{
    "0911/0911.0935_arXiv.txt": {
        "abstract": "The energy and momentum deposited by the radiation from accretion flows onto the supermassive black holes (BHs) that reside at the centres of virtually all galaxies can halt or even reverse gas inflow, providing a natural mechanism for supermassive BHs to regulate their growth and to couple their properties to those of their host galaxies. However, it remains unclear whether this self-regulation occurs on the scale at which the BH is gravitationally dominant, on that of the stellar bulge, the galaxy, or that of the entire dark matter halo. To answer this question, we use self-consistent simulations of the co-evolution of the BH and galaxy populations that reproduce the observed correlations between the masses of the BHs and the properties of their host galaxies. We first confirm unambiguously that the BHs regulate their growth: the amount of energy that the BHs inject into their surroundings remains unchanged when the fraction of the accreted rest mass energy that is injected, is varied by four orders of magnitude. The BHs simply adjust their masses so as to inject the same amount of energy. We then use simulations with artificially reduced star formation rates to demonstrate explicitly that BH mass is not set by the stellar mass. Instead, we find that it is determined by the mass of the dark matter halo with a secondary dependence on the halo concentration, of the form that would be expected if the halo binding energy were the fundamental property that controls the mass of the BH. We predict that the black hole mass, $\\mbh$, scales with halo mass as $\\mbh \\propto \\mhalo^\\alpha$, with $\\alpha \\approx 1.55 \\pm 0.05$ and that the scatter around the mean relation in part reflects the scatter in the halo concentration-mass relation. ",
        "introduction": "Almost all massive galaxies are thought to contain a central supermassive black hole (BH) and the properties of these BHs are tightly correlated with those of the galaxies in which they reside \\citep[e.g.][]{magg98,ferr00,gebh00,trem02,hari04,hopk07b,ho08}.  It is known that most of the mass of the BHs is assembled via luminous accretion of matter \\citep{solt82}.  The energy emitted by this process provides a natural mechanism by which BHs can couple their properties to those of their host galaxies.  Analytic \\citep[e.g.][]{silk98,haeh98,fabi99,adam01,king03,wyit03,murr05,merl08}, semi-analytic \\citep[e.g.][]{kauf00,catt01,gran04,bowe06} and hydrodynamical \\citep[e.g.][]{spri05,dima05,robe06,sija07,hopk07,dima08,okam08,boot09} studies have used this coupling between the energy emitted by luminous accretion and the gas local to the BH to investigate the origin of the observed correlation between BH and galaxy properties, and the buildup of the supermassive BH population.  BHs are expected to regulate the rate at which they accrete gas down to the scale on which they are gravitationally dominant. For example, gas flowing in through an accretion disk can become so hot that its thermal emission becomes energetically important. Scattering of the photons emitted by the accreting matter by free electrons gives rise to the so-called Eddington limit. If the accretion rate exceeds this limit, which is inversely proportional to the assumed radiative efficiency of the accretion disk, then the radiative force exceeds the gravitational attraction of the BH and the inflow is quenched, at least within the region that is optically thin to the radiation.  However, observations indicate that the time-averaged accretion rate is far below Eddington \\citep{koll06}, suggesting the presence of processes acting on larger scales. Indeed, the existence of tight correlations between the mass of the BH and the properties of the stellar bulge indicates that self-regulation may happen on the scale of the bulge \\citep[$\\sim1$~kpc; ][]{adam01,hopk07}, far exceeding the radius within which the BH is gravitationally dominant. However, since galaxy-wide processes such as galaxy mergers can trigger gas flows into the bulge \\citep{sand88,miho94}, it is conceivable that BHs could regulate their growth on the scale of the entire galaxy \\citep[$\\sim10$~kpc; ][]{haeh98,fabi99,wyit03} or even on that of the DM haloes hosting the galaxies \\citep[$\\sim10^2$~kpc; ][]{silk98,ferr02}. Finally, it is possible, perhaps even likely, that self-regulation takes place simultaneously on multiple scales. For example, frequent, short, Eddington-limited outbursts may be able to regulate the inflow of gas on the scale of the bulge averaged over much longer time scales.  In this paper we investigate, using self-consistent simulations of the co-evolution of the BH and galaxy populations, on what scale the self-regulation of BHs takes place.  In Sec.~\\ref{sec:method} we describe the numerical techniques and simulation set employed in this study.  In Sec.~\\ref{sec:results} we demonstrate that BH self-regulation takes place on the scale of the DM halo, and that the BH mass is determined by the binding energy of the DM halo rather than by the stellar mass of the host galaxy. Throughout we assume a flat $\\Lambda$CDM cosmology with the cosmological parameters: $\\{\\Omega_{\\rm m},\\Omega_{\\rm b},\\Omega_\\Lambda,\\sigma_8,n_{\\rm s},h\\}=\\{0.238,0.0418,0.762,0.74,0.951,0.73\\}$, as determined from the WMAP 3-year data \\citep{sper07}. ",
        "conclusions": "\\label{sec:results} It is instructive to first consider under what conditions BHs can regulate their growth. To regulate its growth on a mass scale $\\msr$, a BH of mass $\\mbh$ must be able to inject energy (or momentum) at a rate that is sufficient to counteract the force of gravity on the scale $\\msr$, averaged over the dynamical time associated with this scale. The mass $\\msr$ could, for example, correspond to that of the BH, the stellar bulge, or the dark matter (DM) halo. If the BH cannot inject energy sufficiently rapidly, then gravity will win and its mass will increase. Provided that the maximum rate at which it can inject energy increases with $\\mbh$ (as is for example the case for Bondi-Hoyle and Eddington-limited accretion with a constant radiative efficiency) and provided that this rate increases sufficiently rapidly to counteract the growth of $\\msr$, the BH will ultimately reach the critical mass $m_{\\rm BH,crit}(\\msr)$ required to halt the inflow on the scale $\\msr$.  If, on the other hand, $\\mbh \\gg m_{\\rm BH,crit}(\\msr)$, then the BH will quickly quench the accretion flow and its mass will consequently remain nearly unchanged. The BH will in that case return to the equilibrium value $m_{\\rm BH,crit}(\\msr)$ on the time scale which characterises the growth of $\\msr$.  If the BH regulates its growth on the mass scale $\\msr$ and if $\\mbh \\ll \\msr$, then the critical rate of energy injection required for self-regulation is independent of the mass of the BH. It then follows from Eq.~\\ref{eq:edot} that $\\dot{m}_{\\rm BH}\\propto \\epsf^{-1}$, which implies \\begin{equation} \\left (\\mbh - \\mseed\\right ) \\propto \\epsf^{-1}\\,, \\label{eq:prop} \\end{equation} where $\\mseed$ is the initial mass of the BH. Hence, if the self-gravity of BHs is negligible on the maximum scale on which they regulate their growth and if $\\mbh \\gg \\mseed$, then we expect $\\mbh\\propto \\epsf^{-1}$.  \\begin{figure} \\begin{center} \\includegraphics[width=8.3cm,clip]{massscale.eps} \\end{center} \\vspace{-0.4cm} \\caption{\\label{fig:massscale}Predicted redshift zero global BH mass density (black diamonds) and normalization of the $\\mbh-\\sigma$ relation (black plus signs) as a function of the assumed efficiency of BH feedback, $\\epsf$. Both quantities are normalized to their values in the simulation with $\\epsf=0.15$, which reproduces the observed relations between the mass of the BH and properties of the stellar bulge.  Each point represents a different simulation. For $10^{-4}<\\epsf<1$ all data points track the dotted black line, which is a power-law with index minus one. This implies that in this regime BH mass is inversely proportional to $\\mbh$, and thus that the BHs inject energy into their surroundings at a rate that is independent of $\\epsf$, as expected for self-regulated growth on scales that are sufficiently large for the gravity of the BH to be unimportant. The red data points show results from simulations with a mass resolution that is 8 times worse than the fiducial simulation. The blue data points correspond to simulations with 64 times better resolution than our fiducial resolution, but show results for redshift 2 rather than zero. The agreement between the black, red and blue points confirms numerical convergence and demonstrates that the BHs are already self-regulating at redshift 2.} \\vspace{-0.4cm} \\end{figure}  \\begin{figure*} \\begin{center} \\includegraphics[width=16.6cm,clip]{changing_baryons.eps} \\end{center} \\vspace{-0.8cm} \\caption{\\label{fig:changing_baryons}Median $\\mbh-\\mhalo$ (\\emph{left panel}) and $\\mbh-\\ms$ (\\emph{right panel}) relations for all BHs more massive than $10m_{\\rm seed}$. The black curves correspond to a simulation using our fiducial star formation law and the red, dashed curves show the result for a run in which the star formation efficiency was decreased by a factor of 100.  In order to isolate the effect of stellar mass, we turned off supernova feedback in both runs.  The BH scaling relations therefore differ somewhat from those predicted by our fiducial model, which does include supernova feedback.  Baryons dominate the gravitational potential in the central regions of the galaxy when we use our fiducial star formation law, but DM dominates everywhere in the run with the reduced star formation efficiency.  While the $\\mbh-\\ms$ relation is strongly affected by the change in the star formation efficiency, the relation between BH and halo mass remains invariant. This demonstrates that the BH mass is insensitive to the mass distribution on scales where the stellar mass dominates, and must instead be determined by the mass distribution on larger ($\\gg 10\\,{\\rm kpc}$) scales.} \\vspace{-0.4cm} \\end{figure*}  The black diamonds plotted in Fig.~\\ref{fig:massscale} show the predicted global mass density in BHs at redshift zero as a function of $\\epsf$, the efficiency with which BHs couple energy into the ISM, normalised to the density obtained for $\\epsf=0.15$. Similarly, the black plus signs indicate the normalisation of the $\\mbh-\\sigma$ relation divided by that for the $\\epsf=0.15$ run.  The feedback efficiency, $\\epsf$, is varied, in factors of 4, from $\\epsf=9.2\\times10^{-6}$ to $\\epsf=9.6$, which implies that the fraction of the accreted rest mass energy that is injected (i.e.\\ $\\epsr\\epsf$) varies from $9.2\\times 10^{-7}$ to 0.96.  BH mass is clearly inversely proportional to the assumed feedback efficiency for $10^{-4}<\\epsf< 1$.  For $\\epsf>1$ the trend breaks down because the BH masses remain similar to the assumed seed mass, in accord with Eq.~\\ref{eq:prop}. If we had used a lower seed mass, then the trend would have extended to greater values of $\\epsf$. The deviation from inverse proportionality that sets in below $\\epsf =10^{-4}$ is more interesting. Such low values yield BH masses that are more than $0.15/10^{-4}\\sim 10^{3}$ times greater than observed, in which case they are no longer negligible compared to the masses of their host galaxies. In that case the critical rate of energy deposition will no longer be independent of $\\mbh$ and we do not expect Eq.~\\ref{eq:prop} to hold.  We have thus confirmed that feedback enables BHs to regulate their growth. Moreover, we demonstrated that this self-regulation takes places on scales over which the gravitational influence of the BHs is negligible, provided that the fraction of the accreted rest mass energy that is coupled back into the interstellar medium is $\\gtrsim 10^{-5}$.  To test whether it is the stellar or the dark matter distribution that determines the mass of BHs, we compare the BH masses in two simulations that are identical except for the assumed efficiency of star formation.  One uses our fiducial star formation law, but in the other simulation we reduced its amplitude by a factor of 100, making the gas consumption time scale much longer than the age of the Universe. Because changing the amount of stars would imply changing the rate of injection of supernova energy, which could affect the efficiency of BH feedback, we neglected feedback from star formation in both runs. In the simulation with \\lq normal\\rq\\ star formation the central regions of the galaxies are dominated gravitationally by the baryonic component of the galaxy, whereas in the simulation with reduced star formation the DM dominates everywhere. Fig.~\\ref{fig:changing_baryons} shows the $\\mbh-\\mhalo$ and $\\mbh-\\ms$ relations at redshift 0. While the two runs produce nearly identical BH masses for a fixed halo mass, the $\\mbh-\\ms$ relation is shifted to lower stellar masses by more than an order of magnitude in the model with reduced star formation.  The insensitivity of the relation between $\\mbh$ and $\\mhalo$ to the assumed star formation efficiency demonstrates that the BH mass is not set by the gravitational potential on the scale of the galaxy.  We have verified that the same result holds at redshift two for the simulations with 64 times better mass resolution.  Clearly, stellar mass does not significantly influence the relation between the mass of the BH and that of its host halo. This implies that BH self-regulation occurs on the scale of DM haloes.  If the rate by which the BHs inject energy is independent of the assumed feedback efficiency, then we expect the same to be true for the factor by which BH feedback suppresses star formation. This is confirmed by comparison of the global SFRs in runs with different values of $\\epsf$ (see Fig.~6 of BS09).  \\begin{figure} \\begin{center} \\includegraphics[width=8.3cm,clip]{mhalombh.eps} \\end{center} \\vspace{-0.4cm} \\caption{\\label{fig:mhalombh}The relation between BH mass and DM halo mass for all BHs that belong to central galaxies and have masses greater than $10m_{\\rm seed}$. The DM halo mass, $m_{200}$, is defined as the mass enclosed within a sphere, centred on the potential minimum of the DM halo, that has a mean internal density of 200 times the critical density of the Universe.  The grey pixels show the results from our fiducial simulation ($\\epsf=0.15$), with the colour of each pixel set by the logarithm of the number of BHs in that pixel.  The solid, red line shows the observational determination of the $\\mbh-\\mhalo$ relation (Bandara et al. 2009) and has a slope of 1.55. The dotted, red lines show the 1$\\sigma$ errors on the observations.  The simulation agrees very well with the observed relation.  The value of the slope and the scatter (which correlates with the concentration of the DM halo) suggest that the halo binding energy, rather than mass, determines the masses of BHs.} \\vspace{-0.4cm} \\end{figure}  Fig.~\\ref{fig:mhalombh} compares the predicted $\\log_{10}\\mbh-\\log_{10}\\mhalo$ relation with observation \\citep{band09}.  The agreement is striking. The slope and normalization of the observed $\\log_{10}(\\mbh/\\msun)-\\log_{10}(\\mhalo/10^{13}\\,\\msun)$ relation are $1.55\\pm0.31$ and $8.18\\pm0.11$ respectively, whereas the simulation predicts $1.55\\pm0.05$ and $8.01\\pm0.04$.  Note that the simulation was only tuned to match the normalization of the relations between $\\mbh$ and the galaxy stellar properties.  If the energy injected by a BH is proportional to the halo gravitational binding energy, then, for isothermal models \\citep{silk98}, $\\mbh\\propto\\mhalo^{5/3}$.  Here we extend these models to the more realistic universal halo density profile \\citep{nava97}, whose shape is specified by a concentration parameter, $c$ (we assumed $c\\propto v_{\\rm max}^2/v_{\\rm v}^2$, where $v_{\\rm max}$ and $v_{{\\rm v}}$ are the maximum halo circular velocity and the circular velocity at the virial radius respectively).  It is known that concentration decreases with increasing halo mass, $c\\propto\\mhalo^{-0.1}$ \\citep{bulo01,duff08}, which then affects BH mass through the dependence of halo binding energy on concentration. If the total energy injected by a BH of a given mass is proportional to the energy required to unbind gas from a DM halo \\citep{loka01} out to some fraction of the virial radius, $\\re/r_{\\rm v}$ then \\begin{align} \\mbh\\propto&\\Bigg(\\frac{c}{\\big(\\ln(1+c)-c/(1+c)\\big)^2}\\Bigg)\\times \\nonumber\\\\ &\\Bigg(1-\\frac{1}{(1+c\\frac{\\re}{r_{\\rm v}})^2}-\\frac{2\\ln(1+c\\frac{\\re}{r_{\\rm v}})}{1+c\\frac{\\re}{r_{\\rm v}}}\\Bigg)m_{\\rm v}^{5/3}\\,. \\end{align} Inserting $c\\propto m_{\\rm v}^{-0.1}$ and computing the logarithmic derivative with respect to $m_{\\rm v}$ in the mass range $10^{10}\\,\\msun < m_{\\rm v} < 10^{14}\\,\\msun$, we find that the slope is a weak function of $\\re/r_{\\rm v}$ that varies from 1.50 at $\\re=10^{-1}r_{\\rm v}$ to 1.61 at $\\re=r_{\\rm v}$.  The close match between theory, simulation and observation suggests that the halo binding energy, rather than halo mass, determines the mass of the BH.  The residuals from the $\\mbh-\\mhalo$ relation ($\\Delta\\log_{10}\\mbh$) are correlated with halo concentration (Spearman rank correlation coefficient $\\rho=0.29$, probability of significance $P= 0.9998$) as would be expected if $\\mbh$ is sensitive to the halo binding energy. The residuals are also correlated with galaxy stellar mass, though much less strongly ($\\rho=0.09$; $P=0.96$).  Taken together, these correlations tell us that, at a given halo mass, galaxies with BHs more massive than the average will also contain a larger than average amount of stars, and are hosted by more concentrated haloes. This suggests that the galaxy stellar mass is also determined by the halo binding energy.  Thus, outliers in the $\\mbh-\\mhalo$ relation may still lie close to the mean $\\mbh-\\ms$ relation.  Furthermore, higher concentrations imply earlier formation times and spheroidal components do indeed typically host old stellar populations.  In addition to the \\lq quasar mode\\rq\\ of feedback discussed in this work, it has recently become clear that a second \\lq radio mode\\rq\\ may be required to quench cooling flows in galaxy groups and clusters \\citep[see e.g.][for a review]{catt09}.  Although we do not explicitly include a \\lq radio mode\\rq\\ in the current work, the AGN feedback prescription explored here is capable of suppressing cooling flows, at least on group scales, providing excellent matches to observed group density and temperature profiles as well as galaxy stellar masses and age distributions \\citep{mcca09}. It is known that BHs obtain most of their mass in the \\lq quasar mode\\rq\\ \\citep{solt82} so any discussion of what detemines the masses of BHs must focus primarily on this mode of accretion.  Finally, the ability of a BH to quench cooling flows in the \\lq radio mode\\rq\\ is expected to be closely related the virial properties of the hot halo \\citep{catt09} and would therefore provide an additional link between BHs and DM haloes over and above what we discuss here and so serve to make any fundamental connection between BH mass and the properties of the DM halo even stronger.  We conclude that our simulation results suggest that in order to effectively halt BH (and galaxy) growth, gas must not return to the galaxy on a short timescale. This requires that the BH injects enough energy to eject gas out to scales where the DM halo potential is dominant. The mass of the BH is therefore determined primarily by the mass of the DM halo with a secondary dependence on halo concentration, of the form that would be expected if the BH mass were controlled by the halo binding energy.  The tight correlation between $\\mbh$ and $\\ms$ is then a consequence of the more fundamental relations between halo binding energy and both $\\mbh$ and $\\ms$.  \\vspace{-0.4cm}"
    },
    "0911/0911.2792_arXiv.txt": {
        "abstract": "In cold dark matter cosmologies, small mass halos outnumber larger mass halos at any redshift. However, the lower bound for the mass of a galaxy is unknown, as are the typical luminosity of the smallest galaxies and their numbers in the universe. The answers depend on the extent to which star formation in the first population of small mass halos may be suppressed by radiative feedback loops operating over cosmological distance scales. If early populations of dwarf galaxies did form in significant number, their relics should be found today in the Local Group. These relics have been termed ``fossils of the first galaxies''.  This paper is a review that summarizes our ongoing efforts to simulate and identify these fossils around the Milky Way and Andromeda.  It is widely believed that reionization of the intergalactic medium would have stopped star formation in the fossils of the first galaxies. Thus, they should be among the oldest objects in the Universe. However, here we dispute this idea and discuss a physical mechanism whereby relatively recent episodes of gas accretion and star formation would be produced in some fossils of the first galaxies. We argue that fossils may be characterized either by a single old population of stars or by a bimodal star formation history.  We also propose that the same mechanism could turn small mass dark halos formed before reionization into gas-rich but starless ``dark galaxies''.  We believe that current observational data supports the thesis that a fraction of the new ultra-faint dwarfs recently discovered in the Local Group are in fact fossils of the first galaxies. ",
        "introduction": "There are many questions that remain open in cosmology with regard to the mass, number and properties of the smallest galaxies in the universe. Have we already discovered the smallest galaxies in the universe or we are still missing an elusive but large population of ultra-faint dwarf galaxies?  In cold dark matter (CDM) cosmologies most of the dark halos that formed before reionization had masses smaller than $10^8-10^9$~M$_{\\odot}$ \\citep[\\eg,][]{GnedinO:97}. The small mass halos that survived tidal destruction to the modern epoch, were they able to form stars, would constitute a sub-population of dwarf satellites orbiting larger halos. Small mass dark halos significantly outnumber more massive galaxies like the Milky Way and can be located in the voids between luminous galaxies \\citep[\\eg,][]{Hoeft:06, Ricotti:09}.  However, until recently (\\ie, before 2005) observations did not show a large number of satellites around massive galaxies like the Milky Way and Andromeda. This became known as the ``missing galactic satellite problem, \\citep{Klypin:99, Moore:99}. The voids between bright galaxies appear to be devoid of dwarf galaxies \\citep{Karachentsev:04, Karachentsevetal06, Tullyetal06}. While the abundance of dwarfs in large voids may not pose a problem to CDM cosmology, as shown by \\cite{TinkerC:09}, it is unclear whether the predictions of the number of faint dwarfs in the Local Group is consistent with both the number of observed Milky Way dwarf satellites and the number of relatively isolated dwarfs in the local voids.  Historically, the discrepancy between observation and theory on the number of dwarf galaxies has been interpreted in two ways: 1) as a problem with the CDM paradigm that could be solved by a modification of the dark matter properties -- for instance by introducing warm dark matter \\citep[\\eg,][]{BodeO:01} -- or 2) as an indicator of feedback processes that are exceptionally efficient in preventing star formation in small mass halos, which remain mostly dark \\citep[\\eg,][]{HaimanAR:00}.  The recently discovered population of ultra-faint dwarfs \\citep{Belokurovetal06a, Belokurovetal07, Irwinetal07, Walshetal07, Willmanetal05ApJ, Willmanetal05AJ, Zuckeretal06a, Zuckeretal06b, Ibataetal07, Majewskietal07, Martinetal06} in combination with a proper treatment of observational incompleteness \\citep{SimonGeha07, Koposov:08, Tollerudetal08, Gehaetal08} has increased the estimated number of Milky Way satellites to a level that can be more easily reconciled with theoretical expectations. For instance, the suppression of dwarf galaxy formation due to intergalactic medium (IGM) reheating during reionization \\citep{Babul:92, Efstathiou:92b, ShapiroG:94, HaimanRL:96, Thoul:96, Quinn:96, Weinberg:97, NavarroS:97, Bullock:00, Gnedin:00, Somerville:02, Dijkstra:04, Shapiro:04, Hoeft:06, OkamotoGT08, Ricotti:09}, in conjunction with a strong suppression of star formation in small mass pre-reionization dwarfs, may be sufficient to explain the observed number of Milky Way satellites. In the near future we can hope to answer perhaps a more interesting question: what is the minimum mass that a galaxy can have? This is a non trivial and fundamental question in cosmology. Answering it requires a better understanding of the feedback mechanisms that regulate the formation of the first galaxies before reionization and the details of the process of reionization feedback itself.  The formation of the first dwarf galaxies - before reionization - is self-regulated on cosmological distance scales. This means that the fate of small mass halos (\\ie, whether they remain dark or form stars) depends on local and global feedback effects. This type of galaxy feedback differs from the more familiar model operating in normal galaxies (\\eg, SN explosions, AGN feedback, etc), where the feedback is responsible for regulating the star formation rate within the galaxy itself but does not impact star formation in other distant galaxies. Rather, before reionization, each proto-dwarf galaxy reacts to the existence of all the others.  Different theoretical assumptions and models for the cosmological self-regulation mechanisms will, of course, produce different predictions for the number and luminosity of the first dwarf galaxies \\citep{HaimanAR:00, RicottiGnedinShull02b, WiseA:08, OShea:08, RicottiGnedinShull08}.  We now introduce the basic concepts on how feedback-regulated galaxy formation operates in the early universe (\\ie, before reionization):\\\\ A cooling mechanism for the gas is required in order to initiate star formation in dark halos. In proto-galaxies that form after reionization this is initially provided by hydrogen Lyman-alpha emission. This cooling is efficient at gas temperatures of 20,000 K but becomes negligible below $T \\sim 10,000$~K. Later, as the temperature drops below 10,000 K, the cooling is typically provided by metal line cooling. In the first galaxies, however, both these cooling mechanisms may be absent. This is because the first dwarf galaxies differ when compared to present-day galaxies in two respects: 1) they lack important coolants -- such as carbon and oxygen -- because the gas is nearly primordial in composition, and 2) due to the smaller typical masses of the first dark halos, the gas initially has a temperature that is too low to cool by Lyman-alpha emission.  The gas in small mass halos with circular velocity $v_{\\rm vir}=(GM_{\\rm tot}/r_{\\rm vir})^{1/2} \\simlt 20$~km~s$^{-1}$, where $r_{\\rm vir}$ is the virial radius -- roughly corresponding to a mass $M_{\\rm tot} \\simlt 10^8$ M$_{\\odot}$ at the typical redshift of virialization -- has a temperature at virialization $T \\simlt 10,000$ K. Hence, if the gas has primordial composition it is unable to cool by Lyman-alpha emission and initiate star formation unless it can form a sufficient amount of primordial H$_2$ (an abundance $x_{H_2} \\simgt 10^{-4}$ is required).  Because molecular hydrogen is easily destroyed by far-ultraviolet (FUV) radiation in the Lyman-Werner bands ($11.3<h\\nu<13.6$~eV) emitted by the first stars, it is widely believed that the majority of galaxies with $v_{\\rm vir}<20$~km~s$^{-1}$ remain dark \\citep[\\eg,][]{HaimanAR:00}. However, several studies show that even if the FUV radiation background is strong, a small amount of H$_2$ can always form, particularly in relatively massive halos with virial temperature of several thousands of degrees \\citep{WiseA:07, OShea:08}. Thus, negative feedback from FUV radiation may only delay star formation in the most massive pre-reionization dwarfs rather than fully suppress it \\citep{Machacek:00, Machacek:03}. We will argue later that hydrogen ionizing radiation ($h\\nu > 13.6$~eV) in the extreme-ultraviolet (EUV) plays a far more important role in regulating the formation of the first galaxies than FUV radiation. Thus, in our opinion, models that do not include 3D radiative transfer of H and He ionizing radiation cannot capture the most relevant feedback mechanism that regulates galaxy formation in the early universe \\citep{RicottiGnedinShull02a, RicottiGnedinShull02b}.  After reionization, the formation of dwarf galaxies with $v_{\\rm vir}<20$~km~s$^{-1}$ is strongly inhibited by the increase in the Jeans mass in the IGM.  Thus, according to this model, reionization feedback and negative feedback due to H$_2$ photo-dissociation by the FUV background (important before reionization) determine the mass of the smallest galactic building blocks.  The resultant circular velocity of the smallest galactic building blocks is $v_{\\rm vir} \\sim 20$ km~s$^{-1}$, roughly corresponding to masses $M_{\\rm tot} \\sim 10^8-10^9$~M$_\\odot$. If this is what really happens in the early universe, the ``missing Galactic satellite problem'' can be considered qualitatively solved because the predicted number of Milky Way satellites with $v_{\\rm vir}>20$ km~s$^{-1}$ is already comparable to the estimated number of observed satellites after applying completeness corrections (although this model may still have problems reproducing the observations in detail).  However, as briefly mentioned above, we have argued for some time that most simulations of the first galaxies cannot capture the main feedback mechanism operating in the early universe because they do not include a key physical ingredient: radiative transfer of H and He ionizing radiation. Our simulations of the formation of the first galaxies are to date the only simulations of a cosmologically representative volume of the universe (at $z \\sim 10$) that include 3D radiative transfer of H and He ionizing radiation \\citep{RicottiGnedinShull02a, RicottiGnedinShull02b, RicottiGS:08}. Figure~\\ref{fig:pfr2} shows the evolution of ionized bubbles around the first galaxies in a cubic volume of 1.5~Mpc in size at redshifts $z=21.2, 17.2, 15.7, 13.3$ from one of our simulations. The results suggest that negative feedback from the FUV background is not the dominant feedback mechanism that regulates galaxy formation before reionization. Rather, ``positive feedback'' on H$_2$ formation from ionizing radiation \\citep{HaimanRL:96, RicottiGS:01} dominates over the negative feedback of H$_2$ dissociating radiation. Hence, a strong suppression of galaxy formation in halos with $v_{\\rm vir}<20$~km~s$^{-1}$ does not take place. In this latter case, some galactic satellites would be the fossil remnants of the first galaxies. Comparisons of simulated pre-reionization fossils to dwarf spheroidals in the Local Group show remarkable agreement in properties \\citep[][hereafter RG05]{RicottiGnedin05}. Based on the results of the simulations, we also suggested the existence of the ultra-faint population before it was discovered about a year later \\citep[see RG05,][]{BovillR:09}.  \\begin{figure*}[pth] \\epsscale{1.1} \\plotone{4tile1.eps} \\caption{3D rendering of cosmological \\HII~ regions (fully ionized gas is red and partially ionized gas is blue) around the first galaxies in a box of 1.5~Mpc. The four boxes show a time sequence at redshifts $z=21.2, 17.2, 15.7, 13.3$ for the simulation S2 from \\cite{RicottiGnedinShull02b}. The rendering shows several tens of small size \\HII~ regions around the first galaxies (there are a few hundreds of galaxies in this volume). A movie of the same simulation shows that the \\HII~ regions are short lived: they form, expand to a size comparable to the large scale filamentary structure of the IGM and recombine, promoting the formation of molecular hydrogen inside the relic \\HII~ regions.} \\label{fig:pfr2} \\end{figure*}  \\subsection{Definition of ``pre-reionization fossils'' }  Throughout this paper we define ``pre-reionization fossils'' as the dwarfs hosted in halos with a maximum circular velocity remaining below 20~km~s$^{-1}$ at all times during their evolution: $v_{\\rm max}(t)<20$~km~s$^{-1}$. It will become clear in this paper that this definition is {\\it not} directly related to the ability of fossils to retain gas and form stars after reionization: in \\S~\\ref{sec:infall} we describe a mechanism in which small mass halos with $v_{\\rm max}(t)< 20$~km~s$^{-1}$ are able to have a late phase of gas accretion and possibly star formation.  Our definition of fossil reflects the {\\it special cooling mechanisms and feedback processes} that regulate star formation and the number of luminous halos with $v_{\\rm max}(t)<20$~km~s$^{-1}$, before and {\\it after} reionization. In proto-fossil galaxies -- even adopting the most conservative assumption of maximum efficiency of shock heating of the gas during virialization -- the gas is heated to a temperature below $T \\sim 10,000$~K. Thus the gas cannot cool by Lyman-alpha emission, a very efficient coolant. The cooling of the gas is dominated either by H$_2$ roto-vibrational line emission or by metal cooling, important if the metallicity exceeds $Z \\sim 10^{-3}$~Z$_\\odot$ \\citep[\\eg,][]{BrommF:01, SantoroS:06, RicottiGS:08, Smithetal:09}. These coolants are much less efficient than Lyman-alpha emission. Moreover, H$_2$ abundance and cooling is modulated and often suppressed by the FUV and EUV radiation fields. The FUV radiation in the H$_2$ Lyman-Werner bands and hard ultraviolet radiation have large mean free paths with respect to the typical distances between galaxies, thus their feedback is global in nature. Qualitatively, this explains why the first galaxies have low luminosities and low surface-brightness, similar to dwarf spheroidal (dSph) galaxies in the Local Group \\citep[RG05,][]{BovillR:09, SalvadoriF:09}.  Simulations also show that stars in the the first galaxies do not form in a disk but in a spheroid \\citep[RG05,][]{RicottiGS:08}. A thin galactic disk is not formed because of the high merger rates and the low masses of dark halos in the early universe. Roughly, pre-reionization fossils have a mass at virialization $M_{\\rm tot}<10^8$~M$_\\odot$, assuming they form at $z_{\\rm vir} \\sim 10$, but their mass may increase by up to one order of magnitude by $z=0$ due to secondary infall \\citep[RGS02a, b,][]{ RicottiGnedinShull08}. Secondary infall does not affect $v_{\\rm max}$, which remains roughly constant after virialization.  \\subsection{Pre-reionization fossils and reionization}  The critical value of $v_{\\rm max, crit}$ for which dwarf galaxy formation is suppressed by reionization feedback is close to the 20~km~s$^{-1}$ value that defines a fossil, but it is {\\it not necessarily the same value}. Indeed, it can be significantly larger than 20~km~s$^{-1}$ if the IGM is heated to $T \\gg 10,000$~K \\citep{RicottiGS:00}. Thus, we expect that the virialization of new ``pre-reionization fossils'' is strongly suppressed after reionization due to IGM reheating (\\ie, they mostly form before reionization). However, pre-reionization fossils and dark halos with $v_{\\rm max}<20$~km~s$^{-1}$ that virialized before reionization may accrete gas and, in certain cases, form new stars after reionization at redshifts $z<1-2$ \\citep{Ricotti:09}.  Unfortunately, the value of $v_{\\rm max, crit}$ is uncertain due to our poor understanding of the thermal history of the IGM \\citep{RicottiGS:00}. The uncertainty surrounding the IGM equation of state may partially explain the differences found in literature on the values for $v_{\\rm max, crit}$ and the different levels of suppression of star formation as a function of the halo mass after reionization \\citep[\\eg,][]{Weinberg:97, Gnedin:00, Hoeft:06, OkamotoGT08}. Regardless of assumptions for the reionization feedback model, one should bear in mind that no halo with $v_{\\rm max} <20$~km~s$^{-1}$ can form stars after reionization unless the gas in those halos has been significantly pre-enriched with metals. For instance, the model by \\cite{Koposov:09} assumes star formation after reionization in halos as small as $v_{\\rm max}\\sim 10$~km~s$^{-1}$. With this assumption they find that their model is consistent with observations of ultra-faint dwarfs but claim that fossils are not needed to explain the data. However, star formation in such small halos can only take place in a gas that was pre-enriched with metals, suggesting the existence of older populations of stars in those halos. Indeed, according to our definition, the smallest post-reionization dwarfs with $10~{\\rm km~s}^{-1}< v_{\\rm max} < 20$~km~s$^{-1}$ in the \\cite{Koposov:09} model are ``fossils''. As stated above, fossils may also be able to form stars after reionization due to a late phase of cold gas accretion from the IGM \\citep{Ricotti:09}.  \\subsection{Identification of pre-reionization fossils in observations}  Pre-reionization fossils are not easily identifiable because $v_{\\rm max}$ cannot be measured directly from observations. Understanding the star formation history of dwarf galaxies may help in this respect, as fossils likely show some degree of suppression of their star formation rate occurring about 12.5 Gyr ago due to reionization. However, their identification based on their star formation history may be complicated because some pre-reionization fossils in the last 10~Gyr may have had a late phase of gas accretion and star formation. The caveat is that star formation histories cannot be measured with accuracy better than to within 1-2~Gyr and the accuracy becomes increasingly poorer for old stellar populations. Thus, it is impossible to prove whether an old population of stars formed before reionization (which happened about 1~Gyr after the Big Bang) or at $z \\sim 3$, when the Milky Way was assembled. Nevertheless, ultra-faint dwarfs that show some degree of bimodality in their star formation history are candidates for being pre-reionization fossils.  According to results by RG05 and \\cite{BovillR:09}, Willman~1, Bootes~II, Segue~1 and Segue~2 do not lie on the luminosity-surface brightness relationship of simulated pre-reionization fossils. This result is based on the assumption that fossil properties are not modified by tides.  Their surface brightness is larger than the model predictions for objects with such low-luminosity. An as yet undiscovered population of ultra-faints with lower surface brightness is instead predicted by our simulations. It is likely that the properties of the lowest luminosity ultra-faints may have been modified by tidal forces due to their proximity to the Milky Way disk.  Although it is difficult to identify individual fossils, statistical arguments suggest that at least some ultra-faint dwarf galaxies are pre-reionization fossils. This is because the number of satellites from N-body simulations with $v_{\\rm max}(t)>20$~km~s$^{-1}$ is substantially smaller than the estimated number of observed satellites after completeness corrections. Admittedly the current theoretical and observational uncertainties on the number of satellites are still large. However, if the estimated number (after completeness corrections) of ultra-faint dwarfs increases further, the existence of pre-reionization fossils will be inescapably proven. This is especially the case if a population of ultra-faint dwarfs with luminosities similar to Willman~1, Bootes~II, Segue~1 and Segue~2 but surface brightness below the current sensitivity limit of the SDSS -- as predicted by our simulations -- is discovered.  The possibility of identifying the fossils of the first galaxies in our own backyard is very exciting. It would greatly improve our understanding of the physics involved in self-regulating the formation of the first galaxies before reionization. Clearly, even the launch of the James Webb Space Telescope ({\\it JWST}), would not yield the wealth of observational data on the formation of the first galaxies that could be obtained by studying ultra-faint galaxies in the Local Group.  The rest of the paper is organized as follows. In \\S~\\ref{sec:obs} we briefly review and discuss observational data on Galactic satellites, in \\S~\\ref{sec:theo} we summarize the results of simulations of the formation of the first galaxies in a cosmological volume and the effect of reionization feedback on galaxy formation. In \\S~\\ref{sec:infall} we discuss a recently proposed model for ``late gas accretion'' from the IGM onto small mass halos. In \\S~\\ref{sec:comp} we compare the theoretical properties of simulated pre-reionization fossils to observations. In \\S~\\ref{sec:disc} we compare different ideas for the origin of classical and ultra-faint dwarf spheroidals. We present our conclusions in \\S~\\ref{sec:conc}. ",
        "conclusions": "\\label{sec:conc}  We have summarized our work on the formation of the first galaxies before reionization (\\ie, pre-reionization dwarfs) and the quest to identify the fossils of these first galaxies in the Local Group.  The definition of a pre-reionization fossil is not directly related to the suppression of star formation experienced by these galaxies due to reionization feedback. Indeed, we discussed how pre-reionization fossils may experience a late phase of gas accretion and possibly star formation at redshift $z<1-2$. Most importantly, fossils are a population of dwarf galaxies whose formation (\\ie, the fraction of halos that are luminous) is self-regulated on cosmological distance scales by radiative processes. Their existence is not certain due to a possible strong negative feedback that may prevent the majority of these halos from ever forming stars. In addition, if negative feedback heavily suppresses the number and luminosity of these first galaxies, more massive halos with $v_{\\rm max}>20$~km~s$^{-1}$ will evolve differently because of the lower level of metal pre-enrichment of the IGM.  To summarize, the critical circular velocity $v_{\\rm max} \\sim 20$~km~s$^{-1}$ that we adopt 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. However, it is also close to the critical value for continued gas accretion after IGM reionization \\citep{Gnedin-filteringmass00, Hoeft:06, OkamotoGT08}.  The number of Milky Way and M31 satellites provides an indirect test of galaxy formation and the importance of positive and negative feedback in the early universe. This test, although the uncertainties are large, supports the idea that a fraction of the new ultra-faint dwarfs are fossils.  The good agreement of the SDSS and new M31 ultra-faint dwarf properties with predictions of our simulations (RG05, GK06, Bovill \\& Ricotti 2009) does not prove the primordial origin of the new ultra-faint dwarfs, but it supports this possibility.  More theoretical work and more observational data are needed to prove that some dwarfs in the Local Group are true fossils of the first galaxies. Future theoretical work should focus on improving the accuracy of predictions on the properties of dwarf galaxies formed before reionization and their evolution to the present day. Modeling the evolution of the baryonic component after reionization in dwarf satellites and in the Milky Way -- Andromeda system may be necessary to make robust predictions. More observational data will certainly be available in the near future. A large number of surveys, both at optical and radio wavelengths will be online in the near future (\\eg, {\\it Pan-STARRS, LSST, ALMA, EVLA, JWST, SKA} to mention a few). Different survey strategies may be used to find and characterize fossil dwarf galaxies. A deep pencil beam survey would be useful to find the faintest dwarf satellites of the Milky Way and determine more precisely their Galactocentric distribution. A shallower all sky survey could be used to quantify the degree of anisotropy in the distribution of satellites around the Milky Way.  The star formation history of the dwarf galaxies is not strongly discriminatory because fossil galaxies may have a late phase of gas accretion and star formation during the last $9-10$~Gyrs \\citep{Ricotti:09}. The distinction between fossils and non-fossil galaxies may be quite elusive but it is nevertheless important to understand galaxy formation and feedback in the early universe. Arguments based on counting the number of dwarfs in the Local universe are among the more solid arguments that could be used to prove the existence of fossil galaxies (see Table~\\ref{tab:count}).  Future tests may be provided by deep surveys looking for ultra-faint galaxies in the local voids or looking for gas in dark galaxies (\\ie, dark halos that have been able to accrete gas from the IGM at $z<1-2$). Ultra-faint dwarfs should be present in the voids if dwarf galaxies formed in large numbers before reionization (Bovill \\& Ricotti, in preparation). If pre-reionization dwarfs never formed due to dominant negative feedback in the early universe, it is possible that a faint (in H$\\alpha$ and 21~cm emission) population of dark galaxies exists in the outer parts of the Local Group. Hence, another way to detect fossil galaxies in the outer parts of the Milky Way or outside the super-galactic plane would be to search for neutral or ionized gas that they may have accreted from the IGM. Future radio telescopes (\\eg, {\\it ALMA, EVLA, SKA}) may be able to detect neutral hydrogen in dark galaxies or in ultra-faint dwarfs. Ionized gas in the outer parts of dark halos may be observed in absorption along the line of sight of distant quasars (for instance in \\mbox{\\ion{O}{6}} or \\mbox{\\ion{O}{4}} with {\\it COS} on the {\\it HST}). However, the probability that a line of sight toward a quasar intersects the ionized gas collected from the IGM by dark or fossil galaxies might be small. Additional theoretical work is required to address these issues."
    },
    "0911/0911.2271_arXiv.txt": {
        "abstract": "Protoplanetary disks are composed primarily of gas (99\\% of the mass). Nevertheless, relatively few observational constraints exist for the gas in disks. In this review, I discuss several observational diagnostics in the UV, optical, near-IR, mid-IR, and (sub)-mm wavelengths that have been employed to study the gas in the disks of young stellar objects. I concentrate in diagnostics that probe the inner 20 AU of the disk, the region where planets are expected to form. I discuss the potential and limitations of each gas tracer and present prospects for future research. ",
        "introduction": "\\label{intro} At the time when giant planets form, the mass (99\\%) of the protoplanetary disk is dominated by gas. Dust is a minor constituent of the mass of the disk. Still, it dominates the opacity; consequently, it is much easier to observe. Therefore, most of the observational constraints of disks have been deduced from the study of dust emission (e.g., see chapter by Henning et al.). In contrast, direct observational constraints of the gas in the disk are relatively scarce. Nonetheless, to obtain direct information from the gas content of the disk is crucial for answering fundamental questions in planet formation such as: how long does the protoplanetary disk last?, how much material is available for forming giant planets?, how do the density and the temperature of the disk vary as a function of the radius?, what are the dynamics of the disk?.  Although we have learned important insights from disks from dust observations (see for example the reviews by Henning et al. 2006 and Natta et al. 2007), dust presents several limitations: (i) dust spectral features are broad; consequently, dust emission does not provide kinematical information; (ii) dust properties are expected to change during the planet formation process; therefore, quantities such as the gas-to-dust ratio (needed to derive the disk mass from dust continuum emission in the (sub)-mm) are expected to strongly vary with respect to the conditions of the Interstellar Medium (ISM). In addition, dust signatures are strongly related to dust size. In particular, as soon a dust particle reaches a size close to a decimeter it becomes practically \"invisible\". For example, a crucial quantity such as the disk dissipation time scale is normally deduced from the decline of the fraction of the sources presenting near-IR (JHKL) excess in clusters of increasing ages (e.g., Haisch et al. 2001). In reality, the time scale that we want to constrain is not the time scale in which the small warm dust particles disappear from the surface layers in the inner disk, but the time scale in which the gas in the disk disappears.  In summary, independent information of the gas is crucial to understand protoplanetary disk structure and {\\it combined studies of gas and dust in the disk are needed}. In particular, we are interested in deriving constraints for the inner disk, the region where the planets are expected to form (R$<$20 AU). The observational study of the gas in the inner disk is a relatively new emerging topic. Only with the advent of ground-based high-spectral resolution infrared spectrographs and space-born infrared spectrographs has this research become possible.  \\begin{figure} \\includegraphics[width=0.5\\textwidth]{disk_emission.jpg} \\caption{Line shape, inclination and the emitting region of a disk in Keplerian rotation. The double peaked profile indicates emission from a rotating disk. The line width and double peak separation depend on the inclination of the disk. The line profile wings depend on the innermost radius of the emission. In this toy model, it is assumed that the intensity is constant as a function of the radius and a spectral resolution of 1 km/s.} \\label{fig:1}       % \\end{figure}  The main tool used to study the gas in the disk is molecular spectroscopy. The gas in the disk is heated by collisions with dust, shocks or UV or  X-rays from the central star. The heated molecules populate different rotational and vibrational levels and when they de-excite emission lines are produced. The key is that these emission lines have the imprint of the physical conditions of the gas where they originated. Important constraints can be derived from emission lines: (i) line ratios of different transitions constrain the excitation mechanism (shocks, UV or X-rays) and the temperature of the gas responsible of the emission; (ii) if the emitting gas is optically thin, then the measured line fluxes are proportionally related to the amount of molecules present; therefore, line fluxes can be used to derive column densities and gas masses; (iii) line shapes and spatial extended provide constraints on the size and geometry of emitting region as well as the dynamics of the emitting gas. Figure 1 displays a pedagogical example of how a line profile changes as a function of the inclination and size of the emitting region in the disk. The line width and double peak separation depend on the inclination . The line profile wings depend on the innermost radius of the emission. For the interested reader, recent reviews  by Najita et al. (2000 \\& 2007a), Blake (2003) and Carr (2005) cover the subject of observational studies of gas in disks.  \\begin{figure} \\includegraphics[width=0.5\\textwidth]{disk_structure.jpg} \\label{fig:2} \\caption{Cartoon of the structure of an optically thick disk. The disk has a hot inner region (R$<$20 AU) and a cold outer region (R$>$20 AU). Near-IR and mid-IR diagnostics probe the inner disk, (sub)-mm diagnostics probe the outer disk. The disk has a vertical structure: a dense mid-plane and a hotter less dense surface layer. Near-IR and mid-IR gas and dust emission features originate in the optically thin surface layer of the inner disk. Near-IR and mid-IR continuum originates from the optically thick interior layer. At (sub)-mm wavelengths we observe line emission from optically thick CO gas and optically thin emission of cold dust located in the outer disk.} \\end{figure} ",
        "conclusions": "In this contribution I have discussed several observational diagnostics that allow to study the gas in protoplanetary disks highlighting the strengths and limitation of each diagnostic. Here, I summarize important general points that one should bear in mind about observational constraints of gas in disks. \\begin{itemize} \\item Disks have a radial temperature structure; therefore, different gas diagnostics are employed for probing different regions of the disk. UV probes the hottest gas in the innermost disk (R$<<$1 AU, T$>$2000 K), NIR and MIR probes hot and warm gas in the inner disk (R$\\sim$0.1-20 AU, T$\\sim$100-2000 K), (sub)-mm probes cold gas in the outer disk (R$>20$ AU, T$\\sim$20-100 K). Note that disks have typical sizes of several hundreds of AU. \\item To observe a line from gas (in emission or absorption) the dust in the medium where the line is produced must be optically thin at the observed wavelength, or the gas should be at a higher temperature than the dust. \\item An optically thick disk has a vertical temperature structure. The surface layer is hotter than the interior mid-plane layer. The dust in the surface layer is optically thin, the dust in the interior layer (of the inner disk) is optically thick. In optically thick disks, hot and warm gas emission lines from the UV to the IR are produced in the surface layer. Note that the amount of mass in the surface layer is much smaller than the amount of mass in the interior layer. Gas lines in the IR probe only a limited amount of material. H$_2$ lines in the IR are optically thin. CO lines in the IR can be optically thick. \\item Sub-mm lines trace only the cold outer regions of the disk at R$>$20 AU. Planets usually are not expected to form at this distances. Still, most of the mass of the disk is at R$>$20 AU. \\item CO in the (sub)-mm is not a reliable disk gas mass tracer because (i) CO freezes onto dust grain surfaces, (ii) the CO/H$_2$ conversion factor is unknown in disks (iii) the emission is likely optically thick. \\item Be aware that so far we do not measure directly the mass of gas in disks. We deduce the disk mass from dust continuum emission in the (sub)-mm from cold dust in the outer disk (R$>$20 AU). The deduced mass depends on the dust opacity, dust temperature and the dust-to-gas ratio assumed. The disk mass observed is from material located at R$>$50 AU, and disk evolutionary models are typically of disks of R$<$50 AU. The mass and surface density profiles of gas at R$<$20 AU are poorly constrained from observations. \\item Modeling is required to interprete line observations. Conventional assumptions are that the observed gas is at LTE and is optically thin. The rotational diagrams employed for deducing the temperature of the observed gas make use those two assumptions. To deduce the inner most radius of the observed emission (i.e., R$_{\\rm in}$) from line profile modeling, it is required to assume the disk inclination and the dependence of the intensity as a function of radius ($I(R)$). Inclinations are typically deduced from (sub)-mm imaging and extended to the inner disk. On the other hand, if we want to deduce independently the inclination (e.g., to search for evidence of warps) we need to assume (or calculate) R$_{\\rm in}$ and $I(R)$. \\item Most of the gas tracers (e.g., H$_2$, [Ne II]) can be produced by shocked emission from outflows as well. High spectral/spatial resolution is required for determine whether the emission is observed at the velocity of the central star and whether it is spatially extended or not. \\item Gas emission could be excited by collisions with dust grains, shocks and UV or X-rays. Line ratios are generally employed to discriminate between the different excitation mechanisms. \\item Gas could be present in a disk even when dust emission is weak or not present. Gas emission lines have been observed in the inner regions of transitional disks and in several WTTS. \\item The study of  gas in disks acquired momentum with the advent of high-resolution ground IR spectrographs and high-sensitivity IR spectrographs in space. \\item A new chapter in gas studies was recently written with the detection of emission of simple molecules in the inner disks. \\end{itemize}  At the beginning of the paper I highlighted some fundamental science questions that required direct observational constraints from the gas in the disk. To conclude, I would like to address them again from the actual state of the observational study of gas.  {\\it (i) How long does the protoplanetary disk last? } In general terms, disk gas diagnostics are not detected when the signatures of the dust in the disk have disappeared. Thus, as a first approximation the gas in the disk disappears almost simultaneously with the signatures from small dust particles. Nonetheless, we should be aware of a few details: (a) with the exception of H$\\alpha$ emission, there are many sources for which we know that they have a disk, but given sensitivity issues, we are not able to detect any gas tracer; (b) gas has been observed in several WTTS stars (e.g., H$_2$ in the UV and NIR), and in transitional disks (e.g., H$_2$ in the near-IR, CO at 4.7 $\\mu$m, [Ne II] at 12 $\\mu$m); (c) unbiased surveys for hot/warm gas in a large sample of WTTS have yet not been performed; (d) surveys of H$\\alpha$ at high spectral resolution need to be done in larger samples of WTTS in star-forming regions of different ages. If the accretion rate and spectral resolution are low, one can easily miss objects accreting gas at low accretion rates. In summary, an estimate of disk gas lifetime independent of dust is still required.  {\\it (ii) How much material is available for forming giant planets? } On this aspect our hopes are fainter. Several surveys for H$_2$ emission in the MIR showed that the emission is observed only in a handful of objects, and that in these objects the gas is heated by an additional (to dust collisions) heating mechanism (UV, X-ray excitation). Although these results confirm our two-layer picture of the disk, in practice these results mean that we are blind to the disk's mid-plane and that we are unable to know directly how much mass is in the giant planet forming region of the disk. The good news are two fold: (a) in objects with data at several wavelengths, an educated guess can be found from modeling the surface density (see next point); (b) at some point the disk should become optically thin -by dust evolution- then we will be able to measure the gas mass.  {\\it (iii) How do the density and temperature of the disk vary as a function of the radius? } There is a growing number of sources (e.g., TW Hya, AB Aur) for which we start to have information from several gas tracers at wavelengths from the UV to the (sub)-mm. In such sources, we can attempt to constrain the disk structure employing a disk model aimed to fit all the available data. Therefore, at least for a handful of objects, by modeling multiwavelength spectroscopy, we have the potential to constrain the density and temperature as a function of radius. However, we should always keep in mind that we observe emission from the disk's surface layer and that the mid-plane temperature and density will be deduced from the model. Hence, the final gas mass deduced will be model dependent.  {\\it (iv) What are the dynamics of the disk? } In this aspect we had good news. To detect gas lines in the IR we commonly use the highest spectral resolution available. At the present, the resolution is sufficiently high to spectrally resolve the lines. This allows the modeling of the line shape; therefore, constraints in the dynamics of the emitting gas can be derived. Recently, with the combination of AO and high-resolution spectroscopy, it has been possible to spatially resolve gas disk emission directly or to spatially resolve the spectroastrometric signal with spectra taken at different slit orientations. So far the observed gas is consistent with gas in Keplerian rotation. This kind of studies should be extended in the future to a larger number of systems."
    },
    "0911/0911.5451_arXiv.txt": {
        "abstract": "We present the results of a search for new members of the Taurus star-forming region using data from the {\\it Spitzer Space Telescope} and the {\\it XMM-Newton Observatory}. We have obtained optical and near-infrared spectra of 44 sources that exhibit red {\\it Spitzer} colors that are indicative of stars with circumstellar disks and 51 candidate young stars that were identified by Scelsi and coworkers using {\\it XMM-Newton}. We also performed spectroscopy on four possible companions to members of Taurus that were reported by Kraus and Hillenbrand.  Through these spectra, we have demonstrated the youth and membership of 41 sources, 10 of which were independently confirmed as young stars by Scelsi and coworkers.  Five of the new Taurus members are likely to be brown dwarfs based on their late spectral types ($>$M6).  One of the brown dwarfs has a spectral type of L0, making it the first known L-type member of Taurus and the least massive known member of the region ($M\\sim4$-7~$M_{\\rm Jup}$). Another brown dwarf exhibits a flat infrared spectral energy distribution, which indicates that it could be in the protostellar class~I stage (star+disk+envelope).  Upon inspection of archival images from various observatories, we find that one of the new young stars has a large edge-on disk ($r=2\\farcs5=350$~AU).  The scattered light from this disk has undergone significant variability on a time scale of days in optical images from the Canada-France-Hawaii Telescope.  Using the updated census of Taurus, we have measured the initial mass function for the fields observed by {\\it XMM-Newton}. The resulting mass function is similar to previous ones that we have reported for Taurus, showing a surplus of stars at spectral types of K7-M1 (0.6-0.8~$M_\\odot$) relative to other nearby star-forming regions like IC~348, Chamaeleon~I, and the Orion Nebula Cluster. ",
        "introduction": "\\label{sec:intro}      The Taurus complex of dark clouds is one of the best sites for studying the formation of stars in a quiescent, relatively isolated environment.  It is among the nearest star-forming regions ($d=140$~pc) and exhibits a very low stellar density ($n\\sim1$-10~pc$^{-3}$).  Although the individual clouds are sparsely populated, the cloud complex as a whole contains more than 300 known members.  Working toward a complete census of Taurus is important for the identification of rare objects (e.g., edge-on disks, transitional disks, protostars) as well as the statistical characterization of the stellar population (e.g., disk fraction, initial mass function, spatial distribution). A variety of methods have been employed in surveys for new members of Taurus \\citep{ken08}. Two of these techniques, mid-infrared (IR) imaging and X-ray imaging, are highly complementary. Mid-IR observations can identify stars that have circumstellar disks and can penetrate the high levels of extinction that surround stars at the earliest evolutionary stages while X-ray data can uncover the diskless members of a young stellar population.    Because of their excellent sensitivities and large fields of view, the {\\it Spitzer Space Telescope} \\citep{wer04} and the {\\it XMM-Newton Observatory} \\citep{jan01} are the best available telescopes for wide-field imaging surveys at mid-IR and X-ray wavelengths, respectively. The unique capabilities of these facilities have been applied to the Taurus star-forming region through the Taurus {\\it Spitzer} Legacy Survey (D. Padgett, in preparation) and the {\\it XMM-Newton} Extended Survey of the Taurus Molecular Cloud \\citep[XEST,][]{gud07}. \\citet{luh06tau2,luh09fu} and \\citet{sce07,sce08} have used the data from these surveys to search for new members of Taurus. We have continued those efforts by performing spectroscopy on IR sources that we have identified in the {\\it Spitzer} images and X-ray sources that were reported by \\citet{sce07}. In this paper, we describe the selection of these candidate members of Taurus (\\S~\\ref{sec:select}) and measure their spectral types with optical and IR spectra (\\S~\\ref{sec:spec}). We then characterize the stellar parameters of the new members and discuss other notable properties of these objects (\\S~\\ref{sec:prop}). Finally, we use our updated census of the stellar population in Taurus to measure the initial mass function (IMF) within the fields observed by XEST (\\S~\\ref{sec:imf}). ",
        "conclusions": "We have presented a survey for new members of the Taurus star-forming region in which we obtained spectra of candidate members appearing in images from the {\\it Spitzer Space Telescope} (46~deg$^{2}$) and the {\\it XMM-Newton Observatory} (5~deg$^{2}$).  Using the mid-IR data from {\\it Spitzer}, we identified 44 sources that could be young stars with disks, 24 of which were confirmed as members by our spectroscopy. We also performed spectroscopy on 51 candidates detected in X-rays by the XEST program \\citep{gud07,sce07}, demonstrating the youth and membership of 16 sources. Ten of these new X-ray members were independently confirmed through spectroscopy by \\citet{sce08}. In addition, to the sources from {\\it Spitzer} and {\\it XMM-Newton}, we observed four candidate companions to known members of Taurus that were found by \\citet{kra07} through analysis of 2MASS data, one of which we have classified as a young star.  Our survey has uncovered several rare types of sources that are valuable for studies of various aspects of star and planet formation. They consist of a wide binary brown dwarf that is forming in isolation \\citep{luh09fu}, the first known L-type member of Taurus ($M\\sim4$-7~$M_{\\rm Jup}$), a highly reddened brown dwarf that may be in the class~I stage (star+disk+envelope), a disk that appears to have an inner hole (i.e., transitional disk), and a large edge-on disk ($r=2\\farcs5=350$~AU).  The companion identified by \\citet{kra07} also may comprise the primary in wide, low-mass binary system (M4.5+M8.5, $a=12\\arcsec=1700$~AU).  Meanwhile, the {\\it Spitzer} and {\\it XMM-Newton} data in conjunction with previous optical and near-IR surveys provide relatively well-defined completeness limits for the current census of Taurus, enabling a better characterization of the stellar population.  For instance, we have estimated the IMF within the fields observed by XEST, arriving at a distribution that reaches a maximum near 0.8~$M_\\odot$, which agrees with our previous measurements for Taurus. Thus, the IMF in Taurus continues to appear anomalous compared to other nearby star-forming clusters, which peak at 0.1-0.2~$M_\\odot$.  The disk fraction in the XEST fields and the spatial distribution of the SED classes are investigated by \\citet{luh09tau}.  The completeness of the census of Taurus remains poorly determined among class~I sources at low masses and class~III sources outside of the XEST fields, which focused on the denser stellar aggregates. Future surveys can address these shortcomings through spectroscopy of red, faint sources detected by {\\it Spitzer} and measurements of variability and proper motions with wide-field, multi-epoch imaging (e.g., Panoramic Survey Telescope and Rapid Response System, Large Synoptic Survey Telescope)."
    },
    "0911/0911.3041_arXiv.txt": {
        "abstract": "Semi-empirical method of calculation of quenching factors for scintillators is described. It is based on classical Birks formula with the total stopping powers for electrons and ions which are calculated with the ESTAR and SRIM codes, respectively. Method has only one fitting parameter (the Birks factor $kB$) which can have different values for the same material in different conditions of measurements and data treatment. A hypothesis is used that, once the $kB$ value is obtained by fitting data for particles of one kind and in some energy region (e.g. for a few MeV $\\alpha$ particles from internal contamination of a detector), it can be applied to calculate quenching factors for particles of another kind and for another energies (e.g. for low energy nuclear recoils) if all data are measured in the same experimental conditions and are treated in the same way. Applicability of the method is demonstrated on many examples including materials with different mechanisms of scintillation: organic scintillators (solid C$_8$H$_8$, and liquid C$_{16}$H$_{18}$, C$_9$H$_{12}$); crystal scintillators (pure CdWO$_4$, PbWO$_4$, ZnWO$_4$, CaWO$_4$, CeF$_3$, and doped CaF$_2$(Eu), CsI(Tl), CsI(Na), NaI(Tl)); liquid noble gases (LXe). Estimations of quenching factors for nuclear recoils are also given for some scintillators where experimental data are absent (CdWO$_4$, PbWO$_4$, CeF$_3$, Bi$_4$Ge$_3$O$_{12}$, LiF, ZnSe). ",
        "introduction": "In accordance with our current understanding of astronomical observations, usual matter constitutes only $\\simeq4$\\% of the Universe; the main components are dark matter ($\\simeq23$\\%) and dark energy ($\\simeq73$\\%) \\cite{Ber05}. Various extensions of the Standard Model propose many candidates on the role of dark matter (DM) particles \\cite{Ste09} which are neutral and only weakly interact with matter (Weakly Interacting Massive Particles, WIMPs). One of the approaches to discover these particles is to detect scattering of WIMPs on atomic nuclei in sensitive detectors placed deep underground and measured in extra low background conditions \\cite{Spo07}. Taking into account likely mass range and velocities of WIMPs, energies of nuclear recoils are expected below $\\simeq100$ keV with character interaction rates of $1-10^{-6}$ events kg$^{-1}$ d$^{-1}$. Many searches of WIMPs with semiconductor, scintillator and bolometer detectors to-date gave only negative results (see \\cite{Spo07} and references therein); instead positive evidence for DM particles (WIMPs are a subclass; other candidates and other kinds of interactions are also available) in the galactic halo has been pointed out by DAMA experiments by exploiting the DM annual modulation signature with NaI(Tl) scintillators during more than 10 years long measurements \\cite{Ber08a}.  For a long time it is known that amount of light produced in scintillating material by highly ionizing particles is lower than that produced by electrons of the same energy \\cite{Bir64}. Thus, in a scintillator calibrated with electron and/or $\\gamma$ sources (which is an usual practice), signals from ions will be seen at lower energies (sometimes up to $\\simeq 40$ times) than their real values. Evidently knowledge of these transformation coefficients -- quenching factors -- is extremely important in searches for WIMPs and in predictions where the WIMPs signal should be expected. Many experimental efforts were devoted to measurements, sometimes very sophisticated, of quenching factors at low energies in different detectors (see e.g. recent works \\cite{Bav08,Cal08,Cha08,Cor08,Kob08,Nik08,Sor09} and further references).  Quenching factors are also needed in measurements and interpretation of signals from $\\alpha$ particles in scintillators. As examples, we can mention here recent experiments on searches (and first observations) of extremely rare $\\alpha$ decays ($T_{1/2}=10^{18}-10^{19}$ yr): $^{180}$W in CdWO$_4$ \\cite{Dan03} and CaWO$_4$ \\cite{Zde05} crystal scintillators, and $^{151}$Eu in CaF$_2$(Eu) \\cite{Bel07}.  While few approaches in calculation of quenching factors are known \\cite{Bir51,Mur61,Lin63,Hit05}, satisfactory theory able to exactly predict (and very often even to describe already measured) quenching factors for all detectors and particles still is absent. For example, in the Lindhard's approach \\cite{Lin63}\\footnote{Good description is given in more accessible source \\cite{Mei08}.} it is possible to calculate quenching of ions with atomic number $Z$ in scintillator only with the same $Z$ number; in addition, this theory predicts decrease of quenching factors at low energies, very often in contradiction with experimental data. Hitachi's model \\cite{Hit05} gives better description and for wider data range, however it is not easy to reproduce these calculations independently.  Below we describe rather simple method of calculation of quenching factors for different ions (from protons to heavy recoils), based on semi-empirical approach of Birks \\cite{Bir51} and using available in Internet software for calculation of stopping powers for electrons and ions (ESTAR \\cite{ESTAR} and SRIM \\cite{SRIM} codes, respectively). It employs only one parameter ($kB$ Birks factor) which could be found by fitting experimental data measured for particles of one kind in some energy region (e.g. for $\\alpha$ particles from external sources or internal contamination of a detector by U/Th chains, $^{147}$Sm, $^{190}$Pt, etc.) but afterwards can be used to calculate quenching factors for other particles and in other energy regions (e.g. for nuclear recoils at low energies). Summary of the method is given in section 2. Calculations with this method are demonstrated in section 3 for number of scintillators: organic scintillators (solid C$_8$H$_8$, and liquid C$_{16}$H$_{18}$, C$_9$H$_{12}$); crystal scintillators (pure CdWO$_4$, PbWO$_4$, ZnWO$_4$, CaWO$_4$, CeF$_3$, and doped CaF$_2$(Eu), CsI(Tl), CsI(Na), NaI(Tl)); liquid noble gases (LXe). Estimations of quenching factors for nuclear recoils are also given for some scintillators where experimental data are absent (CdWO$_4$, PbWO$_4$, CeF$_3$, Bi$_4$Ge$_3$O$_{12}$, LiF, ZnSe). Section 4 gives conclusions. ",
        "conclusions": "Semi-empirical and quite simple in realization method of calculation of quenching factors for scintillators was described in this work. It is based on the classical Birks formula with the {\\em total} stopping powers for electrons and ions, and has only one parameter: the Birks factor $kB$. Value of this factor for a given scintillating material can be different in different conditions of measurements and data treatment. However, if experimental conditions and treatment of data are fixed, hypothesis that $kB$ has the same value for particles of different kinds gives reliable results. Once the $kB$ is found by fitting quenching factors for particles of one kind and in some range of energies (e.g. for $\\alpha$ particles from internal contamination of a detector by U/Th chains and/or by $^{147}$Sm, $^{190}$Pt with energies of a few MeV), it can be used to calculate quenching factors for particles of another kinds and for another energies of interest (e.g. for low energy nuclear recoils). Many examples were given for materials which, furthermore, have different mechanisms of scintillation: organic scintillators (solid C$_8$H$_8$, and liquid C$_{16}$H$_{18}$, C$_9$H$_{12}$); crystal scintillators (pure CdWO$_4$, PbWO$_4$, ZnWO$_4$, CaWO$_4$, CeF$_3$, and doped CaF$_2$(Eu), CsI(Tl), CsI(Na), NaI(Tl)); and liquid noble gases (LXe). It was demonstrated for many cases that the method allows not only to {\\em describe} measured data for ions of one kind in a reliable way but also to {\\em predict} behaviour of quenching factors for other particles which sometimes is immediately confirmed by already existing experimental data -- sometimes worse, sometimes better, and sometimes very good, but at least in a rough agreement. Some predictions (e.g. for LNe, LiF and others) could be checked in near future.  Stopping powers for electrons and ions are calculated with the ESTAR and SRIM codes, respectively, which in fact present to-date state-of-art software in this field. It is easy to use these programs and they are publicly available; this makes $Q_i$ calculations quite simple.  Calculations with the SRIM package have some tendency to overestimate quenching factors for $\\alpha$ particles at energies around $\\simeq2$ MeV and underestimate them at high energies ($>8$ MeV) as can be seen in Fig. 4a for CdWO$_4$, Fig. 5a for CaF$_2$(Eu), Fig. 8 for CaWO$_4$, and Fig. 15 for CeF$_3$. At the same time, calculation of the stopping powers for $\\alpha$ particles with the ASTAR package gave better description of $Q_\\alpha$ in CaF$_2$(Eu) scintillator at lower energies (see Fig. 5a). For some other materials difference between ASTAR and SRIM calculations was not big (see Fig. 2 for CsI(Tl) and Fig. 3a for C$_8$H$_8$). Evidently $Q_i$ values will depend on how one calculates stopping powers for ions and electrons, and it is a pity that stopping powers could not be computed in framework of the same package for any particle (SRIM calculates SP for ions in any substance but does not calculate SP for electrons; and STAR gives SP for electrons in any material but SP for ions are possible only for protons and $\\alpha$ particles and for a limited list of materials).  Quenching factors calculated in the presented approach in general increase at low energies, and this encourages experimental searches for dark matter particles. Estimations of quenching factors for nuclear recoils are given for some scintillators where experimental data are absent (CdWO$_4$, PbWO$_4$, CeF$_3$, Bi$_4$Ge$_3$O$_{12}$, LiF, ZnSe)."
    },
    "0911/0911.4457_arXiv.txt": {
        "abstract": "{Supernova remnants are believed to be a major source of energetic particles (\\emph{cosmic rays}) on the Galactic scale. Since their progenitors, namely the most massive stars, are commonly found clustered in \\emph{OB~associations}, one has to consider the possibility of collective effects in the acceleration process.} {We investigate the shape of the spectrum of high-energy protons produced inside the \\emph{superbubbles} blown around clusters of massive stars.} {We embed simple semi-analytical models of particle acceleration and transport inside Monte Carlo simulations of OB~associations timelines. We consider regular acceleration (Fermi~1 process) at the shock front of supernova remnants, as well as stochastic reacceleration (Fermi~2 process) and escape (controlled by magnetic turbulence) occurring between the shocks. In this first attempt, we limit ourselves to linear acceleration by strong shocks and neglect proton energy losses.} {We observe that particle spectra, although highly variable, have a distinctive shape because of the competition between acceleration and escape: they are harder at the lowest energies (index $s<4$) and softer at the highest energies ($s>4$). The momentum at which this spectral break occurs depends on the various bubble parameters, but all their effects can be summarized by a single dimensionless parameter, which we evaluate for a selection of massive star regions in the Galaxy and the LMC.} {The behaviour of a superbubble in terms of particle acceleration critically depends on the magnetic turbulence: if~B is low then the superbubble is simply the host of a collection of individual supernovae shocks, but if~B is high enough (and the turbulence index is not too high), then the superbubble acts as a global accelerator, producing distinctive spectra, that are potentially very hard over a wide range of energies, which has important implications on the high-energy emission from these objects.} ",
        "introduction": "\\label{sec:introduction}  Superbubbles are hot and tenuous large structures that are formed around OB~associations by the powerful winds and the explosions of massive stars \\citep{Higdon2005a}. They are the major hosts of supernovae in the Galaxy, and thus major candidates for the production of energetic particles (e.g., \\citealt{Montmerle1979a}, \\citealt{Bykov2001b}, \\citealt{Butt2009a}, and references therein). Supernovae are indeed believed to be the main contributors of Galactic cosmic rays (along with pulsars and micro-quasars), by means of the \\emph{diffusive shock acceleration} process (a 1st-order, regular Fermi process) occurring at the remnant's blast wave as it goes through the interstellar medium \\citep{Drury1983a,Malkov2001c}.  Supernovae in superbubbles are correlated in space and time, hence the need to investigate acceleration by multiple shocks \\citep{Parizot2004a}. \\citet{Klepach2000a} developed a semi-analytical model of test-particle acceleration by multiple spherical shocks (either wind termination shocks, or supernova shocks plus wind external shocks), based on the limiting assumption of small shocks filling factors. \\citet{Ferrand2008a} performed direct numerical simulations of repeated acceleration by successive planar shocks in the non-linear regime (that is, taking into account the back-reaction of energetic particles on the shocks). However, to ascertain the particle spectrum produced inside the superbubble as a whole, one must also consider important physics occurring \\emph{between} the shocks. Since the bubble interior is probably magnetized and turbulent, we need to evaluate gains and losses caused by the acceleration by waves (a 2nd-order, stochastic Fermi process) and escape from the bubble.  In this study, we combine the effects of regular acceleration (occurring quite discreetly, at shock fronts) and stochastic acceleration and escape (occurring continuously, between shocks), to determine the typical spectra that we can expect inside superbubbles over the lifetime of an OB~cluster. We choose to treat regular acceleration as simply as we can, and concentrate on modeling the relevant scales of stochastic acceleration and escape inside superbubbles. We present our model in Sect.~\\ref{sec:model}, give our general results in Sect.~\\ref{sec:results}, and present specific applications in Sect.~\\ref{sec:application}. Finally we discuss the limitations of our approach in Sect.~\\ref{sec:limitations} and provide our conclusions in Sect.~\\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  Our main conclusions are as follows: \\begin{enumerate} \\item Cosmic-ray spectra inside superbubbles are highly variable: at a given time they depend on the particular history of a given cluster. \\item Nevertheless, spectra follow a distinctive overall trend, produced by a competition between (re-)acceleration by regular and stochastic Fermi processes and escape: they are harder at lower energies ($s<4$) and softer at higher energies ($s>4$), shapes that are in agreement with the results of~\\citet{Bykov2001b} based on different assumptions\\footnote{ \\cite{Bykov2001b} considers acceleration of particles by large-scale motions of the magnetized plasma inside the superbubble, which depends on the ratio $D_u/D_x$ where $D_x$ is the space diffusion coefficient, controlled by magnetic fluctuations at small scales, and $D_u=UL$ describes the effect of large scale turbulence, where $U$ is the average turbulent speed and $L$ is the average size between turbulence sources.}. \\item The momentum at which this spectral break occurs critically depends on the bubble parameters: it increases when the magnetic field value and acceleration region size increase, and decreases when the density and the turbulence external scale increase, all these effects being summarized by the single dimensionless parameter $\\theta^{\\star}$ defined by Eq.~(\\ref{eq:Theta_def}). \\item For reasonable values of superbubble parameters, very hard spectra ($s<3$) can be obtained over a wide range of energies, provided that superbubbles are highly magnetized and turbulent (which is a debated issue). \\end{enumerate}  These results have important implications for the chemistry inside superbubbles and the high-energy emission from these objects. For instance, in the superbubble Perseus~OB2 there is observational evidence of intense spallation activity \\citep{Knauth2000a} attributed to a high density of low-energy cosmic rays, but EGRET has not detected $\\pi^{0}$-decay radiation, which places strong limits on the density of high-energy cosmic rays. This is consistent with the shape of the spectra obtained in this work. We are thus looking forward to seeing how new instruments such as Fermi and AGILE will perform on extended sources such as massive star forming regions, which have recently been established as very high-energy sources. In that respect, we make a final comment that the high intermittency of predicted spectra might explain the puzzling fact that some objects are detected while others remain unseen."
    },
    "0911/0911.1791_arXiv.txt": {
        "abstract": "We present results from the first systematic search for outlying HII regions, as part of a sample of 96  emission-line point sources (referred to as ELdots - emission-line dots) derived from the NOAO Survey for Ionization in Neutral Gas Galaxies (SINGG). Our automated ELdot-finder searches SINGG narrow-band and continuum images for high equivalent width point sources outside the optical radius of the target galaxy ($>$ 2 $\\times$ r$_{25}$ in the R-band). Follow-up longslit spectroscopy and deep  GALEX images (exposure time $>$ 1000 s) distinguish outlying HII regions from background galaxies whose strong emission lines ([OIII], H$\\beta$ or [OII]) have been redshifted into the SINGG bandpass.  We find that these deep GALEX images can serve as a substitute for spectroscopic follow-up because outlying HII regions separate cleanly from background galaxies in color-color space. We identify seven SINGG systems with outlying massive star formation that span a large range in H$\\alpha$ luminosities corresponding to a few O stars in the most nearby cases, and unresolved dwarf satellite companion galaxies in the most distant cases. Six of these seven systems feature galaxies with nearby companions or interacting galaxies. Furthermore, our results indicate that some outlying HII regions are linked to the extended-UV disks discovered by GALEX, representing emission from the most massive O stars among a more abundant population of lower mass (or older) star clusters.  The overall frequency of  outlying HII regions in this sample of gas-rich galaxies is 8 - 11\\% when we correct for background emission-line galaxy contamination ($\\sim75$\\% of ELdots). ",
        "introduction": "\\label{intro} In the last 10 years, deep H$\\alpha$ imaging has confirmed the occurrence of recent star formation far beyond the main optical bodies of  galaxies, in both discrete knots with no underlying continuum emission and faint outer spiral arms (eg. \\nocite{ferguson98, emma,bfq97} Bland-Hawthorn, Freeman, \\& Quinn 1997; Ferguson et al. 1998; Ryan-Weber et al. 2004).  More recently, the {\\it{Galaxy Evolution Explorer}} (GALEX) satellite has broadened our picture of outer-galaxy star formation by revealing that over $\\sim30$\\% of spiral galaxies possess UV-bright extensions of their optical disks, implying that low-intensity star formation in the outskirts of galaxies is not particularly rare \\citep{Thilker07}.  This mode of star formation in the outer disks of galaxies may be a natural extension of inside-out disk formation, in which galactic stellar disks grow gradually with time following the condensation and assembly of their gaseous disks (\\nocite{whiteandfrenk91, bouwens97} e.g. White \\& Frenk 1991; Bouwens, Cayon, \\& Silk 1997).  Moreover, extended-UV (XUV) emission and outlying HII regions sometimes are associated with previous or ongoing galaxy interactions \\citep{Thilker07, werk08}.  Probing the nature and frequency of outer-galaxy star formation can facilitate our understanding of the disk building process, and the ways in which new stellar populations emerge via galaxy interactions.  Star formation in galactic outskirts also has relevance to discussions involving star formation gas density thresholds, as it often occurs in a low-density environment (e.g. \\nocite{gerhard,mendes04, emma}  Gerhard et al. 2002; Mendes de Oliveira et al. 2004; Ryan-Weber et al. 2004). One way to increase our understanding of star formation ``thresholds\" is to assemble a sample of massive stars forming in low column-density gas outside the influence of the existing stellar population in a galactic disk.    The \\ha~ and UV radiation in the extended disks trace different ranges of stellar mass: UV radiation follows primarily the O and B stars, while  \\ha~ emission traces just the O stars. Studies of spiral galaxies with XUV disks \\citep{thilker05, gildepaz05} have indicated that the spatial extent of star formation in outer disks is underestimated by looking for HII regions alone (via \\ha~ imaging), as \\ha~emission tends to be even less widespread than predicted by typical stellar population synthesis models that incorporate a power-law IMF with a slope of $\\alpha\\sim2.35$. The reasons for this relative lack of \\ha~emission at large galactocentric radii could be many, including but not limited to stochastic fluctuations in the IMF at low stellar luminosities \\citep{boissier06}, a top-light IMF in the remotest reaches of galaxies \\citep{meurer09}, or a large fraction of escaping ionizing photons from the less dense outer-galactic regions.  The extent of this underestimation, specifically at large projected galactocentric distances, remains unquantified, as no comprehensive systematic study of \\ha~emission in the outer reaches of galaxies has yet emerged.   Here, we return to  \\ha~imaging to find star formation in the outskirts of galaxies. As a star formation tracer that is sensitive to  only the highest mass stars, \\ha~provides an important complement to the recent GALEX results.  We perform an automated search to find outlying compact sources of net line emission using the imaging data from the Survey for Ionization in Neutral Gas Galaxies (SINGG; Meurer et al 2006; hereafter M06\\nocite{SINGG}). The large angular area outside the optical radius of galaxies in each 14.7\\'~SINGG field presents the opportunity to search for outlying HII regions at large radii and perform a blind search for background emission-line sources. We initially refer to both types of  sources as ``ELdots\" for their appearance as emission line dots in the images.  In all cases, ELdots exhibit strong emission lines in the SINGG narrowband filter, and are point sources well outside the broadband optical emission of nearby galaxies (r$>$2 $\\times$ r$_{25}$). ELdots can be outlying \\ha-emitting HII regions at a similar velocity to the HIPASS source galaxy or background  galaxies emitting a different line ([OIII], [OII], or H$\\beta$) that is redshifted into the narrow filter passband used for the observations. To distinguish between these options, we present follow-up spectroscopy and deep archival GALEX images for a subsample of ELdots.  In this study, we are primarily interested in the \\ha-emitting ELdots, although we tabulate the properties of the background galaxies as well. While a number of previous works have discovered ``intergalactic\" HII regions, this is the first systematic search of an unbiased sample of gas rich galaxies for such emission. Furthermore, although many studies have chosen ``intergalactic\" as a modifier (see \\nocite{boquien09} Boquien et al. 2009 for a detailed description of the terminology), we opt instead for ``outlying.\" These HII regions are distinct from the central galactic star formation, and more sparse, yet they often lie in extended neutral gas and/or an extended UV component associated with the galaxy. Outlying HII regions (abbreviated outer-HIIs) provide a unique laboratory to understand star formation under relatively extreme conditions (e.g. low neutral gas column density, weak galaxy potential), and may shed light on the full extent of stellar disks \\citep{blandhawthorn05, irwin05}.  The paper proceeds as follows: in Section 2, we describe the SINGG sample; in Section 3, we describe our ELdot sample selection; in Section 4, we present spectroscopic observations of a subsample of ELdots; in Section 5 we examine the ELdots in deep GALEX images, where available from the archives; in Section 6, we discuss the properties of our ELdot sample;  in Section 7, we discuss extended star formation in SINGG; and in Section 8 we summarize and present the key findings of this study. ",
        "conclusions": "\\label{conc}  ELdots are emission-line point sources (``dots\") well outside the main optical R-band emission of a galaxy, defined as 2 $\\times$ r$_{25}$ in this work. Using an automated finder, we catalogue a total of 96 ELdots in 50 of 89 SINGG systems.   We classify ELdots as either outlying HII regions (outer-HIIs) or background galaxies by spectroscopy and GALEX morphology and colors. This study highlights four key results: \\begin{itemize}  \\item{Follow-up GALEX data, combined with SINGG photometric properties, can help distinguish between outer-HIIs and background galaxy ELdots, nearly eliminating the need for spectroscopic follow-up. The distribution of ELdot FUV$-$R colors is bimodal, revealing that outer-HIIs are considerably bluer than higher-redshift emission-line background galaxies. The very blue outer-HII colors are due primarily to the extremely low R-band emission from these objects, indicating new stars forming where there is little or no underlying older stellar population. Outer-HIIs also tend to have higher emission-line EWs than background galaxy ELdots and are likely to be found closer to a potential host galaxy than the background galaxies, which are uniformly distributed across the SINGG search area.}  \\item{Through our systematic search of a sample of gas-rich galaxies, we have confirmed that massive star formation in the far outskirts of galaxies tends to be associated with galaxy interactions and nearby companions. This result agrees with other recent findings of star formation outside the main optical bodies of galaxies. }   \\item{ In the cases where long-exposure GALEX data is available, we find that outer-HIIs appear to be associated with XUV emission, specifically of Type-1 morphology.  }   \\item{ We find that the frequency of massive star formation in the far outskirts of galaxies, as traced by \\ha~emission in the full SR1 sample, is between 8 and 11\\%. Outer-HII regions in this full sample span a wide range of H$\\alpha$ luminosities, from 10$^{36.6}$ ergs s$^{-1}$ to 10$^{38.5}$ ergs s$^{-1}$, and have physical sizes from 70 pc to 500 pc, representing OB associations ionized by a few O stars in the most nearby systems, and dwarf galaxies with hundreds of O stars in the most distant systems. When we isolate our sample of outer-HIIs to a nearby sample with D $<$ 30 Mpc, we can compare our frequency statistics more directly with those of XUV galaxies. These outer-HIIs are more homogeneous in nature, similar in size and luminosity to Galactic OB associations. Where \\cite{Thilker07} find XUV emission to be present in $\\sim30$\\% of spiral galaxies in the local universe, we find the frequency of outlying HII regions in gas-rich galaxies more nearby than 30 Mpc to be between 6 and 10\\%. }    \\end{itemize}"
    },
    "0911/0911.4382_arXiv.txt": {
        "abstract": "We present a direct $N$-body simulation modeling the evolution of the old (7 Gyr) open cluster NGC 188.  This is the first $N$-body open cluster simulation whose initial binary population is directly defined by observations of a specific open cluster: M35 (150 Myr).  We compare the simulated color-magnitude diagram at 7 Gyr to that of NGC 188, and discuss the blue stragglers produced in the simulation.  We compare the solar-type main sequence binary period and eccentricity distributions of the simulation to detailed observations of similar binaries in NGC 188. These results demonstrate the importance of detailed observations in guiding $N$-body open cluster simulations.  Finally, we discuss the implications of a few discrepancies between the NGC 188 model and observations and suggest a few methods for bringing $N$-body open cluster simulations into better agreement with observations. ",
        "introduction": "\\label{intro}  Recently, sophisticated $N$-body simulations have incorporated stellar dynamics and stellar evolution self consistently, gaining the ability to accurately model open cluster sized populations ($N \\sim 10^4$) of single and binary stars through many billions of years. (e.g. \\texttt{NBODY6}, \\citealt{aarseth:03}, with stellar and binary evolution included by \\citealt{hurley:00,hurley:02}).  Concurrently, large observational surveys of open clusters, like the WIYN Open Cluster Study \\citep[WOCS;][]{mathieu:00}, are coming to maturity, providing comprehensive databases of open cluster characteristics, including detailed information on their binary populations.  These observations, and specifically those of the binaries, provide important tests and guidance for $N$-body simulations. In this paper, we describe our first step in utilizing this wealth of data to help guide an $N$-body simulation with the aim of creating a realistic model of NGC 188. ",
        "conclusions": "\\label{disc}  By utilizing the observed M35 binaries as a proxy for our primordial binary population, we have managed to nearly reproduce the NGC 188 binaries within our $N$-body simulation at 7 Gyr. However, as noted in Section~\\ref{comp188} the current NGC 188 model is deficient in BSs.  Interestingly, the \\citet{hurley:05} model of M67 was able to produce the large number of observed BSs (though not their binary frequency), but was unable to reproduce the observed binary population.  The quantity of BSs produced in the M67 model was largely dependant on the initial binary population.  Hurley et al.~found that a large fraction of short-period binaries was necessary to create the observed numbers of BSs in M67 (and NGC 188). These short-period binaries not only have a higher likelihood of creating BSs in isolation (e.g., through mass transfer resulting from Roche lobe overflow), but dynamical interactions involving short-period binaries have a higher probability of resulting in collisions and mergers \\citep{fregeau:04}. Observations of M35 binaries are very different from the flat period distribution and large binary frequency used in the M67 model (e.g., Figure~\\ref{M35fig}). The important next step will be to match both the binary population and the BSs population within one cohesive simulation. We are currently exploring a number of possible methods for addressing this issue, the most promising of which appears to be the inclusion of primordial triples (e.g., Kozai induced mergers; \\citealt{perets:09}, larger cross section for dynamical interactions, etc.).  \\begin{figure}% \\begin{center} \\includegraphics[width=0.8\\linewidth]{Geller_f4.eps} \\caption{\\label{NGC188fig3} Eccentricity plotted against the logarithm of the period ($e$ - log $P$) at 7 Gyr for the observed (top) and simulated (bottom) solar-type MS NGC 188 binaries.  The solid gray lines represent the best-fit circularization functions from \\citet{meibom:05} for each sample, respectively.  The tidal processes in the simulation result in a circularization period that is significantly lower than in NGC 188.  Also, the simulation produces a population of circular binaries with $P >> P_{circ}$ that is not observed in NGC 188 (or the Galactic field).  Most of these circular binaries in the model have previously gone through a stage of mass transfer.} \\end{center} \\end{figure}  The tidal circularization period is one of our best observational tools to study the effects of tides on a binary population.  In Figure~\\ref{NGC188fig3} and Section~\\ref{comp188} we show that the circularization period at 7 Gyr in the simulation is significantly lower than what is observed in NGC 188.  This suggests that the tidal energy dissipation rate may be \\textit{underestimated} in the $N$-body model (at least for solar-type MS binaries).  We may also (or alternatively) require a form of pre-MS tidal circularization \\citep[e.g.][]{kroupa:95}.  As pointed out by \\citet{mathieu:08}, tidal processes have a significant impact on dynamical interactions between single and binary stars, and the energy dissipation rates can be the deciding factor between a close fly-by, tidal capture, or even a merger or collision that could result in a BS.  WOCS data \\citep[e.g.,][]{meibom:05} provide the ideal foundation from which we will improve the tidal physics within the $N$-body model.  Finally, the simulation has an excess of circular binaries with periods beyond the circularization period. The majority of these binaries have undergone some form of mass transfer, which has quickly circularized the orbit within the simulation. However, similar populations of circular binaries at $P >> P_{circ}$ are not observed in open clusters or in the Galactic field \\citep[e.g.,][]{meibom:05,duquennoy:91}.  This may suggest that mass transfer, especially in initially eccentric binaries, does not necessarily lead to circular binaries \\citep[e.g.,][]{sepinsky:07,bonacic:08}, or that these cases of mass transfer happen much more frequently within the simulation than in reality.  Correctly modeling the products of mass transfer and common-envelope evolution is critical for accurately reproducing BSs as well as the binary population within an $N$-body simulation.  We will present further details about this NGC 188 model in future papers.  Additionally, we will explore the various methods mentioned above for bringing the NGC 188 model, and $N$-body simulations in general, in better agreement with observations.  This synergy between observations and simulations will continue to refine the $N$-body method and reveal further insights into the origins of BSs, the evolution of a binary population, and the dynamical evolution of star clusters.  \\vspace{1em} \\footnotesize \\flushleft This work was funded in part by a National Science Foundation (NSF) East Asia and Pacific Summer Institutes (EAPSI) fellowship, the Wisconsin Space Grant Consortium, and NSF grant AST-0406615. \\vspace{-1em}"
    },
    "0911/0911.0381_arXiv.txt": {
        "abstract": "{We report  the results  of an  optical campaign carried  out by  the XMM-Newton Survey Science Centre with the  specific goal of identifying the brightest X-ray sources in  the XMM-Newton  Galactic Plane  Survey of Hands  et al.  (2004).  In addition to  photometric and spectroscopic observations obtained  at the ESO-VLT and  ESO-3.6m, we  used cross-correlations  with the  2XMMi, USNO-B1.0,  2MASS and GLIMPSE  catalogues  to progress  the  identification  process.  Active  coronae account for 16 of the 30  positively or tentatively identified X-ray sources and exhibit the  softest X-ray spectra.  Many  of the identified  hard X-ray sources are  associated with  massive stars,  possibly in binary systems and emitting at  intermediate X-ray luminosities  of 10$^{32-34}$\\ergs.   Among  these are  a  very absorbed  likely hyper-luminous star with X-ray/optical  spectra and luminosities comparable with those of $\\eta$  Carina, a new X-ray selected WN8 Wolf-Rayet  star in which most of the  X-ray emission probably  arises from wind  collision in a binary,  a new Be/X-ray  star belonging  to the  growing class  of $\\gamma$-Cas  analogs  and a possible supergiant  X-ray binary of  the kind discovered recently  by INTEGRAL. One of  the sources,  XGPS-25, has a  counterpart of moderate  optical luminosity which exhibits He{\\sc II} $\\lambda$4686 and Bowen C{\\sc III}-N{\\sc III} emission lines suggesting that this may be  a quiescent or  X-ray shielded Low Mass X-ray Binary, although its X-ray properties might also be consistent with a rare kind of cataclysmic variable (CV).   We also report the  discovery of three  new CVs, one  of which  is a  likely  magnetic system displaying strong  X-ray  variability.   The  soft  (0.4--2.0\\,keV)  band  \\lnls\\  curve  is completely dominated by active stars in  the flux range of 1$\\times$10$^{-13}$ to 1$\\times$10$^{-14}$\\ergscm . Several active coronae are also detected above 2\\,keV  suggesting  that  the   population  of  RS  CVn  binaries  significantly contributes  to the  hard X-ray  source  population. In  total, we  are able  to identify a large fraction of the hard (2--10\\,keV) X-ray sources in the flux range of 1$\\times$10$^{-12}$ to 1$\\times$10$^{-13}$\\ergscm\\ with Galactic objects  at a rate consistent  with that expected for the Galactic contribution  only.} ",
        "introduction": "The brightest X-ray sources discovered by the early collimator based experiments in   operation  in  the   1970s  were   very  luminous   ($\\sim$  10$^{38}$\\ergs ) accretion-powered  X-ray binaries located  in our  Galaxy \\citep[see e.g. the fourth Uhuru catalogue of X-ray sources;][] {forman1978}.  The dramatic improvement  in  sensitivity  and   spatial  resolution  afforded  by  the  next generation of instruments exploiting X-ray imaging revealed a Galactic landscape with much larger numbers of lower  luminosity systems powered by a wide range of physical processes.   For example, observations  first with Einstein,  and later with ROSAT, showed  that low Galactic latitude regions are  crowded with a large number  of low-luminosity sources  in  the X-ray  bands explored  by these  two missions (0.5--4.5\\,keV  and 0.2--2.4\\,keV  for Einstein and  ROSAT respectively). These  soft sources  are predominantly  identified with  active  stellar coronae \\citep[e.g.][]{hertz1984,motch1997}.  The  first imaging hard  (2--10\\,keV) X-ray survey  of  the  Galactic  Plane  was  performed  a  few  years  later  with ASCA \\citep{sugizaki2001}.   ASCA  revealed  the  presence of  a  genuinely  Galactic population of  low-luminosity non-coronal  X-ray sources which emit  hard X-rays and   are   therefore  detectable   through   large   columns  of   interstellar absorption. However, the still relatively large positional errors affecting ASCA sources combined with the effects of the strong interstellar extinction in the Galactic Plane prevented their optical identification  in most cases. The prospects for building a detailed  census of the low to  intermediate X-ray luminosity sources has  now  greatly  improved  with  the  launch of  the  Chandra  and  XMM-Newton observatories. Thanks to their unprecedented sensitivity and spatial resolution, these observatories are  able to study the properties of  the faint X-ray source populations  throughout the  Galaxy and  most  notably those  present in  large numbers in the  central parts of the Galaxy \\citep{wang2002,muno2003}.   The fields-of-view of the Chandra and  XMM-Newton X-ray cameras  are such that many  serendipitous sources are  seen in  low-latitude fields  in addition  to the  primary targets  and the preliminary results of  such surveys have already emerged in  the context of the Champlane and XMM-Newton SSC surveys \\citep{grindlay2003,motch2003,hands2004,grindlay2005,motch2006}. A number of dedicated Galactic surveys involving either single deep observations or a mosaic of shallower  exposures have also been carried out or are in progress, \\citep[e..g. the Galactic Centre region,][]{muno2003,wijnands2006,koenig2008,revnivtsev2009,muno2009}.  The nature of the intermediate- to low-luminosity Galactic hard X-ray sources discovered in the Galactic Plane surveys quoted above is only partially known. Cataclysmic variables (CVs) and quiescent X-ray binaries (both low- and high-mass types) account for a significant fraction, however, the exact nature of the majority of the sources detected in the central regions of the Galaxy remains doubtful in the absence of optical or infrared identifications,  which in most cases are hindered by the large interstellar absorption.  The observation  of intermediate-  to low-luminosity Galactic X-ray sources can address many issues related to binary-star evolution and to accretion physics at low rates. In addition surveys which extend the search boundaries have the potential to improve our knowledge of source population statistics and to unveil rare types of compact X-ray-emitting systems which have a low space density in the local Galactic neighbourhood. For instance,  the determination of the space  density and  scale  height of CVs is relevant to estimates of the Galactic novae rate and connects to the origin of low-mass X-ray  binaries and  type Ia  supernovae. Low  X-ray luminosity  but  long lived evolutionary stages of classical low and  high mass X-ray binaries could also be found  in  relatively large  numbers  if  the  predictions of  some  evolutionary scenarios are  correct  \\citep[see e.g.][]{pfahl2002,willems2003}. Massive X-ray  binaries are good proxies  of relatively recent star-bursts and if detected in sufficient numbers might give insight into the recent star-formation history of specific regions of our Galaxy.  Moreover, no  X-ray telescope  planned in the  foreseeable future will  have the necessary  spatial   resolution  and  sensitivity   to  study  in   detail  the distribution  of  these low-  and  intermediate-luminosity  sources in  external galaxies.  Consequently,  X-ray surveys of our  own Galaxy and its Magellanic satellites will remain for  a long time the only means to  study how the various X-ray emitting  populations relate to the  main Galactic structures  such as the thin disc, the thick disc and the bulge.  However, such  studies presuppose the knowledge  of the nature of  the bulk of the X-ray sources  surveyed, a prerequisite which can  only be fulfilled through multi-wavelength spectroscopic identifications. Nevertheless, the statistically controlled cross-correlation with large catalogues such as  the 2MASS catalogue and/or with the recently made available  source lists from the GLIMPSE surveys, provide useful information  on the  intensity and  shape  of the  part of  the spectral  energy distribution in the observing windows least affected by line-of-sight absorption in the interstellar medium.  The work  presented here was  carried out in  the framework of a  Galactic Plane Survey programme instigated by the  XMM-Newton Survey Science Centre (SSC).  The long term  objective of  this project  is to gather  a representative  sample of identified  low-latitude X-ray  sources,  which  eventually could  be  used as  a template  for identifying  and classifying  in a  statistical manner  the entire catalogue  of  serendipitous  XMM-Newton  sources  detected  in  the  Milky  Way region.  Companion programs  are being  carried out  at high  Galactic latitudes using three separate samples of sources selected at faint \\citep{furusawa2008}, medium \\citep{barcons2002} and bright \\citep{dellacecca2004} fluxes.  In the first part of this paper, we recall the main properties of the  \\xmm Galactic Plane Survey (XGPS) of \\cite{hands2004} and describe the way a sub-sample of X-ray sources were selected for optical follow-up at the telescope. The following section explains the method applied to cross-correlate the XGPS source list with the USNO-B1.0 and 2MASS catalogues and how identification probabilities were assigned to each optical or near infrared candidate. After a description of the observing procedures, the next section presents in detail the results of our optical spectroscopic identification campaign. The last sections discuss the \\lnls\\ curves and the statistics of the various kinds of X-ray emitters identified in the soft and hard bands as well as their contributions to the genuine Galactic population of low to medium X-ray luminosity sources. ",
        "conclusions": "We have successfully identified at the telescope a large number of the brightest X-ray sources detected in the XGPS survey. In addition, we were able to increase the number of identifications by cross-correlating X-ray positions with the USNO-B1.0 and 2MASS catalogues thanks to a carefully calibrated statistical method. The bulk of these statistical identifications with relatively bright optical or near infrared objects are likely stellar coronae based on their soft X-ray spectra. However, our optical campaign has demonstrated that some of them, in particular those exhibiting the hardest X-ray energy distributions, might be massive stars exhibiting particularly hard X-ray emission. Our optical spectroscopic follow-up observations have also confirmed the variety of the astrophysical objects contributing to the low to medium X-ray luminosity population of Galactic X-ray sources. We spectroscopically identify a total of 16 solar-type active stars, mostly among the softest X-ray sources with a high fraction of late Me stars. Furthermore, two early B type stars appear as the probable counterparts of some of our bright XGPS sources. We have discovered a total of three new cataclysmic variables with optical magnitudes in the range of $\\sim$ 23-22 and 0.5--10\\,keV fluxes from 0.7 to 4.7$\\times$10$^{-13}$\\ergscm . Four sources can qualify as massive X-ray binary candidates on the basis of hard X-ray emission and of the presence in the error circle of an early type star exhibiting \\Halp\\ emission. The optical counterpart of XGPS-3 is likely to be an hyper-luminous star possibly as bright as $\\eta$ Carinae. We also report the discovery of an X-ray selected Wolf-Rayet star. The brightest source of our sample, XGPS-1 is a known transient Low-Mass X-ray Binary exhibiting a bursting behaviour. The bright optical magnitude and unusual Bowen C{\\sc III}-N{\\sc III} emission of XGPS-25 might be the signature of a Low-Lass X-ray Binary in quiescence, although the observational picture could also be consistent with a rare type of CV. The main findings of our work can be summarized as follows:  \\begin{enumerate}  \\item In the 2--10\\,keV flux range from 10$^{-12}$ to 10$^{-13}$\\ergscm, we positively identify a large fraction of the hard X-ray sources  with Galactic  objects  at a  rate  consistent with  that expected for the Galactic contribution only. We note, however, that the  diversity of the observed X-ray spectra makes the count to flux conversion somewhat uncertain. Moreover, different normalisations of the extragalactic \\lnls\\ curves are proposed in the literature. Taken together, these two issues do not allow a precise estimate to be made of the fraction of Galactic sources. However, based on hardness ratio considerations, we find that about half of the hard XGPS sources are likely Galactic objects, in rough agreement with the number estimated by subtracting the extragalactic ASCA \\lnls\\ curve from the relation obtained using the whole XGPS sample. At the faintest hard X-ray flux considered in our study of \\Fx\\ = 5$\\times$10$^{-14}$\\ergscm , we still identify spectroscopically or statistically between 30\\% and 50\\% of the 'Galactic fraction' of the hard X-ray sources. Since most of the identifications at the faint flux end rely on coincidence in position with a bright optical or infrared candidate, the nature of this optically bright population remains somewhat uncertain. However, drawing on the spectroscopic identifications of the optically bright X-ray brightest objects, we expect them for the most part to be made of particularly active late type stars and of early type stars.  \\item The soft X-ray band is completely dominated by coronally emitting stars with CVs contributing only a small percentage of the identifications. We identify all sources down to \\Fx$\\sim$3$\\times$10$^{-14}$\\ergscm . With a mean integrated Galactic \\nh\\ of $\\sim$ 9$\\times$10$ ^{22}$ cm$^{-2}$ the background of extragalactic sources is completely screened out by line-of-sight absorption. Although the small number statistics does not allow us to draw robust conclusions, it seems that a larger number of faint Me stars is present in the XGPS survey than was identified in the ROSAT Galactic Plane Survey. Such an evolution is in rough agreement with the prediction of X-ray stellar population models.  \\item A few relatively bright active coronae do contribute significantly above 2\\,keV. This probably indicates unusually high temperatures for the hottest of the two thin thermal components generally needed to represent the X-ray spectra of active stars \\citep[see a recent review in][]{guedel2009}. The observation of open clusters and field stars of different ages has established clear relations between stellar age, X-ray luminosity and temperature. For instance, the young stars in Orion (age $<$ 1\\,Myr) require on average $kT_1\\,\\sim$ 0.8\\,keV and  $kT_2\\,\\sim$ 2.9\\,keV, while those of the $\\sim$ 160\\,Myr old stars in NGC 2516 have already dropped to $kT_1\\,\\sim$ 0.5\\,keV and $kT_2\\,\\sim$ 1.7\\,keV  \\citep[see][ and references therein]{sung2008}. In contrast, the population of X-ray bright stars identified in the high galactic latitude XMM-Newton Bright Serendipitous Survey by \\cite{xbss2007} is well characterized by $kT_1\\,\\sim$ 0.3\\,keV and $kT_2\\,\\sim$ 1.0\\,keV, indicating a significantly older age on average. The hard X-ray coronae unveiled in the XGPS are thus likely to be very young stars in agreement with the predictions of X-ray stellar populations models. However, a fraction of these hard X-ray emitting stars may also be active binaries, mainly of the RS CVn type. Close binaries in which rotation and coronal activity is maintained by tidal synchronisation do amount to about one third of the ROSAT/Tycho-2 sample of stars \\citep{guillout2009}.  \\item The hard X-ray source population has a strong component made of  early-type  stars   exhibiting  enhanced  emission   from  a   hard  X-ray component. The physical mechanism responsible for the hard X-ray excess is unclear but could be related to wind collisions in the case of the Wolf-Rayet star counterpart of XGPS-14 and for the likely hyper-luminous star associated with XGPS-3, which shares several similarities with $\\eta$ Carinae. XGPS-36 is another example of an early Be star containing a well developed circumstellar disc and exhibiting thin thermal X-ray emission with typical temperatures of $\\sim$ 10\\,keV, much hotter than usually measured for normal OB stars in which shocks in the high-velocity radiation-pressure driven wind are responsible for the X-ray emission. The origin of the relatively modest X-ray emission (\\Lx\\ $\\sim$ 10$^{32-33}$\\ergs ) remains a matter of lively debate. Two distinct mechanisms have been proposed, mainly in the context of $\\gamma$-Cas, namely accretion on a companion star, most probably a white dwarf, or magnetic interaction between the early-type star and its decretion disc. XGPS-36 is the eighth identified member of the growing  class of $\\gamma$-Cas analogs, and the fourth found so far in XMM-Newton observations of the Galactic Plane.  \\end{enumerate}  Our optical spectroscopic campaign has successfully unveiled the nature of the Galactic sources in the hard X-ray flux interval between 10$^{-12}$ and 10$^{-13}$\\ergscm. This \\Fx\\ range corresponds to the faintest flux limit covered by the ASCA Galactic Plane Survey but still remains between one and two orders of magnitudes above the sensitivity reached by the deepest Chandra pointings. Our identified sample of hard X-ray sources mainly consists of massive stars, possible X-ray binary candidates and CVs with X-ray luminosities in the range of \\Lx $\\sim$ 10$^{32}$ to 10$^{34}$\\ergs\\ and located within a few kpc from the Sun. At first glance, the nature of the ``local'' population of hard X-ray sources does not differ from that inferred from infrared follow-up imaging of deep Chandra pointings in the direction of the Galactic Centre \\citep[see e.g.][and references therein]{mauerhan2009}. However, the relative numbers of the various classes might well be quite different. Our work also illustrates the difficulty of optically identifying sources in this \\Fx\\ range, especially the cataclysmic variables as well as the importance of having arcsecond positions so as to alleviate the confusion issue."
    },
    "0911/0911.2498_arXiv.txt": {
        "abstract": "In this work we present the first non-linear, non-Gaussian full Bayesian large scale structure analysis of the cosmic density field conducted so far. The density inference is based on the Sloan Digital Sky Survey data release 7, which covers the northern galactic cap. We employ a novel Bayesian sampling algorithm, which enables us to explore the extremely high dimensional non-Gaussian, non-linear log-normal Poissonian posterior of the three dimensional density field conditional on the data. These techniques are efficiently implemented in the HADES computer algorithm and permit the precise recovery of poorly sampled objects and non-linear density fields. The non-linear density inference is performed on a 750 Mpc cube with roughly 3 Mpc grid-resolution, while accounting for systematic effects, introduced by survey geometry and selection function of the SDSS, and the correct treatment of a Poissonian shot noise contribution. Our high resolution results represent remarkably well the cosmic web structure of the cosmic density field. Filaments, voids and clusters are clearly visible. Further, we also conduct a dynamical web classification, and estimated the web type posterior distribution conditional on the SDSS data. ",
        "introduction": "Observations of the large scale structure have always attracted enormous interest, since they contain a wealth of information on the origin and evolution of our Universe. The details of structure formation are very complicated and involve many different physical disciplines ranging from quantum field theory, general relativity or modified gravity to the dynamics of collisionless matter and the behavior of the baryonic sector. Throughout cosmic history the interplay of these different physical phenomena therefore has left its imprints in the matter distribution surrounding us. Probes of the large scale structure, such as large galaxy surveys, hence enable us to test current physical and cosmological theories and will generally further our understanding of the Universe.  Especially a cosmographical description of the matter distribution will permit us to study details of structure formation mechanisms and the clustering behavior of galaxies as well as it will provide information on the initial fluctuations and large scale cosmic flows. For this reason, several different methods to recover the three dimensional density or velocity field from galaxy observations have been developed and applied to existing galaxy surveys \\citep[][]{1993PhRvE..47..704E,HOFFMAN1994,1994ASPC...67..171L,1994ApJ...423L..93L,1995A&AS..109...71Z,1995MNRAS.272..885F,1995ApJ...449..446Z,1997MNRAS.287..425W,1999ApJ...520..413Z,2001misk.conf..268V,2006MNRAS.373...45E,ERDOGDU2004,KITAURA2009B}. In particular, recently \\citet{KITAURA2009} presented a high resolution three dimensional Wiener reconstruction of the Sloan Digital Sky Survey data release 6 data, which demonstrated the feasibility of high precision density field inference from galaxy redshift surveys. These three dimensional density maps are interesting for a variety of different scientific applications, such as studying the dependence of galaxy properties on their cosmic environment, increasing the detectability of the integrated Sachs-Wolfe effect in the CMB or performing constrained simulations \\citep[see e.g.][]{BISTOLAS1998,LEE2008,LEELI2008,FROMMERT2008,KLYPIN2003,LIBESKIND2009,2009MNRAS.397.2070M}.  However, modern precision cosmology demands an increasing control of observational systematic and statistical uncertainties, and the means to propagate them to any finally inferred quantity in order not to draw wrong conclusion on the theoretical model to be tested. For this reason, here we present the first application of the new Bayesian large scale structure inference computer algorithm HADES (HAmiltonian Density Estimation and Sampling) to data \\citep[see][for a description of the algorithm]{JASCHE2009B}. HADES performs a full scale non-linear, non-Gaussian Markov Chain Monte Carlo analysis by drawing samples from the lognormal Poissonian posterior of the three dimensional density field conditional on the data. This extremely high dimensional posterior distribution, consisting of \\(\\sim 10^6\\) or more free parameters, is explored via a numerically efficient Hamiltonian sampling scheme which suppresses the random walk behavior of conventional Metropolis Hastings algorithms by following persistent trajectories through the parameter space \\citep[][]{DUANE1987,NEAL1993,NEAL1996}. The advantages of this method are manyfold. Beside correcting observational systematics introduced by survey geometry and selection effects, the exact treatment of the non-Gaussian behavior and structure of the Poissonian shot noise contribution of discrete galaxy distributions, permits very accurate recovery of poorly sampled objects, such as voids and filaments. In addition, the lognormal prior has been demonstrated to be an adequate statistical description for the present density field and hence enables us to infer the cosmic density field deep into the non-linear regime \\citep[see e.g.][]{HUBBLE1934,PEEBLES1980,COLES1991,GAZTANAGA1993,KAYO2001}. The important thing to remark about HADES is, that it does not only yield a single estimate, such as a mean, mode or variance, in fact it provides a sampled representation of the full non-Gaussian density posterior. This posterior encodes the full non-linear and non-Gaussian observational uncertainties, which can easily be propagated to any finally inferred quantity.  The application of HADES to Sloan Digital Sky Survey (SDSS) data therefore is the first non-linear, non-Gaussian full Bayesian large scale structure analysis conducted so far \\citep[SDSS;][]{YORK2000}. In particular, we applied our method to the recent SDSS data release 7 (DR7) data \\citep[DR7;][]{SDSS7}, and produced about 3TB of valuable scientific information in the form of \\(40000\\) high resolution non-linear density samples. The density inference is conducted on an equidistant cubic grid with side length \\(750\\) Mpc consisting of \\(256^3\\) volume elements. The recovered density field clearly reveals the cosmic web structure, consisting of voids, filaments and clusters, of the large scale structure surrounding us.  These results provide the basis for forthcoming investigations on the clustering behavior of galaxies in relation to their large-scale environment. Such analyses yield valuable information about the formation and evolution of galaxies. In example, it has been known since long that physical properties such as morphological type, color, luminosity, spin parameter, star formation rate, concentration parameter, etc., are functions of the cosmic environment \\citep[see e.g.][]{DRESSLER1980,POSTMAN1984,WHITMORE1993,LEWIS2002,GOMEZ2003,GOTO2003,ROJAS2005,KUEHN2005,BLANTON2005,BERNARDI2006,CHOI2007,PARK2007,LEE2008,LEELI2008}.  In this work we already conduct a preliminary examination of the dependence of stellar mass \\(M_{\\star}\\) and \\(g-r\\) color of galaxies on their large-scale environment. However, more thorough investigations will be presented in following works. Analyzing galaxy properties in the large-scale environment also requires to classify the large scale structure into different cosmic web types. We do so by following the dynamic cosmic web type classification procedure as proposed by \\citet{HAHN2007} with the extension of \\citet{FORERO2009}. The application of this procedure to our results yields the cosmic web type posterior, which provides the probability of finding a certain web type (void, sheet, filament, halo) at a given position in the volume conditional on the SDSS data. This permits simple propagation of all observational uncertainties to the final analysis of galaxy properties.  The paper is structured as follows. We start by a brief review of the methodology in section \\ref{Methodology}, particularly describing the lognormal Poissonian posterior and the Bayesian computer algorithm HADES. Additionally, here we describe the dynamic web classification procedure as mentioned above. In section \\ref{DATA} we give a description of the SDSS DR7 data and present necessary data preparation steps required to apply the analysis procedure. Specifically, we describe the preparation of the linear observation response operator and the creation of the three dimensional data cube. In the following section \\ref{RESULTS} we present the results obtained from the non-linear, non-Gaussian sampling procedure. We start by analyzing the convergence behavior of the Markov chain via a Gelman \\& Rubin diagnostic, followed by a discussion of the properties of individual Hamiltonian samples. Here we also provide estimates for the ensemble mean density field and according variance. These fields depict remarkable well the properties of the cosmic web consisting of voids, filaments and halos. Based on these results we perform a simple cosmic web classification in section \\ref{web_classification_data}. In section \\ref{GALAXY_PROPERTIES}, we present a preliminary examination on the correlation between the large-scale environment of galaxies and their physical properties. In particular, here we study the stellar mass and \\(g-r\\) color of galaxies in relation with the density contrast \\(\\delta\\). We conclude the paper in section \\ref{SUMMARYANDCONCLUSION} by summarizing and discussing the results. ",
        "conclusions": "\\label{SUMMARYANDCONCLUSION} In this work we present the first application of the non-linear, non-Gaussian Bayesian large scale structure inference algorithm HADES to SDSS DR7 data.  HADES is a numerically efficient implementation of a Hamiltonian Markov Chain sampler, which performs sampling in extremely high parameter spaces usually consisting of \\(\\sim 10^7\\) or more free parameters. In particular, HADES explores the lognormal Poissonian density posterior, which permits precision recovery of poorly sampled objects and density field inference deep into the non-linear regime \\citep[][]{JASCHE2009}.  The large scale structure inference was conducted on a cubic equidistant grid with sidelength of \\(750\\) Mpc consisting of \\(256^3\\) voxels, yielding a grid resolution of about \\(3\\) Mpc. The large scale structure inference procedure correctly accounts for the survey geometry, completeness and radial selection effects as well as for the correct treatment of Poissonian noise. The analysis yielded  about 3 TB of valuable scientific information in the form of full three dimensional density samples of the lognormal Poissonian density posterior. This set of density samples is thus a sampled representation of the full non-Gaussian density posterior distribution and therefore encodes all observational systematics and statistical uncertainties. Hence, all uncertainties and systematics can seamlessly be propagated to any finally inferred quantity, by simply applying the according inference procedure to the set of samples. In this fashion, the results permit us to make precise and quantitative statements about the large scale density field and any derived quantity.  We stress that our Hamiltonian samples are not the result of a filtering procedure. A filter generally suppresses the power of the signal in low signal-to-noise regions and therefore does not yield a physical meaningful density, since it lacks power in poorly or unobserved regions. However, each Hamiltonian density sample represents a complete physical matter field realization conditional on the observations, in the sense that it possesses correct physical power throughout the entire volume. Already visual inspection of these density samples shows a homogeneous distribution of power throughout the entire volume. This fact was emphasized by the demonstration of power-spectra measured from these density samples, which show no sign of being affected by lack of power or artificial mode coupling nor do they show any sign of being affected by an adaptive smoothing kernel as would be expected for filter applications. It should be noted that this fact marks the crucial difference of our method to previous filter based density estimation procedures.  In section \\ref{MEANANDVARIANCE}, we estimated the ensemble mean and the according variance from the \\(40000\\) density samples. The estimated ensemble mean nicely depicts the cosmic web consisting of filaments, voids and clusters extracted from the SDSS data. It is clear, that the ensemble mean represents the mean estimated from the lognormal Poissonian posterior conditional on the SDSS data. Therefore, it encodes the observational uncertainties and systematics. This can be seen by the fact, that the ensemble mean approaches cosmic mean density in poorly or not observed regions. Further, we plotted the according variance, which demonstrates that the non-Gaussian behavior and structure of the Poissonian shot noise was correctly taken into account in our analysis. Especially, the expected correlation between high mean density and high variance regions was clearly visible. We also estimated the cumulative probabilities for the density amplitude at each volume element, and demonstrated that the recovered density fields truly cover the broad range from linear to non-linear density amplitudes.  To characterize the environment of our galaxy sample, but also to demonstrate the advantages of the Hamiltonian samples, we performed an example cosmic web type classification in section \\ref{web_classification_data}. In particular, we followed the dynamical cosmic web classification approach of \\citet{HAHN2007} with the extensions proposed by \\citet{FORERO2009}. This procedure involves the calculation of the cosmic deformation tensor and its eigenvalues. We demonstrated that this procedure can easily be applied to the set of samples, since they represent full physical matter field realizations. As a byproduct of this procedure we estimated the ensemble mean for the three eigenvalues of the cosmic deformation tensor. Further, we classified the individual volume elements as one of the four different web types void, sheet, filament and halo. The classification into four discrete web types enabled us to explicitely estimate the cosmic web posterior, which provides the probability of finding a specific web type at a given point in the volume conditional on the SDSS data. This result is especially appealing from a Bayesian point of view, since it emphasizes the fact, that the result of a Bayesian method is a complete probability distribution rather than just a single estimate. Here we saw, that especially voids are a sensitive measure for the cosmic web. Of course, it is possible to repeat the cosmic web classification in a similar manner with any other classification procedure.  In the following section \\ref{GALAXY_PROPERTIES}, we presented a preliminary examination of the correlation between the large-scale environment and physical properties of galaxies. In particular, we considered the stellar mass and $g-r$ color of galaxies in relation to the density contrast \\(\\delta\\). A qualitative analysis revealed that there exist correlation between these galaxy properties and the large scale structure. In particular, massive galaxies are more likely to be found in massive structures, while low mass galaxies reside in void like structures. The plots also demonstrate the different clustering behavior of red and blue galaxies. Also note, that these observed trends are consistent with previous works \\citep[][]{LEE2008,LEELI2008,LI2006}. However, more work is required in order to provide quantitative statements. This will be done in a forthcoming publication.  The results presented in this work will be valuable for many subsequent scientific analyses of the dependence of galaxy properties on their cosmic environment. Herefore, particularly the Hamiltonian samples allow for a more intuitive handling of observational data, since they can be understood as full matter field realizations or different multiverses consistent with our data of the Universe we live in. Beside providing quantitative characterizations of the large scale structure, the results also give us an intuitive understanding of the three dimensional matter distribution in our cosmic neighborhood. We intend to make our data publically available to the community.  Future applications will also take into account non-linear bias models and peculiar velocity sampling procedures, to provide even more accurate density analyses.  We hope that this work demonstrates the potential of Bayesian large scale structure inference and its contribution to current and future precision analyses of our Universe."
    },
    "0911/0911.2983_arXiv.txt": {
        "abstract": "We study the \\Mbh{}--\\Mhost{} relation as a function of Cosmic Time in a sample of 96 quasars from $z=3$ to the present epoch. In this paper we describe the sample, the data sources and the new spectroscopic observations. We then illustrate how we derive \\Mbh{} from single-epoch spectra, pointing out the uncertainties in the procedure. In a companion paper, we address the dependence of the ratio between the black hole mass and the host galaxy luminosity and mass on Cosmic Time. ",
        "introduction": "The discovery of tight relations between the mass of massive black holes, \\Mbh{}, and the large scale properties of the galaxies where they reside \\citep[see][for a review]{ferrarese06} is one of the most intriguing results in astrophysics of the last decade, given the consequences in the frame of galaxy formation and evolution. When and how these relations are set are still open questions.  Measuring \\Mbh{} of quiescent massive black holes is extremely challenging even in nearby galaxies, and has been done successfully only in few tens of cases \\citep{ferrarese06,pastorini07}. By contrast, an estimate of \\Mbh{} in Type-I Active Galactic Nuclei (AGN) is possible by assuming that the gas emitting broad lines is in virial equilibrium \\citep[][but see also Marconi et al. 2008 and references therein for possible effects of non-virial components]{peterson00}, and that the line width traces the black hole potential well. Based on reverberation mapping studies \\citep{blandford82}, \\citet{kaspi00} found a correlation between the characteristic size of the broad line region (hereafter, BLR) and the continuum luminosity of the AGN. This allows an estimate of the black hole mass from single epoch low-resolution spectra.  In order to sample the black hole -- host galaxy relations through a wide range of Cosmic Ages, one has to focus on the brightest AGN. Quasars have been detected up to $z\\gsim6$ \\citep[][]{fan04}. Large field surveys such as the SDSS allowed a detailed spectroscopic study of $\\sim60,000$ quasars with $z<4$ \\citep{shen08}. The drawback is that the typical nuclear-to-host luminosity ratio in quasars is such that the light from the host galaxy is outshone by the nuclear component. This usually prevents the detection of stellar features in the spectra of bright quasars. Only through the excellent resolution of the HST, together with state-of-art observing techniques in the NIR at ground-based telescopes, the detection of the extended emission of the host galaxies of few hundreds of quasars up to $z\\lsim3$ became possible \\citep[see][and references therein]{kotilainen09}.  In this project we focus on quasars in the redshift range $0<z<3$ with known host galaxy luminosity, in order to study the evolution of the \\Mbh{}--\\Mhost{} relation. This will shed some light on the joint evolution of black holes and galaxies up to and immediately beyond the crucial age of maximum quasar activity \\citep{dunlop90} and star formation \\citep{madau98}. Here we present the sample, the new spectroscopic observations and we infer the black hole masses. In a companion paper \\citep{paperII}, we address the topic of the \\Mbh--\\Mhost{} relation as a function of Cosmic Time.   Throughout the paper, we adopt a concordance cosmology with $H_0=70$ km/s/Mpc, $\\Omega_m=0.3$, $\\Omega_\\Lambda=0.7$. We converted the results of other authors to this cosmology when adopting their relations and data.   \\section[]{The sample}\\label{sec_sample}  Up to now, few hundreds of quasar host galaxies have been resolved at $z<3$. In order to minimize the uncertainties due to colour and filter corrections, we select objects observed with filters roughly sampling the rest-frame $R$-band. Furthermore, as the \\Mbh{}--\\Mhost{} relation is sensitive to the spheroidal rather than the total mass of the host galaxy \\citep{marconi03,haring04}, we consider only those targets with host galaxies classified as elliptical. Our sample consists of: \\begin{itemize} \\item[i--] 43 low redshift ($z\\lsim0.5$) quasars imaged with the HST-Wide Field Camera \\citep{bahcall97,hooper97,boyce98,kirhakos99,hamilton02,pagani03, dunlop03,floyd04,labita06,kim08a}. They represent the bulk of our knowledge of low-$z$ quasar host galaxies. UV spectra of 28 of these quasars are taken from the HST Faint Object Spectrograph archive \\citep[see][]{labita06}. For 36 objects we collect optical spectra from the Sloan Digital Sky Survey \\citep[][13 quasars]{adelman07} and through on-purpose observations taken at the Asiago Telescope \\citep[23 quasars; see][]{decarli08a}. \\item[ii--] 60 mid- and high-redshift ($z>0.5$) quasars imaged in the NIR through ground-based observations under optimal seeing conditions performed by our group \\citep[50 objects; see][]{kotilainen98,kotilainen00, falomo04,falomo05,hyvonen07a,hyvonen07b,kotilainen07,falomo08,kotilainen09, decarli09a} or from HST-based compilations \\citep[10 sources from][]{kukula01,ridgway01}. We note that all these studies have very high host galaxy detection rates ($>$$85$ per cent at $z<2$, $>$$60$ per cent beyond $z=2$). Optical spectra were collected at the Nordic Optical Telescope and the ESO/3.6m telescope (see section \\ref{sec_observations}). \\end{itemize} The 5 unresolved quasars in \\citet{kotilainen09} were spectroscopically observed when the analysis of the imaging data was not complete yet. They are not included in the study of the evolution of the \\Mbh{}--\\Mhost{} relation. Other two objects were dropped, as we could not estimate \\Mbh{} from our spectra (see Appendix \\ref{sec_appendix}). Therefore the final sample consists of 96 quasars. According to the \\citet{vcv06} catalogue, 48 of them are radio loud quasars (RLQs) and 48 are radio quiet (RQQs). We remark that our sample is approximately twice as large as those of \\citet{peng06a,peng06b} and \\citet{mclure06} and represents the largest dataset ever considered in the study of the evolution of the \\Mbh{}--\\Mhost{} relation. The distribution of our targets in the ($z$,$M_V$) plane is shown in Figure \\ref{fig_distr_z}. Table \\ref{tab_sample} lists the main properties of each quasar in our sample.  \\begin{figure} \\begin{center} \\includegraphics[width=0.49\\textwidth]{fig_zMabs.ps}\\\\ \\caption{The distribution of the quasars in our sample in the ($z$,$M_V$) plane. $M_V$ is the total rest-frame $V$-band absolute magnitude of the quasars, estimated from the $V$-band apparent magnitudes available in the \\citet{vcv06} catalogue and $k$-corrected assuming the quasar template by \\citet{francis91}. Filled symbols refer to the mid- and high-$z$ data for which we present optical spectroscopy in this paper. Empty symbols mark the $z\\lsim0.5$ data from \\citet{labita06} and \\citet{decarli08a}. Circles are radio loud quasars, triangles are radio quiet. The broad emission lines observed in the various redshift windows are also labelled. We note that the $-27<M_V<-26$ luminosity range (long-dashed lines), as well as the $M_*(z)>M_V>M_*(z)-1$ \\citep[where $M_*(z)$ is the characteristic luminosity of the quasar luminosity function by][plotted in short-dashed lines]{boyle00} are well sampled at any redshift bin. }\\label{fig_distr_z} \\end{center} \\end{figure}    \\section[]{New observations, data reduction}\\label{sec_observations}  Spectra of the $z>0.5$ quasars were collected in 5 observing runs at the European Southern Observatory (ESO) $3.6$m telescope in La Silla (Chile) and at the $2.56$m Nordic Optical Telescope (NOT) in La Palma (Spain). Table \\ref{tab_runs} lists the dates of the observations and the number of spectra collected in each run, while table \\ref{tab_journal} summarizes the journal of observations.  The ESO Faint Object Spectrograph and Camera (v.2, EFOSC2; see Buzzoni et al. 1984) and its twin NOT instrument, the Andalucia Faint Object Spectrograph and Camera (ALFOSC), were mounted in long-slit spectroscopy configuration. EFOSC2 observations were carried out with grism \\#4, yielding a spectral resolution of $R\\sim 400$ ($1.2\"$ slit) in the spectral range $4100$--$7500$ \\AA{} ($\\Delta\\lambda/$pxl=$3.36$ \\AA/pxl). For NOT observations, we used ALFOSC grisms \\#6 and \\#7, which allow the observation of the $3500$--$5530$\\footnote{The nominal observed range of ALFOSC grism \\#6 is larger blue-wards, but the sensitivity is so low that we decided to drop the observed spectra at wavelengths below $3500$ \\AA.} and $3800$--$6840$ \\AA{} windows with spectral resolutions $R\\sim490$ and $650$ with the $1.0\"$ slit ($\\Delta\\lambda/$pxl$\\approx1.5$ \\AA/pxl). At the central wavelength, the instrumental resolutions are $12.6$ \\AA{} (EFOSC2+grism \\#4) and $8.2$ \\AA{} (ALFOSC+grism \\#6 and \\#7). \\begin{table} \\begin{center} \\caption{List of the observing runs. (1) Run ID. (2) Telescope (3) Dates of observations. (4) Number of observed objects per run. } \\label{tab_runs} \\begin{tabular}{cccc} \\hline Run & Telescope  & Nights            & N.obj.\\\\ (1) &  (2)       &  (3)              & (4)   \\\\ \\hline E77 & ESO/$3.6$m & Sep.30-Oct.1 2005 &  6 \\\\ E78 & ESO/$3.6$m & Mar.23-25    2007 & 12 \\\\ E79 & ESO/$3.6$m & Sep.8-12     2007 & 22 \\\\ N35 & NOT        & Apr.9-10     2007 &  2 \\\\ N36 & NOT        & Oct.17-19    2007 & 18 \\\\ \\hline \\end{tabular} \\end{center} \\end{table}  \\begin{figure*} \\begin{center} \\includegraphics[angle=-90, width=0.9\\textwidth]{fig_example.ps}\\\\ \\caption{The combined spectrum of the quasars considered in the present study, normalized to the continuum value at 2250 \\AA{}. The median spectrum is plotted with a solid line. Main emission lines are labelled. The shaded regions mark the intervals used in the fit of the continuum. The resulting fit is plotted as a dash line. }\\label{fig_cont} \\end{center} \\end{figure*}   Standard IRAF tools were used in the data reduction. The \\texttt{ccdred} package was employed to perform bias subtraction, flat field correction, image alignment and combination. Cosmic rays were eliminated by combining 3 or more exposures of the same objects, and applying \\texttt{crreject} algorithm while averaging. When only one or two bidimensional spectra were available, we applied the \\texttt{cosmicrays} task in the \\texttt{crutils} package. In these cases, in order to prevent the task from altering the narrow component of emission lines, we masked the central region of our bidimensional spectra. The spectra extraction, the background subtraction and the calibrations both in wavelength and in flux were performed with \\texttt{doslit} task in \\texttt{specred} package, using He-Ar, Th-Ar and He-Ne lamps and spectrophotometric standard stars as reference. Wavelength calibration residuals are around $0.5$, $0.35$ and $0.03$ \\AA{} in the three adopted setups (sub-pixel), thus implying a negligible ($<1$ per cent) error on redshift estimates. Absolute flux calibration of spectra was corrected through the photometry of field stars, as described in \\citet{decarli08a}. This procedure yields uncertainties in the flux calibration as low as $0.05$ mag \\citep[see, e.g., ][]{kotilainen09}, and commonly around $0.1$ mag. Photometric accuracy of each target is reported in table \\ref{tab_journal}. Galactic extinction was accounted for according to \\citet{schlegel98}, assuming $R_V = 3.1$. We shifted the spectra to the rest frame, according to the catalogue $z$, remembering that quasar lines with different ionization potentials may present slightly different shifts \\citep[e.g.][]{bonning07}. Average signal to noise ratio of our spectra is $\\sim$30. The composite spectrum, obtained by median averaging rest frame individual observations of these new data, is presented in Figure \\ref{fig_cont}. The whole dataset is available electronically at \\texttt{www.dfm.uninsubria.it/astro/caqos/}.   \\section[]{Data analysis}\\label{sec_analysis}  We focus our attention on the analysis required to estimate \\Mbh{} from single epoch observations of the rest-frame UV spectra of quasars. Applied to the gas in the BLR of Type-1 AGN, the virial paradigm yields: \\begin{equation}\\label{eq_virial} {\\cal M}_{\\rm BH} = G^{-1} R_{\\rm BLR} v_{\\rm BLR}^2 \\end{equation} where $R_{\\rm BLR}$ is the characteristic radius of the broad line emission, and $v_{\\rm BLR}$ is the velocity of the emitting clouds at $R_{\\rm BLR}$.  The cloud velocity can be inferred from the width of the broad emission lines, e.g.: \\begin{equation}\\label{eq_def_f} v_{\\rm BLR} = f \\cdot {\\rm FWHM} \\end{equation} where $f$ is a geometrical factor around unity which accounts for the de-projection of $v_{\\rm BLR}$ from the line-of-sight, and FWHM is the Full Width at Half Maximum of the line profile (see McGill et al. 2008 and Decarli et al. 2008 for discussions on other line width parametrizations). On the other hand, the BLR size cannot be directly measured from single epoch spectra. Our estimates of the broad line region size rely on the discovery that $R_{\\rm BLR}$ scales with a certain power of the continuum luminosity of the AGN, $\\lambda L_\\lambda$ \\citep[see][]{kaspi00}, as expected from simple photoionization models.   \\subsection[]{The continuum luminosity and the BLR size}  Quasar UV--optical spectra are characterized by the superposition of the following components: \\begin{itemize} \\item[-] A power-law-like continuum from the nucleus; \\item[-] Broad lines emitted within the BH influence radius; \\item[-] Narrow emission lines from the quasar host galaxy and the AGN Narrow Line Region; \\item[-] The star light continuum from the host galaxy; \\item[-] A pseudo-continuum due to the blending of several \\FeII{} and \\FeIII{} multiplets. \\end{itemize} In order to infer the BLR radius, we have to isolate the first component from the others. Our spectra cover the rest-frame UV range of bright AGN, where the flux from the host galaxy star light is always negligible. The contamination of both broad and narrow emission lines is usually avoided simply by fitting the power-law continuum to the observed spectra in a number of wavelength windows free of strong features. Here we adopted the intervals: 1351--1362, 1452--1480, 1680--1710, 1796--1834, 1970--2010, 2188--2243, 2950--2990, 3046--3090 \\AA{} (see Figure \\ref{fig_cont}). The fitted parameters, together with the derived monochromatic specific fluxes and luminosities are reported in table \\ref{tab_fits}. We note that the luminosity estimates obtained through the fit of the power-law component and those obtained from a direct measure of the continuum at 1350 and 3000 \\AA{} are practically equivalent in our datasets, the differences being $<10$ per cent on average.  Concerning luminosity--radius relations, simple photoionization models with a constant ionization parameter predict that a ionizing source emitting isotropically affects a region with a characteristic radius scaling as the square root of the source luminosity. This dependence has been confirmed in several reverberation mapping studies focused on \\Hb{}. Time-lag data of the rest-frame UV lines are still limited in number, therefore any available relation for \\Mgii{} and \\Civ{} is affected by poor statistics \\citep[e.g., see figure 6 in][]{kaspi07}. For consistency with our low-redshift studies, we will refer to the the relations provided by \\citet{mclure02} for \\Mgii{}\\footnote{\\citet{mclure02} set their relation by comparing the \\Hb{} time lags to the continuum luminosity around 3000 \\AA{}, given the fact that \\Mgii{} and \\Hb{} have similar ionizing potentials, hence they are emitted in the same region. This assumption allows us to adopt the same geometrical factor for the \\Mgii{} BLR as the one derived in \\citet{decarli08a} for \\Hb{}. Note that \\citet{mclure02} fitted the line profile of \\Mgii{} into a broad and a narrow component, and used only the former to infer \\Mbh{} \\citep[although there is no evidence of significant narrow components in \\Mgii{} lines: see, e.g.,][]{shen08}. As a consequence, the values of \\Mbh{} reported in their study are systematically larger than ours, but our conclusions are unchanged as we adopt our own internal calibration of $f$.}: \\begin{equation} \\frac{R_{\\rm BLR}({\\rm MgII})}{10~ {\\rm lt-days}}=(2.52\\pm0.3) \\left[\\frac{\\lambda L_{\\lambda}(3000 \\AA)}{10^{44} {\\rm erg/s}}\\right]^{0.47\\pm0.05} \\end{equation} and by \\citet{kaspi07} for \\Civ{}: \\begin{equation} \\frac{R_{\\rm BLR}({\\rm CIV})}{10~ {\\rm lt-days}}=(0.24\\pm0.06) \\left[\\frac{\\lambda L_{\\lambda}(1350 \\AA)}{10^{43} {\\rm erg/s}}\\right]^{0.55\\pm0.04} \\end{equation} All these relations are based on low-redshift ($z<0.3$) objects, with the only exception of S5 0836+71 \\citep[$z=2.17$; see][]{kaspi07}. Throughout this work, we assume that no significant Cosmic evolution of the radius--luminosity relations occurs in the sampled redshift range.  The \\Civ{} relation is poorly constrained, as mentioned above, especially in the bright side where most of our objects reside. Assuming a different slope of the luminosity--time lag relation, e.g., $0.5$, affects our radius estimates at \\lLl $\\sim$ $10^{47}$ erg/s up to $0.2$ dex (i.e., a factor $\\approx1.7$). Nevertheless, we stress here that in terms of \\Mbh{} estimates such a difference would result in a global offset of all values and in a re-definition of the geometrical factor $f$, with little effect on the evolution of the \\Mbh{}--\\Mhost{} relation.  The scatter around the luminosity--radius relations dominates the uncertainties in the radius estimates, contributing up to a factor $\\sim2$, while the uncertainties in the luminosity estimates never exceed 10 per cent.    \\begin{figure} \\begin{center} \\includegraphics[width=0.49\\textwidth]{fig_fit2.ps}\\\\ \\caption{Examples of the line fit, for \\Civ{} in quasar Q2225-403 A (\\emph{upper panel}) and for \\Mgii{} in quasar Q2225-403 B (\\emph{lower panel}). The rest-frame observed spectra of the two quasars are shown in dotted lines. Dotted vertical lines mark the regions used for the underlying continuum estimates and the fitted wavelength range. The thin, solid lines show the continua underlying broad emission lines. The bold, solid lines are the broad emission line fits. In the bottom side of each panel, we plot the fitted \\FeII{} template and the broad line model (dotted lines) together with the fit residual (dashed line). }\\label{fig_fit} \\end{center} \\end{figure}  \\subsection[]{Emission line widths and cloud velocities}  Measuring the width of broad lines consists of three key steps: the definition of the local continuum, the removal of contaminating spectral features (in particular the \\FeII{} multiplets), and the fit with a certain analytical function.  We define underlying continua by matching the fluxes in the windows 2230--2250 and 3020--3050 \\AA{} for \\Mgii{} and 1465--1485 and 1685--1705 \\AA{} for \\Civ{}, where no significant feature is observed (see Figure \\ref{fig_fit}). To evaluate the effects of \\FeII{} contamination, we first fitted the \\Mgii{} and \\Civ{} lines in the 2720--2880 and 1490--1570 \\AA{} windows, where the \\FeII{} contribution is less relevant. Then, we compared these estimates of the line widths with those obtained by fitting the 2450--3020 and 1490--1620 \\AA{} regions with a superposition of a broad line model plus a template reproducing the \\FeII{} emission \\citep{vestergaard01}. In the latter approach the \\FeII{} emission observed in narrow line Type-1 AGN (usually, IZw001) is taken as representative for all quasars, after being properly broadened. The validity of this assumption has been widely discussed in the literature, both from a theoretical and an observational point of view \\citep[e.g.,][]{phillips78,boroson92,marziani03,tsuzuki06,mcgill08}. Since most \\FeII{} lines are blended with each others, their width is hardly constrained. We therefore fixed their broadening to the width of \\Civ{} and \\Mgii{} lines. Figure \\ref{fig_fit} shows examples of fits on \\Civ{} and \\Mgii{}. Line width estimates obtained with or without the \\FeII{} modelling are usually consistent within 20 per cent, but \\Mgii{} shows few deviations larger than 30 per cent, especially when the \\FeII{} emission is strong with respect to the line flux. Hereafter, we will refer to the line models obtained by template fitting the \\FeII{} emission.  \\begin{figure*} \\begin{center} \\includegraphics[width=0.4\\textwidth]{fig_lf_hb.ps} \\includegraphics[width=0.4\\textwidth]{fig_lf_c4_lz.ps}\\\\ \\includegraphics[width=0.4\\textwidth]{fig_lf_m2.ps} \\includegraphics[width=0.4\\textwidth]{fig_lf_c4_hz.ps}\\\\ \\caption{The distribution of our data in the (\\lLl,FWHM) plane, for \\Hb{}, \\Civ{} at low-$z$ (upper panels), \\Mgii{} and \\Civ{} at high-$z$ (bottom panels). The (normalized) projections along each axis are also shown. The loci with constant \\Mbh{} (dot-dashed lines) and \\Edd{} (dashed lines) are plotted. Concerning the luminosities, low-$z$ \\Civ{} data show on average The values of \\Hb{} FWHM show a wider distribution with respect to \\Mgii{}. data are log-normally distributed both in the \\lLl{} and in the FWHM values. \\Civ{} data show less spread in the sampled \\lLl{} space, and a flatter distribution in the line width with respect to the \\Mgii{}. }\\label{fig_labita} \\end{center} \\end{figure*} Concerning the fitted function, a single gaussian usually does not provide satisfactory fits, especially for \\Civ{} \\citep[see][]{decarli08a,shen08,bon09}. We performed our analysis both with the sum of 2 gaussian profiles with the same peak (hereafter, G2) and the Gauss-Hermite series \\citep[GH;][]{vandermarel93}, truncated at the fourth order. Reduced-$\\chi^2$ values obtained from these functions are similarly good, suggesting that both functions can describe line profiles reasonably well. No significant offset is observed among the line width estimates based on the two functions, the residuals lying within 20 per cent \\citep[see also Figure A1 in][]{decarli08a}. We thus conclude that the two functions are equivalent to our purposes. We note that, since G2 fits are necessarily symmetric, while GH fits are not, the consistency between the two estimates of the line width suggests that the FWHM is poorly sensitive to line profile asymmetries. Throughout the paper, we will refer to the FWHM estimates based on GH fits.  Summarizing, we estimate that the typical uncertainties in the line width estimates due to the adopted fit procedures lie within 20 per cent, and usually even lower (if some modelling of the \\FeII{} emission is adopted). Reduced-$\\chi^2$ maps suggest that, given the high signal to noise ratio of our data (exceeding 20 in all but few cases), formal uncertainties in the parameter estimates contribute to $\\lsim 10$ per cent of the FWHM value. Therefore, we conclude that typical uncertainties in the FWHM estimates are around 20 per cent.    \\section[]{Virial estimates of BH masses}\\label{sec_results}  In this section, we match the new estimates of the continuum luminosity, FWHM and \\Mbh{} with those from our previous low-$z$ studies \\citep{labita06,decarli08a}.  In Figure \\ref{fig_labita}, we plot the distribution of our data in the (\\lLl,FWHM) plane \\citep[see a similar approach in][]{fine08,labita09a}. This allows us to monitor how observable quantities (here, the continuum luminosity of the quasar and the line widths) affect our estimates of \\Mbh{} and the Eddington ratio, \\Edd, derived \\emph{assuming} equation \\ref{eq_virial}, $f({\\rm H\\beta,MgII})=1.6$, $f({\\rm CIV})=2.4$ \\citep[as defined in equation \\ref{eq_def_f} and in][]{decarli08a} and the bolometric corrections given in \\citet{richards06}: $L_{\\rm bol}/$\\lLl{}=$9.26$, $5.15$ and $3.81$ for \\Hb{}, \\Mgii{} and \\Civ{} respectively. Average and rms values of FWHM, \\lLl{}, \\Mbh{} and \\Edd{} are provided in Table \\ref{tab_ave_results}, while data of individual quasars are given in table \\ref{tab_fits}.  The monochromatic luminosity at 5100 \\AA{} is on average $3.5$ times fainter than at 1350 \\AA{} in the same redshift bin. This is consistent with an average power-law index of the quasar continuum $\\alpha$ close to 2 (defined so that $F_\\lambda\\propto\\lambda^{-\\alpha}$). The 3000 \\AA{} luminosities sampled in our study range from few times $10^{44}$ to $10^{47}$ erg/s, the bulk being around $10^{46}$ erg/s. The \\Civ{} data at high-$z$ are more clustered around approximately the same luminosity as \\Mgii{} data, $\\approx 6\\cdot 10^{45}$ erg/s. This thinner distribution is due to the combination of two selection effects: lowest luminosity quasars are missed due to the Malmquist bias (the \\Civ{} line falls in the optical bands at $z>1.6$, see Figure \\ref{fig_distr_z}); highest luminosity quasars were rejected in the studies of the host galaxy luminosities, in order to make the detection of the extend emission around the nuclear source more feasible \\citep[see Figure 1 and ][]{kotilainen09}.  Concerning the line widths, \\Hb{} values show a wider dispersion with respect to the \\Mgii{} line, notwithstanding the two lines have similar ionization potential, thus they are believed to originate in the same regions \\citep[see also][]{shen08,labita09a,labita09b}. Low-redshift data from the \\Civ{} line show a smaller average value with respect to both \\Hb{} and high-$z$ \\Civ{} data, and a significantly narrower distribution \\citep[see the discussion in][]{decarli08a}.  With only few exceptions, all our data reside in the locus defined by $10^{8} \\gsim$ \\Mbh{}/\\Msun{} $\\gsim 10^{10}$ and $0.01 \\gsim$ \\Edd{} $\\gsim 1$.  \\begin{table} \\begin{center} \\caption{Average and rms values of FWHM (2), \\lLl{} (3), \\Mbh{} (4) and \\Edd{} (5) for all the subsamples in this study. } \\label{tab_ave_results} \\begin{tabular}{ccccc} \\hline Line             &$\\langle$log FWHM$\\rangle$&$\\langle$log \\lLl$\\rangle$&$\\langle$log \\Mbh{}$\\rangle$&$\\langle$log \\Edd{}$\\rangle$ \\\\ &  [km/s]          & [erg/s]          & [\\Msun]          &                  \\\\ (1)              &  (2)             &  (3)             & (4)              & (5)              \\\\ \\hline \\multicolumn{5}{l}{{\\it Low-$z$}} \\\\ \\Hb{}            & $3.66 \\pm 0.19$ & $45.01 \\pm 0.51$ & $9.01 \\pm 0.47$ & $-1.13 \\pm 0.45$ \\\\ \\Civ{}           & $3.60 \\pm 0.11$ & $45.55 \\pm 0.56$ & $9.06 \\pm 0.38$ & $-1.03 \\pm 0.34$ \\\\ \\hline \\multicolumn{5}{l}{{\\it High-$z$}} \\\\ \\Mgii{}          & $3.63 \\pm 0.19$ & $45.78 \\pm 0.54$ & $9.21 \\pm 0.49$ & $-0.82 \\pm 0.46$ \\\\ \\Civ{}           & $3.66 \\pm 0.19$ & $45.74 \\pm 0.36$ & $9.28 \\pm 0.47$ & $-1.06 \\pm 0.41$ \\\\ \\hline \\end{tabular} \\end{center} \\end{table}     In Figure \\ref{fig_shen} we compare the distributions of \\Mbh{} and \\Edd{} from our dataset with the huge quasar sample from \\citet[][S08]{shen08}\\footnote{From Shen et al. (2008) we take the values of $z$, FWHM and \\lLl{} of the objects in the same redshift ranges than our study. All the derived quantities (e.g., \\Mbh{} and \\Edd{}) are re-computed following the recipes discussed in this work.}. We dropped our low-$z$, \\Civ{}-based data from this comparison as they do not have any counter-part in the S08 data. The distributions of black hole masses and Eddington ratios computed from \\Hb{} in our study are in substantial agreement with those by S08. On the other hand, the black hole masses in our sample have increasingly smaller values than those from S08. We interpret this trend as the superposition of a number of effects: the luminosity cut adopted by \\citet{kotilainen09} for $z>2$ objects, the evolution of the luminosity and mass functions of active black holes through redshift and the occurrence of the Malmquist bias \\citep[on this topic, see][]{labita09a,labita09b}. On average, the Eddington ratios sampled in our study do not show significant differences with respect to those in S08, nor evolution from $z=0$ to $z=3$.  \\begin{figure} \\begin{center} \\includegraphics[width=0.49\\textwidth]{fig_shen2b.ps}\\\\ \\caption{ The distributions of \\Mbh{} and \\Edd{} as derived from \\Civ{} (\\emph{upper panels}), \\Mgii{} (\\emph{central panels}) and \\Hb{} (\\emph{bottom panels}). Low-$z$ \\Civ{}-based estimates are not included, as no counter-part is available in the \\citet{shen08} sample. Shaded histograms refer to our objects, while empty histograms refer to the dataset of \\citet{shen08}. }\\label{fig_shen} \\end{center} \\end{figure}  The relatively small statistics and the luminosity-based selection of the targets in our sample hinder the study of the dependence of \\Mbh{} on the redshift. We note that, on average, the higher is the redshift, the more massive are the black holes (see Figure \\ref{fig_mbh_z}). The linear best fit is: \\begin{equation} \\log {\\cal M}_{\\rm BH}/{\\rm M}_{\\odot} = (0.19 \\pm 0.06) z + (8.98 \\pm 0.06) \\end{equation} This trend is in qualitative agreement with the one found by \\citet{labita09b} with a procedure aimed to minimize the Malmquist bias. However, the occurrence of selection biases cannot be ruled out in the present work.  \\begin{figure} \\begin{center} \\includegraphics[angle=-90, width=0.49\\textwidth]{fig_mbh_z.ps}\\\\ \\caption{The redshift dependence of \\Mbh{} for the quasars in our sample. Circles, triangles and squares refer to \\Civ{}, \\Mgii{} and \\Hb{}-based \\Mbh{} estimates. The best linear fit is shown as a solid line. The dashed line refers to the redshift evolution of the maximum \\Mbh{} as derived by \\citet{labita09b}. }\\label{fig_mbh_z} \\end{center} \\end{figure} ",
        "conclusions": "We start from a list of 96 quasars with known host galaxy luminosity. New, high-quality (S/N$\\sim$30) spectra of the mid- and high-redshift targets ($z>0.5$) are presented here and matched with those of low-$z$ quasars we collected in our previous studies. We analyse the continuum luminosity and profiles of broad emission lines in order to infer black hole masses. In particular, the \\Civ{} line is studied both at low and high redshift, thus avoiding a systematic error related to the adopted emission line in the estimate of \\Mbh{}. We found that high redshift quasars in our sample do have, on average, larger BH masses than local ones, but we note that this result is potentially affected by the luminosity selection and the Malmquist bias. In a accompanying  Paper II, we study the ratio between \\Mbh{} and the luminosity and mass of the host galaxies of the quasars in our sample, and its evolution as a function of the redshift."
    },
    "0911/0911.1949_arXiv.txt": {
        "abstract": "{} {In this paper, we explore the diagnostic power of the far-IR fine-structure lines of [O{\\sc i}] 63.2$\\,\\mu$m, 145.5$\\,\\mu$m, [C{\\sc ii}] 157.7$\\,\\mu$m, as well as the radio and sub-mm lines of CO J=1-0, 2-1 and 3-2 in application to disks around Herbig Ae stars. We aim at understanding where the lines originate from, how the line formation process is affected by density, temperature and chemical abundance in the disk, and to what extent non-LTE effects are important.  The ultimate aim is to provide a robust way to determine the gas mass of protoplanetary disks from line observations.} {We use the recently developed disk code {{\\sc ProDiMo}} to calculate the physico-chemical structure of protoplanetary disks and apply the Monte-Carlo line radiative transfer code {{\\sc Ratran}} to predict observable line profiles and fluxes. We consider a series of Herbig Ae type disk models ranging from $10^{-6}$~M$_\\odot$ to $2.2\\!\\times\\!10^{-2}$~M$_\\odot$ (between 0.5 and 700\\,AU) to discuss the dependency of the line fluxes and ratios on disk mass for otherwise fixed disk parameters. This paper prepares for a more thorough multi-parameter analysis related to the Herschel open time key program {{\\sc Gasps}}.} {We find the [C{\\sc ii}]\\,157.7\\,$\\mu$m line to originate in LTE from the surface layers of the disk, where $\\Tg\\neq\\Td$. The total emission is dominated by surface area and hence depends strongly on disk outer radius. The [O{\\sc i}] lines can be very bright ($>10^{-16}$~W/m$^2$) and form in slightly deeper and closer regions under non-LTE conditions. For low-mass models, the [O{\\sc i}] lines come preferentially from the central regions of the disk, and the peak separation widens. The high-excitation [O{\\sc i}]\\,145.5\\,$\\mu$m line, which has a larger critical density, decreases more rapidly with disk mass than the 63.2\\,$\\mu$m line. Therefore, the [O{\\sc i}] 63.2\\,$\\mu$m/145.5\\,$\\mu$m ratio is a promising disk mass indicator, especially as it is independent of disk outer radius for $R_{\\rm out}>200$~AU. CO is abundant only in deeper layers $A_V\\!\\ga\\!0.05$. For too low disk masses ($M_{\\rm disk}\\!\\la\\!10^{-4}$~M$_\\odot$) the dust starts to become transparent, and CO is almost completely photo-dissociated. For masses larger than that the lines are an excellent independent tracer of disk outer radius and can break the outer radius degeneracy in the [O{\\sc i}]\\,63.2\\,$\\mu$m/[C\\,{\\sc ii}]157.7\\,$\\mu$m line ratio.} {The far-IR fine-structure lines of [C{\\sc ii}] and [O{\\sc i}] observable with Herschel provide a promising tool to measure the disk gas mass, although they are mainly generated in the atomic surface layers.  In spatially unresolved observations, none of these lines carry much information about the inner, possibly hot regions $<\\!30\\,$AU.} ",
        "introduction": "Observations of gas in protoplanetary disks are intrinsically difficult to interpret as they reflect the interplay between a complex chemical and thermal disk structure, statistical equilibrium and optical depth effects. This is particularly true if non-thermal excitation such as fluorescence or photodissociation dominate the statistical equilibrium.  The first studies of gas in protoplanetary disks concentrated on the rotational transitions of abundant molecules such as CO, HCN and HCO$^+$ \\citep[e.g.][]{Beckwith1986, Koerner1993, Dutrey1997, vanZadelhoff2001, Thi2004}. Those lines originate in the outer regions of disks, $r>100$~AU, where densities are at most $n\\sim10^7$~cm$^{-3}$. The interpretation of those lines was mainly based on tools and expertise developed for molecular clouds. Using the CO J=3-2 line, \\citet{Dent2005} inferred for a sample of Herbig Ae and Vega-type stars a trend of disk outer radius with age; on average, the outer disk radius in the 7-20~Myr range is three times smaller than that in the $<7$~Myr range. Also, the disk radii inferred from the dust spectral energy distribution (SED) are generally smaller than those derived from the gas line \\citep{Isella2007,Pietu2005}. \\citet{Hughes2008b} suggest a soft outer edge as a solution to the discrepancy. Comparison of CO J=3-2 maps of four disks to different types of disk models strongly supports a soft edge in favor of a sharp cutoff. \\citet{Pietu2007} use the CO and HCO$^+$ lines to probe the radial and vertical temperature profile of the disk. Simple power law disk models and LTE radiative transfer provides best matching results for radial temperature gradients around $r^{-0.5}$. The $^{12}$CO J=2-1 line, $^{13}$CO J=2-1 and $^{13}$CO J=1-0 lines  used in their analysis probe subsequently deeper layers and reveal a vertical temperature gradient ranging from 50~K in the higher layers to below the freeze-out temperature of CO in the midplane. This confirms earlier findings by \\citet{Dartois2003}.  However, disk masses derived from the CO lines are in general lower than disk masses derived from dust observations \\citep[e.g.][]{Zuckerman1995, Thi2001}. Possible explanations include CO ice formation in the cold midplane and photodissociation in the upper tenuous disk layers. Any single gas tracer alone can only provide gas masses of the species and volume from which it originates, the same way as dust masses derived from a single photometric measurement are only sensitive to grains of a particular size range, namely those grain sizes that dominate the emission at that photometric wavelength. Hence individual gas tracers are more valuable for probing the physical conditions of the volume where they arise than the total disk mass. A combination of suitable gas tracers can then allow us to characterize the gas properties in protoplanetary disks and study it during the planet formation process.  This paper aims at exploring the diagnostic power of the fine structure lines of [O\\,{\\sc i}] and [C\\,{\\sc ii}] in the framework of upcoming Herschel observations. Earlier modeling of these lines indicated that they should be detectable down to disk masses of $10^{-5}$~M$_\\odot$ of gas, so also in the very gas-poor debris disks \\citep{Kamp2005}. More recent work by \\citet{Meijerink2008} presents fine structure lines from the inner 40~AU disk of an X-ray irradiated T Tauri disk; the models indicate that the [O\\,{\\sc i}] emission originates over a wide range of radii and depth and is sensitive to the X-ray luminosity. However, most of the C\\,{\\sc ii} and O\\,{\\sc i} line emission comes from larger radii where the gas temperature is dominated by UV heating processes. \\citet{Jonkheid2007} find from thermo-chemical models of UV dominated Herbig Ae disks that the [O\\,{\\sc i}] lines are generally a factor 10 stronger that the [C\\,{\\sc ii}] line. We started to explore the origin of the fine structure lines in \\citet{Woitke2009a} and find that the [C\\,{\\sc ii}]~157.7~$\\mu$m line probes the upper flared surface layers of the outer disk while the [O\\,{\\sc i}]~63.2~$\\mu$m line originates from the thermally decoupled surface layers inward of about 100~AU, above $A_V \\approx 0.1$. The latter line is very sensitive to the gas temperature and might be used to distinguish between hot ($T_{\\rm gas} \\approx 1000$~K) and cold ($T_{\\rm gas} = T_{\\rm dust}$) disk atmospheres. Since the fine structure lines generally originate from a wider radial and vertical range than for example the $^{12}$CO rotational lines, they are potentially more suitable gas mass tracers.  We use in this paper the disk modeling code {{\\sc ProDiMo}} presented in \\citet{Woitke2009a} to study the gas line emission from disks around Herbig Ae stars. The main focus are the fine-structure lines of C\\,{\\sc ii} and O\\,{\\sc i} which will be observed for a large sample of disks during the {{\\sc Herschel}} open time Key Program {{\\sc Gasps}} (Gas evolution in protoplanetary systems: http://www.laeff.inta.es/projects/herschel). The disk parameters were chosen to resemble the disk around MWC480. According to previous work by \\citet{Mannings1997b}, \\citet{Thi2001} and \\citet{Pietu2007}, the disk around this star extends from $\\approx 0.5$~AU to 700~AU. We choose here a surface density profile $\\Sigma \\sim r^{-1.0}$. The central star is an A2e Herbig star with a mass of 2.2~M$_\\odot$ and an effective temperature of $8500$~K.  Section~\\ref{prodimo} gives a short summary of the disk modeling approach. The line radiative transfer method, re-gridding and the atomic input data are described in Sect.~\\ref{LineRadTrans}. We then briefly discuss some basic properties of the Herbig Ae disk models (Sect.~\\ref{resultsdisks}) before we present the fine-structure lines (Sect.~\\ref{resultslines}) and conclude with a discussion of the diagnostic strength of fine structure line ratios and a comparison to previous ISO and submm observations (Sect.~\\ref{discussion}). ",
        "conclusions": "We used here a limited series of Herbig Ae disk models computed with the new thermo-chemical disk code called {\\sc ProDiMo} to study the origin and diagnostic value of the gas line tracers [C\\,{\\sc ii}], [O\\,{\\sc i}] and CO. We do not include in this study effects of X-ray irradiation, grain settling or mixing. Even though X-rays are generally of minor importance for Herbig stars, grain settling and large scale mixing processes could affect the conclusions. The effect of grain settling will be included in a forthcoming larger model grid, but mixing processes need to be addressed with dynamical time-dependant models. The Monte-Carlo radiative transfer code {\\sc Ratran} is used to compute line profiles and integrated emission from various gas lines.  The main results are:  \\begin{itemize} \\item The [C\\,{\\sc ii}] line originates in the disk surface layer where gas and dust temperatures are decoupled. The total line strength is dominated by emission from the disk outer radius. Thus the 157.7~$\\mu$m line probes mainly the disk extension and outer disk gas temperature. The line forms in LTE. \\item The [O\\,{\\sc i}] lines originate also in the disk surface layer even though somewhat deeper than the [C\\,{\\sc ii}] line. The main contribution comes from radii between 30 and 100~AU. The [O\\,{\\sc i}] lines form partially under NLTE conditions. Differences in line emission from escape probability and Monte Carlo techniques are smaller than 10\\%. \\item CO submm lines are optically thick down to very low disk masses of $<10^{-4}$~M$_\\odot$ and form mostly in LTE. $T_{\\rm gas}=T_{\\rm dust}$ is not a valid approximation for these lines. Differences in line emission from escape probability and Monte Carlo techniques are smaller than 3\\%, except in the case of very optically thin disk models ($10^{-5}$ and $10^{-6}$~M$_\\odot$). \\item The [O\\,{\\sc i}]\\,63/145 $\\mu$m and [O\\,{\\sc i}]\\,63/[C\\,{\\sc ii}]\\,158 $\\mu$m line ratios trace disk mass in the regime between $10^{-2}$ and $10^{-6}$~M$_\\odot$. Since the [C\\,{\\sc ii}]\\,158 $\\mu$m line is very sensitive to the outer disk radius, the [O\\,{\\sc i}]\\,63/[C\\,{\\sc ii}]\\,158 $\\mu$m is degenerate in that respect and its use requires additional constraints from ancilliary gas and/or dust observations. The sensitivity of these two line ratios to the dust grain sizes underlines the importance of using SED constraints along with the gas modeling to mitigate the uncertainty of dust properties. \\item A combination of the [O\\,{\\sc i}]\\,63/145 $\\mu$m and [O\\,{\\sc i}]\\,63/[C\\,{\\sc ii}]\\,158 $\\mu$m line ratios can be used to diminish the degeneracy caused by an unknown outer disk radius. \\item Neither total CO submm line fluxes nor line ratios can be used to measure the disk mass. However, the low rotational lines studied here provide an excellent tool to measure the disk outer radius and can thus help to mitigate the degeneracy between gas mass and outer radius found for the [O\\,{\\sc i}]\\,63/[C\\,{\\sc ii}]\\,158 $\\mu$m line ratio. \\end{itemize}"
    },
    "0911/0911.4794_arXiv.txt": {
        "abstract": "{} {We seek the long-term variations close to the length of a solar cycle  in the mean meridional motion of sunspot groups (a proxy of the meridional plasma flow).} {Using the largest set of available reliable sunspot group data, the combined  Greenwich and Solar Optical Observation Network sunspot group data during the period 1879\\,--\\,2008, we determined  variations in the mean meridional motion of the sunspot groups in the  Sun's whole northern and southern hemispheres and also in different $10^\\circ$ latitude intervals. We determined the variations from   the yearly  data and  for the sake of better statistics  by binning the data into 3\\,--\\,4 year moving time intervals (MTIs)  successively shifted by one year. We  determined the periodicities in the mean meridional  motion from the  fast Fourier transform (FFT) power spectrum analysis.   The values of the periodicities are determined from the maximum entropy method (MEM) and the temporal dependencies of the periodicities are determined from the Morlet-wavelet analyses.} {We find that the mean  meridional motion of the spot groups varies considerably on a time  scale  of  about 5\\,--\\,20 years. The maximum amplitude of the variation is about  10\\,--\\,15 m s$^{-1}$   in both the northern and the southern hemispheres. Variation in the mean motion is considerably different during different solar cycles. At the maximum epoch (year 2000) of the current cycle~23, the mean  motion is relatively strong in the past  100 years and northbound in both the northern and the southern hemispheres. This abnormal behavior of the mean  motion   may be related to the low strength and the long duration of the current cycle, and also to the  violation of the Gnevyshev and Ohl rule by the cycles pair 22,23. The power spectral analyses suggest the existence of $\\approx$ 3.2- and $\\approx$ 4.3-year periodicities in the mean  motion of the spot groups  in the southern hemisphere, whereas a  13\\,--\\,16 year  periodicity is found to exist in the mean  motion of the northern hemisphere. There is  strong evidence for a latitude-time dependency in the periodicities of the mean motion. The north-south  difference in  the mean  motion also varies by about 10 m s$^{-1}$.  During the recent cycles, the north-south difference is negligibly  small. Approximate  12- and 22-year periodicities  are  found to  exist in the north-south difference. The implications of all these results are briefly discussed.} {} ",
        "introduction": "The study of the  variations in the meridional flows is important  for understanding the underlying mechanism of the solar cycle (Babcock 1961;  Ulrich \\& Boyden 2005). Surface Doppler measurements and helioseismology measurements of the surface and the subsurface flows (e.g., Ulrich \\& Boyden 2005; Gonz\\'alez Hern\\'andez et al. 2008) suggest  poleward flows of about 10\\,--\\,20 m s$^{-1}$ and a considerable north-south asymmetry in the meridional flow. Surface Doppler measurements are available  for about 3 cycles (Ulrich \\& Boyden 2005). Doppler measurements suffer from errors  caused by   the scatter light and the B-angle influences (Beckers 2007). The motions of many magnetic tracers, particularly sunspots, have been  used for a long time  as a proxy of the fluid motions to study the solar rotational and  the meridional flows (Schr\\\"oter 1985; Javaraiah \\& Gokhale 2002). The sunspot data have been available for more than 100 years. However, the derived rates of rotation and  meridional flows depend on the method of the selection of the spots or the spot groups (e.g., Howard et al. 1984;  Balthasar et al. 1986; Zappal\\'a \\& Zuccarello 1991; Zuccarello 1993; Javaraiah \\& Gokhale 1997a; Javaraiah 1999; Hiremath 2002; Sivaraman et al. 2003).  Proper motions and evolutionary factors of the spot groups may also   influence the derived rates of the flows to some extent. The proper motions are random in nature, hence their effect can be reduced with the use of a large data set. Recently, Ru\\v{z}djak et al. (2005) have found that evolutionary factors of the sunspot groups in the determination of the positions of the spot groups is small in the estimated mean meridional motion of the spot groups.  A number of scientists studied the solar cycle variations of the mean meridional motion of the spot groups (see Javaraiah \\& Ulrich 2006). Since  yearly data are inadequate  for this purpose, particularly around a solar cycle minimum, some authors have used the superposed epoch analysis of the data during a few or all   cycles for which the data were available and studied the mean solar cycle variation (Balthasar et al. 1986;  Howard \\& Gilman  1986). Recently, by using the same method,  Javaraiah \\& Ulrich (2006) analyzed the combined  Greenwich and Solar Optical Observation Network sunspot group data during the period 1879\\,--\\,2004 and determined the mean solar cycle variation in the meridional motions of the sunspot groups. In that early paper, the spot group data during the cycles 12\\,--\\,20 were superposed according to the minima of these cycles. This  yielded an average solar cycle variation of the meridional motion over about 5\\,--\\,9 cycles, which suggests that only around the end of a solar cycle the meridional motion is considerably significant from zero and it is poleward in both the northern and the southern hemispheres. However, during some individual cycles the variations may be considerably different from  this average solar cycle variation. Javaraiah \\& Ulrich (2006) also determined the cycle-to-cycle variation of the mean (over the duration of a cycle)  meridional motion of the spot groups during cycles 12\\,--\\,23 and  found  the existence of  a weak long-period cycle (Gleissberg cycle) in the cycle-to-cycle variation of the mean motion. Since the meridional speed is negligibly small or zero in some phases  of a cycle or even with opposite signs during different phases of some  cycles, hence,  the motions are washed out in the average over the cycle. In order to get rid of this problem in the study of long-term variations in the mean meridional motion, it is necessary to analyse the spot or the spot group data  in the intervals considerably shorter than the length of a solar cycle.  In the present paper  we have analyzed  the annual spot group data during 1879\\,--\\,2008 and determined the variations in the mean meridional motions of the spot groups in the northern and the southern hemispheres. As expected the statistics is poor in case of the results derived from the annual data, particularly at the  cycles minima. We  have taken some additional precautions, so it is possible  to see the patterns of the variations when they are close to  a solar cycle length, even in the annual time series. However, we have also determined the  variation in the  mean meridional motion of the spot groups by binning the spot group data into the moving-time intervals (MTIs) of lengths (3\\,--\\,4 years) considerably greater than a year, but reasonably smaller than the length of a cycle; $i.e.$, less than the half of the length of a cycle. In such a time-series which comprised the longer time intervals, it is relatively easy to detect the long-term variations near the length of a solar cycle (Javaraiah \\& Gokhale 1995, 1997b). In addition, the sizes of these series reasonably large.  Hence, it enabled us to find the  periodicities that approach the length of a solar cycle from the power spectral analysis.  In the next section we  describe the data and analysis. In Sect.~3 we present the variations in the mean meridional motion of the spot groups  in the  whole northern  and the southern hemispheres, as well as in different $10^\\circ$ latitude intervals--and the corresponding differences between the whole northern and the southern hemispheres--during the period 1879\\,--\\,2008 and point out their important features. In the same section, we show  the periodicities in the mean meridional motion of the spot groups from the traditional FFT analyses.   From the MEM analyses, we  determined the values of the periodicities, and from the Morlet-wavelet analyses we  determined the temporal dependencies of the detected periodicities. In Sect.~4 we  summarize the results and the conclusions, and briefly  discuss the implications of these  results for understanding the solar long-term variability.  \\vspace{0.3cm} ",
        "conclusions": "From the analysis of the largest available reliable  sunspot group data; $i.e.$, the combined Greenwich and SOON sunspot group data during the period 1879\\,--\\,2008, we find the following. \\begin{enumerate} \\item The mean meridional motion  of the sunspot groups varies considerably on  5\\,--\\,20 year timescales during the period 1879\\,--\\,2008. The maximum amplitude of the variation is 10\\,--\\,15 m s$^{-1}$. \\item The pattern  and amplitude  of the solar cycle variation in the mean motion are significantly different during the different cycles. During the maximum epoch (year 2000) of the current cycle, the mean  motion is relatively stronger than in the past  $\\approx$100 years, and it is northbound   in both the northern and the southern  hemispheres. \\item The north-south difference (north-south asymmetry) in the mean meridional motion of the spot groups also varies with a maximum amplitude  of about 10 m s$^{-1}$. The north-south difference was considerably  larger during the early cycles (with a strong contribution from the high latitudes). It is negligible during the recent cycles. \\item Power spectral analyses suggest that  $\\approx$ 3.2- and $\\approx$ 4.3-year periodicities exist in the mean meridional motion of the spot groups of the southern hemisphere, whereas a  13\\,--\\,16 year  periodicity is  found in the mean motion of the spot groups of  the northern hemisphere. \\item The $\\approx$ 12- and $\\approx$ 22-year periodicities are found to exist in the north-south difference of the mean  motion. \\item There is a considerable latitude-time dependence in the periodicities of the mean meridional motion of the spot groups. There is a strong suggestion that,  in the $10^\\circ$\\,--\\,$20^\\circ$  latitude-interval of the northern hemisphere, a periodicity slowly evolved from $\\approx$ 16 year to $\\approx$ 10 year, over the period 1880\\,--\\,2007, and it evolved in the  opposite way,  $\\approx$ 10 year to $\\approx$ 16 year,  in $20^\\circ$\\,--\\,$30^\\circ$ latitude interval. \\end{enumerate}   The behavior of the mean motion of the spot groups in cycle~23 is similar to that of  cycle~14, which is a low-amplitude (lowest in the last century) and considerably long-duration cycle. Cycle~23 is also a relatively  low-amplitude and long-duration cycle, so that the result above (conclusion (2)) may be a part of a real long-term behavior in the mean meridional motion of the spot groups; $i.e.$, most probably it is not an artifact of the  differences (if any) between the  Greenwich and the SOON datasets,  within the continuous time series of the combined dataset used here.   Most of the helioseismic measurements of the  meridional flows during the current sunspot cycle~23 suggest  an increase in the amplitudes of the surface and the  subsurface poleward meridional flows   with a decrease in magnetic activity (Gonz\\'alez Hern\\'andez et al. 2008), whereas we find a strong  northbound mean meridional motion of the spot groups  during the maximum of cycle~23. A reason for this discrepancy may be that sunspot motions  may not  represent the Sun's plasma motions (D'Silva \\& Howard 1994), or the motions of the magnetic structures  may represent the motions of the deeper layers of the Sun's convection zone where these structures are anchored (Javaraiah \\& Gokhale 1997a;  Hiremath 2002; Sivaraman et al. 2003; Meunier 2005). In addition, the mean meridional motion  of the sunspot groups may  only represent the mean solar meridional plasma motion at  low and middle latitudes, because sunspots data are confined to  only these latitudes. The magnetic structures of the only large spot groups during their initial days might be anchored near the base of the convection zone (Javaraiah \\& Gokhale 1997a;  Hiremath 2002; Sivaraman et al. 2003), hence might have largely equatorward motions (Javaraiah 1999). While rising through the convection zone, the magnetic structures of  the large spot groups may be fragmented into the smaller structures (Javaraiah 2003; Sch\\\"ussler \\& Rempel 2005). The small structures may move  mainly toward the poles (\\v{S}vanda et al. 2007). However, as can be seen in Figs.~1\\,--\\,3 there are also equatorward motions (may be due to an   effect of the reverse meridional flows), mainly near minima of the cycles where a  spot group is relatively   small.  Meridional flows can transport magnetic flux and cause cancellation/enhancement of magnetic flux, and it is  believed that poleward meridional flows play a major role  in the polarity reversals of the polar magnetic fields ($e.g.$, Wang 2004). The $\\approx$ 12-year and the $\\approx$ 22-year periodicities of the north-south difference in the  mean meridional motion of the spot groups may have a close relationship with the 11-year solar activity  (the emerging  magnetic flux) cycle and the 22-year solar magnetic cycle, respectively. Many of the other  periodicities found here  also exist in several activity phenomena (Knaack et al. 2005; Song et al. 2009, and references therein) and solar differential rotation determined from sunspot data (Javaraiah \\& Gokhale 1995, 1997b; Javaraiah \\& Komm 1999; Braj\\v{s}a et al. 2006). They may be closely related to the Rossby type waves that were discussed by Ward (1965) and others (e.g., Knaack et al. 2005; Chowdhury et al. 2009).    According to the well known Gnevyshev and Ohl rule (G-O rule), an odd cycle is stronger than its immediately  preceding even cycle (Gnevyshev \\& Ohl 1948). Cycle pair~22,23 violated the G-O rule. The duration  of the current cycle~23 is very long,  and during the declining phase of this cycle,  the activity in the southern hemisphere  is considerably stronger  than   in the northern hemisphere. All these properties of the cycle~23 could be strongly  related to the large and northbound   mean meridional motion of the  spot groups during this cycle. As already mentioned above, motions of  magnetic structures such as  sunspots mimic the motions of deeper layers of the Sun (see also Javaraiah \\& Gokhale 2002). Therefore, the magnetic flux cancellation/enhancement due to the mean meridional motion of sunspot groups  may take place in the subsurface layers of the Sun. By considering the  poleward  meridional plasma flows detected by surface  Doppler measurements and by helioseismology, the deeper counter-motion (suggested by the mean meridional motion of spot groups) might amplify the action of the near-surface dynamo in the southern hemisphere during cycle~23 for  causing stronger magnetic activity on this hemisphere."
    },
    "0911/0911.4107_arXiv.txt": {
        "abstract": "From high-resolution images of 23 Seyfert-1 galaxies at z=0.36 and z=0.57 obtained with the Near Infrared Camera and Multi-Object Spectrometer on board the {\\it Hubble Space Telescope} (HST), we determine host-galaxy morphology, nuclear luminosity, total host-galaxy luminosity and spheroid luminosity. Keck spectroscopy is used to estimate black hole mass (\\mbh). We study the cosmic evolution of the \\mbh-spheroid luminosity (\\ls) relation. In combination with our previous work, totaling 40 Seyfert-1 galaxies, the covered range in BH mass is substantially increased, allowing us to determine for the first time intrinsic scatter and correct evolutionary trends for selection effects.  We re-analyze archival HST images of 19 local reverberation-mapped active galaxies to match the procedure adopted at intermediate redshift. Correcting spheroid luminosity for passive luminosity evolution and taking into account selection effects, we determine that at fixed present-day V-band spheroid luminosity, \\mbh/\\ls $\\propto$ $(1+z)^{2.8 \\pm 1.2}$. When including a sample of 44 quasars out to $z=4.5$ taken from the literature, with luminosity and BH mass corrected to a self-consistent calibration, we extend the BH mass range to over two orders of magnitude, resulting in \\mbh/\\ls $\\propto$ $(1+z)^{1.4 \\pm 0.2}$. The intrinsic scatter of the relation, assumed constant with redshift, is 0.3$\\pm$0.1 dex ($<$0.6 dex at 95\\% CL). The evolutionary trend suggests that BH growth precedes spheroid assembly. Interestingly, the \\mbh-{\\it total host-galaxy} luminosity relation is apparently non-evolving. It hints at either a more fundamental relation or that the spheroid grows by a redistribution of stars. However, the high-z sample does not follow this relation, indicating that major mergers may play the dominant role in growing spheroids above $z \\simeq 1$. ",
        "introduction": "\\label{sec:intro}  Supermassive black holes (BHs) seem to be ubiquitous in the center of spheroids -- elliptical galaxies and classical bulges of spirals \\citep[e.g.,][]{kor95,fer05}. In the local Universe, tight empirical relations have been found between the mass of the BH (\\mbh) and the properties of the spheroid, i.e.~stellar velocity dispersion $\\sigma$ \\citep{fer00,geb00}, stellar mass \\citep[e.g.,][]{mar03}, and luminosity \\citep[e.g.,][]{har04}.  The tightness of these relations is surprising, given the very different scales involved -- from accretion onto the BH ($\\mu$pc scale), the dynamical sphere of influence of the BH (pc scale) to the size of the spheroid (kpc scale) -- and poses a challenge to any theoretical model explaining their origin.  In general, the correlations are believed to indicate a close connection between galaxy formation and evolution and the growth of the BH. A variety of theoretical models have been developed to explain the observed relations, involving galaxy mergers and nuclear feedback through quenching of star formation \\citep[e.g.,][]{kau00,vol03,cio07,hop07,dim08,hop09b}.  Measuring the evolution with redshift of these correlations constrains theoretical interpretations and provides important insights into their origin \\citep[e.g.,][]{cro06,rob06,hop07}. For quiescent galaxies, the biggest challenge is to measure the BH mass, given the pc-scale sphere of influence of the BH which needs to be resolved spatially through either gas or stellar dynamics (see \\citealt{gue09} and \\citealt{gra08} for a recent compilation and references therein; for a review see \\citealt{fer05} and references therein) or from X-ray spectroscopy probing the existence of a central temperature peak of the interstellar medium \\citep[e.g.,][]{bri99,hum06,hum08}. With current technology, direct quiescent black hole mass measurements are thus limited to nearby galaxies.  For active galaxies, for which nuclear luminosity is comparable to or larger than that of the host galaxy, the situation is virtually the opposite.  Estimating BH masses within a factor of 2-3 is fairly straightforward through empirically calibrated relations based on spectroscopic data measuring the kinematics of the broad-line region (BLR) \\citep[e.g.,][]{wan99,woo02,ves02,ves06,mcg08}. Unfortunately, the active galactic nuclei (AGN) often outshines the host galaxy, making it difficult to disentangle nuclear and host-galaxy light for an accurate measurement of the spheroid luminosity. Also, measuring $\\sigma$ from stellar absorption lines is hampered by the contaminating AGN continuum and emission lines. Different groups have tackled these problems in distinct ways, e.g.~by using the [\\ion{O}{3}] emission line width as surrogate of $\\sigma$ \\citep[e.g.,][]{shi03}, or by using gravitational lensing to super-resolve the host galaxies of quasars \\citep[e.g.,][]{pen06a,pen06b}.  Our group \\citep{tre04,woo06,tre07,woo08} has focused on Seyfert-1 galaxies - for which the nucleus is not as bright as for quasars - at moderate redshifts ($z=0.36$ and $z=0.57$, corresponding to look-back times of $\\sim$4-6 Gyrs). The non-negligible stellar light produces strong enough absorption lines to measure $\\sigma$ from unresolved spectra, as shown by \\citet{tre04} and \\citet[][hereafter Paper I \\& III]{woo06, woo08}. At the same time, high resolution {\\it Hubble Space Telescope} (HST) imaging allows for an accurate determination of the AGN luminosity (for an unbiased estimate of nuclear luminosity and hence \\mbh) and spheroid luminosity (to create the \\mbh-\\ls~relation; \\citealt[][hereafter Paper II]{tre07}).  We are thus able to simultaneously study both the \\mbh-$\\sigma$ and \\mbh-\\ls~relations, allowing us to distinguish mechanisms causing evolution in $\\sigma$ (e.g., dissipational merger events) and \\ls~(e.g.~through passive evolution due to aging of the stellar population, or dissipationless mergers).  Results presented in Paper I, II, and III suggest an offset with respect to the local relationships, which cannot be accounted for by known systematic uncertainties. At a given \\mbh, in the range 10$^8$-10$^9$ M$_{\\odot}$, spheroids had smaller velocity dispersion and spheroid mass 6 Gyrs ago ($z\\sim0.57)$, consistent with recent growth and evolution of intermediate-mass spheroids. Paper II concludes that the distant spheroids have to grow by $\\sim$60\\% in stellar mass ($\\Delta \\log M_{\\rm sph}$ = 0.20 $\\pm$ 0.14) at fixed black hole mass in the next 4 billion years to obey the local scaling relations if no significant BH growth is assumed, consistent with the relatively low Eddington ratios.  Indeed, the HST images reveal a large fraction of merging or interacting systems, suggesting that gas rich mergers will be responsible for the spheroid growth.  Although tantalizing, the results presented in our previous papers suffer from several limitations. Samples were small, and the local comparison sample of Seyferts measured in a self-consistent manner was even smaller than the distant sample, thus contributing substantially to the overall error budget. The limited range in black hole mass was insufficient to determine independently the offset of the scaling relation and its scatter, while taking into account selection effects. If the \\mbh-$\\sigma$ and \\mbh-\\ls~relations of active galaxies were not as tight as those for quiescent ones, selection effects could be mimicking evolutionary trends \\citep{tre07,lau07,pen07}.  To overcome these limitations, we have now doubled the sample size (from 20 in Paper II to 40 total here) and expanded the covered range of BH masses to lower masses (from $\\log$ \\mbh/$M_{\\odot} = 8-8.8$ in Paper II to $\\log$ \\mbh/$M_{\\odot} = 7.5-8.8$ here). We focus on the resulting BH mass - spheroid luminosity relation. The BH mass - $\\sigma$ relation will be presented in a separate paper (Woo et al.\\ 2009, in preparation).  We also analyze archival HST images of the sample of local Seyferts with reverberation-mapped (RM) \\mbh~in the same way as our intermediate-z objects, to eliminate possible systematic offsets. Finally, we combine our results with data compiled from the literature and treated in a self-consistent manner to extend the redshift range over which we study evolution. For conciseness the three samples will be referred to as ``intermediate-redshift'' sample, ``local'' sample, and ``high-redshift'' sample, respectively.  The paper is organized as follows. We summarize the properties of our intermediate-redshift Seyfert sample, observations, data reduction, and analysis in \\S~\\ref{sec:sample},~\\ref{sec:obs}, and~\\ref{sec:surface}. \\S~\\ref{sec:quan} summarizes the derived quantities, including the derivation of \\mbh~from Keck spectra. In \\S~\\ref{sec:comp}, we describe the local comparison sample consisting of reverberation-mapped AGNs, re-analyzed here, as well as the high-redshift comparison sample taken from the literature, calibrated for consistency with the other samples. We present our results in \\S~\\ref{sec:res}, including host-galaxy morphology and merger rates, the evolution of the \\mbh-\\ls~relation, a full discussion and treatment of selection effects, and a relation between BH mass and host-galaxy luminosity. We discuss the possible implications of our findings for the origin and evolution of the BH mass scaling relation in \\S~\\ref{sec:disc}.  A summary is given in \\S~\\ref{sec:sum}. In Appendix~\\ref{sec:mc}, we describe Monte Carlo simulations used to probe our analysis and determine errors. Appendix~\\ref{sec:sersic} discusses the choice of the S{\\'e}rsic index in the adapted 2D surface-brightness fitting procedure. Details on the re-analysis of the HST images of the local RM AGNs are given in Appendix~\\ref{sec:rm}. Throughout the paper, we assume a Hubble constant of $H_0$ = 70\\,km\\,s$^{-1}$\\,Mpc$^{-1}$, $\\Omega_{\\Lambda}$ = 0.7 and $\\Omega_{\\rm M}$ = 0.3. Magnitudes are given in the AB system \\citep{oke74}. ",
        "conclusions": "\\label{sec:disc}  \\subsection{The role of mergers} \\label{ssec:mergerd}  Theoretical studies generally invoke mergers to explain the observed scaling relations between BH mass and host-galaxy spheroid properties -- a promising way to grow both spheroid and BH.  In a simple scenario, spheroids grow by (i) the merging of the progenitor bulges (assuming that both progenitors have a spheroidal component), (ii) merger-triggered starbursts in the cold galactic disk, and (iii) by transforming stellar disks into stellar spheroids \\citep[e.g.,][]{bar92,mih94,cox04}, thus increasing the spheroid luminosity and stellar velocity dispersion.  The fueling of the BH, on the other hand, is triggered by the merger event as the gas loses angular momentum, spirals inward and eventually gets accreted onto the BH, giving rise to the bright AGN or 'quasar' period in the evolution of galaxies \\citep[e.g.,][]{kau00,dim05}. Eventually, if BHs are present in the center of both progenitor galaxies, they may coalesce.  In such a simple scenario, an evolution in the BH mass - spheroid luminosity relation is not necessarily expected: Both spheroid and BH grow from the same gas reservoir, and bulge stars added to the final spheroid followed the BH mass - spheroid luminosity relation prior to merging, so the relation will be preserved when the BHs coalesce.  However, while mergers provide a way to grow both spheroids and BHs, they may do so on very different timescales. Moreover, the merger history of galaxies varies, depending e.g.~on formation time and environment. Different types of merger, for example with a different relative role of dissipation \\citep[e.g.,][]{hop09a} have different effects on the growth of spheroid and BH: For a gas-rich major merger between an elliptical galaxy and a spiral galaxy - the latter without a (massive) BH -- the bulge grows more efficiently than the BH by the disruption of the stellar disk \\citep{cro06}.  In general, our images of the intermediate-z Seyfert galaxies support the merger scenario (see Fig.~\\ref{nicmos} and \\S~\\ref{ssec:merger}).  However, objects with evidence for merger/interaction do not form any particular outliers in the BH mass - spheroid luminosity relation (Fig.~\\ref{bhlv}). This may not be too surprising: For those objects for which we still see two separate galaxies in the process of merging, we fitted both separately and the bulge luminosity of the AGN host has not yet increased from the process of merging. Other objects with signs of interaction may be in a later evolutionary stage where the bulge luminosity has already increased and thus, the object falls closer to the local relation. Finally, mergers between similar objects would only move the system parallel to the local relation. In general, the effect of mergers on the measured bulge luminosity of an object depends on the type of the merger, the evolutionary stage of the merger, and the timescales involved to grow spheroid and BH. Such a detailed comparison of merger type and age with theoretical predictions is beyond the scope of this paper, given the small sample of merging objects and the limited information at hand.  Note that the fraction of apparently disturbed systems we find is not higher than that of a comparison sample of inactive galaxies at the same redshift (\\S\\ref{ssec:merger}). Thus, from our images alone, we cannot infer a causal link between a merger/interaction event and the AGN activity we observe. Instead, ``normal'' galaxies may have the same merger history, and ongoing interactions are not necessarily predictive of AGN activity. The role of mergers for the fueling of AGNs is debated in the literature \\citep[e.g.][]{san88,hec84,hut88,dis95,bah97,mcl99,can01,dun03,flo04,can07,urr08,ben08,vei09,tal09}. While there is little doubt that mergers are helpful, they are certainly not a sufficient condition, considering the numerous inactive interacting galaxies.\\footnote{However, this might also be due to the timescales involved, with the signs of interaction outliving the AGN activity (see also the case of present-day Type II AGNs; \\citealt[e.g.][]{cho09}).} Also, mergers may be necessary for the high-luminosity QSOs only while for Seyfert galaxies, secular evolution through processes such as bar instabilities may be the dominant effect in the evolution of these galaxies. We will come back to this issue in \\S~\\ref{ssec:mbhtotr}.  \\subsection{BH Mass - Spheroid Luminosity Relation} \\label{ssec:mbhlr}  Combining results of low-z, intermediate-z and high-z AGNs, treated in a self-consistent manner, we can estimate the intrinsic scatter of the \\mbh-\\ls~scaling relation and correct evolutionary trends for selection effects. We discuss scatter and evolution in the next two subsections.  \\subsubsection{Scatter of  \\mbh-\\ls} \\label{ssec:scatter}  The intrinsic scatter we find (0.3$\\pm$0.1 dex; $<$0.6 dex at 95\\% CL) is non-negligible. However, we assume the intrinsic scatter of the \\mbh-\\ls~relation to be non evolving. While it would be desirable to directly study the evolution of the scatter with redshift, this requires a larger sample than the one we have at hand. Actually, we might expect a larger intrinsic scatter at higher redshifts, given the different ways and timescales involved when growing spheroids and BHs through mergers. Indeed, for the local Universe, the observed tightness in the relations has been a challenge for theoretical studies.  It has been explained by self-regulated models of BH growth \\citep{hop09b} in which the energetic feedback of the AGN eventually halts accretion, preventing the BH from further growth and quenches star formation \\citep[e.g.,][]{cio97,cio01,sil98,mur05,dim05,saz05,hop05,spr05,dim08}.  Also, a significant fraction of the host galaxies of both our local RM AGN sample ($\\sim$9/19) and our intermediate-z sample ($>$15/40) are prominent late-type spirals of type Sa or later which have been found to have a larger intrinsic scatter than elliptical galaxies \\citep[e.g.,][for the \\mbh-$\\sigma$ relation: 0.44 dex when including all galaxies vs. 0.31 for elliptical galaxies only]{gue09}. As already discussed in paper II, the intermediate-z late-type spirals may eventually fall on the local relation later, through merging, in line with ``downsizing'' \\citep[e.g.~][]{cow96,bri00,kod04,bel05,noe07}: Less massive, blue galaxies merge at later times and arrive at the local relation by becoming larger, bulge-dominated red galaxies.  Also, at least some spiral galaxies may not have classical bulges, but pseudobulges which are characterized by surface-brightness profiles closer to exponential profiles, ongoing star formation or starbursts, and nuclear bars or spirals. It is generally believed that they have evolved secularly through dissipative processes rather than being formed by mergers \\citep[see e.g. review in][]{kor04}. BHs have been found to reside in galaxies without classical bulges which may not follow the same scaling relations \\citep[e.g.,][]{gre08}.  \\subsubsection{Evolution of \\mbh-\\ls~with redshift} \\label{ssec:evor}  To generalize our results and to facilitate comparison with theoretical and observational works, it is useful to estimate the evolution of the \\mbh - spheroid {\\it stellar mass} relation.  We can convert the observed evolution of \\mbh~- spheroid luminosity into that between \\mbh~and spheroid mass, if we assume that -- after correction for luminosity evolution -- the mass-to-light ratio does not change from sample to sample\\footnote{Unfortunately, spatially resolved color information for a more sophisticated estimation of the stellar mass of the bulge is currently not available.}.  Under this assumption, an offset of $\\Delta$\\mbh~at fixed \\ls~equals that at fixed $M_{\\rm star}$ and thus, \\mbh/\\ms~$\\propto$$(1+z)^{1.4 \\pm 0.2}$.  We are now in a position to make a broad range of comparisons. In the literature, the BH mass evolution is discussed quite controversially. \\citet{shi03} study the \\mbh-$\\sigma$ relation out to $z=3.3$, estimating \\mbh~from H$\\beta$ and $\\sigma$ from [\\ion{O}{3}] and find that the QSOs and their host galaxies follow the local relation. (Note, however, that using [\\ion{O}{3}] as a surrogate for $\\sigma$ can be problematic as [\\ion{O}{3}] is known to often have an outflow component; for a discussion see e.g. \\citet{gre05,kom07}.) A similar conclusion has been reached by \\citet{shen08} who study over 900 broad-line AGNs out to $z \\simeq 0.4$ from SDSS. \\cite{ade05} use the correlation length of 79 quasar hosts at $z \\sim 2-3$ to estimate the virial mass of the halo and the \\ion{C}{4} line width and UV flux at 1350\\AA~to estimate \\mbh. When comparing the resulting \\mbh~-$M_{\\rm halo}$ relation to the local one \\citep{fer02}, they find no evidence for evolution. In particular, they can rule out evolution of the form \\ms/\\ms~$\\propto$$(1+z)^{2.5}$ with $z=2.5$ at 90\\% CL, given their error bars.  Other observational studies find the same trend in evolution as we do, i.e.~that BHs are too massive for a given bulge mass or velocity dispersion at higher redshifts \\citep{wal04,shi06,mcl06,pen06b,sal07,wei07,rie08,rie09,gu09}. \\citet{mcl06}, for example, study radio-loud AGN ($0 < z < 2$) and find \\mbh/\\ms$\\propto$$(1+z)^{2.07 \\pm 0.76}$. \\citet{pen06b}, whose data, treated in a consistent manner to match our data set, are included in this study, rule out pure luminosity evolution and find that the ratio between \\mbh~and $M_{\\rm sph}$ was $\\sim$four times larger at $z \\sim 2-3$ than today. For 89 broad-line AGNs between $1<z<2.2$ in the zCOSMOS survey, \\citet{mer09} find \\mbh/\\ms$\\propto$$(1+z)^{0.74 \\pm 0.12}$. However, this fit refers to the total host galaxy instead of the spheroid component alone. At least some galaxies will have a non-negligible disk fraction, which, when taken into account, would result in a larger offset \\citep[see also][]{jah09}. For a sample of $\\sim$ 100 quasars selected to reside in elliptical hosts, \\citet{dec09} estimate that \\mbh/\\ms~was $\\sim$8 times larger at $z \\sim 3$ than today, i.e.~\\mbh/\\ms$\\propto$$(1+z)^{1.5}$. The evolution we find is also consistent with our previous results, within their much larger errors: \\mbh/\\ms$\\propto$$(1+z)^{1.5 \\pm 1}$ from the \\mbh-\\ls~relation in Paper II, whose data are included here, and $\\Delta \\log$ \\mbh$\\propto$$(1+z)^{3.1 \\pm 1.5}$ from the \\mbh-$\\sigma$ relation in Paper III.  From a theoretical perspective, the discussion on the evolution of the BH mass scaling relations is not any less controversial. \\citet{sha09}, for example, use the local velocity dispersion function (VDF) of spheroids, together with their inferred age distributions, to predict the VDF at higher redshifts. Using  the \\mbh-$\\sigma$ relation with a normalization allowed to evolve with redshift ($\\propto (1+z)^{\\delta}$), they infer the BH mass density and compare it to the accumulated BH mass density derived from the time integral of the AGN LF. They find a mild redshift evolution ($\\delta < 0.35$), excluding $\\delta > 1.3$ at more than 99\\% CL (with the possibility of a stronger evolution for the more massive BHs). Another study using fully cosmological hydrodynamic simulations of $\\Lambda$CDM following the growth of galaxies and supermassive BHs, as well as their associated feedback processes, finds only limited evolution in \\mbh~with a steepening at z=2-4 \\citep{dim08}. \\citet{mer04} expect a weak evolution of \\mbh/\\ms~$\\propto$ (1+z)$^{0.4-0.6}$, when fitting the total stellar mass and star formation rate density as a function of redshift and comparing that to the hard X-ray selected quasar luminosity function, assuming that BHs only grow through accretion. Such a slope is in agreement with work by \\citet{hop09a} who combine prior observational constraints in halo occupation models with libraries of high-resolution hydrodynamic simulations of galaxy mergers. Using semi-analytic models, \\citet{cro06} predicts an evolution of \\mbh/\\ms~$\\propto$ (1+z)$^{0.4-1.2}$. A more rapid evolution is predicted by \\citet{wyi03} who assume a self-regulated BH growth model and find \\mbh/\\ms~$\\propto$ (1+z)$^{1.5}$, similar to our observational result.  However, the great advantage of the study presented here are the high-quality images at hand, allowing for a detailed bulge-to-disk decomposition of the host galaxy of the low- and intermediate-z Seyfert-1 galaxies. Combining data from a large sample of active galaxies, covering a redshift range from the local Universe out to z=4.5, all treated in a consistent manner, results in smaller error bars on the predicted evolution  than previous studies. Moreover, it allows, for the first time, to correct evolutionary trends for selection effects. The evolution we find (\\mbh/\\ms$\\propto$$(1+z)^{1.4 \\pm 0.2}$) is indicative of BH growth preceding host-spheroid assembly.  Still, we did not take into account that the evolution may depend on BH mass (see also Paper III). Indeed, there are theoretical predictions that objects with higher BH (or bulge) masses evolve faster \\citep[e.g.,][]{hop09a}. For example, \\citet{dim08} find that when restricting their fits to objects with $M_{\\star}\\geq 5\\times 10^{10}\\ M_{\\odot}$, the relation has a slope of $\\sim 1.9$ at z=3-4 and $\\sim 1.5$ at z=2.  Unfortunately, our sample is too small to allow us to address this possibility. There may indeed be some evidence that the offset in BH mass is larger for objects with more massive BHs (Fig.~\\ref{bhlv}, upper left panel and Fig.~\\ref{offset}, right panel).  \\subsection{BH Mass - Galaxy Luminosity Relation} \\label{ssec:mbhtotr}  A different scenario seems to emerge when considering the relation between \\mbh~and {\\it total host-galaxy} luminosity (Fig.~\\ref{bhlv}, lower left panel).  This relation is almost non-evolving within the last six billion years.\\footnote{Note that there is insufficient information to constrain the intrinsic scatter.} Recently, \\citet{jah09} found qualitatively similar results for a small sample of ten AGNs at redshifts between 1 $< z <$ 2: They derive host-galaxy masses from colors based on ACS and NICMOS imaging, finding that they lie on the \\mbh-$M_{\\star, bulge}$ relation in the local Universe \\citep{har04}.  Such a non-evolving \\mbh-\\lh~relation can be interpreted twofold.  (a) The amount by which some of the more distant objects have to grow their spheroid is already contained within the galaxy itself, and the growth can be achieved by the redistribution of stars, i.e.~transforming disk stars into bulge stars. Such a redistribution can be the result of mergers or secular evolution, e.g.~bar instabilities \\citep[e.g.,][]{com81,vanb98,avi05,deb06} and torque-driven accretion \\citep[see e.g. review in][]{kor04} which may coincidentally be also the triggering mechanism for the BH activity we observe \\citep[e.g.,][]{shl93,ath03,dum07,haa09}. While not every object in the intermediate-z sample will experience a major merger in the last 4-6 billion years, secular evolution is a promising alternative way to grow the spheroidal components in these objects. But even if they do experience a major merger (as indeed evidenced for at least some objects in our sample), the role of the merger depends on the merger type as discussed above (e.g.~a merger between similar objects will simply move the system along the local relation). There may again be a dependency on BH mass: For the low-mass objects, the offset becomes almost negative, indicating that in the low-mass range, either the BH is, at the same time, still growing by a non-negligible amount (consistent with the higher Eddington ratio in the low-mass regime\\footnote{Of course, we may be biased against low-mass objects with low Eddington ratios.}) or that not all of the stellar mass will end up in the spheroid component. Indeed, for local RM AGNs, at least 6/19 objects reside in late-type host galaxies (preferentially those with lower BH masses).  (b) The relation between BH mass and host-galaxy luminosity (or mass) may be the more fundamental one. Indeed, this is predicted by \\citet{pen07}: In his thought experiment, he shows that a tight linear relation between \\mbh~and host-galaxy mass can evolve -- if the galaxy mass function declines with increasing mass -- due to ``a central-limit-like tendency for galaxy mergers, which is much stronger for major mergers than for minor mergers, and a convergence toward a linear relation that is due mainly to minor mergers''. Also, it is possible that BHs in late-type galaxies or galaxies without classical bulges, while not following the same \\mbh~scaling relations as spheroids (see discussion in \\S~\\ref{ssec:scatter}), they instead obey a more fundamental relation between BH mass and host-galaxy mass.  However, the relation between host-galaxy luminosity and \\mbh~seems to exist only up to z$\\lesssim$1: The offset for the high-z comparison sample does not decrease as the luminosity given by \\citet{pen06b} is already the total host-galaxy luminosity \\citep[the same is true also for the results by][]{mer09,dec09}. Along the line of argument of (a) above, the growth of the spheroid above a redshift of $z \\gtrsim 1$ cannot simply be achieved through secular evolution (with quasars being predominantly hosted by ellipticals), but instead, major mergers are needed. A major merger is more likely to happen for the high-z sample given the longer time span. Or, following (b), a relation between BH mass and host galaxy is already at place at $z \\lesssim 1$, but still evolving at earlier times. However, we cannot exclude that part of the difference is due to the difference in BH mass between the samples, with the high-z objects generally having larger BH masses.  In the end, the discussion boils down to the following question: What is the dominant mechanism that grows spheroids, and does it depend on spheroid mass and/or redshift? This is debated controversially in the theoretical literature. For example, based on their semi-analytic models, \\citet{par09} find that the majority of ellipticals and spirals never experience a major merger but rather, that they acquire their spheroid stellar mass through minor mergers or disc instabilities. \\citet{hop09c}, on the other hand, combine empirically constrained halo occupation distributions with high-resolution merger simulations, and find that major mergers dominate the formation of $\\sim$$L_{\\star}$ bulges and systems with higher B/T, but that lower-mass or lower B/T systems are preferentially formed by minor mergers. They predict that the major merger rate increases with redshift. Qualitatively, we can reconcile such a scenario with our results: Higher-mass objects and those at higher redshifts (i.e.~the majority of the high-z sample) form their spheroids preferentially through major mergers and are thus still evolving toward a \\mbh-\\lh~relation, while lower-mass and lower-z objects (i.e.~our intermediate-z sample) grow their spheroids through minor mergers or disk instabilities that redistribute the stars and thus, they fall on the \\mbh-\\lh~relation."
    },
    "0911/0911.4325_arXiv.txt": {
        "abstract": "Just comparing with the scenario that the $(3+1)$-dimensional ``real world'' of the Calabi-Yau compactification has a tremendous landscape, we conjecture that a $(4+1)$-dimensional holographic theory may also hold a landscape of its vacua. Unlike the traditional studies of the AdS/CFT phenomenology where the vacua are always constructive, we discuss the proper holographic vacua and their flux compactification, starting from some general compact Einstein manifolds. The proper vacua should be restricted by (i) a consistent worldsheet theory that possesses the superconformal symmetry, (ii) some definite symmetries to keep/break the corresponding symmetries of the dual field theory, (iii) certain brane/flux configurations to cancel anomalies, and (iv) stabilities. We consider diverse fundamental parameters of the dual field theory, fixed by some special vacuum moduli.  In an opposite way, if some field theory such as QCD holds an AdS dual, it may also possesses various fundamental parameters by an ``landscape'' of its vacuum. Different vacua may be adjacent with each other, and divided by domain walls. If the size of a single vacuum region is smaller than the visible universe, it may be testable. We discuss the consequences of this conjecture in the astrophysical environments, include but not limit to: (i) consistency with the critical energy density of the universe, (ii) the behaviors of cosmic rays, (iii) the stability and abundance of deuterons and other nuclei in the big-bang nucleosynthesis and the star burning scenarios, and (iv) the existence of strange/charm stars.  \\vspace{0.5cm} \\noindent \\emph{PACS}: 11.25.Wx, 98.80.Cq, 11.25.Tq, 11.25.Mj ",
        "introduction": "}  The anti-de Sitter/conformal field theory (AdS/CFT) correspondence, one of the most ambitious scenarios in string phenomenology, conjectures that a type IIB superstring theory on $\\mathrm{AdS}_5 \\times S^5$ is equivalent with a $\\mathcal{N} = 4$ $U(N_c)$ super Yang-Mills (SYM) theory in four-dimensions~\\cite{Maldacena:1997re}, or more generally a gravity theory on $\\mathrm{AdS}_{p+2} \\times \\mathcal{M}_q$ is dual to a $(p+1)$-dimensional boundary CFT~\\cite{Witten:1998qj,Aharony:1999ti}. The idea of ``holography''~\\cite{Susskind:1998dq} also pushes the applications of AdS/CFT to more realistic environments, such as QCD~\\cite{Peeters:2007ab,Mateos:2007ay}, or condensed matter systems~\\cite{Herzog:2009xv}.  Nevertheless, the compact dimensions (thus the ``vacua'') and their stabilization in AdS/CFT models, has seldom been studied systematically. On the one hand, theoretical researches always study some specific vacuum by constructive methods. For example, they break boundary supersymmetry by quotient spaces $S^5 / \\Gamma$~\\cite{Kachru:1998ys,Lawrence:1998ja}, or the conifold construction~\\cite{Klebanov:1998hh,Acharya:1998db,Morrison:1998cs}. On the other hand, phenomenological models which aim to approach the ``real world'' physics, always neglect the discussions of the compact dimensions directly. However, although difficult, the study of AdS/CFT vacua has no alternative but within the framework of flux compactification~\\cite{Grana:2005jc,Douglas:2006es}. Some founding works in this direction can be found in~\\cite{Aharony:2008wz,Polchinski:2009ch}.  The original studies of flux compactification, always aim to the Calabi-Yau threefolds. The main reason is that $\\mathrm{CY}_3$, which possesses a special holonomy of $SU(3) \\subset SO(6)$, can reduce the ten-dimension critical superstring theory to some four-dimensional effective field theory which possesses $\\mathcal{N} = 1$ supersymmetry. One of the properties of this scenario beyond one's expectation, is the tremendously abundant vacua, which mainly rise from the not-very-small Betti numbers $b_2$ and $b_3$ of $\\mathrm{CY}_3$, and the various possible fluxes wrapped on it; this set of vacua is always called a ``string landscape''~\\cite{Susskind:2003kw}. Different vacua in the landscape hold different fundamental parameters. It was argued that the number of consistent quasirealistic flux vacua may be greater than $10^{500}$~\\cite{Douglas:2003um}, and models has been constructed to solve the cosmological constant problem using this property~\\cite{Bousso:2000xa}.  In this paper, we conjecture that as an analog, a holographic theory may also hold a landscape of vacua. We verify this hypothesis in two different ways, from top-down and from bottom-up. Along the first root, we discuss properties of the compact manifolds, and the restrictions of them from physical purposes. For the uncompactified dimensions to be AdS, the compact manifold should be Einstein; thus, most of our discussions are within the framework of Einstein manifolds~\\cite{Besse:1987}. After then, we consider the possibilities and phenomenologies of various AdS vacua, especially the properties of domain walls separate them. We also discuss the possibilities for a non-conformal boundary field theory to hold a landscape. Along the opposite root, we studied the consequences of our conjecture, if QCD (as a non-conformal boundary field theory) holds a landscape of vacua. In this case, different vacua of QCD should possess different fundamental parameters, such as the quark masses $m_q$, the running coupling constant $\\alpha_S$, or the CP violating phase $\\theta_\\mathrm{QCD}$. Another vacuum with parameters different from ours may be testable; and we estimate this possibility within several astrophysical environments.  The organization of this paper is as follows. In Sec.~\\ref{sec:Einstein_manifold} we discuss several mathematical and physical issues that relate to the Einstein manifolds. We consider the symmetric conditions of the string worldsheet and the dual field theories, and the properties of wrapped branes and fluxes. We also consider the stability conditions of vacua topologies. In Sec.~\\ref{sec:holo_phenomenology}, we discuss the theoretical issues to approach a QCD landscape. We consider the possibilities to break CFT, the fundamental parameters a vacuum should determine, and the deduced parameters that may relate to applications/observations. In Sec.~\\ref{sec:astronomy}, we consider how a QCD landscape affects astrophysical observations. The applications are abundant, but the studies in our paper are only tentative. We summarize our results in Sec.~\\ref{sec:conclution}. Some mathematical supplements relatively independent to the main text are gathered in Appendix~\\ref{app:mathematics_addons}; and the validity of the orders of magnitudes estimations used in this paper, are reconsidered more carefully in Appendix~\\ref{app:order_estimations}. We gather these materials together, rather than write two separate papers from either the theoretical or the astrophysical aspects, because we think neither one alone is enough to make our conjecture reasonable; however, the two roots can in fact be read separately. We always denote the indices of the extended dimensions by $\\mu, \\nu$, of the compact dimensions by $m, n$, and of the entire target space by $M, N$. We set $\\hbar = k = c = 1$ for simplification throughout this paper. ",
        "conclusions": "}  In this paper, we discussed the possibilities that whether QCD, or more generally some holographic $(3+1)$-dimensional gauge theories, can possess a ``landscape'' of its vacua. We first limited our discussions to some boundary CFTs, for which the compact manifolds are Einstein. An Einstein 5- or 6-manifold $\\mathcal{M}_q$, which needs not to be homogeneous, may have some not-very-small Betti numbers $b_m$. Therefore, if fluxes wrap different patterns on cycles of it, the moduli are also different; a landscape of the holographic vacua should rise for this reason. We examined several relevant issues for this conjecture from the theoretical viewpoint. The geometry of $\\mathrm{AdS}_5 \\times \\mathcal{M}_q$ should also possess some symmetric conditions, such as the worldsheet superconformal symmetry, or definite supersymmetries of the dual field theories. However, it seems that even the $\\mathrm{AdS}_5 \\times S^5$ violates the one-loop worldsheet conformal symmetry. The isometry group of $\\mathcal{M}_q$ decides the R-symmetry of the holographic theory, which is not imposed in our case. In addition, we considered the anomaly cancelation, the no-go theorem, the brane/flux configurations, and the vacua stabilities. We focused on the $(2+1)$-dimensional domain walls within $\\mathrm{Mink}_4$, and also the $(3+1)$-dimensional domain walls filling the radial $\\mathrm{AdS}_4$ inside $\\mathrm{AdS}_5$. We considered the possibilities to break the $\\mathrm{AdS}_5$ geometry, or equivalently the conformal symmetry of the boundary theory, and argued that the confinement-deconfinement phase transition may not affect the vacua configurations. After then, we applied our conjecture of the ``holographic landscape'' directly to QCD, for which the fundamental parameters such as $m_q$, $\\alpha_S$ or $\\theta_\\mathrm{QCD}$ should depend on the moduli of the compact dimensions. We studied the properties of the domain walls, such as their masses, thicknesses, and their dynamical properties; they may be both relativistic and non-relativistic. We discussed how this ``QCD landscape'' affects nuclear physics; they may alter the hadronic masses, the reaction rates, and the stabilities of the nuclei.  In an opposite way, we studied how can a ``QCD landscape'' affects the astrophysical observations. As domain walls may be non-relativistic, another vacuum of QCD may be within the visible universe; if they are not, we can also consider the properties of the other multiverses of QCD, just as what is done in~\\cite{Bousso:2008bu,Bousso:2009gx}. We first considered whether the mass contribution of the domain walls exceeds the critical density of the universe. Most of the case, domain walls are really dangerous; however, it is not enough reasonable that they should be completely ruled out. We then discussed the properties of cosmic rays affecting by this landscape, include the GZK cutoff, the neutron lifetime, and an alternative explanation of the coincidence of the cosmic ray anisotropic spatial distribution with the supergalactic plane. The GZK cutoff depends on, but is not very sensitive to, a different QCD vacua; the variation of the neutron lifetime seems helpless to explain the observations. The alternative explanation may loose the constraints for the origins of the UHECRs; however, as too ambitious the scenario is, it needs to be studied more seriously. We also considered how the QCD landscape affects BBN and the star burning scenarios. As only the abundance of deuterium ${}^2\\mathrm{D}$ can be measured in far away part of the universe, only it can restrain the QCD landscape; the constraint is really loose. A different QCD vacuum can also affect the stellar evolution; for example, the time spent for the proton-proton chain reaction depends on the deuterium mass $m_\\mathrm{D}$ very sensitively. The Chandrasekhar mass limits of white dwarfs and neutron stars are also altered, but not very sensitive to the QCD landscape. In addition, whereas the ``strange quark matter'' may be the true matter ground state of our QCD~\\cite{Witten:1984rs}, charm or $(ud)$ QGP may be the ground states of other QCDs; thus, as an alternative, charm stars or $(ud)$ quark stars may exist in those QCDs. However, the mass-radius relationships and the surface properties of those objects seem really similar to our strange stars.  Recently, Denef and Hartnoll discussed the landscape in condensed matter systems, which results a statistical distribution of critical superconducting temperatures~\\cite{Denef:2009tp}. In here, we compare briefly the original ``string landscape'', their ``atomic landscape'', and our ``QCD landscape''. First, whereas the relativistic quantum critical theories are conformal, our QCD landscape should break conformal symmetry by some mechanics. Second, the original Calabi-Yau compactification possesses a $\\mathcal{N} = 1$ supersymmetry, the ``atomic landscape'' in M-theory by the Sasaki-Einstein 7-manifolds holds a $\\mathcal{N} = 2$ supersymmetry; however, it seems that our ``QCD landscape'' is not restricted by supersymmetry conditions. Third, as the ``atomic landscape'' discussed in~\\cite{Denef:2009tp} is limited to the Freund-Rubin compactification, only the background electromagnetic four-flux contributes to the vacua field equations; thus their statistics of the landscape rises from the different topologies of the compact dimensions. However, the original Calabi-Yau vacua, or the QCD vacua discussed in this paper, are more abundant because of different flux configurations."
    },
    "0911/0911.3396_arXiv.txt": {
        "abstract": "We discuss the universal relation between density and size of observed Dark Matter halos that was recently shown to hold on a wide range of scales, from dwarf galaxies to galaxy clusters. Predictions of $\\Lambda$CDM N-body simulations are consistent with this relation. We demonstrate that this property of $\\Lambda$CDM can be understood analytically in the secondary infall model. Qualitative understanding given by this model provides a new way to predict which deviations from $\\Lambda$CDM or large-scale modifications of gravity can affect universal behavior and, therefore, to constrain them observationally. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0911/0911.2997_arXiv.txt": {
        "abstract": "By means of a semiclassical analysis we show that the trace anomaly does not affect the cosmological constant. We calculate the evolution of the Hubble parameter in quasi de Sitter spacetime, where the Hubble parameter varies slowly in time, and in FLRW spacetimes. We show dynamically that a Universe consisting of matter with a constant equation of state, a cosmological constant and the quantum trace anomaly evolves either to the classical de Sitter attractor or to a quantum trace anomaly driven one. There is no dynamical effect that influences the effective value of the cosmological constant. ",
        "introduction": "T^{\\mu}_{\\phantom{\\mu}\\mu} = 0 \\,. \\end{equation} As an explicit example, consider a conformally coupled scalar field. In quantum field theory the stress-tensor is promoted to an operator. A careful renormalisation procedure renders its expectation value $\\langle \\hat{T}^{\\mu}_{\\phantom{\\mu}\\nu}\\rangle$ finite. However, inevitably, the renormalisation procedure results in general in a non-vanishing trace of the renormalised stress-energy tensor: \\begin{equation}\\label{intro2} \\langle \\hat{T}^{\\mu}_{\\phantom{\\mu}\\mu}\\rangle \\neq 0 \\,. \\end{equation} Classical conformal invariance cannot be preserved at the quantum level. In short, this is the trace anomaly. An important question that immediately arises is how this non-zero trace affects the evolution of the Universe through the Einstein field equations. In particular, one could wonder whether it influences the effective value of the cosmological constant.  It has been argued (see \\cite{Antoniadis:2006wq} and references therein) that the trace anomaly could potentially provide us with a dynamical explanation of the cosmological constant problem. In brief, the line of reasoning is as follows. Firstly, a new conformal degree of freedom is introduced as $g_{\\mu\\nu}(x)=\\exp[2\\sigma(x)]\\overline{g}_{\\mu\\nu}(x)$. Furthermore, the trace anomaly stems from a non-local effective action that generates the conformal anomaly by variation with respect to the metric \\cite{Riegert:1984kt}. The authors of \\cite{Antoniadis:2006wq} then argue that the new conformal field should dynamically screen the cosmological constant, thus solving the cosmological constant problem.  The proposal advocated in \\cite{Antoniadis:2006wq} is very appealing. Before studying the effect of this new conformal degree of freedom, we feel that firstly a proper complete analysis of the dynamics resulting from the effective action of the trace anomaly should be performed. This is what we pursue in this contribution.  We argue in favour of a semiclassical approach to examining the connection between the cosmological constant and the trace anomaly: we take expectation values of inhomogeneous quantum fluctuations with respect to a certain state to study its effect on the background spacetime. Phase transitions aside, quantum fluctuations affect the background \\textit{homogenously}. Moreover, in accordance with the Cosmological Principle, inhomogeneous fluctuations of the metric tensor and in particular of the conformal part of the metric tensor are observed to be small at the largest scales, comparable to the Hubble radius (and expected to be small also in the early Universe). We do not consider a new, conformal degree of freedom. Hence, we do certainly not exclude any possible effect this (inhomogeneous) conformal degree of freedom might have on the cosmological constant. However, it is plausible that in order to address the link between the cosmological constant and the trace anomaly, a semiclassical analysis suffices.  A final note is in order regarding an application of the conformal anomaly: trace anomaly induced inflation. In the absence of a cosmological constant, the trace anomaly could provide us with an effective cosmological constant \\cite{Starobinsky:1980te, Hawking:2000bb, Pelinson:2002ef,Shapiro:2002nz, Shapiro:2003gm, Pelinson:2003gn, Shapiro:2008sf}. If one includes a cosmological constant, the theory of anomaly induced inflation is plagued by instabilities, which we will also come to address. Improving on e.g. \\cite{Pelinson:2002ef, Brevik:2006nh}, we incorporate matter with a constant equation of state in the Einstein field equations. ",
        "conclusions": "We have studied the dynamics of the Hubble parameter both in quasi de Sitter and in FLRW spacetimes including matter, a cosmological term and the trace anomaly. We have seen that there is no dynamical effect that influences the effective value of the cosmological constant, i.e.: the classical de Sitter attractor. Based on our semiclassical analysis we thus conclude that the trace anomaly does \\textit{not} solve the cosmological constant problem.  Ostrogradsky's theorem merits another remark. Clearly, including the higher derivative contributions in FLRW spacetimes modifies the dynamics of the Hubble parameter significantly: attractors, that were stable in the absence of higher derivatives, under certain conditions destabilise. We do not know which of the two approaches is correct. Discarding these higher derivatives and studying the trace anomaly in quasi de Sitter spacetime would seem plausible.  Finally, one could wonder whether the quantum anomaly driven attractor is physical. The quantum attractor is of the order of the Planck mass $M_{\\mathrm{pl}}$, so only when the matter in the early Universe is sufficiently dense, $H \\simeq \\mathcal{O}( M_{\\mathrm{pl}})$. We then expect to evolve towards the quantum attractor. However, at these early times we also expect perturbative general relativity to break down. Hence, this attractor may be seriously affected by quantum fluctuations or it might even not be there."
    },
    "0911/0911.0395_arXiv.txt": {
        "abstract": "{} {We report on the third phase of our study of the neutrino-cooled hyperaccreting torus around a black hole that powers the jet in Gamma Ray Bursts. We focus on the influence of the black hole spin on the properties of the torus.} {The structure of a stationary torus around the Kerr black hole is solved numerically. We take into account the detailed treatment of the microphysics in the nuclear equation of state that includes the neutrino trapping effect. } {We find, that in the case of rapidly rotating black holes, the thermal instability discussed in our previous work is enhanced and develops for much lower accretion rates. We also find the important role of the energy transfer from the rotating black hole to the torus, via the magnetic coupling. } {} ",
        "introduction": "Observed Gamma Ray bursts (GRBs) form at least two classes of events. Long GRBs ($T_{90} > $ 2 s) are the the result of a collapsar in the core of an exploding massive star, as supported by the detections of their afterglows. Shorter bursts ($T_{90} < $ 2 s) are likely to be due to the coalescence of a binary system containing two compact stars, as observationally supported by the localization of their afterglows at the outskirts of their host galaxies (see e.g. the review of Zhang 2007). Moreover, it is possible that a third class of events exists, namely a fraction of the shortest bursts ($T_{90} < $ 0.1 s), that can be produced via the evaporation of the primordial black holes (Cline et al. 2005).   Both the two main classes of bursts involve a stage of a newly formed black hole surrounded by an accreting torus, responsible for the production of a collimated jet. This jet is a source of a highly beamed emission observed in the gamma ray band. The accreting torus is either fed by the material from the stellar envelope, which part falls back after the hypernova explosion, or consumes the material from the disrupted neutron star debris. The torus accretion proceeds with highly hyper-Eddington rates, on the order of a solar mass per second. Such hyper-accreting tori have been already discussed in a number of works (e.g. Popham et al. 1999; Di Matteo et al. 2002; Kohri et al. 2005; Chen \\& Beloborodov 2007; Janiuk et al. 2004; 2007; Lei et al. 2009).  In the present paper we expand on our previous work and we further investigate the properties and evolution of such hot and dense accretion tori. At the extreme densities and temperatures, determined by the hyper-Eddington accretion rate, the torus is cooled mainly by neutrino emission produced primarily by electron and positron capture on nucleons ($\\beta$ reactions). The model we develop here was first described in Janiuk et al. (2004) for the case of a non-rotating black hole and the nuclear matter with a simplified equation of state. We further developed this model in Janiuk et al. (2007), considering the more elaborate equation of state to account for the microphysics of the accreting plasma. In the latter work, we solved for the disk structure and its time evolution by introducing the new EOS which includes photodisintegration of helium, the condition of $\\beta$ equilibrium, and neutrino opacities. We self-consistently calculated the chemical equilibrium in the gas consisting of helium, free protons, neutrons, and electron-positron pairs and compute the chemical potentials of the species, as well as the electron fraction throughout the disk. We found an important property of our solutions: for very large accretion rates, the torus becomes viscously and thermally unstable in a narrow range of radii near the central black hole. The instability results as the intrinsic property of the torus, when the helium is fully photodisintegrated in the neutrino opaque plasma.  In the present work, we further expand our model to account for the black hole rotation. The rotating black hole should naturally be produced in the center of a collapsar or after the merger event, when the collapsing material possesses a large amount of angular momentum. As a consequence, the black hole spin is a plausible mechanism that helps to launch the jets emitted in a narrow cone along the rotation axis. As shown in Janiuk et al. (2008), especially for the long duration GRBs the BH spin is required to sustain the central engine activity and the jet production for a sufficiently long time. Because these long GRBs often exhibit very variable temporal structure (e.g. Beloborodov, Stern \\& Svensson 2000), it seems very important to study a possible instability mechanism that may account for the variability. We check here, for what black hole spin the instability may operate near the central black hole, depending on the accretion rate and viscosity in the disk. We also check what effects may be imposed on the torus, and its unstable strip, by the additional heating via the magnetic fields, threaded by the rotating black hole. The content of the article is as follows. First, we discuss the basic assumptions and present the main equations of the model. Second, we introduce the changes and corrections to the model that are the result of a black hole rotation with an arbitrary spin. Third, we introduce the energy extraction from the rotating black hole via the magnetic field, as an additional physical process that may operate in the gamma ray burst central engine. In Section \\ref{sec:results} we present the results of our calculations, and in Section \\ref{sec:diss} we discuss the results and conclude. ",
        "conclusions": "\\label{sec:diss}  The importance of black hole spin in the accretion disk systems and the formation of jets, also in the presence of the magnetic fields, has been studied by many authors. It has been discussed e.g. in the review by Blandford (1999). In the standard accretion disk theory (Shakura \\& Sunyaev 1973) it is assumed that there is no coupling between the black hole and the disk. However, if the magnetic field is considered, such a coupling may be present. Contrary to the Blandford-Znajek process (in which the magnetic filed lines threading the black hole may close on a very remote load, far away from the black hole), the magnetic field lines coupled to the accretion disk are closed in the vicinity of the inner radius. In this process, the energy and angular momentum are extracted from the black hole and transferred to the disk, if the hole rotates faster than the disk (MacDonald \\& Thorne 1982; Li 2000). Such coupling may have much higher efficiency in extracting the black hole rotation energy than the Blandford-Znajek mechanism, as was first quantitatively estimated by Li (2002).  Let us also remind here after Li (2002), that in this model the torque produced by the magnetic coupling with the black hole propagates only outwards, and no torque is imposed at the inner boundary of the disk. The non-zero torque at the inner boundary is possible if the disk is magnetically connected with the plasma in the transition region between the disk and the black hole (e.g. Krolik 1999), which we do not consider here. Another important assumption behind the present model is that about the quasi steady-state of the disk and magnetic field configuration. Under this assumption, the magnetic field lines are not tangled by the rotation and MHD instabilities of the Balbus \\& Hawley (1991) are not taken onto account. In our present work we hold this assumption, after Li (2002) and Wang et al. (2002; 2003a,b). The evolution of the coronal magnetic field threading the differentially rotating disk was studied however in the literature, e.g. by Lovelace et al. (2002) in the context of Poynting flux dominated jets in quasars, and in Uzdensky (2004; 2005) by numerically solving the Grad-Shafranov equation in the Schwarzschild and Kerr metric.  In the context of Gamma Ray Bursts, the black hole spin interaction with the magnetized torus was studied by van Putten et al. (2004) These authors take into account the dissipation of energy in the torus not only via the thermal and neutrino emission, but also through the gravitational waves and winds. They estimate the lifetime of the central engine powered by rapidly spinning black hole, whose spin-down is governed by the surrounding magnetic field coupled to the inner torus. In the present model, we neglect the gravitational waves emission and we concentrate on the structure of the neutrino-cooled disk. We restrict ourselves to the steady-state solution, however we note that the lifetime of the long gamma ray burst central engine is defined by the period of the rapid spin of the black hole, which may decrease with time during the engine activity (see also Janiuk et al. 2008).   In this article, we study the model of a stationary neutrino cooled disk around a rotating Kerr black hole. We confirmed the presence of the thermal-viscous instability in the region of the large neutrino opacity and efficient helium photodisintegration. The instability was first studied in Janiuk et al. (2007) in the time-dependent model for the Schwarzschild black hole. It was found to produce dramatic changes in the density and temperature profiles, as well as a rapid time variability of the disk emission. Here, we quantitatively described the role of the black hole spin and viscosity in the disk. We found that in the Kerr black hole disks the large spin parameter results in a larger size of an unstable strip, which appears even for a moderate accretion rate. The critical accretion rate is anticorrelated with the BH spin required for the onset of the instability. In the extreme BH case, the instability close to the inner edge may be effective even for $\\dot M=0.5 M_{\\odot}$ s$^{-1}$. On the other hand, the more viscous is the disk, the smaller in size is the unstable strip and the weaker should be the instability.  The time dependent calculations were not performed in the present work, because they would require to compute time evolution of the BH spin and a moving radial grid, which is intended to be studied in the future work. However, we studied the influence of the black hole rotation on the disk structure considering our stationary model. This was done by means of the magnetic transfer of energy from the rotating black hole to the disk, and vice versa.  The first important finding in this work is that thermal instability occurs for quite low accretion rates, on the order of a fraction of solar mass per second, provided the black hole rotates very fast. This can be relevant for the origin of the variable energy input to the GRB jets not only in case of the short GRBs, but also for long ones. The latter are presumably powered by the accretion of the fallback material from the massive star envelope onto the newly born black hole after the hypernova explosion (e.g. Mac Fadyen \\& Woosley 1999). In these models, the accretion rate is found to be substantially lower than in the merger of two neutron stars, and therefore the event can last for much longer time to power the GRB for several hundreds of seconds. On the other hand, the massive star that is a progenitor of a hypernova is a rotating star and therefore a natural consequence of the explosion should be the rapidly rotating black hole in the center.  As was recently shown by Janiuk, Moderski \\& Proga (2008), the large black hole spin is a key ingredient for the occurrence of the longest duration GRBs, and is required for the GRB central engine to operate for a time of the order of hundreds of seconds. The present study leads us therefore to a general conclusion that the longest GRBs, that are powered first by the neutrino annihilation and later by the black hole spin and require $a>0.9$, should be very variable at the initial phase of the prompt emission. Later on, when the accretion rate drops, the variability may be suppressed and what is observed is a smooth tail in the gamma ray lightcurve. On the other hand, the period of the highly variable GRB emission may be somewhat extended if the black hole is being spun up by the accretion. The issue discussed in the literature was whether in such a case the spin equilibrium value can be achieved at $a=0.998$ (Thorne 1974), or a less value (Gammie, Shapiro \\& McKinney 2004).  In our study, we also considered the heating of the innermost regions of the torus due to the magnetic coupling with the rotating black hole. We found that this process can have some influence on the thermal instability, because the size and position of the unstable region are somewhat changed. Also, the profiles of density and temperature in the unstable region are modified. The magnetic coupling enhances the change in density profile, however it weakens the change in temperature profile. The energy transfer to the disk from the rotating black hole does not suppress the thermal instability. This finding also supports our conclusion that the variable and long duration GRBs are powered by the rapidly spinning black holes."
    },
    "0911/0911.2673_arXiv.txt": {
        "abstract": "UV radiation from massive stars is thought to be the dominant heating mechanism of the nuclear ISM in the late stages of evolution of starburst galaxies, creating large photodissociation regions (PDRs) and driving a very specific chemistry. We report the first detection of PDR molecular tracers, namely HOC$^+$, and CO$^+$, and confirm the detection of the also PDR tracer HCO towards the starburst galaxy NGC\\,253, claimed to be mainly dominated by shock heating and in an earlier stage of evolution than M\\,82, the prototypical extragalactic PDR. Our CO$^+$ detection suffers from significant blending to a group of transitions of $^{13}$CH$_3$OH, tentatively detected for the first time in the extragalactic interstellar medium. These species are efficiently formed in the highly UV irradiated outer layers of molecular clouds, as observed in the late stage nuclear starburst in M\\,82. The molecular abundance ratios we derive for these molecules are very similar to those found in M\\,82. This strongly supports the idea that these molecules are tracing the PDR component associated with the starburst in the nuclear region of NGC\\,253. The presence of large abundances of PDR molecules in the ISM of NGC\\,253, which is dominated by shock chemistry, clearly illustrates the potential of chemical complexity studies to establish the evolutionary state of starbursts in galaxies. A comparison with the predictions of chemical models for PDRs shows that the observed molecular ratios are tracing the outer layers of UV illuminated clouds up to two magnitudes of visual extinction. We combine the column densities of PDR tracers reported in this paper with those of easily photodissociated species, such as HNCO, to derive the fraction of material in the well shielded core relative to the UV pervaded envelopes. Chemical models, which include grain formation and photodissociation of HNCO, support the scenario of a photo-dominated chemistry as an explanation to the abundances of the observed species. From this comparison we conclude that the molecular clouds in NGC\\,253 are more massive and with larger column densities than those in M\\,82, as expected from the evolutionary stage of the starbursts in both galaxies. ",
        "introduction": "Intense UV radiation from massive stars is one of the main mechanisms responsible for the heating of the interestelar medium in the nuclear region of starburst galaxies. This mechanism is particularly important in the latest stages of starburst (SB) galaxies where the newly formed massive star clusters are responsible for creating large photodissociation regions (PDRs). This is the case for the prototypical SB galaxy M\\,82, where the large observed abundances of molecular species such as HCO, HOC$^+$, CO$^+$, and H$_3$O$^+$ are claimed to be probes of the high ionization rates in large PDRs formed as a consequence of its extended evolved nuclear starburst \\citep{Burillo02,Fuente06,VdTak08}.   Observational evidences point to a significant enhancement in the abundance of HOC$^+$ in regions with large ionization fractions. The abundance ratio [HCO$^+$]/[HOC$^+$]$=270$ is found in the prototypical Galactic PDRs of the Orion Bar \\citep{Apponi99}. Similar or even lower abundance ratios are observed in the PDRs NGC\\,7023 \\citep[50-120,][]{Fuente03}, Sgr\\,B2(OH) and NGC\\,2024 \\citep[360-900,][]{Ziurys95,Apponi97}, and the Horsehead \\citep[75-200][]{Goico09}, as well as in diffuse clouds \\citep[70-120,][]{Liszt04}. This is in contrast with the much larger ratios of $\\gg 1000$ found in dense molecular clouds well shielded from the UV radiation. However, these low HCO$^+$/HOC$^+$ ratios are not found in other galactic PDRs. Large values of this ratio of $\\gtrsim2000$ are found in the PDRs M17-SW, S140, and NGC2023 \\citep{Apponi99,Savage04}. The HCO molecule has also been observed to be a particularly good tracer of the PDR interfaces. Low ratios of [HCO$^+$]/[HCO]$\\sim2.5-30$ are found in prototypical Galactic PDRs \\citep{Schene88,Schilke01}. The large HCO abundance ($>10^{-10}$) altogether with the low ratio [HCO$^+$]/[HCO]$\\sim1$ in the Horsehead PDR is claimed to be a diagnostic for an ongoing FUV-dominated photochemistry \\citep{Gerin09}. CO$^+$ is also claimed to be particularly prominent in the chemical modeling of PDRs and high abundances of this molecule appear to be correlated to similar enhancements of HOC$^+$ \\citep{Sternberg95,Savage04}. [CO$^+$]/[HOC$^+$] ratios in the range of 1-10 are observed in a number of PDRs \\citep{Savage04}, but only of $\\gtrsim 0.1.$ towards the Horsehead PDR \\citep{Goico09}.  As mentioned above, this set of PDR probes has been extensively studied towards M\\,82. However, no such complete studies have been carried out towards other prototypical galaxies, but for the detection of HCO and HOC$^+$ towards NGC\\,1068 \\citep{Usero04} and H$_3$O$^+$ in Arp\\,220 \\citep{VdTak08}. M\\,82 and NGC\\,253 are the brightest prototypes of nearby SB galaxies, at a similar distance and showing very similar IR luminosities and star formation rates of about $\\sim3\\,M_\\odot \\rm yr^{-1}$ \\citep{Ott05,Minh07}.\\ However, both galaxies show very different chemical composition. The chemistry and to a large extend the heating in the central region of NGC\\,253 is believed to be dominated by large scale low velocity shocks \\citep{Martin06b}. The similar chemical composition found in the nuclear region of NGC\\,253 to that in Galactic star forming molecular complexes points to an earlier evolutionary stage of the starburst in this galaxy than that in M\\,82 \\citep{Martin03,Martin05,Martin06b}.  Furthermore, our recent observations of the PDR component as traced by the easily photodissociated HNCO molecule towards a sample of galaxies \\citep{Martin08} showed the non-detection of HNCO in M\\,82, at a very low abundance limit. This low HNCO abundance supports the scenario that the PDR chemistry dominates the molecular composition of the ISM in this galaxy. However, from the HNCO measured abundance in NGC\\,253, it would be placed in an intermediate stage of evolution where photodissociation should be starting to play a significant role in driving a UV-dominated chemistry which has not been yet identified towards this galaxy. The presence of a significant PDR component in NGC\\,253 claimed from the HNCO abundance is also inferred from the similar intensity of the atomic fine structure line intensities from PDR tracers like CII and OI \\citep{Carral94,Lord96} observed in both M\\,82 and NGC\\,253.  In this paper we present the first detection of PDR molecular tracers HOC$^+$ and CO$^+$, and confirm the detection HCO \\citep[tentatively detected by][]{Sage95} in the central region of NGC\\,253 which allows the evaluation of the influence of the photodissociation radiation in the nuclear ISM of this SB galaxy. The results presented here support the scenario of the presence of a significant PDR component and clearly show the potential of molecular complexity in estimating the contribution of the different heating mechanisms of the ISM in the nuclei of galaxies. ",
        "conclusions": "The comparison of model predictions with the observations presented show that the abundance of the species observed in this work towards NGC\\,253, namely HCO$^+$, CO$^+$, and HCO, are most efficiently formed in the outer region of the molecular clouds where the gas is highly irradiated by the incident UV photons from massive stars. The high molecular abundances derived for these species in NGC\\,253 suggest that the PDR component in this galaxy is  similar to that found in M\\,82, claimed to be the prototype of extragalactic PDR. The abundance ratios found for this limited sample of galaxies are of the same order as those observed towards galactic PDRs, which stress the importance of photo-dominated chemistry in galaxy nuclei. Large amounts of molecular material are affected by photodissociation not only in NGC\\,253, but also towards the star forming regions around the Seyfert 2 nuclei in NGC\\,4945 and NGC\\,1068. This is consistent with the HNCO/CS ratio in these galaxies which suggest that a fraction of HNCO has been photodissociated in PDRs. The combination of the observations of HCO, HOC$^+$ and CO$^+$ with that of HNCO seems to confirm that their abundances reflect the evolutionary stage of the starbursts in these galaxies. Although photodissociation is the most likely scenario for the enhancement of the observed reactive ion in starburst environments, X-ray dominated chemistry has been claimed to be responsible for the high abundances observed around AGNs in circunnuclear disk of NGC\\,1068 \\citep[HOC$^+$][]{Usero04} and towards the ultra luminous infrared galaxy Arp\\,220 \\citep[$\\rm H_3O^+$][]{VdTak08}.  Therefore, M\\,82 is still outstanding not only as a PDR dominated galaxy, but by the underabundance of complex molecules such as CH$_3$OH, HNCO or SiO \\citep{Mauers93,Martin06a,Martin06b}, evidence for the lack of large amounts of dense molecular material which would potentially fuel its nuclear starburst as compared to other starburst galaxies like NGC\\,253 \\citep{Martin09}. Our data in combination with the HNCO abundances \\citep{Martin09} indicate that the molecular clouds in M\\,82 are different from those in NGC253. Although having a similar overall PDR component, the clouds in NGC\\,253 have to be more massive and have larger column densities those in M\\,82."
    },
    "0911/0911.0506_arXiv.txt": {
        "abstract": "Using archival data from the Hubble Space Telescope, we study the quantitative morphological evolution of spectroscopically confirmed bright galaxies in the core regions of nine clusters ranging in redshift from $z = 0.31$ to $z = 0.84$.  We use morphological parameters derived from two dimensional bulge-disk decomposition to study the evolution.  We find an increase in the mean bulge-to-total luminosity ratio $B/T$ as the Universe evolves. We also find a corresponding increase in the fraction of early type galaxies and in the mean S\\'ersic index. We discuss these results and their implications to physical mechanisms for evolution of galaxy morphology. ",
        "introduction": "Clusters serve as laboratories for investigating the dependence of galaxy morphology on the density of the environment.  \\citet{dre80} found that the fraction of early type galaxies increases with local density of galaxies.  He also found that around 80\\% of galaxies in nearby clusters are of the early type.  In recent years, deep, high resolution imaging by the \\textit{Hubble Space Telescope (HST)} has helped to extend the study of galaxy morphology as a function of the environment to $z \\sim 1$.  The observed fraction of different morphological types can be explained in terms of both 'nature' and 'nurture' scenarios.  The former argues that a galaxy's morphological type is determined by initial conditions at formation \\citep[e.g. ][]{egg62}, while the latter depends on the influence of the environment and of secular evolution for determining the final morphological type \\citep[e.g. ][]{too77}. Numerical simulations play an important role in investigating the relative importance of the two scenarios under different physical conditions.  Following Dressler's pioneering work, many observational studies have been done to measure and understand the morphology density relation (MDR), and the dependence of morphological type on distance of the galaxy from the cluster centre \\citep{dre97,got03,smi05,pos05,hol07,tre03,cap07}.  \\citet{smi05} found that the early type fraction is constant in low density environments over the last 10 Gyr, but there is significant evolution in this fraction in higher density regions. According to \\citet{smi05}, this suggests that most of the ellipticals in clusters formed at high redshift, and the increase in the fraction of early type galaxies is because of the physical processes in dense regions which transform disk galaxies with ongoing star formation to early types. Dissipationless merging of cluster galaxies may also be responsible for this increase.  There is increasing evidence to show that massive ellipticals formed by dissipationless (dry) merger of two or more systems.  Such ellipicals must have formed at later times than their low luminosity counterparts \\citep{luc06}.  \\citet{dok05} analysed tidal debris of elliptical galaxies and concluded that $\\sim70\\%$ of the bulge dominated galaxies have experienced a merger.  The analysis of nearby bulge dominated galaxies has shown that the gas to stellar mass ratio is very small and these mergers are mostly 'dry'.  Using a semi-analytic model for galaxy formation, \\citet{naa06} found that both the photometric and kinematic properties of massive elliptical galaxies are in agreement with the scenario where massive elliptical galaxies are produced by mergers of lower mass ellipticals.  They suggested that the merger of two spiral galaxies alone cannot reproduce the observed properties, and that the large remnant mass ($> 6\\times 10^{11} M_\\odot$) implies that they must have undergone elliptical - elliptical mergers.  They found that this process is independent of the environment and redshift, which means that dry mergers can occur at low redshift as well.  By analysing spirals from cluster and group environments at intermediate redshift ($z\\sim 0.5$), \\citet{mor07} found that the Tully-Fisher relation shows larger scatter for cluster spirals than for those in the field.  They also found that the central surface mass density of spirals in clusters is small beyond the cluster virial radius, and argued that these observations provide evidence for merger/harrasment.  Mergers are not common in clusters, as the velocity dispersion of virialized clusters is large.  So if ellipticals in clusters have formed by mergers, that most likely happened during the early stage of cluster collapse \\citep{roo82}. There is some observational evidence to support this idea; \\citet{dok99} showed that there is a large fraction of ongoing mergers in a cluster at $z = 0.83$.  The observation of the unvirialized cluster $\\textrm{RX J}0848+4453$ at $z=1.27$ revealed many ongoing dissipationless mergers of galaxies \\citep{dok01}.  These observations go against the view of monolithic collapse where all the ellipticals formed at the same time, at very high redshift.  In this letter we report on the evolution of galaxies in the core region of clusters, measured using bulge-disk decomposition.  We show that in the case of brightest cluster galaxies, the fraction of galaxies with $B/T>0.4$ and $n>2.5$ evolved significantly over the redshift range 0.31 to 0.83.  Throughout the paper we use the standard concordance cosmology with $\\Omega_\\Lambda = 0.73$, $\\Omega_m = 0.27$ and $H_0 = 71$ km s$^{-1}$ Mpc$^{-1}$. ",
        "conclusions": ""
    },
    "0911/0911.3922_arXiv.txt": {
        "abstract": "We present a Nobeyama 45\\,m Radio Telescope map and Australia Telescope Compact Array pointed observations of \\dia\\, 1-0 emission towards the clustered, low mass star forming Oph B Core within the Ophiuchus molecular cloud. We compare these data with previously published results of high resolution \\amm\\, (1,1) and (2,2) observations in Oph B. We use 3D {\\sc clumpfind} to identify emission features in the single-dish \\dia\\, map, and find that the \\dia\\, `clumps' match well similar features previously identified in \\amm\\, (1,1) emission, but are frequently offset to clumps identified at similar resolution in 850\\,\\micron\\, continuum emission. Wide line widths in the Oph B2 sub-Core indicate non-thermal motions dominate the Core kinematics, and remain transonic at densities $n \\sim 3 \\times 10^5$\\,\\cc\\, with large scatter and no trend with $N(\\mbox{H$_2$})$. In contrast, non-thermal motions in Oph B1 and B3 are subsonic with little variation, but also show no trend with H$_2$ column density. Over all Oph B, non-thermal \\dia\\, line widths are substantially narrower than those traced by \\amm, making it unlikely \\amm\\, and \\dia\\, trace the same material, but the $v_{\\mbox{\\tiny{LSR}}}$ of both species agree well. We find evidence for accretion in Oph B1 from the surrounding ambient gas. The \\amm\\,/\\,\\dia\\, abundance ratio is larger towards starless Oph B1 than towards protostellar Oph B2, similar to recent observational results in other star-forming regions. The interferometer observations reveal small-scale structure in \\dia\\, 1-0 emission, which are again offset from continuum emission. No interferometric \\dia\\, emission peaks were found to be coincident with continuum clumps. In particular, the $\\sim 1$\\,M$_\\odot$ B2-MM8 clump is associated with a \\dia\\, emission minimum and surrounded by a broken ring-like \\dia\\, emission structure, suggestive of \\dia\\, depletion. We find a strong general trend of decreasing \\dia\\, abundance with increasing $N(\\mbox{H$_2$})$ in Oph B which matches that found for \\amm. ",
        "introduction": "Stars form out of the gravitational collapse of centrally condensed cores of dense molecular gas. Recent years have seen leaps forward in our understanding of the structure and evolution of isolated, star forming cores. Most star formation, however, occurs in clustered environments \\citep{lada03}. These regions are more complex, with complicated observed geometries, and contain cores which tend to have higher densities and more compact sizes than those found in isolation \\citep{wardthompson07}. It is likely that due to these differences the evolution of filaments and cores in clustered regions proceeds differently than in the isolated cases. Characterizing the physical and chemical structures of these more complicated regions are thus the first steps towards a better understanding of the process of clustered star formation.  It is now clear that molecular cores become extremely chemically differentiated, as many molecules commonly used for tracing molecular gas, such as CO, become severely depleted in the innermost core regions through adsorption onto dust grains \\citep[see, e.g.,][for a review]{difran07}.  In a recent paper \\citep[hereafter \\one]{friesen09}, we studied the dense gas in several cluster forming Cores\\footnotemark\\footnotetext{\\edits{In this paper, we describe Oph A, B, C, etc., as `Cores' since this is how these features were named in DCO$^+$ observations of the L1688 region by \\citet{loren90}. Since then, higher-resolution data such as those described in this paper have revealed substructure in these features that could be themselves precursors to stars, i.e., cores.  To avoid confusion, we refer the larger features in Oph as `Cores' and Core substructure identified by {\\sc clumpfind} as `clumps'.}} in the Ophiuchus molecular cloud (Oph B, C and F) through high resolution observations of \\amm\\, (1,1) and (2,2) emission made at the Green Bank Telescope (GBT), the Australia Telescope Compact Array (ATCA) and the Very Large Array (VLA). The Ophiuchus molecular cloud \\citep[$\\sim 120$\\,pc distant; ][]{loinard08,lombardi08,knude98} is our closest example of ongoing, clustered star formation. We found that the Ophiuchus Cores presented physical characteristics similar to those found in studies of both isolated and clustered environments. In Oph C, for example, gas motions are subsonic (mean $\\sigma_v = 0.16$\\,\\kms), and decrease in magnitude along with the kinetic gas temperature ($T_K$) towards the thermal dust continuum emission peak in a manner reminiscent of findings in the isolated cores L1544 \\citep{crapsi07}; etc. In contrast to Oph C, \\amm\\, line widths in Oph B are dominated by transonic non-thermal motions and the gas temperatures are warmer than typically found in isolated regions ($\\langle T_K \\rangle = 15$\\,K) and nearly constant across the Core. No contrast in any parameters [such as $T_K$, ratio of non-thermal to thermal line width $\\sigma_{\\mbox{\\tiny{NT}}}\\,/\\,c_s$, or \\amm\\, column density $N(\\mbox{\\amm})$] was found between the two main components of Oph B, labelled B1 and B2, despite the presence of several embedded Class I protostars in B2 while B1 appears starless. Additionally, on small scales ($\\sim 15$\\arcsec, or 1800\\,AU at 120\\,pc), significant offsets were found between the peaks of integrated \\amm\\, (1,1) intensity and those of submillimeter continuum emission, as well as between individual `clumps' identified through 3D {\\sc clumpfind} \\citep{williams94} in the \\amm\\, data cube and those found through 2D {\\sc clumpfind} in 850\\,\\micron\\, continuum emission. Finally, evidence was found for a decreasing fractional \\amm\\, abundance with increasing H$_2$ column density (as traced by 850\\,\\micron\\, continuum emission), suggesting that \\amm\\, (1,1) may not be tracing the densest gas in the Oph Cores.  The diazenylium ion, \\dia, has been shown to be a preferential tracer of quiescent dense gas in molecular clouds \\citep{womack92,caselli02,tafalla02,tafalla04,difran04}. The 1-0 rotational transition has a critical density $n_{cr} \\sim 2 \\times 10^5$\\,\\cc, $\\sim 10-100 \\times$ greater than that of \\amm\\, (1,1), and is therefore possibly better suited to probing gas properties at the high densities expected in clustered star forming Cores. Based on both observations and chemical models of starless cores, \\dia\\, also appears resilient to depletion at the high densities ($n \\gtrsim 10^5$\\,\\cc) and cold temperatures ($T \\sim 10$\\,K) characteristic of the later stages of prestellar core evolution \\citep[see, e.g.,][]{tafalla04,bergin02}, and is thus expected to be an excellent tracer of the physical conditions of dense cores and clumps.  Recently, \\citet{andre07} observed the Oph A, B, C, E and F Cores in \\dia\\, 1-0 emission at 26\\arcsec\\, angular resolution using the IRAM 30\\,m telescope. They found that the gas traced by \\dia\\, emission near clumps identified in millimeter continuum emission \\citep{motte98} is characterized, on average, with relatively narrow line widths showing small non-thermal velocity dispersions, in contrast to the large non-thermal motions traced by \\amm\\, emission in \\one. The authors derived virial masses from the \\dia\\, emission, and found they generally agreed within a factor $\\sim 2$ with the masses derived from dust emission, suggesting that the clumps are gravitationally bound and prestellar. The relative motions of the clumps are also small and subvirial, with a crossing time larger than the expected lifetime of the objects, such that clump-clump interactions are not expected to impact their evolution. These results suggest that a dynamic picture of clump evolution involving competitive accretion at the prestellar stage does not accurately describe the star formation process in central Ophiuchus.  In this work, we discuss the results of higher resolution (18\\arcsec, or $\\sim 2200$\\,AU at 120\\,pc) observations of \\dia\\, 1-0 in the Ophiuchus B Core that reveal the distribution, kinematics and abundance pattern of the Core and associated embedded clumps on smaller physical scales. We additionally examine \\dia\\, 1-0 structure at even higher resolution ($8\\arcsec \\times 5\\arcsec$), through observations made with the ATCA, towards five locations in Oph B2 where small-scale structure was found in \\amm\\, (\\one) and in previously unpublished Berkeley-Illinois-Maryland Association (BIMA) \\dia\\, 1-0 observations. We compare these data with our recently published analysis of \\amm\\, emission in Oph B, in particular focusing on the kinematics and relative abundances of the two species to compare with physical and chemical models of dense core formation and collapse.  We discuss the observations and calibration and \\S2. In \\S3, we present the data, and discuss the results of the hyperfine line structure fitting procedure and derivations of the column density, $N(\\mbox{\\dia})$, and fractional \\dia\\, abundance, $X(\\mbox{\\dia})$, in \\S4. In \\S5, we discuss general trends in the data, and compare these results with those found in \\one\\, and with studies of dense cores in isolated environments. We summarize our findings in \\S6. ",
        "conclusions": "\\subsection{General trends} \\label{sec:trends}  We show in Figure \\ref{fig:n2h_trends} the distribution of the ratio of the non-thermal dispersion to the sound speed, $\\sigma_{\\mbox{\\tiny{NT}}}\\,/\\,c_s$, $n_{ex}$, $N(\\mbox{\\dia})$ and $X(\\mbox{\\dia})$ with $N(\\mbox{H$_2$})$ in Oph B, omitting pixels where the 850\\,\\micron\\, continuum flux $S_\\nu  \\leq 0.1$\\,Jy\\,beam$^{-1}$. Data points represent values for 18\\arcsec\\, pixels, i.e., approximately the Nobeyama beam FWHM. We show individually pixels in Oph B1 and B2, and also show peak values for \\dia\\, clumps, identified in \\S\\ref{sec:clumps}, and 850\\,\\micron\\, continuum clumps \\citep{jorgensen08}. Note that no Oph B3 values are plotted due to the lack of continuum emission at the B3 location. The Figure shows that the \\dia\\, clumps do not reside at the highest H$_2$ column densities, but scatter over the range of $N(\\mbox{H$_2$})$ calculated above our submillimeter flux threshold. The submillimeter clumps tend to be found at higher $N(\\mbox{H$_2$})$ values than the \\dia\\, clumps, on average, but also show a spread in peak $N(\\mbox{H$_2$})$.  In \\one, we used the ratio of \\amm\\, (1,1) and (2,2) emission lines in Oph B to determine the kinetic temperature $T_K$ in each pixel, which we can use to calculate $\\sigma_{\\mbox{\\tiny{NT}}}$ and $c_s$ for \\dia\\, emission in Oph B. Given $T_K$, $\\sigma_{\\mbox{\\tiny{NT}}} = \\sqrt{\\sigma_{obs}^2 - k_B T_K\\,/\\,(\\mu_{mol} m_{\\mbox{\\tiny{H}}})}$, where $k_B$ is the Boltzmann constant, $m_{\\mbox{\\tiny{H}}}$ is the mass of the hydrogen atom, $\\mu_{mol}$ is the molecular weight in atomic units ($\\mu_{\\mbox{\\tiny{\\dia}}} = 29.02$) and $\\sigma_{obs} = \\Delta v\\,/\\,(2\\sqrt{2\\ln2})$. We list values and propagated uncertainties for individual \\dia\\, clumps in Table \\ref{tab:peak_dat}, and give the mean, rms, minimum and maximum $\\sigma_{\\mbox{\\tiny{NT}}}\\,/\\,c_s$ in Table \\ref{tab:col} for each of Oph B1, B2 and B3.  We find a mean $\\sigma_{\\mbox{\\tiny{NT}}}\\,/\\,c_s = 1.02$ for \\dia\\, 1-0 emission across Oph B with an rms variation of 0.49, indicating that the non-thermal motions are approximately equal, on average, to the sound speed. When looked at separately, we find a difference in the  $\\sigma_{\\mbox{\\tiny{NT}}}\\,/\\,c_s$ ratio between Oph B1 and B2. In Oph B1, the mean $\\sigma_{\\mbox{\\tiny{NT}}}\\,/\\,c_s = 0.71$, with an rms variation of 0.16. In Oph B2, the mean $\\sigma_{\\mbox{\\tiny{NT}}}\\,/\\,c_s = 1.26$ with an rms variation of 0.4. Oph B3 is dominated by nearly thermal line widths, with the mean $\\sigma_{\\mbox{\\tiny{NT}}}\\,/\\,c_s = 0.39$ and an rms variation of only 0.13.  We find little variation in $\\sigma_{\\mbox{\\tiny{NT}}}\\,/\\,c_s$ with $N(\\mbox{H$_2$})$ in Oph B, as shown in Figure \\ref{fig:n2h_trends}a. The difference in the mean $\\sigma_{\\mbox{\\tiny{NT}}}\\,/\\,c_s$ ratio between B1 and B2 is clear, as is the significantly larger scatter of non-thermal line widths in B2 than in B1. The scatter in $\\sigma_{\\mbox{\\tiny{NT}}}\\,/\\,c_s$ extends to the highest H$_2$ column densities. In B1, all non-thermal line widths $\\sigma_{\\mbox{\\tiny{NT}}} \\leqq c_s$, shown by the dashed line. Similar results are found for most \\dia\\, clumps, and approximately half the continuum clumps. At 18\\arcsec\\, resolution, the observed non-thermal line widths do not extend to arbitrarily low values, as the data show a cutoff in the $\\sigma_{\\mbox{\\tiny{NT}}}\\,/\\,c_s$ ratio at $\\sigma_{\\mbox{\\tiny{NT}}}\\,/\\,c_s \\sim 0.5$ for both B1 and B2, below which no data points are found (but note that in \\S\\ref{sec:analysis-atca}, we found narrow line widths at high resolution). For $T_K = 15$\\,K, the mean temperature found in Oph B, the $\\sigma_{\\mbox{\\tiny{NT}}}\\,/\\,c_s$ cutoff corresponds to $\\sigma_{\\mbox{\\tiny{NT}}} \\sim 0.12$\\,\\kms. Note that the velocity resolution of the \\dia\\, data, $\\Delta v_{res} = 0.05$\\,\\kms, is significantly less than the observed lower $\\sigma_{\\mbox{\\tiny{NT}}}\\,/\\,c_s$ limit.  In \\one, we discussed the possible origins of the large non-thermal motions in Oph B, and based on timescale arguments ruled out primordial motions (i.e., motions inherited from the parent cloud). It is interesting that most of the gas traced by \\dia\\, 1-0 in Oph B2 remains dominated by transonic non-thermal motions, while the line widths in Oph B1 reveal the Core is significantly more quiescent at the $\\sim 10^5$\\,\\cc\\, densities traced by \\dia\\, 1-0 emission. The most obvious source of this difference is the presence of embedded protostars in Oph B2 while Oph B1 is starless. \\citet{kamazaki03} found an outflow in CO 3-2 emission in Oph B2 centred near Elias 33 and Elias 32 and oriented approximately east-west, but were unable to pinpoint which protostar was the driving source. In upcoming results from the JCMT Gould Belt Legacy survey \\citep{gbls}, which mapped Oph B in CO 3-2, $^{13}$CO 3-2 and C$^{18}$O 3-2 at 14\\arcsec\\, resolution, this highly clumped outflow is shown to extend over more than 10\\arcmin\\, (i.e., 0.4\\,pc at 120\\,pc, White {\\it et al.} 2009, in preparation). The outflow axis is also aligned such that the outflow does not appear to be impacting Oph B1 significantly, potentially explaining the difference in non-thermal line widths between the two sub-Cores\\footnotemark\\footnotetext{JCMT Spring 2009 newsletter, http://www.jach.hawaii.edu/JCMT/publications/newsletter/n30/jcmt-n30.pdf}.   \\dia\\, is not generally thought to be a tracer of protostellar outflows, since it is expected to be destroyed quickly through reactions with CO, which evaporates from dust grains in the higher temperature gas near the driving protostar. On small scales in B2, however, some variations in $v_{\\mbox{\\tiny{LSR}}}$ and $\\Delta v$ appear correlated with the presence of nearby Class I protostars (see \\S\\ref{sec:vlsr} and Figure \\ref{fig:b-mm8}). \\citet{chen08} suggest \\dia\\, can be entrained in protostellar jets before a molecular outflow releases CO from grain surfaces which then can destroy \\dia. In a small \\dia\\, 1-0 survey of the Serpens NW cluster, \\citet{williams00} found that the largest $\\sigma_{\\mbox{\\tiny{NT}}}$ values ($\\sigma_{\\mbox{\\tiny{NT}}} > 0.6$\\,\\kms, greater than seen here in Oph B2) occurred in cores containing protostellar sources which power strong outflows. If caused by the substantial protostellar outflow, the large line widths on a global scale in Oph B2 suggest that protostellar outflows are able to inject additional turbulence into the {\\it high density} gas in cluster-forming Cores, thereby increasing the mass at which clumps become gravitationally unstable and altering the consequent fragmentation and evolution of existing embedded clumps. \\citet{difran01} did not see this effect in the \\dia\\, 1-0 observations of the circumstellar dense gas associated with the protostellar, outflow-driving source NGC 1333 IRAS4B, suggesting that the ability of the outflow to inject turbulent energy may be determined by additional factors, such as the collimation or strength of the outflow. A comparison of the outflow properties in these regions would be useful to probe further the impact of outflows on cluster forming dense gas.  Alternatively, the larger line widths in B2 may be due to global infall motions. In \\one, we determined a virial mass $M_{vir} \\sim 8$\\,M$_\\odot$ for Oph B2 based on \\amm\\, line widths, which is a factor $\\sim 5$ less than Core mass estimates based on thermal dust continuum emission (with a factor $\\sim 2$ uncertainty). In contrast, $M/M_{vir} \\sim 1$ in Oph B1. Some evidence for infall has been observed in self-absorbed molecular line tracers in Oph B2 \\citep{andre07,gurney08}, but it is difficult to determine unambiguously given the complex outflow motions also present. Additionally, the consistent line widths found in \\amm\\, gas for both B1 and B2 (discussed further in \\S\\ref{sec:sig} in comparison with the \\dia\\, results) suggest a common source for non-thermal motions at gas densities $n \\sim 10^{3-4}$\\,\\cc, which then impacts differently the gas at higher densities in the two sub-Cores.    Figure \\ref{fig:n2h_trends}b shows the volume density $n_{\\mbox{\\tiny{ex}}}$ as a function of $N(\\mbox{H$_2$})$. Nearly all pixels have $n_{\\mbox{\\tiny{ex}}} \\gtrsim 10^5$\\,\\cc, and we find slightly greater $n_{\\mbox{\\tiny{ex}}}$ at higher column densities. The \\dia\\, and continuum clumps scatter over the full range of density values found in Oph B.  We find a slight trend of increasing \\dia\\, column density, $N(\\mbox{\\dia})$, with $N(\\mbox{H$_2$})$, shown in Figure \\ref{fig:n2h_trends}c, although the mean $N(\\mbox{\\dia})$ changes by less than a factor of $\\sim 2$ over the order of magnitude range in $N(\\mbox{H$_2$})$ found in Oph B. Several pixels in Oph B1, and one \\dia\\, clump, have significantly greater $N(\\mbox{\\dia})$ for their $N(\\mbox{H$_2$})$ values compared with the rest of the data. The small gradient in $N(\\mbox{\\dia})$ with $N(\\mbox{H$_2$})$ leads to a trend of {\\it decreasing} fractional \\dia\\, abundance with increasing H$_2$ column density, shown in Figure \\ref{fig:n2h_trends}d. We find \\dia\\, abundances decrease by an order of magnitude between the limiting $N(\\mbox{H$_2$})$ threshold and the peak $N(\\mbox{H$_2$})$ values. We fit a linear relationship to $\\log X(\\mbox{\\dia})$ versus $\\log N(\\mbox{H$_2$})$ for Oph B1 and B2 separately. In B1 alone, where we have relatively few data points, we find $\\log X(\\mbox{\\dia})$ is consistent within the uncertainties with being constant with $N(\\mbox{H$_2$})$, but in conjunction with the Oph B2 data, we find  \\begin{equation} \\label{eqn:xvsh} \\log X(\\mbox{\\dia}) = (7.1 \\pm 0.9) - (0.74 \\pm 0.04) \\times \\log N(\\mbox{H$_2$}) \\end{equation}  \\noindent using the same H$_2$ column density limits as in Figure \\ref{fig:n2h_trends}.  \\subsection{Small-Scale Features}  We next discuss the physical properties of the small-scale structure present in Oph B, including the \\dia\\, clumps, continuum clumps, and embedded protostars, and describe in detail the compact, thermal object B2-N6.  \\subsubsection{Comparison of \\dia\\, clumps, continuum clumps and protostars} \\label{sec:compare_clumps}  We list in Table \\ref{tab:prot_dat} the physical parameters derived from \\dia\\, 1-0 emission towards 850\\,\\micron\\, continuum clump locations \\citep{jorgensen08} and embedded Class I protostars \\citep{enoch09}. Columns are the same as in Table \\ref{tab:peak_dat}. In Table \\ref{tab:mean_dat}, we list the mean values of each physical parameter towards \\dia\\, clumps, continuum clumps and protostars. Note that we were only able to solve for all parameters listed towards a single protostar, and discuss below only those mean parameters where we obtained values from three or more objects.  We find no difference in the mean $v_{LSR}$ between \\dia\\, clumps, continuum clumps and protostars. There is a clear reduction in $\\Delta v$ towards \\dia\\, clumps relative to continuum clumps (mean $\\Delta v = 0.49$\\,\\kms and 0.68\\,\\kms, respectively), but excluding one continuum clump (162712-24290) reduces the mean $\\Delta v$ to 0.58\\,\\kms. Protostars are associated with wider mean $\\Delta v$, but the larger mean is also driven by large $\\Delta v$ towards a single object. The differences between \\dia\\, and continuum clumps are marginally significant given the variance of $\\Delta v$ across the entire Core is only 0.08\\,\\kms\\, in Oph B1 and 0.21\\,\\kms\\, in Oph B2. We find similar \\dia\\, line opacities and excitation temperatures towards continuum and \\dia\\, clumps. Continuum clumps are associated with larger non-thermal motions, with a mean $\\sigma_{NT} / c_s = 1.24$ which is 1.5 times that associated with \\dia\\, clumps (mean $\\sigma_{NT}/c_s = 0.86$). \\dia\\, column densities are similar towards \\dia\\, and continuum clumps (within $\\sim 25$\\,\\%), but \\dia\\, clumps are associated with smaller $N(\\mbox{H$_2$})$ and hence greater $X(\\mbox{\\dia})$ by a factor $> 2$. In fact, the mean \\dia\\, abundance towards continuum clumps is less than $X(\\mbox{\\dia})$ averaged over the entire Oph B Core. Derived volume densities are the same for \\dia\\, and continuum clumps.  In summary, the \\dia\\, clumps are associated with smaller line widths and subsonic non-thermal motions relative to continuum clumps and protostars, but the difference in mean values is similar to the variance in line widths across the Core. The most significant difference between \\dia\\, and continuum clumps are found in their \\dia\\, abundances, with a mean $X(\\mbox{\\dia})$ towards continuum clumps that is lower than $X(\\mbox{\\dia})$ in \\dia\\, clumps.  \\subsubsection{Oph B2-N6} \\label{sec:n6}  The B2-N6 clump is a pocket of narrow \\dia\\, 1-0 line emission within the larger, more turbulent B2 Core, with a large fraction ($\\gtrsim 50$\\%) of the emission at the clump peak coming from the small size scales probed by the ATCA observations (see Figure \\ref{fig:nh3peak}). Although we do not consider all \\dia\\, clumps to trace underlying physical structure in Oph B, the remarkable properties of B2-N6 suggest that further attention be paid to that particular clump. Line widths measured from the interferometer data are a factor $\\sim 2$ less than those recorded from the single-dish data ($\\Delta v = 0.21 \\pm 0.04$\\,\\kms\\, compared with $0.39 \\pm 0.01$\\,\\kms). The millimeter continuum object B2-MM15, identified through a spatial filtering technique \\citep{motte98}, lies offset by $\\sim 12$\\arcsec\\, from the peak interferometer integrated intensity (see Figure \\ref{fig:nh3peak}).  \\citeauthor{motte98} calculate a mass of only $0.17$\\,M$_\\odot$ associated with the B2-MM15 clump, but based on its small size (unresolved at 15\\arcsec\\, resolution, using an estimate of $r \\sim 800$\\,AU) derive an average density of $n \\sim 1.2 \\times 10^8$\\,\\cc. This density is significantly larger than the densities calculated in \\S\\ref{sec:nex} at 18\\arcsec\\, resolution, but higher sensitivity \\dia\\, 1-0 interferometer observations are needed to probe the clump density on small scales.  We calculated in \\one\\, a gas kinetic temperature $T_K = 13.9\\pm1.1$\\,K for the co-located B2-A7 \\amm\\, (1,1) clump. We find $\\sigma_{\\mbox{\\tiny{NT}}} = 0.06$\\,\\kms\\, for B2-N6, equal to the expected \\dia\\, thermal dispersion, and $\\sigma_{\\mbox{\\tiny{NT}}}/c_s = 0.26$. A similar compact \\dia\\, 1-0 clump with nearly thermal line widths was found in Oph A \\citep[Oph A-N6, ][]{difran04}. While both clumps are associated with larger-scale \\dia\\, emission, these objects are only distinguishable as compact, nearly purely thermal clumps when observed with interferometric resolution and spatial filtering.  \\citeauthor{difran04} proposed Oph A-N6 may be a candidate thermally dominated, critical Bonnor-Ebert sphere embedded within the more turbulent Oph A Core. Such objects have sizes comparable to the cutoff wavelength for MHD waves and are in equilibrium between their self-gravity and internal and external pressures.  \\citet{myers98} called these objects `kernels', and showed that they could exist in cluster-forming cores with FWHM line widths $\\Delta v > 0.9$\\,\\kms\\, and column densities $N(\\mbox{H$_2$}) > 10^{22}$\\,\\cc, with sizes comparable to the spacing of protostars in embedded clusters, suggesting that a population of kernels within a dense Core could form a stellar cluster. The physical conditions in both Oph A and Oph B2 fit these requirements. The B2-N6 clump, however, is significantly smaller (0.004\\,pc compared with a predicted $\\sim 0.03$\\,pc) and less massive (0.17\\,M$_\\odot$ compared with a predicted $\\sim 1$\\,M$_\\odot$) than the objects discussed by \\citeauthor{myers98}.  If we calculate the mass of a critical Bonnor-Ebert sphere given the ATCA \\dia\\, 1-0 line width and a radius of $\\sim 800$\\,AU, where $M/R = 2.4 \\sigma^2/G$ \\citep{bonnor56}, we find $M_{\\mbox{\\tiny{BE}}} = 0.02$\\,M$_\\odot$, or approximately $10\\times$ smaller than the mass determined by \\citeauthor{motte98} We find a virial mass $M_{vir} = 0.04$\\,M$_\\odot$, assuming B2-N6 is a uniform sphere and $M_{vir} = 5 \\sigma^2 R/G$. If the clump mass is accurate, this suggests that while small, B2-N6 may be gravitationally unstable. \\citet{andre07} made a tentative detection of infall motions towards the clump through observations of optically thick lines such as CS, H$_2$CO or HCO$^+$. Thus B2-N6 may be at a very early stage of clump formation, and may eventually form an additional low mass protostar to the three YSOs already associated with dense gas in Oph B2.   \\subsection{Comparison of \\dia\\, and \\amm\\, emission in Oph B}  We next compare the physical properties of the gas in Oph B derived here from \\dia\\, observations with those derived from \\amm\\, observations, described in \\one. In the following discussion, all comparisons are made after convolving the \\amm\\, data to a final angular resolution of 18\\arcsec\\, to match the Nobeyama \\dia\\, observations.  \\subsubsection{$v_{\\mbox{\\tiny{LSR}}}$ and $\\sigma_{\\mbox{\\tiny{NT}}}\\,/\\,c_s$} \\label{sec:sig}  We first look at the velocities and non-thermal line widths of the gas traced by \\dia\\, gas and those determined from \\amm\\, emission in \\one.  Figure \\ref{fig:compare_nh3}a shows the distribution of $v_{\\mbox{\\tiny{LSR}}}$ derived from \\dia\\, and \\amm\\, emission in Oph B. There is no significant difference in the measured line-of-sight velocity between the two dense gas tracers, with the mean $v_{\\mbox{\\tiny{LSR}}} = 3.98$\\,\\kms\\, for \\amm\\, and 4.01\\,\\kms\\, for \\dia.  In Figure \\ref{fig:compare_nh3}b, we show the distribution of $\\sigma_{\\mbox{\\tiny{NT}}}\\,/\\,c_s$ as determined with \\dia\\, or \\amm\\, using the \\amm-derived $T_K$ values. The non-thermal line widths measured in \\dia\\, 1-0 are substantially smaller than those measured in \\amm\\, in \\one, and additionally the variation in the relative magnitude of the non-thermal motions between Oph B1 and B2 was not found in \\amm\\, emission. The mean \\amm\\, $\\sigma_{\\mbox{\\tiny{NT}}}\\,/\\,c_s = 1.64$ in Oph B, and does not vary significantly between B1 ($\\langle \\sigma_{\\mbox{\\tiny{NT}}}\\,/\\,c_s\\rangle = 1.62$) and B2 ($\\langle \\sigma_{\\mbox{\\tiny{NT}}}\\,/\\,c_s\\rangle = 1.68$).   The assumption that $T_K$ is similar for both \\dia\\, and \\amm\\, is reasonable if both molecules are excited in the same material. The good agreements between the locations of \\amm\\, and \\dia\\, clumps, illustrated in Figure \\ref{fig:b_sep}, and the $v_{\\mbox{\\tiny{LSR}}}$, illustrated in Figure \\ref{fig:compare_nh3}a, support this assumption. The substantial offset in the $\\sigma_{\\mbox{\\tiny{NT}}}\\,/\\,c_s$ ratio seen in Figure \\ref{fig:compare_nh3}b, however, between \\amm\\, and \\dia\\, emission indicates significantly different motions are present in the gas traced by \\amm\\, than in the gas traced by \\dia. The gas densities traced by \\dia\\, are $\\sim 1-2$ orders of magnitude greater than those traced by \\amm\\, ($\\sim 10^5$\\,\\cc\\, and $\\sim 10^{3-4}$\\,\\cc, respectively). In fact, if we attempt to derive the gas kinetic temperature by assuming equal non-thermal motions for the \\dia\\, 1-0 and \\amm\\, (1,1) emission (effectively assuming \\dia\\, and \\amm\\, trace the same material), the resulting average value is a highly unlikely $T_K = 190$\\,K for Oph B2. Given the narrow \\dia\\, lines observed in Oph B1, no physical solution for $T_K$ can be found under the assumption of equal $\\sigma_{\\mbox{\\tiny{NT}}}$.  Starless cores are typically found to be well described by gas temperatures that are either constant or decreasing as a function of increasing density \\citep{difran07}. Given that \\dia\\, traces denser gas, we then expect $T_K$ from \\amm\\, measurements to be biased high in starless cores, e.g., gas traced by \\dia\\, should be colder than gas traced by \\amm. The mean $T_K$ found in Oph B in \\one\\, was 15\\,K. With a lower $T_K$, the sound speed would be smaller and the returned $\\sigma_{\\mbox{\\tiny{NT}}}$ would be larger on average. A gas temperature $T_K = 10$\\,K rather than 15\\,K  would increase the average \\dia\\, $\\sigma_{\\mbox{\\tiny{NT}}}\\,/\\,c_s$ ratio in Oph B by $\\sim 20 - 25$\\,\\%. This increase, while significant, is not large enough for the mean \\dia\\, $\\sigma_{\\mbox{\\tiny{NT}}}\\,/\\,c_s$ to match the \\amm\\, results in Oph B.  With both a constant temperature or decreasing temperature with density, non-thermal motions in Oph B1 and B3 would remain subsonic, while Oph B2 would still be dominated by transonic non-thermal motions. Alternatively, it is possible that the two Class I protostars in Oph B2 could raise the temperature of the dense gas above that traced by \\amm\\, (no $T_K$ difference was found between protostars and starless areas in \\one). In this case, the returned \\dia\\, $\\sigma_{\\mbox{\\tiny{NT}}}$ would be smaller, and would further increase the differences in magnitude of non-thermal motions seen between the \\amm\\, and \\dia\\, emission in Oph B1. In B2, a temperature of 20\\,K would decrease the mean $\\sigma_{\\mbox{\\tiny{NT}}}\\,/\\,c_s$ ratio to 1.1. In a study of dust temperatures in the Ophiuchus Cores, \\citet{stamatellos07} showed, however, that embedded protostars in the Cores will only heat very nearby gas and will not raise the mean temperature of the Cores by more than $\\sim 1-2$\\,K, so an average $T_K = 20$\\,K over all Oph B2 is unlikely.  \\subsubsection{Double peaked line profiles} \\label{sec:dblpeak}  In Figure \\ref{fig:compare_peaks}, we show the isolated $F_1F \\rightarrow F'_1F' = 0 1 \\rightarrow 1 2$ component of the \\dia\\, 1-0 emission line towards the three \\dia\\, clumps in Oph B1 (B1-N1, B1-N3, and B1-N4) which show double peaked line profiles, as described in \\S\\ref{sec:analysis}. Figure \\ref{fig:compare_peaks} also shows a single Gaussian line profile overlaid on each \\dia\\, spectrum, which represents the $v_{\\mbox{\\tiny{LSR}}}$ and $\\Delta v$ of the 18 component Gaussian fit to the \\amm\\, (1,1) emission at the clump peak, normalized to a peak amplitude of 1\\,K. We plot the Gaussian fit for clarity since there are no isolated components in the \\amm\\, (1,1) hyperfine structure. Additionally, spectra of C$_2$S $2_1 - 1_0$ emission line at the peak clump locations are shown. Note that over all Oph B, C$_2$S emission was only detected towards the southern tip of Oph B1, and thus no significant C$_2$S emission was found at the peak location of B1-N1. The C$_2$S emission was observed with the GBT at $\\sim 32$\\arcsec\\, spatial resolution and 0.08\\,\\kms\\, spectral resolution, with the observations and analysis described in \\one. Arrows show the locations of the fitted $v_{\\mbox{\\tiny{LSR}}}$ for all species, including both \\dia\\, velocity components.  The clumps which show the double-peaked line profile structure are the clumps with the highest total \\dia\\, 1-0 line opacities based on a single velocity component fit, with $\\tau = 5, 6$ and 9, respectively for B1-N1, B1-N3 and B1-N4. The $F_1F \\rightarrow F'_1F' = 0 1 \\rightarrow 1 2$ component is expected to have an opacity $\\tau_{iso} = 1/9 \\,\\tau$, so $\\tau_{iso} = 0.6, 0.7$, and $\\sim 1$ for the three clumps, based on a single component fit. If the line is optically thick, however, this fit is likely to underestimate the true opacity given the missing, self-absorbed flux. This is suggestive that the lines are indeed self-absorbed. Since the \\amm\\, and C$_2$S emission is optically thin (\\one), however, if the \\dia\\, emission was self-absorbed due to high optical depth we would expect the emission peak of both \\amm\\, and C$_2$S to be found between the two \\dia\\, components in $v_{\\mbox{\\tiny{LSR}}}$, which is not the case. In all three clumps, we find the $v_{\\mbox{\\tiny{LSR}}}$ of the \\amm\\, emission more closely matches the $v_{\\mbox{\\tiny{LSR}}}$ of the red \\dia\\, component, while the C$_2$S emission in two clumps ($v_{\\mbox{\\tiny{LSR}}} =  3.67$\\,\\kms) more closely matches the $v_{\\mbox{\\tiny{LSR}}}$ of the blue \\dia\\, component ($v_{\\mbox{\\tiny{LSR}}} = 3.85$\\,\\kms\\, and 3.62\\,\\kms\\, for B1-N3 and B1-N4, respectively). This behaviour suggests the \\dia\\, 1-0 emission is not self-absorbed towards these positions. A similar offset of C$_2$S emission relative to \\dia\\, was found by \\citet{swift06} towards the starless Core L1551, however the case of L1551 the C$_2$S emission is redshifted with respect to the systemic velocity of the Core rather than blueshifted, as we find in B1.    The GBT \\amm\\, (1,1) observations described in \\one\\, found extensive emission around Oph B, such that it is possible to determine the $v_{\\mbox{\\tiny{LSR}}}$ and $\\Delta v$ of \\amm\\, beyond where continuum emission traced the Oph B Core. We use an intensity threshold of 2\\,K in the \\amm\\, (1,1) main component to separate Core and off-Core gas. The off-Core gas surrounding Oph B1 has a mean $v_{\\mbox{\\tiny{LSR}}} = 3.72$\\,\\kms\\, with an rms variation of 0.13\\,\\kms, which is significantly different from the mean B1 Core $v_{\\mbox{\\tiny{LSR}}} = 3.96$\\,\\kms\\, and rms variation of 0.15\\,\\kms. Both the C$_2$S emission and the blue \\dia\\, component are thus more kinematically similar to the off-Core \\amm\\, gas. The critical densities of C$_2$S $2_1 - 1_0$ and \\amm\\, (1,1) are similar \\citep{suzuki92,rosolowsky08}. Kinematically, we find $\\sigma_{\\mbox{\\tiny{NT}}}\\,/\\,c_s = 0.7$ for the C$_2$S emission if we assume the $T_K$ determined from \\amm\\, emission accurately represents the temperature of the gas traced by C$_2$S. This result closely matches the \\dia\\, results in Oph B1, but the non-thermal motions in the gas traced by \\amm\\, are significantly larger (\\S\\ref{sec:trends} and \\one). It is impossible to match the non-thermal C$_2$S and \\amm\\, motions for a given $T_K$, since the \\amm\\, (1,1) line $\\sigma_{\\mbox{\\tiny{NT}}} > \\sigma_{obs}$ of the C$_2$S.  At densities $n \\gtrsim$ a few $\\times 10^3 - 10^4$\\,\\cc, chemical models predict C$_2$S is quickly depleted from the gas phase \\citep[timescale $t_{dep} \\sim 10^5$\\,yr;][]{millar90,bergin97}, and accordingly the molecule is observed to be a sensitive tracer of depletion in isolated cores \\citep{lai00,tafalla06}. The volume density in Oph B1 is greater than that required for C$_2$S depletion ($n \\gtrsim 10^5 - 10^6$\\,\\cc\\, from \\S\\ref{sec:nex} and continuum observations), so we would not expect to see any C$_2$S emission unless the gas has only been at high density for $t < t_{dep}$. C$_2$S was only detected towards southern B1. An explanation, therefore, for both the presence of C$_2$S in southern B1 and for its velocity offset relative to the dense gas is that gas from the ambient molecular cloud is accreting onto B1, reaching a density high enough to excite the C$_2$S emission line. This material would be chemically `younger' than the rest of the B1 Core gas, such that the C$_2$S has not had enough time to deplete from the gas phase. A similar result was found in a multi-species study of the starless L1498 and L1517B cores, where \\citet{tafalla04} conclude that non-spherical contraction of the cores produced asymmetric distributions of CS and CO, where the CS and CO `hot spots' reveal the distribution of recently accreted, less chemically processed dense gas.  A single 850\\,\\micron\\, continuum clump \\citep[162715-24303;][]{jorgensen08} was identified in southern Oph B1 where we find B1-N3 and B1-N4. The authors found a clump mass $M = 0.2\\,M_\\odot$, which is a factor $\\sim 2$ less than the virial mass $M_{vir} = 0.4\\,M_\\odot$ we calculate based on the mean line width (of the single velocity component fit) of B1-N3 and B1-N4 ($\\Delta v = 0.52$\\,\\kms, or $\\sigma = 0.22$\\,\\kms) and the continuum clump radius. This suggests that the clump is currently stable against collapse. If the clump is gaining mass through ongoing accretion from the ambient gas, however, then eventually the non-thermal support may not be large enough to prevent gravitational collapse, leading to the formation of a low-mass protostar.   \\subsubsection{Opacity and Excitation Temperature}  We show in Figure \\ref{fig:compare_nh3}c and \\ref{fig:compare_nh3}d the distribution of total line opacity $\\tau$ and excitation temperature $T_{ex}$ derived from hyperfine line fitting of \\amm\\, (1,1) and \\dia\\, 1-0 in Oph B. We find higher total opacities in \\dia\\, emission ($\\langle \\tau \\rangle = 2.5$) than in \\amm\\, emission ($\\langle \\tau \\rangle = 1.0$) in Oph B. The \\dia\\, emission is characterized by lower excitation temperatures than those found for \\amm\\, emission in \\one, with a mean $T_{ex} = 7.1$\\,K for \\dia\\, 1-0 compared with $T_{ex} = 9.5$\\,K for \\amm\\, (1,1).  \\subsubsection{Fractional Abundances and Chemical Evolution} \\label{sec:frac}  We next compare the column densities of \\amm\\, and \\dia\\, towards the Oph B Core. We only show results for pixels which have well-determined values for both $N(\\mbox{\\amm})$ and $N(\\mbox{\\dia})$. In \\one, we found a trend of decreasing fractional \\amm\\, abundance with increasing $N(\\mbox{H$_2$})$ for {\\it both} B1 and B2.  After convolving the \\amm\\, data to match the single-dish \\dia\\, resolution, we find the following linear relationships between $\\log X(\\mbox{\\amm})$ and $\\log N(\\mbox{H$_2$})$:  \\begin{eqnarray} \\mbox{B1}&:& \\log X(\\mbox{\\amm}) = (5 \\pm 4) - (0.6 \\pm 0.2) \\log N(\\mbox{H$_2$}) \\\\ \\mbox{B2}&:& \\log X(\\mbox{\\amm}) = (3.5 \\pm 2.2) - (0.5 \\pm 0.1) \\log N(\\mbox{H$_2$}) \\end{eqnarray}  \\noindent The slopes found for Oph B1 and B2 are both negative and similar to that found in Equation \\ref{eqn:xvsh} for the \\dia\\, fractional abundance as a function of $N(\\mbox{H$_2$})$. This result is suggestive of a trend of increasing depletion of both \\amm\\, and \\dia\\, with increasing $N(\\mbox{H$_2$})$ in Oph B.  We show in Figure \\ref{fig:compare_column} the \\amm\\, column density, $N(\\mbox{\\amm})$, versus the \\dia\\, column density, $N(\\mbox{\\dia})$ in Oph B1 and B2, and also the ratio of $N(\\mbox{\\amm})$ to $N(\\mbox{\\dia})$ as a function of $N(\\mbox{H$_2$})$, calculated in \\S\\ref{sec:nh2} (we do not show results in B3 due to the small number of pixels with both well-determined \\amm\\, and \\dia\\, column densities, but note mean values below). Note that the column density ratio, $N(\\mbox{\\amm})\\,/\\,N(\\mbox{\\dia})$, is equivalent to the ratio of fractional abundances, $X(\\mbox{\\amm})\\,/\\,X(\\mbox{\\dia})$. Figure \\ref{fig:compare_column}a shows a significant difference in the $N(\\mbox{\\amm})\\,/\\,N(\\mbox{\\dia})$ ratio between Oph B1 and B2. We find the mean and the standard deviation of this ratio are each twice as large in B1 ($N(\\mbox{\\amm})\\,/\\,N(\\mbox{\\dia}) = 135$, $\\sigma = 52$) as in B2 ($N(\\mbox{\\amm})\\,/\\,N(\\mbox{\\dia}) = 65$, $\\sigma = 31$) and B3 ($N(\\mbox{\\amm})\\,/\\,N(\\mbox{\\dia}) = 80$, $\\sigma = 56$).  In Oph B1 and B2, Figure \\ref{fig:compare_column}a illustrates that this difference is largely driven by lower $N(\\mbox{\\dia})$ values towards B1, while the spread in the $N(\\mbox{\\amm})\\,/\\,N(\\mbox{\\dia})$ appears due to a wider spread of $N(\\mbox{\\amm})$ values in B1. We find no variation in the $N(\\mbox{\\amm})\\,/\\,N(\\mbox{\\dia})$ ratio as a function of $N(\\mbox{H$_2$})$.  In the clustered star forming region IRAS 20293+3952, \\citet{palau07} found strong \\amm\\, and \\dia\\, differentiation associated with the presence or lack of embedded YSOs. Low $N(\\mbox{\\amm})\\,/\\,N(\\mbox{\\dia}) \\sim 50$ values were found near an embedded YSO cluster, while higher $N(\\mbox{\\amm})\\,/\\,N(\\mbox{\\dia}) \\sim 300$ were found towards cores with no associated YSOs. In a study of two dense cores, each containing a starless main body and a YSO offset from the core center, \\citet{hotzel04} find a factor $\\sim 2$ smaller $X(\\mbox{\\amm})\\,/\\,X(\\mbox{\\dia})$ values towards the starless cores compared with \\citeauthor{palau07}, with abundance ratios ($\\sim 140-190$, towards the starless gas, and $\\sim 60-90$ towards the YSOs) similar to those found in this study.  The relative abundance variation found by \\citeauthor{hotzel04} is driven by a varying \\amm\\, abundance while $X(\\mbox{\\dia})$ remains constant over the cores. Although we find both a decrease in \\dia\\, column density as well as an increase in \\amm\\, column density drives the greater fractional \\amm\\, abundance towards Oph B1 (see Figure \\ref{fig:compare_column}), both $X(\\mbox{\\dia})$ and $X(\\mbox{\\amm})$ are greater in B1, by factors of $\\sim 2$ and $\\sim 4$, respectively, compared with B2 given the lower H$_2$ column densities found in B1 in \\S\\ref{sec:nh2}. In a survey of 60 low mass cloud cores, \\citet{caselli02} find correlations between $N(\\mbox{\\amm})$ and $N(\\mbox{\\dia})$ in starless cores, but the relative column density values do not vary significantly between starless and protostellar objects.  Based on this work and the studies described above, it appears that the relative fractional abundance of \\amm\\, to \\dia\\, remains larger towards starless cores than towards protostellar cores by a factor of $\\sim 2 - 6$ in both isolated and clustered star forming regions. \\edits{We find a 1-$\\sigma$ positive trend in the relative \\amm\\, to \\dia\\, abundance in Oph B2 with increasing distance to an embedded protostar, suggestive of a direct impact by the protostars on the relative abundances. In their work on x-ray emission from young protostars in Ophiuchus, \\citet{casanova95} show that protostellar x-ray emission is potentially significant enough to increase the ionization fraction in nearby dense gas, potentially affecting the local gas chemistry. In this region, the complicated line structures found in both species near the protostars make accurate column density measurements difficult, however, with resulting large uncertainties. }  In their models of collapsing prestellar cores, \\citet{aikawa05} found that \\amm\\, can be enhanced relative to \\dia\\, in core centers at high densities ($n = 3 \\times 10^5 - 3 \\times 10^6$\\,\\cc) due to dissociative recombination reactions of \\dia\\, and $e^-$ to form NH and N. NH then reacts with H$_3^+$ and H$_2$ to form \\amm. The \\dia\\, recombination reaction is dominant only where CO is depleted (if CO is abundant, \\dia\\, is destroyed mainly through proton transfer to CO), and occurs when the abundance ratio of CO to electrons, $n(\\mbox{CO})\\,/\\,n_e \\lesssim 10^3$. When a collapsing core reaches higher central densities ($n \\sim 10^7$\\,\\cc), however, \\citeauthor{aikawa05} predict that the $N(\\mbox{\\amm})$ will begin to decrease towards the core centre due to depletion, while $N(\\mbox{\\dia})$ continues to increase, leading to a higher fractional \\dia\\, abundance relative to \\amm\\, at later times. This prediction is in agreement with $X(\\mbox{\\amm})/X(\\mbox{\\dia})$ results since B2 is more evolved than B1, having already formed protostars, but seems inconsistent with the common slope we find for the decrease in $X(\\mbox{\\amm})$ and $X(\\mbox{\\dia})$ with $N(\\mbox{H$_2$})$ in both Cores. Detailed modelling of the physical and chemical structure of the (non-spherical and clumpy) Cores is beyond the scope of this study, but would help to constrain the central density and abundance structure as a function of Core radius needed for a direct comparison with the \\citeauthor{aikawa05} models.   \\subsection{Are \\dia\\, 1-0 and \\amm\\, (1,1) tracing the Oph B Core interior?}  The hyperfine structure of the \\amm\\, (1,1) inversion transition allows the calculation of volume density $n_{ex}$ as in Equation \\ref{eqn:nex}, with the opacity of a typical hyperfine component $\\tau_{\\mbox{\\tiny{\\it hf}}} = 0.233 \\tau$. The resulting mean density $n_{ex, \\mbox{\\tiny{\\amm}}} = 7.8 \\times 10^3$\\,\\cc\\, and $n_{ex, \\mbox{\\tiny{\\amm}}} = 1.7 \\times 10^4$\\,\\cc\\, for Oph B1 and B2, respectively (note that these values are slightly less than reported in \\one\\, due to using the larger 18\\arcsec\\, beam FWHM). These volume densities  are a factor $\\gtrsim 20$ less than those calculated in \\S\\ref{sec:nex} from \\dia\\, 1-0 emission ($\\sim 26$ in B1 and $\\sim 18$ in B2). There is some question whether the \\amm\\, volume densities are accurate, as recent studies \\citep[\\one; ][]{foster09} have found \\amm-derived $n_{\\mbox{\\tiny{ex}}}$ values which are an order of magnitude less ($10^4$\\,\\cc\\, versus $10^5$\\,\\cc\\, and greater) than volume densities determined from continuum emission. One possible source of error is the collisional de-excitation rate coefficient, $\\gamma_{ul}$, which \\citeauthor{foster09} note is reported in various publications with factor of $\\sim 10$ variation. Since $n_{cr} = A_{ul}\\,/\\,\\gamma_{ul}$, the critical density is therefore also suspect within a factor $\\sim 10$.  It is clear that in Oph B \\amm\\, and \\dia\\, are tracing different motions in the gas. The gas traced by \\dia\\, emission is significantly more quiescent, as has been observed in high density gas in other star forming cores. The fact that higher densities are calculated from \\dia\\, emission relative to those calculated from \\amm\\, emission is thus reasonable, and bolsters the hypothesis that \\amm\\, (1,1) emission is not tracing the highest density gas.  \\edits{How well does the \\dia\\, 1-0 emission trace the highest density gas?} In agreement with results in \\one, we find that the clumps found in \\dia\\, 1-0 emission also do not match well continuum clump locations, and significant offsets between \\dia\\, 1-0 and continuum emission are also apparent in the integrated intensity maps. In particular, both the single-dish- and interferometer-integrated \\dia\\, 1-0 intensity towards the continuum clump B2-MM8 peak offset to the clump, and in the interferometer data we see an integrated intensity minimum at the MM8 location. On larger scales, we find a trend of {\\it decreasing} \\dia\\, abundance with increasing H$_2$ column density, described above, with a slope in Oph B2 that agrees within uncertainties with that determined for $N(\\mbox{\\amm})$ versus $N(\\mbox{H$_2$})$ for both Oph B1 and B2 (see \\S\\ref{sec:trends} and \\S\\ref{sec:frac}). \\edits{Furthermore, the mean $X(\\mbox{\\dia})$ for continuum clumps alone in Oph B is lower than the $X(\\mbox{\\dia})$ over the whole core (see \\S\\ref{sec:compare_clumps}). Together, these findings suggest that \\dia\\, 1-0, though doing a better job than \\amm\\, (1,1), may not itself trace well the coldest, densest material in Oph B. Appropriate caution must be made when using even \\dia\\, 1-0 to probe dense gas in some star-forming regions.}"
    },
    "0911/0911.3055_arXiv.txt": {
        "abstract": "ANTARES is a submarine neutrino telescope deployed in the Mediterranean Sea, at a depth of about 2500 m. It consists of a three-dimensional array of photomultiplier tubes that can detect the Cherenkov light induced by charged particles produced in the interactions of neutrinos with the surrounding medium. Down-going muons produced in atmospheric showers are a physical background to the neutrino detection, and are being studied. In this paper the measurement of the Depth Intensity Relation (DIR) of atmospheric muon flux is presented. The data collected in June and July 2007,  when the ANTARES detector was in its 5-line configuration, are used in the analysis. The corresponding livetime is $724\\,h$. A deconvolution method based on a Bayesian approach was developed, which takes into account detector and reconstruction inefficiencies. Comparison with other experimental results and Monte Carlo expectations are presented and discussed. ",
        "introduction": "The largest event source in neutrino telescopes is \\textit{atmospheric muons}, particles created mainly by the decay of $\\pi$ and $K$ mesons originating in the interaction of cosmic rays with atmospheric nuclei. Although ANTARES \\cite{antares_web, antares_daq, antares_OM, antares_OM1, pcoyle} \"looks downwards\" in order to be less sensitive to signals due to downward going atmospheric muons, these represent the most abundant signal due to their high flux. They can be a background source because they can be occasionally wrongly reconstructed as upward going particles mimicking muons from neutrino interactions. On the other hand they can be used to understand the detector response and possible systematic effects. In this scenario the knowledge of the underwater $\\mu$ intensity is very important for any Cherenkov neutrino telescope and the future projects \\cite{km3net, nemo}. Moreover, it would also provide information on the primary cosmic ray flux and on the interaction models.  \\\\ ",
        "conclusions": "{}} The aim of the presented analysis is the measurement of the muon flux at the depth of ANTARES and the derivation of the vertical component of the atmospheric muon flux as a function of the sea depth. The goal is also to assess the performance of ANTARES in detecting muons. The analysis has been performed on a selection of the experimental data of June and July 2007 when the ANTARES detector was in its 5-line configuration.  Several quality cuts have been applied on the reconstructed events in order to improve their purity, in particular concerning the zenith angle reconstruction.  An unfolding algorithm, based on an iterative method, has been applied on the selected experimental data in order to retrieve back the flux of atmospheric muons with $E_\\mu>20\\,GeV$ at the fixed sea depth $h_0=2000\\,m$. The experimental DIR was finally obtained.  The results are in good agreement, within the uncertainties, with the experimental fluxes obtained by other Cherenkov telescopes.       \\begin{figure}[!t] \\begin{center} \\includegraphics[width=8. cm] {bazzotti_fig5} \\caption{\\label{DIRallNEW} \\textbf{PRELIMINARY}. Depth Intensity Relation of atmospheric muons for $E_\\mu>20\\,GeV$, with systematic uncertainties (the statistical uncertainties are negligible). The DIR obtained from other underwater measurements are also shown: Higashi \\cite{d-higa}, Davitaev \\cite{d-davitaev}, Vavilov \\cite{d-vavilov}, Fyodorov \\cite{d-fyodorov}, DUMAND-SPS \\cite{dumand}, BAIKAL NT-36 \\cite{baikal2}, NESTOR \\cite{d-nestor}, AMANDA B-4 \\cite{d-amandab4}, AMANDA-II \\cite{d-amandaII}. The Sinegovskaya parameterization ($E_\\mu>20\\,GeV$) \\cite{sine} and the MUPAGE simulation are superimposed.} \\end{center} \\end{figure}"
    },
    "0911/0911.2786_arXiv.txt": {
        "abstract": "This talk briefly explains how the breaking of a Lorentz-invariant description of nature at tiny space-time intervals might affect the non-Gaussian character of the primordial fluctuations left by inflation.  For example, a model that contains irrelevant operators that only preserve the spatial symmetries along constant-time surfaces can generate a larger non-Gaussian component in the pattern of primordial fluctuations than is ordinarily predicted by inflation.  This property can be useful for constraining models that allow some Lorentz violation at short distances, beyond the constraints possible from the power spectrum alone. ",
        "introduction": "The universe today is filled with a quite bewildering variety of structures, from stars and galaxies at comparatively small scales (when judged by the size of the observable universe) to the vast clusters and networks of galaxies and emptier voids among them at the largest scales.  But the earlier universe differs markedly from its current appearance.  The earliest relics that we can see directly show a remarkable uniformity in how the material was distributed at those times.  For example, the temperature of the radiation that first escaped as the universe cooled from a hot plasma into a transparent gas varies less than one part in one hundred thousand from one place to another.  Moreover, the abundances of the primordially produced lighter elements also seem to be the same everywhere in the universe, suggesting that this uniformity extends to still earlier times.  Yet the universe needed to have had at least {\\it some\\/} variation, even at the very earliest of times, for without any tiny spatial variations, the growth of the structures that we see today could never have begun in the first place.  The theory of inflation provides one elegant mechanism for explaining the origin of these tiny {\\it primordial fluctuations\\/}.  One very appealing property of the inflationary picture is that such fluctuations are unavoidable, being a simple consequence of the two basic ingredients needed in any inflationary model.  A typical model for inflation requires a space-time that is expanding at an accelerating rate and a quantum field whose dynamics are responsible for the expansion.  But a quantum field is always fluctuating.  It would be quite impossible for it to be otherwise, as it would be inconsistent with the Heisenberg principle for a field to be at once stationary and without any spatial fluctuations.  Let us describe this behavior a little more precisely.  If we denote the quantum field by $\\varphi(t,\\vec x)$ and its quantum state by $|0(t)\\rangle$, then while its mean value in this state may vanish on average, \\begin{equation} \\langle 0(t)| \\varphi(t,\\vec x) |0(t)\\rangle = 0 , \\label{onepoint} \\end{equation} the variance of the field never does, \\begin{equation} \\langle 0(t)| \\varphi(t,\\vec x) \\varphi(t,\\vec y) |0(t)\\rangle \\not= 0 . \\label{variance} \\end{equation} Of course, such quantum fluctuations are also occurring at tiny scales today.  The difference between the early universe and the universe today is that, during the former, the space-time expands very rapidly, stretching these tiny fluctuations to very large scales---even as large as the observable universe today.  Because these fluctuations are such a basic consequence of inflation, the pattern that inflation predicts for them is fairly insensitive to the details of a particular model, though measurements have actually reached a precision that can test inflation beyond its most basic predictions and can even exclude individual models.  Nevertheless, the overall inflationary picture is still in a superb agreement with what is inferred from observations.  Inflation predicts \\begin{itemize} \\item that the fluctuations have more or less the same amplitude at all length scales---including those that otherwise would not have been in causal contact without an inflationary period, \\item that the fluctuations are all in phase and are adiabatic, \\item that their pattern is largely a {\\it Gaussian\\/} one, the meaning of which we shall explain later, and \\item that there should also be a background of primordial gravity waves. \\end{itemize} Each of these predictions---except for the last---matches what is seen in all experiments so far.  With this success, it is important to test the underlying assumptions of the inflationary picture further and to shed some light on its otherwise mysterious elements.  What exactly is this field that is producing the inflation and how does it fit into the rest of our picture for particle physics?  How did the universe find itself initially in a spatially flat configuration, at least to a sufficient degree over a sufficiently large patch?  Does the size of the field change rapidly over small enough distances that we ought to worry about quantum gravity?  Which is the correct quantum state $|0(t)\\rangle$ to choose?  And is it possible to have {\\it too much\\/} expansion of the quantum fluctuations for the consistency of our picture?  If these questions seem to be only so much cavilling, it instructive to compare the setting used in inflation with that of ordinary particle physics to realize how much less we know about the former and how much more limited is our ability to test it.  In particle physics, \\begin{itemize} \\item We know the background (flat space). \\item We know all the relevant ingredients up to some energy scale (and any unseen stuff remains largely decoupled from the lower energy world). \\item We have a well developed, {\\it and well tested\\/}, theoretical framework (quantum field theory in flat space). \\item Gravity is entirely negligible. \\item And perhaps most importantly, we are free to test anything we want, in a controlled environment, up to a limiting energy scale. \\end{itemize} Now contrast each of these points with what we assume for the early universe. \\begin{itemize} \\item We probably know the background (it seems to be approaching a spatially independent one at early enough times). \\item We can only guess what are relevant ingredients, and have no way to observe them directly (at least so far, and probably for a long while yet to come). \\item We have a framework, quantum field theory in curved space; and while it is generally self-consistent, it has not been tested. \\item Gravity is an essential ingredient. \\item And experimentally, we can only observe a very limited and indirect set of evidence, and we are not able to test the picture directly (at least not to the standard applied in an accelerator experiment). \\end{itemize} Nothing on this list implies that the inflationary picture is necessarily wrong; but each item does indicate some unresolved piece of the inflationary framework which we would like to understand better.  Because of its remoteness from anything that we can test directly and because it will probably be difficult to find {\\it complementary\\/} observations to those that we already have which can test this picture, it is important to address some of these questions if we are to have a little more confidence that inflation really did occur in our own universe. ",
        "conclusions": ""
    },
    "0911/0911.0921_arXiv.txt": {
        "abstract": "We evaluate the achievable maximum energy of nuclei diffusively accelerated by shock wave in the jet of Cen\\,A, based on an updated model involving the stochastic magnetic fields that are responsible for recent synchrotron X-ray measurements. For the maximum energy analysis, conceivable energy constraints from spatiotemporal scales are systematically considered for the jet-wide including discrete X-ray knots. We find that in the inner region within $\\sim 1\\,{\\rm arcmin}$ from galactic core, which includes knots AX and BX, proton and iron nucleus can be accelerated to $10^{19}-10^{20}$ and $10^{21}~{\\rm eV}$ ($10-100~{\\rm EeV}$ and ${\\rm ZeV}$) ranges, respectively. The upper cutoff energy of the very energetic neutrinos produced via photopion interaction is also provided. These are essential for identifying the acceleration site of the ultra-high energy cosmic ray detected in the Pierre Auger Observatory, which signifies the arrival from nearby galaxies including Cen\\,A. ",
        "introduction": "To date, the 27 ultra-high energy cosmic ray (UHECR) events of the energy exceeding $5.7\\times 10^{19}~{\\rm eV}$ have been detected in state-of-the-art Auger observatory; in particular, the anisotropy and significant correlation with the active galactic nuclei (AGNs) residing within $75~{\\rm Mpc}$, namely GZK horizon \\citep{greisen,zatsepin}, were discovered \\citep{abraham07}. Surprisingly, the results indicate that two of these UHECR events have arrived within $3\\degr$ of Centaurus\\,A (NGC\\,5128), a closest galaxy \\citep[$3.7~{\\rm Mpc}$ from us;][]{ferrarese}. Although the detailed position of the UHECR production site is still unresolved, the galactic core accompanied by a supermassive black hole, bipolar jets, giant radio lobes \\citep{hardcastle09}, and so on, will be enumerated as the favored candidates. It is desired that genuine theoretical survey be expanded for clarifying the particle acceleration mechanism feasible at these sites, in order to provide the physical interpretation of the observed intriguing results, which include the recent detection of high-energy gamma rays \\citep{aharonian}.  When closely looking at the anatomy of the large-scale jets in the Cen\\,A galaxy \\citep[see,][for a review]{israel}, a clue to solve this challenging problem seems to be in the knotty regions resolved in a deep X-ray image \\citep[e.g.,][]{kraft02}. In an inner region near the galactic core, the rugged features are associated with the shocks, which are considered to be formed where a flowing plasma collides with obstacles \\citep{hardcastle07}. According to the latest theory of plasma kinetic transport, it is known that such a colliding plasma is likely unstable for the electromagnetic current filamentation instability, which generates small-scale magnetic fluctuations with the order of plasma skin depth (\\citealt{medvedev}; \\citealt{honda04} and references therein). In the nonlinear phase, the filaments preferentially coalesce one another, self-organizing larger scale filaments with stronger magnetic fields. This reflects a stochastic nature of the inverse energy cascade of plasma turbulence. The magnetic fluctuations are non-perturbative, and in the strong turbulence regime, in that the energy density is comparable to the thermal energy density of plasma bulk \\citep{honda00a,honda00b}.  Indeed, recent high-resolution observations reveal that the jet of Cen\\,A is in part dominated by the filamentary, sometimes edge-brightened features \\citep{hardcastle07}, as inspected in the well-confirmed jet of a nearby galaxy \\citep[Vir\\,A;][]{owen}. Worth noting thing is that inherent inner structure similar to these has been discovered in many other jets (e.g., Cyg\\,A: \\citealt{perley}; 3C\\,353: \\citealt{swain}; 3C\\,273: \\citealt{lobanov}; 3C\\,438: \\citealt{treichel}; Mrk\\,501: \\citealt{piner}). Also, it was recently found that a filamentary jet model naturally provides the comprehensive explanation for the complicated spectral variability observed in a TeV blazar object \\citep[Mrk\\,421;][]{honda08}. These ingredients strongly encourage us to take the inhomogeneous magnetic field effects into account for precisely modeling the Cen\\,A jet as a cosmic-ray accelerator.  For a standard, diffusive shock acceleration (DSA) scenario \\citep[e.g.,][]{blandford87}, the strong turbulence plays an essential role in particle scatterers is expected. In the stochastic medium, the higher energy heavy particles tend to freely meander \\citep{hh05}, albeit electrons (and positrons) are, if anything, likely gyro-trapped in the magnetized filaments, suffering stronger radiative cooling (Section~3.1). From the fact that in the Cen\\,A jet, electron synchrotron emissions appear to be, in part, already diffusive, it is thus inferred that for nuclei (say, proton) the conventional approximation of small-angle resonant scattering will be inadequate for describing the spatial diffusion, and instead, the three-dimensional rms deflection becomes rather feasible. Importantly, the latter facilitates the back and forth of particles across the shock, increasing the efficiency of DSA. Based on this notion, \\citet{hh04} have argued that a nucleus could be accelerated to the UHE range at a bright jet knot of Vir\\,A, though the arrival from the galaxy ($\\sim 4$ times more distant than the distance to Cen\\,A) is not yet signified \\citep[see, e.g.,][for the useful discussions]{stanev}.  Besides, the bound electrons are co-accelerated, to emit the synchrotron photons, which could serve as a target of the accelerated protons. Then, the resonant interaction of the UHE protons with the target photons could be a major neutrino production process via the decay channels triggered by photopionization: $p\\gamma\\rightarrow \\pi^{\\pm}X\\rightarrow\\mu^{\\pm}\\nu_{\\mu} \\rightarrow e^{\\pm}\\nu_{e}\\nu_{\\mu}$ \\citep{romero}. The emitted neutrinos propagate on the straight, unlike the charged particles that suffer, more or less, deflection by intergalactic magnetic fields \\citep[e.g.,][]{vallee}, so that they play the role of a powerful and complementary messenger of the UHECR production, in light of the observability at a forthcoming kilometer neutrino telescope \\citep[e.g.,][]{halzen}. Recently, \\citet{cuoco} have proposed the model spectra of high-energy neutrinos from Cen\\,A, but putting the detailed mechanisms of the particle accelerator into effect has remained unsolved.  In this paper, based on the filament model, we estimate the maximum possible energies of a proton and iron diffusively accelerated in the Cen\\,A jet, and also, the energy of the produced neutrinos, providing the energy-equipartition among the flavors. The conceivable mechanisms of the energy restriction are taken into consideration for the jet wide including bright knots. It is addressed that the energy is limited dominantly by the shock operation time or particle escape loss, rather than radiative loss and nucleus--nucleus collision timescales. As a result, we find that in the inner region, which contains the X-ray knots BX$n$, proton and iron can be accelerated beyond the Auger limit. In particular, the expected highest energy reaches, for iron, $3~{\\rm ZeV}$ and more (around knot BX5); that is to say, the Cen\\,A jet is a \"Zevatron\" worthy of the candidate (in addition to the Vir\\,A, put forth by \\citet{hh04}).  The present analysis virtually provides a substantially extended version of the previous simple analysis by \\citet{romero}, and hence, a particular attention has been paid to highlight the new points including (1) the proposal of the improved theoretical model compatible with updated X-ray measurements (Section~2.1) and extended arguments on the turbulent magnetic field (Section~2.2), (2) the non-resonant diffusion scenario involved in the DSA (Section~3.1), and (3) thorough survey of the temporal (Section~3.2) and spatial (Section~3.3) limits, in order to figure out the maximum energies of proton and iron achievable in the inner region (Section~3.4), and estimate neutrino energy (Section~3.5). The application of the maximum energy analysis to the outer region of the large-scale jet is also provided (Section~4). At last, the discussion on the comparison with the previous relevant results is expanded (Section~5). ",
        "conclusions": "The validity of the filamentary model has been checked by reproducing the complicated blazar variabilities that involve the non-monotonous hysteresis patterns of broadband electromagnetic spectrum in flare phases \\citep{honda08}. By recalling the idea that regards FR-I radio sources as misaligned BL\\,Lac objects \\citep{urry,tsvetanov}, the application of the filamentary model to the Cen\\,A jet will be reasonably justified. As for the particle acceleration in the filamentary medium, we point out that the present DSA mechanism for nuclei, which owes to the (off-resonant) three-dimensional rms diffusion across the shock (Section~3.1), is simpler than the electron acceleration mechanism that is incorporated with the complexity of the energy hierarchy mediated by the transition injection \\citep{hh07}. The acceleration of highest energy particle in the Cen\\,A was first considered by \\citet{romero}. In the early work, they dealt with the conventional gyro-resonant diffusion model even for proton, that is to say, implicitly supposed a large-scale ordered magnetic field (with small-amplitude perturbations). They addressed that a proton could be accelerated to the energy of $2.7\\times 10^{21}~{\\rm eV}$ in situ, which is an order of magnitude larger than the present $E_{p,m}$ values that are in the range of $\\sim 10^{19}-10^{20}~{\\rm eV}$ (Section~3.4; Table~\\ref{tbl:2}). Such a larger value was derived from simply equating a classical acceleration timescale \\citep[e.g.,][]{biermann} with radiative loss timescales, but this fashion now appears to be somewhat optimistic (Section~3.2.2; Figure~2). It can be claimed that taking into account of the shock operation time or particle escape is essential for the maximum energy analysis for the concerned source, to yield the reduced proton energy, which can hardly reach the $10^{21}~{\\rm eV}$ energy range. It is also mentioned that the situation in which the dynamical timescale limits the particle acceleration is analogous to the situation that typically appears in the supernova remnant shock acceleration, in which the adiabatic expansion of spherical shell likely becomes a major loss mechanism \\citep[e.g.,][]{kobayakawa}.  In conclusion, we have evaluated the achievable maximum energies of the proton and iron nucleus diffusively accelerated by the shock in the Cen\\,A jet including X-ray knots. In particular, we have taken into account of the more realistic DSA scenario that relies on the filamentary jet model responsible for the recent X-ray measurements, and elaborated the conceivable energy restriction stemming from spatiotemporal scales. The key finding is that, for the plausible ranges of the shock speed of $\\sim 0.1c$, viewing angle of $\\gtrsim 10\\degr$, and magnetic intensity of $\\lesssim 500~{\\rm\\mu G}$, the uppermost particle energy tends to be inevitably limited by the shock accelerator operation timescale ($<10^{13}~{\\rm s}$), rather than the radiative cooling losses ($\\gtrsim 10^{13}~{\\rm s}$). As a result, it has been demonstrated that there exists the acceleration region (off the galactic core), in which proton and iron can be energized to $\\sim 10^{20}~{\\rm eV}$ and $\\sim 10^{21}~{\\rm eV}$ ranges, respectively. In particular, for the inner region including the X-ray knot BX5, we work out at the maximum energies of $1\\times 10^{20}~{\\rm eV}$ and $3\\times 10^{21}~{\\rm eV}$ for proton and iron, respectively; and estimate the corresponding cutoff energy of the neutrino produced via the photopionization, as $6\\times 10^{18}~{\\rm eV}$. For the stationary shock case, we can expect the enhancement of the maximum energies, and read that the large-scale region up to the projection of about $200\\arcsec$, as well, has a potential to energize proton and iron beyond $10^{20}~{\\rm eV}$ and $10^{21}~{\\rm eV}$, respectively, to yield the $p\\gamma$ neutrino with the energy exceeding $5\\times 10^{18}~{\\rm eV}$. We here manifest that the jet wide of the Cen\\,A galaxy is a promising candidate for the UHE proton accelerator and Zevatron for high-$Z$ nucleus, as desirable to account for the outcome from the Auger observatory \\citep{abraham07}.  The derived large values of particle energies are of those achievable at the acceleration sites. The observable energies ought to be, of course, reduced, due to a major energy loss mechanism including the photopion interaction with cosmic microwave background. However, the Cen\\,A is so nearby that the ballistic transport is anticipated to be not crucially degraded, on account of the weak decay property (such as $-(1/E)(dE/dt)\\sim 10^{-8}~{\\rm yr^{-1}}$ for $E>10^{20}~{\\rm eV}$; \\citet{romero}). We hope, in near future, the application of the present scenario to cosmologically distant (super GZK) AGNs, to solidify the point source scenario for dominant UHECR production, which is responsible for the suppression of cosmic-ray flux above $4\\times 10^{19}~{\\rm eV}$ that has recently been confirmed by the Auger's experiments \\citep{abraham08}, and also, is of interest in conjunction with the detection of super GZK neutrinos \\citep[e.g.,][]{ringwald}.  I am grateful to Y.~S.~Honda for a useful discussion.\\\\"
    },
    "0911/0911.2115_arXiv.txt": {
        "abstract": "We present new and archival multi-frequency radio and X-ray data for Centaurus A obtained over almost 20 years at the VLA and with {\\it Chandra}, with which we measure the X-ray and radio spectral indices of jet knots, flux density variations in the jet knots, polarization variations, and proper motions.  We compare the observed properties with current knot formation models and particle acceleration mechanisms.  We rule out impulsive particle acceleration as a formation mechanism for all of the knots as we detect the same population of knots in all of the observations and we find no evidence of extreme variability in the X-ray knots.  We find the most likely mechanism for all the stationary knots is a collision resulting in a local shock followed by a steady state of prolonged, stable particle acceleration and X-ray synchrotron emission.  In this scenario, the X-ray-only knots have radio counterparts that are too faint to be detected, while the radio-only knots are due to weak shocks where no particles are accelerated to X-ray emitting energies.  Although the base knots are prime candidates for reconfinement shocks, the presence of a moving knot in this vicinity and the fact that there are two base knots are hard to explain in this model.  We detect apparent motion in three knots; however, their velocities and locations provide no conclusive evidence for or against a faster moving `spine' within the jet.  The radio-only knots, both stationary and moving, may be due to compression of the fluid. ",
        "introduction": "\\label{introjlg}  It is generally agreed that the observed emission from Fanaroff-Riley class I \\citep[FR\\,I;][]{fr74jlg} radio jets is due to the synchrotron process at all wavelengths, with similar jet structure observed from the radio through the optical into the X-ray \\citep[e.g.][]{hardcastle02jlg,harris02jlg}.  The jets of FR\\,I radio galaxies are thought to decelerate as they move away from the core, entraining material and expanding into a plume of diffuse matter \\citep[e.g.][]{bicknell84jlg}.  One of the most significant implications of the synchrotron emission model is reflected in the characteristic loss timescales: in a stable environment, the X-ray emitting electrons have lifetimes of the order of tens of years, tracing regions of current, in situ particle acceleration, while the radio emitting electrons last for hundreds of thousands of years, showing the history of particle acceleration in the jet.  In order to investigate these regions of particle acceleration we need data with sensitivity and resolution sufficient to detect jet substructure on spatial scales comparable to the synchrotron loss scales.  This prompts us to look to the two closest bright FR\\,I radio jets: M87 and Centaurus A (NGC\\,5128, hereafter Cen A).  Both of these jets have been detected in multiple frequencies from the radio through to the X-ray \\citep[e.g.][]{feigelson81jlg,kraft02jlg,hardcastle03jlg,hardcastle06jlg,harris02jlg}. The proximity of these radio galaxies, 16.7\\,Mpc and 3.7\\,Mpc respectively \\citep{blakeslee09jlg,mei07jlg,ferrarese07jlg}, make them unique jet laboratories with spatial scales of 77\\,pc and 17\\,pc per arcsec respectively.  The details revealed in the structure of these jets have been the focus of many recent studies \\citep[e.g.][]{biretta99jlg,hardcastle03jlg,kataoka06jlg,cheung07jlg}. Within the smooth surface brightness observed in both of these jets are clumps of bright material -- the knots -- embedded in diffuse material, all emitting via synchrotron emission.  The precise mechanisms causing the particle acceleration responsible for the diffuse structure and the knots are still unknown.  The most surprising result of recent observations of these two systems was the radio-to-X-ray synchrotron flare of HST-1 in M87. In 2002, the X-ray flux of HST-1 increased by a factor of 2 in only 116 days \\citep{harris03jlg}, implying a change within an emitting volume with a characteristic size less than 0.1\\,pc for a stationary source (much less than the size of HST-1, $\\sim 3$\\,pc).  The X-ray brightness then faded in the following months only to flare again, peaking in 2005.  At its brightest, this flare was higher than its 2001 level by a factor of $\\sim50$.  The UV and radio light-curves were found to vary in step with the X-ray up to this peak \\citep{perlman03jlg,harris06jlg}, but the subsequent decrease appeared to drop off faster in the X-ray than in either the optical or the UV, which drop off in step \\citep{harris09jlg}.  In additional to this spectral variability, it was established with the {\\it Hubble Space Telescope (HST)} and the NRAO Very Long Baseline Array (VLBA) that some of the knots in M87 move superluminally, including subregions of the HST-1 knot \\citep{biretta99jlg,cheung07jlg}.  Together, this suggests that we are observing synchrotron losses in addition to either beaming or compression/rarefaction of the fluid.  Cen A is a factor of 4.5 closer than M87 so we can resolve more details in the complicated fine structure of the jet.  The originally identified features, named A-G by \\citet{feigelson81jlg}, have since been resolved into at least 40 individual knots \\citep{kraft02jlg,hardcastle03jlg} with additional emission from diffuse material.  Some of the diffuse emission has been described as downstream `tails' of emission from the knots \\citep{hardcastle03jlg} or as evidence for limb-brightening of the jet \\citep{kraft00jlg}.  In 2003, Hardcastle et al. presented 8.4\\,GHz radio observations from the NRAO Very Large Array (VLA) of Cen A.  That work used archival data from 1991 and new observations from 2002 to study the jet knots and investigated the offsets and relationships between the radio knots and their X-ray counterparts and vice versa.  They found that only some of the radio knots appeared to have X-ray counterparts, leaving many as `radio-only' knots and `X-ray-only' knots.  They also considered the temporal changes in the radio knots, specifically their proper motions, finding that some of the radio knots were moving.  These moving knots had comparatively little X-ray emission suggesting that high-energy particle acceleration is less efficient in these regions than in the jet as a whole.  Some of the current models explaining the presence of knots within the generally smooth diffuse material of the jet include compressions in the fluid flow, collisions with obstacles in the galaxy causing local shocks, reconfinement of the jet or some other jet-wide process, and magnetic reconnection.  \\citet{hardcastle03jlg} ruled out simple compression of the fluid as a mechanism for producing X-ray bright, radio faint compact knots in favor of in situ particle acceleration associated with local shocks; however, compression could still play a part in the other knots.  They concluded that the most likely model to describe the majority of these knots is an interaction between the jet fluid and an obstacle such as a molecular cloud or a high mass-loss star.  By exploring the temporal behavior of the X-ray and radio emission, we can understand the evolution of the knots and constrain the various models of particle acceleration used to describe the jet features.  In this work we use Chandra and VLA data spread over almost 20 years to measure the X-ray and radio spectral indices knots, flux density variations, polarization variations, and the proper motions of the jet knots in Cen A. Our aims are to detect variability in the radio and X-ray properties of the knots, either extreme variability similar to that of HST-1 in M87 or more subtle changes, and to compare these properties. which will allow us to constrain the knot formation processes at work in the jet of Cen A.  The details of our radio and X-ray data reduction are discussed in Section~\\ref{sec:data}.  In Section~\\ref{sec:res} we discuss the details of our analysis methods and the global results for the knot population, highlighting particularly interesting features.  In Section~\\ref{sec:dis}, we compare these knot properties with the predictions of various models for the formation of knots, their particle acceleration and the jet structure.  Finally we outline the most likely processes for forming knots in Cen A in Section~\\ref{sec:con}. ",
        "conclusions": "\\label{sec:con}  Our results can be summarized as follows:  \\begin{itemize}  \\item We rule out impulsive particle acceleration in the knots of Cen A as we detect no extreme variability in the X-ray knots, in contrast to what is seen in knot HST-1 in M87.  We see essentially the same distribution of X-ray knots in our most recent observation as was seen in the earliest {\\it Chandra} observations in 1999. This would not be the case if the knots were impulsive as they would fade due to synchrotron losses indicating long-lived particle acceleration in the knots of Cen A.  \\item For those radio knots with X-ray counterparts, the most likely formation mechanism is a collision between the jet and an obstacle, resulting in a local shock.  We see no significant variability in many of these knots, suggesting a long-lived, stable stage of particle acceleration during the interaction between the jet and the obstacle.  \\item The formation of knots at the point where the inner hundred-parsec-scale jet broadens abruptly suggests that these base knots (A1A and A1C) may be reconfinement shocks; however, this is complicated by the presence of a radio-only knot (A1B) moving downstream between the possible confinement-shock knots.  \\item We detect a factor of 3 increase in radio flux density of the counterjet knot SJ1.  This knot lies only 17\\,pc from the nucleus so is unresolved in the X-ray; however, it was still increasing in flux in the most recent observation (Dec 2008) so we plan to continue to monitor its radio behavior with the VLA.  \\item We detect proper motions in three of our radio knots; two of which have no compact X-ray counterparts and a third which has only diffuse X-ray emission associated with it.  Studies of the distribution of the moving knots are inconclusive due to the low number of well-established proper motions; however, that the direction of motion of the knots may not be directly parallel to the jet axis which appears to varies along the jet.  Their motions are all downstream and they show no dependence on the position of the knot within the jet.  \\end{itemize}  The most likely cause of knots in the jet is collisions; if the X-ray-only knots have faint radio counterparts and the radio-only knots are seen only during the latter stages of the collision when the interaction is weaker, then only the moving knots are a separate population.  These may include all the radio-only knots but our proper motion measurements are inconclusive for many of these.  We argue that these moving knots are due to compressions in the fluid flow that do not result in particle acceleration to X-ray emitting energies.  It is possible, however, that the X-ray-only knots are also a separate population with flatter X-ray to radio spectra than those with counterparts, in which case we currently have no model for their formation.  Compared to other FR\\,I jets, Cen A is atypical, with an obscuring dust lane extending out to 1\\,kpc from the core which greatly affect the jet and its knots.  Other galaxies where dust has been detected, such as 3C31 and 3C449, have much smaller disks, which cannot affect even the innermost regions of the observed X-ray jet.  If we can attribute the knot dominated particle acceleration of the inner kpc to the presence of this disk then we can postulate that the X-ray jet emission seen in other FR\\,I galaxies should be comparable to the dominant diffuse particle acceleration that dominates farther out in the Cen A jet.  We would then predict that knot-dominated structure will not be seen in other FR\\,I galaxies.  \\vspace{-10pt}"
    },
    "0911/0911.0326_arXiv.txt": {
        "abstract": "The study of scattering and extinction properties of possible nanodiamond grains in the ISM are reported. Calculations using Discrete Dipole Approximation (DDA) for varying ellipsoidal shapes and sizes from $2.5$ to $10 ~nm$ are considered. Nanodiamonds show negligible extinction from IR to near-UV and very sharp far-UV rise. Comparison with observations rule out possibility of independent nanodiamond dust but point towards possibility of nanodiamonds as a component in the ISM. Radiation induced transformations may lead to carbonaceous grains with different core and mantles. So calculations are also performed for a core-mantle target model with nanodiamond core in graphite mantles. The graphite extinction features get modified with the peak at 2175 \\AA{} being lowered, broadened, blue shifted and accompanied by enhanced extinction in the far-UV. Such variations in the 2175 \\AA{} band and simultaneous far-UV rise are observed along some sources. A three component dust model incorporating silicate, graphite and graphite with nanodiamond core is also considered. The model extinction compares very well with the average galactic extinction in the complete range from $0.2$ to $10 ~\\mu m^{-1}$. The best fit requires small size and small number of nanodiamonds. ",
        "introduction": "To account for the ultraviolet extinction in the diffuse interstellar medium \\citet{saslaw69} first proposed the possibility of diamonds in Interstellar Medium (ISM). The 3.43 and 3.53 $\\mu m $ emission bands in circumstellar media of Ae/Be Herbig stars HD~97048 and Elias~1 show convincing presence of nanodiamonds \\citep{guillois99}. % \\citet{kerckhoven02} attribute these bands to C-H stretching in hydrogenated nanodiamond, as distinct from the 3.3 $\\mu$m PAH feature. The 3.47 $\\mu m $ feature in absorption toward a number of proto-stars is attributed to the tertiary C-H stretching mode in diamond like structures \\citep{allamandola92}. Spectra of diamondoid molecules \\citep{pirali07} show that nanodiamonds a few nanometre in size could be responsible for the 3.53 and 3.43 $\\mu$m emission lines and smaller diamondoids give the 3.47 $\\mu$m absorption feature.  Cosmic nanodiamonds are also detected in primitive carbonaceous meteorites of presolar origin \\citep{lewis87, daulton96, dai02, garai06}. In fact, nanodiamonds are considered to be the most abundant presolar component in meteorites \\citep{anders93, zinner98}. \\citet{jones04} studied $C-H$ stretching mode of nanodiamonds extracted from Orgueil meteorite and observed the above mentioned IR bands. If the ISM nanodiamonds are not hydrogenated their detection is hard \\citep{kerckhoven02} and they could be more abundant than estimated by the observations of infrared $C-H$ bands \\citep{dedeigo07}.  It is, therefore, important to incorporate nanodiamonds in UV extinction models for which their independent scattering and extinction properties need to be understood. The extinction properties of nanodiamond dust have been reported \\citep{mutschke04} for spherical shape and the far-UV break of spectral energy in quasars is attributed to nanodiamond dust \\citep{binette05, binette06}. In this communication extinction properties of nanodiamond grains of different ellipsoidal shapes and sizes are reported. To study the effect of nanodiamond within graphite extinction efficiencies are also reported for nanodiamond-graphite as core-mantle grains. An extinction curve modeling is proposed including nanodiamond as a component. ",
        "conclusions": "Various experimental, theoretical and observational studies suggest the possibility of carbon in diamond form in the ISM. Besides enhanced far-UV extinction \\citep{binette05, binette06}, luminescence from diamond nano-crystals could be responsible for the Extended Red Emission (ERE) \\citep{chang06}. The study of extinction properties of possible nanodiamond grains of different shapes and sizes is done to understand their role in far-UV extinction and in modification of the overall extinction curve.  The extinction due to nanodiamonds of all shapes and sizes is essentially similar with negligible extinction from IR to near UV range and sharp rise in extinction in the far-UV. The extinction due to non-spherical shape is higher than that due to spherical. For nanosized particles the extinction efficiency increases linearly with effective radius.  Considering nanodiamond-graphite core-mantle target, modification in the 2175 \\AA{} peak and graphite extinction is studied. In general the 2175 \\AA{} peak gets lowered, broadened and there is enhanced far-UV extinction. Increase in \\% of $ sp^3 $ character in graphite gives rise to higher far-UV extinction. The three component modeling of average galactic extinction gives best results for small sized and low \\% of nanodiamond. This enables us to explain far-UV rise and simultaneous modification of the 2175 \\AA{} bump without putting constraints on the interstellar abundances.  Extinction from specific objects that show enhanced far-UV rise can be attempted by incorporating a nanodiamond component. This will enhance understanding of radiation induced transformations of carbonaceous matter in the ISM."
    },
    "0911/0911.3148_arXiv.txt": {
        "abstract": "Radial velocities measured from near-infrared spectra are a potentially powerful tool to search for planets around cool stars and sub-stellar objects. However, no technique currently exists that yields near-infrared radial velocity precision comparable to that routinely obtained in the visible. We are carrying out a near-infrared radial velocity planet search program targeting a sample of the lowest-mass M dwarfs using the CRIRES instrument on the VLT. In this first paper in a planned series about the project, we describe a method for measuring high-precision relative radial velocities of these stars from $K$-band spectra. The method makes use of a glass cell filled with ammonia gas to calibrate the spectrograph response similar to the ``iodine cell'' technique that has been used very successfully in the visible. Stellar spectra are obtained through the ammonia cell and modeled as the product of a Doppler-shifted template spectrum of the object and a spectrum of the cell, convolved with a variable instrumental profile model. A complicating factor is that a significant number of telluric absorption lines are present in the spectral regions containing useful stellar and ammonia lines. The telluric lines are modeled simultaneously as well using spectrum synthesis with a time-resolved model of the atmosphere over the observatory. The free parameters in the complete model are the wavelength scale of the spectrum, the instrumental profile, adjustments to the water and methane abundances in the atmospheric model, telluric spectrum Doppler shift, and stellar Doppler shift. Tests of the method based on the analysis of hundreds of spectra obtained for late M dwarfs over six months demonstrate that precisions of $\\sim$\\,5\\,m\\,s$^{-1}$ are obtainable over long timescales, and precisions of better than 3\\,m\\,s$^{-1}$ can be obtained over timescales up to a week. The obtained precision is comparable to the predicted photon-limited errors, but primarily limited over long timescales by the imperfect modeling of the telluric lines. ",
        "introduction": "\\subsection{Motivation for NIR radial velocities} The vast majority of known exoplanets have been detected with high-precision (i.e. $\\sigma$ $<$ 10\\,m\\,s$^{-1}$) time-series stellar radial velocity measurements. Invariably, these measurements have been made using spectra obtained in the visible wavelength region. This is natural for two reasons. First, arguably the most interesting targets for planet searches are solar-type stars (broadly defined here as dwarfs with F, G, and K spectral types), and these stars are brightest at wavelengths shorter than 1\\,$\\mu$m. The spectra of solar-type stars at these wavelengths are also rich in deep and sharp spectral lines suitable for Doppler shift measurements. The second reason is that visible wavelength spectrograph technology is much more advanced relative to that required for spectrographs operating in other wavelength regions. When fed by a moderately sized telescope, visible wavelength spectrographs using a cross-dispersed echelle design and CCD detectors can yield high-resolution and high signal-to-noise (S/N) spectra with a large wavelength coverage for a large number of solar-type stars. The more advanced technology in the visible also includes well-established methods for high-precision calibration. The combination of currently obtainable data quality and calibration together make it possible to reach radial velocity precisions of $\\sim$\\,1\\,m\\,s$^{-1}$ for many solar-type stars, which is sufficient to detect planets down to a few Earth masses in short-period orbits around them \\citep[e.g.][]{howard09,bouchy09}.  Despite the obvious utility of radial velocity measurements in the visible, measurements at other wavelengths are also desirable. The near-infrared (NIR, i.e. 1 -- 5\\,$\\mu$m) in particular is an interesting wavelength region for radial velocity exoplanet studies. In contrast to solar-type stars, low-mass objects that are cooler than $\\sim$\\,4000\\,K are brightest at wavelengths of 1\\,$\\mu$m and longer. Furthermore, stars cooler than $\\sim$\\,3200\\,K are so faint at visible wavelengths that high-precision radial velocity measurements are impossible for all but a few very nearby examples even with the best instruments on 8-10\\,m class telescopes. For example, the sample of 40 stars in a long-term radial velocity planet search program specifically targeting low-mass stars that utilized UVES \\citep{dekker00} at the VLT only included two objects with masses below 0.2\\,M$_{\\sun}$, and only three with masses below 0.35\\,M$_{\\sun}$ \\citep{zechmeister09}.  Very low-mass stars are an interesting sample of potential planet hosts notwithstanding their neglected status. They are useful for probing the correlation between gas giant planet frequency and stellar mass \\citep{endl06, johnson07}, are the most numerous stars in the Galaxy \\citep{reid02, henry06, covey08}, and exhibit a larger dynamical response to orbiting planets and have closer-in habitable zones than higher-mass stars \\citep{kasting93}. The latter point means that much less precision is needed to detect potentially habitable planets around low-mass stars with the radial velocity method than is needed for solar-type stars. The ubiquity of mid to late M dwarfs and the closeness of their habitable zones suggests that the closest stars hosting transiting habitable planets are likely of this stellar type \\citep{deming09}. Furthermore, \\citet{deming09} have estimated that basic atmospheric characterization of a few such planets could be obtained with the \\textit{James Webb Space Telescope} using the techniques of transit and occultation spectroscopy. Therefore, very low-mass stars are an important sample for the possible future study of the atmosphere of a habitable planet. However, we first must be able to use the radial velocity method to identify or confirm (in the case the planet is first identified by a transit detection), and measure the masses of low-mass planets around very low-mass stars before these investigations can be carried out.  As faint as the lowest-mass stars are at visible wavelengths, their extreme redness means they are reasonably bright at NIR wavelengths. For example, an M6 star (M$_{\\star}$ $\\approx$ 0.1\\,M$_{\\sun}$) at 10\\,pc will have a $V$-band magnitude of $\\sim$\\,15.5, which would make it a very challenging target for high-precision radial velocity measurements in this wavelength region with current telescopes and instruments. However, the same star will have a $K$-band magnitude of $\\sim$\\,8.5, which is bright enough that spectra with S/N suitable for high-precision radial velocity measurements could be easily obtained using existing high-resolution NIR spectrographs on 8-10\\,m class telescopes - especially those equipped with adaptive optics (AO) systems. Therefore, if the issues and advantages of low-mass stars described above are to be studied and exploited using the radial velocity method, then the NIR spectral region offers the only possible option without building larger telescopes.  Radial velocities measured from NIR spectra also likely offer the advantage that activity induced ``jitter'' is reduced relative to measurements in the visible. Magnetic activity plays a role in limiting the effectiveness of radial velocity planet searches because it results in inhomogeneities on the stellar surfaces (spots and plage). Visible spectra are thus modulated by stellar rotation and variations in the activity. These intrinsic variations mimic real radial velocity changes and are very confusing for planet searches \\citep[e.g.][]{queloz01,paulson04,wright05}. It is thought that radial velocities measured from NIR spectra will exhibit reduced jitter relative to the visible because of the smaller contrast between the cool inhomogeneities caused by activity and the average stellar surface at the longer wavelengths. Significant evidence exists to support the existence of this effect, and this evidence is discussed in the following sub-section.  The issue of activity induced radial velocity jitter is particularly relevant for low-mass stars because it is well established that a much higher fraction of late-type stars are active than earlier-type stars. In the most comprehensive survey so far, \\citet{west04} found that the fraction of active stars increased from $<$ 10\\% at spectral type M3, to 50\\% at M5, and ultimately 75\\% at the end of the main sequence. As high as these activity fractions are, they are likely only lower limits for nearby stars because the \\citet{west04} survey included stars orbiting perpendicular to the Galactic plane, and such stars are potentially very old. Therefore, the fraction of active stars in a radial velocity planet search of M dwarfs will likely be even higher than the \\citet{west04} statistics suggest because they will be drawn from an intrinsically younger sample. This has serious implications for planet searches because it means that the majority of the lowest-mass stars will exhibit substantial activity induced jitter in radial velocities measured from visible wavelength spectra even if their faintness can be overcome.  \\subsection{Previous work in the area} The advantages offered by NIR radial velocities have been identified before and there has been some previous work to realize them. \\citet{martin06} used NIRSPEC \\citep{mclean98} on Keck to make NIR radial velocity measurements of the brown dwarf LP\\,944-20. Their measurements over six nights had an rms dispersion of 360\\,m\\,s$^{-1}$. This result was inconsistent with previously measured visible wavelength radial velocities, which exhibited a coherent periodic signal with an amplitude of 3.5\\,km\\,s$^{-1}$. This result is generally interpreted as confirmation that radial velocity jitter is indeed reduced at NIR wavelengths for active cool dwarfs. However, it should be noted that the NIR and visible data sets were not contemporaneous, and the possibility that the star was experiencing a period of reduced activity during the NIR measurements can not be ruled out.  \\citet{zapatero09} have also reported radial velocity measurements obtained from NIRSPEC data, in this case for the very low-mass star VB\\,10. Their data were obtained over seven years and have an estimated precision of $\\sim$\\,300\\,m\\,s$^{-1}$. These measurements seem to support the astrometric detection of a giant planet around VB\\,10 claimed by \\citet{pravdo09}. However, our measurements of VB\\,10, which are presented in paper II and exhibit a dispersion of only 10\\,m\\,s$^{-1}$, are inconsistent with the \\citet{zapatero09} measurements and also appear to completely rule out the existence of the proposed giant planet around this star \\citep{bean10}. In addition, \\citet{anglada10} did not detect the expected reflex motion of VB\\,10 in their visible wavelength radial velocities, which have a typical precision $\\sim$\\,200\\,m\\,s$^{-1}$.  \\citet{blake07} reported NIR radial velocity measurements with PHOENIX \\citep{hinkle03} on Gemini South. They targeted L dwarfs and reached a precision of 300\\,m\\,s$^{-1}$ over five nights. \\citet{prato08} observed a sample of young stars and some late-type radial velocity standard stars with CSHELL \\citep{greene93} on the IRTF for seven nights. They reached precisions of $\\sim$\\,100\\,m\\,s$^{-1}$ for the standard stars, and 100 -- 300\\,m\\,s$^{-1}$ for the young stars. The measurements for the young stars again indicated no coherence with significant periodicities seen in visible radial velocities, but also were not contemporaneous. Similarly, \\citet{huelamo08} monitored the young star TW Hya over six nights using CRIRES \\citep{kaeufl04} on the VLT. Their measurements yielded an rms dispersion of 35\\,m\\,s$^{-1}$, which was inconsistent with a coherent signal having an amplitude of $\\sim$\\,300\\,m\\,s$^{-1}$ seen in visible radial velocities obtained previously \\citep{setiawan08} and contemporaneously. \\citet{seifahrt08} also used CRIRES to obtain continuous observations of an M giant over five hours and reached a precision of $\\sim$\\,20\\,m\\,s$^{-1}$.  In addition to these stellar observations, there has been some work on measuring NIR radial velocities for the Sun. \\citet{deming87} and \\citet{deming94} presented data from a long-term monitoring program using a Fourier Transform Spectrometer (FTS). Their $K$-band measurements had a precision of better than 5\\,m\\,s$^{-1}$. More recently, \\citet{ramsey08} used a laboratory grating spectrograph to measure the signature of the Earth's rotational velocity in $Y$-band spectra of integrated Sunlight. They maintained a precision of $\\sim$\\,10\\,m\\,s$^{-1}$ over a few hours in two separate experiments.  \\subsection{Challenges of NIR radial velocities} Despite the attention paid to obtaining NIR radial velocities of cool stars, no previous work has achieved a long-term precision on a star other than the Sun within an order of magnitude of the precision that is routinely obtained in the visible. One reason for this is simply that enough signal was not obtained to probe the true limit in some cases. However, based on photon statistics, the observations of \\citet{prato08} and \\citet{huelamo08} should have at least achieved precisions around a factor of two better than they did. Another issue is that no studies with an external check on the obtained precision have been reported continuing for longer than seven days. Therefore, the obtainable long-term precision in NIR radial velocities has not been systematically developed or evaluated.  The main issue that has limited the pursuit of high-precision NIR radial velocities is the lack of a suitable calibration method. For example, useful lines from available emission lamps \\citep[e.g. ThAR,][]{lovis07} are infrequent in the NIR relative to the visible. This is a problem for the existing high-resolution NIR spectrographs because they generally have relatively short wavelength coverage for single exposures. Laser frequency combs have the potential to provide calibration lines in nearly any spectral region conceivable for stellar radial velocity work \\citep{murphy07,li08,steinmetz08}. However, the technique is still not developed enough for long-term use at an astronomical observatory. Furthermore, none of the existing NIR spectrographs can be fed simultaneously with light from a calibration source like a lamp or laser comb system, and are not stabilized to the level necessary to do without simultaneous calibration. In addition, there is also the issue that current NIR spectrographs experience significant variations in illumination on short timescales. This means that even if these instruments could be simultaneously fed with light from a lamp, the calibration would not track all the necessary effects for high-precision radial velocity measurements. Therefore, emission lamps and laser combs are not useful for high-precision calibration in currently available NIR spectrographs, but could be for future instruments designed for larger wavelength coverage and excellent stability \\citep{reiners10}.  All of the previous studies discussed above that targeted stars other than the Sun utilized the telluric spectrum imprinted on the data for \\textit{in situ} calibration, although the lines and methods utilized varied somewhat. In their study, \\citet{seifahrt08} demonstrated that telluric N$_{2}$O lines near 4.1\\,$\\mu$m were stable to a level of 10\\,m\\,s$^{-1}$ over five hours. On the other hand, \\citet{deming87} showed that the telluric methane lines in the $K$-band varied with a semi-amplitude of 20\\,m\\,s$^{-1}$ depending on the hour angle due to the prevailing winds over the observatory. In addition to the possible variability of the telluric lines, another issue with using telluric lines as a radial velocity fiducial is that measurements obtained over a full observing season could be affected by systematics because the stellar lines will move relative to the telluric lines due to the changing barycentric velocity of the observatory along the line of sight. This will lead to more or less blending of features, which could influence the measurements.  The well known ``iodine cell'' method \\citep{butler96} has been used to measure radial velocities to precisions of a few m\\,s$^{-1}$ with numerous visible spectrographs that are not highly stabilized. The general gas cell technique involves placing a cell in front of the spectrograph during observations so that a standard spectrum is imprinted on each obtained stellar spectrum. The cell's spectrum recorded during each exposure can then be used to precisely determine the spectrograph response at the exact moment of the observations and, thus, calibrate the simultaneously obtained stellar spectra for radial velocity measurements.  The success of the iodine cell in the visible implies that a gas cell could be a useful way to calibrate currently available NIR spectrographs. The main constraint for the gas cell method is that the gas in question should exhibit numerous sharp and deep lines in a wavelength interval where the stars targeted for radial velocity measurements have lines as well \\citep{campbell79}. Ideally, this region would also be free of atmospheric lines. However, this is a particularly challenging requirement in the NIR due to the prevalence of strong atmospheric lines even in the traditional windows between water absorption bands.  \\citet{deming87} and \\citet{deming94} used a cell filled with N$_{2}$O to calibrate their FTS measurements of integrated sunlight from 2.0 -- 2.5\\,$\\mu$m. In addition, \\citet{mahadevan09} considered a number of options for NIR gas cells and concluded that the gases H$^{13}$C$^{14}$N, $^{12}$C$_{2}$H$_{2}$, $^{12}$CO, and $^{13}$CO together could provide useful calibration in the $H$-band. However, there have been no reported uses of a gas cell for high-precision NIR radial velocity measurements for stars other than the Sun, and no such measurements obtained using a grating spectrograph.  We have recently developed a new gas cell for simultaneous calibration of spectra obtained in the $K$-band, and are currently using it to conduct a radial velocity search for planets around very low-mass stars with the CRIRES instrument at the VLT. In this paper, which is the first in a planned series about our planet search, we describe a method for measuring high-precision relative radial velocities of cool stars using observations made with the cell, and we present extensive data sets demonstrating the performance of the method over long timescales. The paper is laid out as follows. In \\S2 we describe the new cell and its implementation in CRIRES. We describe the typical observing procedure and data reduction used in \\S3. We detail the radial velocity measurement algorithm in \\S4. In \\S5, we present tests of the obtained radial velocity precision. We conclude in \\S6 with a discussion and outlook for future planet searches utilizing NIR radial velocities. ",
        "conclusions": "The general technique of measuring radial velocities from NIR spectra has recently been receiving a lot of attention as a possible important direction for the observational study of exoplanets for the reasons described in \\S1.1 \\citep[e.g.][]{lunine08}. However, up to now the feasibility of obtaining NIR radial velocity precision similar to that routinely obtained in the visible was largely unknown due to uncertainty about a number of technical issues. The main open question has been how to calibrate NIR spectra, with a second issue being whether the quality of NIR detectors is sufficient.  We have described a new gas cell suitable for calibrating spectra of cool dwarfs in the $K$-band, and demonstrated long-term radial velocity precisions of $\\sim$\\,5\\,m\\,s$^{-1}$ from data obtained with CRIRES at the VLT when using the cell. It is worth noting that the CRIRES detectors have severe cosmetic blemishes relative to CCDs. However, the effects seem to be stable and we have been able to calibrate them out to an acceptable level. Furthermore, the detectors were not even the highest-quality NIR detectors that were available at the time CRIRES was built \\citep[ca. 2002,][]{dorn04}, and are certainly not competitive with the best detectors available as of this writing. Therefore, the issue of the quality of NIR detectors requires attention, but will not likely be a limiting factor for new instruments aiming at $\\sim$\\,1\\,m\\,s$^{-1}$ NIR radial velocity precision.  Our ammonia gas cell has proven to be useful for NIR calibration down to a few m\\,s$^{-1}$ precision despite not being temperature stabilized. Therefore, it could be useful for high-precision radial velocity measurements with other instruments not originally designed for excellent stability. The precision afforded by the cell and the measurement method we have described is already an order of magnitude better than the best previously obtained long-term NIR precision on a star other than the Sun. As a result, it now enables the possibility of searching for planets around a much larger number of very low-mass stars than was feasible before. For example, our planet search project includes 31 objects with estimated masses below 0.2\\,M$_{\\odot}$, 22 of which have estimated masses below 0.15\\,M$_{\\odot}$. A previous high-precision radial velocity planet search utilizing a visible wavelength spectrograph on an identical telescope (UVES on the UT2 telescope of the VLT) was only able to target two objects with estimated masses below 0.2\\,M$_{\\odot}$, and only one of these two has an estimated mass below 0.15\\,M$_{\\odot}$ \\citep{zechmeister09}.  In addition to enabling the search for planets around more low-mass stars, the gas cell and radial velocity measurement algorithm we have developed also opens up a new frontier on the search for potentially habitable planets. The orbital period range for a planet in the habitable zone around a star with a mass M = 0.10\\,M$_{\\odot}$ would be 3 -- 21 days \\citep{selsis07}. The 5\\,m\\,s$^{-1}$ precision obtainable with our method corresponds to the velocity semi-amplitudes induced by a 2.5\\,M$_{\\oplus}$ planet or a 4.6\\,M$_{\\oplus}$ planet on the inner and outer edges of the habitable zone respectively. Therefore, it should be possible to detect Super-Earth type planets in the habitable zones of very low-mass stars with a reasonable expenditure of observing time on current facilities.  Altogether our results have shown that obtaining high-precision NIR radial velocities is possible, and we see no technical reason why a level of precision of 1\\,m\\,s$^{-1}$ could not be achievable with a highly stabilized NIR grating spectrograph. However, the design of such a next generation instrument will have to carefully weigh the availability of calibration against the availability of information (i.e. spectral lines) from the stars being targeted like we have done with our observing program.  As discussed above, we are carrying out a planet search using CRIRES with the ammonia cell and this paper is the first in a planned series from the project. The next papers in the series will present high-precision NIR radial velocities of our M dwarf planet search targets. The data will allow us to detect planets around these stars with similar masses and orbital separations as those that can be detected around solar-type stars using state-of-the-art visible wavelength radial velocities."
    },
    "0911/0911.4768_arXiv.txt": {
        "abstract": "We study the halo bispectrum from non-Gaussian initial conditions. Based on a set of large $N$-body simulations starting from initial density fields with local type non-Gaussianity, we find that the halo bispectrum exhibits a strong dependence on the shape and scale of Fourier space triangles near squeezed configurations at large scales. The amplitude of the halo bispectrum roughly scales as $\\fnl^2$. The resultant scaling on the triangular shape is consistent with that predicted by Jeong \\& Komatsu based on perturbation theory. We systematically investigate this dependence with varying redshifts and halo mass thresholds. It is shown that the $\\fnl$ dependence of the halo bispectrum is stronger for more massive haloes at higher redshifts. This feature can be a useful discriminator of inflation scenarios in future deep and wide galaxy redshift surveys. ",
        "introduction": "The standard cosmological model has successfully explained the observed statistical properties of the cosmic microwave background (CMB) radiation and the large scale structure (LSS) traced by galaxies (e.g., \\cite{WMAP5,Tegmark2006}). The model usually assumes that the primordial density/temperature/curvature fluctuations follow Gaussian statistics. Recently, however, possible deviations from standard Gaussian statistics has attracted great attention with rapid progress in observational techniques. It offers an opportunity to access cosmological information beyond traditional power spectrum analysis. Many recent works have discussed the constraints and future detectability of possible deviations from Gaussianity through the observations of CMB and LSS (e.g., \\cite{Komatsu2001,Carbone2008}).  According to the inflationary scenarios, primordial curvature perturbations are generated during the accelerated phase of cosmic expansion.  The simplest inflation models, in which the inflation takes place with the slow-roll single scalar field that has a canonical kinetic structure, generally predicts a nearly scale-invariant spectrum of curvature perturbations, and a small departure from Gaussianity.  On the other hand, a variety of inflation models that produce a large non-Gaussianity has been recently proposed (see \\cite{Bartolo2004} for a review).  Among these, the models with large non-Gaussianity generated by non-linear dynamics on super-horizon scales can have a generic prediction for non-Gaussian properties.  Denoting a Gaussian field by $\\Phi_{\\rm G}(\\bfx)$, the Bardeen's curvature perturbation during the matter era is characterized as \\cite{Komatsu2001}: \\begin{eqnarray} \\Phi(\\bfx) = \\Phi_{\\rm G}(\\bfx)+\\fnl\\left\\{ \\Phi_{\\rm G}^2(\\bfx)-\\langle\\Phi_{\\rm G}^2(\\bfx)\\rangle\\right\\} +\\cdots. \\label{eq:fnllocal} \\end{eqnarray} This type of non-Gaussianity, described as a local function of the Gaussian field, is often called {\\it local type} non-Gaussianity, and the leading-order coefficient $\\fnl$, which controls the strength of non-Gaussianity, has information on the generation mechanisms for non-Gaussian fluctuations.  Although the current constraint on the parameter $\\fnl$ from CMB data is $-9<\\fnl<111$ ($95\\%$C.L.) \\cite{WMAP5} and it is still consistent with Gaussian ($\\fnl=0$), the constraint will be tightened by the on-going CMB experiment, Planck \\cite{Planck}. As standard inflation models generally predict $|\\fnl|\\ll1$, detection of large non-Gaussianity immediately implies the non-standard mechanism for generation of primordial curvature perturbations.   In this paper, we focus on how this type of non-Gaussianity alters the statistical properties of LSS.  Let us first define the power spectrum of mass density fluctuations assuming isotropy and homogeneity: \\begin{eqnarray} \\langle \\deltam(\\bfk_1)\\deltam(\\bfk_2)\\rangle \\equiv (2\\pi)^3\\delta_{\\rm D}(\\bfk_1+\\bfk_2)\\Pm(k_1), \\label{eq:pow_def} \\end{eqnarray} where $\\deltam(\\bfk)$ denotes the Fourier transform of the density contrast, while $\\delta_{\\rm D}(\\bfk)$ represents the Dirac delta function. If the density field follows Gaussian statistics, its power spectrum determines all the statistical properties. Next we define the bispectrum of a mass density field: \\begin{eqnarray} \\langle \\deltam(\\bfk_1)\\deltam(\\bfk_2)\\deltam(\\bfk_3)\\rangle \\equiv (2\\pi)^3\\delta_{\\rm D} (\\bfk_1+\\bfk_2+\\bfk_3)\\Bm(\\bfk_1,\\bfk_2,\\bfk_3). \\label{eq:bis_def} \\end{eqnarray} Since this is the lowest-order non-vanishing quantity in the presence of non-Gaussianity, the bispectrum is naively expected as the best measure for non-Gaussianity (e.g., \\cite{Scoccimarro2000,Verde2000,Scoccimarro2004,Sefusatti2007}).  Recently, however, the {\\it galaxy} or {\\it halo} power spectrum has been reconsidered in the presence of local and/or equilateral type primordial non-Gaussianity both analytically and numerically (e.g., \\cite{Dalal2008,Afshordi2008,Taruya2008,Matarrese2008,Pillepich2008,Desjacques2009a,Desjacques2009b,Grossi2009,Slosar2008,McDonald2008,Sefusatti2009,Giannantonio2009}). The matter power spectrum or bispectrum is not a direct observable, and the real measurement of LSS gives {\\it galaxy} power spectrum or bispectrum, as defined similarly to equations (\\ref{eq:pow_def}) and (\\ref{eq:bis_def}).  Since galaxies are biased tracers of the dark matter distribution, the information of $\\fnl$ is imprinted in a different manner: a new contribution coming from the primordial non-Gaussianity may dominate over the Gaussian term in the galaxy power spectrum, $\\Pg(k)$, at very large scales ($k\\simlt0.01h$Mpc$^{-1}$).  This new contribution, which is sometimes referred to as the ``scale-dependent bias'' of the galaxy power spectrum, may be a powerful indicator to constrain $\\fnl$.  Indeed, it has been recently applied to the clustering statistics of SDSS LRG and quasar samples, and the tight constraints on $\\fnl$ are comparable to those obtained from CMB measurements have been obtained \\cite{Slosar2008}.  The purpose of this paper is to examine the bispectrum of biased tracers in detail. While the matter bispectrum in the presence of primordial non-Gaussianity has been studied in the literature using both perturbation theory and numerical simulations, the galaxy bispectrum may significantly differ from the matter bispectrum in the presence of primordial non-Gaussianity, just like the difference in the power spectra.  Since the local type non-Gaussianity can be straightforwardly implemented within $N$-body simulations, numerical study on the bispectrum for the dark matter haloes is the first important step toward a practical understanding of the galaxy bispectrum.  Incidentally, Jeong and Komatsu (2009) recently proposed a new parametrized model for the halo/galaxy bispectrum (\\cite{Jeong2009}, see also \\cite{Sefusatti2009}) based on the peak bias model \\cite{Matarrese1986} and the local bias model \\cite{Fry1993}. They found that the formula for the galaxy bispectrum used in \\cite{Sefusatti2007} was missing important contributions from the scale dependent bias effects and they discovered new terms that are important at ``squeezed'' configurations where $k_1, k_2\\gg k_3$. It was argued that these new contributions enable us to put stronger constraints on $f_{\\rm NL}$ than those obtained in \\cite{Sefusatti2007}.  It is of great importance to confirm the scale-dependent bias effects in the bispectrum by $N$-body simulations.  The rest of this paper is organized as follows: we first review the analytical models of the power spectrum and the bispectrum of biased tracers in section \\ref{sec:theory}. We then describe the setup and initial conditions for $N$-body simulations in section \\ref{sec:simulation}. As a first check of our simulations, in section \\ref{sec:pow_results}, we compute the matter and halo power spectra, and the results are compared with previous works.  Section \\ref{sec:bis_results} gives the main results of this paper, in which the simulation results for the matter and halo bispectra are presented and compared with predictions from analytic models, particularly focusing on their scale dependence.  The dependence of the halo bispectrum on the halo mass threshold and redshift is also investigated in detail. Section \\ref{sec:discussion} discusses the future prospects for detecting the primordial non-Gaussianity using the scale-dependent properties of the halo/galaxy bispectrum. Finally, section \\ref{sec:conclusion} is devoted to conclusions and discussions. ",
        "conclusions": "\\label{sec:conclusion} In this paper, we have studied the clustering properties of dark matter haloes from cosmological $N$-body simulations in the presence of local-type primordial non-Gaussianity. We found that the halo bispectrum measured from $N$-body simulations exhibits a strong $f_{\\rm NL}$ dependence which becomes most prominent for squeezed configurations at large scales. In particular, for realistic values of $|\\fnl|\\simlt100$, the dependence of the halo bispectrum on $\\fnl$ is well characterized by the polynomial expansions of $\\fnl$ up to second order.  Since the quadratic dependence on $\\fnl$ does not appear in the matter bispectrum at the lowest order in perturbation theory, this would be a clear indicator for the existence of primordial non-Gaussianity of the local type.  We have investigated the shape and scale dependence of the halo bispectrum arising from the $\\fnl^2$ term, and the simulation results are compared with theoretical predictions based on the local bias model. For the isosceles triangles characterized by $\\alpha\\equiv k_1 / k_3$ and $k\\equiv k_1=k_2$, the dependence of the halo bispectrum on $\\alpha$ measured from $N$-body simulations is found to be consistent with theoretical predictions by \\cite{Jeong2009}.  We also examined the dependence of the halo bispectrum on minimum halo mass and redshift, and showed that the amplitude of the halo bispectrum is more significant for more massive haloes at higher redshifts.  Thus, the strong dependence of the halo/galaxy bispectrum on $\\alpha$ makes the detection of primordial non-Gaussianity much more promising in future surveys. As a preliminary investigation, we have evaluated the signal-to-noise ratio for the scale dependence of the bispectrum in three representative surveys, and found that even with the very limited number of configurations of bispectrum it is possible to detect primordial non-Gaussianity if $\\fnl$ is several dozen. Thus, the detectability of primordial non-Gaussianity is expected to be greatly improved if we use all configurations of the bispectrum.  We leave the following tasks as a future work: (i) Study the effects of redshift-space distortions. Since we focus on very large scales, we expect that these effects are accurately described by linear theory, i.e., they just enhance the amplitude of the bispectrum in a scale independent way. (ii) Construct more elaborate theoretical models that are applicable to a wider range of triangular configurations and compare them with $N$-body simulations (iii) Run higher resolution simulations where we can populate haloes with galaxies and measure the galaxy bispectrum directly from simulations. These tasks are clearly very important to exploit future surveys.   \\ack We thank T.~Sousbie, R.~Nichol, E.~Komatsu, D.~Jeong, and Y.~Suto for useful discussions and comments.  T.~N. is supported by a Grant-in-Aid for Japan Society for the Promotion of Science (JSPS) Fellows (DC1: 19-7066). A.~T. is supported by a Grant-in-Aid for Scientific Research from JSPS (No. 21740168).  K.~K. is supported by the European Research Council, Research Councils UK and the UK's Science \\& Technology Facilities Council (STFC).  C.~S. was funded by a STFC PhD studentship.  Numerical computations were in part carried out on Cray XT4 at Center for Computational Astrophysics, CfCA, of National Astronomical Observatory of Japan.  We are also grateful for the computational time provided by the U.K. National Grid Service (NGS).  This work was supported in part by Grant-in-Aid for Scientific Research on Priority Areas No.~467 ``Probing the Dark Energy through an Extremely Wide and Deep Survey with Subaru Telescope'', and JSPS Core-to-Core Program ``International Research Network for Dark Energy''. \\appendix"
    },
    "0911/0911.3654_arXiv.txt": {
        "abstract": "\\renewcommand{\\thefootnote}{\\fnsymbol{footnote}}   We present kinematic and metallicity profiles for the M\\,31 dwarf elliptical (dE) satellite galaxies NGC~147 and NGC~185.  The profiles represent the most extensive spectroscopic radial coverage for any dE galaxy, extending to a projected distance of eight half-light radii ($8r_{\\rm eff} \\sim 14'$).  We achieve this coverage via Keck/DEIMOS multislit spectroscopic observations of 520 and 442~member red giant branch stars in NGC~147 and NGC~185, respectively.  In contrast to previous studies, we find that both dEs have significant internal rotation.  We measure a maximum rotational velocity of $17\\pm 2$\\kms\\ for NGC~147 and $15\\pm5$\\kms\\ for NGC~185.  While both rotation profiles suggest a flattening in the outer regions, there is no indication that we have reached the radius of maximum rotation velocity.  The velocity dispersions decrease gently with radius with an average dispersion of $16\\pm 1$\\kms\\ for NGC~147 and $24\\pm1$\\kms\\ for NGC~185.  The average metallicity for NGC~147 is [Fe/H] = $-1.1\\pm0.1$ and for NGC~185 is [Fe/H] = $-1.3\\pm0.1$; both dEs have internal metallicity dispersions of 0.5\\,dex, but show no evidence for a radial metallicity gradient.  We construct two-integral axisymmetric dynamical models and find that the observed kinematical profiles cannot be explained without modest amounts of non-baryonic dark matter.  We measure central mass-to-light ratios of $M/L_V = 4.2\\pm0.6$ and $M/L_V = 4.6\\pm0.6$ for NGC~147 and NGC~185, respectively.   Both dE galaxies are consistent with being primarily flattened by their rotational motions, although some anisotropic velocity dispersion is needed to fully explain their observed shapes.  The velocity profiles of all three Local Group dEs (NGC~147, NGC~185 and NGC~205) suggest that rotation is more prevalent in the dE galaxy class than previously assumed, but is often manifest only at several times the effective radius.  Since all dEs outside the Local Group have been probed to only inside the effective radius, this opens the door for formation mechanisms in which dEs are transformed or stripped versions of gas-rich rotating progenitor galaxies. ",
        "introduction": "\\renewcommand{\\thefootnote}{\\fnsymbol{footnote}}   Dwarf elliptical (dE) galaxies are characterized by low surface brightness ($\\mu_{\\rm eff,V} > 22$ mag arcsec$^{-2}$), little to no gas, and old to intermediate age stellar populations.  These galaxies are spatially clustered and account for more than 75\\% of objects brighter than $M_V < -14$ in nearby galaxy clusters \\citep{ferguson91a,ferguson94a}.  Since there are few, if any, isolated dE galaxies, most proposed dE galaxy formation scenarios are environmentally-driven, such as via gas stripping or gravitational harassment of gas-rich spiral or dwarf irregular galaxies \\citep[e.g.\\,][]{dekel86a, moore98a,Aguerri09a}.  Recent evidence for embedded stellar disks and bars in some dEs support these scenarios \\citep{jerjen00a,geha05a,lisker06a,chilingarian07a}. Since gas-rich systems have significant rotation velocities, a key test of this hypothesis class is comparing the kinematics of dE and gas-rich galaxies.  Internal kinematic measurements for dEs in the Fornax and Virgo Cluster have revealed an unexpected dichotomy: roughly half of the observed dEs have significant rotation about the major-axis, while the remaining galaxies show no detectable major-axis rotation \\citep{pedraz02a,geha03a, vanzee04a, deRijcke05a}.  dE galaxies with no rotation present a challenge to any scenario in which dE progenitors are gas-rich and thus rotating \\citep{moore98a}.  Ram pressure stripping alone does not affect the kinematics of the collisionless stellar component.  Gravitational processes such as galaxy harassment can reduce the amount of rotational angular momentum, converting this into random velocity dispersion, however, simulations are unable to completely remove rotational support \\citep{mayer01a,mastropietro05a}.  Because dEs have low surface brightnesses, integrated light (long-slit) kinematic observations of dEs in both galaxy clusters and the Local Group are currently limited to within the half-light effective radius ($r_{\\rm eff}$).  Thus, it is possible that significant rotational angular momentum lies at larger distances than has so far been explored.  Addressing this question requires kinematics tracers at larger radius.  While the dynamics of dE globular cluster systems suggest that there may be significant rotation at large radius \\citep{beasley06a,beasley09a}, the small number of clusters limits the usefulness of this tracer. The dEs in the Local Group share the same general properties as the more numerous dEs in galaxy clusters, however, their proximity allows us to resolve individual stars \\citep{geha06a}, and thus trace kinematics profiles down to arbitrarily low surface brightness and large radius.  The Local Group contains three dE galaxies, NGC~205, NGC~147 and NGC~185, which in addition to the compact elliptical M32, are the brightest satellites around the spiral galaxy M\\,31.   We follow  the terminology conventions of,  e.g.~\\citet{bender92a}, although alternatives are in use. While the differing terminologies can be confusing, the key point is that NGC~205, NGC~147, and NGC~185 on the one hand, and M\\,32 on the other hand, should not be grouped in a single class. The former three galaxies have fundamental plane properties that indicate similarity to galaxies in the dwarf spheroidal class (dSph). By contrast, galaxies such as M\\,32 have properties that indicate similarity to (giant) elliptical galaxies. This likely indicates a difference in formation history, as discussed more fully in, e.g., \\citet{kormendy09a}.  The prototype of the dE galaxy class is NGC~205 which is in very close projection to M\\,31 ($40' = 8$\\,kpc).  Both photometric and kinematic evidence suggests that NGC~205 is tidally interacting with M\\,31 \\citep{choi02a,geha06a}, and may in fact be on its first orbital approach \\citep{howley08a}.  In contrast, NGC~147 and NGC~185 lie at a projected distance of $7^{\\circ}$\\,($\\sim$\\,$150$\\,kpc) from M\\,31.   While the metal-poor stellar halo of M\\,31 extends to this distance \\citep{guhathakurta05a,kalirai06a},  its gravitational influence of M\\,31 is not expected to influence the kinematics of these dE satellites (see \\S\\,\\ref{subsec_bound}).    NGC~147 and NGC~185 are separated from each other by merely $58'$, leading to speculation that they may be a bound galaxy pair \\citep{baade44a, vandenbergh98a}.   An accurate distance estimate suggests that these two galaxies are physically separated by over 60\\,kpc \\citep{mcconnachie05a} and may not be bound (however, see \\S\\,\\ref{subsec_bound}).  The large observed distances between these two dEs and their parent M\\,31, combined with their relative proximity to the Sun, make NGC~147 and NGC~185 the best available targets for studying the properties of dE galaxies to large radius.   NGC~147 and NGC~185 share several similar fundamental properties, such as absolute luminosity ($M_V\\sim -15.5$), half-light radius ($r_{\\rm eff} \\sim 2' = 0.3$\\,kpc), and central velocity dispersion (for exact values see Table~1).  However, they differ markedly in many aspects.  NGC~185 contains some gas, dust and evidence for recent star formation confined to its center \\citep{martinez-delgado99a}, while NGC~147 is devoid of gas or dust and shows no sign of recent star formation activity \\citep{young97a, sage98a}.  An average metallicity of [Fe/H] = $-1.4$ is inferred for NGC~185 \\citep{martinez-delgado99a}, while a more metal rich population of [Fe/H] = $-1.0$ is inferred for NGC~147 \\citep{han97a, davidge05a, goncalves07a}.  HST/WFPC2 imaging of both galaxies also implies the presence of intermediate-age stars, with NGC~147 having a more significant contribution than NGC~185 \\citep{butler05a}.  Thus, while both dEs are dominated by old to intermediate-age stars, the metallicity and age mixture of these components are different in each galaxy.   The first spectroscopic long-slit observations of Local Group dEs by \\citet{bender91a} suggested that while NGC~185 appeared supported entirely by random motions, NGC~147 had a significant rotational component of $6.5\\pm1.1$\\kms.  Somewhat deeper observations by \\citet{simien02a} extended the kinematic profiles out $\\sim2'$ (1$r_{\\rm eff}$), but did not measure a significant rotation for either galaxy.  In contrast, both studies measured a rotational component in NGC~205 (within the tidal radius), later confirmed by \\citet{geha06a} to be $11\\pm5$\\kms.  Using the Simien \\& Prugniel data, \\citet{derijcke06a} constructed dynamical models for all three Local Group dEs, concluding that, within the half-light radii, all three galaxies have mass-to-light ratios above that expected from the stellar populations alone.   In this paper, we present accurate radial velocities out to 8\\,$r_{\\rm eff}$ in the Local Group dE galaxies NGC~147 and NGC~185, based on Keck/DEIMOS multi-slit spectroscopy for 520 and 442~red giant branch (RGB) stars, respectively.  The paper is organized as follows: in \\S\\,\\ref{sec_data} we discuss target selection for our DEIMOS slitmasks, the observing procedure and data reduction.  In \\S\\,\\ref{sec_results}, we discuss the velocity, velocity dispersion and metallicity profiles.  In \\S\\,\\ref{sec_model}, we construct two-integral dynamical models for these galaxies and discuss the best fit mass distribution and dark matter content.  Finally, in \\S\\,\\ref{sec_disc}, we discuss the implications of these results for the formation of dE galaxies.  Throughout this paper we adopt distance moduli determined by \\citet{mcconnachie05a} via the tip of the RGB method.  The distance modulus for NGC~147 is $(m -M)_0 = 24.43 \\pm 0.04$ ($675\\pm 27$\\,kpc) and NGC~185 is $(m -M)_0 = 24.23 \\pm 0.03$ ($616\\pm 26$\\,kpc).  This places NGC~147 and NGC~185 at projected distances of 140 and 185\\,kpc from their parent galaxy M\\,31, respectively.    \\begin{figure*}[t!] \\plotone{dEfig_dss.eps} \\caption{Palomar Sky Survey images of the two Local Group dE galaxies, NGC~147 ({\\it left}) and NGC~185 ({\\it right}). Images are $30'\\times 30'$; North is up, East is to the left. The major-axis of each galaxy is indicated by the solid line.  The spatial position of Keck/DEIMOS spectroscopically confirmed member stars are shown as blue circles, non-members are plotted as open squares.  \\label{fig_dss}} \\end{figure*} ",
        "conclusions": "\\label{sec_disc}  We have presented mean velocity, velocity dispersion and metallicity profiles for the Local Group dE galaxies NGC~147 and NGC~185 based on Keck/DEIMOS spectroscopic observations of 520 and 442~member RGB stars in each dE, respectively.  The profiles represent the most extensive spectroscopic radial coverage for any dE galaxy, extending to a projected distance of eight half-light radii ($14' \\sim 8r_{\\rm eff}$).  Contrary to previous results, we find that both dEs have significant rotation velocities which contributes to the observed flattened shapes of each galaxy.  Our two-integral dynamical modeling suggests that the observed rotation velocities cannot fully explained the observed shapes and that some anisotropic velocity dispersion is required.  Our modeling estimates $M/L$ ratios of $M/L_V = 4.2\\pm0.6$ and $M/L_V = 4.6\\pm0.6$ for NGC~147 and NGC~185, respectively, which is in excess of that expected from stellar populations alone.  Thus, some dark matter is required to fully explain the observe kinematics.  The mean velocity profiles of NGC~147 and NGC~185 suggest that rotation may be far more prevalent in dE galaxies than previously assumed.  We conclude that these Local Group dEs are rotationally-supported, however, if placed at the distance of the Virgo or Fornax Clusters, integrated-light observations such as those of \\citet{geha03a} would have concluded that these dEs have little to no rotation.    It is not guaranteed that the Local Group dE galaxies  had similar formation mechanisms as those residing in the much denser environments of, e.g., the Virgo or Fornax clusters. However, this is certainly plausible, given that these galaxies share similar fundamental plane properties, as defined by the photometric and kinematic properties inside $r_{\\rm eff}$ (see e.g., Fig 7).  Our results then indicate that cluster dEs may also have significant rotation velocities that are manifest only at larger radii than current observations allow.  If true, this significantly modifies the observational constraints under which dEs can form.    The observations presented here open the door for formation mechanisms in which dEs are transformed or stripped versions of gas-rich rotating progenitor galaxies."
    },
    "0911/0911.2652_arXiv.txt": {
        "abstract": "An artificial neural network (ANN) is investigated as a tool for estimating rate coefficients for the collisional excitation of molecules.  The performance of such a tool can be evaluated by testing it on a dataset of collisionally-induced transitions for which rate coefficients are already known: the network is trained on a subset of that dataset and tested on the remainder.  Results obtained by this method are typically accurate to within a factor $\\sim 2.1$ (median value) for transitions with low excitation rates and $\\sim 1.7$ for those with medium or high excitation rates, although 4$\\%$ of the ANN outputs are discrepant by a factor of 10 more.  The results suggest that ANNs will be valuable in extrapolating a dataset of collisional rate coefficients to include high-lying transitions that have not yet been calculated.  For the asymmetric top molecules considered in this paper, the favored architecture is a cascade-correlation network that creates 16 hidden neurons during the course of training, with 3 input neurons to characterize the nature of the transition and one output neuron to provide the logarithm of the rate coefficient. ",
        "introduction": "The interpretation of astrophysical spectra rests heavily upon the availability of relevant atomic and molecular data.  Of critical importance are (1) the energies of the various accessible quantum states, (2) the rates of spontaneous radiative decay between states, and (3) the state-to-state rate coefficients for excitation and deexcitation in inelastic collisions.  The last of these have proven quite elusive; in many circumstances, our ability to interpret astrophysical spectra is severely limited by the availability of collisional rate coefficients.  As a result of new spectroscopic facilities such as the {\\it Spitzer Space Telescope}, the {\\it Herschel Space Observatory}, and the planned Atacama Large Millimetric Array (ALMA), the need for fundamental molecular data has become all the more acute.  In modeling the emission from interstellar molecular clouds, in particular, there is a critical need for rate coefficients for the rotational excitation of molecules in collisions with H$_2$.  The theoretical calculation of such data is extremely time consuming, involving the computation of a potential energy surface for the system in question, followed by extensive scattering calculations.  Even for those molecules for which such calculations have been carried out, the available data are often limited to collisionally-excited collisions involving a relatively small set of low-lying states.  Submillimeter and infrared spectra of molecular clouds often reveal the presence of higher-lying transitions for which collisional rate coefficients have not been computed; here, some kind of extrapolation method is needed.  For the case of linear molecules, the infinite order sudden (IOS) and related approximations provide an obvious method of extrapolation; for asymmetric molecules, there is no simple analogous method.  Given a limited set of transitions of an asymmetric molecule for which collisional rate coefficients have been computed, previous efforts at extrapolation (e.g.\\ Neufeld \\& Melnick 1987, Faure \\& Josselin 2008) have typically adopted an expression with one or more adjustable parameters for how the collisional rate coefficient depends upon the relevant parameters for each transition (energy difference, change in rotational quantum number, etc.)  The adjustable parameters are then chosen so as to optimize the fit to those transitions for which collisional rate coefficients are known, and then the expression is used to estimate the values for higher-lying transitions where calculated results are unavailable.  For example, the recent study by Faure \\& Josselin (2008) considered the rate coefficients for the excitation of H$_2$O by H$_2$, which will be of crucial importance to the interpretation of many {\\it Spitzer} and {\\it Herschel} spectra.  Here, Faure \\& Josselin extrapolated a set of rate coefficients obtained from quasi-classical trajectory calculations for transitions among the lowest 45 rotational states of water. Each such state is characterized by an energy, $E$, and rotational quantum numbers $J$, $K_a$ and $K_c$ that describe the total angular momentum and its projection onto the principal axes of the molecule. For transitions with $\\Delta K_a \\le 2$, $\\Delta K_c \\le 2$ and $\\Delta J \\le 2$ -- termed ``high-propensity\" transitions -- Faure \\& Josselin assumed that rate coefficients depend solely upon $\\Delta K_a$, $\\Delta K_c$ and $\\Delta J$.  For all other transitions (``low propensity transitions\"), they assumed the rate coefficients to be {\\it independent} of $\\Delta K_a$, $\\Delta K_c$ and $\\Delta J$ and proportional to $\\exp (-\\Theta \\Delta E/kT)$, where $\\Theta$ is adjusted to optimize the fit to the available data.  As an alternative to the methods adopted previously, I have explored the use of an artificial neural network (ANN) to accomplish the extrapolation.  In contrast to previous techniques, the use of an ANN does not make any {\\it a priori} assumption about how the transitions might be categorized or how the rate coefficients might depend upon the parameters for each transition. The calculational method is discussed in \\S2 (with the more technical details given in the Appendix).  The performance of the ANN in reproducing a known set of collisional rate coefficients in presented in \\S3.  A discussion follows in \\S4. ",
        "conclusions": "1.  Artificial neural networks (ANN) can be successfully trained to determine rate coefficients for the collisional excitation of astrophysical molecules.  2.  The type of ANN favored here is a cascade correlation network which creates 16 hidden neurons during the course of training.  There are three input neurons that inform the network of $\\Delta E$, $\\Delta J$ and $\\Delta \\tau$ for a given transition, and one output neuron that provides the logarithm of the rate coefficient for collisional de-excitation.  3.  The performance of such an ANN can be evaluated by training it on a subset of the set of transitions for which collisional rate coefficients are known, and then testing it on the remainder of that set.  Since the motivation for this study is to develop a method for estimating the rate coefficients for high-lying transitions for which values have not been computed, the training set involved the lowest-lying transitions among the available data.  4.  The performance of the ANN increases slightly if the results from $\\sim 10$ realizations are averaged.  5.  In the primary test case considered here, the excitation of ortho-H$_2$ by H$_2$ ($J=0$), the available data -- consisting of 990 rate coefficients -- can be modeled equally well provided that the network is trained on at least $20\\%$ of the available data.  6.  The performance of the ANN is summarized in Table 1, which lists the typical output errors in every case considered.  Here the RMS and median values of $\\rm log_{10}(ANN\\,\\,output/actual$ value) are listed for the entire test set and separately for low, medium and high values of the actual rate.  Overall, for ANN with 16 hidden neurons and 3 input neurons that are trained on at least 20$\\%$ of the available data, the typical RMS value of $\\rm log_{10}(ANN\\,\\,output/actual\\,\\,value)$ is $\\sim 0.4$; for transitions with rate coefficients $\\le 10^{-13} \\rm cm^{-3}\\,s^{-1}$, the typical RMS value is $\\sim 0.5$.  The corresponding median values of $\\vert \\rm log_{10}(ANN\\,\\,output/actual\\,\\,value) \\vert$ are $\\sim 0.25$ and $\\sim 0.33$.  These median values imply that one-half of the ANN outputs are typically accurate to within factors of 1.8 (or $\\sim 2.1$ in the case of rate coefficients $\\le 10^{-13} \\rm cm^{-3}\\,s^{-1}$.)   However, the results shown in Figure 4 indicate that some of the ANN outputs are highly discrepant (with 4$\\%$ in error by more than a factor 10). It should be noted that the errors given here apply to individual state-to-state rate coefficients.  When these are used to model molecular line emissions, the errors in the resultant line strengths can be expected to be smaller than in the state-to-state rate coefficients because a given radiative transitions may be pumped by a combination of several different collisionally-induced transitions.  \\vskip 0.5 true in  \\centerline{\\bf Appendix: Calculational details}  The calculation made use of the Fast Artificial Neural Network (FANN) Library, version 2.0.0, an open source software package developed by Steffen Nissen and Evan Nemerson and available at http://fann.sourceforge.net.  The cascade training option was adopted for this problem, and was found to result in rapid training.  In this algorithm (Fahlman \\& Lebiere 1990), the training starts with an empty network devoid of hidden neurons.   Hidden neurons are trained and added to the network one at a time; each new hidden neuron is chosen from a selection of candidate neurons with different initial weights and one of several possible activation functions.  In using the FANN library, the input neuron values were rescaled to \\re{small values (relative to unity) to ensure that the input neurons did not saturate}: energies (given as wavenumber in units of cm$^{-1}$) were therefore divided by $10^5$ and rotational quantum numbers were divided by 100.  \\re{No other preprocessing was performed.  With these rescalings, the input values lay in the ranges [0, 0.14], [--0.03, 0.10], and [--0.16, 0.18] respectively for $\\Delta E$, $\\Delta J$, and $\\Delta \\tau$.} The desired output neuron values were the logarithm \\re{to base 10} of the rate coefficient in units of $\\rm cm^{3}\\,s^{-1}$, \\re{and covered the range [--15.3, --9.9]}  The cascade algorithm was used with all the default settings except for one: a linear function was selected as the activation function for the output layer.  This choice was necessitated by the fact that the output neuron values lay outside the range of the other activation functions\\footnote{\\re{An alternative solution would have been to rescale the output neuron values.}}.  The desired error was set to a very small value to ensure that the ANN made use of the maximum number of hidden neurons specified for each test. \\re{The default settings for the FANN software make six activation functions available to the algorithm, all of which are continuous functions with a domain of [$-\\infty,+\\infty$] and a range of either [0, 1] or [--1, 1].  The network is optimized using the i-RPROP algorithm of Igel \\& H\\\"usken (2000), a variant of the RPROP algorithm introduced by Riedmiller \\& Braun (1993).}  Even if the settings and training set are unchanged, the results are slightly different every time the ANN is run.  This is because the initial weights are randomly assigned at the start of training.  I therefore averaged the output values over a set of $N_r$ runs.  Figure 7 shows how the performance of the network depends upon $N_r$, for a network with 3 input neurons and 16 hidden neurons and with $N_{train}=30$.  As in Figures 3 -- 5, black curves apply to the entire test set, while blue, green and red histograms apply to transitions with low, medium and high actual rate coefficients; for this purpose, ``low\" applies to value smaller than $10^{-13} \\rm \\, cm^3\\, s^{-1}$, ``medium\" to value in the range $10^{-13} - 10^{-11} \\rm \\, cm^3\\, s^{-1}$, and ``high\"\u00a0to a value greater than $10^{-11} \\rm \\, cm^3\\, s^{-1}$. The magenta curve refers to the training set.  Clearly, the performance on the test set is improved if multiple runs are averaged, but for $N_r$ greater than a few, no further improvements are realized.  Based upon the behavior exhibited in Figure 7, I selected $N_r = 10$ as the standard parameter for the calculations that yielded the results discussed in \\S3.  Figures 8 and 9 show how the performance depends upon the final number of hidden neurons.  The color coding in Figure 9 is the same as that in Figure 7.   Clearly, if enough hidden neurons are provided, the fit to the training set becomes extremely accurate (lower right panel in Figure 8), with the RMS error for 64 neurons being less than 3$\\%$.  However, the performance on the {\\it test} set worsens somewhat once the number of hidden neurons exceeds 16.  This behavior indicates that the network is ``overfitting\" the data (e.g.\\ Smith 1993).  The following analogy is presented by the use of polynomial fitting functions.  Given a function of one variable that has been evaluated at $N$ values, a perfect polynomial fit to those values can always be obtained with a polynomial of order $N-1$.  However, such a polynomial may show strong wiggles between (and beyond) the points where the function was evaluated (particularly if the function evaluations have errors).  A lower order polynomial may fit the data points imperfectly but may be preferable for the purposes of interpolation and particularly extrapolation. Based upon the behavior exhibited in Figure 8 and 9, I selected $N_{hidden} = 16$ as the standard parameter for the calculations that yielded the results discussed in \\S3."
    },
    "0911/0911.2187_arXiv.txt": {
        "abstract": " ",
        "introduction": "Whenever bending of light by gravitational fields is detected in astronomical observations it is generally very weak. Even in the spectacular formation of giant arcs, typical of the so-called strong lensing, photons are deflected by no more than a few arcseconds. This leads to the conclusion that the first order Einstein formula for the deflection angle is sufficient to explain all observed phenomenology% . On the other hand, the existence of very compact objects, such as neutron stars or black holes, is now well-established through many independent astrophysical observations. In the neighbourhood of such objects, electromagnetic radiation is generated and travels through very strong gravitational fields. In such extreme cases, the calculation of the deflection of photons needs to be pushed beyond the first order approximation on which the Einstein formula relies. It must be said that no direct observation of gravitational lensing by black holes or other compact objects has ever been performed up to date. Yet, some peculiarities of the emission spectra of black holes, such as the shape of the $K_\\alpha$ line of the iron, have been interpreted as the final effect of the strong deflection of photons emitted by the accretion disk.  The aim of this article is to review the theoretical aspects of gravitational lensing by black holes, and discuss the perspectives for realistic observations. We will first treat lensing by spherically symmetric black holes, in which the formation of infinite sequences of higher order images emerges in the clearest way. We will then consider the effects of the spin of the black hole, with the formation of giant higher order caustics and multiple images. Finally, we will consider the perspectives for observations of black hole lensing, from the detection of secondary images of stellar sources to the interpretation of iron K-$\\alpha$ lines and direct imaging of the shadow of the black hole. ",
        "conclusions": ""
    },
    "0911/0911.0711_arXiv.txt": {
        "abstract": "We investigate how the divergence-free property of magnetic fields can be exploited to resolve the azimuthal ambiguity  present in solar vector magnetogram data, by using line-of-sight and horizontal heliographic derivative information as approximated from discrete measurements. Using synthetic data we test several methods that each make different assumptions about how the divergence-free property can be used to resolve the ambiguity. We find that the most robust algorithm involves the minimisation of the absolute value of the divergence summed over the entire field of view. Away from disk centre this method requires the sign and magnitude of the line-of-sight derivatives of all three components of the magnetic field vector. ",
        "introduction": "Measurements of the vector magnetic field in the solar atmosphere are crucial for understanding a wide variety of physical phenomena. Using the linear polarisation of magnetically sensitive spectral lines to infer the component of the field transverse to the line-of-sight results in an ambiguity of 180$^\\circ$ in its direction \\cite{1969PhDT3H}. This ambiguity must be resolved in order to completely determine the magnetic field vector. Several methods are currently in use to resolve the ambiguity (for overviews see \\opencite{metcalfetal06}; \\opencite{lekaetal2008}). In general, each method involves an assumption or approximation that may not be appropriate for observations of solar magnetic fields.  A promising approach, which can potentially avoid unrealistic assumptions, is to exploit the divergence-free property of magnetic fields (\\myeg \\opencite{wuai1990}; \\opencite{1993AA...278..279C}; \\opencite{1993AA...279..214L}; \\opencite{1999AA...347.1005B}; \\opencite{lietal07}; \\opencite{cb2007}). One challenge in calculating the divergence is that it requires information about the variation of the field in the direction perpendicular to the solar surface. To this end it is possible to derive from observations the variation of the field along the  line-of-sight direction (\\myeg \\opencite{ruizcobodeltooroiniesta92}; \\opencite{1994AA...291..622C}; \\opencite{1995ApJ...439..474M}; \\opencite{1996SoPh..164..169D}; \\opencite{1996SoPh..169...79L}; \\opencite{1998ApJ...494..453W}, \\citeyear{2001ApJ...547.1130W}; \\opencite{2000ApJ...530..977S}; \\opencite{eibeetal2002}; \\opencite{lekametcalf2003}; \\opencite{2005ApJ...631L.167S}, \\citeyear{2007ApJS..169..439S}). The line-of-sight direction is not perpendicular to the solar surface, except at disk centre; therefore some care is required to correctly treat the geometrical perspective.  In \\citeauthor{cb2007}~(\\citeyear{cb2007}, henceforth paper~I) we showed that the  ambiguity  can be resolved with line-of-sight and horizontal heliographic derivative information by using the divergence-free property  without additional assumptions about the nature of the field. To resolve the ambiguity away from disk centre we demonstrated that it is sufficient to determine four pieces of information: {\\it i)}   the sign of the line-of-sight derivative of the line-of-sight component of the magnetic field; {\\it ii)}  the sign and magnitude of the line-of-sight derivative of the magnitude of the transverse component of the field; {\\it iii)} the sign and magnitude of the line-of-sight derivative of the azimuthal angle; and {\\it iv)}  the sign and magnitude of the horizontal heliographic derivatives of the magnitude of the transverse component of the field and the azimuthal angle. Paper~I was a theoretical examination in which we assumed that all  necessary derivatives could be determined exactly at a given location and that the solar surface could be represented by the heliographic plane.  The aim of this paper is to investigate how the divergence-free property can be used to resolve the ambiguity for data where the magnetic field is measured at discrete spatial locations. Using synthetic data we test two methods that have previously appeared in the literature (\\myeg \\opencite{wuai1990}; \\opencite{1993AA...278..279C}; \\opencite{1993AA...279..214L}; \\opencite{1999AA...347.1005B}; \\opencite{lietal07}); both of these are based on the divergence but they  make different assumptions about how to use it to resolve the ambiguity. We show that both of these methods can produce significant errors. For this reason, we present a more robust algorithm, the global minimisation method.  The outline of this paper is as follows. In Section~\\ref{sec_divb} we describe the derivation of the divergence-free condition for any position on the solar disk in terms of observable quantities and discuss the consequences for ambiguity resolution. In Section~\\ref{sec_synth} we review the synthetic data and metrics that will be used to test the performance of the ambiguity resolution algorithms. In Section~\\ref{algorithms} we test the algorithms and the validity of the assumptions made. In Section~\\ref{sec_conc} we draw conclusions. ",
        "conclusions": "\\label{sec_conc}  We investigate how the divergence-free property of magnetic fields can be exploited to resolve the azimuthal ambiguity that is present in solar vector magnetogram data by using line-of-sight and horizontal heliographic derivative information approximated from discrete measurements. Using synthetic data we test several methods that each make different assumptions about how to use the divergence to resolve the ambiguity. We find that all of the assumptions considered can be violated when the divergence is approximated from discrete measurements (depending on the nature of the field and the resolution of the observations), but some assumptions are more robust than others.  The \\citeauthor{wuai1990} criterion expresses the divergence-free condition as an inequality (\\myeg \\opencite{wuai1990}; \\opencite{1993AA...278..279C}; \\opencite{1993AA...279..214L}; \\opencite{lietal07}). We find that methods based on this criterion are very sensitive to the smoothness of the initial configuration of azimuthal angles and the order in which pixels are examined. In addition, the underlying assumption can be incorrect over a significant fraction of the field of view when derivatives are approximated from discrete measurements. The sequential minimisation method (\\opencite{1999AA...347.1005B}) assumes that the magnitude of the divergence is minimised over the pixels used to approximate the divergence at each point. Our tests indicate that this approach is more robust than the \\citeauthor{wuai1990} criterion but can produce errors when the underlying assumption is violated. Moreover, the algorithm can propagate erroneous solutions into regions that do not violate the underlying assumption.  We find that the most promising algorithm is the global minimisation method which assumes that the correct configuration of azimuthal angles corresponds to the minimum of the magnitude of the divergence summed over the entire field of view. This method relies on the  magnitude of the divergence. Therefore, with this method, the {\\it sign and magnitude} of the line-of-sight derivatives of all three components of the magnetic field vector are required to resolve the azimuthal ambiguity away from disk centre.  For the synthetic data employed here, the field is sampled at discrete spatial locations where the exact value of the field is provided. Thus, the effects of noise and unresolved structure in both the horizontal heliographic and line-of-sight directions are neglected. The performance of the global minimisation method is very promising, but before this approach can be reliably applied to real magnetogram data its performance should be tested with more realistic synthetic data that include the effects of noise and unresolved structure (\\myeg \\opencite{lekaetal2008}).  \\begin{acks} The authors would like to thank Yuhong Fan for providing one of the synthetic data sets used in this investigation. This work was supported by funding from NASA under contracts NNH05CC75C and NNH09CE60C. \\end{acks}"
    },
    "0911/0911.2008_arXiv.txt": {
        "abstract": "Analytical relations are derived for the amplitude of astrometric, photometric and radial velocity perturbations caused by a single rotating spot. The relative power of the star spot jitter is estimated and compared with the available data for $\\kappa^1$ Ceti and HD 166435, as well as with numerical simulations for $\\kappa^1$ Ceti and the Sun. A Sun-like star inclined at $i=90\\degr$ at 10 pc is predicted to have a RMS jitter of $0.087$ \\uas\\ in its astrometric position along the equator, and $0.38$ m s$^{-1}$ in radial velocities. If the presence of spots due to stellar activity is the ultimate limiting factor for planet detection, the sensitivity of SIM Lite to Earth-like planets in habitable zones is about an order of magnitude higher that the sensitivity of prospective ultra-precise radial velocity observations of nearby stars. ",
        "introduction": "With the anticipated launch of the SIM Lite mission in the near future, we are embarking on a long and exciting journey of exoplanet detection by astrometric means. One of the main goals of this mission is the detection of habitable Earth-like planets around nearby stars \\citep{unw}. To date, most of the exoplanet discoveries have been made by the Doppler-shift technique, while the astrometric method has been limited to the use of the FGS on the Hubble Space Telescope \\citep{ben} and to ground-based CCD observations of low-mass stars with giant, super-Jupiter,  planetary companions \\citep{sha}. In achieving the strategic goal of confident detection of rocky, Earth-sized planets in the habitable zone, the prospective astrometric and spectroscopic ultra-precise measurements will encounter a number of limitations of technical and astrophysical nature.  For the Doppler-shift technique, many of these limitations will be dealt with by further improvement in the instrumentation or refinement of the observational procedure \\citep{may}. However, the presence of astrophysical noise due to stellar magnetic activity emerges as the ultimate bound on the sensitivity of planet detection techniques, and the only remedy suggested thus far, is selection of particularly inactive, slowly rotating stars. Indeed, a very small fraction of stars in the high-precision HARPS program of exoplanet search exhibit radial velocity (RV) scatter of less than 0.5 m s$^{-1}$. Although this type of variability is probably driven by the rotation of bright and dark structures on the surface (star spots and plage areas), the frequency power spectrum of such perturbations can be fairly flat, extending to frequencies much higher or lower than the rotation, as shown by \\citet{cat} for the Sun. Arguments have been presented \\citep[e.g., ][]{eri} that star spots also result in very large astrometric noise of $\\sim 10$ $\\mu$AU, which should thwart discovery of habitable Earth analogs. The aim of this paper is to quantify the effects of rotating spots in astrometric photometric and RV measurements more accurately, taking into account the limb darkening, geometric projection and differential rotation, and to assess the expected vulnerability of the RV and astrometric methods to such perturbations. We do this by direct analysis as well as by numerical simulation, and support our findings by the data for the Sun and two rapidly rotating stars. ",
        "conclusions": "\\label{con.sec} Our results for the Sun are in good agreement with the approximate relations by \\citet{eri}, who estimated a positional standard deviation of 0.7 $\\mu$AU. At the same time, their conclusion that \"for most spectral types the astrometric jitter is expected to be of the order of 10 micro-AU or greater\" is misleading, because it is largely based on overestimated values of photometric variability from ground-based observations, and it does not differentiate the luminosity classes of giants and dwarfs. It can not be concluded that the Sun is exceptionally inactive compared to its peers just because the ultra-precise solar irradiance data, such as PMOD or SOHO, reveal a much smaller scatter than the inferior photometric data for other stars. We investigated the indices of chromospheric activity ($\\log R^\\prime_{\\rm HK}$) and available rotation periods for some 80 SIM targets, and found that half of them should rotate with the same rate as the Sun, or slower.   \\citet{hal} presented a detailed study of variability of solar-type stars and its relation to the index of chromospheric activity $\\log R'_{\\rm HK}$, based on 14 years of photometric and spectroscopic observations. They found that the Sun at $\\log R'_{\\rm HK}=-4.96$ is not more variable than its F-G peers at the low end of the activity distribution. Given that most Solar-type stars exhibit similar low levels of chromospheric activity \\citep{gra}, we expect that finding stars with levels of jitter similar to, or lower than the Sun, should not be a problem. Low-jitter, stable stars are common and plentiful, which augurs well for the prospects of finding small, rocky planets with {\\it Kepler} \\citep{bat}.  \\begin{deluxetable}{lrrr} \\tabletypesize{\\scriptsize} \\tablecaption{Observable signals and star spot jitters for an Earth-like planet orbiting a typical dwarf star at 10 pc. \\label{snr.tab}} \\tablewidth{0pt} \\startdata Star type \\dotfill & Sun & F5V & K5V \\\\ Rotation period, d \\dotfill & 25.4 & 18 & 30 \\\\ Astrometric signal, \\uas\\ \\dotfill & 0.30 & 0.23 & 0.45 \\\\ RV signal, m s$^{-1}$ \\dotfill & 0.089 & 0.078 & 0.109 \\\\ Astrometric jitter, \\uas\\ \\dotfill & 0.087 & 0.113 & 0.063\\\\ RV jitter, m s$^{-1}$ \\dotfill & 0.38 & 0.69 & 0.23  \\\\ Astrometric SNR \\dotfill & 3.4 & 2.0 & 7.1 \\\\ RV SNR \\dotfill & 0.23 & 0.11 & 0.47 \\\\ \\enddata \\end{deluxetable}  The SIM Lite Astrometric Observatory (formerly known as the Space Interferometry Mission) will achieve a single-measurement accuracy of 1 \\uas\\ or better in the differential regime of observation \\citep{unw, sim}. Several previous studies have addressed SIM's exoplanet detection and orbital characterization capabilities \\citep[][and references therein]{catpa}. The \"Tier 1\" program includes $\\sim 60$ nearby stars for which the astrometric signature of a terrestrial habitable planet is large enough to be confidently measured by SIM. The recently completed double-blind test \\citep{tra} demonstrated that Earth-like planets around nearby stars can be discovered and measured even in complex planetary systems. In this paper, we are concerned with the more general question of the ultimate limit to planet detectability set by the activity-related jitter. Assuming that the instrumentation progresses to levels where observational noise becomes insignificant, which technique holds the best prospects for detection of habitable Earth-like planets? Table \\ref{snr.tab} summarizes the expected RMS jitter and the exoplanet signal for three typical nearby stars. In all cases, the solar spot filling factor ($r^2$) is assumed. The signal-to-noise ratios (SNR) per observation include only the star spot jitter. Note that the astrometric SNR in this case is independent of the distance to the host star, because both the signal and the star spot jitter are inversely proportional to distance. As a measure of the relative sensitivity of the two methods, the ratio of the SNR values for a given star is independent of the planet mass or the filling factor to first-order approximation. The physical radius of the star and the period of rotation are the two parameters with the largest impact on the relative sensitivity, but their combined effect is rather modest for normal stars, as is seen in the Table. Therefore, in the ultimate limit of exoplanet detection defined by intrinsic astrophysical perturbations, the astrometric method is at least an order of magnitude more sensitive than the Doppler technique for most nearby solar-type stars."
    },
    "0911/0911.1301_arXiv.txt": {
        "abstract": "High contrast coronagraphic imaging of the immediate surrounding of stars requires exquisite control of low-order wavefront aberrations, such as tip-tilt (pointing) and focus. We propose an accurate, efficient and easy to implement technique to measure such aberrations in coronagraphs which use a focal plane mask to block starlight. The Coronagraphic Low Order Wavefront Sensor (CLOWFS) produces a defocused image of a reflective focal plane ring to measure low order aberrations. Even for small levels of wavefront aberration, the proposed scheme produces large intensity signals which can be easily measured, and therefore does not require highly accurate calibration of either the detector or optical elements. The CLOWFS achieves nearly optimal sensitivity and is immune from non-common path errors. This technique is especially well suited for high performance low inner working angle (IWA) coronagraphs. On phase-induced amplitude apodization (PIAA) type coronagraphs, it can unambiguously recover aberrations which originate from either side of the beam shaping introduced by the PIAA optics. We show that the proposed CLOWFS can measure sub-milliarcsecond telescope pointing errors several orders of magnitude faster than would be possible in the coronagraphic science focal plane alone, and can also accurately calibrate residual coronagraphic leaks due to residual low order aberrations. We have demonstrated $\\approx 10^{-3} \\lambda/D$ pointing stability in a laboratory demonstration of the CLOWFS on a PIAA type coronagraph. ",
        "introduction": "High performance coronagraphs with small inner working angle (IWA) are unavoidably very sensitive to small pointing errors and other low order aberrations \\citep{lloy05,shak05,siva05,beli06,guyo06}. This property is due to the fact that the wavefront of a source at small angular distance (typically between 1 and 2 $\\lambda/D$) from the optical axis is \"similar\" (in the linear algebra sense of the term) to an on-axis wavefront with a small ($\\ll \\lambda/D$) pointing error. If the coronagraph must \"transmit\" the former, it will also transmit a significant part of the latter, and therefore be extremely sensitive to pointing errors. This behavior is indeed verified by performance comparison between coronagraph concepts \\citep{guyo06}.  \\begin{figure*}[htb] \\includegraphics[scale=0.35]{f1.eps} % \\caption{\\label{fig:lowfsprinciple} Optical layout of a coronagraphic low order wavefront sensor system, shown here with a PIAA coronagraph. See text for details.} \\end{figure*}  For high performance coronagraphs, such as the Phase-Induced Amplitude Apodization (PIAA) coronagraph \\citep{guyo03,guyo05} used as example in this paper, the coronagraph should ideally be designed to balance IWA against stellar angular diameter, which sets a fundamental limit on the achievable coronagraph performance. Such coronagraphs are therefore pushed to become sensitive to pointing errors corresponding to the angular size of nearby stars, roughly 1 milliarcsecond (mas), and are also highly sensitive to other low order wavefront errors such as focus and astigmatism. Milliarcsecond-level pointing error can increase stellar leakage in the coronagraph to the point where a planet would be lost in photon noise. Even smaller errors can create, if not independantly measured, a signal which is very similar to a planet's image.    A robust and accurate measurement of low order aberrations (especially tip-tilt errors, which are easily generated by telescope pointing errors and vibrations) is therefore essential for high contrast coronagraphic observations at small angular separation. The science focal plane after the coronagraph is unfortunately \"blind\" to small levels of low order aberrations, which can only be seen when already too large to maintain high contrast in the coronagraphic science image. A better option is to monitor pointing errors by using starlight which would otherwise be rejected by the coronagraph. This scheme was successfully implemented on the LYOT project coronagraph \\citep{oppe04,digb06}. Alternatively, the measurement could be performed independently from the coronagraph optical train (for example, the wavefront sensor in the adaptive optics system upstream of the coronagraph). We propose in this paper an improved solution to obtain accurate measurement of several low-order aberrations including pointing: the coronagraphic low-order wavefront sensor (CLOWFS). The wavefront control requirements for a PIAA coronagraph are first clearly defined in \\S\\ref{sec:requ}. The CLOWFS principle is presented in \\S\\ref{sec:principle}. Wavefront reconstruction algorithms and CLOWFS sensitivity are discussed in \\S\\ref{sec:wfreconstr}. The aberration sensitivity of a PIAA coronagraph equipped with a CLOWFS is discussed in \\S\\ref{sec:aberrsens}, and the results of a laboratory demonstration on a PIAA Coronagraph system are shown in \\S\\ref{sec:labexp}. ",
        "conclusions": "The CLOWFS design presented in this paper can efficiently measure low order aberrations \"for free\", as it uses light that would otherwise be discarded by the coronagraph. Both the hardware configuration and software algorithms presented are easy to implement and their performance is robust against calibration errors, chromaticity, non-common path errors and small errors/aberrations in the optical components. The CLOWFS pointing measurement can also lead to improved astrometric accuracy for the position of faint companions, as the star position on the image is usually difficult to measure in coronagraphic images \\citep{digb06}.   In this paper, we have studied a CLOWFS design on a low-IWA PIAA coronagraph. Although coronagraphs with larger IWA can tolerate larger amount of low order aberrations (see for example Kuchner \\& Traub 2002; Kuchner et al. 2005), they require these aberration to be measured ahead of the coronagraphed beam because they are \"blind\" to aberrations until they are large enough to produce significant coronagraph leaks. The CLOWFS would therefore be very useful to any high contrast coronagraph, and the optical design presented in this paper can readily be used on any coronagraph where a focal plane mask physically blocks starlight.  The CLOWFS can also be applied to phase mask coronagraphs. For such coronagraphs (see for example Roddier \\& Roddier 1997; Rouan et al. 2000; Palacios 2005), where starlight is diffracted outside the pupil by the focal plane mask, a modified Lyot stop is placed on the pupil plane to reflect starlight to the CLOWFS. In addition, using a pattern matching algorithm, the CLOWFS can estimate low-order wavefront aberrations accurately and quickly even in non-linear region. These new CLOWFS designs and their performances will be presented in an upcoming paper."
    },
    "0911/0911.1592_arXiv.txt": {
        "abstract": "Stacking analysis is a means of detecting faint sources using a priori position information to estimate an aggregate signal from individually undetected objects.  Confusion severely limits the effectiveness of stacking in deep surveys with limited angular resolution, particularly at far infrared to submillimeter wavelengths, and causes a bias in stacking results.  Deblending corrects measured fluxes for confusion from adjacent sources; however, we find that standard deblending methods only reduce the bias by roughly a factor of two while tripling the variance.  We present an improved algorithm for simultaneous stacking and deblending that greatly reduces bias in the flux estimate with nearly minimum variance.  When confusion from neighboring sources is the dominant error, our method improves upon RMS error by at least a factor of three and as much as an order of magnitude compared to other algorithms.  This improvement will be useful for Herschel and other telescopes working in a source confused, low signal to noise regime. ",
        "introduction": "At every wavelength of astronomy, there is a need to extract faint signals buried in noisy images.  For example, a certain class of objects may be lurking undetected below the noise threshold of a data set at a particular wavelength.  However, if accurate positions for the objects are known a priori from detections at other wavelengths, then this prior information may be exploited to break through the noise floor and achieve a detection, at least in a statistical sense, for the objects in question.  Stacking is an effective analysis technique in cases where noisy images at one wavelength are complemented by catalogs of known objects that have been detected in the same region of sky at other wavelengths.  The noisy data at known object positions are averaged, and the resulting aggregate flux measurement has an enhanced signal to noise ratio due to averaging, which reduces the noise by a factor $1/\\sqrt{R}$, where $R$ is the number of target objects in the stack.  As an early application of this method in x-ray astronomy, stacking analysis was used to study x-ray emission from stars \\citep{1985ApJ...289..279C}; subsequently, stacking became a standard x-ray analysis technique that, among others, has been applied to normal galaxies \\citep{2001AJ....122....1B}, Lyman Break Galaxies (LBGs) \\citep{2001ApJ...558L...5B} and radio sources \\citep{2003MNRAS.345..939G}.  At optical to near infrared wavelengths stacking has also become widely adopted, and has been used for example to study galactic halos \\citep{Zibetti2004} and intergalactic stars \\citep{Zibetti2005}, as well as the cosmic infrared background \\citep{2006A&A...451..417D}, star forming galaxies \\citep{2004ApJ...610..745L, 2007ApJ...671..278G} and Extremely Red Objects (EROs) \\citep{1997AJ....113..474H}.  At submillimeter wavelengths, stacking has been employed to study the submillimeter background \\citep{2000MNRAS.318..535P, 2004ApJS..154..118S, 2006ApJ...647...74W, 2006ApJ...644..769D, 2009arXiv0904.1205M, 2009arXiv0904.0028G}, optical/IR color selected galaxies \\citep{2004ApJ...605..645W, 2005ApJ...631L..13D, 2007MNRAS.381.1154T, 2009MNRAS.394....3D, 2009arXiv0904.0028G}, and radio detected galaxies \\citep{2008MNRAS.385.2225S}.  Similarly, at radio wavelengths stacking has been widely adopted, for example to study LBGs \\citep{2007ApJ...660L..77I, 2008ApJ...689..883C}, star forming BzK galaxies \\citep{2005ApJ...631L..13D, 2007MNRAS.381.1154T, 2009MNRAS.394....3D}, Distant Red Galaxies (DRGs) \\citep{2007ApJ...660L..77I}, EROs \\citep{2009MNRAS.394....3D} and quasars \\citep{2007ApJ...654...99W}.  It has been widely recognized that confusion noise limits astronomical measurements particularly at long wavelengths \\citep{1974ApJ...188..279C, 2001AJ....121.1207H} and confusion poses special complications for stacking analyses.  When multiple sources crowd a resolution element of the data, then blending of sources must be accounted for in the stacking procedure.  Blending is a particular problem in the submillimeter regime, where at present images suffer from poor resolution compared to many other astronomical observations.  Methods have been used for deblending in submillimeter data, however a systematic assessment of the effectiveness of these techniques has not been previously reported.  This paper addresses the problem of deblending the various unresolved sources of submillimeter emission in order to effectively stack objects in the confusion dominated, low signal to noise, regime.  We present an algorithm that accomplishes simultaneous stacking and deblending of these data and demonstrate its effectiveness through Monte Carlo simulation.  We also demonstrate that the standard approaches to deblending are in fact subject to substantial bias and scatter when applied to real data. ",
        "conclusions": "The simultaneous stacking and deblending algorithm presented here demonstrated more than a factor of three improvement in RMS error over other deblending algorithms at all signal to noise ratios, and by an order of magnitude at large signal to noise ratios.  This improvement translates into a greater sensitivity for stacking analyses that will be applicable to present and future far infrared and submillimeter surveys.  For example, the Herschel SPIRE instrument will reach 1$\\sigma$ survey depths of 3 mJy in 10-16 hours  at 18-36$''$ spatial resolution \\citep{2008SPIE.7010E...4G}.  At this depth, and with 1,000 known source positions, stacking analyses can reach a depth of 0.1 mJy when limited only by statistical error.  This depth would be sufficient to detect stacking signals from color selected galaxies such as BzKs, DRGs and EROs that have been detected in stacking at 870 $\\mu$m with fluxes in the 0.2-0.4 mJy range with the ground based LABOCA instrument \\citep{2009arXiv0904.0028G}.  However, with the factor of three poorer sensitivity caused by standard deblending algorithms, these objects would remain out of reach to stacking analysis with Herschel, or their reported stacked fluxes would be significantly biased."
    },
    "0911/0911.1889_arXiv.txt": {
        "abstract": "We present new measurements of the energy spectra of cosmic-ray (CR) nuclei from the second flight of  the balloon-borne experiment Cosmic Ray Energetics And Mass (CREAM). The instrument  included different particle detectors to provide  redundant charge identification and  measure the energy of CRs up to several hundred TeV. The measured individual energy spectra of C, O, Ne, Mg, Si, and Fe are presented up to $\\sim 10^{14}$ eV. The spectral shape looks nearly the same for these primary elements and it can be fitted to an $E^{-2.66 \\pm 0.04}$ power law in energy. Moreover, a new measurement of the absolute intensity of nitrogen in the 100-800 GeV/$n$ energy range with smaller errors than previous observations, clearly indicates a hardening of the  spectrum at high energy. The relative abundance of N/O at the top of the atmosphere is measured to be $0.080 \\pm 0.025\\, $(stat.)$\\, \\pm 0.025\\,$(sys.) at $\\sim\\,$800 GeV/$n$, in good agreement with a recent result from the first CREAM flight. ",
        "introduction": "Experimental studies of charged cosmic rays (CRs) are focused on the  understanding of the  acceleration mechanism of high-energy CRs, identification of their sources, and clarification of their interactions with the interstellar medium (ISM). The origin of CRs is still under debate, although direct measurements  with instruments on stratospheric balloons or in space have provided  information on their elemental composition and energy spectra, and indirect detection by ground-based experiments has extended  the measurements up to the end of the observed all-particle spectrum. \\\\ It is generally accepted that CRs are accelerated in blast waves of supernova remnants (SNRs), which are the only galactic candidates known with sufficient energy output to sustain the CR flux \\citep{Bell, hillas}. Recent observation of emission of TeV gamma-rays from SNR RX J1713.7-3946 by the HESS Cherenkov telescope array proved that high-energy charged particles are accelerated in SNR shocks  up to energies beyond 100 TeV \\citep{HESS}. This result is consistent with  expectations of the class of theoretical models that predict the existence of a rigidity-dependent limit, above which the diffusive shock acceleration becomes inefficient. The maximum energy attainable by a nucleus of charge $Z$ may range from $Z\\times$10$^{14}$ eV to $\\sim$ $Z\\times$10$^{17}$ eV depending on the model and types of supernovae considered \\citep{LC, Berezhko, parizot}. In this scenario, the energy spectra of elements exhibit a $Z$-dependent cutoff. As a consequence, the CR elemental composition is expected to change as a function of energy, marked by a  depletion of low-$Z$ nuclei in the several hundred TeV region. This mechanism has been proposed as a possible explanation of the steepening (``knee'') in the CR energy spectrum (described by a power law: $dN$/$dE \\propto E^{-\\gamma}$) with a change of the spectral index from $\\gamma\\approx\\,$2.7 to $\\gamma\\approx\\,$3.1  observed at energies around 4 PeV. An alternative approach is adopted by models that relate the knee to leakage of CRs from the Galaxy. In this case, the knee is expected to occur at lower energies for light elements  as compared to heavy nuclei, due to the rigidity-dependence  of the Larmor radius of  CR particles  propagating in the galactic magnetic field  \\citep{horandel2}.\\\\ A detailed understanding of the propagation of CRs during their wanderings inside the Galaxy and their interactions with the ISM is needed to infer  the injection spectra  of individual elements at the source from the observed spectra at Earth. CR nuclei can be divided into primaries (i.e., H, He, C, O, Fe, etc.) that come from CR sources, and secondaries (such as Li, Be, B, Sc, Ti, V, etc.), which are  products of the interactions of primary nuclei with the ISM. The amount of material traversed by CRs between injection and observation can be derived from the measured ratio of secondary-to-primary  nuclei, such as the boron-to-carbon ratio (B/C) or the ratio of sub-iron elements to iron. Likewise, their confinement time in the Galaxy can be determined through measurements of long-lived radioactive secondary nuclei \\citep{longair, yanasak}. Observations from space-based experiments like CRN on the Spacelab2 mission of the Space Shuttle \\citep{CRN2} and  C2 \\citep{HEAO} and Heavy Nuclei Experiment (HNE) \\citep{binns} onboard the {\\em HEAO 3} satellite have shown that the energy spectra of the secondary nuclei  are steeper than those of the  primaries. The escape pathlength of CRs from the Galaxy depends on the  magnetic rigidity as $R^{-\\delta}$, with the parameter $\\delta\\approx\\,$0.6. This implies that the spectra observed at Earth are steeper  than the injection spectra, i.e., the spectral index at the source is smaller by the value $\\delta$. \\\\ Those pioneering measurements were statistically limited at energies of order $\\sim$ 10$^{11}$ eV. Accurate direct measurements of the energy spectra of individual elements into the knee region are needed to discriminate between different astrophysical models proposed to explain the acceleration and propagation mechanisms. Recently, a new generation of balloon-borne experiments has performed accurate measurements of CRs in the previously experimentally unexplored TeV region \\citep{ATIC, TRACER}. Among these, the Cosmic Ray Energetics And Mass (CREAM) experiment was designed to  measure directly the elemental composition  and the energy spectra of CRs from hydrogen to iron over the energy range 10$^{11}$-10$^{15}$ eV in a series of flights \\citep{seo}. Since 2004, four instruments were  successfully flown on long-duration balloons in Antarctica. The instrument configurations varied slightly in each mission due to various detector upgrades. The first flight produced important results extending the measurement of the relative abundances of CR secondary nuclei B and N close to the highest  energies ($\\sim$ 1.5 TeV/$n$) allowed by the irreducible background generated by the residual atmospheric overburden at balloon altitudes. The data clearly indicate that the escape pathlength of CRs from the Galaxy has a power-law rigidity-dependence, with the parameter $\\delta$ in the range 0.5-0.6 \\citep{Ahn2008}.  In this work, we present new  measurements of the energy spectra of the major primary CR nuclei from carbon to iron, up to a few TeV/$n$ made with the  second CREAM flight (CREAM-II). A new measurement of the nitrogen intensity with unprecedented statistics is also presented up to 800 GeV/$n$. The measured N/O ratio confirms the results of CREAM-I. The procedure used to analyze the data and reconstruct the CR energy spectra is described in detail together with an assessment of the systematic uncertainties.\\\\ ",
        "conclusions": "CREAM-II  carried out measurements of high-$Z$ CR nuclei with an excellent charge resolution and  a reliable energy determination. The use of a dual layer of silicon sensors allowed us to unambiguously identify each individual element up to iron and to reduce the background of misidentified nuclei. We demonstrated the feasibility of measuring the particle energy up to hundreds of TeV by using  a thin ionization sampling calorimeter, with sufficient resolution to reconstruct accurately the CR spectra. \\\\ The energy spectra of the major primary heavy nuclei from C to Fe were measured up to $\\sim\\,$10$^{14}$ eV and found to agree well with earlier direct measurements. All the spectra follow a power law in energy with a remarkably similar spectral index $\\bar{\\gamma} = -2.66 \\pm 0.04$, as  is plausible if they have the same origin and  share the same acceleration and propagation processes. \\\\ A new measurement of the nitrogen intensity in an energy region so far experimentally unexplored indicates a harder  spectrum than at lower energies, supporting the idea that nitrogen has  secondary as well as primary contributions. These results provide new clues for understanding the CR acceleration and propagation mechanisms, but further accurate observations with high statistics beyond 10$^{14}$ eV are needed to finally unravel the mystery of CR origin."
    },
    "0911/0911.0288_arXiv.txt": {
        "abstract": "We consider the problem of characterisation of burst sources detected with the Laser Interferometer Space Antenna (LISA) using the multi-modal nested sampling algorithm, {\\sc MultiNest}. We use {\\sc MultiNest} as a tool to search for modelled bursts from cosmic string cusps, and compute the Bayesian evidence associated with the cosmic string model. As an alternative burst model, we consider sine-Gaussian burst signals, and show how the evidence ratio can be used to choose between these two alternatives. We present results from an application of {\\sc MultiNest} to the last round of the Mock LISA Data Challenge, in which we were able to successfully detect and characterise all three of the cosmic string burst sources present in the release data set. We also present results of independent trials and show that {\\sc MultiNest} can detect cosmic string signals with signal-to-noise ratio (SNR) as low as $\\sim 7$ and sine-Gaussian signals with SNR as low as $\\sim8$. In both cases, we show that the threshold at which the sources become detectable coincides with the SNR at which the evidence ratio begins to favour the correct model over the alternative. ",
        "introduction": "The development of data analysis algorithms for gravitational wave detectors on the ground and for the proposed space-based gravitational wave detector, the Laser Interferometer Space Antenna (LISA)~\\cite{lisa}, is an active area of current research. LISA data analysis activities are being encouraged by the Mock LISA Data Challenges~\\cite{mldc} (MLDCs). Prior to the most recent round, round 3, which finished in April 2009, the MLDCs had included data sets containing individual and multiple white-dwarf binaries, a realisation of the whole galaxy of compact binaries, single and multiple non-spinning supermassive black hole (SMBH) binaries, either isolated or on top of a galactic confusion background and isolated extreme-mass-ratio inspiral (EMRI) sources in purely instrumental noise. Round 3 included a realisation of the galactic binaries that included chirping systems, a single data set containing multiple overlapping EMRIs and data sets containing three new types of source --- inspirals of spinning supermassive black hole binaries, bursts from cosmic string cusps and a stochastic background of gravitational radiation~\\cite{mldc}. Several of these sources, including the EMRI signals, the spinning black hole binaries and the cosmic string bursts, have highly multi-modal likelihood surfaces which can cause problems for algorithms such as Markov Chain Monte Carlo (MCMC), as these may become stuck in secondary maxima rather than finding the primary mode~\\cite{EtfAG,neilemri,BBGPa,BBGPb}. Recently~\\cite{MNnospin} we applied a new algorithm to the problem of gravitational wave data analysis, {\\sc MultiNest}~\\cite{feroz08,multinest}, which is a nested sampling algorithm, optimized for problems with highly multi-modal posteriors. We demonstrated that {\\sc MultiNest} could be used as a search tool and to recover posterior probability distributions for the case of data sets containing multiple signals from non-spinning SMBH binary inspirals. We used the algorithm as a search tool in MLDC round 3 to analyse two of the data sets --- the spinning SMBH inspirals (challenge 3.2) and the cosmic string bursts (challenge 3.4). We will discuss the second of these applications in this paper.  The nested sampling algorithm~\\cite{Skilling04} was developed as a tool for evaluating the Bayesian evidence. It employs a set of live points, each of which represents a particular set of parameters in the multi-dimensional search space. These points move as the algorithm progresses, and climb together through nested contours of increasing likelihood.  At each step the algorithm works by finding a point of higher likelihood than the lowest likelihood point in the live point set and then replacing the lowest likelihood point with the new point. The difficulty is to sample, efficiently, new points of higher likelihood from the remaining prior volume. {\\sc MultiNest} achieves this using an ellipsoidal rejection sampling scheme% . It has demonstrated orders of magnitude improvement over standard methods in several applications in cosmology and particle physics~\\cite{2008arXiv08100781F,2008arXiv08111199F,Feroz:2008wr,2008JHEP12024T}. It also proved to be very effective in a gravitational wave context for the test case mentioned above~\\cite{MNnospin}.  {\\sc MultiNest} can be used as a search tool and returns the posterior probability distributions as a by-product. We employed the algorithm successfully to analyse MLDC challenge 3.4, accurately finding all three of the cosmic string burst sources present in the data set. This will be discussed in more detail in Section~\\ref{sec:mldcres}. The cosmic string cusp is a special type of burst source, since the waveform is known and can be modelled, so a matched filtering search is possible and that is what {\\sc MultiNest} does in computing the likelihoods across the parameter space. However, there may be bursts of gravitational waves in the LISA data stream from other sources and the question arises as to whether we would be able to detect them and if we can distinguish cosmic string bursts from bursts due to these other sources. The evidence that {\\sc MultiNest} computes is one tool that can be used to address model selection. Evidence has been used for gravitational wave model selection to test the theory of relativity in ground based observations~\\cite{Veitch:2008wd}, and, in a LISA context, to distinguish between an empty data set and one containing a signal~\\cite{littenberg09}. Calculation of an evidence ratio requires a second model for the burst as an alternative to compare against. We chose to use a sine-Gaussian waveform as a generic alternative burst model, as this is one of the burst models commonly used in LIGO data analysis (see for instance~\\cite{LSCUsingSG}). We find that the evidence is a powerful tool for characterising bursts --- for the majority of detectable bursts, the evidence ratio strongly favours the true model over the alternative. While the sine-Gaussian model does not necessarily well describe all possible un-modelled bursts, these results suggest that the evidence can be used to correctly identify any cosmic string bursts that are present in the LISA data stream.  This paper is organised as follows. In Section~\\ref{sec:methods} we describe the methods that we employ in this analysis ---  we describe Bayesian inference, nested sampling, the {\\sc MultiNest} algorithm, the two burst waveform models we use and the detector noise spectral density. In Section~\\ref{sec:res} we present our results, including a description of our methods and results for the analysis of MLDC challenge 3.4 (in Section~\\ref{sec:mldcres}); an estimate of the signal-to-noise ratio (SNR) required for detection of cosmic string and sine-Gaussian bursts using the {\\sc MultiNest} algorithm; and a calculation of the evidence ratio of the two models for data sets containing each type of source. We finish in Section~\\ref{sec:discuss} with a summary and discussion of directions for further research. ",
        "conclusions": "\\label{sec:discuss} We have considered the use of the multi-modal nested sampling algorithm \\MN~for detection and characterisation of cosmic string burst sources in LISA data. As a search tool, the algorithm was able successfully to find the three cosmic string bursts that were present in the MLDC challenge data set. These sources, and the five sources in the MLDC training data, were correctly identified in the sense that the full signal-to-noise ratio of the injected source was recovered, and a posterior distribution for the parameters obtained. The maximum likelihood and maximum a-posteriori parameters were not particularly close to the true parameters of the injected signals, but this was a consequence of the intrinsic degeneracies in the cosmic string model parameter space and in all cases the true parameters were consistent with the recovered posterior distributions.  In controlled studies, we found that the SNR threshold required for detection of the cosmic string bursts was $\\sim 7$--$11$, depending on the burst parameters. Bursts with a low break-frequency require a higher SNR to detect than those with high break frequencies. We also explored the detection of sine-Gaussian bursts and in that case the SNR required for detection was slightly higher, being typically $8$--$11$, with sources having frequency close to Nyquist being more difficult to detect.  \\MN~is designed to evaluate the evidence of the data under a certain hypothesis, and this can be used to compare possible models for the burst sources. LISA may detect bursts from several different sources, and it is important for scientific interpretation that the nature of the burst be correctly identified. We used the Bayesian evidence as a tool to choose between two different models for a LISA burst source --- the cosmic string model and the sine-Gaussian model, which was chosen to represent a generic burst. The Bayesian evidence works very well as a discriminator between these two models. The evidence ratio begins to clearly favour the correct model over the alternative at the same SNR that the sources become loud enough to detect in the first place.  The usefulness of \\MN~as a search tool in this problem is a further illustration of the potential utility of this algorithm for LISA data analysis, as previously demonstrated in a search for non-spinning SMBH binaries~\\cite{MNnospin}. Other algorithms based on Markov Chain Monte Carlo techniques have also been applied to the search for cosmic strings~\\cite{CornishCS}. Both approaches performed equally well as search tools in the last round of the MLDC. We are now exploring the application of \\MN~to searches for other LISA sources, including spinning massive black hole binaries and extreme-mass-ratio inspirals. \\MN~was not designed primarily as a search algorithm, but as a tool for evidence evaluation and this work has demonstrated the utility of the Bayesian evidence as a tool for model selection in a LISA context. Other problems where the evidence ratio approach could be applied include choosing between relativity and alternative theories of gravity as explanations for the gravitational waves observed by LISA, or choosing between different models for a gravitational wave background present in the LISA data set. The Bayesian evidence was previously used in a LIGO context as a tool to choose between alternative theories of gravity~\\cite{Veitch:2008wd} and in a LISA context to distinguish a data set containing a source from one containing purely instrumental noise~\\cite{littenberg09}. \\MN~provides a more efficient way to compute the evidence and this should be explored in more detail in the future.  In the context of interpretation of LISA burst events, what we have considered here is only part of the picture. We have shown that we are able to correctly choose between two particular models for a burst, and this can easily be extended to include other burst models. However, LISA might also detect bursts from unmodelled sources. In that case, algorithms such as \\MN~which rely on matched filtering would find the best fit parameters within the model space, but a higher intrinsic SNR of the source would be required for detection. In such a situation, we would like to be able to say that the source was probably not from a model of particular type, e.g., not a cosmic string burst. There are several clues which would provide an indication that this was the case. The sine-Gaussian model is sufficiently generic that we would expect it, in general, to provide a better match to unmodelled bursts than the cosmic string model, which has a very specific form. Therefore, we could say that if the evidence ratio favoured the cosmic string model over the sine-Gaussian model it was highly likely that the burst was in fact a cosmic string and not something else. Similarly, if we found that several of the alternative models had almost equal evidence, but the SNR was quite high, it would be indicative that the burst was not described by any of the models. We have seen that at relatively moderate SNRs, when the signal is described by one of the models, the evidence clearly favours the true model over an alternative. If we found that two models gave almost equally good descriptions of the source, it would suggest that the burst was not fully described by either of them. A third clue would come from the shape of the posterior for the source parameters. The cosmic string waveform space contains many degeneracies, but these can be characterised theoretically for a given choice of source parameters. If the signal was not from a cosmic string, we might find that the structure of the posterior was modified. Finally, some techniques have been developed for the Bayesian reconstruction of generic bursts~\\cite{roever09} which could also be applied in a LISA context. The usefulness of these various approaches can be explored further by analysing data sets into which numerical supernovae burst waveforms have been injected. While the necessary mass of the progenitor is probably unphysically high for a supernova to produce a burst in the LISA frequency band, such waveforms provide examples of unmodelled burst signals on which to test analysis techniques. The final LISA analysis will employ a family of burst models to characterize any detected events. The work described here demonstrates that the Bayesian evidence will be a useful tool for choosing between such models, and \\MN~is a useful tool for computing those evidences."
    },
    "0911/0911.5522_arXiv.txt": {
        "abstract": "Comet 133P/Elst-Pizarro is the first-known and currently best-characterised member of the main-belt comets, a recently-identified class of objects that exhibit cometary activity but which are dynamically indistinguishable from main-belt asteroids.  We report here on the results of a multi-year monitoring campaign from 2003 to 2008, and present observations of the return of activity in 2007.  We find a pattern of activity consistent with the seasonal activity modulation hypothesis proposed by Hsieh et al. (2004, AJ, 127, 2997).  Additionally, recomputation of phase function parameters using data in which 133P was inactive yields new IAU parameters of $H_R=15.49\\pm0.05$~mag and $G_R=0.04\\pm0.05$, and linear parameters of $m_R(1,1,0)=15.80\\pm0.05$~mag and $\\beta=0.041\\pm0.005$~mag~deg$^{-1}$. Comparison between predicted magnitudes using these new parameters and the comet's actual brightnesses during its 2002 and 2007 active periods reveals the presence of unresolved coma during both episodes, on the order of $\\sim$0.20 of the nucleus cross-section in 2002 and $\\sim$0.25 in 2007. Multifilter observations during 133P's 2007 active outburst yield mean nucleus colours of $B-V=0.65\\pm0.03$, $V-R=0.36\\pm0.01$, and $R-I=0.32\\pm0.01$, with no indication of significant rotational variation, and similar colours for the trail. Finally, while 133P's trail appears shorter and weaker in 2007 than in 2002, other measures of activity strength such as dust velocity and coma contamination of nucleus photometry are found to remain approximately constant.  We attribute changes in trail strength to the timing of observations and projection effects, thus finding no evidence of any substantial decrease in activity strength between 2002 and 2007. ",
        "introduction": "Discovered on 1996 August 7 \\citep{els96}, Comet 133P/Elst-Pizarro (also designated 7968 Elst-Pizarro; hereafter 133P) orbits in the main asteroid belt ($a=3.156$~AU, $e=0.165$, $i=1.39^{\\circ}$).  It has a Tisserand parameter (with respect to Jupiter) of $T_J=3.184$, while classical comets have $T_J<3$ \\citep{vag73,kre80}. In 2005, two more objects displaying cometary activity that are likewise dynamically indistinguishable from main-belt asteroids were identified: P/2005 U1 (Read) \\citep{rea05} and 176P/LINEAR (also known as asteroid 118401 (1999 RE$_{70}$)) \\citep{hsi06c}.  Their discoveries led to the designation of a new cometary class --- the main-belt comets (MBCs) --- among which 133P is also classified \\citep{hsi06b}.  A fourth MBC, P/2008 R1 (Garradd), has also since been discovered \\citep{gar08,jew09}.  Despite the initial excitement over the discovery of the cometary nature of 133P in 1996, no physical studies or monitoring reports were published in the refereed literature until the comet's activity was re-observed in 2002 \\citep{hsi04,low05}.  Consequently, little is known about 133P's active behaviour in that intervening period.  Since knowledge of the timing of active episodes can constrain hypotheses concerning the source of the activity, we report the results of our own monitoring campaign, which began following 133P's active outburst in 2002 and which culminated in observations of renewed activity in 133P in 2007. ",
        "conclusions": "}  \\subsection{Monitoring Campaign\\label{monitoring}}  For all monitoring observations, individual $R$-band images (aligned on the object's photocenter using linear interpolation) from each night were combined into single composite images (Fig.~\\ref{images_133p}).  For reference, we also show composite images from 133P's 2002 active phase \\citep[Figs.~\\ref{images_133p}a--\\ref{images_133p}d;][]{hsi04}. Activity is marginally visible in images from 2007 May 19, 2007 August 18, and 2007 September 12 (Figs.~\\ref{images_133p}p, \\ref{images_133p}r, \\ref{images_133p}s), while the comet's characteristic dust trail is clearly visible in the image from 2007 July 17 (Fig.~\\ref{images_133p}q). We find no evidence of activity in images from 2003 September 22 through 2007 March 21 (Figs.~\\ref{images_133p}e--\\ref{images_133p}o) and from 2008 July 1 (Fig.~\\ref{images_133p}t). In all images, even those obtained while 133P was active, the FWHM of the object's surface brightness profile is consistent with the typical FWHM seeing at the time of night when those images were obtained, implying that little or no coma is present.  In Figure~\\ref{actv133p}, we mark the positions where we observed 133P to be active or where others reported it to be active, as well as positions where we observed it to be inactive, on a plan view of its orbit. The figure shows that reports of activity in 133P are approximately confined to the quadrant following perihelion, with the earliest detection of activity occurring shortly before perihelion at a true anomaly of $\\nu\\approx350^{\\circ}$ and the latest detection occurring at $\\nu\\approx90^{\\circ}$. This activity profile is consistent with the hypothesis of seasonal activity modulation described in \\citet{hsi04} and \\citet{hsi06a}, whereby 133P's activity is driven by the sublimation of a localised patch of exposed volatile material confined to either the ``northern'' or ``southern'' hemisphere of the body. Assuming non-zero obliquity, activity then only occurs during the portion of the orbit when that active site receives enough solar heating to drive sublimation, i.e., during that hemisphere's ``summer''. We note that our observations of 133P on 2008 July 1 at the NTT showed it to be inactive despite the object being observed to be active at nearly the same orbital position in 2002. We attribute this discrepancy to a combination of the low signal-to-noise of this observation and the expected extremely weak activity of 133P at that point in its orbit \\citep{hsi04}.   \\subsection{Photometric Activity Detection and Measurement\\label{activitydetection}}  When no coma is clearly visible for an object, an alternate method for detecting activity is examination of its photometric behaviour: i.e., determining whether it is consistent with an inactive object of a fixed size, or whether it shows anomalous brightening over a certain portion of its orbit.  This type of analysis led to the discovery of activity in 95P/(2060) Chiron \\citep{tho88,bus88,mee89,har90}.  In applying this approach to 133P, we recall that \\citet{hsi04} originally derived linear and IAU $H$,$G$ phase function solutions for 133P using data taken in 2002 when the object was visibly emitting dust.  In the case of that data set, 133P's activity was judged to contribute negligibly to nucleus photometry (as no significant coma was detected) and was thus assumed to affect phase function derivations similarly negligibly. Having since accumulated a substantial set of observations while 133P was entirely inactive, though, we can now assess the validity of this neglect by deriving new phase function solutions and comparing the results to those of \\citet{hsi04}.  We caution that, unlike the data used by \\citet{hsi04}, the photometric data used in this follow-up analysis (2003 Sep 22 to 2007 Mar 21) all consist of ``snapshot observations'', which are short sequences of exposures at unknown rotational phases, instead of full lightcurves.  This caveat is significant because rotation of the body is expected to cause deviations in measured brightness by as much as 0.2~mag from the comet's true mean brightness at a given time \\citep{hsi04}.  Given a sufficiently large data set, however, we expect that the average of these fluctuations will approach zero, allowing us to derive reasonably accurate phase function solutions without necessarily knowing the rotational phase at which each individual photometry point was obtained. Nonetheless, the lack of rotational phase information for our snapshot observations remains a source of uncertainty.  We compute the reduced magnitude, $m_R(1,1,\\alpha)$, of 133P at the time of each observation using \\begin{equation} m_R(1,1,\\alpha) = m_{mid}(R,\\Delta,\\alpha) - 5\\log(R\\Delta) \\end{equation} where $R$ is the heliocentric distance of the object in AU, $\\Delta$ is the object's geocentric distance in AU, and $m_{mid}(R,\\Delta,\\alpha)$ is the estimated $R$-band magnitude at the midpoint of the full photometric range of the rotational lightcurve (Table~\\ref{obs_elstpiz}). For observations of full lightcurves, $m_{mid}$ is determined by simply plotting the data and locating the midpoint between the maximum and minimum values of the lightcurve.  For snapshot observations, $m_{mid}$ is generally taken to be the mean of the available photometry data with large error bars applied to reflect rotational phase uncertainties, assuming a full possible photometric range of 0.40~mag. We fit reduced magnitude values to both a linear phase function and an IAU phase function, finding best-fit values of $m_R(1,1,0)=15.80\\pm0.07$~mag and $\\beta=0.041\\pm0.005$~mag~deg$^{-1}$ for the linear phase function where \\begin{equation} m_R(1,1,\\alpha) = m_R(1,1,0) + \\beta\\alpha . \\end{equation} We also find best-fit values of $H_R=15.49\\pm0.05$~mag and $G_R=0.04\\pm0.05$ for the IAU phase function as defined in \\citet{bow89}. Photometry obtained at phase angles of $\\alpha<5^{\\circ}$, where an opposition surge effect is expected, are included in the derivation of the IAU phase function but omitted from the derivation of the linear phase function. We plot our best-fit solutions in Figure~\\ref{phaselaws}. A modest amount of scatter around our solutions is present, as expected, but in all cases the deviations from the best-fit phase functions are consistent with expected brightness fluctuations due to 133P's rotation. Due to the uncertainty of the active status of 133P on 2008 July 01 (\\S\\ref{monitoring}), the photometry from that night is plotted but was not included in the computation of the best-fit phase functions.  While the slope parameters of both newly-derived functions are consistent with the parameters computed in \\citet{hsi04} ($\\beta=0.044\\pm0.007$~mag~deg$^{-1}$; $G_R=0.026\\pm0.1$), both newly-derived absolute magnitudes are $\\sim$0.2~mag fainter than their previously derived values ($m_R(1,1,0)=15.61\\pm0.01$~mag; $H_R=15.3\\pm0.1$~mag), strongly suggesting that the previously derived parameters were affected by contamination by 133P's dust emission. This contamination is assumed to consist of a combination of coma and the portion of the dust trail (as projected in the plane of the sky) contained within the seeing disc. This suggestion of dust contamination is reinforced by Figure~\\ref{phaselaws} where we note that photometry from both 133P's 2002 and 2007 active phases are consistently brighter than expected from our new phase function solutions.  Because most of the data points from 2002 and 2007 are mean magnitudes derived from fully sampled lightcurves, brightness fluctuations due to rotation cannot account for the discrepancies.  Assuming that the discrepancy between an observed magnitude, $m_{mid}$, and expected magnitude, $m_{exp}$, is due to dust contamination, the scattering surface area of the dust, $A_d$, is given by \\begin{equation} A_d = A_n\\left({A_d\\over A_n}\\right) = A_n\\left(10^{0.4(m_{exp} - m_{mid})} - 1\\right) \\label{eqadust} \\end{equation} where $A_n=\\pi r_e^2=1.13\\times10^7$~m$^2$ is the scattering cross-section of the nucleus \\citep{hsi09}, and albedos of the nucleus and dust are assumed to be equal. Assuming optically thin dust, the total dust mass, $M_d$, can then be estimated from \\begin{equation} M_d \\sim {4\\over 3}\\pi\\rho a_d^3 \\left({A_d\\over\\pi a_d^2}\\right) \\label{eqmdust} \\end{equation} where we adopt typical dust grain radii of $a_d=10$~$\\mu$m and a bulk grain density of $\\rho_d=1300$~kg~m$^{-3}$ \\citep[{\\it cf}.][]{hsi04}.  For reference, we also compute $Af\\rho$ \\citep[{\\it cf}.][]{ahe84} for each set of observations where the parameter is given by \\begin{equation} Af\\rho = {(2R\\Delta)^2\\over \\rho} 10^{0.4[m_{\\odot}-m_R(R,\\Delta,0)]} \\label{eqafrho} \\end{equation} where $R$ is in AU, $\\Delta$ is in cm, $\\rho$ is the physical radius in cm of a $4\\farcs0$-radius photometry aperture at the distance of the comet, and $m_R(R,\\Delta,0)$ is the phase-angle-corrected $R$-band magnitude of the comet measured using a $4\\farcs0$-radius aperture, which we calculate using \\begin{equation} m_R(R,\\Delta,0) = m_{mid}(R,\\Delta,\\alpha) + 2.5\\log\\left[(1-G)\\Phi_1(\\alpha)+G\\Phi_2(\\alpha)\\right] \\end{equation} where $\\Phi_1$ and $\\Phi_2$ are given by \\begin{equation} \\Phi_1=\\exp\\left[-3.33\\left(\\tan{\\alpha\\over 2}\\right)^{0.63}\\right] \\end{equation} \\begin{equation} \\Phi_2=\\exp\\left[-1.87\\left(\\tan{\\alpha\\over 2}\\right)^{1.22}\\right] \\end{equation} \\citep{bow89}.  Using Equations~\\ref{eqadust}, \\ref{eqmdust}, and \\ref{eqafrho}, we compute $A_d$, $M_d$, and $Af\\rho$ for each set of observations from 2002 and 2007 during which 133P was observed to be active, and tabulate the results in Table~\\ref{dustcontrib}. We find that for data from 2002, dust contamination is approximately constant with a scattering surface area of $\\sim$0.20$A_n$ and a dust mass of $M_d\\sim4\\times10^4$~kg contained within $\\sim$3~arcsec ($\\sim$4500~km in August-November; $\\sim$6500~km in December) photometry apertures. The relatively constant amount of dust over this time period explains why we were able to derive reasonably accurate slope parameters for 133P from our 2002 data despite arriving at incorrect results for the comet's absolute magnitude due to the dust contamination.  In data from July 2007, we find that 133P's inferred dust coma has a strength comparable to that observed in 2002, having a scattering surface area equivalent to $\\sim$0.25$A_n$ and a dust mass of $M_d\\sim5\\times10^4$~kg contained within $\\sim$4~arcsec ($\\sim$4700~km) photometry apertures. The slightly larger amount of inferred dust in 2007 could indicate a higher rate of dust production, but could also be due to different viewing geometries, given that 133P was close to opposition when observed in July 2007.  At this position, the antisolar vector for 133P points very nearly directly behind the object as seen from Earth, causing more of the dust trail to be located within the seeing disc of the comet as projected on the sky. Given the limitations of our observations, however, we are unable to disentangle this possible projection effect from any intrinsic increase in dust production. Additionally, 133P's apparent brightness could also have been enhanced by an opposition surge effect from the dust in its coma, though we unfortunately lack observational constraints for quantifying this effect.  Given these various possible contributing factors to 133P's enhanced brightness on 2007 July 17 and 20, we are unable to determine whether 133P was more active on these dates compared to 2002 August 19 through 2002 November 7.  We can conclude, however, that coma contamination is present in nucleus photometry performed for 133P during both observing periods, and that the measured magnitude enhancements suggest at least comparable levels of activity in each case.  The remainder of our photometry from 133P's 2007 active phase is derived from incomplete lightcurve information, and as such, coma estimates at these times have much larger uncertainties than at other times. We find no definitive evidence of a coma on 2007 May 19, but find that the inferred comae on 2007 Aug 18 and 2007 Sep 12 are far stronger ($\\sim$0.65$A_n$) than in any other observations, a rather unexpected discovery given the minimal amount of time elapsed since our 2007 July observations.  We suggest that the large inferred dust contribution to nucleus photometry in August and September could be at least partly due to geometric effects.  As can be seen in Figures~\\ref{images_133p}q-\\ref{images_133p}s, the orientation of the projection of the dust trail appears to change over this period of time. We caution that poor seeing during our August and September observations and the small aperture (1.3~m) of the telescope used to obtain these data mean that the observed morphology (namely, the near-disappearance of the dust trail) cannot be considered entirely reliable. If the observed morphology is believed, however, much of the precipitous increase in 133P's apparent coma strength between July and August could be due to the dust trail becoming almost directly aligned behind the nucleus in August and September, thus becoming unavoidably included within our photometry apertures.  We can account for this viewing geometry effect by integrating the scattering surface area of the visible dust trail measured in July data (discussed below; \\S\\ref{dusttrail}) and then assuming that it all falls within the photometry aperture used to measure the nucleus magnitudes in August and September. The net increase in dust scattering surface area implied by photometry between 2007 July 20 and 2007 August 18 is $\\sim$0.40$A_n$. The integrated scattering surface area of the dust trail on 2007 July 20 over the first 30~arcsec from the nucleus (the trail becomes too faint to measure reliably beyond this point), however, is $\\sim$0.20$A_n$, accounting for only about half of the observed increase in dust contamination between July and August.  The remainder of the observed increase could be partly due to distant material in the dust trail that was too diffuse to detect in trail form in July data, but nevertheless contributed positively to nucleus photometry when projected directly behind the nucleus in August and September. It seems unlikely, however, that half of the dust in the trail could go undetected in our July data, and as such, we surmise that at least part of the increase must in fact be due to a real increase in dust production, which of course would certainly be plausible at this early stage in 133P's active phase.   \\subsection{The Lightcurve Revisited\\label{ltcurverevisit}}  \\subsubsection{Search for Rotational Colour Variations\\label{colorvariations}}  During our 2007 NTT run when 133P was active, we observed the comet in continuously cycling filters ($VRI$ on 2007 July 17 and $BVRI$ on 2007 July 20).  Observations were made in this way to allow us to obtain deep imaging of 133P in multiple filters and also construct simultaneous lightcurves in each filter. These lightcurves then allowed us to search for surface colour inhomogeneities that, for example, may constrain the position of the localised active site hypothesized by \\citet{hsi04}. These lightcurves, phased to a rotational period of $P_{rot}=3.471$~hr \\citep{hsi04}, are plotted in Figure~\\ref{nttltcurves}. To then assess colour variation as a function of rotational phase for each filter pair, we use linear interpolation to obtain the magnitudes of the object in the second filter at times of observations in the first filter and then plot the differences (Fig.~\\ref{nttcolors}), again phased to $P_{rot}=3.471$~hr.  We find mean nucleus colours of $B-V=0.65\\pm0.03$~mag, $V-R=0.36\\pm0.01$~mag, and $R-I=0.32\\pm0.01$~mag.  These values are somewhat different from the mean colours found for 133P by \\citet{hsi04}, but are within the range of individual values measured in that work.  We regard the colour measurements presented here to be more accurate since our repeated multifilter observations of 133P here allowed us to account for both rotational magnitude variations (via lightcurve interpolation) and minor extinction variability (using field stars as references for making differential photometric corrections). The single sets of multifilter observations used to make 133P's previous colour measurements did not permit either of these corrective measures.  Upon examining individual colour measurements, we find no conclusive evidence of rotational colour inhomogeneity.  We find maximum colour variations of only $\\Delta(B-V)=0.11\\pm0.13$~mag, $\\Delta(V-R)=0.06\\pm0.07$~mag, and $\\Delta(R-I)=0.08\\pm0.08$~mag, where the non-systematic distribution of even these small variations indicates that they are most likely due to ordinary measurement uncertainties.  We note that this result does not rule out the possibility that 133P's active area exhibits a different colour signature than inactive surface material.  First, the coma that is likely present (\\S\\ref{activitydetection}) should act to obscure colour variations on the nucleus surface, with the precise amount of obscuration varying with rotational phase as the ratio of the nucleus's scattering cross-section to the coma's cross-section changes.  Furthermore, under the seasonal heating hypothesis \\citep{hsi04,hsi06a}, the active site is in fact expected to be illuminated by the Sun at all rotational phases when near perihelion (assumed to be close to solstice) when these observations were made. The nucleus orientation at this time allows the active site to receive maximal solar heating but also means that the active site is always in the line of sight as viewed from Earth. We suggest that more favourable conditions for detecting colour inhomogeneities will occur around 133P's next pre-perihelion equinox (likely near $\\nu\\sim270^{\\circ}$).  Based on prior observations (\\S\\ref{monitoring}), the nucleus should be largely coma-free over this portion of the orbit, and at equinox, the active site should pass into and out of the line of sight as the nucleus rotates, maximising any colour variations.  We therefore encourage additional rotationally-resolved colour measurements of 133P between late 2011 and early 2012.   \\subsubsection{Implications for 133P's Pole Orientation}  For reference, we remove the estimated dust contamination from both our 2002 and 2007 lightcurve data, and overplot the two sets of lightcurves (Fig.~\\ref{ltcurves0207}). Each of the two sets of data are phased self-consistently to $P_{rot}=3.471$~hr, though given the great difficulty of phasing data together that are separated by almost 5 years to such a short rotational period, the 2002 and 2007 data are simply aligned by eye.  Due to the two-peaked nature of 133P's lightcurve, though, there is an ambiguity in performing this alignment.  In one case (Fig.~\\ref{ltcurves0207}a), the data can be aligned such that the lightcurve shape and photometric range appear largely unchanged between the two observation epochs. In the second case (Fig.~\\ref{ltcurves0207}b), the data can be aligned such that the photometric range of lightcurve appears to decline to $\\Delta m_R\\sim0.25$~mag in 2007 from $\\Delta m_R\\sim0.35$~mag in 2002.  In the latter case, it should be recalled that the coma contribution to the data plotted has already been subtracted, and as such the change in photometric range cannot be attributed to differences in the amount of coma. Unfortunately, due to the incomplete sampling of the lightcurve in 2007, it is not possible to resolve the ambiguity between these two cases.  This ambiguity is significant because of the implications of photometric range behaviour for the orientation of 133P's rotational pole. To gain more insight as to how the photometric range of 133P should change depending on pole orientation and observing geometry, we simulate its lightcurve behaviour using the model presented in \\citet{lac07}. We assume a simple prolate ellipsoidal shape for the nucleus of 133P and render it at various observing geometries and rotational phases.  At each rotational phase, the light reflected back to the observer is integrated to generate lightcurve points. The 2002 September coma-corrected photometric range for 133P was measured to be $\\Delta m_R=0.35$~mag, and so we use a nucleus axis ratio of $a/b=10^{0.4\\Delta m_R}=1.39$ (it should be noted that this is a lower limit due to the unknown projection angle at the time). We use a Lommel-Seeliger ``lunar'' scattering function \\citep[{\\it cf}.][]{fai05} which has no free parameters and is appropriate for simulating the low albedo \\citep[$p_R=0.05\\pm0.02$;][]{hsi09} surface of 133P.  To simplify the geometry, we neglect the small orbital inclination ($i=1.4\\degr$) of 133P and assume that it is coplanar with the Earth (i.e., $i=0\\degr$).  The seasonal heating hypothesis implies that 133P is at solstice when close to perihelion, i.e. has a true anomaly at solstice of $\\nu_\\mathrm{sol}\\approx0\\degr$, and also requires that the object have non-zero obliquity ($\\varepsilon\\neq0\\degr$). In principle, $\\nu_\\mathrm{sol}$ could potentially have any value from $\\nu_\\mathrm{sol}\\approx0\\degr$ to $\\nu_\\mathrm{sol}\\approx45\\degr$, since the temperature of the hemisphere where 133P's active site is located will begin to rise due to solar heating before the spin axis direction is actually aligned with the Sun.  The seasonal heating hypothesis is inconsistent, however, for pole orientations for which $\\nu_\\mathrm{sol}\\approx90\\degr$, or $\\varepsilon=0\\degr$. We simulate the lightcurve behaviour of 133P for $\\nu_\\mathrm{sol}=0\\degr$, $\\nu_\\mathrm{sol}=40\\degr$ and $\\nu_\\mathrm{sol}=90\\degr$. The first and third pole orientations are limiting cases that are consistent and inconsistent with the seasonal heating hypothesis, respectively. The intermediate geometry, in which solstice is reached approximately half-way through the active portion of the orbit, is meant to test how sensitive we are to the exact longitude of the pole.  Because we assume zero orbital inclination, each case sets the ecliptic longitude of pole, and the ecliptic latitude is defined by the choice of obliquity. We simulate obliquities of $\\varepsilon=0\\degr$, $\\varepsilon=10\\degr$, $\\varepsilon=20\\degr$ and $\\varepsilon=30\\degr$. Only $\\varepsilon=0\\degr$ is inconsistent with the seasonal hypothesis. Rendered samples of 133P, where we assume $\\varepsilon=30\\degr$, are shown in Figures~\\ref{ltcurve_nu0}, \\ref{ltcurve_nu40} and \\ref{ltcurve_nu90}.  Figure~\\ref{rangevsgeom} shows the expected photometric range in 2002 September and 2007 July for each pole orientation. As expected, the photometric range changes in opposite directions for $\\nu_\\mathrm{sol}=0\\degr$ and $\\nu_\\mathrm{sol}=90\\degr$, whereas the intermediate pole orientation ($\\nu_\\mathrm{sol}=40\\degr$) produces only a small change between the two epochs. The absolute value of $\\Delta m_R$ in the figure depends on the assumed axis ratio and is unimportant in this analysis, in which we are primarily concerned with relative changes. The key feature is the variation of the range between the two epochs.  The two possible scenarios indicated by the data ({\\it cf}. Fig.~\\ref{ltcurves0207}) are where (a) both the 2002 and 2007 photometric ranges are similar ($\\Delta m_R\\sim0.35$ mag), or (b) the 2002 photometric range ($\\Delta m_R\\sim 0.35$ mag) is larger than the 2007 range ($\\Delta m_R\\sim0.25$ mag). Inspection of Figure~\\ref{rangevsgeom} shows that the first scenario is consistent with low obliquity ($\\varepsilon\\lesssim10\\degr$) and any of the considered pole orientations. The second scenario is only consistent with a solstice around $\\nu_\\mathrm{sol}=0\\degr$ and significant obliquity ($\\varepsilon\\gtrsim20\\degr$). Both scenarios rule out a pole orientation where $\\nu_\\mathrm{sol}=90\\degr$ if there is also significant obliquity.  Clearly, additional and more complete lightcurve observations at different points in 133P's orbit are needed to clarify how 133P's photometric range varies with orbit position, constrain the object's pole orientation, and determine whether the seasonal heating hypothesis remains plausible. Given our current data, we can neither confirm nor reject the plausibility of seasonal activity modulation as described by \\citet{hsi04}.  While the pattern of activity of 133P along its orbit appears consistent with the seasonal heating hypothesis, the discovery of an incompatible pole solution could indicate that activity is in fact modulated by factors other than obliquity, {\\it e.g.}, shadowing of the active site by crater walls or other local topographic features. In Figure~\\ref{rangevsorbit}, we use our model to forecast the photometric range behaviour of 133P over 1.5 orbits from its perihelion passage in 2007 August and to its aphelion passage in 2016 January. We plot solutions for four pole positions, two consistent with the seasonal heating hypothesis ( $\\varepsilon=20\\degr$, $\\nu_\\mathrm{sol}=0\\degr$, and $\\varepsilon=20\\degr$, $\\nu_\\mathrm{sol}=40\\degr$) and two inconsistent with that hypothesis ($\\varepsilon=20\\degr$, $\\nu_\\mathrm{sol}=90\\degr$, and $\\varepsilon=0\\degr$). The observability of 133P during this period is also indicated in the figure, and should assist in planning observations that are best-suited for discriminating between the various pole orientations that we consider here.   \\subsection{The Dust Trail Revisited\\label{dusttrail}}  To produce deep composite images from our 2007 NTT data, we use linear interpolation to shift the multiple images obtained in each filter to align the photocenters of the nucleus in each image, and sum the resulting shifted images. To measure the surface brightness profiles of the dust trail in these composite images, we then rotate the images to make the trail horizontal in the image frames, and measure the net flux in rectangular apertures placed along the length of the trail \\citep[{\\it cf}.][]{hsi04}. The dimensions of these equally-sized apertures are set to lengths (along the direction of the trail) of 5 pixels, and widths (perpendicular to the trail) of 6 pixels (approximately equal to the FWHM of the trail cross-section on each night).  The net fluxes in these apertures are then converted to net fluxes per linear arcsec (as measured along the length of the trail) and normalised with respect to the total net flux of the nucleus.  We plot the resulting surface brightness profiles for both 2007 Jul 17 and 2007 Jul 20 in Figure~\\ref{trailcolors07}.  From these plots, we see that the trail profile does not change significantly between the two nights.  We also note that there are minimal differences in the profiles of the trail as observed in different filters, indicating that the colours of the dust along the trail are consistently similar to those of the nucleus.  To quantify this observation, we measure the surface brightness of the trail as observed on 2007 July 20 in each filter in a single aperture approximately 5 arcsec (15 pixels, or $\\sim$6000~km) in length and 1 arcsec (3 pixels or $\\sim$1200~km) in width placed along the trail. Seeking to minimise the effect of the nucleus on our trail photometry, we place the nearest edge of this aperture $\\sim$3.0~arcsec from the nucleus photocentre. We find surface brightnesses of $\\Sigma_B=24.88\\pm0.17$~mag~arcsec$^{-2}$, $\\Sigma_V=24.35\\pm0.05$~mag~arcsec$^{-2}$, $\\Sigma_R=24.04\\pm0.05$~mag~arcsec$^{-2}$, and $\\Sigma_I=23.70\\pm0.07$~mag~arcsec$^{-2}$, giving colours of $B-V=0.53\\pm0.18$~mag~arcsec$^{-2}$, $V-R=0.31\\pm0.07$~mag~arcsec$^{-2}$, and $R-I=0.34\\pm0.09$~mag~arcsec$^{-2}$, consistent with the colours of the nucleus found in \\S\\ref{ltcurverevisit}.  We also wish to know how trail morphology changes between 133P's 2002 and 2007 active episodes.  The most obvious difference between the two observing epochs is that the dust trail of 133P is significantly shorter in our 2007 data than in 2002 (despite composite images from each epoch being of approximately equivalent effective exposure time), extending only $\\sim$30 arcsec from the nucleus in 2007 observations, compared to nearly 3 arcmin in 2002 \\citep{hsi04}. In terms of trail width, the observed mean FWHM of the trail on 2002 Sep 07 over the first 10~arcsec of the trail, as measured from the edge of the nucleus's seeing disc (taken to be 2.5$\\times$ the FWHM seeing), was measured to be $\\theta_o=1\\farcs3$.  This observed value corresponds to an intrinsic FWHM of $\\theta_i=0\\farcs9$ ($\\sim$1300~km in the plane of the sky), which is computed using \\begin{equation} \\theta_i = (\\theta_o^2 - \\theta_s^2)^{1/2} \\label{intrinsicwidth} \\end{equation} where the FWHM seeing was $\\theta_s=0\\farcs9$ on 2002 Sep 07. For comparison, the observed FWHM of the trail on 2007 July 17 was $\\theta_o=1\\farcs9$, corresponding to $\\theta_i=1\\farcs3$ ($\\sim$1500~km in the plane of the sky), where $\\theta_s=1\\farcs4$. Given that viewing geometries (parametrized by out-of-plane viewing angles, $\\alpha_{pl}$) in 2002 and 2007 were comparable, we therefore find that the computed intrinsic width of the dust trail is approximately equal in both our 2002 and 2007 observations. As \\citet{hsi04} found the primary factor controlling 133P's trail width to be particle ejection velocity, this result suggests that sublimation took place with comparable intensity in both 2002 and 2007.  In order to further compare 133P's activity level in 2002 and 2007, we measure the profile of 133P's trail in $R$-band data from 2002 September 07 using the procedure described above, i.e., using rectangular apertures placed along the length of the trail with lengths of 5 pixels and widths of 6 pixels each.  We then compare the resulting profile to the mean $R$-band trail profile from 2007 (Fig.~\\ref{trailcomparison}), finding that the trail is noticeably weaker in 2007 than it was in 2002.  The difference in trail strength in 2002 and 2007 could be due to several factors.  The simplest explanation is that the activity was actually weaker in 2007 due to depletion of exposed volatile material on 133P by the previous outburst.  This explanation, however, is at odds with our findings of comparable dust ejection velocities for the two observing epochs (above), and comparable dust enhancement of the nucleus brightness (\\S\\ref{activitydetection}).  A more likely explanation is that by the time our 2007 NTT observations were made, 133P was no more than 4 months into its current active phase, whereas it had been active for about a year by the time it was observed on 2002 September 07.  Thus, 133P may simply have not yet reached its peak level of activity by the time we observed it with the NTT in 2007. The position of 133P near opposition on 2007 July 17 and 20 also meant that a dust trail pointed in the antisolar direction would be highly projected in the sky, which would additionally explain why the trail appeared to be so much shorter in 2007 than in 2002."
    },
    "0911/0911.0131_arXiv.txt": {
        "abstract": "We propose here a robust scheme to infer the physical parameters of compact stars from their f-mode gravitational wave signals. We first show that the frequency and the damping rate of f-mode oscillation of compact stars can be expressed in terms of universal functions of stellar mass and moment of inertia. By employing the universality in the f-mode one can then infer accurate values of the mass, the moment of inertia and the radius of a compact star. In addition, we demonstrate that our new scheme works well for both realistic neutron stars and quark stars, and hence provides a unifying way to infer the physical parameters of compact stars. ",
        "introduction": "Introduction} Ever since the first prediction of neutron stars (NS) by \\citet{Baade1934a}, NS have become a major test bed of various theories of dense matter and nuclear matter. The extreme high pressure and density achievable inside NS offers an ideal and unique environment to infer the equation of state (EOS) of nuclear matter from astronomical observations. On the other hand, the possible existence of quark stars (QS) would provide a strong support to a new form of matter, namely, the quark matter \\citep{Itoh70, Bodmer71, Witten}. Moreover, the size and the mass of QS could also cast light on the properties of quark matter. Thus, extracting useful information and constraints about nuclear and particle physics from astronomical observations of compact stars like NS and QS has become  an active direction in astrophysics \\citep[see, e.g.,][for reviews] {lattimer2007nso,haensel2007nse}.  Since the density in the core of a NS (or QS) can be several times the normal nuclear density, the EOS there is rather uncertain. Various physical phenomena including pion condensation, hyperon formation and deconfined quark matter have been considered by theoretical nuclear physicists and sophisticated techniques of many-body calculations have been developed to yield a bunch of EOS intended for the description of the deep interior of compact stars. It is standard practice for astrophysicists to construct compact stars with these EOS for nuclear matter and establish various relationships among different  physical quantities which could be obtained from astronomical observations. Motivated by the possibilities of constraining EOS of nuclear matter and extracting physical parameters of compact stars as well, researchers in astrophysics and nuclear physics have been actively seeking for EOS-dependent and EOS-independent relationships  by examining the physical characteristics of compact stars constructed with different EOS \\citep[see, e.g.,][]{Lattimer:2001,Bejger:2002p8392,Lattimer:2005p7082, bejger2005cdm,lattimer2007nso,haensel2007nse}. For example, \\citet{Lattimer:2001} compared the structure of compact stars of various kinds and discovered several empirical relationships connecting different physical characteristics of a star, such as the mass, the radius, the moment of inertia, and the mass distribution function as well. Such relationships can be applied to infer the physical attributes of a compact star, including its EOS, from astronomical observations.   The main objective of the present paper is to unveil the universality embedded in the pulsation frequencies of non-rotating compact stars and to exploit these findings to infer the physical characteristics of NS or QS (e.g., mass, radius, moment of inertia and EOS) from their oscillation spectra. The pulsations of compact stars are an interesting and popular topic in its own right, for gravitational waves (GW) can be generated in such processes \\citep[see][for a seminal exposition of the topic]{Thorne}. It is indisputable that the detection of GW would be a major milestone in general relativity and astrophysics. Several Earth-based GW interferometric detectors such as LIGO, VIRGO, GEO600 and TAMA300 have been operating. While the current detectors are still not sensitive enough to detect GW directly, interesting upper limits have been placed on either the GW strains or event rate for several potential astrophysical sources \\citep[see, e.g., ][for the results obtained from the latest science runs of LIGO]{LIGO_S4_BHringdown,LIGO_S5_LowMass, LIGO_S5_pulsars}.  Owing to the energy carried away by GW, pulsations of compact stars are damped harmonic oscillations, which are analyzed in terms of quasi-normal modes (QNM) \\citep[see, e.g.][]{Press_1971,Leaver_1986,rmp,Kokkotas_rev}. Each QNM is characterized by a complex eigenfrequency $\\omega=\\omega_\\r+ i\\omega_\\i$ and has a time dependence $\\exp(i\\omega t)$, displaying exponential decay with a damping time $\\tau\\equiv 1/\\omega_\\i$. Several attempts have been made to relate $\\omega_\\r$ and $\\omega_\\i$ (or $\\tau$) to the mass $M$ and the radius $R$ of a compact star and some universal relationships which are, to certain degree of accuracy, EOS-independent have been found \\citep{Andersson1996,Andersson1998,Ferrari,Ferrari_prd,Tsui05:1}. However, most of  these relationships fail to describe QNM of QS, which are stiff and self-bound.  On the other hand, \\citet{Lattimer:2005p7082} found an empirical relation approximately expressing the moment of inertia in terms of the mass and the radius. Such a relationship is universal as long as the EOS under consideration is not too soft (e.g., hyperon matter) or too stiff (e.g., quark matter). They proposed to use the universality to estimate the radius of a compact star from its mass and moment of inertia, with the latter two quantities being measurable for double pulsar systems by considering the spin-orbit coupling effect in general relativity \\citep{Lattimer:2005p7082}.    Motivated by the discovery of \\citet{Lattimer:2005p7082}, we propose here to investigate the correspondence between the QNM frequency (both $\\omega_\\r$ and $\\omega_\\i$) of the (fundamental) f-mode oscillations, $M$ and the moment of inertia $I$ of compact stars. We show that there exist EOS-independent universal relationships between them, namely equations (\\ref{lau_r}) and (\\ref{lau_i}), which are accurate up to a few percent or better and work nicely for QS as well. Based on equations (\\ref{lau_r}) and (\\ref{lau_i}), a feasible scheme is in turn established to determine the mass and the moment of inertia of a compact star once its f-mode frequency is found from GW observations. As the radius $R$ of a NS (or QS) is approximately expressible in terms of the mass and the moment of inertia, our scheme can also yield good estimate of the radius and hence imposing constraints on the EOS \\citep{Bejger:2002p8392,Lattimer:2005p7082}.  The outline of the paper is as follows. Section 2 is a brief review on the the behavior of f-mode oscillations of compact stars \\citep{Tsui05:1} and the universal relationships among the mass, the radius and the moment of inertia of compact stars \\citep{Bejger:2002p8392,Lattimer:2005p7082}.  In Section 3  we introduce an accurate universal relation between the scaled frequency $M \\omega$ and the physical quantity $\\sqrt{M^3/I}$.  In Section 4 we establish a feasible and robust scheme to apply our findings reported in Section 3 to invert the mass, the moment of inertia, and in turn the radius of a compact star if its f-mode GW signal is detected. We finally summarize our paper in Section 5 with a brief discussion. Unless otherwise noted, we adopt geometric units with $G=c=1$ throughout the whole paper, and use kilometers as the unit of lengths. ",
        "conclusions": "Conclusion and discussions} In this paper we study the universality embedded in f-mode pulsations of compact stars and propose to use it to determine the physical parameters (including mass, radius and moment of inertia) of a compact star from which the f-mode gravitational wave signal is detected. We first establish a pair of empirical equations (\\ref{lau_r}) and (\\ref{lau_i}) that can predict the frequency and the damping rate of f-modes from $M$ and $I$ with good accuracy for both NS and QS. We then apply such discovery to develop an inversion scheme to infer the mass, radius and moment of inertia of a compact star from its f-mode. In particular, the scheme is shown to work well for both NS and QS.    The universality discovered in the present paper takes moment of inertia into the consideration, replacing the role of radius used in previous attempts \\citep{Andersson1996,Andersson1998,Ferrari,Ferrari_prd,Tsui05:1}. It is worth noting that the radius of a star is sensitive to the low-density part of its EOS, whereas f-mode frequencies are expected to be dominated by dynamical behavior in the high-density regime. It is therefore physically appealing to use moment of inertia, which measures global mass distribution, to study f-mode oscillations. As clearly shown in Table~\\ref{uni_quality}, the replacement of $R$ with $I$ indeed leads to a much improved universal behavior. As it is arguable that the accuracy of the inversion scheme directly reflects the quality of the universal behavior upon which the scheme is based, the universality discovered in the present paper leads to an accurate inversion scheme.  Finally, it is worthy of mentioning that the first gravitational-wave search sensitive to the f-modes has recently been carried out by the LIGO detectors \\citep{abbott-2008sgr}. While the current detectors have not yet detected gravitational waves directly, it is plausible that the advanced LIGO detectors, which have a factor of 10 improvement on the sensitivity, might detect the f-mode gravitational wave signals from compact stars. The inversion scheme proposed in this work will provide a unifying way to infer the mass, radius and moment of inertia of both NS and QS accurately. Of course in reality the situation is often complicated by the presence of magnetic field and rotation, in the future we shall consider the impacts of these factors on our proposal."
    },
    "0911/0911.0313_arXiv.txt": {
        "abstract": "{In this paper we present the first system test in which we demonstrate the concept of using an array of Distributed Read Out Imaging Devices (DROIDs) for optical photon detection.} {After the successful S-Cam 3 detector the next step in the development of a cryogenic optical photon counting imaging spectrometer under the S-Cam project is to increase the field of view using DROIDs. With this modification the field of view of the camera has been increased by a factor of 5 in area, while keeping the number of readout channels the same.} {The test has been performed using the flexible S-Cam 3 system and exchanging the 10x12 Superconducting Tunnel Junction array for a 3x20 DROID array. The extra data reduction needed with DROIDs is performed offline.} {We show that, although the responsivity (number of tunnelled quasiparticles per unit of absorbed photon energy, e-/eV) of the current array is too low for direct astronomical applications, the imaging quality is already good enough for pattern detection, and will improve further with increasing responsivity.} {The obtained knowledge can be used to optimise the system for the use of DROIDs.} ",
        "introduction": "With the S-Cam project the Advanced Studies \\& Technology Preparation Division of the European Space Agency is developing a series of prototype cryogenic detectors to be used as optical photon counting imaging spectrometers for ground-based astronomy. S-Cam uses Superconducting Tunnel Junctions (STJs) (\\cite{Friedrich:2006}; \\cite{Prober:2006}; \\cite{Peacock:1996}) as its detector technology. The merit of this and other cryogenic detectors (\\cite{Romani:1999}) is that they combine single photon detection with sub-microsecond time resolution and intrinsic wavelength resolution, imaging and good detection efficiency in a single device.  STJs consist of 2 superconducting layers separated by a thin insulating layer acting as a tunnel barrier. With the absorption of a photon in the superconducting layer a large quantity (several thousands) of Cooper pairs are broken into quasiparticles which can tunnel across the barrier and, under the influence of an applied bias voltage, produce a measurable current pulse. The number of created quasiparticles is given by: $N(E_{0})=\\frac{E_{0}}{\\varepsilon}$, with $N(E_{0})$ the number of created quasiparticles, $E_{0}$ the energy of the absorbed photon and $\\varepsilon=1.75\\Delta_{g}$ the mean energy needed to create a quasiparticle (\\cite{Kurakado:1981}) with $\\Delta_{g}$ the energy gap of the superconducting material. As shown the number of created quasiparticles, and hence the amplitude of corresponding tunnel current, is proportional to the energy of the absorbed photon, thus providing the detector with its spectrographic capabilities. The theoretical limit for the intrinsic energy resolution is given by: $\\Delta E= 2.355\\sqrt{\\varepsilon E_{0} (F+G)}$, where $F$ is the Fano factor (\\cite{Fano:1947}), equal to $F=0.2$ (\\cite{Kurakado:1981}; \\cite{Rando:1991}), and $G=1+\\frac{1}{<n>}$ (\\cite{Mears:1993}) accounts for the statistical variations in the tunnel process, with $<n>$ the average number of tunnels of a single quasiparticle. The energy gap, $\\Delta_{g}$, of the superconducting material is proportional to its critical temperature, $T_{c}$), the temperature at which the phase changes from superconducting to normal metal. For a BCS type superconductor (usually an elemental superconducting material which follows the theory developed by \\cite{Bardeen:1957}), $\\Delta_{g}=1.764k_{b}T_{c}$. A lower energy gap of the superconducting material will therefore increase the number of created quasiparticles and provide better spectrographic capabilities, but it also puts increasing constraints on the operating temperature ($T_{op}$). This needs to be well below the critical temperature of the superconducting layer ($T_{op}\\approx0.1T_{c}$) in order to sufficiently reduce the thermally excited quasiparticle population. For a more  extended overview of the STJ technology the reader is referred to \\cite{Peacock:1996}.  Each STJ needs to be read out using a dedicated electronics chain which limits the maximum number of pixels that can be read out in a practical application (\\cite{Martin:2006}). To overcome this limitation the Distributed Read Out Imaging Device (DROID) (\\cite{Kraus:1989}) is being developed. A DROID consists of a superconducting absorber strip with STJs on either end, see fig. \\ref{fig:Figure 1}. The photon is absorbed in the absorber and the created quasiparticles diffuse towards the STJs where they tunnel. The sum of the tunnel signals of both STJs is a measure for the energy of the absorbed photon while the ratio is a measure for the absorption position. Depending on the position resolution of the DROID it can replace a number of single STJs and reduce the number of read out channels for a given sensitive area (\\cite{Hijmering:2008}).  Within the S-Cam project three prototype cameras have already successfully been used on telescopes such as the William Herschel Telescope (WHT) on La Palma and the Optical Ground Station (OGS) on Tenerife (\\cite{Martin:2004}). S-Cam 1 (\\cite{Verhoeve:2002}) and 2 (\\cite{Rando:2000}) were based on a $6\\times6$ pixel array ($25\\times25 \\mu m^{2}$ pixels) with a wavelength resolving power of 6. S-Cam 3 (\\cite{Martin:2007}; \\cite{Martin:2006}) was based on a $10\\times12$ pixel array ($35\\times35 \\mu m^{2}$ pixels), increasing the field of view on the WHT from $4\\arcsec\\times4\\arcsec$ to $10\\arcsec\\times12\\arcsec$. Also the covered wavelength range, operating temperature and resolving power ($\\sim14@500nm$) has been enhanced with S-Cam 3. The applicability of this type of detector has been proven in different observation campaigns in which several types of astronomical objects have been observed. The high time resolution spectrally resolved S-Cam data has provided strong constraints on the geometry of eclipsing binaries (\\cite{Perryman:2001}; \\cite{de Bruijne1:2002}; \\cite{Martin:2003}). Precise timing of the Crab pulsar light curve has shown that the optical pulses lead the radio pulses by $273\\pm65\\mu s$ (\\cite{Perryman:1999}; \\cite{Oosterbroek:2006}). The spectral information provided by the STJs has enabled the direct determination of quasar redshifts (\\cite{de Bruijne2:2002}) and stellar temperatures (\\cite{Reynolds:2003}). The next step is to increase the field of view further with the use of DROIDs. Here we present the results of the first system test using a $3\\times20$ DROID array as a detector. ",
        "conclusions": "We have successfully demonstrated operation of an array of DROIDs as a photon counting and imaging detector in an astronomical instrument. Although the responsivity of the array was too low for practical use, the resolving power and imaging capabilities, which will improve further with increasing responsivity, are already adequate. From this first system test we have obtained a good understanding on how to further optimize the system for photon detection with DROIDs."
    },
    "0911/0911.2316_arXiv.txt": {
        "abstract": "We have resimulated the six galaxy-sized haloes of the Aquarius Project including metal-dependent cooling, star formation and supernova feedback. This allows us to study not only how dark matter haloes respond to galaxy formation, but also how this response is affected by details of halo assembly history. In agreement with previous work, we find baryon condensation to lead to increased dark matter concentration. Dark matter density profiles differ substantially in shape from halo to halo when baryons are included, but in all cases the velocity dispersion decreases monotonically with radius. Some haloes show an approximately constant dark matter velocity anisotropy with $ \\beta \\approx 0.1-02$, while others retain the anisotropy structure of their baryon-free versions. Most of our haloes become approximately oblate in their inner regions, although a few retain the shape of their dissipationless counterparts.  Pseudo-phase-space densities are described by a power law in radius of altered slope when baryons are included. The shape and concentration of the dark matter density profiles are not well reproduced by published adiabatic contraction models.  The significant spread we find in the density and kinematic structure of our haloes appears related to differences in their formation histories.  Such differences already affect the final structure in baryon-free simulations, but they are reinforced by the inclusion of baryons, and new features are produced. The details of galaxy formation need to be better understood before the inner dark matter structure of galaxies can be used to constrain cosmological models or the nature of dark matter. ",
        "introduction": "Dissipationless cosmological simulations have contributed substantially to our understanding of the structure and evolution of cold dark matter (CDM) haloes, showing them to have triaxial shapes (Frenk et al. 1998; Dubinski \\& Carlberg 1991; Jing \\& Suto 2002; Hayashi, Navarro \\& Springel 2007) and density profiles with inner cusps (Dubinski \\& Carlberg 1991, Navarro, Frenk \\& White 1996; Moore et al. 1999; Diemand et al. 2005) and an approximately universal shape, independent of mass or cosmological parameters (Navarro, Frenk \\& White 1997).  This universal cuspy profile appears in conflict with several observations, for example, the slow increase of rotation velocity with radius at the centre of low surface brightness galaxies (Flores \\& Primack 1994; McGaugh \\& de Blok 1998; Moore et al 1999; Salucci, Yegorova \\& Drory 2008) and the relatively weak concentration of galaxy clusters inferred from strong lensing (Sand et al. 2008). Additional evidence supporting central dark matter densities lower than predicted by CDM models comes from the difficulty in simultaneously matching observed luminosity functions and the zero-point of the Tully-Fisher relation (e.g.  Dutton, van den Bosch \\& Courteau 2008).  Recently, Navarro et al. (2008, hereafter N08) analysed in detail the density profiles of very high-resolution dark matter-only simulations of six Milky Way mass haloes from the Aquarius Project (Springel et al. 2008). They found the innermost cusps in these haloes to be weaker than claimed in some earlier work, and showed how halo-to-halo variations in radial structure reflect the detailed formation histories of individual haloes (see also Vogelsberger et al. 2009).  Galaxy formation might reinforce such history-specific features in the dark matter distribution, since baryons are subject to dissipative processes such as cooling, star formation and feedback in addition to gravity.  The condensation of baryons within dark matter haloes modifies both their dynamics and their structure, but exactly how this happens is still quite uncertain. The simplest model assumes that the dark halo is compressed radially and adiabatically by the added mass at its centre (Blumenthal et al. 1986, hereafter B86; see Eggen et al. 1962, Zeldovich et al. 1980 and Barnes \\& White 1984 for previous applications of this formalism). Recent simulations have found the very simple adiabatic compression (AC) scheme of B86 to overestimate the effect (e.g. Gnedin et al. 2004 and Abadi et al. 2009, hereafter G04 and A09, respectively; Pedrosa, Tissera \\&   Scannapieco 2010) but it is nevertheless often used in simplified modelling of galaxy formation (e.g. Mo, Mao \\& White 1998).  More sophisticated AC models, based, for example, on the formalism of Young (1980) have claimed to reproduce better the contraction and the final shape of haloes (G04; Sellwood \\& McGaugh 2005). The common outcome of these schemes is nevertheless an increase in the dark matter density in the central regions, exacerbating the problem of reconciling $\\Lambda$CDM models with observation.    Simulations of galaxy formation which follow both dark matter and baryons in their proper cosmological context are the primary tool for studying how galaxy assembly affects dark matter haloes.  There are many papers devoted to this subject. Improving numerical algorithms and increasing computer power have led to continual progress in understanding the complex interplay between these two components.  Previous analyses of the evolution of dark haloes in hydrodynamical simulations have reported an increase both of the central mass concentration and of the central velocity dispersion, which no longer shows the 'temperature inversion' characteristic of dissipationless CDM haloes (e.g. Katz \\& Gunn 1991; Evrard, Summers \\& Davis 1994; Navarro \\& White 1994; Tissera \\& Dominguez-Tenreiro 1998; O\\~norbe et al. 2008; Romano-Diaz et al. 2008). Recently, Pedrosa, Tissera \\& Scannapieco (2009, 2010) investigated how baryonic assembly history affects the dark matter distribution. They found the SN feedback process to play a key role by regulating star formation activity and ejecting material not only from the main system, but also from infalling satellites.  It seems clear that the final distribution of baryons at the centre of a halo is insufficient to determine halo response to galaxy assembly, thus contradicting the AC hypothesis.   Typical dark matter haloes have long been reported to be triaxial with major to minor axis ratios often exceeding two in CDM-only simulations (e.g. Barnes \\& Efstathiou 1987; Frenk et al. 1988; Dubinski \\& Carlberg 1991; Jing \\& Suto 2002; Hayashi, Navarro \\& Springel 2007). However, their shape changes, becoming more nearly oblate as baryons condense within them (Katz \\& Gunn 1991; Evrard et al. 1994; Tissera \\& Dom\\'{\\i}nguez-Teneiro 1998; Kazantzidis et al 2004; Debattista et al. 2008; A09). These results suggest the need for a comprehensive analysis of halo structure in simulations which include both the appropriate cosmological context and a realistic description of the physics of baryon condensation. It is clearly necessary to analyse a number of different galaxy-sized haloes in order to explore how differing assembly histories affect the structure of the final haloes.  In this paper, we study a set of high-resolution resimulations of the six galaxy-mass haloes of the Aquarius Project (Springel et al. 2008). These were carried out with a version of {\\small GADGET-3} which includes a multi-phase treatment of metal-dependent cooling, star formation and SN feedback (Scannapieco et al. 2005, 2006). The original haloes were selected from a cosmological CDM-only simulation with no restriction on merger history, except that implied by eliminating objects with high-mass close neighbours. This Aquarius halo set is well-suited to study how formation history and baryonic condensation together determine the final structure of dark matter haloes.  Our ability to isolate these effects is aided by comparing results from our hydrodynamical simulations (hereafter SPH runs) with corresponding results from the CDM-only simulations (hereafter DM runs) as reported by N08. Properties of the galaxies in these haloes are studied in Scannapieco et al. (2009) and will be further analysed in a forthcoming paper.  This paper is organized as follows. In Section 2, we describe the numerical experiments. In Section 3, we analyse the dark matter density profiles. Section 4, describes the velocity dispersion structure of our haloes. In Section 5, we study the effects of baryons on the pseudo-phase-space density profile, while Section 6 discusses halo shapes. Section 7 compares the change in halo circular velocity profile between the SPH and the DM runs with the predictions of AC models. Finally, in Section 8, we summarize our main findings. ",
        "conclusions": "We have analysed the dark matter distributions in six galaxy-sized haloes belonging to the Aquarius Project, comparing results from the original dark matter only simulations to those from re-simulations including baryonic processes (S09).  In agreement with previous work, we find that dark matter haloes become more concentrated when baryons condense at their centres, but that the characteristics of the contraction do not correlate in a simple way with the total amount of baryons.  Our main contribution here is to analyse similar mass haloes with a variety of formation histories using high resolution simulations with and without baryons.  Our results show that the response of haloes to the presence of baryons is sensitive to the details of halo assembly.  This set of six galaxy-sized haloes provides the opportunity to study which properties can be considered common to such haloes and which depend significantly on their particular formation history.     Our findings can be summarized as follows: \\begin{itemize}  \\item In the regions dominated by baryons, haloes become significantly more concentrated than their dark matter only counterparts. The level of concentration varies significantly from object to object, however, it is not simply related to the total baryonic mass accumulated in the central galaxy.  For $r < r_{-2}$, the dark matter density profiles of many of our simulations are nearly isothermal except, possibly, very close to the centre.  \\item The velocity dispersion structure is modified in all haloes, with velocity dispersion increasing monotonically to small radii in all cases. The temperature inversion of NFW profiles is no longer present once the galaxy has formed.  In some systems, the tangential dispersion increases more than the radial dispersion (but not all), causing them to become more nearly isotropic.   \\item The $\\beta$-$\\gamma$ relation proposed by Hansen \\& Moore (2006) is obeyed at most over a restricted radial range for some of our haloes. It works best in the central region of those systems where the isothermal behaviour extends over a relatively small range.  The departures from their predictions become large when the profile is isothermal or has constant anisotropy over an extended radial range.  \\item Pseudo-phase-space density  no longer follows the same power law in radius as in Bertschinger's (1985) similarity solution once galaxy formation is included. The profile differs from one halo to the next and in no halo is it consistent with the purely adiabatic contraction of the dark matter only case.  \\item As in previous work, the condensation of baryons makes the central regions of all our haloes less aspherical and more nearly oblate, although in two cases the changes are small. One of these has no significant disk component and the other has a bulge with substantial nett rotation.  \\item None of the simple adiabatic contraction models proposed in earlier work is able to describe how the radial density profiles of  our haloes are modified by baryon condensation. The scheme suggested by Abadi et al (2009; see also Pedrosa et al. (2010)) is the most successful of those we consider, although it can significantly over- or underestimate the effects in individual cases. For the same haloes, the central mass is  overestimated by more than $50\\%$.   The circular velocity curves of our galaxy formation simulations are not as centrally peaked as predicted by such adiabatic contraction models, but are still more rapidly rising  than in many observed spirals.  Since none of our simulated galaxies is disk-dominated it is unclear whether this is a problem.  \\end{itemize}   Our analysis show that haloes of similar virial mass respond in different ways to baryon condensation depending on the details of their assembly history. Consequently, a much more detailed understanding of galaxy formation is needed before we can make reliable predictions for the detailed structure of the dark matter haloes surrounding galaxies."
    },
    "0911/0911.0852_arXiv.txt": {
        "abstract": "We present archival high spatial resolution VLA and VLBA data of the nuclei of seven of the nearest and brightest Seyfert galaxies in the Southern Hemisphere. At VLA resolution ($\\sim 0.1$ arcsec), the nucleus of the Seyfert galaxies is unresolved, with the exception of MCG-5-23-16 and NGC\\,7469 showing a core-jet structure. Three Seyfert nuclei are surrounded by diffuse radio emission related to star-forming regions. VLBA observations with parsec-scale resolution pointed out that in MRK\\,1239 the nucleus is clearly resolved in two components separated by $\\sim$ 30 pc, while the nucleus of NGC\\,3783 is unresolved. Further comparison between VLA and VLBA data of these two sources shows that the flux density at parsec scales is only 20\\% of that measured by the VLA. This suggests that the radio emission is not concentrated in a single central component, as in elliptical radio galaxies, and an additional low-surface brightness component must be present. A comparison between Seyfert nuclei with different radio spectra points out that the ``presence'' of undetected flux on milli-arcsecond scale is common in steep-spectrum objects, while in flat-spectrum objects essentially all the radio emission is recovered. In the steep-spectrum objects, the nature of this ``missing'' flux is likely due to non-thermal AGN-related radiation, perhaps from a jet that gets disrupted in Seyfert galaxies because of the denser environment of their spiral hosts. ",
        "introduction": "Only a small fraction ($\\sim$ 10\\%) of the population of active galactic nuclei (AGN) possess a powerful radio emission ($L_{\\rm 1.4\\,GHz} > 10^{23}$ W/Hz), as found in radio galaxies/quasars and blazars. Seyfert galaxies are part of the ``radio-quiet'' AGN population, with radio luminosity $L_{\\rm 1.4\\,GHz} \\leq 10^{20 - 23}$ W/Hz. Despite their weak radio emission, Seyfert nuclei are very nearby, allowing the study in detail of the radio properties of their central engine.\\\\ Radio observations with arcsecond resolution of several Seyfert samples \\citep[see e.g.][]{ulvestad84, morganti99, thean00} showed that a large fraction of Seyferts have resolved structures, with hint of jets and/or extended emission, the latter usually related to star-forming regions. Several objects, such as NGC\\,1052 \\citep{wrobler84}, NGC\\,1068 \\citep{ulvestad87}, NGC\\,7674 \\citep{momjian03}, and MRK\\,3 \\citep{kukula99} have been found to display a radio morphology with core, collimated jets and hot spots, similar to those found in radio-loud galaxies. However, powerful radio sources have linear structures reaching hundreds of kpc or even Mpc scales, while in Seyferts the radio emission is confined to a few kpc or even sub-kpc scales.\\\\ When observed with milli-arcsecond resolution, the pc-scale structure of Seyfert nuclei is usually resolved in several components \\citep[e.g. NGC\\,3079,][]{trotter98}, resembling a jet structure \\citep[e.g. NGC\\,4151,][]{ulvestad98, nagar01} and sometimes with the presence of extended emission \\citep[e.g. NGC\\,5793,][]{haghi00}. The comparison between arcsec and milli-arcsec radio properties pinpointed a frequent misalignment between pc and kpc-scale jets, suggesting either a change in jet ejection axis, or a bending due to pressure gradients in the ambient medium \\citep{middelberg04}. \\\\ An intriguing characteristic shown by a large number of Seyfert nuclei is that the radio emission arising from their pc-scale structure is often much fainter than that derived from observations with lower resolution, even in the case the nucleus is unresolved. This result suggests that in Seyfert nuclei the radio emission is not concentrated in the central region, as found in powerful radio galaxies, but it extends on scales of tens or hundreds pc \\citep[see e.g.][]{sadler95}. However, not all the Seyfert nuclei have missing flux on parsec scales, as in the case of MRK\\,530 \\citep{lal04}, indicating that the radio emission mainly arises from the central compact component, without evidence of extended, low-surface brightness features.\\\\ In this paper, we present the results of multi-frequency archival VLA and/or VLBA data of a sample of some of the nearest and brightest Seyfert galaxies, taken from the infrared high spatial resolution studies conducted by \\citet{prieto09} and \\citet{reunanen09}. For these sources, no complete and/or unambiguous information on their radio properties could be found in the literature. The comparison between these data with those at pc-scale resolution available from the literature allows a better determination of the physical conditions of the radio emission at kpc and pc scales.  \\begin{figure*} \\begin{center} \\special{psfile=mn09_1328_fig1a.eps voffset=-175 hoffset=0 vscale=80 hscale=80} \\special{psfile=mn09_1328_fig1b.eps voffset=-175 hoffset=180 vscale=80 hscale=80} \\special{psfile=mn09_1328_fig1c.eps voffset=-175 hoffset=360 vscale=80 hscale=80} \\special{psfile=mn09_1328_fig1d.eps voffset=-345 hoffset=0 vscale=80 hscale=80} \\special{psfile=mn09_1328_fig1e.eps voffset=-345 hoffset=180 vscale=80 hscale=80} \\special{psfile=mn09_1328_fig1f.eps voffset=-345 hoffset=360 vscale=80 hscale=80} \\special{psfile=mn09_1328_fig1g.eps voffset=-515 hoffset=0 vscale=80 hscale=80} \\vspace{18cm} \\caption{VLT-VISIR images at 11.8 $\\mu$m of the Seyfert nuclei studied in this paper. Images adapted from \\citet{reunanen09}. Contours are at 3$\\sigma$, 5, 7, 11 and 19 levels. The triangle in NGC\\,7582 represents the ionization cone.} \\label{ir} \\end{center} \\end{figure*}    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. The spectral index is defined as $S {\\rm (\\nu)} \\propto \\nu^{- \\alpha}$. ",
        "conclusions": "We presented the analysis of multi-frequency archival VLA and VLBA data of 7 close and bright Seyfert nuclei. The conclusions we can draw from this investigation are:\\\\  \\begin{itemize}   \\item At VLA resolution, FWHM $\\sim 0^{\\prime\\prime}.1$, the nucleus of the Seyfert galaxies is unresolved, with the only exception of MGC-5-23-16 and NGC\\,7469 which show a core-jet structure. At VLBA resolution, equivalent to a few parsecs, the nucleus of MRK\\,1239 is resolved in two components, while in NGC\\,3783 it is still unresolved. \\\\  \\item The Seyfert galaxies in this study with known circumnuclear star-forming regions in the IR, namely NGC\\,1097, NGC\\,7469 and NGC\\,7582 \\citep{reunanen09} present a radio counterpart for a few ($\\sim$10\\%) of these regions. Most are not detected in radio, indicating that those detected may be characterised by higher supernovae rate. Their steep radio spectral index is in line with this idea. Conversely, none of the other Seyfert nuclei present any circumnuclear radio emission. The only exception is NGC\\,5506 that shows a radio halo surrounding the nucleus. \\\\  \\item A comparison between arcsecond and milli-arsecond resolution in MRK\\,1239 and NGC\\,3783, pointed out that almost 80\\% of the radio flux density detected in VLA observations is not recovered at pc-scale resolution. This suggests the presence of a diffuse component on scales of a few tens of parsecs, undetected with the VLBA. Similar situation is found in NGC\\,1068, NGC\\,4151, NGC\\,5506 and NGC\\,7469. The nature of this ``missing'' flux components is likely due to synchrotron AGN-related emission. This difference between the flux density measured on arcsecond and milli-arcsecond resolution images is not found in the case of elliptical radio galaxies, but it appears to be a common phenomenon in Seyfert galaxies with a steep spectrum, mostly hosted in spirals. If of synchrotron origin, this emission may be spilt off from a jet that may get easily distorted and/or disrupted by the dense interstellar medium in the nucleus of spirals.\\\\  \\item A comparison between Seyfert nuclei with different spectral properties points out that in flat-spectrum nuclei, almost all the flux density is recovered on milli-arcsecond scale. This indicates that in flat-spectrum objects the radio emission is essentially concentrated in the compact core, without evidence of jet-like structure even on milli-arcsecond scale, while in steep-spectrum objects a significant fraction of the radio emission arises from low-surface brightness, extended features.  \\end{itemize}"
    },
    "0911/0911.2585_arXiv.txt": {
        "abstract": "A solution of the large discrepancy existing between inclusive and exclusive measurements of the ${}^8{\\rm Li}+{}^4{\\rm He}\\to{}^{11}{\\rm B}+n$ reaction cross section at $E_{cm} <3$ MeV is evaluated. This problem has profound astrophysical relevance for this reaction is of great interest in Big-Bang and r-process nucleosynthesis. By means of a novel technique, a comprehensive study of  all existing ${}^8{\\rm Li}+{}^4{\\rm He}\\to{}^{11}{\\rm B}+n$ cross section data is carried out, setting up a consistent picture in which all the inclusive measurements provide the reliable value of the cross section.  New unambiguous signatures of the strong branch pattern non-uniformities, near the threshold of higher ${}^{11}{\\rm B}$ excited levels, are presented and their possible origin, in terms of the cluster structure of the involved excited states of ${}^{11}{\\rm B}$ and ${}^{12}{\\rm B}$ nuclei, is discussed. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0911/0911.4195_arXiv.txt": {
        "abstract": "A set of HI sources extracted from the north Galactic polar region by the ongoing ALFALFA survey has properties that are consistent with the interpretation that they are associated with isolated minihalos in the outskirts of the Local Group (LG). Unlike objects detected by previous surveys, such as the Compact High Velocity Clouds of Braun \\& Burton (1999), the HI clouds found by ALFALFA do not violate any structural requirements or halo scaling laws of the $\\Lambda$CDM structure paradigm, nor would they have been detected by extant HI surveys of nearby galaxy groups other than the LG. At a distance of $d$ Mpc, their HI masses range between $5\\times 10^4d^2$ and $10^6d^2$ \\msun ~and their HI radii between $<0.4d$ and $1.6d$ kpc. If they are parts of gravitationally bound halos, the total masses would be on order of $10^8$--$10^9$ \\msun, their baryonic content would be signifcantly smaller than the cosmic fraction of 0.16 and present in a ionized gas phase of mass well exceeding that of the neutral phase. This study does not however prove that the minihalo interpretation is unique. Among possible alternatives would be that the clouds are shreds of the Leading Arm of the Magellanic Stream. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0911/0911.1315_arXiv.txt": {
        "abstract": "WMAP5 and related data have greatly restricted the range of acceptable cosmologies, by providing precise likelihood ellypses on the the $w_0$--$w_a$ plane. We discuss first how such ellypses can be numerically rebuilt, and present then a map of constant--$w$ models whose spectra, at various redshift, are expected to coincide with acceptable models within $\\sim 1\\, \\%$. ",
        "introduction": "\\label{intro}  One of the main puzzles of cosmology is why a model as $\\Lambda$CDM, implying so many conceptual problems, apparently fits all linear data in such unrivalled fashion \\citep{astier,riess,komatsu}.  It is then important that the fine tuning paradox of $\\Lambda$CDM is eased in cosmologies where Dark Energy (DE) is a self--interacting scalar field $\\phi$ (dDE cosmologies), with no likelihood downgrade \\citep{colombo,lavacca}. Unfortunately, however, only cosmological observables can provide information on the form of the self--interaction potential, even though several researchers incline to priviledge potentials allowing tracking solutions.  When aiming to obtain information on the physical potential, the basic observable is however the evolution of the DE scale parameter, $w(a)$. Here $a = 1/(1+z)$ is the scale factor of a spatially flat metric \\begin{equation} ds^2 = c^2 dt^2 - a^2(t) d\\ell^2~, \\label{metric} \\end{equation} while $z$ is the redshift.  The analysis of available data, made by the WMAP team \\citep{komatsu}, was able to constrain the coefficients $w_0$ and $w_a$ in the expression \\begin{equation} w(a) = w_0 + (1-a)\\, w_a~, \\label{wa} \\end{equation} putting again into evidence that a model with $w \\equiv -1$ is close to top likelihood, but also stressing a preference for the phantom area ($w < -1$), which is hardly consistent with current tracking potentials.  The main tool, to go beyond these constraints on the $w(a)$ law, will be tomographic shear analyses \\citep{refregier}, able to reconstruct the spectrum of density fluctuations at various $z$'s, with a precision approaching 1$\\, \\%$ \\citep{huterer}.  This work is therefore focused on the relevance of spectral predictions, and aims at providing a tool to ease the determination of model spectra. Within this context, it is essential to outline the spectral equivalence criterion (SEC).  It has been noted since a few years \\citep{white,linderW} that the density fluctuation spectra $P(k) = \\langle |\\delta \\rho/\\rho |^2 \\rangle$, up to $k=3\\, h$Mpc$^{-1}$, essentially depend on the distance from the LSB (Last Scattering Band). This was first verified at $z=0$ \\citep{linder} and then tested at higher $z$ \\citep{CaMaB}. $N$--body simulations show that the SEC is however true when model parameters are suitably tuned.  Let us be more specific on this point: When a given $w(a)$ is considered, it is easy to determine the comoving distance of the LSB. Keeping then the same values of $\\Omega_{b,c}$ and $h$, we can seek a constant--$w$ cosmology with the same distance from the LSB; \\citep{linder} tested that the spectra of these two cosmologies coincide (within $1\\, \\%$). They also saw that, when $z \\neq 0$ values are explored, spectral discrepancies are mostly greater, in the range of a few percents. \\begin{figure}% \\vskip -2.4truecm \\begin{center} \\includegraphics[width=10.cm]{wmap5.w.eps} \\end{center} \\vskip -.6truecm \\caption{1-- and 2--$\\sigma$ curves, yielding the marginalized model likelihood on the $w_0$--$w_a$ plane, as obtained from the Lambda NASA site. The reproduction device is detailed in the text.  Also 0.5-- and 1.5--$\\sigma$ curves are provided. Crosses indicates models for which the SEC was explicitly tested at various $z$'s. } \\vskip -.3truecm \\label{elly} \\end{figure}  The SEC however works also at $z \\neq 0$, provided that the distance between such $z$ and the LSB is evaluated and a constant--$w$ auxiliary model is defined, with an equal distance between $z$ and the LSB. The assigned cosmology and the auxiliary model must also have equal $\\Omega_{b,c}$, $h$ {\\it at $z=0$}, while $\\sigma_8(z=0)$ must be tuned in order that, at $z$, the {\\it r.m.s.}  density fluctuations of the two models, on the $8\\, h^{-1}$Mpc scale, coincide \\citep{CaMaB}.  Let us outline, in particular, that the SEC does not require that the auxiliary model shares the values of $\\Omega_{b,c}$ and $h$ at the assigned $z \\neq 0$. In fact, by multiplying the critical density definition by $\\Omega_{c,b}$, one has \\begin{equation} \\omega_{c,b} \\propto \\Omega_{c,b} H^2 = (8\\pi G/3) \\Omega_{c,b} \\rho_{cr} = (8\\pi G/3) \\rho_{c,b} \\propto a^{-3} \\end{equation} so that the assigned cosmology and the auxiliary model share the values of $\\omega_{c,b}$ at any $z$, and this is enough.  This comes as no surprise, however: most linear feature, {\\it e.g.} BAO's, essentially depend just on $\\omega_{b,c}$.  This recipe, however, prescribes a different constant--$w$ at any $z$ and is somehow curious that, starting from the assigned $w(a)$, one builds a $w_{eff}(a)$ law, completely different from it. In a similar way, one can build a $\\sigma_{8;(z=0)}$ dependence from $z$, in order that auxiliary models be fairly normalized at any $z$. Examples of $w_{eff}(z) $ and $\\sigma_{8;(z=0)}(z)$ are provided by \\citep{CaMaB}.  The scope of this work, instead, amounts to applying the SEC to models with $w(a)$ of the form (\\ref{wa}) and consistent with data, exploring decreasing likelihood ellypsoids.  The plan of the paper is as follows: In the next Section we shall discuss how one can rebuild the likelihood ellypsoids on the $w_0$--$w_a$ plane.  This Section defines the parameter $t$ used in the plots which are one of the results of this work. In Section 3 we shall then discuss such plots. Finally, Section 4 will be devoted to drawing our conclusions. ",
        "conclusions": "\\label{conclusions}  When tomographic cosmic shear data will become available, an inspection on DE nature will surely start from comparing them with the predictions of constant--$w$ cosmologies.  As soon as data will become more refined, it will be possible to bin them, discriminating among the $w$ values best fitting data at various redshift.  It is quite possible that these inspections yield $w$ values compatible with a constant, all through the redshift range explored. As it is possible that such value is compatible with $-1$, so vanifying the efforts to improve our knowledge on DE nature.  Let us suppose that, instead, data analysis is consistent with the same $\\omega_{b,c}$ values in all bins, but require different $w$'s in different bins. \\begin{figure}% \\begin{center} \\includegraphics[width=10cm]{tw2.s1.eps} \\end{center} \\caption{Bias in observational $w$ values (named here $w_{eff}$), in respect to the physical value of the state parameter, at each redshift. The points in the curves of the upper frame refer to different models at 1--$\\sigma$ from the top--likelihood cosmology, discriminated~by different $t$ values. In the lower frame the ratio $w_{eff}/w$ is explicitly shown as a function of $t$~.} \\label{tw21} \\end{figure} \\begin{figure}% \\begin{center} \\includegraphics[width=10cm]{tw2.s2.eps} \\end{center} \\caption{As Figure \\ref{tw21} for models at 2--$\\sigma$ from the top--likelihood cosmology.} \\label{tw22} \\end{figure} We wish to outline here a major danger that data analysis could meet in this welcome case. As a matter of fact, one could be tempted to conclude that the $w(z)$ dependence found is the physical scale dependence of the DE state parameter. Unfortunately, this could be badly untrue.  In fact, cosmic shear spectra can be easily translated into fluctuation spectra, so that observational values of $w(z)$ would correspond to $w_{eff}(z)$, not to the physical $w(z).$  How different the two behaviors can be is already represented by Figures 2--5 in this paper. But Figures \\ref{tw21} and \\ref{tw22} further illustrate this point. They show what is $w_{eff}$ when $w$ has a given value, for models laying on the 1--$\\sigma$ or 2--$\\sigma$ curves.  Different colors in each plot correspond to different $z$'s. In the lower frame, the $w_{eff}/w$ ratio is also shown, as a function of $t$.  Even in the most furtunate cases, provided that the {\\it true} cosmology is not a constant--$w$ one, discrepancies are hardly below 10$\\, \\%$ and are greatest at $z$ approaching zero.  Another point concerning model fitting is that a systematic mapping of (reasonable) dDE cosmologies, in order to fit future data, is apparently unnecessary. At each $z$, in fact, there will be a constant--$w$ model able to fit any dDE cosmology. dDE will then be revealed by the $w$ dependence on $z$, as above outlined. At present, well approximated analytical expressions of $P(k)$, at different $z$, are provided by the so--called {\\it Halofit} formulae, holding for $\\Lambda$CDM cosmologies \\citep{smith}. Some attempt to generalize {\\it Halofit} to constant $w \\neq -1$ were also performed \\citep{mcdonald}, but they do not cover the desired parameter ranges. Our conclusion is that it will be important to extend {\\it Halofit} to constant--$w$ cosmologies, for the whole range of (reasonable) cosmological parameters; this will enable us to fit future cosmic shear data without any substantial restriction on the $w(a)$ behavior."
    },
    "0911/0911.1409_arXiv.txt": {
        "abstract": "We observed CO $J=3-2$ emission from the ``water fountain\" sources, which exhibit high-velocity collimated stellar jets traced by  \\h2o maser emission, with the Atacama Submillimeter Telescope Experiment (ASTE) 10~m telescope. We detected the CO emission from two sources,  \\i1634 and  \\j1828. The  \\i1634 CO emission exhibits a spectrum that is well fit to a Gaussian profile, rather than to a parabolic profile, with a velocity width (FWHM) of $158 \\pm 6$ \\kms \\ and an intensity peak at  \\vlsr$=50 \\pm2$ \\kms. The mass loss rate of the star is estimated to be $  \\sim 2.9\\times 10^{-5}M_{ \\odot}$~yr$^{-1}$. Our morpho-kinematic models suggest that the CO emission is optically thin and associated with a bipolar outflow rather than with a (cold and relatively small) torus. The  \\j1828 CO emission has a velocity width (FWHM) of $3.0 \\pm 0.2$ \\kms, smaller than typically seen in AGB envelopes. The narrow velocity width of the CO emission suggests that it originates from either an interstellar molecular cloud or a slowly-rotating circumstellar envelope that harbors the \\h2o maser source. ",
        "introduction": "The process in which mass outflow shapes planetary nebulae (PNe) is still unresolved and needs to be elucidated. One of the striking features of some PNe is the bipolarity in their morphological and kinematical structures, some of which clearly exhibit high velocity bipolar jets ($V_{\\mbox jet}>$100 \\kms). Such bipolar jets have been found at the beginning of the post asymptotic giant branch (AGB) phase and even at the end of the AGB phase (e.g. \\cite{sah98,ima02,ima07b}). In particular, a small fraction of stellar jets have velocities higher than a typical expansion velocity of a circumstellar envelope (CSE) found in 1612~MHz OH maser emission (\\vjet $ \\gtrsim$30\\kms), and are traced by  \\h2o \\ maser emission. Such high velocity jets are called ``water fountains\" \\citep{ima07b}. As of 2009, 13 such objects have been identified \\citep{wal09,sua09}. Some spatio-kinematical structures of the jets have been elucidated by very long baseline interferometry (VLBI) observations of the  \\h2o masers  \\citep{ima02,ima05,ima07,ima07b, bob07,cla09}.  Interestingly, all of the estimated apparent dynamical ages of the jets are about 100~yr or shorter, which is consistent with the rare detection of water fountain sources. This implies that these jets should have just been launched at the final evolutionary stage of the dying stars. Note that the  \\h2o maser emission is excited when the bipolar tips of the stellar jet strike into the ambient CSE that was produced from the previous spherically symmetric stellar mass loss. Therefore, in order to more reliably estimate the stellar mass loss rate, the dynamical age, and the whole morphological and kinematical structure of the jet, a mapping observation of thermal emission such as CO emission is   crucial.  The first detection of CO emission from a water fountain source was reported with the Arizona Radio Observatory 10~m telescope towards  \\i1634  \\citep{he08}. Usually it is difficult to detect such CO emission towards the water fountain sources because they are located close to the Galactic plane with heavy contamination from the interstellar CO emission or they are too distant ($D \\gtrsim$2~kpc). Here we present results of the first systematic CO $J=3-2$ emission observations towards the water fountain sources with the Atacama Submillimeter Telescope Experiment (ASTE) 10~m telescope. In addition, we also show results of CO $J=1-0$ emission observations towards W~43A and  \\j1828 with the Nobeyama 45~m telescope. In section 2 and 3, we describe in detail the observations and results, respectively. In section 4, we discuss the mass loss rate and the morphology and kinematics of  \\i1634, whose CO $J=3-2$ emission is detected. ",
        "conclusions": "Using ASTE, we have surveyed CO $J=3-2$ emission from nine water fountain sources and detected the CO emission towards  I16342 and  \\j1828. I16342 may have a mass loss rate of $ \\dot{M}_{ \\mbox{ \\scriptsize{gas}}}  \\approx  5.6\\times10^{-5} \\:M_{ \\odot}$~yr$^{-1}$, with most emission attributed to a bipolar outflow rather than a high velocity expanding torus or a biconical flow with a wide opening angle. On the other hand, to further clarify the intrinsic CO detection from \\j1828 with a very narrow velocity width ($\\sim 3$\\kms), we need interferometric observations of this object. Because the water fountain sources are located at large distances, higher sensitivity is essential in future observations.  \\bigskip  We deeply appreciate members of the ASTE team for their careful observations, preparation and kind operation support. The ASTE project is driven by Nobeyama Radio Observatory (NRO), a branch of National Astronomical Observatory of Japan (NAOJ), in collaboration with University of Chile, and Japanese institutes including University of Tokyo, Nagoya University, Osaka Prefecture University, Ibaraki University, and Hokkaido University. Observations with ASTE were in part carried out remotely from Japan by using NTT's GEMnet2 and its partner R\\&E (Research and Education) networks, which are based on AccessNova collaboration of University of Chile, NTT Laboratories, and NAOJ. We also acknowledge the referee, Albert Zijlstra, for careful reading and providing several fruitful comments for paper improvement.  HI and SD have been financially supported by Grant-in-Aid for Scientific Research from Japan Society for Promotion Science (20540234). JN was supported by the Research Grants Council of the Hong Kong under grants HKU703308P, and by the financial support from the Seed Funding Programme for Basic Research in HKU (200802159006)."
    },
    "0911/0911.1912_arXiv.txt": {
        "abstract": "\\normalsize  The inner 10~pc of our galaxy contains many counterpart candidates of the very high energy (VHE; $> 100$~GeV) \\gr\\ point source \\hgcfull.  Within the point spread function of the \\hess\\ measurement, at least three objects are capable of accelerating particles to very high energies and beyond, and of providing the observed \\gr\\ flux. Previous attempts to address this source confusion were hampered by the fact that the projected distances between those objects were of the order of the error circle radius of the emission centroid (34'', dominated by the pointing uncertainty of the \\hess\\ instrument). Here we present \\hess\\ data of the Galactic Centre region, recorded with an improved control of the instrument pointing compared to \\hess\\ standard pointing procedures. Stars observed during \\gr\\ observations by optical guiding cameras mounted on each \\hess\\ telescope are used for off-line pointing calibration, thereby decreasing the systematic pointing uncertainties from 20'' to 6'' per axis. The position of \\hgc\\ is obtained by fitting a multi-Gaussian profile to the background-subtracted \\gr\\ count map. A spatial comparison of the best-fit position of \\hgc\\ with the position and morphology of candidate counterparts is performed.  The position is, within a total error circle radius of 13'', coincident with the position of the supermassive black hole \\astar\\ and the recently discovered pulsar wind nebula candidate \\pwn. It is significantly displaced from the centroid of the supernova remnant \\aeast, excluding this object with high probability as the dominant source of the VHE \\gr\\ emission. \\vspace{5mm}\\\\ Key words: Galaxy: centre -- ISM: individual: Sgr~A~East -- ISM: individual: Sgr~A* -- ISM: individual: G~359.95-0.04 -- gamma-rays: observations ",
        "introduction": "\\label{Introduction}  Since the discovery of the strong compact radio source \\astar\\ \\citep{Balick:1974aa}, the Galactic Centre (GC), as the closest galactic nucleus, has served as a unique laboratory for investigating the astrophysics of galactic nuclei in general. The radio picture \\citep{LaRosa00} of the central few 100~pc around the centre of the Milky Way exhibits a complex and very active region, with numerous sources of non-thermal radiation, making this region a prime target for observations at very high energies (VHE; $> 100$~GeV).  Indeed several Imaging Atmospheric Cherenkov Telescopes (IACTs) have detected a source of VHE \\grs\\ in the direction of the GC \\citep{Aharonian:2004wa,Albert:2005kh,Kosack:2004ri,Tsuchiya:2004wv}. The \\hess\\ instrument \\citep[see][and references therein]{crab} provides the to date most precise VHE data on this source, henceforth called \\hgcfull.  As shown with deep observations in 2004, \\hgc\\ is a point source for \\hess\\ (rms spatial extension $< 1\\farcm 2$ at 95\\% CL), and is within $7''\\pm 14''_{\\mathrm{stat}}\\pm 28''_{\\mathrm{sys}}$ positionally coincident with the bright radio source \\astar\\ \\citep{Aharonian:2006wh}. The measured energy spectrum does not fit Dark Matter (DM) model spectra -- at least for the most popular models of DM annihilation --, ruling out the bulk of the TeV emission soley to be of a DM origin \\citep{Aharonian:2006wh}.  Of all possible astrophysics counterparts, the $3\\times 10^6\\ \\mathrm{M}_{\\astrosun}$ supermassive black hole (SMBH) coincident with the \\astar\\ radio position is a compelling candidate. Various models predict VHE emission from this object, produced either close to the SMBH itself \\citep{Aharonian:2005ti}, within an ${\\cal O}(10)$~pc zone around \\astar\\ due to the interaction of run-away protons with the ambient medium \\citep{Aharonian:2005b,Liu2006a,Wang:2009}, or by electrons accelerated in termination shocks driven by winds emerging from within a couple of Schwarzschild radii \\citep{Atoyan2004}. \\astar\\ is a source of bright and frequent X-ray and infrared flares. Detection of quasi-periodic oscillations (QPOs) on time scales of 100-2250~s has been claimed \\citep[e.g.][]{Baganoff2001,Genzel:2003,Porquet2003}. Recently, however, observations with the Keck II telescope could not confirm the existence of such QPOs \\citep{Meyer:2008}. No hint for variability, flaring activity, or QPOs has been found in the VHE \\gr\\ lightcurve in 93~h live time of \\hess\\ data collected during the years 2004-2006 \\citep{GCSpectrum}. Moreover, during a campaign of simultaneous \\hess\\ and Chandra observations of \\astar\\ in 2005, a major X-ray flare of 1600~s duration was observed. Although the X-ray flux increased to $\\approx 9$ times the quiescent level, no evidence for flaring activity was detected in the VHE lightcurve \\citep{Aharonian:2008yb}. This result makes it highly unlikely that X-ray and VHE emission originate from the same source region, and puts constraints on models predicting correlated flaring.  Besides \\astar\\ and its immediate vicinity, there are at least two other production site candidates for VHE emission. The first one is the radio-bright, shell-like supernova remnant (SNR) \\aeast, which surrounds partially \\astar. SNRs have been shown to be efficient particle accelerators \\citep[see e.g.][]{Helder:2009fm}, and the presence of an ${\\cal O}$(mG) magnetic field \\citep{Yusef96} makes \\aeast\\ a compelling candidate for particle acceleration to very high energies \\citep{Crocker2005}. The second one is the recently detected pulsar wind nebula (PWN) candidate \\pwn\\ \\citep{Wang:2005ya}. Despite of its faint X-ray flux, it may plausibly emit TeV \\grs\\ at an energy flux level compatible with the \\hess\\ observations \\citep{Hinton:2006zk}, assuming that \\pwn\\ is located at the same distance as \\astar.  A firm identification of \\hgc\\ is particularly hampered by the -- compared to radio or X-ray instruments -- modest angular resolution of the current generation of Cherenkov telescopes ($\\leq 5'$ for a single \\gr\\ at TeV energies), which gives rise to source confusion in this densely populated region of the galaxy. Adopting a distance to the GC of 8.33~kpc \\citep{Gillessen:2009}, the \\hess\\ source size upper limit encloses a region of about 2.9~pc radius. Comparing this number to the projected distance of \\astar\\ to the radio maximum of \\aeast\\ and the X-ray maximum of \\pwn\\ (3.7~pc and 0.4~pc, respectively), it becomes clear that a precise position measurement of the centre-of-gravity of \\hgc\\ can help to shed light on the nature of this source.  Although previous \\hess\\ position measurements have been unprecedentedly precise, the relatively large -- compared to statistical errors -- systematic errors due to pointing uncertainties of the \\hess\\ array rendered the identification of the major contributing source of the VHE emission, and especially a clear statement on the role of \\aeast, difficult. In this paper a refined measurement of \\hgc's emission centroid is reported. Using improved telescope pointing control, the systematic error of the measurement is decreased by a factor of three compared to previous results, and the total error on the centroid position is reduced to $13''$ (68\\% containment radius), compared to $34''$ in \\cite{Aharonian:2006wh}. ",
        "conclusions": "\\begin{figure} \\includegraphics[width=0.49\\textwidth]{./figures/Fig2.eps} \\caption{90~cm VLA radio flux density map \\citep{LaRosa00} of the innermost 20~pc of the GC, showing emission from the SNR \\aeast. Black contours denote radio flux levels of 2, 4, and 6~Jy/beam. The centre of the SNR \\citep{Green:2009qf} is marked by the white square, and the positions of \\astar\\ \\citep{saga_radio} and \\pwn\\ \\citep[head position,][]{Wang:2005ya} are given by the cross hairs and the black triangle, respectively. The 68\\% CL total error contour of the best-fit centroid position of \\hgc\\ is given by the white circle. The dashed white circle shows the same contour for the previously reported \\hess\\ measurement \\citep{Aharonian:2006wh}. The white and black dashed-dotted lines show the 95\\% CL upper limit contour of the source extension and the 68\\% containment region of the \\hess\\ PSF, respectively. The white stars marked \\emph{A} and \\emph{B} denote the position of the radio maximum and the best-fit position for the radio emission after smoothing with the PSF of the \\hess\\ instrument, respectively.  } \\label{fig:RadioMap} \\end{figure}  Fig.~\\ref{fig:RadioMap} shows a VLA 90~cm image of the innermost 20~pc region of the GC, centred on \\astar. The shell-like radio structure of the SNR \\aeast\\ is clearly visible. The best-fit position of \\hgc\\ lies in a region where the radio emission is comparatively low, and is shown as a 68\\%~CL total error contour, computed from the summed (in quadrature) statistical and systematic best-fit position errors. As can be seen from the figure, the centroid of the VHE source is coincident with the positions of \\astar\\ and \\pwn, but inconsistent with the regions of intense radio emission from \\aeast. Two rather independent approaches to derive a quantitative statement about the compatibility of \\hgc's best-fit position with \\aeast\\ are detailed in the following.  The white star labelled $A$ in Fig.~\\ref{fig:RadioMap} denotes the position of \\aeast's radio maximum. Comparing the 68\\%~CL radius of the observed VHE centroid to the angular distance between $A$ and the best-fit position, a coincidence of the two positions is ruled out with 7.1 standard deviations. By the same arguments, VHE point emission from the centre of the SNR \\citep{Green:2009qf}, indicated as a white square in Fig.~\\ref{fig:RadioMap}, is ruled out with 4.7 standard deviations.  Instead of point emission, extended emission can be considered, e.g. by assuming that the hypothetical VHE emission from \\aeast\\ follows closely the morphology of the radio flux. The centroid of such emission would be detected at the coordinates marked $B$ in Fig.~\\ref{fig:RadioMap}. This position was derived by fitting the radio map -- smoothed with the \\hess\\ PSF -- with the technique used above for the VHE \\gr\\ data. Following the methods used for position $A$, the radio fit position is 5.4 standard deviations away from the best-fit VHE centroid position.  The above results are obtained assuming that the VHE emission and radio morphology are correlated. Since the best-fit position does not coincide with a region of intense radio emission (see Fig.~ \\ref{fig:RadioMap}), relaxing this assumption leads to more conservative estimates of the association probability.  A priori, it would appear conservative to assume that the centroid of TeV emission associated with \\aeast\\ might appear anywhere within the boundaries, with equal probability. With this assumption one can calculate the probability that the VHE emission is produced inside \\aeast, but is only by chance positionally coincident with \\astar\\ and \\pwn, which themselves are plausible emitters of VHE radiation and thus viable counterpart candidates of \\hgc. Defining the 2~Jy/beam radio contour of \\aeast\\ as the SNR boundary, which encloses the best-fit position of the emission centroid, a chance probability of $9\\times 10^{-5}$ is derived (corresponding to 3.9 standard deviations). This number does slightly change depending on the assumed size of the SNR boundary.  It is clear, however, that even with this conservative approach an association of \\aeast\\ with the observed VHE \\gr\\ emission is rather unlikely.  The exclusion of \\aeast\\ as the main contributor to the VHE emission is a major step towards an identification of \\hgc.  Despite the fact that \\hgc\\ is a non-variable \\gr\\ source, both \\pwn\\ and \\astar\\ are compelling counterpart candidates, as models exist (see section \\ref{Introduction}), which can explain a steady $\\gamma$-ray flux and variable X-ray emission from \\astar\\ at the same time.  More information is needed to discriminate between these two objects. Due to their enhanced angular resolution and sensitivity, and their extended energy range, proposed future VHE \\gr\\ observatories such as CTA or AGIS could shed light on the open question of which of these sources dominates the production of the VHE \\gr\\ emission from the gravitational centre of our galaxy."
    },
    "0911/0911.2700_arXiv.txt": {
        "abstract": "We make an inventory of the baryonic and gravitating mass in structures ranging from the smallest galaxies to rich clusters of galaxies.  We find that the fraction of baryons converted to stars reaches a maximum between ${M}_{500} = 10^{12}$ and $10^{13}\\;\\mathrm{M}_{\\sun}$, suggesting that star formation is most efficient in bright galaxies in groups. The fraction of baryons detected in all forms deviates monotonically from the cosmic baryon fraction as a function of mass. On the largest scales of clusters, most of the expected baryons are detected, while in the smallest dwarf galaxies, fewer than 1\\% are detected. Where these missing baryons reside is unclear. ",
        "introduction": "The early universe was a highly uniform and homogeneous mix of dark and baryonic matter with baryon fraction ${f_b = 0.17 \\pm 0.01}$ \\citep{WMAP5} . If this primordial mix persists as individual gravitationally bound structures emerge during the course of cosmic evolution, then the baryonic mass of any given object would be $M_b = f_b M_{tot}$.  Indeed, combining this with the baryon density constraint from big bang nucleosynthesis \\citep{BBN} provides one important argument that the density parameter is less than unity \\citep{White93,OS}.  Here we reverse the logic and ask what fraction of the expected baryons are actually detected.  We present an inventory of the baryonic and gravitating masses of cosmic structures spanning a dozen decades in detected baryonic mass.  The mass of baryons known in each system correlates well with the total mass, but not as a simple proportion.  The implication is that most of the baryons associated with individual dark matter halos are now missing. While some are in the intergalactic medium \\citep{OVI}, a complete accounting of where the baryons now reside, and how they relate to their parent structures, remains wanting. ",
        "conclusions": "The stellar fraction reaches a maximum between ${M}_{500} = 10^{12}$ and $10^{13}\\;\\mathrm{M}_{\\sun}$ (Fig.~\\ref{fdfstV}). This is broadly consistent with previous results based on counting statistics \\citep{yang}. However, there may be some offset in the peak mass scale relating to the long standing dichotomy between Tully-Fisher and luminosity function based normalizations of the halo mass function.  Notably, the efficiency of star formation appears to increase monotonically with mass for individual galaxies \\citep{baldry}.  After the transition to cluster halos containing many galaxies, the efficiency declines again.  Perhaps the most striking aspect of Fig.~\\ref{fdfstV} is that the fraction of detected baryons falls short of the cosmic fraction at all scales.  In no system is it unity, as would be expected if we had a complete accounting of all the baryons associated with a given bound system.  Where are all these missing baryons?  An obvious possibility is that the missing baryons are present, and simply inhabit their dark matter halos in some undetected form.  Indeed, one might expect that not all baryons would have time to cool into the observed cold gas and stellar component of galaxies.  In this case, many baryons might remain mixed in with the dark halo.  The notion that we are simply not seeing many or even most of the baryons in individual galaxies is profoundly unsatisfactory.  Direct searches for hot halo gas have turned up nothing substantial: halo baryon reservoirs fail to explain the observed deficit by two orders of magnitude \\citep{bregman,andbreg}. In clusters, the hot gas is detected, and constitutes the majority of the baryons. Clusters fall short of the cosmic baryon fraction by a modest amount \\citep{mccarthy,giodini} which might be readily explicable \\citep{crain}.  However, the fraction of missing baryons is highly significant on the scales of individual galaxies. In dwarfs, fewer than 1\\% of the baryons expected from the cosmic fraction are detected.  The detected baryon fraction varies systematically with scale. It is a matter of taste whether one chooses to describe this scale as one of circular velocity, mass, or potential well depth.  It is tempting to attribute this correlation to feedback processes being more effective in objects with smaller potential wells \\citep{DS}.  However, the details are heinously complicated \\citep{MayerMoore,crazy} and not well understood. We would naively expect feedback to be a messy process resulting in lots of scatter in any correlations that might result.  Instead, the observed relation is remarkably tight. In principle, the entire range $0 \\le f_d \\le 1$ is accessible at each mass, yet only a very particular value is observed.  Moreover, the current potential well is the result of the hierarchical assembly of many smaller building blocks.  Left to itself, a small dark matter halo will have a small $f_d$.  If incorporated into a larger halo, $f_d$ goes up.  How does each building block know to bring the right amount of baryons to the final halo?  We do not see a satisfactory solution to this missing baryon problem at present. Considerable work remains to be done to obtain a complete understanding of the universe and its contents."
    },
    "0911/0911.2470_arXiv.txt": {
        "abstract": "Classical Cepheid variable stars have been important indicators of extragalactic distance and Galactic evolution for over a century. The \\spitzer{} Space Telescope has opened the possibility of extending the study of Cepheids into the mid- and far-infrared, where interstellar extinction is reduced. We have obtained photometry from images of a sample of Galactic Cepheids with the IRAC and MIPS instruments on \\spitzer. Here we present the first mid-infrared period--luminosity relations for Classical Cepheids in the Galaxy, and the first ever Cepheid period--luminosity relations at 24 and 70~\\micron. We compare these relations with theoretical predictions, and with period--luminosity relations obtained in recent studies of the Large Magellanic Cloud. We find a significant period--color relation for the $[3.6]-[8.0]$ IRAC color. Other mid-infrared colors for both Cepheids and non-variable supergiants are strongly affected by variable molecular spectral features, in particular deep CO absorption bands. We do not find strong evidence for mid-infrared excess caused by warm ($\\sim 500$~K) circumstellar dust. We discuss the possibility that recent detections with near-infrared interferometers of circumstellar shells around $\\delta$~Cep, $\\ell$~Car, Polaris, Y~Oph and RS~Pup may be a signature of shocked gas emission in a dust-poor wind associated to pulsation-driven mass loss. ",
        "introduction": "Although it was over 100 years ago that Henrietta Leavitt discovered the Cepheid Period--Luminosity relation (PL, \\citealt{leavitt1908}), or ``Leavitt Law'', few tools in astronomy have such enduring importance. Classical Cepheid variable stars are fundamental calibrators of the extragalactic distance scale and in addition their observed properties are a benchmark for stellar evolution models of intermediate mass stars. It is now widely accepted that distances derived from SN Ia are still significantly influenced by the classical Cepheid calibration: \\citet{riess2005} found a change of 15\\% in the value of $H_0$ when only modern, high-quality SN Ia data and HST Cepheids were used. Recent works have extended optical and near-infrared PL and period--color (PC) relations to the mid-infrared (mid-IR), where there is less interstellar extinction, with observations of Large Magellanic Cloud (LMC) Cepheids (\\citealt{freedman2008, ngeow2008, ngeow2009, madore2009a}) with the \\spitzer{} Space Telescope \\citep{werner2004}.  Evolutionary and pulsation properties of intermediate-mass stars in the core He-burning phase, like Cepheids, play a crucial role in several long-standing astrophysical problems. They are transition objects between stellar structures ending up their evolution either as white dwarfs or as core collapse supernovae. Therefore, they are not only the most popular primary distance indicators, but are also crucial to understanding the chemical evolution of stellar systems hosting a substantial fraction of young stars, e.g the Galactic disk and the dwarf irregular galaxies in the Local Group and in the Local Volume ($d \\la 10$~Mpc).  Although these objects are fundamental for stellar evolution \\citep{bono2000, beaulieu2001}, stellar pulsation \\citep{bono1999, marconi2005}, and Galactic chemical evolution models \\citep{pedicelli2009, spitoni2009}, current predictions are still hampered by several problems. The most outstanding issue is the ``Cepheid mass discrepancy'' between pulsation masses of classical Cepheids and their evolutionary masses. Evidence was brought forward more than 30 years ago by \\citet{fricke1972} who found that pulsation masses were from 1.5 to 2 times smaller than the evolutionary masses. This conundrum was partially solved \\citep{moskalik1992} by the new sets of radiative opacities released by the Opacity Project \\citep{seaton1994} and by OPAL \\citep{rogers1992}. However, several recent investigations focused on Galactic Cepheids \\citep{bono2001a, caputo2005} and Magellanic Cloud Cepheids \\citep{beaulieu2001, bono2002, keller2006} suggest that such a discrepancy still amounts to 10--15\\%. Measured masses of Galactic binary Cepheids (e.g. \\citealt{evans2008}) are also smaller than predicted by evolutionary models neglecting core convective overshooting during central hydrogen burning phases.  The relative importance of the main factors affecting Cepheid model mass estimates (extra-mixing, rotation, radiative opacity, mass loss and binarity) is still debated. Even though commonly used semi-empirical relations \\citep{reimers1975, dejager1997} do not predict enough mass loss to solve the Cepheid mass discrepancy problem, mass loss may indeed be the key culprit among the physical mechanisms suggested to explain the mass discrepancy problem. The semi-empirical mass loss relation derived by \\citet{reimers1975} is clearly inadequate to correctly estimate the mass loss for certain evolutionary phases of giant stars (see e.g. \\citealt{willson2000}). A plausible increase in the typically adopted Reimers wind free parameter ($0.2 \\le \\eta \\le 0.4$) does not account for the entire range in mass covered by cluster horizontal branch (central helium burning) stars \\citep{yong2000, castellani2005, serenelli2005}. There is no good reason why a similar discrepancy should not affect intermediate-mass stars.  Mass loss is often betrayed by the presence of dust shells, detectable as infrared excesses, or by stellar winds which produce blue-shifted absorption dips in the ultraviolet. Empirical estimates of mass loss rates based on infrared (IRAS) and ultraviolet (IUE spectra) observations for a large sample of Galactic Cepheids suggest mass-loss rates ranging from 10$^{-10}$ to 10$^{-7}$~\\msun~yr$^{-1}$ \\citep{deasy1988}. However, evidence for mass loss rates high enough to affect evolution is very rare, and it is not clear that mass loss is a wide-spread phenomenon. \\citet{mcalary1986} used IRAS photometry \\citep{beichman1985} and found evidence of very cool dust ($T_d \\la 50$~K) around two classical Cepheids (\\rspup{} and SU~Cas) known to be associated with reflection nebul\\ae{}. Due to the very large IRAS beamsize (as large as $\\sim 5$~arcmin), however, it was very difficult to separate local dust emission from ``Galactic cirrus'' background emission. More recently, $K$ band near-IR interferometric observations \\citep{merand2006, merand2007, kervella2006, kervella2008, kervella2009} have detected circumstellar emission around five Classical Cepheids: \\deltacep{}, $\\ell$~Car, Polaris, Y~Oph and RS~Pup. While the nature of the material responsible for this emission remains mysterious, this is a tantalizing suggestion that mass loss activity may be present around these nearby Cepheids. A similar conclusion was also reached by \\citet{neilson2009}, that using OGLE (Optical Gravitational Lensing Experiment; \\citealt{udalski1999}) and \\spitzer{} data for Magellanic Cepheids, found evidence of a wide range of mass loss rates.  To investigate this possibility, we have obtained \\spitzer{} observations of a sample of Galactic Classical Cepheids. All stars were observed with both the \\spitzer{} Infrared Array Camera (IRAC, \\citealt{fazio2004}) and the Multiband Infrared Photometer for \\spitzer{} (MIPS, \\citealt{rieke2004}). The aim of the IRAC observations was to characterize Galactic Cepheid colors in the mid-IR, where the emission from the photosphere is still dominant, and search for infrared excess related to warm ($\\sim 500$~K) circumstellar dust. The MIPS observations were intended to investigate the presence of extended emission from cool ($\\la 100$~K) dust, taking advantage of the higher angular resolution of \\spitzer{} (5~arcsec at 24~\\micron). In this paper we present the photometry of our Cepheid sample in the IRAC and MIPS bands, discussing their PL and period-color (PC) relations. The results of our search for extended emission in IRAC and MIPS images will be presented in a separate paper (Barmby et al. in preparation).  The criteria for our sample selection are laid out in section~\\ref{sec-sample}, and the observations are described in section~\\ref{sec-obs}. The techniques adopted to measure the source photometry in all IRAC and MIPS bands are discussed in detail in section~\\ref{sec-phot}. In section~\\ref{sec-leavitt} we derive the Leavitt Law and PC relations in the IRAC and MIPS bands, and we compare them with similar relations obtained for the LMC. We study the intrinsic mid-IR colors of Cepheids in section~\\ref{sec-excess}, where we set limits on infrared excess at \\spitzer{} wavelengths. Our results are discussed in section~\\ref{sec-discussion} and summarized in section~\\ref{sec-summary}. ",
        "conclusions": "\\label{sec-discussion}  The reliability of Classical Cepheids as standard candles is of paramount importance for astronomy. The recent advances in infrared space astronomy have shifted the focus of obtaining accurate PL relations to wavelengths longer than the visible, where interstellar extinction is reduced. \\citet{madore2009b} and \\citet{freedman2009} have demonstrated how PL relations obtained at IRAC wavelengths can be effectively used to measure the distance of nearby galaxies. This work is a further step in the characterization of PL relations in all \\spitzer{} photometric bands, including two MIPS bands at 24 and 70~\\micron, by using a sample of Classical Cepheids in the Galaxy.  In order to provide a detailed comparison between theory and observations we adopted the large set of nonlinear, convective models computed by \\citet{bono1999, marconi2005} and by \\citet{fiorentino2007}. For Galactic Cepheids we chose a scaled-solar chemical composition (helium, $Y=0.28$; metals, $Z=0.02$) and accounted for fundamental mode pulsators. Moreover, we covered a broad range of stellar masses ($3.5 \\le M/M_\\odot \\le 11.5$), and to account for current uncertainties affecting the size of the helium core \\citep{bono2006} we adopted two different mass-luminosity relations based on canonical and non-canonical (the latter incorporating a range of main sequence core convective overshoot) evolutionary models. Theoretical predictions were transformed into the observational plane using scaled-solar atmospheres based on {\\tt ATLAS9} models \\citep{castelli2003}, and multiplied with the IRAC band-passes. The slopes and zero points of these model Leavitt laws are listed in the last two columns of Table~\\ref{tab-PL}. While the zero points are in general agreement with all our fits, the slopes are significantly shallower than our best fits obtained with both ``new'' and ``old'' distances (the disagreement is however reduced in the fits adopting the ``new'' distances). The PL relations derived with the exclusive use of the astrometric distances (that do not rely on the $p(P)$ relation), are however in excellent agreement with the theoretical relations. This result may suggest that the period dependence of the $p-$factor currently adopted in IRSB distances may still need further refinement. Because of the small number of Cepheids and the limited range of periods for which accurate parallaxes are available, we could not attempt to invert the problem and estimate the $p(P)$ dependence with our data.  Comparison between our Galactic PL relations and the relations derived for the LMC by \\citet{madore2009a} and \\citet{ngeow2009} leads to contradictory results. The PL relations obtained with both the ``new'' and ``old'' IRSB distances are steeper than the LMC relations, while the Galactic PL relations we obtain with the astrometric distances are significantly shallower. To add to the confusion, it should be noted that the slopes derived by \\citet{ngeow2009} and \\citet{madore2009a} disagree by more than their respective uncertainties, despite using stars from the same galaxy\\footnote{The difference between the two LMC results can be an issue of crowding and binarity in the \\citet{ngeow2009} sample (selected without individual image inspection, and containing a larger fraction of fainter short period Cepheids) that could result in shallower slopes.}. The contradictory results in our Galactic PL relation slopes shows once more the effects of the systematics introduced by the IRSB distances and their dependence on the $p(P)$ relation. Based on the more reliable astrometric distances we should conclude that the slope of the PL relation could be shallower at Galactic metallicity than in the LMC, but even this result is not statistically significant, due to the larger uncertainty in the PL slope. In conclusion, the insufficient reliability of the IRSB distances, and the small number of stars for which astrometric distances are available, prevents us from resolving the dependence of the PL relations from metallicity. This problem has already been noticed at optical and near-IR wavelengths by other authors (see e.g. \\citealt{storm2004, fouque2007, romaniello2008}). \\citet{riess2009} also noted that the precision of current datasets does not allow measurement of the dependence of the PL relation on metallicity; those authors measured statistically consistent slopes at 1.6 micron for the Milky Way, the LMC and other galaxies.  We also cannot determine a reliable wavelength dependence of the PL slope, as all our values are within 1$\\sigma$ from the average slope value. This is consistent with the expected flattening of the PL slope wavelength dependence in the mid-IR shown by the models. If we compare our slope values with Figure~4 in \\citet{freedman2008}, however, we note a similar trend. Our result for 4.5~\\micron{} is marginally higher (by $\\sim 0.1$) than the slope at the 3.6 and 8.0~\\micron, as also shown by the LMC Cepheid fit (and by the models). This anomaly in the 4.5~\\micron{} (and, to a minor extent, 5.8~\\micron) slope is related to the presence of the variable CO band, which increases the amplitude variation in these bandpasses.  For the first time we have measured the PL relation in the MIPS 24 and 70~\\micron{} bands. The 24~\\micron{} relation, in particular, is important for the determination of Cepheid distances in the Galactic plane from datasets such as MIPSGAL \\citep{carey2009}, where high extinction from ISM dust may complicate their detection at IRAC bands.  The determination of a PC relation is in principle not limited by the uncertainty in the distances affecting our PL relations. Figure~\\ref{fig-pc}, however, clearly shows that such a relation can be derived with sufficient accuracy only in the $[3.6]-[8.0]$ color. Other colors show a scatter that is increasingly larger with the period, as high as $\\sim 0.1$~mag, one order of magnitude larger than our color accuracy of $\\sim 0.02$. The fact that this scatter was not found by \\citet{ngeow2009} may again be a consequence of the difference in the median period between our Galactic and their LMC sample.  Based on our data and on numerical modeling of one representative Cepheid ($\\zeta$~Gem), we conclude that the color scatter affecting long period Cepheids is a consequence of variable stellar CO absorption, rather than infrared excess from circumstellar dust. Circumstellar dust at $\\sim 500$~K (the temperature where dust thermal emission is maximum in the IRAC bands) would produce the largest excess in the $[3.6]-[8.0]$ color, which we do not observe, with the possible exception of a weak detection ($\\sim 0.1$~mag) around Polaris. Our $\\zeta$~Gem models show that scatter in the IRAC 4.5 and 5.8~\\micron{} bands can be induced by variations in effective temperature and/or the propagation of shocks through the Cepheid atmospheres. A more comprehensive modeling effort, including Cepheids with different periods and pulsation modes, is however required to confirm and quantify this effect.  Not finding an infrared excess in most, if not all, stars is somewhat unexpected, given the detection of circumstellar shells by near-IR interferometers. Figure~8 in \\citet{merand2007} shows a number of targets, including the long period Cepheids Y~Oph and $\\ell$~Car (both part of our sample), with a significant (4--5\\%) excess in the K band. We do not measure any excess of this magnitude in these or other stars, which seems to rule out the presence of circumstellar dust in the shells detected around these and other Cepheids. The interferometric determination of the radius of these circumstellar shells, however, indicate values as low as $\\sim 2$ stellar radii \\citep{merand2006}. For a Cepheid with $T_{\\textrm{eff}} \\sim 5$,000--6,000~K, dust equilibrium temperature at $\\sim 2$~R$_*$ would be $T_d \\simeq T_*/\\sqrt{R_d/R_*} \\la 3$,500~K. This is well above the sublimation temperature of any known astronomical dust.  Our results can be reconciled with the interferometric detection of circumstellar shells close to the star if the emission is not due to dust, but rather to some strong molecular line emission. Given that the observations in \\citet{merand2006, merand2007} and \\citet{kervella2006} were made in the $K$ band, a possible candidate for this emission is shocked $H_2$. The presence of $H_2$ shocked emission lines could further contribute to the scatter in the IRAC 4.5~\\micron{} band, for certain densities and shock velocities \\citep{smith2006}. This hypothesis supports the idea that Cepheid stars indeed have a strong stellar wind associated with a pulsation driven mass-loss mechanism, as suggested by \\citet{merand2007}. Unlike the case of mass loss in red giants, AGB and supergiant stars (showing a dust-driven wind), this stellar wind would be largely dustless, due to the higher temperature of the star. This hypothesis needs to be tested by near-IR multi-epoch spectral monitoring of these stars.  This hypothesis does not imply that a Cepheid stellar wind is completely devoid of dust, that may be condensing at larger radii, in quantities below the detection limits of our IRAC observations (see $\\ell$~Car mid-IR detection of mid-IR circumstellar emission with the VLT MIDI and VISIR instruments, \\citealt{kervella2009}). Cold dust may also be collected by the outflow, at much larger distances (thousands of AU), from the ISM, as in the case of the well known nebula around RS~Pup \\citep{kervella2009}. We find evidence supporting this mechanism in our IRAC and MIPS images. Extended emission at 5.8, 8.0, 24 and 70~\\micron{} around $\\delta$~Cep indicate that this star may be losing mass due to a strong wind pushing into the local interstellar medium, which is leading to the formation of a 70~\\micron{} bow shock detected at large distance from the star ($\\sim 10$,000~AU). A detailed analysis of this phenomenon is being published elsewhere (Marengo et al., in preparation).  In a more comprehensive paper \\citep{barmby2009} we will discuss the presence of spatially resolved extended emission at 24 and 70~\\micron{} around more targets, and quantify the occurrence of mass loss in the Cepheid phase that can be inferred from our imaging data.      We have derived the PL and PC relations in \\spitzer/IRAC bands for a sample of Galactic Cepheids. These relations are critically dependent on the choice of the period dependence of the $p$-factors used for the distance determination with the IRSB method, even though the uncertainties in our fits prevent the assessment of the best $p(P)$ relation choice. The best agreement with theoretical PL relations, however, is obtained when only distances obtained by astrometric methods are used. We do not detect statistically significant variations between the slope and zero points of the PL relations between our Galactic sample and LMC relations obtained by \\citet{madore2009a}, despite the difference in metallicity. We find that the intrinsic variations in the 4.5 and 5.8~\\micron{} fluxes are larger for long period Cepheids. These variations (of the order of $\\sim 0.1$~mag) are related to deep CO absorption, dependent on the stellar $T_{eff}$ and hydrodynamic effects associated to the stellar pulsations. We do not find significant infrared excess related to warm circumstellar dust, except for a weak excess detected at IRAC wavelengths for Polaris, and at 70~\\micron{} for SZ~Tau. This may rule out the presence of extensive dust driven mass loss in the Cepheid phase, but leaves open the possibility of pulsation-driven mass loss from a dust-poor wind, as suggested by recent interferometric observations."
    },
    "0911/0911.4019.txt": {
        "abstract": "{}{}{}{}{} % 5 {} token are mandatory  \\abstract %  % context heading (optional) %  % {} leave it empty if necessary %   {XXX} %  % aims heading (mandatory) %   {YYY} %  % methods heading (mandatory) %   {ZZZ} %  % results heading (mandatory) %   {ZZZ} %  % conclusions heading (optional), leave it empty if necessary {The planetary nebula \\TS\\ (also called \\png\\ or \\sbs), with its record-holding low oxygen abundance and its double degenerate close binary core (period 3.9 h),  is an exceptional object located in the Galactic halo.  We have secured observational data in a complete wavelength range in order to pin down the abundances of half a dozen elements in the nebula.  The abundances are obtained via detailed photoionization modelling taking into account all the observational constraints (including geometry and aperture effects) using the pseudo-3D photoionization code Cloudy$\\_$3D. The spectral energy distribution of the ionizing radiation is taken from appropriate model atmospheres. Incidentally, from the new observational constraints, we find that both stellar components  contribute to the ionization: the ``cool'' one provides the bulk of hydrogen ionization, and the ``hot'' one  is responsible for the presence of the most highly charged ions, which explains why previous attempts to model the nebula experienced difficulties.  The nebular abundances of C, N, O, and Ne are found to be respectively,  1/3.5, 1/4.2, 1/70, and   1/11 of the Solar value, with uncertainties of a factor 2. Thus the extreme O deficiency of this object is confirmed. The abundances of S and Ar are less than  1/30 of  Solar.  The  abundance of He relative to H is  0.089$\\pm$0.009.  Standard models of stellar evolution and nucleosynthesis  cannot explain the  abundance pattern observed in the nebula.  To obtain an extreme oxygen deficiency in a star whose progenitor has an initial mass of about 1\\,\\msun\\   requires an additional mixing process, which can be induced by stellar rotation and/or  by the presence of the close companion.  We have computed a  stellar model with initial mass of 1\\,\\msun, appropriate metallicity, and initial rotation of 100\\kms, and find that rotation greatly improves the agreement between  the predicted and observed abundances. } ",
        "introduction": "\\label{sec:introduction}  \\sbs\\ was discovered in the second Byurakan Sky Survey  and first classified as a cataclysmic variable \\citep{1999PASP..111.1099S}. \\cite{2001A&A...370..456T}  discussed in detail the nature of the object and came to the conclusion that, in fact, it is a  planetary nebula (PN). The  object was renamed  \\png, following the nomenclature for Galactic PNe from the Strasbourg-ESO catalogue of Galactic Planetary Nebulae \\citep{1992secg.book.....A}.  For the sake of brevity, we will refer to it as \\TS\\ in the rest of the paper. This PN is special in at least three important aspects.  First of all, its oxygen abundance is very low, significantly lower than in any other PN known up to now \\citep{2001A&A...370..456T, 2002A&A...395..929R, 2002AJ....124.3340J, 2005A&A...430..187P}. Second, its nucleus is a spectroscopic binary, with a  period  of only a few hours \\citep{2004ApJ...616..485T}. Third, it appears, from estimates of the nature and masses of the two stellar components, that \\TS\\  could turn into a double degenerate Type Ia Supernova \\citep{2004ApJ...616..485T}. Each of these aspects, even taken alone, makes  \\TS\\ an exceptional object.  In this paper, we reexamine the  chemical composition of \\TS. Briefly, the story of the determination of the chemical composition of this object is the following. Tovmassian et al. (2001) had optical spectra of \\TS\\ in the range 3900--7000\\,\\AA\\  obtained with 2\\,m class telescopes which showed no lines from heavy elements except a very weak \\Oiii, with an intensity a few percent of H$\\beta$. A coarse photoionization  analysis  suggested an oxygen abundance smaller than 1/100 of Solar. Note that standard empirical methods for abundance determinations in PNe cannot be used for \\TS, since the electron temperature cannot be determined directly from observations. To go further  in the abundance determination of \\TS\\ required an estimate of the  effective temperature of the central star. One way is to obtain a good blue spectrum of the PN, and use  the   \\Nev/ \\Neiii\\ ratio (or a limit on it) as a constraint.  \\cite{2002A&A...395..929R} at the Canada-France-Hawaii Telescope (CFHT) and \\cite{2002AJ....124.3340J} at the Multiple Mirror Telescope (MMT) secured deep blue spectra in order to detect these lines. \\cite{2002AJ....124.3340J} detected the \\Nev\\ line at a level of ~0.8 H$\\beta$.  \\cite{2002A&A...395..929R} found only an upper limit of 0.1 H$\\beta$! Concerning  the \\Neiii\\ line, {\\cite{2002AJ....124.3340J} measured an intensity about 10 times larger than \\cite{2002A&A...395..929R}. The two papers appeared within a few days of each other on astro-ph, revealing this big conflict in the observations. The two groups conducted independent photoionization analyses, and both concluded that the O/H ratio is less than 1/100 of Solar (the main reason for their similar result for the oxygen abundance was the similar \\Nev/\\Neiii\\ ratio used by both studies).  \\cite{2005A&A...430..187P} merged and discussed the two observational data sets and conducted their own photoionization analysis. They concluded that the O/H ratio of \\TS\\ lies between 1/30 - 1/15 of Solar (still holding the record for the most oxygen poor planetary nebula but much higher than previously published). However,  \\cite{2005A&A...430..187P} neglected to consider observations of \\TS\\ made with the \\textit{Hubble Space Telescope} (HST) and  the \\textit{Far Ultraviolet Spectroscopic Explorer} (FUSE). As a result, some of their ``predicted'' line intensities  are in conflict with what is actually observed in the UV. HST observations were obtained in 2003, and presented in a short, preliminary version by {\\cite{2006IAUS..234..431J}. Those authors quoted an oxygen abundance of 1/30 - 1/40 of Solar, and carbon and nitrogen abundances roughly 1/10 of Solar.  Before embarking on a new determination of abundances, we have chosen to gather the best possible observations at all wavelength ranges. These data provide many more constraints than were available in any previous study. In order to make the best use of the large amount of data obtained with different telescopes, we use a pseudo-3D photoionization code, Cloudy\\_3D, which is able to account for the nebular geometry as we see it now, and with which we can properly take into account the aperture effects. This code is based on CLOUDY \\citep{1998PASP..110..761F} and was written by \\cite{2006IAUS..234..467M}.  The paper is organized as follows. Section 2  presents the new observational material: several optical spectra, HST imaging and spectroscopy, infrared spectroscopy with the \\emph{Spitzer} Telescope, and mentions our X-ray observations with XMM. Section 3 summarizes other data that we  used as constraints for the photoionization modelling. Section 4 describes our modelling strategy, and presents our ``reference model''. Section 5 evaluates the error bars on the derived elemental abundances, taking into account  observational uncertainties in emission-line fluxes, uncertainties in model input parameters as well as uncertainties arising from an imperfect description of the physical processes included in the models. In Section 6, we compare the chemical composition of \\TS\\ with that of other PNe in the Galactic halo and discuss it in terms of stellar  nucleosynthesis in the Asymptotic Giant Branch (AGB) phase. Finally, Section 7 summarizes our main findings. %__________________________________________________________________ ",
        "conclusions": ""
    },
    "0911/0911.3689_arXiv.txt": {
        "abstract": " ",
        "introduction": "Much progress has been achieved in modeling the central stars of planetary nebulae, seen through the progressive refinement by many authors (eg. Harm and Schwartzschild 1975; Sch\\\"{o}nberner 1979, 1981, 1983; Iben 1984; Wood and Faulkner 1986). Together with dynamical models for the evolution of the ejected nebula (Mathews 1966; Sofia and Hunter 1968; Kwok, Purton and Fitzgerald 1978; Kwok 1982), most authors are now modeling PN evolution with a more holistic approach (Okorokv et al. 1985; Schmidt-Voigt and K\\\"{o}ppen 1987a; Stasinska 1989; 2008, Dopita and Meatheringham 1990, 1992; Guenther et al. 2003; Schwarz \\& Monteiro 2006; Holovatyy et al. 2008). This increasingly places more dependence on our understanding of the photoionisation of the nebula and may affect how we interpret diagnostic tools such as excitation class which has been established to give some indication of the central star's effective temperature.  Most meaningful physical parameters that can be derived from a PN's observed emission lines, including ionized and total nebular mass and the brightness and evolutionary state of the central star, depend on accurate distances (Ciardullo et al. 1999). This is difficult in our own Galaxy due to inherent problems with variable extinction and lack of central star homoegeneity (Terzian, 1997). The well determined 50.6~Kpc LMC distance (e.g. Keller \\& Wood 2006), modest, 35~degree inclination angle and disk thickness (only $\\sim500$~pc, van der Marel \\& Cioni 2001), mean that LMC PNe are effectively co-located. Since dimming of their light by intervening gas and dust is low and uniform (e.g. Kaler \\& Jacoby, 1990), we can better estimate absolute nebula luminosity and size. This allows us to test the various methods for deriving excitation class and place this diagnostic on a better footing.  We have constructed the most complete, least biased and homogeneous census of a PNe population ever compiled for a single Galaxy via discoveries from our deep AAO/UKST H$\\alpha$ multi-exposure stack of the LMC's central 25~deg$^{2}$ (Reid \\& Parker 2006a,b; Parker et al. 2005). PNe were confirmed by 2dF spectroscopy which gave 460 new PNe and independently recovered all 169 previously known PNe in the central area. % These data have led to significant advances in our understanding of the PN luminosity function (PNLF) (Reid \\& Parker, 2010) and physical parameters such as temperatures, densities, nebula masses and abundances (Reid 2007). Perfect for excitation studies, the LMC sample provides a large range of intrinsic flux intensities for individual emission lines (eg. 8.8E-12 to 3.7E-16 ergs cm$^{2}$ s$^{-1}$ for \\OIII\\,5007) and interesting differences in line ratios.    Despite all the advantages of using PNe in the Magellanic Clouds for the modeling of dynamical evolution, without deep HST imaging the central stars cannot be directly observed. In order to understand the evolution of these PNe, properties of the central star such as mass and luminosity must be modeled using properties of the nebula. Fortunately, much work has been conducted in this area and it has been convincingly argued that the strength of emission lines in the nebula are uniquely determined by the properties of the central star (Dopita et al. 1987, 1988, 1990, 1992).  Much of this central star modeling depends on some evaluation of the so-called excitation class parameter which is determined from ratios of specific emission lines, measured from the ionised nebula. In this paper, we use the large, new population of PNe in the LMC to compare the three main methods currently used for determining excitation class. Each of these different methods is somewhat grounded in the physics of low to high ionisation states but produce differing results. At last, however, with a large, more representative population of PNe in one system, covering a large evolutionary range, we can observe the different effects produced by these differing excitation schemes.  Having plotted, assessed and compared the existing excitation schemes, each is found to be reflecting different aspects of the nebula-central star relationship, none of which is fully adequate. We suggest a new method for the determination of excitation class. This method provides a smooth transition between low and high excitation regimes and allows both the helium-to-hydrogen and the oxygen-to-hydrogen ratios to contribute a weighted proportion toward the excitation estimate.  Each method is then assessed and compared by constructing diagrams, based on the standard H-R diagram. In this case, excitation class substitutes for central star temperature and the H$\\beta$ luminosity of the nebula represents a rough estimate of stellar luminosity. The diagram constructed using the most correct method should provide an evolutionary map, over which previously determined evolutionary tracks should show some degree of unity if excitation class has any true intrinsic value. ",
        "conclusions": "The utility of the commonly used excitation class estimator for PNe is investigated and appraised using a large, new sample of PNe in the LMC.  The importance of determining the best form for deriving the excitation class is going to be an important part of our on-going study of LMC PNe. The large number of PNe now available over a wide luminosity and evolutionary range has permitted the excitation class parameter to be re-investigated and re-worked. The 3 main methods current in the literature for determining excitation class have been carefully plotted, compared and evaluated. We advise against using the Ex$_\\mathrm{neb}$ scheme in its current form due to the distortion it creates at the medium-to-high excitation region. The Ex$_{\\ast}$ scheme also appears to create some anomalies at the medium excitation levels where the input equation is re-worked. Using (\\OIII\\,$\\lambda4959+\\lambda5007)/\\mathrm{H\\beta}$ for all the PNe produces a smooth transition but suffers for not applying the \\HeII\\,$\\lambda$4686~ratio.  The new Ex$_\\rho$ method offered in this paper appears to improve excitation and stellar temperature estimates by allowing both \\HeII\\,$\\lambda$4686 and \\OIII\\,$\\lambda$5007\\ to become constant variables against the H$\\beta$ intensity. It also minimises any sudden increase in PNe at the point where the \\HeII\\ line is introduced into the derived equation for high-excitation PNe. The method also correlates well with the best published Zanstra temperatures, placing this new excitation class parameter on a far more solid footing.  This presentation has been prepared in order to lay the groundwork for further research using central star temperatures, electron densities, nebula mass and expansion velocities and new photoionisation codes. The disparity between the current methods for assigning excitation classes has been highlighted in this paper. Our suggested Ex$_{\\rho}$ method appears to produce encouraging results as it allows excitation to increase using both the \\OIII/H$\\beta$ and \\HeII\\,$\\lambda$4686/H$\\beta$ CS temperature indicators.  It is clear that excitation is a quality that remains key to the ionised nebulae well into the ageing process and must be correctly modeled. Guided by these results and comparisons, we will continue to model the excitation class based on the ratios of the same ionic species, for example the ratios of O$^{+}$ to O$^{++}$ and He$^{+}$ to He$^{++}$. This method, used already by some authors (e.g. Martin \\& Viegas 2002) is expected to be a good CS temperature indicator. Our preliminary results using the LMC PNe are encouraging. With new photoionisation models, we intend refining the input parameters and anticipate that new evolutionary tracks and diagrams describing the AGB to WD stages will be produced."
    },
    "0911/0911.1339_arXiv.txt": {
        "abstract": "We report the detection of a 115 day periodicity in {\\it SWIFT}/XRT monitoring data from the ultraluminous X-ray source (ULX) NGC 5408 X-1.  Our ongoing campaign samples its X-ray flux approximately twice weekly and has now achieved a temporal baseline of $\\approx 485$ days. Periodogram analysis reveals a significant periodicity with a period of $115.5 \\pm 4$ days. The modulation is detected with a significance of $3.2 \\times 10^{-4}$. The fractional modulation amplitude decreases with increasing energy, ranging from $0.13 \\pm 0.02$ above 1 keV to $0.24 \\pm 0.02$ below 1 keV.  The shape of the profile evolves as well, becoming less sharply peaked at higher energies. The periodogram analysis is consistent with a periodic process, however, continued monitoring is required to confirm the coherent nature of the modulation.  Spectral analysis indicates that NGC 5408 X-1 can reach 0.3 - 10 keV luminosities of $\\approx 2 \\times 10^{40}$ ergs s$^{-1}$. We suggest that, like the 62 day period of the ULX in M82 (X41.4+60), the periodicity detected in NGC 5408 X-1 represents the orbital period of the black hole binary containing the ULX.  If this is true then the secondary can only be a giant or supergiant star. ",
        "introduction": "The existence of black holes in the mass range from $10^{2} - 10^{4}$ $M_{\\odot}$--intermediate mass black holes (IMBH)--is still not widely accepted.  It has been argued based on their extreme luminosities that some of the bright X-ray sources found in nearby galaxies, the ultraluminous X-ray sources (ULXs), may be IMBHs (Colbert \\& Mushotzky 1999), \u0153but it has also been suggested that these objects may appear luminous due to beaming of their X-ray radiation (King et al. 2001). As yet there has been no direct measurement of the mass of a ULX.  At present the best IMBH candidates include X41.4+60, the brightest ULX in the starburst galaxy M82 (also referred to as M82 X-1, Kaaret, Feng \\& Gorski 2009; Strohmayer \\& Mushotzky 2003), and the ULX NGC 5408 X-1 (Strohmayer et al. 2007; Kaaret \\& Corbel 2009).  Very recently, Farrell et al. (2009) have reported the identification of an X-ray source (2XMM J011028.1-460421) very near the absorption line galaxay ESO 243-49. They argue for a physical association with this galaxy which at a redshift of $z = 0.0224$ implies a luminosity in excess of $10^{42}$ erg s$^{-1}$, and suggest it is an IMBH in excess of 500 $M_{\\odot}$. If the association with ESO 243-49 is correct then this object is the most luminous ULX currently known.  X41.4+60 is also amongst the most luminous ULXs, on occasion having a luminosity upwards of $10^{41}$ ergs s$^{-1}$ (Kaaret et al. 2009). This object also shows a 62 day periodic modulation in its X-ray flux, which has been proposed to be the orbital period of the binary system containing the black hole (Kaaret, Simet \\& Lang 2006; Kaaret \\& Feng 2007).  Both NGC 5408 X-1 and X41.4+60 show quasiperiodic oscillations (QPOs) in their X-ray fluxes, and these ULX QPOs appear at systematically lower frequencies than the analogous QPOs observed in Galactic black hole binary systems (Strohmayer \\& Mushotzky 2003; Strohmayer et al. 2007).  Indeed, Strohmayer \\& Mushotzky (2009) have recently shown that the QPO properties in NGC 5408 X-1 correlate with the energy spectrum and X-ray flux in a manner fully consistent with the behavior observed for so-called Type C QPOs in Galactic systems.  The Type C QPOs in stellar mass black holes are strong (fractional rms amplitude $\\sim 15\\%$), relatively coherent ($\\nu_{qpo} / \\Delta\\nu > 10$) oscillations with characteristic frequencies of $\\sim 1 - 10$ Hz that vary in frequency in correlation with source flux and spectral index (see Sobczak et al 2000; Vignarca et al. 2003; Casella, Belloni \\& Stella 2005). They are associated with an energy spectral state in which approximately half or more of the flux is carried by a power-law component with slope $\\Gamma \\sim 2 - 2.5$, commonly referred to as the Steep Power Law state (SPL, McClintock \\& Remillard 2006), or the Hard Intermediate State (HIMS, Belloni 2006). The precise origin of these QPOs is still uncertain, but their phenomenology has been used to estimate the masses of black holes (see Shaposhnikov \\& Titarchuk 2009).  The Type C QPOs in NGC 5408 X-1 identified by Strohmayer \\& Mushotzky (2009) have frequencies of a few tens of mHz, a factor of 100 lower than the Galactic black holes.  Scaling the observed QPO frequencies in NGC 5408 X-1 to those observed in Galactic systems of known mass suggests that NGC 5408 X-1 is an IMBH with a mass greater than 1000 $M_{\\odot}$ (Strohmayer \\& Mushotzky 2009).  Moreover, recent optical spectroscopy reported by Kaaret \\& Corbel (2009) indicates that the optical counterpart to NGC 5408 X-1 is associated with reprocessed emission from an accretion disk as well as a strongly photoionized optical nebula. This, combined with the presence of a powerful radio nebula also appears consistent with the presence of an IMBH (Lang et al. 2007; Soria et al. 2006; Kaaret et al. 2003).  While it is likely that many of the ULXs are accreting black hole binary systems, very little is known about the nature of these putative binaries.  If some of the brighter systems are indeed IMBHs accreting via Roche lobe overflow, then their orbital periods could be as long as a few hundred days (Portegies Zwart et al. 2004).  With the exception of X41.4+60 noted above, very few objects have been monitored with a cadence and duration that would enable sensitive searches for X-ray periods in the range of 10s to 100s of days.  The {\\it Swift} observatory uniquely provides both the X-ray imaging sensitivity and scheduling flexibility to enable such searches. Starting with Observing Cycle 4 {\\it Swift} has been monitoring NGC 5408 X-1 several times per week as part of an approved program.  These observations have continued into Cycle 5 and are presently ongoing. Here we present results from this campaign which provide evidence for a 115 day periodicity in NGC 5408 X-1 which could very likely be the orbital period of a Roche lobe overflow binary containing an IMBH. ",
        "conclusions": "A number of accreting binaries show X-ray modulations at their orbital periods. Among black hole binaries Cyg X-1 and Cyg X-3 are well known examples (Wen et al. 1999; Elsner 1980), and more recent detections include LMC X-3 (Boyd et al. 2001), 1E 1740.7-2942 and GRS 1758-258 (Smith et al. 2002). Many neutron star binaries also show X-ray modulations at the binary orbital period.  The orbital modulations in accreting binaries have typically been attributed to periodic obscuration produced by a vertically and azimuthally structured accretion disk or scattering in an extending accretion disk corona or stellar wind from the donor (Parmar \\& White 1988; Wen et al. 1999). A recent example of an X-ray modulation at the putative orbital period in a neutron star system is that of GX 13+1 (Corbet 2003). Thus, if ULXs are indeed accreting black hole binaries, then it is not unexpected to detect X-ray modulations at their orbital periods.  While their have been several claims of detection of orbital modulations in ULXs (see, Kaaret, Simet \\& Lang 2006 for a brief summary), the most compelling detection to date is that of the 62 day modulation from the M82 ULX X41.4+60 (Kaaret, Simet \\& Lang 2006; Kaaret \\& Feng 2007).  The average modulation in X41.4+60 as observed with the RXTE/PCA is roughly sinusoidal with an amplitude of $\\approx 20\\%$ (Kaaret \\& Feng 2007). This is qualitatively consistent with the $0.2 < E < 8$ keV modulation amplitude and profile that we report here for NGC 5408 X-1, and suggests that similar physical processes may produce the observed modulation in both systems.  Superorbital periods are known in both neutron star and black hole binaries. Well known examples include the 35 day modulation in the neutron star system Her X-1, and the 164 day period in the putative black hole binary SS 433 (Wijers \\& Pringle 1999).  These variations have generally been ascribed to accretion disk precession (Ogilvie \\& Dubus 2001).  Kaaret \\& Feng (2007) argued that the 62 day period in X41.4+60 was unlikely to be a superorbital modulation because the observed period is only consistent with the observed superorbital periods of neutron star systems, but that the extreme luminosity of X41.4+60 comfortably rules out an accreting neutron star (see their Figure 4 for a distribution of observed superorbital periods). Similar arguments would appear valid for the 115 day modulation in NGC 5408 X-1.  The 115 day period is shorter than all known superorbital periods for black hole candidate binaries, and its peak luminosity is well in excess of the Eddington limit for a neutron star ($\\approx 10^{38}$ ergs s$^{-1}$).  Using X-ray timing measurements with {\\it XMM-Newton}, Strohmayer et al. (2009) have recently shown that the strong QPO observed in NGC 5408 X-1 varies in frequency and amplitude with changes in the X-ray flux and energy spectrum in a manner that closely mimics the correlated temporal and spectral variations observed in stellar mass black holes in the so-called Intermediate State (also known as the Steep Power-law State).  Moreover, recent optical spectroscopy reported by Kaaret \\& Corbel (2009) indicates that its optical counterpart has a significant contribution from reprocessing of the X-ray luminosity in the outer parts of the disk.  The so-called ``slim disks'' that may exist at super-Eddington accretion rates preferentially radiate more flux out along the disk axis, and therefore have relatively less available for reprocessing (Abramowicz et al 1988). These models thus have difficulty accounting for the observed optical to X-ray flux ratio in NGC 5408 X-1.  Further, modeling of the observed nebular optical emission lines supports the conclusion that NGC 5408 X-1 radiates $\\sim 10^{40}$ ergs s$^{-1}$ in an approximately isotropic manner (Kaaret \\& Corbel 2009). These observations support the notion that NGC 5408 X-1 harbors a ``standard'' thin, optically thick accretion disk similar to those inferred to exist in Galactic black hole binaries.  The 115 day modulation we have found in NGC 5408 X-1 shows interesting energy dependent effects that would appear consistent with orbital modulation.  A system which shows qualitatively similar effects to those seen in NGC 5408 X-1 is Cyg X-1. Specifically, the orbital modulation in Cyg X-1 has been modeled in the context of absorption and scattering of X-rays in the partially ionized wind from the companion star (Wen et al. 1999). This model predicts a decrease in modulation amplitude with increasing energy, and a minimum in the hardness ratio at the maximum of the modulation profile (see Wen et al. 1999, Figure 3), both of which are evident in the NGC 5408 X-1 data. While the shapes of the modulation profiles in Cyg X-1 do not exactly match the results we find for NGC 5408 X-1, the behavior of the $E > 1$ profile (see Figure 3) appears qualitatively similar.  An important distinction would seem to be that in the case of NGC 5408 X-1 we can directly observe the thermal disk flux whereas the orbital modulation measurements for Cyg X-1 concern the low-hard state, when presumably its thermal flux is weak or absent. Moreover, any disk flux appears largely shortward of the RXTE/PCA energy band.  Cyg X-1 is also a wind accreting system, whereas wind accretion would likely be unable to account for the high luminosity of NGC 5408 X-1.  If the 115 day period is indeed the orbital period of a Roche lobe filling binary, then the mean density of the secondary is constrained to be $\\rho_{mean} \\approx 0.2 (P_{\\rm days})^{-2}$ g cm$^{-3}$ (Frank et al. 2002). For a period of 115.5 days this gives $\\rho_{mean} = 1.50 \\times 10^{-5}$ g cm$^{-3}$, and the companion would have to be a giant or supergiant star.  Recent observations suggest that a substantial fraction of the optical counterpart is due to reprocessing of the X-ray flux in an accretion disk, and while the companion star still has not been observed directly, the recent optical measurements suggest that it is probably a giant in the 3 - 5 $M_{\\odot}$ range, and with a spectral type of B or later (Kaaret \\& Corbel 2009). Binary evolution calculations for IMBHs and massive companions appear consistent with the notion that NGC 5408 X-1 contains a $\\sim 1000 M_{\\odot}$ black hole with a 3 - 5 $M_{\\odot}$ companion (Li 2004). Continued monitoring of NGC 5408 X-1 with {\\it Swift} is important in order to confirm the orbital nature of the modulation."
    },
    "0911/0911.4036_arXiv.txt": {
        "abstract": "{} {Our goal is to understand the nature of blazars and the mechanisms for the generation of high-energy $\\gamma$-rays, through the investigation of the blazar 3C 66A.} {We model the high energy spectrum of 3C 66A, which has been observed recently with the Fermi-LAT and VERITAS telescope. The spectrum has a hard change from the energy range of 0.2-100 GeV to 200-500 GeV in recent almost contemporaneous observations of two telescopes. } { The de-absorbed VERITAS spectrum greatly depends on the redshift, which is highly uncertain. If z=0.444 is adopted, we are able to use the SSC model to produce the Fermi-LAT component and the EC model to the VERITAS component. However, if z=0.1, the intrinsic VERITAS spectrum will be softer, there will be a smooth link between the Fermi-LAT and VERITAS spectra which can be explained using a SSC model. } {} ",
        "introduction": "Blazars are a peculiar class of active galactic nuclei (AGN) and their jets point at small angles with respect to our line of sight. Many of them have been observed at all wavelengths, from radio to very high energy (VHE) $\\gamma$-rays. Their spectral energy distribution (SED) consists of two bumps which are attributed to the synchrotron and the inverse Compton (IC) emission of ultrarelativistic particles. The different soft photon sources deduce synchrotron self-Compton (SSC) and external Compton (EC) models to produce high energy emission. In the SSC model, the soft photons are provided by the synchrotron emission of the same electrons (\\cite{Marscher80}; \\cite{Ghisellini89}; \\cite{Marscher96}). However in the EC model, the soft photons mostly come from the outside of the jet, such as outer disk, broad-line region (BLR) clouds (etc. \\cite{Dermer93}; \\cite{Sikora94}).  In 3C 66A, \\cite{Miller78} have given the redshift z=0.444 by a weak Mg II emission line detection, but it is very uncertain (\\cite{Bramel05}). When 3C 66A locates at z=0.444, its TeV photons will suffer the strong pair production absorption of the Extragalactic Background light (EBL). After corrected by the EBL absorption, TeV emission presents an inverted intrinsic spectrum (see the Acciari et al. (2009) Fig. 2, the de-absorbed photon spectral index is calculated to be $1.1 \\pm 0.4$). In this paper, we take z = 0.444 to analyze TeV emissive mechanism, and will discuss the behavior of the de-absorbed VERITAS spectrum in different redshifts. Generally, 3C 66A is classified as a low-frequency peaked BL Lac object (LBL). The peak of the low-frequency component of LBLs usually lies in the IR or optical regime, whereas the peak of high-energy component locates at several GeV. The luminosity of $\\gamma$-ray is typically comparable to or slightly higher than the luminosity of the synchrotron component. Such as, \\cite{Joshi07} have argued that for 3C 66A the peak of low-frequency component locates at the optical regime, the peak of high-frequency component reaches multi MeV-GeV range. However, \\cite{Perri03} have revealed that the synchrotron peak locates in between $10^{15}$ and $10^{16}$ Hz, then 3C 66A is classified as an intermediate-frequency peaked BL Lac (IBL). From the X-ray spectrum with the photon spectral index $\\Gamma \\sim 2.5$ (Bottcher et al. 2005; Donato et al. 2005; Foschini et al. 2006), which might be the tail of the synchrotron emission, 3C 66A is considered as an IBL in this paper. 3C 66A is observed in radio, IR, optical, X-rays and $\\gamma$-rays, and shows large luminosity variations. As described in \\cite{Bottcher05}, the object exhibits several outbursts in the optical band and the variations of $\\Delta$m $\\sim$ 0.3-0.5 over a timescale of several days. Until now, the majority of BL Lacs detected at VHE (very high energy: E $>$ 100 GeV) are HBLs (high-frequency peaked BL Lacs). Only IBL W Comae (\\cite{Acciari08}), LBL BL Lacertae (\\cite{Albert07}) and 3C 279 (\\cite{Albert08}) display the potential to enlarge the extragalactic TeV source. For 3C 66A, the Crimean Astrophysical Observatory have reported a 5.1$\\sigma$ detection above 900 GeV(Stepanyan et al. 2002).   Recently, VERITAS have carried out 14 hours' observations for 3C 66A from September 2007 through January 2008 (hereafter, the 2007-2008 season), and from September through November 2008 (hereafter, the 2008-2009 season) a further 46 hours' data have been taken (\\cite{Acciari09}). Due to the limited spatial resolution of Cherenkov telescopes it is difficult to accurately identify the emission region. The radio galaxy 3C 66B lies in the same view field of 3C 66A at a separation of 0.12$^{\\circ}$ and is also a plausible source of VHE radiation (\\cite{Tavecchio08}). The recent detection by MAGIC favored 3C 66B as VHE source and excluded 3C 66A at an 85\\% confidence level (\\cite{Aliu09}). However , VERITAS have found that 3C 66A lies 0.01$^{\\circ}$ from the fit position while 3C 66B lies 0.13$^{\\circ}$ away, and 3C 66A is VHE source. If 3C 66A has a redshift of z = 0.444, its de-absorbed spectral index is 1.1$\\pm $0.4 showing very hard intrinsic spectrum (\\cite{Acciari09}). In first three months, the Fermi-LAT Gamma-ray Space Telescope have observed 3C 66A (\\cite{Abdo09}), almost at contemporaneous observation with VERITAS in the 2008-2009 season. However, very soft spectrum with the spectral index of 1.97$\\pm $0.04 appears in the Fermi-LAT observing energy range. The $\\gamma$-ray spectrum suddenly hardens from 0.2-100 GeV to 200-500 GeV and challenges the one-zone homogeneous SSC model.     In Section 2 we present the jet models for application to 3C 66A. We use the observed data to constrain the model parameters in Section 3. We finish with discussions and conclusions in Section 4. Throughout this paper, we use a soft cosmology with a deceleration factor $q_0 = 0.5$ and a Hubble constant $H_0 = 75 km s^{-1} Mpc^{-1}$. ",
        "conclusions": "The redshift of 3C 66A has an uncertain value (\\cite{Bramel05}), and is usually adopted as z = 0.444. However the redshift is crucial in constructing intrinsic high energy spectrum because of the EBL absorption (\\cite{Hauser01}). This absorption decreases the observed flux and softens the observed spectrum. If the redshift is less than 0.444, such as just $\\geq$ 0.096 suggested by \\cite{Finke08}, the intrinsic spectrum in the VERITAS energy range will be softer. There will be a smooth link spectrum between the Fermi-LAT and the VERITAS energy ranges. In the Fig.~\\ref{fig.2}, we generate the intrinsic spectra of VERITAS observations corrected by EBL absorption according to \\cite{Franceschini08} model, assuming the source to be at the different redshifts z=0.03, 0.1, 0.3, and 0.5. It is shown that the de-absorbed spectrum strongly depends on the redshift. When z=0.3, the de-absorbed spectrum has a little inverted, but it becomes an inverted spectrum in z=0.5. If the redshift is less than 0.1, the de-absorbed spectrum will present the usual SED of an LBL such as W Comae (\\cite{Acciari08}). In the Fig.~\\ref{fig.3}, we show the de-absorbed SEDs under z=0.1 and the modeling. A smooth spectrum can link the Fermi-LAT and VERITAS data and be explained with a SSC model, in which $U_B=9.95 \\times 10^{-5} erg \\cdot cm^{-3}$ and $U_e=2.31 \\times 10^{-3} erg \\cdot cm^{-3}$.   In fact the bulk motion of the emitting blob affects the observed SED (e.g., \\cite{Dermer95}; Georganopoulos et al. (2001)). The peak frequencies are given by $\\nu_s\\approx 3.7 \\times 10^6 \\gamma _{peak}^2 \\delta B /(1+z)$ , $\\nu_{SSC}\\approx \\frac{4}{3}\\gamma _{peak}^2 \\nu_{s}$, and $\\nu_{EC}\\approx \\frac{4}{3} \\gamma _{peak}^2 \\nu_{ext} \\Gamma \\delta /(1+z)\\label{nu_ec}$. We can see that $\\nu_{EC}$ would be larger than $\\nu_{SSC}$ provided the blob has larger $\\delta$ or $\\Gamma$ (Usually the viewing angle $\\theta\\sim 1/\\Gamma$ is assumed, $\\delta\\sim \\Gamma$.) From the ratio of peak luminosity, $\\frac{L_{EC}}{L_{SSC}} \\approx \\frac{L_{EC}^{'}}{L_{SSC}^{'}} = \\frac{\\dot\\gamma_{peak,EC}}{\\dot\\gamma_{peak,SSC}} \\approx \\frac{U^{'}_{ext}}{U^{'}_{syn}} \\approx \\delta^4 \\Gamma^2 \\frac{\\xi L_{ext}}{L_s} \\frac{R^2}{R^2_{ext}}$ ($U_{ext}$ is amplified by a factor of $\\Gamma^2$, see the Sikora et al. (1994) and Dermer (1995)), where $\\xi$ is the reproduce fraction of the $L_{ext}$, we show that $\\frac{L_{EC}}{L_{SSC}}$ is strongly affected by the bulk motion of the blob. % In the Fig. 1, using $U_B^{'}=4.6 \\times 10^{-5} erg \\cdot cm^{-3}$, $U_e^{'}=1.67 \\times 10^{-3} erg \\cdot cm^{-3}$, and $U_{ext}^{'} =3.5 \\times 10^{-6} erg \\cdot cm^{-3}$ we can model the SED. $U_{ext}^{'}$ is obviously lower than $U_B^{'}$, however, the EC luminosity is comparable with the synchrotron ones (see the Fig. 1). In the observer frame, the beaming factor is different for EC ($\\delta ^{4+2\\alpha}$ (\\cite{Dermer95}; Georganopoulos et al. 2001)), synchrotron and SSC emission ($\\delta ^{3+ \\alpha}$). The difference of EC and synchrotron luminosity is reasonable by considering their beaming factor.     The EC emission is less clear for the BL Lac objects. The lack of strong emission lines and UV excesses suggests that the external photon density $\\xi L_{ext}$ is very low, while the Lorentz factor of BL Lac objects is typically smaller than that of quasars (\\cite{piner08}). Their high energy emission strongly favors the SSC mechanism over the EC mechanism. But, 3C 66A might be an exception and reveal an existence of larger bulk velocity in the high energy emissive region. Therefore, the high energy emission caused by EC mechanism is likely observed in the IBL. This possibility needs future Fermi-LAT and VERTAS observations for 3C 66A and a precise redshift determination."
    },
    "0911/0911.3774.txt": {
        "abstract": "The velocity of the inner ejecta of stripped-envelope core-collapse supernovae (CC-SNe) is studied by means of an analysis of their nebular spectra. Stripped-envelope CC-SNe are the result of the explosion of bare cores of massive stars ($\\geq 8$ M$_{\\odot}$), and their late-time spectra are typically dominated by a strong [O {\\sc i}] $\\lambda\\lambda$6300, 6363 emission line produced by the innermost, slow-moving ejecta which are not visible at earlier times as they are located below the photosphere. A characteristic velocity of the inner ejecta is obtained for a sample of 56 stripped-envelope CC-SNe of different spectral types (IIb, Ib, Ic) using direct measurements of the line width as well as spectral fitting. For most SNe, this value shows a small scatter around 4500 km s$^{-1}$. Observations ($< 100$ days) of stripped-envelope CC-SNe have revealed a subclass of very energetic SNe, termed broad-lined SNe (BL-SNe) or hypernovae, which are characterised by broad absorption lines in the early-time spectra, indicative of outer ejecta moving at very high velocity ($v \\geq 0.1 c$). SNe identified as BL in the early phase show large variations of core velocities at late phases, with some having much higher and some having similar velocities with respect to regular CC-SNe. This might indicate asphericity of the inner ejecta of BL-SNe, a possibility we investigate using synthetic three-dimensional nebular spectra. ",
        "introduction": "\\label{intro}  Massive stars ($>8$ M$_\\odot$) collapse when the nuclear fuel in their central regions is consumed, producing a core-collapse supernova (CC-SN) and forming a black hole or a neutron star. CC-SNe with a H-rich spectrum are classified as Type II. If the envelope was stripped to some degree prior to the explosion, the SNe are classified as Type IIb (strong He lines, and weak but clear H), Type Ib (strong He lines but no H), and Type Ic (no He or H lines) \\citep{1997ARA&A..35..309F}.  Some CC-SNe, called broad-lined SNe (BL-SNe), exhibit very broad absorption lines in their early phase, resulting from the presence of sufficiently massive ejecta expanding at high velocities. BL-SNe seem to be preferentially of Type Ic [two exceptions are the Type IIb SN 2003bg \\citep{2009arXiv0908.1783H} and the Type Ib SN 2008D \\citep{2008Sci...321.1185M,2009ApJ...702..226M}]; see also SN 1987K \\citep{1988AJ.....96.1941F}, which is mentioned by \\citep{2009arXiv0908.1783H}. Some BL-SNe can reach kinetic energies of $\\ge$ 10$^{52}$ ergs. They are sometimes called hypernovae, and can be associated with long-duration gamma-ray bursts (GRBs) \\citep[see][and references therein]{2006ARA&A..44..507W}. However not all BL-SNe are associated with GRBs.  An important question in the context of CC-SNe is how the gravitational energy is converted into outward motion of the ejecta during the collapse; see \\citep{2007PhR...442...38J} for a recent review. In GRB scenarios a relativistic outflow is launched by the central engine and deposits some fraction of its energy into the SN ejecta, probably preferentially along the polar axis, which might cause strong asymmetries \\citep[e.g.,][]{2002ApJ...565..405M}. The nearest, best-studied GRB-SNe are SN 1998bw / GRB 980425 \\citep{1999A&AS..138..465G}, SN 2003dh / GRB 030329 \\citep{2004cetd.conf..351M}, SN 2003lw / GRB 031203 \\citep{2004ApJ...609L...5M}, and SN 2006aj / GRB/XRF 060218 \\citep{2006Natur.442.1011P}, although it is not fully established that the GRBs (or X-ray flashes) accompanying nearby CC-SNe can be compared directly to high-redshift GRBs. CC-SNe may be characterised by asphericities although a jet does not necessarily form \\citep{2003ApJ...584..971B,2007ApJ...655..416B,2004ApJ...608..391K,2006MNRAS.370..501M,2009ApJ...691.1360T}.  Extremely massive stars ($> 100 $ M$_\\odot$) are thought to end their lives as pair-instability SNe (PI SNe). A star with enough mass to form a He core with more than 40 M$_\\odot$ will suffer electron-positron pair instability, leading to rapid collapse. This triggers explosive oxygen burning, leading to the complete disruption of the star \\citep{1967PhRvL..18..379B,2005IAUS..228..297H}. This process can produce large amounts of $^{56}$Ni. The ejecta mass and the explosion kinetic energy are high, but the ejecta velocities are moderate \\citep{2005ApJ...633.1031S}.  Stripped CC-SNe, which lost part of their envelope before collapse, offer a clearer view of their inner ejecta than SNe which have retained it. Thus, in this paper we exclusively address stripped CC-SNe (Types IIb, Ib, Ic); we do not include Type II SNe, despite using the term ``CC-SN.''  Asphericities in their inner and outer ejecta \\citep[e.g.,][]{2001ApJ...559.1047M} are evident in at least some CC-SNe. Two main indicators are velocity differences of Fe and lighter-element lines, and polarisation measurements \\citep[e.g.,][]{1991A&A...246..481H}.  Independent of their type, with time SNe become increasingly transparent to optical light, as the ejecta thin out. At late times ($> 200$ days after the explosion), the innermost layers of the SN can be observed. This epoch is called the nebular phase, because the spectrum turns from being absorption dominated to emission, mostly in forbidden lines.  In this phase the radiated energy of a SN is provided by the decay of radioactive $^{56}$Co (which is produced by the earlier decay of $^{56}$Ni). Decaying $^{56}$Co emits $\\gamma$-rays and positrons which are absorbed by the SN ejecta. As the deposition rate of $\\gamma$-rays and positrons depends on the density and $^{56}$Ni distribution, the inner parts of the SN dominate the nebular spectra.  Recently, several authors \\citep{2008Sci...319.1220M,2008ApJ...687L...9M,2009arXiv0904.4632T} have studied nebular spectra of CC-SNe. They concentrated on the shape of the [O {\\sc i}] $\\lambda\\lambda$6300, 6364 doublet (which is produced by much of the mass) and concluded that torus-shaped oxygen distributions might cause the double-peaked [O {\\sc i}] profile observed in many CC-SNe nebular spectra. However, there is ongoing discussion regarding whether geometry is the dominant reason for this type of line profile \\citep[][also see Appendix A]{2009arXiv0904.4256M}.  Several authors have modelled nebular-phase spectra of SNe to derive quantities such as the $^{56}$Ni mass \\citep[e.g.,][]{2004ApJ...614..858M,2006A&A...460..793S,2006MNRAS.369.1939S,2007ApJ...658L...5M}, ejecta velocities \\citep[e.g.,][]{2007Sci...315..825M}, asphericities \\citep[e.g.,][]{2005Sci...308.1284M,2007ApJ...670..592M,2008Sci...319.1220M}, and elemental abundances \\citep[e.g.,][]{2007ApJ...658L...5M, 2007ApJ...666.1069M}. The nebular phase is especially suitable for studying the core of SNe. If different explosion scenarios are involved for different types of CC-SNe, one might expect the largest, most revealing differences to be in the central region of the explosion. Therefore, in contrast to the standard classification of BL-SNe, which is based primarily on early-time spectroscopy and describes the velocity field of the outer SN layers, here we focus on the centre of the explosion. We have modelled the nebular spectra of over 50 SNe, the largest sample of CC-SNe so far, and obtained a statistically significant representation of their core velocities.  We describe our data set in Section 2 and the modelling procedure in Section 3. In Section 4 we test the reliability of the modelling approach. Results are discussed in Section 5. ",
        "conclusions": "\\label{conc}  We estimated $^{56}$Ni masses for 29 SNe for which $^{56}$Ni masses were not previously known. Their average agrees with the average of 18 CC-SN $^{56}$Ni masses estimated before rather well, however there might be a small systematic over-estimate of the $^{56}$Ni mass in this work. Individual estimates might be quite erroneous given the poor quality of the data set. We then measured characteristic velocities for 56 CC-SN cores, which range from 3000 km s$^{-1}$ to 7000 km s$^{-1}$ ($v_{\\alpha}$). Several BL-SNe with high-velocity outer ejecta have high-velocity cores as well, but some BL-SNe do not. We have shown that this might be due to ejecta asphericities, which are expected theoretically and might have been detected by different methods before. We found that the average core velocities of SNe~IIb are slightly lower than core velocities of SNe~Ib and SNe~Ic. SNe~Ib and Ic have very similar average core velocities. SNe~IIb show much less variance ($\\sigma_v \\approx 400$ km s$^{-1}$) of their core velocities than SNe~Ib and Ic ($\\sigma_v \\approx 850$ km s$^{-1}$).  There seems to be no strong dependence of core velocity on $^{56}$Ni mass. We also checked for correlations between total ejecta mass or kinetic energy and the core velocity and found none. There is only a weak correlation between the ratio of total kinetic energy to total mass and the core velocity (BL-SNe have the highest core velocities on average, but the scatter is so large that there is almost no predictive power). Therefore, although the total mass of a SN might be estimated (if the spectrum is flux calibrated), it is not possible to estimate the SN total kinetic energy from a nebular spectrum with good accuracy, since the core velocity seems to correlate with the outer ejecta velocities only weakly. We will further study the relation between outer and inner ejecta velocities in future work.  The uncertainties in our estimates are rather large. They are caused by the general properties of our method (background subtraction, reddening, time evolution of line width, uncertainty on epoch, $^{56}$Ni distribution, and $^{56}$Ni mass) and are difficult to improve upon when only limited high-quality data are available that cover different CC-SNe at several different epochs. Therefore more and better data is needed."
    },
    "0911/0911.4493_arXiv.txt": {
        "abstract": "We present recent results of 3D magnetohydrodynamic simulations of neutron stars with small misalignment angles, as regards the features in lightcurves produced by regular movements of the hot spots during accretion onto the star. In particular, we show that the variation of position of the hot spot created by the infalling matter, as observed in 3D simulations, can produce high frequency Quasi Periodic Oscillations with frequencies associated with the inner zone of the disk. Previously reported simulations showed that the usual assumption of a fixed hot spot near the polar region is valid only for misalignment angles $\\mis$ relatively large. Otherwise, two phenomena challenge the assumption: one is the presence of Rayleigh-Taylor instabilities at the disk-magnetospheric boundary, which produce tongues of accreting matter that can reach the star almost anywhere between the equator and the polar region; the other one is the motion of the hot spot around the magnetic pole during stable accretion. In this paper we start by showing that both phenomena are capable of producing short-term oscillations in the lightcurves. We then use Monte Carlo techniques to produce model lightcurves based on the features of the movements observed, and we show that the main features of kHz QPOs can be reproduced. Finally, we show the behavior of the frequencies of the moving spots as the mass accretion rate changes, and propose a mechanism for the production of double QPO peaks. ",
        "introduction": "\\label{sec:intro}  Quasi Periodic Oscillations (QPO) have been observed in many X-ray sources for the last 24 years \\citep{vanderKlis:1985p2318,Lamb:1985p2327}. They appear as broad peaks in the power spectra, with a behavior dependent on the source state, the mass accretion rate, and other physical properties of the sources. They are present in both neutron star (hereafter NS) and black hole (BH) sources, sometimes with peculiar similarities between the two classes of objects \\citep[see][for a review]{2006csxs.book...39V}.  Of particular interest are the kHz QPOs: they were first discovered in the Low Mass X-ray Binary (LMXB) and Z source Sco-X1 \\citep{1996ApJ...469L...1V}, to be then observed in the spectra of nearly all Z and atoll sources, as well as several weak LMXBs \\citep[see for example][]{Strohmayer:1996p2317,Wijnands:1997}. They have frequencies between 300Hz and 1.2kHz \\citep[see][]{2006csxs.book...39V}. The peaks move with timescales of hours/days and their quality factor $Q=\\nu/\\Delta\\nu$ (ratio between the frequency and the width of the peak, a measure of coherence) is not constant, but can be related with the QPO frequency \\citep{Mendez:2001p4640, diSalvo:2003p5128,Barret:2006p7645} and with the mass accretion rate \\citep[see for example][]{Mendez:2006p5133}.  The kHz QPOs usually appear as a pair of peaks (the {\\em twin peaks}), labeled ``upper'' and ``lower'' (and their frequencies, accordingly, $\\nu_u$ and $\\nu_l$).  While on timescales of hours the two frequencies change, their difference $\\Delta\\nu=\\nu_u-\\nu_l$ tends to be almost costant. As with the discovery of burst oscillations \\citep{Strohmayer:1996p2317} and Accreting Millisecond X-ray Pulsars (hereafter AMXP) \\citep {Wijnands:1998p1343} the spin frequencies $\\spinf$ of more and more accreting neutron stars were measured, an interesting behavior was observed: $\\Delta\\nu$ was usually found to be close to either $\\spinf$ or $0.5\\spinf$. Recent works \\citep[see for example][]{Mendez:2007p5756} call this assumption into question. In particular, it can be shown that $\\Delta \\nu$ could in principle be independent from $\\spinf$.  It has been shown that the kHz QPOs frequency is correlated, at least on a given system and for short timescales, with the X-ray count rate \\citep{1999ApJ...511L..49M} and thus also, probably, with the mass accretion rate.  Many models have been developed to explain the kHz QPO phenomenon. They involve a number of different mechanisms, like for example relativistic precession frequencies \\citep{Stella:1998p4748}, relativistic resonances \\citep{Kluzniak:2001p4749}, beats between the keplerian frequency at a particular radius in the disk and the frequency of the star \\citep{Alpar:1985p2315,Lamb:1985p2327,Miller:1998p4107,Lamb:2004p6829}, oscillations in the comptonizing medium \\citep{Lee:1998p4811}. Recently, \\citet{Lovelace:2007p4213} and \\citet{Lovelace:2009p4323}  found and analyzed a radially localized Rossby wave instability which exists in the inner region of a disk around a magnetized star where the angular rotation rate of the matter decreases approaching the star.  This instability may explain the twin kilo-Hertz QPOs observed in some LMXBs. The importance for QPOs of the disk angular rotation rate having a maximum was discussed earlier by \\citet{Alpar:2005,Alpar:2008}  Most of the models of kHz QPOs, based on the fact that QPOs also appear in black hole systems and there seems to be a link between $\\nu_l$ with frequencies observed in black hole systems \\citep{Psaltis:1999p4878}, try to explain the behavior of kHz QPOs in both the neutron star and the black hole cases, excluding the possibility of a surface emission. \\citet{2005AN....326..812G} claim that, according to spectral emission properties of QPOs, in Atoll and Z sources periodic and quasi-periodic emission is more likely to be produced on the surface of the star.  During the past years MHD simulations have been amongst the most fruitful methods to investigate accreting magnetized stars \\citep[see for example][]{Goodson:1997p5065, Romanova:2002p4187, Romanova:2003, Romanova:2004p4095, Kulkarni:2005p4093, Bessolaz:2008p4909, Long:2008p4237, Romanova:2008}. 3D simulations have given insight into the way matter accretes, showing in more detail trajectories, shapes of magnetic field lines, shapes and movements of hot spots on the surface of the star \\citep[see for example][]{Romanova:2003,Romanova:2004p4095}.  Simulations show that matter may accrete in several different regimes: \\begin{itemize} \\item the {\\em stable accretion} regime, when matter is stopped by the magnetic field and conveyed to the magnetic poles via the so called {\\em funnel flow} \\citep{Romanova:2002p4187,Romanova:2003,Kulkarni:2005p4093}; \\item the {\\em unstable} regime, when Rayleigh-Taylor instabilities form at the boundary between the disk and the magnetosphere \\citep{Romanova:2006p2479,Romanova:2008,Kulkarni:2008}, and accretion takes place through tongues of matter from the inner edge of the disk to random places between the equator and the magnetic poles; \\item the {\\em magnetic boundary layer} regime, during which the disk is very close to the surface of the star and the accretion takes place via two big antipodal equatorial tongues which rotate with the frequency of the inner disk \\citep{Romanova:2009p6603}. \\end{itemize}  One of the interesting phenomena observed in 3D simulations is the movement of spots around the magnetic pole of the accreting star, both in stable \\citep[][]{Romanova:2003,Romanova:2004p4095,Romanova:2006p3862} and in unstable regime \\citep{Kulkarni:2009}. It was noticed that in some cases, such as in the magnetic boundary layer regime, the moving spots  produce QPO features \\citep{Romanova:2009p6603}. Usually, models of NS surface emission (like in the case of Accreting Millisecond Pulsars) consider funneling of matter onto the polar caps as a static process, with a fixed polar cap and thus a fixed emission spot \\\\citep[e.g.][]{Beloborodov:2002p4173,Poutanen:2003p4169}. 3D simulations show that this may not be the case in particular for small misalignment angles.  Recent works \\citep{Lamb:2009A,Lamb:2009B} investigate the wandering of hot spots at which matter accretes in sources with a small misalignment angle, finding that it might affect the coherence of the pulsed signal from neutron stars and even destroy it, being the origin of the transient emission in AMXPs. The movements hypothesized in those works, though, are different from the movements observed in our 3D simulations, where the motion of the hotspots is a rotational motion around the magnetic pole or in the equatorial region, on timescales of milliseconds. Moreover, the dimension of the hotspot in those two works (a radius of $\\sim 25^{\\circ}$) does not find confirmation in our simulations. It corresponds to the dimension of the whole ``hot ring'' where our hotspots (much smaller) move (see \\fref{fig:sf-stab-join}).  In this paper we show that the observed movements carry with them the information of disk frequencies near the inner disk, and their effect is the production of quasi-periodic oscillations at the relative frequencies. We perform a series of 3D MHD simulations (mostly for the cases of very small $\\mis$) and analyze the rotation of the spots and the related frequencies in detail. MHD simulations can only follow the motion of the disk on very short timescales compared to the disk evolution timescale and the observational time. For this reason, we use the features of the moving spots produced in simulations to generate long lightcurves by means of Monte Carlo techniques, showing that the motion observed produces QPOs. Then, we show that the phenomenon follows a predictable behavior as the mass accretion rate changes.  We start in \\sref{sec:model} with a general description of the 3D MHD model used for this work. Then, we show the way spots move in our simulations in \\sref{sec:3d}. We then show with a Monte Carlo-generated lightcurve our model of surface QPO production in \\sref{sec:mc}. Finally, in \\sref{sec:mdot} we show how the angular velocities of hot spots correlate with the mass accretion rate. ",
        "conclusions": "We developed a model of quasi-periodic oscillations based on the spot movements on the surface of the accreting neutron star with a slightly misaligned dipole magnetic field.  3D MHD simulations have shown that spots may rotate faster/slower than the star in cases of stable or unstable accretion  \\citep[see also][]{Romanova:2003,Romanova:2004p4095,Kulkarni:2008,Kulkarni:2009} and in the magnetic boundary layer regime \\citep{Romanova:2009p6603}. We have shown that in all above cases, the spot movements develop ordered oscillations in the light-curves. Modeling these oscillations as short-lived sinusoidal trains, we showed that if they are powerful enough to be seen, their fingerprint is a QPO. Given the range of values of the frequencies involved, we hypothesize that kHz QPOs could be produced in this way.  There are several features of QPOs that can give confirmation to this model. First of all, the general trend with the mass accretion rate, as shown in \\sref{sec:mdot}, is as it is expected. Though it is not generally true that luminosity and mass accretion rate are proportional to each other, we can assume it to be true on a given source and for a short timescale. This is the case of the simulations in this paper, where only one parameter (the mass accretion rate) changed between the various simulations.  The rms amplitude of real life QPOs has a strong dependence on energy \\citep {Berger:1996p4473}, suggesting that the mechanism of emission is different from that of the background. In our opinion, this is because the QPOs are produced on the surface, or perhaps just above the surface by means of shocks in the accretion column similarly to what happens in AMXPs \\citep[see for example][]{Gierlinski:2002p5829},  while most of the background is produced by the disk. The rms amplitude is also generally higher in atoll sources (up to 20\\%) than in Z sources (2-5\\%) \\citep{Jonker:2001p4470}. The reason for this difference is probably related to the higher accretion rate of Z sources.  In some cases  3D simulations show that  two frequencies appear at the same time. This numerical discovery is important for possible interpretation of two QPO peaks observed in a number of LMXBs. In fact, this mechanism for the production of multiple QPOs is able to explain more than the mere appearance of two peaks in the Power Spectrum. When double QPOs appear in observations, they normally present different behavior as regards their $Q$ factor \\citep{diSalvo:2003p5128,Barret:2006p7645}, and hence (according to our assumption) their coherence time with the upper peak generally less coherent; in our model, this is due to the fact that the lower one comes from a smooth motion around the magnetic pole, while the upper one is produced by instability tongues, shorter-lasting and thus with a lower coherence. The term ``unstable'' should not mislead the reader though: the unstable regime is very common in simulations, expecially for fast rotators, also at small values of the mass accretion rate, as shown in \\fref{fig:mdot18}. A higher coherence of the upper kHz QPO, in our model, could in principle be produced by a very high accretion rate and the onset of the magnetic boundary layer regime, in which instabilities produce long-lasting opposite tongues in the equatorial plane. During this regime though, it is very probable that the zone around the surface of the neutron star is optically thick. Just as observations show, multiple QPOs can appear simultaneously or at different times.  As regards the models predicting the upper limit of $\\nu_u$ as the Keplerian frequency at the ISCO \\citep{Miller:1998p4107}, it would be interesting to investigate the matter further. For the values of mass, magnetic field and radius used in the present job, the ISCO is never reached by the inner disk during QPO production with the two channels described because of the onset of the magnetic boundary layer regime. Further investigation should include the use of more extreme values of the mass in order to study the behavior of the inner disk when the ISCO is considerably far from the surface.   Another interesting point is the $\\nu-L_x$ correlation observed in many papers \\citep[for example][]{Yu:1997p4754,Mendez:1998p6564}. The correlation only holds for a given source and on time scales shorter than a few hours. The main reason why it does not hold across sources can easily be found in the magnetic field: for a given value of the mass accretion rate, different values of the magnetic field correspond to a different size of the magnetosphere and, as a consequence, to different values of the frequency of the moving spots. Again, this is evident from our simulations: the exact same frequency in figure \\ref{fig:mdot15} is obtained for $\\mdot\\approx4\\cdot 10^{-11}\\msun y^{-1}$ in a system with $B=10^8$G and for $\\mdot=10^{-9}\\msun y^{-1}$ in a system with $B=5\\cdot10^8$G. But the simulations presented here show only one of the many possible configurations of the disk around the star. For a given source, it is possible that changes in the structure (density distribution for example) of the disk produce a different behavior of the moving hot spots, even with the same mass accretion rate. Besides, the relationship between luminosity and mass accretion rate is not clear. This point will be investigated more deeply in future papers, by simulating systems similar to the one shown here but with different initial disk conditions.  The results presented here are based on simulations of neutron stars with very small misalignment angles. An equivalent discussion can be made at slightly larger misalignment angles, where the behavior of the spots does not change considerably (see \\fref{fig:spotvsmis}, but also \\citealt{Kulkarni:2008,Kulkarni:2009,Romanova:2008}), and obtain almost the same features in the power spectrum. In this case though, from simple geometrical considerations it is reasonable to expect some kind of modulation at the frequency of the star. The least we can expect is a broad Lorentzian-like feature at the frequency of the star \\citep{Burderi:1993p328,Lazzati:1997p8008,Menna:2002p949} just from the action of the rotating spot in the polar region, due to the rotation of the star. Its quality factor $Q$ would be comparable to that of the moving hotspots ($Q_*/Q_{spot} \\approx {\\spinf/\\nu_{spot}}$, according to the considerations of \\sref{sec:q}). The observation of such a feature, even very dim, in observations where the lower kHz QPO (the one from the funnel flow in this model) is present, would give interesting evidence in favor of this model.  As a final remark, we notice that the difference between the frequencies of the two QPOs is $200-300$Hz, close to the frequency of the star. It is generally known that for relatively slow rotators the frequency difference of QPOs is similar to the frequency of the star. Until about two years ago, it seemed that the $\\Delta\\nu$ was approximately equal to $\\spinf$ for slow rotators and $0.5\\spinf$ for fast rotators. \\citet{Mendez:2007p5756}, analyzing the matter in detail, find that this distinction is not so clear, and that data can be interpreted and fitted equally well by values of $\\Delta\\nu$ around $300$Hz for all sources, with the caveat \\citep{vanStraaten:2005p4594} of a few AMXPs showing all variability components (and thus also $\\Delta\\nu$) shifted down by a factor of $\\sim 1.5$. The range of values of $\\Delta\\nu$ found in this work are compatible with both hypotheses, because the simulated system is a slow rotator and $\\Delta\\nu$ is both near the frequency of the star and $300$Hz. We are planning to investigate the matter with simulations of faster rotators. Nevertheless, this mechanism for the production of double QPOs does not need the two frequencies to be origined from one another, giving a fixed value of $\\delta\\nu$, which was the main caveat of a somewhat similar model, the sonic point beat frequency model by \\citet{Miller:1998p4107}.  From what we observe in this work, the two QPOs are two separate effects of the interaction between the rotating magnetosphere and the disk. The funnel flow comes from a zone just behind the inner radius where the magnetic field is strong enough to capture matter and make it drift towards the pole. The instabilities form right at the inner radius, where magnetic field energy density and ram pressure are equal (the green line in \\fref{fig:sf-stab-join}). The results of this paper can give insight into the matter: the frequencies found in observations are a probe of the interaction between the magnetic field and the disk."
    },
    "0911/0911.1878_arXiv.txt": {
        "abstract": "The virtual observatory (VO) is a collection of interoperable data archives, tools and applications that together form an environment in which original astronomical research can be carried out. The VO is opening up new ways of exploiting the huge amount of data provided by the ever-growing number of ground-based and space facilities, as well as by computer simulations. This presentation summarises a variety of scientific results spanning various fields of astronomy, obtained thanks to the VO, after highlighting its structure, infrastructure and various capabilities. ",
        "introduction": "\\label{sec:intro}  A Virtual Observatory (VO) is a collection of inter-operating data archives and software tools which utilize the Internet to form an environment in which original research can be conducted. The VO is an international astronomical community-based initiative, whose main goal is to allow transparent and distributed access to data available worldwide. This is achieved by developing and applying common standards and by ensuring the interoperability between the various data collections, tools and services.  This allows scientists to discover, access, analyze, and combine observational and model data from heterogeneous data collections in a coherent and user-friendly manner. ",
        "conclusions": "\\label{conclude}  The scientific usage of the VO tools and infrastructure is now taking up, and this is reflected by the increasing number of refereed and other publications making use of the VO capabilities. The EURO-VO and other national VO initiative are here to help interested groups to carry out VO-enabled research and provide support for projects that want to publish their data in the VO. Dedicated initiatives such a workshops and Schools are organised world-wide in order to make astronomers aware of the grand potential of the VO and the scientific possibilities it offers."
    },
    "0911/0911.3006.txt": {
        "abstract": "We present Integral Field Spectroscopy (IFS) of NGC~595, one of the most luminous \\hii\\ regions in M33. This type of observations allows us to study the variation of the principal emission-line ratios across the surface of the nebula. At each position of the field of view, we fit the main emission-line features of the spectrum within the spectral range 3650-6990~\\AA, and create maps of the principal emission-line ratios for the total surface of the region. The extinction map derived from the Balmer decrement and the {\\it absorbed} \\ha\\ luminosity show good spatial correlation with the 24~\\mi\\ emission from {\\it Spitzer}. We also show here the capability of the IFS to study the existence of Wolf-Rayet (WR) stars, identifying the previously catalogued WR stars and detecting a new candidate towards the north of the region. The ionization structure of the region nicely follows the \\ha\\ shell morphology and is clearly related to the location of the central ionizing stars. The electron density distribution does not show strong variations within the \\hii\\ region nor any trend with the \\ha\\ emission distribution. We study the behaviour within the \\hii\\ region of several classical emission-line ratios proposed as metallicity calibrators: while [\\nii]/\\ha\\  and [\\nii]/[\\oiii] show important variations, the \\rdostres\\ index is substantially constant across the surface of the nebula, despite the strong variation of the ionization parameter as a function of the radial distance from the ionizing stars. These results show the reliability in using the R$_{\\rm 23}$ index to characterize the metallicity of \\hii\\ regions even when only a fraction of the total area is covered by the observations. ",
        "introduction": "Emission-line spectra of Galactic and extragalactic HII regions are normally obtained to characterize their physical properties, to extract information of the stellar populations that ionize them \\citep{Osterbrock:2006p498} as well as to study the variation of the physical properties of the star-forming regions across the galaxy disc (e.g. \\citealt*{McCall:1985p496}; \\citealt{Vilchez:1988p505}; \\citealt{VilaCostas:1992p506}). The \\hii\\ regions are far from the idealized spherical ionized gas clouds with homogenous physical properties, as has clearly been shown in detailed studies of the nearest Galactic and extragalactic \\hii\\ regions: e. g. Orion Nebula (\\citealt{Baldwin:1991p116}; \\citealt{Kennicutt:2000p110}; \\citealt{Sanchez:2007p501}), 30~Doradus (\\citealt{Mathis:1985p115}; \\citealt{Kennicutt:2000p110}) and NGC~604 (\\citealt{GonzalezDelgado:2000p486}; \\citealt{MaizApellaniz:2004p493}), among others. The derived physical properties, generally obtained from long-slit spectroscopic measurements centred on the most intense knots, are not necessarily representative of the conditions in all the locations within the regions (\\citealt{Oey:2000p124}; \\citealt{Oey:2000p121}). The variations of the physical properties are especially seen in star-forming regions where the action of the massive stars strongly influence the surrounding medium (e.g. \\citealt*{Kobulnicky:1996p489}, 1997). Moreover, in some cases the gas within the \\hii\\ regions is ionized by multiple star clusters which are not always centrally located. This will produce different temperature and ionization structures than the normally assumed configurations from spherical models (\\citealt*{Ercolano:2007p482}; \\citealt{Jamet:2008p570}).  In order to study the variations of the physical properties within the \\hii\\ regions, the usual technique carried out up to now consists of a set of spectroscopic observations with slits located at different positions in the region and covering as much surface as possible. This technique has limitations however: it requires a large amount of observation time and does not cover, in general, the complete face of the region; thus interpolations of the observational parameters in the gap between slits need to be made. Integral Field Spectroscopy (IFS) overcomes these limitations and offers the opportunity of extracting spectroscopic information at different positions across a continuous field of view. Therefore, this kind of surveys allows the study of the variation of the physical properties within star-forming regions.  We present here IFS observations of NGC~595, the second most luminous \\hii\\ region in M33. NGC~595 presents an angular size of $\\sim$1 arcmin, which at the distance of M33 (840~kpc; \\citealt*{Freedman:1991p485}) translates into a linear physical size of $\\sim$200~pc. This makes the region particularly suitable for being studied with the IFS technique. Although the current IFS facilities usually have a relatively small field of view, the proximity of this region (which implies a relatively high surface brightness) makes possible to map it in a reasonable amount of observing time.   NGC~595 has an \\ha\\ shell morphology that shows the action of the stellar winds of the massive stars located in its interior. The stellar content studied by \\citet*{Malumuth:1996p494} using optical photometry reveals the existence of $\\sim$250 OB-type stars,  $\\sim$13 supergiants and 10 candidate Wolf-Rayet (WR) stars. These authors derived an age of 4.5 $\\pm$ 1.0~Myr for the stellar cluster, which is consistent with the age predicted from the synthesis of integrated spectra in the far-ultraviolet (FUV) wavelength range \\citep{Pellerin:2006p499}. Recently, \\citet{Drissen:2008p481} spectroscopically confirmed nine of the WR candidates previously identified by \\citet*{Drissen:1993p105} using photometric observations and discarded two possible WR candidates from their sample. The physical properties for NGC~595 have been derived with long-slit spectroscopic observations located at the most intense knots within the region \\citep{Vilchez:1988p505} and recently with echelle spectroscopy \\citep{Esteban:2009p563}. \\citet{Vilchez:1988p505} obtained a temperature of T$_{\\rm e} \\sim$ 8000~K, an electron density consistent with the low density limit and a metallicity of 12 + log[O/H] = 8.44 $\\pm$ 0.09 (Z = 0.6\\zsun\\footnote{We assume solar abundance from \\citet*{Asplund:2005p404}, 12 + log[O/H] = 8.66}), which was recently confirmed by \\citet{Esteban:2009p563}. The extinction within the region has been studied previously by other authors: \\citet*{Bosch:2002p478} obtained the extinction using the \\ha/\\hb\\ emission-line ratio and \\citet*{Viallefond:1983p504} derived the extinction with the \\ha\\ and thermal radio emission ratio. \\citet{Relano:2009p558} revised the extinctions using radio-to-\\ha\\ emission and derived new values using the 24~\\mi-to-\\ha\\ integrated emission ratio, finding consistent results from both methods. They also found differences in the emission distributions at several wavelength ranges: while the infrared emission at 24~\\mi\\ and \\ha\\ are spatially correlated with each other, the FUV emission is located within the observed \\ha\\ shell structure of the region. They suggest that the dust emitting at 24~\\mi\\ will probably be mixed with the ionized gas of the region and will be heated by the central ionizing stars.  In this paper, we analyse IFS observations covering the whole face of NGC~595. We are able to perform a detailed analysis of the principal emission-line ratios which are widely used to describe the main properties of \\hii\\ regions such as extinction and density structure, ionization parameter, metallicity, existence of shocks and evolutionary state. With these observations, we are in a position to test the influence of the geometry and the stellar distribution within the region on the integrated fluxes of the emission lines, as has been predicted by the models. The comparison of our results with those obtained from long-slit spectroscopy makes it possible to estimate the bias of the observations when only one part of the region, generally the most intense, is covered. Finally, the power of the IFS and the quality of the observations presented here allow us to analyse the WR stellar content of the total surface of the region and to make a complete census of this stellar population in NGC~595. ",
        "conclusions": "We present IFS of NGC~595, one of the most luminous \\hii\\ regions in M33, covering an unprecedented area in the disc of the spiral galaxies of the Local Group, a $\\sim$174$\\times$340~parsec$^2$ field of view representing the complete surface of the region. Taking advantage of the power of these observations, we are able to identify and catalogue the WR population of the region and to make the best census of WR stars in NGC 595 up till now. We have analysed the variations within the region of the main emission-line ratios that describe the physical properties of the region. The analysis of the observations yields the following results.  \\begin{itemize}  \\item We present the properties of NGC~595 derived from the integrated spectrum of the region. The electron density and the metallicity are consistent with previously reported values in the literature using long-slit spectroscopy, covering just the most intense knots in the region.  \\item The extinction map obtained from the \\ha/\\hb\\ emission-line ratio at each \\emph{spaxel} in the field of view presents a concentrated distribution with the maximum located at the centre of the \\ha\\ shell structure. We have compared this map with the {\\it Spitzer} 24 and 8~\\mi\\  bands using elliptical radial profiles. The 24~\\mi\\ emission and the \\ha/\\hb\\ radial profiles are similar, with their maxima located at the same distance from the ionizing stars. The maximum of the 8~\\mi\\ emission radial profile shows a displacement of $\\sim$4-5\\arcsec\\ (15-18~pc) with respect to the maximum of the \\ha/\\hb\\ radial profile. This shows that the extinction suffered by the gas is produced by dust emitting at 24~\\mi, which is probably mixed with the ionized gas within the region, while the dust emitting at 8~\\mi\\ is more related to the outer parts of the region where the molecular cloud is located.  \\item The [\\sii]$\\lambda$6717/[\\sii]$\\lambda$6731 emission-line ratio map does not show any structure, implying that the electron density is quite constant within the region, despite the pronounced \\ha\\ shell morphology of NGC~595. The values of [\\sii]$\\lambda$6717/[\\sii]$\\lambda$6731 are consistent with the low electron density limit.  \\item We have produced BPT diagrams using the emission-line ratios at each \\emph{spaxel} in the field of view. We find that the \\emph{spaxels} with a low \\ha\\ flux cover the total excitation range of the region. The location of the \\emph{spaxels} with a high \\ha\\ flux in the [\\oiii]/\\hb-[\\sii]/\\ha\\ diagram is displaced with respect to the position given in the diagram for the integrated values. This shows the limitations of the long-slit observations, which usually cover the most intense \\ha\\ knot in the regions, to obtain a reliable value for the [\\sii]/\\ha\\ emission-line ratio.  \\item The study of the WR population for NGC~\u00ca595 has been completed with these observations. We have covered the total surface of the region and found a new WR candidate whose spectrum is characteristic of a WN star. The WR candidate, situated far away from the other WR stars in the region, is close to the \\ha\\ shell but corresponds to a zone of low \\ha\\ emission.  \\item  We have analysed the behavior of the most used metallicity calibrators: \\rdostres\\ shows small variations of 0.13~dex despite the spatial distribution of the stars and gas in the region and the strong trend of [\\oiii]/[\\oii] to decrease to the outer parts of the region. This shows the robustness of \\rdostres\\ to estimate the metallicity of the region; in spite of the \\hii\\ region morphology, far away from the classical Str\\\"omgren sphere picture, the \\rdostres\\ parameter varies slightly within the region and thus can be used to obtain a representative value of the \\hii\\ region metallicity. In contrast, other parameters such as [\\nii]/\\ha\\ and [\\nii]/[\\oiii] show strong variations within the region (up to one order of magnitude in the case of [\\nii]/[\\oiii]), which show their strong dependence on the ionization parameter. \\end{itemize}  We show in this paper the power of the IFS, which overcomes the limitations of long-slit observations in \\hii\\ regions. The results presented here show that for the emission lines dominated by the contribution of the most intense knots, long-slit observations are representative of the integrated emission of \\hii\\ regions. The situation is different in the case of lower excitation emission lines, as it happens in the [\\sii] emission lines, for which there is a significant bias for long-slit observations. This bias has only been possible to quantify using IFS observations with a complete \\hii\\ region coverage.  The physical properties can vary within \\hii\\ regions and the emission-line ratio maps derived from IFS observations allow us to make a detailed study of these variations. For NGC~595 we find that the ionization parameter shows strong variations with the radial distance from the stars; the ionization structure is clearly depicted with the corresponding emission-line diagnostics and the reddening map presents a non-uniform distribution with a maximum located at the center of the observed \\ha\\ shell structure. The electron density map, however, does not present significant variations within the region.  A change of one order of magnitude in [\\oii]/[\\oiii] within NGC~595 allows us to study the dependence of the ionization parameter on the most widely used metallicity calibrators, since other properties such as density or the ionizing spectrum do not change within this \\hii\\ region. The main result is that \\rdostres\\ does not depend significantly on the ionization parameter, but other emission-line ratios such as [\\nii]/\\ha\\ and [\\nii]/[\\oiii] show important variations within the region.  Finally, the capability of the IFS data to make complete censuses of the WR population in stellar clusters is very nicely shown here: while the hunting of WR stars are based on identification in broad-band images with follow-up spectroscopic observations, we are able to perform such an analysis in only one night of observation, covering the whole surface of the region and allowing us to identify WR stars that are not located close to the centre of the cluster."
    },
    "0911/0911.1080_arXiv.txt": {
        "abstract": "From observations collected with the ESPaDOnS spectropolarimeter at the Canada-France-Hawaii Telescope (CFHT), we report the detection of Zeeman signatures on the low-mass classical T~Tauri star (cTTS) V2247~Oph.  Profile distortions and circular polarisation signatures detected in photospheric lines can be interpreted as caused by cool spots and magnetic regions at the surface of the star.  The large-scale field is of moderate strength and highly complex;  moreover, both the spot distribution and the magnetic field show significant variability on a timescale of only one week, as a likely result of strong differential rotation.  Both properties make V2247~Oph very different from the (more massive) prototypical cTTS BP~Tau;  we speculate that this difference reflects the lower mass of V2247~Oph.  During our observations, V2247~Oph was in a low-accretion state, with emission lines showing only weak levels of circular polarisation;  we nevertheless find that excess emission apparently concentrates in a mid-latitude region of strong radial field, suggesting that it is the footpoint of an accretion funnel.  The weaker and more complex field that we report on V2247~Oph may share similarities with those of very-low-mass late-M dwarfs and potentially explain why low-mass cTTSs rotate on average faster than intermediate mass ones.  These surprising results need confirmation from new independent data sets on V2247~Oph and other similar low-mass cTTSs. ",
        "introduction": "\\label{sec:int}  Whereas our understanding of most phases of stellar evolution made considerable progress throughout the twentieth century, stellar formation remained rather enigmatic and poorly constrained by observations until about two decades ago with the advent of the most powerful ground-based and space telescopes.  One of the major discoveries obtained then is that protostellar accretion discs are often associated with extremely powerful and highly collimated jets escaping the systems along their rotation axis \\citep[e.g.,][]{Snell80}.  This finding has revolutionized the field of stellar formation;  in particular, it provided us with strong evidence that magnetic fields are playing an important role throughout stellar formation.  Magnetic fields are expected to modify significantly the contraction of molecular clouds into protostellar cores and discs \\citep[e.g.,][for a review]{Andre08}. They are also expected to influence the next formation stage, at least for stars with masses lower than about 2.5~\\msun;  these protostars (called classical T~Tauri stars or cTTSs) apparently host magnetic fields strong enough to disrupt the central regions of their accretion discs, connect the protostars to the inner disc rim through discrete accretion funnels and possibly even slow down the rotation of protostars through the resulting star/disc magnetic torque \\citep[e.g.,][for a review]{Bouvier07}.  While magnetic fields have been repeatedly reported over the last decade at the surface of a number of prototypical TTSs \\citep[e.g.,][]{Johns99b, Johns07}, it is only recently that their large-scale topologies can be investigated into some details through phase-resolved spectropolarimetric observations \\citep{Donati07, Donati08b, Hussain09}.  This new option offers a direct opportunity to pin down the role of magnetic fields in this crucial phase of the formation process.  In particular, it straightforwardly allows studying how magnetic topologies of cTTSs depend on mass and rotation rate, giving hints on the potential origin of such fields;  moreover, it provides material for quantitatively studying how protostars magnetically connect and interact with their accretion discs.  Up to now, magnetic topologies of 4 cTTSs have been investigated with this method.  Among these 4, only BP~Tau ($\\mstar\\simeq0.7$~\\msun) is fully convective and hosts a rather simple, dominantly dipolar, large-scale magnetic field, the 3 others being partly convective (as a result of their relatively higher mass, in excess of 1.3~\\msun) and hosting a more complex magnetic topology.  In particular, this transition from simple to complex fields for stars on either sides of the full convection limit is reminiscent of the magnetic properties of low-mass main-sequence dwarfs \\citep{Morin08b, Donati08d};  in addition to suggesting that magnetic fields of cTTSs are likely of dynamo origin, it also gives potential hints on why magnetospheric accretion and rotation properties of cTTSs are discrepant on both sides of the full convection limit (e.g., with higher mass cTTSs rotating more quickly in average).  Extending this magnetic survey to a larger sample of cTTSs is needed to go further;  this is the exact purpose of the MaPP (Magnetic Protostars and Planets) program, in the framework of the international MagIcS research initiative (aimed at qualifying the magnetic properties of stars throughout the HR diagram). This survey has been allocated 690~hr of observing time on the 3.6~m Canada-France-Hawaii Telescope (CFHT) over 9 successive semesters (from 2008b to 2012b). The present paper concentrates on the low-mass\\footnote{Throughout this paper, cTTSs with masses lower than 0.5~\\msun, ranging from 0.5 to 1~\\msun\\ and larger than 1~\\msun\\ are respectively called low-mass, intermediate-mass and high-mass cTTSs. } cTTS V2247~Oph (spectral type M1) in the $\\rho$~Oph star forming region (Lynds~1688 dark cloud); with a mass of only about 0.35~\\msun\\ (see Sec.~\\ref{sec:v22}), it is less massive than all cTTSs yet magnetically imaged and thus appears as an ideal candidate to expand our sample towards cooler cTTSs. ",
        "conclusions": "\\label{sec:dis}  We presented in this paper the first magnetic map for the low-mass cTTS V2247~Oph. This is the lowest mass cTTS to be magnetically imaged yet;  despite it being in a low-state of accretion (with $\\log \\Mdot \\simeq -9.8$ with \\Mdot\\ expressed in \\mspy), this study brings a number of new and unexpected results.  The brightness distribution we derive at the surface of V2247~Oph (from the rotational modulation of Stokes $I$ photospheric LSD profiles) is relatively simple, with dark spots covering about 3.5\\% of the total surface;  in particular, it does not feature a highly-contrasted dark polar spot as seen, e.g., on V2129~Oph \\citep{Donati07}, BP~Tau \\citep{Donati08b} or other cTTSs \\citep{Hussain09}. It is however roughly similar to what is found on low-mass fully-convective dwarfs where large-scale spots are usually low-contrast and cover only a small fraction of the stellar surface \\citep[e.g.,][]{Morin08a}.  We also detect clear Zeeman signatures from V2247~Oph, tracing a rather complex and moderately strong multipolar large-scale magnetic topology with an average strength of 300~G and a mixed amount of poloidal and toroidal field;  the poloidal field is mostly non-axisymmetric and features a weak (tilted) dipole component of about 80--100~G.  Again, this is fairly different from what is seen on intermediate-mass cTTSs like BP~Tau, where the field is far simpler, predominantly poloidal, axisymmetric with a dipole component exceeding 1~kG \\citep{Donati08b}. This is also in contrast to low-mass fully-convective mid-M dwarfs \\citep[e.g.,][]{Morin08b} that almost always exhibit very intense, mainly potential large-scale fields with a simple topology roughly aligned with the rotation axis. New results however indicate that a significant fraction of very-low-mass dwarfs (late-M dwarfs with $\\mstar<0.2$~\\msun, i.e., below about half the mass threshold for full convection on the main sequence) are found to host moderate and complex non-axisymmetric large-scale magnetic fields (Morin et al 2009, in preparation), very different from the strong and simple fields of mid-M dwarfs.  Although V2247~Oph is significantly more massive (at 0.35~\\msun) than late-M dwarfs, the difference in evolutionary stage may easily compensate for the discrepancy;  the mass of V2247~Oph is indeed also lower than half the mass threshold for full convection at an age of about 1~Myr. Further observations are of course needed to confirm whether this analogy truly holds;  if so, it would bring additional evidence that magnetic fields of intermediate- and low-mass cTTSs are similar in nature to those of mid-M and late-M dwarfs and are thus produced through dynamo processes (rather than being fossil leftovers from an earlier formation stage).  In addition, we find that the spot distribution and the magnetic field of V2247~Oph evolve on a very short timescale, of order one week;  in both cases, this evolution is compatible with surface differential rotation shearing the photosphere 5 to 6 times faster than on the Sun, with the equator lapping the pole by one complete rotation every $19\\pm2$~d.  This is again fairly unexpected since all fully-convective stars have been reported to show little to no differential rotation up to now \\citep[e.g.,][]{Donati06a, Morin08a, Morin08b}.  Only late F stars with very shallow convective zones are yet known to exhibit a similar degree of photospheric shear \\citep[e.g.,][]{Donati08c}.  Further confirmation from new data sets on V2247~Oph (and other similar low-mass cTTSs) are thus needed to validate this surprising result.  The amount of differential rotation we estimate is compatible with the range of photometric periods measured on V2247~Oph over the last 2 decades \\citep[from 3.4 to 3.6~d, implying $\\dom>0.1$~\\rpd,][]{Grankin08}.  Given the differential rotation law we derive, these periods suggest that spots preferentially cluster at low to intermediate latitudes (15--40\\degr) on V2247~Oph; this is at least qualitatively compatible with the fact that no high-contrast polar spot is present at the surface. This is the first time that differential rotation is detected on a cTTS;  the recent claim \\citep{Herbst06} that the high-mass cTTS HBC~338 (spectral type G9) is differentially rotating (given the observed variations of the photometric period) needs confirmation, cTTSs with massive accretion discs being prone to photometric perturbations (including sudden changes of the light-curve period) likely caused by accretion \\citep[e.g.,][]{Simon90, Donati08b} rather than resulting from differential rotation.  Finally, we find that V2247~Oph features a region of excess IRT (and \\hal\\ and \\hbe) emission at phase 0.45 and intermediate latitude. There is no obvious correlation between the location of this emission region and the cool spots mapped from the distorted LSD profiles of photospheric lines;  however, we find that it roughly coincides with one of the strongest field region detected on V2247~Oph, the positive radial field spot whose magnetic flux reaches about 0.5~kG.  We speculate that this spot may be the footpoint of an accretion funnel linking the star to its accretion disc at the time of our observations.  If confirmed, we expect such accretion funnels to be rather short lived, with differential rotation continuously distorting the magnetic connections between the star and the disc as well as the stellar field itself.  In particular, this effect may be partly responsible for the sporadic accretion episodes observed on V2247~Oph, with more intense accretion episodes occurring when the magnetic topology linking the star to the disc is favouring accretion.  In addition to the differences in their magnetic topologies,  V2247~Oph and BP~Tau show quite distinct rotational properties;  as emphasised in Sec.~\\ref{sec:v22}, each of them belongs to a distinct population of young stars \\citep{Lamm05}, V2247~Oph to the low-mass cTTSs with short rotation period (of about half a week in average) while BP~Tau to the intermediate-mass cTTSs with long rotation periods (of about one week in average). These two issues may relate together, i.e., the magnetic topology may have an impact on the rotation period through a coupling mechanism such as, e.g., the well-known disc-locking picture \\citep{Camenzind90, Konigl91}. In this model, the rotation period at the surface of the star settles close to the Keplerian period of the point in the disc (called the Alfven radius and noted \\rmag) where the ram pressure of the average accretion flow roughly equals the magnetic pressure of the largest-scale (i.e., the dipole component) poloidal field \\citep[e.g.,][]{Bessolaz08}.  Using Eq.~1 of \\citet{Bessolaz08} (with $B_*$ set to the average dipole field strength at the rotation equator, i.e., about 40~G for V2247~Oph), we estimate $\\rmag\\simeq3.9$~\\rstar; we can also work out that the radius at which the Keplerian period matches the average rotation period (called the corotation radius and noted \\rcor) is equal to $\\rcor\\simeq3.4$~\\rstar\\ for V2247~Oph.  Both estimates roughly agree with each other, suggesting a reasonable agreement with the disc-locking model predictions.  Comparing directly with BP~Tau however \\citep[and assuming a logarithmic mass accretion rate of --7.5 taken from the literature as in][]{Donati08b}, the model predicts that both stars should have similar \\rmag\\ and thus behave similarly with respect to the disc-locking mechanism, the 15-fold weaker dipole field of V2247~Oph being almost compensated by the 200-fold weaker mass accretion rate;  re-estimating $\\log \\Mdot$ for BP~Tau using the same procedure as for V2247~Oph, we actually find a much smaller value, equal to $-8.6\\pm0.4$. While it is not clear whether this new estimate is more accurate than the one used previously, it is at least more consistent with the one we derived for V2247~Oph when it comes to comparing both stars; in this case, we obtain that $\\rmag\\simeq7.1$~\\rstar\\ for BP~Tau (taking $B_*\\simeq600$~G) in reasonable agreement with $\\rcor\\simeq7.4$~\\rstar.  Obviously, more stars of different masses and rotation periods, and with consistent estimates of $B_*$ and \\Mdot, are needed to assess the validity of the disc-locking picture; if confirmed, it may ultimately explain why low-mass cTTSs rotate in average faster than their intermediate mass equivalents. The MaPP program should bring new results along this line in the near future."
    },
    "0911/0911.0423_arXiv.txt": {
        "abstract": "We present a sample of \\noptall\\ optically selected \\bl\\ candidates from the SDSS DR7 spectroscopic database encompassing  8250~deg$^2$ of sky; our sample constitutes one of the largest uniform BL Lac samples yet derived. Each \\bl\\ candidate has a high-quality SDSS spectrum from which we determine spectroscopic redshifts for $\\sim$60\\% of the objects. Redshift lower limits are estimated for the remaining objects utilizing the lack of host galaxy flux contamination in their optical spectra; we find that objects lacking spectroscopic redshifts are likely at systematically higher redshifts.  Approximately 80\\% of our \\bl\\ candidates match to a radio source in FIRST/NVSS, and $\\sim$40\\% match to a ROSAT X-ray source.  The homogeneous multiwavelength coverage allows subdivision of the sample into \\noptrl\\ radio-loud \\bl\\ candidates and \\noptrq\\ weak-featured radio-quiet objects. The radio-loud objects broadly support the standard paradigm unifying \\bl\\ objects with beamed radio galaxies. We propose that the majority of the radio-quiet objects may be lower-redshift ($z<2.2$) analogs to high-redshift weak line quasars (i.e., AGN with unusually anemic broad emission line regions).  These would constitute the largest sample of such objects, being of similar size and complementary in redshift to the samples of high-redshift weak line quasars previously discovered by the SDSS.  However, some fraction of the weak-featured radio-quiet objects may instead populate a rare and extreme radio-weak tail of the much larger radio-loud \\bl\\ population.  Serendipitous discoveries of unusual white dwarfs, high-redshift weak line quasars, and broad absorption line quasars with extreme continuum dropoffs blueward of rest-frame 2800~\\AA\\ are also briefly described. ",
        "introduction": "BL~Lacertae objects compose an especially rare subclass of Active Galactic Nuclei (AGNs) traditionally interpreted as  low-luminosity radio galaxies with a relativistic jet pointed toward the observer \\citep[e.g., see][]{blandford78, urry95}.  Emission from the relativistically boosted jet dominates the observed flux, yielding featureless (or nearly featureless) optical spectra, strong radio and X-ray emission, strong polarization, flat radio spectra, and rapid variability across the entire electromagnetic spectrum \\citep[e.g., see][]{kollgaard94,perlman01}.  Given their odd spectral characteristics, and the need to obtain optical spectra to confirm a candidate as a \\bl\\ object, the most efficient \\bl\\ selection algorithms tend to search for optical counterparts to strong radio and/or X-ray sources \\citep[e.g., see][]{stickel91,stocke91,perlman98,laurent99,landt01,anderson07,padovani07,piranomonte07,turriziani07,plotkin08}.  Optically selected \\bl\\ samples have historically been extraordinarily difficult to assemble, but they are highly desired.  Systematic searches based only on optical characteristics would allow the exploration of \\bl\\ properties with minimal concern for  biases introduced by radio and X-ray surveys (which generally tend to have shallower flux sensitivities than optical surveys.)  Also, optically selected samples might reveal new populations of \\bl\\ objects, and they might even lead to serendipitous discoveries of other, perhaps new, classes of featureless objects. Unfortunately, the odd spectral characteristics of \\bl\\ objects conspire to make selection based solely on optical colors ineffective.  For example, only 4 \\bl\\ objects were identified in the Palomar-Green Survey for ultraviolet-excess objects \\citep[PG,][]{green86}; another 3 PG sources initially misclassified as DC (i.e., featureless) white dwarfs were identified as \\bl\\ objects only after the {\\it ROSAT} All Sky Survey \\citep[RASS,][]{voges99,voges00} came online \\citep{fleming93}.  Attempts have been made to exploit other distinctive \\bl\\ properties, like polarization and flux variability, to derive optically selected samples.   Such techniques, however, require observing large areas of the sky at multiple epochs just to assemble manageable lists of candidates suitable for spectroscopic follow-up.  For example, it would seem that the strong polarization of \\bl\\ objects ($P\\gtrsim2\\%$) should allow for efficient identification of \\bl\\ candidates.  However, most attempts along these lines have so far been unsuccessful \\citep{impey82_optsel,borra84,jannuzi93}, probably because the polarized flux is variable with a low duty cycle: for example, X-ray selected \\bl\\ objects from the {\\it Einstein} Observatory Extended Medium-Sensitivity Survey \\citep[EMSS,][]{stocke91}  spend over half their time with $P<4\\%$ \\citep[][]{jannuzi94}.  Not until we entered the age of large-scale spectroscopic digital sky surveys has optical selection of \\bl\\ objects been a realistic endeavor.  Traditionally, the primary limitation of efficient \\bl\\ recovery algorithms is the large number of false positives returned, even when radio and/or X-ray information is consulted.  Optical spectroscopy is then required to make firm \\bl\\ identifications.  However, with a digital survey such as the Sloan Digital Sky Survey \\citep[SDSS,][]{york00}, this task is less daunting because the requisite spectra already exist.  Furthermore, the most relevant spectra can be searched within databases, allowing for automated removal of many of the contaminants that previously had to be culled manually (after a significant investment in telescope time.)  The first successful attempt at assembling an optically selected sample is the 2QZ \\bl\\ survey \\citep{londish02,londish07}, derived from  the 2-degree field (2dF) and the 6-degree field (6dF) quasi-stellar object Redshift Surveys \\citep{croom04}.  They searched over $\\sim$10$^3$~deg$^2$ and recovered 7 confident \\bl\\ objects.  All 7 objects were additionally identified with radio and X-ray sources (post-selection), and they also showed optical flux variations in a photometric monitoring campaign over 2002--2004 \\citep{nesci05}.   An additional object ($z=0.494$) is identified as a radio-quiet/weak \\bl\\ candidate \\citep{londish04}.  This source lacks both a radio and an X-ray detection (with a radio to optical flux ratio limit placing it firmly in the radio-quiet regime\\footnote{Radio-quiet quasars are commonly defined as having radio to optical flux ratios $R<10$ \\citep[e.g., see][]{kellermann89, stocke92}.   Throughout this paper, we adopt the approximately similar $\\alpha_{ro}<0.2$, where $\\alpha_{ro}$ is the radio-to-optical broad-band spectral index: $\\alpha_{ro}=-\\log(L_o/L_r)/5.08$, where $L_o$ and $L_r$ are optical and radio specific luminosities at rest-frames 5000~\\AA\\ and 5~GHz, respectively \\citep{stocke85}.}),  % but its optical properties are similar to (radio-loud) \\bl\\ objects.  The recovery of only a single extragalactic object with a featureless spectrum lacking radio emission by the  2QZ survey confirms the \\citet{stocke90} result that radio-quiet \\bl\\ objects must be extremely rare if they even exist at all \\citep[also see][]{jannuzi93}.  Searching 2860~deg$^2$, \\citet{collinge05}, hereafter \\citetalias{collinge05}, assembled a larger optically selected sample from the SDSS.  Their sample contains 386 nearly featureless spectra, including 240 {\\it probable} \\bl\\ candidates and 146 {\\it possible} \\bl\\ candidates (the latter are most likely DC white dwarfs, proper motions are consulted for their classification.)  The vast majority of the {\\it probable} candidates either have firm extragalactic redshifts, or they match to a radio source in the Faint Images of the Radio Sky at Twenty-cm survey \\citep[FIRST,][]{becker95} and/or the NRAO VLA Sky Survey \\citep[NVSS,][]{condon98},  or to a RASS X-ray source (correlations to multiwavelength catalogs are performed post-selection.)  \\citetalias{collinge05} also present a list of 27 intriguing radio-quiet/weak \\bl\\ candidates.   These objects have nearly featureless optical spectra, and their radio fluxes (or limits in the cases of radio non-detections) place their radio to optical flux ratios within or tantalizingly close to the radio-quiet regime.   All but one of these also lack X-ray detections in RASS (but their extant X-ray limits are not sensitive enough to declare them X-ray weak.)  Some of these objects are at high-redshift ($z>2.2$), and those high-redshift objects may be alternately described as members of a population of high-redshift weak line quasars (WLQs) discovered by the SDSS \\citep[e.g., see][]{fan99,anderson01,shemmer06,shemmer09,diamond09}.  Three of these high-redshift objects were detected in follow-on {\\it Chandra} X-ray observations by \\citet{shemmer09}, and they are X-ray weaker than typical radio-loud \\bl\\ objects.   Other radio-quiet \\bl\\ candidates in \\citetalias{collinge05} have lower redshifts, and it is unclear if they are more closely related to \\bl\\ objects or to WLQs.  Regardless of their true nature, \\citetalias{collinge05} demonstrate that optical selection of \\bl\\ objects from the SDSS is efficient,  producing a large enough sample to reveal especially rare sub-populations of objects with featureless spectra in numbers comparable to entire venerable radio and X-ray selected \\bl\\ samples.  Here, we expand on \\citetalias{collinge05} and continue the search for \\bl\\ objects in the SDSS (which now covers almost three times the sky area as the \\citetalias{collinge05} sample).  We recover \\noptall\\ objects.  Our automated selection algorithm is described in \\S \\ref{sec:ch4_sampsel}.  We visually inspect $\\sim$23,000 spectra, and we discuss the removal of various classes of contaminants (including stars, post-starburst galaxies, quasars, etc.)\\ in \\S \\ref{sec:ch4_visinspec}; we also present lists of serendipitously discovered unusual white dwarfs, WLQs, and unusual broad absorption line quasars.  In \\S \\ref{sec:ch4_postvisinspec} we describe the removal of contaminants that survive visual inspection (including featureless white dwarfs and large red galaxies.)  The final sample is presented in \\S \\ref{sec:ch4_finalsamp}, and we describe a spectral decomposition of the \\bl\\ host galaxy from the AGN.  We use this decomposition to measure optical spectral indices for the AGN, to estimate lower redshift limits by assuming \\bl\\ host galaxies are standard candles, and to calculate decomposed AGN component optical fluxes and luminosities.  We also correlate the \\noptall\\ \\bl\\ candidates to the FIRST and NVSS radio surveys, and to the RASS X-ray survey.   The sample is discussed in \\S \\ref{sec:ch4_discussion}, which focuses on a particularly interesting subset of \\noptrq\\ radio-quiet objects with weak spectral features.   Finally, our main conclusions are summarized in \\S \\ref{sec:ch4_summary}.  Throughout, we use {\\it quasar} to refer to any AGN (regardless of luminosity or radio loudness and including \\bl\\ objects), and we adopt $H_0= 71$~km s$^{-1}$~Mpc$^{-1}$, $\\Omega_m=0.27$, and $\\Omega_{\\Lambda}=0.73$. ",
        "conclusions": "  1.\\ Among the tens of thousands of contaminants removed during visual inspection, we note potentially interesting serendipitous discoveries, including 17 unusual white dwarfs,  9 unusual BALQSOs with extreme continuum dropoffs blueward of \\ion{Mg}{2}, and \\tabnwlq\\ high-redshift WLQs (with reliable $z>2.2$ and Ly$\\alpha$+\\ion{N}{5} $REW<10$~\\AA).  Six of these WLQs are additionally classified as \\bl\\ candidates because their emission features have $REW<5$~\\AA.  2.\\ Approximately 60\\% of our \\bl\\ candidates have spectroscopic redshifts (or limits from \\ion{Mg}{2} absorption).  For objects lacking spectroscopic redshifts, we assume \\bl\\ host galaxies are standard candles, and we place redshift limits to those objects utilizing the fact that we do not detect host galaxy contamination to their SDSS spectra.  We estimate their ``host galaxy'' redshift limits to be accurate to $\\sigma_z = \\pm0.064$.  We thus have redshift information (either from weak spectral features or limits placed from the lack of a host galaxy detection) for {\\it every} single object in our sample.  We infer the objects lacking spectroscopically-derived redshifts are at systematically higher redshift.  We expect this result, since the objects with the weakest spectral features are arguably among the most highly beamed, which requires a low probability geometry.  So, we are more likely to find the most highly beamed objects at higher redshift, where more volume is probed.  However, there may also be some examples of highly beamed nearby objects in this sample as well.  3.\\ Each object emerges with homogeneous multiwavelength data coverage, and we include radio/X-ray fluxes (or limits) for all objects.   Around 80\\% of the \\bl\\ candidates presented here match to a radio source in FIRST or NVSS, and we place upper limits on radio flux densities for the rest (except for \\noptallNoRadioInfo\\ objects outside the FIRST footprint and not detected by NVSS.)   Approximately 40\\% match to an X-ray source in RASS, and we place upper X-ray flux limits for the rest.  4.\\ We use each object's high-quality SDSS spectrum to perform spectral decompositions of each object's host galaxy and nucleus, and we report AGN component optical fluxes and luminosities.  5.\\  Based on radio fluxes and limits from the FIRST/NVSS radio surveys, we subdivide our sample into \\noptrl\\ radio-loud \\bl\\ candidates and \\noptrq\\ radio-quiet weak-featured spectra.   There is significant overlap between the optical properties of the radio-loud and radio-quiet objects.  However, the radio-quiet objects tend to be at higher redshift, and they also have a smaller X-ray detection rate.  We argue that  the $\\alpha_{ro}<0.2$ objects are a mix of different populations: some of these objects (especially ones lacking spectroscopic redshifts or with small redshifts) probably populate the radio-weak tail of the much larger radio-loud \\bl\\ population.  Other sources (especially the ones at higher redshifts) are likely WLQs.  The very high-redshift nature of published SDSS WLQs is a selection effect (because they are identified by the presence of the Ly$\\alpha$ forest), and  this study provides a systematic search of the SDSS database capable of revealing lower-redshift analogs in large numbers.  Even just the $z<2.2$ radio-quiet objects constitute, to our knowledge, the largest sample of such weak-featured radio-quiet objects.  The sample size is larger than many entire venerable \\bl\\ samples, and this subset is of similar size and complementary to the number of known high-redshift SDSS WLQs.  6.\\ The radio-loud objects in general support the standard unification paradigm, which maintains that radio and X-ray selected \\bl\\ objects populate extreme ends of a single, larger, and continuous \\bl\\ population.  We recover a mix of LBLs, IBLs, and XBLs, although the SDSS appears biased toward finding HBLs.  However, we note that the large observed scatter in \\bl\\ properties (e.g., the distribution of their AGN optical spectral indices) is probably due to innate differences as well.  Even just the \\noptrl\\ radio-loud objects constitute one of the largest samples of \\bl\\ objects derived from a single selection technique.  7.\\   It will be interesting to assemble light curves for these \\bl\\ candidates once large-scale time-domain surveys, like the  Large Synoptic Survey Telescope \\citep[LSST,][]{ivezic08_ph} and  the Panoramic Survey Telescope \\& Rapid Response System (Pan-STARRS\\footnote{\\url{http://pan-starrs.ifa.hawaii.edu}}), come online.   It may even be possible to use flux variability from those surveys as a \\bl\\ selection criterion.  For example, \\citet{bauer09} recently found $\\sim$3000 objects in the Palomar-QUEST survey identified by their optical flux variability; many of these objects may be blazars.    Time-domain \\bl\\ searches over huge areas of the sky  requires telescope access that is unrealistic with standard non-survey telescopes, but well within the capabilities of LSST and Pan-STARRS."
    },
    "0911/0911.1205_arXiv.txt": {
        "abstract": "{Is the Doppler interpretation of galaxy redshifts in a Friedmann-Lema\\^{\\i}tre-Robertson-Walker (FLRW) model valid in the context of the approach to comoving spatial sections pioneered by de~Sitter, Friedmann, Lema\\^{\\i}tre and Robertson, i.e. according to which the 3-manifold of comoving space is characterised by both its curvature and topology?} {Holonomy transformations for flat, spherical and hyperbolic FLRW spatial sections are \\postrefereechanges{proposed}.} {\\postrefereechanges{By quotienting a simply-connected FLRW spatial section by an appropriate group of holonomy transformations, the} Doppler interpretation in a non-expanding Minkowski space-time, obtained via four-velocity parallel transport along a photon path, is found to imply that an inertial observer is receding from herself at a speed greater than zero, implying contradictory world-lines.  The contradiction \\postrefereechanges{in the multiply-connected case} occurs for arbitrary redshifts in the flat and spherical cases, and for certain large redshifts in the hyperbolic case.} {The link between the Doppler interpretation of redshifts and cosmic topology can be understood physically as the link between parallel transport along a photon path and the fact that the comoving spatial geodesic corresponding to a photon's path \\postrefereechanges{can be} a closed loop in an FLRW model of any curvature. Closed comoving spatial loops are fundamental to cosmic topology.} ",
        "introduction": "Much debate has recently taken place regarding the interpretation of the redshifts of comoving galaxies in cosmological models as a special-relativistic Doppler effect in the absence of the concept of expanding space \\citep{Chod07a,Barnes06,Chod07b,Abram07,Francis07,Lewis08,BunnHogg09,Peacock08,Abram09,Chod08,Faraoni09}. Here, it is shown that in the flat and spherical Friedmann-Lema\\^{\\i}tre-Robertson-Walker (FLRW) models, where galaxies are comoving massless test objects in the standard comoving-coordinate system, the Doppler interpretation in non-expanding Minkowski space-time leads to a contradiction for galaxies at distances from the observer that are arbitrarily small, provided that these galaxies are considered to be comoving.  The corresponding (weaker) argument in hyperbolic FLRW models is also presented.  In \\SSS\\ref{s-cos-top}, some mathematical properties of the comoving spatial sections in FLRW models are recalled. The methods of using these properties in flat, positively curved, and negatively curved FLRW models are presented in \\SSS\\ref{s-meth-zero}, \\SSS\\ref{s-meth-pos}, and \\SSS\\ref{s-meth-neg}, respectively. The consequences are described in \\SSS\\ref{s-results}. The way that density perturbations modify these consequences is discussed in \\SSS\\ref{s-almostflrw}. The failure of the Doppler interpretation to arbitrarily small separations for non-negatively curved exact-FLRW models may seem to be inconsistent with the Minkowski nature of FLRW models in the limit towards a point in space-time. This paradox is explained in \\SSS\\ref{s-Mink-limit}. A definition of ``expansion of space'' and related statements from the literature are briefly discussed in \\SSS\\ref{s-expanding}, and the notion of an expanding Minkowski space-time is mentioned in \\SSS\\ref{s-expanding-Mink}. Conclusions are presented in \\SSS\\ref{s-conclu}.  For clarity, the terms ``hyperbolic'' and ``spherical'' are used for negatively and positively curved FLRW models, respectively, rather than the ambiguous terms ``open'' and ``closed''.  The term ``galaxy'' is used for an external massless galaxy located at some non-zero distance from the observer in the covering space. The observer should be considered to be located in our (massless) Galaxy. The ``galaxy'' and the ``Galaxy'' are massless in order not to violate homogeneity. \\postrefereechanges{A ``comoving spatial geodesic'' is distinct from a space-time geodesic. The former refers to a geodesic (a curve that minimises the metric distance between any close pair of points on that curve) of a comoving spatial section (where the metric only has a spatial component). It can also be thought of as a space-time geodesic (e.g. the path of a photon) projected to a comoving spatial section by ignoring the cosmological time coordinate. For example, in a positively curved FLRW model, a comoving spatial geodesic is an arc that is part of a great circle of the hypersphere $S^3$, since great circles are straight lines.}  \\postrefereechanges{For a short introduction to the topology of FLRW models and observational approaches to measuring it, see \\cite{Rouk00BASI}.  For reviews, see \\cite{LaLu95,Lum98,Stark98,LR99,BR99,RG04}. Analyses of WMAP data presently include analyses suggesting that sub-matter-horizon cosmic topology has been detected \\citep[][and references therein]{Gundermann2005,Caillerie07,RBG08,Aurich09a}.  and analyses disfavouring it \\citep[][and references therein]{KeyCSS06,NJ07}. Key terminology regarding the comoving spatial section, which is a 3-manifold, includes the following, using the 3-torus as an example.  The manifold $M = T^3$ can be thought of as a cube, the ``fundamental domain'', with identified faces. It can also be thought of as a tiling of Euclidean 3-space, $\\mathbb{R}^3$, called the ``covering space'', $\\widetilde{M}$, which can also be informally called the ``apparent space'', since it contains many copies of any individual physical object. A mapping $f$ from one copy of the cube in $\\widetilde{M}$ to another copy is a (non-arbitrary) isometry of $\\widetilde{M}$, i.e. it preserves metric distances, and is called a ``holonomy transformation''. The set of holonomy transformations for a given manifold $M$ is the group $\\Gamma$.  This group has the same structure as the group of all possible closed comoving spatial paths in $M$ that cannot be continuously transformed into one another. The latter group is called the first homotopy group, $\\pi_1(M)$. Switching from $\\widetilde{M}$ to $M = \\widetilde{M}/\\Gamma$ can be described as ``quotienting'' $\\widetilde{M}$ by $\\Gamma$.} ",
        "conclusions": "\\label{s-conclu}  Cosmic topology is not just a luxury referred to by de~Sitter, Friedmann, Lema\\^{\\i}tre and Robertson \\citep[and ][in the pre-relativistic era]{Schw00}\\footnote{English translation: \\protect\\citet{Schw98}.}. Apart from providing a competitor to the infinite flat model for understanding WMAP observational data, cosmic topology can also help to improve physical insight into fundamentals of FLRW models that initially appear to be unrelated to global topology. \\postrefereechanges{ If the parallel-transported four-velocity Doppler interpretation of galaxy redshifts depends only on the metric (including $a(t)$), then it} leads to the physical contradiction of an inertial observer in a non-expanding Minkowski space-time receding from herself. This contradiction can be resolved, but only at the cost of introducing the notion of expanding space, in which case the motivation for a Doppler interpretation is weakened, or by explicitly stating that the comoving distance to the galaxy must be less than twice the injectivity radius of comovinig space.  The link between the Doppler interpretation of redshifts and cosmic topology can be summarised physically as follows. \\begin{list}{(\\roman{enumi})}{\\usecounter{enumi}} \\item The relativistic Doppler interpretation of a cosmological redshift is obtained by parallel transporting the emitting galaxy's four-velocity along a photon path to the observer. \\item \\postrefereechanges{For an FLRW model of any of the three curvatures (with the restrictions as detailed above), the comoving spatial geodesic of the photon's path can be considered to be a closed path by changing from a simply-connected FLRW model to a multiply-connected FLRW model, without any modification of the metric. This should not affect the Doppler interpretation if that interpretation depends only on the metric.} \\item Closed comoving spatial paths constitute the mathematical foundation of cosmic \\postrefereechanges{topology \\citep[e.g., \\SSS{}3.4,][]{LaLu95}}. \\end{list} Hence, it is unsurprising that insistence on the possibility of interpreting galaxy redshifts as a relativistic Doppler effect leads to the properties of cosmic topology, and in turn requires the physical concept of expanding space, i.e. epochs \\postrefereechanges{when $\\dot{a}>0$} (Defn~\\ref{d-eos})."
    },
    "0911/0911.3346_arXiv.txt": {
        "abstract": "We construct a new Hartree-Fock-Bogoliubov (HFB) mass model, labeled HFB-18, with a generalized Skyrme force. The additional terms that we have introduced into the force are density-dependent generalizations of the usual $t_1$ and $t_2$ terms, and are chosen in such a way as to avoid the high-density ferromagnetic instability of neutron stars that is a general feature of conventional Skyrme forces, and in particular of the Skyrme forces underlying all the HFB mass models that we have developed in the past. The remaining parameters of the model are then fitted to essentially all the available mass data, an rms deviation $\\sigma$ of 0.585 MeV being obtained. The new model thus gives almost as good a mass fit as our best-fit model HFB-17 ($\\sigma$ = 0.581 MeV), and has the advantage of avoiding the ferromagnetic collapse of neutron stars. ",
        "introduction": "Astrophysical considerations require that one have available nuclear-mass models as rigorously based as possible. In this way one might hope to be able to extrapolate from the mass data, which cluster fairly closely to the stability line, out towards the neutron drip line, and make reliable estimates of the masses of nuclei so neutron rich that there is no hope of measuring  them in the foreseeable future; such nuclei are found in the outer crusts of neutron stars, and also play a vital role in the r-process of nucleosynthesis. To this end we have developed a series of nuclear-mass models based on the Hartree-Fock-Bogoliubov (HFB) method with Skyrme and contact-pairing forces, together with phenomenological Wigner terms and correction terms for the spurious collective energy; all the model parameters are fitted to essentially all the experimental mass data (see Ref.~\\cite{gcp09} and references quoted therein). To make the extrapolations to neutron-rich nuclei as reliable as possible, the underlying Skyrme force in model HFB-9 \\cite{sg05} and all later models was constrained to fit the zero-temperature equation of state (EoS) of neutron matter (NeuM), as calculated by Friedman and Pandharipande~\\cite{fp81} (FP) for realistic two- and three-nucleon forces. (We have so far been unable to obtain mass fits as good as our published ones when constraining to the more complete, and slightly stiffer, realistic neutron-matter EoS labeled A18 + $\\delta\\,v$ + UIX$^*$~\\cite{apr98}. Our findings are consistent with a recent analysis of the $\\pi^-/\\pi^+$ ratio in central heavy-ion collisions indicating that this EoS is too stiff~\\cite{xiao09}.)  Because of the neutron-matter constraint our models can be used to extrapolate beyond the neutron drip line and calculate with some confidence the EoS of the inner crust of neutron stars~\\cite{onsi08} (throughout this paper we assume zero temperature). Since the good agreement of our forces with the FP calculation \\cite{fp81} of neutron matter extends to the highest density encountered in neutron stars it might be thought that our inner-crust EoS is continuous with that of the homogeneous core, the transition between the two regions taking place at around $0.5\\rho_0$, where $\\rho_0$ ($\\simeq$ 0.16 nucleons.fm$^{-3}$) is the equilibrium density of symmetric nuclear matter (SNM). However, in fitting our forces to the neutron matter of FP \\cite{fp81}, we {\\it assume} that our ground state is spin unpolarized, but in fact the underlying Skyrme forces of all our models, like all conventional Skyrme forces of the form (\\ref{1}), predict that beyond a certain supernuclear density the ground state of NeuM becomes ferromagnetic, i.e., at least partially polarized~\\cite{vida84,kw94,mar02}. In the case of the Skyrme force BSk17, the force underlying our best-fit model, HFB-17~\\cite{gcp09}, complete polarization sets in at a density of $\\rho_{frmg} = 1.24\\,\\rho_0$. On the other hand, microscopic calculations using different realistic forces and different methods~\\cite{fan01, bomb06, sam07, bord08,vid02} all predict no such polarization, at least up to about $5\\rho_0$. Moreover, in the case of all our own previously published HFB forces the predicted ferromagnetic state is unstable against collapse, the energy becoming more and more negative as the density increases (see, for example, the lowest curve in Fig.~\\ref{fig1}, constructed for force BSk17). This predicted ferromagnetic collapse sets in at densities that are certainly encountered in the cores of all neutron stars, and is contradicted by the very existence of neutron stars.  Actually, the core of neutron stars does not consist of pure NeuM but rather of so-called neutron-star matter (N*M), which just below the crust is composed of neutrons with a weak admixture of proton-electron pairs in beta equilibrium (for simplicity we will assume that these are the only particles present also at higher densities). But if NeuM were indeed unstable against collapse then N*M would be likewise, since, at the very least, it would always be energetically advantageous for N*M to spontaneously transform into NeuM through electron capture beyond some critical density. Thus the stability of NeuM against collapse is a necessary condition for the existence of neutron stars. But even if no such instability is implied, there is a general tendency for Skyrme forces of the conventional form (\\ref{1}) to lead to the ground state of N*M being polarized. However, calculations based on realistic forces show the ground state of N*M, like that of NeuM, to be unpolarized at all densities~\\cite{bord08b}.  Our purpose here is to show that by adding suitable terms to the Skyrme force it is possible to eliminate the anomalous prediction of a ferromagnetic transition in neutron stars, with essentially no impact on the high-quality fits to the mass data that we have previously obtained with conventional Skyrme forces. (An alternative approach to this problem has been followed by Margueron {\\it et al.} \\cite{marg09a, marg09}.) In Section II we discuss in more detail the nature of the spurious transition to a ferromagnetic state in NeuM and N*M associated with conventional Skyrme forces of the form (\\ref{1}), considering not only our own BSk17 \\cite{gcp09}, but also the widely used SLy4 \\cite{cha98}, which was specifically constructed for neutron-star calculations. Section III shows how extra terms in the Skyrme force can stop this spurious transition, while in Section IV we describe the new mass fit. Our conclusions are summarized in Section V, and in the Appendix we present the full formalism for the generalized form of Skyrme force used here. ",
        "conclusions": "We have extended our earlier Skyrme-HFB mass models by the inclusion of terms that are density-dependent generalizations of the usual $t_1$ and $t_2$ terms. We have shown that these new terms can be chosen in such a way as to prevent the high-density ferromagnetic collapse of neutron stars that was a general feature of our previous HFB mass models, without compromising the excellent fit to the mass data that we obtained in the past, and without relaxing any of the previously imposed constraints of conformity to reality. The mass predictions made by the new model, HFB-18, are, in fact, very similar to those made by the preceding model, HFB-17~\\cite{gcp09}. These two models not only give better fits to the mass data than does any other published model except that of Duflo and Zucker~\\cite{dz95}, but they are also by far the most microscopically founded models, and in particular their underlying interactions (BSk17 and BSk18, respectively) have been fitted to realistic calculations of both the EoS and the $^1S_0$ gap of neutron matter. They can thus be expected to make more reliable predictions of the highly neutron-rich nuclei that appear in the outer crust of neutron stars and that are involved in the r-process. Moreover, these mass models can be used to extrapolate beyond the drip line to the inner crust of neutron stars, using the respective interactions to calculate the EoS in this region. Our confidence in this extrapolation derives not only from the fit of the interactions to neutron matter but also from the precision fit to masses, which means that the presence of protons and the existence of inhomogeneities in the inner crust are well represented.  Finally, pursuing our aim of a unified treatment of the different regions of a neutron star, we can use these effective interactions to calculate the EoS of the homogeneous core (at least in its outer parts where no complications from the possible appearance of hyperons and other particles arise). Of course, for pure NeuM nothing new can be obtained in this way beyond what has already been given by the realistic calculations to which our effective interactions have been fitted, but these interactions can then be reliably used to calculate N*M, which is not treated in the realistic calculations of FP~\\cite{fp81}. Another important application of these effective interactions, which would hardly be practical with realistic forces, is to make a detailed study of the transition between the inner crust and the fluid core.  We have seen that for N*M, as for masses, the two interactions BSk17 and BSk18 give very similar results, provided we assume that the ground state for the former is unpolarized, which in fact is not the case. Here lies the basic difference between the two interactions: with BSk18 the ground state of NeuM and N*M is indeed unpolarized at all densities prevailing in neutron stars. This elimination of the spurious polarization is the essential contribution of this paper, made possible by the introduction of the terms in $t_4$ and $t_5$. However, we have so far made only a partial search in the new, extended parameter space, and there remains the possibility not only of improving the fit to the mass data still further but also of imposing further realistic constraints on the Skyrme force.  \\appendix"
    },
    "0911/0911.4085_arXiv.txt": {
        "abstract": "Coronal emission line intensities are commonly used to measure electron temperatures using emission measure and/or line ratio methods.  In the presence of systematic errors in atomic excitation calculations and data noise, the information on underlying temperature distributions is fundamentally limited. Increasing the number of emission lines used does not necessarily improve the ability to discriminate between different kinds of temperature distributions. ",
        "introduction": "Spectroscopic measurements of the temperature of coronal plasma have been made for decades \\citep[e.g.][]{Seaton1962c,Noci2003}. \\citet{Feldman+others1999,Feldman+Landi2008} and \\citet{Landi+Feldman2008} presented intriguing  evidence that electron temperatures of solar coronal plasma, measured from emission line spectra obtained high in the coroma, are clumped into several peaks, and are not broadly distributed.   Narrow distributions of plasma temperature have important implications for the energy balance of the corona. In the methods used, frequency integrated line intensities $I$ are assumed to be simple functions $G(T)$ of the logarithm of the electron temperature, $T$, because of the dominance of two body collisional processes which lead to the well-known ``coronal approximation'' \\citep[e.g.][]{Woolley+Allen1948,Seaton1964a}.   Two approaches were used.  In one they solved for the differential emission measure $\\xi(T)$ which is an optimal solution to the inverse problem: \\begin{equation} I_i = \\int G_i(T) \\xi(T) dT,\\ \\ \\ \\ i=1\\ldots n \\end{equation} where there is a set of $n$ different emission lines.  The other method sought the single logarithmic temperature $T_0$ such that $\\xi(T) \\propto \\delta(T-T_0)$, by plotting $I_i/G_i(T)$ and identifying intersection points for the observed lines.  These approaches are, in fact, formally equivalent \\citep{McIntosh+Brown+Judge1997}. In both cases, information must be added to obtain both $\\xi(T)$ and $T_0$: in the first case the problem has to be ``regularized'' as it is ill-posed \\citep[e.g.][]{Craig+Brown1986}, one searches for example for the  ``least structured'' solution $\\xi(T)$ that is compatible with the data. In the second case it is assumed that the plasma is indeed approximately isothermal. In the analysis of \\citet{Feldman+Landi2008}, the individual peaks in the $\\xi(T)$ functions have widths of 0.1-0.2 in $T$, and the curves of $I_i/G_i(T)$ intersect one another within similar margins.  Some general questions arise.  Given typical uncertainties in observable and model parameters, are the observed data compatible with different $\\xi(T)$ distributions?  What, then, is the accuracy of the derived temperatures?   In this paper these questions are addressed by asking, how broad can $\\xi(T)$ functions be to be incompatible with observed data. By how much can one change the temperature of an isothermal plasma before the differences in line intensities become significant? In the problem at hand there are unavoidably large and systematic uncertainties in the $G(T)$ functions. These uncertainties limit the information which can be extracted from emission line spectra. It is found that acceptable widths $\\width$ of the $\\xi(T)$ functions exceed the precision by which the lines can in principle determine that two different plasmas have slightly different temperatures.  Thus, these widths set the lower limit to the ability of emission line techniques to diagnose electron temperatures. ",
        "conclusions": "The largest values of $\\width$ compatible with observations yield the accuracy with which emission lines can measure electron temperatures.  The widths of the distributions compatible with known sources of uncertainties set a natural limit on the sharpness of a detectable peak in the underlying emission measure distributions. Typically the full width, $2\\width$, is 0.2 to 0.3 in the logarithmic electron temperature.   The precision is an order of magnitude better, being estimated using the sensitivity of the spectra to two isothermal plasmas differing in logarithmic temperature by $\\tau$. Although  this implies that two strictly isothermal plasmas {\\em could} be differentiated by using line ratios to a precision of order $\\epsilon_i \\widi/2 \\sim 0.015$ in the logarithmic electron temperature, we will perhaps never know if such isothermal plasmas exist with a width narrower than $\\width$, based upon these data.  These properties prompt the following comments:  \\begin{enumerate} \\item Error bars quoted below $\\sim0.1$ in logarithmic temperatures are not credible using these techniques. Landi and Feldman quote error bars of 0.04 and 0.05 in their work which seeks to provide spectroscopic evidence that the coronal temperature is ``quantized''. \\item The widths (FWHM) of peaks in the $\\xi(T)$ functions found by \\citet[][their figure 5]{Landi+Feldman2008} are between 0.1 and 0.2. These widths are characteristic of the limit with which emission lines can determine the isothermality of the emitting plasma, and are therefore determined more by regularization than by the data themselves. Their analysis is consistent with isothermal plasma, but is limited by the uncertainties discussed here. \\item  Coronal temperatures are controlled by (unknown) heating mechanisms and cooling by heat conduction, flows, radiation.  The scalings of \\citet{Rosner+Tucker+Vaiana1978}, based on simple energy balance considerations,  show that a given coronal electron temperature $T_e$ in a loop of length $L$ requires an energy flux density $$ {\\cal F} \\propto (T_eL)^{7/2} $$ to sustain it agains conductive and radiation losses.  Therefore, for a given $L$, if $T_e$ is uncertain to $\\pm0.15$ in its logarithm, the absolute value of log$_{10}{\\cal F}$ can be determined from a given length loop to an accuracy of $\\sim \\pm0.5$.  The energy flux is constrained only to somewhere within a range spanning a factor of ten using emission lines. Conversely, variations of the energy flux ${\\cal F}$ dissipated in coronal loops by $\\lta 0.5$ in the logarithm cannot be distinguished through temperature dependent emission line methods. \\item  {\\em Relative} changes in ${\\cal F}$ could, however, be detected to within about $\\pm 12\\%$, but again this precision depends on the plasma being nearly isothermal (i.e. actual $\\width \\sim \\delt$), which is neither known from the data nor expected from first principles. \\end{enumerate}  The analytical results offer some hope that carefully selected sets of emission lines can be used to reduce the widths of distributions compatible with the data.  Specifically, using lines within the same element removes errors arising from incorrect abundances, and can include lines with large values of $|x_i|$.  In this case just a few lines with different values of $x_i$ can reduce $\\width$ via equation~(\\pref{eqn:bigx})).  Using lines from within the same ion stage yields significantly  smaller values of $\\epsilon_i \\sim 0.1$, but while the precision improves linearly with $\\epsilon_i$ the accuracy improves only quadratically."
    },
    "0911/0911.3170_arXiv.txt": {
        "abstract": "\\begin{list}{ } {\\rightmargin 1in} \\baselineskip = 11pt \\parindent=1pc {\\small We survey the basic principles of atmospheric dynamics relevant to explaining existing and future observations of exoplanets, both gas giant and terrestrial. Given the paucity of data on exoplanet atmospheres, our approach is to emphasize fundamental principles and insights gained from Solar-System studies that are likely to be generalizable to exoplanets.  We begin by presenting the hierarchy of basic equations used in atmospheric dynamics, including the Navier-Stokes, primitive, shallow-water, and two-dimensional nondivergent models.  We  then survey key concepts in atmospheric dynamics, including the importance of planetary rotation, the concept of balance, and simple scaling arguments to show how turbulent interactions generally produce large-scale east-west banding on rotating planets.  We next turn to issues specific to giant planets, including their expected interior and atmospheric thermal structures, the implications for their wind patterns, and mechanisms to pump their east-west jets.  Hot Jupiter atmospheric dynamics are given particular attention, as these close-in planets have been the subject of most of the concrete developments in the study of exoplanetary atmospheres. We then turn to the basic elements of circulation on terrestrial planets as inferred from Solar-System studies, including Hadley cells, jet streams, processes that govern the large-scale horizontal temperature contrasts, and climate, and we discuss how these insights may apply to terrestrial exoplanets.   Although exoplanets surely possess a greater diversity of circulation regimes than seen on the planets in our Solar System, our guiding philosophy is that the multi-decade study of Solar-System planets reviewed here provides a foundation upon which our understanding of more exotic exoplanetary meteorology must build. \\\\~\\\\~\\\\~}%  \\end{list} ",
        "introduction": "\\label{Intro}  The study of atmospheric circulation and climate began hundreds of years ago with attempts to understand the processes that determine the distribution of surface winds on the Earth \\citep[e.g.,][]{hadley-1735}.  As theories of Earth's general circulation became more sophisticated \\citep[e.g.,][]{lorenz-1967}, the characterization of Mars, Venus, Jupiter, and other Solar-System planets by spacecraft starting in the 1960s demonstrated that the climate and circulation of other atmospheres differ, sometimes radically, from that of Earth. Exoplanets, occupying a far greater range of physical and orbital characteristics than planets in our Solar System, likewise plausibly span an even greater diversity of circulation and climate regimes.  This diversity provides a motivation for extending the theory of atmospheric circulation beyond our terrestrial experience.  Despite continuing questions, our understanding of the circulation of the modern Earth atmosphere is now well developed \\citep[see, e.g.,][]{held-2000, schneider-2006, vallis-2006}, but attempts to unravel the atmospheric dynamics of Venus, Jupiter, and other Solar-System planets remain ongoing, and the study of atmospheric circulation of exoplanets is in its infancy.  For exoplanets, driving questions fall into several overlapping categories.  First, we wish to understand and explain new observations constraining atmospheric structure, such as light curves, photometry, and spectra obtained with the {\\it Spitzer, } {\\it Hubble}, or {\\it James Webb Space Telescopes (JWST)}, thus helping to characterize specific exoplanets as remote worlds.   Second, we wish to extend the theory of atmospheric circulation to the wide range of planetary parameters encompassed by exoplanets. Existing theory was primarily developed for conditions relevant to Earth, and our understanding of how atmospheric circulation depends on atmospheric mass, composition, stellar flux, planetary rotation rate, orbital eccentricity, and other parameters remains rudimentary. Significant progress is possible with theoretical, numerical, and laboratory investigations that span a wider range of planetary parameters.  Third, we wish to understand the conditions under which planets are habitable, and answering this question requires addressing the intertwined issues of atmospheric circulation and climate.   \\begin{figure*} \\epsscale{1.0} \\plotone{coupled-dynamics-radiation-crop2.jpg} \\caption{\\small Atmospheric circulation results from a coupled interaction between radiation and hydrodynamics:  horizontal temperature and pressure contrasts generate winds, which drive the atmosphere away from local radiative equilibrium.  This in turn allows the spatially variable thermodynamic (radiative) heating and cooling that maintains the horizontal temperature and pressure contrasts.} \\label{flow-chart} \\end{figure*}  What drives atmospheric circulation? Horizontal temperature contrasts imply the existence of horizontal pressure contrasts, which drive winds. The winds in turn push the atmosphere away from radiative equilibrium by transporting heat from hot regions to cold regions (e.g., from the equator to the poles on Earth).  This deviation from radiative equilibrium allows net radiative heating and cooling to occur, thus helping to maintain the horizontal temperature and pressure contrasts that drive the winds (see Fig.~\\ref{flow-chart}). Spatial contrasts in thermodynamic heating/cooling thus fundamentally drive the circulation, yet it is the existence of the circulation that allows these heating/cooling patterns to exist. (In the absence of a circulation, the atmosphere would relax into a radiative-equilibrium state with a net heating rate of zero.)  The atmospheric circulation is thus a coupled radiation-hydrodynamics problem.  On the Earth, for example (see Fig.~\\ref{earth-radiation-balance}), the equator and poles are not in radiative equilibrium.  The equator is subject to net heating, the poles to net cooling, and it is the mean latitudinal heat transport that is both responsible for and driven by these net imbalances.   \\begin{figure*} \\epsscale{1.0} \\plotone{earth-radiation-balance.jpg} \\caption{\\small Earth's energy balance.  The Earth absorbs more sunlight at the equator than the poles (blue curve, denoted ``shortwave'').  The Earth also radiates more infrared energy to space at the equator than the poles (red curve, denoted ``longwave''). However, because the atmospheric/oceanic circulation act to mute the latitudinal temperature contrasts (relative to radiative equilibrium), the longwave radiation exhibits less latitudinal variation than the shortwave absorption.  Thus, the circulation leads to net heating at the equator and net cooling at the poles, which in turn drives the circulation. Data are an annual-average for 1987 obtained from the NASA Earth Radiation Budget Experiment (ERBE) project.  Copyright M. Pidwirny, www.physicalgeography.net, used with permission.} \\label{earth-radiation-balance} \\end{figure*}  The mean climate (e.g., the global-mean surface temperature of a planet) depends foremost on the absorbed stellar flux and the atmosphere's need to reradiate that energy to space.  Yet even the global-mean climate is strongly affected by the atmospheric mass, composition, and circulation.  On a terrestrial planet, for example, the circulation helps to control the distribution of clouds and surface ice, which in turn determine the planetary albedo and the mean surface temperature. In some cases, a planetary climate can have multiple equilibria (e.g., a warm, ice-free state or a cold, ice-covered ``snowball Earth'' state), and in such cases the circulation plays an important role in determining the relative stability of these equilibria.  Understanding the atmosphere/climate system is challenging because of its nonlinearity, which involves multiple positive and negative feedbacks between radiation, clouds, dynamics, surface processes, planetary interior, and life (if any).  The inherent nonlinearity of fluid motion further implies that even atmospheric-circulation models neglecting the radiative, cloud, and surface/interior components can exhibit a large variety of behaviors.   From the perspective of studying the atmospheric circulation, transiting exoplanets are particularly intriguing because they allow constraints on key planetary attributes that are a prerequisite to characterizing an atmosphere's circulation regime.  When combined with Doppler velocity data, transit observations permit a direct measurement of the exoplanet's radius, mass and thus surface gravity\\footnote{Combining Doppler velocity and transit measurements lifts the mass-inclination degeneracy.}.  With the additional expectation that close-in exoplanets are tidally locked if on a circular orbit, or pseudo-synchronized\\footnote{Pseudo-synchronization refers to a state of tidal synchronization achieved only at periastron passage (=closest approach), as expected from the strong dependence of tides with orbital separation.} if on an eccentric orbit, the planetary rotation rate is thus indirectly known as well.  Knowledge of the radius, surface gravity, rotation rate and external irradiation conditions for several exoplanets, together with the availability of direct observational constraints on their emission, absorption and reflection properties, opens the way for the development of comparative atmospheric science beyond the reach of our own Solar System.    The need to interpret these astronomical data reliably, by accounting for the effects of atmospheric circulation and understanding its consequences for the resulting planetary emission, absorption and reflection properties, is the central theme of this chapter. Tidally locked close-in exoplanets, for example, are subject to an unusual situation of permanent day/night radiative forcing, which does not exist in our Solar System\\footnote{Venus may provide a partial analogy, which has not been fully exploited yet.}.  To address the new regimes of forcings and responses of these exoplanetary atmospheres, a discussion of fundamental principles of atmospheric fluid dynamics and how they are implemented in multi-dimensional, coupled radiation-hydrodynamics numerical models of the GCM (General Circulation Model) type is required.   Contemplating the wide diversity of exoplanets raises a number of fundamental questions.  What determines the mean wind speeds, direction, and 3D flow geometry in atmospheres?  What controls the equator-to-pole and day-night temperature differences?  What controls the frequencies and spatial scales of temporal variability? What role does the circulation play in controlling the mean climate (e.g., global-mean surface temperature, composition) of an atmosphere? How do these answers depend on parameters such as the planetary rotation rate, gravity, atmospheric mass and composition, and stellar flux?  And, finally, what are the implications for observations and habitability of exoplanets?   At present, only partial answers to these questions exist \\citep[see reviews by][]{showman-etal-2008b, cho-2008}. With upcoming observations of exoplanets, constraints from Solar-System atmospheres, and careful theoretical work, significant progress is possible over the next decade. While a rich variety of atmospheric flow behaviors is realized in the Solar System alone---and an even wider diversity is possible on exoplanets---the fundamental physical principles obeyed by all planetary atmospheres are nonetheless universal. With this unifying notion in mind, this chapter provides a basic description of atmospheric circulation principles developed on the basis of extensive Solar-System studies and discusses the prospects for using these principles to better understand physical conditions in the atmospheres of remote worlds.   The plan of this chapter is as follows. In Section~\\ref{equations}, we introduce several of the equation sets that are used to investigate atmospheric circulation at varying levels of complexity. This is followed (Section~\\ref{basic-concepts}) by a tutorial on basic ideas in atmospheric dynamics, including atmospheric energetics, timescale arguments, force balances relevant to the large-scale circulation, the important role of rotation in generating east-west banding, and the role of waves and eddies in shaping the circulation.  In Section~\\ref{giants}, we survey the atmospheric dynamics of giant planets, beginning with generic arguments to constrain the thermal and dynamical structure and proceeding to specific models for understanding the circulation of our ``local'' giant planets (Jupiter, Saturn, Uranus, Neptune) as well as hot Jupiters and hot Neptunes.\\footnote{The terms hot Jupiter and hot Neptune refer to giant exoplanets with masses comparable to those of Jupiter and Neptune, respectively, with orbital semi-major axes less than $\\sim0.1\\,$AU, leading to high temperatures.}  In Section~\\ref{terrestrial}, we turn to the climate and circulation of terrestrial exoplanets. Observational constraints in this area do not yet exist, and so our goal is simply to summarize basic concepts that we expect to become relevant as this field expands over the next decade.  This includes a description of climate feedbacks (Section~\\ref{climate}), global circulation regimes (Section~\\ref{regimes}), Hadley-cell dynamics (Section~\\ref{hadley}), the dynamics of the so-called midlatitude ``baroclinic'' zones where baroclinic instabilities dominate (Section~\\ref{baroclinic-zone}), the slowly rotating regime relevant to Venus and Titan (Section~\\ref{slowly-rotating}), and finally a survey of how the circulation responds to the unusual forcing associated with synchronous rotation, extreme obliquities, or extreme orbital eccentricities (Section~\\ref{unusual-forcing}). The latter topics, while perhaps the most relevant, are the least understood theoretically.  In Section~\\ref{highlights} we summarize recent highlights, both observational and theoretical, and in Section~\\ref{future} we finish with a survey of future prospects.     \\bigskip ",
        "conclusions": ""
    },
    "0911/0911.2495_arXiv.txt": {
        "abstract": "We give the evolution and constraint equations on an asymmetrically embedded brane in the form of average and difference equations. ",
        "introduction": "In a recent paper \\cite{3+1+1} we have developed the covariant gravitational dynamics in a 3+1+1 dimensional space-time in the spirit of the general relativistic 3+1 covariant cosmology \\cite{3+1}, generalizing both previous approaches to 5-dimensional (5d) gravitational dynamics \\cite{brane-dynamics}% , brane 3+1 covariant cosmology \\cite{Maartens-eqs} and 2+1+1 covariant dynamics \\cite{2+1+1}. Such a formalism may turn useful in discussing perturbations on the brane \\cite{branepert}. The singled-out directions are the normal $n^{a}$ to the hypersurface representing the time-evolving 3-dimensional space (the brane with metric $h_{ab}$ and tension $\\lambda $) and a temporal direction $u^{a}$ tangent to this hypersurface. All employed gravitational variables are projected to the brane. They consist of kinematical variables $(\\Theta ,\\sigma _{ab},\\omega _{ab},A_{a},K_{a},L_{a})$ related to the vector $u^{a}$; analogous quantities, carrying an overhat, related to $n^{a}$; two kinematical scalars $(K,\\hat{K})$ related to both; finally gravito-electro-magnetic quantities $(\\mathcal{E},\\mathcal{H}_{k},% \\mathcal{F}_{kl},\\mathcal{E}_{k},\\widehat{\\mathcal{E}}_{k},\\mathcal{E}_{kl},% \\widehat{\\mathcal{E}}_{kl},\\mathcal{H}_{kl},\\widehat{\\mathcal{H}}_{kl})$. For details on the definitions of these quantities see Ref. \\cite{3+1+1}. The matter on the brane has been decomposed as $T_{ab}=\\rho u_{a}u_{b}\\!+\\!q_{(a}u_{b)}\\!+\\!ph_{ab}\\!+\\!\\pi _{ab},$ while possible non-standard model fields nesting in the 5d space-\\negthinspace time as $% \\widetilde{T}_{ab}=\\widetilde{\\rho }u_{a}u_{b}\\!+\\!2\\widetilde{q}% _{(a}u_{b)}\\!+\\!2\\widetilde{q}u_{(a}n_{b)}\\!+\\!\\widetilde{p}h_{ab}\\!+\\!% \\widetilde{\\pi }n_{a}n_{b}\\!+\\!2\\widetilde{\\pi }_{(a}n_{b)}\\!+\\!\\widetilde{% \\pi }_{ab}.$ The gravitational coupling constants $\\widetilde\\kappa^2$ and $% \\kappa^2=\\widetilde\\kappa^4\\lambda/6$ act in 5d and on the brane, respectively.  The generic form of the evolution and constraint equations in the 5-dimensional spacetime were given as Appendix C of Ref. \\cite{3+1+1} and they were specified on a $Z_{2}$-symmetrically embedded brane in Subsection IV.D. The embedding however is not necessarely symmetric: an asymmetric embedding is known to generate late time acceleration in a cosmological setup \\cite{Induced}. Therefore here we generalize the formalism to an asymmetrically embedded brane. Following the recipe presented in Subsection IV.D of Ref. \\cite{3+1+1}, we give here the evolution and constraint equations obtained both as averages and differences across the brane. We denote the average over the two sides of the brane of any quantity by angle brackets and the jump by $\\Delta $. For the extrinsic curvature components $% \\mathcal{K}\\equiv (\\widehat{\\Theta },\\widehat{\\sigma }_{ab},\\widehat{K},% \\widehat{K}_{a})$ the latter is directly related to the brane matter variables and brane tension cf. the Israel-Lanczos condition, Eqs. (76)-(79) of Ref. \\cite{3+1+1}. Angle brackets on indices indicate projection to the 3-space, symmetrization and trace-free character. ",
        "conclusions": "Both the average and the difference equations reduce to the corresponding equations given in Subsection IV.D of Ref \\cite{3+1+1}, by taking into account that in the particular case of a symmetric embedding for quantities defined with an odd (even) number of $n^{a}$, the conditions $\\Delta f=2f,\\langle f\\rangle =0$ ($\\Delta f=0,\\langle f\\rangle =f$) hold. In particular, the extrinsic curvature components $\\mathcal{K}$ belong to the first group.  \\textit{Acknowledgements}: This work was supported by the Pol\\'{a}nyi and Sun Programs of the Hungarian National Office for Research and Technology (NKTH), and by the Hungarian Scientific Research Fund (OTKA) grant 69036. We acknowledge financial support from the organizers of the Grassmannian Conference in Fundamental Cosmology (Grasscosmofun'09)."
    },
    "0911/0911.0035_arXiv.txt": {
        "abstract": "We use {\\it Spitzer} data to infer that the small infrared excess of V819 Tau, a weak-lined T Tauri star in Taurus, is real and not attributable to a ``companion'' 10\\arcsec\\ to the south. We do not confirm the mid-infrared excess in HBC 427 and V410 X-ray 3, which are also non-accreting T Tauri stars in the same region; instead, for the former object, the excess arises from a red companion 9\\arcsec\\ to the east. A single-temperature blackbody fit to the continuum excess of V819 Tau implies a dust temperature of 143 K; however, a better fit is achieved when the weak 10 and 20 $\\mu$m silicate emission features are also included. We infer a disk of sub-$\\mu$m silicate grains between about 1 AU and several 100 AU with a constant surface density distribution. The mid-infrared excess of V819 Tau can be successfully modeled with dust composed mostly of small amorphous olivine grains at a temperature of 85 K, and most of the excess emission is optically thin. The disk could still be primordial, but gas-poor and therefore short-lived, or already at the debris disk stage, which would make it one of the youngest debris disk systems known. ",
        "introduction": "The evolution and dissipation of protoplanetary disks has been an active area of research for the last few decades; especially the advent of sensitive near- and mid-infrared observations has allowed us to explore inner disk regions (out to a few AU), encompassing the radii where planets are thought to form. The role of planet formation in disk dissipation has been all but proven; while disks likely evolve from a flared, optically thick configuration to a flat, settled disk via grain growth and settling, it is not clear to what extent the eventual dissipation of the remaining dust and gas and the formation of planets are linked.  A key evolutionary stage in disk dissipation is the transition from the classical T Tauri stage, when a pre-main-sequence star is surrounded by an accreting disk, to the weak-lined T Tauri phase, when accretion ends and the disk disappears. This transitional stage is believed to be short ($\\lesssim$ 10$^5$ years), since only few stars with vanishing infrared excesses have been observed \\citep[e.g.,][]{skrutskie90,simon95}. The Infrared Spectrograph\\footnote{The IRS was a collaborative venture between Cornell University and Ball Aerospace Corporation funded by NASA through the Jet Propulsion Laboratory and the Ames Research Center.} \\citep[IRS;][]{houck04} on board the {\\it Spitzer Space Telescope} \\citep{werner04} has provided new insights into the study of transitional disks. For example, in the nearby and well-studied Taurus star-forming region \\citep{kenyon95}, out of a total of 111 T Tauri stars, five objects were identified whose inner disk regions are cleared to different degrees, but outer, optically thick disks remain, delimited by a well-defined inner disk rim \\citep{dalessio05,calvet05,furlan06,espaillat07}. Planet formation, photoevaporation, or inner disk draining induced by the magneto-rotational instability may play a role in removing the inner disk \\citep{clarke01,quillen04,alexander07,chiang07}, but also close binary companions can clear inner regions due to orbital resonances \\citep{artymowicz94,ireland08}.  V819 Tau is a weak-lined T Tauri star with a spectral type of K7 in the Taurus star-forming region \\citep{herbig88}. We presented its {\\it Spitzer} IRS spectrum in \\citet{furlan06} as part of the sample of Class III objects in Taurus, noting that it showed an infrared excess beyond about 12 $\\mu$m (see Figure \\ref{V819Tau_IRS}). However, we could not confirm whether this excess was real due to the presence of a ``companion'' 10\\arcsec\\ to the south that is detected in 2MASS images \\citep{skrutskie06}. Two other weak-lined T Tauri stars in Taurus, HBC 427 and V410 X-ray 3, with spectral types of K5 and M6, respectively \\citep{steffen01,strom94}, were also introduced in \\citet{furlan06} as objects with uncertain infrared excesses. While HBC 427 has a ``companion'' star about 15\\arcsec\\ to the southeast that is seen in 2MASS images, V410 X-ray 3 appears to be single. In the meantime, we obtained IRS peak-up images of V819 Tau and HBC 427, and we re-reduced the IRS spectra of all three objects. We do not confirm the mid-infrared excess in HBC 427 and V410 X-ray 3, but, together with Multiband Imaging Photometer for {\\it Spitzer} \\citep[MIPS;][]{rieke04} 24 $\\mu$m images from the {\\it Spitzer} archive, we establish that the infrared excess is intrinsic to V819 Tau, and we apply simple models to derive the distribution of dust around this T Tauri star.  \\begin{figure} \\centering \\includegraphics[angle=90, scale=0.35]{f1.eps} \\caption{The spectral energy distribution of V819 Tau; the optical photometry is from \\citet{kenyon95}, JHK$_s$ photometry from 2MASS \\citep{skrutskie06}, IRAC fluxes from \\citet{luhman06}, and the 450 $\\mu$m upper limit from \\citet{andrews05}. The IRS spectrum (also shown with error bars in the figure inset), IRS peak-up fluxes, and MIPS fluxes are from this work. The data were dereddened using Mathis's reddening law for $R_V$=3.1 \\citep{mathis90} and $A_V$=1.7 \\citep{furlan06}. The stellar photosphere is a scaled NextGen model with $T_{eff}$=4000 K and $log(g)$=3.5 (see text for details). \\label{V819Tau_IRS}} \\end{figure}  \\newpage ",
        "conclusions": "Of the three Class III objects in Taurus that showed tentative mid-infrared excess emission in \\citet{furlan06}, we could confirm an infrared excess only in V819 Tau. While V819 Tau is a single star, HBC 427 and V410 X-ray 3 are close binaries: HBC 427 is a spectroscopic binary with a semi-major axis of about 0.03\\arcsec, comprised of a K5 and an M2 star \\citep{steffen01}; V410 X-ray 3 consists of an M6 and an M7.7 star separated by $\\sim$ 0.05\\arcsec\\ \\citep{kraus06}. Neither object shows signs of accretion \\citep{strom94, kenyon98, mohanty05}, and therefore both dust and gas have already dissipated in their inner disks. HBC 427 is also not detected at sub-millimeter wavelengths, implying an upper limit for the disk mass of 0.0007 \\Msun \\citep{andrews05}. It is likely that interaction between the binary and the disk in these systems is the cause for the absence of disk material \\citep{jensen94}.  V819 Tau is also not accreting any more; its H${\\alpha}$ 10\\% width amounts to 166 km s$^{-1}$ \\citep{nguyen09}, which is considerably less than the lower limit of 270 km s$^{-1}$ that characterizes accretors \\citep{white03}. Also the low H${\\alpha}$ equivalent width of 1.7-3.2 {\\AA} \\citep{strom94, kenyon98} confirms its nature as a weak-lined T Tauri star. \\citet{white01} determined an upper limit of $1.4 \\times 10^{-9}$ {\\Msun} yr$^{-1}$ for the mass accretion rate based on the lack of $U$-band excess. No H$_2$ emission from warm inner disk regions was detected \\citep{bary03}, implying that both the gas and dust have been removed there. Our data indicate that there are no small, warm dust grains within approximately 1 AU from the star.  One process that can generate cleared inner disk regions is photoevaporation \\citep{clarke01}; it sets in once the mass accretion rate has dropped to levels below a few 10$^{-10}$ {\\Msun} yr$^{-1}$, and it acts by eroding the disk in a photoevaporative flow beyond a certain radius, which, for V819 Tau, amounts to 7.1 AU. This value is larger than the inner radius of our best-fit optically thin disk models, but it is remarkably close to the value of 6.8 AU we estimated by assuming the excess emission to arise from an optically thick, blackbody surface. However, it seems that the dust emission is optically thin, given the presence of 10 and 20 $\\mu$m silicate emission features and the low infrared excess luminosity ($L_{IR}/L_{bol} \\sim $ 10$^{-3}$). Moreover, our models require small dust grains at distances of at least $\\sim$ 1 AU from the star. Thus, photoevaporation cannot explain the inner disk clearing of V819 Tau; a change in grain opacity caused by grain growth to sizes well beyond 1 $\\mu$m in this region, or the gravitational influence of another body, such as a planet, are possible explanations for the inner disk clearing.  The question also arises whether the material we detect around V819 Tau is primordial or already second-generation dust. Detection of gas emission from the inner disk would support the primordial nature of the disk; so far, the absence of accretion signatures and H$_2$ emission from the inner disk suggest that the gas has already dissipated. Without the gas, dust particles would be subject to collisions, Poynting-Robertson (PR) drag, radiation pressure, and corpuscular stellar wind pressure \\citep[e.g.,][]{wyatt99, chen06}, which, in the region from $\\sim$ 1 to 100 AU from V819 Tau, would limit the lifetime of sub-$\\mu$m grains to $\\lesssim$ 10$^5$ years.  Therefore, in the absence of gas, dust would have to be continuously replenished, since the age of V819 Tau is about 2 Myr. V819 Tau would be one of the youngest debris disk systems; its infrared excess would lie on the higher side of what is usually measured in debris disks \\citep[e.g.,][]{chen06}, but still within the observed range. Collisions among planetesimals, as are thought to occur in debris disks, typically produce larger grains ($\\gtrsim$ 10 $\\mu$m), since the mid-infrared spectra of the majority of debris disks are featureless \\citep{chen06,carpenter09}. The fact that small, amorphous silicate grains dominate the optically thin emission of V819 Tau, with just a minor possible presence of crystalline silicates, suggests that the dust we observe consists of mostly pristine material. Colliding or evaporating comet nuclei, which likely form in the outer disk regions and therefore incorporate more unprocessed silicates into ice, could be the source of this pristine dust.  It is also possible that we are just observing ``true'' primordial, unprocessed dust from large disk radii; this would imply that we are witnessing a rare stage in the life of a protoplanetary disk, when the gas has already been mostly removed from the inner disk, and the remaining dust is dissipated on a timescale of 10$^5$ years or less. Then, not only could the disk structure of V819 Tau be described as transitional due to the lack of infrared excess below $\\sim$ 8 $\\mu$m and the clear presence of such an excess beyond 12 $\\mu$m, but also the evolutionary stage of this object could be considered transitional, i.e., having just transitioned from a primordial, optically thick disk to an optically thin one. Even though a few percent of the T Tauri population in Taurus can be considered as transitional (defined as disks with substantial inner disk clearings and thus low near-infrared excesses; \\citealt{furlan09}), V819 Tau would stand out by its very low infrared excess. It would be the only object in Taurus observed so far at an advanced transitional stage, when the disk material is already optically thin. Given the small amount of dust left in the disk, planet formation would have to be already completed in this disk; this conclusion would also apply if the dust were second-generation and therefore the result of dynamically perturbed, remnant planetary building blocks."
    },
    "0911/0911.0729_arXiv.txt": {
        "abstract": "One of the main science goal of the future European Extremely Large Telescope will be to understand the mass assembly process in galaxies as a function of cosmic time. To this aim, a multi-object, AO-assisted integral field spectrograph will be required to map the physical and chemical properties of very distant galaxies. In this paper, we examine the ability of such an instrument to obtain spatially resolved spectroscopy of a large sample of massive ($0.1 \\leq M_{stellar} \\leq 5\\times10^{11}M_{\\odot}$) galaxies at $2\\leq z < 6$, selected from future large area optical-near IR surveys. We produced a set of about one thousand numerical simulations of 3D observations using reasonable assumptions about the site, telescope, and instrument, and about the physics of distant galaxies. These data-cubes were analysed as real data to produce realistic kinematic measurements of very distant galaxies. We then studied how sensible the scientific goals are to the observational (i.e., site-, telescope-, and instrument-related) and physical (i.e., galaxy-related) parameters. We specifically investigated the impact of AO performance on the science goal. We did not identify any breaking points with respect to the parameters (e.g., the telescope diameter), with the exception of the telescope thermal background, which strongly limits the performance in the highest (z$>$5) redshift bin. We find that a survey of $N_{gal}$ galaxies that fulfil the range of science goals can be achieved with a $\\sim$90 nights program on the E-ELT, provided a multiplex capability $M \\sim N_{gal}/8$. ",
        "introduction": "Over the last decade, the synergy of 8-10 meter class telescopes with HST has strongly re-invigorated the field of galaxy formation and evolution by unveiling very distant galaxies up to z$\\sim$6 (e.g., \\citealt{cuby03,rhoads03,bremer04,bouwens06}), by allowing the first determination of the global star formation history since redshift z$\\sim$6 (e.g., \\citealt{hopkins04}), and by providing the first insights on the stellar mass assembly history out to z$\\sim$5 (e.g., \\citealt{drory05,pozzetti07,marchesini07,perez07}). Despite these recent progresses, the outstanding question remains on how and when galaxies assembled their baryonic mass across cosmic time. The CDM standard model has provided a satisfactory scenario describing the hierarchical assembly of dark matter halos, in a bottom-up sequence which is now well-established over the whole mass structure spectrum. In contrast, little progress has been made in the physical understanding of the formation and evolution of the baryonic component because the conversion of baryons into stars is a complex, poorly understood process.   As a result, all intellectual advances in galaxy formation and evolution over the last decade have been essentially empirical, often based on phenomenological (or semi-analytical) models, which heavily rely on observations to describe, with simplistic rules, such processes as star formation efficiency, energy feedback from star formation and AGN, chemical evolution, angular momentum transfer in merging, etc. Cornerstones observations in this empirical framework are the total and stellar mass of galaxies and their physical properties, including the age and metallicities of their underlying stellar populations, dust extinction, star formation rate, and structural/morphological parameters. The study of well-established scaling relations involving a number of these physical parameters (e.g. mass-metallicity, fundamental plane, colour-magnitude, morphology-density) are essential for understanding the physical processes driving galaxy evolution. However, with the current generation of 10m-class telescopes, we have been able to construct for example the fundamental plane of early-type galaxies, or to measure the Tully-Fisher relation of late-type galaxies over a wide range of masses only at low and intermediate redshifts (z$<$1), whereas only the brightest or most massive galaxies have been accessible at z $>$ 1, and a direct measurement of masses has been almost completely out of reach at z $>$ 2. Thus, our ability to explore the evolution and origin of the aforementioned scaling relations has rapidly reached the limit of 10m-class telescopes.   Hence, most of the outstanding questions arisen from recent observational galaxy evolution studies which have pushed the 10m-class telescopes to their limits call for an Extremely Large Telescope (ELT), specifically to extend the spectroscopic limit by at least two magnitudes with near-diffraction limit angular resolution. IFU spectrographs on 8-10m class telescopes (e.g., VLT/SINFONI and Keck/Osiris) are now routinely deriving the spatially-resolved kinematics of massive (i.e., with stellar masses larger than 10$^{10}M_\\odot$), distant galaxies, from z$\\sim$0.4 to z$\\sim$3 \\citep{flores06,puech06,yang08,forster06,wright07,shapiro08,bournaud08,genzel08,vanstarkenburg08,law09,wright09,epinat09,forsterschreiber09,lemoine09}. Such studies have brought new and essential insights into galaxy evolution processes, such as the fraction of rotators as a function of redshift \\citep{yang08,forsterschreiber09}, a better understanding of the evolution of the Tully-Fisher relation \\citep{puech08,cresci09,puech09}, or the first glimpse into the evolution of the angular momentum \\citep{puech07,bouche07}. However, several uncertainties and limitations remain, especially at the highest redshifts (i.e., z$>$1). For instance, z$\\sim$2 surveys were drawn from different selection criteria (e.g., BX, BM, or BzK galaxies), and this might bias the resulting sample in a subtle and quite uncontrolled way. Besides, samples in the NIR are selected in optimal atmospheric windows, free of OH sky lines and maximising the atmospheric throughput, which intrinsically limits the redshift range and size of the resulting sample. The only way to overcome these issues is to move to telescopes with larger collecting areas.    ESO is currently developing a Phase B study for a European 42-meter telescope \\citep{gilmozzi08,spyromilio08}. As part of the project development, the Design Reference Mission (DRM) aims at producing a set of observing proposals and corresponding simulated data which together provide the project with a tool to assess the extent to which the E-ELT addresses key scientific questions and assist in critical trade-off decisions\\footnote{see http://www.eso.org/sci/facilities/eelt/science/drm/}. In the frame of the DRM effort, simulations of 3D high-z galaxy observations were undertaken. These observations are expected to yield direct kinematics of stars and gas in the first generation of massive galaxies (in the range 0.1 $\\leq$ M $\\leq$ 5$\\times$10$^{11}$M$_\\odot$), as well as their stellar population properties. This will allow one to derive dynamical masses, ages, metallicities, star-formation rates, dust extinction maps, to investigate the presence of disk and spheroidal components and the importance of dynamical processes (e.g. merging, in/outflows) which govern galaxy evolution. These data will also allow one to study the onset of well known scaling relations at lower redshifts, and to witness the gradual shift of star formation from the most massive galaxies in the highest density regions to less massive galaxies in the field.   To assess the science achievement of an NIR IFS on the E-ELT, as well as to better understand the interactions between the telescope, site, and instrument, we have produced an extensive set of $\\sim$1000 simulations of observations of very distant galaxies. It is important to realize that if one wants to study what kind of instrument is needed to understand the galaxy mass assembly process, then one needs to explore a huge parameter space, i.e., from relaxed smooth or clumpy rotating disks to major and minor mergers, taking care of spanning all possible mass ratios, viewing angles, merger geometry and so on, which would make such a direct approach very difficult, and probably even impossible, on the practical side. Moreover, one has to recognise that our knowledge and understanding of the physics of distant galaxies is still incomplete, which necessitates to extrapolate some of their characteristics. Therefore, the simulations presented in this paper do not intend to be exhaustive in any sense. Rather, we aim at exploring this very large parameter space using reasonable assumptions and guesses. Indeed, the DRM exercise consisted in a broad exploration of the parameter space in realistic observing conditions, in order to identify possible limitations and/or breaking points, as well as interactions between the telescope, site, and instrument, which could potentially impact the telescope design.   This paper is the second of a series that explores the performances of a NIR IFS on the E-ELT. In the first paper, we presented our methodology to simulate realistic observations of distant galaxies \\citep{puech08b}. We also produced first simulations and explored performance using a few scientifically-motivated cases. They illustrated the concept of ``scale-coupling'', i.e., the relationship between the IFU pixel scale and the size of the kinematic features that need to be recovered by 3D spectroscopy in order to understand the nature of the galaxy and its substructure. In \\cite{puech08b}, we focused on the largest spatial scales, which are of particular interest because they carry most of the kinematic information useful to reveal the process underlying galaxy dynamics, i.e., whether a given galaxy is in a coherent and stable dynamical state (e.g., rotation), or, on the contrary, out of equilibrium (e.g., subsequently to a merger).   In this paper, we present an updated version of the simulation pipeline. The main improvement is related to the inclusion of thermal emissions from both the IFS and the telescope. This allowed us to explore a wider parameter space, and especially areas where observations are limited by the thermal emission rather than by the sky background (e.g., relatively faint targets in the K band). We also significantly extended the range of morpho-kinematic templates and of physical parameters (e.g., radius, mass, and velocity as a function of redshift) considered in the simulations. This paper is organised as follows. In Sect. 2, we present our methodology and the pipeline used for simulations. In Sect. 3, we detail the scientific and observational inputs used in the simulations. In Sect. 4, we present the results of the simulations, which are discussed in Sect. 5. A conclusion is drawn in Sect 6. Throughout, we adopt $H_0=70$ km/s/Mpc, $\\Omega _M=0.3$, and $\\Omega _\\Lambda=0.7$, and the $AB$ magnitude system. ",
        "conclusions": "We have conducted simulations of the ``Physics and mass assembly of galaxies out of z$\\sim$6'' science case for the E-ELT, exploring a wide range of observational and physical parameters. We have defined figures of merit for this science case despite the inherent complexity of the science goals, derived empirical scaling relations between the signal-to-noise ratio of kinematic and intensity maps and the main telescope and instrument parameters, as well as a relation between the limit in stellar mass that can be reached for a given signal-to-noise ratio as a function of redshift. We specifically investigated the impact of AO performance on the science goal. We did not identify any breaking points with respect to all parameters (e.g., the telescope diameter), with the exception of the telescope thermal background, which strongly limits the performance in the highest (z$>$5) redshift bin. We find that the full range of science goals can be achieved with a $\\sim$100 nights program on the E-ELT, provided a high multiplex advantage $M\\sim N_{gal}/8$. We stress that several assumptions and guided guesses had to be made on both the observational conditions and physical characteristics of distant galaxies under study. This introduces an inherent uncertainty, which can be mitigated with future simulations as the telescope design will be consolidated and more details about the physics of high-z galaxies will become available."
    },
    "0911/0911.2212_arXiv.txt": {
        "abstract": "We have used the Gemini Near-infrared Integral Field Spectrograph (NIFS) to map the gas kinematics of the  inner $\\sim$\\, 200$\\times$500\\,pc of the Seyfert galaxy NGC\\,4151 in the Z, J, H and K bands at a resolving power $\\ge$\\,5000 and spatial resolution of $\\sim$\\,8\\,pc. The ionised gas emission is most extended along the known ionisation bi-cone at position angle PA=60--240$\\degr$, but is observed also along its equatorial plane. This indicates that the AGN ionizes gas beyond the borders of the bi-cone, within a sphere with $\\approx$\\,1\\arcsec\\ radius around the nucleus. The ionised gas has three kinematic components: (1) one observed at the systemic velocity and interpreted as originating in the galaxy disk; (2) one outflowing along the bi-cone, with line-of-sight velocities between $-600$ and 600\\,\\kms\\ and strongest emission at $\\pm\\,(100-300)$\\,\\kms;  (3) and another component due to the interaction of the radio jet with ambient gas.  The radio jet (at PA=75--255$\\degr$) is not aligned with the NLR, and produces flux enhancements  mostly observed at the systemic velocity, suggesting that the jet is launched close to the plane of the  galaxy ($\\sim$ plane of the sky).  The mass outflow rate, estimated to be $\\approx$\\,1\\,M$_\\odot$\\,yr$^{-1}$ along each cone, exceeds the inferred black hole accretion rate by a factor of $\\sim$\\,100. This can be understood if the Narrow-Line-Region (NLR) is formed mostly by entrained gas from the circumnuclear interstellar medium by an outflow probably originating in the accretion disk. This flow represents feedback from the AGN, estimated to release a kinetic power of $\\dot{E}\\approx\\,2.4\\,\\times10^{41}$\\,erg\\,s$^{-1}$, which is only $\\sim$\\,0.3\\% of the bolometric luminosity of the AGN.  There is no evidence in our data for the gradual acceleration followed by gradual deceleration proposed by previous modelling of the [OIII] emitting gas. Our data allow the possibility that the NLR clouds are accelerated close to the nucleus (within 0$\\farcs$1 -- $\\approx$\\,6\\,pc) after which the flow moves at essentially constant velocity ($\\approx\\,600$\\kms), being consistent with NIR emission arising predominantly from the interaction of the outflow with gas in the galactic disk.  The molecular gas exhiibits distinct kinematics relative to the ionised gas. Its emission arises in extended regions approximately perpendicular to the axis of the bi-cone and along the axis of the galaxy's stellar bar, avoiding the innermost ionised regions. It does not show an outflowing component, being observed only at velocities very close to systemic, and is thus consistent with an origin in the galaxy plane. This hot molecular gas may only  be the tracer of a larger reservoir of colder gas which represents the AGN feeding. ",
        "introduction": "This is a continuing study of the narrow-line region (hereafter NLR) of the Seyfert galaxy NGC\\,4151 using data obtained with the Gemini Near-infrared Integral Field Spectrograph (NIFS). The data comprise spectra of the inner$~\\approx\\,200\\,\\times\\,500\\,$pc$^2$, at a spatial resolution of$~\\approx\\,8\\,$pc at the galaxy, covering the wavelength range$~0.95$--$2.51\\,\\mu$m at a spectral resolving power R$\\,\\geq 5000$. In a previous paper \\citep[][hereafter Paper\\,I]{sb09} we have used these data to map the NLR intensity distributions of $14$ emission lines, as well as their ratios. The main results were the distinct flux distributions and physical properties observed for the ionised, molecular and coronal gas. The ionised gas is co-spatial with the ionisation bi-cone observed in the optical \\foiii\\ emission at position angle (hereafter PA)$~60\\degr$ \\citep{evans93,hutchings98}, and seems to trace the outflow along the bi-cone \\citep{das05,crenshaw00a,hutchings99}. In the inner region, the NIR ionised-gas emission extends beyond the borders of the cone, which does not have a sharp apex as noted also by  \\citet{kra08}. The \\hmol\\ molecular gas intensity distribution, on the other hand, avoids the region of the bi-cone and seems to originate in the galaxy disc, while the coronal gas emission is barely resolved. In another recent work \\citep{riffel09b}, we have studied the unresolved nuclear continuum showing that its origin is emission by hot dust within $\\approx$\\,4\\,pc from the nucleus -- as expected from the dusty torus postulated by the unified model \\citep{antmi85}. The origin of this structure is probably a dusty wind that originates in the outer parts of the accretion disk. This wind is probably clumpy, as suggested by recent models \\citep{elitzur06}, and necessary in order to allow the escape of  radiation along the equatorial plane of the bi-cone in order to ionize the gas and produce the observed flux maps of Paper\\,I.  In the present paper we use the NIFS data to map the NLR kinematics of NGC\\,4151. Although many papers have been devoted to such a study, most of them are based on long-slit spectroscopy obtained in the optical with the Hubble Space Telescope \\citep[e.g.][]{winge97,hutchings98,crenshaw00a,das05}. The present study was performed in the near-IR, a waveband which is less affected by dust, and using integral field spectroscopy, which allows a full two-dimensional (hereafter 2D) coverage of the NLR kinematics. With our data and analysis we aim to examine previous claims of acceleration of the gas along the NLR, quantify the mass outflow rate and the corresponding feedback power, as well as investigate the origin of the NLR gas. As NGC\\,4151 harbors the closest bright AGN, its NLR  is one of the best suited for this type of study.  Our approach for the analysis of the NLR kinematics of NGC\\,4151 is as follows. First we map the centroid velocity and velocity dispersion of the ionised, molecular and coronal gas, obtained from fits to the emission-line profiles and compare our results with those of previous studies. Then we use an alternative approach to map the NLR kinematics, made possible by integral field spectrographs: a ``velocity tomography'' of the emitting gas, obtained by slicing the line profiles in velocity bins of$~60$\\kms, producing channel maps which provide a clearer view of the velocity distribution in the different gas phases. We have grouped the channel maps together in a sequence of velocity bins, generating movies which can be recovered at the authors' website. \\footnote{http://www.if.ufrgs.br/$\\sim$thaisa/ifu\\_movies/ngc4151}   We adopt in this paper the same distance to NGC\\,4151 as in Paper\\,I:$~13.3\\,$Mpc, corresponding to a scale at the galaxy of$~65\\,$pc\\,arcsec$^{-1}$ \\citep{mundell03}. The paper is organised as follows. In \\S\\,\\ref{data} we provide a quick description of the data, in \\S\\,\\ref{results} we present the centroid velocity and velocity dispersion measurements, in \\S\\,\\ref{tomography} we present the emission-line ``tomography'', in \\S\\,\\ref{discussion} we discuss our results, in \\S\\,\\ref{feed} we estimate the mass outflow rate along the NLR and compare with the mass accretion rate and in \\S\\,\\ref{conclusion} we present our conclusions. ",
        "conclusions": "\\label{conclusion}  We have measured the kinematics of the ionised and molecular gas surrounding the nucleus of  NGC\\,4151 using integral-field spectroscopy obtained with the NIFS instrument at the Gemini North telescope. The main results of this paper are:  \\begin{itemize}  \\item The ionised-gas kinematics are dominated by three velocity components: (1) extended emission at systemic velocity observed in a circular region around the nucleus;  (2) a component outflowing along the bi-cone (PA$=60\\deg$); (3) another component due to the interaction of the radio jet with the galactic disk.  \\item Regarding the extended emission at systemic velocity: The origin of this component is gas from the galaxy plane ionised by the AGN up to at least 1 arcsec (65\\,pc) from the nucleus, revealing that there is escape of ionizing radiation even along the equatorial plane of the bi-cone.  \\item Regarding the outflowing component: Channel maps extracted along the emission-line profiles show high-velocity gas all the way from the nucleus to the border of the emitting region. This seems to contradict  previous studies showing acceleration along the NLR. Our data suggest that the NLR clouds are accelerated very close to the nucleus (within $\\approx$ 10\\,pc), after which the flow moves at essentially constant velocity.  \\item We do not detect NIR counterparts to the highest-velocity \\foiii\\ emission seen by previous authors. It may originate in very low-density gas, which could be ablated material flowing off NLR clouds. This emission would not be characteristic of the bulk of the NLR mass.  \\item The origin of the bulk of NLR mass seems to be  entrained gas from the galaxy plane. This interpretation is supported by the large mass-outflow rate in ionised gas of $\\sim$\\,1\\,M$_\\odot$\\,yr$^{-1}$ obtained for the outflowing component in each cone, which is $\\sim$100 times the nuclear mass accretion rate.  \\item The estimated kinetic power of the outflow is $\\approx2.4\\times10^{41}$\\,erg\\,s$^{-1}$, which is $\\sim$\\,0.3\\% of L$_{bol}$, and is $\\approx$\\,10 times larger than values obtained previously by our group for other nearby Seyfert galaxies.  \\item Regarding the third kinematic component: Our data clearly show the interaction of the radio jet with ambient gas at systemic velocity. This has not been seen previously and indicates that the jet is launched close to the plane of the galaxy, pushing the gas encountered on its way.  \\item  Molecular hydrogen emission arises in extended regions along the axis of the galaxy's stellar bar, avoiding the region of the bi-cone. The \\hmol\\ velocities are close to systemic with a small rotation component, supporting an origin in the galaxy plane. The emitting molecular gas is probably just the ``hot skin'' of a larger reservoir that may have been built up by the previously observed HI inflow along the bar of the galaxy, and is probably the source of the AGN feeding."
    },
    "0911/0911.5574_arXiv.txt": {
        "abstract": "{ We report an indication  ($3.22 \\sigma$) of $\\approx 1860$ Hz quasi-periodic oscillations from a neutron star low-mass X-ray binary 4U 1636-536. If confirmed, this will be by far the highest frequency feature observed from an accreting neutron star system, and hence could be very useful to understand such systems. This plausible timing feature was observed simultaneously with lower ($\\approx 585$ Hz) and upper ($\\approx 904$ Hz) kilohertz quasi-periodic oscillations. The two kilohertz quasi-periodic oscillation frequencies had the ratio of $\\approx 1.5$, and the frequency of the alleged $\\approx 1860$ Hz feature was close to the triple and the double of these frequencies. This can be useful to constrain the models of all the three features. In particular, the $\\approx 1860$ Hz feature could be (1) from a new and heretofore unknown class of quasi-periodic oscillations, or (2) the first observed overtone of lower or upper kilohertz quasi-periodic oscillations. Finally we note that, although the relatively low significance of the $\\approx 1860$ Hz feature argues for caution, even a $3.22 \\sigma$ feature at such a uniquely high frequency should be interesting enough to spur a systematic search in the archival data, as well as to scientifically motivate sufficiently large timing instruments for the next generation X-ray missions. ",
        "introduction": "Quasi-periodic oscillations (QPOs) are observed from many neutron star low-mass X-ray binary (LMXB) systems \\citep{vanderKlis2006}. There are several kinds of QPOs, such as millihertz QPO, horizontal branch QPO, normal/flaring branch QPO, hectohertz QPO, lower kilohertz (kHz) QPO, upper kHz QPO, etc. \\citep{vanderKlis2006}. In this paper we will mostly concentrate on high frequency QPOs, that can currently be observed only with the proportional counter array (PCA) instrument of the {\\it Rossi X-ray Timing Explorer} ({\\it RXTE}) space mission. KHz QPOs are such high frequency QPOs, and often they appear as a pair. For a given source, their frequencies normally move up and down together in the $\\approx 200-1200$ Hz range in correlation with the source state. KHz QPOs are scientifically very important for the following reasons: (1) their high frequencies indicate that they originate from regions close to the neutron stars, and hence could be used as a tool to probe the strong gravity region around these stars, as well as the neutron star properties, and (2) they have been observed from many sources, and repeatedly from a given source, with high significance. Indeed, there are several models that involve relativistic orbital frequencies and the neutron star spin frequency, as well as beating and resonances among them \\citep{StellaVietri1998, Lamb+1985, Miller+1998, AbramowiczKluzniak2001, Torok+2008, Wijnands+2003, LambMiller2003, Zhang2004, Mukhopadhyay2009}. However, although there are many kHz QPO models available in the literature, none of them can explain all the major properties of this timing feature. Therefore, these QPOs cannot yet be used as a reliable tool to probe the strong gravity region, or to measure the neutron star parameters.  An intriguing aspect of the kHz QPOs is that so far they have never been observed with a frequency greater than 1330 Hz \\citep{vanderKlis2006}, while PCA can easily detect a feature with a much higher frequency. In fact, even this 1330 Hz QPO was not confirmed by a later analysis \\citep{Boutelieretal2009b}. This implies that no timing feature has ever been observed with $> 1330$ Hz from an accreting neutron star system. In this paper, we report the indication of an $\\approx 1860$ Hz QPO from the persistent neutron star LMXB 4U 1636-536 (see \\citet{Altamirano+08} for the other timing features). This plausible QPO was observed simultaneously with the pair of kHz QPOs, which suggests that the $\\approx 1860$ Hz QPO could be from a new and previously unknown class of very high frequency QPOs. Note that kHz QPOs from 4U 1636-536 were first discovered by \\citet{vanderKlisetal1996, Zhangetal1996a, Zhangetal1996b}, and later were reported by many authors \\citep{Wijnandsetal1997, Vaughanetal1997, Zhangetal1997, Vaughanetal1998, Psaltisetal1998, Mendezetal1998, Mendez1998, MisraShanthi2004, Kaaret+1999, Psaltisetal1999, Markwardtetal1999, Jonkeretal2000, Fordetal2000, Mendez2002, Jonkeretal2002, DiSalvoetal2003, MisraShanthi2004, Jonkeretal2005, Barretetal2007, Bellonietal2007, Torok+2008}. ",
        "conclusions": "\\label{Discussion}  In this paper, we report the detection of a pair of kHz QPOs from a neutron star LMXB 4U 1636-536. The ratio of these QPO frequencies is roughly 1.5, which is somewhat consistent with the resonance model described in \\citet{Abramowiczetal2003} (but see \\citet{Boutelieretal2009a}; see also \\citet{Zhangetal2006}). It is generally observed that the separation of the twin kHz QPO frequencies clusters around the neutron star spin frequency, or half of that. Models have been proposed in order to explain this aspect. However, for the kHz QPOs reported in this paper, the separation is $\\approx 319$ Hz, which is quite different from the half of the stellar spin frequency (see also \\citet{MendezBelloni2007, Yinetal2007}). Note that the neutron star of 4U 1636-536 spins with a rate of 582 Hz \\citep{StrohmayerMarkwardt2002}.  Apart from the  kHz QPOs, we have found an indication of a QPO at $\\approx 1860$ Hz. This is by far the highest frequency feature observed from an accreting neutron star system, and hence could be extremely useful to understand such systems. Moreover, this plausible QPO was observed simultaneously with the lower and the upper kHz QPOs, which means that the $\\approx 1860$ Hz QPO could be from a new class of QPOs.  We will now briefly discuss the plausible origin of this QPO using general arguments, and without going into specifics. In some of the models, the lower and upper kHz QPO and some other QPO frequencies are thought to be one or more of the following frequencies (for equatorial circular orbits in Kerr spacetime; \\citet{vanderKlis2006}): (1) Keplerian orbital frequency $\\nu_{\\phi} = \\nu_{K}(1+j(r_g/r)^{3/2})^{-1} = (2\\pi)^{-1}(GM/r^3)^{1/2}(1+j(r_g/r)^{3/2})^{-1}$; (2) radial epicyclic frequency $\\nu_r = \\nu_{\\phi}(1-6(r_g/r)+8j(r_g/r)^{3/2}-3j^2(r_g/r)^2)^{1/2}$; (3) vertical epicyclic frequency $\\nu_{\\theta} = \\nu_{\\phi}(1-4j(r_g/r)^{3/2}+3j^2(r_g/r)^2)^{1/2}$; (4) periastron precession frequency $\\nu_{\\rm peri} = \\nu_{\\phi} - \\nu_r$; and (5) nodal precession frequency $\\nu_{\\rm nodal} = \\nu_{\\phi} - \\nu_{\\theta}$. Here the last four frequencies are for infinitesimally tilted and eccentric orbits, $M$ is the neutron star mass, $r_g = GM/c^2$, $j = Jc/GM^2$, $J$ is the total angular momentum of the neutron star, and $r$ is the distance of an orbit from the centre of the neutron star. For a sample of radio pulsars, $M$ was found to be distributed around $1.35 M_\\odot$ in a narrow Gaussian ($\\sigma = 0.04 M_\\odot$; \\citet{ThorsettChakrabarty1999}). Therefore, considering $M = 1.35 M_\\odot$, we find that only $\\nu_{\\phi}$ or $\\nu_{\\theta}$ can be the frequency of the plausible $\\approx 1860$ Hz QPO, and that too for relatively large angular momentum parameter values (as far as a neutron star is concerned) and for orbits very close to the innermost stable circular orbit (ISCO; see Fig.~\\ref{freq1}). As the neutron star mass increases due to accretion, $M$ is expected to be greater than $1.35 M_\\odot$ for accreting neutron stars (such as the one in 4U 1636-536). Moreover, since $\\nu_{K} = (c^3/2\\pi G)(r/r_g)^{-3/2}(1/M)$, $\\nu_{K}$ decreases as $M$ increases for a given $r/r_g$. Therefore, it is unlikely that any of the above mentioned five frequencies can be the frequency of the plausible $\\approx 1860$ Hz QPO. Besides, since the spin frequency of the neutron star in 4U 1636-536 is $\\nu_{\\rm spin} = 582$ Hz, the beating of $\\nu_{\\rm spin}$ with any of $\\nu_{\\phi}$, $\\nu_r$ or $\\nu_{\\theta}$ cannot explain the frequency of this QPO. However, the $\\approx 1860$ Hz QPO could be an overtone of any of $\\nu_{\\phi}$, $\\nu_r$, $\\nu_{\\theta}$ or $\\nu_{\\rm peri}$, or beating between an overtone and a fundamental of these frequencies. This plausible QPO could also be an overtone of one of the kHz QPOs (see \\S~\\ref{DataAnalysisandResults}).  Although it is unlikely that the $\\approx 1860$ Hz QPO frequency is a Keplerian frequency $\\nu_{\\phi}$, it is still instructive to examine what constraints $\\nu_{\\phi} = 1860$ Hz can impose on the neutron star parameter values. This is because, the measurement of neutron star parameters provides the only way to constrain the theoretically proposed equation of state (EoS) models of neutron star cores, and hence to understand the nature of supranuclear core matter \\citep{LattimerPrakash2007, Bhattacharyyaetal2000, Bhattacharyyaetal2001a}. Various authors suggested ways to constrain the neutron star parameters assuming one of the kHz QPO frequencies as a Keplerian frequency \\citep{Kaaretetal1997, Miller+1998, Zhangetal2007, Zhang2009}. For example, this assumption, and the following two reasonable conditions can constrain the mass ($M$) and the radius ($R$) of a neutron star \\citep{Miller+1998}: \\begin{eqnarray} R \\le r, \\label{freq4} \\end{eqnarray} where $r$ is the radius of the orbit associated with the kHz QPO via the expression of $\\nu_{\\phi}$; and \\begin{eqnarray} r_{\\rm ISCO} \\le r, \\label{freq5} \\end{eqnarray} where $r_{\\rm ISCO}$ is the radius of the ISCO. This is because the first condition gives a mass-dependent upper limit on $R$ via the expression of $\\nu_{\\phi}$; and the second condition gives an upper limit on $M$: $M < c^3/(2\\pi 6^{3/2}G\\nu_\\phi|_r)$ (for Schwarzschild spacetime). Therefore, if 1860 Hz were a Keplerian frequency, then the upper limit of neutron star mass and radius would be $\\approx 1.2 M_\\odot$ and $\\approx 10.5$ km respectively. Such upper limits would support the strange star EoS models \\citep{Chengetal1998, Bhattacharyyaetal2001b}, although the less exotic neutron star EoS models could not be completely ruled out.  \\begin{figure}[h] \\centering \\includegraphics[width=9.0cm, angle=0]{fig3.ps} \\begin{minipage}[]{85mm} \\caption{Radial profiles of various frequencies (colour coded) of equatorial circular orbits in Kerr spacetime (see \\S~\\ref{Discussion}). Two angular momentum parameters ($j = 0.0$ (solid) and $j = 0.3$ (dotted)), and neutron star mass $= 1.35 M_\\odot$ are used. Dashed horizontal line exhibits the 1860 Hz frequency. This figure shows that these frequencies, that are often used to explain kHz and some other QPOs, cannot possibly explain a $\\approx 1860$ Hz QPO. }\\end{minipage} \\label{freq1} \\end{figure}  The uniquely high frequency of the plausible $\\approx 1860$ Hz QPO and its coexistence with both the kHz QPOs make it extremely interesting. Furthermore, although no such high frequency signal was claimed previously, \\citet{vanderKlis2006} mentioned, ``the distinct impression of the observers is that there is still much hiding below the formal detection levels\". It should, therefore, be worthwhile to report this feature, which will spur a systematic search of such high frequency features in the archival {\\it RXTE} PCA data, as well as in the large area xenon proportional counters (LAXPC) instrument data of the upcoming {\\it Astrosat} space mission. Moreover, next generation X-ray timing instruments, such as the high timing resolution spectrometer (HTRS; proposed for the {\\it International X-ray Observatory}) and the proposed Si pixel detector of the {\\it Advanced X-ray Timing Array}, will have much better capability to detect weak QPOs (see, for example, Fig.~2 of \\citet{Barret+08}; \\citet{Chakrabarty+08}). The plausible $\\approx 1860$ Hz QPO will therefore motivate the future X-ray timing instruments."
    },
    "0911/0911.2448_arXiv.txt": {
        "abstract": "We compare the relative merits of AGN selection at X-ray and mid-infrared wavelengths using data from moderately deep fields observed by both {\\em Chandra} and {\\em Spitzer}. The X-ray-selected AGN sample and associated photometric and spectroscopic optical follow-up are drawn from a subset of fields studied as part of the Serendipitous Extragalactic X-ray Source Identification (SEXSI) program.   Mid-infrared data in these fields are derived from targeted and archival {\\em Spitzer} imaging, and mid-infrared AGN selection is accomplished primarily through application of the IRAC color-color AGN `wedge' selection technique.  Nearly all X-ray sources in these fields which exhibit clear spectroscopic signatures of AGN activity have mid-infrared colors consistent with IRAC AGN selection.  These are predominantly the most luminous X-ray sources.  X-ray sources that lack high-ionization and/or broad lines in their optical spectra are far less likely to be selected as AGN by mid-infrared color selection techniques.  The fraction of X-ray sources identified as AGN in the mid-infrared increases monotonically as the X-ray luminosity increases.  Conversely, only 22\\%\\ of mid-infrared-selected AGN are detected at X-ray energies in the moderately deep ($\\langle t_{\\rm exp} \\rangle \\approx 100$~ks) SEXSI {\\em Chandra} data.  We hypothesize that IRAC sources with AGN colors that lack X-ray detections are predominantly high-luminosity AGN that are obscured and/or lie at high redshift.  A stacking analysis of X-ray-undetected sources shows that objects in the mid-infrared AGN selection wedge have average X-ray fluxes in the $2 - 8$~keV band three times higher than sources that fall outside the wedge. Their X-ray spectra are also harder.  The hardness ratio of the wedge-selected stack is consistent with moderate intrinsic obscuration, but is not suggestive of a highly obscured, Compton-thick source population. It is evident from this comparative study that in order to create a complete, unbiased census of supermassive black hole growth and evolution,  a combination of sensitive infrared, X-ray and hard X-ray selection is required. We conclude by discussing what samples will be provided by upcoming survey missions such as {\\em WISE}, {\\em eROSITA}, and {\\em NuSTAR}. ",
        "introduction": "\\label{sexsiconclusion_sec:intro}  The tight correlation of nuclear black hole mass with the velocity dispersion and mass of the galactic bulge \\citep[e.g.,][]{Magorrian:98,Ferrarese:00,Tremaine:02} implies that the growth and evolution of galaxies is closely linked to the growth and evolution of the supermassive black holes which reside in (at least) all massive galaxies \\citep[e.g.,][]{Kauffmann:00,Heckman:08}. Indeed, recent theoretical work suggests that feedback from active galactic nuclei (AGN) plays a dominant role in establishing the present-day appearances of galaxies, providing a natural, physical explanation for both cosmic downsizing and the possibly related bimodality in local galaxy properties \\citep[e.g.,][]{Scannapieco:05, Hopkins:08, Cattaneo:09}. However, obtaining an unbiased census of black holes in the universe remains challenging, hampering our ability to fully probe this connection.  Most surveys for active galaxies are severely biased towards unobscured ({\\em type~1}) AGN since nuclear emission in such sources dominates over host galaxy light at most wavelengths, making type~1 AGN both more readily identifiable and easier to follow up spectroscopically.  However, both unified AGN models \\citep[e.g.,][]{Antonucci:93,Urry:95} and the shape of the X-ray background suggest a population of obscured ({\\em type~2}) sources which outnumber the type~1 AGN by up to a factor of ten \\citep[e.g.,][]{Comastri:95, Treister:04, Gilli:07}. Determining the ratio of unobscured to obscured AGN as a function of luminosity and redshift can directly constrain the growth history of supermassive black holes, and help to quantify their influence on the evolution of their host galaxies \\citep[e.g.,][]{Ueda:03,Barger:05,Hasinger:05,Hopkins:07}. Since different search techniques for obscured AGN suffer from different biases, the problem is best addressed through multiwavelength studies.  The hard (\\hardrange) X-ray and mid-infrared wavebands provide powerful, complementary methods for identifying and studying AGN over a wide range of intrinsic obscuration. Radiation seen from active galaxies at X-ray energies is primarily due to direct emission from accretion processes near the central supermassive black hole, with higher energy photons less susceptible to absorption, while mid-infrared radiation (rest-frame $\\lambda_0 \\simgt 2~\\mu$m) is typically dominated by emission from dusty obscuring material surrounding the AGN central engine, and is likewise relatively immune to absorption \\citep[e.g.,][]{Krolik:99}. Hard X-ray and mid-infrared surveys of AGN will therefore sample a wide range of intrinsic absorption and, in this regard, will be both less biased and more complete than surveys in the optical \\citep[e.g.,][]{Richards:06} and soft X-ray \\citep[$E \\simlt 2.4$~keV; e.g.,][]{Hasinger:98, Schmidt:98} bands, both of which suffer from strong attenuation by even moderate (\\nh$\\simgt 10^{21}$~cm$^{-2}$) columns of dusty, obscuring material.  However, AGN samples selected in the \\hardrange\\ X-ray and infrared bands will also suffer incompleteness. Examples exist of AGN identified from optical spectroscopy which remain undetected even in the deepest X-ray images yet obtained \\citep[e.g.,][]{Steidel:02}; similarly, there are examples of optically selected AGN not identified in infrared surveys (e.g., \\citealt{Stern:05b}). Conversely, X-ray missions have identified AGN whose optical spectra are devoid of AGN signatures (e.g., X-ray bright, optically normal galaxies, or ``XBONGs''; \\citealt{Comastri:02}).   Therefore a complete census will necessarily require combining selection techniques from several spectral regimes.  The current generation of X-ray telescopes has dramatically advanced our knowledge of the AGN population (see~\\citealt{Brandt:05}\\ for a review). The {\\em Chandra X-ray Observatory} \\citep{Weisskopf:96} provides a large collecting area, a moderate field of view (FOV), and exquisite angular resolution ($<1$\\arcsec) from 0.5 to 8~keV.  These capabilities have allowed {\\em Chandra} extragalactic surveys to efficiently select and optically identify large samples of AGN in the 2 -- 8~keV range. Observing in this hard X-ray energy band means that even sources shrouded by considerable obscuring column densities (up to $10^{24}$ cm$^{-2}$) can be detected at low redshift, and sources with even higher column densities can be detected out to $z \\simgt 2$.  Previous sensitive X-ray telescopes were primarily restricted to energies below $\\sim 2$~keV and thus missed many of the obscured AGN, a source population that has long been theorized to explain the mismatch in spectral shape between the \\hardrange\\ X-ray background ($\\Gamma \\approx 1.4$) and the unobscured active galaxies which dominate the source counts at soft X-ray energies ($\\Gamma\\approx 1.9$; \\citealt{Nandra:94}).  The 2003 launch of the {\\em Spitzer Space Telescope} \\citep{Werner:04} opened a new era in mid-infrared observations, providing orders of magnitude improvement in sensitivity in the $3.6-160~\\mu$m band.  The increased sensitivity, combined with the large FOV, allows for the first time efficient long-wavelength survey capabilities.  The primary imaging cameras on {\\it Spitzer} are the Infrared Array Camera \\citep[IRAC;][]{Fazio:04}, providing simultaneous imaging at 3.6, 4.5, 5.8, and 8~$\\mu$m, and the Multiband Imaging Photometer for {\\em Spitzer} \\citep[MIPS;][]{Rieke:04}, providing simultaneous imaging at 24, 70, and 160~$\\mu$m.  Several methods have been developed to select AGN based on their {\\em Spitzer} colors~\\citep{Lacy:04,Stern:05b,Alonso-Herrero:06}. These methods exploit the difference in the typical spectral energy distribution (SED) of AGN compared to `normal' galaxies. The near-infrared emission of both typical galaxies and vigorously star-forming galaxies is primarily produced by a thermal stellar population, resulting in SEDs peaked near the rest-frame 1.6~$\\mu$m ``bump.'' This bump, which is caused by the minimum in the opacity of the H$^{-}$ ion near 1.6~$\\mu$m \\citep[e.g.,][]{John:88}, is a feature of almost all stellar populations \\citep[e.g.,][]{Wright:94,Simpson:99}, whereas AGN-dominated SEDs have a non-thermal, roughly power-law shape (for $\\lambda \\simlt 10~\\mu$m). At longer wavelengths contributions from stellar blackbody emission is low, while radiation from the obscuring dust near the central engine provides strong, relatively isotropic emission which is quite distinct from stellar light (although the shape of the AGN SED does somewhat depend on viewing angle, e.g., \\citealt[][]{Elitzur:08}). \\citet{Stern:05b}, for example, suggested a selection technique exploiting the fact that, for AGN, the long-wavelength side of the 1.6~$\\mu$m stellar peak does not decline; the technique uses an empirically determined `wedge' in IRAC color-color space ([3.6]$-$[4.5] versus [5.8]$-$[8.0]) that preferentially contains AGN as compared to normal galaxies or Galactic stars.  Combining a sample of 10,000 $R<21.5$ spectroscopically identified sources from the AGN and Galaxy Evolution Survey (AGES; Kochanek et al., in prep.) and mid-infrared observations from the IRAC Shallow Survey \\citep{Eisenhardt:04}, \\citet{Stern:05b} defined mid-infrared AGN selection criteria which robustly identify broad- and narrow-lined AGN, with only 18\\% sample contamination from galaxies (17\\%) and stars (1\\%).  The true sample contamination is likely lower, since many of the spectroscopically normal galaxies may harbor active nuclei (i.e., are mid-infrared versions of XBONGs).  Working from the full spectroscopically defined AGES sample, \\citet{Stern:05b} found that the wedge selects 91\\% of the broad-lined AGN, 40\\% of the narrow-lined AGN, and fewer than 3\\% of the normal galaxies.  This paper presents an exploration of the relative strengths of {\\em Chandra} and {\\em Spitzer} as black-hole finders, using a subset of six fields from the Serendipitous Extragalactic X-ray Source Identification program (SEXSI; \\citealt{Harrison:03}, \\citealt{Eckart:05}, and \\citealt{Eckart:06}) for which we have obtained mid-infrared coverage with {\\em Spitzer}. In addition to examining the properties of $\\sim 250$ hard X-ray-selected AGN, we extend our sensitivity to X-ray emission using a stacking analysis on sources with different mid-infrared characteristics. This enables us to make a comprehensive comparison of the relative characteristics of AGN samples selected from their X-ray and mid-infrared properties.   We organize the paper as follows: \\S\\ref{sexsiconclusion_sec:experiment} introduces the SEXSI program and presents the complementary mid-infrared observations and the resulting source catalog; \\S\\ref{sexsiconclusion_sec:xrayselected} discusses the mid-infrared properties of the \\hardrange\\ SEXSI sources; \\S\\ref{sexsiconclusion_sec:stacking} presents the mean X-ray properties of X-ray-non-detected {\\em Spitzer} sources; \\S\\ref{sexsiconclusion_sec:discussion} provides a discussion; and \\S\\ref{sexsiconclusion_sec:conclusions} summarizes our conclusions. Luminosity calculations assume a standard cosmology used in our previous work \\citep{Eckart:06}, $\\Omega_0 = 0.3$, $\\Lambda = 0.7$, and $H_0 = 65~ {\\rm km}~ {\\rm s}^{-1}~ {\\rm Mpc}^{-1}$. ",
        "conclusions": "\\label{sexsiconclusion_sec:conclusions}  We have presented an exploration of the relative efficiency of selecting AGN in the X-ray with {\\em Chandra} and in the mid-infrared with {\\em Spitzer}.  In this section we summarize our comparison of the selection techniques and make predictions about the AGN samples that will be selected with future X-ray and mid-infrared surveys.  \\subsection{Summary} \\label{sexsiconclusion_sec:midIRsummary}   Two-thirds of the X-ray-selected AGN are also selected by the \\citet{Stern:05b}\\ {\\em Spitzer} AGN selection criteria. Nearly all of the SEXSI X-ray sources that are classified as either broad-line or narrow-line AGN (i.e., those that exhibit high-ionization optical/UV emission lines in their optical spectra) are consistent with wedge selection within photometric uncertainties.   A large fraction of the low luminosity ($L_{\\rm 2-10~keV} < 10^{43.5} \\lumin$), hard X-ray (\\hardrange) AGN identified by {\\em Chandra} will be missed by {\\em Spitzer} selection techniques. These low X-ray luminosity sources tend to lack the high-ionization emission lines that allow AGN identification via optical spectroscopy; thus, the sources that are missed by the {\\em Spitzer} color-color AGN selection will also remain unidentified in optical spectroscopic surveys. These sources tend to lie at lower redshift and have {\\em higher} fluxes in the infrared, on average. This suggests that many of these non-wedge-selected X-ray sources have mid-infrared fluxes that are dominated by stellar (starburst) emission as opposed to AGN activity; indeed, most have optical spectra of ELGs, consistent with the conclusion that emission related to star formation outshines the emission from the active nucleus in some parts of the spectrum (i.e., in the mid-infrared and optical), while the X-ray luminosities of these sources are dominated by X-ray emission from the active nuclei. In addition, we find that while a fair fraction ($\\sim 63$\\%) of the optically bright ($R<21$) wedge-selected {\\em Spitzer} AGN are also identified by {\\em Chandra}, over 80 percent of the optically faint wedge sources are not detected in the X-ray.  It is likely that these fainter sources are AGN that are heavily obscured and/or lie at high redshift. We also find that X-ray-selected high-redshift, type-2 quasars are also selected via the $3.6~\\mu$m$-24~\\mu$m selection criteria proposed by \\citet{Martinez-Sansigre:05}; the radio selection criteria used by those authors to eliminate starburst contaminants is not necessary when the sources have X-ray properties indicative of AGN activity.   To explore the mean X-ray properties of the X-ray undetected, IRAC-selected AGN candidates, we have stacked the  {\\em Chandra} data extracted from the positions of IRAC sources. The stacked wedge sources show significant X-ray signals in the full, soft, and hard X-ray bands. The average hardness ratio of the stacked spectrum of sources inside the wedge is higher than that of sources outside of the wedge, and the average $2-8$~keV flux is three times larger.  The hardness ratio of the wedge-selected stacked spectrum is consistent with moderate intrinsic obscuration (\\nh$\\sim 10^{22}-10^{23.7}$~cm$^{-2}$), but is not suggestive of a highly obscured, Compton-thick source population.   \\subsection{Conclusions} \\label{sexsiconclusion_sec:finalconclusions}  The results of our study illustrate the challenges of identifying obscured and low-luminosity AGN  using any one technique.  A complete understanding of AGN activity and black hole growth will require multiwavelength data sets: low-luminosity AGN in galaxies with active star formation cannot be selected using mid-infrared color techniques, but are effectively found at low redshift with the current-generation of X-ray telescopes;  obscured high-redshift objects are beyond the sensitivity threshold of {\\em Chandra} and {\\em XMM}, but they are effectively identified in the mid-infrared if the AGN luminosity is at the high end of the luminosity distribution.  The results from current X-ray and mid-infrared satellites enable us to estimate the progress possible with upcoming space missions that plan to conduct all-sky or pencil beam surveys. The {\\em Wide-field Infrared Survey Explorer} ({\\em WISE;} \\citealt{Mainzer:05}), scheduled for launch at the end of 2009, will provide full-sky mid-infrared imaging reaching a depth of 120~$\\mu$Jy at 3.3~$\\mu$m (similar to IRAC channel~1) and 160~$\\mu$Jy at 4.7~$\\mu$m (similar to IRAC channel~2). At these shallow depths, almost every source redder than [3.3]$-$[4.7] $\\approx 0.5$ will be an AGN; high-redshift galaxies are too faint to make this flux cut, and the coolest brown dwarfs are too rare. The primary contaminant will be actively star-forming galaxies at redshifts of a few tenths, which should be readily identifiable from their optical morphologies and colors in relatively shallow imaging such as is available from the Sloan Digital Sky Survey. By applying a flux cut at the {\\em WISE} depths to the Bo\\\"{o}tes IRAC survey \\citep{Eisenhardt:04,Ashby:09}, we predict that {\\em WISE} should detect approximately 85 AGN candidates per square degree, of which $\\approx 15\\%$ will be low-redshift star-forming galaxy contaminants. {\\em WISE} will therefore provide an unprecedented sample of high-luminosity obscured AGN across a range of redshifts up to  $z \\sim 3$ identified based on mid-infrared color selection alone.    {\\em WISE} will also detect numerous other low-redshift ($z < 1$), low-luminosity ($L_{\\rm 2-10~keV} \\simlt 10^{43} \\lumin$) active galaxies, but the significant number of open symbols above the {\\em WISE} threshold shown in Figure~\\ref{sexsiconclusion_fig:f36vsfx} shows that mid-infrared colors alone are insufficient to identify them as AGN.   The {\\em extended R\\\"ontgen Survey with an Imaging Telescope Array} ({\\em eROSITA;} \\citealt{Predehl:06}) X-ray mission, planned for launch in 2011,  will provide full-sky hard X-ray images to a depth of $f_{\\rm 2-10~keV} = 1.5 \\times 10^{-13}~\\fluxu$. Using the sample of XBo\\\"{o}tes sources \\citep{Kenter:05} with fluxes above the {\\em eROSITA} hard-band sensitivity limit, we predict that {\\em eROSITA} should detect $\\approx 4$ hard X-ray sources per square degree. The mission will detect many more AGN in the soft X-ray band because the effective area is much larger at energies below 2~keV, reaching soft-band depths of $f_{\\rm 0.5-2~keV} = 9 \\times 10^{-15}~\\fluxu$.  Scaling from the XBo\\\"{o}tes \\softrange\\ sample we predict that {\\em eROSITA} will detect $\\approx 120$ sources per square degree in the soft band. Figure~\\ref{sexsiconclusion_fig:f36vsfx} shows that {\\em eROSITA} will identify a large population of AGN below the {\\em WISE} threshold, most of which we project will be at $z \\simgt 1$. The majority of these sources will be unobscured BLAGN, but a significant population will be narrow-line, obscured AGN (NLAGN and ELG). However, {\\em eROSITA} does not reach sufficient depth to detect many of the low-luminosity, low-redshift AGN detected, but not identified as such, by {\\em WISE}.   Clearly {\\em WISE} and {\\em eROSITA} combined will miss (or mis-identify) low-luminosity obscured AGN at all redshifts, most of which will also not be found in optical and soft X-ray surveys. At low redshift only a sensitive X-ray mission extending to E$_x > 10$~keV will effectively uncover this population, which will be visible only in reprocessed light, if at all, at E$_x < 10$~keV. The  {\\em Nuclear Spectroscopic Telescope Array} ({\\em NuSTAR;} \\citealt{Harrison:05})\\footnote{\\tt http://www.nustar.caltech.edu.}, scheduled for launch in 2011, will make major advances in studying faint extragalactic sources at energies above 10~keV. {\\em NuSTAR} will be the first focusing, high-energy X-ray satellite and is designed to make targeted observations of the hard X-ray sky from $6-79$~keV. {\\em NuSTAR}, which has a field-of-view comparable to {\\em Chandra}, plans to cover  the Bo\\\"{o}tes and Great Observatories Origins Deep Survey (GOODS) fields, and should resolve $\\simgt 50$\\% of the 30~keV X-ray background, reaching $10-40$~keV flux limits of $< 10^{-14}~\\fluxu$. Many of the resolved sources are expected to be low-luminosity, nearby AGN not identified as such in current mid-infrared or \\hardrange\\ X-ray surveys, and some {\\em NuSTAR} sources may be AGN with hard spectra owing to low-efficiency accretion. For studying the high-redshift, low-luminosity population, the {\\em International X-ray Observatory} ({\\em IXO}) will reach \\hardrange\\ flux limits of $\\sim 5 \\times 10^{-17}~\\fluxu$ in a megasecond exposure with its  wide-field imager. It could survey up to 10~deg$^2$ to a depth of $10^{-16}~\\fluxu$ and 0.25~deg$^2$ to a depth of $10^{-17}~\\fluxu$, uncovering the largest population of high-redshift AGN yet detected.  Predictions for the AGN populations that will be uncovered by these future missions must rely on careful studies of the AGN identified by existing X-ray and infrared surveys. As illustrated in this paper, an understanding of the AGN selection techniques of the various missions, and their relative strengths and weaknesses, is crucial as datasets are synthesized to obtain an understanding of black hole growth and evolution in the universe.  Finally, we attempt to put the results of our study in the context of understanding the cosmic history of black hole growth and AGN accretion modes. The primary conclusions of this paper are that heavily obscured luminous AGN are often missed by X-ray selection, while low-luminosity AGN are often missed by mid-infrared selection.  For the high-luminosity sources, various work has shown that, at least for the optically-selected unobscured sources, most are emitting at close to their Eddington limit \\citep[e.g.,][]{Kollmeier:06}.  Multiwavelength observations are consistent with the heavily-obscured high-luminosity sources having similar intrinsic SEDs, simply with their UV, optical, and X-ray emission suffering from absorption.  As for the low-luminosity AGN, there are two possibilities:  they could be lower mass black holes again emitting at close to their Eddington limits, or they could be average-sized black holes, with masses similar to the high-luminosity sources, but simply emitting at lower Eddington ratios.  Using a new set of empirical, low-resolution SED templates for AGN and galaxies, \\citet{Assef:09} shows that the likelihood of a source to be selected as an AGN using the \\citet{Stern:05b} mid-infrared color criteria is primarily a function of how strong the AGN is relative to the host galaxy --- e.g., the Eddington ratio assuming that, to first order, the host galaxy mid-infrared emission scales with the galaxy total stellar mass which scales with the nuclear supermassive black hole mass.  Thus, the fact that we find few low-luminosity sources using the mid-infrared criteria suggests that the low-luminosity sources are {\\em not} emitting at close to their Eddington ratio.  This is consistent with the results of \\citet{Babic:07}\\ who find that low-luminosity X-ray sources in the {\\em Chandra} Deep Field-South have a wide range of Eddington ratios, $10^{-5} \\simlt \\log(L_{\\rm bol}/L_{\\rm Edd}) \\simlt 1$.  One caveat on this interpretation is that it assumes that the AGN SED is independent of both Eddington ratio and luminosity --- the former of which has recently been questioned by \\citet{Vasudevan:07}.  Future work, in particular, higher energy observations with {\\em NuSTAR}, will quantify the extent to which the intrinsic AGN SEDs depend on both luminosity and Eddington ratio."
    },
    "0911/0911.1161_arXiv.txt": {
        "abstract": "We analyze a solar active region observed by the {\\it Hinode} \\ion{Ca}{2}~H line using the time-distance helioseismology technique, and infer wave-speed perturbation structures and flow fields beneath the active region with a high spatial resolution. The general subsurface wave-speed structure is similar to the previous results obtained from {\\it SOHO}/MDI observations. The general subsurface flow structure is also similar, and the downward flows beneath the sunspot and the mass circulations around the sunspot are clearly resolved. Below the sunspot, some organized divergent flow cells are observed, and these structures may indicate the existence of mesoscale convective motions. Near the light bridge inside the sunspot, hotter plasma is found beneath, and flows divergent from this area are observed. The {\\it Hinode} data also allow us to investigate potential uncertainties caused by the use of phase-speed filter for short travel distances. Comparing the measurements with and without the phase-speed filtering, we find out that inside the sunspot, mean acoustic travel times are in basic agreement, but the values are underestimated by a factor of $20-40\\%$ inside the sunspot umbra for measurements with the filtering. The initial acoustic tomography results from {\\it Hinode} show a great potential of using high-resolution observations for probing the internal structure and dynamics of sunspots. ",
        "introduction": "Deriving subsurface structures and flow fields of solar active regions is one of the major topics for local helioseismology studies. Using \\ion{Ca}{2}~K observations made at the geographic South Pole in early 1990s, \\citet{duv96} found the first evidence of downdrafts below active regions through measuring acoustic travel times by employing the time-distance helioseismology technique. Later inversions \\citep{kos96} using these travel time measurements found strong converging downflows and areas of the excess of sound-speed beneath growing active regions. With the availability of {\\it Solar and Heliospheric Observatory} / Michelson Doppler Imager \\citep[{\\it SOHO}/MDI;][]{sch95} Doppler observations that are seeing-free and more suitable for local helioseismology studies, more investigations of the sunspot's subsurface structures and flow fields have been carried out by use of various theoretical models: ray-path approximation \\citep[e.g.,][]{kos00, zha01}, Fresnel-zone approximation \\citep[e.g.,][]{jen01, cou04}, and Born approximation \\citep[e.g.,][]{cou06}. Despite the use of different approaches of calculating travel time sensitivity kernels, the results remain largely the same, i.e., for active region subsurface structures, a negative sound-speed perturbation was found up to 5 Mm below the photosphere, and a positive perturbation was seen below that depth. For subsurface flow fields, converging downflows were inferred from the photosphere to about 5 Mm beneath it, and divergent flows were found below this layer. Another local helioseismology technique, ring-diagram analysis, although could not provide subsurface flow maps with such a high spatial resolution, gave results that were in general agreement \\citep{hab04, kom05}. Some direct comparisons of these two techniques were also performed \\citep{hin04}.  However, on the other hand, these results were not uncontroversial. By measuring acoustic phase shifts in sunspot penumbra using MDI Doppler observations when sunspots were located near the solar disk limb, \\citet{sch05} demonstrated that the measured acoustic travel times varied around sunspot penumbra when the sunspot was located near the limb, and believed this effect was caused by the inclined magnetic field. While acknowledging the existence of such an effect in oscillation signals observed in Dopplergrams, \\citet{zha06} found out that such an effect did not exist in the oscillations observed in MDI continuum intensity. Additionally, \\citet{raj06} showed that the phase-speed filtering in the time-distance acoustic travel time measurements would also introduce errors in active regions where oscillation amplitudes were suppressed. Using numerical simulations \\citet{par08} confirmed this effect, and concluded that this effect led to underestimation of the mean travel times at short distances. Some recent attempts of using different filtering procedures, e.g., ridge filtering \\citep{bra08}, showed that different filtering might give different measurements. Thus, it is certainly worthwhile examining how filters would change measurement results. On the other hand, it is also arguable whether those time-distance observed acoustic travel time variations in solar magnetic areas could be interpreted as interior sound speed anomalies, or were caused by some surface magnetism effects such as showerglass effect \\citep{lin05a, lin05b}. Furthermore, it is also not clear how the interaction of acoustic waves and magnetic field would effect the phase of these waves, hence the travel time variations in the measured acoustic travel times. It would be very difficult to interpret the measurements if acoustic waves experience some phase shifts at the boundary of unmagnetized and magnetized areas \\citep{cal09}.  High spatial resolution observation of sunspots by the Solar Optical Telescope \\citep[SOT;][]{tsu08} onboard the Japanese solar spacecraft {\\it Hinode} \\citep{kos07} not only gives us another possibility of investigating subsurface properties of solar active regions in addition to the already existing helioseismology instruments, but also provides an unprecedented high spatial resolution that has enabled some studies that were improbable using these existing instruments. In this paper, we analyze an active region observed by {\\it Hinode} utilizing the time-distance helioseismology technique, and infer the wave-speed profiles and flow fields beneath this region, and present these results in \\S3. We also analyze the acoustic signals without applying any filters except filtering out solar convection and $f$-modes, and investigate how the phase-speed filtering effects our measurements. Such analyses can hardly be carried out by use of MDI high resolution observations, and these results are presented in \\S4. In \\S5, we discuss and summarize our results. ",
        "conclusions": "\\subsection{Subsurface Structure and Flow Field} By a time-distance analysis of unprecedented high spatial resolution observations from the {\\it Hinode}/SOT, we have investigated high resolution wave-speed structures and mass flows beneath active region AR10953. For the subsurface wave-speed structure, the inverted results are remarkably similar to previous results based on MDI Dopplergrams with different inversion technique as well as different sensitivity kernels. For subsurface flow fields, the general picture is similar to what has been obtained from MDI observations, i.e., converging downward flows near the surface below the sunspot area. Despite the similarities, the current picture also has some differences compared with the earlier one. One difference is that the downward flows beneath the sunspot are more prominent in the flow fields inferred from this study. Another remarkable result from this study is the mass circulation outside the sunspot, which seems to keep mass conservative and is much more clear than the previous MDI inversions. It is widely believed in theory and in numerical simulations that such downdraft and converging flows near the sunspot surface play important roles in keeping sunspots stable \\citep{par79, hur00, hur08}.  However, it is still not quite clear how the overall flow structures around the sunspot's surface and interior look like. It is already well known the existence of penumbral Evershed outflows, outgoing moat flows beyond sunspot's penumbra, and inflows from the inner penumbra and umbra measured by tracking penumbra/umbra dots \\citep[e.g.,][] {sob09}. For the sunspot's interior, there were reports of outflows from f-mode analysis \\citep{giz00}, inflows at the depth of 1.5 - 5 Mm (Zhao et al. 2001), and large scale inflows around active regions up to 10 Mm in depth from ring-diagram analysis \\citep{hab04}. Based on all these observations, \\citet{giz03} proposed a schematic flow structure of the sunspot, and \\citet{hin09} recently proposed a similar one. Both flow structures show outflows near the surface, and inflows below that, with transition depth of these opposite flows undetermined. Here, based on all previous results, as well as on the recent numerical simulation of Evershed flows \\citep{kit09}, we also propose a schematic flow structure of a sunspot (Figure~\\ref{flow_profile}), slightly different from the plots of two previous studies, but essentially similar. However, it is also acknowledged that this flow structure is not consistent with what was recently found by \\citet{giz09} and recent numerical simulations by \\citet{rem09}. Thus, further observational and theoretical studies are required to determine the subsurface dynamics of sunspots.  The high resolution subsurface wave-speed maps (Figure~\\ref{cs_horiz}) and flow fields (Figure~\\ref{hr_flows}), present us an unprecedented opportunity to see the sunspot's subsurface structures and dynamics with many details. These results reveal the complexity of the subsurface perturbations and dynamics and their relationship with the sunspot structure. The subsurface image at the depth of $0 - 1.5$ Mm displays a larger wave-speed perturbation, presumably hotter in temperature, near the light bridge area than other areas inside the sunspot, where the wave-speed perturbations are largely negative. This is in agreement with the surface observations that light bridges are formed by upwelling of hot plasma \\citep{kat07}. We also see clear divergent flows from the light bridge in high resolution subsurface flow map at the same depth (Figure~\\ref{hr_flows}). This is consistent with the plasma upwelling that occurs in this area, but not in good agreement with the proper motion tracking results of this region \\citep{lou08}.  It is also interesting to see quite a few divergent flow cells inside both the sunspot umbra and penumbra. The cells are bigger than granules and smaller than supergranules, and indicate the existence of a mesoscale convection inside the sunspot. The discovery of these motions is intriguing because they may be a signature of the downdraft vortex rings around magnetic flux bundles, as suggested by \\citet{par92}. The Parker's conjecture was that the observed clustering of magnetic flux bundles in sunspot is at least in part a consequence of hydrodynamic attraction of the downdraft vortex rings. Thus, it is very important to continue the investigations of the subsurface dynamics of sunspots, both observationally by high resolution helioseismology and theoretically by realistic MHD simulations.  \\subsection{Analysis without Phase-Speed Filtering} As introduced in \\S1, the phase-speed filtering may introduce some errors in active regions due to the smaller oscillation amplitude in these areas \\citep{raj06}, and using ridge filtering may give different results as using phase-speed filtering in active regions \\citep{bra08}. All of these prompt us a reexamination of the use of phase-speed filtering and an evaluation of how phase-speed filtering alter measured acoustic travel times, in particular, inside active regions.  The {\\it Hinode} high resolution observation has helped us to overcome an obstacle that acoustic travel times could hardly be measured in short distances inside active regions if phase-speed filtering was not applied. Our measured distance-dependent acoustic travel time differences relative to the solar quiet region obtained without using the phase-speed filtering, as shown in Figure~\\ref{td_diff}, have evidently demonstrated that for short distances, the acoustic travel time is longer inside both sunspot umbra and penumbra than the quiet region, and for distances longer than $\\sim 13$ Mm, the mean travel time is shorter. This is generally in agreement with the measurements using phase-speed filters, except that the ones with filters would underestimate the values by $20-40\\%$ in sunspot umbra. It is believed that the underestimation is caused by a combination of oscillation amplitude suppression inside active regions and the use of phase-speed filtering. Therefore, we believe that the phase-speed filtering does not change qualitatively the measured travel times, therefore the inferred interior properties of active regions, but would change the inferred values quantitatively.  The {\\it Hinode} data also enable us to construct acoustic travel time maps without using filters, and this gives us an opportunity to assess the correctness and accuracy of various filters. Further investigations are particularly important in this regard.  \\subsection{Summary} We summarize our results as follows: \\begin{enumerate} \\item By analyzing high resolution {\\it Hinode} \\ion{Ca}{2}~H observations using time-distance helioseismology, we have derived wave-speed structures beneath solar active regions. The inferred structures are remarkably similar to results obtained by various authors analyzing MDI Dopplergram observations, although the instrument, the spectrum line used to observe, the type of observation (intensity and Doppler velocity), the inversion technique, and the sensitivity kernels are mostly different. \\item The subsurface flow field structure is generally in agreement with previous MDI result, but it is clear that the downward flows near the sunspot surface is more prominent, and the mass circulation around the sunspot is more clear and seems self-consistent, although no mass conservation constraint is applied in the inversion. \\item Our analysis without using phase-speed filtering has convincingly demonstrated that the phase-speed filtering does not change travel time measurements qualitatively, but may underestimate travel time anomalies inside active regions. \\item High spatial resolution subsurface flow fields reveal quite a few organized flow structures inside sunspot umbra and penumbra, which are perhaps corresponding to some convection structures or downdraft vortex rings around magnetic flux bundles. In the light bridge area, subsurface hotter temperature is found, and plasma flows divergent from the light bridge is also seen. These initial results provide a basis for further development of high-resolution acoustic tomography of sunspots, and comparison with numerical simulations. \\end{enumerate}"
    },
    "0911/0911.1357_arXiv.txt": {
        "abstract": "The equivalent width (EW) of the \\lya{} line is directly related to the ratio of star formation rates determined from Ly$\\alpha$ flux and UV flux density [SFR(\\lya{})/SFR(UV)]. We use published data --in the literature EW and SFR(\\lya{})/SFR(UV) are treated as independent quantities-- to show that the predicted relation holds for the vast majority of observed \\lya{} emitting galaxies (LAEs). We show that the relation between EW and SFR(\\lya{})/SFR(UV) applies irrespective of a galaxy's ``true'' underlying star formation rate, and that its only source of scatter is the variation in the spectral slope of the UV continuum between individual galaxies. The derived relation, when combined with the observed EW distribution, implies that the ratio SFR(UV)/SFR(Ly$\\alpha$) is described well by a log-normal distribution with a standard deviation of $\\sim 0.3-0.35$. This result is useful when modelling the statistical properties of LAEs. We further discuss why the relation between EW and SFR(\\lya{})/SFR(UV) may help identifying galaxies with unusual stellar populations. ",
        "introduction": "\\label{sec:intro} Lyman $\\alpha$ (hereafter Ly$\\alpha$) equivalent width represents a fundamental quantity in the study of Ly$\\alpha$ emitting galaxies. The equivalent width (EW) of the \\lya{} line emitted by galaxies is a sensitive indicator of the initial mass function (IMF) or gas metallicity from which stars form \\citep[e.g.][]{S02,S03}. The existence of large equivalent width (rest-frame $\\mathrm{EW} \\ga 240$\\,\\AA{}) \\lya{} emitters has led to speculation on whether population III formation from pristine gas may actually have been observed \\citep{MR02,J06,DW07}. However, this interpretation depends sensitively on the details of radiative transfer through both the intergalactic medium \\citep[IGM, see e.g.][for a discussion]{S08}, and the interstellar medium \\citep[ISM, e.g.][]{N90,H06,Fi}. It is safe to conclude that at present, no truly convincing candidates for population III galaxies exist. The search for large equivalent width \\lya{} emitters -- and population III galaxy formation -- is a key science driver for the observational community, while modelling \\lya{} radiative transfer in large EW emitters is a challenge for theorists.  The main goal of this paper is to draw attention to the fact that \\lya{} EW is frequently discussed independently from the quantity SFR(\\lya{})/SFR(UV) \\citep[e.g.][]{Ajiki03,Fujita03,Venemans04,Shima06,Gronwall07,Tapken07,P08}. This quantity denotes the ratio of the star formation rates derived from the observed \\lya{} flux and rest-frame UV flux density. One can consider the quantity SFR(\\lya{})/SFR(UV) as an ``alternative'' measurement of EW, because the quantities EW and SFR(\\lya{})/SFR(UV) are directly related (see \\S~\\ref{sec:basis} of this paper, and e.g. \\citealp{DW07}, \\citealp{Rauch08},  \\citealp{Dayal08}, \\citealp{Nagamine08}, \\citealp{Nilsson09}).  Since EW represents a fundamental property of \\lya{} emitting galaxies, a more detailed investigation of its relation to the ratio SFR(\\lya{})/SFR(UV) is warranted.  We derive the relation between EW and SFR(\\lya{})/SFR(UV), and compare with observations in \\S~\\ref{sec:result}. We discuss the implications of our results in \\S~\\ref{sec:conc}. Throughout this paper we denote the rest frame equivalent width by REW, and the observed equivalent width by OEW. The two are related by REW=OEW$/(1+z)$. When we write ``EW'' this refers to both OEW and REW. ",
        "conclusions": "\\label{sec:conc} We investigate the relation between REW and the ratio of star formation rates derived from \\lya{} flux, and rest-frame UV flux density [SFR(\\lya{})/SFR(UV)] (Eq~\\ref{eq:ratio2}). This relation derives directly from the definition of equivalent width and star formation rate conversion factors, and its only source of scatter is the variation in the slope of the UV continuum at $\\lambda_{{\\rm Ly}\\alpha}< \\lambda <\\lambda_{\\rm UV}$ between individual galaxies. The correlation exists regardless of the assumed star formation calibrators (which themselves depend on the assumed IMF and gas metallicity), or the true star formation rates of these galaxies.  Despite their fundamentally tight relation, \\lya{} REW and SFR(\\lya{})/SFR(UV) are often discussed as independent quantities. We investigate their correlation in existing data, and find the vast majority of galaxies to be consistent with the predicted relation (see Figs~\\ref{fig:corr1}). The existence of the relation has interesting applications, which are discussed next.  \\subsection{An Empirically Constrained Ly$\\alpha$ Based Star Formation Indicator}  A SFR derived from the UV flux density is likely more reliable than a SFR derived from Ly$\\alpha$ flux. Ly$\\alpha$ scatters through the ISM and IGM which makes it hard to determine the amount of extinction. Our relation can be used to derive a more accurate Ly$\\alpha$ based star formation calibrator, if we require that --statistically-- the Ly$\\alpha$ derived SFR should be equal to the UV derived SFR. We introduce the constant $\\mathcal{M}$ such that  \\begin{equation} {\\rm SFR}_{\\mathcal{M}}({\\rm Ly}\\alpha)\\equiv \\mathcal{M}\\times {\\rm SFR(Ly}\\alpha{\\rm )}\\equiv {\\rm SFR(UV)}. \\label{eq:m} \\end{equation} The constant $\\mathcal{M}$ ensures that the SFR derived for a certain galaxy from its measured Ly$\\alpha$ flux is equal to that derived from its UV flux density. According to Eq~\\ref{eq:ratio2} $\\mathcal{M}=[C({\\rm REW}/{\\rm REW}_{\\rm c})]^{-1}$; this implies that a galaxy with an unusually large REW has $\\mathcal{M}\\ll 1$. Without this correction one would overestimate the SFR from the Ly$\\alpha$ flux alone.  In the most general case, the probability $P(\\mathcal{M})d\\mathcal{M}$ that $\\mathcal{M}$ lies in the range $\\mathcal{M}\\pm d\\mathcal{M}/2$ is given by (see Appendix~\\ref{app:m})  \\begin{eqnarray} P(\\mathcal{M})d\\mathcal{M}=d\\mathcal{M}\\hs\\mathcal{N}\\int_{-\\infty}^{\\infty} P(\\beta) P({\\rm REW}_{\\mathcal M})\\frac{{\\rm REW}_{\\mathcal M}(\\beta)}{\\mathcal{M}}d\\beta, \\label{eq:pm} \\end{eqnarray} with $\\mathcal{N}$ the normalization factor which ensures that $\\int_{0}^{\\infty} P(\\mathcal{M})d\\mathcal{M}=1$, and REW$_{\\mathcal{M}}(\\beta)\\equiv {\\rm REW}_{\\rm c}/(\\mathcal{M}C)$. Furthermore, $P(\\mathrm{REW})d\\mathrm{REW}$ denotes the probability that a LAE has an observed REW in the range REW$\\pm d$REW/2; $P(\\beta)d\\beta$ denotes the probability that a LAE has an observed $\\beta$ in the range $\\beta \\pm d\\beta/2$.  Figure~\\ref{fig:pdf} shows the probability distribution $P(\\mathcal{M})$ obtained from Eq~\\ref{eq:pm} ({\\it black solid line}). In this calculation we assume: ({\\it i}) $P(\\beta)d\\beta$ is a Gaussian with $\\bar{\\beta}=1.2$, and $\\sigma_{\\beta}=0.9$. This choice ensures that $\\sim 90\\%$ of the galaxies have $0 \\la \\beta \\la 2.4$ \\citep[cf.][]{Tapken07}; ({\\it ii}) $P(\\mathrm{REW})d\\mathrm{REW}$ is an exponential with a scale length of REW$_{\\rm L}$=76 \\AA\\hs\\citep[][]{Gronwall07}, based on observed Ly$\\alpha$ emitting galaxies at $z=3.1$ with REW$> 20$ \\AA, i.e $P$(REW)dREW$\\propto$ exp(-REW/REW$_{\\rm L})$ for REW$>$REW$_{\\rm min}\\sim$ 20 \\AA, and P(REW)$=0$ otherwise. In Appendix~\\ref{app:dos} we show that the precise shape of the distribution is not very sensitive to the assumed probability density functions (PDFs) of $\\beta$ and REW.  \\begin{figure} \\vbox{\\centerline{\\epsfig{file=fig3.eps,angle=270,width=8.0truecm}}} \\caption[]{The {\\it black solid line} shows the probability distribution for $\\mathcal{M}\\equiv$SFR(UV)/SFR(Ly$\\alpha$) as given by Eq~\\ref{eq:ratio2}, in which $P$(REW) and $P$($\\beta$) are determined from the observations (see text). The {\\it red dotted lines} mark the $68\\%$ confidence interval centered on the median $\\mathcal{M}=1.13$ (indicated by the {\\it blue dashed line}). We thus find that$(68\\%$ of LAEs have $\\mathcal{M}=1.1^{+1.4}_{-0.7}$. The data points denote the directly measured distribution of SFR(UV)/SFR(Ly$\\alpha$) collected from the papers used in our analysis. The errorbars denote Poisson uncertainties in these data. The observed data distribution agrees well with our derived PDF. Furthermore, the derived PDF is reproduced quite well by a log-normal distribution which is indicated by the {\\it grey dot-dashed line}, which is easily adopted when modeling statistical properties of LAEs.} \\label{fig:pdf} \\end{figure} Figure~\\ref{fig:pdf} shows that $P(\\mathcal{M})$ peaks at $\\mathcal{M}\\sim 0.6$. The function $P(\\mathcal{M})$ is highly asymmetric and its median is $\\mathcal{M}=1.13$, which is indicated as the {\\it blue dashed line}. That is, SFR(Ly$\\alpha$) $<0.88$ SFR(UV) for $50\\%$ of the galaxies (similarly, SFR(Ly$\\alpha$) $<$ SFR(UV) for $\\sim 57\\%$ of the galaxies). The {\\it red dotted lines} mark the $68\\%$ confidence interval centered on the median value. We therefore find that for $68\\%$ of LAEs $\\mathcal{M}=1.1^{+1.4}_{-0.7}$. In other words, the Ly$\\alpha$ derived SFR lies within a factor of $\\sim 2.5$ from the UV derived SFR for 68\\% of LAEs. Overplotted as the data points is the measured distribution of SFR(UV)/SFR(Ly$\\alpha$) collected from the papers that were used in our analysis. The errorbars denote Poisson uncertainties. This observed distribution agrees quite well with our derived PDF. The {\\it grey dot-dashed line} indicates a log-normal distribution \\begin{eqnarray} P(\\mathcal{M})d\\mathcal{M}=\\frac{1}{\\sigma\\sqrt{2\\pi}}{\\rm exp}\\Big{[} -\\frac{1}{2}\\Big{(}\\frac{[\\log \\mathcal{M}]-x}{\\sigma}\\Big{)}^2\\Big{]}\\frac{d\\mathcal{M}}{\\mathcal{M}\\ln 10}, \\label{eq:log} \\end{eqnarray} with $x=0.04$ and $\\sigma=0.35$. This log-normal distribution provides a decent fit to the derived and observed distribution, and is easily adopted when modeling statistical properties of LAEs.  \\citet{Ouchi08} conclude that between $z=3$ and $z=6$ the observed REW PDF is consistent with no redshift evolution. This implies that our derived PDF for $\\mathcal{M}$ also applies at redshifts greater than $z=3$. On the other hand, the measured scalelength of the exponential REW distribution at $z=2.3$ is REW$_{\\rm L}=48$ \\AA\\hs\\citep{Nilsson09}, causing the PDF to broaden and shift to larger $\\mathcal{M}$ (see Appendix \\ref{app:dos}). Table~\\ref{table:param} summarizes the redshift dependence of the parameters describing the log-normal distribution for $\\mathcal{M}$. In both redshift bins the standard deviation  of the log-normal distribution is very similar, with $\\sigma \\sim 0.3-0.35$. The abrupt change of the value of $x$ at $z=3$ is clearly a crude approximation of the real redshift evolution of the $\\mathcal{M}$-PDF. We have chosen this parametrization because it corresponds to the simplest description that is consistent with existing data.  \\begin{table} \\begin{minipage}{8cm} \\centering \\caption{Fit Parameters for the log-normal PDF for $P(\\log \\mathcal{M})d\\mathcal{M}$.} \\begin{tabular}{l c c} \\hline\\hline & x & $\\sigma$\\\\ $z \\gsim 3$\\hs\\footnote{We assume that the REW-PDF, and hence the $\\mathcal{M}$-PDF, does not to evolve between $z=3-6$ (Ouchi et al. 2008). We take the observed $z=3$ REW-PDF from Gronwall et al. (2007) for $z\\gsim 3$, and the observed $z=2.3$ REW-PDF from Nilsson et al. (2009) for $z\\lsim3$.} & 0.04 & 0.35 \\\\ $z \\lsim 3$ & 0.20& 0.31 \\\\ \\hline\\hline \\end{tabular} \\label{table:param} \\end{minipage} \\end{table}  In theoretical work the inverse problem often arises in which a Ly$\\alpha$ luminosity must be obtained from a physical SFR. In such a case one may write $L_{{\\rm Ly}\\alpha}=\\frac{1.1\\times 10^{42}\\hs{\\rm erg}\\hs{\\rm s}^{-1}}{\\mathcal{M}}\\Big{(}\\frac{{\\rm SFR}}{M_{\\odot}\\hs{\\rm yr}^{-1}}\\Big{)}$ and adopt our PDF for $\\mathcal{M}$. This allows one to assign an empirically calibrated, variable Ly$\\alpha$ luminosity to galaxies of a given SFR. This prescription is clearly not perfect, because not all galaxies that are actively forming stars show a Ly$\\alpha$ emission line. For example, only 20-25$\\%$ of Lyman Break galaxies (LBGs) at $z=3$ have a Ly$\\alpha$ line with REW$\\gsim 20\\%$ \\AA\\hs and would classify as a LAE \\citep[e.g.][]{Shapley03}, and the connection between LBGs and LAEs is not well understood. Nevertheless, our suggested prescription for assigning Ly$\\alpha$ luminosities to galaxies  of a given SFR is more realistic than the often used one-to-one relation $L_{{\\rm Ly}\\alpha}=1.1 \\times 10^{42}\\hs{\\rm erg}\\hs{\\rm s}^{-1}\\hs {\\rm SFR}/(M_{\\odot}\\hs{\\rm yr}^{-1})$.  \\subsection{Outliers in the SFR(Ly$\\alpha$)/SFR(UV)-REW Plane: Signposts for Unusual Galaxies?}  ``Normal'' star forming galaxies can emit a maximum REW$_{\\rm max}=240$\\,\\AA. Galaxies with REW$>$REW$_{\\rm max}$ may signal the presence of a galaxy that contains population III stars \\citep[e.g.][]{MR02}. Using Eq~\\ref{eq:ratio2} this corresponds to [SFR(\\lya{})/SFR(UV)]$_{\\rm max} =2.8^{+0.4}_{-0.5}$. The majority of observed galaxies that have quoted uncertainties on SFR(\\lya{})/SFR(UV) are consistent with this upper limit.  Objects where nebular emission dominates, such as galaxies containing population III stars \\citep{S02} or cooling clouds \\citep{D09}, may be dominated by the two-photon continuum at 1216\\,\\AA{} $< \\lambda <$ 1600\\,\\AA{}. This results in unusually negative values for $\\beta$. These objects may have been identified in the spectrum itself, because deviations from the correlation are caused by unusual spectral slopes. However, reliable measurements of the continuum just redward of the Ly$\\alpha$ line and at $\\lambda_{\\rm UV}=1400$ \\AA\\hs provide a long baseline in wavelength, which may more clearly reveal the presence of a continuum dominated by two-photon emission. This suggests that outliers in the REW - SFR(Ly$\\alpha$)/SFR(UV) plane may provide a more sensitive probe to cooling clouds or primordial galaxies than the spectrum alone. That additional probe is important especially at high redshifts, where the IGM may transmit only a small fraction of the \\lya{} emitted by galaxies \\citep{igm}. In this case even those star forming galaxies containing population III stars may have REW$<$REW$_{\\rm max}$. This strongly suggests that the determined REW alone is not enough to identify a population III galaxy.  {\\bf Acknowledgements} M.D. is supported by Harvard University funds. E.W. acknowledges the Smithsonian Institution for the support of his postdoctoral fellowship. We thank an anonymous referee for helpful, constructive comments that improved the content of this paper."
    },
    "0911/0911.2244.txt": {
        "abstract": " ",
        "introduction": "In the early nineties, cosmic shear was predicted to be a promising way to study the distribution of matter in the Universe \\cite{1991MNRAS.251..600B,1992ApJ...388..272K,1991ApJ...380....1M} and since its first detection \\cite{2000MNRAS.318..625B,2000astro.ph..3338K,2000A&A...358...30V,2000Natur.405..143W} it has shown to be a precious mean of investigations of the large-scale structure of the Universe, enabling us to explore dark energy properties or uncover signatures of mode coupling effects \\cite{2008ARNPS..58...99H,1999ARA&A..37..127M,2008PhR...462...67M}.  So far cosmic shear surveys have covered only a limited field in the sky. For instance, the CFHTLS,\\footnote{http://www.cfht.hawaii.edu/Science/CFHTLS} which has produced  very promising results over the last years  (see e.g.~\\cite{2008A&A...479....9F} for a recent account of these observations), is limited to about a 170 squared degree range. With the demonstration of the robustness of cosmic shear observations, (nearly) full-sky surveys such as Pan-STARRS, DES, LSST, JDEM, or Euclid are under preparation. They will open the way to new types of studies. Akin to CMB observations, such surveys will be an excellent tool to explore the physics of the Universe at scales comparable to the Hubble radius, therefore testing genuinely general relativistic effects.\\footnote{See \\cite{2009arXiv0907.0707Y} for a recent account of these effects on galaxy clustering observations.} In particular, the study of mode couplings, already well established on Newtonian scales, can be extended at these very large scales therefore testing the details of our understanding of the origin and formation of the large-scale structure.   Such an investigation requires that we know the types of mode couplings that are expected to be seen at such large scales. Calculations have been undertaken to predict the nonlinear growth of metric and density fluctuations after modes reenter the Hubble radius \\cite{Bartolo:2005kv,Boubekeur:2008kn,Fitzpatrick:2009ci}. In the context of the CMB anisotropies, progress has been recently made in understanding the effect of these nonlinearities, from concentrating on the large angular scales \\cite{Pyne:1995bs,Bartolo:2004ty,2009JCAP...08..029B} to the details of the physics of recombination (see for instance \\cite{2006JCAP...06..024B,2009CQGra..26f5006P}).   So far the investigations of mode couplings in weak lensing  were limited to small angular scales, corresponding to scales much smaller than the angular diameter distance at the source. Accidentally, this distance roughly corresponds to the Hubble radius at the source. Thus, on these scales one can consistently neglect general relativistic effects that are suppressed by the ratio between the scale probed and the Hubble scale. On small angular scales the dominant contribution to the cosmic shear comes from fluctuations of the gravitational potential transverse to the line of sight. Perturbations along the line of sight average out and do not yield appreciable effects. In this regime the dominant geometrical mode couplings were identified more than a decade ago in \\cite{1997A&A...322....1B}. They include the Born correction and the lens-lens coupling. In the so-called Born approximation one integrates the lensing distortion over an unperturbed photon path. One can consider the correction due to the fact that the photon path is perturbed. The lens-lens coupling consists in the correction due to the deformation of a distant lens caused by a foreground one. The consequences of these effects have been extensively described  in the literature and they have been found to have an impact on both the shear power spectrum and higher-order statistical observables such as the bispectrum \\cite{Cooray:2002mj,2003MNRAS.344..857T,Dodelson:2005zj,2006JCAP...03..007S,Hilbert:2008kb,Krause:2009yr}. As the shape distortion probes the reduced shear rather than the shear itself, there is another correction associated to the nonlinear conversion between these two quantities \\cite{Dodelson:2005rf,Schneider:1997ge,Krause:2009yr}. Finally, another nonlinear effect is the source-lens clustering, due to the fact that the source of the lensed light is itself a perturbed field with specific clustering properties correlated with the lens \\cite{1998A&A...338..375B,2002A&A...389..729S}. For the current surveys restricted to a limited angular field all these types of couplings are undoubtedly the dominant ones.    In view of full-sky surveys one needs to go beyond the small-angle approximation and probe scales of the order of the angular diameter distance to the source. In this case fluctuations along the line of sight are not negligible and terms other than those described above may become important. Furthermore, as we are probing scales comparable to the Hubble size, one needs to undertake a full general relativistic treatment. This is important in order to compute accessible higher-order observables in full or almost full-sky surveys. In particular, in such surveys, it becomes necessary if one wants to compute the lensing bispectrum in the squeezed configuration, when one of the scales probed is taken to be much larger than the other two.    We present here the exhaustive calculation of the weak lensing cosmic shear at second order including all general relativistic contributions and without relying on the small-angle or thin-lens approximations. However, we do not include in our study the effect of source-lens clustering \\cite{1998A&A...338..375B,2002A&A...389..729S} and other intrinsic effects in the alignment and ellipticity of galaxies. We work in the so-called generalized Poisson gauge without specifying the matter content of the Universe. We assume that there are no primordial vector and tensor perturbations. However, we will take into account vector and tensor components of the metric generated at second order from scalar fluctuations. In this gauge we derive the reduced -- i.e.~observable -- shear by solving the Sachs equation  \\cite{1961RSPSA.264..309S} which describes the distortion of the cross section of an infinitesimal bundle of light rays in the geometric-optics limit. The advantage of using the Sachs equation instead of  the geodesic equation is that it deals only with physically observable quantities. As the resolution of the Sachs equation is extremely tedious and involves a large number of terms we will develop tests that allow us to check the validity of its solution. In particular, some of the contributions to the second-order shear that we compute -- and that are usually neglected in the small-angle approximation -- can be compared to those expected from the lensing shear at linear order in a universe with spatial curvature.  Using the solution of the Sachs equation we will compute the reduced shear by adding the nonlinear corrections coming from the relation between this quantity and the shear itself. Finally, as in the Poisson gauge hypersurfaces of constant redshift are inhomogeneous, we take into account the corrections due to the inhomogeneity of the redshift of the source \\cite{2008PhRvD..78l3530B}. The final observable reduced shear field that we obtain is a gauge invariant quantity, although  its separate contributions are not necessarily so. As it is customary for CMB polarization, we express the reduced shear in terms of spin-2 operators on the sky. In particular, angular gradients on the sky will be written in terms of spin raising and lowering operators, whose eigenfunctions are the well-known spin-weighted spherical harmonics.    The plan of the paper is the following. In Sec.~\\ref{sec:weak} we give the outline of our calculation. In particular, we describe the Sachs equation and how the transverse size of a propagating beam can be related to the observable reduced shear including the effects of the source redshift inhomogeneities. We will solve the Sachs equation at first order in Sec.~\\ref{sec:linear} while the full second-order calculation will be presented in Sec.~\\ref{sec:shear_2nd}. We discuss and comment on our results in Sec.~\\ref{sec:conclusion} in the context of cosmic shear surveys regarding the generation of $B$ modes and the expected contributions to the cosmic shear bispectrum.      %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% ",
        "conclusions": "\\label{sec:conclusion} In this article we have derived the expression of the reduced cosmic shear up to second order in the perturbations with full-sky validity. Our main result is summarized in eq.~(\\ref{final_result}). As it is expressed in terms of spin-2 operators on the sphere it can be decomposed as sum of spin-weighted spherical harmonics on the sky. Indeed, this description ensures that our observable has a genuine spin-2 behavior on the celestial sphere.  Our result is written in terms of the metric perturbations in the generalized Poisson gauge. These are the scalar potentials $\\phi$ and $\\psi$ and the vector and tensor components of the metric generated at second order, respectively $\\omega_i$ and $h_{ij}$. Let us first comment on the first two terms on the right-hand side of eq.~(\\ref{final_result}). Remarkably, the contribution from scalar perturbations from the sum of these two terms can be expressed in terms of the Weyl potential $\\Psi = (\\phi+\\psi)/2$ only. As explained, this is due to the fact that null geodesics are conformally invariant. These two terms contain the well-known second-order corrections due to lens-lens coupling and departure from the Born approximation, which dominate in the small-angle approximation. On larger angular scales new couplings become important. These are an intrinsic contribution which is a purely general relativistic effect at second order, a coupling between the gravitational potential at the source with the lens and corrections due to couplings between the lens and the photon time-delay. We have checked that these contributions can be independently reconstructed from the calculation of the shear at first order in a universe with a radially dependent spatial curvature. Other checks, such as the invariance under a homogeneous time shift and a conformal transformation can be used to verify the validity of these new corrections. Another scalar correction appears in the form of products of two spin-1 fields and comes from the couplings between two photon deflections. Finally, besides the scalar contributions, the shear gets a contribution from  spin-2 quantities defined from the vector and tensor components of the metric generated at second-order. Note that the separation between all these contributions is not gauge invariant.   In Poisson gauge, the correction due to the coupling between the photon redshift perturbation and the lens cannot be written in terms of $\\Psi$ only. Indeed, the integrated contribution to the photon redshift -- the integrated Sachs-Wolfe effect -- is a time integral over $\\dot \\Psi$ but the intrinsic contributions -- Sachs-Wolfe and Doppler effects -- are expressed in terms of the Newtonian gravitational potential and the velocity along the line of sight and do not depend on $\\Psi$ only.   We are now in the position to explore the phenomenological consequences of these results in view of the future (partially) full-sky lensing surveys. In particular, the new corrections that we have computed should become relevant in deriving the lensing bispectrum on large angular scales. For instance, to compute the bispectrum in the squeezed limit one needs to take one of the three modes to be much smaller than the other two, corresponding to angular scales comparable to the depth of the survey. As the lensing is a cumulative effect integrated along the line of sight, it is difficult, at this stage, to precisely guess the relative importance of the various contributions. In particular, although the lens-lens coupling terms are a priori larger by a factor $\\sim l^2$ compared to the others, they may be damped by geometrical factors. We leave these investigations for the future."
    },
    "0911/0911.4955_arXiv.txt": {
        "abstract": "We constrain the cosmological density of cosmic string loops using two observational signatures -- gravitational microlensing and the Kaiser-Stebbins effect. Photometry from RXTE and CoRoT space missions and pulsar timing from Parkes Pulsar Timing Array, Arecibo and Green Bank radio telescopes allow us to probe cosmic strings in a wide range of  tensions $G\\mu/c^2=10^{-16}\\div10^{-10}$. We find that pulsar timing data provide the most stringent constraints on the abundance of light strings at the level $\\Omega_s \\sim 10^{-3}$. Future observational facilities such as the Square Kilometer Array will allow one to improve these constraints by orders of magnitude. ",
        "introduction": "} Cosmic strings are now a widely recognized part of cosmological theory. Cosmic stings appear naturally in a multitude of inflationary models as topological defects from the early Universe (e.g. \\cite{Allen1990,Copeland2009}, for more  studies see references in \\cite{Chernoff2009}). Similar objects commonly referred to as cosmic superstrings can also be produced in fundametal string and M-theories \\cite{daviskibble, polchinski}. In the present study we will not differentiate between the two classes, because their observational signatures considered in this paper are the same.  The key parameter of a cosmic string is its tension $\\mu$, which is assumed to be related to the effective energy scale of the string-producing theory $\\Lambda$ by \\cite{Vilenkin1994} $$ \\frac{G\\mu}{c^2}\\sim\\frac{\\Lambda^2}{M_\\mathrm{Planck}^2}. $$ Earliest theories of string formation placed time of their generation to the Grand Unification Theory epoch and therefore their tension seemed to be of order $10^{-6}$. Initially, possible tensions of string were constrained from both sides: $10^{-11}<G\\mu/c^2<10^{-6}$, but eventually the lower bound was removed and  strings with arbitrarily low tension are theoretically allowed now (e.g.~\\cite{Firouzjahi2005}). Cosmic strings with low tension can solve some astrophysical problems: e.g., recently, cosmic strings with tensions about $G\\mu/c^2\\sim10^{-12}$ were proposed as prominent source of high-energy cosmic rays \\cite{Vachaspati2009}.  Simulations \\citep{Allen1990, Vanchurin} suggest the energy fraction in strings has a scaling behavior: their density $\\Omega_s$ (in units of critical density $\\rho_0=3H_0^2/8\\pi G$) does not depend on cosmological time. Recent works suggest that strings contribute a subdominant fraction to the energy balance of the Universe \\cite{Sakellariadou2006}. This question has not been completely resolved yet. We will treat $\\Omega_s$  essentially as a free parameter and will try to constrain it observationally.  There have been a number of attempts to limit the density of strings from various observational perspectives. Heavy enough strings, if present, would make distinctive imprints on the cosmic microwave background and shapes of distant lensed galaxies; these techniques indicate absence of strings with $G\\mu/c^2\\ge 10^{-7}$ in the Universe \\cite{Morganson2009,Fraisse2008, sazhin}. More stringent, though more model-dependent constraints come from pulsar timing (PT): cosmic strings emit gravitational waves and the corresponding background can be detected by usual methods of PT\\cite{Damour2005}. This method rules out any significant presence of cosmic strings with tensions $\\mu$ at $G\\mu/c^2<10^{-9}-10^{-8}$ \\cite{DePies2007}. Ultimately, planned mission \\textit{LISA} are expected to test the presence of gravitational wave background from lighter strings with tensions down to $G\\mu/c^2\\sim 10^{-14}-10^{-16}$.  String network has a complicated structure with a combination of long straight segments and a population of string loops of various lengths $L$ that were formed in the interconnections between straight strings. The loops oscillate relativistically ($\\beta\\sim\\mathcal{O}(1)$) with an amplitude of order $L$ and period $T=L/2c$, emitting gravitational waves and eventually decaying. Only sufficiently long loops survive by the present time.  Recent simulations \\cite{Chernoff2009} show that the surviving large-scale ($L\\sim1\\,\\mathrm{pc}$ and above) loops of light cosmic strings experience considerable clustering, which closely follows that of the dark matter,  albeit with a somewhat lower amplitude. In central parts of large galaxies, such as the Milky Way, the loops' density can be enhanced by up to $10^5$ relative to its average cosmological value. This density enhancement significantly boosts the detection rates of experiments sensitive to the local population of lenses \\cite{Chernoff2009, chetye}.  In this paper, we investigate how the clustering affects the prospects of detection of local cosmic strings via two observational signatures -- lensing on the string as seen in the photometry of background objects and Kaiser-Stebbins effect affecting the timing of millisecond pulsars.  The paper is organized as follows. In the next Section~\\ref{depthsection}, we discuss the expected rate of the effects and lay out a simple formalism with which one can interpret the non-detection of the effect in a given experiment. Then in Sections~\\ref{Eventphenomena} and~\\ref{Pulsartiming} we consider lensing and Kaiser-Stebbins effect in a greater detail and calculate the constraints on the string loop population from existing observations. The final Section discusses the obtained results and offers suggestions for further research. ",
        "conclusions": "\\begin{figure} \\center \\includegraphics[width=86mm]{combined2.jpg} \\caption{Combined constraints on the average density of cosmic string loops based on \\corot, {\\it RXTE} and {\\it PPTA} data at P=95\\% level. One can see that pulsar timing currently provides the strongest constraints down to $G\\mu/c^2\\sim 10^{-14}$ while lighter strings are constrained by the available data on Sco~X-1; \\corot data do not significantly constrain the population of loops at present.} \\label{combinedconstraintfig} \\end{figure}  Figure~\\ref{combinedconstraintfig} presents our final result including constrains from pulsar timing, X-ray data for Sco~X-1 and precision photometry from \\corot. One can see that the most stringent limits come from pulsar timing except for the lightest strings where competitive constraints are provided by {\\it RXTE} data. Presently available \\corot data do not have much constraining power in the range of tensions we consider. Existing observations allow one to limit the average density of cosmic string loops down to $\\Omega_s\\sim10^{-3}$ at $G\\mu/c^2=10^{-14}$. For larger tensions the limits become eventually weaker, proportionally to $\\mu^{-1}$ for $G\\mu/c^2 \\le 10^{-13}$ and $\\mu^{-1.5}$  for $G\\mu/c^2 > 10^{-13}$ because of lesser enhancement of density  of heavy strings.  These results clearly demonstrate that the density enhancement in the Galaxy improves the limits on cosmic string abundance that can be set by test sensitive to the local population of loops as suggested by~\\cite{Chernoff2009}. However, the same approach can also be applied to extragalactic sources, most notably quasars. The probability of lensing by string loops for these sources  is lower due to absence of enhancement in the extragalactic case. On the other hand, the huge distance to the source  can easily compensate for the relative deficit of string density. Moreover, these observations are insensitive to theoretical uncertainty in predicting $\\eta$.  As an example we consider the 'Einstein cross' quasar QSO~2237+0305. The Optical Gravitational Lensing Experiment ({\\it OGLE}) data for this source currently span more than 7 years of high-quality photometry sampled every few days~\\cite{ogle2237} and no obvious signature of string lensing are seen in these light curves\\footnote{The light curve for image~B does show an approximately twofold increase in brightness starting near $JD\\sim2453500$ but present data is far from being sufficient to draw any firm conclusions. Unfortunately, no observations can resolve this source at present to the level needed to rule out the possibility of lensing by cosmic string.}. This non-detection can be used to place some limits on the density of cosmic strings in the range of tensions $3\\cdot10^{-12}<G\\mu/c^2<3\\cdot10^{-10}$ (corresponding to the strings with crossing time equal to the sampling interval and data span, respectively).  Using (\\ref{Omegaconstrain}) with enhancement  factor $\\eta=1$, $D\\sim2\\,\\mathrm{Gpc}$, $T\\sim7.5\\,\\rm years$ yields \\bea \\label{huchra_2} \\Omega_s\\sim4\\cdot10^{-2}\\left(\\frac{G\\mu/c^2}{3\\cdot10^{-13}}\\right) \\ena These limits are inferior to constraints from pulsar timing by about an order of magnitude.  Future observational projects will noticeably improve the limits established in this paper. In \\cite{Kuijken2008} it was proposed to use upcoming large-scale observational surveys of quasar variability to search for cosmic strings. In line with this paper, we might assume that a large number $N\\sim10^3$ of quasars is monitored for $T\\sim10\\,\\mathrm{years}$ on a daily basis. Non-detection of string crossings in such survey would limit the average cosmic string loop density with $G\\mu/c^2\\ge 10^{-13})$ at the level of \\bea \\label{qso_survey_constraint} \\Omega_s= 2\\cdot10^{-5}\\left(\\frac{G\\mu/c^2}{10^{-13}}\\right) \\left(\\frac{1\\,\\mathrm{Gpc}}{D}\\right)\\left(\\frac{10^3}{N}\\right)\\left(\\frac{10~\\rm yr}{T}\\right), \\ena which is orders of magnitudes better that constraints presented in this paper --  essentially as a by-product of AGN long-time variability study.  Other avenues to better constraining $\\Omega$ at lower tensions exist. With the eventual arrival of the Square Kilometer Array (\\textit{SKA}) we can expect further enhancement  in the sensitivity due to the increase of number of pulsars and timing of more distant pulsars~\\cite{ska}. Pulsar timing array at \\textit{SKA} will consist of $\\sim$100 pulsars, that would be timed with precision better than 100 ns. Exact limits will depend on distances  of pulsars, but we can forecast that accessible range and density estimates will improve by orders of magnitude."
    },
    "0911/0911.3058_arXiv.txt": {
        "abstract": "We have performed a series of $N$-body simulations to model the Arches cluster. Our aim is to find the best fitting model for the Arches cluster by comparing our simulations with observational data and to constrain the parameters for the initial conditions of the cluster. By neglecting the Galactic potential and stellar evolution, we are able to efficiently search through a large parameter space to determine e.g. the IMF, size, and mass of the cluster. We find, that the cluster's observed present-day mass function can be well explained with an initial Salpeter IMF. The lower mass-limit of the IMF cannot be well constrained from our models. In our best models, the total mass and the virial radius of the cluster are initially $(5.1 \\pm 0.8)\\cdot10^4\\,\\msun$ and $0.76 \\pm 0.12\\,\\pc$, respectively. The concentration parameter of the initial King model is $w_0 = 3-5$. ",
        "introduction": "The Arches cluster is one of only a few young and massive starburst clusters in the Milky Way. Its location at a projected distance of less than $30\\,\\pc$ from the Galactic center and an age of only $\\sim 2.5\\,\\myr$ \\citep{FNG02,NFH04} make this cluster a unique object for studying star formation and dynamical processes in the center of galaxies.  The observed present-day mass of the Arches cluster has been estimated with $\\sim 1-2 \\cdot 10^4\\,\\msun$ \\citep{FKM99,ESM09}.  With this mass a cluster will not survive long in the Galactic center environment and evaporate on a time scale maybe as fast as $\\sim 10\\,\\myr$ \\citep{KML99,PZMM02}. The initial mass of the cluster has been determined from $N$-body simulations, however, different results have been obtained by different authors: \\citet{KFL00} found that their best model for the Arches cluster had a total mass of about $2\\cdot 10^4\\,\\msun$; \\citet{PZMM02}, on the other hand, came to the conclusion that the cluster was initially more massive than $\\sim4\\cdot 10^4\\,\\msun$.  The initial mass function (IMF), a key aspect of star formation, seems to be uniform throughout the universe \\citep{Kroupa02}. This universal IMF can be described by the power-law found by \\citet{S55} for stars in the solar neighborhood and is valid from $0.5\\,\\msun$ to the highest masses. Below $0.5\\,\\msun$, the IMF is significantly flattened \\citep[e.g.][]{Kroupa02}.  Determining the IMF of young clusters from observations is not a straight-forward process. Uncertainties can arise from the measured luminosities, the estimated age of the cluster, the completeness of the observed sample, and the stellar evolution models. In addition, the non-linear dynamical evolution of the cluster has to be taken into account as shown in Fig.~\\ref{fig:IMFMS}: as the star cluster evolves, more massive stars (star symbols) will move towards the cluster center and low-mass stars (points) in the opposite direction (indicated by the arrows in the left image). If the detection of cluster members is radially limited (dashed circle), it will result in an observed mass function (MF) in the mass-segregated cluster (right image) that is different from the IMF. This effect is visualized in Fig.~\\ref{fig:IMFMS} by the ratio of low- to high-mass stars inside the dashes circles before and after mass segregation.  \\begin{figure} \\includegraphics[width=84mm,angle=0]{plots/IMFandMassSeg.eps} \\caption{Schematical view on the cluster mass function and evolution. Low- and high-mass stars are shown by points and star symbols, respectively. The dashed circle indicates an radial observational limit. In the left image, the cluster is at $t=0$ with arrows denoting the effects of dynamical evolution.  The right images shows the mass-segregated cluster.} \\label{fig:IMFMS} \\end{figure}   In case of the Arches cluster, observations have revealed that the slope of the observed mass function (MF) is significantly flattened with $\\Gamma \\approx -0.9 \\pm 0.15$ with respect to the standard Salpeter IMF ($\\Gamma = -1.35$) \\citep{SBG05,SGB02,FKM99}, and therefore the Arches cluster has been regarded as an argument against the universality of the IMF. More recently, however, \\citet{ESM09} derived a slope of $\\Gamma = -1.1 \\pm 0.2$ in $R < 0.4\\,\\pc$ and concluded that a standard Salpeter IMF cannot be ruled out for Arches. In addition to the radial variation in $A_V$, these authors also accounted for differential extinction variations, which can severely affect the incompleteness and may have biased the earlier results. Large uncertainties in the slope still remain, revealing the necessity to compare the observed cluster MF with simulations.  In addition to the flattened slope, there has been some debate whether the IMF of Arches is truncated at the low-mass end as the result of the extreme conditions at the Galactic center where the cluster has formed. Possible evidence for this was reported by \\citet{SBG05}, who determined a low- and intermediate-mass depleted MF in the cluster core ($R < 0.2\\,\\pc$) with a turn-over at $6-7\\,\\msun$. This truncation in the MF was not seen by \\citet{KFK06}. They only found a local bump in the MF at $\\sim6\\,\\msun$.  Even if the MF is truncated at the low-mass end, it remains unclear whether this would be the result of a truncated IMF or a dynamical effect such as tidal stripping of low-mass stars.     Several studies have been done to determine the IMF of the Arches cluster using numerical simulations, again coming to different",
        "conclusions": "\\label{sec:conclusions}  \\begin{figure*} \\includegraphics[width=84mm,angle=0]{plots/Arches_all_gray.eps} \\includegraphics[width=84mm,angle=0]{plots/Arches_IKW03S05N150_Seed010_Rvir070_gray.eps} \\caption{The observed cluster (left) and a snapshot of one of the best fitting models (right) in comparison. The images are centered on the center of density and the circles indicate a radius of $0.4\\,\\pc$. Gray-scale and point size represent stellar masses. In the right panel, stars have been removed randomly to mimic incompleteness.} \\label{fig:bfmsnapshot} \\end{figure*}  A number of models in Tab.~\\ref{table:bestfit} can be considered best fit models, for example model IKW03S05 with $\\rvir = 0.70$ and $\\NMS = 150$. A snapshot of this model at $t=2.5\\,\\myr$ is shown in Fig.~\\ref{fig:bfmsnapshot}. Stellar masses are indicated by the point size and gray-scale and the center of density of the cluster is located at the origin. The dashed circle has a radius of $0.4\\,\\pc$. In the left panel, the NACO data is plotted for comparision in the same way as the simulation data in the right panel. In the latter, stars have been randomly removed to mimic incompleteness. The probability that a star is removed is given by a function fitted to the data shown in Fig.~\\ref{fig:inc} and depends on mass and position of the star.  At $2.5\\,\\myr$, the cluster has a total mass of $\\sim2\\cdot10^4\\,\\msun$ inside a radius of $0.4\\,\\pc$ and twice that mass inside the tidal radius of $1\\,\\pc$. On average, our best models predict that the Arches cluster is more massive than observed, even if we take into account mass loss by stellar evolution. About $6\\cdot10^3\\,\\msun$ of the total mass are in stars below $1\\,\\msun$. Our results are in good agreement with the previous findings of \\citet{PZMM02}.  Generally, only models with the Salpeter IMF produce acceptable fits and we therefore conclude that the IMF in the Arches cluster is normal despite the extreme environment in which the cluster is formed. However, we cannot rule out a turn-over as suggested by \\citep{SBG05} as our results are not very sensitive to the lower mass limit of the IMF. Our best fitting models can have any lower mass limit in the the range of $0.5 - 4.0\\,\\msun$ that we investigated.  \\begin{figure} \\includegraphics[width=84mm,angle=0]{plots/imf_IKW03S05N150R070.eps} \\caption{The IMF of the best fitting model for all stars and the stars in two radial bins.} \\label{fig:bfmimf} \\end{figure}  Fig.~\\ref{fig:bfmimf} shows the MF of one of our best fit models. A bin size of $\\Delta log(m/\\msun) = 0.2$ was used to obtain the individual star counts and the data was also averaged over the ten different realization of this model. The top line shows the MF for all stars which is simply the Salpeter IMF used in the initial conditions of this model. The other two lines show the MF in two radial bins (the same used by \\citet{SBG05} and \\citet{ESM09}). The MF is noticeably flattened for massive stars ($m \\gap 10\\,\\msun$) in the inner radial bin ($R < 0.2\\,\\pc$), while it remains unchanged in the outer radial bin ($0.2 < R < 0.4\\,\\pc$). The flattening of the MF in the cluster core has been reported in all the previous observational studies, though with varying $\\Gamma$-values determined for the slope. Our best fit models are in very good agreement, however, with the latest results from \\citet{ESM09}. Stars with $m>10\\,\\msun$ can be found in our models as far as $10\\,\\pc$ from the cluster center. However, we find that in order to measure the correct slope of the MF, we only need the stars inside the tidal radius of $1\\,\\pc$.  No turn-over at the low mass end of the IMF can be seen. This means that the turn-over seen by \\citet{SBG05} cannot be explained by mass segregation (unlike the flat IMF at the high-mass end). Deriving the IMF from our simulation data is not hampered by incompleteness, selection effects, and mass determination from observed luminosities. All these effects may bias the observed star counts to produce a turn-over, however, the turn-over appears at $\\sim6\\,\\msun$ where the data is still 50\\% complete. So, either the IMF is indeed truncated or another effect has to be considered. One possible explanation is that tidal stripping preferably removes low-mass stars from the cluster. This effect is not included in our current simulations but will be tested in a follow-up paper. Alternatively, \\citet{ESM09} pointed out that local variations in extinction can account for (some of) the flattening of the obserevd MF. The same effect could cause lower-mass stars to be lost preferentially. Then the completeness function, which only takes into account crowding and sensitivity effect, would not be sufficient to correct the low-mass end of the IMF.  \\begin{figure} \\includegraphics[width=84mm,angle=0]{plots/rcore.eps} \\caption{The core radius evolution with time of model IKW03S05 with $\\rvir=0.7\\,\\pc$ and $\\NMS=150$.  The black line with errorbars shows the average of ten random realizations of the same model (grey lines). Core collapse happens at $\\sim4\\,\\myr$. } \\label{fig:rcore} \\end{figure}  In Fig.~\\ref{fig:rcore} we show the core radius evolution with time for one of our best fit models. At its current age of $3.5\\,\\myr$ the cluster is a little more than halfway to core collapse which will occur at $\\sim4\\,\\myr$. This result is in agreement with the findings of \\citet{PZGC07}.      We have performed a large number of $N$-body simulations in order to find the best fitting model for the Arches cluster. The available observational data has been used to constrain the free parameters in our model. In a systematic analysis, we compared the total mass, the cumulative mass profile, and the present-day MF and defined fitness parameters for each of the three observables.  The main conclusion from our analysis is that the Arches cluster, despite of being born in an extreme environment, has formed with a standard Salpeter IMF. Due to mass segregation, the slope of the observed MF is flattened inside a radius of $0.4\\,\\pc$. This radius was imposed a limit from the obversavational data used, and we estimate that a limiting radius of $\\sim1\\,\\pc$ would be required for the observed MF to match the underlying IMF.  We conclude from our best fitting models that the Arches cluster was born with an initial virial radius of $0.76 \\pm 0.12\\,\\pc$ and an initial total mass of about $(5.1 \\pm 0.8) \\cdot 10^4\\,\\msun$ assuming a lower mass limit of $0.5\\,\\msun$ for the Salpeter IMF. The lower mass limit cannot be constrained well from our models, giving rise to additional uncertainties in determining the initial cluster mass. The King model concentration parameter of the best fitting models is $W_0 = 3-5$, however, the $W_0 = 7$ produced also reasonable fits so that this parameter is not very well constrained.  We neglected the Galactic potential and also stellar evolution for the simulations in this paper. These processes will be included in a following paper to get a more realistic model of the Arches cluster, however this first step was needed in order to reduce the number of free parameters for the these (computationally more expensive) simulations, which further constrain the dynamical evolution and tidal effects acting on the Arches cluster, and thereby the initial conditions of this nearby starburst, such as the cluster mass, the orbital motion and the IMF at the birth of the cluster.    {\\bf Acknowledgments} SH and SPZ are grateful for the support from the NWO Computational Science STARE project \\#643.200.503 and NWO grant \\#639.073.803.  AS acknowledges funding from the German Science Foundation (DFG) Emmy-Noether-Programme under grant \\mbox{STO 496-3/1}."
    },
    "0911/0911.1788_arXiv.txt": {
        "abstract": "This is the fourth of a series of papers in which we derive simultaneous constraints on cosmological parameters and X-ray scaling relations using observations of the growth of massive, X-ray flux-selected galaxy clusters. Our data set consists of 238 clusters drawn from the \\ROSAT{} All-Sky Survey, and incorporates extensive follow-up observations using the \\Chandra{} X-ray Observatory. Here we examine the constraints on neutrino properties that are enabled by the precise and robust constraint on the amplitude of the matter power spectrum at low redshift available from our data. In combination with cluster gas mass fraction, cosmic microwave background, supernova and baryon acoustic oscillation data, and incorporating conservative allowances for systematic uncertainties, we limit the species-summed neutrino mass, \\Mnu{}, to $<0.33\\eV$ at 95.4 per cent confidence in a spatially flat, cosmological constant (\\LCDM) model. In a flat \\LCDM{} model where the effective number of neutrino species, \\Neff{}, is allowed to vary, we find $\\Neff=3.4_{-0.5}^{+0.6}$ (68.3 per cent confidence, incorporating a direct constraint on the Hubble parameter from Cepheid and supernova data). We also obtain results with additional degrees of freedom in the cosmological model, in the form of global spatial curvature (\\Omegak) and a primordial spectrum of tensor perturbations ($r$ and \\nt{}). The results are not immune to these generalizations; however, in the most general case we consider, in which \\Mnu{}, \\Neff{}, curvature and tensors are all free, we still obtain $\\Mnu<0.70\\eV$ and $\\Neff=3.7 \\pm 0.7$ (at, respectively, the same confidence levels as above). These results agree well with recent work using independent data, and highlight the importance of measuring cosmic structure and expansion at low as well as high ($z \\sim 1100$) redshifts. Although our cluster data extend to redshift $z=0.5$, the direct effect of neutrino mass on the growth of structure at late times is not yet detected at a significant level. ",
        "introduction": "\\label{sec:introduction}  Observations of neutrino flavor oscillation have conclusively shown that the neutrino mass eigenstates are non-degenerate \\citep[e.g.][]{Fukuda98,Fukuda02,Ahmad02,Ahn03,Ahn06,Eguchi03,Sanchez03,Giacomelli04,Aharmim05}. Although, these observations can place tight constraints on the differences in squared mass, measuring the absolute mass scale remains challenging. Current laboratory efforts focus on tritium beta decay \\citep[e.g.][]{Lobashev03,Kraus05} and neutrinoless double beta decay \\citep[e.g.][]{Aalseth99,Klapdor01,Arnaboldi05,Arnold05}; the latter, if observed, would additionally indicate that neutrinos are Majorana rather than Dirac fermions. The squared mass differences measured from flavor oscillations place a lower bound on the sum of the three masses, $\\Mnu=\\sum_i m_i$, at $\\sim 0.056$ (0.095)$\\eV/c^2$ in the normal (inverted) hierarchy, while current tritium beta decay results provide an upper bound on the mass of the electron neutrino at $\\sim 2\\eV/c^2$ (thus on \\Mnu{} at $\\sim 6\\eV/c^2$).\\footnote{Henceforth, we set $c=1$.}  Because neutrinos play a prominent role in the early Universe, cosmological observations are also sensitive to their properties (for a review, see \\citealt{Lesgourgues06}; see also \\secref{sec:background}). The primary effect of non-zero neutrino mass on cosmological observables is to suppress the formation of cosmic structure on intermediate and small scales. Comparison of the Cosmic Microwave Background (CMB), which reflects large-scale structure at early times, with measurements of the intermediate- or small-scale structure in the local Universe thus provides a way to constrain the absolute mass scale of the neutrino \\citep[e.g.][]{Fukugita00}.  This approach has the disadvantage that any neutrino properties inferred are at some level sensitive to our incomplete understanding of cosmology, in particular dark energy and inflation. It has long been recognized, however, that using complementary measurements of structure, as described above, significantly reduces the sensitivity of the results to such necessary assumptions (e.g. \\citealt*{Allen03a}; \\citealt{Tegmark04}). Nevertheless, it is imperative that any systematic uncertainties affecting the measurements of cosmic structure be properly accounted for in order to obtain robust results.  The amount of structure in the local Universe is generally described by the parameter $\\sigma_8$, defined as the present day root-mean-square fluctuation of the linearly evolved matter density field, smoothed by a spherical top-hat window of comoving radius $8h^{-1}\\Mpc$; here $h=H_0/100\\km\\second^{-1}\\Mpc^{-1}$ is the normalized Hubble parameter. At present, the most robust observation for measuring this quantity is arguably the abundance of massive galaxy clusters;\\footnote{We note that current cluster measurements do not constrain $\\sigma_8$ independent of spectral index of the power spectrum, \\ns{}. However, the best fitting value of $\\sigma_8$ does not vary rapidly with \\ns{}; moreover, \\ns{} is well constrained by CMB data in all cosmological models considered here.} recent advances in cluster simulation \\citep*[e.g.][]{Nagai07}, comparisons of different mass measurement techniques \\citep[e.g.][]{Bradac08,Mahdavi08,Newman09}, and the availability of robust mass proxies (e.g. \\citealt{Allen04,Allen08}; \\citealt*{Kravtsov06}) have significantly reduced systematic uncertainties associated with the measurement of cluster masses. Consequently, recent estimates of $\\sigma_8$ based on independent analyses of galaxy cluster data (both optically and X-ray selected clusters) are in very good agreement [$\\sigma_8 \\sim 0.8$; \\citealt{Mantz08,Mantz09} (hereafter \\cosmopaper{}); \\citealt{Henry09,Vikhlinin09a,Rozo10}]. Recent analyses of cosmic shear data are also in good agreement \\citep{Benjamin07,Fu08}, and provide compatible constraints on the neutrino mass to our own \\citep{Tereno09}. Galaxy redshift surveys have also produced comparable results \\citep*{Thomas09}.  The number of neutrino species participating in weak interactions is known to be three to high precision \\citep{Amsler08}. However, the possibility remains that additional ``sterile'' species exist. Results consistent with the existence of a fourth neutrino were reported by the Liquid Scintillator Neutrino Detector \\citep[LSND;][]{Aguilar01}; these results are disfavored by the MiniBooNE experiment \\citep{MiniBooNE09,MiniBooNE09a}, although the interpretation remains somewhat ambiguous at present. This is another question that cosmological observations can address. In particular, the synthesis of light elements is sensitive to the number of relativistic species present in the early Universe, since these determine the expansion rate at that time; however, observations of primordial deuterium and helium abundances are challenging and are subject to large systematic uncertainties. An independent probe is provided by the CMB, in combination with other cosmological data, as described in \\secref{sec:background}. In this work, we consider only the case where the neutrino species (whatever their number) have approximately degenerate mass; in this case, \\Mnu{} and the effective number of neutrino species, \\Neff{}, are sufficient to describe the cosmological effect of neutrinos. More general mass splittings, and in particular the case favored by the initial LSND results, in which a sterile neutrino is significantly more massive than the other species, require a more complete treatment \\citep[as in, e.g.,][]{Crotty04}.  In this paper, we apply the statistically rigorous analysis method of \\cosmopaper{} and the X-ray flux-limited cluster samples and follow-up observations described in \\citet[][hereafter \\scalingpaper{}]{Mantz09a} to the problem of inferring neutrino properties from cosmological data. These data (collectively referred to here as the cluster X-ray Luminosity Function, or XLF) provide a robust means to measure $\\sigma_8$; our analysis method includes generous systematic allowances and accounts fully for all parameter degeneracies. In obtaining our results, we also incorporate CMB data and measurements of cosmic distance in the form of cluster gas mass fractions (\\fgas{}), type Ia supernova (SNIa) fluxes and Baryon Acoustic Oscillation (BAO) data (see \\secref{sec:data} and references therein). We note that \\citet{Reid09a} obtain very similar results to our own by importance sampling 5-year {\\it Wilkinson Microwave Anisotropy Probe} (\\WMAP{}) results using a prior based on the analysis of optically selected clusters by \\citet{Rozo10}.  The basic cosmology that we consider in this paper is the spatially flat, cosmological constant (\\LCDM{}) model parametrized by the mean baryon density, \\Omegab{}; the mean total matter density including baryons, neutrinos and cold dark matter (CDM), \\Omegam{}; the Hubble parameter, $H_0$; the normalization of the matter power spectrum, $\\sigma_8$; the spectral index of the primordial scalar power spectrum, \\ns{}; and the optical depth to reionization, $\\tau$. Here the mean densities refer to the present day (redshift $z=0$), since their values at other times are then determined by the Friedmann equation, and $\\sigma_8^2$ is the $z=0$ variance in the matter density field at scales of $8h^{-1}\\Mpc$, as defined above. This simple and commonly used model assumes $\\Mnu=0$ and $\\Neff=3.046$, the predicted value for the three weakly interacting neutrinos.  In addition to freeing \\Mnu{} and \\Neff{}, we will consider the effect of marginalizing over other parameters with which they are degenerate. However, since the flat \\LCDM{} model is known to provide a good fit to currently available cosmological data, we are conservative in incorporating these additional degrees of freedom. In particular, we consider generalizing the description of dark energy, either by allowing global spatial curvature or by varying the dark energy equation of state in a flat universe, and including the effects of primordial tensor perturbations. The former case is parameterized by the effective curvature energy density, \\Omegak{}, or the equation of state parameter, $w$, while in the latter case we simultaneously marginalize over the tensor-to-scalar ratio, $r$ (defined with respect to wavenumber $k=0.002h\\Mpc^{-1}$, as in, e.g.,  \\citealt{Komatsu09}), and the tensor power spectral index, \\nt{}.  In our results, we will consistently quote one-sided limits (upper bounds) on \\Mnu{} and $r$ at the 95.4 per cent confidence level and two-sided constraints on all other parameters at 68.3 per cent confidence. ",
        "conclusions": "\\label{sec:conclusion}  We have applied measurements of the growth of massive galaxy clusters (detailed in Papers~I and II), in combination with other cosmological data, to the problem of constraining the species-summed neutrino mass, \\Mnu{}, and effective number, \\Neff{}.  Our results show that a robust measurement of $\\sigma_8$ from clusters significantly improves limits on \\Mnu{} from current data, and reduces their sensitivity to assumptions about the cosmological model. In a simple \\LCDM+\\Mnu{} cosmology, the addition of the XLF data improves the 95.4 per cent confidence upper limit on \\Mnu{} to 0.33\\eV{}, compared with 0.61\\eV{} from the combination of CMB, \\fgas, SNIa and BAO data. In a more general model, marginalized over spatial curvature, primordial tensors and the effective number of neutrinos, and incorporating a prior on the Hubble parameter, the XLF data improve the limit from 1.75\\eV{} to 0.7\\eV{}. The results indicate that this improvement is due entirely to the ability of the cluster data to constrain $\\sigma_8$ (at $z=0$); while measurements of the ($z$-dependent) growth of structure in principle contain additional information about \\Mnu{}, current data are not sufficient to exploit it.  The cluster data additionally improve constraints on the effective number of relativistic species, from $\\Neff=3.6_{-0.6}^{+0.7}$ to $\\Neff=3.4_{-0.5}^{+0.6}$ in a simple \\LCDM+\\Neff{} model (68.3 per cent confidence). For the more general model, including curvature, tensors and neutrino mass, only a marginal improvement is observed.  The results obtained here are compatible with other recent estimates based on galaxy clusters \\citep{Reid09a,Vikhlinin09a}, reflecting the good agreement in recent $\\sigma_8$ constraints based on X-ray and optically selected clusters \\citep[\\cosmopaper{};][]{Henry09,Rozo10}. Although current cosmological data are not sufficient to rule out the existence of sterile neutrino species or distinguish between the normal and inverted mass hierarchies, they continue to provide some of the tightest constraints available, in particular on the mass scale.  We also consider the prospects for further improvement, finding that a few per cent level measurement of the Hubble parameter will significantly improve the constraints on models where \\Neff{} is free. A similarly improved determination of $\\sigma_8$, combined with improved CMB data from {\\it Planck}, will further tighten limits on the neutrino mass scale,\\footnote{We note that the 7-year \\WMAP{} results, which were released while this work was in revision, already offer some improvement. Specifically, \\citet{Komatsu10} find $\\Mnu<0.58\\eV$ by combining 7-year \\WMAP{} data with the \\citet{Riess09} $H_0$ prior and the BAO results of \\citet{Percival10}, a noticeable improvement over the limit $\\Mnu<0.66\\eV$ obtained by \\citet{Sekiguchi09} from the same auxiliary data sets combined with WMAP5.} perhaps providing the first hints of non-zero mass from cosmological data. More precise measurements of the growth of structure at late times, and the extension of such data to higher redshifts, will provide a powerful, new mechanism to constrain neutrino mass."
    },
    "0911/0911.5625_arXiv.txt": {
        "abstract": "Comet 17P/Holmes was observed for linear polarisation using the optical polarimeter mounted on the 1.2m telescope atop Gurushikhar peak near Mt. Abu during the period November-December 2007. Observations  were  conducted through the IHW narrow band (continuum) filters.  During the observing run the phase angle was near $13^{\\circ}$  at which the comet showed negative polarisation. On the basis of the observed polarisation data we find comet 17P/Holmes to be a typical comet with usual dust characteristics. We note that radial rate of change of brightness in coma in red band is higher than that in blue band; it has decreased by a factor of 3.6 and 2.5 respectively in red and blue bands during the November - December run, indicating relative increase in the abundance of smaller dust particles out ward. Radial brightness variation seen near the nucleus on November 6 is indicative of the presence of a blob or shocked region beyond 10\" from the nucleus which has gradually smoothened by December 13. The brightness distribution is found steeper during November 5-7 as compared to on December 13. ",
        "introduction": "Photopolarimetric observations of comet 17P/Holmes were made during the period November 5-7 and on December 13, 2007 with a two channel photo polarimeter \\citep{Deshpande1985,Joshi1987}, which has been fully automated recently \\citep{ganesh2008p}, mounted on the 1.2m telescope of Mt. Abu Observatory operated by Physical Research Laboratory (PRL), Ahmedabad. {\\bf PRL polarimeter works on rapid modulation principle. A rotating  super-achromatic half-wave plate in front of a fixed Wollaston prism modulates the polarized component of the incident light at a frequency of 10Hz. Another identical Panchratnam half-wave plate is put between the rotating half-wave plate and the Wollaston prism to eliminate any wavelength dependence of optic axis of half-wave plate. Rotating half plate rotates at discrete steps with 1 ms sampling time. One full rotation of half-wave plate is completed in 96 step and thus one modulation cycle involves 24 steps (i.e. 24ms) and the data are folded and accumulated in 24 bins. With the rapid modulation, the atmospheric scintillation effects (eg. sky transparency fluctuations) are eliminated. \\\\}  The instrument is equipped with IAU's International Halley Watch ($IHW$)) continuum filters (3650/80\\AA, 4845/65\\AA, 6840/90\\AA~) \\citep{Osborn1990} and Johnson-Cousins' BVR broad band filters. The $IHW$ filters acquired for Comet Halley have been used for these observations.  These filters have been in regular use for observations of several other comets \\citep{Joshi1987,Sen1991,Ganesh1998,Joshi2002, Joshi2003}. The filters have been carefully stored in dry atmospheric conditions to preserve their transmission characteristics. The observations made with the same set of filters facilitate better comparison with other comets observed  earlier,  hence their continued use is justified.\\\\  {\\bf The online data reduction performed after each integration invokes a least square fit to the counts (comet-sky) obtained from the two photo-multiplier tubes as a function of the rotating half-wave plate position. The mean error in polarisation is estimated from actual deviation of the counts from the fitted curve. To further reduce the error, several such measurements were taken and averaged. The error bars represent this error and instrumental error (obtained by observing several zero polarisation standards during each night and  found to be $<0.03\\%$).  Regarding the sky polarisation, on an average it was $\\sim 5-6\\% (\\pm 3\\%)$ but the sky as seen through the 26\" aperture is $\\sim 4$ mag fainter compared to the comet observed through the same aperture. These numbers were consistent during the observing run. Hence, the error communicated due to the sky to the comet data is negligible.} Nonetheless, to take care of the sky polarisation, observations were made alternately centred on the photo Centre of the comet and on a region of the sky more than 30 arcmin away from the comet (along the anti-tail direction).  All the observations were made under dark sky conditions. The errors in the position angle are obtained using the equation 8.5.4 given by \\citet{Serkowski1974}.\\\\  The observations were taken with apertures (non-metallic diaphragms centred on the photo-centre of the comet) of different sizes- 10\", 20\", 26\" and 54\" with the projected diameter varying from 11750 to 69925 km (see Table \\ref{obstab})  to study the behaviour of the dust as a function of radial distance from the comet nucleus. All the observations were made under dark sky conditions, and the  comet was much brighter than the sky resulting in negligible contribution by the sky to the observed degree of polarisation of the comet. Nonetheless, to take care of the sky polarisation, observations were made alternately centred on the photo Centre of the comet and on a region of the sky more than 30 arcmin away from the comet  (along the anti-tail direction).\\\\  Polarisation standard 9 Gem was observed to calibrate the observed position angle. Comet's IHW magnitudes were obtained using the observations of solar type stars, namely HD29461, HD76151.  Polarisation values, corrected position angle and IHW magnitudes in continuum bands are given in Table \\ref{obstab}.  \\begin{table*} \\caption{Polarisation observations of comet 17P/Holmes. Listed entries are Julian date(JDT), Heliocentric range(r), Geocentric range($\\Delta$), phase angle($\\alpha$), filter, aperture(arcsec), Diameter(km), total integration time(IT sec), degree of polarisation($P\\%$), error in polarisation(${\\epsilon}_P\\%$), position angle($\\theta^\\circ$) in equatorial plane, magnitude at the time of observations. JDT=JD-2454400} \\begin{tabular}{|l|l|l|l|l|l|l|l|l|l|l|l|l|} \\hline Date&JDT& r&${\\Delta}$&${\\alpha}$& Filter(\\AA)& Ap(\")&Dia(km)& IT(sec)& $P\\%$ & ${\\epsilon}_p\\%$ &${\\theta^\\circ}$ & mag \\\\ \\hline Nov 5&10.32687 & 2.485602 & 1.620252  & 13.9 & 6840 &  10 & 11751 & 400 &  -1.17 &  0.50 & 19 & 10.05\\\\ &10.33699 & 2.485642 & 1.620252  & 13.9 & 4845 &  10 & 11751 & 300 &  -1.70 &  0.51 &  4 & 11.09\\\\ &10.34764 & 2.485684 & 1.620251  & 13.9 & 3650 &  10 & 11751 & 500 &  -2.72 &  2.14 & 81 & 12.45\\\\ &10.35990 & 2.485732 & 1.620251  & 13.8 & 6840 &  26 & 30553 & 400 &  -1.22 &  0.21 & 24 &  8.24\\\\ &10.36958 & 2.485770 & 1.620252  & 13.8 & 4845 &  26 & 30553 & 400 &  -1.09 &  0.18 & 22 &  9.28\\\\ &10.38014 & 2.485791 & 1.620252  & 13.8 & 3650 &  20 & 23502 & 100 &  -1.25 &  0.76 & 86 & 10.58\\\\ Nov 6&11.28874 & 2.489386 & 1.620319  & 13.6 & 6840 &  26 & 30554 & 500 &  -1.60 &  0.27 & 14 &  8.66\\\\ &11.30002 & 2.489430 & 1.620321  & 13.6 & 4845 &  26 & 30554 & 500 &  -0.95 &  0.19 & 20 &  9.72\\\\ &11.31650 & 2.489495 & 1.620323  & 13.6 & 3650 &  26 & 30554 & 900 &  -1.51 &  0.70 & 32 & 10.98\\\\ &11.33392 & 2.489564 & 1.620327  & 13.6 & 3650 &  10 & 11751 & 400 &  -2.60 &  2.15 & 15 & 12.56\\\\ &11.34823 & 2.489620 & 1.620330  & 13.6 & 4845 &  10 & 11751 & 800 &  -0.96 &  0.43 & 22 & 11.25\\\\ &11.36806 & 2.489698 & 1.620334  & 13.6 & 6840 &  10 & 11751 & 700 &  -1.10 &  0.54 & 43 & 10.16\\\\ &11.38504 & 2.489765 & 1.620338  & 13.6 & 6840 &  20 & 23503 & 600 &  -1.49 &  0.29 & 20 &  9.01\\\\ &11.40166 & 2.489831 & 1.620343  & 13.6 & 4845 &  20 & 23503 & 400 &  -1.36 &  0.27 & 20 & 10.05\\\\ &11.41631 & 2.489888 & 1.620348  & 13.6 & 6840 &  54 & 63460 & 400 &  -1.09 &  0.14 & 13 &  7.21\\\\ Nov 7&12.30185 & 2.493380 & 1.620604  & 13.4 & 6840 &  26 & 30559 & 600 &  -1.60 &  0.23 & 15 &  8.67\\\\ &12.32505 & 2.493472 & 1.620613  & 13.4 & 4845 &  26 & 30560 & 600 &  -1.17 &  0.16 & 17 &  9.73\\\\ &12.33843 & 2.493524 & 1.620618  & 13.4 & 3650 &  26 & 30560 & 500 &  -1.02 &  0.71 & 97 & 11.03\\\\ Dec 13&48.25186 & 2.639940 & 1.784706  & 12.9 & 6840 &  10 & 12944 & 600 &  -3.96 &  1.77 & 97 & 11.69\\\\ &48.26690 & 2.640003 & 1.784838  & 12.9 & 4845 &  10 & 12944 & 600 &  -1.39 &  0.76 & 96 & 12.16\\\\ &48.28473 & 2.640077 & 1.784996  & 12.9 & 3650 &  10 & 12946 & 600 &  -2.88 &  3.46 &142 & 12.98\\\\ &48.29981 & 2.640141 & 1.785130  & 12.9 & 3650 &  26 & 33662 & 300 &  -3.20 &  4.12 &131 & 12.72\\\\ &48.31233 & 2.640193 & 1.785241  & 12.9 & 3650 &  20 & 25895 & 300 &  -3.52 &  3.55 &138 & 12.84\\\\ &48.32681 & 2.640254 & 1.785370  & 12.9 & 3650 &  54 & 69923 & 400 &  -4.56 &  3.53 &107 & 12.49\\\\ &48.33570 & 2.640291 & 1.785449  & 12.9 & 4845 &  54 & 69926 & 300 &  -1.79 &  0.88 &114 & 11.55\\\\ &48.34532 & 2.640331 & 1.785535  & 12.9 & 4845 &  26 & 33669 & 400 &  -1.97 &  0.83 &110 & 11.92\\\\ &48.35515 & 2.640372 & 1.785623  & 12.9 & 4845 &  20 & 25901 & 400 &  -1.60 &  0.79 & 64 & 12.02\\\\  \\hline \\end{tabular} \\label{obstab} \\end{table*} ",
        "conclusions": "  \\begin{itemize}  \\item Like other comets, 17P/Holmes shows negative polarisation at  low phase angle (ie below $20^\\circ$). Based on the polarisation values, we believe that 17P/Holmes is not unusual comet as claimed by \\citet{rosenbush2009} and the dust seems to be of the same nature as found in other comets. The discrepancy could possibly be due to the gas emission contamination in the broad band observations.\\\\  \\item Though there is indication that the value of polarisation decreases towards longer wavelength, we infer that, within the errors, the wavelength dependence of polarisation as observed through different sized apertures is typical of comets.\\\\  \\item  Radial distribution of brightness in the coma shows higher variation with time in the red wave band as compared to that in the blue wave band; it has decreased by a factor of 3.6 and 2.5 respectively in red and blue bands indicating increase in the relative abundance of smaller dust particles away from the nucleus.\\\\  \\item Radial brightness distribution shows the non-uniform distribution of material near the comet nucleus on November 6 which has gradually smoothened by December 13. Also the brightness distribution has become steeper from November 5-7 run to December 13 run.  \\end{itemize}"
    },
    "0911/0911.2469_arXiv.txt": {
        "abstract": "The horizontal branch (HB) morphology of globular clusters (GCs) is most strongly influenced by metallicity. The second parameter phenomenon, first described in the 1960's, acknowledges that metallicity alone is not enough to describe the HB morphology of all GCs. In particular, astronomers noticed that the outer Galactic halo contains GCs with redder HBs at a given metallicity than are found inside the Solar circle. Thus, at least a second parameter was required to characterize HB morphology. While the term `second parameter' has since come to be used in a broader context, its identity with respect to the original problem has not been conclusively determined. Here we analyze the median color difference between the HB and the red giant branch (RGB), hereafter denoted $\\dvi$, measured from Hubble Space Telescope (HST) Advanced Camera for Surveys (ACS) photometry of 60 GCs within $\\sim$20 kpc of the Galactic Center. Analysis of this homogeneous data set reveals that, after the influence of metallicity has been removed from the data, the correlation between $\\dvi$ and age is stronger than that of any other parameter considered. Expanding the sample to include HST ACS and Wide Field Planetary Camera 2 (WFPC2) photometry of the 6 most distant Galactic GCs lends additional support to the correlation between $\\dvi$ and age. This result is robust with respect to the adopted metallicity scale and the method of age determination, but must bear the caveat that high quality, detailed abundance information is not available for a significant fraction of the sample. Furthermore, when a subset of GCs with similar metallicities and ages are considered, a correlation between $\\dvi$ and central luminosity density is exposed. With respect to the existence of GCs with anomalously red HBs at a given metallicity, we conclude that age is the second parameter and central density is most likely the third. Important problems related to HB morphology in GCs, notably multi-modal distributions and faint blue tails, remain to be explained. ",
        "introduction": "}  It has been clear for decades that metallicity is the most influential factor governing the HB morphologies of Galactic GCs: metallicity is the first parameter.  The earliest photographic color-magnitude diagrams (CMDs) of GCs by, e.g.\\,\\citet{ar52,sa53}, revealed that HB stars in metal-rich GCs tend to lie on the red side of the RR Lyrae instability strip while HB stars in metal-poor GCs lie primarily on the blue side. As more and more CMDs were assembled, exceptions to this rule were uncovered. \\citet{sa60} noted that M~13 and M~22 display HB morphologies appropriate for metal-poor GCs, despite the fact both GCs were of intermediate metallicity, and suggested that a difference in age might be responsible. \\citet{sa67} noticed that the HB of NGC~7006 is redder than either M~13 or M~3, despite the fact that all three GCs appeared to have very similar metallicities. Such anomalies suggested the need for a second parameter that could account for differences in HB morphology not obviously caused by metallicity. \\citet{va65} summarized the problem, stating that metallicity is not sufficient to explain the extent of observed HB morphologies but that differences in age or He enrichment could explain the observed variations.  Both of these suggestions remain valid up to the present time. \\citet{va67} analyzed the integrated colors of 49 GCs with $UBV$ photometry and concluded that `at least two parameters (one of which is metal abundance) are required to describe globular clusters.'  The second parameter phenomenon took on a greater significance with the seminal work of \\citet{sz78}.  Searle \\& Zinn recognized that GCs with unusually red HBs are relatively rare in the inner regions of the Galactic halo (Galactocentric radius, $\\rgc \\la 8$ kpc) but become increasingly common at greater $\\rgc$ (see their Figure 10). Searle \\& Zinn used this fact to argue that the inner halo was assembled early and in a fairly short time while the outer halo was assembled over an extended period of time. Current galaxy formation scenarios envision the outer regions of the Galaxy as the accumulated debris of the many accretion events that shaped the early evolution of the Galaxy. In that context, understanding the origin and existence of age and metallicity gradients in the Galactic GC population is as relevant now as it was at the time when Searle \\& Zinn first introduced their halo formation scenario.  Subsequent efforts relating to the second parameter problem fall into two general categories: those that attempt to measure the age difference between two (or a few) GCs with similar metallicities but markedly different HB morphologies and those that investigate the ensemble properties of a large sample of GCs.  One canonical pair in the former category is NGC~288, with a blue HB, and NGC~362, with a red HB. The first CCD-based, differential photometric study of these clusters was conducted by \\citet{bo89}. Bolte concluded that NGC~288 is $\\sim$3 Gyr older than NGC~362. The subsequent work of \\citet{gr90}, \\citet{sa90}, and \\citet{va90} echoed this result. By contrast, \\citet{vd90} reached a very different conclusion. Using the brightest RGB star in each GC, \\citet{vd90} corrected the photometry of NGC~288 and NGC~362 for their relative distances and found their main sequence turnoffs to be roughly coincident, thereby suggesting a negligible age difference. \\citet{be01} found NGC~288 to be 2$\\pm$1 Gyr older than NGC~362 using three different techniques; in a follow-up, \\citet{ca01} found that an age difference of 2 Gyr was plausible if both GCs have $\\feh \\sim -1.2$ but that their synthetic HB models were unable to match the detailed HB morphology of either GC using canonical assumptions of average mass loss and dispersion on the HB. In considering differences between second parameter pairs, it is important to understand just how similar the abundances are in both clusters. \\citet{sh00} performed a detailed comparison of abundances in red giants with 13 stars in NGC~288 and 12 in NGC~362. These authors concluded that NGC~288 has a lower $\\feh$ than NGC~362 by 0.06 dex, that the average $\\afe$ ratios are very similar, and that both GCs exhibit variations among O, Na, and Al.  The other canonical second parameter pair is M~13, with a blue HB, and M~3, with an intermediate HB.  M~13 and M~3 are $\\sim$0.2 dex more metal poor than NGC~288 and NGC~362. \\citet{va90} found evidence for an age difference but the quality of their CMDs did not allow a more definitive statement. \\citet{ca95} estimated an upper limit to the age difference of $\\sim$3 Gyr and concluded that it was insufficient to explain the HB morphologies, assuming both GCs have the same chemical composition. Alternatively, \\citet{jo98} suggested the difference in their HB morphologies was due to a difference in their He abundances. Johnson \\& Bolte suggested that M~13 had a He mass fraction $\\sim$0.05 greater than M~3. However, \\citet{sw98} showed that a difference of $\\Delta$Y $\\sim 0.05$ at constant $\\feh$ would make the level of the HB brighter by $\\sim$0.2 mag. Such a difference between M~3 and M~13 was ruled out by the photometry of \\citet{re01}, who also concluded that M~13 is older than M~3 by $1.7\\pm0.7$ Gyr. \\citet{sn04} compared spectra of 28 red giants in M~3 with 35 in M~13 and found the two GCs' mean $\\feh$ values the same within their 1-$\\sigma$ error bars. However, \\citet{sn04} reported differences in light element abundance variations, particularly O, for which M~13 displays a larger range of variation than M~3 by $\\sim$0.5 dex.  A discussion of the search for the second parameter in GC-to-GC comparisons would not be complete without including the work of \\citet{st99} and \\citet{do08b}. These investigations used HST WFPC2 photometry to measure the ages of the outer halo GCs Palomar~3, Palomar~4, and Eridanus \\citep{st99} and AM-1 and Palomar~14 \\citep{do08b} relative to inner halo GCs with similar metallicities, M3 and M5\\footnote{There is substantial evidence that M~5 is actually an outer halo GC currently near its perigalacticon, see \\citet{sc96} and \\citet{di99}.}. Each of the five outer halo GCs has a redder HB morphology than its comparison inner halo GC. Both studies concluded that the outer halo GCs are $\\sim$1.5-2 Gyr younger than M3 and M5, provided the chemical compositions of the outer halo GCs are comparable to their inner halo counterparts. A recent study by \\citet{ko09} found the abundances of Pal~3 are essentially the same as inner halo GCs of similar metallicity and thus the relative age comparison with M~3 by \\citet{st99} was justified. Overall, the outer halo GCs' chemical compositions remain poorly understood by comparison with the inner halo, see e.g.\\,\\citet{pvi05}. \\citet{ca00} found the reported age differences between Pal~4 and Eridanus, with red HBs, and M~5, with an intermediate HB, too small to explain the difference in HB morphologies unless all three GCs are younger than 10 Gyr, assuming standard assumptions of mass loss on the RGB.  \\citet{ca01b} concluded that it was possible to explain the difference in HB morphology between Pal~3 and M~3 if the former has less HB mass dispersion and is younger than the latter. These examples indicate that if age is the second parameter, then it is our lack of understanding of mass loss that confuses GC-to-GC comparisons as first pointed out by \\citet{ro73}. The review by \\citet{ca05} includes a thorough discussion of different mass loss prescriptions and their efficacy. Given the HB morphologies and relatively young ages of the most distant outer halo GCs\\footnote{The exception in the outer halo is NGC~2419. A number of dedicated HST photometric studies have focused on this massive, distant GC. \\citet{ha97} and \\citet{sa08} both found NGC~2419 to be an old, metal poor GC with a blue HB.}, it is important to include them when considering the properties of the entire Galactic GC population, especially in the context of Searle \\& Zinn's halo formation scenario.  The second category of second parameter studies are those that consider the properties of a large sample of Galactic GCs, e.g.\\,\\citet{sz78}. In the first CCD-era study, \\citet{sa89} compiled properties for 31 Galactic GCs. Among other things, \\citet{sa89} examined the variation of HB morphology with age for GCs in a narrow range of metal abundance and showed that GCs with red HBs are significantly younger than those with blue HBs by as much as 5 Gyr. The subsequent studies of \\citet{ch92}, \\citet{sa95}, and \\citet{ch96} updated and reaffirmed this result. \\citet{ro99} found that the GCs with $\\rgc > 8$ kpc are, on average, younger than those with $\\rgc < 8$ kpc, which is consistent with age as the second parameter. By contrast, \\citet{ri96} examined the CMDs of 36 GCs, found an age dispersion of $\\sim$1 Gyr with no significant age gradient in the Galactic halo, and concluded that this age range was too small for metallicity and age alone to explain HB morphology in GCs.  Theoretical efforts have provided further insights into the complexities of HB morphology. In particular, several studies have applied the synthetic HB model developed by \\citet{ro73}. \\citet{ldz90}, \\citet{ca93}, and \\citet{ldz94} used synthetic HB models to explore the interplay of HB morphology, metallicity, and age in the CMD. These studies demonstrated considerable degeneracy in the HB morphology--metallicity diagram and, in particular, \\citet{ca93} argued that unless absolute values of the He abundance, $\\afe$ ratios, and RGB mass loss were known, HB morphology did not constitute a reliable age constraint.  Nevertheless, \\citet{ldz94} concluded that there was evidence for an age dispersion of $\\la 5$ Gyr among Galactic GCs of similar metallicity but markedly different HB morphologies. A number of theoretical studies have also focused on second parameter pairs or triads but, unfortunately, the lack of a firm theoretical understanding of mass loss along the RGB has impaired the these efforts, as first pointed out by \\citet{ro73}. See \\citet{ca05} for a recent review.  As it pertains to HB morphology, the use of the term 'second parameter' has expanded to cover a broad range of complex behaviors. In particular, this includes the faint extent of the blue HB tail in some GCs. \\citet{fp93} analyzed 53 GCs and found that the length of the HB in the CMD--and the extent of the blue tail--is correlated with central density.  \\citet{bu97} also reported a link between an extended blue tail and central density in GCs and concluded that `environment is ``a'' second parameter'. \\citet{sm04} showed that there is a correlation between central density and HB morphology for intermediate metallicity GCs ($-1.7 < \\feh < -1.3$) where the second parameter effect is most pronounced. \\citet{rb06} analyzed the properties of 54 Galactic GCs with homogeneous photometry from an HST WFPC2 Snapshot Survey \\citep{pi02} and concluded that more massive GCs tend to have more extended blue HBs. \\citet{rb06} recognized a link between the effective temperature of the hottest HB star in a GC and its mass (as inferred from its integrated luminosity) and suggested that self-pollution could explain the existence of faint blue tails in preferentially massive GCs.  Evidence abounds for the presence of chemical abundance variations in all GCs \\citep{gr04,car09a,car09b} and the possible correlation between the degree of abundance variations and the faint extent of the blue HB \\citep{car07}. However, it is unlikely that the existence of faint, blue tails in the HBs of some GCs is directly related to the appearance of anomalously red HBs in others. It is also unlikely that chemical abundance variations will unduly affect age estimates of most GCs, provided the total metal content is constant across all stars \\citep{pie09}. For cases in which the total metal abundance is not constant, or there are distinct stellar populations present in the CMD, age estimates are necessarily more complicated and careful analysis of each of these GCs is needed. See \\citet{pi09} for a recent summary of GCs with multiple stellar populations visible in the CMD, several of which were discovered with photometry from the ACS Survey of Galactic GCs. The complex issue of multiple stellar populations in GCs remains poorly understood and the extent to which multiple-population GCs permeate the Galactic GC population is unknown at present.  To summarize, the presence of metal poor GCs with red HBs that primarily reside in the outer Galactic halo is well-known observationally. Although the present study is focused on the Galactic GCs, including some of those associated with the Sagittarius dwarf, there is ample evidence to suggest that metal poor GCs with red HBs are also found in the Magellanic Clouds \\citep{jo99,gl08} and Fornax \\citep{bu99}. \\citet{mg04} summarize our current knowledge of the GC populations in the Galaxy and its satellites in the HB morphology--metallicity diagram. Despite an abundance of observational evidence, no consensus has been reached as to what parameter(s) are responsible for the appearance of relatively red HBs in metal poor GCs. Age is frequently offered as a candidate but considerable doubt still remains because the age difference claimed--or required by theoretical studies of HB morphology--is too large to satisfy the observations.  The existence of a homogeneous database of deep, high quality photometry from the ACS Survey of Galactic GCs \\citep[Paper I in this series]{sa07} has motivated a re-examination of HB morphology and its relation to a variety of GC properties. The paper is organized as follows: $\\S$\\ref{data} describes the data sample;  $\\S$\\ref{hbm} describes the methods that were employed to determine the HB morphologies; $\\S$\\ref{gcages} describes the sources of GC ages and provides some discussion of complicating factors in age determination; $\\S$\\ref{results} presents the analysis performed on the assembled data and discusses the important results; and, finally, $\\S$\\ref{conclusion} provides a summary of the salient points. ",
        "conclusions": "} HB morphologies characterized by $\\dvi$ of 66 Galactic GCs were examined to determine the sensitivity of HB morphology to a variety of different factors. Deep, homogeneous photometry from the ACS Survey of Galactic GCs accounts for 60 of these while the remaining 6 GCs are the most distant Galactic GCs known. The complete sample is the largest examined to date and consists solely of high-quality HST photometry. It spans the full range of $\\rgc$ and almost the entire range of metallicity present in the Galactic GC population. $\\dvi$ values and two independently measured sets of ages were joined with other quantities from the literature to assess the relative importance of these quantities on HB morphology.  The data were split into two groups, roughly equal in number, based on $\\rgc$. The tight relationship between metallicity and HB morphology in the inner halo group ($\\rgc < 8$ kpc) was characterized by a fitting function and this trend was subtracted off of the outer halo group. The difference between fit and measured $\\dvi$ was then compared to a variety of parameters, of which only age showed a significant correlation. The age correlation does not rely on the presence of the 6 most distant GCs in the analysis, though their presence does strengthen the result.  Hence we conclude that, after metallicity, age has the most influence on $\\dvi$.  The age spread among the bulk of GCs in the Galactic halo was found to be 2-2.5 Gyr, though there are a few younger outliers such as Pal~12 and Ter~7. Further analysis, in which both metallicity and age were restricted, provided strong evidence that central luminosity density ($\\rho_0$) is the third most influential parameter on $\\dvi$."
    },
    "0911/0911.3659_arXiv.txt": {
        "abstract": "We derive a relation for the steepening of blazar $\\gamma$-ray spectra between  the multi-GeV  {\\it Fermi}  energy range  and the  TeV energy range observed by atmospheric  \\v{C}erenkov telescopes.  The change in spectral  index   is  produced  by  two  effects:   (1)  an  intrinsic steepening,  independent  of  redshift,  owing to  the  properties  of emission and  absorption in the  source, and (2)  a redshift-dependent steepening produced  by intergalactic pair  production interactions of blazar $\\gamma$-rays  with low  energy photons of  the ``intergalactic background light''  (IBL). Given this relation, with  good enough data on the mean \\gray SED of  TeV Selected BL Lacs, the redshift evolution of the IBL  can, in principle, be determined  independently of stellar evolution models.   We apply our relation  to the results  of new {\\it Fermi} observations of TeV selected blazars. ",
        "introduction": "Stecker \\&  Scully (2006) (SS06) derived a  simple analytic expression for the change in spectral index of a TeV \\gray source in the redshift range  between 0.05  and  0.4.  They  showed  that the  change in  the spectral  index caused  by  intergalactic absorption  is  given by  an approximately  linear  relation   in  redshift,  {\\it  i.e.},  $\\Delta \\Gamma_{a} \\simeq C + Dz$.  The purpose of this letter  is to generalze this relation by including the effect of  intrinsic steepening in the source  spectra between the mutli-GeV  energy range  observed by  {\\it  Fermi} and  the TeV  range observed by atmospheric  \\v{C}herenkov telescopes.  Our general result is  roughly independent  of the  specific model  of  the intergalactic background light (IBL)  used, because it only depends  on the shape of the average galaxy  spectral energy distribution (SED) on  the near IR side of the  starlight peak that determines the  absorption in the TeV energy range.  We compare our  relation for the specific baseline  and fast evolution models of Stecker, Malkan \\& Scully (2006) (SMS06) with the results of recent {\\it Fermi} observations of 13 TeV selected BL Lac AGN. We also show  how it  can  be  used to  independently  determine the  redshift evolution of the IBL. ",
        "conclusions": "We have  derived a simple  analytic approximation for  determining the steepening in the spectra of TeV  selected BL Lac AGN and compared our results with recent observational data  on 13 TeV selected BL Lac AGN. Our relation is in excellent agreement with the observational data and provides a framework for understanding and interpreting both these and observational data.  SS06 have  shown that  the effect of  intergalactic absorption  on the spectra of AGN in the energy range 0.2 TeV $ < E_{\\gamma} <$ 2 TeV and the redshift range $0.05 < z < 0.4$ is a simple power-law to power-law steepening.  Absorption in this  energy range is primarily produced by interactions with near  infrared photons from on the  low energy sides of  the  starlight peaks  in  galaxy SEDs.   The  {\\it  shape} of  the resulting  peak in  the  intergalactic SED  produces an  approximately logarithmic  energy dependence for  the function  $\\tau (E_{\\gamma})$. This energy  dependence should be roughly  the same for  all models of the IBL.   Therefore, our general  result of a linear  $\\Delta\\Gamma = (C+K) + Dz$ relationship is  roughly independent of the specific IBL model used  because  it only depends on the shape of the average  galaxy SED on the near IR side of the IBL starlight  peak. The parameters $C,  D$ and $K$  will be different for the  different models  since the absolute  value of  $\\tau$ varies from one IBL model to another.  Given our derived relation between $\\Delta \\Gamma$ and redshift, and with good enough data on the mean \\gray SED of TeV Selected BL Lacs used to determine $K$, the redshift evolution of the IBL can, in principle, be determined independently of stellar evolution models."
    },
    "0911/0911.3974_arXiv.txt": {
        "abstract": "The SPace Infrared telescope for Cosmology and Astrophysics (SPICA) is a proposed mid-to-far infrared (4-200 $\\mu$m) astronomy mission, scheduled for launch in 2017.  A single, 3.5m aperture telescope would provide superior image quality at 5-200 $\\mu$m, and its very cold ($\\sim$5 K) instrumentation would provide superior sensitivity in the 25-200 $\\mu$m wavelength regimes. This would provide a breakthrough opportunity for studies of exoplanets, protoplanetary and debris disk, and small solar system bodies. This paper summarizes the potential scientific impacts for the proposed instrumentation. ",
        "introduction": "The SPace Infrared telescope for Cosmology and Astrophysics (SPICA) is a proposed mission for mid-to-far infrared (MIR/FIR) astronomy, consisting of a single 3.5-m aperture space telescope with cooled ($\\sim$5 K) instrumentation (see Nakagawa 2008). This mission would provide a significant step forward in detection sensitivity in the 4 to 200 $\\mu$m wavelength regime, which would revolutionize our understanding of how galaxies, stars and planets form, and how interactions between complex astrophysical processes have ultimately led to the formation of our own Solar system and the emergence of life on Earth.  The role of SPICA would complement other forthcoming space infrared missions; Herschel and JWST.  Herschel will provide improved capabilities at wavelengths longer than 57 $\\mu$m, while its relatively warm temperature will produce a modest improvement over the Spitzer Space Telescope.  JWST will provide the best performance at wavelengths shorter than 25 $\\mu$m, and will have the highest angular resolution. SPICA would (1) achieve the best sensitivity at 25--200 $\\mu$m, due to its cold mirror ($\\sim$5 K); (2) achieve a clean point-spread function at 5--200 $\\mu$m due to its non-segmented mirror, thereby providing the best performance for high-contrast coronagraphy; and (3) fill the gap in wavelength coverage between Herschel and JWST.  SPICA would offer a unique opportunity for studying exoplanets, planet formation, circumstellar disks, and small bodies in the solar system. Table 1 summarizes the instruments proposed to date. Those selected for the mission well be chosen based on technical constraints (weight, volume, heat dissipation etc.) and scientific impact.  In \\S 2, we describe potential scientific achievements in the above areas of research.  In \\S 3 we briefly describe the project schedule, including the selection process for the instruments. ",
        "conclusions": "SPICA would provide significant advances in the study of exoplanets, protoplanetary disks, debris disks and small bodies in the solar system.  Its imaging capability would be able to detect debris disks with an unprecedented sensitivity. These capabilities would allow to measure the total flux of small bodies in the solar system, useful for the study of their size distribution and albedos. SPICA's spectral capability would probe the gas and dust associated with circumstellar disks, and the atmospheres of the transiting planets. Its coronagraphic capability would lead to the discovery of a number of exoplanets, and the subsequent investigation of their atmospheres and the geometry and distribution of ice in the spatially resolved circumstellar disks. Possible achievements with the individual instruments are summarized in Fig 4.  The conceptual design of SPICA is currently underway, and we will begin the definition phase of the mission in 2009. This phase, which will include the final decisions for the instruments, will be completed in 2011. \\\\ \\\\ {\\it Acknowledgement} --- We are grateful to Dr. Yuri Aikawa and Akira Kouchi for useful discussions. We thank Dr. Jennifer Karr for reading our manuscript, and the editor and anonymous referee for useful comments."
    },
    "0911/0911.3145_arXiv.txt": {
        "abstract": "From the analysis of measurements of the linear polarisation of visible light coming from quasars, the existence of large-scale coherent orientations of quasar polarisation vectors in some regions of the sky has been reported. Here, we show that this can be explained by the mixing of the incoming photons with nearly massless pseudoscalar (axion-like) particles in extragalactic magnetic fields. We present a new treatment in terms of wave packets and discuss its implications for the circular polarisation. ",
        "introduction": "In this work~\\cite{paper}, we are looking at the effect that axion-photon mixing can have on the polarisation of light emitted by distant astronomical sources. In particular, the observation of redshift-dependent large-scale coherent orientations of AGN polarisation vectors~\\cite{hutsemekers} can, at least qualitatively, be reproduced as a result of such a mixing of incoming photons with extremely light axion-like particles as they travel inside external magnetic fields. These observations (see Figure~\\ref{fig:hutsemekers}) are based on good-quality measurements in visible light of the linear polarisation for a sample of 355 quasars from different groups, working on different instruments.  \\begin{center} \\begin{figure}[h] \\includegraphics[trim = 10mm 10mm 0mm 0mm, clip, width=0.8\\textwidth]{fig_map_nc.ps} \\caption{Maps of the same region of the sky in right ascension ($x$ axis) and declination ($y$ axis), both given in degrees, for AGN characterised by different redshifts, $z$. These polarisation vectors are thus for objects approximately along the same line of sight but at different distances from us. In the low-$z$ case, the average direction is $\\bar{\\theta}\\approx 79^\\circ$ while at high-$z$, it would be $\\bar{\\theta}\\approx 8^\\circ$ ($\\theta$ being counted from North to East)~\\cite{hutsemekers}. } \\label{fig:hutsemekers} \\end{figure} \\end{center}  This has been discussed in terms of axion-photon mixing by several authors, in the case of plane waves, and a prediction from this mixing, in this particular case and in general, is a circular polarisation comparable to the linear one and, hence, observable~\\cite{planewaves,reviewHLPW}. Here, we present a treatment in which light is described by wave packets and show that the circular polarisation can be suppressed with respect to that which is predicted in the plane wave case. ",
        "conclusions": "We have reviewed axion-photon mixing in the case of plane waves and have briefly presented our new formalism in terms of wave packets. The main consequence of this new treatment is the net decrease of circular polarisation with respect to what is predicted using plane waves. From this we conclude that the lack of circular polarisation in the light from AGN does not rule out the ALP-photon mixing.    \\begin{theacknowledgments} A.~P. would like to thank the IISN for funding and to acknowledge constructive discussions on physical and numerical matters with Fredrik Sandin and Davide Mancusi. D.~H. is senior research associate FNRS. \\end{theacknowledgments}"
    },
    "0911/0911.3235_arXiv.txt": {
        "abstract": "Motivated by recent observations from Pamela, Fermi and H.E.S.S., we consider dark matter decays in the framework of supersymmetric SU(5) grand unification theories. An SU(5) singlet $S$ is assumed to be the main component of dark matters, which decays into visible particles through dimension six operators suppressed by the grand unification scale. Under certain conditions, $S$ decays dominantly into a pair of sleptons with universal coupling for all generations. Subsequently, electrons and positrons are produced from cascade decays of these sleptons. These cascade decay chains smooth the $e^{+}+e^{-}$ spectrum, which permit naturally a good fit to the Fermi LAT data. The observed positron fraction upturn by PAMELA can be reproduced simultaneously. We have also calculated diffuse gamma-ray spectra due to the $e^{\\pm}$ excesses and compared them with the preliminary Fermi LAT data from 0.1 GeV to 10 GeV in the region $0^{\\circ}\\leq l\\leq360^{\\circ},10^{\\circ}\\leq|b|\\leq20^{\\circ}$. The photon spectrum of energy above 100 GeV, mainly from final state radiations, may be checked in the near future. ",
        "introduction": "Electron, proton, photon, neutrino and their antiparticles are stable, at least on the cosmological time scale. Detection of these particles from cosmic rays provides an interesting window to look into the deep universe. Recently, the PAMELA experiment reported a significant excess in the positron fraction $e^{+}/(e^{+}+e^{-})$ between $10$ GeV and $100$ GeV \\cite{Adriani:2008zr}. On the other hand, the measured antiproton to proton flux ratio appears to be consistent with predictions \\cite{Adriani:2008zq}. More recently, the Fermi LAT collaboration observed a smooth $e^{+}+e^{-}$ spectrum with high accuracy. It is found to be falling as $E^{-3.0}$ from $20$ GeV to $1$ TeV \\cite{Abdo:2009zk}, much harder than the predictions of conventional models. The H.E.S.S. collaboration measured the $e^{+}+e^{-}$ spectrum from $600$ GeV up to several TeV \\cite{Aharonian:2009ah}, which is consistent with the Fermi data in overlapping regions and steepens at about $1$ TeV towards higher energy.  These excesses of electrons and positrons could be due to unidentified astrophysical sources, e.g., nearby pulsars or supernova remnants \\cite{Stawarz:2009ig,Piran:2009tx,Grasso:2009ma}. However, an explanation via dark matter (DM) annihilation or decay is, arguably, a much more interesting possibility, at least from the perspective of particle physics. The electron and positron spectra alone, even with higher precision and broader energy range, cannot decisively decide which explanation is more plausible \\cite{Malyshev:2009tw}. Hopefully, the energy spectrum and the angular dependence of cosmic gamma rays \\cite{Bertone:2008xr,Zhang:2008tb,Bergstrom:2008ag,Ibarra:2009nw,Zhang:2009kp}, to be measured by the Fermi LAT in the near future, may provide a more definite answer. For the DM interpretation, the mass of the DM should be around several TeV, to provide the $e^{\\pm}$ excesses from $20$ GeV to $1$ TeV and steepen sharply above $1$ TeV. Furthermore, traditional WIMP DM candidates usually produce extra antiprotons. As Pamela does not observe any deviation on antiproton spectrum from the anticipation, WIMP DMs are now disfavored as potential sources of the observed cosmic-ray excesses. Still, there are plenty of freedoms for both DM annihilation and decay to reproduce the experimental $e^{\\pm}$ spectra reasonably \\cite{Meade:2009iu,Bergstrom:2009fa}. For DM annihilation, a large boost factor in the order of $10^{2}$ to $10^{3}$ is needed for the theory to be consistent with the relic abundance measured by the WMAP \\cite{Dunkley:2008ie}. As the clumpiness property of the DM distribution falls far short of such a large factor, one usually resorts to nonperturbative Sommerfeld \\cite{Hisano:2004ds,Hisano:2006nn,Cirelli:2007xd,ArkaniHamed:2008qn} or Breit\\textendash{}Wigner \\cite{Feldman:2008xs,Ibe:2008ye,Guo:2009aj} enhancement in model buildings. For DM decays, the lifetime should typically be around the order of $10^{26}s$ to fit the $e^{\\pm}$ data\\cite{Meade:2009iu,Shirai:2009fq,Mardon:2009gw}, which is much longer than the lifetime of the universe. Therefore the DM decay rates will not affect the relic abundance appreciably.  The energetic $e^{\\pm}$ flux produced from DM annihilations/decays would inevitably emit gamma rays. These gamma rays depend on the DM density as $\\rho^{2}$ for annihilations and $\\rho$ for decays. This will lead to different angular dependence of the gamma ray spectrum, which may be measurable in the near future to differentiate these two scenarios. The gamma ray spectrum can also be used to differentiate DM explanations from astrophysical ones.  In this paper, we will focused on DM decays. Notice that a lot of suppression will be needed for a TeV scale particle to have a lifetime $\\sim10^{26}s$. If it decays via dimension four operators, tremendous fine tunings will be needed. If it decays via dimension six operators, it still needs to be suppressed by a scale $\\sim10^{16}$ GeV, which turns out to coincide with the grand unification theory (GUT) scale \\cite{Langacker:1991an,Amaldi:1991cn,Ellis:1990wk}. In the same spirit of Refs. \\cite{Arvanitaki:2008hq,Arvanitaki:2009yb}, we will take a singlet as the dark matter candidate and provide a detailed analysis in the frame of supersymmetric SU(5) GUT. To be consistent with the Pamela antiproton measurement, squark masses are assumed to be heavier than that of the SU(5) singlet $S$, so the $S$ decay would be quark phobic. The $S$ then decays dominantly into slepton pairs with a universal coupling for all generations. These sleptons decay quickly into leptons and lightest supersymmetric particles (LSPs), if R-parity is conserved. In this framework, we have obtained a reasonable fit to all $e^{\\pm}$ data from Pamela, Fermi and H.E.S.S..  The $e^{\\pm}$ fluxes from $S$ decays are inevitably accompanied by hard photons: coming from final state radiations (FSR) of cascade decays, including $S\\to\\tilde{\\tau}\\to\\tau\\to\\pi^{0}\\to2\\gamma$ and the inverse Compton scattering (ICS) on the interstellar radiation field (ISRF). The gamma ray fluxes could have Galactic and extragalactic origins. We have calculated all these gamma ray spectra and compared them with the recent Fermi LAT measurement in the region $0^{\\circ}\\leq l\\leq360^{\\circ},10^{\\circ}\\leq|b|\\leq20^{\\circ}$ \\cite{Porter:2009sg}.  This paper is organized as follows. The supersymmetric SU(5) model plus a singlet $S$ is presented in Section II, where we have also discussed the possible decay channels of $S$ in some detail. In Section III, a reasonable fit is obtained to reproduce the observed $e^{\\pm}$ fluxes, by tuning relevant parameters in the model. Section IV is devoted to the study of gamma-ray spectra from $e^{\\pm}$ excesses. Finally we conclude with a summary in section V. The component field structure of the dimension six effective operators will be presented in the Appendix. In this paper, we have used the Navarro-Frenk-White (NFW) halo model \\cite{Navarro:1996gj} for DM distribution and the MED propagation model \\cite{Donato:2003xg,Delahaye:2007fr}. For other halo and propagation models, the conclusions are similar. In addition, all computations on astrophysical effects are performed semi-analytically instead of using the GALPROP program% \\footnote{Web page: http://galprop.stanford.edu/web\\_galprop/galprop\\_home.html% }. ",
        "conclusions": "In this paper we have studied the DM decay in supersymmetric SU(5) models. An SU(5) singlet $S$, instead of LSP, is assumed to be the dominant component of DM. With R-parity conservation and a spontaneously broken $Z_{2}$ symmetry, the singlet $S$ can decay into visible particles through dimension six effective operators suppressed by the GUT scale. Assuming the squarks to be heavier than $S$, $S$ decays dominantly into a pair of sleptons through the effective operator $\\widetilde{s}^{*}(\\widetilde{l_{L}}^{*}\\square\\widetilde{l_{L}}+\\widetilde{l_{R}}\\square\\widetilde{l_{R}}^{*})$. Typically, the lifetime of $S$ is around $10^{26}$ s, much longer than the age of the Universe. Since the decay products of $S$ do not contain any quarks, our model is consistent with the Pamela antiproton measurement automatically. For illustration, we have chosen $M^{DM}=6.5$ TeV, $M_{GUT}=10^{16}$ GeV, $M_{\\widetilde{e}}=380$ GeV, $M_{\\tilde{\\mu}}=370$ GeV, $M_{\\widetilde{\\tau}}=330$ GeV and $M_{LSP}=300$ GeV. With this parameter set, we have shown that the $S$ decays can account for the PAMELA, H.E.S.S. and Fermi LAT $e^{\\pm}$ excesses. Numerically, we have adopted the NFW profile for the dark matter distribution and the MED propagation model for the cosmic ray propagation. To interpret these results, one should keep in mind that there exist substantial astrophysical uncertainties about $e^{\\pm}$ background, $e^{\\pm}$ propagation and the DM distribution.  When $e^{\\pm}$ propagate from the decay position to the Earth, hard photons are emitted inevitable due to inverse Compton scattering and final state radiation. Future measurements of the diffuse gamma ray may distinguish DM explanations from astrophysical explanations by looking into the energy and angular distributions. We have calculated the gamma ray spectra in our model. The predicted photon spectrum are compared to the preliminary Fermi LAT measurement from $0.1$ GeV to $10$ GeV in the region $0^{\\circ}\\leq l\\leq360^{\\circ},~10^{\\circ}\\leq|b|\\leq20^{\\circ}$, which seems to be consistent with each other. The total gamma ray spectrum are dominated by photons from Galactic final state radiation for the photon energy above $100$ GeV, which may be tested by Fermi LAT in the near future."
    },
    "0911/0911.0416_arXiv.txt": {
        "abstract": "We present three-dimensional simulations on a new mechanism for the detonation of a sub-Chandrasekhar CO white dwarf in a dynamically unstable system where the secondary is either a pure He white dwarf or a He/CO hybrid. For dynamically unstable systems where the accretion stream directly impacts the surface of the primary, the final tens of orbits can have mass accretion rates that range from $10^{-5}$ to $10^{-3} M_{\\odot}$ s$^{-1}$, leading to the rapid accumulation of helium on the surface of the primary. After $\\sim 10^{-2}$ $M_{\\odot}$ of helium has been accreted, the ram pressure of the hot helium torus can deflect the accretion stream such that the stream no longer directly impacts the surface. The velocity difference between the stream and the torus produces shearing which seeds large-scale Kelvin-Helmholtz instabilities along the interface between the two regions. These instabilities eventually grow into dense knots of material that periodically strike the surface of the primary, adiabatically compressing the underlying helium torus. If the temperature of the compressed material is raised above a critical temperature, the timescale for triple-$\\alpha$ reactions becomes comparable to the dynamical timescale, leading to the detonation of the primary's helium envelope. This detonation drives shockwaves into the primary which tend to concentrate at one or more focal points within the primary's CO core. If a relatively small amount of mass is raised above a critical temperature and density at these focal points, the CO core may itself be detonated. ",
        "introduction": "White dwarfs (WDs), the end point of stellar evolution for most stars, are extremely common, with about $10^{10}$ of them residing within the Milky Way \\citep{Napiwotzki:2009p3570}. WDs are frequently observed in binary systems with normal stellar companions, and, less frequently, with compact stellar companions. Double degenerate (DD) systems are those in which the companion is another WD. Typically, DD systems are formed via common envelope evolution \\citep{Nelemans:2001p1711}. Often, the final result of this evolutionary process is a binary consisting of a carbon-oxygen (CO) WD and a lower-mass helium WD companion \\citep{Napiwotzki:2007p1710}.  For decades, Type Ia supernovae (SNe) have been employed as standard candles. Still, even the preferred mechanism exhibits substantial variability \\citep{Timmes:2003p2073, Kasen:2009p3427}. Complicating this issue further is that WDs may not need to exceed the Chandrasekhar limit to be capable of exploding; other mechanisms include collisions between WDs in dense stellar environments such as the cores of globular clusters \\citep{Rosswog:2009p3556}, tidal encounters with moderately massive black holes \\citep{Rosswog:2008p3059, Rosswog:2009p3553, RamirezRuiz:2009p3071}, or as we are introducing in this paper, the accretion of dense material from a He WD companion.  At the distance of tens of WD radii, gravitational radiation can bring two WDs close enough for the less-massive WD (secondary) to overfill its Roche lobe and begin transferring mass to the more-massive WD (primary). Because of the mass-radius relationship of WDs, mass transfer is often unstable and can eventually lead to a merger \\citep{Marsh:2004p1880, Gokhale:2007p2850}. For binary mass ratios close to unity the circularization radius $R_{\\rm h}$ drops below the radius of the primary $R_{1}$, and thus the accretion stream will directly impact the primary's surface.  Most studies of dynamically unstable DD systems focus on the evolution of the post-merger object and, assuming that the final object has a mass larger than the Chandrasekhar limit, on how the merged remnant may eventually lead to a Type Ia SNe \\citep{Livio:2000p3540, Yoon:2007p1003}. These models all assume that the rapid accretion that precedes coalescence is uneventful. In this \\textit{Letter}, we present three-dimensional hydrodynamics simulations that demonstrate that explosive phenomena are an inevitable results of the high accretion rates characteristic of the final stages of a merger in a DD system. In some cases, these explosive phenomena may lead to the complete detonation of the CO primary.  \\begin{figure*}[tb] \\centering\\includegraphics[width=0.5\\linewidth,clip=true,angle=-90]{fig1.eps} \\caption{Setup of run Ba in FLASH, with several annotated regions of interest. The color scheme shows $\\log T$ through a slice of the orbital plane. The cyan circle shows the spherical outflow boundary condition centered about the secondary, while the cyan wedge shows a cross-section of the cone used as the mass inflow boundary. The dashed contour shows the system's Roche surface. The white boxes with labels are described in Section \\ref{accinsta}.} \\label{slicefig} \\end{figure*} ",
        "conclusions": "\\label{results}  Surface detonations are present in both of the longer runs with a 0.9 $M_{\\odot}$ CO primary, although run B detonates much later in its evolution than run C. The geometry and evolution of the detonations are similar in both runs. A small region of the surface helium is heated to a temperature of $\\sim 2 \\times 10^{9}$ K, leading to a detonation front that expands outwards from the ignition site along the primary's surface (Figure \\ref{volfig}, second column). Because the helium layer is toroidal rather than spherical, the front runs out of fuel as it propagates towards the primary's poles. Consequently, the detonation splits into two fronts that run clockwise and counterclockwise around the equator of the primary (Figure \\ref{volfig}, third column). These detonation fronts eventually run into each other along a longitudinal line that is opposite to the original ignition site (Figure \\ref{volfig}, fourth column). The highest temperatures are produced in this convergence region.  For run A, the primary's surface gravity is too low to compress the helium layer above the critical temperature necessary for thermonuclear runway. No detonation was observed in run Ba, despite being initialized using a checkpoint from run B just prior to the observed detonation in B. This is because of the stochastic nature of the ignition mechanism --- the particular dense knot of material that led to the surface detonation in B has a different shape in Ba, and did not have a favorable geometry for igniting the helium torus. And because Ba has twice the linear resolution of run B, we could only afford to evolve it for a short period of time.  \\begin{figure}[tb] \\centering\\includegraphics[width=\\linewidth,clip=true]{fig4.eps} \\caption{Maximum temperature $T$ in a mass shell as a function of interior mass $M\\left(< r\\right)$ and time $t$ for run B. Each shell lies on a surface of equal gravitational potential. The shading is logarithmically binned, with red corresponding to $T_{\\max}$ and blue corresponding to $T_{\\min} = T_{\\max}/10^{1.5}$. The dashed line shows the region where $X_{\\rm He} > 0.02$, while the dotted line shows where $X_{\\rm Ne} > 10^{-3}$. For convenience, we include a zoomed view of the surface detonation as an inset.} \\label{radialfig} \\end{figure}  If a helium surface detonation occurs as the result of stream instabilities in a DD system, the resulting transient event could potential resemble a dim Type Ia SNe \\citep{Bildsten:2007p2743, Perets:2009p3232}. The results of our simulations shows a large degree of variability which ultimately depends on the geometry of the dense knots when they strike the surface of the primary in the impact zone. However, if the post-detonation temperature is not much larger than $2 \\times 10^{9}$ K or the geometry of the detonation fronts are not favorable, it is possible to only synthesize intermediate-mass $\\alpha$-elements (Table \\ref{outcomes}).  The prospect of igniting the CO core itself is attractive because the typical densities of moderate-mass WDs ($M \\sim 1.0 M_{\\odot}$) are $\\sim 10^{6}-10^{7}$ g cm$^{-3}$, comparable to the densities found in the DDT for Type Ia SNe \\citep{Khokhlov:1997p3443}. \\cite{Fink:2007p1886} (hereafter FHR) show that when $0.1 M_{\\odot}$ of helium is ignited at the He/CO interface of a moderate-mass WD, strong shocks are launched deep into the CO core, which converge at a focusing point. This focusing produces a region where the temperature and density meet the criteria for explosively igniting carbon \\citep{Niemeyer:1997p1901, Ropke:2007p3332}, thus leading to a full detonation of the CO core.  The linear resolution of our simulations is substantially coarser than the highest-resolution two-dimensional simulations of FHR, and we are certainly under-resolving the true density and temperature peaks. However, FHR's results appear to be rather optimistic when applied to our model for the following reasons. First, FHR assumes perfect mirror symmetry, which dramatically increases the amount of focusing by forcing the shocks to converge at a single point rather than a locus of points. Second, FHR does not use a nuclear network and makes the simplifying assumption that the entire He envelope is burned to Ni, which naturally overestimates the energy injected into the CO core when compared to the more realistic case where burning is incomplete.  Because of these differences, the conditions for double detonation are still not met even in our run with the most favorable shock geometry and strongest surface detonation (run C). In the focusing region of run C, the temperature and density are $2.5 \\times 10^{8}$ K and $2.0 \\times 10^{7}$ g cm$^{-3}$, respectively. Conditions for CO core detonation may be more propitious in systems containing a slightly more massive primary, but ultimately the successful detonation of the core depends on the precise distribution of temperature and density within the high-pressure regions of the convergence zone \\citep{Seitenzahl:2009p1881}.  The mechanism reported here for the detonation of a sub-Chandrasekhar CO WD in a dynamically unstable binary is not tied to a particular mass scale and therefore allows for considerably more diversity. As outlined above, mass transfer between a pure He WD or a He/CO hybrid and a CO WD before the merger provides a novel pathway to ignite CO WDs. Even if a critical amount of mass is not raised above the conditions required for CO detonation, a peculiar underluminous optical transient should signal the last few orbits of a merging system."
    },
    "0911/0911.5233_arXiv.txt": {
        "abstract": "In recent years we have been witnessing the discovery of one extrasolar gas giant after another. Now the time has come to detect more low-mass planets like Super-Earths and Earth-like objects. An  interesting question to ask is: where should we look for them? We have explored here the possibility of finding Super-Earths in the close vicinity of gas giants, as a result of the early evolution  of  planetary systems. For this purpose, we have considered a young planetary system containing a Super-Earth and a gas giant, both embedded in a protoplanetary disc. We have shown that, if the Super-Earth is on the internal orbit relative to the gas giant, the planets can easily become locked in a mean motion resonance. This is no longer true, however, if the Super-Earth is on the  external orbit. In this case we have obtained that the low-mass planet is captured in a trap at the outer edge of the gap opened by the giant planet and no first  order mean motion commensurabilities are expected. Our investigations might be particularly useful for the observational TTV (Transit Timing Variation) technique. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0911/0911.5696_arXiv.txt": {
        "abstract": "The Transit Timing Variation (TTV) method relies on monitoring changes in timing of transits of known exoplanets. Non-transiting planets in the system can be inferred from TTVs by their gravitational interaction with the transiting planet. The TTV method is sensitive to low-mass planets that cannot be detected by other means. Here we describe a fast algorithm that can be used to determine the mass and orbit of the non-transiting planets from the TTV data. We apply our code, {\\tt ttvim.f}, to a wide variety of planetary systems to test the uniqueness of the TTV inversion problem and its dependence on the precision of TTV observations. We find that planetary parameters, including the mass and mutual orbital inclination of planets, can be determined from the TTV datasets that should become available in near future. Unlike the radial velocity technique, the TTV method can therefore be used to characterize the inclination distribution of multi-planet systems. ",
        "introduction": "In Nesvorn\\'y \\& Morbidelli (2008) and Nesvorn\\'y (2009) (hereafter NM08 and N09) we developed and tested a fast inversion method that can be used to characterize planetary systems from the observed Transit Timing Variations (TTVs; Agol et al. 2005; Holman \\& Murray 2005). See NM08 and N09 for a technical description of the method. Here we use this new method to solve the TTV inversion  problem for an arbitrary planetary system. The results provide a baseline for studies of real exoplanetary systems for which TTVs will be detected. Examples of past work that would greatly benefit from the application of the fast inversion algorithm discussed here include Steffen \\& Agol (2005), Agol \\& Steffen (2007), Miller-Ricci et al. (2008) and Gibson et al. (2009).  In \\S2, we briefly describe the TTV inversion method. In \\S3, we apply it to a case with coplanar planetary orbits. Inclined planetary orbits are discussed in \\S4. We show, for example, that the mutual inclination of planetary orbits can be determined from TTVs. This important parameter, which may be used to test planet-migration theories (e.g., Rasio \\& Ford 1996, Goldreich \\& Sari 2003), is not typically available from other existing planet-detection methods. ",
        "conclusions": "The method developed here can be used to analyze TTVs found for any of the potentially hundreds of planets expected to be discovered by {\\it Kepler} (Beatty \\& Gaudi 2008). {\\it Kepler} should be able to detect transit timing variations of only a few seconds (Holman \\& Murray 2005), which should easily exist in many systems, extrapolating from the radial velocity planets (Agol et al. 2005, Fabrycky 2008).  Perhaps the most interesting result that comes out of this work is that the shape of the TTV signal is generally sensitive to the orbital inclination of the non-transiting planetary companion. Thus, the TTV method can provide means of determining mutual inclinations in systems in which at least one planet is transiting. This parameter cannot be determined by other planet-detection methods.  TTVIM algorithm can be easily extended to incorporate uncertainties in the transiting planet's parameters. This can be done by sampling dimensions that correspond to the additional parameters. For example, in N09 we extended the NM08 method to the case with $e_1\\neq0$. This may be especially relevant to the transiting planets that will be discovered by Kepler because these planets are expected to have wider orbits, which are less susceptible to the circularizing effects of tides. The low CPU cost of the TTVIM algorithm is the key element which will make such studies possible."
    },
    "0911/0911.4659_arXiv.txt": {
        "abstract": "{ We report on the results of new simulations of near-infrared (NIR) observations of the Sagittarius~A* (Sgr~A*) counterpart associated with the super-massive black hole at the Galactic Center.}{Our goal is to investigate and understand the physical processes behind the variability associated with the NIR flaring emission from Sgr~A*. } { The observations have been carried out using the  NACO adaptive optics (AO) instrument at the European Southern Observatory's Very Large Telescope and CIAO NIR camera on the Subaru telescope (13 June 2004, 30 July 2005, 1 June 2006, 15 May 2007, 17 May 2007  and 28 May 2008). We used a model of synchrotron emission from relativistic electrons in the inner parts of an accretion disk. The relativistic simulations have been carried out using the Karas-Yaqoob (KY) ray-tracing code.} {We probe the existence of a correlation between the modulations of the observed flux density light curves and changes in polarimetric data. Furthermore, we confirm that the same correlation is also predicted by the hot spot model. Correlations between intensity and polarimetric parameters of the observed light curves as well as a comparison of predicted and observed light curve features through a pattern recognition algorithm result in the detection of a signature of orbiting matter under the influence of strong gravity. This pattern is detected statistically significant against randomly polarized red noise. Expected results from future observations of VLT interferometry like GRAVITY experiment are also discussed.}{ The observed correlations between flux modulations and changes in linear polarization degree and angle can be a sign that the NIR flares have properties that are not expected from purely random red-noise. We find that the geometric shape of the emission region plays a major role in the predictions of the model. From fully relativistic simulations of a spiral shape emitting region, we conclude that the observed swings of the polarization angle during NIR flares support the idea of compact orbiting spots instead of extended patterns. The effects of gravitational shearing, fast synchrotron cooling of the components and confusion from a variable accretion disk have been taken into account. Simulated centroids of NIR images lead us to the conclusion that a clear observation of the position wander of the center of NIR images with future infrared interferometers will prove the existence of orbiting hot spots in the vicinity of our Galactic super-massive black hole.} ",
        "introduction": "\\begin{figure*}[!htb] \\begin{minipage}{\\textwidth} \\centering{\\includegraphics[width=\\textwidth]{12473f1.jpg}} \\end{minipage} \\caption{Sgr A* as it was observed in  NIR $L^{'}$-band (3.8 $\\mu$m) on 3 June 2008 between 05:29:00 - 09:42:00 ~(UT time). ({\\it a-f}) show the observed images of Sgr A* after 0, 50, 91, 135, 155 and 212 minutes off the start of observation. The images show that when Sgr A* is in its flaring state the flux changes up to 100\\% in time intervals  of the order of only tens of minutes.} \\label{label} \\end{figure*}  The nearest super-massive black hole candidate ($\\sim~4\\times10^6 M_\\odot$) lies at the center of our galaxy, as inferred from motions of stars near the Galactic Center (Eckart \\& Genzel 1996, 1997; Eckart et al. 2002; Sch\\\"{o}del et al. 2002; Eisenhauer et al. 2003; Ghez et al. 2000, 2005, 2008; Gillessen et al. 2009). With a luminosity of $10^{-9}-10^{-10} L_{Edd}$, where $L_{Edd}$ is its limiting Eddington luminosity, Sagittarius A*, the radio source associated with this SMBH, is one of the most extreme sub-Eddington sources accessible to observations. However, X-ray and near-infrared (NIR) flares are routinely detected with high spatial and spectral resolution observations (Baganoff et al. 2001; Porquet et al. 2003, 2008; Genzel et al. 2003; Eckart et al. 2004, 2006a-c, 2008a-c; Meyer et al. 2006a,b, 2007; Yusef-Zadeh et al. 2006a,b, 2007, 2008). These short bursts of increased radiation last normally for about 100 minutes and occur four to five times a day (see Fig. 1 for a typical behavior of Sgr~A* during a flaring phase in NIR bandwidth).   Recent NIR and X-ray observations have revealed the non-thermal nature of  high frequency radiation from Sgr~A* (Eckart et al. 2006a-c, 2008a-c; Gillessen et al. 2006; Hornstein et al. 2007). Sgr~A* is probably visible in the NIR regime only during its flaring state. The short time scale variabilities seen during several observed NIR and X-ray flares argue for an emitting region not bigger than about ten Schwarzschild radii ($r_s=\\frac{2GM}{c^2}=2r_g=1.2\\times10^{12}\\big(\\frac{M}{4\\times10^6 M_{\\odot} }\\big)$ cm) of the associated super-massive black hole (Baganoff et al. 2001; Genzel et al. 2003). We have scaled the relevant physical distances according to the gravitational radius ($r_g$) throughout this paper. The NIR flares are highly polarized and normally have X-ray counterparts, which strongly suggests a synchrotron-self-Compton (SSC) or inverse Compton emission as the responsible radiation mechanism (Eckart et al. 2004, 2006a,b; Yuan et al. 2004; Liu et al. 2006). Several observations have already confirmed the existence of a time lag between the simultaneous NIR/X-ray flares and the flares in the lower frequencies. This is interpreted as a sign for cooling down via adiabatic expansion (Eckart et al. 2006a, 2008b,c; Yusef-Zadeh et al. 2006a,b, 2007, 2008; Marrone et al. 2008, Zamaninasab et al. 2008a).  \\begin{figure*}[!t] \\begin{minipage}{1.\\textwidth} \\centering{\\includegraphics[width=\\textwidth]{12473f2.jpg}} \\end{minipage} \\caption{Our sample of light curves of Sgr~A* flares observed in NIR $K_s$ band (2.2$\\mu$m) polarimetry mode. The events were observed on  2004 June 13 (a), 2005 July 30 (b), 2006 June 1 (c), 2007 May 15 (d), 2007 May 17 (e) and  2008 May 28 (f). In each panel, the top shows the de-reddened flux density measured in mJy (black), the middle shows the polarization angle (blue) and the bottom shows the degree of linear polarization (red). The gaps in the light curves are due to the sky background measurements.} \\label{events} \\end{figure*}   \\begin{table*}[btp] \\begin{center} \\begin{tabular}{lcccccccr}\\hline \\hline Date (Telescope) & Spectral domain & UT start & UT stop & Max. flux  & Min. flux& Average flux & Average polarization  \\\\ &  & time & time &  & & sampling rate & sampling rate  \\\\ \\hline \\hline \\\\ 2004 June 13 (NACO)& 2.2$\\mu$m & 07:20:02 & 09:15:08  & 5.19 mJy & 2.21 mJy & 1.2 min &  3 min \\\\ \\\\ &        &          &       &       \\\\ 2005 July 30 (NACO)& 2.2$\\mu$m & 02:07:36 & 07:03:39 &  8.19 mJy & 1.23 mJy & 1.2 min & 3 min \\\\ &        &          &       &      \\\\ &        &          &       &       \\\\ 2006 June 1 (NACO)& 2.2$\\mu$m & 04:26:03 & 10:44:27 & 19.33 mJy & 0.72 mJy & 1.5 min & 2 min \\\\ &        &          &       &      \\\\ \\\\ 2007 May 15 (NACO)& 2.2$\\mu$m & 09:08:14 & 09:42:12 & 22.27 mJy & 1.70 mJy & 1.5 min & 2 min \\\\ &        &          &       &      \\\\ \\\\ 2007 May 17 (NACO)& 2.2$\\mu$m & 04:42:14 & 09:34:40 & 11.87 mJy & 1.86 mJy & 1.5 min & 2 min \\\\ &        &          &       &      \\\\ \\\\ 2008 May 28 (CIAO)& 2.15$\\mu$m & 09:22:51 & 13:00:37 & 7.70 mJy & 0.97 mJy & 3.3 min & 10 min \\\\ &        &          &       &      \\\\ \\hline \\hline \\end{tabular} \\end{center} \\caption{Observations log.} \\label{table1} \\end{table*}  The other feature related to these NIR/X-ray flares are the claimed quasi-periodic oscillations (QPOs) with a period of $20\\pm5$ minutes, which have been reported  in several of these events (Genzel et al. 2003; Belanger et al. 2006; Eckart et al. 2006b,c; 2008a; Meyer et al. 2006a,b, Hamaus et al. (2009)). Short periods of increased radiation (the so called \"NIR flares\", normally around 100 minutes) seem to be accompanied by QPOs. All the studies mentioned above probed this 20$\\pm$5 minutes quasi-periodicity, by performing a sliding window analysis with window lengths of the order of the flaring time. Recently, Do et al. (2009) argued that they did not find any significant periodicity at any time scale while probing their sample of observations for a periodic signal. Their method is based on the Lomb-Scargle periodogram analysis of a sample of six light curves and comparing them with several thousands of artificial light curves with the same underlying red-noise. One must note that the suggested QPOs are transient phenomena, lasting for only very few cycles (50-100 minutes). This kind of behavior, along with the inevitable uncertainty  in the red noise power law index determination, makes a clear and unambiguous detection ($>5\\sigma$) of a periodic signal very difficult. Whenever a flare of Sgr~A* was observed with polarimetry, it was found that it is accompanied by significant polarization which varies on similarly short timescales as the light curve itself. By carrying out an analysis only on the total flux, some pieces of the observed information are ignored. It is already known that polarimetric data have been shown to be able to reveal substructure in flares, even when the light curve appears largely featureless (e.g. see Fig. 4 in Eckart et al. 2006b). The other main advantage of polarimetric observations is that, in addition to the flux density light curve, one can analyze the changes in the observed degree of polarization and the changes in polarization angle during flaring time as well.  The claimed quasi-periodicity has been interpreted as being related to the orbital time scale of the matter in the inner parts of the accretion disk. According to well-known observed high frequency quasi-periodic oscillations (HFQPOs) in X-ray light curves of stellar mass black holes and binaries (Nowak \\& Lehr 1998), this interpretation is of special interest since it shows a way to better understand  the behavior of accretion disks for a wide range of black hole masses. The recent unambiguous discovery of a one hour (quasi-)periodicity in the X-ray emission light curve of the active galaxy RE J1034+396 provides further support to this idea and extends the similarity between stellar-mass and super-massive black holes to a new territory (Gierli\\'{n}ski et al. 2008).  Although the origin of the observed QPOs in the sources associated with black holes is still a matter of debate, several magnetohydrodynamic (MHD) simulations confirmed that it could be related to instabilities in the inner parts of the accretion disks, very close to the marginally stable orbit of the black hole ($r_{mso}$), and also possibly connected with the so-called \"stress edge\" (Hawley 1991; Chan et al. 2009b). If the flux modulations are related to a single azimuthal compact over-dense region (hereafter: \"hot spot\"), orbiting with the same speed as the underlying accretion disk, one can constrain the spin of the black hole by connecting the observed time scales of QPOs to the orbital time scale of matter around the black hole: $T = 2.07 (r^{\\frac{3}{2}}+a)(\\frac{M}{4\\times10^6M_\\odot})$~min (Bardeen et al. 1972, where $ -1\\leq a\\leq1$ is the black hole dimensionless spin parameter and $r$ is the distance of the spot from the black hole). The characteristic behavior of general relativistic flux modulations produced via the orbiting hot spots have been discussed in several papers (see e.g. Cunningham \\& Bardeen 1973; Abramowicz et al. 1991; Karas \\& Bao 1992; Hollywood et al. 1995; Dov\\v{c}iak 2004, 2007). In this paper, we have used a spotted accretion disk scenario to model the observed patterns of our sample of NIR light curves.  Several authors have proposed different models in order to explain the flaring activity of Sgr~A*. These models cover a wide range of hypotheses like disk-star interactions (Nayakshin et al. 2004), stochastic acceleration of electrons in the inner region of the disk (Liu et al. 2006), sudden changes in the accretion rate of the black hole (Liu et al. 2002), heating of electrons close to the core of a jet (Markoff et al. 2001; Yuan et al. 2002), trapped oscillatory modes in the inner regions of the accretion disk in the form of spiral patterns or Rossby waves (Tagger \\& Melia 2006; Falanga et al. 2007; Karas et al. 2008), non-axisymmetric density perturbations which emerge as the disk evolves in time (Chan et al. 2009b), non-Keplerian orbiting spots falling inward inside the plunging region created via magnetic reconnections (Falanga et al. 2008), and also  comet like objects trapped and tidally disrupted by the black hole (\\v{C}ade\\v{z} et al. 2006; Kosti\\'{c} et al. 2009).  Observational data render some of these models unlikely. The star-disk interaction model is unable to produce the repeated flux modulations and also the high degree of polarization since it mainly deals with thermal emission. Tidal disruption of  comet-like objects are also unable to reproduce the observed rate of  flares per  day, since the estimated capture rate of such objects for the Sgr~A* environment is at least one order of magnitude lower (\\v{C}ade\\v{z} et al. 2006). Nevertheless, several viable models exist and make different predictions that can be distinguished observationally. For example, one important characteristic prediction of hot spot models is about the wobbling of the center of the images (Broderick \\& Loeb 2006a,b; Paumard et al. 2006; Zamaninasab et al. 2008b; Hamaus et al. 2009). Significant  effort has been already devoted to measure this possible position wander of Sgr~A* in the mm, sub-mm and NIR regimes (Eisenhauer 2005b, 2008, Gillesen 2006, Reid et al. 2008, Doelleman et al. 2008). In this paper we discuss how NIR polarimetry and the next generations of VLT interferometry (VLTI) and Very Long Baseline Interferometry (VLBI) experiments can provide data to support or reject certain models for the accretion flow/outflow related to Sgr~A*. Obtaining accurate data on the accretion flow of Sgr~A* can lead us to a better understanding of the physics of strong gravitational regimes, formation of black holes and their possible relation to the galaxy formation process in a cosmological context.  In Sect. 2 we present a complete sample of NIR light curves observed in  the polarimetry mode. A brief description about the details of the observation and data reduction methods is provided. We discuss  the quasi-periodicity detection methods and present the results of a correlation analysis between the flux and polarimetric parameters. A general description of our model setup and results of simulations are discussed in Sect. 3. We show how NIR polarimetry can be used as a way to constrain physical parameters of the emitting region (like its geometrical shape) in Sect. 4 and Sect. 5. In Sect. 6 we mainly discuss the predictions that the future NIR interferometer (GRAVITY) is expected to reveal and how different assumptions in the model parameters can modify the results. In Sect. 7 we summarize the main results of the paper and draw our conclusions. ",
        "conclusions": "We used a sample of NIR flares of Sgr~A* observed in polarimetry mode to study the nature of the observed variability. Using the z-transformed discrete correlation function algorithm, we found a significant correlation between changes in the measured polarimetric data and  total flux densities. This provides  evidence that the variations probably  originate from inner parts of an accretion disk while a strong gravity's effects are manifested inside them.  In order to obtain this information polarization data is indispensable since the corresponding signals are very difficult or even impossible to extract from simple total intensity light curves only (Do et al. 2008). The presence of significant signals from orbiting matter then calls for detailed modeling of the NIR polarized light curves in order to analyze the distribution of the emitting material and the magnetic field structure within the accretion disk (e.g. Eckart et al. 2006a, 2008a, Meyer et al. 2006ab,2007).  In addition, the pattern recognition algorithm we employed in this paper is an efficient tool to search for flare events that carry the signature of strong gravity in light curves which are significantly longer than the orbiting time scale over which such an event can typically occur. Other methods that make simultaneous use of the entire light curve (like e.g. the Lomb-Scargle algorithm) tend to dilute these signals and lower the possibility for significant detections dramatically (Do et al. 2008, Meyer et al. 2009).  In order to constrain the physical properties of the emitting region we employed a relativistic disk model with azimuthal over-densities of relativistic electrons. A combination of a synchrotron mechanism and relativistic amplifications allows us to fit the real observed data and make predictions about astrometric parameters of the accretion disk around Sgr~A*. The modeled light curves show the same correlation between the flux and polarimetric data as the one deduced from observations.  Simulations have been carried out in a way that they cover a wide range of parameters, including the effects of gravitational shearing inside the accretion disk, the heating and cooling time scales, the inclination and the spin of the black hole. It is discussed how the observed swings in the polarization angle support the idea of a compact source for the emission, instead of  radially extended spiral shapes. Furthermore, we present a model in which the observed NIR polarization angle can lead to confining the expected region for the expected outflow/wind from Sgr~A*. The model also predicts that when observations will be able to resolve the position of such an outflow, the magnetic field structure inside the accretion disk could be confined.  Finally, the centroid paths of the NIR images are discussed. In comparison with the results by Broderick \\& Loeb (2006a,b) and Hamaus et al. (2009), we have shown that the geometrical structure of the emitting region (elongation of the spot according to the Keplerian shearing, multi-component structures, spiral arms, confusion caused by the radiation from the hot torus) can affect the expected centroid tracks. While all the mentioned geometries are able to fit the observed fluxes, we show how the future NIR interferometer GRAVITY on the VLT can break these degeneracies. The results of simulations propose that focusing GRAVITY observations on the polarimetry mode could reveal a clear centriod track of the spot(s). We conclude that even though a non-detection of centroid shifts can not rule out the multi-components model or spiral arms scenarios, a clear position wander in the center of NIR images during the flares will  support the idea of  bright long-lived spots orbiting occasionally around the central black hole. This possible detection opens a new window for testing the physics very close to the edge of infinity."
    },
    "0911/0911.3401_arXiv.txt": {
        "abstract": "Previous solar system constraints of the Brans-Dicke (BD) parameter $\\omega$ have either ignored the effects of the scalar field potential (mass terms) or assumed a highly massive scalar field. Here, we interpolate between the above two assumptions and derive the solar system constraints on the BD parameter $\\omega$ for {\\it any} field mass. We show that for $\\omega=O(1)$ the solar system constraints relax for a field mass  $m \\gsim 20 \\times m_{AU}= 20\\times 10^{-27}GeV$. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0911/0911.1859_arXiv.txt": {
        "abstract": "{We here propose the time drift of subtended angles as a new possible cosmological probe. In particular, with the coming era of microarcsecond astrometry, our proposal can be used to measure the Hubble expansion rate of our universe in a direct way.} \\begin{document} ",
        "introduction": "Research in cosmology has become extraordinarily lively in the past quarter century. In particular, the use of Type Ia supernovae as standard candles lead to the discovery that the expansion of our universe is accelerating, implying that the energy of the universe may be dominated by some sort of exotic matter, with a ratio of pressure to density less than one third, dubbed dark energy\\cite{Weinberg}. However, besides the cosmological constant, there are various other dark energy models devised to explain the mysterious accelerating expansion. On the other hand, the current accelerating expansion of our universe can also be accounted for by either the modification of Einstein gravity or the violation of Copernican principle. With the other cosmological probes, some of them has been ruled out, but there still remain a number of theoretical models surviving such observational tests\\cite{FTH}.  In this era of empirical cosmology, it is thus significant to propose some new observational programs that may shed light on our understanding of the universe in a way independent from other known cosmological probes. By comparing and combining results from very different methods of determining cosmological parameters, we may both obtain stronger constraints than any method alone would impose, and test techniques against one another to identify signatures of systematic effects. In particular, in a field so afflicted by systematic errors as cosmology, having many independent but complementary techniques is the best way to ensure that we are on the right track to explore the genuine mechanism underlying the evolution of our universe. We here contribute to such a theme by proposing the time drift of subtended angles as a new cosmological probe. As depicted in Fig.\\ref{td}, the ordinary cosmological probes are usually related to the sky survey along the past light cone of today. However, we have another way to acquire the evolution information of our universe by measuring the time drift of cosmological objects. The classical exemplification is Sandage-Loeb test\\cite{Sandage,Loeb}.  In next section we shall popularize the spirit of time drift by confining ourselves onto the case of subtended angles and argue that the time drift of subtended angle may provide us with a new cosmological probe, just like Sandage-Loeb test. Some discussions will be presented in the last section. \\begin{figure} \\includegraphics[width=12cm,height=10cm] {td.eps}\\\\ \\caption{Time drift versus sky survey in the redshift space.}\\label{td} \\end{figure} ",
        "conclusions": "We have argued that the time drift of subtended angles may be used as a new promising cosmological probe to measure the Hubble expansion rate of our universe in a direct way.  However, so far the analysis of its feasibility rests on an order of magnitude estimate. It is thus important to see whether the realistic data will distort such an estimate dramatically. On the other hand, although it is hard to imagine that the virialized systems will undertake a size evolution to the same order as the Hubble expansion, it may be expected that at least a few of them will violate this general belief. Nevertheless we may employ the other side of the same coin by instead using Eq.(\\ref{coin}) to determine the size evolution rate for those exceptional objects and non-virialized ones if we know the evolution law of our universe from other cosmological probes, and vice versa, since the subtended angle by definition entangles the whole universe with those luminous objects in it.  It is interesting to ask whether we can disentangle them intrinsically. One way is to combine it with the time drift of the redshift difference from the far and near points of the object to obtain an evolution free cosmological probe, reminiscent of Alcock-Paczynski test\\cite{AP}. Another way out is to go directly for the time drift of angle diameter distance, which by expression is immune to the size evolution effect totally.  All of these issues are worthy of further investigation but beyond the scope of this paper, we expect to report them elsewhere in the near future."
    },
    "0911/0911.1297_arXiv.txt": {
        "abstract": "We study relativistic stars in the context of  scalar tensor theories of gravity that try to account for the observed cosmic acceleration and satisfy the local gravity constraints via the chameleon mechanism.  More specifically, we consider two types of models: scalar tensor theories with an inverse power law potential and f(R) theories. Using a relaxation algorithm, we construct numerically static relativistic stars, both for constant energy density configurations and for a polytropic equation of state. We can reach a gravitational potential up to $\\Phi\\sim 0.3$ at the surface of the star, even in f(R) theories with an ``unprotected\" curvature singularity. However, we find static configurations only if the pressure does not exceed one third of the energy density, except possibly in a limited region of the star (otherwise, one expects tachyonic instabilities to develop). This constraint is  satisfied by realistic equations of state for neutron stars. ",
        "introduction": "One of the most challenging tasks for cosmology and fundamental physics today is to try to understand the apparent acceleration of the Universe. Beyond the minimal assumption of a pure cosmological constant, two main approaches have been explored. The first one consists in assuming some unknown form of matter, called {\\it dark energy}, characterized by an equation of state  $P\\simeq -\\rho$. The second approach is more radical, as it tries to explain the present observations as the manifestation of a modified theory of gravity, which mimicks general relativity on solar system scales, but significantly deviates from it on cosmological scales.  It turns out that it is rather difficult to construct a theory of gravity that, while being  internally consistent, can account for the cosmological observations and be compatible with the present gravity constraints deduced from laboratory experiments and from  solar  and astrophysical systems. A class which has attracted a lot of attention is the so-called $f(R)$ gravity theories where the standard Einstein-Hilbert gravitational Lagrangian, proportional to the scalar curvature $R$, is replaced by a function of $R$ while the matter part of the Lagrangian is left unchanged (see e.g. \\cite{review} for a recent review). After several detours, it has been realized that viable $f(R)$ theories must satisfy stringent conditions in order to avoid instabilities and to satisfy the present laboratory and solar system constraints, and a few models have been carefully constructed to meet these requirements \\cite{Hu:2007nk,Appleby:2007vb,Starobinsky:2007hu}.  To explore the full viability of these theories, it is important to go one step further by studying their behaviour  in the  strong gravity regime, such as reigns in the core of the most relativistic stars, namely neutron stars. In this context, it has been claimed in \\cite{Frolov,KM1} that  very relativistic stars do not exist  in the models  \\cite{Hu:2007nk,Appleby:2007vb,Starobinsky:2007hu} because of the presence of an easily accessible singularity. In a recent work \\cite{BL1}, we have shown that this claim does not hold by constructing numerically static relativistic stars. Note that  relativistic stars have also been studied in \\cite{Kainulainen:2007bt}  but within  different $f(R)$ models.  $f(R)$ models can also be seen as a subclass of scalar-tensor theories. In particular, viable $f(R)$ models rely on the so-called chameleon mechanism. For this reason,  it is interesting to extend the study of relativistic stars to  chameleon models. In this work, we show that the behaviour of chameleon and $f(R)$ models is quite similar. In particular, in both cases, the scalar field in the innermost part of the star sticks to the minimum of its effective potential  (if  this minimum exists). If the equation of state is such that  $\\rho-3P<0$, which occurs in the central part of  highly relativistic constant energy density stars, then there is no minimum. As we show here, it is nevertheless possible to construct numerically static stars, up to some critical value of the central energy density. We believe that this is due to tachyonic instabilities, associated with a negative effective squared mass, which develop and prevent the existence of a static star configuration. However, this problem does not apply to  realistic neutron stars: although there is a large uncertainty on the equation of state deep inside a neutron star, the equations of state that have been proposed in the literature verify $\\rho-3P>0$ throughout the star. To construct   a simple approximation of a realistic neutron star, we  have used a polytropic equation of state.  { Although, according to our previous work \\cite{BL1} and the present one, the singularity of the models \\cite{Hu:2007nk,Appleby:2007vb,Starobinsky:2007hu} does not seem so far to be an obstacle  for relativistic stars, it appears to be problematic for cosmology~\\cite{Frolov}.} This has motivated the construction of regularized versions of these models by adding for instance an extra $R^2$ term~ \\cite{Starobinsky:2007hu,Capozziello:2009hc,Dev:2008rx,Abdalla:2004sw,Thongkool:2009js}. A more sophisticated model was also proposed recently in~\\cite{Appleby:2009uf}.  We will consider these ``cured''  $f(R)$ theories and compare their behaviour in relativistic stars with their ``singular'' counterparts.   This paper is organized as follows. In the next section, we derive the main equations governing static and spherically  symmetric configurations in scalar-tensor theories. Section III is devoted to relativistic stars in chameleon models. We then consider, in Section IV, the $f(R)$ models, including the ``cured'' models that have recently been introduced to solve the cosmological singularity problem of some of these models. We finally conclude in Section V. ",
        "conclusions": "We have studied in this work  relativistic stars in scalar tensor and $f(R)$ theories that use the chameleon mechanism. The behaviour of the scalar field is extremely similar in the two types of models. The main properties are the following.  Deep inside the star, the scalar field follows very closely the minimum of its effective potential (if it exists). As a consequence, if the minimum changes as a function of the radius, the scalar field will closely follow it. If $\\trho-3\\tP<0$, which can occur for instance in the central region of very compact stars with constant energy density, there is no minimum for the effective potential. It is however possible to find numerical solutions for constant energy stars with both $w<1/3$  and $w>1/3$ regions, but only up to some maximum gravitational potential. We have  given qualitative arguments to show that one expects tachyonic instabilities if most of the star is characterized by $w>1/3$.  As already emphasized in our previous work \\cite{BL1}, our results invalidate the claim that highly relativistic stars cannot be constructed in $f(R)$ theories. Numerically, the construction of a relativistic stars in $f(R)$ gravity is a  challenging task because  the scalar field value is extremely close to the singularity in the center of the star.   We have also constructed relativistic stars in the so-called ``cured'' $f(R)$ models, which have been advocated to solve a cosmological singularity problem.  We have also illustrated how the screening effect of the chameleon mechanism manifests itself for relativistic  stars, by computing numerically the effective coupling of the star as well as the post-Newtonian parameter $\\tilde\\gamma$.  \\vskip 1cm {\\bf Note:} while this paper was being completed, another paper~\\cite{Cooney:2009rr} on relativistic stars in $f(R)$ theories appeared on the arXiv, although in a different context.  \\vskip 1cm"
    },
    "0911/0911.5361_arXiv.txt": {
        "abstract": "The misalignment between the orbital plane of a transiting exoplanet and the spin axis of its host star provides important insights into the system's dynamical history. The amplitude and asymmetry of the radial-velocity distortion during a planetary transit (the Rossiter-McLaughlin effect) depend on the projected stellar rotation rate $v\\sin I$ and misalignment angle $\\lambda$, where the stellar rotation axis is inclined at angle $I$  to the line of sight. The parameters derived from modelling the R-M effect have, however, been found to be prone to systematic errors arising from the time-variable asymmetry of the stellar spectral lines during transit. Here we present a direct method for isolating the component of the starlight blocked by a planet as it transits the host star, and apply it to spectra of the bright transiting planet HD 189733b. We model the global shape of the stellar cross-correlation function as the convolution of a limb-darkened rotation profile and a gaussian representing the Doppler core of the average photospheric line profile. The light blocked by the planet during the transit is a gaussian of the same intrinsic width, whose trajectory across the line profile yields a precise measure of the misalignment angle and an independent measure of $v\\sin I$. We show that even when $v\\sin I$ is less than the width of the intrinsic line profile, the travelling Doppler ``shadow'' cast by the planet creates an identifiable distortion in the line profiles which is amenable to direct modelling. Direct measurement of the trajectory of the missing starlight yields self-consistent measures of the projected stellar rotation rate, the intrinsic  width of the mean local photospheric line profile, the projected spin-orbit misalignment angle, and the system's centre-of-mass velocity. Combined with the photometric rotation period, the results give a geometrical measure of the stellar radius which agrees closely with values obtained from high-precision transit photometry if a small amount of differential rotation is present in the stellar photosphere. ",
        "introduction": "The recent discoveries of strong projected spin-orbit misalignments in the transiting exoplanets XO-3b (\\citealt{hebrard2008_xo3_rm}, \\citealt{winn2009xo3_rm}), HD 80606b (\\citealt{winn2009hd80606rm}, \\citealt{pont2009hd80606rm}) and WASP-14b \\citep{johnson2009wasp-14rm}, and retrograde orbital motion in WASP-17b \\citep{anderson2009wasp-17} and HAT-P-7b (\\citealt{winn2009hat-p-7rm}; \\citealt{narita2009hat-p-7rm}) indicate violent dynamical histories for a significant fraction of the population of hot-Jupiter planets.  These misalignments are measured by obtaining densely-sampled radial-velocity observations during transits. The starlight blocked by the planet during a transit possesses the radial velocity of the obscured part of the stellar surface. The radial-velocity centroid of the light emanating from the visible parts of the stellar disk exhibits a reflex motion disturbance --the Rossiter-McLaughlin (RM) effect -- whose form depends on the projected stellar equatorial velocity  $v\\sin I$, the impact parameter $b$ of the planet's path across the stellar disc, the projected misalignment angle $\\lambda$ between the orbital axis and the stellar spin axis, and the relative sizes of the star and planet.  Detailed semi-analytic formulations for determining the anomalous departure of the line centroid from the stellar reflex orbit have been developed by \\citet{ohta2005rm}, \\citet{gimenez2006rm} and \\citet{hirano2009rm}. The process of measuring and interpreting this reflex velocity shift is not, however, straightforward. If the stellar rotation profile is even partially resolved by the spectrograph, the missing light within the planet's silhouette introduces an asymmetry into the line spread function. For instruments such as SOPHIE and HARPS, radial velocities are derived using a gaussian fit to a cross-correlation function (CCF) computed from the stellar spectrum and a weighted line mask (\\citealt{baranne96elodie}; \\citealt{pepe2002harps}). Such methods work extremely well for the symmetric and time-invariant profile shapes encountered outside transit. When applied to profiles with a time-varying degree of intrinsic asymmetry, however, they yield velocity measures that depart systematically from the velocity centroid of the visible starlight.  \\citet{winn2005rm} overcame this difficulty for their iodine-cell observations of HD 209458b using a direct modelling approach which they have employed in subsequent papers. They built models of both the unobscured stellar profile and the light blocked by the planet into their analysis of the line-spread function,  deriving semi-empirical corrections to the model velocities to enable meaningful comparison with the data. Nonetheless, their study of the RM effect in HD 189733 \\citep{winn2006hd189733} shows a clear pattern of correlated radial-velocity residuals during the transit. In a more recent HARPS study of the RM effect during a transit of HD 189733b, \\citet{triaud2009rm} noted an almost identical systematic pattern of correlated residuals between the radial velocities of the gaussian fits to the CCFs, and the line centroid velocities of the best-fitting model computed using the  formulation of \\citet{gimenez2006rm}.  In developing an analytic model of this velocity anomaly for cross-correlation spectra, \\citet{hirano2009rm} have shown that such errors are exacerbated as $v\\sin I$ increases and the asymmetry becomes more extreme. If uncorrected, the anomaly can lead to significant over-estimation of the value of $v\\sin I$, particularly for rapidly-rotating planet-host stars. The motivation for the present study is to develop a comprehensive model of  the changes in CCF  morphology that occur before, during and after a transit, in terms of  the projected stellar equatorial rotation speed, the spin-orbit misalignment angle, the intrinsic width of the combined stellar CCF and instrumental broadening function, and the stellar limb-darkening coefficient. In principle this approach is similar  to that of \\citet{winn2005rm}, but we use it to decompose the observed CCF into its various components and track them directly. A similar methodology has already been used to model the RM effect in the binary stars V1143 Cyg \\citep{albrecht2007} and DI Her \\citet{albrecht2009}.  To illustrate the effectiveness of this approach as an alternative to gaussian-fitting in the study of transiting exoplanets, we re-analyse the same HARPS observations of HD 189733b described by \\citet{triaud2009rm}. We show that the spectral signature of the light blocked by the planet is clearly discernible in the residuals when a model of the unobscured starlight is subtracted from the data, and track it directly to derive new measurements of the projected stellar rotation velocity and spin-orbit misalignment angle. ",
        "conclusions": "In principle we could determine a lower limit on the stellar radius from $v\\sin I$ and the photometric rotation period, which was determined by   \\citet{henry2008hd189733} to be $11.953\\pm 0.009$ days. The very close alignment of the projected orbital and spin axes in the plane of the sky makes it highly probable that the inclination $I$ of the stellar spin axis to the line of sight should be very close to the inclination $i=85.5$ degrees of the planet's orbit. Assuming $\\sin I = \\sin i = 0.997$, our value for $v\\sin I$ yields a direct estimate of the stellar radius $R_*=0.732\\pm 0.007$ R$_\\odot$, independently of the usual assumptions involving isochrone fits to the stellar density and effective temperature. This is significantly smaller than the $R_*=0.77\\pm 0.01$ R$_\\odot$ determined by \\citet{triaud2009rm}, and marginally less than the $R_*=0.75\\pm 0.01$ R$_\\odot$ found by \\citet{pont2007hst}.  We suspect that the discrepancy arises from differential rotation of the stellar photosphere. \\citet{gaudi2007rm} pointed out that under some circumstances, the RM effect could be used to measure stellar differential rotation. As with most previous studies of this kind, we have measured $v\\sin I$ under the assumption that the star rotates as a solid body. In all stars for which differential rotation has been measured directly, however, a solar-like pattern is found: the equator rotates faster than the poles, with the departure from the equatorial spin rate having a sine-squared dependence on latitude. The typical difference in rotational angular frequency $\\Delta\\Omega = \\Omega_{\\rm equator}-\\Omega_{\\rm pole}$ is found to be of order 0.04 radian per day for rapidly-rotating stars K dwarfs similar to HD 189733, and only weakly dependent on rotation rate \\citep{barnes2005diffrot}. This is very similar to the solar value of $\\Delta\\Omega$.  If the starspots giving rise to the photometric modulation signal are concentrated at the same stellar latitude as the planet's path across the stellar disc, we should derive the correct value of the stellar radius. Unfortunately we have no way of determining the dominant latitude from which the modulation signal originates. If the orbital and spin axes are closely aligned, the impact parameter of the planet's path across the star carries it across a stellar latitude of 39$^\\circ$ in the hemisphere facing away from the observer. Although \\citet{pont2007hst}  noted the passage of the planet across one or two isolated dark spots, these spots are significantly foreshortened. It is reasonable to expect the modulation signal to arise mainly from spots at intermediate to low latitude in the observer's hemisphere. A difference in latitude of 20 or 30 degrees is sufficient to give a difference in angular velocity $\\delta\\Omega\\simeq 0.015$ radians per day between the main spot belts and the latitudes traversed by the planet. This discrepancy yields a stellar radius that is too small by about 3 percent, which is sufficient to reconcile our radius estimate with those of \\citet{pont2007hst} and \\citet{triaud2009rm}.  We conclude that direct modelling of the cross-correlation function during a planetary transit is feasible even for a host star such as HD~189733 whose $v\\sin I$ is comparable to the intrinsic line width. The profile decomposition method described here yields values and error estimates for the stellar spin rate and orbital obliquity that agree closely with the results of previous studies of this system. The parameter values are free of the systematic errors that occur when velocity measurements are derived from gaussian fits to line profiles with inherent time-variable asymmetry. This circumvents the need for semi-empirical corrections when modelling the Rossiter-McLaughlin effect. The method makes the fullest use of all the information present in the cross-correlation function: $v\\sin I$ is tightly constrained by the shape of the stellar rotation profile as well as by the spectral signature of the starlight blocked by the planet. The stellar radius obtained from the photometric rotation period and the stellar $v\\sin I$ agrees well with values obtained from high-precision transit photometry, if a modest amount of differential rotation is present.  The method described here could in principle be extended to observations for which an iodine cell is used to track the spectrograph PSF and wavelength scale. The analysis of such observations requires a reference spectrum of the host star, taken with the same instrument without the iodine cell. An artificial reference spectrum could be generated by convolving a limb-darkened rotation profile with the the spectrum of a narrow-lined star of the same spectral type. The Doppler signature of the light blocked by the planet could be mimicked by scaling, shifting and subtracting the same narrow-lined spectrum from the broadened stellar spectrum; the result would then be used as the reference spectrum. A very similar procedure was described by \\citet{winn2005rm} for calibrating the departure of the velocities measured during a transit of HD 209458b from the predictions of the \\citet{ohta2005rm} model, using high-resolution solar spectra."
    },
    "0911/0911.0896_arXiv.txt": {
        "abstract": " ",
        "introduction": " ",
        "conclusions": "\\label{sec5} In the last decade ground based as well as satellite observations probed with greater accuracy the nature of the magnetic field of the sun. The SOHO \\cite{soho}, TRACE \\cite{trace}, GONG \\cite{gong} experiments \\footnote{ SOHO stands for Solar and Heliospheric Observatory; TRACE stands for  TRansition Region and Coronal Explorer; GONG stands for Global Oscillation Network Group.} helped in deepening solar observations. In spite of these direct probes, it is fair to say that the 22-year magnetic cycle of the Sun cannot  be claimed to be fully understood in terms of the so-called $\\alpha\\omega$ mechanism. The GONG observations give, for instance, the profile of the angular velocity of the sun not only on the surface but also in the interior of the sun. In spite of this valuable determination other effects may play a role so that it is fair to say that the situation is not clear \\cite{sun}.  The sun is the closest and better observed astrophysical object and still the features of dynamo amplification are under debate. In the case of large-scale magnetism we face a similar situation.  It seems that we will probably never be able to know the initial conditions of the galactic dynamo as accurately as we pretend to understand the initial conditions of  solar dynamos where, still, crucial puzzles remain.  The hopes of clarifying this problem in the near future might not be so forlorn.  Indeed, as repeatedly argued,  CMB anisotropies and polarization \\cite{mga,dyn} are a powerful window on pre-decoupling physics and, consequently, on pre-decoupling magnetic fields. Waiting for more direct observational evidences it is plausible to speculate that a moderate dynamo action combined with compressional amplification could indeed bridge the regime of the initial conditions with the observed large-scale magnetic fields.  Even in this simplified framework important theoretical problems persist and they have to do with the way the Universe becomes a good conductor at the end of inflation.  The early variation of the gauge couplings is a potential candidate for producing large-scale magnetic seeds for galaxies and clusters.  However, more effort is certainly required especially in modeling the finite conductivity effects and the self-consistent evolution of the whole system of equations. \\newpage"
    },
    "0911/0911.4967.txt": {
        "abstract": "{ We consider the application of quantum corrections computed using renormalization group arguments in the astrophysical domain and show that, for the most natural interpretation of the renormalization group scale parameter, a gravitational coupling parameter $G$ varying $10^{-7}$ of its value across a galaxy (which is roughly a variation of  $10^{-12}$ per light-year) is sufficient to generate galaxy rotation curves in  agreement with the observations. The quality of the resulting fit is similar to the Isothermal profile quality once both the shape of the rotation curve and the mass-to-light ratios are considered for evaluation. In order to perform the analysis, we use recent high quality data from nine regular disk galaxies. For the sake of comparison, the same set of data is modeled also for the Modified Newtonian Dynamics (MOND) and for the recently proposed Scalar Tensor Vector Gravity (STVG). At face value, the model based on quantum corrections clearly leads to better fits than these two alternative theories.} ",
        "introduction": "\\label{intro}  The dynamics of many galaxies looks like as if the main part of their masses is distributed in a different way than their light is, with  heavier densities than those expected from the emitted radiation by the stars and gas. Curiously, although galaxies are far from being ``hairless\", like fundamental particles or black holes, these masses discrepancies follow considerably strong patterns, including a series of correlations with the luminous matter, e.g. \\cite{urc, btf}.  Usually, in order to solve this missing matter problem, a qualitatively new kind of matter, dark matter (DM), is evoked. Among all kinds of dark matter, current cosmological observations (like the large scale structure and the microwave background anisotropies) favor the collisionless and ``cold'' type of dark matter (CDM). However the CDM paradigm faces some difficulties at galactic scales, like the angular momentum issues on galaxy formation, the discrepancies on the number of satellite galaxies,  and the cuspy dark matter density profile (see \\cite{primack} for a recent review). The CDM paradigm has not yet yielded a satisfactory dark matter profile for spiral galaxies from the simulations of cosmological evolution. The problem with the cuspy dark matter distribution in galaxies was already pointed in \\cite{flores, moore1994}, and in particular both the Navarro-Frenk-White \\cite{nfw} and the Moore \\cite{moore} dark matter profiles  have a cusp on the dark matter density at the galactic center whose dynamical consequences do not match the observations; the discrepancies with the observations becoming particularly clear for the Low Surface Brightness galaxies. The problematics associated with this cuspy dark matter profile, and possible solutions, were systematically analyzed in many references, including \\cite{blokrubin, spano,  Gentile, ThingsRot, ohthings, Donato}, see \\cite{corecuspblok} for a recent review.   It is well known that the parameters of the dark matter profiles that best fit the observed rotation curves  (like the Isothermal \\cite{BBS} or the Burkert\\footnote{Considering the CDM simulations, the Burkert profile has merits over the Isothermal one, in particular for its large $R$ behavior and its finite total mass; see also \\cite{burkertCDM}.  } \\cite{burkert} ones, which have a core  and were proposed on phenomenological grounds) display many correlations with the baryonic matter behavior, e.g. \\cite{mcgaughdmlsb, corecor, Donato, BBS, burkert}, see also \\cite{Cthesis}. These correlations between a galaxy rotation curve and its baryonic matter suggest another interpretation to the missing mass problem in galaxies, namely  a change of the gravitational law itself instead of introducing DM.  This has lead to diverse proposals, including modifications on the law of inertia  \\cite{mond} and its extensions to General Relativity (e.g., \\cite{revteves}), proposals with extra-dimensions \\cite{extrad}, actions of General Relativity with an extra affine connection or a new graviton \\cite{ebi}, conformal gravity \\cite{Mannheim}, non-symmetric metrics \\cite{BrownMoffat}, among others possibilities. MOND \\cite{mond, sandersmcgaugh} is by far the most studied example of the last approach. Although it seems to be consistent with  the general features of the galactic dynamics, its concordance with observations (if the original recipe is applied) started to be questioned in the light of the higher precision recent observations and the improvements on the  stellar mass-to-light ratios constraints, e.g. \\cite{Gentile, m33m31, tevesgalaxy, mondmilky, mondearlydisk}.   One can suppose that the deviation from the gravitational (either Einstein or Newton) law is due to the quantum effects, such as semiclassical corrections, effects of quantum gravity, consequences of string theory physics, extra dimensions, branes or some other (maybe yet unknown) model of ``quantum gravity''.  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% \\vskip 1mm    The corrections to the Newton law is a relatively common feature of the different models of ``quantum gravity'', including higher derivative quantum gravity \\cite{frts82}, different versions of the effective low-energy quantum theory of gravitational field \\cite{AntMot,don,Wood}, so-called non-perturbative quantum gravity based on the hypothesis of the existence of the non-Gaussian UV fixed point \\cite{Reuter} and, finally, in the semiclassical approach to quantum gravity \\cite{book} (see also references therein). The application of these corrections has been elaborated recently in the cosmological \\cite{nova,CCfit,Gruni} and astrophysical \\cite{Gruni} areas. One can see the recent paper \\cite{DCCR} for the detailed discussion of the quantum field theory backgrounds of these quantum effects. Let us emphasize, from the very beginning, that in the most cases the present day state of art in all mentioned approaches to quantum gravity does not enable one to really calculate the relevant quantum contributions to the Newton law in a unique and consistent way. At the same time, we have many reasons to believe that these quantum corrections can be non-zero and, in principle, may have a measurable effect. Typically, the theoretical estimate for the quantum contribution to the gravitational law involves more or less strong arbitrariness, related to the dependence of the quantum corrections on some dimensional parameter (scale parameter in the renormalization group-based approaches, for instance) from one side, and to the physical interpretation of this parameter from the other side.  In the present paper we will consider a relatively general model of quantum corrections which is mainly controlled by covariance. Furthermore we will use a new identification of the scale parameter $\\mu$ in the astrophysical setting (rotation curves problem), which we believe to be the more natural compared to the one \\ ($\\mu \\sim 1/r$) \\ considered before in \\cite{Bertol,Gruni,Reuter} (see also \\cite{Mazzitelli-94}). It turns out that this new identification of scale leads to the surprisingly good result for the rotation curves of the galaxies, which is very close in quality to the output of the mainstream approach based on the DM paradigm.  The previous papers on the subject were based on the renormalization group equations coming from the higher-derivative quantum gravity \\cite{Bertol} (unfortunately, the proper renormalization group equations \\cite{frts82} are ambiguous in this case, making their application to astrophysics doubtful \\cite{nova}), on the assumption of the Appelquist and Carazzone decoupling of the quantum corrections in the low-energy domain \\cite{babic,nova,CCfit} and on the assumption of the asymptotic safety and existence of the non-Gaussian fixed point in the four-dimensional quantum gravity \\cite{Reuter}. It is remarkable that all those, in fact rather different approaches, converge in predicting qualitatively similar logarithmic running of the effective Newton constant with the scale $\\mu$. It was shown in \\cite{Gruni} (see also previous qualitatively similar consideration in \\cite{Bertol}) that this logarithmic behaviour, together with the mentioned above identification of $\\mu^{-1}$ to the distance from the center of the galaxy, $r$, can partially explain the main features of the rotation curves without invoking the DM. The astrophysical applications of \\cite{Bertol}, \\cite{Gruni} and \\cite{Reuter} were performed for the point-like model of the galaxy, which is not supposed to be realistic. It was pointed out in \\cite{Gruni} that it would be very interesting to make a more detailed investigation, taking the case of a plane or thin disk distribution of mass in the galaxy. We present the corresponding analysis and show that it meets certain obstacles at both theoretical and phenomenological levels.  In order to make a conclusive consideration of the astrophysical application of the renormalization group method, one has to ask whether the $\\mu \\sim 1/r$ identification is the unique possible one, or there are some alternative options. One can remember that the identification of scale in the cosmological setting is a nontrivial issue, which was treated differently by different authors in different papers \\cite{CCcosm,nova,babic,CCfit,Reuter-t,Reuter}. Furthermore, was an interesting attempt to construct the regular scale-fitting procedure \\cite{Guberina-scale}. The output of this procedure was the most natural identification of $\\mu$ with the Hubble parameter $H$, which is the energy characteristic of the external metric leg of the Feynman diagram, corresponding to the quantum correction to the expansion of the universe. Indeed, similar energy characteristic can be constructed in the case of the galaxies rotation curves, and it is indeed different from $1/r$. Moreover, this ``right'' setting of the energy scale leads to the immediate and dramatic improvement of the phenomenological analysis results. After all, we gain a chance to explain the rotation curves on the basis of quantum corrections to the classical action of gravity, without using DM.  The paper is organized as follows. In section 2 we discuss very general features of quantum contributions and establish their likely form, which is based on the general covariance and on the proper existence of these contributions, following \\cite{Gruni}. As we have already mentioned above, this existence can not be either proved or disproved at the present-day state of art in this area \\cite{DCCR} and it is legitimate to apply a phenomenological approach. Let us note that the covariance enables one to put serious restrictions on the possible forms of quantum corrections in the cosmological setting, the astrophysical ones can be deduced starting from the cosmological setting result and the assumption of a unique effective action for gravity.  The main new element in section 2 is the new definition of the energy scale, which is appropriate for the application to the rotation curves problem. In section 3 we present the numerical results of our model applied to the data of nine regular galaxies divided into two samples of data. For comparison purposes, our results are shown together with the numerical results of three other proposals, obtained from exactly the same data and procedures. These three proposals are: dark matter (modeled with the Isothermal profile \\cite{BBS}), MOND \\cite{mond} and STVG \\cite{stvg, BrownMoffat}. The first is the leading candidate for the missing matter problem in galaxies and all the universe altogether, MOND is the paradigmatic model for galaxy rotation curves without dark matter and STVG is a recently proposed model for astrophysics and cosmology without dark matter that also uses the variation of the gravitational coupling parameter $G$ to fit galaxy rotation curves. Finally, in section 4 we draw our Conclusions and discuss the perspectives of the theory of quantum corrections (especially the ones based on the renormalization group) as a potential competitor of the standard $\\Lambda$CDM model, at both astrophysical and cosmological scales.  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% ",
        "conclusions": "\\label{conclusions}  We  presented a model, motivated by quantum corrections to the Einstein-Hilbert action, that introduces small inhomogeneities in the gravitational coupling across a galaxy (of about 1 part in $10^7$) and can generate galaxy rotation curves in agreement with the observational data, without the introduction of dark matter as a new kind of matter. Considering the samples of  galaxies evaluated in this paper, the quality of the rotation curve fits and the associated mass-to-light ratios from our model is a bit lower, or about the same, than  the Isothermal profile quality, but with one less free parameter. We also compared the results of our model with MOND \\cite{mond} and STVG\\cite{stvg,BrownMoffat}, and at face value our model yielded clearly better results.  Our results can be seen as a next step compared to the previous models motivated by renormalization group effects in gravity, \\cite{Gruni, Reuter:2007de}. Their original analyses could only yield a rough estimate on the galaxy rotation curves, since they were restricted to modeling a galaxy as a single point. Trying to extend this approach to real galaxies, we have shown that the phenomenologically optimized proper scale for the renormalization group phenomenology is not of a geometric type, like the inverse of the distance, but is proportional to the Newtonian potential with null boundary condition at infinity. This setting is shown to be in agreement with  both  theoretical expectations and with observations.    The essential feature for the rotation curve fittings is the formula (\\ref{vsvg}), which is by itself a remarkably simple formula that provides a very efficient description of galaxy rotation curves. We have shown that  it can be derived from the assumption of a gravitational coupling parameter with a very small departure from the constant $G_0$, such that, in the weak field limit, it depends on the  logarithm of the Newtonian potential, e.g. $G(\\Phi_\\Newt) = G_0 \\[ 1 - \\beta \\, \\ln (\\Phi_\\Newt / \\Phi_0) \\]$.  In the last equation, $\\beta$ is a positive and very small ($\\sim 10^{-7}$) effective parameter that necessarily depends on the distribution of mass of the system. The precise relation between $\\beta$ (or $\\alpha  \\nu$ using the theory of quantum corrections) with the galaxy parameters is yet to be unveiled. To this end, from a phenomenological approach, a larger sample of galaxies would be necessary to find meaningful results. On the other hand, using the Tully-Fisher law, it is not hard to guess that $\\beta$, for disk galaxies, should scale with the mass approximately as $\\sqrt{M}$, in which $M$ is the baryonic mass of the galaxy. The results for $\\alpha$ in tables \\ref{resultsA} and \\ref{resultsB}, clearly show that it indeed increases with the galactic mass. In order to avoid premature statements, we think that the nature of $\\beta$ (or $\\alpha$) should be explored in more detail in the forthcoming papers.  In the present paper we have addressed the issue of generating galaxies rotation curves by developing and extending the proposal of \\cite{Gruni}. It turns out that, at least for the sample galaxies we dealt with here, one can explain these curves without invoking the dark matter concept. Does it mean that we can really be free of dark matter in constructing a realistic cosmological scenario? The answer to this question is probably negative. It is well known that there many other issues associated with dark matter, or the lack of it, that are necessary to take into account. In particular, the results on the density perturbations and related issues such as cosmic microwave background radiation (CMBR) and the large scale structure (LSS) data, baryon acoustic oscillations, big bang nucleosynthesis, gravitational lensing and others (see, e.g. \\cite{Ross} for a recent review). In all these observational and experimental issues the standard way of explaining the date is to assume the existence of  dark matter. At the same time, one can not underestimate the fact that the rotation curves may be explained without dark matter effect. One can imagine, for instance, the scenario with essentially smaller amount of dark matter which has slightly different set of properties compared to the usual CDM model. In this case the rotation curves will be explained by summing up the effect of quantum corrections and the one of the dark matter content. The preliminary analysis shows this possibility can not be ruled out \\cite{AToJFa}. Finally, there is a chance to address all mentioned issues trading a large amount of cold dark matter content by another one, which can have qualitatively other origin, and possibly invoking the quantum corrections to the cosmological perturbations spectrum (see, e.g., \\cite{den, GCC}).   \\vspace{.2in} {\\it Note added.} After the first version of this work was finished, we became aware, through \\cite{brownstein}, and contrary to the expectations in \\cite{stvg}, that there are small and perhaps sensible differences between the fits of the rotation curves using either the STVG or the MSTG formulations. If precision on the nomenclature and results is at stake, our results on the model developed by Moffat and collaborators apply to the older formulation, namely MSTG.    \\vspace{.2in}"
    },
    "0911/0911.1224_arXiv.txt": {
        "abstract": "\\noindent We study the linear differential equation $\\dot{x} = Lx$ in $1$:$1$ resonance. That is, $x \\in \\fR^4$ and $L$ is $4 \\times 4$ matrix with a semi-simple double pair of imaginary eigenvalues $(i\\beta,-i\\beta,i\\beta,-i\\beta)$. We wish to find all perturbations of this linear system such that the perturbed system is stable. Since linear differential equations are in one to one correspondence with linear maps we translate this problem to $\\gl(4,\\fR)$. In this setting our aim is to determine the stability domain and the singularities of its boundary. The dimension of $\\gl(4,\\fR)$ is $16$, therefore we first reduce the dimension as far as possible. Here we use a versal unfolding of $L$ ie a transverse section of the orbit of $L$ under the adjoint action of $\\Gl(4,\\fR)$. Repeating a similar procedure in the versal unfolding we are able to reduce the dimension to $4$. A $3$-sphere in this $4$-dimensional space contains all information about the neighborhood of $L$ in $\\gl(4,\\fR)$. Considering the $3$-sphere as two $3$-discs glued smoothly along their common boundary we find that the boundary of the stability domain is contained in two right conoids, one in each $3$-disc. The singularities of this surface are transverse self-intersections, Whitney umbrellas and an intersection of self-intersections where the surface has a self-tangency. A Whitney stratification of the $3$-sphere such that the eigenvalue configurations of corresponding matrices are constant on strata allows us to describe the neighborhood of $L$ and in particular identify the stability domain. ",
        "introduction": "\\label{sec:intro}    \\subsection{Setting}\\label{sec:introset} Suppose that the ordinary differential equation \\begin{equation}\\label{eq:orig} \\dot{x} = F_{\\mu}(x) \\end{equation} has a stationary point at $x=0$ for all $\\mu$. In this equation $x \\in \\fR^n$, $\\mu \\in \\fR^p$ and $F_{\\mu}(x)\\frac{\\partial}{\\partial x}$ is a vector field on $\\fR^n$ smoothly depending on $x$ and $\\mu$. Furthermore suppose that $A: \\fR^p \\to \\gl(n,\\fR): \\mu \\mapsto A(\\mu) = DF_{\\mu}(0)$ is the linear part of the vector field at $x=0$. We are interested in the case that the phase space is $4$-dimensional and the linear part $A(\\mu)$ has a double semi-simple pair of complex conjugate imaginary eigenvalues $(i\\beta,-i\\beta,i\\beta,-i\\beta)$, $\\beta \\neq 0$, at $\\mu = 0$. A system with such a linear part is also said to be in 1:1-\\emph{resonance}. With $I$ denoting a two by two identity matrix and $\\beta=1$ we have \\begin{equation}\\label{eq:l} A(0) = L = \\begin{pmatrix} 0 & I\\\\-I & 0 \\end{pmatrix}. \\end{equation}  Our aim is to give a description of a small neighborhood of $L$ in $\\gl(4,\\fR)$ and in particular find the stability domain and its boundary. On the latter we expect bifurcations of the stationary point of the original differential equation \\eqref{eq:orig}. Thus this study fits in a much larger study of local bifurcations of systems having a stationary point with zero or imaginary eigenvalues in the linear part.  We give a short list of generic bifurcations with increasing number of zero or imaginary eigenvalues in the linear part. We do not specify non-degeneracy conditions, but instead refer to the literature. In dimension one, only one eigenvalue can be zero and we have \\emph{transcritical} (TC) and \\emph{saddle-node} (SN) bifurcations. However in the presence of symmetry the \\emph{pitchfork} (PF) bifurcation occurs. In dimension two there can be two non semi-simple zero eigenvalues denoted by $0^2$, then the system has a \\emph{Bogdanov-Takens} (BT) bifurcation. But there can also be a pair of imaginary eigenvalues, then the system undergoes a \\emph{Hopf} bifurcation. If the two dimensional system is Hamiltonian and has two zero eigenvalues $0^2$, then there is in general a \\emph{Hamiltonian saddle-node} (SSN) bifurcation.  When the dimension gets larger, results get sparser because the codimensions of the bifurcations readily increase. In dimension three, the simplest case is the \\emph{Hopf-saddle-node} (HSN) bifurcation when there is one zero eigenvalue and an imaginary pair, denoted $0\\beta$. Other cases are systems with linearizations having eigenvalues $0^3$ or $00^2$. The bifurcations for these cases are not yet well studied. The simplest bifurcation for general systems in dimension 4 is the \\emph{Hopf-Hopf} (HH) bifurcation where the linearization at the stationary point has two pairs of \\emph{non-resonant} imaginary eigenvalues $\\beta_1\\beta_2$. Other results for $4$-dimensional systems are known for special cases only. For Hamiltonian systems there is a bifurcation at $k$:$l$-resonance of imaginary eigenvalues $\\beta_1\\beta_2$. A special place is taken by the 1:$\\pm 1$-resonances, where the sign in 1:$\\pm 1$ designates \\emph{symplectic signature} of eigenvalues. At $1$:$-1$-resonance there is a so called \\emph{Hamiltonian-Hopf} (SH) bifurcation. But at $1$:$1$-resonance there is another bifurcation that has not been analyzed completely. The $1$:$1$-resonance in reversible systems is very similar to the $1$:$-1$-resonance in Hamiltonian systems, note that there is no signature for imaginary eigenvalues in reversible systems. However, there is in reversible systems a \\emph{reversible sign} for real and zero eigenvalues, see \\cite{hvn96}. The case $0_+^4$ has been studied, but the case $0_-^4$ has not, as far as we know. We summarize the local bifurcations (with lowest codimension) occurring in systems up to dimension 4 in table \\ref{tab:bifs}.  \\begin{table}[htbp] \\begin{center} \\begin{tabular}{|c|c|l|c|l|l|}\\hline dim & evc & bif & codim & comments & references\\\\\\hline 1   & $0$ & SN  & 1 & & \\cite{gh, kuz}\\\\ &     & TC  & 1 & ``$x=0$ stationary for all $\\mu$'' & \\cite{gh, kuz}\\\\ &     & PF  & 2 & & \\cite{gh, kuz}\\\\ &     & PF  & 1 & $\\fZ_2$-symmetric & \\cite{gh, kuz}\\\\\\hline 2  & $0^2$   & BT  & 2 & & \\cite{drsz}\\\\ & $\\beta$ & H   & 1 & & \\cite{gh, kuz}\\\\ & $0^2$   & SSN & 1 & symplectic & \\cite{wig}\\\\\\hline 3  & $0\\beta$ & HSN & 2 & & \\cite{bv, gh, gpd}\\\\\\hline 4  & $\\beta_1\\beta_2$ & HH & 2 & $\\beta_1 : \\beta_2$ irrational & \\cite{gh, kuz}\\\\ & $\\beta^2$        & SH & 1 & symplectic 1:$-1$-resonance & \\cite{mee}\\\\ & $\\beta_1\\beta_2$ & nn & 1 & symplectic $1$:$2$-resonance  & \\cite{dui}\\\\ & $\\beta_1\\beta_2$ & nn & 3 & symplectic $1$:$3$-resonance  & \\cite{dui}\\\\ & $\\beta_1\\beta_2$ & nn & 2 & symplectic $k$:$l$-resonance, $k$:$l \\neq 1$:$2$, $k$:$l \\neq 1$:$3$  & \\cite{dui}\\\\ & $\\beta^2$        & nn & 1 & reversible 1:1-resonance  & \\cite{msv}\\\\ & $\\beta_+\\beta_+$ & nn & nk & symplectic 1:1-resonance & \\cite{hh}\\\\ & $\\beta\\beta$     & nn & nk & 1:1-resonance &\\\\ & $0_+^4$          & nn & 2 & reversible, reversible sign +1 & \\cite{ioo}\\\\\\hline \\end{tabular} \\end{center} \\caption{\\textit{Low codimension bifurcations of stationary points in systems up to dimension 4. We do not list the non-degeneracy conditions. When such conditions are violated but higher order conditions are met, we generally get a higher codimension bifurcation for the same eigenvalue configuration. The abbreviations have the following meaning dim: dimension of phase space; evc: eigenvalue configuration (see section \\ref{sec:stabdombkar}); bif: name of bifurcation; codim: codimension of bifurcation; SN: saddle-node; TC: transcritical; PF: pitchfork; BT: Bogdanov-Takens; SSN: Hamiltonian saddle-node; HSN: Hopf-saddle-node; HH: Hopf-Hopf; SH: Hamiltonian Hopf; nn: no name; nk: not known.}\\label{tab:bifs}} \\end{table}  Here we do not even attempt to describe the bifurcation of system \\eqref{eq:orig}, whose linear part at $x=0$ and $\\mu=0$ is in 1:1-resonance. Instead we focus on the linearized system $\\dot{x}=A(\\mu)x$ near $\\mu=0$. As mentioned before we will describe the neighborhood of $A(0)=L$ in $\\gl(4,\\fR)$, in particular the stability domain and the singularities on its boundary. Moreover it is of both theoretical and practical interest to know how linear Hamiltonian, reversible and equivariant subsystems appear as subspaces in $\\gl(4,\\fR)$ and how they intersect the stability domain and its boundary near $L$. This last issue will be treated in more detail in \\cite{hkb}.  The \\emph{stability domain} in $\\gl(4,\\fR)$ is defined as follows. \\begin{equation}\\label{eq:stabdom} \\stabdom = \\{A \\in \\gl(4,\\fR)\\;|\\; \\text{if $\\lambda$ is an eigenvalue of $A$ then $\\Re(\\lambda) < 0$}\\} \\end{equation} Then the boundary of the stability domain $\\dstabdom$, is characterized by vanishing real parts of one or more eigenvalues of $A \\in \\gl(4,\\fR)$. By the implicit function theorem, $\\dstabdom$ as a hyper surface in $\\gl(4,\\fR)$, is smooth in points where the corresponding matrix has a simple zero eigenvalue or a simple complex conjugate pair with vanishing real part. At points where multiple eigenvalues have vanishing real parts we may expect singularities. A simple example being two pairs of purely imaginary eigenvalues $\\pm i \\beta_1$ and $\\pm i \\beta_2$ with $\\beta_1 \\neq \\beta_2$. Then $\\dstabdom$ has generically a transverse self-intersection. We will only consider the stability domain and its boundary in a small neighborhood of $L$ and in figure \\ref{fig:eigcfgs} we list the possible \\emph{eigenvalue configurations}. For an informal definition of eigenvalue configuration see section \\ref{sec:statres}, a precise definition will be given in section \\ref{sec:stabdom}. We use the following coding: $\\gamma$ for a pair of complex conjugate eigenvalues; $\\beta$ for a pair of complex conjugate imaginary eigenvalues; $\\alpha$ for a real eigenvalue; $0$ for a zero eigenvalue; $\\gamma \\gamma$ for semi-simple double eigenvalues and $\\gamma^2$ for double eigenvalues with nilpotent part of height 2. A subscript $\\pm$ denotes the sign of the real part and an optional index is used to denote different eigenvalues. \\begin{figure}[htbp] \\setlength{\\unitlength}{1mm} \\begin{picture}(100,25)(0,0) \\put(10, 15){\\includegraphics{eigcfgs.ps}} \\put(2,   1){$\\beta\\beta$} \\put(24,  1){$\\beta^2$} \\put(42,  1){$\\beta_1\\beta_2$} \\put(60,  1){$\\beta\\gamma_-$} \\put(82,  1){$\\gamma_{-1}\\gamma_{-2}$} \\put(104, 1){$\\gamma_-\\gamma_+$} \\put(124, 1){$\\beta\\gamma_+$} \\put(142, 1){$\\gamma_{+1}\\gamma_{+2}$} \\end{picture} \\caption{\\textit{Eigenvalue configurations near $L$. $\\beta\\beta$ is the eigenvalue configuration of $L$. On $\\dstabdom$ we have $\\beta^2$, $\\beta_1\\beta_2$ or $\\beta\\gamma_-$, on $\\stabdom$ we only have $\\gamma_{-1}\\gamma_{-2}$ and elsewhere we have $\\gamma_-\\gamma_+$, $\\beta\\gamma_+$ or $\\gamma_{+1}\\gamma_{+2}$.}\\label{fig:eigcfgs}} \\end{figure} \\begin{remark}\\label{rem:randcfgs} There are many more eigenvalue configurations on $\\dstabdom$, including zero or real eigenvalues, for example $\\alpha_{-1}\\alpha_{-2}\\beta$, $0\\alpha_-\\beta$ and $0^4$. In points of $\\dstabdom$ where the corresponding matrix has eigenvalue configuration $\\alpha_{-1}\\alpha_{-2}\\beta$, $\\dstabdom$ is smooth, for $0\\alpha_-\\beta$ it has a self-intersection and for $0^4$, $\\dstabdom$ will be more singular. In a study of $0^4$ the present study of $L$ will appear as a sub case. However, these eigenvalue configurations do not occur on an arbitrary small neighborhood of $L$. \\end{remark}  The question of singularities on the boundary of the stability domain has been taken up earlier, see for example a discussion of the \\emph{decrement diagram} in \\cite{arn1}. There a classification is guided by codimension that is by the number of parameters in a family of matrices $A(\\mu)$ where $A(0)$ is the central singularity. Also see \\cite{le82,ms99} for an elaboration on this idea with examples. Here we wish to view a study of $L$ in a classification guided by dimension of the phase space. This will lead to high codimensions. Indeed, in studying $L$ we have to consider an eight parameter family. See section \\ref{sec:methrcu} how we reduce such families. In the end it turns out that we have to study a three parameter family. In this family we do find most of the singularities listed in \\cite{arn1} for generic three parameter families.    \\subsection{Motivation and main questions}\\label{sec:intromoq} The main motivation for this study comes from a wide variety of applications where the question of stability of a system near 1:1-resonance turns up in various forms. For example, double semi-simple imaginary eigenvalues are natural in the spectra of rotationally and spherically symmetrical models of solids and fluids \\cite{ki09}, ranging from car brakes \\cite{ki08}, rotating shafts \\cite{nn98}, and computer hard discs \\cite{cb92} to rotating elastic Earth \\cite{rv09}, and from vortex tubes \\cite{hf08} to magneto-hydrodynamics \\cite{kgs09}. Double semi-simple eigenvalues are also characteristic of optimal structures and are responsible for high sensitivity of the latter to small imperfections \\cite{slo94}. Many \\emph{dissipation-induced instabilities} in water wave models can be traced back to the occurrence of double imaginary eigenvalues, see \\cite{br97} and references therein, but also \\cite{mk91,bkmr94,mo95,km07,ki07,bmr08}. An early observation of friction induced instability can be found in \\cite{zie}, which has been related to a singularity on the boundary of the stability domain by \\cite{bot}. For more applications and references see \\cite{hkb}.  The questions in the applications mentioned above in many instances boil down to questions about the stability domain and its boundary in $\\gl(4,\\fR)$ near a matrix with double semi-simple imaginary eigenvalues. The main questions we wish to address here are the following. \\begin{enumerate}\\itemsep 0pt \\item What are the open domains in parameter space with constant eigenvalue configuration? \\item What are the singularities on the boundaries of these domains? \\item In particular, we address the above two questions for the stability domain.  \\commentaar{volgende naar artikel B}  \\end{enumerate} We will consider these questions for a general system. But in many applications such a system can be considered as small, dissipative perturbations of a Hamiltonian system in $1:\\pm1$-resonance. However, also reversibility and equivariance with respect to a circle group play a prominent role. In all our constructions to study the neighborhood of a matrix with double semi-simple imaginary eigenvalues, we take care to carry them out in such a way that Hamiltonian, reversible and equivariant subsystems can be recognized easily.    \\subsection{Organization}\\label{sec:introorg} In section \\ref{sec:statres} we give an overview of the results, that is a description of the domains with constant eigenvalue configuration near $L$ in $\\gl(4,\\fR)$. In particular we present the stability domain and its singularities. The methods we use will be outlined in section \\ref{sec:methods}. We give a short overview of the centralizer unfolding of a matrix $A$ in $\\gl(n,\\fR)$, a family describing a neighborhood of $A$. We also present a method to reduce the number of parameters in this family leading to a reduced centralizer unfolding with the same properties but easier to analyze. This will be applied in section \\ref{sec:unfol} to the matrix $L$ introduced in section \\ref{sec:introset}. The resulting reduced centralizer unfolding of $L$ is further analyzed in section \\ref{sec:stabdom} where we characterize the stability domain and its singularities. Finally in section \\ref{sec:evcnbhdl} we give a description of a small neighborhood of $L$ in $\\gl(4,\\fR)$. ",
        "conclusions": ""
    },
    "0911/0911.3221_arXiv.txt": {
        "abstract": "{Information on physical characteristics of astrometric radio sources, magnitude and redshift in the first place, is of great importance for many astronomical studies.  However, data usually used in radio astrometry is incomplete and often outdated.} {Our purpose is to study the optical characteristics of more than 4000 radio sources observed by the astrometric VLBI technique since 1979. Also we studied an effect of the asymmetry in the distribution of the reference radio sources on the correlation matrices between vector spherical harmonics of the first and second degrees.} {The radio source characteristics were mainly taken from the NASA/IPAC Extragalactic Database (NED).  Characteristics of the gravitational lenses were checked with the CfA-Arizona Space Telescope LEns Survey.  SIMBAD and HyperLeda databases was also used to clarify the characteristics of some objects.  Also we simulated and investigated a list of 4000 radio sources evenly distributed around the celestial sphere.  We estimated the correlation matrices between the vector spherical harmonics using the real as well as modelled distribution of the radio sources.} {A new list of physical characteristics of 4261 astrometric radio sources, including all 717 ICRF-Ext.2 sources has been compiled. Comparison of our data of optical characteristics with the official International Earth Rotation and Reference Systems Service (IERS) list showed significant discrepancies for about half of 667 common sources. Finally, we found that asymmetry in the radio sources distribution between hemispheres could cause significant correlation between the vector spherical harmonics, especially if the case of sparse distribution of the sources with high redshift. We also identified radio sources having many-year observation history and lack redshift. This sources should be urgently observed at large optical telescopes.} {The list of optical characteristics created in this paper is recommended for use as a supplement material for the next International Celestial Reference Frame (ICRF) realization. It can be also effectively used for cosmological studies and planning of observing programs both in radio and optics.} ",
        "introduction": "Information on physical characteristics of the astrometric radio sources is important for planning of VLBI experiments and analysis of VLBI data to do a research in cosmology, kinematics of the Universe, etc.  In particular, the primary mainspring to this work was a support of the investigation of the systematic effects in apparent motion of the astrometric radio sources observed by VLBI (Gwinn et al. 1997, MacMillan 2005, Titov 2008a, Titov 2008b).  The official list of the physical characteristics of the ICRF radio sources is supported by the IERS (International Earth Rotation and Reference Systems Service) ICRS (International Celestial Reference System) Product Center (Archinal et al. 1997). The latest version of the IERS list is available in the Internet\\footnote{http://hpiers.obspm.fr/icrs-pc/info/car\\_physique\\_ext1}. However this list has some deficiencies: \\begin{itemize} \\item Not all the sources observed in the framework of geodetic and astrometric experiments are included in the IERS list. \\item The characteristics of some sources in the IERS list are outdated or doubtful. \\end{itemize}  To overcome this problems, we performed a compilation of new list of the physical characteristics of astrometric radio sources using the latest information. Hereafter this list is referred to as OCARS (Optical Characteristics of Astrometric Radio Sources).  The list of radio sources with their positions was originally taken from the Goddard VLBI astrometric catalogues\\footnote{http://vlbi.gsfc.nasa.gov/solutions/astro}, version 2007c, with addition of two absent ICRF-Ext.2 (Fey et al. 2004) sources 1039-474 and 1329-665 {(\\bf hereafter we will use 8-char IERS designation HHMMsDDd which is an abridged version of the IAU-compliant name 'IERS BHHMMsDDd')}. In the last version, the source list was updated using the 2009a\\_astro.cat catalogue computed by Leonid Petrov\\footnote{http://astrogeo.org/vlbi/solutions}. It gives 4261 radio sources in total.  At this stage mainly the NASA/IPAC Extragalactic Database\\footnote{http://nedwww.ipac.caltech.edu/} (NED) was scoured. Characteristics of the gravitational lenses were checked with the CfA-Arizona Space Telescope LEns Survey\\footnote{http://cfa-www.harvard.edu/glensdata/} (CASTLES). Several sources were checked with SIMBAD\\footnote{http://simbad.u-strasbg.fr/} the HyperLeda\\footnote{http://leda.univ-lyon1.fr/} databases. In the OCARS list we have included only the optical characteristics of astrometric radio sources: source type, redshift and visual magnitude. The flux parameters are not included in our list because they are available from other centers directly working on correlation and primary processing of the VLBI observations.  The OCARS was preliminary presented in (Malkin \\& Titov 2008). In this paper we investigated statistical properties of the list in more detail and study their impact on the kinematic analysis of radio source motions.  Analysis of the radio source apparent motion revealed some statistically significant systematics described by the vector spherical harmonics of the first and second orders (dipole and quadrupole effects, respectively) (Titov 2008a, Titov 2008b). The dipole effect could be caused by the Galactocentric acceleration of the Solar system (Gwinn et al. 1997; Sovers et al. 1998; Kovalevsky 2003; Klioner 2003; Kopeikin \\& Makarov 2006) or a hypothetic acceleration of the Galaxy relative to the reference quasars. The quadrupole harmonic, considered in details by Kristian \\& Sachs (1966), could be caused either of the primordial gravitational waves or anisotropic expansion of the Universe. This result was confirmed by Ellis et al. (1985) although they stated ``the major problem is that neither the distortion nor the proper motions are likely to be measurable in practice in the foreseeable future''. In this case the quadrupole effect should be redshift-dependent, and the apparent proper motion will increase with redshift. However, Pyne et al. (1996) and Gwinn et al. (1997) also discussed the gravitational waves with the wavelength shorter than the Hubble length. Thus, the proper motion, induced by the short-wavelength gravitational waves, also might be constant over all redshifts.  Due to asymmetry of the astrometric radio source distribution around the sky, the correlation between the vector spherical harmonic components is not zero. Therefore, we studied the effect of the asymmetry using the real uneven and simulated even distribution of the sources.  The OCARS list can be used as a supplement material for the second realization of the International Celestial Reference Frame (ICRF2), as well as a database for kinematic studies of the Universe and other related works, including scheduling of dedicated IVS (International VLBI Service for Geodesy and Astrometry, Schl\\\"ueter \\& Behrend 2007) programs. ",
        "conclusions": "To conclude, it is necessary to note that this `historic' deficit of the radio sources (and, additionally, the radio sources with measured redshift) might cause problems with further investigation of the hardly detectable systematic effects in the proper motion of the reference radio sources. Therefore, large observational projects for spectroscopy of the astrometric radio sources in the Southern hemisphere are very important.  Nonetheless, some observations in the North hemisphere also need to be undertaken.  Independently, the MASIV scintillation survey (Lovell et al. 2003, 2009) also demonstrates a highly significant dramatic decrease in the numbers of scintillators for redshifts in excess of z = 2. The lack of scintillation at high redshifts is clear evidence for an increase in the source angular sizes with increasing redshift. Such an increase may be cosmological in origin or may be a propagation effect of inter-galactic scattering (Lovell et al. 2009).  To observe and study radio source in both frequency range (optical and radio) characteristics such as visual magnitude, redshift in optic and flux density in several radio bands have to be measured. We also need to be sure that the same source is observed by the optical and radio instruments. Due to possible misalignment between optical and radio positions the physical characteristics might help to solve the problem of identification.  To chase this aim, a new list of optical characteristics of 4261 astrometric radio sources, OCARS, including all 717 ICRF-Ext.2 sources has been compiled. The OCARS list includes source type, redshift and visual magnitude (when available). Detailed comments are provided when necessary, which is especially useful in understanding of incomplete, contradictory and controversial astrophysical data. The OCARS may serve to various VLBI tasks, for instance:  \\begin{itemize} \\item As a supplement material for the second ICRF realization ICRF2 (Ma 2008). \\item As a database for VLBI data analysis. \\item For planning of IVS and other observing programs, in order to enrich the observational history of the sources with reliable determined redshift. \\item For future link between optical (GAIA) and radio (ICRF) celestial reference frames, once the GAIA optical position of about 100,000 quasars will be available. \\end{itemize}  We performed a detailed comparison of the OCARS list with the official IERS list, and found many discrepancies for about a half of common sources.  This comparison showed that the IERS list seems to be outdated and should be used with care. Besides, the IERS list, being intended to provide physical characteristics of the IERS sources only, contains only a small fraction of the whole set of astrometric radio sources used nowadays.  We also compared the OCARS with the newest Large Quasar Astrometric Catalogue LQAC, and found discrepancies which worth further investigating. Most of discrepancies seems to be a result of different object identification.  This is only the first stage of our work. We are planning the following steps: \\begin{itemize} \\item To continue searching for the missing and checking out the ambiguous characteristics through literature and astronomical databases. \\item To organize photometric and spectroscopy observations of astrometric radio sources with large optical telescopes.  In particular, such an observational program started at Pulkovo Observatory in 2008. Observations are being made on the 6-m BTA telescope of the Special Astrophysical Observatory in North Caucasus. \\end{itemize}  The list of optical characteristics of astrometric radio sources presented in this paper is publicly available at \\mbox{\\tt www.gao.spb.ru/english/as/ac\\_vlbi/sou\\_car.dat} and is updated once in several months."
    },
    "0911/0911.3198_arXiv.txt": {
        "abstract": "We propose to use alternative cosmic tracers to measure the dark energy equation of state and the matter content of the Universe [w$(z)$ \\& $\\Omega_m$]. Our proposed method consists of two components: (a) tracing the Hubble relation using HII galaxies which can be detected up to very large redshifts, $z\\sim 4$, as an alternative to supernovae type Ia, and (b) measuring the clustering pattern of X-ray selected AGN at a median redshift of $\\sim 1$. Each component of the method can in itself provide interesting constraints on the cosmological parameters, especially under our anticipation that we will reduce the corresponding random and systematic errors significantly. However, by joining their likelihood functions we will be able to put stringent cosmological constraints and break the known degeneracies between the {\\em dark energy} equation of state (whether it is constant or variable) and the matter content of the universe and provide a powerful and alternative route to measure the contribution to the global dynamics and the equation of state of {\\em dark energy}. A preliminary joint analysis of X-ray selected AGN clustering (based on the largest to-date XMM survey; the 2XMM) and the currently largest SNIa sample, the {\\em Constitution} set (Hicken et al.), using as priors a flat universe and the WMAP5 normalization of the power-spectrum, provides: $\\Omega_{\\rm m}=0.27\\pm 0.02$ and w$=-0.96\\pm 0.07$. Equivalent and consistent results are provided by the joint analysis of X-ray selected AGN clustering and the latest Baryonic Acoustic Oscillation measures, providing: $\\Omega_{\\rm m}=0.27\\pm 0.02$ and w$=-0.97\\pm 0.04$. ",
        "introduction": "We live in a very exciting period for our understanding of the Cosmos. Over the past decade the accumulation and detailed analyses of high quality cosmological data (eg., supernovae type Ia, CMB temperature fluctuations, galaxy clustering, high-z clusters of galaxies, etc.) have strongly suggested that we live in a flat and accelerating universe, which contains at least some sort of cold dark matter to explain the clustering of extragalactic sources, and an extra component which acts as having a negative pressure, as for example the energy of the vacuum (or in a more general setting the so called {\\em dark energy}), to explain the observed accelerated cosmic expansion  (eg. Riess, et al. 1998; 2004; 2007, Perlmutter et al. 1999; Spergel et al. 2003, 2007, Tonry et al. 2003; Schuecker et al. 2003; Tegmark et al. 2004; Seljak et al. 2004; Allen et al. 2004; Basilakos \\& Plionis 2005; 2006; 2009; Blake et al. 2007; Wood-Vasey et al. 2007, Davis et al. 2007; Kowalski et al. 2008, Komatsu et al. 2008; Hicken et al. 2009, etc).  Due to the absence of a well-motivated fundamental theory, there have been many theoretical speculations regarding the nature of the exotic {\\em dark energy}, on whether it is a cosmological constant, a scalar or vector fields which provide a time varying dark-energy equation of state, usually parametrized by: \\begin{equation} p_Q= {\\rm w}(z) \\rho_Q\\;, \\end{equation} with $p_Q$ and $\\rho_Q$ the pressure and density of the exotic dark energy fluid and \\be {\\rm w}(z)={\\rm w}_0 + {\\rm w}_1 f(z) \\;, \\label{eqstatez} \\ee with w$_0=$w$(0)$ and $f(z)$ an increasing function of redshift [ eg., $f(z)=z/(1+z)$] (see Peebles \\& Ratra 2003 and references therein, Chevalier \\& Polarski 2001, Linder 2003, Dicus \\& Repko 2004; Wang \\& Mukherjee 2006). Of course, the equation of state could be such that w does not evolve cosmologically.  Two very extensive recent reports have identified {\\em dark energy} as a top priority for future research: \"Report of the Dark Energy Task Force (advising DOE, NASA and NSF) by Albrecht et al. (2006), and ``Report of the ESA/ESO Working Group on Fundamental Cosmology'', by Peacock et al. (2006).  It is clear that one of the most important questions in Cosmology and cosmic structure formation is related to the nature of {\\em dark energy} (as well as whether it is the sole interpretation of the observed accelerated expansion of the Universe) and its interpretation within a fundamental physical theory. To this end a large number of very expensive experiments are planned and are at various stages of development, among which the {\\em Dark Energy Survey} (DES: {\\tt http://www.darkenergysurvey.org/}), the {\\em Joint Dark Energy Mission} (JDEM: {\\tt http://jdem.gsfc.nasa.gov/}), {\\em HETDEX} ({\\tt http://www.as.utexas.edu/hetdex/}), Pan-STARRS: {\\tt http://pan-starrs.ifa.hawaii.edu}, etc.  Therefore, the paramount importance of the detection and quantification of {\\em dark energy} for our understanding of the cosmos and for fundamental theories implies that the results of the different experiments should not only be scrutinized, but alternative, even higher-risk, methods to measure {\\em dark energy} should be developed and applied as well. It is within this paradigm that our current work falls. Indeed, we wish to constrain the {\\em dark energy} equation of state using, individually and in combination, the Hubble relation and large-scale structure (clustering) methods, but utilizing alternative cosmic tracers for both of these components.  From one side we wish to trace the Hubble function using HII galaxies, which can be observed at higher redshifts than those sampled by current SNIa surveys and thus at distances where the Hubble function is more sensitive to the cosmological parameters. The HII galaxies can be used as standard candles (Melnick, Terlevich \\& Terlevich 2000, Melnick 2003; Siegel et al. 2005; Plionis et al. 2009) due to the correlation between their velocity dispersion, metallicity and $H_{\\beta}$ luminosity (Melnick 1978, Terlevich \\& Melnick 1981, Melnick, Terlevich \\& Moles 1988), once we reduce significantly their distance modulus uncertainties, which at present are unacceptably large for precision cosmology ($\\sigma_\\mu\\simeq 0.52$ mag; Melnick, Terlevich \\& Terlevich 2000). Furthermore, the use of such an alternative high-$z$ tracer will enable us to check the SNIa based results and lift any doubts that arise from the fact that they are the only tracers of the Hubble relation used to-date (for possible usage of GRBs see for example, Ghirlanda et al. 2006; Basilakos \\& Perivolaropoulos 2008)\\footnote{ GRBs appear to be anything but standard candles, having a very wide range of isotropic equivalent luminosities and energy outputs. Nevertheless, correlations between various properties of the prompt emission and in some cases also the afterglow emission have been used to determine their distances. A serious problem that hampers a straight forward use of GRBs as Cosmological probes is the intrinsic faintness of the nearby events, a fact which introduces a bias towards low (or high) values of GRB observables and therefore the extrapolation of their correlations to low-$z$ events is faced with serious problems. One might also expect a significant evolution of the intrinsic properties of GRBs with redshift (also between intermediate and high redshifts) which can be hard to disentangle from cosmological effects. Finally, even if a reliable scaling relation can be identified and used, the scatter in the resulting luminosity and thus distance modulus is still fairly large.}.  From the other side we wish to use X-ray selected AGN at a median redshift of $\\sim 1$, which is roughly the peak of their redshift distribution (see Basilakos et al. 2004; 2005, Miyaji et al. 2007), in order to determine their clustering pattern and compare it with that predicted by different cosmological models.  Although each of the previously discussed components of our project (Hubble relation using HII galaxies and angular/spatial clustering of X-ray AGN) will provide interesting and relatively stringent constraints on the cosmological parameters, especially under our anticipation that we will reduce significantly the corresponding random and systematic errors, it is the combined likelihood of these two type of analyses that enables us to break the known degeneracies between cosmological parameters and determine with great accuracy the {\\em dark energy} equation of state (see Basilakos \\& Plionis 2005; 2006; 2009).  Below we present the basic methodology of each of the two main components of our proposal, necessary in order to constrain the {\\em dark energy} equation of state. ",
        "conclusions": ""
    },
    "0911/0911.3151_arXiv.txt": {
        "abstract": "Stellar population synthesis (SPS) provides the link between the stellar and dust content of galaxies and their observed spectral energy distributions.  In the present work we perform a comprehensive calibration of our own flexible SPS (FSPS) model against a suite of data.  These data include ultraviolet, optical, and near--IR photometry, surface brightness fluctuations, and integrated spectra of star clusters in the Magellanic Clouds (MCs), M87, M31, and the Milky Way (MW), and photometry and spectral indices of both quiescent and post--starburst galaxies at $z\\sim0$. Several public SPS models are intercompared, including the models of Bruzual \\& Charlot (BC03), Maraston (M05) and FSPS.  The relative strengths and weaknesses of these models are evaluated, with the",
        "introduction": "\\label{s:intro}  The spectral energy distribution (SED) of a galaxy contains a wealth of information regarding its star formation history, dust content, and chemical abundance pattern.  These properties provide essential clues to the physical processes governing the formation and evolution of galaxies from high redshift to the present.  It is therefore highly desirable to have a robust method for extracting the physical properties of galaxies from their SEDs.  The process of translating observed SEDs into physical properties is, unfortunately, very challenging because it requires 1) an accurate understanding of all phases of stellar evolution, 2) well--calibrated stellar spectral libraries for converting stellar evolution calculations into measurable fluxes, 3) an initial mass function (IMF), specifying the weight given to each stellar mass, 4) detailed knowledge of the star--dust geometry in conjunction with an appropriate extinction curve; i.e., knowledge of the physical conditions of the interstellar medium (ISM).  Each of these requirements depend on chemical composition, further compounding the problem.  Combining these ingredients in order to predict the spectrum of a galaxy is known as stellar population synthesis (SPS), and has an extensive history \\citep[e.g.,][]{Tinsley76, Tinsley80, Bruzual83, Renzini86, Buzzoni89, Bruzual93, Worthey94, Maraston98, Leitherer99, Fioc97, Vazdekis99, Yi03, Bruzual03, Jimenez04, LeBorgne04, Maraston05, Schiavon07, Coelho07, Conroy09a, Molla09, Kotulla09}.  In the past two decades enormous progress has been made on each of the requirements mentioned above.  Yet, substantial uncertainties remain. A non--exhaustive list includes 1) the treatment of the core convective boundary in main sequence stars (i.e., convective core overshooting); 2) metallicity--dependent mass--loss along the red giant branch (RGB) and, relatedly, the morphology of the horizontal branch (HB); 3) the treatment of the thermally--pulsating asymptotic giant branch (TP--AGB) phase; 4) the spectral libraries, especially for M giants, TP--AGB stars, and stars at non--solar metallicities \\citep[e.g.,][]{Martins07}; 4) blue straggler (BS) stars and their ubiquity \\citep[e.g.,][]{Preston00, Li08a}; 5) the effects of binarism, which might be especially relevant for massive star evolution \\citep[e.g.,][]{Eldridge08} and BS stars; 6) the importance of rotation on massive star evolution \\citep[e.g.,][]{Meynet00}; 7) non--solar abundance ratios, which effect not only the stellar evolution calculations but also the stellar spectra \\citep[e.g.,][]{Coelho07}; 8) the unknown dependence of the IMF on ISM properties such as metallicity and pressure.  These uncertainties can in many cases dramatically impact the ability to convert observables into physical properties and vice--versa \\citep[e.g.,][]{Charlot96a, Charlot96b, Yi03a, Gallart05, Maraston06, LeeHC07, Conroy09a, Muzzin09, Conroy09c}.  Thankfully, there exists a wide array of data capable of constraining SPS models.  By far the most common type of object used for comparison is the star cluster\\footnote{Herein both open and globular clusters will be referred to as star clusters.}.  The approximately uniform age and metallicity of the stars within star clusters and the lack of internal reddening (except perhaps for very young clusters) affords direct comparison between them and the most basic ingredients in SPS --- the stellar evolution calculations and stellar spectral libraries. Comparisons to star clusters are either made in CMD space or in an integrated sense (i.e., the light from all the cluster stars are added together).  The former technique provides a sensitive probe of the main--sequence including the turn-off point, sub--giant and red giant branches \\citep[e.g.,][]{Worthey94, Bruzual03, An09}, while the latter technique provides a more robust measure of the rarer brightest stars that dominate the integrated light from the cluster.  It is the latter method of comparison that is more relevant if the goal of SPS modeling is to understand the integrated light from galaxies.  For this reason, the integrated colors, surface brightness fluctuations (SBFs), and spectra of star clusters have been used extensively to compare and calibrate SPS models \\citep[e.g.,][]{Bruzual03, Gonzalez04, Maraston05, Cohen07, Cordier07, Pessev08, Marigo08, Koleva08, LeeHC09b}.  There are two significant sources of concern when attempting to calibrate SPS models with star cluster data.  The first concern is that the brightest stars, which dominate the integrated light, are rare, and thus stochastic effects must be carefully modeled. Recently, this issue has been addressed observationally by stacking clusters in bins of age in order to synthesize `superclusters' that are relatively unaffected by stochastic effects \\citep{Pessev06, Gonzalez04}.  The second concern is less tractable, even in principle, and arises from the fact that the primary sources of star clusters for which ages and metallicities can be reliably estimated are the Milky Way (MW), M31, and the Large and Small Magellanic Clouds (LMC and SMC, respectively).  This fact imposes severe restrictions on the region of the age--metallicity parameter space that can be constrained with star cluster data.  For example, one of the most important regions of this parameter space for studying galaxies --- old and metal rich --- contains very few star clusters (except in the bulge of the MW, although high extinction diminishes the utility of star clusters there).  In Local Group star clusters, a bin in age contains clusters of a relatively narrow range in metallicity owing to the simple fact that the metallicity of galaxies tends to increase with age.  Because of these issues, entire galaxies are often also used to assess the accuracy of SPS models \\citep[e.g.,][]{Worthey94, Bruzual03, Maraston06, Eminian08}.  Galaxies are not subject to the two concerns mentioned above, but using them to constrain SPS models can be very difficult because 1) galaxies contain stars of a range of ages and metallicities; and 2) starlight in galaxies is attenuated by interstellar dust.  Nevertheless, if subsamples of galaxies are carefully constructed with known and regular properties, then galaxies can be used to assess the reliability of SPS models.  The present work seeks to provide a comprehensive comparison between SPS models and observed star clusters and galaxies.  Such a comparison is timely because of an abundance of new, high--quality data, including near--IR photometry and SBFs of star clusters in the LMC and SMC, UV photometry of clusters in M31 and M87, CMDs and integrated spectra of clusters in the MW, and optical and near--IR photometry and spectral indices for large samples of low--redshift galaxies.  We will consider not only our own SPS model \\citep{Conroy09a} but also the commonly used models of \\citet[][BC03]{Bruzual03} and \\citet[][M05]{Maraston05}.  Such a detailed comparison between these popular models has not been undertaken until now, and will thus provide a much--needed evaluation of the relative strengths and weaknesses of these models.  In addition, the flexible nature of our own model will be exploited to explore the extent to which various uncertain phases of stellar evolution, including thermally--pulsating asymptotic giant branch stars (TP--AGB), blue HB stars, and post--AGB stars, can be constrained by the data.  This task is critical if we are to have confidence in the derived physical properties of galaxies, such as stellar masses and star formation rates.  We proceed as follows. $\\S$\\ref{s:data} contains a detailed description of the star cluster and galaxy data used to constrain the models, which are themselves described in $\\S$\\ref{s:model}.  In $\\S$\\ref{s:calib} we present an extensive comparison between models and data.  A discussion and summary are provided in $\\S$\\ref{s:disc} and $\\S$\\ref{s:sum}, respectively.  The zero points of the star cluster UBVRIJHK magnitudes are in the Vega system; all other magnitudes are quoted in the $AB$ system \\citep{Oke83}.  Where necessary, we adopt a flat $\\Lambda$CDM cosmology with $(\\Omega_m, \\Omega_\\Lambda,h)=(0.26,0.74,0.72)$. ",
        "conclusions": "\\label{s:sum}  We now summarize our principle results.  \\begin{itemize}  \\item  The latest stellar models from the Padova group cannot fit the near--IR photometry and surface brightness fluctuations of star clusters in the MCs.  With our flexible SPS (FSPS) code we have modified the TP--AGB phase in the Padova isochrones to produce better fits to the data.  The M05 models also fail to reproduce the data, both because the colors are too red and the age--dependence is incorrect.  The BC03 and the FSPS model using the BaSTI isochrones fare well, although both models are somewhat too blue in the near--IR.  A significant limiting factor in more accurate calibrations is the difficulty in constructing reliable photometry of MC clusters, owing to the fact that TP--AGB stars are luminous and rare.  \\item  At low metallicities (log$(Z/Z_\\Sol)<-0.5$) and old ages, all models predict similar $UBVRIJHK$ colors (the typical variation between models is $\\approx0.05$ mags).  When compared to MW and M31 star clusters, the models are 0.2 mags too blue in $V-K$ at log$(Z/Z_\\Sol)>-1.0$, and 0.1 mags too red in $J-K$ at all $Z$.  The origin of the discrepancies in the near--IR is unclear.  Decreasing the temperature of the RGB by $\\sim200$K alleviates much of the disagreement.  If this modification is isolated to the brightest stars, then such a modification could improve agreement with integrated photometry without compromising agreement in CMD--space. The FSPS model with the BaSTI isochrones performs worst at low metallicities and old ages of the models considered (although the disagreement is not dramatic), while the FSPS model with the Padova isochrones performs well.  Multi--color CMDs of two metal--rich clusters, NGC 6791 and NGC 188, are well--fit by FSPS.  Spectral indices of MW star clusters are generally well--fit by both FSPS and BC03, although the models predict D$_n4000$ strengths too large and H$\\delta_A$ strengths too weak compared to two log$(Z/Z_\\Sol)\\approx0.0$ clusters.  \\item  The ultraviolet photometry of the MW, M87, and M31 star clusters can be well--fit by FSPS because FSPS contains flexible treatments of the post--AGB and horizontal branch evolutionary phases.  This flexibility is essential given the substantial scatter in the ultraviolet data at fixed metallicity, and given our inadequate theoretical understanding of these phases.  The M05 and BC03 models perform less well because of their incomplete treatment of these advanced evolutionary phases.  \\item  SPS models are compared to $ugrzYJHK$ photometry of massive red sequence galaxies.  The FSPS, BC03, and M05 models fare well in most respects, with a few exceptions: all models are too blue in $J-H$; the M05 model is somewhat too blue in $Y-K$; all models are far too red in $g-r$ and $u-g$.  These conclusions hold under the assumption that red sequence galaxies are composed exclusively of old metal--rich stars with canonical evolutionary phases.  The disagreement in the $ugr$ colors can be alleviated if some combination of young ($\\sim1$ Gyr) stars, metal--poor stars, and blue straggler stars are added in small amounts.  \\item  Optical spectral indices of massive galaxies are generally well--fit by the FSPS and BC03 models, with the important exceptions that both models predict D$_n4000$ in excess of observations and neither fit the observed locus of galaxies in the D$_n4000-$H$\\delta_A$ plane. These conclusions hold for the same assumptions noted above, and, as above, small additions of some combination of young stars, metal--poor stars, blue straggler and horizontal branch stars can remedy this disagreement.  \\item  The FSPS and BC03 models adequately describe the $grizYK$ colors of K+A, or `post--starburst' galaxies, while the M05 model performs less well.  Such galaxies contain a large proportion of intermediate age ($0.5<t<2$ Gyr) stars and thus provide unique constraints on the importance of TP--AGB stars in galaxies.  \\end{itemize}"
    },
    "0911/0911.4771_arXiv.txt": {
        "abstract": "{% We give a pedagogical introduction to two aspects of magnetic fields in the early universe. We first focus on how to formulate electrodynamics in curved space time, defining appropriate magnetic and electric fields and writing Maxwell equations in terms of these fields. We then specialize to the case of magnetohydrodynamics in the expanding universe. We emphasize the usefulness of tetrads in this context. We then review the generation of magnetic fields during the inflationary era, deriving in detail the predicted magnetic and electric spectra for some models. We discuss potential problems arising from back reaction effects and from the large variation of the coupling constants required for such field generation. } ",
        "introduction": "The universe is magnetized, from scales of planets to the largest collapsed objects like galaxy clusters. The origin and evolution of these magnetic fields is a subject of intense study. Much of the work on magnetic field origin centers on the idea that small seed magnetic fields are generated by purely astrophysical batteries and are subsequently amplified to the observed levels by a dynamo (cf. Brandenburg and Subramanian, 2005 for a review). An interesting alternative is that the observed large-scale magnetic fields are partially a relic field from the early Universe, which have been further amplified by motions.  We provide here a pedagogical review of two aspects of magnetohydrodynamics (MHD) in the early universe. In the first part of the review we focus on how to formulate electrodynamics in curved space time, especially how to define magnetic and electric fields and write Maxwell equations in terms of these fields. This issue may perhaps be well known to relativists but may not be so familiar to practitioners of MHD. The perspective provided is that of the author.  In the second part, we review, again in a pedagogical manner, generation of magnetic fields during the inflationary era. It is well known that scalar (density or potential) perturbations and gravitational waves (or tensor perturbations) can be generated during inflation. Could magnetic field perturbations also be generated? Indeed, a large number of mechanisms whereby magnetic fields are generated in the early universe have been discussed in the literature (Turner \\& Widrow 1988; Ratra 1992; Widrow 2002; Giovannini 2007). These generically involve the breaking of the conformal invariance of the electromagnetic action, and the predicted amplitudes are rather model dependent. Nevertheless, if a primordial magnetic field is generated, with a present-day strength of $B \\sim 10^{-9}$~G and coherent on Mpc scales, it can strongly influence early galaxy formation (cf. Sethi and Subramanian, 2005), or induce signals on the cosmic microwave backround radiation (CMBR) (cf. Subramanian, 2006 for a review) including CMBR non-Gaussianity (cf. Seshadri and Subramanian, 2009; Caprini et al., 2009). An even weaker field, sheared and amplified due to flux freezing, during galaxy and cluster formation may kick-start the dynamo. It is then worth considering if one can generate such primordial fields. Our discussion of inflationary generation of magnetic fields follows the standard literature, but again the perspective offered is that of the author. We hope that our pedagogical discussion of these aspects of MHD in the early universe will be useful to those entering the subject. ",
        "conclusions": "We end with a few comments. We have emphasized the use of tetrads in defining properly behaved magnetic and electric fields. Such a procedure is also followed when using with the $3+1$ formalism in the context of black-hole electrodynamics (cf. Macdonald et al., 1986). Thus it seems to us somewhat surprising why this does not usually get mentioned when dealing with MHD in the context of cosmology. Regarding magnetic field generation during inflation, we have a paradigm, but no compelling model as yet. Clearly more work is required to find such a model, which at the same time avoids the problems of back reaction, the strong coupling problem mentioned by DMR, or the strong decay of coupling constants.  Another possibility is generation of primordial fields is during a later phase transition, like the Electro-Weak or the quark-hadron transition. Here the main problem is that the generated field typically has a tiny correlation scale, less than the Hubble radius $H^{-1}$, at the epoch of the phase transition, unless magnetic helicity is also generated. Interestingly there are several ideas for such helicity generation during the Electro-weak phase transition (cf. Vachaspati, 2001; Diaz-Gil et al., 2008; Copi et al., 2008). Magnetic energy decay conserving helcity can then lead to an inverse cascade and larger coherence scales, as first emphasized in the early universe context by Brandenburg, Enquist and Olesen (1996) (see also Christensson, Hindmarsh and Brandenburg, 2001; Banerjee and Jedamjik, 2005), but that is a story for another review."
    },
    "0911/0911.0687_arXiv.txt": {
        "abstract": "\\noindent Upcoming gravitational wave (GW) detectors might detect a stochastic background of GWs potentially arising from many possible sources, including bubble collisions from a strongly first-order electroweak phase transition.  We investigate whether it is possible to connect, via a semi-analytical approximation to the tunneling rate of scalar fields with quartic potentials, the GW signal through detonations with the parameters entering the potential that drives the electroweak phase transition.  To this end, we consider a finite temperature effective potential similar in form to the Higgs potential in the Standard Model (SM).  In the context of a semi-analytic approximation to the three dimensional Euclidean action, we derive a general approximate form for the tunneling temperature and the relevant GW parameters.  We explore the GW signal across the parameter space describing the potential which drives the phase transition.  We comment on the potential detectability of a GW signal with future experiments, and physical relevance of the associated potential parameters in the context of theories which have effective potentials similar in form to that of the SM.  In particular we consider singlet, triplet, higher dimensional operators, and top-flavor extensions to the Higgs sector of the SM.  We find that the addition of a temperature independent cubic term in the potential, arising from a gauge singlet for instance, can greatly enhance the GW power.  The other parameters have milder, but potentially noticeable, effects. ",
        "introduction": "Gravitational wave (GW) interferometers currently taking data, and under development or construction, have as a possible target a stochastic background of GWs of cosmological origin (for reviews, see\\cite{thorne,allen:1996vm,Maggiore:1999vm,Maggiore:2000gv}, for recent results, see \\cite{ligonature}). Since GWs propagate freely through the Universe after being produced, their detection provides a powerful diagnostics for the physics of the Early Universe. Several mechanisms that might generate such GWs have been discussed, including quantum fluctuations during or shortly after inflation \\cite{Turner:1990rc,hogan,Witten:1984rs}, cosmic strings \\cite{Steinhardt:1981ct}, cosmological magnetic fields \\cite{Gyulassy:1983rq,VanHove:1983rk,Enqvist:1991xw,Huet:1992ex}, plasma turbulence \\cite{landau,Abney:1993ku,weinberggr,Turner:1992tz} and bubble wall collisions during first-order phase transitions \\cite{Kosowsky:1992rz,Kosowsky:1991ua,Kosowsky:1992vn,Coleman:1977py,Callan:1977pt,Kamionkowski:1992dc}.  Specifically, the space interferometer LISA, expected to fly, or to be nearing completion, over the next decade will have a sensitivity peak for GWs with a frequency between $10^{-4}$ and $1$ Hz \\cite{Maggiore:2000gv}. Quite fortunately, this range corresponds to the frequency range expected today, after redshifting, from GWs produced at a temperature $T \\sim 100\\ {\\rm GeV}\\sim E_{\\rm EW}$, where the latter symbol indicates the electroweak scale. Intriguingly enough, the production of GWs at $T\\sim E_{\\rm EW}$ is indeed expected if the electroweak phase transition is strongly first order. In turn, this is a necessary condition for the success of scenarios where the baryon asymmetry of the Universe is produced at the electroweak phase transition (electroweak baryogenesis, EWB; for a review, see \\cite{Riotto:1999yt}).  However, there is a generic tension between the requirement of a small bubble wall velocity needed in the context of electroweak baryogenesis models (see e.g.~\\cite{bauref1,bauref2,bauref3}) and the super-sonic bubble velocity values (detonation) under investigation here.  Depending on the sources of net baryon number, super-sonic bubble wall velocities might be compatible with a baryon asymmetry produced at the EWPT, although this is generally not the case.  We are at the dawn of the exciting period of exploring the electroweak scale with the possibility to connect such diverse experimental endeavors as GW detection and the Large Hadron Collider.  These grand experimental enterprises may also help answer the question of how the baryon asymmetry arose in the Early Universe, which is all the more exciting.   In its minimal setup, the strength of the electroweak phase transition in the Standard Model only depends only on the mass of the Higgs bosons (for a review, see \\cite{tempreview}).  For the Higgs mass range compatible with searches at LEP-II, non-perturbative lattice computations indicate that there is no phase transition at all, but rather a smooth crossover \\cite{Kajantie:1996mn}.  No GW production is thus expected at the EW phase transition, if the electroweak sector corresponds to the minimal Standard Model.  In addition, if this is the case, EWB cannot be the mechanism underlying the production of the baryon asymmetry in the Universe.  Fortunately, the strength of the EW phase transition is actually strongly model-dependent and parameter-dependent (for a review, see \\cite{Rubakov:1996vz}), and numerous extensions of the Standard Model do predict a strongly first order EW phase transition. If this is further extended with the other necessary ingredients for EWB, the EW scale might indeed be responsible for generating the observed baryon-antibaryon asymmetry.  If the EW phase transition is strongly first order, the universe finds itself trapped, at $T\\lesssim E_{\\rm EW}$, in a metastable EW unbroken phase where the vacuum state of the universe is a false vacuum (i.e.~it is not the lowest energy state) as the universe cools down and its temperature decreases.  A potential energy barrier exists between the false and the true (lowest energy) electroweak symmetry breaking vacuum.  Quantum mechanical tunneling produces bubbles of true vacuum (broken EW phase), which then expand, collide, and combine to fill the universe with the true vacuum.  GWs can be abundantly produced at the EW phase transition, primarily through bubble collisions \\cite{Witten:1984rs,hogan,Turner:1990rc,Kosowsky:1992rz,Kosowsky:1991ua,Turner:1992tz,Kosowsky:1992vn,Kamionkowski:1993fg}.  This mechanism has been extensively investigated, analytically and numerically, and in relation to the effective potential that drives the transition itself.  In particular, the possibility of a GW signal from the EW phase transition has attracted, not surprisingly, a great deal of attention.  In view of the above mentioned large model-dependence, though, most of the existing literature has been devoted to either special cases and particular corners of parameter space (such as the light stop scenario in the context of the minimal supersymmetric extension of the Standard Model, MSSM; for a review, see \\cite{Quiros:2000wk}), and/or to accurate but numerical studies only \\cite{Ashoorioon:2009nf,gwbubble}.  Alternatively, model independent results on the detectability of a GW signal from a strongly first order phase transition have been derived assuming a few relevant parameters for the GW dynamics could be computed from the effective potential that drives the particular phase transition under consideration (for recent related work, see e.g.~\\cite{Grojean:2006bp,Caprini:2007xq,Kahniashvili:2008pf}).  The scope of the present study is to use semi-analytical results of the tunneling rate of scalar fields with quartic potentials \\cite{se3approx} to predict {\\em analytically} the strength of the EW phase transition and the GW signal in extensions of the Standard Model EW sector that can be characterized with the dynamics of a single order parameter, a scalar field $\\phi$.  These models include simple generalizations of the SM effective potential, in appropriate dynamical regimes.  The main result of this paper is a closed analytic formula for two parameters describing the GW signal as a function of parameters appearing in the effective potential of the scalar field driving the EW phase transition.   The rest of this paper is organized as follows: in Section \\ref{sec:backgrnd} we summarize the relevant physics and definitions of the quantities we will be studying, in particular an approximation for the three dimensional Euclidean action, which is the basis for computing the finite temperature tunneling rate.  Then, in Section \\ref{sec:approx}, we derive an approximation for the tunneling temperature and present exact and approximate formulas for the GW parameters for any effective potential similar in form to the SM case.  Following that, in Section \\ref{sec:param} we constrain some parameters in order to reproduce the usual electroweak symmetry breaking pattern, and plot the effect of the various parameters in the potential on the parameters driving the GW signal.  We examine the physical relevance of the parameter space and the detectability of the GW when varying various parameters beyond their SM values.  In Section \\ref{sec:models} we describe specific examples where our formalism can be applied in certain regimes, including $SU(2)$ triplet and singlet extensions to the Higgs sector, top-flavor models, and models encompassing higher dimensional operators. % ",
        "conclusions": "Upcoming experiments may soon observe the first GW signals, including those from early universe processes.  Such experiments will deliver the very exciting prospect of observing GWs from the electroweak phase transition, if it is strongly first-order.  While there have been many studies on this topic, they largely involve numerical computations of the phase transition temperature and GW parameters.  In this work we have presented an analysis of a generic effective potential which is motivated by the SM Higgs potential and an additional contribution from a gauge singlet (or other possible origins) in the form of a temperature-independent cubic coupling.  By approximating the tunneling temperature to be very close to the critical temperature (based on the general form of the action), we have derived an expression for the tunneling temperature based on the parameters of the effective potential.  Using this result, the GW parameters also have expressions depending on the parameters of the potential.  Once these expressions were found, the parameter space of the potential was explored starting from the SM values.  While all the parameters can have noticeable effects, $e$, motivated from singlet models, easily has the greatest effect.  Following this, we showed how this parameter arises from a decoupling analysis of a simple SM plus singlet model.  Other models can also be analyzed with our effective potential.  The addition of an additional $SU(2)$ triplet affects $\\lambda$, while higher dimensional operators can also affect $E$.  Increasing $E$ by an order of magnitude from its SM value could also produce a strongly first-order phase transition.  The topflavor model also has an effective potential of the same form as the SM for its earlier phase transition, which can also be analyzed through this work.  In this case, the GW signal would not be from the (later) electroweak phase transition.  Finally, although the effective potential studied has clear origins in the SM and singlet extensions, the potential is indeed very general.  A low energy, effective theory attempting to model the electroweak phase transition will likely closely model this form of the potential.  Additionally, the potential is a high-temperature approximation to a one-loop, finite temperature, effective potential of a scalar field with gauge bosons, fermions, and singlets.  So in this sense its form is also generic and expected to be common in a phase transition with a scalar field order parameter.  The potential applications of the results presented include applying it to other models for the electroweak (or other) phase transition.  By approximating or otherwise finding a way to put an effective potential in the form we analyzed, the tunneling temperature and GW spectrum is now easily obtained through the above formulae.  % Besides being significant itself, the detection of a stochastic background of GWs may provide insight into the specifics of models of the electroweak phase transition, and the fundamental mechanism underlying the generation of the baryon asymmetry of the universe."
    },
    "0911/0911.5588_arXiv.txt": {
        "abstract": "The magnetic star HD 21699 possesses a unique magnetic field structure where the magnetic dipole is displaced from the centre by 0.4 $\\pm$ 0.1 of the stellar radius (perpendicularly to the magnetic axis), as a result, the magnetic poles are situated close to one another on the stellar surface with an angular separation of 55\\grad and not 180\\grad  as seen in the case of a centred dipole. Respectively, the two magnetic poles form a large \"magnetic spot\". High-resolution spectra were obtained allowing He I and Si II abundance variations to be studied as a function of rotational phase. The results show that the helium abundance is concentrated in one hemisphere of the star, near the magnetic poles and it is comparatively weaker in another hemisphere, where magnetic field lines are horizontal with respect to the stellar surface. At the same time, the silicon abundance is greatest between longitudes of 180 - 320\\grad, the same place where the helium abundance is the weakest. These abundance variations (with rotational phase) support predictions made by the theory of atomic diffusion in the presence of a magnetic field. Simultaneously, these result support the possibility of the formation of unusual structures in stellar magnetic fields. Analysis of vertical stratification of the silicon and helium abundances shows that the boundaries of an abundance jump (in the two step model) are similar for each element; $\\tau_{5000}$ = 0.8-1.2 for helium and 0.5-1.3 for silicon. The elemental abundances in the layers of effective formation of selected absorption lines for various phases are also correlated with the excitation energies of low transition levels: abundances are enhanced for higher excitation energy and higher optical depth within the applied model atmosphere. ",
        "introduction": "HD~21699 (HR~1063) was initially classified by Roman and Morgan (1950) as a B8IIIvar star. Molnar (1972) later performed an analysis of its spectra and determined that helium was extremely deficient (by factor of 5), as a result he classified this star as a He-weak. Shore et al. (1987) refer to HD~21699 as a He-weak silicon star (sn class - with broad and diffuse He I lines). The Sn class is primarily attributed by Morgan (1977) to the silicon and He-weak star HD 5737. In the classification of Preston (1974) both objects are categorised as chemically peculiar (CP2) stars. Because of the extremely reduced helium abundance, it is possible that HD~21699 has been given an incorrect spectral classification: the MK classification gives Sp = B8, whereas Abt et al. (2002) suggest a spectral class of B8IIIpMn that corresponds to \\Tef = 12000 K, while Glagolevskij (2002) derived \\Tef = 16100 K from the analysis of colour indices. Using the spectra obtained with the 6-m telescope, the following parameters were derived by Glagolevskij et al. (2006) by an analysis of the $H_{\\delta}$ line: \\Tef = 16000 K, lg g = 4.15, \\Vt = 0.8 km\\,$s^{-1}$.   The period of axial rotation for HD~21699 is P = $2^d.4765$ (Brown et al. 1985). Using the relation between the equatorial velocity (km\\,$s^{-1}$), period (days), stellar radius in solar radii ($V =  50.613* R/P$) and the measured value of \\vsini =  35 km\\,$s^{-1}$ yields an inclination angle of $i$ = $32^{\\circ}$ % (Glagolevskij \\& Chuntonov 2007). The positive magnetic field maximum occurs at the phase $\\phi= 0$, and the negative extremum occurs at $\\phi=  0.4$, according to the ephemerides of Glagolevskij \\& Chuntonov (2007) for initial phase; (JD = 2445595.529 + $2^d.49246$)         \\vskip2mm ",
        "conclusions": "For the first time photospheric chemical element abundances of HD 21699 were obtained during the whole rotational period. { Magnetic splitting of spectral lines and elements stratification in the atmosphere were taken into account with line profile modelling. Helium abundance is weakened for the whole stellar surface, which is usual for He-weak stars. Nevertheless, helium has higher abundance in the region of the \"magnetic spot\" due to the influence of stellar wind, which elevates the helium to the outer layers of the stellar atmosphere, and helium abundance should increase with optical depth in the atmospheres of He-weak stars. (Vauclair et al. 1991).   Silicon is accumulated in the part of the star where the magnetic field lines are predominantly horizontal, as it was predicted by Vauclair et al. (1979), Alecian and Vauclair (1981)  and Meggessier (1984). Silicon abundance determined from lines with low excitation energies (8-10eV) appears to be lower in the area of the magnetic poles and is approximately solar in the region with horizontal magnetic lines. The lines with high excitation energies (12-16eV) also show enhancement of silicon abundance for the region with horizontal magnetic lines, but this abundance is significantly higher than solar everywhere on the stellar surface. Silicon abundance is lower in the outer parts of the atmosphere and higher in the deeper layers.   The optical depth of the abundance jump (in two-steps model) is near to  $\\tau_{5000}$ = 1. In the cases of He I and Si II, we can see a tendency to change $\\tau_{5000}$ with the phase contrary to abundances, which could be the result of competition between the stellar wind, diffusion, and the local magnetic field geometry that would lead to accumulations of helium and silicon at different depths in the atmosphere.   Since IUE observations of the C IV resonance doublet show existence of magnetically controlled stellar wind in HD 21699 (Brown et al. 1985), one should expect some variations in Balmer lines, similar to that seen in H-alpha in 36 Lyn (Wade et al. 2006) and in HD 37479 (Smith et al. 2006) and in many Balmer lines in HD 37479 (Smith and Bohlender, 2007). Unfortunately, due to the incomplete number of observed phases it is difficult to find similar changes in Balmer lines in HD 21699."
    },
    "0911/0911.0014_arXiv.txt": {
        "abstract": "Leading models of galaxy formation require large-scale energetic outflows to regulate the growth of distant galaxies and their central black holes. However, current observational support for this hypothesis at high redshift is mostly limited to rare $z>2$ radio galaxies. Here we present Gemini-North NIFS Intregral Field Unit (IFU) observations of the [O~{\\sc iii}]$\\lambda$5007 emission from a $z\\approx$~2 ultraluminous infrared galaxy ($L_{\\rm IR}>10^{12}$~$L_{\\odot}$) with an optically identified Active Galactic Nucleus (AGN). The spatial extent ($\\approx$~4--8~kpc) of the high velocity and broad [O~{\\sc iii}] emission are consistent with that found in $z>2$ radio galaxies, indicating the presence of a large-scale energetic outflow in a galaxy population potentially orders of magnitude more common than distant radio galaxies. The low radio luminosity of this system indicates that radio-bright jets are unlikely to be responsible for driving the outflow. However, the estimated energy input required to produce the large-scale outflow signatures (of order $\\approx10^{59}$~ergs over $\\approx$~30~Myrs) could be delivered by a wind radiatively driven by the AGN and/or supernovae winds from intense star formation. The energy injection required to drive the outflow is comparable to the estimated binding energy of the galaxy spheroid, suggesting that it can have a significant impact on the evolution of the galaxy. We argue that the outflow observed in this system is likely to be comparatively typical of the high-redshift ULIRG population and discuss the implications of these observations for galaxy formation models. ",
        "introduction": "The most successful current models of galaxy formation invoke large-scale energetic outflows to explain many of the properties of local massive galaxies and the intragalactic medium (IGM; i.e.,\\ red optical colours of massive galaxies, steep optical galaxy luminosity function, the black-hole--spheroid mass relationship, metal-enrichment of the IGM; e.g.,\\ Silk \\& Rees 1998; Fabian 1999; King 2003; Benson et~al. 2003; Granato et~al. 2004; Di Matteo et~al. 2005; Springel et~al. 2005; Bower et~al. 2006; Croton et~al. 2006; Hopkins et~al. 2006). A key attribute of these simulated outflows is the injection of significant amounts of kinetic energy into the interstellar medium (ISM), which can inhibit and terminate star formation by either heating the ISM or ejecting the gas out of the gravitational potential of the host galaxy. Large-scale outflows can be powered by star formation and/or Active Galactic Nuclei (AGN) activity (e.g.,\\ Heckman et~al. 1990; Crenshaw et~al. 2003; Veilleux et~al. 2005), although most galaxy formation models predict that only AGN-driven outflows will be sufficiently energetic to have significant impact on the formation and evolution of massive galaxies. These large-scale outflows are expected to be particularly effective at $z\\approx$~2, when star formation in the most massive galaxies was in significant decline (e.g.,\\ Juneau et~al. 2005; Panter et~al. 2007; P{\\'e}rez-Gonz{\\'a}lez et~al. 2008; Damen et~al. 2009).  Broadly speaking, AGN-driven outflows are kinematically energetic ``winds'' or ``jets'' that are initially launched close to the central black hole. The two main catalysts that have been proposed to power an AGN-driven outflow are (1) a radio jet or lobe, and (2) a radiatively driven wind. A radio-jet/lobe driven outflow will be most prevalent in radio-loud AGN but might only occur in a minority of the AGN population.\\footnote{The typical definition of a radio-loud AGN is ${\\nu}{L_{\\rm 5GHz}}$/${\\nu}{L_{\\rm 440nm}}\\simgt$~10 (e.g.,\\ Kellerman et al. 1989).} For example, $\\approx$~4--24\\% of the optically bright (radiatively strong) AGN population are radio loud, which appears to be a function of both redshift and AGN luminosity (e.g.,\\ Hooper et~al. 1995; Jiang et~al. 2007), although this may indicate the ``duty cycle'' of radio AGN activity. The fraction of optically faint (radiatively weak) radio-loud AGN fraction is a strong function of stellar mass (e.g.,\\ Best et~al. 2005; Smol{\\v c}i{\\'c} et~al. 2009). By comparison, a radiatively driven wind would be most effective in luminous AGNs with high mass-accretion rates, since the wind is directly driven by radiation pressure from the accretion disk.  High signal-to-noise ratio X-ray and ultra-violet absorption-line spectroscopy have shown that high-velocity outflows are present in many AGNs, and may be a ubiqitious property of the AGN population at both low and high redshift (e.g.,\\ Crenshaw et~al. 1999: Chartas et~al. 2002, 2007a,b; Laor \\& Brandt 2002; Pounds et~al. 2003; Reeves et~al. 2003; Porquet et~al. 2004; Ganguly \\& Brotherton 2008; Gibson et~al. 2009). For example, at least $\\approx$~60\\% of unobscured AGNs show evidence for high-velocity outflows, and the maximum measured outflow velocity is a strong function of the AGN luminosity (e.g.,\\ Crenshaw et~al. 1999; Ganguly \\& Brotherton 2008). The estimated energy injection required to produce the X-ray absorption features (which sometimes suggest outflow velocities exceeding $v\\approx$~0.1$c$) is typically $\\approx$~0.1--1 of the AGN bolometric luminosity (e.g.,\\ Pounds et~al. 2003; Reeves et~al. 2003; Chartas et~al. 2007a; Pounds \\& Reeves 2009). However, X-ray variability studies and consideration of the required energy input to produce the high-velocity features, indicate that these outflow signatures must be produced in the vicinity of the accretion disc on $<$~1~pc scales (e.g.,\\ Crenshaw et~al. 2003; King \\& Pounds 2003). Therefore, while it is clear that energetic outflows are present close to the accreting black hole of AGNs, it is far from clear what impact these outflows will have on the gas and star formation in the host galaxy on $\\approx$~1--10~kpc scales. To directly test the impact that outflows have on the formation and evolution of galaxies, it is necessary to identify large-scale energetic outflows, particularly in high-redshift ($z\\simgt2$) massive galaxies where they are predicted to have been most prevalent.  Optical spectroscopy of distant galaxies have revealed blueshifted absorption/emission lines in many systems, a signature of outflowing gas (e.g.,\\ Pettini et~al. 2001; Shapley et~al. 2003; Tremonti et~al. 2007). However, the most direct insight into the identification and interpretation of large-scale outflows is provided by spatially resolved spectroscopy (integral-field unit; IFU observations), which provide direct information about both the velocity {\\it and} extent of any outflowing gas (e.g.,\\ Swinbank et~al. 2005, 2006; Nesvadba et~al. 2006). Using rest-frame optical IFU observations of distant radio-loud AGNs (i.e.,\\ high-redshift radio galaxies, hereafter referred to as HzRGs), Nesvadba et~al. (2006, 2007a, 2008) showed that large-scale energetic outflows are present in at least a fraction of the high-redshift galaxy population. The key diagnostic that revealed an AGN-driven outflow in these systems was the presence of kinematically complex and extended [O~{\\sc iii}]$\\lambda$5007 emission;\\footnote{Emission from [O~{\\sc iii}] can be produced by a number of processes, including photo-ionisation and shocks (Osterbrock 1989) but it cannot be produced in dense environments, such as the broad-line region of AGN, without being collisionally de-excited (e.g.,\\ Osterbrock 1989; Robson 1996). The detection of broad (FWHM$\\simgt$~500--1000~km$^{-1}$) [O~{\\sc iii}] emission therefore indicates the presence of off-nuclear kinematic components.} see Holt et~al. (2008) for similar constraints on $z\\simlt0.6$ radio galaxies. Nesvadba et~al. (2008) showed that the extent of the broad [O~{\\sc iii}] emission is similar to that of the radio emission, providing evidence for a causal connection between the radio emission and a large-scale outflow. The estimated kinetic energy required to drive the outflows in these HzRGs is $\\approx$~1--40\\% of that potentially provided by the radio jet, indicating that they could be powered by mechanical energy from the radio jet. However, HzRGs are rare systems (co-moving space densities at $z\\approx$~2 of $\\Phi\\approx10^{-8}$--$10^{-7}$~Mpc$^{-3}$; e.g.,\\ Willott et~al. 1998), possibly representing the most massive galaxies in the most biased regions of the distant Universe (e.g.,\\ Stevens et~al. 2003; Seymour et~al. 2007). To fully assess the global impact of large-scale energetic outflows on the formation and evolution of distant galaxies, and to distinguish between the different mechanisms proposed to drive outflows, it is therefore necessary to search for large-scale energetic outflows in the more typical radio-quiet AGN population.   In this paper we present Gemini-North NIFS IFU observations of the [O~{\\sc iii}] emission from a $z\\approx$~2 quasar (SMM J123716.01+620323.3, hereafter referred to as SMM~J1237+6203) that is hosted in an Ultraluminous Infrared (IR) Galaxy (ULIRG; $L_{\\rm IR}>10^{12}$~$L_{\\odot}$; e.g.,\\ Sanders \\& Mirabel 1996). ULIRGs (both at low and high redshift) are rapidly evolving galaxies undergoing intense dust-obscured star formation and AGN activity (e.g.,\\ Sanders \\& Mirabel 1996; Genzel et~al. 1998; Page et~al. 2001; Alexander et~al. 2005b; Stevens et~al. 2005; Schweitzer et~al. 2006; Coppin et~al. 2008; Lutz et~al. 2008; Pope et~al. 2008). Ultra-luminous infrared quasars are often believed to be a key phase in the evolution of a ULIRG, where energetic outflows are starting to shut down star formation in the host galaxy, before revealing an unobscured quasar (e.g.,\\ Sanders et~al. 1988; Canalizo \\& Stockton 2001; Granato et~al. 2004; Hopkins et~al. 2005a; Kawakatu et~al. 2006). It is possible that every massive galaxy in the local Universe underwent a ULIRG phase at some time over the past $\\approx$~13~Gyrs during which the galaxy spheroid and black hole were predominantly grown (e.g.,\\ Swinbank et~al. 2006; Alexander et~al. 2008; Tacconi et~al. 2008). The [O~{\\sc iii}] luminosity of SMM~J1237+6203 is comparable to the HzRGs studied by Nesvadba et~al. (2006, 2007a, 2008), suggesting similar intrinsic AGN luminosities. However, importantly, SMM~J1237+6203 is $\\approx$~3--4 orders of magnitude fainter at radio luminosities than the HzRGs. We identify the signatures of a large-scale energetic outflow in this ultraluminous infrared quasar, showing that these features are not restricted to the distant radio-loud AGN population. We estimate the energy input required to produce these features, assess what could be responsible for driving this large-scale outflow, and discuss the implications for galaxy formation models. We have adopted $H_{0}=71\\kms$, $\\Omega_{M}=0.27$ and $\\Omega_{\\Lambda}=0.73$; in this cosmology 1$''$ corresponds to 8.5\\,kpc at $z=2.0$. ",
        "conclusions": "We have presented Gemini-North NIFS IFU observations of the [O~{\\sc iii}]$\\lambda$5007 emission in a $z\\approx$~2 ultraluminous infrared galaxy hosting an optically identified AGN (SMM~J1237+6203). Our main findings are the following:  \\begin{itemize}  \\smallskip \\item The spatial extent ($\\approx$~4--8~kpc) of the high velocity (up-to $v\\approx$~300~km~s$^{-1}$) and broad (up-to $\\approx$~900~km~s$^{-1}$) [O~{\\sc iii}] emission is consistent with that found in $z>2$ radio galaxies, suggesting the presence of a large-scale energetic outflow in a galaxy population potentially orders of magnitude more common than $z>2$ radio galaxies. See \\S3, \\S4.1, \\& \\S4.3.  \\smallskip \\item The estimated energy required to produce these large-scale outflow features ($\\approx$~[0.6--3]~$\\times10^{44}$~erg~s$^{-1}$) could be provided by a wind radiatively driven by the quasar ($\\approx$~[0.3--3]~$\\times10^{45}$~erg~s$^{-1}$) and/or supernovae winds from intense star formation ($\\approx3\\times10^{44}$~erg~s$^{-1}$), so long as the wind-gas coupling efficiencies are comparatively high (of order $\\approx$~10--100\\%). However, the low radio luminosity of this system indicates that radio-bright jets are unlikely to be responsible for driving the outflow, in contrast to $z>2$ radio galaxies. See \\S4.1 and \\S4.2.  \\smallskip \\item The estimated energy injection required to drive the large-scale outflow is comparable to the estimated binding energy of the galaxy spheroid in SMM~J1237+6203, suggesting that the outflow can have a significant impact on the evolution of the galaxy. However, on the basis of the maximum energy input into the outflow and the measured outflow velocities, it seems unlikely that the majority of the ISM will be expelled from the host galaxy. This may mean that another process (or longer lived AGN activity) is required to keep the gas unbound and prevent future star formation from occuring. See \\S4.3.  \\end{itemize}  Finally, we conclude that SMM~J1237+6203 is the first high-redshift ULIRG with spatially extended broad [O~{\\sc III}] emission to be mapped with IFU observations, revealing the signatures expected for a large-scale energetic outflow. However, since $\\approx$~50\\% of nearby ULIRGs hosting optical AGN activity also have broad [O~{\\sc iii}] emission, we expect SMM~J1237+6203 to be comparatively typical of the high-redshift ULIRG population. High signal-to-noise ratio (and high spatial resolution) IFU observations of further high-redshift ULIRGs will provide more definitive constraints."
    },
    "0911/0911.2011_arXiv.txt": {
        "abstract": "We describe the first axisymmetric numerical code based on the generalized harmonic formulation of the Einstein equations, which is regular at the axis. We test the code by investigating gravitational collapse of distributions of complex scalar field in a Kaluza-Klein spacetime. One of the key issues of the harmonic formulation is the choice of the gauge source functions, and we conclude that a damped wave gauge is remarkably robust in this case. Our preliminary study indicates that evolution of regular initial data leads to formation both of black holes with spherical and cylindrical horizon topologies. Intriguingly, we find evidence that near threshold for black hole formation the number of outcomes proliferates. Specifically, the collapsing matter splits into individual pulses, two of which travel in the opposite directions along the compact dimension and one which is ejected radially from the axis. Depending on the initial conditions, a curvature singularity develops inside the pulses. ",
        "introduction": "\\label{sec_intro}   In general, a detailed investigation of fully nonlinear gravitational dynamics is impossible by other than numerical means. Luckily, the numerical methods have recently reached the level of maturity that finally allows addressing many long-standing puzzles. Perhaps the most remarkable is the progress achieved in solving a general-relativistic two-body problem---the coalescence of black holes \\cite{FP2,FP1,FP3,RIT,nasa,AEI,CaltechCornell}. Driven by the gravitational wave detection prospects, the problem of the collision of two black holes or neutron stars continues to be the central front where an overwhelming majority of numerical relativity research is done. Fortunately, the computational methods employed there are portable and---as demonstrated below---can readily be applied on other problems of interest.    The success of the numerical simulations was backed up by a parallel development of the software and the hardware, which provided the necessary computational resources. The rapid hardware evolution combined with the persisting regularity problems in axial symmetry eventually led to a direct transition from highly symmetrical spherical configurations to fully general 3D situations without any symmetries at all, essentially bypassing the intermediate axisymmetric case. However, here we argue that important theoretical and practical reasons exist to explore axisymmetry better, and we describe a new regular numerical code that, we believe, will be capable of achieving this.~\\footnote{See \\cite{Garfinkle_axi,Choptuik_axi,Rinne_axi_1,Rinne_axi_2} for alternative approaches.}  Before describing any concrete setup we would point out one important possible use of an axisymmetric code, specifically, that it can be regarded as an efficient ``calibration tool'' for more general 3D codes. \\footnote{and as a probe of reliability of the cartoon methods \\cite{cartoon_method,FP1} used to effectively simulate axisymmetric spacetimes in 3D.}  Indeed, we expect that an intrinsically axisymmetric code applied to, say, the head-on collision of two black holes would be capable of following the evolution and the resulting gravitational radiation more accurately compared to Cartesian 3D numerical implementations, both because of explicit use of the symmetry and since higher numerical resolution can be employed for given hardware resources.  Several interesting open problems arise in axisymmetric gravitational collapse situations.  In particular, it remains unclear whether or not the weak cosmic censorship is violated in collapse of prolate Brill waves \\cite{BrillWaves,Rinne_axi_1}. An independent observation of universality in critical collapse of gravitational waves \\cite{AbrahamsEvans} is pending, as well as further investigation of the non-spherical unstable mode that apparently shows up at threshold for black hole formation in axisymmetric collapse of a scalar field \\cite{crit_collapse_axi}. A basic problem of mathematical relativity concerning the stability of black holes with respect to nonlinear axisymmetric perturbations, can be equivalently addressed in a collapse situation by computing how fast a newly formed black hole radiates away higher multipole moments.   In addition, we shall mention other, rarely cited in the context of numerical relativity, axisymmetric systems which are of great interest in theoretical work on higher-dimensional gravity.  One of the most basic motivations for studying higher-dimensional spacetimes relies on the observation that the Einstein equations, describing classical General Relativity (GR), are independent of the spacetime dimension. Nevertheless, certain properties of the solutions to these equations vary dramatically with the dimension.  A striking example is that axisymmetric black holes in dimensions greater than four do not necessarily have spherical horizons, but also admit horizons of toroidal topology \\cite{BlackRing}. Moreover, and in sharp contrast to their four-dimensional counterparts, higher dimensional black holes do not respect the Kerr limit \\cite{MP}, and they are unstable in certain range of parameters \\cite{GL}. One of the fundamental unresolved puzzles \\cite{HM,Choptuik_BS} is whether or not the instability leads to fragmentation of the horizon and exposition of the inner singularity, hence violating the cosmic censorship hypothesis.  Since it is unlikely that any of these these problems can be addressed analytically in a systematic manner, we turn to computational methods.  However, solving the Einstein equations numerically is notoriously difficult and depends crucially on the way these equations are formulated and evolved.  In this paper we focus on the generalized harmonic (GH) formulation \\cite{Friedrich85,Friedrich96,Garfinkle2001} that has recently gained popularity because of its great success in the simulations of black hole binaries \\cite{FP2,FP3,CaltechCornell,FP1,Lindblom_etal}.  In a nutshell, the GH approach is a way to write the field equations such that the resulting system is manifestly hyperbolic, taking the form of a set of quasilinear wave equations for the metric components. As the name suggests, the GH method generalizes the harmonic approach, which achieves strong hyperbolicity by choosing the harmonic spacetime coordinates. In the GH approach much of the coordinate freedom is regained through the introduction of certain gauge source functions, which also maintains the desirable property of strong hyperbolicity of the field equations. In fact, the source functions can be thought of as representing the coordinate freedom of the Einstein equations, and when constructing solutions of the equations, via an initial value approach, for example, they must be completely specified in some fashion. Choosing the source functions in a controlled way is a key issue of the GH technique and after testing several recent prescription, we conclude that a damped wave gauge \\cite{LS} is remarkably robust in collapse situations.  Applications of the GH approach in spherically-symmetric situations were studied in \\cite{SorkinChoptuik} and here we employ the method in axisymmetry. In both cases it is natural to use coordinates in which the symmetries of the spacetime are explicit. However, these coordinates, are formally singular: at the origin in spherical symmetry and on the axis in axial symmetry.  Thus, the field equations have to be regularized in numerical implementations; here we describe a regularization procedure that is compatible with the GH formulation.  Our numerical implementation of the GH system is a free evolution code that advances initial data by solving a set of wave equations. In addition, there are also constraint equations that must be satisfied during the evolution.  Although the constraints are consistently preserved in the GH approach in the continuum limit, in numerical computations at finite resolution, constraint violations generically develop.  In order to maintain stability these deviations must be damped and we discuss an effective method that achieves this.   Having in mind implications for higher-dimensional GR we test our new code by studying gravitational collapse of a complex scalar field in a $D$-dimensional Kaluza-Klein (KK) spacetime.  This background, which has a single compact extra dimension curled into a circle, is a classical example of a higher-dimensional compactified spacetime that in certain limits can appear four dimensional, for example when the size of the compact dimension is small.\\footnote{In this paper, however, the size of the KK circle is arbitrary.} Assuming spherical symmetry in the infinite $(D-2)$-dimensional portion of the space makes the problem 2+1 that depends on time and two spatial coordinates: one in the radial direction, and one along the KK circle.  We perform a series of numerical simulations where the initial distribution of the scalar matter is freely specified and the outcome of the evolution depends on the ``strength'' of the initial data as well as on its topology.  The weak data correspond to the dispersion of relatively dilute pulses, while a typical strong data configuration leads to black hole formation or nearly does so. Our preliminary results indicate a wide range of the black hole forming scenarios, including how many holes form and of what topology.  For instance, a static distribution of matter centered at the axis and localized in the KK direction with the energy-density above certain threshold collapses to form a black hole with a quasi-cylindrical horizon, smeared along the extra-dimension.  Data with the energy-density below this threshold evolve to form a quasi-spherical horizon centered around the initial matter distribution. By further reducing the density we find that resulting black holes become progressively smaller, and at some critical density a pulse of matter is emitted radially away from the axis and a curvature singularity develops inside it. We find that for slightly lower initial densities the evolving matter splits into several pulses and two of them individually collapse to form the singularities moving apart along the KK dimension. Finally, when the initial matter distribution becomes dilute below a certain limit, no black holes or curvature singularities are created.  In the next section we describe the class of effectively 2+1 dimensional models where our code can currently be applied. In Sec. \\ref{sec_GH} we present the basic formulas of the GH formulation and discuss the constraints and a method to damp their violations. We describe an axis regularization in Sec. \\ref{sec_axis} and boundary conditions  in Sec. \\ref{sec_eqs}. Although, in this work we integrate full $D$-dimensional equations, in the Appendix we describe an alternative approach---one that uses a dimensional reduction on the symmetry and integrates the reduced 2+1 equations---and compare its performance with ours.  Coordinate conditions and the initial data problem are formulated in Secs. \\ref{sec_coord_condition} and \\ref{sec_id}, respectively. We use several diagnostics to probe the spacetimes that we construct, including computation of asymptotic measurables and apparent horizons, and describe that in Sec. \\ref{sec_diagnostics}.  After elaborating on our numerical algorithm in Sec. \\ref{sec_numercal_approach} we test its performance in Sec. \\ref{sec_num_experements}, giving detailed accounts of various aspects, such as specific coordinate choices, constraint damping, numerical dissipation, and convergence. We conclude in Sec. \\ref{sec_conclusion}, outline possibilities of improvement, and discuss some future prospects. ",
        "conclusions": ""
    },
    "0911/0911.1077_arXiv.txt": {
        "abstract": "We discuss the use of SPICA to study the cosmic history of star formation and accretion by supermassive black holes. The cooling lines, in particular the high-$J$ rotational lines of CO, provide a clear-cut and unique diagnostic for separating the contributions of star formation and AGN accretion to the total infrared luminosity of active, gas-rich galaxies. We briefly review existing efforts for studying high-$J$ CO emission from galaxies at low and high redshift. We finally comment on the detectability of cooling radiation from primordial (very low metallicity) galaxies containing an accreting supermassive black hole with SPICA/SAFARI. ",
        "introduction": "For several decades, our only direct probes of the high-redshift universe were QSOs and radio galaxies. While significant insight into the properties and evolution of these populations was obtained, almost nothing was known about the population of normal galaxies at high $z$. This has changed radically over the last decade. Deep imaging and spectroscopy from the ground and from space have revealed a rich and diverse population of high-$z$ galaxies, and enormous progress has been made in establishing a complete inventory of the high-$z$ universe, in terms of mass and metallicity budget, energy output and evolution over cosmic time. In addition, first investigations of the genesis of the fundamental scaling relations of galaxies (e.g., the fundamental plane and the Tully-Fisher relation) are being undertaken, and will provide insight into the origin of the galaxy population itself.  One of the most remarkable scaling relations is the $M_\\bullet-\\sigma$ relation, i.e., the relation between the mass of a nuclear supermassive black hole (SMBH) and the velocity dispersion of the spheroid (elliptical or bulge) hosting the black hole \\citep{Magorrianetal98,FerrareseMerritt00}. The fact that the black hole mass and total galaxy mass are closely related is truly remarkable, given that there is a factor of $\\sim10^8$ between the AU-size Schwarzschild radius of the black hole and the kpc-size dimension of the galaxy. This is generally interpreted as evidence that the formation and growth of the SMBH is directly related to the formation process of the stellar population, e.g., in a violent burst of star formation. The picture of the simultaneous build-up of a SMBH with an extreme burst of star formation provides a new context for the luminous and ultraluminous infrared galaxies (LIRGs and ULIRGs), long suggested to be the birthplaces of SMBHs powering active galactic nuclei (AGNs). Studies of local ULIRGs demonstrate that they plausibly evolve into moderate mass field ellipticals \\citep{Genzeletal01}. The giant ellipticals in the local universe would then be the result of similar but much more extreme events at higher redshifts. Indeed, the population of submillimetre galaxies (SMGs) provides the more luminous high redshift counterpart of the local LIRGs and ULIRGs and may be responsible for the formation of local massive spheroids and for generating QSO activity at high redshift \\citep{Blainetal99}.  \\begin{figure}[th] \\begin{center} \\includegraphics[width=8cm]{vanderWerfP_fig1.eps} \\end{center} \\caption{CO emergent intensity distributions for a PDR and an XDR with {\\it identical\\/} impinging energy density, as described in Section~\\ref{vanderWerfP.SecDiag}. Note that high-$J$ CO emission is expected in both cases, but vastly more prominently in the XDR \\citep{SpaansMeijerink08}.} \\label{vanderWerfP_fig1} \\end{figure} ",
        "conclusions": ""
    },
    "0911/0911.1241_arXiv.txt": {
        "abstract": "We provide an analytical expression for the trispectrum of the Cosmic Microwave Background (CMB) temperature anisotropies induced by cosmic strings. Our result is derived for the small angular scales under the assumption that the temperature anisotropy is induced by the Gott--Kaiser--Stebbins effect. The trispectrum is predicted to decay with a non-integer power-law exponent $\\ell^{-\\rho}$ with $6<\\rho<7$, depending on the string microstructure, and thus on the string model. For Nambu--Goto strings, this exponent is related to the string mean square velocity and the loop distribution function. We then explore two classes of wavenumber configuration in Fourier space, the kite and trapezium quadrilaterals. The trispectrum can be of any sign and appears to be strongly enhanced for all squeezed quadrilaterals. ",
        "introduction": "Although cosmic strings may be of various early universe origins~\\cite{Kibble:1976sj,Dabholkar:1990yf, Hindmarsh:1994re, VilShe94, Kofman:1994rk,Yokoyama:1989pa, Sakellariadou:2006qs, Copeland:2003bj, Sarangi:2002yt, Dvali:2003zj}, being line-like gravitational objects, they induce temperature discontinuities in the CMB through the Gott--Kaiser--Stebbins (GKS) effect \\cite{Gott:1984ef,Kaiser:1984iv}. Direct searches for such discontinuities have been performed without success but do provide upper limits to the string tension $U$~\\cite{Jeong:2004ut, Jeong:2007, Christiansen:2008vi}. On the other hand, if cosmic strings are added to the standard power-law $\\Lambda$CDM model~\\cite{Bouchet:2000hd}, it has been shown in Refs.~\\cite{Battye:2006pk,Bevis:2007gh} that the CMB data are fitted even better if the fraction of the temperature power spectrum due to strings is about $10\\%$ (at $\\ell=10$). Such a fraction of string would even dominate the primary anisotropies of inflationary origin for $\\ell \\gtrsim 3000$~\\cite{Fraisse:2007nu}. With the advent of the arc-minute resolution CMB experiments and the soon incoming Planck satellite data, it is therefore crucial to develop reliable tests for strings~\\cite{Seljak:2006}, as to understand the non-Gaussian signals. The probability distribution of the fluctuations due to the GKS effect is known to be skewed and has a less steep decay than Gaussian~\\cite{Fraisse:2007nu}, a feature which can be explained in a simple model of kinked string~\\cite{Takahashi:2008ui}. In Ref.~\\cite{Hindmarsh:2009qk}, we have studied the temperature bispectrum induced by cosmic strings both analytically and numerically by using Nambu--Goto string simulations. We found good agreement between the analytical and numerical bispectrum for both the overall amplitude and the geometrical factor associated with various triangle configurations of the wavevectors. This agreement suggests that our analytical assumptions are capturing the relevant non-Gaussian features of a string network and could be used to derive other statistical properties. In this paper, we present new results concerning the trispectrum, i.e. the four-point function of the temperature anisotropy~\\cite{Hu:2001fa}. As pointed out in Ref.~\\cite{Hindmarsh:2009qk}, the bispectrum is generated only when the background spacetime breaks the time reversal symmetry. Because our universe is expanding, the time reversal symmetry is indeed broken and we get a non-vanishing string bispectrum. On the other hand, the trispectrum can be generated even in Minkowski spacetime and one may naively expect a stronger non-Gaussian signal than for the bispectrum. Motivated by this observation, we provide in this paper an analytical derivation of the string trispectrum and study its dependency for various quadrilateral configurations in Fourier space. Given the fact that analytical predictions and numerical results exhibit good agreement both for the power spectrum~\\cite{Hindmarsh:1993pu, Fraisse:2007nu} and the bispectrum~\\cite{Hindmarsh:2009qk}, we expect our result here to agree as well with the numerics. Performing such a comparison would however require a significant amount of computing resources which motivated us to leave it for a future work. Our main result can be summarized by Eq.~(\\ref{eq:finaltri}). Interestingly, the power-law behaviour of the trispectrum exhibits a non-integer exponent which can be related to the small scale behaviour of the string tangent vector correlator. In the framework of Nambu--Goto strings, this exponent is related to the mean square string velocity~\\cite{Polchinski:2006ee, Copeland:2009dk} and to the scaling loop distribution function~\\cite{Ringeval:2005kr, Rocha:2007ni}.  The paper is organised as follows. In the next section we briefly recall the assumptions at the basis of our analytical approach and derive the trispectrum in Sec.~\\ref{sec:trispectrum}. As an illustration, we apply our result to some specific quadrilaterals in Sec.~\\ref{sec:quad} and exhibit some configurations that leads to a divergent trispectrum. They should provide the cleanest way to look for a non-Gaussian string signal. ",
        "conclusions": "\\label{sec:concl}  In this paper, we have analytically derived the CMB temperature trispectrum induced by cosmic strings using the string correlation functions in the Gaussian approximation. The trispectrum generically decays with a non-integer power-law behaviour at small angular scales which depends on the string microstructure through the behaviour of the tangent vector correlator on small distances. Its eventual detection and measurement may therefore help to distinguish between different string models. We have also found that the trispectrum diverges, in the framework of our approximations, on all squeezed configurations whose measurements remain however limited by the finite experimental resolution. In fact, such a non-integer power-law is linked to the existence of a ``flat direction'' at leading order and the four-point function ends up being sensitive to the next-to-leading order string tangent vector correlator. This situation is also present in the n-point function and we do expect all of the higher n-point function to exhibit non-integer power-law behaviours. Since this situation was not encountered for the two- and three-point functions, the next step will be to compare our results here with the trispectrum computed from CMB maps obtained by string network simulations.  Finally, let us notice that we have not attempted to make any comparison with a CMB trispectrum produced by primordial non-Gaussianities of inflationary origin. The situation is nearly the same as it is for the string bispectrum~\\cite{Hindmarsh:2009qk}. The so-called $\\tauNL$ and $\\gNL$ parameters quantify the amplitude of the primordial four-point function of the curvature perturbation on super-Hubble scales. As a result, the induced trispectrum of the CMB temperature fluctuations strongly depends on the CMB transfer functions and exhibits damped oscillations with respect to the multipole moments. Here, we have direcly derived the CMB temperature trispectrum produced by the strings and it would therefore make no sense to find an associated $\\tauNL$ and $\\gNL$. An alternative approach might be to estimate what values $\\tauNL$ and $\\gNL$ would assume in a primordial-type oriented data analysis if the non-Gaussianities were actually due to strings. This could be done with a Fisher matrix analysis for a given experiment but we leave this question for a forthcoming work."
    },
    "0911/0911.1131_arXiv.txt": {
        "abstract": "We use a novel method to predict the contribution of normal star-forming galaxies, merger-induced bursts, and obscured AGN, to IR luminosity functions (LFs) and global SFR densities. We use empirical halo occupation constraints to populate halos with galaxies and determine the distribution of normal and merging galaxies. Each system can then be associated with high-resolution hydrodynamic simulations. We predict the distribution of observed luminosities and SFRs, from different galaxy classes, as a function of redshift from $z=0-6$. We provide fitting functions for the predicted LFs, quantify the uncertainties, and compare with observations. At all redshifts, `normal' galaxies dominate the LF at moderate luminosities $\\sim L_{\\ast}$ (the `knee'). Merger-induced bursts increasingly dominate at $L\\gg L_{\\ast}$; at the most extreme luminosities, AGN are important. However, all populations increase in luminosity at higher redshifts, owing to increasing gas fractions. Thus the `transition luminosity' between normal and merger-dominated sources increases from the LIRG-ULIRG threshold at $z\\sim0$ to bright Hyper-LIRG thresholds at $z\\sim2$. The transition to dominance by obscured AGN evolves similarly, at factor of several higher $L_{\\rm IR}$. At all redshifts, non-merging systems dominate the total luminosity/SFR density, with merger-induced bursts constituting $\\sim5-10\\%$ and AGN $\\sim1-5\\%$. Bursts contribute little to scatter in the SFR-stellar mass relation. In fact, many systems identified as `ongoing' mergers will be forming stars in their `normal' (non-burst) mode. Counting this as `merger-induced' star formation leads to a stronger apparent redshift evolution in the contribution of mergers to the SFR density. We quantify how the evolution in LFs depends on evolution in galaxy gas fractions, merger rates, and possible evolution in the Schmidt-Kennicutt relation. We discuss areas where more detailed study, with full radiative transfer treatment of complex three-dimensional clumpy geometries in mixed AGN-star forming systems, is necessary. ",
        "introduction": "\\label{sec:intro}  Understanding the global star-formation history of the Universe remains an important unresolved goal in cosmology. In recent years, observations of the properties of galaxies in the infrared, at redshifts $z=0-3$, have begun to shed light on this history, but have also revealed a number of intriguing questions.  Of particular interest are the roles of mergers and AGN in driving star formation and/or the infrared luminosities of massive systems. A wide range of observed phenomena support the view that gas-rich\\footnote{ By ``gas,'' we refer specifically to cold, star-forming gas in galaxy disks, as opposed to hot, virialized gas.} mergers are important to galaxy evolution; but it is less clear what their role is in the global star formation process and buildup of stellar mass in the Universe. In the local Universe, the population of star-forming galaxies appears to transition from ``quiescent'' (non-disturbed)\\footnote{In this paper, we use the term ``quiescent'' to refer to star-forming systems that are not strongly disturbed in e.g.\\ major mergers and are forming stars in similar fashion to most ``normal'' disks. We do {\\em not} mean non-star forming systems, as the term is used in some literature.} disks -- which dominate the {\\em total} star formation rate/IR luminosity density -- at the luminous infrared galaxy (LIRG) threshold $10^{11}\\,\\lsun$ ($\\dot{M}_{\\ast}\\sim 10-20\\,\\msun\\,{\\rm yr^{-1}}$) to clearly merging, violently disturbed systems at a few times this luminosity. The most intense starbursts at $z=0$, ultraluminous infrared galaxies (ULIRGs; $L_{\\rm IR}>10^{12}\\,\\lsun$), are invariably associated with mergers \\citep[e.g.][]{joseph85,sanders96:ulirgs.mergers, evans:ulirgs.are.mergers}, with dense gas in their centers providing material to feed black hole (BH) growth and to boost the concentration and central phase space density of merging spirals to match those of ellipticals \\citep{hernquist:phasespace,robertson:fp}. Various studies have shown that the mass involved in these starburst events is critical to explain the relations between spirals, mergers, and ellipticals, and has a dramatic impact on the properties of merger remnants \\citep[e.g.,][]{LakeDressler86,Doyon94,ShierFischer98,James99, Genzel01,tacconi:ulirgs.sb.profiles,dasyra:mass.ratio.conditions,dasyra:pg.qso.dynamics, rj:profiles,rothberg.joseph:kinematics,hopkins:cusps.ell,hopkins:cores}.  At high redshifts, the role of mergers is less clear. It is clear that LIRGs and ULIRGs increase in relative importance with redshift, with LIRGs dominating the star formation rate/IR luminosity densities at $z\\sim1$ and ULIRGs dominating at $z\\sim2$ \\citep[e.g.][]{lefloch:ir.lfs,perezgonzalez:ir.lfs,caputi:ir.lfs,magnelli:z1.ir.lfs}. This, together with the fact that merger rates are expected and observed to increase with redshift \\citep[by a factor $\\sim10$ from $z=0-2$; see e.g.][and references therein]{hopkins:merger.rates} has led to speculation that the merger rate evolution may in fact drive the observed evolution in the cosmic SFR density, which rises rapidly from $z\\sim0-2$ and then turns over, declining more slowly \\citep[e.g.][and references therein]{hopkinsbeacom:sfh}.  However, many LIRGs at $z\\sim1$, and potentially ULIRGs at $z\\sim2$, appear to be ``normal'' galaxies, without dramatic morphological disturbances associated with the local starburst population or large apparent AGN contributions \\citep{yan:z2.sf.seds,sajina:pah.qso.vs.sf, dey:2008.dog.population,melbourne:2008.dog.morph.smooth, dasyra:highz.ulirg.imaging.not.major}. At the same time, even more luminous systems appear, including large populations of Hyper-LIRG (HyLIRG; $L_{\\rm IR}>10^{13}\\,\\lsun$) and bright sub-millimeter galaxies \\citep[e.g.][]{chapman:submm.lfs,younger:highz.smgs, younger:sma.hylirg.obs,casey:highz.ulirg.pops}. These systems exhibit many of the traits more commonly associated with merger-driven starbursts, including morphological disturbances, and may be linked to the emergence of massive, quenched (non star-forming), compact ellipticals at times as early as $z\\sim2-4$ \\citep{papovich:highz.sb.gal.timescales, younger:smg.sizes,tacconi:smg.maximal.sb.sizes, schinnerer:submm.merger.w.compact.mol.gas, chapman:submm.halo.clustering,tacconi:smg.mgr.lifetime.to.quiescent}. But reproducing their abundance and luminosities remains a challenge for current models of galaxy formation \\citep{baugh:sam, swinbank:smg.counts.vs.durham, narayanan:smg.modeling, younger:warm.ulirg.evol}.  In a related vein, observations of a tight correlation between the masses of super-massive BHs and their host spheroid properties \\citep{Gebhardt00,FM00,magorrian,novak:scatter,hopkins:bhfp.obs} suggest a tight coupling between BH growth and star formation, perhaps in particular to the mergers believed to drive the formation of the most massive bulges. Considering the energy output required to form the BH population \\citep[e.g.][]{soltan82}, or the observed bolometric quasar energy density as a function of redshift \\citep[see][and references therein]{hopkins:bol.qlf}, it is clear that the {\\em bolometric} output of quasars and AGN is at least roughly comparable to the total infrared luminosity density of the Universe at most redshifts ($z\\sim1-3$) -- although the measurements above suggest it is still a factor $\\sim2-3$ lower. Some recent observations have suggested that the population of very luminous, highly obscured (Compton-thick) quasars may be considerably larger than previous estimates, in which case the heavily-obscured AGN population could represent a large fraction of the total IR luminosity density at high redshifts \\citep{hickox:bootes.obscured.agn,daddi:2007.high.compton.thick.pops, treister:compton.thick.fractions}. This would have dramatic implications not just for BH populations and e.g.\\ the implied radiative efficiencies of BH accretion, but also for the implied total star formation rate density. Some apparent discrepancies between e.g.\\ the total mass density observed in old stars and the implied star formation rate density have been cited as possible evidence of a time-dependent stellar initial mass function (IMF); but a rising contribution from obscured AGN at high redshifts could mimic this effect \\citep{hopkinsbeacom:sfh,dave:imf.evol}.  In particular, there are long-standing questions of what powers the most luminous infrared sources, for example, ULIRGs and sub-millimeter galaxies. This debate extends to the discovery of these objects \\citep[see e.g.][]{soifer84b,soifer:iras,scoville86, sargent87,sanders88:warm.ulirgs, solomon.downes:ulirg.ism}, and has persisted despite the addition of millimeter spectroscopy and observations in a large number of independent wavebands \\citep[for a review of the debate, see both][]{sanders:agn.vs.sf.in.ulirgs, joseph:sb.vs.agn.power.ulirgs}. Although some evidence suggests that they are primarily powered by star formation \\citep{farrah:qso.vs.sf.sed.fitting, lutz:pah.qso.vs.sf.local,sajina:pah.qso.vs.sf, pope:2008.pah.agn.dont.dominate.smgs,pope:2008.dog.sfr.properties, watabe:highz.ulirg.sb.vs.agn,nardini:2009.agn.vs.sb.contrib.in.ulirgs}, the constraints and correlations typically invoked have inherent factor $\\sim2$ uncertainties, and thus could easily accommodate comparable power input from star formation and AGN. Moreover, a sufficiently obscured AGN, in a medium with the right optical depth properties, is indistinguishable from star formation by the usual indicators (e.g.\\ PAH strengths, emission region sizes, or any other infrared spectral or morphological criteria). Hence even at $z=0$, debate surrounds the power source of many bright infrared systems, and there exist a number of examples of systems classified as ``star formation dominated'' by all of these metrics that later revealed Compton-thick AGN whose longer-wavelength emission has been fully re-processed, even into ``cool'' dust \\citep[see e.g.][and references therein]{alexander:2008.compton.thick.z2.qsos}. In a bolometric sense, the most luminous galaxies observed (with $L\\gtrsim 10^{14}\\,\\lsun$ or $\\gg 10^{47}\\,{\\rm erg\\,s^{-1}}$) are the most luminous quasars; although the contribution of these systems to the infrared remains highly uncertain. This may be important for resolving the theoretical difficulties in modeling these bright systems.  There has been important theoretical progress in modeling these processes in an {\\em a priori} manner \\citep[see e.g.][]{baugh:sam, hopkins:groups.qso,narayanan:smg.modeling,younger:warm.ulirg.evol}. However, two basic limitations remain. In direct cosmological hydrodynamic simulations, as well as semi-analytic models of galaxy formation, it is well known that it remains challenging to accurately reproduce global quantities such as the galaxy mass function and the distribution of sizes, gas fractions, and hence star formation rates, especially the distributions of star-forming gas and their relations to whether or not galaxies are ``quenched'' \\citep[recently, see e.g.][]{weinmann:group.cat.vs.sam, maller:sph.merger.rates,kimm:passive.sats.vs.centrals,fontanot:downsizing.vs.sams}. This makes it difficult to determine whether discrepancies between such models and the observations owe to their treatment of star formation, or to discrepancies in these quantities. Moreover, it is difficult to disentangle the effects of these different properties on the distribution of star formation rates. In addition, for merger-induced starburst and AGN activity, although it may be possible to roughly estimate some global quantities (e.g.\\ the total mass involved in a starburst) from simple analytic motivations or low-resolution cosmological simulations (several $\\sim$kpc typical), it is not straightforward to estimate the chaotic, time-dependent behavior of full lightcurves needed to estimate the distribution of time spent at different, rapidly varying luminosities without high-resolution simulations of individual systems. Since the number density of the most bright systems is exponentially declining, fluctuations and features in the starburst/AGN fueling history on small time and spatial scales ($\\Delta t\\sim10^{7}\\,{\\rm yr}$, $R\\lesssim 100\\,$pc) can be critical for correct estimates of their contributions to bright populations.  In this paper, we present theoretical predictions for the distribution of galaxy star formation rates and infrared luminosities, as a function of galaxy mass and redshift, using a novel methodology that can circumvent some of these obstacles. We combine a halo-occupation based approach, in which we take galaxy properties as fixed from observations at each epoch, and then apply rules for the distribution of star formation rates/infrared luminosities in ``quiescent'' systems, merger-induced starbursts, and obscured AGN, calculated from a large suite of high-resolution hydrodynamic simulations of individual galaxies and galaxy mergers. We use this to independently estimate the contributions of ``normal'' galaxies, mergers, and AGN to the luminosity functions. The comparisons we make are approximate -- we do not include full time-dependent radiative transfer in simulations (the subject of future work, in progress), and so focus on integral quantities such as the total IR luminosity and SFR distributions, that are less sensitive to issues of e.g.\\ the exact dust distribution, temperature, and other properties. We also explicitly separate the contributions of AGN and star formation, but stress that, in real systems where the two are comparable, their additive effects are non-linear, and will require further study. We show how adding or removing components of the model taken from observations such as e.g.\\ the distribution of galaxy sizes and gas fractions affects these consequences. We compare to observations of all quantities, where available, and find reasonable agreement but with some interesting apparent discrepancies at high redshifts. We also show how these populations relate to the scatter in star formation rates at fixed galaxy masses, and in a global sense to the total star formation rate density, the star formation rate density in mergers, and the fraction of the inferred star formation rate density which might really be driven by obscured AGN activity. Readers interested primarily in the comparison of predictions and observations may wish to skip directly to \\S~\\ref{sec:lfs}.  Throughout, we adopt a $\\Omega_{\\rm M}=0.3$, $\\Omega_{\\Lambda}=0.7$, $h=0.7$ cosmology and a \\citet{chabrier:imf} stellar IMF (discussed further below), but these choices do not affect our conclusions. ",
        "conclusions": "\\label{sec:discussion}  We present a simple model for the distribution of star formation rates and infrared luminosities owing to ``normal'' star-forming disks, merger-induced starbursts, and AGN. Comparing this with observations, we find reasonable agreement at $z\\sim0-3$. At all redshifts, we find that the low-luminosity population is dominated by disks, whereas the high-luminosity population becomes progressively more dominated by merger-induced bursts and then, ultimately, obscured AGN.  The threshold for this transition is always at high luminosities and low space densities. At higher redshifts, gas fractions in all systems increase -- hence, specific star formation rates even in ``typical'' systems are much higher at high-$z$. As a consequence, the luminosity threshold for the disk-merger transition increases with redshift, from between LIRGs and ULIRGs in the local Universe, to ULIRG luminosities at $z\\sim1$, to HyLIRG luminosities at $z\\sim2-4$. Similarly, the threshold for AGN dominance, at the bright ULIRG range at low redshifts, rises to HyLIRG luminosities at $z\\sim1-2$ and to the very brightest HyLIRGs at $z>2$.  We provide simple fitting functions for each of these quantities, and show how they depend on different parameters in the model. Most critically, it is the evolution in galaxy gas fractions that drives most of the evolution in the IR LFs. Observations have shown that typical gas fractions in massive, star-forming galaxies increase rapidly with redshift \\citep[see e.g.][]{erb:lbg.gasmasses, bouche:z2.kennicutt,puech:tf.evol, mannucci:z3.gal.gfs.tf, forsterschreiber:z2.sf.gal.spectroscopy}. This naturally follows from the facts that cooling rates onto galaxies at high redshift are much higher than at low redshift, and there has simply been less time to process gas into stars (higher cosmic densities may also make stellar and AGN feedback relatively less efficient). A higher gas fraction by a factor of $\\sim3$, as implied by observations of Milky-Way mass disks at $z\\sim2-3$, leads to a factor $\\sim5-6$ higher star formation rate, according to the Kennicutt-Schmidt relation. At high masses/luminosities, i.e.\\ where the number density of systems is falling exponentially, such a systematic increase in luminosity more than offsets the declining space density of massive galaxies with redshift.  Other parameters have little effect. Changes in galaxy sizes and/or structural parameters, while potentially very important for e.g.\\ how gas fractions are maintained, make little difference to the SFR distributions given some gas fraction and stellar mass distributions. Allowing for a more steep index in the \\citet{kennicutt98} relation may help to account for the most luminous observed systems such as bright sub-millimeter galaxies at the Hyper-LIRG threshold at $z\\sim2-4$ \\citep{chapman:submm.lfs}. However, the number counts of such objects remains quite uncertain, and recently it has been suggested that the average counts might be much lower with cosmic variance still a concern \\citep{austermann:2009.aztec.submm.source.counts}. Larger samples and better calibration of bolometric corrections -- in particular, real knowledge of the appropriate dust temperatures for conversion to total-IR luminosities, which requires sampling both sides of the cold dust peak \\citep[see e.g.][]{younger:mm.obs.z2.ulirgs} -- will be needed for better understanding of extreme systems.  Our simple model succeeds reasonably well at explaining the observed global SFR density. At all redshifts, normal systems dominate the global SFR density. Obscured AGN contribute little, $\\lesssim5\\%$, to the total IR luminosity density. Substantial bias to IR-based SFR density estimates from obscured AGN would require an undiscovered population of heavily, isotropically obscured sources with luminosity densities $\\sim5$ times what is currently suggested (which would be in conflict with relic BH mass densities).  The contribution of merger-induced bursts is similarly small, $\\sim5-10\\%$ at most redshifts. This owes both to the physics discussed above -- disks also rapidly increase their SFR density with redshift owing to higher gas fractions -- and also to the fact that increasing disk gas fractions arbitrarily will not continue to increase the merger-induced burst contribution arbitrarily. Rather, as discussed in detail in \\citet{hopkins:disk.survival,hopkins:disk.survival.cosmo}, at high gas fractions angular momentum loss in mergers becomes less efficient -- thus for an otherwise identical merger with a much larger gas fraction, the fraction of gas funneled into the nuclear starburst, relative to the total available, will be less (and the fraction that remains in an extended disk distribution to continue ``normal'' mode star formation will be larger), even if the absolute mass in the burst is larger. Similarly, at a given mass, the distribution of SFRs at all redshifts is dominated by normal-mode star formation -- merger-induced bursts are important only in the high-SFR tail ($\\sim2-3\\,\\sigma$) of the distribution. These trends explain a number of recent observations that similarly indicate a small effect of merger enhancements to star formation rates, and that show that most star formation by number and luminosity density appears to follow a simple trend or ``main sequence'' as a function of galaxy stellar mass and redshift \\citep{blain:ir.lf.synthesis.model, noeske:2007.sfh.part1, noeske:sfh,papovich:ssfr, bell:morphology.vs.sfr, jogee:merger.density.08, robaina:2009.sf.fraction.from.mergers}.  Support for these fractions also comes from completely independent sources. Recently, a number of high-resolution studies of spheroid formation via galaxy mergers have shown that properties such as the surface brightness profiles, sizes, concentrations, kinematics, and isophotal shapes of spheroids are very sensitive to the mass fractions formed in such bursts, which produce dense, disky, nuclear mass concentrations, versus the mass in a more extended envelope formed via the violent relaxation of the pre-burst stellar disks \\citep[see e.g.][]{cox:kinematics, naab:gas,robertson:fp,burkert:anisotropy, jesseit:merger.rem.spin.vs.gas,hopkins:cusps.mergers, hopkins:cusps.fp,hopkins:cusps.evol}. In particular, typical $\\sim L_{\\ast}$ early-type galaxies, which dominate the spheroid stellar mass density, have properties that are reproduced accurately by simulations if and only if this burst fraction is $\\sim10\\%$ \\citep[for details, see][]{hopkins:cusps.ell,hopkins:cores}. Independent analysis of their stellar population properties leads to similar conclusions \\citep{mcdermid:sauron.profiles,sanchezblazquez:ssp.gradients, reda:ssp.gradients,foster:metallicity.gradients.prep}.  There is an important technical distinction between the total SFR density in ``ongoing'' or recent mergers and that actually {\\em induced} by the merger (we present predictions for both). The former includes systems in their ``normal'' star-forming mode, observable for $\\sim$Gyr as perturbed or in pairs; the latter reflects specifically the $\\sim10^{8}\\,$yr event where gravitational torques drive a nuclear starburst. Under some circumstances, especially in very gas-rich mergers, the sum of this ``normal mode'' star formation over a long $\\sim$Gyr duration yields significantly more total stellar mass formed than in the burst itself. The ``ongoing'' merger SFR density must, of course, rise with the observed merger fraction, reaching $\\sim20\\%$ at $z=2$ and as high as $\\sim20-50\\%$ at $z>4$; however we stress that these high fractions reflect predominantly the ``normal'' modes of star formation simply present in systems that may be on their way to merging.  Finally, we caution that the above comparisons are approximate, and intended as a broad means of comparing the primary drivers of star formation and their contributions relative to observed IR luminosity functions and SFR distributions. We have ignored a number of potentially important effects: for example, obscuration is a strong function of time in a merger, and may affect various luminosities and morphological stages differently. Moreover, our simple linear addition of the star formation contribution of mergers to the IR LF and the AGN contribution is only technically correct if one or the other dominates the IR luminosity at a given time in the merger; however, there are clearly times during the final merger stages when the contributions are comparable. Resolving these issues requires detailed, time-dependent radiative transfer solutions through high-resolution simulations that properly sample the merger and quiescent galaxy parameter space at each redshift, and is outside the scope of this work \\citep[although an important subject for future, more detailed study; see, e.g.][]{li:radiative.transfer,narayanan:smg.modeling, narayanan:molecular.gas.in.smgs,younger:warm.ulirg.evol}. It would be a mistake, therefore, to read too much into e.g.\\ the detailed predictions for sub-millimeter galaxies or other extreme populations in Figure~\\ref{fig:ir.lfs} that may have complex dust geometries and/or a non-trivial mix of contributions from all of ``normal'' and ``burst'' mode star formation as well as AGN. However, most of our predicted qualitative trends, including the evolution of the luminosity density (and approximate relative contribution of mergers) and the shift in where quiescent or merger-driven populations dominate the bright IR LF, should be robust."
    },
    "0911/0911.4839_arXiv.txt": {
        "abstract": "{There is mounting evidence for an extra-planar gas layer around the Milky Way disk, similar to the anomalous \\hi~ gas detected in a few other galaxies. As much as 10\\% of the gas may be in this phase.}  {We analyze \\hi~ clouds located in the disk-halo interface outside the solar circle to probe the properties of the extra-planar \\hi~ gas, which is following Galactic rotation.} {We use the Leiden/Argentine/Bonn (LAB) 21-cm line survey to search for \\hi~ clouds which take part in the rotation of the Galactic plane, but are located above the disk layer. Selected regions are mapped with the Effelsberg 100-m telescope. Two of the \\hi~ halo clouds are studied in detail for their small scale structure using the Westerbork Synthesis Radio Telescope (WSRT)\\thanks{The Westerbork Synthesis Radio Telescope is operated by ASTRON (Netherlands Foundation for Research in Astronomy) with support from the Netherlands Foundation for Scientific Research NWO.} and the NRAO Very Large Array (VLA)\\thanks{The National Radio Astronomy Observatory is a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc. }.} {Data from the 100m telescope allow for the parameterization of 25 distinct \\hi~ halo clouds at Galactocentric radii 10 kpc $<R<$15 kpc and heights 1 kpc$ <z<$ 5 kpc. The clouds have a median temperature of 620 K,  column densities of \\NH$\\sim10^{19}$cm$^{-2}$, and most of them are surrounded by an extended envelope of warmer \\hi~ gas. Interferometer observations for two selected regions resolve the \\hi~ clouds into several arc-minute sized cores. These cores show narrow line widths (\\FWHM$ \\sim 3$ \\kms), they have volume densities of $n>1.3$ cm$^{-3}$, masses up to 24 \\msun, and are on average in pressure equilibrium with the surrounding envelopes. Pressures and densities fall within the expectations from theoretical phase diagrams ($P$ vs $\\langle n_{h} \\rangle$). The \\hi~ cores tend to be unstable if one assumes a thermally bistable medium, but are in better agreement with models that predict thermal fragmentation driven by a turbulent flow.  } {} ",
        "introduction": "\\label{sec_1}  The \\hi~ gas is a major constituent of the interstellar medium (ISM), and it is a well known property of this gas that it settles in the Galactic plane. From the very first \\hi~ observations it is also known that the disk gas is co-rotating with the stellar disk. Taking both properties together one may use the gas distribution to describe the morphology of the Galactic disk.  As yet, there is no sharp boundary for the disk emission. \\citet{Oort1962dm1} was the first who mentioned this fact.  ``Well outside the real disk one still finds neutral hydrogen with an average density of between 5 and 10 per cent of the intensities one observes in the plane''. Oort was referring to observations with the Dwingeloo telescope, and \\citet{Shane1967IAUS} described this gas later as a ``galactic envelope'', a smooth envelope of neutral hydrogen surrounding the spiral structure, following the same Galactic rotation as the gas in the plane. Further discussion of this envelope was given by \\citet{Takakubo1967BAN} and by \\citet{Shane1971AAS}, but there was some concern about a possible contamination by stray radiation from the antenna diagram of the Dwingeloo telescope.  The extra-planar gas component was also visible in the \\citet{WW1973AAS} survey, and \\citet{Lockman1984} studied this feature in some more detail. Supplementing observations were made with the NRAO 300-foot telescope and came also from the NRAO 140-foot survey by \\citet{Burton1983AAS}. Lockman argued that his analysis was not affected by stray radiation. He found for Galactocentric radii $4 \\la R \\la 8$ kpc that 13\\% of the \\hi~ gas is located outside the disk, extending to $z$-distances of 1 kpc or more and termed this component an ``\\hi~ halo''.  The Bell Labs survey \\citep{Stark1992ApJS} is only little affected by instrumental effects, and \\citet{Lockman1991} used this survey to analyze the nature of the vertical \\hi~ gas distribution in the direction of the Galactic poles. They proposed for the \\hi~ gas a decomposition in several layered structures, corresponding to distinct different isothermal cloud populations. The scale height for each component results from the pressure balance of a cloudy turbulent medium against the gravitational potential of the Milky Way. The concept of a layered structure of the \\hi~ contains essentially three components: a cold neutral medium (CNM), a warm neutral medium (WNM) and an extra-planar component \\citep{Dickey1990ARAAD} which is often called ``Lockman Layer''.  The layer concept is based on the average emission from the extra-planar \\hi~ gas layer which is very faint. The clumpy nature of the \\hi~ gas implies then that extra-planar \\hi~ clouds must have a low volume filling factor. The layer concept describes therefore an ensemble of \\hi~ clouds or the probability distribution of such objects. First indications for a population of such clumps were found by \\citet{Simonson1971AA}. These early data came from the Dwingeloo telescope. Almost three decades later the Leiden/Dwingeloo survey (LDS,\\citep{Hartmann1997}) became available and provided a much improved database, more sensitive and essentially free of stray radiation. Channel maps show numerous clumps and filaments that are detached from the disk, and \\citet{KalberlaWestphalen1998} argued for an extra-planar gas layer which can be characterized by a distribution with a velocity dispersion of $\\sigma = 60$ \\kms, considerably larger than the dispersion suggested by \\citet{Lockman1991}.  The first high resolution data of the extra-planar gas layer at a beam-width of 9\\arcmin~ have been taken by \\citet{Lockman2002} with the Robert C. Byrd Green Bank Telescope (GBT). These observations demonstrated convincingly the nature of the extra-planar gas layer as a population of cold clumps with a typical mass of 50 \\msun. Many of these clumps appear to be surrounded by warmer envelopes. A larger sample of clouds in the lower halo was studied by \\citet{Ford2008ApJ} with the Parkes Telescope. These clumps are somewhat larger and more massive than the sample detected by \\citet{Lockman2002}. This cloud population, located close to $R \\sim 3.8 $ kpc, is interpreted as originating from a Galactic fountain. Such a model would also explain the high kinetic energy which is needed for individual clouds to reach large $z$-distances. Support for such an interpretation comes from \\citet{Stil2006}. They found fast moving clumps in the Galactic plane with velocity vectors located {\\it within} the Galactic plane, analogous to fast velocities {\\it perpendicular} to the plane as suggested as an explanation for the extra-planar gas layer.  So far we discussed predominantly Galactocentric distances $R \\la 8.5$ kpc, where 8.5 kpc is the I.A.U Sun--Galactic center distance, since most of the observations are in this range. For a more general description of this phenomenon, in particular for the question whether the extra-planar gas layer is caused by a fountain flow, objects at larger distances are needed. \\citet{Kalberla2008AA} argue that extra-planar gas is present even at $R \\ga 35$ kpc. Gas at such distances can hardly originate from fountain events.  Direct evidence for a population of extra-planar \\hi~ clouds outside the Solar circle was first given by \\citet{Snezana2006}. Arecibo data in the direction towards the anti-center suggest that these clouds are not restricted to the inner part of the Milky Way disk, which is similar to preliminary results with the Effelsberg telescope reported by \\citet{Kalberla2005ASPC}. In the following we intend to explore the extra-planar gas layer of the outer part of the Milky Way in some more detail. Our results are based on single-dish observations with the Effelsberg 100-m radio telescope and on interferometer observations with the VLA and the WSRT array.  This paper is organized as follows. In Sect. \\ref{sec_2} we explain our selection criteria for targets that have been mapped with the 100-m telescope. Our observations are described in Sect. \\ref{sec_3}. We discuss the properties of the \\hi~ cloud sample detected by us with the 100-m telescope, and also the results from two targeted interferometer observations in Sect. \\ref{sec_4}. We find evidence for a multi-phase structure and compare in Sect. \\ref{sec_6} the derived physical parameters with theoretical models. Sect. \\ref{sec_7} gives our summary and conclusion. ",
        "conclusions": "\\label{sec_7}  We discussed a population of \\hi~ clouds residing in the lower halo of the Milky Way, co-rotating with the Galactic disk. The sample was observed with the 100-m Effelsberg telescope. Search criteria were angular sizes $s$, the brightness temperatures \\TB and line width \\FWHM~ which are considered to be typical for halo clumps. The sample includes \\hi~ clumps with the following properties:  \\begin{itemize}  \\item they reside in the outer galaxy with  Galactocentric radii $R$ $ 10 < R < 15 $ kpc. \\item they belong to the lower halo ($ 0.9 < z < 5.4$ kpc).  \\item the gas is cold, with a median \\TKIN $ \\sim 600$ K and a line width \\FWHM=5.3\\kms.  \\item the sample shows a prominent two-component structure. Cold \\hi~ cores are surrounded by an extended component with broad line emission. \\end{itemize}  Two of the most prominent \\hi~ clouds were observed using synthesis arrays, the WSRT and the VLA. These high-resolution observations resolve the clouds into a conglomeration of arc-minute sized \\hi~ cores. These cores are embedded in a more diffuse medium which is detectable only with single dish telescopes. The cores contain a significant fraction of the \\hi~ mass and tend to be in pressure equilibrium with the surrounding envelopes.  Taking into account the influence of turbulence onto  the line widths \\FWHM, the median line width values of 3.3 \\kms and 4.3 \\kms observed at the cores implies that the \\hi~ gas is very cold.   Estimating densities and pressures for clumps and surrounding envelopes, we find some scattering but also a reasonable agreement with models which predict pressure equilibrium and a multi-phase structure caused by thermal instabilities \\citep{Wolfire2003}. The clumps tend to populate unstable regions in the phase diagrams, in agreement with recent predictions of turbulence driven instabilities \\citep{Audit2005,Gazol2005}.  Comparing samples observed with big single dish telescopes e.g. GBT \\citep{Lockman2002}, Effelsberg (this work), and Parkes \\citep{Ford2008ApJ}, we find similar column densities, peak temperatures, line widths and masses. Our interferometer observations imply that some of the derived parameters may be heavily biased if the small scale structure observed by us may be considered as typical for \\hi~ halo clumps."
    },
    "0911/0911.3905_arXiv.txt": {
        "abstract": "{} {We aim to demonstrate the usability of mm-wavelength imaging data obtained from the APEX-SZ bolometer array to derive the radial temperature profile of the hot intra-cluster gas out to radius $r_{500}$ and beyond. The goal is to study the physical properties of the intra-cluster gas by using a non-parametric de-projection method that is, aside from the assumption of spherical symmetry, free from modeling bias.} {We use publicly available X-ray spectroscopic-imaging data in the 0.7--2 keV energy band from the \\textit{XMM-Newton} observatory and our Sunyaev-Zel'dovich Effect (SZE) imaging data from the APEX-SZ experiment at 150 GHz to de-project the density and temperature profiles for a well-studied relaxed cluster, Abell 2204. We derive the gas density, temperature and entropy profiles assuming spherical symmetry, and obtain the total mass profile under the assumption of hydrostatic equilibrium. For comparison with X-ray spectroscopic temperature models, a re-analysis of recent \\textit{Chandra} observation is done with the latest calibration updates. We compare the results with that from an unrelaxed cluster, Abell 2163, to illustrate some differences between relaxed and merging systems.} {Using the non-parametric modeling, we demonstrate a decrease of gas temperature in the cluster outskirts, and also measure gas entropy profiles, both of which are done for the first time independently of X-ray spectroscopy using the SZE and X-ray imaging data. The gas entropy measurement in the central 100 kpc shows the usability of APEX-SZ data for inferring cluster dynamical states with this method. The contribution of the SZE systematic uncertainties in measuring $T_e$ at large radii is shown to be small compared to  \\textit{XMM-Newton} and \\textit{Chandra} systematic spectroscopic errors. The total mass profile obtained using the hydrostatic equilibrium assumption is in agreement with the published X-ray and weak lensing results; the upper limit on $M_{200}$ derived from the non-parametric method is consistent with the NFW model prediction from weak lensing analysis.} {} ",
        "introduction": "Current cosmological models are built upon two complementary approaches of astronomical observation: the statistical study of the ensemble properties in a large sample of objects (i.e. from surveys) and the detailed analysis of the individual objects for gaining better understanding of the physical processes affecting those ensemble properties. This is particularly important in the study of galaxy clusters, where extraction of cosmological parameters from large survey samples (X-ray, optical, or in the radio/mm wavebands) relies critically on our understanding of different mass observables, which depends on the detailed physical processes affecting constituent gas and galaxies.  Accurately determining the thermodynamic state of the intra-cluster medium (ICM) out to a large radius is critical for understanding the link between cluster mass and observables. For over a decade, observations of the thermal Sunyaev-Zel'dovich Effect (tSZE, hereafter simply SZE; Sunyaev \\& Zel'dovich 1970, Birkinshaw 1999) have been considered as a promising complement to X-ray observations for modeling the ICM in galaxy clusters, yet only recently has it been possible to make meaningful de-projections of gas temperature and density profiles using  SZE imaging data from multi-pixel bolometer arrays, in combination with X-ray data. The APEX-SZ experiment (Dobbs et al. 2006, Halverson et al. 2009) employs one of the first such powerful multi-pixel Transition-Edge Sensor (TES) bolometer cameras, and a joint analysis of the ICM properties using SZE and X-ray data has been presented by Nord et al. (2009, Hereafter NBP09) for the massive cluster Abell 2163.   In this paper we use the de-projection method used in NBP09 on the prototypical relaxed cluster Abell 2204. Our non-parametric analysis does not rely on any prior physical models in the construction of temperature and density profiles (apart from the assumption of spherical symmetry), hence the results are not based on parametric model fits. We measure the ICM entropy profile, as well as demonstrate the decrease of the ICM temperature in the cluster outskirts, first time from an SZE imaging data and independently from the X-ray spectroscopy. The derived ICM and cluster properties are compared with available X-ray and lensing results to highlight the level of accuracy of this independent method.  Joint SZE/X-ray de-projection analysis is expected to become a standard tool in the near future for understanding the ICM physical state, as large numbers of resolved SZE maps will be available from the new generation SZE experiments. Our  analysis assumes the gas to be in thermal equilibrium to model its physical properties, but presence of multi-phase ICM due to gas clumping will drive the electron temperature lower than the ion temperature in the electron-ion plasma (Evrard et al. 1996, Nagai et al. 2000). Recent hydro-simulations by Rudd \\& Nagai (2009) have shown, with a limited sample of halo models, that this deviation is small (about 5\\%) near $r_{200}$ for a relaxed cluster. Joint SZE/X-ray analysis using interferometric measurement of the SZE with OVRO/BIMA (Reese et al. 2002) has already shown that clumping effects are not large in the cluster interior (within $r_{500}$). Jia et al. (2008) have demonstrated the effect of the gas clumping on SZE and X-ray derived gas temperatures, and also found that these two quantities are in very good agreement within $r_{500}$ for the massive relaxed cluster RXC J2228.6$+$2036. But at large radii the gas should get clumpier, due to the onset of filamentary structures. One vital goal for sensitive imaging of the SZE signal using wide-field, multi-pixel bolometer cameras, and its combination with the X-ray and weak-lensing measurements, will be to provide an ultimate tool for measuring the gas clumping and thermodynamic state near the cluster virial radius, to give a dynamic view on the growth of clusters through accretion.   \\subsection{Previous SZE/X-ray joint modeling}  Due to the unavailability of resolved SZE images most of previous SZE/X-ray joint analysis studies have been limited to analytical or numerically simulated cluster models with idealized noise properties. Zaroubi et al. (2001) considered a method for reconstructing the triaxial structure of clusters based on Fourier slice theorem and applied it to a set of cluster simulations. Lee \\& Suto (2004) also considered de-projection method combining SZE and X-ray data and applied to analytical cluster models. Puchwein \\& Bartelman (2006) have employed the Richardson-Lucy de-projection technique to reconstruct the ICM and probe the dynamical state of clusters from simulations, and Ameglio et al. (2007) used a joint SZE/X-ray likelihood function maximization using a Monte Carlo Markov Chain (MCMC) for a similar objective.  Modeling ICM properties from real SZE observations has been limited mainly to isothermal $\\beta$-models (Cavaliere \\& Fusco-Femiano 1978).  Holzapfel et al. (1997), Hughes \\& Birkinshaw (1998) used isothermal models to constrain the Hubble parameter from observations of the clusters Abell 2163 and CL 0016$+$16, respectively, and later Reese et al. (2002) extended this analysis to a sample of 18 clusters detected by OVRO/BIMA.  De Filippis et al. (2005) used published SZE decrement values and X-ray imaging data to constrain the triaxial structure of clusters using isothermal $\\beta$-models. Zhang \\& Wu (2000) similarly used the $\\beta$-model to combine SZE and X-ray data to derive central gas temperature in clusters. A more detailed parametric modeling has been done by Mahdavi et al. (2007) for the cluster Abell 478, using simultaneous fits to the X-ray, lensing and SZE data assuming parametric models for dark matter, gas and stellar mass distribution, and hydrostatic equilibrium.  Yoshikawa \\& Suto (1999) first used Abel's integral inversion technique, originally proposed by Silk \\& White (1978), for a non-parametric reconstruction of radial density and temperature profiles using analytical and simulated cluster models. More recently Yuan et al. (2008) has extended this method for the most X-ray luminous cluster RXC J1347.5-1145 using published $\\beta$-model fit values from SZE and X-ray measurements. Extrapolation of the density and temperature profiles to the cluster outskirts based on such parametric modeling can be problematic, in particular for clusters with a very peaked central emission such as RXC J1347.5-1145. Additionally, this cluster is considered to be a merging system (Cohen \\& Kneib 2002) where the assumptions of spherical symmetry and hydrostatic equilibrium may not be valid. The nearest approach to non-parametric modeling was made by Kitayama et al. (2004) for the same cluster, RXC J1347.5-1145, using a beta-model density profile to fit the X-ray surface brightness and obtaining fitted temperature values separately in each radial bin from their SZE imaging data. The small extent of their SZE map (less than 2 arcmin) limited the temperature modeling again to the cluster core region.  \\subsection{Scope of the present work}  In this paper we apply the non-parametric ICM modeling based on Abel's integral inversion technique, as presented in NBP09, to the well studied and dynamically relaxed galaxy cluster Abell 2204  ($z=0.1523$, $L_{\\mathrm{X}}=26.9\\times 10^{44}$ $h_{50}^{-2}$ ~erg s$^{-1}$ in the $0.1-2.4$ keV band, $T_{\\mathrm{X}} = 7.21 \\pm 0.25$ keV; Reiprich \\& B\\\"ohringer 2002). The only assumptions in this analysis are spherical symmetry for reconstructing temperature and density profiles, and hydrostatic equilibrium (HSE) for reconstructing the total mass profile. The primary aim is to confirm the validity of this method for modeling the ICM distribution and cluster mass -- and compare the results with those obtained from deep X-ray spectroscopic and weak lensing data -- in a cluster where the assumptions of spherical symmetry and HSE are generally accepted to be valid.  We compute the \\textit{Chandra} spectral temperature profile with the latest calibration updates and  compare it with the SZE-derived temperature profile. In contrast to the X-ray spectroscopic  measurements from \\textit{Chandra}, the SZE-derived ICM temperature measurements near the cluster virial radius are constrained primarily by the statistical uncertainties in the SZE data. This fact demonstrates the potential for stacking the SZE signal of several relaxed clusters to put tighter constraints on the slope of the gas temperature profile in the cluster outskirts (Basu et al., in preparation). For a single cluster (Abell 2204), the depth in the APEX-SZ map allows us to model the temperature profile with meaningful errors up to $\\sim80\\%$ of the cluster virial radius (which we take to be $r_{200}$, the radius within which the mean total density is 200 times the critical density).  From density and temperature profiles we derive other physical properties like total gravitational mass, gas mass fraction and the gas entropy index. The total mass modeling provides a quantitative comparison with the published X-ray and lensing results. The modeling of the gas entropy profile from SZE/X-ray imaging data is a first, and we compare the central entropy values of two clusters with different morphologies, A2204 and A2163 (APEX-SZ analysis of the latter was presented in NBP09). This comparison shows how the gas entropy in the cluster core derived from SZE/X-ray joint modeling can be used to infer the dynamical state of clusters without the need for X-ray spectroscopy. A further comparison of the baryonic fraction of the ICM between A2204 and two other dynamically complex clusters detected by APEX-SZ (Bullet and A2163) illustrates a statistically significant difference of $f_{\\mathrm{gas}}$ inside $r_{2500}$.  All the scientific results in this paper are computed from the radial profiles of two observables: the SZE temperature decrement at 150 GHz, and the X-ray surface brightness in the $0.7-2$ keV band of XMM-Newton. In \\S2 and \\S3, we describe the map making and radial profile extraction steps from the X-ray and SZE data, and discuss the different systematic uncertainties associated with each profile. \\S4 describes Abel's integral inversion method and presents our primary results in the form of the radial density and temperature profiles. In \\S5 we present the other derived quantities like gas entropy and the total cluster mass profiles, and list the conclusions in \\S6.  We use the currently favored $\\Lambda$CDM cosmology with the following parameters:  $\\Omega_m=0.27$, $\\Omega_{\\Lambda}=1-\\Omega_m=0.73$, and the Hubble constant $H_0=70$ km s$^{-1}$ Mpc$^{-1}$. At redshift of $z=0.1523$, the angular diameter distance of Abell 2204 is 541.6 Mpc. To put the radial profiles in perspective using the  characteristic cluster radii, we adopted the maximum likelihood NFW fit parameters from Corless et al. (2009), $M_{200} = 7.1\\times 10^{14} M_{\\odot}$ and $c=4.5$, which gives  $r_{200}=1.76$ Mpc ($11.2'$), $r_{500}=1.16$ Mpc ($7.3'$) and $r_{2500}=0.51$ Mpc ($3.2'$). ",
        "conclusions": "\\begin{enumerate}  \\item We describe the detailed application of a direct, nonparametric de-projection method of cluster density and temperature profiles, using APEX-SZ and XMM-Newton data.  The method was presented in NBP09, the current paper builds upon the previous work by applying this technique to the well-studied relaxed cluster Abell 2204.  \\item Analysis of both SZE and X-ray data have been done from their raw data sets, to create images and radial profile. In particular, we describe the creation of a set of SZE radial profiles, all consistent with the APEX-SZ measurement, to characterize the statistical uncertainties on the bin values and minimize the numerical errors in Abel's de-projection method. Our final results are dominated by the statistical uncertainties in the SZE data, the signal at $r_{200}$ is essentially an upper limit for A2204. We describe the different sources of systematic uncertainties and include them in the analysis.  \\item The decreasing gas temperature in the cluster outskirts is demonstrated for the first time from SZE measurements, using a broad re-binning of the APEX-SZ (and X-ray) data. The temperature drop can be confirmed to $98\\%$ confidence level. We also compare the direct de-projected pressure profile with some parametric models, and show that the Nagai profile is adequate for modeling the gas pressure, within the current statistical uncertainties in APEX-SZ imaging of a single cluster.  \\item We re-perform the X-ray spectral analysis for the ICM temperature profile from publicly available \\textit{Chandra} data, primarily to find the changes from the recent calibration updates (CALDB 4), but also to show the effect of systematic uncertainties due to the background modeling in the X-ray spectral analysis. A comparison with the projected temperature profile obtained from SZE data confirms that our SZE derived temperature values are much less affected by systematic uncertainties at large radii, in comparison with \\textit{Chandra} and \\textit{XMM-Newton}. Precise comparison between the SZE and X-ray spectroscopic measurements of the gas temperature in the cluster outskirts will be a promising method to constrain gas clumping and non-LTE effects.  \\item The integrated total mass profile is computed assuming hydrostatic equilibrium for the cluster gas. The mass profile is in excellent agreement with the recent X-ray and weak lensing analyses. Our model prediction for $M_{500}$ is $(2.6\\pm 2.2) \\times 10^{14}~h^{-1} M_{\\mathrm{\\odot}}$. This is somewhat lower than the X-ray and lensing results but consistent within $1\\sigma$ errors. The upper limit on $M_{200}$ from our analysis is in good agreement with the published NFW model fit from weak lensing analysis of A2204.  \\item The ICM mass fraction as function of radius is computed using the non-parametric modeling, and found to be below 0.1 within $r_{500}$. The low $f_{\\mathrm{gas}}$ value in A2204 in the cluster center is in contrast with the previous APEX-SZ measurement of this ratio for Abell 2163 and the Bullet cluster.  \\item We compute the ICM entropy profile from SZE/X-ray joint analysis and confirm the general agreement with the self-similar cluster model predictions within the present statistical uncertainties. The significance of the APEX-SZ measurement of A2204 is not sufficiently high at $r\\gtrsim r_{500}$ to constrain the slope of the entropy profile in the cluster outskirts.  \\item We compare the entropy profiles of Abell 2204 (relaxed) and Abell 2163 (merging system), using the same non-parametric SZE/X-ray de-projection and radial binning, and find a clear entropy difference in their central 200 kpc. This corresponds to the different dynamical states of these two clusters and seen for the first time from SZE derived $T_{\\mathrm{gas}}$ measurement.  \\end{enumerate}"
    },
    "0911/0911.1307_arXiv.txt": {
        "abstract": "The Phase-Induced Amplitude Apodization (PIAA) coronagraph is a high performance coronagraph concept able to work at small angular separation with little loss in throughput. We present results obtained with a laboratory PIAA system including active wavefront control. The system has a 94.3\\% throughput (excluding coating losses) and operates in air with monochromatic light.  Our testbed achieved a 2.27 $10^{-7}$ raw contrast between 1.65 $\\lambda/D$ (inner working angle of the coronagraph configuration tested) and 4.4 $\\lambda/D$ (outer working angle). Through careful calibration, we were able to separate this residual light into a dynamic coherent component (turbulence, vibrations) at 4.5 $10^{-8}$ contrast and a static incoherent component (ghosts and/or polarization missmatch) at 1.6 $10^{-7}$ contrast. Pointing errors are controlled at the $10^{-3}$ $\\lambda/D$ level using a dedicated low order wavefront sensor.  While not sufficient for direct imaging of Earth-like planets from space, the  2.27 $10^{-7}$ raw contrast achieved already exceeds requirements for a ground-based Extreme Adaptive Optics system aimed at direct detection of more massive exoplanets. We show that over a 4hr long period, averaged wavefront errors have been controlled to the 3.5 $10^{-9}$ contrast level. This result is particularly encouraging for ground based Extreme-AO systems relying on long term stability and absence of static wavefront errors to recover planets much fainter than the fast boiling speckle halo. ",
        "introduction": "\\label{sec:intro} An imaging system aimed at dection or characterization (spectroscopy) of exoplanets must overcome the large contrast beween the planet and its star. This is particularly challenging for Earth-like planets, where the contrast is $\\approx 10^{-10}$ in the visible and the angular separation is 0.1\\arcsec for a system at 10pc. Many coronagraph concepts have recently been proposed to overcome this challenge (see review by \\cite{guyo06}). Among the approaches suggested, Phase-Induced Amplitude Apodization (PIAA) coronagraphy is particularly attractive. In a PIAA coronagraph, aspheric optics (mirrors or lenses) apodize the telescope beam with no loss in throughput. A PIAA coronagraph combines high throughput, small inner working angle (2 $\\lambda/D$ for $10^{-10}$ contrast), low chromaticity (when mirrors are used), full 360 degree discovery space, and full $1 \\lambda/D$ angular resolution. Angular resolution (size of the planet's PSF in the image) is a critical performance parameter as exoplanet imaging sensitivity, even if speckles have been perfectly removed, is usually background-limited due to sky or thermal emission (near-IR and mid-IR imaging from the ground) or zodiacal and exozodiacal light (direct imaging of Earth-like planets in the visible from space). The PIAA concept, orginally formulated by \\cite{guyo03}, has since been studied in depth in several subsequent publications \\citep{trau03,guyo05,vand05,gali05,mart06,vand06,pluz06,guyo06,beli06,guyo09,lozi09}, which the reader can refer to for detailed technical information.  In the first laboratory demonstration of the PIAA concept \\citep{gali05}, lossless beam apodization was demonstrated, and the field aberrations introduced by the PIAA optics were confirmed experimentally. In this first prototype, the PIAA acrilic optics lacked surface accuracy required for high contrast imaging, and since this experiment did not include active wavefront control, the high contrast imaging potential of the technique could not be demonstrated. In the present paper, we report on results obtained with a new system which includes reflective PIAA optics and wavefront control. Our prototype combines the main elements/subsystems envisionned for a successful PIAA imaging coronagraph instrument, with the exception of corrective optics required to remove the strong off-axis aberrations introduced by the PIAA optics. This last subsystem has been designed and built for another testbed, and its laboratory performance is reported in a separate paper \\citep{lozi09}.  The overall system architecture adopted for our test is presented and justified in \\S\\ref{sec:archi}. The design of the main components of the coronagraphs (PIAA mirrors, masks) is also described in this section. Wavefront control and calibration are discussed in \\S\\ref{sec:wfc}. Laboratory results are presented in \\S\\ref{sec:labresults}.   \\begin{figure*}[htb] \\includegraphics[scale=0.6]{f1.eps} \\caption{\\label{fig:optlayout} Optical layout of the laboratory PIAA coronagraph system. The grey shaded area shows the rigid PIAA bench on which the two PIAA mirrors are mounted. The light source is at the upper corner of the figure. The focal plane mask (near the bottom, center) separates light into the imaging channel and the coronagraphic low order wavefront sensor (CLOWFS) channel.} \\end{figure*} ",
        "conclusions": "The results obtained in this experiment are especially encouraging for ground-based coronagraphy. The 2 $10^{-7}$ raw contrast we have achieved in the 1.65 to 4.4 $\\lambda/D$ angular separation range already exceeds by two orders of magnitudes the raw contrast that can be hoped for in even a theoretically ideal Extreme-AO system \\citep{guyo05}. We note that with a more careful optical design and anti-reflection coated optics, our experiment could probably have reached 5 $10^{-8}$ raw contrast (level of coherent light leak in the current experiment). More importantly, we have demonstrated that with the coronagraph + wavefront control architecture adopted in our experiment, static coherent speckles can be pushed down very low (3.5 $10^{-9}$) in long exposures. Our system successfully removed long term correlations in the coherent speckles, and their averaged level in long exposure was reduced with a $1/\\sqrt{T}$ law. The combination of a high performance PIAA coronagraph and a focal-plane based wavefront control therefore appears extremely attractive for ground-based Extreme-AO. In that regard, our experiment has been a successful validation of the key technologies and control algorithms of the Subaru Coronagraphic Extreme-AO (SCExAO) system currently in assembly. The major differences between the SCExAO PIAA coronagraph and our laboratory prototype are (1) the need to design and operate a PIAA coronagraph on a centrally obscured pupil with thick spider vanes and (2) the need for corrective optics to recover a wide field of view. These two requirements have been validated in a separate laboratory experiment using the final SCExAO coronagraph optics \\citep{lozi09}.  Our experiment was limited at the 2 $10^{-7}$ contrast by an optical ghost and at the 5 $10^{-8}$ contrast by turbulence or vibrations. The PIAA coronagraph could therefore not be tested to the contrast level required for direct imaging of Earth-like planets from space (approximately $10^{-10}$), although several key concepts were demonstrated, including simultaneous operation of a low-order wavefront sensor using starlight in the PSF core and high-order wavefront sensor using scattered light in the science focal plane. New calibration schemes which will be very useful for high contrast coronagraphy were also developed and validated, such as the system response matrix optimization loop, which can executed as a background task to fine-tune the system.  PIAA coronagraph technologies for high contrast space applications are now being actively developed and tested at NASA Ames Research Center and NASA Jet Propulsion Laboratory. A new set of PIAA mirrors was recently manufactured to higher surface accuracy than the ones used for this demonstration, and is being integrated within the High Contrast Imaging Testbed (HCIT) vacuum chamber at NASA Jet Propulsion Laboratory. We note that the HCIT chamber has already demonstrated stability to the $10^{-10}$ contrast with a Lyot-type coronagraph \\citep{trau07}. The experiment described in this paper served as a precursor to this new step, which is aimed at reaching higher contrast (minimum goal of $10^{-9}$) in broadband light using a two-DM wavefront correction. In parallel to this effort, a highly flexible high stability air testbed at NASA Ames Research Center is coming online to further explore technology and architecture trades for PIAA systems.  In addition to pushing the contrast further, future laboratory demonstration of the PIAA coronagraph will explore chromaticity and wavefront control issues unique to a PIAA coronagraph architecture. The monochromatic experiment described in this paper did not address these important points, and should therefore be considered as only a first step toward validation of the PIAA coronagraph technique for high contrast space-based exoplanet imaging mission."
    },
    "0911/0911.4714_arXiv.txt": {
        "abstract": "s{ We report on the observations  of  cosmic rays with energies $ \\geq $ $10^{18}$ eV from Jan  2004 to April 2009 by the Pierre Auger Observatory. During this period the Observatory has grown from about 300 surface detectors to about 1600 upon its completion in November 2008. The 1600 surface detectors are overlooked by 24 fluorescence telescopes. We report on measurements of the cosmic ray spectrum, the arrival directions and the elongation rate. We also report limits for the photon and neutrino components of this cosmic radiation.}  \\begin{figure} \\begin{minipage}[b]{0.47\\linewidth} \\includegraphics  [width=3.0in] {blois1.jpg} \\label{fig:blois1.jpg} \\end{minipage} \\hfill \\begin{minipage}[b]{0.47\\linewidth} \\includegraphics  [width=3.0in] {blois2.jpg} \\label{fig:blois2.jpg} \\end{minipage}\\\\ \\begin{minipage}[t]{0.47\\linewidth} \\caption{\\it Cosmic ray spectrum (data compiled by Simon Swordy, University of Chicago).} \\end{minipage} \\hfill \\begin{minipage}[t]{0.47\\linewidth} \\caption{\\it GZK horizons for protons and nuclei (figure courtesy Denis Allard).} \\end{minipage} \\end{figure} ",
        "introduction": "The cosmic radiation discovered by Hess \\cite{Hess} extends from very low energies $\\leq$ 10$^{6}$ eV to $\\geq$ 10$^{20}$ eV. The latter energy is equal to 16 joules - a macroscopic energy in a microscopic particle as the cosmic rays are principally atomic nuclei ranging from protons to iron. Figure 1 shows the full cosmic ray spectrum. It is roughly a  power law falling by 30 orders of magnitude in flux over 10 orders of magnitude increase in energy. The upper end of the spectrum represents a  mystery as there is no clear understanding of how Nature can accelerate atomic nuclei to such high energies.  The study of this category of cosmic rays is a scientific imperative and Nature provides two important analytical tools for the investigation.  First, the very highest energy cosmic rays must come from nearby. Consequently one can expect that there are a small number of sources that can contribute to the flux of the highest energy cosmic rays. Protons interact with the cosmic microwave background (CMB) losing energy while producing pions. This is the famous GZK effect. Complex nuclei are photo disintegrated by the CMB.  The result of these interactions is that half of the cosmic rays with  energy $\\geq$ 6x10$^{19}$ eV must come from distances less than 70 Mpc \\cite{Allard}. In Figure 2 we show the expected distribution of  distances for several nuclear species on the basis of a uniform source distribution. It is noteworthy that for distances $\\geq$ 50 Mpc only protons and iron nuclei survive.  In composition analysis at these high energies the assumption of two components is more than just an ansatz - it is a reasonable assumption.  Second, the higher energies and shorter distances will reduce the effects of the random magnetic fields which at lower energies  decouple the observed arrival direction from the true direction of the source.  Thus it is quite possible that the arrival directions for energies $\\geq$ 6x10$^{19}$ will correlate with the distribution of extragalactic objects located within 100 Mpc.    \\begin{figure} \\begin{minipage}[b]{0.47\\linewidth} \\includegraphics  [width=2.0in] {blois3.jpg} \\label{fig:blois3.jpg} \\end{minipage} \\hfill \\begin{minipage}[b]{0.47\\linewidth} \\includegraphics  [width=3.0in] {blois4.jpg} \\label{fig:blois4.jpg} \\end{minipage}\\\\ \\begin{minipage}[t]{0.47\\linewidth} \\caption{\\it Cartoon showing the two techniques for detection of air showers (figure courtesy Enrique Zas .} \\end{minipage} \\hfill \\begin{minipage}[t]{0.47\\linewidth} \\caption{\\it Cosmic ray spectra from HiRes and AGASA circa 2004   .} \\end{minipage} \\end{figure} ",
        "conclusions": ""
    },
    "0911/0911.2768_arXiv.txt": {
        "abstract": "The  preliminary results of a three-site CCD photometric campaign are reported. The  $\\delta$ Scuti variable V650 Tauri  belonging to the Pleiades cluster was observed photometrically for 14 days on three continents during 2008 November.  An overall run of 164 hr of data was collected. At least five significant frequencies for V650 Tauri have been detected. ",
        "introduction": "$\\delta$ Scuti variables are stars with masses between 1.5 and 2.5 $M_{\\odot}$ located at the intersection of the classical Cepheid instability strip with the main sequence. These variables are thought to be excellent laboratories for probing the internal structure of intermediate mass stars. Intents of modelling $\\delta$ Scuti stars belonging to open clusters have been performed recently (e.g. [1], [2], [3]). Although the constraints imposed by the cluster parameters have proved to be very useful when modelling an ensemble of $\\delta$ Scuti stars, more detected frequencies in individual stars would improve current seismic studies.   \\bigskip The target star V650 Tau (HD 23643, $V=7^{\\rm m}.79$, A7)  was identified as a short-period pulsating variable by Breger (1972). Intensive observations performed by the STEPHI network in November 1990, revealed four frequency peaks in V650 Tau [4]. One-site CCD photometric observations carried out by [5] in November-December 1993, confirmed the results obtained by the STEPHI campaign. Since then, no new observations of V650 Tauri have been performed.  \\bigskip The present paper provides preliminary observational results of a three-site campaign on V650 Tauri in 2008. ",
        "conclusions": ""
    },
    "0911/0911.2597_arXiv.txt": {
        "abstract": "It is commonly stated that we have entered the era of precision cosmology in which a number of important observations have reached a degree of precision, and a level of agreement with theory, that is comparable with many Earth-based physics experiments. One of the consequences is the need to examine at what point our usual, well-worn assumption of homogeneity associated to the use of perturbation theory begins to compromise the accuracy of our models. It is now a widely accepted fact that the effect of the inhomogeneities observed in the Universe cannot be ignored when one wants to construct an accurate cosmological model. Well-established physics can explain several of the observed phenomena without introducing highly speculative elements, like dark matter, dark energy, exponential expansion at densities never attained in any experiment (i.e. inflation), and the like. Two main classes of methods are currently used to deal with these issues. Averaging, sometimes linked to fitting procedures a la Stoegger and Ellis, provide us with one promising way of solving the problem. Another approach is the use of exact inhomogeneous solutions of General Relativity. This will be developed here. ",
        "introduction": "\\label{int}  The observation of its structures show that our Universe is not homogeneous. We see voids, groups of galaxies, clusters, superclusters, walls, filaments, etc. However, it is usually argued in the literature that the Universe should be nearly homogeneous at large scales which is supposed to validate the use of Friedmannian models. But how large these scales are and what nearly does imply is never precisely stated.  It has however become, during the last few years, as a widely accepted fact, that the effect of the inhomogeneities cannot be ignored when one wants to construct an accurate cosmological model up to the regions where structures start forming and their evolution becomes non-linear. Three different methods have been proposed to deal with this issue:  \\begin{enumerate}  \\item Linear perturbation theory. However, this method is only valid when {\\it both} the curvature and density contrasts remain small, which is not the case in the non-linear regime of structure formation and where the SNe Ia are observed.  \\item Averaging methods `{\\`a} la Buchert', promising, but needing to be improved (see \\cite{TB08} and references therein).  \\item Exact inhomogeneous solutions, valid at all scales and exact perturbations of the Friedmann background which they can reproduce as a limit with any precision.  \\end{enumerate}  The use of such exact solutions shows that well-established physics can explain several of the phenomena observed in astrophysics and cosmology without introducing highly speculative elements, like dark matter, dark energy, exponential expansion at densities never attained in any experiment (inflation), and the like. Here, we foccuss on the application of a couple of exact solutions of general relativity to structure formation and evolution and to the reproduction of cosmological data. ",
        "conclusions": "The increasing precision of observational data implies that FLRW models must now be considered just a zeroth order approximation, and linear perturbation theory a first order approximation whose domain of validity is an early, nearly homogeneous Universe.  In the nonlinear regime, which was entered since structures formed, there is no escape from the use of exact methods (or of averaging schemes aiming at investigating this issue from the standpoint of backreaction).  In the era of `precision cosmology', the effect of the inhomogeneities on the determination of the cosmological models cannot be ignored. Inhomogeneous models constitute an exact perturbation of the Friedmann background and can reproduce it as a limit with any precision. This is the reason why they are fully adapted for the purpose of studying astrophysical and cosmological effects and for constructing precise models of universe.  While using L--T models with a central observer to represent our `local' Universe averaged over angles around us, a giant void is not mandatory to explain away dark energy. A giant overdensity can also do the job. However while neither the void nor the overdensity are directly observable, the giant void alone can be tested with observations of the density function on our past light cone. The giant hump will need more and more precise data to be constrained.  Exact inhomogeneous solutions can be employed not only for studying the geometry and dynamics of the Universe, but also to investigate the formation and the evolution of structures. They give enhanced formation efficiency and might therefore help solving the problem of structure formation pertaining to the standard model.  While the L--T models have been mostly used up to now for modeling the inhomogeneities of the Universe, the need of getting rid of spherical symmetry for this purpose will lead the cosmological community to consider other solutions, and among them QSS models, for the future developments of inhomogeneous cosmology."
    },
    "0911/0911.4843_arXiv.txt": {
        "abstract": "Previous velocity images which reveal flows of ionized gas along the most prominent cometary tail (from Knot 38) in the Helix planetary nebula are compared with that taken at optical wavelengths with the Hubble Space Telescope and with an image in the emission from molecular hydrogen. The flows from the second most prominent tail from Knot 14 are also considered. The kinematics of the tail from the more complex Knot 32, shown here for the first time, also reveals an acceleration away from the central star.  All of the tails are explained as accelerating ionized flows of ablated material driven by the previous, mildly supersonic, AGB wind from the central star. The longest tail of ionized gas, even though formed by this mechanism in a very clumpy medium, as revealed by the emission from molecular hydrogen, appears to be a coherent outflowing feature. ",
        "introduction": "Dense, neutral (\\citealt{mea92}; \\citealt{hug92}) knots with ionized cometary tails are found in the central regions of planetary nebulae (PNe); the most famous being those in the Helix nebula for, at a distance of only 213 $+$30/$-$16 pc \\citep{har07}, they were easily detected and resolved by early ground--based observations (Baade and reported by \\citealt{vor68}). The incredibly clumpy nature of the neutral material in the disk of the Helix nebula (when modelled as a bi--polar PNe viewed along its axis, \\citealt{mea98}; \\citealt{mea05}) has now been revealed in spectacular fashion in the imagery in the H$_{2}$ emission line by \\citet{mat09}. Previously, \\citet{mei05} had estimated that there were 23,000 cometary knots and that inevitably tail interactions must be occurring. \\citet{mat09} show conclusively that it is an inner region in the Helix disk, towards the central star, where tails are observed from neutral globules surrounded by an outer clumpy region free of such tails.  \\citet{dys03} reviewed the two broad models for the creation of the cometary tails; either they are shadowed from the ionizing radiation of the central star, and their surfaces photo-ionized by scattered Lyman photons in the nebula, or they are dynamically produced as the particle winds from the central star swept past the slowly expanding system of dense globules.  A critical distinction between these models arises if the cometary tails can be seen to be flowing along their lengths away from their parent globules. Only the `dynamic' model would have this effect. In fact, a numerical simulation of the flow of a moderately supersonic particle wind past the ionized head of a globule \\citep{dys06} predicted not only the creation of a cometary tail but a moderate acceleration of ionized material parallel to the tail surface and away from the globule. In this initial model the neutral material away from the globule had a smooth density distribution.  The most comprehensive optical study of the kinematics of the Helix system of globules remains that described in \\citet{mea98}. Here, spatially resolved \\nii\\ profiles, from 300 separate long-slit (163\\arcsec\\ long) positions, were obtained with the Manchester echelle spectrometer (MES -- \\citealt{mea84} but now with a CCD as the detector) on the Anglo--Australian telescope in exceptional `seeing' conditions, over three separate `blocks' of globules and their tails in the nebular core. Furthermore, a kinematical `case study' of the globule which is apparently the 2nd closest to the nebular core (Knot 14), and a length of its tail, was carried out in a range of emission lines. The principal outcome was to show that the system of central globules is concentrated in a disk expanding itself at 14\\kms. Flows parallel to the surfaces of the cometary tails of two of the most prominent globules (Knots 38 and 14) were also detected and even an acceleration of this flow along the length of the longest tail (from Knot 38) suggested. \\citet{ode07} challenged these latter assertions without making any further kinematical observations but simply because they claim that more recent HST images reveal confusing minor globules in the longest tail (Knot 38). With the dismissal of these kinematical effects they proceeded to support the shadowing theories of the creation of the tails. The aim of the present paper is to reassess the strength of the deductions from the \\citet{mea98} kinematical data set in the light of the subsequent HST and very latest H$_{2}$ imagery by \\citet{mat09}. In particular, to consider if the evidence in \\citet{mea98} for flows along the tail surfaces of Knots 38 and 14, away from the central star, is now invalidated by this HST imagery alone as suggested by \\citet{ode07}. Furthermore, the flow behind Knot 32, which was not considered hitherto, is now presented to strengthen the original suggestion that accelerating flows in the cometary tails could be ubiquitous though in a clumpy medium. \\section[]{Knots 38 and 14} Knot 38 has the longest (62\\arcsec) cometary tail and is the apparently closest knot to the central star of NGC 7293, while Knot 14 is the next closest. The heads of both have arcs of \\oiii\\ emission facing this star which indicates that they protrude into the hard radiation field of the central volume of the bi--polar, ellipsoidally--shaped, nebula that produces the characteristic helical appearance in emission lines of lower ionization species (\\citealt{mea98}; \\citealt{mea05}). The HST archival images (PI NAME: Meixner, PID:9700) of Knots 38 (HST ACS/WFC J8KR14040) and 14 (HST ACS/WFC J8KR0840) are presented respectively in Figs.  1a and 2.  These should be compared with those taken with the New Technology Telescope (NNT - Chile) and presented in Meaburn et al (1998). All are in the light of the \\ha\\ plus \\NII\\ nebular emission lines.  The uniquely long ionized tail of Knot 38 in Fig. 1 is prominent in both the HST and former \\citep{mea98} NTT images and appears to be a coherent structure in both i.e. it is not simply a consequence of chance superposition of a large number of fore-- or background tails along the same sightlines. Minor ionized knots appear along the tail in the image in Fig. 1a but it is the H$_{2}$ 2.12$\\mu$m image \\citep{mat09} shown in Fig. 1b that emphasises that the tail of ionized gas, being so long, is formed around an internal core composed of a large number of minor clumps of neutral material.  With this as a starting point the kinematics along this tail should be re--considered by examining the \\nii\\ line profiles presented in \\citet{mea98}. The centroids of these profiles are shown in fig. 13 of that paper to change along the 62\\arcsec\\ length of this tail by a radial velocity difference (from central knot to tail end) by 10\\kms. The velocity images in figs. 4 and 12 of the same paper confirm this systematic radial velocity change in a different way. Unfortunately, the four images in different heliocentric radial velocity (\\vhel) ranges became jumbled in the production of fig. 12 in the \\citet{mea98} paper. These should be \\vhel\\ $= -$31 to $-$27 (top right), $-$25 to $-$21 (top left), $-$20 to $-$15 (bottom right) and $-$14 to $-$10 (bottom left) and all \\kms. When these four images are considered with this correction it is clear that the head of Knot 38 appears alone in the top right frame then progressively the tail appears in the subsequent three frames towards more positive velocities as the image of the head declines. The acceleration along the tail length, if regarded as a coherent feature starting at Knot 38, is very clear and not dominated by the minor confusing knots apparent in Fig. 1a. Those areas free of these along the tail length of Knot 38 in Fig. 1 clearly show this radial velocity change. The appearance of the neutral material in the tail of Knot 38 and seen in Fig. 1b suggests that the ionized outflow is in a sheath around a clumpy neutral core.  Similarly, the \\nii\\ line profiles up to 5\\arcsec\\ from the head of Knot 14 (fig. 9 of \\citealt{mea98}) show a systematic change of radial velocity of their centroids of -6 \\kms. This was not covered by the H$_{2}$ imagery of \\citep{mat09}. The tail from Knot 14  was modelled kinematically as a flow parallel to the globule and tail surfaces. The HST  images in Fig. 2 show that there are no significant minor ionized knots along this small length of the tail. Again an accelerating flow away from the central star is indicated. If tilted at 25\\degree\\ to the plane of the sky this change of outflow velocity amounts to 14 \\kms\\ compared to the apex of the head of the cometary knot. ",
        "conclusions": ""
    },
    "0911/0911.4591_arXiv.txt": {
        "abstract": "Apart from the 11-year solar cycle, another periodicity around 155-160 days was discovered during solar cycle 21 in high energy solar flares, and its presence in sunspot areas and strong magnetic flux has been also reported.  This periodicity has an elusive and enigmatic character, since it usually appears only near the maxima of solar cycles, and seems to be related with a periodic emergence of strong magnetic flux at the solar surface.  Therefore, it is probably connected with the tachocline, a thin layer located near the base of the solar convection zone, where strong dynamo magnetic field is stored. We study the dynamics of Rossby waves in the tachocline in the presence of a toroidal magnetic field and latitudinal differential rotation. Our analysis shows that the magnetic Rossby waves are generally unstable and that the growth rates are sensitive to the magnetic field strength and to the latitudinal differential rotation parameters. Variation of the differential rotation and the magnetic field strength throughout the solar cycle enhance the growth rate of a particular harmonic in the upper part of the tachocline around the maximum of the solar cycle. This harmonic is symmetric with respect to the equator and has a period of 155-160 days. A rapid increase of the wave amplitude could give place to a magnetic flux emergence leading to observed periodicities in solar activity indicators related with magnetic flux. ",
        "introduction": "During solar cycle 21, a short periodicity between 152--158 days was discovered in $\\gamma$ ray flares \\citep {Rieger84}, X ray flares \\citep{Rieger84, Dennis85, Bai87, Kile91, Dimitropoulou08}, flares producing energetic interplanetary electrons \\citep{Droge90}, type II and IV radio bursts \\citep{Verma91}, and microwave flares \\citep{Bai85, Kile91}.  However, this periodicity was absent during solar cycle 22 \\citep{Kile91, Bai92a, ozguc94}.  The periodicity has also been detected in indicators of solar activity (sunspot blocking function, sunspot areas, ``active'' sunspot groups, group sunspot numbers) which suggest that it is associated preferentially with photospheric regions of compact magnetic field structures \\citep{Lean89, Lean90, pap90, carball90, Bou92, carball92, Verma92, oliver98, ballester99, Krivo02}.  Probably, the most important, and enigmatic, feature of the periodicity is that it appears during epochs of maximum activity and that it occurs in episodes of 1 to 3 years.  \\citet{rabin91} performed a study of the magnetic flux variations during solar cycle 21 which reveals the existence of quasi-periodic pulses or episodes of enhanced magnetic activity. The duration of the pulses is $\\approx$ 5 rotations during the years around maximum activity, the epoch in which the flare periodicity appears, and the comparison with magnetic field maps indicates that those pulses of activity correspond to the occurrence of complex active regions containing large sunspots \\citep{Bai87a}.  \\citet{ballester02, ballester04} analyzed several data sets of, or strongly related to, photospheric magnetic flux to point out that the appearance of the near 160-day periodicity in different manifestations of solar activity during solar cycle 21 has its underlying cause in the appearance of the periodicity in the magnetic flux linked to regions of strong magnetic field.  They also showed that during solar cycle 22 the periodicity does not appear in the photospheric magnetic flux records and, as a consequence, the periodicity did not appear in other solar activity indicators, while during solar cycle 23 it appeared in the photospheric magnetic flux but not in other solar activity indicators.  Several mechanisms have been put forward in order to explain the existence of this periodicity.  \\citet{wolff83} linked it to the interaction of rotating features (active longitude bands) resulting from g-modes with $l=2$ and $l=3$.  \\citet{Bai87b} suggested that the cause of this periodicity must be a mechanism that causes active regions to be more flare productive.  Later, \\citet{Bai87} concluded that it cannot be due to the interaction of \\lq \\lq hot spots\", i.e. regions where flare activity is higher than elsewhere \\citep{Bai87a, Bai88}, rotating at different rates and that the cause must be a mechanism involving the whole Sun.  \\citet{Ichi85} suggested that it is related to the timescale for storage and/or escape of magnetic fields in the solar convection zone.  \\citet{B90}, taking into account the possible intermittency of the periodicity, suggested that this behavior could be simulated with a damped, periodically forced non-linear oscillator, which shows periodic behavior for some values of the parameters and chaotic behavior for other values. \\citet{wolff92} argued that such periodicity can be understood in terms of the normal modes of oscillation of a nearly spherical, slowly rotating star, when two r-modes (inertial modes) couple with an interior g-mode beat.  This suggestion seems to agree qualitatively with the fact that the periodicity is stronger around the activity maximum.  \\citet{Bai91} and \\citet{stu92} proposed that the Sun contains a \\lq \\lq clock\", modeled by an oblique rotator or oscillator, with a period of 25.8 days and suggested that the periodicity of 154 days is just a subharmonic of that fundamental period.  Later, \\citet{Bai93} modified the earlier period to the value 25.50 days, but that model seems to be very constrained by helioseismological data about the rotation of the Sun's interior. \\citet{lou00} suggested that such periodicities can be related to large-scale equatorially trapped Rossby-type waves showing that, for typical solar parameters, the periods of these waves (with n = 1 and m even) are in good agreement with the observed ones. Moreover, \\citet{lou00} has also pointed out that such waves can give rise to detectable features, such as surface elevations in the photosphere. Coincidently, \\citet{kuhn00} have reported observations made with MDI onboard SOHO and claim to have detected a regular structure of 100-m \\lq \\lq hills'', uniformly spaced over the surface of the Sun with a characteristic separation of 90,000 km.  They suggest that this structure is the surface manifestation of Rossby waves, or r-modes oscillations. Finally, \\citet{Dimitropoulou08} have linked the found periodicities in different classes (B, C, M, X) of X-ray flares with the theoretical periods derived by \\citet{lou00}, pointing out that odd m periodicities are also frequent and significant.  On the other hand, most of the proposed mechanisms to explain solar flares, specially the most energetic ones, accept as a prerequisite the emergence of magnetic flux \\citep{priest90, forbes91} which, by reconnection with the ambient field, triggers the destabilization of active regions.  Based on this mechanism, Carbonell \\& Ballester (1990, 1992) suggested that the periodic increase in the occurrence rate of energetic flares is related to a periodic emergence of magnetic flux through the photosphere. Later, Oliver et al.  (1998) showed that during solar cycle 21 there was a perfect time correlation between the intervals of occurrence of the periodicity in sunspot areas and energetic flares, and \\citet{ballester02} clearly pointed out that in cycle 21, and during the time interval in which the periodicity appeared, there was a perfect time and frequency coincidence between the impulses of high-energy flares and those corresponding to strong photospheric magnetic flux. The efficiency of the reconnection mechanism depends on the geometry of the two flux systems \\citep{Galsgaard07} and recent high resolution observations performed by \\citet{zuca08} have confirmed the suitability of the mentioned mechanism for flare production.  Emerged magnetic flux is probably connected to deeper regions, namely to the tachocline, which is a thin, transition layer between differentially rotating convection zone and rigidly rotating radiative envelope. The tachocline may prevent the spreading of the solar angular momentum from the convection zone to the interior \\citep{spi92,gough98,gough07,gar07} and probably it is the place, where the large-scale magnetic field which governs the solar activity is generated/amplified.  The observed periodicity of 155--160 days in the emerging flux is in the range of Rossby wave spectrum. Therefore, we suggest that the periodicity is connected to the Rossby wave activity in the tachocline. Rossby waves are well studied in the geophysical context \\citep{gill82,ped87}, however, the presence of magnetic fields significantly modifies their dynamics \\citep{zaqa07,zaqa09}. On the other hand, the differential rotation, which is inevitably present in the tachocline, may lead to the instability of particular harmonics of magnetic Rossby waves. It has been shown that the joint action of toroidal magnetic field and the differential rotation generally leads to tachocline instabilities \\citep{gilman97,cal03,dik05,gil07,gil007}. However, the stability analysis usually has been performed in an inertial frame, which complicates to extract the information about unstable Rossby modes. Therefore, it is of paramount importance to perform the stability analysis in a rotating frame. Another important point is that the consideration of a rotating frame may tighten the stability criteria as it has been suggested by \\citet{hug01}. The difference between the present analysis and that by \\citet{hug01} is the inclusion of rotation which allows us to obtain Rossby wave solutions.  In this paper, we use a rotating spherical coordinate system to study the linear stability of magnetic Rossby waves in the solar tachocline taking into account the latitudinal differential rotation and the toroidal magnetic field. We perform a two dimensional analysis, which can be followed in the future by more sophisticated shallow water considerations \\citep{gil00}. We first derive the analytical conditions of instability similar to \\citet{dahlburg98} and \\citet{hug01}. Then, we perform  a detailed stability analysis using Legendre polynomial expansions \\citep{longuet68} to obtain the spectrum of unstable harmonics of magnetic Rossby waves. ",
        "conclusions": "In summary, we have shown that the destabilization of magnetic Rossby waves in the solar tachocline is produced by the joint effect of the latitudinal differential rotation and the toroidal magnetic field. The frequencies and growth rates of unstable harmonics depend on the combination of the differential rotation parameters and the magnetic field strength. The possible increase of latitudinal differential rotation at the solar maximum may trigger the instability of symmetric harmonic with period of 155-160 days in the upper part of the tachocline. This instability has a direct correlation with magnetic flux emergence, therefore the periodicity also appears in solar activity indicators related with magnetic flux. Later on, and probably via reconnection, this periodic magnetic flux emergence triggers the observed periodicity in solar flares.    The magnetic Rossby wave theory opens a new research area about the activity on the Sun and other stars, and magnetic Rossby waves can be of paramount importance for observed intermediate periodicities in solar and stellar activity \\citep{massi98, massi05}.   {\\bf Acknowledgements} The authors acknowledge the financial support provided by MICINN and FEDER funds under grant AYA2006-07637.  Also, the Conselleria d'Economia, Hisenda i Innovaci\\'o of the Government of the Balearic Islands is gratefully acknowledged for the funding provided under grant PCTIB2005GC3-03. T. V. Z. acknowledges financial support from the Austrian Fond zur F\\\"orderung der wissenschaftlichen Forschung (under project P21197-N16), the Georgian National Science Foundation (under grant GNSF/ST06/4-098) and the Universitat de les Illes Balears. Wavelet software was provided by C. Torrence and G. Compo \\footnote{The software is available at http://paos.colorado.edu/research/wavelets}.  \\appendix"
    },
    "0911/0911.4072_arXiv.txt": {
        "abstract": "{} {Our main aim is to study the influence of the initial conditions of a cloud in the intermediate/high-mass star formation process.} {We observed with the VLA, PdBI, and SMA the centimeter and millimeter continuum, \\nth\\,(1--0), and CO\\,(2--1) emission associated with a dusty cloud harboring a nascent cluster with intermediate-mass protostars.} {At centimeter wavelengths we found a strong source, tracing a \\uchii\\ region, at the eastern edge of the dusty cloud, with a shell-like structure, and with the near-infrared counterpart falling in the center of the shell. This is presumably the most massive source of the forming cluster. About $15''$ to the west of the \\uchii\\ region and well embedded in the dusty cloud, we detected a strong millimeter source, MM1, associated with centimeter and near-infrared emission. MM1 seems to be driving a prominent high-velocity CO bipolar outflow elongated in the northeast-southwest direction, and is embedded in a ridge of dense gas traced by \\nth, elongated roughly in the same direction as the outflow. We estimated that MM1 is an intermediate-mass source in the Class 0/I phase. About $15''$ to the south of MM1, and still more deeply embedded in the dusty cloud, we detected a compact millimeter source, MM2, with neither centimeter nor near-infrared emission, but with water maser emission. MM2 is associated with a clump of \\nth, whose kinematics reveal a clear velocity gradient and additionally we found signposts of infall motions. MM2, being deeply embedded within the dusty cloud, with an associated water maser but no hints of CO outflow emission, is an intriguing object, presumably of intermediate mass.} {The \\uchii\\ region is found at the border of a dusty cloud which is currently undergoing active star formation. Two intermediate-mass protostars in the dusty cloud seem to have formed after the \\uchii\\ region and have different properties related to the outflow phenomenon. Thus, a single cloud with similar dust emission and similar dense gas column densities seems to be forming objects with different properties, suggesting that the initial conditions in the cloud are not determining all the star formation process.} ",
        "introduction": "}  It is well established that intermediate (2--8~\\mo) and high-mass ($\\gtrsim8$~\\mo) stars form in cluster environments (\\eg\\ Kurtz \\et\\ \\cite{kurtz2000}; Evans \\et\\ \\cite{evans2009}), and that at the very first stages of their formation, they are embedded within dense and cold gas and dust from their original natal cloud. However, it is not clear to what extent the initial conditions of the natal cloud are determining the star formation story of the cloud and the properties of the nascent stars of low, intermediate, and high mass, forming within it. Regarding the star formation story, there is increasing evidence that a single cloud can undergo different episodes of star formation, as suggested by studies of deeply embedded clusters observed at spatial scales similar to the cluster member separation ($\\sim5000$~AU). In these studies, the intermediate/high-mass young stellar objects (YSOs) in the forming cluster seem to be in different evolutionary stages (judging from the peak of their spectral energy distributions: Beuther \\et\\ \\cite{beuther2007}; Palau \\et\\ \\cite{palau2007a},b; Leurini \\et\\ \\cite{leurini2007}; Williams \\et\\ \\cite{williams2009}).   Concerning the properties of the intermediate/high-mass YSOs forming within the cloud, a few studies toward massive star-forming regions have revealed that two objects formed in the same environment may have very different properties in the ejection of matter, with one YSO driving a highly collimated outflow nearby another YSO driving an almost spherical mass ejection (\\eg\\ Torrelles \\et\\ \\cite{torrelles2001}, \\cite{torrelles2003}; Zapata \\et\\ \\cite{zapata2008}), indicating that the initial conditions in the cloud may not be the only agent determining the formation and evolution of the intermediate/high-mass members in a forming cluster. While most of these cases have been found at very high spatial scales ($\\sim1$~AU), a broad observational base in other regions and at other spatial scales is required to properly understand the role of initial conditions of the cloud in the cluster formation process and ultimately in the formation of intermediate/high-mass stars.  In this paper we show a high angular resolution study of a deeply embedded cluster whose intermediate-mass protostars show differences not only in their spectral energy distributions, but also in the ejection phenomena associated with the YSOs. The region, IRAS\\,00117+6412, was selected from the list of Molinari \\et\\ (\\cite{molinari1996}) in a search for deeply embedded clusters which are luminous ($>1000$~\\lo) and nearby (distance $<3$~kpc). The IRAS source has a bolometric luminosity of 1400~\\lo\\ and is located at a distance of 1.8~kpc (Molinari \\et\\ \\cite{molinari1996}). The millimeter single-dish image reported by S\\'anchez-Monge \\et\\ (\\cite{sanchezmonge2008}) shows strong emission with some substructure, tracing a dusty cloud, and the centimeter emission reveals an ultra-compact \\hii\\ (\\uchii) region, associated with the brightest 2MASS source of the field at the eastern border of the dusty cloud. The dusty cloud is associated with an embedded cluster reported by Kumar \\et\\ (\\cite{kumar2006}), two H$_2$O masers spots (Cesaroni \\et\\ \\cite{cesaroni1988}; Wouterloot \\et\\ \\cite{wouterloot1993}), and CO\\,(2--1) bipolar outflow emission (Zhang \\et\\ \\cite{zhang2005}; Kim \\& Kurtz \\cite{kimkurtz2006}). All this is suggestive of the dusty cloud harboring a \\uchii\\ region and a deeply embedded stellar cluster forming in its surroundings.  We conducted high-sensitivity radio interferometric observations in order to study the different millimeter sources embedded in the dusty cloud. To properly characterize the protostars, we also studied the distribution of dense gas, outflow and ionized gas emission. In the present paper we show the first results obtained from \\nth\\ dense gas and CO outflow emission, focusing mainly on the intermediate/high-mass content of the protocluster. ",
        "conclusions": "In this paper we study with high angular resolution the centimeter, and millimeter continuum, and \\nth\\,(1--0), and CO\\,(2--1) emission of the intermediate-mass YSOs forming within a dusty cloud, with the goal of assessing the role of the initial conditions in the star formation process in clusters. Our conclusions can be summarized as follows:  \\begin{enumerate}  \\item A \\uchii\\ region is found at the eastern border of the dusty cloud, with a shell-like structure and a flat spectral index, $-0.03\\pm0.08$. The estimated age and mass of the underlying star is $\\sim1$~Myr and $\\sim6$~\\mo.  \\item Deeply embedded within the dusty cloud, we have discovered a millimeter source, MM1, associated with a 2MASS infrared source, which is driving a CO\\,(2--1) powerful and collimated high-velocity outflow, oriented in the southwest-northeast direction. The mass derived from the millimeter continuum emission for MM1 is $\\sim3$~\\mo. MM1 is embedded within a ridge of dense gas as traced by \\nth, which seems to be rotating roughly along the outflow axis. MM1 is associated with centimeter emission, whose spectral index is compatible with an ionized wind, and at 1.2~mm splits up into different subcomponents when observed with an angular resolution of $\\lesssim1''$. From the derived outflow momentum rate, we estimated a luminosity for MM1 of 400--600~\\lo. Thus, MM1 seems to be a Class 0/I intermediate-mass YSO.  \\item About $\\sim15''$ to the south of MM1, our observations revealed a dust compact condensation, MM2, lying in a dark infrared region, associated with water maser emission and a dense core traced by \\nth\\ emission. The mass from the dust emission is $\\sim1.7$~\\mo, and the \\nth\\ excitation temperature and column density are similar to the ones derived for MM1. The dense core in MM2 is rotating in the same sense as the ridge associated with MM1 and seems to be undergoing infall. We modeled the MM2 dense core as a disk-like structure with an inner radius of $\\sim1''$ and an outer radius of $\\sim5''$, with a rotation velocity in the outer radius of $\\sim0.6\\pm0.2$~\\kms, and an infall velocity at the same radius of $\\sim0.17\\pm0.02$~\\kms. The non-detection of CO at any velocity toward MM2 makes this object intriguing.  \\item Although MM1 and MM2 formed within the same cloud and have similar dust and dense gas emission, their properties, specially concerning the ejection phenomenon, seem to be different and could be indicating that the initial conditions in a cloud forming a cluster are not the only agent determining the properties of the members of the cluster.  \\end{enumerate}"
    },
    "0911/0911.3652_arXiv.txt": {
        "abstract": "We investigate the nature and spatial variations of turbulence in the Small Magellanic Cloud (SMC) by applying several statistical methods on  the neutral hydrogen (HI) column density image of the SMC and a database of isothermal numerical simulations. By using the 3rd and 4th statistical moments we derive the spatial distribution of the sonic Mach number (${\\cal M}_s$) across the SMC. We find that about 90\\% of the HI in the SMC is subsonic or transonic. However, edges of the SMC `bar' have ${\\cal M}_s\\sim4$ and may be tracing shearing or turbulent flows. Using numerical simulations we also investigate how the slope of the spatial power spectrum depends on both sonic and Alfv\\'en Mach numbers. This allows us to gauge the Alfv\\'en Mach number of the SMC and conclude that its gas pressure dominates over the magnetic pressure. The super-Alfv\\'enic nature of the HI gas in the SMC is also highlighted by the bispectrum, a three-point correlation function which characterizes the level of non-Gaussianity in wave modes. We find that the bispectrum of the SMC HI column density displays similar large-scale  correlations as numerical simulations, however it has localized enhancements of correlations. In addition, we find a break in correlations at a scale of $\\sim160$ pc. This may be caused by numerous expanding shells of a similar size. ",
        "introduction": "\\label{intro}  In the recent decade, many advances in both observations and computational models have provided new insights into the workings and evolution of the interstellar medium (ISM). The emerging picture is that interstellar turbulence plays the key role in ISM structure formation and evolution (McKee \\& Ostriker 2007). In the Galaxy and the Magellanic Clouds, the ISM is turbulent on scales ranging from less than a parsec to a few kiloparsecs (Crovisier \\& Dickey 1983; Stanimirovic et al. 1999; Deshpande et al. 2000; Dickey et al. 2001; Elmegreen et al. 2001; Elmegreen \\& Scalo 2004). Although the observational evidence for the importance of turbulence in the ISM is overwhelming, many questions remain open. For example, what are the dominant energy sources and physical processes that convert kinetic energy into turbulence (Burkert 2006)? At what scales and through which modes is turbulent energy dissipated (Heyer \\& Zweibel 2004)? How do the level and type of turbulence depend on properties of the interstellar gas (e.g. presence/absence of star formation, presence/absence of tidal effects, or the strength of magnetic field)? Since no complete theory of astrophysical turbulence exists, studying its effects on the multiphase ISM can be challenging and calls for a combination of numerical and observational efforts.   Statistical studies have proved to be important in characterizing the magnetized turbulent ISM  \\cite[]{Lazarian09}, however the interpretation of results is not always straight forward. Several statistical methods have been extensively used for both observational and synthetic data. These statistics include probability density functions (PDFs), wavelets, the principal component analysis,  higher order moments, Velocity Coordinate Spectrum (VCS), and Velocity Channel Analysis (VCA), to name just a few (Gill \\& Henriksen 1990; Brunt \\& Heyer 2002; Kowal, Lazarian \\& Pogosyan 2000, 2004, 2006;  Lazarian \\& Beresnyak 2007).   Most of these statistical methods require large datasets with a large spatial or velocity dynamic range, and produce a single, mostly one-dimensional, measure. This results in the lack of spatial information about turbulent properties across a given interstellar cloud, or a galaxy, making a connection with underlying physical properties highly difficult.   In this paper we explore a new method for obtaining spatial information about the level and nature of ISM turbulence on the neutral hydrogen (HI) observations of the Small Magellanic Cloud (SMC). The SMC, a dwarf irregular galaxy in the Local Group, has a highly gas rich ISM environment (see, Stanimirovic et al. 1999, henceforth known as SX99), and is an excellent candidate for ISM studies. Being nearby (60 kpc, \\cite{West91}), the SMC is distant enough for all its objects to be treated as having roughly the same distance, unlike the Milky Way where distance determination is relatively uncertain.     The HI observations of the SMC, obtained using the Australia Telescope Compact Array (ATCA) and the Parkes telescope (SX99), have been used for several investigations,  including the HI spatial power spectrum and the kinematic study of HI, which revealed the existence of many expanding shells of gas and three supergiant shells. The power law index of HI density and velocity distributions was derived in SX99 and \\cite{Stan01}, while the Genus statistic in \\cite{Chep08} revealed spatial variations of HI morphology. Because the SX99 SMC data set is  well studied, it is a perfect candidate to investigate new statistical methods. We can acquire new information, but also test and confirm past results, as well as validate the promise of these statistical tools for further use in other observational studies.      In this study, we investigate turbulent properties of the HI in the SMC by applying the higher order moments on the HI column density image. We then use a database of MHD simulations to bootstrap the spatial distribution of the sonic Mach number across the SMC. The crucial aspect of our approach is the confluence of observations and numerical simulations: only by combining the two we can retrieve the spatial variations of turbulent properties. This is the reason why we oscillate between observational and synthetic data in this paper. We also investigate whether and how interstellar shocks leave footprints on the HI gas by employing the bispectrum, a three point statistical measure, on the SMC HI column density image. Again, to interpret our results we apply the same statistics on the database of MHD simulated column density images.   In particular, the paper is organized as follows. We start with \\S~\\ref{sec:background} by providing a brief summary of previous work regarding the statistical methods used in our study. In \\S~\\ref{sec:data} we describe the SMC HI column density map and the database of numerical simulations of compressible MHD turbulence used for the comparison with observations. In \\S~4 we introduce higher order moments and their dependence on the sonic and Alfvenic Mach numbers. We then apply higher order moments on the SMC HI observations to derive an image of the sonic Mach number across the SMC, in \\S~5. We  compare our results with an observational estimate of the sonic Mach number of the cold neutral medium (CNM) in the SMC, based on a comparison of the spin and kinetic temperature of HI absorption profiles, in \\S~\\ref{ratio}. In \\S~\\ref{ps} we show how the power-law slope of the spatial power spectrum depends on the sonic Mach number and use this to gauge the Alfvenic Mach number of the HI gas seen in emission. In \\S~\\ref{sec:bispectrum} we present an analysis of the bispectrum of the SMC, as well as a brief discussion of the noise and windowing effects. In \\S~\\ref{sec:Discussion} we provide a discussion of our results, followed by our conclusions. ",
        "conclusions": "\\label{sec:Discussion}  \\subsection{Turbulence properties of the HI gas in the SMC}  Using 3rd and 4th statistical moments of the HI column density image and boothstreping turbulent information from a database of isothermal MHD simulations, we have mapped spatial variations of the sonic Mach number across the SMC. While most of the HI seen in emission in the SMC appears to be subsonic or transonic, several supersonic regions have emerged from our study.  It is interesting that these regions do not correlate well with the most recent sites of star formation and seem to point out to large scale shearing or tidal flows. Commonly, it is believed that supernovae and superbubbles are the main drivers of galactic turbulence (McCray \\& Snow 1979), with a typical size of $\\sim100$ pc. While we do not have high enough resolution to see changes on such small scales ($\\sim10'$) in our derived map of ${\\cal M}_s$, most of the star-forming bar of the SMC appears to have subsonic or transonic properties when viewed at resolution of 30$'$.   The most turbulent regions in the SMC may be tracing some kind of shearing flows between the SMC bar and the surrounding HI. This suggests that SMC's chaotic history with the LMC  and our own Milky Way has probably left strong turbulent imprints on the HI gas. The lack of a turnover in the HI spatial power spectrum on the largest observed scales is also indicative of the fact that turbulent energy injection happens largely on scales larger than the size of the SMC  \\cite[]{Stan01}. Similarly, Goldman (2000) suggested that the HI turbulence in the SMC was induced by large-scale flows from tidal interactions with the Milky Way and the LMC about $2\\times10^8$ yrs ago. Such large-scale bulk flows could have generated turbulence through shear instabilities, and this turbulence has not have had enough time to decay.    Most of the HI in the SMC has a sonic Mach number of 1-2. This is on average at least two times smaller than what we inferred from HI absorption observations for the CNM in the SMC, ${\\cal M}_s\\sim$3.5-4. Similarly, for the CNM in the Milky Way Heiles \\& Troland (2003) found ${\\cal M}_s \\sim3$ with a large dispersion. A sonic Mach number of about 4-5 is commonly assumed for cold gas in molecular clouds (Federrath et al. 2009). For example, Heyer et al. (2006) measured from CO observations ${\\cal M}_s=4.2\\pm0.17$ for the Rosette molecular cloud, and ${\\cal M}_s=4.7\\pm0.12$ for G216-2.5. On the other hand, Hill et al. (2008) found that the distribution of the warm ionized medium (WIM) in the Milky Way can be best fit by models for mildly supersonic turbulence with ${\\cal M}_s\\sim1.4-2.4$. Turbulent properties of the HI in emission in the SMC are therefore closer to properties of the WIM in the Milky Way than properties of the CNM. This may suggest a large fraction of warm relative to cold HI being traced in HI emission.      In the Milky Way, the HI is known to consist of at least two components with different temperature: the WNM with $T_{warm}=6000$~K and the CNM with $T_{cold}=70$~K. In addition, there could be a substantial amount of gas at intermediate temperatures \\cite[]{Heiles03}. Due to its lower metallicity, the HI in the SMC has different properties. Dickey et al. (2000) found $T_{cold}=40$~K, in agreement with theoretical expectations by Wolfire et al. (1995) whereby the existence of the two-phase medium is possible only at higher pressures compared with the range that applies for solar neighborhood conditions. They also estimated the fraction of cold HI in the SMC to be $\\la15$\\%. This is lower than $\\sim25$\\% found for the Milky Way.  As our simulations are isothermal it is obvious to wonder how does the multi-phase HI affect our statistics and conclusions. We investigate this in Figure~\\ref{fig:ex} by producing a simulated data cube from a weighted combination of two cubes, one subsonic (${\\cal M}_{s}$=0.7) and one supersonic. The subsonic cube represents contribution from warmer gas, while the supersonic cube represents the cold gas. We combined the two cubes with different emphasis on warm vs cold gas, obtained the column density image of the resultant cube, and calculated its moments.    Figure~\\ref{fig:ex} shows that for the case when the supersonic cube has ${\\cal M}_{s}$=2.0 skewness of the final cube with up to 50\\% of subsonic gas will be biased towards supersonic gas and will appear dominated by cold gas. If we increase the sonic Mach number of the supersonic cube to 4.0, the dominance of the higher turbulence is even more pronounced. A cube with up to 60-70\\% of subsonic gas and 25\\% of supersonic gas, will still have high skewness biased by the supersonic contribution. Considering that the HI column density image results in the mean ${\\cal M}_{s}=1-2$, significantly lower than what is expected for the cold HI, the fraction of the CNM along any LOS is most likely $\\la25$\\%. This supports the Dickey et al. (2000) estimate of the fraction of cold HI in the SMC being about 15\\%.     \\begin{figure*}[tbh]\\centering \\includegraphics[scale=.7]{f16a.eps} \\caption{Skewness vs. percent of gas that is subsonic for column density of a $512^3$cube weighted with supersonic and subsonic gas.  Supersonic gas generally dominates the skewness of the column density PDF.} \\label{fig:ex} \\end{figure*}     \\subsection{Is the HI in the SMC sub-Alfv\\'enic or Super-Alfv\\'enic?}   Two different statistical approaches in our study suggest that the HI gas in the SMC seen in emission is super-Alfv\\'enic. As we have shown in Figure ~\\ref{fig:power}, in addition to the sonic Mach number the spectral slope of the spatial power spectrum is sensitive to the Alf\\'venic Mach number for ${\\cal M}_{s}<2$. The sub-Alfv\\'enic models generally show steeper slopes due to large scale influence of magnetic fields. Thus, if one independently knows the sonic Mach number, it is possible to estimate the Alf\\'venic one using just the column density data. While the dependence of the spectral slope on the Alf\\'venic Mach number has not received much attention in the past, it is somewhat expected. Essentially, magnetization decreases compression in the shocks. Strong magnetic forces mix up density clumps preventing formation of isolated peaks, which results in a steeper spectrum. In addition, in the sub-Alfv\\'enic case we expect oblique shocks to be disrupted by Alfv\\'en shearing, which in turn, produces more small scale shocks \\cite[]{beresnyak05}.  Another indication that the HI in the SMC is super-Alfv\\'enic comes from the bispectrum. The very sharp decrease in the bispectral amplitudes from large to small scales observed for the SMC is the closest to the trend found for simulated data for the case of ${\\cal M}_{s}\\sim2$ and  ${\\cal M}_{A}\\sim2$. Detailed comparison between simulated and observed bispectra awaits future work, however this qualitative comparison is certainly encouraging.    Assuming on average ${\\cal M}_{s}\\sim1-2$, the power spectrum slope suggests a super-Alfv\\'enic HI in the SMC with ${\\cal M}_{A}\\sim1-3$. This is generally in agreement with the observationally inferred strength of the magnetic field by Mao et al. (2008). Using their estimate for $B_{\\rm ext}=2$ $\\mu$G, a radius of the SMC of 2 kpc, the total hydrogen mass of $4.2\\times10^{9}$ M$_{\\odot}$, and a typical velocity dispersion of 20 \\kms (Stanimirovic et al. 2004), we estimate ${\\cal M}_{A}\\sim3$. As the Alf\\'venic Mach number shows the nature of the interplay between gas pressure and magnetic fields, it appears that the gas pressure in the SMC dominates over the magnetic pressure.      \\subsection{Intermittency in mode correlations? }  Our bispectrum analysis of the SMC HI data was the first attempt to apply bispectrum on observed astrophysical data. While more detailed comparison between observations and simulations awaits future work, we clearly see trends in the bispectral amplitudes similar to what was found for simulations of supersonic MHD turbulence. The most interesting finding is, however, the effect of small-scale variations in the $k_1=k_2$ correlations and a strong break in correlations at a scale of 160 pc. Such small-scale variations, or jumps, have not been seen in the bispectrum of simulated data. We can speculate about several possible scenarios that could explain their existence. The jumps could be caused by the energy injection due to processes other than turbulence affecting specific spatial scales. Alternatively, the jumps may be marking the presence of colder or multi-phase gas. Similarly, the observed break in the bispectrum at about 160 pc is intriguing. As we already pointed out, it is interesting that most expanding shells in the SMC (more than 500 were cataloged so far) have a diameter of $\\sim120$ pc. The break could be due to the lack of correlations on scales similar to the distance between two shell centers. Obviously this will require further studies.    \\subsection{Limitations of the present study}  A natural question to ask is how results presented in this paper depend on the resolution of numerical simulations. For example, Kritsuk et al. (2007) investigated how resolution of numerical simulations affects the power spectrum of density.  These authors found that the spectral index estimates based on low resolution simulations bear large uncertainties due to the bottle neck contamination, and that the power spectra of $512^3$ simulations are substantially shallower then  models with resolution of $1024^3$. However, while Kritsuk et al. (2007) only examined hydrodynamic turbulence, Beresnyak et al. (2008) showed that the slopes were very different between the MHD and pure hydrodynamic cases. For instance, the slopes for hydrodynamic simulations showed a pronounced and well defined bottleneck effect, while the MHD slopes were much less affected. This is indicative of MHD turbulence being less local than the Kolmogorov turbulence, and suggests that our simulations will be less affected by resolution. In addition, Kritsuk et al. (2007) found a difference in the slope between hydrodynamic $512^3$ and $1024^3$ simulations to be 0.17. This would result in a change of $d{\\cal M}_{s}\\sim0.5$ only and will not change our interpretation. We also add that in the case of higher statistical moments and the bispectrum BFKL confirmed trends noticed by KLB at lower resolution of 128$^3$.  Another issue that should be further addressed and that could affect our results is the type of numerical forcing of turbulence. Federrath et al. (2009) recently investigated the effects of solenoidal vs. compressive (divergence-free vs. curl-free) forcing on a variety of statistics including PDFs and higher order moments. They found that both types of driving mechanisms are compatible with observations of molecular clouds however, depending on the data studied, one type could be superior then the other in terms of the statistics and reproduced observables. This implies that different regions in  the SMC may exhibit statistical signatures of either compressive or solenoidally driven turbulence.   \\subsection{Summary}  \\label{sec:conclusions}  We have investigated a new method for constraining turbulent properties of the ISM, specifically the sonic Mach number, by using the HI column density image and a database of numerical simulations with a range of sonic and Alf\\'venic Mach numbers. By applying the 3rd and 4th statistical moments on both observed and simulated data we have derived the spatial distribution of the sonic Mach number across the SMC with angular resolution of 30$'$. To provide an estimate of the Alf\\'venic Mach number we used two approaches: the spatial power spectrum and the bispectrum. Using the database of numerical simulations we have shown that the spatial power spectrum varies with both the sonic and Alf\\'venic Mach numbers. If the sonic number is known the Alf\\'venic number can be constrained from this dependence. The bispectrum shows the level of correlation between turbulent eddies of different size and depends greatly on the sonic Mach number, and somewhat on the Alf\\'venic Mach number. By comparing the bispectra of observations and simulations we have gauged the importance of magnetic fields relative to the gas pressure in the SMC. The following results were discussed in the paper.    \\begin{itemize} \\item Skewness and kurtosis of the HI column density generally correlate well and are within the range expected from MHD simulations. This suggests that departures from Gaussianity could be interpreted as being governed by MHD turbulence.  \\item Most of the HI in the  SMC bar and the Eastern Wing is subsonic or transonic with ${\\cal M}_{s}\\sim0-2$. Sites of most recent star formation have ${\\cal M}_{s}\\sim1$. Regions with the highest skewness and kurtosis, which could be interpreted as having ${\\cal M}_{s}\\sim4$, correspond to the edges of the bar. The most turbulent regions are most likely tracing tidal or shearing flows. The fraction of the SMC with different turbulent properties is: 10\\% with ${\\cal M}_{s}>2$, 80\\% with $0<{\\cal M}_{s}<2$, and about 10\\% with very low values of ${\\cal M}_{s}$.   \\item Using HI absorption profiles from Dickey et al. (2000) we have estimated that the CNM in the SMC is highly supersonic with ${\\cal M}_{s}=3.5-4$. This is at least a factor of two higher than what we measured from the higher statistical moments for the HI gas seen in emission. One possible reason for this discrepancy could be that HI emission is dominated by warm gas and the fraction of the CNM in the SMC is $\\la20$\\%.   \\item The slope of the spatial power spectrum and the bispectrum suggest that the HI in the SMC is super-Alf\\'venic with ${\\cal M}_{A}\\sim1-3$. This is implies that the gas pressure dominates over the magnetic pressure.    \\item The bispectrum of the HI column density shows large scale wave correlations suggesting a large scale energy injection mechanism. Contrary to simulations which show a smooth decrease of wave-wave correlations from large to small scales, the SMC bispectrum shows localized enhancements of correlations and at least one prominent break at $\\sim160$ pc. We speculate that the multi-phase medium, and/or energy injection by processes other than turbulence, could be responsible for the correlation jumps. The break on the other hand appears at a scale similar to the diameter of the majority of expanding shells in the SMC.   \\end{itemize}"
    },
    "0911/0911.4134_arXiv.txt": {
        "abstract": "In this paper we examine the contribution of galaxies with different infrared (IR) spectral energy distributions (SEDs) to the comoving infrared luminosity density, a proxy for the comoving star formation rate (SFR) density. We characterise galaxies as having either a {\\em cold} or {\\em hot} IR SED depending upon whether the rest-frame wavelength of their peak IR energy output is above or below $90\\,\\um$. Our work is based on a far-IR selected sample both in the local Universe and at high redshift, the former consisting of {\\it IRAS} $60\\,\\um$-selected galaxies at $z< 0.07$ and the latter of {\\it Spitzer} $70\\,\\um$ selected galaxies across $0.1< z\\le 1$. We find that the total IR luminosity densities for each redshift/luminosity bin agree well with results derived from other deep mid/far-IR surveys. At $z<0.07$ we observe the previously known results: that moderate luminosity galaxies ($L_{IR}<10^{11}\\,L_\\odot$) dominate the total luminosity density and that the fraction of cold galaxies decreases with increasing luminosity, becoming negligible at the highest luminosities. Conversely, above $z=0.1$ we find that luminous IR galaxies ($L_{IR}>10^{11}\\,L_\\odot$), the majority of which are cold, dominate the IR luminosity density. We therefore infer that cold galaxies dominate the IR luminosity density across the whole $0< z< 1$ range, hence appear to be the main driver behind the increase in SFR density up to $z\\sim1$ whereas local luminous galaxies are not, on the whole, representative of the high redshift population. ",
        "introduction": "The rise in the comoving star formation rate (SFR) density up to $z\\sim1$ and its subsequent flattening \\citep{Lilly:96, Madau:96} has now been well studied at several wavelengths \\citep[e.g.][and references therein]{Bunker:04, Hopkins:06}. While the global picture to $z=1$ has been well constrained by observation, the details of how or why this change occurs remain largely unknown. There now is evidence that star formation depends on both galaxy mass \\citep[][]{Feulner:05, Juneau:05} and environment \\citep[][]{Lewis:02,Elbaz:07} presenting a more complicated picture of galaxy evolution than simple evolution of the luminosity function.  The infrared (IR), and particularly the far-IR, is one of the most powerful tracers of star formation as IR luminosity directly scales with SFR \\citep[][and references therein]{Kennicutt:98b} and far-IR emission originates from regions of cold dust and gas that constitute the fuel for an on-going burst of star formation. Another advantage of selecting sources at long wavelengths is the low contribution by active galactic nuclei (AGN), if present, to the total IR luminosity \\citep{Alexander:03, Clements:08}. Studies of the distant Universe at IR wavelengths remain limited and the most sensitive probe has been surveys with the {\\it Spitzer} $24\\,\\um$ band. However, this wavelength is relatively far from the peak of all but the hottest IR galaxies and progressively shifts to shorter wavelengths at higher redshifts where strong spectral features in the observed frame can also complicate matters. The few studies done with deep {\\it Spitzer} $70\\,\\um$ imaging have found rapid evolution in the total IR luminosity function \\citep{Huynh:07b, Magnelli:09}. When the IR luminosity function is integrated at different redshifts it is found that the IR luminosity density increases rapidly up to $z=1$ in a similar fashion to the SFR density. Nevertheless, the relative contribution by luminosity to the IR luminosity density changes with redshift with starbursts ($10\\le\\log(L_{IR}/L_\\odot)< 11$) dominating locally, but with luminous IR galaxies (LIRGs: $11\\le\\log(L_{IR}/L_\\odot)< 12$) and ultra-luminous IR galaxies (ULIRGs: $12\\le\\log(L_{IR}/L_\\odot)< 13$) becoming increasing contributors at higher redshifts \\citep{LeFloch:05,Magnelli:09}.  In \\citet[][hereafter S09]{Symeonidis:09} we studied the IR spectral energy distributions (SEDs) of a sample of $70\\,\\um$ selected galaxies at $z\\ge0.1$, using a sample of local, $z<0.1$,  {\\it IRAS} $60\\,\\um$ selected galaxies for comparison. We fitted the mid to far-IR photometry of both samples with models from \\citet{Siebenmorgen:07}. We found that the majority of the $70\\,\\um$ sources had IR SEDs which peaked, in $\\nu\\times F_\\nu$, at longer wavelengths than galaxies from the local sample with similar luminosities. In the local sample we observed a shift  of the IR SED peak to shorter wavelengths with increasing luminosity \\citep[as has previously been noted:][]{Sanders:96, Chapman:03a, Rieke:09}, whereas the $70\\,\\um$ sample had a wide range of peak wavelengths the distribution of which varied little with luminosity. The observation that the IR SEDs of luminous, distant galaxies were on average different to their local analogues was in contrast to other recent results, \\citet[e.g.][]{Magnelli:09} who concluded there was no significant change in IR SED of luminous galaxies with redshift. Models of galaxy evolution in the IR have tended to assume that the SEDs of high redshift luminous sources follow the luminosity trend seen in local sources \\citep[][]{Lagache:05,Pearson:05,LeBorgne:09,RowanRobinson:09}. However, as shown in S09, the range of IR SEDs for luminous galaxies is much wider at high redshifts than seen locally.  In this paper we examine the contribution of galaxies with different IR SEDs to the comoving IR luminosity density (IRLD). We will compare our results to earlier work \\citep{LeFloch:05} by examining the contribution to the IRLD by luminosity, but this time with a sample selected at $70\\,\\um$ rather than $24\\,\\um$. This selection enables us to use a more robust selection of IR luminous galaxies as $70\\,\\um$ lies closer to the peak of typical IR SEDs. Hence, the total IR luminosity can also be estimated more accurately, especially with constraints (mainly detections) from $160\\,\\um$ data. We examine the IRLD in more detail by focusing on the contribution by IR SED type within each bin. By obtaining an estimate of the peak of the IR SED in rest-frame $\\nu\\times F_\\nu$, we can characterise these IR bright sources as being either {\\em cold} or {\\em hot} depending on whether the SED peaks above or below $90\\,\\um$ ($32.2\\,$K for a blackbody). We choose this wavelength for two reasons. Firstly, local IR galaxies appear to have a warm IR component ($\\bar{\\lambda}\\sim 60\\,\\um$) associated with dust around young star forming regions and a cooler ``cirrus'' component ($\\lambda\\ge\\,100\\,\\um$) associated with more extended dust heated by the interstellar radiation field \\citep{Lonsdale:87}. Secondly, this $90\\,\\um$ cut also marks a divide between the SEDs of most local ULIRGs and those of cold SMGs seen at high redshift, most of which would be classified as cold by our definition \\citep{Chapman:05}. We note that there is a linear relationship between the {\\it IRAS} colours often used to define the IR galaxies as `cold' or `hot' \\citep[e.g.][]{Chapman:03a} and the peak wavelength used in this work.  We present our far-IR galaxy sample in \\S2, describe our determination of the comoving IRLD in \\S3. We present our results in \\S4 and discuss them in \\S5. Throughout we use the current `concordance' cosmology: $\\Omega_M = 1 - \\Omega_{\\Lambda} = 0.3$, $\\Omega_0 = 1$, and $H_0 = 70\\, \\kmpspMpc$ \\citep{Spergel:03}.     \\begin{figure} \\centerline{ \\psfig{file=fig1.ps,width=8.5cm,angle=270} } \\caption{The redshift/total ($8-1000\\,\\um$) IR luminosity  distribution of the local ($z<0.07$) BGS sample and our {\\it Spitzer} $70\\,\\um$ sample ($z>0.1$) with sources characterised as cold or hot depending on whether their IR SED peaks above or below $90\\,\\um$ (rest-frame in $\\nu\\times F_\\nu$). The dotted lines indicate the binning chosen to investigate the comoving IRLD by IR luminosity and SED.} \\label{fig:lz} \\end{figure} ",
        "conclusions": "In S09 we demonstrated that the IR SEDs of individual LIRGs and ULIRGs at $0.1\\le z\\le 1.0$ peak at longer wavelengths, and span a wider range, than local galaxies at similar luminosities. Other changes in the properties of star forming galaxies have occured since $z\\sim1$. Star formation shifts from massive galaxies at high redshifts to progressively less massive galaxies at lower redshifts \\citep[e.g.][]{Cowie:96,Juneau:05}. A shift is also seen in environment of the most strongly star forming galaxies with star formation tending to occur in over-dense environments at $z\\sim 1$ \\citep{Elbaz:07}, a reverse of the trend seen locally. While the dependence of star formation on stellar mass and environment inform us about the star formation history and potential or current galaxy mergers, information about the IR SED is a direct measure of the physical conditions in galaxies underlying major episodes of star formation.   In this paper we have expanded the results of S09 by deriving IR luminosity densities at different luminosities and epochs. We observe that the total IR luminosity densities derived from a $70\\,\\um$ sample are broadly consistent with results derived from shorter wavelengths (Fig~\\ref{fig:lumden}). The low IR luminosity density of the $z\\sim0.2$ starburst bin compared to other results is likely due to the effect of sample variance over the small cosmic volume probed. We also separate our sample into galaxies that can be characterised as having cold or hot SEDs and study their contribution to each redshift/luminosity bin. We find two striking new results. Firstly, we observed a rapid change in the make up of LIRG and ULIRG bins: locally they are dominated by hot galaxies, but at $z>0.2$ they are dominated by cold galaxies. Secondly, cold galaxies appear to be the dominant contributor to the total IR luminosity density over $0\\le z\\le 1.0$; when starburst galaxies dominate at $z<0.2$ we find that they are mainly cold and by the time LIRGs and ULIRGs begin to dominate at higher redshifts they are also mainly cold. Hence, a cold mode of star formation has dominated galaxy evolution since $z\\sim1$ and is likely responsible for the increase in the star formation rate density up to $z=1$.  While Fig.~\\ref{fig:lumden} does dramatically show the rise of cold galaxies in the LIRG and ULIRG bins at higher redshifts we note that hot galaxies are still found at those redshifts. In fact the results presented here do not rule out a rise with redshift in the number (and luminosity density) of hot galaxies. Therefore, the distribution of types of SEDs remains quite wide at high redshift. Fig.~\\ref{fig:frac} most strikingly demonstrates the evolution of the relative contribution of cold galaxies to each redshift/luminosity bin. The rise in the contribution of cold galaxies to the LIRG and ULIRG bins is most rapid between $z\\sim0$ and $z\\sim0.3$. We now consider if AGN could be responsible for the change in the fraction of cold galaxies. As a powerful AGN  would make the IR SED hotter the presence of an AGN can not explain the prevalence of cold galaxies at high redshift. AGN could, however, be the cause of the high fraction of hot galaxies locally. Studies have shown that the incidence of AGN are low in local LIRGs, $10-20\\%$, (Petric et al., 2009, in press), but while it is higher in ULRIGs, $\\sim40\\%$ \\citep{Farrah:07} the $8-1000\\,\\mu$m IR luminosity remains dominated by star formation. We note that at higher redshift the incidence of AGN, and their contribution to total IR luminosity, also remains low \\citep{Alexander:03, Clements:08}. Hence, AGN are not a significant contributor to the IR luminosity of the sources studied here and therefore are not responsible for any of the trends with redshift.  The prominence of cold galaxies discovered here might suggest that there could be a substantial contribution to the IR emission of luminous galaxies at high redshift from a cold cirrus component akin to that observed in local IR galaxies. Such an interpretation is in agreement with the observation that many of the distant, cold, SMGs are extended on scales of $\\sim10\\,$kpc \\citep{Chapman:04} in contrast with the compact ULIRGs seen locally, and is also consistent with the change in the mean IR SED of luminous galaxies as inferred from the observed $70\\,\\um$ to radio flux density ratio of high redshift star forming galaxies \\citep{Seymour:09}. In addition, changes in dust properties (such as opacity, grain distribution etc.) with redshift are also a possible cause of the larger spread in SED types at high redshift.  Most galaxy evolution models \\citep[e.g.][]{Lagache:05,Pearson:05,LeBorgne:09,RowanRobinson:09} assume that the IR SED is only dependent on IR luminosity, i.e. they typically use certain templates for star forming galaxies of a given luminosity, often using local galaxies like Arp 220 (which can be characterised as having a hot IR SED by our definition) for the most luminous population. This choice was often necessary as there were few constraints on any change with redshift until now. Models based on cooler SEDs could , for example, reconcile the flat or decreasing star formation rate density above $z=1$ and the far-IR and sub-millimetre source counts.  In conclusion, we have found that the star formation history of the Universe is dominated by cold galaxies across $0<z<1$ and that the hot luminous IR galaxies we see locally are not representative of the luminous galaxies that make up the bulk of the star formation at $0.1<z<1$."
    },
    "0911/0911.1527_arXiv.txt": {
        "abstract": "We investigate the RR Lyrae population in NGC 147, a dwarf satellite galaxy of M31 (Andromeda). We used both Thuan-Gunn $g$-band ground-based photometry from the literature and Hubble Space Telescope Wide Field Planetary Camera 2 archival data in the F555W and F814W passbands to investigate the pulsation properties of RR Lyrae variable candidates in NGC 147. These datasets represent the two extreme cases often found in RR Lyrae studies with respect to the phase coverage of the observations and the quality of the photometric measurements. Extensive artificial variable star tests for both cases were performed. We conclude that neither dataset is sufficient to confidently determine the pulsation properties of the NGC 147 RR Lyraes. Thus, while we can assert that NGC 147 contains RR Lyrae variables, and therefore a population older than $\\sim$10 Gyr, it is not possible at this time to use the pulsation properties of these RR Lyraes to study other aspects of this old population. Our results provide a good reference for gauging the completeness of RR Lyrae variable detection in future studies. ",
        "introduction": "Local Group dwarf galaxies exhibit diversity in their star formation histories (Da Costa \\& Mould 1998; Grebel 1999). Most of these galaxies show some evidence for intermediate age stars and some have even younger populations along with a significant amount of gas and dust. The broad red giant branch (RGB) morphology shown in their color-magnitude diagrams (CMDs) reflects a wide range of metallicities and/or ages, implying that these galaxies have experienced complex star formation histories. Despite their different and complex star formation histories, RR Lyrae variables are believed to exist in most if not all dwarf galaxies in the Local Group. The existence of RR Lyrae (RRL) populations in a given stellar system indicates that the system is older that $\\sim$10 Gyr. Furthermore, their pulsation properties and absolute magnitudes are correlated with their metallicities (Bono et al. 2003; Bono et al. 2007; Kunder \\& Chaboyer 2009). Therefore, detailed investigations of the physical properties of RRL variable stars are crucial to improve our understanding of the early stages of star formation in Local Group dwarf galaxies.  NGC 147 is one of the dwarf spheroidal (dSph) satellites of the Andromeda galaxy (M31). In terms of its morphology, NGC 147 is a typical dwarf elliptical or spheroidal galaxy. However, it is distinct from other dwarf galaxies in the Local Group, with the possible exception of Tucana (Fraternali et al. 2009), because NGC 147 is dominated by an old stellar population (Hodge 1989). There is a slight indication of intermediate age stars (e.g. extended asymptotic giant branch stars) concentrated in the central regions (Han et al. 1997), but their signature is weaker than the intermediate age stellar populations found in other Local Group dwarf galaxies. Furthermore, NGC 147 appears to show a complete deficit of gas and dust.  The presence of an RR Lyrae population in NGC 147 was reported by Saha et al. (1990, hereafter S90). They used Thuan-Gunn $g$-band ground based photometry obtained from the Hale 5m telescope equipped with the 4-shooter CCD system to identify RR Lyrae variables in this galaxy. They found 32 RR Lyrae candidates and determined the distance modulus of NGC 147 to be $(m-M)_0$= 23.92 based on their mean apparent $g$-band magnitudes. Their periods and amplitudes were estimated by using a prototype of the phase dispersion minimization method developed by Lafler \\& Kinman (1965), the so-called L-K method. Saha et al. (1990) did not present a comparison of the pulsation properties of these RR Lyrae candidates with those in other Local Group galaxies. One reason for this may be that the large photometric errors ( $>$ 0.15 mag) at the level of the horizontal branch (HB) of NGC 147 produced significant errors in the period determinations - especially for the shorter period RR Lyraes (P $<$ 0.4 d). However, the effect of the photometric errors on the period determination was not addressed because the L-K method ignores the photometric errors in the process of identifying the optimal period.  Saha et al's (1990) paper appears to be the only previous study of the RR Lyrae population in NGC 147. We are motivated to revisit the properties of these stars for two reasons. First, we have had good success in using a light curve template-fitting algorithm (Layden \\& Sarajedini 2000) to determine the properties of RR Lyraes such as periods, amplitudes, and mean magnitudes. This method includes the photometric errors in the analysis and has been streamlined and redesigned to be more user-friendly by Mancone \\& Sarajedini (2008). This will allow us to refine the determination of the RRL periods and place better constraints on the total number of such variables in NGC 147. Second, there are archival imaging data from the Hubble Space Telescope (HST) Wide Field Planetary Camera 2 (WFPC2) for NGC 147 that may allow us to update the list of RR Lyraes published by S90. Although the WFPC2 data provide accurate photometry at the level of the HB in NGC 147, they exhibit poor phase coverage with a time baseline of $\\sim$0.4 day. However, it is still useful because it could help to identify the shorter period, lower amplitude RR Lyrae stars that may not have been detected by S90.  The next section describes the observational datasets that we will analyze. Sections 2 and 3 make it clear that neither the S90 data nor the WFPC2 data are  ideal for the purpose of studying the RR Lyraes in NGC 147, but they do complement each other nicely. As discussed in Section 4, it is important to carry out simulations to fully understand the biases and caveats inherent in the results obtained from each of the datasets. The conclusions are presented in Section 5 as well as the case for future work to better characeterize the variable stars in NGC 147 and its M31 dwarf satellite cousin NGC 185. ",
        "conclusions": "  1. The $g$-band photometry from the work of Saha et al. (1990) likely possesses an adequate observational baseline and available epochs. However, our simulations showed that the photometric errors at the level of the horizontal branch significantly hinder the accurate determination of the pulsation periods of the RR Lyrae candidates in NGC 147.  2. Our template light curve fitting technique (FITLC) detected 36 probable RR Lyrae candidates from HST/WFPC2 archival data. However, our simulations reveal that the short observational baseline and small number of observations severely affect the accurate characterization of RR Lyrae periods longer than $\\sim$0.4 days, which are essentially the ab-type RR Lyraes.  3. The $g$-band photometry and the WFPC2 archival data analyzed herein present two extreme cases often found in period finding studies - good phase coverage but with large photometric errors, and high quality photometry with poor phase coverage. Our investigation of these two extreme cases not only provides a good reference for interpreting the pulsation properties of RR Lyrae variables in other similar situations, but also calls attention to a strong need for new high quality time-series observations of NGC 147. Thus, while we can confidently assert that NGC 147 contains RR Lyrae variables, and therefore a population older than $\\sim$10 Gyr, it is not possible at this time to use the pulsation properties of these RR Lyraes to study other aspects of this old population."
    },
    "0911/0911.2712_arXiv.txt": {
        "abstract": "We present the first comprehensive study of short-timescale chromospheric H$\\alpha$ variability in M dwarfs using the individual 15 min spectroscopic exposures for $52,392$ objects from the Sloan Digital Sky Survey.  Our sample contains about $10^3-10^4$ objects per spectral type bin in the range M0--M9, with a total of about $206,000$ spectra and a typical number of 3 exposures per object (ranging up to a maximum of $30$ exposures).  Using this extensive data set we find that about $16\\%$ of the sources exhibit H$\\alpha$ emission in at least one exposure, and of those about $45\\%$ exhibit H$\\alpha$ emission in all of the available exposures.  As in previous studies of H$\\alpha$ {\\it activity} ($L_{\\rm H\\alpha}/L_{\\rm bol}$) we find a rapid increase in the fraction of active objects from M0--M6. However, we find a subsequent decline in later spectral types that we attribute to our use of a spectral type dependent equivalent width threshold.  Similarly, we find saturated activity at a level of $L_{\\rm H\\alpha}/L_{\\rm bol}\\approx 10^{-3.6}$ for spectral types M0--M5, followed by a decline to about $10^{-4.3}$ in the range M7--M9.  Within the sample of objects with H$\\alpha$ emission, only 26\\% are consistent with non-variable emission, independent of spectral type.  The H$\\alpha$ {\\it variability}, quantified in terms of the ratio of maximum to minimum H$\\alpha$ equivalent width ($R_{\\rm EW}$), and the ratio of the standard deviation to the mean ($\\sigma_{\\rm EW}/\\langle {\\rm EW}\\rangle$), exhibits a rapid rise from M0 to M5, followed by a plateau and a possible decline in M9 objects.  In particular, $R_{\\rm EW}$ increases from a median value of about 1.8 for M0--M3 to about 2.5 for M7--M9, and variability with $R_{\\rm EW}\\gtrsim 10$ is only observed in objects later than M5.  For the combined sample we find that the $R_{\\rm EW}$ values follow an exponential distribution with $N(R_{\\rm EW})\\propto {\\rm exp}[-(R_{\\rm EW}-1)/2]$; for M5--M9 objects the characteristic scale is $R_{\\rm EW}-1\\approx 2.7$, indicative of stronger variability.  In addition, we find that objects with persistent H$\\alpha$ emission exhibit smaller values of $R_{\\rm EW}$ than those with intermittent H$\\alpha$ emission.  Based on these results we conclude that H$\\alpha$ variability in M dwarfs on timescales of 15 min to 1 hr increases with later spectral type, and that the variability is larger for intermittent sources.  Future studies using this large sample will address the variability timescales, the variability of other chromospheric emission lines (e.g., H$\\beta$, \\ion{Ca}{2} H\\&K), and the origin of the highest amplitude events. ",
        "introduction": "The study of magnetic activity in fully convective low mass stars and brown dwarfs (spectral types late-M, L, and T) has progressed in recent years from the question of whether these objects produce stable fields to a quantitative investigation of how the fields are produced and dissipated.  Unlike the solar-type $\\alpha\\Omega$ dynamo, which operates in the transition region between the radiative and convective zones (the tachocline; \\citealt{par55}), objects later than spectral type M3 can only support a convective dynamo.  Numerical simulations of such a mechanism are still at an early stage, but they suggest that large-scale axisymmetric fields can indeed be generated in fully convective objects, at least for conditions that roughly correspond to mid-M dwarfs ($M\\sim 0.3$ M$_\\odot$; \\citealt{bro08}).  Thus, observational constraints on the scale, geometry, and dissipation of the fields are essential.  Several observational techniques are now being used to address these questions, including Zeeman measurements in Stokes $I$ and $V$, which probe the strength of the integrated surface fields and their large-scale topology, respectively \\citep{rb07,dmp+08,mdp+08}, and activity indicators such as radio, X-ray, and H$\\alpha$ emission, which trace the dissipation of the field and hence its strength and geometry (e.g., \\citealt{whw+04,ber06,bbf+09}).  The activity indicators also provide insight into magnetic heating, and their temporal variability can potentially trace the field properties on small scales that are inaccessible to the Zeeman measurements.  The current results from Zeeman measurements point to a transition from mainly toroidal and non-axisymmetric fields in M0--M3 dwarfs to predominantly poloidal axisymmetric fields in mid-M dwarfs \\citep{dmp+08,mdp+08}, with field strength of $\\sim 0.1-3$ kG \\citep{rb07,dmp+08,mdp+08}.  Studies of activity indicators point to a rapid decline in X-ray and H$\\alpha$ activity (i.e., $L_{X,{\\rm H\\alpha}}/L_{\\rm bol}$) at about spectral type M7, and uniform radio luminosity (i.e., increasing $L_{\\rm rad}/L_{\\rm bol}$) at least to spectral type L3 (\\citealt{bbf+09} and references therein).  The disparate trends may be due to a decoupling of the magnetic fields from the increasingly neutral atmospheres, or to a shift in the magnetic field configuration.  While X-ray and radio observations provide powerful insight into the nature of the magnetic fields, H$\\alpha$ chromospheric emission, which traces gas at temperatures of $\\sim 10^4$ K, is more easily accessible for large samples as a by-product of standard optical spectroscopic observations.  In recent years, large spectroscopic samples of M dwarfs have become available through dedicated studies and large-scale surveys such as the Sloan Digital Sky Survey (SDSS).  These extensive samples have led to several important results concerning chromospheric activity.  First, the fraction of objects that exhibit H$\\alpha$ emission increases rapidly from $\\sim 5\\%$ in the K5--M3 dwarfs to a peak of $\\sim 80-100\\%$ around spectral type M7, followed by a subsequent decline to a few percent in the L dwarfs \\citep{gmr+00,whw+04,whb+08}.  Second, while the level of activity increases with both rotation and youth in F--K stars, it reaches a saturated value of $L_{\\rm H\\alpha}/L_{\\rm bol}\\approx 10^{-3.6}$ in M0--M6 dwarfs, followed by a rapid decline to $L_{\\rm H\\alpha}/L_{\\rm bol}\\approx 10^{-5}$ by spectral type L0 \\citep{hgr96,gmr+00,whw+04}, and a breakdown of the rotation-activity relation \\citep{mb02}. Finally, a small fraction ($\\lesssim 5\\%$) of late-M and L dwarfs have been serendipitously observed to exhibit H$\\alpha$ flares that reach the saturated emission levels found in the early-M dwarfs \\citep{lkc+03}.  To uniformly address the latter point --- H$\\alpha$ variability --- we recently carried out spectroscopic monitoring observations of about 40 M4--M8 dwarfs with known H$\\alpha$ emission \\citep{lbk09}.  With observations of about 1 hr per source and a time resolution of $5-10$ min, we found that about $80\\%$ of these objects exhibit H$\\alpha$ variability on a wide range of timescales, ranging in amplitude from tens of percent up to a factor of about 5.  Indeed, the timescale distribution for variability ``events'' is nearly flat from 10 min to 1 hr, with fewer events on timescales below 10 min.  The variability amplitudes follow an exponential distribution with a characteristic scale of ${\\rm Max(EW)/Min(EW)}-1\\approx 0.7$.  Finally, we found tentative evidence for increased variability with later spectral type.  Here, we extend our study of H$\\alpha$ variability by three orders of magnitude using SDSS time-resolved spectroscopic data for a sample of 52,392 M0--M9 dwarfs (Knapp et al. 2009 in preparation; hereafter Paper I).  For the first time we take advantage of the individual 15 min exposures, with a typical number of 3 exposures per source.  Using these data we can thus probe H$\\alpha$ variability on timescale of $15$ min to $\\sim 1$ hr.  In this paper we focus on the H$\\alpha$ variability amplitudes and their relation to the H$\\alpha$ activity level ($L_{\\rm H\\alpha}/L_{\\rm bol}$) and spectral type.  We also re-investigate the relation between H$\\alpha$ activity and spectral type using the individual 15 min spectra, rather than the SDSS pipeline-combined spectra, which were used in previous work \\citep{whw+04,whb+08}.  In \\S\\ref{sec:obs} we define our sample selection and describe the analysis method for measuring H$\\alpha$ equivalent widths.  The results for H$\\alpha$ variability and activity are described in \\S\\ref{sec:results}.  Future papers will focus on the variability timescales, the highest amplitude events, the variability of other chromospheric emission lines (e.g., higher-order Balmer lines, \\ion{Ca}{2} H\\&K, \\ion{He}{1} lines), and the relation between chromospheric variability and various source properties, such as age. ",
        "conclusions": "\\label{sec:conc}  We study the variability of the H$\\alpha$ chromospheric emission line in M dwarfs using the individual 15 min spectroscopic exposures from SDSS for an unprecedentedly large sample of 52,392 stars.  The typical timescale probed by these data is 15 min to 1 hr.  About $16\\%$ of the sources exhibit H$\\alpha$ emission in at least one exposure, and of those about $45\\%$ have H$\\alpha$ emission in all of the available exposures.  Only $26\\%$ of all the objects with H$\\alpha$ emission are consistent with steady emission, spread relatively uniformly from M0 to M9.  The detection fraction as a function of spectral type exhibits the known trend of a sharp increase from a few percent in M0--M3 to tens of percent in later objects.  However, unlike the results of previous studies based on SDSS pipeline-combined spectra and a fixed equivalent width threshold of 1 \\AA\\ \\citep{whw+04,whb+08}, we find a peak detection fraction of $\\sim 40\\%$, with a steady decline beyond M6.  We attribute this trend to our use of a spectral type dependent detection threshold of $\\sim 3-6$ \\AA\\ ($3\\sigma$) in M5--M9 objects.  The H$\\alpha$ variability, quantified in terms of $R_{\\rm EW}$, exhibits a substantial increase with later spectral type.  For M dwarfs as a whole, the distribution of $R_{\\rm EW}$ values appears to follow an exponential with $N(R_{\\rm EW})\\propto{\\rm exp}[-(R_{\\rm EW}-1)/2]$; the characteristic scale is about 2.7 for M5--M9.  We also find that objects with partial detections exhibit a wider distribution of $R_{\\rm EW}$ values than those with persistent H$\\alpha$ emission. These results indicate that H$\\alpha$ variability increases with later spectral type, even as H$\\alpha$ activity declines, and that M dwarfs with intermittent H$\\alpha$ activity are more variable than those with persistent H$\\alpha$ emission.  In the context of chromospheric heating, these trends suggest that the magnetic energy input typically varies by a factor of about 2 on timescales of $\\sim 15$ min to $\\sim 1$ hr.  Moreover, the increased variability as a function of later spectral type, with a particularly large increase beyond spectral types M3--M4, hints at an overall shift in the magnetic field, leading to less uniform field dissipation. Taken at face value, this conclusion suggests that the relative contribution from small scales, which are likely to dissipate on short timescales, increases for later M dwarfs.  This possible trend is in conflict with preliminary results from Zeeman measurements of a few objects in the range M0--M5, which point to an increase in the relative contribution from large-scale fields \\citep{dmp+08,mdp+08,rb09}.  Alternatively, the increase in H$\\alpha$ variability, particularly for sources with intermittent emission, may be due to an increase in the stochastic dissipation of the field in the presence of increasingly neutral atmospheres.  In this scenario, heating of the bulk chromospheric plasma is suppressed by the decoupling of the field from the atmosphere (leading to a decline in $L_{\\rm H\\alpha}/L_{\\rm bol}$), but small-scale stochastic heating still takes place (leading to an increase in the H$\\alpha$ variability).  The role of these two scenarios may become clearer as samples of late-M dwarfs with Zeeman measurements become available."
    },
    "0911/0911.0521_arXiv.txt": {
        "abstract": "We study, within an effective field theory framework, $O(E^{2}/\\Mpl^{2})$ Planck-scale suppressed Lorentz invariance violation (LV) effects in the neutrino sector, whose size we parameterize by a dimensionless parameter $\\eta_{\\nu}$. We find deviations from predictions of Lorentz invariant physics in the cosmogenic neutrino spectrum. For positive $O(1)$ coefficients no neutrino will survive above $10^{19}~\\eV$. The existence of this cutoff generates a bump in the neutrino spectrum at energies of $10^{17}~\\eV$. Although at present no constraint can be cast, as current experiments do not have enough sensitivity to detect ultra-high-energy neutrinos, we show that experiments in construction or being planned have the potential to cast limits as strong as $\\eta_{\\nu} \\lesssim 10^{-4}$ on the neutrino LV parameter, depending on how LV is distributed among neutrino mass states. Constraints on $\\eta_{\\nu} < 0$ can in principle be obtained with this strategy, but they require a more detailed modeling of how LV affects the neutrino sector. ",
        "introduction": "Over the last fifteen years there has been consistent theoretical interest in possible small deviations from the exact local Lorentz Invariance (LI) of general relativity as well as a flourishing of observational tests. The theoretical interest is driven primarily by ideas in the Quantum Gravity (QG) community that Lorentz invariance may not be an exact local symmetry of the vacuum. The possibility of outright Lorentz symmetry violation (LV) or a different realization of the symmetry than in special relativity has arisen in string theory~\\cite{KS89, Ellis:2008gg}, Loop QG~\\cite{LoopQG,Rovelli:2002vp,Alfaro:2002ya}, non-commutative geometry~\\cite{Carroll:2001ws, Lukierski:1993wx, AmelinoCamelia:1999pm, Chaichian:2004za}, space-time foam~\\cite{AmelinoCamelia:1997gz}, some brane-world backgrounds~\\cite{Burgess:2002tb}, and condensed matter analogues of ``emergent gravity''~\\cite{Analogues}. As well, allowing the fundamental theory of gravity to be non-relativistic in the ultraviolet can make gravity renormalizable while avoiding some of the other pathologies that plague renormalizable gravitational actions with higher derivative terms~\\cite{Horava:2009uw}.  Constructing useful tests for the various active models and ideas is therefore vital.  Since there are so many theoretical models around, a good approach is to work within a calculable framework where all possible LV terms are parameterized.  Each theory then picks out a certain combination of terms which can be constrained. In this vein, a standard method is to simply analyze a Lagrangian containing the standard model fields and all LV operators of interest that can be constructed by coupling the standard model fields to new LV tensor fields that have non-zero vacuum expectation values\\footnote{There are other approaches to either violate or modify Lorentz invariance, that do not necessarily yield a low energy EFT (see \\cite{AmelinoCamelia:2008qg} and refs therein). However, these models do not easily lend themselves to particle physics constraints as the dynamics of particles is less well understood and hence we do not consider them here. In particular, we remark here that ideas of deformation, rather than breaking, of the Lorentz symmetry (see, e.g., \\cite{AmelinoCamelia:2000mn}) do not have an ordinary-EFT formulation, hence they cannot be tested with the arguments presented in this work.}.  All renormalizable LV operators that can be added to the standard model in this way are known as the Standard Model Extension (SME)~\\cite{Colladay:1998fq}. These operators all have mass dimension three or four and can be further classified by their behavior under CPT. Higher mass dimension operators can be systematically explored as well, which is useful in case the naive EFT hierarchy breaks down due to other new physics (e.g.~SUSY) or quantum gravity introducing a custodial mechanism for the renormalizable operators in the infrared.  The CPT odd dimension five kinetic terms for QED coupled to a non-zero vector field were written down in~\\cite{Myers:2003fd} while the full set of dimension five operators with a vector were analyzed in~\\cite{Bolokhov:2007yc}. The dimension five and six CPT even kinetic terms for QED for particles coupled to a non-zero background vector, which we are primarily interested in here, were partially analyzed in~\\cite{Mattingly:2008pw}.  The full set of dimension five and six operators for QED has recently been introduced in~\\cite{Kostelecky:2009zp}. It is notable that SUSY forbids renormalizable operators for matter coupled to non-zero vectors~\\cite{GrootNibbelink:2004za} but permits certain nonrenormalizable operators at mass dimension five and six.  Many of the parameterized LV operators have been very tightly constrained via direct observations (see~\\cite{Mattingly:2005re,Jacobson:2005bg, AmelinoCamelia:2008qg, Liberati:2009pf} for reviews). In particular, the dimension five and six CPT even operators for LV with a vector field have been recently directly\\footnote{Notice that all these operators can be indirectly constrained by EFT arguments~\\cite{Collins:2004bp}, as higher dimension LV operators induce large renormalizable ones if we assume no other relevant physics enters between the TeV and $\\Mpl$ energies.  SUSY, however, is an example of new relevant physics that can change this.} constrained in the hadronic sector by exploiting ultra-high energy cosmic-ray observations \\cite{Maccione:2009ju} performed by the Pierre Auger Observatory (PAO). Indeed, the construction and successful operation of this instrument has brought UHECRs to the interest of a wide community of scientists and it is expected to allow, in the near future, the assessment of several problems of UHECR physics and also to test fundamental physics (in particular Lorentz invariance in the QED sector) with unprecedented precision \\cite{Galaverni:2007tq,Maccione:2008iw,Galaverni:2008yj}.  The UHECR constraints \\cite{Kifune:1999ex,Aloisio:2000cm,AmelinoCamelia:2000zs,Stecker:2004xm,GonzalezMestres:2009di,Scully:2008jp,Maccione:2009ju,Stecker:2009hj} rely on the behavior of particle reaction thresholds with LV, which are one of the best methods in the EFT approach to constrain nonrenormalizable LV operators.  Many LV operators give modified dispersion relations for free particles, where the energy as a function of momentum deviates slightly from the special relativistic form. For threshold reactions what matters is not the size of the LV correction to the energy compared to the absolute energy of the particle, but instead the size of the LV correction to the mass of the particles in the reaction.  Hence the LV terms usually become important when their size becomes comparable to the mass of the heaviest particle.  If the LV term scales with energy as $E^n$, then this critical energy is $E_{cr} \\sim \\left(m^{2}\\Mpl^{n-2}\\right)^{1/n}$ \\cite{Jacobson:2002hd}. According to this reasoning, the larger the particle mass the higher is the energy at which threshold LV effects come into play. This is why $\\gtrsim \\TeV$ electrons and positrons, but not protons, can be used to constrain $n=3$ LV \\cite{Maccione:2007yc}, and why UHE protons are needed to obtain constraints on hadronic LV with $n=4$ scaling (which corresponds to CPT even mass dimension five and six operators).  From this point of view, neutrinos, with their tiny mass of order $m_{\\nu} \\simeq 0.01~\\eV$ \\cite{pdg}, are in principle the most suited particles to provide strong constraints on LV, at least for reactions involving {\\em only} neutrinos.  One such reaction is that of neutrino oscillation.  Indeed, for a decade neutrino oscillations have proven to be excellent tests of the SME and other LV models~\\cite{Coleman:1998ti,Kostelecky:2009zp,GonzalezGarcia:2005xw,Diaz:2009qk,Yang:2009ge} as when the LV corrections are near the neutrino mass, the oscillation pattern can change as a function of energy, direction and mass. ICECUBE may even be able to probe dimension six operators with time of flight techniques with TeV neutrinos from distant Gamma Ray Bursts~\\cite{GonzalezGarcia:2006na}. UHECR experiments have also the capability of placing constraints on SME parameters by exploiting neutrino oscillations \\cite{Bhattacharya:2009tx}\\footnote{We notice that in the same work \\cite{Bhattacharya:2009tx} the process of neutrino decay is discussed, but in a different context than what we consider here.}. One can construct complementary and in some measures even more sensitive neutrino tests of higher dimension operators by leveraging observables which deviate more strongly from their special relativistic values as the neutrino travel distance increases. One such observable is found to be related to UHE neutrino spectrum observations.  Despite the threshold being low for LV effects to kick in, neutrinos with ultra-high energy are necessary to achieve a signal, as they interact so weakly that the phase space for a LV reaction must be huge to generate an appreciable rate. This requirement implies that, since the LV terms and hence the phase space grow with energy, very large energies are needed. Indeed, for renormalizable operators, where the phase space does not grow quickly enough, reactions of the type we consider here never achieve the necessary rate given current bounds on their coefficients, even for UHE neutrinos over cosmological distances~\\cite{Coleman:1997xq,Coleman:1998ti}. However, for the CPT even non-renormalizable operators the situation is different, and we find that for energies of order $10^{17}~\\eV$ there can be a significant modification of the spectrum. Next generation neutrino detectors such as ANITA \\cite{Gorham:2008yk} and SuperEUSO \\cite{Petrolini:2009cg,Santangelo:2009et} are sensitive to neutrinos of energies $>10^{19}~\\eV$. Further experiments, like the planned ARIANNA \\cite{Barwick:2006tg,Barwick:2009zz} and IceRay \\cite{Allison:2009rz}, will cover the range $10^{17}\\div 10^{20}~\\eV$.  In the present work we study how limits on the absolute scale of non-renormalizable CPT even LV neutrino parameters can be obtained from UHE neutrino observations.  In particular we will determine the energy scale where the neutrino spectrum might begin to deviate as a function of the size of the LV coefficients.  With three neutrino species as well as assorted light leptons possibly involved, the spectrum as a function of various LV parameters can only be computed by a detailed parameterized numerical search.  At this stage, where LV in this sector is only speculative, we feel that such a search is unwarranted.  However, we will show the ``best case'' scenario for a LV neutrino signal with dimension six operators as well as a scenario where a signal in the UHE spectrum is much more subtle even though LV is still relatively strong.  This paper is structured as follows. In section~\\ref{sec:theo} we describe the theoretical LV framework in which we will derive the LV effects on the UHE neutrino spectrum. In sections~\\ref{sec:nuLV} and \\ref{sec:nusplitting} we will give general information on the standard understanding of UHE neutrino generation by UHECRs, and we will present and detail the main LV reactions possibly affecting their spectrum. Section~\\ref{sec:bestcase} is devoted to present results in our test cases. In Section \\ref{sec:otherprocesses} we discuss the possible role of other processes than neutrino splitting.  In section~\\ref{sec:conclusions} we will report our final remarks and conclusions. ",
        "conclusions": "\\label{sec:conclusions}  In this work we have investigated possible signals of higher dimension LV CPT even operators in the UHE neutrino spectrum, in particular the effect of the neutrino splitting on the UHE neutrino spectrum. This process provides a clean test as it does not involve other LV operators apart from the neutrinos' ones.  In addition, since the dominant neutrino field is left-handed, there are no complications due to different LV for different chirality of fermion, as there are in the UHECR case, and one ends up with one LV coefficient for each neutrino mass eigenstate. In the flavor blind scenario, where every mass eigenstate undergoes splitting approximately at the same energy, there is both a precocious fall off of the neutrino flux at UHE as well as a significant excess in the UHE neutrino flux at energies as low as $10^{17.5}$ eV. Noticeably, this kind of energies are well within reach of current and future UHECR experiments and about order of magnitude below those so far used for LV tests with hadronic and electromagnetic UHECRs.  According to our study, existing or planned UHE neutrino experiments have the potential to probe LV in the neutrino sector for coefficients $\\eta \\gtrsim 10^{-4}$. However, we have discovered a serious difficulty with deriving constraints in the absence of a positive LV signal event. The distribution in mass eigenstates is roughly equal for UHE neutrinos in realistic mixing scenarios. Although it might seem somewhat unnatural that LV has different effects on different mass states of the same particle field, if this is the case then it is possible that LV can exist/be strong for one mass eigenstate yet be almost invisible for UHE neutrino detectors. In particular, the observation of a neutrino flux up to some maximal energy does not imply a firm conclusion on LV, at least with present accuracy. However, let us note that if a neutrino is ever detected at energy $E_{\\nu}$, we can rule out a flavor blind $\\eta_{\\nu} \\gtrsim (E_{\\nu}/10^{18.8}~\\eV)^{-13/4}$ according to Eq.~(\\ref{eq:constraint_naive}). On the other hand, the bump feature can lead to constraints on $\\eta_{\\nu} > 0$, as this bump should be observable in any such scenario of LV in the neutrino sector. In this case, to obtain a $O(1)$ constraint on LV requires at least a 50\\% accuracy in the determination of the neutrino spectrum in the energy range $10^{17}\\div10^{19}~\\eV$, which can be achieved by future experiments.    \\ack We thank J.~Kelley and B.~McElrath for useful discussions. This work was supported by the Deutsche Forschungsgemeinschaft through the collaborative research centre SFB 676 ``Particles, Strings and the Early Universe: The Structure of Matter and Space-Time''. LM acknowledges support from the State of Hamburg, through the Collaborative Research program ``Connecting Particles with the Cosmos'' within the framework of the LandesExzellenzInitiative (LEXI)."
    },
    "0911/0911.0698_arXiv.txt": {
        "abstract": "Following an initial explosion that might be launched either by magnetic interactions or neutrinos, a rotating magnetar radiating according to the classic dipole formula could power a very luminous supernova. While some $^{56}$Ni might be produced in the initial explosion, the peak of the light curve in a Type I supernova would not be directly related to its mass. In fact, the peak luminosity would be most sensitive to the dipole field strength of the magnetar. The tail of the light curve could resemble radioactive decay for some time but, assuming complete trapping of the pulsar emission, would eventually be brighter. Depending on the initial explosion energy, both high and moderate velocities could accompany a very luminous light curve. ",
        "introduction": "\\lSect{intro}  The role of rotation in powering the explosion of supernovae has long been debated \\citep[e.g.,][]{Hoy46,Leb70,Ost71,Aki03}.  Most recent work has focused on the possibility that a rotating neutron star could, by way of a magnetic interaction, be the energy source for {\\sl exploding} a massive star. While that issue is far from resolved, and recent headway has been made in exploding these same stars using neutrino transport, it is certain that a large fraction of supernova explosions produce rotating neutron stars and that those neutron stars frequently have large magnetic fields. Magnetars are a class of neutron stars with field strengths 10$^{14}$ to 10$^{15}$ Gauss and more. They may constitute 10\\% of the neutron star birthrate \\citep{Kou98}. It is likely that fast rotation in the collapsing iron core is responsible for creating the large magnetic field \\citep{Dun92}, so it is reasonable to expect that the birth of rapidly rotating, highly magnetic neutron stars is commonplace.  It is also generally assumed that these rapidly rotating neutron stars are magnetically braked by dipole emission early on, accounting for the slow periods observed in anomalous x-ray pulsars and soft gamma-ray repeaters \\citep{Dun92,Kou98}. Since the initial rotational energy of such stars at birth must have been large and 1000 years later is small, where did the difference go? Might it have been emitted in some observable form?  As a rough approximation, assume that the neutron star radiates its rotational energy away at a rate given by the traditional dipole formula for pulsars.  For a typical moment of inertia of 10$^{45}$ g cm$^2$, the rotational energy of a neutron star with period, P$_{\\rm ms}$, in milliseconds is \\begin{equation} \\begin{split} E &= \\frac{1}{2} I \\omega^2 \\\\ &\\approx 2 \\times 10^{52} {\\rm P}_{\\rm ms}^{-2} \\ {\\rm erg}. \\end{split} \\lEq{pulsener} \\end{equation} The approximate energy loss for dipole radiation is given by the Larmor formula \\citep[e.g.,][]{Lan80}, \\begin{equation} \\begin{split} \\frac{d E}{d t} &= \\frac{2}{3 c^3} \\left(B R^3 \\ {\\rm Sin} \\, \\alpha \\right)^2 \\left(\\frac{2 \\pi}{{\\rm P}}\\right)^4 \\\\ &\\approx 10^{49} B_{15}^2 {\\rm P}_{\\rm ms}^{-4} \\ {\\rm erg \\ s^{-1}}. \\end{split} \\lEq{dipole} \\end{equation} Here $B_{15}$ is the surface dipole field in 10$^{15}$ Gauss, $R \\approx 10^6$ cm is the neutron star radius, and $\\alpha$ is the inclination angle between the magnetic and rotational axes, taken arbitrarily to be 30 degrees.  At face value, these simple formulae suggest large luminosities for magnetars during the first few weeks of their lives. For example, the magnetar in SGR 1627-41 is estimated to have a current spin-down luminosity of $\\sim 4 \\times 10^{34}$ erg s$^{-1}$ at an age of 2.2 ky implying a magnetic field $B \\, {\\rm Sin} \\, \\alpha \\ \\sim 2 \\times 10^{14}$ Gauss \\citep{Esp09}. Extrapolated back to when the magnetar was 10 days old, \\eq{dipole} implies a luminosity of $2 \\times 10^{44}$ erg s$^{-1}$.  Provided the initial rotation rate was rapid compared with its value at 10 days, this result is independent of the initial rotation rate.  Here we explore the observational characteristics of Type Ib/c supernovae in which pulsars with magnetar-like characteristics have been embedded. The brightest supernovae actually come from neutron stars with fields that, for a magnetar, are relatively modest, $\\sim 10^{14}$ Gauss. Larger fields imply that most of the magnetar's rotational energy is dissipated early on, adding to the kinetic energy and mixing, but not appreciably affecting the luminosity.  These luminous magnetar-powered supernovae might easily be confused with other forms of ``hypernovae'' where the luminosity has a radioactive origin. Alternatively, the lack of such emission constrains the existence of any rapidly rotating magnetar. This could be an important constraint in the context of gamma-ray bursts where the the possibility of a magnetar power source is currently debated. ",
        "conclusions": "Pulsars with magnetar-like magnetic fields can contribute appreciably to the light curves of Type Ib and Ic supernovae. In fact, it would be surprising if they never did. For moderate field strengths, the light curve is very luminous and long lasting and might be confused with those of pair-instability supernovae or circumstellar interaction. Indeed, \\citet{Mae07} have suggested a pulsar as the possible source powering the second (principal) maximum of the light curve of SN 2005bf.  If magnetars are the central engine that powers the long-soft class of gamma-ray bursts \\citep[e.g.,][]{Woo06}, then it is reasonable to expect that they may contribute to the light curves of the supernovae that accompany them. On the one hand, this might facilitate the magnetar paradigm because the production of the necessary large amount of $^{56}$Ni has proven problematic \\citep{Buc09}. Having an alternate explanation that involves magnetar energy input would solve this problem. On the other hand, if a magnetar contribution to the light curve can be ruled out based on the spectrum and late time light curve, one must wonder how the magnetar is so effectively concealed. Is the rotational energy extracted with such high efficiency in the first few seconds that the magnetar forever afterwards rotates slowly, or do rotating magnetars not emit according to the popular dipole formula during the first year?  Perhaps they do not. The absence of a pulsar contribution to the luminosity of typical supernovae is easy to understand. A pulsar born with a 14 ms period (10$^{50}$ erg), would only have a pulsar luminosity of $6 \\times 10^{39}$ erg s during the time when it is bright. This is small compared with the 10$^{42}$ - 10$^{43}$ erg s$^{-1}$ resulting from shock energy released by recombination (Type II supernova) or radioactivity (Type I supernova). But the extremely faint optical luminosity of SN 1987A, $< 8 \\times 10^{33}$ erg s$^{-1}$ seventeen years after its birth \\citep{Gra05} is very difficult to reconcile with any model with a young active pulsar. While dust extinction could be considerable, the lack of a point source in SN 1987A remains a mystery."
    },
    "0911/0911.2181_arXiv.txt": {
        "abstract": "{ Neutron-induced reaction rates, including fission and neutron capture, are calculated in the temperature range $10^{8} \\leq T({\\rm K})\\leq 10^{10}$ within the framework of the statistical model for targets with  the  atomic number   $84\\leq Z\\leq 118$ (from Po to Uuo) from the neutron to the proton drip-line.  Four sets of rates have been calculated, utilizing - where possible - consistent nuclear data for neutron separation energies and fission barriers from Thomas-Fermi (TF), Extended Thomas-Fermi plus Strutinsky Integral (ETFSI), Finite-Range Droplet Model (FRDM) and Hartree-Fock-Bogolyubov (HFB) predictions.  Tables of calculated values as well as analytic seven parameter fits in the standard REACLIB format are supplied$^0$. We also discuss the sensitivity of the rates to the input, aiming at a better understanding of the variations introduced by the nuclear input. } ",
        "introduction": "\\label{sec:intro} Investigations of nucleosynthesis processes make use of reaction networks including thousands of nuclei and tens of thousands of reactions. Most of these reactions occur far from stability and thus cannot yet be directly studied in the laboratory. In addition most of the nuclear properties including reaction rates,   which are  also required for the calculation of cross sections and astrophysical reaction rates, are not experimentally known either. Therefore, predictions based on theoretical models are necessary. While close to stability partial experimental information is available, relying fully on theoretical information leads to relatively large variations  in computed cross sections far from stability. This is especially true for the region of fissionable nuclei, which is the focus of the present investigation.  In the past, a series of efforts were applied to calculate   neutron-capture rates for r-process nucleosynthesis and other astrophysical applications    \\citep[e.g.,][and references therein]{A72,HWF76,WFH78,sar82,thi87,cowan91,rafkt00,most05,gori08}.     Fission has often been neglected in astrophysical calculations. In early applications to astrophysical  nucleosynthesis, usually only one mode was considered, beta-delayed fission \\citep{TMK83} or a phenomenological model of spontaneous fission   \\citep{gori99,frei99,cowan99}. However, it was shown recently that neutron-induced fission is more important  than beta-delayed fission  in r-process nucleosynthesis \\citep{pan2002,panfkt04,fiss07}. Thus, the need to provide a compilation of neutron-induced fission rates is obvious.   Initial investigations have been undertaken by \\cite{pan05} and \\cite{talys09}. Here we present extended calculations of   neutron-induced fission rates for different predictions of masses and fission barriers. The present work also completes existing nuclear neutron-capture  rate sets  by extending the works of \\cite{rafkt00} and \\cite{pan05} to the region  $84\\leq Z\\leq 118$  in order to provide the necessary input for nucleosynthesis studies under high neutron densities. As in \\citet{pan05}, the statistical model approach of Wolfenstein-Hauser-Feshbach \\citep{Wol51,HF52} for compound nuclear reactions was used, but employing more recent and complete data and predictions for masses, spins, and fission barriers. \\footnotetext{Tables 3-18 with these data are only available in electronic form at the CDS via anonymous ftp to cdsarc.u-strasbg.fr (130.79.128.5) or via http://cdsweb.u-strasbg.fr/cgi-bin/qcat?J/A+A/}    Nuclear mass and fission barrier predictions have a strong model dependence, and none of the existing models can reproduce all experimentally known data. Moreover, the fission process itself is complicated, and extended calculations for neutron-induced fission across the nuclear chart  have to be done carefully. Here, we aim to provide rates for studying the endpoint of the r-process and the possible production of super-heavy elements. By comparing rates obtained with different  choices of mass and fission barrier predictions we attempt to give a measure of the involved variations. Astrophysical models, providing the nucleosynthesis conditions, bear large  variations in themselves. This is especially true for the r-process, for which the astrophysical site is still unknown  despite decades of study. For a realistic and exhaustive exploration of synthesis conditions, simulations do not only have to vary astrophysical parameters, but also have to include a variation range of involved reaction rates   given by different mass and fission barrier models.       Our paper is structured as follows. In Sect. \\ref{sec:statmod} we  briefly describe the statistical model used in the calculations as well as the nuclear input data and give a comparison of cross sections or rates for  a number of experimentally known nuclei with existing experimental information and other theoretical models. These methods are then applied to supplement the rate sets of \\citet{rafkt00} of $(n,\\gamma)$-rates for chemical elements with $Z>83$ and predict neutron-induced fission cross sections and rates (where available in comparison to experiments). Section \\ref{sec:rates} presents these results and shows the sensitivity with respect to mass models and fission barriers employed. Rate fits for utilization in astrophysical calculations are discussed in Sect. \\ref{sec:rates_2}. In Sect. \\ref{sec:yields} we give a brief discussion and some examples of the mass distribution of fission fragments, which will be provided in an extended way in a forthcoming paper. The final Sect. \\ref{sec:conclusion} contains conclusions and a summary. The explanation of the tables and their structure are given in Appendix {\\bf A}. The complete tables of reaction rates and their fits are found   at CDS in electronic form. \\begin{figure*} \\begin{center} \\hspace{3.5mm}\\includegraphics*[width=.32\\textwidth]{11967fg01a.eps} \\hspace{0mm}\\includegraphics*[width=.32\\textwidth]{11967fg01b.eps} \\hspace{0mm} \\includegraphics*[width=.325\\textwidth]{11967fg01c.eps} \\hspace*{2mm}\\includegraphics*[width=.335\\textwidth]{11967fg01d.eps} \\hspace{2mm}\\includegraphics*[width=.305\\textwidth]{11967fg01e.eps} \\hspace{3mm} \\includegraphics*[width=.305\\textwidth]{11967fg01f.eps} \\hspace*{0mm} \\includegraphics*[width=.345\\textwidth]{11967fg01g.eps} \\hspace{1mm} \\includegraphics*[width=.305\\textwidth]{11967fg01h.eps} \\includegraphics*[width=.325\\textwidth]{11967fg01i.eps} \\end{center} \\caption{  Present predictions  of energy-dependent $(n,f)$ cross sections $\\sigma_{nf}(E)$ for some target nuclei  of U, Np and Pu calculated in the framework of different mass and fission barrier predictions (ETFSI, TF, HFB-14) and experimental data, marked $B_{\\rm f}^{\\rm exp}$ as well. Experimentally measured cross-sections  were used after JENDL-3.3 \\citep{jendl05}, averaged by the code JANIS \\citet{janis07}, displayed by a black line. All the predictions are given for a ground-state population.   Our previous results \\citep{pan05} are shown as well. } \\label{ex_jeja} \\end{figure*} ",
        "conclusions": "\\label{sec:conclusion}   We provide predictions of neutron-induced fission rates and $(n,\\gamma)$-rates for a wide range of astrophysical temperatures ($10^{8} \\leq T({\\rm K})\\leq 10^{10}$) and targets (proton- to neutron- drip-line for  $84 \\leq Z \\leq 118$, i.e. from Po to Uuo ) in the framework of the Wolfenstein-Hauser-Feshbach model, making use of a variety of different mass and fission barrier predictions \\citep{mysw99,mamdo01,etfsi,frdm,homo80,talys09}. The astrophysical (stellar) reaction rates were fitted   as in previous works \\citep{thi87,rafkt00}  in the common REACLIB seven parameter form, and these parameters are also tabulated. This provides the basis for r-process nucleosynthesis calculations where the abundance predictions for the highest mass numbers as well as the effect of fission cycling are strongly dependent on the interplay of neutron capture and fission.   In order to give an impression of the reliability of the results, we compared them  with experiment and with available independent predictions before exploring the currently unreachable regions of the nuclear chart with a variety of theoretical predictions for nuclear masses and fission barriers (FRDM, ETFSI, TF, HFB). An extended comparison of neutron-induced fission rates with experiment and with available independent predictions was done. The  dependence of rates  on nuclear input data, most of all fission barriers, is   high. Astrophysical nucleosynthesis yield predictions, especially in the transuranium region, should  take into account these large differences in order to explore the variations  involved. For this reason   extended tables for neutron-induced fission rates as well neutron capture rates are presented for different mass and fission barrier predictions in  fitted form for nucleosynthesis calculations. Their structure is given in the Appendix {\\bf A} (note that the full rate and fit tables are available at the CDS). Given that  fission predictions far from stability have not been tested yet, and even close to stability none of the existing models has yet been proven to be superior (see Fig. \\ref{U_Pu_sigv_nf}).   Nucleosynthesis calculations should probably continue to use a variety of these models. A further requirement for   nucleosynthesis modeling in the region of fissioning nuclei is the  knowledge  of the   mass distribution of fission products. This work is in progress (see Sect. \\ref{sec:yields})."
    },
    "0911/0911.0968_arXiv.txt": {
        "abstract": "{} { Our primary goal is to search for planets around intermediate mass stars. We are also interested in studying the nature of radial velocity (RV) variations of K giant stars. } { We selected about 55 early K giant (K0 - K4) stars brighter than fifth magnitude that were observed  using BOES, a high resolution spectrograph attached to the 1.8 m telescope at BOAO (Bohyunsan Optical Astronomy Observatory). BOES is equipped with $I_2$ absorption cell for high precision RV measurements. } { We detected a periodic radial velocity variations in the K0 III star \\gam1leo with a period of P = 429 days. An orbital fit of the observed RVs yields a period of P = 429 days, a semi-amplitude of K = 208 \\mps, and an eccentricity of $e = 0.14.$ To investigate the nature of the RV variations, we analyzed the photometric, Ca II $\\lambda$ 8662 equivalent width, and line-bisector variations of \\gam1leo. We conclude  that the detected RV variations can be best explained by a planetary companion with an estimated mass of m $\\sin i = 8.78 M_{Jupiter}$ and a semi-major axis of $a = 1.19$ AU, assuming a stellar mass of 1.23 \\Msun. } {} ",
        "introduction": "Since the first exoplanet around a main sequence star was discovered in 1995, more than 300 exoplanets have been detected (http://exoplanet.eu/catalog.php). Among them, the majority of exoplanets candidates has been detected by using radial velocity methods around late F-G and K main-sequence stars.  The existence of planets around intermediate mass early-type stars and their planetary parameters have not been investigated well due to their fast stellar rotational velocities and the lack of an appropriate number of sharp absorption lines in the stellar spectra. When intermediate mass stars evolve toward the red giant stage, they go through the G and K-giant phase where many sharp absorption lines appropriate for high precision RV measurements are available. Therefore, G and K giant stars are suitable targets for detecting exoplanets with the RV technique. There are several ongoing exoplanet survey projects around giant stars (Setiawan et al. 2005, Hatzes et al. 2005, Sato et al. 2007, Johnson et al. 2007, Lovis \\& Mayor 2007, Niedzielski et al. 2007, and Liu et al. 2008). Now, more than 20 exoplanets have been detected around giant stars and from this sample some statistical studies on the planetary systems around giants have been made (Pasquini et al. 2007, Hekker et al. 2008, and Sato et al. 2008).  We started our precise RV survey using the 1.8m telescope at BOAO (Bohyunsan Optical Astronomy Observatory) in 2003 in order to search for exoplanets and to study the pulsations and stellar surface activity of K-giant stars. Our sample consists of 55 K0 - K4  giant  stars. Most of them are brighter than fifth magnitude. Results from  the first 3 years of our survey show that the majority (more than 90\\%) of our sample of K-giants show RV dariations. Among them, we found spectroscopic binaries, pulsating variables (see Kim et al. 2006), and several periodic RV variable stars. Our survey confirmed an exoplanet around $\\beta$ Gem (Han et al. 2008) discovered by Hatzes et al. (2006) and Reffert et al. (2006). In this paper we present a new exoplanet detection around the giant star \\gam1leo. ",
        "conclusions": "There is no doubt that a 429-d periodic variation is present in the RV measurements of {\\gam1leo}. This variation has persisted for almost 5 cycles with no change in phase or amplitude. To clarify the cause of the periodic RV variation, we investigated the possibility of any correlation between the RV variation and other stellar variations such as Hipparcos photometry, EW of Ca II $\\lambda$ 8662 line, line bisectors, and Ca II H line profile. Within the accuracy of our measurements, we found no convincing evidence of such correlation. Though we leave the real cause of the detected 429-d periodic RV variation open, we conclude that the observed periodic RV variation is best explained by an unseen companion with an estimated mass of m $\\sin i = 8.78 \\pm 0.2 M_{Jupiter}$. The non-negligible eccentricity of 0.14 found from the orbital solution also lends some support to a Keplerian orbital motion as the cause of the periodic RV variations. But we would like to emphasize that continued RV observations of \\gam1leo are essential to clarify the true nature of the RV variation. Perhaps the best way to confirm the orbital hypothesis is to detect the astrometric perturbation. For \\gam1leo the expected astrometric perturbation is several tenths of milli arcseconds, which is about the limit of current ground-based astrometric techniques.  The nature of the 1340-d period found in the residual RVs remains unclear. We do not see any evidence for this period in either Hipparcos photometry, Ca II EW, or line bisectors. At face value this would argue for an additional companion, but this is not certain. Continued RV measurements are needed to confirm a possible second companion. We also found a significant power at P = 8.5 days in the residual RVs. The 8.5-d period is close to the estimated fundamental radial pulsation period. To confirm the reality of the 8.5-d period, more observations with better sampling are needed.  It is known that planet harboring giants does not show metal-rich tendency found from dwarf stars. The low metalliicity of \\gam1leo strengthens this trend further. \\gam1leo adds one more exoplanet discovered around a binary star system. Since there are not still many exoplanet found around binary system, we consider our discovery as a valuable addition to the field."
    },
    "0911/0911.0473_arXiv.txt": {
        "abstract": "We investigate the recovery chances of highly spinning waveforms immersed in LIGO S5-like noise by performing a matched filtering with $10^{6}$ randomly chosen spinning waveforms generated with the LAL package. While the masses of the compact binary are reasonably well recovered (slightly overestimated), the same does not hold true for the spins. We show the best fit matches both in the time-domain and the frequency-domain. These encompass some of the spinning characteristics of the signal, but far less than what would be required to identify the astrophysical parameters of the system. An improvement of the matching method is necessary, though may be difficult due to the noisy signal. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0911/0911.5076.txt": {
        "abstract": "% { The Bulge is the least understood major stellar population of the Milky Way. Most of what we know about the formation and evolution of the Bulge comes from bright giant stars. The underlying assumption that giants represent all the stars, and accurately trace the chemical evolution of a stellar population, is under debate. In particular, recent observations of a few microlensed dwarf stars give a very different picture of the evolution of the Bulge from that given by the giant stars. } { We aim to resolve the apparent discrepancy between Bulge metallicity  distributions derived from microlensed dwarf stars and giant stars. Additionally, we aim to put observational constraints on the elemental abundance trends and chemical evolution of the Bulge. } { We perform a detailed elemental abundance analysis of dwarf stars in the Galactic bulge, based on high-resolution spectra that were obtained while the stars were optically magnified during gravitational microlensing events. The analysis method is the same as for a large sample of F and G dwarf stars in the Solar neighbourhood, enabling a fully differential comparison between the Bulge and the local stellar populations in the Galactic disc. } { We present detailed elemental abundances and stellar ages for six new dwarf stars in the Galactic bulge. Combining these with previous events, here re-analysed with the same methods, we study a homogeneous sample of 15 stars, which constitute the largest sample to date of microlensed dwarf stars in the Galactic bulge.  We find that the stars span the full range of metallicities from $\\rm [Fe/H]=-0.72$ to $+0.54$, and an average metallicity of $\\rm\\langle [Fe/H]\\rangle=-0.08\\pm 0.47$, close to the average metallicity based on giant stars in the Bulge. Furthermore, the stars follow well-defined abundance trends, that for $\\rm [Fe/H]<0$ are very similar to those of the local Galactic thick disc. This suggests that the Bulge and the thick disc have had, at least partially, comparable chemical histories. At sub-solar metallicities we find the Bulge dwarf stars to have consistently old ages, while at super-solar metallicities we find a wide range of ages. Using the new age and abundance results from the microlensed dwarf stars we investigate possible formation scenarios for the Bulge. } {} ",
        "introduction": " ",
        "conclusions": ""
    },
    "0911/0911.0190_arXiv.txt": {
        "abstract": "{A hypothetical time-variation of the gravitational constant $G$ would make neutron stars expand or contract, so the matter in their interiors would depart from beta equilibrium. This induces non-equilibrium weak reactions, which release energy that is invested partly in neutrino emission and partly in internal heating. Eventually, the star arrives at a stationary state in which the temperature remains nearly constant, as the forcing through the change of $G$ is balanced by the ongoing reactions. Using the surface temperature of the nearest millisecond pulsar (PSR J0437$-$4715) inferred from ultraviolet observations and results from theoretical modelling of the thermal evolution, we estimate two upper limits for this variation: (1) $|\\dot G/G| < 2 \\times 10^{-10}~\\mathrm{yr}^{-1},$ if the fast, ``direct Urca'' reactions are allowed, and (2) $|\\dot G/G|<4\\times 10^{-12}~\\mathrm{yr}^{-1},$ considering only the slower, ``modified Urca'' reactions. The latter is among the most restrictive upper limits obtained by other methods. ",
        "introduction": "A number of theorists have proposed that the so-called ``fundamental constants\" of Nature may vary with cosmological time, ever since Dirac suggested that the gravitational force may be weakening \\citep{dirac37}. Several experiments aimed at constraining the time variation of $G$ have been conducted, and their results can be grouped by the time scales they probe (see \\citealt{reisenegger07} for a list with references): (1) \\emph{Human lifetime} ($\\sim 10$~yr) experiments rely on real-time monitoring of distances within the Solar System, white-dwarf oscillation periods, and pulse arrival times of isolated and binary pulsars, (2) \\emph{long time} ($\\sim 10^{9-10}$~yr) measurements use stellar astrophysics and paleontology, and (3) \\emph{cosmological} ($\\gtrsim 10^{10}$~yr) experiments use the Big Bang nucleosynthesis and Cosmic Microwave Background anisotropies to compare the value of $G$ in the early Universe to the present value. Generically, they set constraints on $|\\dot{G}/G|$ down to $\\sim 10^{-12}$~yr$^{-1}$, although the comparison between different timescales  depends on the assumed form of the function $G(t)$.  Here, we review our previously introduced method for setting constraints on $\\dot{G}$ \\citep{jofre06}, to our knowledge the only one so far to probe timescales $\\sim 10^{7-9}~\\mathrm{yr}$, intermediate between the \\emph{human} and \\emph{long} timescales mentioned above. It relies on the change in the internal structure of a neutron star induced by a change in $G$, which, together with the slow response timescale of weak interactions, result in internal heating and an increase in the observed surface temperature.  Neutron stars have mean densities exceeding nuclear saturation density, $\\rho_{\\rm nuc}\\sim 3\\times 10^{14}$~g~cm$^{-3}$. In their outer layers, they are composed of heavy atomic nuclei and free electrons, giving way to free neutrons, free protons, muons, and potentially more exotic particles as the density increases inward. A short time after their formation, their internal temperatures drop orders of magnitude below the Fermi energies of free particles ($\\sim 10-100$~MeV), and so their structure is well approximated by zero-temperature models.  Nonetheless, weak interactions still play an important role, as characteristic temperatures $10^{6-8}$~K yield non-negligible neutrino emission. Given that the equation of state of matter above nuclear density still remains poorly known, researchers construct models to calculate the evolution of the thermal content of neutron stars and compare their predictions with X-ray observations, in order to set constraints on the models for dense matter (e.g., \\citealt{yakovlev04}). In these thermal evolution models, the structure of the star is calculated assuming an equation of state, and the evolution of the internal temperature is obtained by considering losses due to different neutrino emission processes from the stellar interior, as well as thermal electromagnetic radiation from the surface.  We have previously explored \\emph{rotochemical heating}, the effect of the progressive loss of rotational support of millisecond pulsars has on their internal structure, which leads to weak interaction processes (beta decays) and net heating \\citep{reisenegger95,reisenegger97,fernandez05}. A time variation of the gravitational constant has an analogous effect: it compresses the star and therefore also causes heating (\\emph{gravitochemical heating}, \\citealt{jofre06}). In what follows, we discuss how this process constrains $\\dot{G}$. ",
        "conclusions": "\\emph{Gravitochemical heating} sets constraints on the time variation of the gravitational constant, using the fact that a non-zero change would generate internal heating in neutron stars, which for nearby cases can be detected with existing telescopes. These constraints are the only ones on timescales $10^{7-9}$~yr. If it could be assured that the observed neutron stars cannot have direct Urca reactions, these constraints would be of the same order as the best ones available on other time scales. This is not the case at the moment, but could become so as neutron star interiors become better understood from other studies. In particular, observed temperatures of other old neutron stars (such as millisecond pulsars) will provide useful information."
    },
    "0911/0911.2689_arXiv.txt": {
        "abstract": "We describe a new algorithm for the ``perfect'' extraction of one-dimensional spectra from two-dimensional (2D) digital images of optical fiber spectrographs, based on accurate 2D forward modeling of the raw pixel data.  The algorithm is correct for arbitrarily complicated 2D point-spread functions (PSFs), as compared to the traditional optimal extraction algorithm, which is only correct for a limited class of separable PSFs.  The algorithm results in statistically independent extracted samples in the 1D spectrum, and preserves the full native resolution of the 2D spectrograph without degradation.  Both the statistical errors and the 1D resolution of the extracted spectrum are accurately determined, allowing a correct $\\chi^2$ comparison of any model spectrum with the data. Using a model PSF similar to that found in the red channel of the Sloan Digital Sky Survey spectrograph, we compare the performance of our algorithm to that of cross-section based optimal extraction, and also demonstrate that our method allows coaddition and foreground estimation to be carried out as an integral part of the extraction step. This work demonstrates the feasibility of current- and next-generation multi-fiber spectrographs for faint galaxy surveys even in the presence of strong night-sky foregrounds.  We describe the handling of subtleties arising from fiber-to-fiber crosstalk, discuss some of the likely challenges in deploying our method to the analysis of a full-scale survey, and note that our algorithm could be generalized into an optimal method for the rectification and combination of astronomical imaging data. ",
        "introduction": "Optical fibers offer a compelling advantage for astronomical survey spectroscopy.  By allowing the light from targets over a wide field of view on the sky to be rearranged into a compact format and fed to any number of spectrographs in parallel, they can provide a vast multiplex advantage over imaging spectrographs.  For this reason, fiber technology has been adopted for a number of major survey programs such as the Las Campanas Redshift Survey (LCRS: \\citealt{schech_lcrs}), the Two-degree Field Survey (2dF: \\citealt{col_2df}), the Sloan Digital Sky Survey (SDSS: \\citealt{york_sdss}), and the recently initiated Baryon Oscillation Spectroscopic Survey (BOSS: \\citealt{sch_boss}) of the SDSS3 project. The use of fiber spectrographs for faint-object spectroscopy has however been restrained by concerns over throughput and systematic limitations on the quality of subtraction of night-sky emission foregrounds.  The spectroscopic extractions of the SDSS multi-fiber instrument have established a high standard, but significant systematic shortcomings remain. Methods to characterize and partially remove sky-subtraction residuals in SDSS fiber spectra can mitigate the problem somewhat (e.g.\\ \\citealt{bol04, wild05}), but do not substitute for a formally correct sky-subtraction model. The faintest spectroscopic galaxy surveys have tended to make use of multi-slit imaging spectrographs (e.g., \\citealt{cow96, steidel03, dav03, gdds1, lef05}), along with sky-subtraction techniques such as nod-and-shuffle \\citep{cuil94, glaz01} or B-spline-based modeling of the two-dimensional sky spectra in the slits (e.g., \\citealt{kel03}). Nod-and-shuffle techniques in particular are ill-suited to most fiber spectrographs, require at least double the detector area, and furthermore reduce the background-limited signal-to-noise by a factor of $1/\\sqrt{2}$ due to their subtraction of data from data.  This paper outlines an algorithmic framework for the modeling and extraction of optical and near-infrared astronomical fiber spectroscopy to the limit of photon noise and native instrumental resolution. By comparing our method with the standard techniques of optimal extraction currently in wide use, we identify and resolve key systematic barriers to ``perfect'' extraction.  As a result, multifiber spectroscopy emerges as a clear and compelling technique for current and future generations of faint-galaxy spectroscopic surveys, even well below the brightness of the night sky at all wavelengths.  This algorithm will deliver significant benefits to the reanalysis of the original SDSS (hereafter SDSS1) archive and to the ongoing analysis of the BOSS survey, both for core redshift-survey goals and for projects that aim to select rare emission-line objects from within the regions of the spectrum dominated by OH line emission from the night sky (e.g., \\citealt{bol04, wil05, bol06, wil06, bol08}).  In considering this subject, we will make a clear distinction between the problems of ``calibration'' and ``extraction'': calibration is the description of the way in which any set of astronomical and environmental stimuli translate into the responses of the digital detectors (which we assume here to be pixellated charge-coupled devices, or CCDs); extraction is the reconstruction of particular stimuli from particular responses. More strictly speaking, we view an ``extracted spectrum'' not as a model for the flux of the observed source itself, but rather as a properly calibrated one-dimensional compression of the instrumental response to an observation of that source. When executed and reported correctly, an extracted spectrum permits a statistically valid $\\chi^2$ test of an input model spectrum against the extracted pixels. This paper specifically illustrates a method for carrying out this sort of perfect extraction assuming that perfect calibration is available. We do not mean to trivialize the problem of calibration, and in our concluding remarks we will discuss the relationship of our extraction method to current and future spectroscopic calibration regimes.  The paper is organized as follows. Section~\\ref{mod2d} frames the problem of extraction in terms of image modeling, lays out the first part of our algorithm, and compares its performance with that of standard extraction techniques in terms of the quality of their respective models to the raw 2D data.  Section~\\ref{resol} addresses the issue of resolution and covariance in the extracted spectra, and presents the second part of our method, which establishes optimal properties in both these regards. Section~\\ref{test4x3} presents a more realistic demonstration of our algorithm on simulated data, illustrating several additional strengths and subtleties of the method. Finally, Section~\\ref{dc} provides a discussion and conclusions.  We will observe the following conventions in this paper. Without loss of generality, we will assume that spectroscopic traces run roughly parallel to CCD columns (i.e., ``vertically''), with wavelength increasing with row number and cross-sectional position increasing with column number.  We will denote vectors in lowercase bold-face type ($\\mathbf{f}$) and matrices in uppercase bold-face type ($\\mathbf{A}$).  We assume all errors have a Gaussian distribution, and we assume no formal priors on fitted model parameters. ",
        "conclusions": "\\label{dc}  The extraction algorithm that we have described does not come without a price.  In the presence of fiber-to-fiber cross-talk, the standard row-wise optimal extraction will couple extracted amplitudes between fibers in a single row, leading to a banded matrix of dimension equal to the number of fibers that must be inverted; this process must then be repeated for each row in each exposure. Our 2D modeling extraction, on the other hand, couples extracted amplitudes between fibers, wavelengths, and exposures. Thus the matrix to invert for the solution set of spectra has sides of dimension equal to \\begin{equation} N_{\\mbox{spectra}} \\times N_{\\mbox{samples per spectrum}} \\times N_{\\mbox{exposures}}~. \\end{equation} For one SDSS1 pointing, this would correspond to 320 fibers (one of the two spectrographs), approximately 4000 sampling points, and three exposures: i.e., a square matrix nearly 4 million to a side.  With a brute-force approach this matrix could never be stored, let alone inverted, with any hardware of the present or foreseeable future.  The way forward to determining the extracted spectra will no doubt lie in exploiting the sparseness of the inverse covariance matrix to reduce storage and computation, and to apply an iterative method such as the conjugate gradient to solve for the extracted spectra.  To determine the resolution with which to reconvolve the extracted spectrum, which formally requires the inversion of the full inverse covariance matrix $\\mathbf{C}^{-1}$ , the practical solution will be to invert a sufficient subspace of $\\mathbf{C}^{-1}$ surrounding each spectrum (or subsegment of a spectrum) to define an acceptably accurate approximation to the desired resolution matrix. The exact requirements will depend upon the degree of cross-talk between neighboring fibers in the spectrograph under consideration. Even with these strategies, a usable implementation of our algorithm for real multifiber survey data may require high-performance parallel computing, depending on the computational expense of the necessary matrix-element calculations.  Most immediately, we plan to implement the strategy outlined in this paper to the extraction of spectra from the BOSS Survey. We also plan to conduct a reanalysis of the SDSS1 archive, to provide the definitive version of this important spectral database with the best possible extracted resolution, signal-to-noise, and foreground subtraction.  These techniques also offer promise for the upcoming Apache Point Observatory Galaxy Evolution Experiment (APOGEE: \\citealt{ap_apogee}) of the SDSS3, a high-resolution, near-infrared multifiber survey of red giant stars in our Galaxy that aims to constrain their evolutionary history as traced by multiple chemical abundances. This extraction strategy will also form a key part of the technical feasibility of the BigBOSS survey \\citep{sch_bigboss}, which proposes to use a 5000-fiber spectrograph fed by a $3^{\\circ}$ field-of-view focal plane positioner system on a 4m-class telescope to measure the baryon acoustic scale and redshift-space distortions over the redshift range $0.2 < z < 3.5$.  Although the implementations for these different surveys will differ in detail, we believe that the software engine for the core extraction computations can be written in a general-purpose form. The application of this technique to slit spectroscopy may be possible as well, although the preservation of object spatial information by the slit makes the problem substantially more complicated.  In all cases, the full power of this extraction can only be realized with sufficiently accurate calibration.  The current standard calibrations for fiber spectroscopy are exposures of uniform spatial illumination by flat-spectrum incandescent lamps (``flats'') and multi-emission line gas-discharge lamps (``arcs'').  Within the framework of our extraction algorithm, the former will determine the relative sensitivity of the individual fibers and pixels, and the latter will determine the spectrograph PSF shape and position as a function of illuminating wavelength in each fiber.  Assuming there are no systematic offsets between the calibrations and the science exposures, and assuming that the variation of the spectrograph PSF is smooth enough with wavelength to be well-sampled by the arc frames (which are sparsely distributed in wavelength), these calibrations should contain sufficient information to determine the calibration matrix $\\mathbf{A}$. We may proceed by extracting the arcs and flats together, with each one described by a single spectrum projected through an initial guess for $\\mathbf{A}$, and then optimize the parameters of $\\mathbf{A}$ by non-linear iteration so as to improve the quality of the 2D extraction models to convergence. A more direct and ambitious approach to the determination of the calibration matrix would be to illuminate the facility calibration screen with either a high-wattage monochrometer or a tunable laser (c.f.\\ \\citealt{st06, cra09}), and to step the illumination source through wavelength in subsequent exposures so as to map out $\\mathbf{A}$ explicitly. In practice, with the exception of spectrograph systems that are very stable thermally and mechanically, there is likely to be some shifting of the fiber positions and focus on the CCDs between the calibration and science frames. In this case, ``tweaks'' to the parameters of $\\mathbf{A}$ will be derived from the shapes and positions of the fiber traces and night-sky emission lines in the science frames. Finally, we note that the calibration may be significantly improved by incorporating all available knowledge about the optical design of the telescope and instrument, rather than treating the system as a black box to be specified entirely by empirical calibration data.  Putting aside the computational challenges that arise from the consideration of continuously two-dimensional input data, the method we have described can also be generalized into an optimal recipe for the rectification and combination of multiple CCD imaging exposures: i.e., taking the $\\mathbf{f}$ of \\S\\ref{mod2d} to be a \\textit{two-dimensional} image model to be extracted from the data. As with the spectroscopic application, the resulting extractions would have optimal resolution, statistically independent extracted image samples, and a definition in terms of the optimization of a well-motivated scalar objective function describing the quality of a model fit to all of the individual exposures.  The implementation would be non-trivial, but the benefits could be great.  A significant challenge on the calibration side is that the details of the imaging PSF, which must be known, will depend upon the spectral energy distributions of the imaged objects, which will in general be unknown.  In summary, we have demonstrated a new spectrum extraction algorithm for optical and near-infrared astronomical fiber spectroscopy. Given sufficiently accurate calibration, this method can extract spectra to the statistical noise limit in the presence of arbitrarily complicated point-spread functions and arbitrarily high and wavelength-varying foregrounds.  The extracted spectra have optimal and precisely quantified resolution and signal-to-noise, along with statistically uncorrelated pixels, for any number of sub-exposures in combination together. As such, statistically accurate $\\chi^2$ tests may be made between the extracted data and theoretical models of the input object spectrum. This algorithm represents a fundamental improvement upon the current state-of-the-art methods in use for the extraction of fiber spectroscopy, and thus motivates a serious and positive reevaluation of the promise of fiber-fed spectrographs for next-generation ground-based faint-object surveys."
    },
    "0911/0911.2530_arXiv.txt": {
        "abstract": "We present an unbiased deep \\OII emission survey of a cluster XMMXCS~J2215.9-1738 at $z=1.46$, the most distant cluster to date with a detection of extended X-ray emission. With wide-field optical and near-infrared cameras (Suprime-Cam and MOIRCS, respectively) on Subaru telescope, we performed deep imaging with a narrow-band filter $NB912$ ($\\lambda_c$ = 9139\\AA, $\\Delta\\lambda$=134\\AA) as well as broad-band filters ($B$, $z'$, $J$ and $K_s$). From the photometric catalogues, we have identified 44 \\OII emitters in the cluster central region of $6'\\times6'$ down to a dust-free star formation rate of 2.6 \\Msun/yr (3$\\sigma$). Interestingly, it is found that there are many \\OII emitters even in the central high density region. In fact, the fraction of \\OII emitters to the cluster members as well as their star formation rates and equivalent widths stay almost constant with decreasing cluster-centric distance up to the cluster core. Unlike clusters at lower redshifts ($z\\la1$) where star formation activity is mostly quenched in their central regions, this higher redshift 2215 cluster shows its high star formation activity even at its centre, suggesting that we are beginning to enter the formation epoch of some galaxies in the cluster core eventually. Moreover, we find a deficit of galaxies on the red sequence at magnitudes fainter than $\\sim$$M^*+0.5$ on the colour-magnitude diagram. This break magnitude is brighter than that of lower redshift clusters, and it is likely that we are seeing the formation phase of more massive red galaxies in the cluster core at $z\\sim1$. These results may indicate inside-out and down-sizing propagation of star formation activity in the course of cluster evolution. ",
        "introduction": "\\label{sec;intro} Galaxy formation and evolution is strongly dependent on environment and on galaxy mass (e.g., \\citealt{dre97,tan05,coo06,tas09}). Galaxies in high density regions are systematically older than those in lower density regions, and massive galaxies are on average older than lower mass galaxies in stellar ages. It seems that massive galaxies in the high density regions form first and galaxy formation activity is propagated to lower density regions and to lower-mass galaxies with time.  Such dependence on environment and mass can be understood by intrinsic effects and extrinsic effects. Intrinsic effects are those determined by the initial condition of galaxy formation and extrinsic effects are those in effect during the evolution of galaxies after formation.  The environmental dependence can be partly understood by the intrinsic effect called ``galaxy formation bias'' (e.g., \\citealt{cen93}) where high density regions should have started off from the largest initial density fluctuations that collapse first in the universe, and galaxy formation takes place earliest in such regions and subsequent evolution is prompted. In lower density regions, however, galaxy formation is delayed and time scales of star formation and mass assembly are longer.  Likewise, the mass dependence of galaxy formation called ``down-sizing'' (e.g., \\citealt{cow96,kod04a}) may be partly understood by the scaled-down version of the bias on galactic scale, where more massive galaxies today correspond to higher initial density fluctuations on galactic scale and their formation such as star formation and assembly of building blocks take place earlier than less massive galaxies.  Galaxies are also subject to external effects from their surrounding environments, such as galaxy-galaxy interactions/mergers and ram-pressure stripping in dense environments (e.g., \\citealt{aba99,qui00} ). Such interactions may enhance and/or quench the star formation activity in galaxies preferentially in high density regions, which would therefore result in conspicuous dependence of galaxy properties on environments.  However, the relative importance of the intrinsic effects and the extrinsic effects is almost totally unknown yet. One of the most effective methods to verify the existence of the intrinsic effects is to go back in time and directly see the galaxies in the distant universe. By doing so, we may eventually reach the epoch when galaxies are forming rapidly in the biased cluster core, while galaxies are not yet formed or only slowly forming in lower density regions.  Also, as we go back in time, host galaxies of star formation would be shifted to higher mass systems and we may eventually see the active star formation in action in massive galaxies in the cluster core.  In the low redshift universe, star formation activity is a monotonically decreasing function of local density. However, we have not yet known the dependence of star formation activity on environments at high redshifts in detail. According to recent studies, it seems that star formation is in fact biased at high density region at high redshifts. \\citet{elb07} reported for the first time in the GOODS North/South surveys that at $z\\sim1$ the galaxies at denser environment tend to have higher star formation rates (SFRs), in contrast to the local Universe. \\citet{coo08} and \\citet{ide09} also showed similar trends for \\OII emitters at $z\\sim1$ in DEEP2 and COSMOS surveys, respectively. At a slightly lower redshift, $z=0.81$, based on the mid-infrared observation of the RX~J1716.4+6708 cluster, \\citet{koy08} reported that the star formation activity is probably enhanced in the medium density region, such as cluster outskirts or galaxy groups, rather than in the highest density region. \\citet{pog08} also suggested that SFRs of galaxies may have a peak at intermediate densities at $z$=0.4--0.8 based on the EDisCS survey. These observational findings may imply that the environment that hosts active star formation is shifted towards denser regions at higher redshifts.  On the other hand, it is well-known as the ``Butcher-Oemler (B-O) effect'' that more distant clusters show higher fraction of blue galaxies up to $z\\sim0.5$, suggesting the enhancement of star formation activity in higher-$z$ clusters (e.g., \\citealt{bo78,bo84,mar01}). However, the fraction of blue galaxies decreases with cluster centric radius, and red galaxies still dominate in the core regions at $z\\sim0.5$ (\\citealt{kod01,ell01}). Such studies of the B-O effect in clusters have been extended up to $z\\sim1$.  \\citet{pos01} found that the fraction of active galaxies in the central regions of clusters is higher at $z\\sim$0.7--0.9 compared to $z\\sim$0.2--0.5. However, \\citet{nak05} suggested that the fraction of galaxies with strong \\OII emission in the cluster cores at $z<1.0$ is not significantly dependent on redshift. It thus seems that we have not yet reached a bursting phase of massive galaxy formation in cluster cores even at $z\\sim1$, although the global activity of star formation within clusters is already enhanced.  In this paper, we present an \\OII emitter survey of the XMMXCS~J2215.9-1738 cluster (hereafter 2215 cluster) with a narrow-band filter $NB912$ on Suprime-Cam (Figure \\ref{fig;transmission}), and discuss spatial distribution of star formation activity (see also \\citealt{koy09}, which presents our \\Ha emitter survey of the RX~J1716.4+6708 cluster at $z=0.81$). The 2215 cluster is the most distant cluster to date at $z$=1.46 with a detection of extended X-ray emission (\\citealt{sta06}). \\citet{hil07,hil09} have confirmed dozens of the cluster members by spectroscopy, and found that velocity dispersion of the member galaxies is $\\sigma=580\\pm140$ [km s$^{-1}$]. With the $NB912$ filter, we can survey \\OII emissions from the galaxies with the line-of-sight velocities between $-$2794 $<$ $\\Delta V_{\\rm los}$ [km s$^{-1}$] $<$ 1598 with respect to the velocity centre of the cluster. Our $NB912$ filter thus perfectly matches this cluster, and should be able to detect all the \\OII emission lines from the cluster members to a certain flux limit without introducing any bias (Figure \\ref{fig;transmission}). Therefore, our survey is unique, and the 2215 cluster is an ideal target for us to investigate the environmental dependence of star formation activity over a wide range in environment at this high redshift.  The X-ray luminosity and inter-cluster medium temperature for the 2215 cluster are $L_X=4.4^{+0.8}_{-0.6}\\times10^{44}$[erg\\ s$^{-1}$], and $kT=7.4^{+2.6}_{-1.8}$[keV], respectively (\\citealt{sta06}). The luminosity is fainter than what is expected from the temperature compared to the local $L_X-T$ relation.  \\citet{hil07} thus point out that this cluster may experience a merger within the last few Gyr. It is also found that colour-magnitude diagram in this cluster shows red sequence (\\citealt{sta06,hil09}). \\citet{hil09} investigated the morphologies of bright member galaxies, and found that about 60\\% of members are E or S0 galaxies. Even at $z=1.46$, cluster core is already dominated by early-type galaxies, as far as morphology is concerned.  The structure of this paper is as follows. The observations and data reduction are described in \\S~\\ref{sec;obs}. In \\S~\\ref{sec;selection}, we describe how we select the \\OII emitters associated to the cluster at $z=1.46$ from the photometric catalogues. We show the spatial distribution, SFRs, and equivalent widths of the \\OII emitters, and investigate the star forming activity in the 2215 cluster in \\S \\ref{sec;results}.  We also discuss the deficit of faint red galaxies and its evolution. A summary is given in \\S \\ref{sec;summary}. Throughout this paper, magnitudes are in the AB system, and we adopt cosmological parameters of $h=0.7$, $\\Omega_{m0}=0.3$ and $\\Omega_{\\Lambda 0}=0.7$. Vega magnitudes in $J$ and $K_s$, if preferred, can be obtained from our AB magnitudes using the relations: $J$(Vega)=$J$(AB)$-$0.92 and $K_s$(Vega)=$K_s$(AB)$-$1.80.   \\begin{figure} \\begin{center} \\includegraphics[width=70mm]{transmission.eps} \\end{center} \\caption{A transmission curve of the $NB912$ filter ($\\lambda_c$=9139\\AA, $\\Delta\\lambda$=134\\AA). Upper x-axis shows peculiar velocity with respect to the velocity centre of the cluster at $z=1.46$. The histograms show the velocity distribution of the spectroscopically confirmed member galaxies of the 2215 cluster (\\citealt{hil09}). } \\label{fig;transmission} \\end{figure} ",
        "conclusions": "\\label{sec;results} \\subsection{Star formation activity in a cluster core} \\label{sec;sf_activity}   \\begin{figure} \\begin{center} \\includegraphics[width=70mm]{r_fraction_linear_oii.eps} \\end{center} \\caption{Fraction of \\OII emitters to the cluster members as a function of radius from a cluster centre. The contribution of field galaxies is statistically subtracted from our colour-selected cluster member candidates. The error-bars indicate the statistical errors associated to both \\OII emitters and the cluster members. } \\label{fig;oii_fraction} \\end{figure}  \\begin{figure} \\begin{center} \\includegraphics[width=84mm]{comparison_2215-1716.eps} \\end{center} \\caption{Close-up views of the central $2'\\times2'$ regions for XCS2215 cluster at $z=1.46$ (left) and RXJ1716 cluster at $z=0.81$ (right). Blue open squares show \\OII emitters for 2215 cluster and \\Ha emitters for RXJ1716 cluster, respectively. Dots show cluster member candidates selected in \\S\\ref{sec;clmember} for the 2215 cluster, and member galaxies selected by photometric redshifts of $\\Delta z=0.76-0.83$ for RXJ1716 cluster (\\citealt{koy07}), respectively. Both panels show the areas of similar physical scales (0.51Mpc/arcmin at $z=1.46$ and 0.45Mpc/arcmin at $z=0.81$, respectively). } \\label{fig;1716} \\end{figure}  \\begin{figure*} \\includegraphics[width=55mm]{radius_sfr_oii_linear.eps} \\includegraphics[width=55mm]{radius_ssfr_oii_linear.eps} \\includegraphics[width=55mm]{radius_ew_oii_linear.eps} \\caption{ SFRs(left), specific SFRs (SSFRs) (middle) and observed equivalent widths (right) for \\OII emitters as a function of distance from the cluster centre. SFRs are derived and corrected for dust extinction using the relations in \\citet{mou06} (see text). Stellar masses are estimated using the equation of \\citet{dad04}. Equivalent widths are calculated with the equations (\\ref{eq;fline}) and (\\ref{eq;fcont}). Blue open squares show 44 \\OII emitters, and 4 red \\OII emitters are marked by magenta filled square. Typical errors on each value are shown in the upper or the lower right corner on each panel. The errors on SFR and equivalent width are derived from magnitude errors in $NB912$ and $z'$, and $\\sigma(\\Delta \\log({\\rm M_\\star}))=0.20$ is applied to the error in stellar mass (see text). In each diagram, the solid line connects median values in four radius bins, and the errors contain both the typical error on each value and the standard deviation in each bin. } \\label{fig;sfr_ew} \\end{figure*}  Figure \\ref{fig;map_member} shows that there are many \\OII emitters in the central region of the 2215 cluster. If we assume that \\OII lines are emitted from ionized gas in/around the star forming regions, it seems that the 2215 cluster is still actively forming stars even at its core region. This is not the case in lower-$z$ clusters where star forming activity is much lower in the central region. To be more quantitative, we estimate a fraction of \\OII emitters to cluster members as a function of cluster-centric distance in Figure~\\ref{fig;oii_fraction}. Here, the amount of remaining contamination from foreground/background galaxies are estimated from the SDF sample, and it is statistically subtracted from the colour-selected sample of cluster member candidates. Figure \\ref{fig;oii_fraction} suggests that the 2215 cluster maintains a high fraction of star forming galaxies in the central region, $\\sim 30$\\%, even at the most inner part within 0.25~Mpc in physical scale (1'=0.51Mpc at $z=1.46$). This fraction is higher than the \\OII fraction in the cluster cores in \\citet{nak05}, which is $\\la 20$\\% at $z<1.0$.  \\citet{lid08} recently reported a distribution of spectroscopic sample of \\OII emitters in a cluster XMMU~J2235.3-2557 at $z=1.39$, a similar redshift to that of the 2215 cluster. While there are no galaxies in the very core within 90~kpc in radius that show star forming activity, \\OII lines are detected from more than half of galaxies near the centre. This result is similar to ours on the \\OII emitters in 2215 cluster at a similar (slightly higher) redshift of 1.46. It seems likely that distant X-ray detected clusters at $z\\ga1.4$ are actively forming new stars even in the central region within a few hundreds kpc in radius.  We have recently conducted a narrow-band \\Ha emitter survey for RX~J1716.4+6708 cluster at $z$=0.81 (\\citealt{koy09} in preparation). This survey reveals quite a different situation where no \\Ha emitters are observed in the core region within a radius of 0.25~Mpc (Figure \\ref{fig;1716}). A difference in spatial distribution of the emitters is very impressive. The \\Ha emitters in the 1716 cluster are selected with a combination of $NB119$ ($\\lambda_c$ = 11885\\AA, $\\Delta\\lambda$=141\\AA) and $J$ filters on Subaru/MOIRCS. The \\Ha survey reached to the depth of a limiting line flux of $\\sim$4.1$\\times$10$^{-17}$ erg s$^{-1}$ cm$^{-2}$, which corresponds to a dust-free SFR of $\\sim1.0$\\Msun yr$^{-1}$. The \\Ha survey for the 1716 cluster is, thus, more sensitive to star forming galaxies with slightly lower SFR than the \\OII survey for the 2215 cluster, presented in this paper.  Although the lines are different between the two clusters, \\Ha and \\OII lines, considering the fact that \\Ha line is much less affected by dust extinction than \\OII line, the intrinsic difference in spatial distribution of the emitters would be more significant. Furthermore, \\citet{koy09} find that the fraction of star forming galaxies decreases as one goes to denser region, which is different from what we see for the 2215 cluster.  Similar trends are also found in lower redshift clusters at $z$=0.4--0.8 (\\citealt{kod04b,pog08}). \\citet{koy08,koy09} have also unveiled the star forming activities hidden by dust in the 1716 cluster based on the mid-infrared imaging with AKARI. These studies conclude that star formation activity has already been quenched in the cluster core, while it comes to an peak in the medium density regions away from the cluster core.  These facts may imply that we find galaxy formation bias in the highest density region at $z\\sim1.5$.  Recent studies also support the biased star formation in a relatively dense environment at $z>1$ (\\citealt{elb07,coo08,ide09}). The difference of critical environments in star formation at various redshifts may suggest that galaxy formation bias plays an important role in the dependence of galaxy properties on environments.  In such a comparison between different clusters at different redshifts, we must keep in mind that the size and mass of the clusters can be different. In fact, the bolometric X-ray luminosity of 1716 cluster is $\\sim3$ times larger than that of 2215 cluster (\\citealt{ett04,sta06}). Also, the sub-structures of 1716 cluster is more prominent (\\citealt{koy07}). We cannot therefore conclude whether the difference seen in spatial distribution is largely due to time evolution or due to different masses or characteristics. But it may be the case that we are witnessing the evolution in star forming activity in the core of high-$z$ clusters from $z\\sim1.5$ to $z\\sim0.8$.  On the other hand, \\citet{fin05} report a diversity in distribution of \\Ha emitters in the cores of clusters at $z$=0.7--0.8. Two among their three clusters at similar redshifts show a decrease in \\Ha fraction toward the cluster centre, while the other cluster shows an opposite trend. This fact may imply that it is not only the star formation bias that causes the high fraction of \\OII emitters in the core of the 2215 cluster. It is thus crucial to observe more clusters at $z>1.0$, and to evaluate to what extent the presence of star forming galaxies in the cluster core is a general property at $z>1$.  The $L_X-T$ relation of the 2215 cluster suggests that it is possible that it experienced a merger event within the last few Gyr (\\citealt{hil07}). This can be another possible reason for the high fraction of \\OII emitters, since the cluster merging event might cause an enhancement of star formation activity in galaxies in the cluster core. Also, since the distribution of \\OII emitters is projected on celestial sphere, it may be possible that some member galaxies in the outskirt of the cluster happen to be seen in the direction to the cluster centre, which may result in apparent high fraction of \\OII emitters at the centre. However, the number density of galaxies in the outskirt is likely to be much lower than that of galaxies in the core region (Figure \\ref{fig;bzk}(b)), and such projection effect would be small in the direction to the centre.  It is also possible that an active galactic nucleus (AGN) enhances \\OII line flux.  Recent studies suggest that a fraction of AGNs in clusters increases with redshifts.  \\citet{gal09} have found an overdensity of AGNs within a radius of 0.25~Mpc in clusters at $z>0.5$, and that the density of X-ray selected AGNs in clusters at $1.0<z<1.5$ is 0.102 arcmin$^{-2}$, which is $\\sim2$ times larger than that of clusters at $0.5<z<1.0$.  It indicates that it is possible that 3--4 AGNs reside in the central 33.8 arcmin$^{2}$ region of the 2215 cluster. \\citet{mar09} also show a monotonous increase of AGN fraction in clusters from $z\\sim0.2$ to $z\\sim1.0$. It is worth mentioning that \\citet{sta06} report no obvious X-ray point source in the cluster core, although we cannot completely exclude the presence of any point source. We thus consider that AGN contribution is small. \\citet{yan06} suggest that \\OII emitters in the red sequence tend to be AGNs, most commonly LINERs, and that \\OII line can be valid as an indicator of star formation activity only for blue galaxies. We find four red \\OII emitters with $z'-K_s$ colours redder than $-0.088K_s+4.26$ (see \\S\\ref{sec;cmd} and Figure \\ref{fig;cmd_lf}(a)), which are located near the centre (Figure \\ref{fig;map_member}). Even if all these red \\OII emitters are AGNs, the enhancement of galaxy activity in the core would be still valid, since the presence of AGN is also a sign of activity.  This would then provide us with the clues to understanding the influence of AGNs on the evolution of massive galaxies in the cluster core, such as quenching of star formation. It is essential therefore to obtain NIR spectra of these four \\OII emitters in order to constrain the origin of the \\OII lines.   \\subsection{SFR, specific SFR and equivalent width} According to \\citet{ly07}, \\OII line flux [ergs\\ s$^{-1}$ cm$^{-2}$] and continuum flux density [ergs\\ s$^{-1}$ cm$^{-2}$ \\AA$^{-1}$] for 44 \\OII emitters are calculated from flux densities in $NB912$ and $z'$ bands ($f_{NB912}$ and $f_{z'}$), respectively, as follows; \\begin{equation} F({\\rm [O\\,{\\scriptstyle II}]})=f_{NB912}\\Delta_{NB912}\\frac{1-(f_{z'}/f_{NB912})}{1-(\\Delta_{NB912}/\\Delta_{z'})}, \\label{eq;fline} \\end{equation} \\begin{equation} f_{\\lambda,{\\rm cont}}=f_{z'}\\frac{1-(f_{NB912}/f_{z'})(\\Delta_{NB912}/\\Delta_{z'})}{1-(\\Delta_{NB912}/\\Delta_{z'})}, \\label{eq;fcont} \\end{equation} where $\\Delta_{NB912}$ and $\\Delta_{z'}$ indicate FWHMs of the filters, and $\\Delta_{NB912}=134$\\AA\\ and $\\Delta_{z'}=955$\\AA.  In order to estimate dust-corrected SFR of our \\OII emitters, we use the empirical SFR calibration for \\OII parametrized in terms of the $B$-band luminosity developed by Moustakas, Kennicutt \\& Tremonti (2006) for local star forming galaxies. The $B$-band luminosity was regarded as a proxy of stellar mass. They were aware that the mass-to-light ratios of star forming galaxies span a wide range and that the $B$-band luminosity would not be a perfect indicator of stellar mass. However, they empirically found that the amount of dust extinction and metallicity are correlated with the $B$-band luminosity (see Figure 16 in Moustakas et al. (2006)). They also found that the correlation for star forming galaxies at 0.5$<$$z$$<$1.5 is roughly consistent with that for local galaxies. It should be noted, however, that there is a larger uncertainty in the derived SFR for individual galaxies at high redshifts. We apply their correction scheme of dust extinction to our \\OII emitters at $z=1.46$. We infer the rest-frame $B$-band fluxes from the $J$-band fluxes and $z-J$ colours and use them to correct for dust extinction. Figure \\ref{fig;sfr_ew}(a) shows thus derived dust-corrected SFRs of our \\OII emitters as a function of distance from the cluster center.  Stellar masses of the \\OII emitters are estimated from $K_s$ total magnitudes using the equation given in \\citet{dad04}. \\citet{dad04} estimate the stellar masses for $K$-selected galaxies at $z>1.4$ using K20 survey data. Mass to $K$-band luminosity ratio is derived from multi-wavelength data from $U$ to $K$ with \\citet{sal55} IMF (\\citealt{fon04}). The relation is calibrated with $z-K$ colour, which can reduce the dispersion of derived stellar mass to $\\sigma(\\Delta \\log({\\rm M_\\star}))=0.20$. We also derive specific SFR by dividing SFR by stellar mass. Figure \\ref{fig;sfr_ew}(b) shows the specific SFRs for \\OII emitters as a function of radius from a centre.  The observed equivalent width of an \\OII emission is also derived from \\OII line flux and continuum flux density (Equations (\\ref{eq;fline}) and (\\ref{eq;fcont})). Figure \\ref{fig;sfr_ew}(c) shows the observed equivalent widths for the \\OII emitters as a function of radius from a centre.  Figure \\ref{fig;sfr_ew} suggests that star formation activities of the \\OII emitters in denser regions are as active as those in the outskirt of the cluster. The equivalent widths are not correlated with the position of galaxies in the cluster. These facts also suggest active star formation in the central region of the 2215 cluster. On the other hand, the specific SFRs may show a mild correlation in the sense that \\OII emitters at the inner regions tend to have lower specific SFRs. Since the similar correlation is not seen in their SFRs, this is due to the fact that massive galaxies tend to reside near the centre of the cluster. This tendency may be due to the galaxy formation bias that massive galaxies are formed in the cluster core, or may be a result of efficient mass assembly due to merging that can be more effective in high density regions. Figure \\ref{fig;sfr_ew}(b) also shows that the red \\OII emitters tend to have lower specific SFRs. This may be because we underestimate their intrinsic SFRs due to the strong dust extinction for the red galaxies. Otherwise, this may imply that these galaxies are just quenching their star formation activities.  \\subsection{Colour-magnitude diagram} \\label{sec;cmd}   \\begin{figure*} \\includegraphics[width=70mm]{kzkcolor_clmember-Kdet_center.eps} \\includegraphics[width=70mm]{lf_complete_clmemred_center_Ks-det.eps} \\caption{(a) Left panel: Colour-magnitude diagram of $z'-K_s$ vs. $K_s$ for the $K_s$-detected galaxies in the central $2'\\times2'$ region.  Black dots shows cluster member candidates. Red solid line show the expected location of the colour-magnitude relation of passively galaxies galaxies at $z=1.46$ formed at $z_{form}$=5 (\\citealt{kod98}). We define the red sequence galaxies as those falling between the two broken lines, $\\Delta(z'-K_s)=\\pm0.3$ around the colour-magnitude relation.  Long-dashed and dotted lines are 3$\\sigma$ and 5$\\sigma$ confidence levels in colours, respectively. Blue open squares show 18 \\OII emitters in this region. Three dot-dashed lines indicate iso-stellar mass curves for $2\\times10^{11}$\\Msun, $10^{11}$\\Msun and $5\\times10^{10}$\\Msun, respectively, which are drawn based on the equation given in \\citet{dad04}. (b) Right panel: $K_s$-band luminosity function of the red sequence galaxies in the central $2'\\times2'$ region. The contribution of field galaxies, estimated from the SDF sample, is statistically subtracted.  Error-bars show Poisson errors. } \\label{fig;cmd_lf} \\end{figure*}   Figure \\ref{fig;cmd_lf}(a) shows a colour-magnitude diagram of $z'-K_s$ vs. $K_s$ for cluster member candidates in the central $2'\\times2'$ region. Note that the $K_s$-selected galaxy sample is used in this section. The use of the $NB912$-selected sample may cause incompleteness in the number of galaxies at faint $z'$ magnitudes. The solid line in the figure shows the expected location of the colour-magnitude relation (CMR) of passively evolving galaxies at $z=1.46$ formed at $z_{form}=5$ inferred from the \\citet{kod98} model which is calibrated to reproduce the CMR of elliptical galaxies in the Coma cluster at $z=0$; \\begin{equation} z'-K_s=-0.088K_s+4.56. \\label{eq;cmr} \\end{equation} We define the red sequence galaxies as those falling in-between $\\Delta(z'-K_s)=\\pm0.3$ from the predicted CMR as shown by the broken lines in Figure \\ref{fig;cmd_lf}(a).  The number of the red sequence galaxies seems to decrease at magnitude fainter than $K_s\\sim$21.5. The long-dashed and dotted lines show the 5$\\sigma$ and 3$\\sigma$ confidence levels of $z'-K_s$ colours, respectively. Therefore the decrease in the number of the red galaxies is not driven by incompleteness. \\citet{sta06} and \\citet{hil09} also examine colour-magnitude diagrams, and find that there is a well defined CMR of red galaxies in the 2215 cluster. In \\citet{sta06}, however, shallowness of the data does not allow us to discuss the faint end of the red sequence. \\citet{hil09} use deeper data to investigate the CMR, and colour-magnitude diagram of Figure~6 in \\citet{hil09} shows a deficit of red galaxies at magnitudes fainter than $K_s\\sim21.5$. This is consistent with our result.  Since all the cluster member candidates are plotted in Figure \\ref{fig;cmd_lf}(a), some galaxies actually do not belong to the cluster. In order to correctly evaluate the deficit of the red galaxies, we derive the field-corrected $K_s$-band luminosity function in the central $2'\\times2'$ region by statistically subtracting the field contamination. The detection completeness is also corrected for as follows. Artificial objects with Gaussian profile are randomly generated and embedded in the raw $K_s$ image, and detection of these objects is conducted with the same manner as in \\S\\ref{sec;selection_nbcatalogue}. The assumed distribution of brightness of the artificial objects is uniform within each magnitude bin. The completeness thus estimated in each bin is always as high as $\\ga$90\\%.  The resulting luminosity function of the red sequence galaxies is shown in Figure \\ref{fig;cmd_lf}(b). It is clear that the number of red member galaxies decreases at $K_s\\la 21.5$.  This magnitude corresponds to $K_s^*$+0.5 with respect to the passively evolving galaxies at $z=1.46$ \\citep{kod98}. In the 2215 cluster at $z=1.46$, the red sequence is visible only down to $\\sim$$K_s^*$+0.5 and truncated at that magnitude.  Such truncation is seen, if any, at fainter magnitudes ($\\ga M^*$+0.5) in the cores of lower redshift clusters. For example, \\citet{and06} find no deficit of galaxies on red sequence down to $M^*+3.5$ in the MS1054.4-0321 cluster at $z$=0.83. \\citet{koy07} suggest that the build-up of the CMR depends on X-ray luminosity of clusters at $z\\sim0.8$, and find that the CMR is established down to $M^*+2.0$ even for a X-ray fainter cluster RXJ1716+6708 at $z=0.81$. \\citet{lid04} and \\citet{tan08} study CMR for the RDCS~J1252.9-2927 cluster at $z$=1.24. While the red sequence galaxies only appear down to $K\\sim22$ in sub-clumps in the outskirt of the cluster (\\citealt{tan08}), faint red galaxies clearly exist down to $K\\sim24$ in the main cluster (\\citealt{lid04}). \\citet{lid08} also find that there are red faint galaxies down to $K\\sim24.5$ in the core region of the XMMU~J2235.3-2557 cluster at $z$=1.39. We also note that \\citet{kod07} show that the CMR of proto-clusters at much higher redshifts ($2\\la z\\la 3$) becomes less conspicuous, and even the bright-end of the red sequence seems to disappear in proto-clusters at $z\\ga 2.5$.  In Figure \\ref{fig;cmd_lf}(a), we mark the \\OII emitters with blue open squares.  Most of the \\OII emitters have bluer colours as expected since they are likely to be still forming stars.  The galaxies that are fainter in $K_s$ hence less massive galaxies tend to be slightly bluer. This may suggest that star formation activity in more massive galaxies tends to be truncated at earlier times. We draw three iso-stellar mass curves of $2\\times10^{11}$\\Msun, $10^{11}$\\Msun\\ and $5\\times10^{10}$\\Msun, respectively, using the equation given in \\citet{dad04}. If the \\OII emitters cease their star formation, they would move along these curves until they reach on to the red sequence. If an \\OII emitter has a SFR of 50\\Msun/yr, which is close to the median SFR in our \\OII emitter sample, and keeps this rate constant until $z=1.0$, its stellar mass would increase by $\\Delta$M$_\\star$$\\simeq7\\times10^{10}$\\Msun. If we assume that an \\OII emitter gradually becomes red while increasing its stellar mass at a constant SFR, the faint end of the red sequence would be filled up by $z\\sim1.0$. If this is the case, the \\OII emitters with $\\la5\\times10^{10}$\\Msun\\ may be good progenitors of the faint galaxies on the red sequence. However, we do not know yet the mechanisms of changing galaxy colours and quenching their star formation. As described in the last paragraph of \\S\\ref{sec;sf_activity}, contribution of AGNs to the \\OII emitters on the red sequence can be large (\\citealt{yan06}).  Perhaps AGN feedback is one of the key mechanisms to reduce the star formation activities.  Combining with the previous studies of CMR (see above), we come up with the following scenario of formation of red sequence in clusters at $z\\sim1.5$.  The most massive galaxies brighter than $K_s^*$ (i.e., $>10^{11}$\\Msun) are formed at $z>2$, and become red by quenching the star formation in early epoch. Some galaxies on the red sequence may still be keeping residual star formation activities.  On the other hand, less massive galaxies are actively growing at $z<2$, and we hardly see galaxies fainter than $K_s^*$ on the red sequence. Star forming galaxies with $\\sim5\\times10^{10}$\\Msun\\ at $z\\sim1.5$ may evolve into faint red sequence galaxies by $z\\sim1.0$. This suggest down-sizing propagation of star formation in high redshift clusters.   We performed a unique, unbiased \\OII line survey of star forming galaxies in the XMMXCS~J2215.9-1738 cluster at $z=1.46$, which is currently the most distant cluster ever identified with a detection of extended X-ray emission.  We have obtained wide-field optical ($B$, $z'$, $NB912$) and near-infrared ($J$ and $K_s$) data with Suprime-Cam and MOIRCS, respectively. With a combination of $NB912$ narrow-band filter ($\\lambda_c$ = 9139\\AA, FWHM=134\\AA) and the $z'$-band filter, we detect 69 $NB912$ emitters in the central $6'\\times6'$ region where near-infrared data are also available. Among them, 44 emitters are identified as \\OII emitters associated to the cluster based on the $B-z'$ and $z'-K_s$ colours, down to a dust-free star formation rate of 2.6 \\Msun yr$^{-1}$ (3$\\sigma$).  We find that many \\OII emitters reside in the central high density region even within a radius of 0.25 Mpc (physical scale).  We also find that the fraction of \\OII emitters to cluster members remains high up to the core region. This suggests that the 2215 cluster is still actively forming stars even at the central region, in contrast to lower redshift clusters, where old passively evolving elliptical galaxies dominate. This indicates an inside-out propagation of star formation in high redshift clusters, and we may be eventually beginning to enter the epoch of biased galaxy formation in the densest region at $z=1.46$.  SFRs, specific SFRs, and equivalent widths are derived for 44 \\OII emitters.  It is found that the emitters have similar SFRs and equivalent widths irrespective of the location within the cluster. It seems however that the specific SFRs tend to decrease slightly toward the cluster core, probably due to the fact that more massive \\OII emitters exist near the centre. We may be approaching to the formation phase of massive galaxies at the cluster core.  Moreover, the colour-magnitude diagram in the 2215 cluster shows a deficit of red sequence galaxies fainter than $\\sim$$M^*+0.5$, while the red sequence in lower redshift clusters extends to much fainter magnitudes. While some bright \\OII emitters are located on the red sequence, all the faint \\OII emitters with $>$$M^*$ have blue colours. It is likely that those blue \\OII emitters become redder once they truncate their star formation, and they would eventually reach and fill the faint end of the red sequence at lower redshifts. This indicates a down-sizing propagation of star formation in high redshift clusters."
    },
    "0911/0911.0703_arXiv.txt": {
        "abstract": "The combination of high spatial and spectral resolution in optical astronomy enables new observational approaches to many open problems in stellar and circumstellar astrophysics. However, constructing a high-resolution spectrograph for an interferometer is a costly and time-intensive undertaking.  Our aim is to show that, by coupling existing high-resolution spectrographs to existing interferometers, one could observe in the domain of high spectral and spatial resolution, and avoid the construction of a new complex and expensive instrument. We investigate in this article the different challenges which arise from combining an interferometer with a high-resolution spectrograph. The requirements for the different sub-systems are determined, with special attention given to the problems of fringe tracking and dispersion. A concept study for the combination of the VLTI (Very Large Telescope Interferometer) with UVES (UV-Visual Echelle Spectrograph) is carried out, and several other specific instrument pairings are discussed.  We show that the proposed combination of an interferometer with a high-resolution spectrograph is indeed feasible with current technology, for a fraction of the cost of building a whole new spectrograph. The impact on the existing instruments and their ongoing programs would be minimal. ",
        "introduction": "In recent years optical interferometers have proven that they can produce excellent science in the field of stellar and circumstellar astrophysics. Over the same period high-resolution spectrographs have enabled the discovery of the first extra-solar planets, and contributed substantially to great progress in the field of asteroseismology.  A number of current interferometric instruments have some spectroscopic capabilities. For example the mid- and near-infrared instruments MIDI (The Mid-Infrared instrument, at the VLTI) and AMBER (Astronomical Multiple BEam Recombiner, at the VLTI) provide spectral resolutions of up to $R\\sim250$ and $R\\sim12 000$, over bandpasses of $\\sim 5$~$\\mu$m and $\\sim 50$~nm, respectively. At the CHARA array, the Vega (Visible spEctroGraph and polArimeter) project is under construction with a spectral resolution of $R\\sim30 000$ and a bandpass of $\\sim 50$~nm.  For science results obtained with spectrally resolved interferometry see for example \\cite{vakili1998} and \\cite{weigelt2007}. Unfortunately the combination of very high spectral resolution over a bandpass greater than a few tens of nanometer, to enable a real analog to classical Echelle spectroscopy with interferometric spatial resolution is not yet available.  In building a dedicated high-resolution Echelle spectrograph for an existing interferometer such as the VLTI, one would face several challenges. First of all, building a high-resolution spectrograph is a very costly and time-intensive undertaking. In addition it would be hard to justify building such an instrument only for use with an interferometer, as current optical interferometers are still restricted to very bright objects in comparison with single telescopes. Furthermore, high-resolution spectrographs are usually large instruments, while the space available in the beam combining laboratories of interferometers is often limited. Therefore, it is unlikely that such a dedicated instrument will be built in the near future.  In this article we advocate a different approach. By using an existing spectrograph and only building an interface between it and an interferometer at the same site, the combination of high spectral and spatial resolution could be achieved on a much shorter timescale, and for a fraction of the cost of a complete new instrument.  The pre-existing infrastructure would need to consist of two telescopes, delay lines for path compensation, a fringe sensing unit to acquire and stabilize the fringes, and a high-resolution spectrograph on the same site. These conditions are already fulfilled, or will be fulfilled in the very near future, at several observatories. The two most promising sites are: \\begin{itemize}  \\item In the Southern hemisphere at Paranal Observatory with the VLTI in combination with the UVES spectrograph at Unit Telescope 2 (UT2), or the High-Resolution IR Echelle Spectrometer (CRIRES) spectrograph at UT1.  \\item In the Northern hemisphere at Mauna Kea with the Keck Interferometer (KI) and HIgh Resolution Echelle Spectrometer (HIRES) at Keck I telescope or the NIRSPEC spectrograph at Keck II. Also at Mauna Kea, the OHANA (Optical Hawaiian Array for Nanoradian Astronomy) interferometer is currently under development, and will allow the combination of several other pairs of telescopes at the Manual Kea site. \\end{itemize} In each case, three additional hardware components would be required:   \\begin{enumerate}  \\item A beam combiner that accepts two input beams from the telescopes and feeds the outputs carrying the fringe signals (coded as intensity variations) into fiber feeds;  \\item Fibers that connect the interferometer to the spectrograph;  \\item A fiber head that feeds the light from the fibers into the spectrograph.  \\end{enumerate}  If separate telescopes are available for interferometry (as in the case of the Auxiliary Telescopes of the VLTI), the impact on the single-telescope-use of the spectrograph would be minimal, as it could be used in the interferometric mode during times when other instruments are scheduled for use with the main telescope.  The outline of this article is as follows. In Section~\\ref{sect:science} the scientific motivation for the proposed setup will be discussed.  Section~\\ref{instrument} will give an overview of the challenges in designing and building the proposed instrument, and their possible solutions. In Section~\\ref{sect:test_case:_vlti_uves} we will give more information about the proposed combination of the VLTI with the UVES spectrograph, and its expected performance. Partial and preliminary results of this work have been presented in \\cite{quirrenbach2008}. Section \\ref{sect:other_interferometer-spectrograph_pairings} highlights the main points for other possible interferometer-spectrograph pairings, and Section~\\ref{Conclusion} gives our conclusions. ",
        "conclusions": "\\label{Conclusion}  This article shows that the combination of high-resolution spectroscopy with long-baseline interferometry gives access to hitherto unobservable properties of stellar surfaces and circumstellar matter.  Furthermore the article shows how this combination can be achieved without building a major new instrument; only a beam combiner and a fiber link are needed. No instrument components are required which would be expensive or time-consuming to make. The use of external fringe tracking is essential for an instrument like this as it enables long integration times. Time-varying longitudinal dispersion could severely limit the possibilities of such an instrument. Dispersion compensation techniques are investigated in this article and a solution is presented which allows integration times up to a few minutes.  The implementation of this approach is shown for the example combination UVES-VLTI. The resulting instrument would differ from other instruments or efforts taken to achieve high spatial and spectral resolution. It would offer spectral resolution nearly a factor 2 higher than for any other interferometric instrument, and over a wide spectral range of a few hundred rather than a few tens of nanometers. It offers the same possibilities to an astronomer as a high-resolution Echelle spectrograph behind a single telescope does, plus the high spatial resolution due to the interferometer, albeit only for bright targets. However it is important to realize that this concept, the combination of two existing instruments to create a new one, is not limited to a specific location or instrument, but rather is an approach which can be followed at different observatories in both hemispheres.  It is worth pointing out that the measurements taken with the proposed instruments are differential in nature (e.g. change of visibility amplitude and visibility phase over spectral lines), allowing many interferometric calibration problems to be circumvented. Therefore science relying on absolute V$^{2}$ measurements would need additional data form other instruments. However as noted in Section~\\ref{sect:science}, differential measurements contain a wealth of astronomical information in the optical/IR regime."
    },
    "0911/0911.1533_arXiv.txt": {
        "abstract": "We apply ionization balance and MHD calculations to investigate whether magnetic activity moderated by recombination on dust grains can account for the mass accretion rates and the mid-infrared spectra and variability of protostellar disks.  The MHD calculations use the stratified shearing-box approach and include grain settling and the feedback from the changing dust abundance on the resistivity of the gas.  The two-decade spread in accretion rates among Solar-mass T~Tauri stars is too large to result solely from variations in the grain size and stellar X-ray luminosity, but can plausibly be produced by varying these parameters together with the disk magnetic flux.  The diverse shapes and strengths of the mid-infrared silicate bands can come from the coupling of grain settling to the distribution of the magneto-rotational turbulence, through the following three effects.  First, recombination on grains 1~$\\mu$m or smaller yields a magnetically-inactive dead zone extending more than two scale heights from the midplane, while turbulent motions in the magnetically-active disk atmosphere overshoot the dead zone boundary by only about one scale height. Second, grains deep in the dead zone oscillate vertically in wave motions driven by the turbulent layer above, but on average settle at the rates found in laminar flow, so that the interior of the dead zone is a particle sink and the disk atmosphere will become dust-depleted unless resupplied from elsewhere.  Third, with sufficient depletion, the dead zone is thinner and mixing dredges grains off the midplane.  The last of these processes enables evolutionary signatures such as the degree of settling to sometimes decrease with age.  The MHD results also show that the magnetic activity intermittently lifts clouds of small grains into the atmosphere.  Consequently the photosphere height changes by up to one-third over timescales of a few orbits, while the extinction along lines of sight grazing the disk surface varies by factors of two over times down to a tenth of an orbit.  We suggest that the changing shadows cast by the dust clouds on the outer disk are a cause of the daily to monthly mid-infrared variability found in some young stars. ",
        "introduction": "}  Concentration of the primordial solid material is a fundamental requirement for planet formation.  The interstellar medium contains about 1\\% solids by mass in the form of dust grains, and the Sun has a similar abundance of heavy elements, while the terrestrial planets consist almost entirely of refractory material.  Even the most gas-rich Solar system planet, Jupiter, has between 1.5 and 6~times the Solar abundance of heavy elements \\citep{sg04}.  A basic concentration process operating in the Solar nebula and other protostellar disks is the settling of dust particles due to the component of the stellar gravity perpendicular to the disk plane.  The grains settle at different speeds depending on their ratio of mass to area, leading to mutual collisions.  If the disk gas is laminar and the initially sub-micron particles readily stick together, then most of the solid material settles into a thin midplane layer within a few tens of thousands of years at 1 to 10~AU \\citep{w80,nn81,ns86}.  Such rapid loss of the grains conflicts with observations of scattered starlight showing some dust remains suspended in the atmospheres of million-year-old protostellar disks \\citep{bs96,sm03,ws04}.  However mid-infrared colors suggest the dust abundance is reduced near the disk photosphere, consistent with the incorporation of some material into larger bodies in the interior \\citep{fh06}.  The smaller particles can remain aloft if stirred by turbulence in the gas \\citep{dd04b}.  Turbulence is also commonly invoked to drive the observed flow of material onto the central star.  In particular, the turbulence resulting from the magneto-rotational instability \\citep[MRI;][]{bh91,bh98} can provide the necessary outward transfer of orbital angular momentum.  However, the MRI requires a sufficient level of ionization for coupling the disk gas to the magnetic fields. The weak internal heating means thermal ionization is effective only very near the star, while the interstellar cosmic rays and stellar X-rays ionize only the top and bottom surfaces of the disk, failing to penetrate the interior \\citep{g96,ig99}.  Furthermore, the recombination of the free charges on grain surfaces is so efficient that a small mass fraction of sub-micron particles is sufficient to shut off MRI turbulence in much of the region where the planets formed \\citep{sm00}.  As a result, protostellar disks consist of a laminar dead zone sandwiched between two turbulent active layers.  The dead zone extends from 0.1 to 15~AU in the minimum-mass Solar nebula, given micron-sized grains and ionization by the median stellar X-ray luminosity \\citep{td09}.  This paper deals with the consequences of the active and dead layers for dust settling and disk evolution.  We investigate how the gas flows govern the distribution of the grains, and how the grains in turn affect the dead zone size, using the methods outlined in \\S\\ref{sec:methods} to make ionization balance and 3-D MHD calculations of small patches of the disk around a T~Tauri star. Since a wealth of data is now available on the abundance, size and composition of dust in the surface layers of protostellar disks overlying the regions where planets are likely to form \\citep{fh06,sw06,bh08,wl09}, we examine the observational signatures of the movements of the grains, seeking to understand the following issues. \\begin{enumerate} \\item What is the cause of the large measured range in accretion rate at a given stellar mass (\\S\\ref{sec:massflow})?  Solar-mass T~Tauri stars grow at rates between about $10^{-9}$ and $10^{-7}$~M$_\\odot$ \\citep[compiled by][]{ml05}, while simple magnetic accretion models have a fixed column of accreting gas set by the penetration depth of the ionizing radiation, implying a unique mass flow rate \\citep{hd06}. \\item Why do T~Tauri stars show such diverse mid-infrared spectra (\\S\\ref{sec:movements})?  The variety in the silicate band strengths and shapes indicates wide ranges in the abundances of small and crystalline particles.  High crystalline mass fractions appear more often in systems having colors consistent with grains concentrated near the midplane \\citep{fh06,wl09}. \\item What are the origins of the variability over intervals of days to years in the 10-$\\mu$m silicate feature \\citep{wh00,ww04,bl09} and nearby continuum \\citep{la08,mf09} (\\S\\ref{sec:variable})?  The shorter variation timescales are substantially faster than the orbital periods at the AU radii where the feature forms. \\end{enumerate} A summary, discussion and conclusions are in \\S\\ref{sec:conc}. ",
        "conclusions": "}  We have made resistive MHD calculations of a patch of disk in orbit around a young star, treating simultaneously for the first time three processes that together regulate the distribution of magneto-rotational turbulence.  First, the dust drifts through the gas under the gravity and gas drag forces.  Second, the dust alters the resistivity of the gas, and third, the resistivity controls the evolution of the magnetic fields driving the gas flows.  The three processes combine in quite different ways in the disk atmosphere and interior.  In the atmosphere, the high flux of ionizing radiation from the star allows good coupling between gas and magnetic fields even when the dust abundance is high.  Intermittency in the turbulence and in the distribution of dust result from fluctuations in the magnetic field strength.  During intervals of weak fields, the magnetic forces drive turbulence that carries dust particles high in the atmosphere.  The turbulent mixing of the grains is faster than either the settling or the outflows resulting from the open boundaries.  When the fields are stronger, their large pressure prevents magneto-rotational turbulence and reduces the gas density, so that the particles settle quickly.  At the same time in the dead zone near the disk midplane, recombination on grain surfaces reduces the ionization fraction below the level needed to sustain turbulence.  Grains settle at the laminar rate, with superimposed vertical oscillations due to the propagation of waves excited in the turbulent layers.  Particles falling into the laminar layer do not return to the disk atmosphere and will eventually concentrate near the midplane, in a layer whose thickness is not less than the minimum imposed by the vertical oscillations.  A dead zone enriched in solids may provide a favorable environment for the growth of larger bodies.  Many T~Tauri systems show a combination of 10-$\\mu$m silicate emission and shallow millimeter spectral slope that is consistent with micron-sized grains in the disk atmosphere and millimeter-sized or larger particles in the interior, though it should be noted that the two wavelengths probe different distances from the star \\citep{dc06,pp08}.  The turbulent and dead zones exchange dust and gas when turbulent motions overshoot the mutual boundary.  The mixing extends about one scale height into the dead zone, as shown by the penetration of the dust-enriched material in figures~\\ref{fig:xgvsz1um} to~\\ref{fig:xgvsz100um}.  The motions can perhaps be described in terms of Alfv\\'en waves driven by the breakup of channel flows in the disk atmosphere \\citep{si09}.  The net effect is a transfer of dust to the dead zone, because settling proceeds fastest in the disk atmosphere.  During the early stages of the concentration of the solid material, turbulence thus speeds up settling.  The turbulent layer will eventually become dust-free unless replenished through radial transport within the disk or rain-out from an outflow or from the circumstellar envelope.  Well-mixed micron or sub-micron grains with the interstellar mass fraction produce a dead zone reaching more than two scale heights from the midplane at 5~AU in the minimum-mass solar nebula.  However, grains 10~$\\mu$m or larger yield a dead zone ending below one scale height.  Turbulent mixing then penetrates to the midplane, consistent with the finding by \\cite{fp06} that diffusion adequately describes the effects of the gas motions on particles in a dead zone with a small vertical extent.  While we were able to run the case with 1-$\\mu$m grains only about 60~orbits, the continuing loss of dust into the dead zone without resupply will ultimately yield a depletion factor in the surface layers exceeding the value of about ten needed for mixing to reach the midplane.  The rate at which mass flows through the patch of disk toward the star is the product of the turbulent layer depth with the mean accretion stress.  Both these quantities increase with the grain size, the ionizing flux from the star, and the magnetic flux threading the disk. Depletion of the dust in the surface layers will have a similar effect to an increase in the grain size.  The two-decade spread in accretion rate among T~Tauri stars cannot easily be produced by grain growth and the observed spread in stellar X-ray output, but can be understood if in addition the objects have a modest range of disk magnetic field strengths.  The mass accretion rate resulting from the magnetic stresses will vary with radius, leading to the build-up of material in some annuli \\citep{al01,zh09}, an effect not captured in our shearing-box calculations.  Apparently the stellar X-rays both increase the disk accretion rate within 5~AU (figure~\\ref{fig:activeht}) and increase the rate at which material escapes from the disk into a photoevaporating wind at 10-40~AU \\citep{de09}.  Furthermore the ionization of the inner disk could contribute to a wind driven by magnetic activity \\citep{si09}. All these effects tend to reduce the surface density of gas in the planet-forming region faster around stars with high X-ray luminosities.  Whether it is possible in this way to account for the tendency of non-accreting T~Tauri stars to have higher X-ray luminosities \\citep{ns95,sn01,fd03,pk05,tg07} or for an inverse correlation between X-ray luminosity and stellar accretion rate \\citep{de09} remains unclear.  We suggest that the diversity in the mid-infrared silicate bands among classical T~Tauri stars is a consequence of the concentration of solid material associated with planet formation.  The growth of larger bodies requires a loss of dust, leading generally to a declining recombination cross-section.  The smallest grains provide most of the recombination when the particles have a size spectrum appropriate for the interstellar medium \\citep{sm00}, and will dominate still more in the disk atmosphere, where settling removes the larger grains.  The disruption of aggregates through collisions (\\S\\ref{sec:stirring}) thus slows the decline.  However, unless the grains making up the aggregates also fragment so that the minimum particle size becomes smaller, the loss of cross-section is likely to continue, since even 1-$\\mu$m particles show appreciable settling in the turbulent disk atmosphere.  The decreasing recombination rates will allow the active layers to expand into the disk interior, and if the mass column is not too great, turbulent mixing will eventually reach the midplane, returning the second-generation debris from planet formation to the atmosphere.  Questions for future investigation include when and where the mixing reaches the midplane, and whether the observed correlation between the degree of settling and the crystalline mass fraction \\citep{wl09} can be explained by the return to the atmosphere of minerals formed near the midplane at high temperature and pressure in planetesimal interiors, impacts, or shocks.  Any such correlation will be blurred by the fact that the degree of dust depletion required for mixing to reach the midplane depends on the column of gas, the X-ray luminosity of the star and the magnetic flux delivered to the disk from the surroundings.  The mid-infrared variability of protostellar disks has been proposed to arise from the changing illumination of dust in the disk atmosphere.  Explaining timescales shorter than a week is a challenge, since the dust in the disks around young Solar-mass stars extends inward only to about 0.1~AU \\citep{mc03,ab05,a08}, where the orbital period is 12~days.  The results of the MHD calculations demonstrate that magnetic activity moves dust from place to place in the disk atmosphere on several timescales.  The optical depth along rays grazing the surface of the patch of disk frequently varies by a factor two within a tenth of an orbit as regions of enhanced dust density pass by.  Over longer time intervals of a few orbits, the photosphere moves up and down by as much as a scale height due to changes in the balance between settling and stirring.  Clearly the effects of the magnetic activity on the dust distribution should be considered when attempting to understand the infrared variability.  Properly accounting for the shadows cast on the disk surface requires models of the inner edge of the dust distribution, perhaps informed by the time-steady picture that has been developed for Herbig~Ae stars \\citep{dd01,in05}.  Our results suggest that magneto-rotational turbulence near the dust sublimation radius can cast shadows that vary over timescales as short as a day."
    },
    "0911/0911.3536_arXiv.txt": {
        "abstract": "It is widely believed that the cosmological redshift is not a Doppler shift. However, Bunn \\& Hogg have recently pointed out that to settle properly this problem, one has to transport parallely the velocity four-vector of a distant galaxy to the observer's position. Performing such a transport along the null geodesic of photons arriving from the galaxy, they found that the cosmological redshift is purely kinematic. Here we argue that one should rather transport the velocity four-vector along the geodesic connecting the points of intersection of the world-lines of the galaxy and the observer with the hypersurface of constant {\\em cosmic time}. We find that the resulting relation between the transported velocity and the redshift of arriving photons is {\\em not\\/} given by a relativistic Doppler formula. Instead, for small redshifts it coincides with the well known non-relativistic decomposition of the redshift into a Doppler (kinematic) component and a gravitational one. We perform such a decomposition for arbitrary large redshifts and derive a formula for the kinematic component of the cosmological redshift, valid for any FLRW cosmology. In particular, in a universe with $\\Omega_\\mathrm{m} = 0.24$ and $\\Omega_\\Lambda = 0.76$, a quasar at a redshift $6$, at the time of emission of photons reaching us today had the recession velocity $v = 0.997c$. This can be contrasted with $v = 0.96c$, had the redshift been entirely kinematic. Thus, for recession velocities of such high-redshift sources, the effect of deceleration of the early Universe clearly prevails over the effect of its relatively recent acceleration. Last but not least, we show that the so-called {\\em proper\\/} recession velocities of galaxies, commonly used in cosmology, are in fact radial components of the galaxies' four-velocity vectors. As such, they can indeed attain superluminal values, but should not be regarded as real velocities. ",
        "introduction": "\\label{sec:intro} A standard interpretation of the cosmological redshift in the framework of the Friedman-Lema{\\^\\i}tre-Robertson-Walker (FLRW) models is that it is an effect of the expansion of the Universe. This interpretation is obviously correct since $1 + z = a(\\tO)/a(\\tE)$, where $z$ is the value of the redshift, $a(t)$ is the scale factor of the Universe and $\\tE$ and $\\tO$ are respectively the times of emission and observation of a sent photon. In semi-popular literature (e.g. Kaufmann \\& Freedman 1999; Franknoi, Morrison \\& Wolff 2004; Seeds 2007), but also in professional (e.g. Harrison 2000; Abramowicz \\etal\\ 2007), one can often find statements that distant galaxies are `really' at rest and the observed redshift is caused by the `expansion of space'. According to other authors (e.g. Peacock 1999; Whiting 2004; Chodorowski 2007a,b; Bunn \\& Hogg 2009, hereafter BH9), such statements are misleading and cause misunderstandings about the cosmological expansion. A presentation of the ongoing debate in the literature on this issue is beyond the scope of the present work; a (possibly non-exhaustive) list of papers includes Davis \\& Lineweaver (2001), Davis \\& Lineweaver (2004), Barnes \\etal\\ (2006), Francis \\etal\\ (2007), Lewis \\etal\\ (2007), Lewis \\etal\\ (2008), Gr{\\o}n \\& Elgar{\\o}y (2007), Peacock (2008), Abramowicz \\etal\\ (2009), Chodorowski (2008), Cook \\& Burns (2009).  On the other hand, there is broad agreement that the cosmological redshift is not a pure Doppler shift; the gravitational field must also generate a gravitational shift. Gravity can be neglected only locally, in the local inertial frame (LIF) of an observer. It turns out that for small redshifts, the cosmological redshift can be decomposed into a Doppler shift and a Newtonian gravitational one (Bondi 1947). The latter is a shift induced by the Newtonian gravitational potential. Can the cosmological redshift be decomposed into a Doppler shift and a gravitational shift (not necessarily Newtonian) for an {\\em arbitrary\\/} value of the redshift? This is the question which we want to deal with in this Paper. Formally, the answer is no. There is no invariant definition of the recession velocity of a distant galaxy in general relativity (GR). This velocity is a relative velocity of the galaxy and the observer, and in curved spacetime there is no unique way to compare vectors at widely separated points. A natural way to define the recession velocity is to {\\em transport parallely\\/} the velocity four-vector of the distant galaxy to the observer, but the result will depend on the chosen path. (This is just the definition of curvature.) In practice, however, as a `preferred' path one can choose a {\\em geodesic\\/} connecting the galaxy and the observer. Moreover, in FLRW models there is a natural foliation of spacetime, into space-like hypersurfaces of constant {\\em cosmic time}. In our Paper, as the geodesic we will adopt the geodesic lying on such a hypersurface, i.e.\\ connecting the points of intersection of the world-lines of the galaxy and the observer with the hypersurface of constant cosmic time.  In a seminal study of the cosmological redshift, BH9, following Synge~(1960) and Narlikar~(1994), adopted another geodesic for the parallel transport: the null geodesic along which the photon is travelling from the source to the observer. This approach results in one `effective' velocity, while we think it is important to make a distinction between the velocity {\\em at the time of emission\\/} and the velocity {\\em at the time of observation}. These two velocities are obtained by transporting parallely the velocity four-vector of the source respectively on the hypersurface of constant $\\tE$ and $\\tO$. Not surprisingly, our result differs from that obtained by BH9. However, unlike theirs, ours correctly reproduces the small-redshift decomposition of the cosmological redshift into a Doppler component and a gravitational component, mentioned above. Specifically, our decomposition coincides with that of Bondi (1947) for isotropic {\\em and\\/} homogeneous matter distribution.  This Paper is organized as follows. In Section~\\ref{sec:transport} we transport parallely the velocity four-vector of a distant galaxy to the observer. From the transported vector we calculate the recession velocity of the galaxy, which turns out to depend on the galaxy's comoving distance and the assumed background cosmological model. In Section~\\ref{sec:specific} we find specific relations between the cosmological redshift and its Dopplerian component for two particularly simple cosmological models: the empty model and the Einstein-de Sitter model. In Section~\\ref{sec:PoE}, using only the Principle of Equivalence, we derive the small-redshift decomposition of the cosmological redshift and find it to be identical with that obtained in Section~\\ref{sec:specific} using generally relativistic approach. In Section~\\ref{sec:rec} we calculate the recession velocities in the two models mentioned above, as well as for the currently favoured, flat non-zero $\\Lambda$ model with $\\Omega_{\\mathrm{m}} = 0.24$. We find also that the velocities are {\\em subluminal\\/} in {\\em all\\/} FLRW cosmological models and for {\\em all\\/} values of the redshift. A comparison of the results obtained in this Paper with some earlier works on the subject is given in Section~\\ref{sec:disc}. Summary is presented in Section~\\ref{sec:summ}. ",
        "conclusions": "recession velocities of all galaxies inside our particle horizon are subluminal. This fact is perhaps more fundamental than the specific value of recession velocity for a given redshift."
    },
    "0911/0911.4120_arXiv.txt": {
        "abstract": "{Earlier work has suggested that large-scale dynamos can reach and maintain equipartition field strengths on a dynamical time scale only if magnetic helicity of the fluctuating field can be shed from the domain through open boundaries.}% {Our aim is to test this scenario in convection-driven dynamos by comparing results for open and closed boundary conditions.} {Three-dimensional numerical simulations of turbulent compressible convection with shear and rotation are used to study the effects of boundary conditions on the excitation and saturation of large-scale dynamos. Open (vertical-field) and closed (perfect-conductor) boundary conditions are used for the magnetic field. The shear flow is such that the contours of shear are vertical, crossing the outer surface, and are thus ideally suited for driving a shear-induced magnetic helicity flux.}% {We find that for given shear and rotation rate, the growth rate of the magnetic field is larger if open boundary conditions are used. The growth rate first increases for small magnetic Reynolds number, $\\Rm$, but then levels off at an approximately constant value for intermediate values of $\\Rm$. For large enough $\\Rm$, a small-scale dynamo is excited and the growth rate of the field in this regime increases as $\\Rm^{1/2}$. Regarding the nonlinear regime, the saturation level of the energy of the total magnetic field is independent of $\\Rm$ when open boundaries are used. In the case of perfect conductor boundaries, the saturation level first increases as a function of $\\Rm$, but then decreases proportional to $\\Rm^{-1}$ for $\\Rm\\ga 30$, indicative of catastrophic quenching. These results suggest that the shear-induced magnetic helicity flux is efficient in alleviating catastrophic quenching when open boundaries are used. The horizontally averaged mean field is still weakly decreasing as a function of $\\Rm$ even for open boundaries.}% {}% ",
        "introduction": "The classical view of the effect of shear in large-scale dynamos is that it promotes the generation of magnetic fields, and that it is instrumental in exciting oscillatory solutions (e.g.\\ Parker \\cite{P55}; Steenbeck \\& Krause \\cite{SK69}). While these aspects remain unchanged, a much more subtle effect of shear has been discovered in studies of the saturation mechanism of the large-scale dynamo. If the saturation of the dynamo occurs due to magnetic helicity conservation, a fully periodic or closed system experiences catastrophic quenching with the energy of the large-scale magnetic field decreasing in inverse proportion to the magnetic Reynolds number (Vainshtein \\& Cattaneo \\cite{VC92}; Cattaneo \\& Hughes \\cite{CH96}; Brandenburg \\cite{B01}). A way to avoid catastrophic quenching is to allow magnetic helicity fluxes to escape from the system through open boundaries, thus allowing the large-scale fields to saturate at equipartition strength in a dynamical time scale (Blackman \\& Field \\cite{BF00}; Kleeorin et al.\\ \\cite{KMRS00}). One of the most promising mechanisms, introduced by Vishniac \\& Cho (\\cite{VC01}), operates by driving a flux of magnetic helicity along the isocontours of shear. Forced turbulence simulations with shear have confirmed that the quenching of the $\\alpha$-effect is up to 30 times weaker when closed boundaries are replaced by open ones (Brandenburg \\& Sandin \\cite{BS04}).  Recent numerical experiments with forced turbulence (Brandenburg \\cite{B05}; Yousef et al.\\ \\cite{YHSea08a,YHSea08b}; Brandenburg et al.\\ \\cite{BRRK08}) and convection (K\\\"apyl\\\"a et al.\\ \\cite{KKB08}, hereafter Paper~I; Hughes \\& Proctor \\cite{HP09}; see also K\\\"apyl\\\"a et al.\\ \\cite{KKB10}) with imposed large-scale shear flows have established the existence of a robust large-scale dynamo. The origin of this dynamo is still under debate (shear--current versus incoherent $\\alpha$--shear effects). This is discussed in detail elsewhere (e.g.\\ Brandenburg et al.\\ \\cite{BRRK08}; K\\\"apyl\\\"a et al.\\ \\cite{KKB09b}). In forced turbulence simulations with fully periodic boundaries, that do not allow magnetic helicity fluxes, large-scale magnetic fields show slow saturation on a diffusive timescale (Brandenburg \\cite{B01}). Earlier results from convection also suggest that no appreciable large-scale magnetic fields are seen if magnetic he\\-li\\-ci\\-ty fluxes are suppressed (Paper~I).  In most of the simulations of Paper~I the isocontours of shear were vertical and thus intersecting the boundaries. However, the boundary conditions (bc) imposed on the magnetic field on the vertical borders would either allow (open vertical-field bc) or inhibit (closed perfect-conductor bc) a net flux out of the domain. It was shown that when the flux is suppressed, only weak large-scale magnetic fields occur in the saturated state at large $\\Rm$. Conversely, allowing a flux produced a significant large-scale field. Furthermore, using a similar setup as in Tobias et al.\\ (\\cite{TCB08}), where the isocontours of shear are horizontal, a similar effect was found. Thus it would appear that the shear-induced flux plays a critical role in exciting a large-scale dynamo in turbulent convection. This has given rise to the impression that the dynamos seen in Paper~I are solely due to the magnetic helicity flux and thus inherently nonlinear (Hughes \\& Proctor \\cite{HP09}). In the present paper we argue that this interpretation is incorrect and that the lack of appreciable large-scale magnetic fields in some of our earlier cases with vanishing helicity flux is due to the fact that the critical dynamo number for the excitation of the dynamo is larger in that case. We also note that, more recently, large-scale dynamos have been found in turbulent convection without shear (K\\\"apyl\\\"a et al.\\ \\cite{KKB09a}), irrespective of the presence of magnetic helicity fluxes.  The origin of the dynamos reported in Paper~I was interpreted in the context of turbulent dynamo theory in K\\\"apyl\\\"a et al.\\ (\\cite{KKB09b}). An important feature of turbulent dynamos is that according to standard mean-field theory neither the turbulent transport coefficients nor the growth rate of the field should depend on the microscopic resistivity. The former prediction was confirmed by K\\\"apyl\\\"a et al.\\ (\\cite{KKB09b}), but the latter has not yet been studied in detail in direct simulations. This is one of the goals of the present paper.  Another aim is to study the saturation behaviour of the large-scale dynamo with both open and closed boundary conditions. Firstly, in homogeneous systems, i.e.\\ where the turbulent transport coefficients do not vary in space, full saturation is expected to occur on a slow diffusive time scale with closed boundaries whereas with open boundaries the saturation is expected to happen on a dynamical time scale if a helicity flux is present (Brandenburg \\cite{B01}). However, in inhomogeneous systems the situation is likely to be more complex (e.g.\\ Mitra et al.\\ \\cite{MCCTB10}). Furthermore, simulations of helically forced isotropic turbulence have shown that, in the absence of shear, the saturation level of the magnetic field scales as $\\meanv{B}^2/\\Beq^2\\propto\\Rm^{-1}$ for open boundaries (Fig.~2 of Brandenburg \\& Subramanian \\cite{BS05a}). When shear is present, the shear-induced magnetic helicity flux should alleviate the $\\Rm$-dependence of the saturation level. Preliminary results in Paper~I indicate that this might indeed be true, but the results are not fully conclusive because other parameters, such as the Rayleigh number and the fluid Reynolds number changed when $\\Rm$ was varied. In the present paper we perform a more detailed study where only $\\Rm$ is varied. ",
        "conclusions": "We have studied the effects of magnetic boundary conditions on the excitation and saturation of large-scale dynamos driven by turbulent convection, shear, and rotation by means of numerical simulations. We find that the critical magnetic Reynolds number is greater ($\\Rm\\approx5$) in the case of closed boundaries (perfect-conductor) compared to simulations with open (vertical-field) boundaries where the large-scale dynamo is excited already for $\\Rm\\approx1$. A similar result is obtained from mean-field models operating in the same parameter regime. Furthermore, the relevant dynamo number is roughly three times greater with closed boundaries. These effects are likely to explain the weak large-scale dynamo action seen in the perfect-conductor runs in Paper~I.  The measured growth rate of the magnetic field is independent of the microscopic resistivity when $\\Rm$ is sufficiently above the critical value. This is manifested by the approximately constant growth rate in the intermediate $\\Rm$ range in Fig.~\\ref{fig:pgr}. For $\\Rm\\ga30$, the small-scale dynamo is excited, and for $\\Rm\\ga60$ it becomes dominant. The growth rate of the magnetic field is then consistent with $\\lambda\\propto\\Rm^{1/2}$ scaling, which is in accordance with the results of Schekochihin et al.\\ (\\cite{Scheko04}) and Haugen et al.\\ (\\cite{{HBD04}}). The fact that the qualitative behaviour of the growth rate of the magnetic field is similar for both boundary conditions suggests that the origin of the large-scale dynamo is not likely to be a process that is essentially nonlinear, as suggested by Hughes \\& Proctor (\\cite{HP09}), but that it can be understood within the framework of classical kinematic mean-field theory.  In the saturated state the energy of the total magnetic field remains independent of $\\Rm$ for open boundaries and decreases as $\\Rm^{-1}$ for closed boundaries. The latter result is consistent with catastrophic quenching while the former result suggests that magnetic helicity fluxes are efficiently driven out of the system. On the other hand, the energy of the mean field, taken here to be represented by horizontal averages, decreases approximately as $\\Rm^{-0.25}$ for open and as $\\Rm^{-1.6}$ for closed boundaries. It is not yet clear why the mean fields tend to show a steeper decline than the total field. It is possible that this declining trend levels off at a higher $\\Rm$ and that the magnetic Reynolds numbers in our simulations are still not large enough (cf.\\ Brandenburg et al.\\ \\cite{BCC09}). A similarly weak $\\Rm$-dependence has been observed in the cycle period of $\\alpha$-shear dynamos with isotropically forced turbulence (K\\\"apyl\\\"a \\& Brandenburg \\cite{KB09}).  We find that, for intermediate values of $\\Rm$, the large-scale magnetic field saturates on a resistive time scale when closed boundaries are used. However, with the largest $\\Rm$ no clear signs of slow saturation are observed. Earlier results using perfect-conductor boundaries have shown that the mean field can evolve in steps (Brandenburg \\& Dobler \\cite{BD02}; see also Brandenburg et al.\\ \\cite{BKMMT07}) which are associated with a change of the large-scale magnetic field configuration. We have not seen such a behaviour in our current simulations but the existence of such events at a later stage cannot be ruled out. Another possible explanation is that there are magnetic helicity fluxes occurring inside the domain which arise from the spatial gradients of magnetic helicity (e.g.\\ Covas et al.\\ \\cite{CTTB98}; Kleeorin et al.\\ \\cite{KMRS00}; Mitra et al.\\ \\cite{MCCTB10}). However, a quantitative study of these effects requires more detailed knowledge of the helicity fluxes and possibly an anisotropic formulation of the magnetic $\\alpha$ effect. These issues merit further investigation and are beyond the scope of the present paper."
    },
    "0911/0911.1369_arXiv.txt": {
        "abstract": "We investigate a new theory of the origin of the irregular satellites of the giant planets: capture of one member of a $\\sim$100-km binary asteroid after tidal disruption.  The energy loss from disruption is sufficient for capture, but it cannot deliver the bodies directly to the observed orbits of the irregular satellites.  Instead, the long-lived capture orbits subsequently evolve inward due to interactions with a tenuous circumplanetary gas disk.  We focus on the capture by Jupiter, which, due to its large mass, provides the most stringent test of our model.  We investigate the possible fates of disrupted bodies, the differences between prograde and retrograde captures, and the effects of Callisto on captured objects.  We make an impulse approximation and discuss how it allows us to generalize capture results from equal-mass binaries to binaries with arbitrary mass ratios.  We find that at Jupiter, binaries offer an increase of a factor of $\\sim$10 in the capture rate of 100-km objects as compared to single bodies, for objects separated by tens of radii that approach the planet on relatively low-energy trajectories.  These bodies are at risk of collision with Callisto, but may be preserved by gas drag if their pericenters are raised quickly enough.  We conclude that our mechanism is as capable of producing large irregular satellites as previous suggestions, and it avoids several problems faced by alternative models. ",
        "introduction": "\\label{intro}  \\subsection{Previously suggested capture models}  \\label{models}  With discoveries accelerating in the last decade, we now know of over 150 satellites orbiting the giant planets.  About one-third of these are classified as regular, with nearly circular and planar orbits.  It is thought that these satellites are formed by accretion in circumplanetary disks.  The majority of the satellites, however, are irregular and follow distant, highly eccentric and inclined paths.  It is widely believed that irregular satellites originated in heliocentric orbits and were later captured by their planets, but the details of how this occurred are still uncertain.  At least seven different models have been proposed, involving dissipative forces, collisions, resonances, and three-body effects.  Each model has its own strengths and weaknesses.  In one long-standing theory, planetesimals are slowed as they punch through the gas disk surrounding a young, growing planet (Pollack \\etall, 1979).  For this mechanism to be efficient, the gas must be sufficiently dense to capture the planetesimals in one pass.  This is problematic, however, because if the gas disk does not rarefy substantially in $\\sim$100-1000 years, the orbits of the new satellites will decay inward, leading to collisions with the planet and its regular satellites.  Furthermore, the atmospheres of Uranus and Neptune have only a few Earth-masses of hydrogen and helium at present, so their gas disks could not have been as extensive or long-lived as those of Jupiter and Saturn.  A likely outcome of this model, then, is that satellite capture should have been different at Jupiter and Saturn than at Uranus and Neptune; however, current observational estimates suggest equal efficiencies (Jewitt \\& Sheppard, 2005).  With a model similar to that of Pollack \\etall, {\\'C}uk \\& Burns (2004a) found that Jupiter's largest irregular satellite, Himalia, would evolve inward to its current orbit in $10^4-10^6$ years.  This tenuous gas, however, may make capture difficult.  In another model, planetesimals are captured when the mass of the planet increases (Heppenheimer \\& Porco, 1977).  This mass growth causes the planet's escape velocity to increase, rendering a previously free planetesimal bound to the planet.  For this method to be effective, the planet's mass must increase substantially on $\\sim$100-1000-year timescales.  However, in most planet formation models (e.g. Pollack \\etall, 1996), giant planet growth is hypothesized to take place on timescales many orders of magnitude longer than required by this capture scenario.  Furthermore, Uranus and Neptune's gas deficiency implies that their growth was of very short duration.  Thus, our current understanding of planetary formation makes this model improbable.  The observation that the four giant planets contain approximately the same number of irregular satellites (accounting for observational biases; Jewitt \\& Sheppard, 2005) has led to a renewal of interest in capture theories that do not depend strongly on the planet's formation process.  In one such scenario, a planetesimal collides with a current satellite or another planetesimal in the vicinity of the planet, resulting in its capture (Colombo \\& Franklin, 1971).  Though collisions were certainly more common in the early Solar System than they are today, if they resulted in enough energy loss to permit capture, they would likely also have catastrophically disrupted the bodies.  Nevertheless, the fragments might then have become independent satellites.  A fourth suggestion involves the possible instability in the orbits of the outer planets early in the Solar System (e.g., by a 2:1 resonance crossing between Jupiter and Saturn).  Outlining the theory of the Nice model of Solar System evolution, Tsiganis \\etal (2005) have shown that such an event could cause Uranus and Neptune to have many close approaches with each other and with Jupiter and Saturn.  During these encounters, the influence of the massive interloping planet can cause planetesimals to be stabilized as satellites (Nesvorn{\\'y} \\etall, 2007).  This method is promising but has an important disadvantage in that Jupiter (and Saturn, to a lesser extent) sustains very few close encounters relative to the ice giants.  Thus the gas giants are inefficient at capturing satellites in this way (Nesvorn{\\'y} \\etall, 2007).  Astakhov \\etal (2003) examined low-energy orbits that linger near Jupiter and Saturn.  While these bodies are not permanently captured, the authors found that some of them were stable for thousands of years, long enough to allow a weak dissipative force such as gas drag to complete the capture process.  However, the overall percentage of temporary captures that do not escape is small, and many of these bodies are threatened by collision with the planets' large outer satellites (e.g., Callisto and Titan).  Agnor \\& Hamilton (2006a) examined the capture of Triton from an exchange reaction between a binary pair and Neptune.  Their motivation stemmed from the newly-discovered abundance of binaries in small-body populations.  Currently, it is estimated that binaries account for $\\sim$30\\% of Kuiper belt objects (KBOs) with inclinations $<$ 5$\\deg$, $\\sim$5\\% of the rest of the KBOs (Noll \\etall, 2008), and $\\sim$2\\% of large main belt asteroids (diameters $>$ 20 km; percentage increases for smaller objects; Merline \\etall, 2007).  In Agnor \\& Hamilton's capture model, a binary is tidally disrupted and one of its members, Triton, is captured as a satellite.  This process is most effective for large satellites like Triton, with radius 1350 km.  However, the largest of the other irregular satellites are more than 10 times smaller than Triton: Himalia at Jupiter is $\\sim$85 km in radius, Saturn's largest irregular, Phoebe, is $\\sim$110 km, Uranus's Sycorax is $\\sim$80 km, and Neptune's Halimede and Neso are only $\\sim$30 km each.  Capturing these satellites via binary exchange reactions would be significantly more difficult, as we will discuss further below.  Finally, Vokrouhlick{\\'y} \\etal (2008) examined binary exchange reactions during the first 100 Myr after an assumed Jupiter/Saturn 2:1 resonance crossing, using results of the Nice model (Tsiganis \\etall, 2005) to guide their initial conditions.  Because planetesimal speeds relative to the planets are high after the scattering phase of the Nice model, they found that captures from binaries during that time do not match current orbital parameters and occur too infrequently to account for today's populations.   \\subsection{Our model: Capture from 100-km binaries}  \\label{ourmodel}  All of the above models have promising aspects coupled with important limitations.  In this work, we seek to combine the best features of several models into a viable capture scenario.  In particular, we examine binaries (as in Agnor \\& Hamilton, 2006a and Vokrouhlick{\\'y} \\etall, 2008) as a way to augment capture from low-velocity orbits resulting from three-body interactions like those studied by Astakhov \\etal (2003).  While Vokrouhlick{\\'y} \\etal (2008) studied exchange reactions in the context of an assumed initial planetesimal population, we focus on assessing the viability of the mechanism itself.  Our goal is to determine how various parameters of the model affect its plausibility.  We examine its viability at Jupiter, as a number of the above models suggest that capturing at the largest gas giant is especially difficult.  As the largest of the existing irregular satellites are $\\sim$80-110 km, capture of objects in this size-range is particularly interesting. Since it is likely that the irregular satellite population contains collisional families (Nesvorn{\\'y} \\etall, 2003; Sheppard \\& Jewitt, 2003), it may be the case that only the largest objects were captured, while the smaller satellites formed later, via collisions.  For this reason, we focus our investigation on capturing the $\\sim$100-km progenitors.  In order to stabilize and shrink the resulting capture orbits, a dissipation source is required; we suggest a tenuous version of the gas drag originally proposed by Pollack \\etal (1979).  Two of Jupiter's irregular satellites, Pasiphae and Sinope, as well as Saturn's satellite, Siarnaq, and Uranus' Stephano, are found in resonances that seem to require just such a weak dissipative force (Whipple \\& Shelus, 1993; Saha \\& Tremaine, 1993; {\\'C}uk \\& Burns, 2004b).  Furthermore, a tenuous circumplanetary disk is consistent with current theories of late-stage planetary formation.  Jupiter's massive gaseous envelope of hydrogen and helium necessitates that it formed in the Solar System's circumstellar gas disk.  Before the end of its accretion, Jupiter was likely able to open a gap in the local density distribution of the gas (for a review, see e.g. Papaloizou \\etall, 2008).  After gap opening, gas continues to leak into the planet's Hill sphere through the $L_1$ and $L_2$ points, but at a rate much reduced in comparison to the previous epochs.  A tenuous circumplanetary gas disk results (e.g., Lubow \\etall, 1999; D'Angelo \\etall, 2003; Bate \\etall, 2003), from which material may condense and regular satellites may accrete near the planet (e.g., Canup \\& Ward, 2002; Mosquiera \\& Estrada, 2003).  In {\\'C}uk \\& Burns' study (2004a) of the Himalia progenitor's orbital evolution, they considered circumjovian nebular conditions consistent with hydrodynamical simulations of Jupiter's gap opening in a circumstellar gas disk (e.g., Lubow \\etall, 1999) and found that the post-capture timescale for evolving this progenitor to its present orbit to be roughly in the range of $10^4-10^6$ years.  This is similar to the timescale in which extrasolar circumstellar disks transition from optically thick to thin ($\\sim10^5$ years; Skrutskie \\etall, 1990; Silverstone \\etall, 2006; Cieza \\etall, 2007).  The similarity of timescales suggests that satellites captured at the onset of disk dispersal have a good chance of experiencing stabilizing orbital evolution while also avoiding collision with the planet.  The timescale for binary capture is very short compared to evolution timescales from a tenuous gas disk.  Therefore, we focus our study first on characterizing the effectiveness of binary capture in the absence of gas.  In the following sections, we critically evaluate our model for capturing irregular satellites from low-mass ($\\sim$100-km) binaries.  We begin with a closer examination of the three-body capture process and then explore parameter space with a large suite of numerical simulations.  We then discuss the ability of gas drag to stabilize post-capture orbits in Section~\\ref{survivability}. ",
        "conclusions": "The new model discussed here, capture from low-mass binaries with subsequent orbital evolution, has both significant advantages and disadvantages in comparison to other suggested models.  One important advantage is that capture is viable at Jupiter and Saturn, unlike three of the models discussed in Section~\\ref{intro}. For example, Agnor \\& Hamilton's (2006) direct three-body capture model works only for very large bodies in the gas giants' high-approach-speed environments.  Also, the theory of Nesvorn{\\'y} \\etal (2003) requires close approaches among the giant planets, of which Jupiter has very few in the Nice model scenario.  Saturn encounters its outer neighbors more frequently than Jupiter, but it still suffers from few close approaches overall in this model.  In Vokrouhlick{\\'y} \\etal (2007), the capture statistics from planet-binary encounters are low for all planets, but especially Jupiter and Saturn.  This is primarily because of the high relative velocities assumed in their model.  Though these latter two capture mechanisms in their current forms cannot explain the gas giants' irregular satellites, they are worth further study, perhaps in the context of altered versions of the Nice model or other early Solar System models.  Our model also has an important advantage over that of Pollack \\etall, 1979, in that our capture scenario allows for a much weaker gas, since the gas here is needed for shrinking the orbits, not capturing satellites. Also, the gas in our model can persist for much longer than that in Pollack \\etall, as a weaker gas does not have the problem of quickly destroying the captured bodies.  Furthermore, various groups (e.g., Canup \\& Ward, 2002) have proposed that the Galilean satellites formed from a tenuous gas; if true, the gas was most likely even thinner when the irregular satellites were captured.  Thus it is important that our capture mechanism does not rely on a dense gas.  To directly compare our mechanism with that of Astakhov \\etal (2003), we need to consider both relative capture rates (and their survivability) and the prevalence of binaries vs. single bodies.  In our simulations, we find that binary capture can provide a significant advantage over capturing from populations of single bodies for binaries with particular characteristics: high enough masses ($\\gtrsim$ 100 km), optimal separations ($\\sim$10-20 $R_B$), and low incoming energies (corresponding to Jacobi constants $\\gtrsim$ 3.02 and mostly prograde encounters).  However, like Astakhov \\etal (2003), we also find that the probability of captures (from either binaries or single bodies) of avoiding collisions with Callisto is low for readily captured progenitors.  In Section~\\ref{survivability}, we discussed that this problem can be alleviated by altering the capture orbits with the surrounding gas or by capturing binaries with larger-than-optimal separations that do not lead to Callisto-crossing orbits.  So, how common were easily-captured binaries early in Solar System history?  This question is difficult to assess.  Observational surveys of the current population of the cold, classical Kuiper belt (i.e. objects with modest inclinations and eccentricities) find a 30\\% binary fraction among bodies larger than 100 km (Noll \\etall, 2008), many with nearly equal mass components.  Also, recent studies of planetesimal formation have suggested that large, $\\gtrsim$ 100-km bodies may form quickly in the gaseous proto-planetary disk, providing the building blocks of subsequent planet formation (Johansen \\etall, 2007; Cuzzi \\etall, 2008).  Binary formation is likely to be contemporaneous with the formation of these bodies (Nesvorn{\\'y}, 2008).  Further, Morbidelli \\etal (2009) have shown that the size-frequency distribution of asteroids in the main belt is consistent with large initial planetesimals, at least $\\sim$100-km in size.  Together these results indicate that the $\\sim$100-km binary objects considered here may have been quite common as the very last portions of the Solar System's gas disk were being depleted.  Finally, the known irregular satellite population numbers $<$ 100, and many of these are probably members of families -- thus, we need only produce at most a few dozen captures.  While accounting for the origin of such a small population is difficult, our simulations show that it is likely binaries played a role.  We conclude by offering our model as a new idea that alleviates many but not all of the problems faced by previous models, but acknowledge that the without a detailed understanding of the initial population including binary statistics, a firm conclusion is not possible.   \\newpage"
    },
    "0911/0911.5476_arXiv.txt": {
        "abstract": "\\noindent We present new brightness and magnetic images of the weak-line T~Tauri star V410~Tau, made using data from the NARVAL spectropolarimeter at T\\'{e}lescope Bernard Lyot (TBL). The brightness image shows a large polar spot and significant spot coverage at lower latitudes. The magnetic maps show a field that is predominantly dipolar and non-axisymmetric with a strong azimuthal component. The field is 50\\% poloidal and 50\\% toroidal, and there is very little differential rotation apparent from the magnetic images.  A photometric monitoring campaign on this star has previously revealed V--band variability of up to 0.6 magnitudes but in 2009 the lightcurve is much flatter. The Doppler image presented here is consistent with this low variability. Calculating the flux predicted by the mapped spot distribution gives an peak-to-peak variability of 0.04 magnitudes. The reduction in the amplitude of the lightcurve, compared with previous observations, appears to be related to a change in the distribution of the spots, rather than the number or area.  This paper is the first from a Zeeman-Doppler imaging campaign being carried out on V410~Tau between 2009--2012 at TBL. During this time it is expected that the lightcurve will return to a high amplitude state, allowing us to ascertain whether the photometric changes are accompanied by a change in the magnetic field topology. ",
        "introduction": "T~Tauri stars (TTS) are pre-main sequence late-type stars. As they contract towards the main sequence the increasing density in their interiors will lead to the development of a radiative core, if the star is sufficiently massive. A point of uncertainty concerns the evolution of their dynamos. In fully convective stars, or stars with small radiative cores, we expect the dynamo to look fundamentally different to that of the Sun. We presume that the solar dynamo is concentrated in the boundary between the convective and radiative layers, known as the tachocline. As the radiative core in the young star grows the dynamo should change from being distributed throughout the convective zone to being confined to the tachocline. Observations of young stars with a range of ages, masses and rotation rates will help to illuminate this process.  V410~Tau is one of the most well-observed TTS. It is a weak-line T~Tauri star (wTTS) which has already dissipated its disk, hence it does not display signatures of accretion characteristic of classical TTS (cTTS). It has a \\vsini\\ of $74\\pm3$\\kms\\ and a period of 1.872~days \\citep{stelzer03}, and is therefore an example of the young, rapidly rotating stars which are characterised by strong, large scale magnetic fields. Other, more evolved, examples include the 10~Myr old TWA~6 \\citep{skelly1}, Speedy~Mic \\citep[20 Myr,][]{barnes01,barnes05c, duns06} and AB~Dor \\citep[50~Myr,][]{cc02b,donati03b,jeffers07}.  The strong fields in these stars have either been directly measured using Zeeman-Doppler imaging \\citep[ZDI,][]{semel89,jf97,tausco06}, or their presence indicated indirectly e.g. by the large spot coverage on the surface. In the case of V410~Tau this field has revealed itself in the presence of cool spots covering a large fraction of the photosphere. In fact, V410~Tau has exhibited one of the highest optical variabilities observed on a young star, with $\\rm \\Delta V$ reaching 0.6 magnitudes \\citep{stelzer03,grankin08}.  V410~Tau is the target of an ongoing photometric monitoring campaign which began in 1981 \\citep{vrba88, herbst89,petrov94, grankin08}. For most of this time the lightcurve has been smooth, repeating, with a clearly defined period and an amplitude between 0.2 and 0.6 magnitudes.  Two significant decreases in the amplitude of the lightcurve, where $\\Delta \\rm V$ fell below 0.2, have been observed during this campaign. The first of these was between 1981--1983 and the second began in 2005, reaching a minimum in 2007. Additional observations of V410~Tau dating from 1905--1987 are available on photographic plates in the archives of the Sternberg Astronomy Institute. \\cite{sokoloff08} have digitised this data and used it to investigate the variability over a longer period of time. This revealed a third decrease in lightcurve amplitude between 1963--1970, this time in the $B$--band.  These changes in the lightcurves suggest a rearrangement of the cool spots. Given that starspots are a magnetic phenomenon such a change is likely to be caused by a change in the magnetic topology. The fact that this change has been observed three times with intervals of approximately 20~years may be evidence of a stellar cycle, similar to the 22-year solar cycle. The existence of magnetic cycles in active stars has been extensively investigated through changes in their starspot coverage \\cite[see, e.g., ][and references therein]{berd05, strass09}. Possible short-period cycles (of between 5--7 years) have been found on V410~Tau by \\cite{stelzer03} and \\cite{olah09}, using photometric data spanning several decades.  Doppler imaging \\citep{vogt87,strass02} can provide more detail on the changes in spot distribution. Maps of the spot distribution on V410~Tau have been made by \\cite{v41094b, hatzes95b} and \\cite{v41096}. The images have shown decentred polar spots and low latitude spots. Often the lower latitudes features are found to be confined within particular longitude ranges \\citep[e.g.][]{v41096}.  The flattening of the lightcurve suggests that a change in the distribution or number of photospheric spots has occurred, i.e., either there are fewer spots than previously or there is a more even longitudinal distribution of spots. \\cite{grankin08} found a stronger correlation between the non-uniformity of the spot distribution and the amplitude of the lightcurve than the total spot area. This suggests that the reduction in lightcurve amplitude is related to a rearrangement in the spot distribution, rather than a change in the number or area of the spots.  To date, no magnetic images of V410~Tau have been published. This has motivated a long--term campaign, being carried out at T\\'{e}lescope Bernard Lyot (TBL), which aims to produce several spot and magnetic maps of V410~Tau between 2009--2012. We expect that during this time the lightcurve will return to the state observed between 1988--2004. In this paper we present the first set of images. Sec.~2 discusses the evolutionary status of V410~Tau. Sec.~3 describes the observations and data reduction; in Sec.~4 we describe the image production and present the brightness and magnetic images, and a new measurement of the differential rotation. Sec.~5 is a discussion of the implication of the results for the dynamo theory in pre-main sequence stars. ",
        "conclusions": "The Doppler image of the weak-line T~Tauri star V410~Tau in January 2009 shows a large polar spot and a several lower latitude spots spread around the latitude band between $0 - 30^{\\circ}$. This spot distribution can explain the relatively flat V--band lightcurve currently being observed on the star as, unlike during previous Doppler imaging campaigns, the spots are not confined to a particular longitude range. The polar spot on the image is centred on the pole, unlike on previous images where it was shifted from the pole.  The \\ha\\ profile is complex and highly variable. Fast-moving absorption transients in the line may indicate the presence of high-lying circumstellar material, such as slingshot prominences.  The magnetic field is predominately dipolar and non-axisymmetric. It has a complex radial field as well as a strong azimuthal field with large monolithic regions. The field resembles zero-age main sequence stars such as AB~Dor, but unlike that star it has very weak differential rotation. The rotation rate and depth of convection zone appear to be the most important parameters in determining the strength and configuration of the magnetic field  Future Zeeman Doppler images planned as part of a campaign on T\\'{e}lescope Bernard Lyot (TBL) will reveal whether changes in the amplitude of the optical lightcurve are accompanied by significant changes in the spot distribution or magnetic field topology."
    },
    "0911/0911.0779_arXiv.txt": {
        "abstract": "{\\bf star formation; cluster formation} Stellar clusters are born in cold and dusty molecular clouds and the youngest clusters are embedded to various degrees in dusty dark molecular material. Such embedded clusters can be considered protocluster systems. The most deeply buried examples are so heavily obscured by dust that they are only visible at infrared wavelengths. These embedded protoclusters constitute the nearest laboratories for direct astronomical investigation of the physical processes of cluster formation and early evolution. I review the present state of empirical knowledge concerning embedded cluster systems and discuss the implications for understanding their formation and subsequent evolution to produce bound stellar clusters. ",
        "introduction": "The question of the origin of stellar clusters is an old one. As early as 1785, William Herschel first considered this problem in the pages of these {\\it Transactions} when he speculated on the origin of the clusters Messier 80 and Messier 4 in Ophiuchus. More than two centuries later, despite profound advances in astronomical science, we find that the question of the physical origin of stellar clusters remains largely a mystery. This is, in part, due to the fact that cluster formation is a complex physical process that is intimately linked to the process of star formation, for which there is as yet no complete theory. The development of a theoretical understanding of star and cluster formation therefore depends on first acquiring detailed empirical knowledge of these phenomena. This can be a very difficult task. Consider, for example, the globular clusters, the most massive stellar clusters in our own Galaxy, the Milky Way. These stellar systems are more than 12 billion years old and are no longer formed in the Milky Way. Consequently, direct empirical study of their formation process is not possible. The situation is considerably better for Galactic open clusters. These typically span ages from 1 Myr to 1 Gyr and thus must be continually forming in the Milky Way, making direct observational study of their formation process possible, at least in principle. However, such studies have been seriously hindered by the fact that open clusters are born in molecular clouds and, during their formation and early evolution, are completely embedded in molecular gas and dust. They are thus obscured from view at optical wavelengths, where the traditional astronomical techniques are most effective.  Fortunately, molecular clouds are considerably less opaque at infrared wavelengths and the development and deployment of infrared imaging cameras on optical- and infrared-optimized telescopes during the past two decades has provided astronomers with the ability to detect, survey and systematically study the extremely young embedded stellar clusters within nearby molecular clouds. Such studies indicated that embedded clusters are quite numerous and that they account for a significant fraction, if not the majority, of all star formation presently taking place in the Galaxy. Similarly, as discussed in accompanying articles by de Grijs (2010) and Larsen (2010), young clusters also appear to account for a significant fraction of star formation in other galaxies, such as in the very active and luminous starburst galaxies and in the vigorous star-formation episodes which accompany galaxy mergers and close interactions. In our Galaxy, the embedded-cluster phase lasts only 2--4 Myr and the vast majority of embedded clusters which form in molecular clouds dissolve within 10 Myr or less of their birth. High cluster infant mortality has also been inferred for clusters in other galaxies (see, e.g., de Grijs 2010; Larsen 2010). This early mortality of embedded clusters is likely a result of the low star-formation efficiency that characterizes the massive molecular-cloud cores within which the clusters form. The physical reason for the low formation efficiencies is not well understood and the origin of the massive cores themselves is one of the many mysteries of modern astrophysical research.  The embedded clusters are the primary laboratory for research into the question of the physical origin of stellar clusters. At the present time, the answer to this question is shrouded by the dusty veils of the dark molecular clouds in which stars and clusters are born. The ultimate goal of a theoretical understanding of the physical process of cluster formation is far from being realized. The first step towards achieving this goal is to construct a solid empirical foundation upon which a physical theory can eventually be built. This then requires a detailed understanding of the physical nature of embedded clusters and the environments in which they form. In this article, I will review the empirical knowledge that is setting the stage for the eventual development of a theory of cluster formation. I will describe existing knowledge concerning many of the basic physical properties of the embedded-cluster population, including the cluster mass function, cluster birthrate, structural properties, etc., as well as the nature of the gaseous cloud cores in which they form. I will (at the end) briefly speculate on the implications of these results for understanding the origin of stellar clusters. A more theoretical discussion of possible physical mechanisms for forming such clusters is contained in the accompanying article by Clarke (2010).  \\begin{figure} \\vspace{-3.5cm} \\begin{center} \\includegraphics[scale=0.55, angle=90]{lada_rs_f1.ps} \\end{center} \\vskip -1.0in \\caption{{\\it (left)} Optical and {\\it (right)} infrared view of an embedded Trapezium cluster associated with the Great Orion Nebula. North is on the left, west at the top. (From Lada \\& Lada 2003.)} \\end{figure} ",
        "conclusions": "More than two centuries after Herschel (1785) first considered the question, we now understand that clusters are formed in cold, dark molecular clouds. Stellar clusters begin their lives deeply buried in dense molecular gas and dust as embedded infrared protoclusters. Over the past two decades, advances in detector and telescope technology have led to steady advances in our knowledge of embedded protoclusters. These extremely young stellar systems appear to account for a significant fraction of all star formation presently taking place in the Milky Way and, thus, they are tracers of the current epoch of star formation in the Galaxy. Although considerable progress has been made towards understanding the basic physical properties of embedded clusters, a physical theory of cluster formation eludes us. In the Milky Way, the primary mode of cluster formation at the present epoch appears to be one in which relatively compact, massive dense cores in GMCs undergo a process of Jeans-like fragmentation that transforms cold, dense interstellar material into stellar form. Triggering of clouds by shocks or other mechanisms that lead to an increase in external pressure may represent an additional mode of Galactic cluster formation. Although its relative contribution to present-day cluster formation in the Milky Way is unclear, this latter mechanism may be significant in interacting galaxies and nuclear starbursts, and perhaps even in the early history of the Milky Way itself. It is now clear that to develop a predictive theory of cluster formation will require both a better understanding of the process of star formation and a comprehensive understanding of the physical mechanisms that organize the dense material of a GMC into massive, compact cores which, for a typical GMC in the Milky Way, occupy only a small fraction ($< 1$\\%) of its volume and account for only a small fraction (1--10\\%) of its total mass."
    },
    "0911/0911.2310_arXiv.txt": {
        "abstract": "We investigate the global dynamics of the universe within the framework of the Interacting Dark Matter (IDM) scenario. Considering that the dark matter obeys the collisional Boltzmann equation, we can obtain analytical solutions of the global density evolution, which can accommodate an accelerated expansion, equivalent to either the {\\em quintessence} or the standard $\\Lambda$ models. This is possible if there is a disequilibrium between the DM particle creation and annihilation processes with the former process dominating, which creates an effective source term with negative pressure. Comparing the predicted Hubble expansion of one of the IDM models (the simplest) with observational data, we find that the effective annihilation term is quite small, as suggested by various experiments. ",
        "introduction": "The detailed analysis of the available high quality cosmological observations (\\cite{Spergel07,essence,Kowal08,komatsu08} and references therein) have converged during the last decade towards a cosmic expansion history that involves a spatial flat geometry and a recent accelerating expansion of the universe. This expansion has been attributed to an energy component (the so called dark energy) with negative pressure which dominates the universe at late times and causes the observed accelerating expansion. The nature of the dark energy is still a mystery and it is one of the most fundamental current problems in physics and cosmology. Indeed, due to the absence of a physically well-established fundamental theory, there have been many theoretical speculations regarding the nature of the above exotic dark energy (DE) among which a cosmological constant, scalar or vector fields (see \\cite{Weinberg89,Wetterich:1994bg, Caldwell98,Peebles03,Brax:1999gp,Brookfield:2005td, Boehmer:2007qa} and references therein).  Most of the recent papers in this kind of studies are based on the assumption that the DE evolves independently of the dark matter (DM). The unknown nature of both DM and DE implies that we can not preclude future surprises regarding the interactions in the dark sector. This is very important because interactions between the DM and {\\em quintessence} could provide possible solutions to the cosmological coincidence problem. Recently, several papers have been published in this area \\cite{Amm:1999, Bin:2006} proposing that the DE and DM could be coupled, assuming also that there is only one type of non-interacting DM.  However, there are other possibilities. (a) It is plausible that the dark matter is self-interacting (IDM) \\cite{SperStei}, a possibility that has been proposed to solve discrepancies between theoretical predictions and astrophysical observations, among which the gamma-ray and microwave emission from the center of our galaxy (eg. \\cite{92, 90}, \\cite{hoop, regis} and references therein). It has also been shown that some dark matter interactions could provide an accelerated expansion phase of the Universe \\cite{zim, balakin, Lima2008}. (b) The DM could potentially contain more than one particle species, for example a mixture of cold and warm or hot dark matter \\cite{Farrar:2003uw,Gubser:2004uh}, with or without inter-component interactions.  In this work we are not concerned with the viability of the different such possibilities, nor with the properties of interacting DM models. The aim of this work is to investigate only whether there are repercussions of DM self-interactions for the global dynamics of the universe and specifically whether such models can yield an accelerated phase of the cosmic expansion, without the need of the {\\em dark energy}. We note that we do not ``design'' the fluid interactions to produce the desired accelerated cosmic evolution, as in some previous works (eg., \\cite{balakin}), but investigate the circumstances under which the analytical solution space of the collisional Boltzmann equation, in the expanding universe, allows for a late accelerated phase of the universe. ",
        "conclusions": "In this work we investigate analytically the evolution of the global density of the universe in the context of an interacting DM scenario by solving analytically the collisional Boltzmann equation in an expanding Universe. The possible disequilibrium between the DM particle creation and annihilation processes, regardless of its cause and in which the particle creation term dominates, creates an effective source term with negative pressure which (acting as {\\em dark energy}) provides an accelerated expansion phase of the scale factor. Finally, comparing the observed Hubble function of a few high-redshift elliptical galaxies with that predicted by our simplest IDM model ($m=0$), we find that the effective annihilation term is quite small."
    },
    "0911/0911.3316_arXiv.txt": {
        "abstract": "{Recent investigations of the white dwarf (WD) + He star channel of Type Ia supernovae (SNe Ia) imply that this channel can produce SNe Ia with short delay times. The companion stars in this channel would survive and potentially be identifiable.} {In this {\\em Letter}, we study the properties of the companion stars of this channel at the moment of SN explosion, which can be verified by future observations.} {According to SN Ia production regions of the WD + He star channel and three formation channels of WD + He star systems, we performed a detailed binary population synthesis study to obtain the properties of the surviving companions.} {We obtained the distributions of many properties of the companion stars of this channel at the moment of SN explosion. We find that the surviving companion stars have a high spatial velocity ($>$400\\,km/s) after SN explosion, which could be an alternative origin for hypervelocity stars (HVSs), especially for HVSs such as US 708.} {} ",
        "introduction": "\\label{1. Introduction} Type Ia supernovae (SNe Ia) play an important role in the study of cosmic evolution, especially in cosmology. They have been applied successfully in determining cosmological parameters (e.g., $\\Omega$ and $\\Lambda$; Riess et al. 1998; Perlmutter et al. 1999). It is generally believed that SNe Ia are thermonuclear explosions of carbon--oxygen white dwarfs (CO WDs) in binaries (for the review see Nomoto et al. 1997). However, there is still no agreement on the nature of their progenitors (Hillebrandt \\& Niemeyer 2000; Podsiadlowski et al. 2008; Wang et al. 2008), and no SN Ia progenitor system has been conclusively identified from before the explosion.  Over the past few decades, two families of SN Ia progenitor models have been proposed, i.e., the double-degenerate (DD) and single-degenerate (SD) models. Of these two models, the SD model is widely accepted at present (Nomoto et al. 1984). It is suggested that the DD model, which involves the merger of two CO WDs (Iben \\& Tutukov 1984; Webbink 1984; Han 1998), likely leads to an accretion-induced collapse rather than to an SN Ia (Nomoto \\& Iben 1985). For the SD model, the companion is probably a main-sequence (MS) star, a slightly evolved subgiant star (WD + MS channel), or a red-giant star (WD + RG channel) (e.g., Hachisu et al. 1996, 1999a,b; Li $\\&$ van den Heuvel 1997; Langer et al. 2000; Han $\\&$ Podsiadlowski 2004, 2006; Chen $\\&$ Li 2007, 2009; Meng et al. 2009; L\\\"{u} et al. 2009; Wang, Li \\& Han 2009). An explosion following the merger of two WDs would leave no remnant, while the companion star in the SD model would survive and be potentially identifiable (Podsiadlowski 2003). There has been no conclusive proof yet that any individual object is the surviving companion star of an SN Ia. It will be a promising method to test SN Ia progenitor models by identifying their surviving companion stars.  Yoon $\\&$ Langer (2003) followed the evolution of a CO WD + He star system with a $1.0\\,M_{\\odot}$ CO WD and a $1.6\\,M_{\\odot}$ He star in a 0.124\\,d orbit. In this binary, the WD accretes He from the He star and grows in mass to the Chandrasekhar (Ch) mass. SNe Ia from this binary channel can neatly avoid H lines. Recently, Wang et al. (2009a) systematically studied the WD + He star channel of SNe Ia. In the study, they carried out binary evolution calculations of this channel for about 2600 close WD binaries, in which a CO WD accretes material from an He MS star or an He subgiant to increase its mass to the Ch mass. The study shows the parameter spaces for the progenitors of SNe Ia. By using a detailed binary population synthesis (BPS) approach, Wang et al. (2009b) find that the Galactic SN Ia birthrate from this channel is $\\sim$$0.3\\times 10^{-3}\\ {\\rm yr}^{-1}$ and that this channel can produce SNe Ia with short delay times ($\\sim$45$-$140\\,Myr) from the star formation to SN explosion. The companion star in this channel would survive and show distinguishing properties.  In recent years hypervelocity stars (HVSs) have been observed in the halo of the Galaxy. HVSs are stars with a velocity so great that they are able to escape the gravitational pull of the Galaxy. However, the formation of HVSs is still unclear (for a recent review see Tutukov \\& Fedorova 2009). It has been suggested that such HVSs can be formed by the tidal disruption of a binary through interaction with the super-massive black hole (SMBH) at the Galactic center (GC) (Hills 1988; Yu \\& Tremaine 2003). The first three HVSs have only recently been discovered serendipitously (e.g., Brown et al. 2005; Hirsch et al. 2005; Edelmann et al. 2005). Up to now, about 17 HVSs have been discovered in the Galaxy (Brown et al. 2009; Tillich et al. 2009), most of which are B-type stars, probably with masses ranging from 3 to 5\\,$M_\\odot$ (Brown et al. 2005, 2009; Edelmann et al. 2005). One HVS, HE 0437-5439, is known to be an apparently normal early B-type star. Edelmann et al. (2005) suggests that the star could have originated in the Large Magellanic Cloud, because it is much closer to this galaxy (18 kpc) than to the GC (see also Przybilla et al. 2008).  At present, only one HVS, US 708, is a subdwarf O (sdO) star, and Hirsch et al. (2005) speculatS that US 708 is formed by the merger of two He WDs in a close binary induced by the interaction with the SMBH in the GC and then escaped. Recently, Perets (2009) has suggested that US 708 may have been ejected as a binary from a triple disruption by the SMBH, which later on evolved and merged to form an sdO star. Because of the existence of the short orbital periods ($\\sim$1\\,h) for the WD + He star systems, Justham et al. (2009) argues that the WD + He star channel of SNe Ia may provide a natural explanation for stars like US 708.  Han (2008a) obtained the distributions of many properties of the surviving companions from the WD + MS channel of SNe Ia. The properties can be verified by future observations. The purpose of this {\\em Letter} is to investigate the properties of the surviving companions of the WD + He star channel and to explore whether HVSs such as US 708 could have been released from the binaries that produced SNe Ia. In Section 2, we describe the BPS approach and the simulation results for the properties of the surviving companions. Finally, a discussion is given in Section 3. ",
        "conclusions": "\\label{3. Discussion} \\begin{figure} \\includegraphics[width=5.6cm,angle=270]{f1.ps} \\caption{The distribution of properties of companion stars in the plane of ($V_{\\rm orb}$, $M_2^{\\rm SN}$) at the current epoch, where $V_{\\rm orb}$ is the orbital velocity  and $M_2^{\\rm SN}$ the mass at the moment of SN explosion.} \\end{figure}  \\begin{figure} \\includegraphics[width=5.6cm,angle=270]{f2.ps} \\caption{Similar to Fig. 1, but in the plane of ($\\log T_{\\rm eff}$, $\\log g$), where $T_{\\rm eff}$ is the effective temperature of companion stars at the moment  of SN explosion and $\\log g$ the surface gravity. The companion stars are out of thermal equilibrium, e.g., a 1.243$\\,M_\\odot$ He star with 0.274$\\,R_\\odot$ and ($\\log T_{\\rm eff}$, $\\log g$)=(4.70, 5.65) will be a 0.267$\\,R_\\odot$ He star and with ($\\log T_{\\rm eff}$, $\\log g$)=(4.75, 5.67) after the He star back to the thermal equilibrium.} \\end{figure}   \\begin{figure} \\includegraphics[width=8.6cm,angle=0]{f3.ps} \\caption{The distribution of spatial velocity for surviving companion stars of SNe Ia. The dotted line denotes the results of ejecta velocity 11000\\,km/s, while the solid line shows ejecta velocity 13500\\,km/s.} \\end{figure}   \\begin{figure} \\includegraphics[width=5.6cm,angle=270]{f4.ps} \\caption{Similar to Fig. 1, but in the plane of ($\\log P^{\\rm SN}$, $M_2^{\\rm SN}$), where $P^{\\rm SN}$is the orbital period at the moment of SN explosion.} \\end{figure}  \\begin{figure} \\includegraphics[width=5.6cm,angle=270]{f5.ps} \\caption{Similar to Fig. 1, but in the plane of ($V_{\\rm rot}$, $M_2^{\\rm SN}$), where $V_{\\rm rot}$ is the equatorial rotational velocity of companion stars at the moment of SN explosion.} \\end{figure}   Figures 1 and 2 show the distributions of the masses, the orbital velocities, the effective temperatures, and the surface gravities of companion stars at the moment of SN explosion. In Figure 1, the companion star has an orbital velocity of $\\sim$300$-$500\\,${\\rm km/s}$ for a corresponding mass of $\\sim$0.6$-$1.7\\,$M_\\odot$ at the moment of SN explosion. In the three formation channels of the WD + He star systems, SNe Ia mainly come from the He star channel. For this formation channel, the recorded properties at each step show that a primordial binary system with a primary mass $M_{\\rm 1i}\\sim 5.0-8.0\\,M_\\odot$, a secondary mass $M_{\\rm 2i} \\sim 2.0-6.5\\,M_\\odot$, and an orbital period $P_{\\rm i} \\sim 10-40\\,{\\rm d}$ would evolve to a close WD + He star system with a WD mass $M^{\\rm i}_{\\rm WD} \\sim 0.87-1.2\\,M_\\odot$, a He star mass $M^{\\rm i}_{\\rm 2} \\sim 1.0-2.6\\,M_\\odot$, and an orbital period $P^{\\rm i} \\sim 0.04-0.2\\,{\\rm d}$. Finally, the WD + He star system results in an SN Ia explosion and survives a companion star.  However, Figures 1 and 2 are for the moment of SN explosion, which could be modified by the explosion. The SN ejecta will interact with its companion. The companion stars will be stripped of some mass and receive a kick velocity that is perpendicular to the orbital velocity. Adopting a similar method to Meng, Chen \\& Han (2007), we estimated the stripped mass for the companion stars of the WD + He star channel and find that the stripped mass is very low, e.g., a 1.243$\\,M_\\odot$ companion star with 0.274$\\,R_\\odot$ only loses a mass of 0.015$\\,M_\\odot$ in our simulation. We also roughly obtain the kick velocity based on the momentum conservation equation.\\footnote{ However, in reality, the collision between the SN ejecta and its companion cannot be elastic. If we adopt the inelastic collision in our calculations, the kick velocity will only decrease by $\\sim$3\\%$-$5\\%. Thus, the inelastic collision is no significant influence on the kick velocity.} The kick velocity mainly depends on the ratio of separation to the radius of companions at the moment of SN explosion, $A/R_{\\rm 2}^{\\rm SN}$, and the leading head velocity of SN ejecta, which is assumed to be in the range of 11000 to 13500\\,km/s. The velocity of 11000\\,km/s is from the SN ejecta kinetic energy $1.0\\times10^{51}$\\,erg corresponding to the lower limit of normal SN Ia kinetic energy, while the velocity of 13500\\,km/s is from the SN ejecta kinetic energy $1.5\\times10^{51}$\\,erg corresponding to the upper limit of the kinetic energy (Gamezo et al. 2003). We can obtain the spatial velocity by the formula $V_{\\rm 2}^{\\rm SN}=\\sqrt{V_{\\rm kick}^{2}+V_{\\rm orb}^{2}}$, where $V_{\\rm kick}$ and $V_{\\rm orb}$ are the kick velocity and the orbital velocity of the companion star at the moment of SN explosion, respectively. In Figure 3, we show the distribution of the spatial velocity for the surviving companion stars of the WD + He star channel. We see that the surviving companion stars have high spatial velocities ($>$400\\,km/s) that almost entirely exceed the gravitational pull of the Galaxy nearby the sun. Thus, the surviving companion stars from the WD + He star channel could be an alternative origin for HVSs.  US 708 is an extremely He-rich sdO star in the Galaxy halo, with a heliocentric radial velocity of +$708\\pm15$\\,${\\rm km/s}$ (Hirsch et al. 2005). We note that the local velocity relative to the Galatic center may lead to a higher observation velocity for the surviving companion stars, but this may also lead to a lower observation velocity. Considering the local velocity near the sun ($\\sim$220\\,km/s), we find that $\\sim$30\\% of the surviving companion stars may be observed to have velocity $V>700\\,{\\rm km/s}$ for a given SN ejecta velocity 13500\\,km/s. In addition, the asymmetric explosion of SNe Ia may also enhance the velocity of the surviving companions. Thus, a surviving companion star in the WD + He star channel may have a high velocity like US 708 (see also Justham et al. 2008).  The companion stars are out of thermal equilibrium at the moment of SN explosion. For He stars, the equilibrium radii are lower than at the moment of SN explosion. Thus, the surface gravity at equilibrium should be greater than the one in Figure 2; e.g., a 1.243$\\,M_\\odot$ He star with 0.274$\\,R_\\odot$ and ($\\log T_{\\rm eff}$, $\\log g$)=(4.70, 5.65) will be a 0.267$\\,R_\\odot$ He star and with ($\\log T_{\\rm eff}$, $\\log g$)=(4.75, 5.67) after the He star is back to the thermal equilibrium. It is well known that a shock will develop after impact by the ejecta. A large part of the material in the companion's envelope is heated by the shock, and some of the material is vaporized from the surface of the companion stars. Following a similar method to Chen \\& Li (2007), we also estimated the vaporized mass, and find that the mass loss from the companion stars is not significant ($<$5\\%). Such a surviving companion star may be significantly overluminous or underluminous depending on the amount of heating (e.g., Podsiadlowski 2003). Figure 2 could be a starting point for further studies of this kind.  Figure 4 shows the distributions of orbital periods and secondary masses of the WD + He star systems at the moment of SN explosion. The orbital periods and secondary masses of the WD + He star systems at this moment are basic input parameters when one simulates the interaction between SN ejecta and its companion. It is suggested that, for hot stars with radiative envelopes (such as He stars), tidal forces may be inefficient for synchronization (e.g., Zahn 1977). However, the recent study by Charpinet et al. (2008) supports efficient tidal synchronization for hot subdwarf stars. The work by Toledano et al. (2007) also indicates, on the other hand, that even stars with radiative envelopes may have efficient tidal interaction on a time scale comparable to convective envelopes. Thus, we make an assumption that the companion stars co-rotate with their orbits. In Figure 5, we show the distributions of equatorial rotational velocities of the companions stars. We see that the surviving companion stars are fast rotators, so their spectral lines should be broadened noticeably.  The simulation in this {\\em Letter} was made with $\\alpha_{\\rm ce}\\lambda =0.5$. If we adopt a higher value for $\\alpha_{\\rm ce}\\lambda$, e.g., 1.5, the birthrate of SNe Ia would be a little bit higher and the delay time from the star formation to SN explosion longer. This is because binaries emerging from CE ejections tend to have longer orbital periods for a large $\\alpha_{\\rm ce}\\lambda$ and are more likely to be located in the SN Ia production region (Fig. 8 of Wang et al. 2009a). Due to the lack of WD binaries with short orbital periods ($\\log P^{\\rm i} < -1.2$) for a large $\\alpha_{\\rm ce}\\lambda$, the companions with orbital velocity $V_{\\rm orb}>430\\,{\\rm km/s}$ would be noticeably absent in Figure 1.  The distributions are results of the current epoch for a constant SFR. For a single starburst, most of the SN explosions occur between $\\sim$45\\,Myr and $\\sim$140\\,Myr after the starburst; i.e., SNe Ia from the WD + He star channel will be absent in old galaxies. The Galactic SN Ia birthrate from this channel is $\\sim$$0.3\\times 10^{-3}\\ {\\rm yr}^{-1}$ (Wang et al. 2009b). By multiplying the birthrate with a typical MS lifetime of He stars, $\\sim$$10^7$\\,yr, we estimated the current number of this type of HVSs in the Galaxy to be $\\sim$$10^3$. In future investigations, we will employ the Large sky Area Multi-Object fiber Spectral Telescope (LAMOST) to search the HVSs originating from the surviving companion stars of SNe Ia."
    },
    "0911/0911.1639_arXiv.txt": {
        "abstract": "I show that Eddington accretion episodes in AGN are likely to produce winds with velocities $v \\sim 0.1c$ and ionization parameters up to $\\xi \\sim 10^4$~(cgs), implying the presence of resonance lines of helium-- and hydrogenlike iron. These properties are direct consequences of momentum and mass conservation respectively, and agree with recent X--ray observations of fast outflows from AGN. Because the wind is significantly subluminal, it can persist long after the AGN is observed to have become sub--Eddington. The wind creates a strong cooling shock as it interacts with the interstellar medium of the host galaxy, and this cooling region may be observable in an inverse Compton continuum and lower--excitation emission lines associated with lower velocities. The shell of matter swept up by the (`momentum--driven') shocked wind must propagate beyond the black hole's sphere of influence on a timescale $\\la 3\\times 10^5$~yr. Outside this radius the shell stalls unless the black hole mass has reached the value $M_{\\sigma}$ implied by the $M - \\sigma$ relation. If the wind shock did not cool, as suggested here, the resulting (`energy--driven') outflow would imply a far smaller SMBH mass than actually observed. In galaxies with large bulges the black hole may grow somewhat beyond this value, suggesting that the observed $M -\\sigma$ relation may curve upwards at large $M$. Minor accretion events with small gas fractions can produce galaxy--wide outflows with velocities significantly exceeding $\\sigma$, including fossil outflows in galaxies where there is little current AGN activity. Some rare cases may reveal the energy--driven outflows which sweep gas out of the galaxy and establish the black hole--bulge mass relation. However these require the quasar to be at the Eddington luminosity. ",
        "introduction": "Outflows and winds are widely observed in many types of galaxies, from active to normal (e.g. Chartas et al., 2003; Pounds et al., 2003a, 2006; O'Brien et al., 2005, Holt et al., 2008; Krongold et al., 2007; Tremonti et al., 2007). There is often a strong presumption that the central supermassive black hole (SMBH) is implicated. Although other types of driving exist, e.g. by supernovae or starbursts, in many cases these phenomena are themselves associated with accretion episodes on to the SMBH.  From a theoretical viewpoint, outflows driven by black holes offer a simple way of establishing relations between the SMBH and its host galaxy, and hence potential explanations for the $M - \\sigma$ and $M - M_{\\rm bulge}$ relations (Ferrarese \\& Merritt, 2000; Gebhardt et al. 2000; H\\\"aring \\& Rix 2004). Such outflows are plausible, as AGN must feed at high rates to grow the observed SMBH masses, given the short duty cycle implied by the rarity of AGN among all galaxies. There is no obvious reason why these rates should respect the hole's Eddington limit, and outflows are a natural consequence.  However it is unclear just how, if at all, the various types of observed outflows mentioned in the first paragraph fit together in terms of these ideas. This paper aims to clarify this. ",
        "conclusions": "We have shown that Eddington accretion episodes in AGN produce winds with velocities $v \\sim 0.1c$ and ionization parameters $\\xi \\sim 10^4$~(cgs) requiring the presence of resonance lines of helium-- and hydrogenlike iron. These properties follow from momentum and mass conservation, and agree with recent X--ray observations of high--speed outflows from AGN. Because the winds have speeds $\\sim 0.1c$ they can persist long after the AGN have become sub--Eddington.  The wind creates a strong cooling shock as it impacts the interstellar medium of the host galaxy. This cooling region may be observable, as it produces an inverse Compton continuum and lower--excitation emission lines associated with lower velocities. The shell of matter swept up by the ('momentum--driven') shocked wind emerges from the black hole's sphere of influence on a timescale $\\la 3\\times 10^5$~yr. The shell then stalls unless the black hole mass has reached the value $M_{\\sigma}$ implied by the $M - \\sigma$ relation. If the wind shock did not cool, as suggested here, the resulting (`energy--driven') outflow would imply a far smaller SMBH mass than actually observed. In galaxies with large bulges the black hole must grow somewhat beyond this value, suggesting that the observed $M -\\sigma$ relation may curve upwards at large $M$. Galaxy--wide outflows with velocities significantly exceeding $\\sigma$ probably result from minor accretion episodes with low gas fractions. These can appear as fossil outflows in galaxies where there is little current AGN activity. In rare cases it may be possible to observe the energy--driven outflows which sweep gas out of the galaxy and establish the black hole--bulge mass relation. However these require the quasar to be at the Eddington luminosity."
    },
    "0911/0911.3120_arXiv.txt": {
        "abstract": "{The historic plates of the {\\it Carte du Ciel}, an international cooperative project launched in 1887, offer valuable first-epoch material for the determination of absolute proper motions. } {We present the CdC-SF, an astrometric catalogue of positions and proper motions derived from the  \\textit{Carte du Ciel} plates of the San Fernando zone, photographic material with a mean epoch of 1901.4 and a limiting magnitude of V$\\sim$16, covering the declination range of $-10^{\\circ} \\leq \\delta \\leq -2^{\\circ}$. } {Digitization has been made using a conventional flatbed scanner. Special techniques have been developed to handle the combination of plate material and the large distortion introduced by the scanner. The equatorial coordinates are on the ICRS defined by Tycho-2, and proper motions are derived using UCAC2 as second-epoch positions. } { The result is a catalogue with positions and proper motions for 560000 stars, covering 1080 degrees$^2$. The mean positional uncertainty is 0$\\farcs$20 (0$\\farcs$12 for well-measured stars) and the proper-motion uncertainty is 2.0 mas/yr (1.2 mas/yr for well-measured stars). } {The proper motion catalogue CdC-SF is effectively a deeper extension of Hipparcos, in terms of proper motions, to a magnitude of 15. } ",
        "introduction": "The Hipparcos Catalogue (Perryman et al. 1997) contains the positions, parallaxes and proper motions of 118218 stars with great precision (proper motions precision is 1~mas/year), complete up to a magnitude of $V=7.3$. Due to its precision and homogeneity, this catalogue constitutes an important source of information and has provided a useful tool in numerous kinematic studies. However, its relatively bright magnitude limit restricts it to very nearby stars, with the number of stars per unit area also being lower than what is often desired in kinematic analyses of the Galaxy.  For this reason, various efforts have been made to develop astrometric catalogues that are more dense and reach deeper magnitudes, becoming in effect an extension of the Hipparcos catalogue. It is necessary that such catalogues contain proper motions besides positions in order to provide the essential information with the quality and completeness needed.  The Tycho-2 Catalogue (H{\\o}g et al. ~ 2000) can be considered an extension of Hipparcos in the way that it contains positions and proper motions for stars brighter than V$\\sim$12 with errors of 2.5~mas/year. At the present time the Hipparcos and Tycho-2 catalogues are the largest high-quality proper motion surveys.  Other proper motion catalogues available at the present time are UCAC2 catalogues (Zacharias et al. 2004, 2--7 mas/yr depending on the magnitude and current version has systematic errors of the order of 3~mas/yr), NPM (Hanson et al. 2004, 5~mas/yr, $8<B<18$) and SPM (Girard et al. 2004, 4~mas/yr, $5<V<18.5$). For an overview in detail of the current astrometric catalogues, see Girard (2008).  When determining proper motions, a long time-baseline is highly desirable. Historical photographic records provide ideal material for the calculation of proper motions. This is a prime motivation for the digitization and reduction of the collections of photographic plates that were accumulated in observatories at the end of the 19th century.  In 1887, the first great astronomical project of international cooperation in history was planned in Paris. Its objective was the construction of a complete catalogue up to a magnitude of V$\\sim$12, the {\\it Astrographic Catalogue} (AC), and to map the sky up to V$\\sim$15, the {\\it Carte du Ciel} (CdC). A total of twenty observatories around the world participated in the task of taking the necessary observations. This material represents the first existing photographic record of the complete sky and is an extremely important potential resource for determining proper motions due to the long time-baseline relative to modern observations, $\\sim$100 years in most cases.  This study is part of a larger project, the goal of which is to produce an astrometric catalogue with proper motions for a variety of applications in Galactic kinematics, as well as to provide a dense, faint reference frame by recovering and utilizing the historical CdC plates. In order to construct the catalogue, the first necessary step is to digitize these plates. The digitization and measurement process is discussed in detail in a previous paper (Vicente at al. 2007, hereafter Paper~I), where an innovative method of digitization using a commercial flatbed scanner was presented. It was demonstrated that this device, together with the techniques developed to remove the distortion it introduces, is an alternative to more specialized measuring machines that are less available, while yielding similar measuring errors of 0$\\farcs$18 as other groups, deriving astrometry from similar plate material, although from other CdC collections (Rapaport 2006, Ortiz-Gil 1998, Geffert 1996).   The main scientific value of positions at the epoch of 1900 resides in their usefulness in providing positions to serve as the first epoch in proper motion determinations involving the compilation of several catalogues of different epochs.  The current paper explains the various steps and procedures used during the construction of the CdC-SF proper-motion catalogue. The catalogue's properties are described in detail, including an estimation of uncertainties in the positions and proper motions. For the convenience of the reader, some of the figures and tables from Paper~I are reproduced here for their relevance to the catalogue description. ",
        "conclusions": "We have presented the CdC-SF Catalogue, an astrometric catalogue with positions and proper motions for more than 500,000 stars, in the ICRS system, with a mean epoch of 1901.4, to $V=16$, covering 1080 degrees$^2$ in the sky area between $\\alpha=(06^h,14^h),\\delta=(-10.5^{\\circ},-2.5^{\\circ})$. The mean uncertainty in positions is 0$\\farcs$20 (0$\\farcs$12 for well-measured stars), and in proper motions it is 2~mas/yr (1.1~mas/yr for well-measured stars). The catalogue is based on the reduction of the original {\\it Carte du Ciel} photographic plates, San Fernando zone, that have been ``revived'' by digitization with a commercial flatbed scanner.  External comparison of the data has been made with the catalogues Hipparcos and UCAC2, as well as a study of the proper-motion errors using open clusters. The catalogue achieves approximately the same precision in proper motions as the Hipparcos catalogue, but it extends to seven magnitudes fainter.  Finally, the CdC-SF catalogue covers -10$^{\\circ}$ to +60$^{\\circ}$ galactic latitude and will provide an excellent basis for studies of the kinematics of different components of our Galaxy. The positions at the epoch 1900 should prove to be useful for other proper-motion determinations."
    },
    "0911/0911.3911_arXiv.txt": {
        "abstract": "It is well established that a dominant phase in the growth of massive galaxies occurred at high redshift and was heavily obscured by gas and dust. Many studies have explored the stellar growth of massive galaxies but few have combined these constraints with the growth of the supermassive black hole (SMBH; i.e.,\\ identified as AGN activity). In this brief contribution we highlight our work aimed at identifying AGNs in $z\\approx$~2 luminous dust-obscured galaxies. Using both sensitive X-ray and infrared (IR)--submillimeter (submm) observations, we show that AGN activity is common in $z\\approx$~2 dust-obscured systems. With a variety of techniques we have found that the majority of the AGN activity is heavily obscured, and construct diagnostics based on X-ray--IR data to identify some of the most heavily obscured AGNs in the Universe (i.e.,\\ AGNs obscured by Compton-thick material; $N_{\\rm H}>1.5\\times10^{24}$~cm$^{-2}$). On the basis of these techniques we show that SMBH growth was typically heavily obscured ($N_{\\rm H}\\ge10^{23}$~cm$^{-2}$) at $z\\approx$~2, and find that the growth of the SMBH and spheroid was closely connected, even in the most rapidly evolving systems. ",
        "introduction": "%  There is no doubt that Active Galactic Nuclei (AGN) are an important component in the formation and evolution of galaxies. The seminal discovery that massive galaxies in the local Universe harbor a super-massive black hole (SMBH) indicates that they must have hosted AGN activity over the past $\\approx$~13~Gyrs (e.g.,\\ Soltan 1982; Kormendy \\& Richstone 1995). Furthermore, the close relationship between the mass of the SMBH and the spheroid in local galaxies points towards a close connection between the growth of the SMBH and the host galaxy (e.g.,\\ Magorrian et~al. 1998; Ferrarese \\& Merritt 2000). Finally, the most successful galaxy formation and large-scale structure models require AGN outflows to suppress star formation and reproduce the properties of massive galaxies in the local Universe (e.g.,\\ Croton et~al. 2006; Bower et~al. 2006).  Stellar population synthesis modelling of nearby galaxies and studies of the cosmic history of star formation indicate that the bulk of the stellar build up of massive galaxies must have occurred at high redshift ($z\\approx$~2; i.e.,\\ when the Universe was $\\approx$~25\\% the current age; e.g.,\\ Heavens et~al. 2004; Hopkins et~al. 2006). Under the assumption that the growth of the SMBH is concordant with that of the galaxy spheroid, the dominant growth phase of SMBH growth in massive galaxies must also have occurred at high redshift. Currently the most {\\it efficient} identification of AGNs (and therefore SMBH growth) is made with deep X-ray surveys (e.g.,\\ Brandt \\& Hasinger 2005). For example, the deepest X-ray surveys (e.g.,\\ Alexander et~al. 2003; Luo et~al. 2008) are able to detect even moderately luminous AGN activity at $z\\approx$~2 ($L_{\\rm X}\\approx10^{42}$--$10^{43}$~erg~s$^{-1}$; i.e.,\\ $>10$ times below the canonical threshold assumed for quasars). Furthermore, the high rest-frame energies probed by these surveys at $z\\approx$~2 ($\\approx$~1.5--24~keV) means that the selection of AGNs is {\\it almost} obscuration independent. This latter point is important since the dominant growth phase of massive galaxies appears to have been heavily obscured by dust and gas (e.g.,\\ Chapman et~al. 2005; Le Floc'h et~al. 2005).  In this brief contribution we hightlight some of our recent work aimed at identifying obscured AGN activity in distant rapidly evolving galaxies at $z\\approx$~2; see \\S2 \\& \\S3. We discuss these results within the context of the SMBH--sperhoid growth phase of today's massive galaxies ($M_{\\rm GAL}\\approx10^{11}$~$M_{\\odot}$); see \\S4. ",
        "conclusions": ""
    },
    "0911/0911.1063_arXiv.txt": {
        "abstract": "We discuss a possible solution to the cosmological constant problem based on the hypothesis of a fixed point for higher dimensional quantum gravity coupled to a scalar. At the fixed point, which is reached only for an infinite value of the scalar field, dilatation symmetry becomes exact. For this limit we concentrate on the absence of a scalar potential since a polynomial potential is not consistent with dilatation symmetry in higher dimensions. We find generic solutions of the higher dimensional field equations for which the effective four-dimensional cosmological constant vanishes, independently of the parameters of the higher dimensional effective action. Under rather general circumstances these are the only quasistatic stable extrema of the effective action which lead to a finite four-dimensional Planck mass. We discuss the associated higher dimensional self-tuning mechanism for the cosmological constant. If cosmological runaway solutions approach the fixed point as time goes to infinity, the effective dark energy vanishes asymptotically.  In the present cosmological epoch the fixed point is not yet reached completely, resulting in a tiny amount of dark energy, comparable to dark matter. We discuss explicitly higher dimensional geometries which realize such asymptotic solutions for time going to infinity. ",
        "introduction": "\\label{intro} Anomalous dilatation symmetry may be a key ingredient for a dynamical solution of the cosmological constant problem \\cite{CWQ}, \\cite{PSW}. In the asymptotic limit of time $t$ going to infinity, a cosmological runaway solution of the field equations can approach a fixed point. At the fixed point all memory of explicit mass or length scales is lost and the quantum effective action becomes dilatation symmetric. In other words, the dilatation anomaly vanishes when it is evaluated for the field configurations corresponding to the fixed point \\cite{CWCC}. In consequence, a scalar field becomes massless in the asymptotic limit for $t\\to\\infty$, corresponding to the Goldstone boson of spontaneously broken dilatation symmetry.  For the approach to the fixed point at finite $t$ the anomaly is not yet zero, and correspondingly the scalar ``pseudo-Goldstone boson'' still has a small mass, that vanishes only asymptotically. These ideas are realized in practice in quintessence cosmologies, where the ``cosmon''-field plays the role of the pseudo-Goldstone boson of spontaneously broken anomalous dilatation symmetry \\cite{CWQ}, \\cite{CWCC}. The cosmon mass is varying with time and of the order of the Hubble parameter \\cite{CWAA}.  Before discussing cosmological solutions, it is crucial to understand the asymptotic solution towards which the cosmological runaway solution converges asymptotically. In our scenario, this asymptotic solution should correspond to flat Minkowski space and therefore to an ``asymptotically vanishing cosmological constant''. However, a vanishing cosmological constant is not enforced by dilatation symmetry - asymptotic dilatation symmetry could, in principle, also be realized with a non-zero effective cosmological constant.  In two recent papers \\cite{CWCC}, \\cite{CWNL} we have argued that a higher dimensional setting sheds new light on the question why the effective four-dimensional cosmological constant $\\Lambda$ vanishes asymptotically. Actually, such an asymptotic vanishing happens for a large generic class of cosmological runaway solutions without any tuning of parameters. The remaining problem concerns the strict limits on the possible variation of fundamental couplings, which are not obeyed for some classes of such solutions. In ref. \\cite{CWCC} we have discussed several scenarios how to reconcile asymptotically static couplings (with an interesting possible exception in the neutrino sector \\cite{CWN}) and an asymptotically vanishing cosmological constant.  In the present paper we investigate higher dimensional models for which the quantum effective action exhibits an exact dilatation symmetry for a suitable fixed point. This fixed point should become relevant for the asymptotic solution. A dilatation symmetric effective action is also the starting point for approaches where dilatation symmetry is realized as an exact quantum symmetry \\cite{SD}. Thus the basic object of our investigation is the quantum effective action $\\Gamma$ where all quantum fluctuations are already included. The field equations derived from $\\Gamma$ are exact without any further quantum corrections. Our key finding is that in the presence of higher dimensional dilatation symmetry these field equations have generic solutions for which the four-dimensional cosmological constant vanishes, $\\Lambda=0$. They correspond to stable extrema of $\\Gamma$.  We do not postulate here that the effective action of a fundamental theory is dilatation symmetric. In general, it is not, and dilatation anomalies are present. For example, a dilatation anomaly arises if running couplings induce an explicite scale, or if a higher dimensional cosmological constant is present. We only make the hypothesis that $\\Gamma$ has a fixed point for certain asymptotic field values to be specified below. Only in this asymptotic limit the dilatation anomaly vanishes - only the ``fixed point effective action'' is dilatation symmetric.  For a cosmological runaway solution the values of fields are not static but continue to change for all times. We investigate solutions where the fields move towards the fixed point region for $t\\to\\infty$. For $t\\to\\infty$ the field equations derived from the fixed point effective action become accurate. As the runaway solution approaches a fixed point the anomalous parts of the effective action vanish. Since such a fixed point is generally reached only asymptotically for $t\\to\\infty$, it is only in this limit that the field equations exhibit an exact dilatation symmetry.  Our discussion of a dilatation symmetric effective action will therefore be insufficient for the cosmological solutions. The solutions of the field equations derived from the dilatation symmetric effective action account only for the possible asymptotic states for $t\\to\\infty$. Nevertheless, it is a necessary condition in our scenario that the dilatation symmetric effective action has an extremum which describes an acceptable asymptotic state. This solution should have flat four-dimensional space. Furthermore, a simple solution to the problem of time-varying couplings would be a static ratio between the characteristic length scale of internal geometry and the effective four-dimensional Planck length. The latter should lead after dimensional reduction to finite non-zero values of the gauge couplings and other dimensionless couplings and mass ratios in the resulting model of particle physics. Inversely, if a satisfactory asymptotic solution exists, the chances that a runaway solution will approach it for $t\\to\\infty$ are quite high.  Dilatation symmetry is easily realized in models that contain besides the metric also a scalar field $\\xi$. A striking difference between dilatation symmetry in higher dimensions, $d>6$, as compared to four dimensions is the absence of a polynomial potential for a scalar field with a canonical kinetic term \\cite{CWCC}. Indeed, for $d=4~mod~4$ (or $d=2~mod~4$ with a symmetry $\\xi\\to-\\xi$) the most general effective action , which is polynominal in $\\xi$ and consistent with general coordinate transformations (diffeomorphism-symmetry) and global dilatation symmetry, reads \\begin{equation}\\label{AA1a} \\Gamma=\\int_{\\hat x}\\hat g^{1/2}\\Big\\{-\\frac12\\xi^2\\hat R +\\frac\\zeta2\\partial^{\\hat\\mu}\\xi\\partial_{\\hat\\mu}\\xi +F(\\hat R_{\\hat\\mu\\hat\\nu\\hat\\rho\\hat\\sigma})\\Big\\}. \\end{equation} The normalization of $\\xi$ is chosen such that it has a canonical kinetic term, up to the free dimensionless parameter $\\zeta$.   While for $d=4$ a term $\\lambda\\xi^4$ is dilatation invariant (and for $d=6$ a term $\\gamma\\xi^3$), no polynomial $\\xi^n$ with integer $n$ is dilatation invariant for $d>6$.  We may first investigate the case that $F$ contains only polynomials of the curvature tensor $\\hat R_{\\hat\\mu\\hat\\nu\\hat\\rho\\hat\\sigma}$ or derivatives thereof. A typical example for $F$ is \\begin{equation}\\label{AA2a} F=\\tau\\hat R^{\\frac d2}, \\end{equation} and we will discuss more general polynomial forms of $F$ in the next section (cf. ref. \\cite{CWCC}, \\cite{CWNL}.) For odd dimensions no invariant polynomial in the curvature tensor exists and therefore $F=0$. (We disregard invariants involving the $\\epsilon$-tensor.) We will often concentrate on the simplest form of such a fixed point effective action by taking $F=0$ in eq. \\eqref{AA1a}. Nevertheless, we also discuss more general forms. The polynomial form of $F$ is actually not essential for our findings and we will extend our discussion to the most general dilatation symmetric term $F$ that only involves the metric. In contrast, the absence of a dilatation symmetric (non-polynomial) potential $V(\\xi)$ at the fixed point remains important for certain aspects of our discussion. We only briefly comment in sect. \\ref{quasi} on possible modifications if a non-polynominal potential for $\\xi$ would be present at the fixed point. It seems likely that this will not change the existence of extrema of $\\Gamma$ with $\\Lambda=0$.   The key argument for the generic existence of solutions with $\\Lambda=0$ is rather simple and directly related to higher dimensional dilatation symmetry. Let us consider the most general form of a dilatation symmetric quantum effective action \\begin{equation}\\label{2A} \\Gamma=\\int_{\\hat x}\\hat g^{1/2}{\\cal L}. \\end{equation} Dilatation transformations correspond to a rescaling of the metric by a constant factor $\\alpha^2$, and an associated rescaling of $\\xi$, \\begin{eqnarray}\\label{2B} \\hat g_{\\hat\\mu\\hat\\nu} \\to \\alpha^2\\hat g_{\\hat\\mu\\hat\\nu}~,~\\hat g^{1/2}\\to\\alpha^d\\hat g^{1/2},\\nonumber\\\\ \\xi\\to\\alpha^{-\\frac{d-2}{2}}\\xi~,~{\\cal L}\\to\\alpha^{-d}{\\cal L}. \\end{eqnarray} A dilatation symmetric effective action remains invariant under these rescalings. For general field values we may define $\\Gamma_\\kappa[\\hat g_{\\hat\\mu\\hat\\nu},\\xi]=\\Gamma[\\kappa^{-2}\\hat g_{\\hat\\mu\\hat\\nu}, \\kappa^{\\frac{d-2}{2}}\\xi]$ and consider the asymptotic limit $\\kappa\\to\\infty$. Our fixed  point hypothesis states then that $\\Gamma_\\kappa$ becomes dilatation symmetric for $\\kappa\\to\\infty$.  The special role of dilatation symmetry for the problem of the cosmological constant is visible already for the most general form of a dilatation symmetric effective action. We are interested in configurations with a block diagonal metric \\begin{equation}\\label{2c} \\hat g_{\\hat\\mu\\hat\\nu} (x,y)= \\left(\\begin{array}{ccc} \\sigma(y)g^{(4)}_{\\mu\\nu} (x) &,&0\\\\ 0&,&g^{(D)}_{\\alpha\\beta}(y) \\end{array}\\right). \\end{equation} Here $x^\\mu$ denotes the four-dimensional coordinates and $y^\\alpha$ are coordinates of $D$-dimensional internal space, with corresponding metrics $g^{(4)}_{\\mu\\nu}$ and $g^{(D)}_{\\alpha\\beta},d=D+4$. The function $\\sigma(y)$ accounts for a possible warping \\cite{RSW,7A,RDW,RS}. The configuration for the metric is supplemented by a configuration for the scalar field $\\xi(y)$. With $\\hat g^{1/2}=(g^{(4)})^{1/2}\\sigma^2(g^{(D)})^{1/2}$ we define \\begin{equation}\\label{2D} W(x)=\\int_y(g^{(D)}{(y)})^{1/2}\\sigma^2(y){\\cal L}(x,y), \\end{equation} and write \\begin{equation}\\label{2E} \\Gamma=\\int_x(g^{(4)})^{1/2} W. \\end{equation} We observe the scaling under dilatations $(g^{(4)}_{\\mu\\nu}\\to\\alpha^2 g^{(4)}_{\\mu\\nu}~,~g^{(D)}_{\\alpha\\beta}\\to\\alpha^2 g^{(D)}_{\\alpha\\beta})$ \\begin{equation}\\label{2F} W\\to\\alpha^{-4}W. \\end{equation}  The possible extrema of $\\Gamma$ can be classified into two categories, according to the existence of an extremum of $W(x)$ or not. Consider first the case where an extremum of $W(x)$ exists for an appropriate field configuration $\\hat g_{\\hat\\mu\\hat\\nu},\\xi$. This means that a (infinitesimally) close neighboring field configuration does not change the function $W(x)$. We may denote the value of $W(x)$ at its extremum by $W_0(x)$. The scaling property \\eqref{2F} immediately implies $W_0(x)=0$. Indeed, using in eq. \\eqref{2B} $\\alpha=1+\\epsilon$, with infinitesimal $\\epsilon$, defines a neighboring configuration. The combination of scaling and extremum condition, $\\partial_\\epsilon(1+\\epsilon)^{-4}W_0=0$, can be obeyed only for $W_0=0$. With eq. \\eqref{2E} and $W_0=0$ an extremum of $W$ is also an extremum of the effective action $\\Gamma$. Thus, whenever an extremum of $W$ exists, this defines an extremum of $\\Gamma$ obeying the field equations, with a vanishing value of $\\Gamma$ at the extremum, $\\Gamma_0=0$.  For all solutions which admit dimensional reduction to an effective four dimensional ``local'' theory of gravity we will show that $\\Gamma_0=0$ implies that the effective four-dimensional cosmological constant $\\Lambda$ vanishes. The existence of extrema of $W$ is a rather generic feature and we find that within our setting stable extrema in this ``flat phase'' always exist. Only the question remains if this flat phase includes solutions with interesting particle physics. There may exist additional extrema of $\\Gamma$ which are not extrema of $W$ and for which $\\Lambda$ does not vanish. Four dimensional theories with a non-vanishing cosmological constant are possible only for this possible second category of extrema of $\\Gamma$ which are not extrema of $W(x)$. In the absence of a non-polynominal scalar potential we find that all such possible solutions in the non-flat phase are unstable if $\\xi \\neq 0$. This singles out a vanishing four-dimensional cosmological constant.  The  existence of extrema of $W$ is a qualitative feature  that does not depend on the precise values of the parameters characterizing ${\\cal L}$. As an example we may consider the polynomial effective action \\eqref{AA1a}. It is obvious that $\\hat R_{\\hat\\mu\\hat\\nu\\hat\\rho\\hat\\sigma}=0~,~\\xi=\\xi_0=$const. corresponds to such an extremum, with ${\\cal L}_0=0$. The existence of this extremum extends to a very large class of non-polynomial ${\\cal L}$ as well. It is sufficient that the $\\xi$-dependent part of ${\\cal L}$ vanishes for flat space and constant $\\xi$, and that the remaining $\\xi$ independent part $F$ vanishes for $\\hat R_{\\hat\\mu\\hat\\nu\\hat\\rho\\hat\\sigma}=0$.  Furthermore, if \\begin{equation}\\label{2G} W_F(x)=\\int_y(g^{(D)})^{1/2}\\sigma^2 F \\end{equation} admits an extremum with respect to variations of the metric, this must occur for $W_{F,0}(x)=0$ by the same scaling arguments as above. (No polynomial form of $F$ needs to be assumed here.) A dilatation symmetric theory involving only gravity (without the scalar field $\\xi$) will imply a vanishing cosmological  constant if an extremum of $W_F$ exists which leads to acceptable four dimensional gravity (with nonzero and finite effective Planck mass).  Let us next extend the setting by adding a scalar part ${\\cal L}_\\xi$ which is quadratic in $\\xi$, without being necessarily polynomial in the curvature tensor. For a given metric $\\hat g_{\\hat\\mu\\hat\\nu}$ we can find a partial extremum of $\\Gamma_\\xi=\\int_{\\hat x}\\hat g^{1/2}{\\cal L}_\\xi=\\int_x(g^{(4)})^{1/2} W_\\xi$ by solving the higher dimensional field equation for $\\xi$, consistent with an extremum condition in case of singular geometries which will be derived later. This results in $W_\\xi=0$, as expected for possible extrema of a purely quadratic polynomial. A necessary condition for an extremum of $W=W_\\xi+W_F$ remains therefore $W_F=0$. However, only the sum $W_\\xi+W_F$ has to be an extremum with respect to variations of the metric, and not $W_\\xi$ and $W_F$ separately. We discuss an example of this type of extrema in the appendix. We observe that for a quadratic ${\\cal L}_\\xi$ we could ``integrate out'' the scalar field in favor of non-local gravitational interactions. This demonstrates incidentally that our setting is not restricted to a local effective gravitational action.  From these simple observations we conclude that in the presence of dilatation symmetry the vanishing of the cosmological constant $\\Lambda$ is very robust with respect to variations of the precise form of $\\Gamma$. The effective action may be characterized by many parameters, as for example the coefficients of different terms appearing in a polynomial approximation of $F$. A change of the values of these parameters will typically not change the value $\\Lambda=0$, as long as an extremum of $W$ continues to exist and is compatible with effective four dimensional gravity. This feature is a particular consequence of dilatation symmetry - it no longer holds in presence of dilatation anomalies. In view of the robustness of $\\Lambda=0$ we find many analogies to phases in many body theories. For the ``flat phase'' an extremum of $W$ exists and $\\Lambda=0$. Extrema of $\\Gamma$ in the ``non-flat phase'' are not extrema of $W$, and we will typically find $\\Lambda>0$.  For effective four dimensional theories in the flat phase we will show that $W$ can be associated with the effective potential $V$ for four-dimensional scalar fields. The minimum of $V$ occurs then necessarily for $V_0=0$. Furthermore, $V$ has flat directions. One such flat direction corresponds to the dilatations \\eqref{2B} and the corresponding massless scalar field is the dilaton. There may be more flat directions since the extrema of $W(x)$ may occur for a whole class of different configurations $\\hat g_{\\hat\\mu\\hat\\nu},\\xi$. In this case we expect additional massless fields for the dilatation symmetric asymptotic cosmological solution for $t\\to\\infty$. For finite $t$ the presence of dilatation anomalies will induce mass terms for these scalars, which vanish only for $t\\to\\infty$. It is conceivable that such additional light scalar fields (beyond the cosmon) are interesting candidates for dark matter.  Our investigation focuses on two issues. \\begin{itemize} \\item [(i)] The existence of extrema of $W$ and a first discussion of characteristic properties of solutions in the flat phase. \\item [(ii)] The dimensional reduction to effective four dimensional gravity  and the establishment that $\\Lambda=0$ in the flat phase. \\end{itemize} A demonstration that an effective action which admits a flat phase does generically not allow other extrema with arbitrarily small $|\\Lambda|$ is given in ref. \\cite{CWNL}. In particular, there are no continuous families of extrema where $\\Lambda$ appears as a continuous parameter. This issue is important for warped geometries with singularities where the existence of families of solutions of the higher dimensional field equations with continuous $\\Lambda$ is known \\cite{RSW}, \\cite{7A}, \\cite{RDW}.  In this case the extremum conditions for $\\Gamma$ go beyond the higher dimensional field equations \\cite{CWCON}. They precisely select the solutions with $\\Lambda=0$ out of the continuous family of solutions. The particular properties of warped spaces within our general dilatation symmetry setting will be discussed in an accompanying paper \\cite{P2}.  In the course of our discussion we will explicitly address the issues of quantum corrections, ``tuning of the cosmological constant to zero'', and ``naturalness'' of the solutions with $\\Lambda=0$. Some of the general aspects are already discussed in \\cite{CWCC} and not repeated here, such that we concentrate here on our specific setting. We discuss the special role of higher dimensions for the ``self-tuning'' of the cosmological constant to zero. This self-tuning is particular to the case of dilatation symmetry and closely connected to the robustness of the existence of extrema of $W$ under a change of parameters in the effective action. We find rather satisfactory answers to the naturalness problem. Asymptotic dilatation symmetry in higher dimensional theories may indeed provide the key for a solution of the cosmological constant problem.  Our paper is organized as follows. In sect. \\ref{dilatation} we present a first discussion of a polynomial effective action \\eqref{AA1a}. We establish explicitely the existence of extrema of $W$ belonging to the flat phase for a large class of effective actions, and demonstrate the vanishing of the cosmological constant if effective four dimensional gravity exists. We extend this discussion to the most general dilatation symmetric effective action and show that extrema in the flat phase with $\\Lambda=0$  exist whenever an appropriate $d-4$ dimensional functional $\\bar W$ admits an extremum. Section \\ref{quasi} addresses the existence of the two ``phases'' for solutions in an effective four dimensional framework and shows why $\\Lambda=0$ is singled out in the flat phase. We concentrate on ``quasistatic solutions'' for which the four dimensional fields are static and homogeneous, while the metric describes a geometry with maximally four dimensional symmetry, with positive, negative or vanishing $\\Lambda$. It also demonstrates that the non-flat phase with $\\Lambda>0$ can only be realized if certain conditions in parameter space are met. The non-flat phase appears to be less generic than the flat phase. We find that all possible extrema in the non-flat phase for $\\xi \\neq 0$ are unstable.  In sect. \\ref{extended} we discuss an interesting ``extended scaling symmetry'' which becomes realized for the simplest fixed point with $F=0$, or if the contribution of the term $\\sim F$ can be neglected for the extremum condition. Extended scaling symmetry admits quasistatic solutions only in the flat phase. Sect. \\ref{adjustment} deals with the issue of an ``adjustment'' or tuning'' of the cosmological constant. The robustness of $\\Lambda=0$ with respect to changes of parameters in the dilatation symmetric effective action is associated to a mechanism of ``self-adjustment'' or ``self-tuning''. In this respect we highlight the differences between a four dimensional theory with a finite number of scalar fields and a higher dimensional setting with infinitely many effective four dimensional scalar fields. This is an important ingredient for the understanding of the robustness of the flat phase, which would be hard to implement with a finite number of degrees of freedom.  Sect. \\ref{higherorder} investigates the dilatation symmetric effective action \\eqref{AA1a} with a polynomial form of $F$. We display the field equations and discuss simple solutions with a vanishing four-dimensional cosmological constant $\\Lambda$. An appendix is devoted to particular solutions with non-Ricci-flat internal space, which nevertheless results in $\\Lambda=0$. These solutions demonstrate explictly that non-abelian isometries of internal space are possible for extrema in the flat phase.  The final part presents a simple estimate of the dilatation anomaly in sect. \\ref{dilatationanomaly}. In sect. \\ref{cosmologicalrunaway} we show that this leads to cosmological runaway solutions for which $\\Lambda$ vanishes asymptotically, while for finite $t$ a homogeneous dark energy component accounts for quintessence. We present our conclusions in sect. \\ref{conclusions}. ",
        "conclusions": "\\label{conclusions}   In general, the quantum effective action $\\Gamma[\\hat g_{\\hat\\mu\\hat\\nu},\\xi]$ for gravity coupled to a scalar field $\\xi$ has a complicated form. We explore here the possibility of a simple scaling limit for large $\\xi$ and small $\\hat g_{\\hat\\mu\\hat\\nu}$. More precisely, we investigate the scaling with a multiplicative power of a dimensionless parameter $\\kappa$ and define \\begin{equation}\\label{Z1} \\Gamma_\\kappa[\\hat g_{\\hat\\mu\\hat\\nu},\\xi]=\\Gamma[\\kappa^{-2}\\hat g_{\\hat\\mu\\hat\\nu},\\kappa^{\\frac{d-2}{2}}\\xi]. \\end{equation} Our hypothesis is a simple limiting behavior of $\\Gamma_\\kappa$ for $\\kappa\\to\\infty$. In particular, we investigate a possible fixed point \\cite{CWQ}, \\cite{CWNL} \\begin{eqnarray}\\label{Z2} \\lim_{\\kappa\\to\\infty}\\Gamma_\\kappa=\\int_{\\hat x}\\hat g^{1/2} \\left\\{-\\frac12\\xi^2\\hat R+\\frac{\\zeta}{2} \\partial^{\\hat \\mu}\\xi\\partial_{\\hat \\mu}\\xi\\right\\}. \\end{eqnarray} As appropriate for a fixed point, eq. \\eqref{Z2} states that $\\Gamma_\\kappa$ becomes independent of $\\kappa$ for $\\kappa\\to\\infty$. In this limit the effective action is therefore invariant with respect to dilatations.  If a fixed point \\eqref{Z2} exists, it may be used for a definition of quantum gravity, with a non-perturbative renormalization similar to the asymptotic safety scenario \\cite{25}, \\cite{26}. Quantum fluctuations are expected to modify the simple asymptotic form of the effective action by adding for finite $\\kappa$ additional terms to the r.h.s. of eq. \\eqref{Z2}. Such terms typically violate the dilatation symmetry and will then be treated as ``dilatation anomalies''. An example would be a cosmological constant \\begin{equation}\\label{Z3} \\Gamma^{(\\bar\\mu)}=\\int_{\\hat x} \\hat g^{1/2}\\bar\\mu^d. \\end{equation} Its contribution to $\\Gamma_\\kappa$ indeed vanishes for $\\kappa\\to\\infty$ \\begin{equation}\\label{Z4} \\Gamma^{(\\bar\\mu)}_\\kappa=\\kappa^{-d}\\int_{\\hat x}\\hat g^{1/2}\\bar\\mu^d. \\end{equation}  One may imagine a non-linear functional flow equation for the $\\kappa$-dependence of $\\Gamma_\\kappa$ (at fixed $\\hat g_{\\hat\\mu\\hat\\nu},\\xi$ and besides the trivial linear one following from the definition \\eqref{Z1}), \\begin{equation}\\label{Z5} \\kappa\\partial_\\kappa\\Gamma_\\kappa[\\hat g_{\\hat\\mu\\hat\\nu},\\xi]= {\\cal F}[\\hat g_{\\hat\\mu\\hat\\nu},\\xi], \\end{equation} with a functional ${\\cal F}$ typically involving $\\Gamma_\\kappa$ and its functional derivatives. One could then define quantum gravity by a solution of this flow equation with the ``boundary condition'' or ``initial value'' \\eqref{Z2}. The fixed point would be an ``ultraviolet fixed point'' in the sense of $\\kappa\\to\\infty$. Effective actions which are close to this fixed point for finite $\\kappa$ may then be treated in the standard renormalization group formalism for relevant and marginal directions.  The fixed point may become important for late time cosmological solutions. This happens if $\\xi(t)$ increases and $\\hat g^{1/2}(t)$ decreases for increasing time, such that the field equations, which obtain from the functional derivatives of $\\Gamma$ with respect to $\\hat g_{\\hat\\mu\\hat\\nu}$ and $\\xi$, have to be evaluated for values of the fields where the limit $\\Gamma_{\\kappa\\to\\infty}$ becomes relevant. In sect. \\ref{cosmologicalrunaway} we have discussed an explicit example for this type of ``runaway cosmology''. The approach to the fixed point leads to an interesting candidate for dynamical dark energy, where the potential energy of an appropriate scalar field vanishes only asymptotically for $t\\to\\infty$.  The fixed point \\eqref{Z2} may exist for arbitrary dimension $d>2$, including $d=4$. In this paper we are interested in higher dimensional theories. Indeed, for $d>6$ no polynomial potential for $\\xi$ is consistent with dilatation symmetry. This may add to the plausibility of a fixed point effective action without a potential $V(\\xi)$. On the other hand, the vanishing of the higher dimensional cosmological constant is, in general, not sufficient to guarantee a vanishing effective four-dimensional cosmological constant after ``spontaneous compactification'' of the internal dimensions. The curvature of internal space or a non-trivial warping may generate such an effective cosmological constant. One of the important findings of the present paper is a statement about the most general ``compactified'' higher dimensional solutions of the field equations derived from the fixed point action \\eqref{Z2}: all stable solutions with maximal four-dimensional symmetry and finite effective four-dimensional gravitational constant must have a vanishing four-dimensional cosmological constant. This asymptotic vanishing of the cosmological constant can be understood by a self-adjustment mechanism for infinitely many degrees of freedom, as discussed in sect. \\ref{adjustment}.  Consider now stable cosmological runaway solutions which drive the fields into a region where the fixed point effective action \\eqref{Z2} becomes relevant, and which allow for an effective four-dimensional gravity. Such solutions lead to a vanishing four-dimensional cosmological constant for $t\\to\\infty$, thus solving the ``cosmological constant problem''. Furthermore, since the Universe is very old (in units of the Planck time), but not infinitely old, a small amount of homogeneous energy density remains present in the effective four-dimensional Universe. This could give an explanation for the observed dark energy. Due to the huge age of the Universe (in Planck units) we expect that the present geometry of the Universe is very close to the asymptotic geometry, which is a quasi static solution of the field equations derived from eq. \\eqref{Z2}.  It is easy to find candidates for such asymptotic solutions - the simplest being a direct product of four-dimensional flat space and an internal $d-4$-dimensional Ricci-flat space with finite volume, accompanied by a constant value of the scalar field $\\xi$. If internal space admits isometries, the gauge couplings of the resulting four-dimensional gauge interactions become time-independent for large $t$. This solution has a vanishing four-dimensional cosmological constant $\\Lambda$ and exists for all values of $\\zeta$ in eq. \\eqref{Z2}. On the other hand, we have shown that all possible stable quasistatic solutions of the field equations derived from the effective action \\eqref{Z2}, which are consistent with four dimensional gravity, must have a vanishing four-dimensional constant $\\Lambda=0$. Stable asymptotic solutions approaching de Sitter or anti-de Sitter space $(\\Lambda\\neq 0)$ are not possible. This singles out a vanishing cosmological constant for the asymptotic behavior for $t\\to\\infty$ and may be the basis for the solution of the ``cosmological constant problem''.  Furthermore, we have extended our discussion to a more general form of a fixed point effective action. It includes dilatation symmetric higher curvature invariants for the metric. In the absence of a potential for $\\xi$ our conclusions remain essentially unchanged. Finally, we have presented a simple estimate of the dilatation anomaly and discussed the qualitative behavior of the corresponding runaway cosmology. It approaches a zero cosmological constant for infinite time, and has a dynamical dark energy component of similar size as dark matter.  Our main conclusion is simple: if the quantum effective action for large $\\xi$ and small $\\hat g^{1/2}$ shows a limiting fixed point behavior \\eqref{Z2}, or a similar dilatation symmetric  behavior without a potential term for $\\xi$, and if a stable cosmological runaway solution approaches this field region for large $t$, the cosmological constant problem can be solved without any fine tuning of parameters or initial conditions.   \\LARGE"
    },
    "0911/0911.1549_arXiv.txt": {
        "abstract": "The theory of radiative transfer provides the link between the physical conditions in an astrophysical object and the observable radiation which it emits. Thus accurately modelling radiative transfer is often a necessary part of testing theoretical models by comparison with observations.  We describe a new radiative transfer code which employs Monte Carlo methods for the numerical simulation of radiation transport in expanding media. We discuss the application of this code to the calculation of synthetic spectra and light curves for a Type Ia supernova explosion model and describe the sensitivity of the results to certain approximations made in the simulations. ",
        "introduction": "Since almost all we know about astronomical objects is inferred from the radiation which they emit, the theory of radiation transport often has a key role in testing our understanding of astrophysics. Although there are a variety of competitive approaches used for radiative transfer simulations, Monte Carlo methods are particularly well-suited for many modern astrophysical applications. In the Monte Carlo approach, the radiation field is discretized into quanta which represent bundles of photons. By propagating these quanta through a model of an astrophysical object, and simulating their interactions, synthetic spectra and light curves can be obtained. This method has the particular advantage that matter-radiation interactions are always treated locally meaning that multi-dimensionality, time-dependence and large-scale velocity fields can all be incorporated readily. ",
        "conclusions": ""
    },
    "0911/0911.0980_arXiv.txt": {
        "abstract": "We simulate the formation and evolution of galaxies with a self-consistent 3D hydrodynamical model including star formation, supernova feedback, and chemical enrichment. Hypernova feedback plays an essential role not only in solving the [Zn/Fe] problem, but also reproducing the cosmic star formation rate history and the mass-metallicity relations. In a Milky-Way type galaxy, kinematics and chemical abundances are different in bulge, disk, and thick disk because of different star formation histories and the contribution of Type Ia Supernovae. ",
        "introduction": "While the evolution of the dark matter is reasonably well understood, the evolution of the baryonic component is much less certain because of the complexity of the relevant physical processes, such as star formation and feedback. With the commonly employed, schematic star formation criteria alone, the predicted star formation rates (SFRs) are higher than what is compatible with the observed luminosity density.  Thus feedback mechanisms are in general invoked to reheat gas and suppress star formation. Supernovae inject not only thermal energy but also heavy elements into the interstellar medium, which can enhance star formation. Chemical enrichment must be solved as well as energy feedback. ``Feedback'' is also important for solving the angular momentum problem and the missing satellite problem, and for explaining the existence of heavy elements in intracluster medium and intergalactic medium, and the mass-metallicity relation of galaxies (\\cite{kob07}, hereafter K07).  In the next decade, high-resolution multi-object spectroscopy (HERMES) and space astrometry mission (GAIA) will provide kinematics and chemical abundances of a million stars in the Local Group. Since different heavy elements are produced from different supernovae with different timescales, elemental abundance ratios can provide independent information on ``age''. Therefore, stars in a galaxy are fossils to untangle the history of the galaxy. The galactic archeology technique can be used to study the galaxy formation and evolution in general. Metallicities are measured in various objects with different galaxy mass scale and as a function of redshift/time. The internal structure of galaxies have being observed with integral field spectrographs (e.g., the SAURON project, SINFONI on VLT). Chemodynamical simulations can provide useful predictions and physical interpretation of these observations.  \\vspace*{-5mm} ",
        "conclusions": ""
    },
    "0911/0911.2727_arXiv.txt": {
        "abstract": "{ In recent years it has emerged that the high energy behavior of gravity could be governed by an ultraviolet non-Gaussian fixed point of the (dimensionless) Newton's constant, whose behavior at high energy is thus {\\it antiscreened}. This phenomenon has several astrophysical implications. In particular in this article recent works on renormalization group  improved cosmologies based upon a renormalization group  trajectory of Quantum Einstein Gravity with realistic parameter values will be reviewed. It will be argued that quantum  effects can account for the entire entropy of the present Universe in the massless sector and give rise to a phase of inflationary expansion. Moreover the prediction for the final state of the black hole evaporation is a Planck size remnant which is formed in an infinite time. }  \\FullConference{Workshop on Continuum and Lattice Approaches to Quantum Gravity\\\\ September 17-19 2008\\\\ Brighton, University of Sussex, United Kingdom}  \\begin{document} ",
        "introduction": "Cosmology is a natural setting to study quantum gravity, which  may provide answers to fundamental questions as why is the expansion of the universe isotropic, can the initial singularity be avoided, why does the vacuum energy ``gravitate'' so little (Cosmological Constant problem)?    In recent years it has emerged that the asymptotic safety scenario \\cite{wein2,2009AnPhy.324..414C,niedermaier_asymptotic_2006} could provide the right framework to address the above questions. According to this approach the  ultraviolet (UV) behavior of quantum gravity is controlled by a fixed point at a non-zero value of the (dimensionless) coupling constant, so that the dimensionful Newton's constant reduces its strength at higher energies, it is thus {\\it antiscreened}. The non-perturbative renormalization group (RG) equation employed in this investigation predicts that the dimensionless cosmological constant reaches a non-gaussian fixed point (NGFP) in the infinite cutoff limit, so that the full Einstein-Hilbert Lagrangian is renormalizable at a non-perturbative level around this fixed point.  The gravitational {antiscreening} behavior is very similar to the running of the non-Abelian gauge coupling in Yang-Mills Theory, but only after the introduction of the effective average action and its functional renormalization group equation for gravity \\cite{1998PhRvD..57..971R} detailed investigations of the scaling behavior of the Newtons's constant have become possible \\cite{1998PhRvD..57..971R,1998CQGra..15.3449D,2002PhRvD..65b5013L,2002PhRvD..66b5026L,2002CQGra..19..483L, 2002PhRvD..65f5016R,2002PhRvD..66l5001R,1999PThPh.102..181S,2003PhRvD..67h1503P,2004PhRvL..92t1301L, 2005JHEP...02..035B,2006PhRvL..97v1301C,2009GReGr..41..983R,2009PhRvD..79j5005R,2009arXiv0904.2510M,2009arXiv0905.4220M}. The non-perturbative renormalization group equation underlying this approach defines a Wilsonian RG flow on a theory space which consists of all diffeomorphism invariant functionals of the metric $g_{\\mu\\nu}$.  This framework  turned out to be an ideal setting for investigating the asymptotic safety scenario in gravity \\cite{wein2,2009AnPhy.324..414C,niedermaier_asymptotic_2006} and, in fact, substantial evidence was found for the non-perturbative renormalizability of Quantum Einstein Gravity. The theory emerging  from this construction (``QEG\") is not a quantization of classical general relativity. Instead, its bare action corresponds to a nontrivial fixed point of the RG flow and therefore is a {\\it prediction}.  The effective average action \\cite{1998PhRvD..57..971R,2002PhR...363..223B} has crucial advantages as compared to other continuum implementations of the Wilson RG, in particular it is closely related to the standard effective action and defines a family of effective field theories $\\{ \\Gamma_k[g_{\\mu\\nu}], 0 \\leq k < \\infty \\}$ labeled by the coarse graining scale $k$. The latter property opens the door to a rather direct extraction of physical information from the RG flow, at least in single-scale cases: If the physical process or phenomenon under consideration involves only a single typical momentum scale $p_0$ it can be described by a tree-level evaluation of $\\Gamma_k[g_{\\mu\\nu}]$, with $k=p_0$. The precision which can be achieved by this effective field theory description depends on the size of the fluctuations relative to the mean values. If they are large, or if more than one scale is involved, it might be necessary to go  beyond the tree analysis.  The qualitative scale dependence of Newton's constant can be grasped with the help of the following physical argument. Let us imagine that in the large distance limit the leading quantum effects of the geometry are described by quantizing the linear fluctuations of the metric, $g_{\\mu \\nu}$. The resulting theory is a minimallu coupled theory in a curved background spacetime whose elementary quanta, the gravitons, carry energy and momentum. The vacuum of this theory will be populated by virtual graviton pairs, and the problem is to understand how these virtual gravitons respond to the perturbation by an external test body which we immerse in the vacuum. Assuming that also in this situation gravity is universally attractive, the gravitons will be attracted towards the test body. It will thus become ``dressed\" by a cloud of virtual gravitons surrounding it so that its effective mass seen by a distant observer is larger than it would be in absence of any quantum effects. This means that while in QED the quantum fluctuations {\\it screen} external charges, in quantum gravity they have an {\\it antiscreening} effect on external test masses. The consequence of this simple {\\it Gedanken} experiment entails Newton's constant becoming a scale dependent quantity $ G\\!\\left(k \\right)$ which is small at small distances $ r\\sim 1/k$, and which becomes large at larger distances.  In QED the {screening} behavior is well-known but it is interesting to recall how this result is obtained from the \"renormalization group improvement\", a standard device, in particle physics, in order to add the dominant quantum corrections to the Born approximation of a scattering cross section for instance. One starts from the classical potential energy $V_{\\rm cl}(r) = e^2/4\\pi r$ and replaces $e^2$ by the running gauge coupling in the one-loop approximation: \\begin{equation}\\label{uel} e^2(k)=e^2(k_0)[1-b\\ln(k/k_0)]^{-1},\\;\\;\\;\\;\\; b\\equiv e^2(k_0)/6\\pi^2. \\end{equation} The crucial step is to identify the renormalization point $k$ with the inverse of the distance $r$ so that result of this substitution reads \\be\\label{rea} V(r)=-e^2(r_0^{-1})[1+b\\ln(r_0/r) + O(e^4)]/4\\pi r \\ee where the IR reference scale $r_0\\equiv 1/k_0$ has to be kept finite in the massless theory. We emphasize that eq.(\\ref{rea}) is the correct (one-loop, massless) Uehling potential which is usually derived by more conventional perturbative methods \\cite{dittrich_effective_1985}. Obviously the position dependent renormalization group improvement $e^2\\rightarrow e^2(k)$, $k\\propto 1/r$ encapsulates the most important effects which the quantum fluctuations have on the electric field produced by a point charge.  The effective field theory techniques proved useful for an understanding of the scale dependent geometry of the effective QEG spacetimes \\cite{2005JHEP...10..050L,2006JHEP...01..070R,2007JHEP...01..049R}. In particular it has been shown \\cite{2002PhRvD..65b5013L,2005JHEP...10..050L} that these spacetimes have fractal properties, with  a fractal dimension of 2 at small, and 4 at large distances. The same dynamical dimensional reduction was also observed in numerical studies of Lorentzian dynamical triangulations \\cite{2005PhLB..607..205A,2004PhRvL..93m1301A,2005PhRvL..95q1301A} and in \\cite{2006JHEP...11..081C} A.Connes et al. speculated about its possible relevance to the non-commutative geometry of the standard model.  In order to extract all the relevant information from the RG evolution, it is thus necessary to relate the cutoff scale $k$ which corresponds to the resolution of the RG flow,  to the spacetime properties. This  procedure is called ``cutoff identification\" for which the relevant energy scale $k$ is related to a characteristic length scale where the quanta with energy $k$ propagate.  In the case of massless QED the choice $k\\propto 1/r$ was clearly the only possible one, as there are no other relevant scales in the problem. When several scales are present the prescription which emerges from the general theory of the Effective Average Action \\cite{2002PhR...363..223B} is that $\\Gamma_k$ is defined at a scale $k$ which is the {\\it largest} one of the various competing scales in the fluctuation determinant of the Average Action, $\\Gamma^{(2)}_k$, namely \\be\\label{flucd} \\Gamma^{(2)}_k=\\frac{\\delta^2 \\Gamma_k}{\\delta \\Phi^2 } \\ee where $\\Phi$ is the so-called ``blocked''  field \\cite{1995PhRvD..52..969B}.  The  difficulty arises when we decide to apply the same ``recipe'' in gravity by writing \\be\\label{klen} k \\sim 1/\\ell(x^\\mu)), \\;\\;\\;\\;\\;\\;\\;\\;\\;\\;\\;   \\ell=\\ell(g_{\\mu\\nu}) \\ee being $\\ell$ a characteristic length where the fluctuations with energy $k$ propagate.  The reason is that the flow equation is by construction diffeomorphism invariant at any $k$ so that the RG flow itself does not know anything about the background field metric $g_{\\mu\\nu}$ that has been used for projecting on a finite-dimensional subspace of the ``theory space''.  There are two possible strategies to overcome this issue. The first one amounts to choose a fiducial metric which is a {\\it solution} of the Einstein equations and RG-improve it by substituting the Newton constant $G$ with the running $G(k)$ together with a cutoff identification of the type (\\ref{klen}). The limitation of this approach lies in the fact that in general the improved metric may not be a solution of Einstein equation, but one can imagine that this is a sort of ``Thomas-Fermi'' approximation where only the leading quantum corrections are  taken into account \\cite{2000PhRvD..62d3008B,2006PhRvD..73h3005B}. The improved $g_{\\mu\\nu}(k)$ metric represents then a sort of ``emergent'' spacetime description of the effective geometry \\cite{2005LRR.....8...12B,2004CQGra..21.1725M,2007arXiv0712.0810V}  according to the scale dependence of the Newton constant.  A second possibility is to consider the energy scale $k$ associated to the field strength itself rather than to an observational scale $\\ell$. This is motivated by the analogy with the QED (and QCD) case, where higher loop contributions to the Uehling potential are obtained by renormalization group improvement of the QED {\\it action} by using the field strength $(F_{\\mu\\nu}F^{\\mu\\nu})^{1/4}$  as a cutoff instead than $1/r$ \\cite{1973NuPhB..52..483M,1978NuPhB.134..539M,1978NuPhB.143..485P}. In this case the short distance correction to the static potential is obtained from the non-linear differential equations \\begin{eqnarray}\\label{maty} &&\\nabla \\cdot  \\mathbf D = J_0, \\;\\;\\;\\;\\; \\mathbf D = \\mathbf E \\; \\epsilon (E) , \\;\\;\\; \\mathbf E = -\\nabla A^0\\\\[2mm]\\nonumber && \\epsilon(E)=1-\\frac{e^2}{12\\pi^2} \\log (e E/k_0^2)+\\ldots \\end{eqnarray} whose solution reproduces the Uehling potential in the long distance limit, but in general the solutions of Eq.(\\ref{maty})  include higher loop effects due to the non-linearities of the effective action in the short distance limit.  The two approaches discussed above are obviously related, at least in some limit. In the case of Robertson-Walker spaces it will be shown that due to the very high degree of symmetry of the spacetime, the time-scale defined by ``Hubble parameter''  behaves essentially like the characteristic time scale associated to the relevant curvature invariants\\cite{2004ForPh..52..650R,2005JCAP...09..012R,2007JCAP...08..024B}. In the case of spherically symmetric spacetimes \\cite{2000PhRvD..62d3008B,2004JCAP...12..001R,2004PhRvD..70l4028R,2006PhRvD..73h3005B}, near the singularity the proper distance of a radially free falling observer behaves essentially as $1/\\sqrt{\\Psi_2}$, being $\\Psi_2$ the ``Coulombian'' component of the Weyl tensor.  From the above discussion it is then clear that in general there is not a preferred strategy to perform the RG improvement in gravity. In some case it might  be  more interesting to RG improve {\\it  solutions} and to make contact with an emergent spacetime description of the effective geometry. In some other cases it could be more convenient to work with a RG improvement at the level of {\\it field equations} or {\\it actions} \\cite{2004PhRvD..69j4022R,2004PhRvD..70l4028R,2004CQGra..21.5005B,2005IJMPA..20.2358B}.  It is important to remark that it is not surprising  that different cutoff might provide quantitavely different evolutions, as usually the $\\beta$-functions are not ``universal'' quantities. For instance it is well known \\cite{Goldenfeld:1992qy} that different realizations of the block-spin RG transformation applied to the Ising model may provide different values for fixed points, as we are essentially using different type of ``microscopes''\\footnote{From this point of view the criticism expressed in \\cite{2008PhRvL.101h1301W} should not be seriously considered.}. On the other hand truly universal quantities, like the critical exponents, are essentially insensitive to the cutoff choice.   In this review recent results obtained with the RG improvement of Einstein theory will be discussed in the framework QEG will be reviewed, with particular emphasis on recent results obtained in cosmology \\cite{2007JCAP...08..024B}. In particular in Sec.2 the RG evolution of the Newton constant and Cosmological constant describing our Universe are described. In Sec.3 a covariant formalism to improve the Einstein field equation is presented while in Sec.4 the RG improved Robertson-Walker Cosmology is discussed. In Sec.5 the basic mechanism to produce the entropy of the Universe is presented. In Sec.6 the properties of a class of solutions of the RG equations are discussed. In Sec.7 a mechanism to produce a power-law inflation is studied. In Sec.8 and Sec.9 the properties of RG improved Black Hole metric is studied. In Sec.10 the possibility that Quantum Gravity effects are present on Astrophysical distances is reviewed. Sec.11 is devoted to the Conclusions. ",
        "conclusions": "In these notes some important astrophysical consequences of the Asymptotic Safety Scenario have been reviewed.  In particular it was advocated the point of view that the scale dependence of the gravitational parameters has an impact on the physics of the Universe we live in and, in particular, it has been possible to identify known features of the Universe which could possibly  be due to this scale dependence. Three possible candidates for such features are proposed: the entropy carried by the radiation which fills the Universe today, a period of automatic, $\\Lambda$-driven inflation that requires no ad hoc inflaton, and the primordial density perturbations.  Moreover, the impact of the leading quantum gravity effects on the dynamics of the Hawking evaporation process of a black hole have also been investigated. Its spacetime structure is described by a renormalization group improved Vaidya metric. Its event horizon, apparent horizon, and timelike limit surface have been obtained taking the scale dependence of Newton's constant into account. The emergence of a quantum ergosphere is discussed. The final state of the evaporation process is a cold, Planck size remnant.  It would be interesting to investigate the possible astrophysical implications of a population  of stable Planck size mini-black holes produced in the Early Universe or by the interaction of cosmic rays with the interstellar medium. I hope to address this issue in a subsequent publication."
    },
    "0911/0911.2511_arXiv.txt": {
        "abstract": "We compare H$\\alpha$, ultraviolet (UV) and infrared (IR) indicators of star formation rate (SFR) for a well-defined sample of $z=0.84$ emission-line galaxies from the High-$z$ Emission Line Survey (HiZELS).  Using emission-line, optical, IR, radio and X-ray diagnostics, we estimate that 5 -- 11~per~cent of H$\\alpha$ emitters at this redshift are active galactic nuclei.  We detect 35~per~cent of the H$\\alpha$ emitters individually at 24~$\\mu$m, and stack the locations of star-forming emitters on deep 24-$\\mu$m {\\it Spitzer Space Telescope} images in order to calculate the typical SFRs of our H$\\alpha$-emitting galaxies.  These are compared to the observed H$\\alpha$ line fluxes in order to estimate the extinction at $z=0.84$, and we find a significant increase in dust extinction for galaxies with higher SFRs.  We demonstrate that the relationship between SFR and extinction found in the local Universe is also suitable for our high-redshift galaxies, and attribute the overall increase in the typical dust extinction for $z=0.84$ galaxies to an increase in the average SFR, rather than to a change in dust properties at higher redshift.  We calculate the UV extinction, and find a similar dependence on SFR to the H$\\alpha$ results, but no evidence for a 2175~\\AA\\ UV bump in the dust attenuation law for high-redshift star-forming galaxies.  By comparing H$\\alpha$ and UV indicators, we calculate the conversion between the dust attenuation of nebular and stellar radiation, $\\gamma$, and show that $\\gamma=0.50\\pm0.14$.  The extinction / SFR relationship is shown to be applicable to galaxies with a range of morphologies and bulge-to-disk ratios, to both merging and non-merging galaxies, and to galaxies within high- and low-density environments, implying that it is a fundamental property of star-forming regions.  In order to allow future studies to easily correct for a SFR-dependent amount of dust extinction, we present an equation to predict the extinction of a galaxy, based solely upon its observed H$\\alpha$ luminosity, and use this to recalculate the H$\\alpha$ luminosity function and star formation rate density at $z=0.84$. ",
        "introduction": "\\label{sec:introduction} There are many methods of estimating the global star formation rate (SFR) of a galaxy, from measurements of the ultra-violet (UV) emission coming directly from visible young massive stars, to the bolometric infrared (IR) luminosity resulting from dust-reprocessed stellar emission \\citep[see][for a review of various star formation indicators]{Kennicutt98}.  All SFR indicators suffer from different biases and uncertainties, and while SFRs estimated in different ways in the local Universe are now in general agreement, there is much greater scatter in the results at $z\\sim1$ \\citep[see, e.g.][]{Hopkins06}.  Locally one of the best measures of the current rate of star formation within a galaxy comes from observations of the H$\\alpha$ transition of atomic hydrogen.  The luminosity of this recombination line scales directly with the rate of ionising flux from young, massive stars and gives an excellent indication of the presence of ongoing star formation.  However, its usefulness as a quantitative SFR indicator is reduced by the effect of dust obscuration; observed H$\\alpha$ line fluxes need to be corrected for extinction before the intrinsic H$\\alpha$ luminosity can be recovered, and the SFR calculated.  The initial method used to estimate H$\\alpha$ extinction came from a comparison of thermal radio emission, which is unaffected by dust attenuation, and the observed H$\\alpha$ + [N\\,{\\sc ii}] line emission within local galaxies \\citep{Kennicutt83}.  Disentangling thermal and non-thermal radio emission is non-trivial \\citep[see, e.g.][for more details]{Condon92}, as thermal radio emission is only a small fraction of the total radio luminosity of a galaxy \\citep[$\\lesssim10$~per~cent at 1.4~GHz;][]{Condon90}, and so this method is not often used in practise.  A much more common method to estimate extinction in individual galaxies comes from the `Balmer decrement', comparing the observed value of the H$\\alpha$/H$\\beta$ line flux ratio to the intrinsic value when no dust is present.  This allows a calculation of the reddening between the wavelengths of the H$\\alpha$ and H$\\beta$ lines and, with a choice of dust attenuation model \\citep[e.g.][]{Calzetti00}, permits an estimate of the attenuation at a given wavelength \\citep*[see e.g.][]{Moustakas06}.  However, this technique requires sensitive spectroscopy to obtain the line flux ratio with a good signal-to-noise ratio (SNR), which is time-consuming for large samples of galaxies at moderate redshifts.  It is therefore more usual for studies with limited available photometry to simply assume that a single extinction can be applied to all galaxies.  The canonical amount of H$\\alpha$ extinction for a typical galaxy is usually taken to be $A_{\\rm H\\alpha} \\sim1$~mag \\citep[e.g.][]{Kennicutt83,Bell01, Pascual01,Charlot02,Fujita03,Brinchmann04,Ly07,Geach08,Sobral09Halpha}, based upon measurements of extinction in the local Universe.  At higher redshift, the appropriate H$\\alpha$ extinction for a typical galaxy is not well determined.  Significant numbers of extremely dusty galaxies have been discovered at high redshift \\citep[e.g.][]{Chapman05}, and the large dust content implies a greater H$\\alpha$ attenuation.  The global rate of star formation increases out to $z\\sim1$ \\citep[e.g.][]{Hopkins06}, and galaxies with higher SFR are thought to undergo greater extinction \\citep{Hopkins01, Sullivan01, Berta03}, also implying that the typical H$\\alpha$ attenuation will be greater at $z\\sim1$ than in the local Universe. \\citet{Villar08} have estimated that the typical H$\\alpha$ extinction of galaxies within their $z=0.84$ sample is $\\sim1.5$~mag.  The total IR luminosity of a galaxy is a bolometric measurement of the thermal emission from dust grains, which have been heated by the UV and optical radiation produced by young massive stars \\citep[e.g.][]{Kennicutt98}.  For star-forming galaxies with a reasonable dust content, this emission gives a direct measurement of the total stellar power output.  However, since longer wavelength IR emission will more closely trace the cool dust component within a galaxy, which predominantly results from heating by an old stellar component rather than from recent star formation, it could be argued that mid-IR emission might be a better SFR indicator than the bolometric IR luminosity \\citep[e.g.][]{Dopita02}.  \\citet{Rieke09} have demonstrated that the flux density measured in the {\\it Spitzer Space Telescope} 24-$\\mu$m band can be used as a good proxy for SFR for galaxies at $z=0$, and, after applying suitable corrections, has a comparable accuracy to estimating SFRs from the extinction-corrected Pa$\\alpha$ luminosity or total IR luminosity for normal galaxies in the redshift range $0<z<2$.  Observations at 24~$\\mu$m have the additional advantage over longer wavelength IR observations in that they are more sensitive (thus allowing fainter, more distant galaxies to be studied), and have a higher resolution ($\\sim6$~arcsec, making it relatively straightforward to identify counterparts at other wavelengths; the 70- and 160-$\\mu$m {\\it Spitzer} bands only have resolutions of $\\sim20$ and $\\sim40$~arcsec respectively).  These disadvantages will be overcome to some extent with {\\it Herschel}, which should be able to make sensitive $\\sim6$~arcsec resolution images of galaxies at 70~$\\mu$m, close to the peak in luminosity for the spectral energy distribution (SED) of a typical galaxy at $z\\sim0$.  In this work, we compare H$\\alpha$, UV and IR SFR indicators in order to estimate the extinction for a well-defined sample of galaxies at $z=0.84$, using data from the High-$z$ Emission Line Survey \\citep[HiZELS;][]{Geach08, Sobral09Halpha, Sobral09Lyalpha}.  HiZELS is a panoramic extragalactic survey making use of the wide area coverage of the Wide Field Camera on the 3.8-m UK Infrared Telescope (UKIRT).  Narrow-band observations of the Cosmological Evolution Survey field \\citep[COSMOS;][]{Scoville07overview}, and the Subaru-XMM-UKIDSS Ultra Deep Survey field \\citep[UDS;][]{Lawrence07} have been taken in the $J$, $H$ and $K$ bands in order to target H$\\alpha$ emitters at $z=0.84$, $z = 1.47$ and $z=2.23$ respectively; this paper concentrates on the properties of the $z=0.84$ H$\\alpha$ sample, the creation of which is described in full in \\citet{Sobral09Halpha}.  In Section~\\ref{sec:sampleselection} we present details of the HiZELS sample selection, and the available multi-wavelength data in the COSMOS and UDS fields.  In Section~\\ref{sec:agn} we describe various methods which we use to identify active galactic nuclei (AGN) contaminants, and quantify the fraction of H$\\alpha$-emitting sources which may be AGN at $z=0.84$.  We describe our method for stacking galaxies and measuring typical 24-$\\mu$m flux densities in Section~\\ref{sec:results}, and convert our stacking results into SFRs in order to estimate the H$\\alpha$ extinction that is required to make this indicator agree with the IR-based SFRs.  In Section~\\ref{sec:discussion} we discuss our results, investigate the effects that a variable extinction has on the H$\\alpha$ luminosity function at $z=0.84$, and present a method for correcting observed H$\\alpha$ luminosities for a SFR-dependent amount of extinction. Section~\\ref{sec:conclusions} reports our conclusions.  Throughout this work we assume a concordance cosmology of $\\Omega_{\\rm M}=0.3$, $\\Omega_{\\Lambda}=0.7$ and H$_{0}=70$~km~s$^{-1}$~Mpc$^{-1}$. All magnitudes are given in the AB system \\citep{Oke83}. ",
        "conclusions": "\\label{sec:conclusions} We have used a well-defined sample of H$\\alpha$-emitting galaxies at $z=0.84$ to study the link between star formation and dust extinction. The HiZELS sample consisted of 694 H$\\alpha$-emitting galaxies within the redshift range $z=0.845\\pm0.021$ in the well-studied COSMOS and UDS fields, with additional information available in multiple UV, optical and IR bands, along with ancillary radio and X-ray data.  In order to focus on star-forming galaxies, it was necessary to remove AGN contaminants from our sample.  We performed this removal through a number of multi-wavelength classification methods; emission-line diagnostics, deviations away from the IR / radio correlation, X-ray counterparts, mid-IR and mid- to far-IR shapes and full template fitting were all used to identify potential AGN.  We estimate an AGN contamination rate of between 5 and 11~per~cent for H$\\alpha$-selected emission-line sources at $z=0.84$.  No single diagnostic identified more than half of the AGN contaminants, with SED template fitting being the most successful; we suggest that using a combination of template-fitting and colour-colour diagnostics will maximise the success rate of AGN identification in H$\\alpha$ surveys.  In order to constrain the average IR luminosity of H$\\alpha$ emitters at $z=0.84$, we have binned our sample by observed H$\\alpha$ flux, and used the technique of source stacking to calculate typical 24-$\\mu$m flux densities for the H$\\alpha$ emitters.  Several previous studies have found that dust extinction correlates with either luminosity or SFR, and we looked for a relationship between the SFR of a galaxy and its dust attenuation.  We found that individual extinctions vary significantly between galaxies, but confirm that this correlation holds for the average properties of our sample; galaxies with higher SFR also undergo a greater amount of extinction.  We have presented an approximate method of calculating the typical H$\\alpha$ extinction of a galaxy, based solely upon the observed H$\\alpha$ luminosity; $A_{\\rm H\\alpha} = -21.963 + 0.562~{\\rm log}_{10}(L_{\\rm H\\alpha, obs}/{\\rm erg~s^{-1}})$, which is within 10~per~cent of the full numerical solution over the luminosity range $10^{41} - 10^{44}$~erg~s$^{-1}$.  This is found to describe the properties of galaxies with disk, spiral and irregular morphologies, a range of bulge-to-disk ratios, galaxies in high- and low-density environments, and both merging and non-merging galaxies, although a potential deviation in this relationship may be seen for galaxies with low stellar masses.  We emphasize that this is only an average relationship, and that the extinction measured within individual galaxies can vary away from this relationship by up to $\\sim2$~mag.  We have calculated the far-UV extinction of galaxies in a similar manner to the H$\\alpha$ extinction, and found a qualitatively similar correlation between this extinction and the SFR.  We find that the UV dust obscuration is about half of that expected from an extrapolation of the \\citet{Calzetti00} dust attenuation law.  By calculating UV extinctions from flux densities measured in three different bands, we show that, similar to the local Universe, there is no evidence for a 2175~\\AA\\ UV bump in the dust attenuation of high-redshift star-forming galaxies.  The relationship between dust extinction and SFR at $z=0.84$ is the same as is seen at $z\\sim0$; we conclude that there is no evidence for a change in the dust properties of galaxies over this redshift range, and that any average increase in the extinction of galaxies at higher redshift is due to the overall increase in star formation activity out to $z\\sim2$.  Using our relationship between observed H$\\alpha$ luminosity and extinction, we recalculate the H$\\alpha$ luminosity function for galaxies at $z=0.84$.  We find an increase in the characteristic luminosity compared with the value obtained by assuming a single value of extinction is applicable to all sources, implying a greater number density of bright H$\\alpha$ emitters than previous estimates at $z=0.84$."
    },
    "0911/0911.3667_arXiv.txt": {
        "abstract": "We report observations of Lyman Alpha Blob 1 (LAB1) in the SSA 22 protocluster region ($z=3.09$) with the integral-field spectrograph SAURON. We increased the signal-to-noise in the spectra by more than a factor three compared to our previous observations. This allows us to probe the structure of the LAB system in detail, examining its structure in the spatial and wavelength dimensions. We find that the emission from the system comes largely from five distinct blobs. Two of the emission regions are associated with Lyman Break Galaxies, while a third appears to be associated with a heavily obscured submillimeter galaxy. The fourth and fifth components do not appear to be associated with any galaxy despite the deep imaging that is available in this field. If we interpret wavelength shifts in the line centroid as velocity structure in the underlying gas, many of these emission systems show evidence of velocity shear. It remains difficult to distinguish between an underlying rotation of the gas and an outflow driven by the central object.  We have examined all of the line profiles for evidence of strong absorption features. While several systems are better fitted by the inclusion of a weak absorption component, we do not see evidence for a large-scale coherent absorption feature such as that seen in LAB2.  \\date{} ",
        "introduction": "\\label{sec:introduction}  By allowing us to probe the gaseous haloes around $z \\sim 3$ galaxies, large scale Lyman $\\alpha$ nebulae provide a fascinating insight into the formation of high-redshift galaxies. The first and brightest of these haloes were discovered by Steidel et al. (2000)\\nocite{2000ApJ...532..170S} in the SSA 22 protocluster region at $z=3.09$. Subsequently a population of fainter Lyman Alpha Blobs (LABs) was detected in deep narrow-band imaging surveys (e.g. Matsuda et al. 2004\\nocite{2004AJ....128..569M}; Nilsson et al. 2006\\nocite{2006A&A...452L..23N}, Smith \\& Jarvis 2007\\nocite{2007MNRAS.378L..49S}), revealing that LABs have a large spread in properties such as surface brightness and morphology, and it has been suggested that their presence is linked to dense environments (Matsuda et al.  2004\\nocite{2004AJ....128..569M}). Follow-up observations in the optical, near-infrared and particularly the far-infrared suggest that LABs are sites of massive galaxy formation, enhanced by the cluster environment (e.g. Chapman et al. 2004\\nocite{2004ApJ...606...85C}).  This view is supported by the discovery of luminous submillimeter sources in several LABs (e.g. Geach et al. 2005\\nocite{2005MNRAS.363.1398G}). LAB systems like those discovered by Steidel et al. (2000)\\nocite{2000ApJ...532..170S} are also found around high-redshift radio galaxies (e.g. Chambers, Miley \\& van Breugel 1990\\nocite{1990ApJ...363...21C}; Villar-Mart\\'in et al. 2002\\nocite{2002MNRAS.336..436V}). However, the study of such radio loud systems is complicated by the presence of radio jets and lobes.  It is unclear whether these systems are directly comparable to the radio quiet LABs that we discuss in this paper, or that they have a different power source, such as the injection of cosmic rays by the radio source (e.g. Ferland et al. 2009\\nocite{2009MNRAS.392.1475F}).  The origin of radio quiet LABs is still unclear, and three different scenarios have been proposed to explain their existence. One is that the gas in LABs is heated by photo-ionisation, caused by massive stars and/or active galactic nuclei (AGN) (Geach et al. 2009\\nocite{2009ApJ...700....1G}). However, for one third of the LABs in the sample of Matsuda et al. (2004)\\nocite{2004AJ....128..569M}, the observed UV luminosities are too low to produce the observed \\Lya\\ radiation, although putative ionising sources could be obscured along our line of sight. For example, Geach et al. (2009)\\nocite{2009ApJ...700....1G} argue that in a large fraction of LABs there is sufficient UV flux from an obscured AGN component to power the extended line emission, despite large dust covering fractions. While photo-ionisation is likely to play a role in powering LABs, it is nonetheless instructive to consider alternative or additional power sources for the \\Lya\\ emission in LABs, since it is not clear that a single mechanism (such as photo-ionisation) is responsible for all of the observed properties of these objects.  An alternative scenario is that the gas in LABs is excited by cooling flows (e.g. Haiman, Spaans \\& Quataert 2000\\nocite{2000ApJ...537L...5H}; Fardal et al. 2001\\nocite{2001ApJ...562..605F}; Dijkstra \\& Loeb 2009\\nocite{2009arXiv0902.2999D}). Nilsson et al. (2006)\\nocite{2006A&A...452L..23N} argue that the emission in the LAB they detected in the GOODS South field originates from cold accretion onto a dark matter halo. This view is supported by the lack of continuum counterparts, which could photo-ionise the gas, and the absence of a massive starburst in the infrared, that could point to a superwind outflow (see below). Instead, they find a good match between their observed surface brightness profile and the profiles derived from theoretical models for collapsing clouds from Dijkstra, Haiman \\& Spaans (2006)\\nocite{2006ApJ...649...14D}. A similar analysis is given by Smith et al. (2008)\\nocite{2008MNRAS.389..799S} for the LAB presented in Smith \\& Jarvis (2007)\\nocite{2007MNRAS.378L..49S}.  The third proposed origin for extended \\Lya\\ emission haloes is provided by the so-called superwind model (e.g. Taniguchi \\& Shioya 2000\\nocite{2000ApJ...532L..13T}; Ohyama et al. 2003\\nocite{2003ApJ...591L...9O}). After an initial starburst, massive stars die in supernovae. If the resulting supernova remnants overlap, they could form a superbubble (e.g. Heckman, Armus \\& Miley 1990\\nocite{ 1990ApJS...74..833H}), from which a superwind can blow gas into the intergalactic medium if the kinetic energy in the gas is large enough to overcome the gravitational potential. Taniguchi \\& Ohyama (2000)\\nocite{2000ApJ...532L..13T} suggest an evolutionary sequence for elliptical galaxies that includes LABs. During the initial starburst, a galaxy can be enshrouded by gas and dust grains, and is therefore observable as a (dusty) submillimeter source.  The LABs represent the subsequent superwind phase, expelling the gas and dust, and therefore rendering the galaxy fainter in the submillimeter regime. After this phase, the galaxy continues to develop into a normal elliptical galaxy.  The brightest and most extended LAB observed to date is LAB1 (SSA22a-C11), one of the two LABs described by Steidel et al. (2000)\\nocite{2000ApJ...532..170S}. This LAB has a \\Lya\\ luminosity of $1.1 \\times 10^{44}$ ergs s$^{-1}$ at $z=3.1$ (Matsuda et al. 2004\\nocite{2004AJ....128..569M}), and a spatial extent of $\\sim$ 100 kpc. Matsuda et al. (2004)\\nocite{2000ApJ...532..170S} find bubble-like structures in narrowband images of LAB1, in support of the superwind model for this LAB. Bower et al. (2004)\\nocite{2004MNRAS.351...63B} observed LAB1 with the integral-field spectrograph SAURON (Bacon et al. 2001\\nocite{2001MNRAS.326...23B}). They found extensive emission, and a large velocity dispersion for the \\Lya\\ emission line ($\\sim$ 500 km s$^{-1}$). Although interpretation of the spectra of LAB1 is not straightforward, since \\Lya\\ is a resonant line and therefore emission can easily get scattered, they find similarities between LAB1 and the local emission-line halo of NGC 1275 in the Perseus cluster. Longer observations of LAB2 with the same spectrograph suggested the existence of a dense outflowing shell of material around this system (Wilman et al. 2005\\nocite{2005Natur.436..227W}). However, spectra of higher signal-to-noise are needed to find evidence for outflows in LAB1.  We therefore reobserved LAB1 with SAURON, adding signal to the data already published by Bower et al. (2004)\\nocite{2004MNRAS.351...63B}. By increasing the observing time from 9 hours in the original dataset to 23.5 hours in our new datacube, we increased the signal-to-noise in the spectra, and therefore were able to obtain line profiles and kinematic maps of the \\Lya\\ emission in this region. We describe the new observations and data reduction of LAB1 in Section~\\ref{sec:observations} and analyse the spectra in Section~3. In Section 4 we discuss our results and speculate on the structure and origin of LAB1. Throughout this paper, we assume a flat cosmology with $H_0 = 70$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega$=0.3 and $\\Lambda$ = 0.7. In this scenario, 1 arcsec at $z=3.1$ corresponds to 7.5 kpc.         \\begin{figure*} \\begin{center} \\psfig{figure=figures/bigplot_v6.ps,angle=90,width=18cm} \\caption{\\Lya\\ emission in the LAB1 region. Left panel: continuum subtracted \\Lya\\ emission, obtained from collapsing the SAURON spectra over a narrow wavelength range centred on the emission line. Interesting regions are indicated by boxes (see text). Middle panel: {\\it HST}/STIS optical image overlaid with \\Lya\\ contours from the left panel. The faintest contour has a surface brightness of 5.6 $\\times$ 10$^{-18}$ erg s$^{-1}$ cm$^{-2}$ arcsec$^{-2}$ and contours increase in steps of 3.7$\\times$ 10$^{-18}$ erg s$^{-1}$ cm$^{-2}$ arcsec$^{-2}$. Note however the 50 per cent uncertainty in absolute flux calibration (see \\S2). Right panel: same as middle panel, but now a {\\it Spitzer}/IRAC 3.6 $\\mu$m image is displayed. Sources identified in Geach et al. (2007) are indicated with identical nomenclature ($a$-$e$). In all plots, the blue dot denotes the position of the radio source (Chapman et al. 2004). All images are plotted on the same scale and are orientated such that North is up and East to the left.} \\label{fig:lya} \\end{center} \\end{figure*} ",
        "conclusions": "\\label{sec:conclusion}  We have presented very deep integral-field spectrograph observations of the LAB1 \\Lya\\ emission halo. The deep data allow us to study the spatial and velocity structure of this system in unprecedented detail.  We find that the giant halo is made up from a superposition of five distinct blobs, with the emission from these regions accounting for at least 55 per cent of the total diffuse flux. Most of the emission regions are associated with individual galaxies. The regions C15 and C11 are associated with optical Lyman Break Galaxies and region R3 is associated with a bright submillimeter source. While this source is faint in the optical, it is bright at 3.6 $\\mu$m, suggesting that it is a strongly dust-obscured starforming galaxy. None of these sources are detected at X-ray wavelengths, suggesting that it is unlikely that the emission is powered by an AGN. The regions R1 and R2 are not associated with any optical or IR source, and the emission from these blobs plausibly comes from gas associated with the proto-cluster potential.  The integral-field spectra allow us to examine the emission line profiles and velocity structure in each of the blobs. In R1, R3 and C11, we find that the integrated emission line profiles are not adequately fitted by a single component Gaussian, and that a better fit is obtained with either a two component Gaussian or the combination of a Gaussian emission and Voigt absorption line profile. We find that the underlying emission has a velocity centroid that is extremely similar from region to region, reinforcing the idea that the LAB1 system is a virialised group with velocity dispersion $\\sim100$ km s$^{-1}$. However, in contrast to the LAB2 system, we find no evidence for a coherent shell of absorption that covers the entire system. Any absorption features are significantly weaker than those seen in LAB2, so that they might arise in the large-scale structure foreground to the proto-cluster.  C15, C11 and possibly R1 show evidence of coherent velocity shear arising from an outflow or rotation. In C15 the velocity gradient is perpendicular to the morphology of the underlying galaxy, consistent with the pattern expected for an outflowing galactic wind. In the other systems the relation between the velocity field and the underlying galaxy is unclear. The velocity shear is largest in C11 where it is $\\sim 550$ km s$^{-1}$ over approximately 25 kpc.  The implied outflow velocity is comparable with that seen in many other Lyman break systems, and does not suggest that the sources seen in LAB1 are undergoing unusually strong feedback.  The primary motivation for these observations was to discover whether the coherent absorption systems such as seen by Wilman et al. (2005)\\nocite{2005Natur.436..227W} in LAB2 are a ubiquitous feature. These observations have shown that they are not.  The data for LAB2 are best interpreted as a large-scale super-bubble of material that has been expelled by a high power, perhaps explosive feedback event. The two LAB systems both seem to be made up of smaller emission clouds. So why does the absorption pattern in LAB1 and LAB2 differ? One possible answer is that a similar event has occured in LAB1 in the past, but that the shell has now broken up, as it passed through the intergalactic medium of this proto-cluster, or that it is sufficiently blue shifted that it cannot be seen in absorption against the system's \\Lya\\ emission. Another explanation might be that LAB1 is a younger system in which the large scale outflow is yet to develop. This seems unlikely, however, since we see no signs of spectacular outflows associated with the individual emission systems. Now that we have dissected the giant emission halo into a number of smaller systems, these seem quite comparable to the many radio quiet LAB systems identified by Matsuda et al. (2004\\nocite{2004AJ....128..569M}). In many ways, the composite system represents a microcosm of diffuse \\Lya\\ emission in general, with the systems reflecting a diversity of power sources. Geach et al. (2009)\\nocite{2009ApJ...700....1G} favour photo-ionisation as the principle power source for LABs. This fits in well with C11, C15 and R3 (if we allow for the possibility that it is only so strongly obscured along our line of sight). However, this appears to describe the R1/R2 emission less well. Possibly the emission from this part of the system is much more closely related to the emission seen around radio galaxies (Chambers et al. 1990\\nocite{1990ApJ...363...21C}; Villar-Mart\\'in et al. 2002\\nocite{2002MNRAS.336..436V}) and (perhaps) in local cooling flow clusters (Johnstone \\& Fabian 1988\\nocite{1988MNRAS.233..581J})."
    },
    "0911/0911.0791_arXiv.txt": {
        "abstract": "{{\\bf stellar evolution; population synthesis; spectral evolution; simple stellar populations; TP--AGB stars}} In this paper, I review to what extent we can understand the photometric properties of star clusters, and of low-mass, unresolved galaxies, in terms of population synthesis models designed to describe `simple stellar populations' (SSPs), i.e., groups of stars born at the same time, in the same volume of space, and from a gas cloud of homogeneous chemical composition. The photometric properties predicted by these models do not readily match the observations of most star clusters, unless we properly take into account the expected variation in the number of stars occupying sparsely populated evolutionary stages, due to stochastic fluctuations in the stellar initial mass function. In this case, population synthesis models reproduce remarkably well the full ranges of observed integrated colours and absolute magnitudes of star clusters of various ages and metallicities.  The disagreement between the model predictions and observations of cluster colours and magnitudes may indicate problems with or deficiencies in the modelling, and dioes not necessarily tell us that star clusters do not behave like SSPs. Matching the photometric properties of star clusters using SSP models is a necessary (but not sufficient) condition for clusters to be considered simple stellar populations.  Composite models, characterized by complex star-formation histories, also match the observed cluster colours. ",
        "introduction": "Simple stellar populations (SSPs), defined as groups of stars born at the same time, in the same volume of space, and from a gas cloud of homogeneous chemical composition certainly exist in Nature.  A priori, we cannot say that all stellar groups, associations or star clusters are SSPs, however. Most galaxies certainly are not.  Conceptually, SSPs are appealing because they are easy to model theoretically and their temporal evolution can be followed accurately. All stars of an SSP should have the same initial metal content.  At any given time, the stars composing an SSP describe an isochrone in the theoretical Hertzsprung--Russell (HR) diagram, which can easily be transformed to an observational colour--magnitude diagram (CMD). Detailed observations of this kind are available only for resolved stellar populations, either Galactic clusters or star clusters in nearby galaxies, or the stars making up the satellites of the Milky Way. Only in these few cases can we reliably establish the simplicity of a stellar population by inspection of its CMD.  Integrated properties of stellar populations, such as their colours or spectral-energy distributions, are subject to degeneracies (e.g., old, metal-poor populations resemble younger, metal-richer ones) and statistical (stochastic) uncertainties (due to the small number of stars present in low-mass systems) which, in most cases, prevent us from establishing with certainty if we are observing an SSP.  Of course, the question as to whether or not star clusters can be described by SSPs does not apply to those clusters which have been shown explicitly to host stellar populations of a composite nature, e.g., resolved clusters where a double or triple main sequence (MS) has been detected such as NGC 2808 (Piotto \\ea 2007; see also Kalirai \\& Richer 2010; van Loon 2010), or clusters showing evidence of prolonged star formation like $\\omega$ Centauri (Villanova \\ea 2007), nor to dwarf galaxies like Leo A (Cole \\ea 2007) or other galaxies in the Local Group with well-established multiple episodes of star formation (Gallart \\ea 2007). ",
        "conclusions": "The magnitudes and colours predicted by stellar population synthesis models do not readily match observations of unresolved star clusters, which are commonly expected to behave like ideal SSPs. This lack of agreement between the model predictions and observations may indicate problems or deficiencies in the modelling, and does not necessarily tell us that star clusters do not behave like SSPs.  In this review, I have briefly summarized the results of simple simulations which show how the range of colours observed in intermediate-age LMC star clusters can be understood on the basis of current stellar evolution theory, if we properly take into account the expected variation in the number of stars occupying sparsely populated evolutionary stages, due to stochastic fluctuations in the IMF. In this case, population synthesis models reproduce remarkably well the full ranges of observed integrated colours and absolute magnitudes of star clusters of various ages and metallicities. Some young clusters are described by supersolar-metallicity models, which may not be realistic for the LMC.  There is no need to introduce ad hoc assumptions into population synthesis models (Maraston 1998; Maraston \\ea 2001), representing a departure from our current understanding of stellar evolution theory, to explain the observed range of cluster colours and magnitudes. The predicted fluctuations in the integrated photometric properties of simulated clusters increase with decreasing cluster mass, as expected on the basis of the results of Cervi\\~no \\ea (2000, 2001, 2002).  It is worth pointing out that, because of the stochastic nature of the integrated-light properties of star clusters, single clusters may not be taken as reference standards of SSPs of a given age and metallicity. It should be emphasized that matching the photometric properties of star clusters using SSP models is a necessary condition for clusters to be considered simple stellar populations, but not sufficient.  So far, the single MS and unique MS turnoff required by SSPs can be established with certainty only for resolved stellar populations."
    },
    "0911/0911.5457_arXiv.txt": {
        "abstract": "We have analyzed nearly all images of the Taurus star-forming region at 3.6, 4.5, 5.8, 8.0, and 24~\\micron\\ that were obtained during the cryogenic mission of the {\\it Spitzer Space Telescope} (46~deg$^{2}$) and have measured photometry for all known members of the region that are within these data, corresponding to 348 sources, or 99\\% of the known stellar population. By combining these measurements with previous observations with the {\\it Spitzer} Infrared Spectrograph and other facilities, we have classified the members of Taurus according to whether they show evidence of circumstellar disks and envelopes (classes~I, II, and III). Through these classifications, we find that the disk fraction in Taurus, N(II)/N(II+III), is $\\sim75$\\% for solar-mass stars and declines to $\\sim45$\\% for low-mass stars and brown dwarfs (0.01--0.3~$M_\\odot$). This dependence on stellar mass is similar to that measured for Chamaeleon~I, although the disk fraction in Taurus is slightly higher overall, probably because of its younger age (1~Myr vs. 2--3~Myr). In comparison, the disk fraction for solar-mass stars is much lower ($\\sim20$\\%) in IC~348 and $\\sigma$~Ori, which are denser than Taurus and Chamaeleon~I and are roughly coeval with the latter. These data indicate that disk lifetimes for solar-mass stars are longer in star-forming regions that have lower stellar densities. Through an analysis of multiple epochs of {\\it Spitzer} photometry that are available for $\\sim200$ Taurus members, we find that stars with disks exhibit significantly greater mid-infrared variability than diskless stars, which agrees with the results of similar variability measurements for a smaller sample of stars in Chamaeleon~I. The variability fraction for stars with disks is higher in Taurus than in Chamaeleon~I, indicating that the IR variability of disks decreases with age. Finally, we have used our data in Taurus to refine the observational criteria for primordial, evolved, and transitional disks. The ratio of the number of evolved and transitional disks to the number of primordial disks in Taurus is 15/98 for spectral types of K5--M5, indicating a timescale of $0.15\\times\\tau_{\\rm primordial}\\sim0.45$~Myr for the clearing of the inner regions of optically thick disks. After applying the same criteria to older clusters and associations (2--10~Myr) that have been observed with {\\it Spitzer}, we find that the proportions of evolved and transitional disks in those populations are consistent with the measurements in Taurus when their star formation histories are properly taken into account.  {\\bf ERRATUM:} In Table~7, we inadvertently omitted the spectral type bins in which class~II sources were placed in Table~8 based on their bolometric luminosities (applies only to stars that lack spectroscopic classifications). The bins were K6--M3.5 for FT~Tau, DK~Tau~B, and IRAS 04370+2559, M3.5--M6 for IRAS 04200+2759, IT~Tau~B, and ITG~1, and M6--M8 for IRAS 04325+2402~C. In addition, the values of $K_s-[3.6]$ in Table~13 and Figure~26 for spectral types of M4--M9 are incorrect. We present corrected versions of Table~13 and Figure~26. ",
        "introduction": "\\label{sec:intro}      Much of our knowledge of the processes of star and planet formation has been derived from observations of circumstellar disks around newborn stars. The Taurus complex of dark clouds is one of the principle sites for studies of disks. Taurus is among the nearest star-forming regions to the Sun ($d=140$~pc) and contains more than 300 young stars and brown dwarfs \\citep{ken08}. Because of the low stellar density and the absence of photoionizing stars, most members of Taurus are not embedded within bright nebulae, and thus are not subject to high levels of background emission that would hinder infrared (IR) observations of disks. The low density of Taurus also permitted resolved photometry of individual stars with the low-resolution imaging that was offered by the first major far-IR telescope, the {\\it Infrared Astronomical Satellite} \\citep[IRAS,][]{bei86,ken90,ken94}. By combining 12--100~\\micron\\ photometry from IRAS with 1--10~\\micron\\ data from ground-based telescopes, \\citet{kh95} performed the most comprehensive census of disks in a star-forming region at that time. The disk-bearing stars identified in Taurus have served as targets for a multitude of detailed observations of circumstellar disks \\citep[][references therein]{ken08}.  The census of the stellar population of Taurus has expanded significantly since the disk study of \\citet{kh95}, particularly at masses below 0.5~$M_\\odot$. The sensitivities of IR telescopes also have improved dramatically, most notably with the deployment of the {\\it Spitzer Space Telescope} \\citep{wer04}, which is capable of detecting members of Taurus with masses below 0.01~$M_\\odot$. The Taurus stellar population has been imaged extensively by {\\it Spitzer} through pointed observations of individual stars as well as wide-field maps. These images have been used to measure photometry for known members of Taurus \\citep{har05,luh06tau2,gui07} and search for new disk-bearing members \\citep{luh06tau2,luh09fu,luh09tau1,reb10}.  With the recent completion of the cryogenic mission of {\\it Spitzer}, we wish to present a study of disks in Taurus that makes use of all {\\it Spitzer} images of the region, which cover a total area of 46~deg$^{2}$ and encompass 99\\% of the known stellar population. We begin by analyzing all {\\it Spitzer} images of Taurus at 3.6--24~\\micron\\ that are publicly available and measuring photometry for all known members that are detected in these data (\\S~\\ref{sec:images}). We use these measurements in conjunction with previous observations by the {\\it Spitzer} Infrared Spectrograph and other telescopes to classify the spectral energy distributions of all members of Taurus (\\S~\\ref{sec:sed}). Through these classifications, we investigate how the prevalence of disks depends on stellar mass (\\S~\\ref{sec:frac}) and location (\\S~\\ref{sec:spatial}). We then characterize the mid-IR variability of disks in Taurus with the multi-epoch photometry from {\\it Spitzer} (\\S~\\ref{sec:var}). Finally, we use our mid-IR data in Taurus to refine the observational criteria for the advanced stages of disk evolution and we apply these criteria to {\\it Spitzer} measurements in nearby clusters and associations with ages of 2--10~Myr (\\S~\\ref{sec:trans}). ",
        "conclusions": "We have performed a census of the circumstellar disk population of the Taurus star-forming region ($\\tau\\sim1$~Myr) using mid-IR images obtained with the {\\it Spitzer Space Telescope}. The results of this study are summarized as follows:  \\begin{enumerate}  \\item We have analyzed nearly all images of the Taurus cloud complex at 3.6, 4.5, 5.8, 8.0, and 24~\\micron\\ that were collected by {\\it Spitzer} during its cryogen mission. The IRAC and MIPS cameras on board {\\it Spitzer} each covered a total area 46~deg$^{2}$ and encompassed 346 and 299 members of Taurus, respectively, corresponding to 99\\% of the known stellar population. We have presented photometry for all members that were detected in these images.  \\item We have used our mid-IR photometry in conjunction with other available observations to determine the likely evolutionary stages of all members of Taurus (classes~0--III). Stars that exhibit evidence of circumstellar envelopes in previous measurements (e.g., IRS spectra) are assigned to classes~0 and I. We have classified the remaining members of Taurus using their mid-IR SEDs, or H$\\alpha$ emission for the few stars that lack {\\it Spitzer} data.  \\item Based on our classifications, the disk fraction in Taurus, N(II)/N(II+III), is $\\sim75$\\% for solar-mass stars and declines to $\\sim45$\\% for low-mass stars and brown dwarfs (0.01--0.3~$M_\\odot$). A similar dependence on stellar mass has been observed in Chamaeleon~I \\citep{luh08cha1}. In contrast to Taurus and Chamaeleon~I, IC~348 exhibits a disk fraction of only $\\sim20$\\% for solar-mass stars \\citep{lada06,mue07,luh05frac}. Given that IC~348 is roughly coeval with Chamaeleon~I ($\\tau\\sim2$--3~Myr), the comparison of these three regions suggests that disk lifetimes for solar-mass stars are longer in star-forming regions like Taurus and Chamaeleon~I that have lower stellar densities.  \\item As previously observed in Taurus \\citep{har02}, we find that the positions of the class~I and II members closely follow the distribution of dense gas while the class~III stars are more widely distributed. An analysis of the nearest neighbor distances also suggests that class~II sources may have a wider distribution than class~I sources, although this difference is only marginally significant. The median of the nearest neighbor distances for classes~I and II is 0.15~pc, which is twice the average value in nearby star-forming clusters \\citep{gut09}.  \\item Our study has produced multiple epochs of mid-IR photometry for $\\sim200$ members of Taurus. In these data, the mid-IR variability of class~I/II sources is much greater than that of class~III stars, which agrees with similar measurements with {\\it Spitzer} in Chamaeleon~I \\citep{luh08cha1}. The fraction of disk-bearing stars that are variable is higher in Taurus than in Chamaeleon~I, indicating that the variability of disks decreases with age.  \\item We have used our {\\it Spitzer} photometry for the disk population in Taurus to refine the observational criteria for the evolutionary phases of disks. When plotted in terms of $K_s-[5.8]$, $K_s-[8.0]$, and $K_s-[24]$, the members of Taurus appear predominantly within two distinct groups that are well-separated from each other, as found in early mid-IR studies of Taurus \\citep{skr90}. The colors of the stars in the bluer group are consistent with stellar photospheres. We have defined the large, continuous population of much redder sources as primordial disks. Using our models of accretion disks, we have demonstrated that the colors of these primordial disks can be explained in terms of optically thick disks. The sources that fall within the gap between primordial disks and stellar photospheres are defined as evolved disks (weak excess in all bands), transitional disks (weak or no excess at $\\lambda<10$~\\micron, large excess at longer $\\lambda$), and evolved transitional disks or debris disks (only weak 24~\\micron\\ excess). We have identified 19 members of Taurus that are candidates for disks in these later stages of disk evolution, 11 of which have not been previously recognized as such.  \\item We have applied our classification criteria for disks to nearby clusters and associations with ages of 2--10~Myr that have been observed with {\\it Spitzer}. We find that the number of evolved and transitional disks in those regions is consistent with the paucity of such disks in Taurus. Some of the sources that have been classified as evolved and transitional disks in previous {\\it Spitzer} studies have colors that are similar to those of the bluer primordial disks in Taurus (i.e., flatter optically thick disks). The number ratio of evolved and transitional disks to primordial disks in Taurus is 15/98 for spectral types of K5--M5, indicating a timescale of inner disk clearing that is $\\sim15$\\% of the lifetime of primordial disks ($\\sim3$~Myr). The data in Taurus and the older populations that we have examined are inconsistent with notion that the lifetime of the evolved and transitional phases (i.e., disk clearing timescale) is comparable to that of primordial disks.  \\end{enumerate}"
    },
    "0911/0911.0044_arXiv.txt": {
        "abstract": "We examine the colour-magnitude relation of $\\sim 5000$ Brightest Cluster Galaxies (BCGs) in the SDSS. The colour-magnitude and colour-velocity dispersion relations of the BCGs are flatter in slope, by factors of two or more, than those of non-BCG early-type galaxies of similar luminosity ($M_r<-22.5$), while their $g-r$ colours at the half-light radius are $\\sim 0.01$ magnitude redder.  We investigate and compare radial colour gradients (which are usually negative) in these BCGs and a sample of 37000 early-type galaxies, using a gradient estimator based on the ratio of de Vaucouleurs effective radii in $g$ and $r$ passbands, ${{{\\rm r}_{eff}(g)}\\over{{\\rm r}_{eff}(r)}}-1$. The mean colour gradient of BCGs is flatter (by $23\\%$) than for other E/S0 galaxies of similar luminosity.  The colour gradients in early-type galaxies are stronger at intermediate luminosity ($M_r\\simeq -22$) than at low or the highest luminosities, and tend to decrease with increasing velocity dispersion ($\\sigma$). In non-BCG E/S0s, colour gradients increase for larger effective radii (up to 10--12 kpc), and are negatively correlated with $10~{\\rm log}~\\sigma+M_r$, the mass density, and stellar age. However, gradients can be reduced or inverted (positive) for post-starburst galaxies at the youngest ages. In BCGs, these trends are absent and the mean colour gradient remains at a relatively low level ($\\sim 0.08$ by our measure) whatever the other properties of the galaxies. The redder half-light radius colours of the BCGs can be explained by their slightly greater ages combined with flatter radial colour gradients.  We discuss possible explanations in terms of spheroidal galaxies initially having a colour gradient positively correlated with luminosity and positively correlated with large radius and/or low density. Subsequently, elliptical-elliptical `dry' mergers progressively reduce colour gradients, towards a low but non-zero value. This has occurred a greater number of times during the formation histories of the most massive E/S0s, and to by far the greatest degree in the BCGs. ",
        "introduction": "In this paper we investigate colour-magnitude relations, and, to greater depth, the properties of radial colour gradients, for early-type galaxies and in particular Brightest Cluster Galaxies (BCGs). Both can provide information on these objects' formation.  Early-type galaxies (elliptical and S0) with higher luminosities are redder in colour than those of lower luminosity, following a relatively tight and linear  colour-magnitude relation (CMR), or `red sequence'. Gallazzi et al. (2005) performed a detailed analysis of $2\\times 10^5$ galaxy spectra from the Sloan Digital Sky Survey (SDSS), using a set of 5 spectral indices (D4000, $\\rm H\\beta$, $\\rm H_{\\delta A}+H_{\\gamma A}$, $[\\rm Mg_2 Fe]$ and $\\rm [MgFe]^{\\prime}$) to estimate mean stellar ages and metallicities. Gallazzi et al. (2006) and Jimenez et al. (2007), from spectral analysis of the early-type galaxies, conclude that the slope of the red-sequence CMR is primarily caused by an increase in mean stellar metallicity with galaxy mass or velocity dispersion. Mergers may then affect the CMR by building  more luminous galaxies with the same composition as their less massive progenitors, which would tend to flatten the CMR slope at $L>L^*$ luminosities (eg. Skelton, Bell and Somerville 2009).  Age also has a (secondary) role in explaining the colour trend - the most massive E/S0s formed early and rapidly (in $<1$ Gyr), while those of lower luminosity formed stars over longer timescales, giving slightly younger flux-weighted ages. There is also evidence that E/S0 galaxy formation occurred slightly earlier in dense environments (eg. Sheth et al. 2006, Bernardi 2009), while galaxies in cluster centres may have had an unusual and extreme merger history (e.g. Boylan-Kolchin, Ma and Quataert 2006).  Thus merger history and cluster environment may both influence the CMR.  Our studies of the CMR in Roche, Bernardi and Hyde 2009 (hereafter Paper I) suggested (see Section 3.1) that the radial colour gradients in E/S0 galaxies might be correlated with their other properties and it would be fruitful to investigate this further. These colour gradients are of interest in that they are sensitive tracers of evolution processes and vary strongly between the individual E/S0s. Most E/S0 galaxies have negative colour gradients, meaning that the centres are reddest and  colours become bluer outwards. This could be the result of a negative gradient in age, metallicity or dust, or some combination of these.  Tamura et al. (2000) and Tamura and Ohta (2003) estimate the mean colour gradient in E/S0s as ${{\\Delta(B-R)}\\over{\\Delta(\\rm log~r)}}=-0.09$ mag $\\rm dex^{-1}$ (with a scatter of 0.04), and attributed this to a radial gradient of metallicity (rather than age) ${{d(\\rm log~Z)}\\over{d(\\rm log~r)}}= -0.3\\pm 0.1$, approximately constant to $z\\simeq 1$. Wu et al. (2005) similarly estimated ${{d({\\rm log}~Z)}\\over{d(\\rm log~r)}}= -0.25$ with a large scatter $\\sigma=0.19$. Mehlert et al. (2003) found negative metallicity gradients in  Coma-cluster ellipticals, but negligible gradients in age or $\\alpha$/Fe. Michard (2005) concluded colour gradients in 50 nearby ellipticals were primarily due to metallicity gradients and the effects of dust were usually small. La Barbera and Carvalho (2009) compare optical and near-IR colour gradients and again conclude they are produced by negative gradients in metallicity (they find the radial age gradients in E/S0s may be even be positive).  Kobayashi (2004) performed detailed chemodynamical simulations of a set of over 100 model elliptical galaxies, with star-formation and merging, and predicted that spheroidals which formed monolithically, or at least with only minor mergers, would have steeper metallicity gradients (${{d({\\rm log}~Z)}\\over{d(\\rm log~r)}}\\simeq -0.3$ to -0.5) than those formed by major mergers (${{d({\\rm log}~Z)}\\over{d(\\rm log~r)}}\\simeq -0.2)$. Within the simulated galaxy set, merger history was the primary determinant (more so than mass or luminosity) of present-day metallicity/colour gradient. This might account for the observed wide distribution in colour gradients if similar numbers of spheroidals had monolithic, minor-merger or major-merger histories, and might also lead to an environmental dependence.  In this paper we focus especially on Brightest Cluster Galaxies (BCGs), of which several thousand have been identified in the SDSS. Previously, the BCGs in the SDSS have been found to follow a steeper relation of radius to luminosity (${\\rm r}_{eff}\\propto L$) than the other E/S0s (${\\rm r}_{eff}\\propto L ^{0.6}$); (Bernardi et al. 2007; Bernardi 2009). In this paper, we look for systematic differences in their CMR and radial colour gradients with respect to other E/S0 galaxies, especially those in the same luminosity range ($M_r<-22.5$ to $M_r\\simeq -25$). Some models, for example,  predict a much flatter CMR for the BCGs, as a result of formation from a relatively large number of mergers of (already) old and red galaxies (de Lucia and Blaizot 2007). Ko and Im (2005) found indications that the colour gradients of ellipticals in dense cluster environments tended to be less strong than in field environments, which again could be due to (more) mergers.  In Section 2 of this paper we describe the data used and our selection of early-type galaxies and BCGs. In Sections 3 and 4 we investigate and compare their CMR and colour-$\\sigma$ relation. In Section 5 we use an estimator of colour gradient based on the ratio of $g$ and $r$-band effective radii to compare the colour gradients of BCGs and other E/S0s and examine the dependences on luminosity, radius, and density. In Section 6 we look at the influence of stellar age, as estimated from the spectra. We conclude in Section 7 with summary and further discussion.  SDSS magnitudes are given in the AB system where $m_{AB}=-48.60-2.5$ log $F_{\\nu}$ (in ergs $\\rm cm^{-2}s^{-1}Hz^{-1}$); equivalently, $m_{AB}=0$ is 3631 Jy. We assume throughout a spatially flat cosmology with $H_0=70$ km $\\rm s^{-1}Mpc^{-1}$, $\\Omega_{M}=0.27$ and $\\Omega_{\\Lambda}=0.73$, giving the age of the Universe as 13.88 Gyr. ",
        "conclusions": "(i) The colour-magnitude relation (CMR) and colour-velocity dispersion relation (C$\\sigma$R) of BCGs are significantly flatter in slope(${{d(g-r)}\\over{d M_r}}$ and ${{d(g-r)}\\over{d{\\rm log}~\\sigma}}$) than the respective relations for non-BCG E/S0 galaxies of comparable (high) luminosity. The difference in slope is about a factor of two, whether we define the $g-r$ colour using k-corrected model magnitudes or rest-frame magnitudes derived from integrating the SDSS spectra. BCGs are also, on average, $\\sim 0.01$ mag redder in model magnitudes (k-corrected to rest-frame) than E/S0s of the same $M_r$ or $\\sigma$, but with less or no offset in spectra-derived colours.  The hierarchical merging model of Lucia and Blaizot (2007) predicts a flat CMR for BCGs as a result of late assembly from a large number of red progenitors which formed their stars very much earlier ($z\\sim 5$). This model included a very early truncation of star-formation. Skelton, Bell and Somerville (2009) present a simplified model in which dry mergers of already merged E/S0s on a `creation red sequence'  mildly flatten the CMR slope at higher luminosities. There may indeed be some flattening in the bright end of the CMR for all E/S0s (see Paper I), but for BCGs we find much more, although the BCG CMR is not entirely flat. This suggests the BCGs or their progenitors evolve by a less extreme scenario than the first of these models but perhaps with a higher dry merger rate compared to the E/S0s in the second model. In addition to being flattened, the BCG CMR is offset redwards. \\medskip   (ii) As a simple quantifier of radial colour gradient we use  ${{{\\rm r}_{eff}(g)}\\over{{\\rm r}_{eff}(r)}}-1$, with the effective radii ${\\rm r}_{eff}(g)$ and ${\\rm r}_{eff}(r)$, corrected to the rest-frame by linearly interpolating between the radii fitted in the bands $g$, $r$, $i$ and $z$. We show that at least for galaxies with angular $r_{eff}$ radii above 1.8 arcsec, this ratio is well and linearly correlated to ${{d(g-r)}\\over {d(\\rm log~r)}}$ across the observed range of colour gradients and galaxy luminosities, and estimate ${{d(g-r)}\\over {d(\\rm log~r)}}\\simeq -0.87 ({{{\\rm r}_{eff}(g)}\\over{{\\rm r}_{eff}(r)}}-1)$. The large scatter in this correlation, which we initially estimated as 0.122, is reduced to a much more acceptable 0.056 if only galaxies with $r_{eff}$ larger than 1.8 arcsec are included.  After excluding the $r_{eff}<1.8$ arcsec galaxies and a few others which we believed to be spiral types, we found for the E/S0s a mean colour gradient, using our measure, of $0.11115\\pm 0.00061$ with a comparable galaxy-to galaxy scatter 0.1172. For the $M_r<-22.5$ non-BCG the mean is similar at $0.10550\\pm 0.00098$ (scatter 0.1141) and for max-BCGs the mean  is $0.08142\\pm 0.00114$ (scatter 0.0764). Therefore we find at high significance that BCGs have a flatter (by an average of $23\\%$) radial colour gradient than other high luminosity spheroidals.  This again could be due to intensified merging in the cluster-core environment. La Barbera et al. (2005), and similarly Ko and Im (2005), reported mean colour gradients in E/S0s of all luminosities were weaker by almost a factor two in the richest clusters ($N_{gal}>50$) compared to the field, and suggested this was because formation histories in denser cluster environments included more elliptical-elliptical mergers (e.g. Tran et al. 2005), which reduce colour/metallicity gradients (Kobayashi 2004).  The simulations of Boylan-Kolchin, Ma and Quataert (2006) account for the steeper $R_e$--$L$ of BCGs, i.e. their `abnormally' large sizes (e.g. Bernardi 2009), as the result of repeated `dry' (dissipationless, non-star-forming) mergers, and especially of the anisotropic accretion of smaller spheroidals in near-radial, low angular momentum  orbits, which might occur preferentially in a cluster core. A few BCGs have been observed undergoing such mergers (Liu et al. 2008; Tran et al. 2008). Di Matteo et al. (2009) predict from simulations that the metallicity gradient in a  dissipationless (`dry') merger remnant will consistently be $\\simeq 0.6$ the mean of the progenitor galaxies' gradients (but not reduced to zero). \\medskip  (iii) In non-BCG E/S0s, we examine the trends in mean colour gradient as a function of other galaxy properties. We find a dependence of colour gradient on absolute magnitude $M_r$, with a broad peak in colour gradient at $M_r\\simeq - 22$ and a decrease to lower and higher luminosities. The luminosity dependence is relatively mild, less than a factor 1.5, which may explain why it was not seen by La Barbera et al. (2005). However, Spolaor et al. (2009) do find a luminosity dependence similar to ours. Colour gradients tend to decrease with increasing velocity dispersion $\\sigma$, by almost 1/2  between 150 and 300 km $\\rm s^{-1}$, confirming the suggestion in our Paper I that the trends with $M_r$ and $\\sigma$ are different. \\medskip  (iv) Colour gradients in E/S0s show strong correlations with some other galaxy properties. Colour gradients increase, by about a factor of two, with from small to larger effective radius, up to a maximum at 8-12 kpc (depending on luminosity). A positive correlation with radius was prevously reported by Tamura and Ohta (2003) for a small sample of E/S0s in Abell 2199. At even larger radii we find the mean colour gradients decrease, but this is actually because of a reduced gradient in the most luminous ($M_r<-23$) galaxies. If E/S0s are divided by luminosity, within  low/moderate luminosity intervals there is simply a steep increase in colour gradient with $\\rm r_{eff}$.  We find colour gradients to be negatively correlated with $\\rm 10~log~\\sigma+M_r$ (the $\\sigma$ residual relative to $\\langle \\sigma|M_r\\rangle$)  and mass density ($\\sigma^2/\\rm r_{eff}^2$).  Of course, these quantities are related -- a high density implies a high $\\sigma$ at a given dynamic mass (rather than luminosity). These negative correlations do not seem to be caused by a selection effect (from the flux limit of the sample), as they are seen within a wide range of narrow luminosity intervals, being strongest for E/S0s of moderate/high luminosity $-21<M_r<-24$. \\medskip  (v) We also examine the trend in colour gradient with the `mean stellar ages' estimated (Gallazzi et al. 2005, 2006) for most of these galaxies using 5 indices from the SDSS spectra. In non-BCG E/S0s the colour gradients are strongest for relatively young ages of 4--5 Gyr (which may signify the presence of younger stellar populations or more extended star-formation, rather than the galaxies being younger) and decrease by up to a factor of two to the oldest ages of 11 Gyr. This may be primarily because older stellar populations have experienced more mergers and interactions, especially at high redshifts. Early-type galaxies which ceased forming stars and joined the red sequence more recently (e.g. post-mergers of spirals at $z<1$) would tend to have stronger colour gradients, if it is dry mergers (occurring after the end of star-formation) that flatten the gradients.  However, there is a drop in mean colour gradient for the galaxies with mean stellar ages less than 3--4 Gyr (even in BCGs). This is likely due to the inclusion of post-starburst galaxies, which can have blue cores giving inverted (positive) colour gradients (Friaca and Terlevich 1999; Yamauchi and Goto 2005) together with young age indicators like strong Balmer absorption lines. To confirm this we examined the spectra of 40 galaxies with age $<3$ Gyr and inverted colour gradients ($<-0.1$) and found many with enhanced Balmer absorption including 7 with equivalent widths $>5\\rm \\AA$  in both $\\rm H\\beta$ and $\\rm H\\delta$. Suh et al. (2010) found that blue-cored early-types in general had at least mildly enhanced $\\rm H\\beta$ absorption, and showed that the Balmer enhancement and the positive colour gradients could both be accounted for by centrally concentrated, very young stellar populations (age $\\leq \\rm 0.5 Gyr$), comprising only $\\sim0.5$--$2\\%$  of the stellar mass. \\medskip  (vi) Colour gradients in BCGs are consistently lower than in non-BCG spheroidals of the same luminosity, velocity dispersion or radius. Some classes of non-BCG  ellipticals, those with (i) the highest luminosities and largest radii, (ii) the highest $\\sigma$ relative to $M_r$, or (iii) the highest densities, also have mean colour gradients of ${{{\\rm r}_{eff}(g)}\\over{{\\rm r}_{eff}(r)}}\\simeq 0.08$. But within the BCG class mean ${{{\\rm r}_{eff}(g)}\\over {{\\rm r}_{eff}(r)}}-1$ remains at $\\sim 0.08$, almost regardless of the other galaxy properties. The mean colour gradient in BCGs decreases a little for the largest radii and with increasing dynamic mass, but is uncorrelated with mass density, the $10~{\\rm log}~\\sigma+M_r$ residual, or stellar age (except for the drop at $<3$ Gyr).  The process of colour gradient flattening through dry mergers appears to have operated most strongly on (all) BCGs, reducing the mean colour gradients to $\\sim 0.08$, but never much lower than this, whatever the initial value. In the Di Matteo et al. (2009) model of dry mergers, a galaxy with a strong colour gradient will probably merge with one of lower gradient and experience a significant flattening, but a galaxy with already a very flat colour gradient would probably merge with one with higher gradient and experience no further decrease or even a slight increase. With repeated mergers this would cause an exponential decrease in the mean gradient and also a narrowed distribution, which is what we seem to find for the merged classes of galaxies, with mean gradients of $\\simeq 0.08$ consistently. We found the BCGs had ${1\\over 3}$ less galaxy-to-galaxy variation in colour gradient than the E/S0s (this difference might be even greater after taking measurement errors into account).   In addition, the radial mergers thought to occur in BCGs would strongly  increase their radii (and reduce density) while simultanously flattening colour gradients, thus erasing the gradient-radius correlation (and gradient-density anticorrelation) seen in other spheroidals. \\medskip  (vii) We examine colour as a function of stellar age, comparing the BCGs and other E/S0s. The BCGs are on average 0.5 Gyr older at observation than the $M_r<-22.5$ non-BCGs, and their mean lookback time to formation is estimated as 0.68 Gyr greater. This could simply reflect the correlation of earlier formation (by as much as $\\sim 1$ Gyr) with high-density cluster environment (e.g. Sheth et al. 2006; Rogers et al. 2010). The BCGs also have a flatter age-luminosity relation, as might be expected if BCGs (of all masses) form from dry major mergers of red early-type galaxies (with older ages associated with the cluster-core environment), which would increase luminosity without changing the mean stellar age. In the same way, their CMR is flattened. If, on the other hand, the BCGs had formed `monolithically', without major merging, they might show stronger trends in their stellar populations (age, composition) with luminosity/mass, extrapolating the trends seen in less massive E/S0s.  In rest-frame model-magnitude $g-r$, BCGs are slightly redder than non-BCG E/S0s of the same age, whereas in spectra-derived $g-r$, they lie on the same (approximately linear) relation of colour to age. We conclude that the ($\\simeq 0.01$ mag) redder model-magnitude $g-r$ colours of the BCGs in the CMR and C$\\sigma$R can be explained simply by their greater mean ages combined with their flatter colour gradients, which would give them a slightly redder model-magnitude colour for a given central-aperture (e.g. spectra-derived) colour. \\medskip  (viii) We re-examine the correlations of colour gradient with galaxy properties, dividing by age (star-formation redshift). The positive correlation of colour gradient with radius, and the anticorrelation with density and $10~{\\rm log}~\\sigma +M_r$ are strongest in younger galaxies but are seen within all age intervals. With increasing age, these correlations flatten somewhat while the mean colour gradients decrease, as might be expected for sequential dry mergers in the Di Matteo et al. (2009) model.  The decrease in colour gradient with the $\\sigma$ residual $10~{\\rm log}~\\sigma+M_r$ is not solely due to the known correlation of this quantity with stellar age (Forbes and Ponman 1999; Bernardi et al. 2005; Gallazzi et al. 2006), as it is seen within each age interval. Rather, it seems the colour gradient must have independent anticorrelations with the $\\sigma$ to luminosity  ratio and with stellar age.  It seems the process of spheroid formation sets up a  metallicity/colour gradient, with a steepness positively correlated with the mass/luminosity of stars and also with a large radius and low mass density for the galaxy. Subsequent mergers flatten the stronger gradients while adding low-gradient galaxies at the high-luminosity end of the E/S0 luminosity function. As dry mergers tend to increase size this also weakens the gradient-density relation (Figure 23). A comparison with colour gradients in much higher redshift E/S0s would help to confirm this. The mean colour gradient in E/S0s decreases somewhat at $M_{dyn}>10^{11.5}M_{\\odot}$ because most of these higher-mass galaxies are the products of elliptical-elliptical (dry) mergers (Tran et al. 2005; Naab, Khochfar and Burkert 2006). Yet some E/S0s remain today with colour gradients more than twice the observed average, like those in the `monolithic' models of Kobayashi et al. (2004).  The position of BCGs on Figures 21 to 24 is that of galaxies which have experienced more age and/or merger-related colour gradient suppression than even the $z_{form}>2$ E/S0s. On Fig 18 the BCG colour gradient matches that of non-BCGs of mean stellar age $>10$ Gyr, although the BCG mean stellar age is only 7.5 Gyr. The colour gradients in BCGs could be effectively `super-aged' by several Gyr by their environment and their position in the path of infalling cluster spheroidals."
    },
    "0911/0911.5727_arXiv.txt": {
        "abstract": "\\begin{center}{\\bf Abstract}\\end{center} A new accelerating cosmology  driven only by baryons plus cold dark matter (CDM) is proposed in the framework of general relativity. In this model the present accelerating stage of the Universe is powered by the negative pressure describing the gravitationally-induced particle production of cold dark matter particles.  This kind of scenario has only one free parameter and the differential equation governing the evolution of the scale factor is exactly the same of the $\\Lambda$CDM model. For a  spatially flat Universe, as predicted by inflation ($\\Omega_{dm}+\\Omega_{baryon}=1$), it is found that the effectively observed  matter density parameter is  $\\Omega_{meff} = 1- \\alpha$,   where $\\alpha$ is the constant parameter specifying the CDM particle creation rate. The supernovae test based on the Union data (2008) requires $\\alpha\\sim 0.71$ so that $\\Omega_{meff} \\sim 0.29$  as independently derived  from weak gravitational lensing, the large scale structure and other complementary  observations. ",
        "introduction": "It is well known that observations  from Supernovae Type  Ia (SNeIa) provide strong evidence for an expanding accelerating Universe \\cite{Riess07,Union08}. In relativistic cosmology, such a phenomenon is usually explained by the existence of a new dark component (in addition to cold dark matter), an exotic fluid endowed with negative pressure \\cite{review}.  Many candidates for dark energy have been proposed in the literature \\cite{decaying,XM,SF,CGas}, among them: (i) A cosmological constant ($\\Lambda$), (ii) a decaying vacuum energy density or $\\Lambda(t)$-term,  (iii) a relic scalar field slowly rolling down its potential, (iv) the ``X-matter'',  an extra component characterized by equation of state $p_x=\\omega \\rho_x$, where $\\omega$ may be constant or a redshift dependent function, (iv) a Chaplygin-type gas whose equation of state is $p=-A/\\rho^{\\gamma}$, where A and $\\gamma$ are positive parameters.  All these models explain the accelerating stage, and, as such, the space parameter of the basic observational quantities is rather degenerate. Nowadays, the most economical explanation is provided by the flat $\\Lambda$CDM model which has only dynamic free parameter, namely, the vacuum energy density. It seems to be consistent with all the available observations provided that the vacuum energy density is fine tuned to fit the data ($\\Omega_{\\Lambda} \\sim 0.7$).  However, even considering that the addition of these fields explain the late time accelerating stage  and other complementary observations \\cite{CMB,Clusters}, the need of (yet to be observed) dark energy  component with  unusual properties is certainly a severe hindrance.  In general relativistic cosmology, the presence of a negative pressure is the key ingredient to accelerate the expansion. In particular, this means  that cosmological models dominated by pressureless fluid like a CDM component expands in a decelerating way. However, as first discussed by Prigogine and coworkers \\cite{Prigogine} and somewhat clarified by Calv\\~{a}o and collaborators \\cite{LCW}  through a manifestly covariant formulation, the  matter creation process at the expense of the gravitational field is also macroscopically described by a negative pressure. Later on, it was also demonstrated that the matter creation is an irreversible process completely  different from the bulk viscosity description \\cite{LG92} originally proposed by Zeldovich \\cite{Zeld70} to avoid the singularity, as well as to describe phenomenologically the emergence  of particles in the begin of the Universe evolution early (see also \\cite{SLC02} for a more complete discussion comparing particle creation and bulk viscosity).  Microscopically, the  gravitationally induced particle creation mechanism has also been discussed by many authors \\cite{Parker,BirrellD}. A non-stationary gravitational background influences quantum fields in such a way that the frequency becomes time-dependent. In the case of a flat Friedmann-Robertson-Walker (FRW) spacetime described in conformal time coordinates, the key result is that the scalar field obeys the same equation of motion as a massive scalar field in Minkowski spacetime, except that the effective mass becomes time dependent (the dispersion relation in the FRW metric involves the scale factor and its second derivative). When the field is quantized, this leads to particle creation, with the energy for newly created particles being supplied by the classical, time-varying gravitational background.  In this context, we are proposing here a new flat cosmological scenario where the cosmic acceleration is powered uniquely by the  creation of cold dark matter particles.  It will be assumed that the CDM particles are described by a real scalar field  so that only particle creation takes place because in this case it is its own antiparticle. The model can be seen as a workable alternative to the cosmic concordance cosmology because it has only one free parameter and the equation of motion is exactly the same of the $\\Lambda$CDM model.  As we shall see, in the case of a  spatially flat Universe ($\\Omega_{dm}+\\Omega_{bar}=1$), the effectively observed  matter density parameter is  $\\Omega_{eff} = 1- \\alpha$,   where $\\alpha$ is the constant parameter defining the creation rate. The supernovae test requires the central value $\\alpha\\sim 0.71$ so that $\\Omega_{eff} \\sim 0.29$  in accordance with the large scale structure and other complementary  observations. ",
        "conclusions": "A new creation cold dark matter (CCDM) cosmology has been proposed. In this late time CDM dominated model, the  vacuum energy density  parameter is  $\\Omega_{\\Lambda} = 0$, and, therefore, the so-called cosmological constant problem \\cite{weinb, decaying} is absent.  The late time acceleration is  powered here by an irreversible creation of CDM particles and the value of  $H_0$ does not need to be small in order to solve the age problem.  \\begin{table}[ht] \\caption{$\\Lambda$CDM vs. CCDM} \\centering \\begin{tabular}{c c} \\hline\\hline $\\Lambda$CDM & CCDM \\\\ [0.5ex] \\hline $\\Omega_\\Lambda$ & $\\alpha$ \\\\ $\\Omega_m$ & $\\Omega_{meff}\\equiv \\Omega_m - \\alpha$ \\\\ Vacuum DE & Creation of CDM \\\\ Acceleration ($z_t \\approx 0.71$, $k=0$)\\, & \\,acceleration ($z_t \\approx 0.71$, $k=0$) \\\\ [1ex] \\hline \\end{tabular} \\label{tab1} \\end{table}  It is worth noticing the existence of a dynamic equivalence between CCDM and $\\Lambda$CDM cosmologies at the level of the background equations.  Actually, the CCDM scenario can formally be interpreted as a two component fluid mixture: a pressureless matter with density parameter, $\\Omega_{eff} = \\Omega_m -\\alpha$, plus a ``vacuum fluid\"  with $\\rho_v=-p_v = \\alpha \\rho_{co}$, where $\\rho_{co}$ is the critical density parameter.  A simple qualitative  comparison between both approaches is summarized on Table 1. Note that for nonflat CCDM, there are two dynamic free parameters, namely, $\\alpha$ and  $\\Omega_m$ (or $\\Omega_{meff}$), similarly to nonflat $\\Lambda$CDM,  whose dynamic free parameters are  $\\Omega_{\\Lambda}$ and $\\Omega_m$. For the flat CCDM case,  there is just one dynamic free parameter (like in  flat $\\Lambda$CDM model), say, $\\alpha$.  This formal equivalence explains why CCDM scenarios provide  an excellent fit to the observed dimming of distant type Ia supernovae (see  text and Figs. 1a and 1b). As it appears, one may say that $\\Lambda$CDM cosmology is one of the possible effective descriptions of cold dark matter creation scenarios.  On the other hand, since the creation mechanism adopted here is  classically described as an irreversible process \\cite{Prigogine,LCW}, the basic  problem with this new cosmology is related  to  the absence of a consistent approach based in quantum field theory in curved spacetimes.  However,  the  last 30 years thinking about the cosmological constant problem are suggesting that the possible difficulties in searching for a more rigorous quantum approach for matter creation in  the expanding Universe are much smaller than in the so-called $\\Lambda$-problem. Indeed, the basic tools have already been discussed long ago \\cite{Parker,BirrellD}, and, as such, the problem now is reduced  to take into account properly the entropy production rate present in the creation mechanism and the associated creation pressure.  Finally, for those believing that  $\\Lambda$CDM contains all the physics that will be needed to confront  the next generation of cosmological tests we call attention for the model proposed here. It is simple like $\\Lambda$CDM, has the same dynamics, and, more important, it is based only in cold dark matter whose status is relatively higher than any kind of dark energy. Naturally, new constraints on the relevant parameters ($\\alpha$ and $\\Omega_m$) from complementary observations need to be investigated in order to see whether the CCDM model proposed here provides a realistic description of the observed Universe.  In principle, additional tests  measuring  the matter power spectrum, the weak gravitational lensing distortion by foreground galaxies and the cluster mass function may decide between $\\Lambda$CDM and CCDM cosmologies. New bounds on the CCDM parameters coming from background and perturbed cosmological equations will be discussed in a forthcoming communication."
    },
    "0911/0911.0872.txt": {
        "abstract": "An interesting new high-energy pulsar sub-population is emerging following early discoveries of gamma-ray millisecond pulsars (MSPs) by the \\textit{Fermi} Large Area Telescope (LAT). We present results from 3D emission modeling, including the Special Relativistic effects of aberration and time-of-flight delays and also rotational sweepback of B-field lines, in the geometric context of polar cap (PC), outer gap (OG), and two-pole caustic (TPC) pulsar models. In contrast to the general belief that these very old, rapidly-rotating neutron stars (NSs) should have largely pair-starved magnetospheres due to the absence of significant pair production, we find that most of the light curves are best fit by TPC and OG models, which indicates the presence of narrow accelerating gaps limited by robust pair production -- even in these pulsars with very low spin-down luminosities. The gamma-ray pulse shapes and relative phase lags with respect to the radio pulses point to high-altitude emission being dominant for all geometries. We also find exclusive differentiation of the current gamma-ray MSP population into two MSP sub-classes: light curve shapes and lags across wavebands impose either pair-starved PC (PSPC) or TPC / OG-type geometries. In the first case, the radio pulse has a small lag with respect to the single gamma-ray pulse, while the (first) gamma-ray peak usually trails the radio by a large phase offset in the latter case. Finally, we find that the flux correction factor as a function of magnetic inclination and observer angles is typically of order unity for all models. Our calculation of light curves and flux correction factor for the case of MSPs is therefore complementary to the ``ATLAS paper'' of Watters et al.\\ for younger pulsars. ",
        "introduction": "\\label{sec:intro} The field of gamma-ray pulsars has already benefited profoundly from discoveries made during the first year of operation of the \\textit{Fermi} / Large Area Telescope (LAT). These include detections of the radio-quiet gamma-ray pulsar inside the supernova remnant CTA~1 \\citep{Abdo08_CTA1}, the second gamma-ray millisecond pulsar (MSP) \\citep{Abdo09_J0030} following the \\textit{EGRET} $4.9\\sigma$-detection of PSR~J0218+4232 \\citep{Kuiper04}, the~6 high-confidence \\textit{EGRET} pulsars \\citep{Thompson99,Thompson04}, and discovery of~16 radio-quiet pulsars using blind searches \\citep{Abdo09_BS}. In addition, 8~MSPs have now been unveiled \\citep[][see Table~\\ref{tab1}]{Abdo09_MSP}, confirming expectations prior to \\textit{Fermi's} launch in June 2008 \\citep{HUM05,Venter05_Cherenkov}. A \\textit{Fermi} six-month pulsar catalog is expected to be released shortly \\citep{Abdo09_Cat}. \\textit{AGILE} has also reported the discovery of~4 new gamma-ray pulsars, and marginal detection of~4 more \\citep{Halpern08,Pellizzoni09b}, in addition to the detection of~4 of the \\textit{EGRET} pulsars \\citep{Pellizzoni09a}. Except for the detection of the Crab at energies above 25~GeV \\citep{Aliu08}, no other pulsed emission from pulsars has as yet been detected by ground-based Cherenkov telescopes \\citep{Schmidt05,Albert07,Aharonian07,Fuessling08,Kildea08,Konopelko08,Albert08,Celik08,DelosReyes09}.  MSPs are characterized by relatively short periods $P\\lesssim30$~ms and low surface magnetic fields $B_0\\sim10^8-10^9$~G, and appear in the lower left corner of the $P\\dot{P}$-diagram (with $\\dot{P}$ the time-derivative of $P$; see Figure~\\ref{fig:PPdot}, where the newly-discovered \\textit{Fermi} MSPs are indicated by squares). MSPs are thought to have been spun-up to millisecond periods by transfer of mass and angular momentum from a binary companion during an accretion phase \\citep{Alpar82}. This follows an evolutionary phase of cessation of radio emission from their mature pulsar progenitors, after these have spun down to long periods and crossed the ``death line'' for radio emission. These ``radio-silent'' progenitors \\citep{Glendenning00} are thought to reside in the ``death valley'' of the $P\\dot{P}$-diagram, which lies below the inverse Compton scattering (ICS) pair death line \\citep{HM02}.  The standard ``recycling scenario'' \\citep{Bhattacharya91} hypothesizing that MSP birth is connected to low-mass X-ray binaries (LMXRBs) might have been confirmed recently by the detection of radio pulsations from a nearby MSP in an LMXRB system, with an optical companion star \\citep{Archibald09}. Optical observations indicate the presence of an accretion disk within the past decade, but none today, raising the possibility that the radio MSP has ``turned on'' after termination of recent accretion activity, thus providing a link between LMXRBs and the birth of radio MSPs.  High-energy (HE) radiation from pulsars has mainly been explained as originating from two emission regions. Polar cap (PC) models \\citep{Harding78,Daugherty82,Sturner95,DH96} assume extraction of primaries from the stellar surface and magnetic pair production of ensuing HE curvature radiation (CR) or ICS gamma rays, leading to low-altitude pair formation fronts (PFFs) which screen the accelerating electric field \\citep{HM98,HM01,HM02}. These space-charge-limited-flow (SCLF) models have since been extended to allow for the variation of the CR PFF altitude across the PC and therefore acceleration of primaries along the last open magnetic field lines in a slot gap (SG) scenario \\citep{AS79,Arons83,MH03_SG,MH04_SG,Harding08_Crab}. The SG results from the absence of pair creation along these field lines, forming a narrow acceleration gap that extends from the neutron star (NS) surface to near the light cylinder. The SG model is thus a possible physical realization of the two-pole caustic (TPC) geometry \\citep{Dyks03}, developed to explain pulsar HE light curves. On the other hand, outer gap (OG) models \\citep{CHR86a,CHR86b,CR92,CR94,Romani96,Cheng00,Zhang04} assume that HE radiation is produced by photon-photon pair production-induced cascades along the last open field lines above the null-charge surfaces ($\\mathbf{\\Omega}\\cdot\\mathbf{B}=0$, with $\\Omega=2\\pi/P$), where the Goldreich-Julian charge density \\citep{GJ69} changes sign. The pairs screen the accelerating E-field, and limit both the parallel and transverse gap size \\citep{Takata04}. Classical OG models may be categorized as ``one-pole caustic models'', as the assumed geometry prevents observation of radiation from gaps (caustics) associated with both magnetic poles \\citep{Harding05}. More recently, however, \\citet{Hirotani06,Hirotani07} found and applied a 2D, and subsequently a 3D \\citep{Hirotani08} OG solution which extends toward the NS surface, where a small acceleration field extracts ions from the stellar surface in an SCLF-regime (see also \\citet{Takata04,Takata06}, and in particular \\citet{Takata08} for application to Vela). Lastly, \\citet{Takata09} modeled Geminga using an OG residing between a ``critical'' B-field line (perpendicular to the rotational axis at the light cylinder) and the last open field line.  Current models using dipole field structure to model MSPs predict largely unscreened magnetospheres due to the relatively low B-fields inhibiting copious magnetic pair production. Such pulsars may be described by a variation of the PC model (applicable for younger pulsars), which we will refer to as a ``pair-starved polar cap'' (PSPC) model \\citep{MH04_PS,HUM05,MH09_PS}. In a PSPC model, the pair multiplicity is not high enough to screen the accelerating electric field, and charges are continually accelerated up to high altitudes over the full open-field-line region. The formation of a PSPC ``gap'' is furthermore naturally understood in the context of an SG accelerator progressively increasing in size with pulsar age, which, in the limit of no electric field screening, relaxes to a PSPC structure.  Several authors have modeled MSP gamma-ray fluxes, spectra and light curves in both the PSPC \\citep{Frackowiak05a,Frackowiak05b,HUM05,VdeJ05,Venter_phd,Zajczyk08} and OG \\citep{Zhang03,Zhang07} cases. Collective emission from a population of MSPs in globular clusters \\citep{HUM05,Zhang07,BS07,Venter08,Venter09} and in the Galactic Center \\citep{Wang06} have also been considered. \\citet{Watters09} recently calculated beaming patterns and light curves from a population of canonical pulsars with spin-down luminosities $\\dot{E}_{\\rm rot} > 10^{34}$~erg\\,s$^{-1}$ using geometric PC, TPC, and OG models. They obtained predictions of peak multiplicity, peak separation, and flux correction factor $f_\\Omega$ as functions of magnetic inclination and observer angles $\\alpha$ and $\\zeta$, and gap width $w$. The latter factor $f_\\Omega$ is used for converting observed phase-averaged energy flux $G_{\\rm obs}$ to the total radiated (gamma-ray) luminosity $L_\\gamma$, which is important for calculating the efficiency of converting $\\dot{E}_{\\rm rot}$ into $L_\\gamma$. A good example is the inference of the conversion efficiencies of globular-cluster MSPs which may be collectively responsible for the HE radiation observed from 47~Tucanae by \\textit{Fermi}-LAT \\citep{Abdo09_Tuc}.  In this paper, we present results from 3D emission modeling, including Special Relativistic (SR) effects of aberration and time-of-flight delays, and rotational sweepback of B-field lines, in the geometric context of OG, TPC, and PSPC pulsar models. We study the newly-discovered gamma-ray MSP population \\citep{Abdo09_MSP}, and obtain fits for gamma-ray and radio light curves. Our calculation of light curves and flux correction factors $f_\\Omega(\\alpha,\\zeta,P)$ for the case of MSPs is therefore complementary to the work of \\citet{Watters09} which focuses on younger pulsars, although our TPC and OG models include non-zero emission width. Section~\\ref{sec:Model} deals with details of the various models we have applied. We discuss light curves from both observational and theoretical perspectives in Section~\\ref{sec:LCs}, and present our results and conclusions in Sections~\\ref{sec:Results} and~\\ref{sec:Con}. ",
        "conclusions": "\\label{sec:Con} We presented results from 3D emission modeling of gamma-ray and radio radiation in the framework of geometric PC, OG, and TPC pulsar models, and also for the full-radiation PSPC model. We have applied our results to recent measurements of newly-discovered MSPs by \\textit{Fermi}-LAT. In this sense, we present results complementary to those obtained by \\citet{Watters09} for young pulsars.  Previously, it was believed that most MSPs should have unscreened magnetospheres \\citep{HUM05}, as they lie below the predicted CR pair death line on the $P\\dot{P}$-diagram. It was expected that such pair-starved MSPs should have single gamma-ray pulses roughly in phase with the radio \\citep{VdeJ05}. From Figure~\\ref{fig:LC_com1} and~\\ref{fig:LC_com3}, we see the surprising fact that there are indeed MSPs that have double-peaked light curves well fit by TPC / OG models, as are many of the young gamma-ray pulsars. This is interpreted as indicating the operation of a magnetic pair formation mechanism, and copious production of pairs to set up the required emitting gap structure.  New ways of creating pairs in low-$\\dot{E}_{\\rm rot}$ pulsars will have to be found to explain this phenomenon. PSR~J0030+0451 illustrates this point very well in that it has the lowest $\\dot{E}_{\\rm rot}$ of the MSP sample ($3.5\\times10^{33}$~erg s$^{-1}$), therefore lying significantly below the calculated CR death line \\citep[e.g.,][]{HMZ02}, and yet exhibits the sharpest double peaks of the current population, implying emission originating in very thin TPC / OG gaps. The problem may be alleviated somewhat by increasing the stellar compactness $\\kappa^\\prime$ (larger mass or smaller radius), motivated by recent measurements of large MSP masses \\citep[up to $\\sim1.7M_\\odot$; see][and references therein]{Verbiest08,Freire09}. This will boost the GR E-fields, and enhance pair creation probability. Another way to do this would be to increase the magnetic field. B-fields that are larger than those usually inferred using the dipole spin-down model (and having smaller curvature radii) may be present when there are multipolar B-components near the surface (or an offset-dipole geometry).  In fact, offset dipoles have been suggested in modeling the X-ray light curves of MSPs J0437$-$4715 and J0030$+$0451 \\citep{Bogdanov07,Bogdanov09}.  However, detailed investigation of such a scenario and its implications for pair cascades is necessary to place this speculation on sure footing. Another possible origin for higher surface fields is the movement of magnetic poles toward the spin axis during the spin-up phase of an MSP \\citep{Lamb08}. It has been argued that during the spin-up to millisecond periods, the inward motion of the neutron star superfluid vortices produces a strain on the crust, causing the magnetic poles to drift toward the spin-axis \\citep{Ruderman91}. If the two poles are in the same hemisphere prior to spin-up, the poles drift toward each other, producing a nearly orthogonal rotator having the same dipole moment but a surface field that can be orders of magnitude higher \\citep{Chen93}.  We find that there is exclusive differentiation between the TPC / OG models on the one hand, and the PSPC model on the other hand. Six MSPs have gamma-ray light curves which lag the radio and are explained using TPC or OG fits, but not PSPC fits. For the remaining two MSPs, the radio light curves slightly lag the gamma-ray light curves, and these are fit by the PSPC model (and not by the TPC / OG models). It therefore seems that there are two subclasses emerging within the current gamma-ray MSP sample, and it is not obvious which pulsar characteristics provide a means to predict subclass membership. From our model light curve fitting, we furthermore find $(\\alpha,\\zeta)$ values which are in reasonable agreement with values inferred from MSP polarization measurements. Although we find good PSPC fits for the last two MSPs, we caution that the E-field is only approximately known (e.g., it follows from a local electrodynamical model based on a GR dipolar B-field). Future models which take global current flow patterns into account, along with more sophisticated B-field structure, may produce more realistic solutions for the E-field.  Our ability to discriminate between different classes of models derives from the fact that we produced both the gamma-ray and radio curves within the same model. We could then use the shape \\textit{and} relative radio-to-gamma phase lag provided by the data to obtain the best-fit model type for each MSP. The data also enabled us to conclude that the emission, in \\textit{all} models considered, must come from the outer magnetosphere. This has now been observed to be true for the bulk of the gamma-ray pulsar population \\citep{Abdo09_Cat}.  In the case of PSR~J0437-4715 and PSR~J0613-0200, we find that the TPC model predicts a significant precursor to the main gamma-ray peak, while the OG model predicts no such low-level emission. With more statistics, this effect may possibly become a discriminator between the TPC and OG models. (We assumed that the TPC emission region starts at $r_{\\rm em}=R$ when creating our plots. However, the relative intensity of the precursor and low-level emission predicted by the TPC model may be reduced by limiting the emission region's extension, i.e.\\ only collecting photons above a certain radius $r_{\\rm em} \\geq R_{\\rm min} > R$.)  We calculated the flux correction factor in the context of the different models, and found that $f_\\Omega \\sim1$. These values imply a wide beaming angle, and derives from the fact that we obtain best fits for large impact angles. \\citet{Venter_phd} previously found $\\Lambda_{\\rm avg} \\sim 10-30$ (i.e.\\ $f_\\Omega\\sim0.8-2.4$), and $\\Lambda_{\\max} \\sim300$ ($f^{\\max}_\\Omega\\sim24$) for the PSPF model. Now, we find $f_\\Omega$ $\\sim 0.5-2$, and $f^{\\max}_\\Omega\\sim 4$. These results are roughly consistent, with the differences stemming from the following: (i) \\citet{Venter_phd} used energy flux ratios to calculate $\\Lambda$, while we are using photon flux ratios to calculate $f_\\Omega$, assuming that the photon and energy fluxes have similar distributions across $(\\zeta,\\phi)$-space; (ii) \\citet{Venter_phd} only used $E_{||}^{(1)}$ and $E_{||}^{(2)}$ for the E-field, while we now also include the high-altitude solution ($E_{||}^{(3)}$) for the PSPF case. This leads to more intense high-altitude emission, and therefore smaller values of $f_\\Omega$ for off-beam emission.  We lastly remark that the larger radio beam widths of MSPs compared to those of canonical pulsars should lead one to expect relatively few radio-quiet MSPs.  The spectacular data from \\textit{Fermi}-LAT hold the promise of phase-resolved spectroscopy, at least for the brightest pulsars, and will challenge existing pulsar models to reproduce such unprecedented detail. Future work therefore includes using full acceleration and radiation models to study gamma-ray spectra, luminosities, and light curves, in order to constrain fundamental electrodynamical quantities, and possibly providing the opportunity of probing the emission geometry and B-field structure more deeply. Improved understanding of pulsar models will also feed back into more accurate population synthesis models \\citep[e.g.,][]{Story07}. In addition, we hope to obtain better understanding of important quantities such as MSP efficiencies, and whether this quantity is similar for Galactic-Field and globular-cluster MSPs \\citep{Abdo09_Tuc}."
    },
    "0911/0911.0102_arXiv.txt": {
        "abstract": "{We investigate claims according to which the X-ray selection of AGN is not as efficient compared to that based on  [OIII] selection because of the effects of X-ray absorption. We construct the predicted X-ray luminosity function both for all Seyferts as well as separately for Seyfert-1 and Seyfert-2 type galaxies, by combining the optical AGN [OIII] luminosity functions derived in SDSS with the corresponding $\\rm L_X-L_{[OIII]}$ relations. These relations are derived from   {\\it XMM-Newton} observations of all Seyfert galaxies in the Palomar spectroscopic sample of nearby galaxies after correction for X-ray absorption and optical reddening. We compare the predicted X-ray  luminosity functions with those  actually observed in the local Universe by {\\it HEAO-1}, {\\it RXTE} as well as {\\it INTEGRAL}. The last luminosity function is derived in the 17-60 keV region and thus is not affected by absorption even in the case of Compton-thick sources. In the common luminosity regions, the optically and X-ray selected  Seyfert galaxies show reasonable agreement. We thus find no  evidence  that the [OIII] selection provides a more robust tracer of powerful AGN compared to the X-ray. Still, the optical selection probes less luminous Seyferts compared to the current X-ray surveys. These low luminosity levels, are populated by a large number of X-ray unobscured Seyfert-2 galaxies. \\keywords {X-rays: general; X-rays: diffuse background; X-rays: galaxies}} ",
        "introduction": "Since the detection of the nearby AGN 3C273 in X-rays by {\\it UHURU} (Kellogg et al. 1971), X-ray observations have been considered to be the primary tool for selecting AGN. Recently, the deepest ever observations  in the Chandra Deep Field North and South (Alexander et al. 2003, Giaconni et al. 2002, Luo et al. 2008) have resolved 80-90\\% of the extragalactic X-ray light, the X-ray background, in the 2-10 keV band. These observations reveal a sky density of about 5000 sources per square degree (Bauer et al. 2004), the vast majority of which are AGN (for a review see Brandt \\&  Hasinger 2005). These surveys allowed the derivation of the luminosity function and probed with good accuracy the accretion history of the Universe (Ueda et al. 2003, La Franca et al. 2005, Barger et al. 2005).   In the local Universe, the X-ray luminosity function has been derived from the  wide-angle surveys of {\\it RXTE} and {\\it HEAO-1}. Sazonov \\& Revnitsev (2004) derived the luminosity function for bright AGN, detected at fluxes $>2.5\\times 10^{-11}$ \\funits in the 3-20 keV band, in the {\\it RXTE} slew survey. Shinozaki et al. (2006) derived the luminosity function of AGN  using the all-sky HEAO-1 in the 2-10 keV band. Even these hard X-ray surveys may be missing a number of extremely obscured AGN. In the case of Compton-thick AGN, at column densities $>10^{24}$ \\cunits (equivalent to $A_V\\sim 450$ using the Galactic dust-to-gas ratio), a large fraction of the intrinsic flux will be absorbed. The {\\it SWIFT}  (Gehrels et al. 2004) and the {\\it INTEGRAL} missions \\ (Winkler et al. 2003) which carry ultra-hard X-ray detectors ($>$15 keV), albeit  with limited imaging capabilities, probed energies which are  immune to X-ray obscuration up to column densities of $10^{25}$ \\cunits. These missions helped towards further constraining the number density of such heavily absorbed sources at very bright fluxes, $\\rm f_{17-60~keV} >10^{-11}$ \\funits in the local Universe, $z<0.1$ (Beckmann et al. 2006, Bassani et al. 2006,  Winter et al. 2008, Winter et al. 2009). The luminosity function in these energies has been derived by Sazonov et al. (2007), Paltani et al. (2008) and Tueller et al. (2009). It appears that the fraction of Compton-thick sources is small and thus the {\\it RXTE} and   {\\it HEAO-1} luminosity functions are little  affected.   In the optical, QSOs have been traditionally selected using colours (Schmidt \\& Green 1983, Marshall et al. 1987, Boyle et al. 2000). The advent of the 2dF (Croom et al. 2004) and the SDSS (York et al. 2000) surveys have generated vast samples of QSOs providing a leap forward in the study of their luminosity function. The derived sky density of QSOs (few hundred per square degree) is at least an order of magnitude lower than that derived from X-ray surveys. This is because the colour selection of AGN requires that the nuclear optical luminosity is much higher than that of the host galaxy for the AGN to be detected (typically $M_B<-23$). Therefore the optical selection based on colours is biased against low luminosity AGN in contrast to the X-ray selection.  The limitations of colour optical selection techniques can be circumvented by selecting AGN via their emission lines. Such methods can extend the optical luminosity function to low luminosities. Ho et al. (1997) carried out a spectroscopic survey of about 500 nearby galaxies selected from the revised Shapley-Ames catalogue (Sandage \\& Tammann 1981). They identified Seyfert emission-line characteristics in 52 galaxies. Ulvestad \\& Ho (2001) derived the optical B-band luminosity function from their sample extending the AGN luminosity function to $\\rm M_B\\sim -17$ (see also Georgantopoulos et al. 1999). A limitation in the above works is that the B-band is strongly affected by the host galaxy light. Hao et al. (2005a) extended significantly these results using the SDSS survey to select a sample of about 3000 AGN in the redshift range $0<z<0.15$.   Hao et al. (2005b) derive the $Ha$ and $[OIII]$ $\\lambda5007$ luminosity function separately for Seyfert-1 and Seyfert-2 galaxies. The [OIII] line is considered a much better proxy of the nuclear power as compared to the B-band luminosity, although in the case of  strong star-forming emission some contamination of the [OIII] emission is expected.   The question which arises is whether the optical emission lines or the X-ray  selection methods are more efficient for finding low-luminosity AGN. This can be addressed by  comparing the X-ray and optical AGN luminosity functions.   Heckman et al. (2005) addressed  this issue by converting the optical [OIII] Seyfert (combined type 1 and type 2) luminosity function of Hao et al. (2005b) to X-ray wavelengths using the relation between X-ray and [OIII] luminosity. Comparison with the {\\it RXTE} X-ray luminosity function reveals that the latter  lies consistently below the optical one. As these authors have not taken into account, intentionally, the X-ray absorption, they conclude that the mismatch between the two luminosity functions can be more naturally  attributed to absorption in X-ray wavelengths. The overall conclusion from this work is that the optical [OIII] emission may provide a more efficient method for picking AGN. In this paper we attempt to  further address this issue and to understand the reason for the disagreement found by Heckman et al. (2005). We derive the X-ray luminosity function again from the [OIII] luminosity function  of Hao et al. by combining it with the $\\rm L_X-L_{[OIII]}$  relation  derived from {\\it XMM-Newton} observations (Akylas \\& Georgantopoulos 2009) of all the Seyferts in the Ho et al. sample.   Our analysis involves the following improvements: \\begin{itemize} \\item{The available  {\\it XMM-Newton} spectra allow for correcting for the X-ray absorption}. \\item{We produce individually the luminosity function of Seyfert-1 and Seyfert-2 galaxies}. \\item{We derive the X-ray luminosity function from the bi-variate optical/X-ray luminosity function taking fully into consideration the $\\rm L_X-L_{[OIII]}$ relation and its dispersion}. \\end{itemize} We adopt $\\rm H_o=  75  \\,  km  \\,  s^{-1}  \\,  Mpc^{-1}$,  $\\rm \\Omega_{M}  =  0.3$, $\\Omega_\\Lambda = 0.7$ throughout the paper.   \\begin{table*} \\begin{center} \\caption{LogX-ray vs Log[OIII] luminosity relation in optically selected AGN} \\label{lxlo_ho} \\begin{tabular}{lccccc} \\hline Sample & Number & Slope &  Intercept & $\\rm <log(L_x/L_{[OIII]})>$ & Dispersion  \\\\ \\hline A. Seyfert-1 (Ho \\& Heckman)& 28 & $0.84\\pm0.09$  & $8.27\\pm 3.87$ & 1.66 & 0.60 \\\\ B. Seyfert-1 (Ho)           &  8 & $0.66\\pm0.22$ & $15.31\\pm8.84$ & 1.84 & 0.82 \\\\ C. Seyfert-1 (Heckman)      & 20 & $1.04\\pm0.17$ & $-0.17\\pm7.34$ & 1.59 & 0.48 \\\\ D. Seyfert-2  (Ho luminous+weak)      & 23 & $0.83\\pm0.15$ & $7.3\\pm6.1$ & 0.87 & 0.80 \\\\ E. Seyfert-2 (Ho luminous)  & 11  & $0.64\\pm0.17$ & $15.34\\pm7.15$ & 1.16 & 0.77 \\\\ F. Seyfert-2 (Ho weak)      & 12 & $0.59\\pm0.29$ & $16.18\\pm 11.48$ & 0.61 & 0.78 \\\\ G. Seyfert 1+ 2 (luminous)     & 17 & $0.65\\pm0.14$ & $15.35\\pm5.90$ & 1.44 & 0.84 \\\\ \\hline \\end{tabular} \\end{center} \\end{table*} ",
        "conclusions": "We derived the predicted X-ray luminosity function for Seyfert galaxies in the local Universe, by combining the optical SDSS [OIII] Seyfert luminosity functions with the corresponding  {\\lxlo} relation. These relations have been derived  using {\\it XMM-Newton} observations (Akylas \\& Georgantopoulos 2009) of the local, optically selected AGN sample of Ho et al. (1997). This sample covers a comparable luminosity range with the SDSS Seyfert sample. We have corrected the X-ray luminosity for the effects of absorption. Our analysis above shows that the predicted X-ray luminosity function is in reasonable agreement with the observed X-ray  Seyfert luminosity functions derived in the 2-10 keV and 2-20 keV bands by {\\it HEAO-1} and {\\it RXTE} respectively. Most importantly, it is in agreement with the  ultra-hard 17-60 keV {\\it INTEGRAL} luminosity function (Sazonov et al. 2007). As the {\\it INTEGRAL} luminosity function is practically immune to X-ray absorption, this suggests that absorption played little role in the optical/X-ray luminosity function discrepancy reported by Heckman et al. (2005). This is also independently supported by the {\\it XMM-Newton} observations of the Palomar optically selected Seyfert sample (Akylas \\& Georgantopoulos 2009). The fraction of Compton-thick AGN in this optically selected sample is small.  In addition we examined separately, the Seyfert-1 and Seyfert-2 luminosity function. The Seyfert-1 luminosity function is in rough agreement or even somewhat above the observed X-ray luminosity function. The predicted Seyfert-2 luminosity function agrees quite well with the optical luminosity function again disfavouring the absorption hypothesis. Indeed, if absorption were the problem, then our predicted luminosity function, which takes X-ray absorption into account, would be well above that of {\\it RXTE}.  Having addressed this matter, we need to understand  why Heckman et al. (2005) found that the {\\it RXTE} X-ray luminosity function lies below the optical one by about a factor of 3. The comparison was based on the fact that the two luminosity functions have the same slope. However,  this is true only for the  bright part of the X-ray luminosity function (i.e. for L$_{\\star}(3-20~ \\rm keV)>3\\times 10^{43}$ \\lunits or equivalently L$_{\\star}(2-10 \\rm keV)>2.3 \\times 10^{43}$ \\lunits). Using $\\rm (L_X/L_{[OIII]}) \\approx 100 $ the corresponding optical luminosities must be  greater than $5.7 \\times 10^7 L_{\\odot}$. Since the upper end of the optical luminosity function of Hao et al. (2005b) is $\\sim 3 \\times 10^8 L_{\\odot}$  the  overlap is very limited (see Fig.5 in Heckman et al.) and thus the comparison is not quite robust. Moreover, the effect of the dispersion plays a critical role as shown here.  Although, the densities of the [OIII] and the X-ray selected AGN  are comparable, this does not mean that  these methods favour the selection of the same objects. The optical selection favours X-ray weak Seyfert-2 (see Fig. 1). It has been proposed (see Ho 2008 for a review) that a large number of sources at low luminosities appear as Seyfert-2 possibly because no  Broad-Line-Region is formed at low accretion rates (Nicastro 2000) or low luminosities (Elitzur  \\& Shlosman 2006).  If the above models hold true, these 'naked' Seyfert-2 galaxies should present no hidden Broad Lines in spectropolarimetric observations (e.g. Tran 2003). Interestingly, Spinoglio et al. (2009) find that the  ratio of the X-ray to 12$\\mu m$ luminosity, $L_X/L_{12\\mu m}$, is much lower in the Seyfert-2 with no hidden Broad-Line-Region, suggesting that these are weak X-ray emitters. There is however, one caveat in these 'naked' X-ray weak Seyfert-2 interpretation. The [OIII] (or the mid-IR) luminosity is believed to be a good proxy of the ionizing nuclear luminosity and thus to the X-ray luminosity. But in these Seyfert-2 sources the $L_X/L_{[OIII]}$ ratio is low, implying that only the X-ray luminosity is weak. One explanation could be that star-formation is contributing a large part of the  [OIII] emission.  The future missions {\\it NUSTAR} and {\\it ASTRO-H} having superb imaging capabilities ($\\sim$ 1arcmin) at ultra-hard energies (10-70 keV), will provide the opportunity to obtain the least unbiased AGN samples, reaching  flux levels at least two orders of magnitude fainter than {\\it SWIFT} and {\\it INTEGRAL}."
    },
    "0911/0911.3101_arXiv.txt": {
        "abstract": "s{ \\ps\\ was successfully launched on May 14th, 2009, from the Kourou space port, in French Guyana. After recalling the objectives that we set out - back in 1996 - to fulfill with this project, I recall some of the technological breakthroughs which needed to be made and report on the exciting scientific outlook of the project in light of the knowledge we now have of the actual performances of the two on-board instruments. I also include one of our more recent results even though it was not yet available at the time of the conference. }   The measurement goals of \\plancks may be stated rather simply: to build an experiment able to perform the ``ultimate'' measurement of the primary CMB temperature anisotropies, which requires:\\begin{itemize} \\item full sky coverage and a good enough angular resolution in order to completely mine all scales at which the Cosmic Microwave background (CMB) primary anisotropies contain information ($\\simgt 5$ minutes of arc) \\item a final sensitivity essentially limited by the ability to remove the astrophysical foregrounds, implying a large frequency coverage from 30 GHz to 1 THz (provided by the two instruments: HFI and LFI), with sensitivity at each of the 9 survey frequencies in line with the role of each map in determining the CMB properties . \\end {itemize} For the measurement of the polarisation of the CMB anisotropies, \\plancks goal was ``only'' to get the best polarisation performances with the technology available at the design time  This is on these simple but ambitious goals and the proposed way of reaching them that, after 3 years of preparatory work, the project was selected by the European Space Agency (ESA), as the $3^{rd}$ Medium size mission of its Horizon 2000+ program. This selection occurred in march 1996, \\ie contemporaneously with that of WMAP by NASA, which rather proposed reaching earlier less ambitious goals with already existing technology.  \\newcommand{\\hf}{\\hfill} \\begin{table}[htb] \\caption{Summary of \\plancks performance goals for the required 14 months of routine operations, which allows nearly all detectors to map the entire sky twice. Requirements on sensitivity are simply two times worse than the stated goals. Central band frequencies, $\\nu$, are in Gigahertz, the FWHM angular sizes are in arc minute, and the (sky-averaged) sensitivities $c^X_{noise}$, with $X = T, Q$ or $U$, are expressed in $\\mu \\mathrm{K.deg}$; this number indicates the rms detector noise, expressed as a equivalent temperature fluctuation in $\\mu \\mathrm{K}$, whcih is expected once it is averaged in a pixel of 1 degree of linear size).} \\vspace{0.4cm} \\begin{center} \\scalebox{1.}{ \\begin{tabular}{|c||c|c|c||c|c|c|c|c|c|} \\hline \\hline & \\multicolumn{3}{|c||}{\\lfi\\ goals}&\\multicolumn{6}{|c|}{\\hfis goals} \\\\ \\hline\\hline $\\nu$ \\hf [GHz] & \\hf 30 & \\hf 44 & \\hf 70 & \\hf 100 & \\hf 143 & \\hf 217 & \\hf 353 & \\hf 545 & \\hf 857 \\\\ FWHM \\hf [arcmin] & \\hf 33 & \\hf 24 & \\hf 14 & \\hf 9.5 & \\hf 7.1 & \\hf 5.0 & \\hf 5.0 & \\hf 5.0 & \\hf 5.0 \\\\ $c^T_{noise}$ \\hf [$\\mu \\mathrm{K.deg}$]& \\hf 3.0 & \\hf 3.0 & \\hf 3.0 & \\hf 1.1 & \\hf 1.4 & \\hf 2.2 & \\hf 6.8 & \\hf - & \\hf - \\\\ $c^{Q or U}_{noise}$ \\hf [$\\mu \\mathrm{K.deg}$]& \\hf 4.5 & \\hf 4.6 & \\hf 4.6 & \\hf 1.8 & \\hf 1.4 & \\hf 2.4 & \\hf 7.3 & \\hf & \\hf \\\\ \\hline \\end{tabular} } \\end{center} \\label{tab:perfs} \\end{table}  Table~\\ref{tab:perfs} summarises the main performance goals of \\ps, expressed for instance as the average detector noise within a square patch of 1 degree of linear size, $c_{noise}$, for the 14 months baseline duration of the mission, which would allow covering twice all the sky by nearly all the detectors. It is interesting to note that if we take the noise performance figure for the average of the central CMB frequencies (the 100-143-217 GHz \\hfis\\ channels, assuming all the other channels are devoted to foregrounds removal), one finds $0.5\\ \\mathrm{\\mu K.deg}$ in temperature and $1\\ \\mathrm{\\mu K.deg}$ for the Q \\& U Stokes parameters. The magnitude of this step forward, if achieved, may be judged by comparing with the WMAP sensitivity which is given in Table~\\ref{tab:WMAPperfs}.  The aggregate sensitivity of the WMAP 60 \\& 90 GHz channels is $\\sim 10.8\\ \\mathrm{\\mu K.deg}$ in a year, which would imply about $(10.8/.5)^2\\sim$ 460 years of operations to reach the baseline \\ps\\ sensitivity. In other words, the error bars from noise in the angular power spectrum should be at least hundred times smaller for \\plancks than for \\map (with an even larger difference at smaller scales which will be much better known from \\plancks\\ thanks to the twice higher angular resolution of HFI).  \\begin{table}[htb] \\caption{Summary of WMAP in-flight performance per full year of operations. Same units than in Table~\\ref{tab:perfs}. } \\vspace{0.4cm} \\begin{center} { % \\begin{tabular}{|c||c|c|c|c|c|} \\hline \\hline \\multicolumn{6}{|c|}{\\maps (in flight) } \\\\ \\hline\\hline $\\nu$ \\hf [GHz] & \\hf 23 & \\hf 33 & \\hf 41 & \\hf 61 & \\hf 94 \\\\ FWHM \\hf [arc min] & \\hf 49.2 & \\hf 37.2 & \\hf 29.4 & \\hf 19.8 & \\hf 12.6 \\\\ $c^T_{noise}$ \\hf [$\\mu \\mathrm{K.deg}$] & \\hf 12.6 & \\hf 12.9 & \\hf 13.3 & \\hf 15.6 & \\hf 15.0 \\\\ \\hline \\end{tabular} } % \\end{center} \\label{tab:WMAPperfs} \\end{table}  \\begin{figure}[htbp] \\begin{center} \\psfig{file=Planck-build-up.jpg, width=\\textwidth} \\end{center} \\caption[]{Planck build-up from the inside out. Going from left to right, one sees on the top row (t1) the HFI instrument with its 52 detector horns poking out of it's outer shell at 4\\,K. (t2) HFI is surrounded by the 11 larger horns from LFI. HFI and LFI together form (t3) the focal plane assembly from which (t4) the electrical signal departs (though a bunch of wave guides for the LFI and a harness of wires for HFI) to connect to the warm electronics parts of the detection chain which are located within the service module at $\\sim$ 300\\,K. (t5) The (top) cold and warm (bottom) parts are separated by three thermally isolating V-grooves which allow radiating to space heat from the spacecraft sideways and quite efficiently. The third (top) V-grooves operating temperature is about 40\\,K. On the middle row, one sees (m1) the beds of the sorption cooler and it's piping around the V-groove, bringing the overall focal-plane structure to LFI's operational temperature of $\\sim 18$\\, K. (m2) The back-to-back (to damp the first harmonics of the vibrations) compressors of the 4\\,K cooler allow bringing HFI outer shell to 4\\,K, while (m3) isotopes from the He3 tank and the three He4 tanks are brought to the mixing pipes within HFI to cool filters (within the horns) to 1.6\\,K and the bolometer plate to 0.1\\,K, before (m4) being released to space. (m5) The passive cooling and the three active stage constitute this complex but powerful cooling chain in space. On the last bottom row, one can also see (b1) some of the electronic boxes in the service module (SVM) which in addition to the warm part of the electronic and cooling chains also contain all ``services'' needed for transmitting data, reconstructing the spacecraft attitude, powering the whole satellite... The bottom of the SVM is covered with solar panels, while supporting struts begin on its top which allows (b3) positioning the secondary and primary reflectors. The top part is surrounded by a large baffle to shield at best the focal plane from stray-light. The back view (b5) allows distinguishing in the back the supporting structure of the primary mirror, and the wave guides from LFI. The spin axis of Planck (vertical on these plots) is meant to be close to the sun-earth line, with the solar panel near perpendicular to that line and the rotation of the line-of-sight (at 1 rpm) causing the detectors to survey circles on the sky with an opening angle around 85 degrees. Copyright ESA.} \\label{fig:build} \\end{figure}  We proposed to achieve the ambitious sensitivity goals of \\plancks with a small number of detectors, limited principally by the photon noise of the background (for the CMB ones), in each frequency band. This implied to achieve several technological feats never achieved in space before (see in particular Lamarre et al. 2003, in New Astronomy Reviews, 47, pp. 1017): \\begin{itemize} \\item sensitive \\& fast bolometers with a Noise Equivalent Power $< 2 \\times 10^{-17}\\ \\mathrm{W/Hz}^{1/2}$ and time constants typically smaller than about 5 milliseconds (which thus requires cooling them down to $\\sim 100$ mK, and build them with a very low heat capacity \\& charged particles sensitivity) \\item total power read out electronics with very low noise, $< 6\\ \\mathrm{nV/Hz}^{1/2}$ from 10 mHz (1 rpm) to 100 Hz (\\ie from the largest to the smallest angular scales to measure at the \\ps\\ scanning speed) \\item excellent temperature stability, from 10 mHz to 100 Hz (cf. Lamarre et al. 04), such that the induced variation be a small fraction of the detector temporal noise:  \\begin{itemize} \\item better than $ 10\\ \\mathrm{\\mu K/Hz}^{1/2}$ for the 4K box (assuming 30\\% emissivity) \\item better than $30\\ \\mathrm{\\mu K/Hz}^{1/2}$ on the 1.6K filter plate (assuming a 20\\% emissivity) \\item better than $20\\ \\mathrm{nK/Hz}^{1/2}$ for the detector plate (a damping factor $\\sim 5000$ needed) \\end{itemize} \\item very low noise HEMT amplifiers (therefore cooled to 20\\,K) \\& very stable cold reference loads (at 4\\,K) \\end{itemize} In addition, \\plancks requires:\\begin{itemize} \\item a low emissivity telescope with very low side lobes (\\ie strongly under-illuminated) \\item no windows, and minimum warm surfaces between the detectors and the telescope \\item a quite complex cryogenic cooling chain (cf. figure~\\ref{fig:build}) which begins by reaching $\\sim 40K$ via passive cooling, by radiating about 2 Watts to space, followed by three active stages, at 20\\,K, 4\\,K, and 0.1\\,K: \\begin{itemize} \\item 20\\,K for the LFI, with a large cooling power, $\\sim 0.7$ Watts (provided by $H_2$ Joule-Thomson sorption pumps developed by JPL, USA) \\item 4\\,K, 1.6\\,K and 100\\,mK for the HFI (the 15 milli-Watts cooling power at 4K is provided by mechanical pumps provided by the RAL, UK, in order to perform a Joule-Thomson expansion of He; the 1.6K stage has a pre-cooling power of about 0.5 milli-Watts, thanks to another Joule-Thomson expansion, while the final dilution fridge of He$^3$ \\& He$^4$, from a French collaboration between Air Liquide, the CRTBT, can deal with 0.2  micro-Watts at 0.1\\,K). \\item a thermal architecture optimised to damp thermal fluctuations (active+passive) \\end{itemize} \\end{itemize} Furthermore, a tight control of vibrations is needed, in particular since the dilution cooler does not tolerate micro-vibrations at sub-mg level. And as little as $7\\times 10^{10}$ He atoms accumulated on the dilution heat exchanger (an He pressure typically at the $1\\times 10^{-10}$ mb level) would be too much.  These top-level design goals have now been turned into real instruments, which went through several qualification models. Before delivering the actual flight model of both instruments to industry for integration with the satellite, both instrumental consortia organised extensive calibrations campaigns starting at the individual components levels, then at the sub-systems levels (e.g. individual photometric pixels), then at instrument level. For HFI, the detector-level tests were done mainly at JPL in the USA, and the pixel level tests were performed in Cardiff in the UK, while the flight instrument calibration was performed at the Institut d'Astrophysique Spatiale in Orsay, France from April till the end of July 2006. During that period, we obtained in particular 19 days of scientific data at normal operating conditions. We could then confirm that HFI satisfies all our {\\em requirements}, and for the most part actually reaches or exceeds the more ambitious design {\\em goals}, in particular concerning the sensitivity, and speed of the bolometers, the very low noise of the read out electronics and the overall thermal stability. In addition, the total optical efficiency has been verified to be satisfactory, optical cross-talk appears negligible, as well as the Current cross-talk and the cross-talk in intensity is weak. Main beams are well-defined and are quite well described by the models, polarisation measurements confirm expectations, etc. The LFI instrument also went through detailed testing around the same time and it does reach most of its ambitious requirements.  \\begin{figure}[htbp] \\begin{center} \\psfig{file=csl.jpg, width=\\textwidth} \\end{center} \\caption[]{Planck on May 2008, hanging outside the vacuum cryogenic chamber at CSL before completing it's first and last full thermal vacuum test. Copyright ESA.} \\label{fig:CSL} \\end{figure}  The integration of the LFI and HFI instruments was performed at Thales premises in Cannes in November 2006 and within a year, by December 2007, the full satellite was ready for vibration testing. Planck was then flown from Cannes to ESA's ESTEC centre (in Noordwijk, Holland) where among other things it went through load balancing on April 7th, before travelling again to the ``Centre Spatial de Li\\'eges'' (CSL) in april 2008. Figure~\\ref{fig:CSL} is a picture of Planck hanging outside the vacuum cryogenic chamber at CSL, before the start of the first (and last) full thermal test with all elements of the cryogenic chain present and operating. This ultimate system-level (ground) test demonstrated in particular the following: \\begin{itemize}  \\item the dilution system can work with the minimal Helium 3 and 4 flux, which should allow 30 months of survey duration (nominal duration being 14 months!)  \\item the extremely demanding temperature stability required (at 1/5 of the detection noise) has been verified  \\item bolometers sensitivities in flight conditions are indeed centered around the goal, as shown in figure~\\ref{fig:bolos}. \\end{itemize}  \\begin{figure}\\begin{center} \\psfig{figure=Bolos.jpg,width=0.6\\textwidth} \\end{center} \\caption[]{Measured values of the Noise Equivalent Power of HFI detectors during the ground test in flight conditions at CSL on May 2008. One sees that the median value is at the goal level.} \\label{fig:bolos} \\end{figure}  \\ps\\ was then shipped to Kourou, and after a few more nerve \u2013wracking delays, we finally lost sight of Planck for ever (when it was covered by the SYLDA support system on the top of which laid Herschel for a joint launch). Launch was on May 14th, and it was essentially perfect. After separating from Herschel, \\ps\\ was set in rotation and started its to the L2 Lagrange point of the sun-earth system, at 1.5 million kilometres away from earth, \\ie about 1\\% further away from the sun than the earth. The final injection in the L2 orbit was at the end of June, shortly after the end of the Blois meeting (see figure~\\ref{fig:trajcool}-a), at the same time than the cooling sequence ended successfully. Indeed, figure~\\ref{fig:trajcool}-b shows how the various thermal stages reached their operating temperature, cooling of coarse from the outside-in, and closely following the predicted pattern.  \\begin{figure}\\begin{center} \\psfig{figure=trajectory.jpg,width=0.45\\textwidth} \\psfig{figure=cooling.jpg,width=0.54\\textwidth} \\end{center} \\caption[]{a) Spacecraft trajectory to and on the L2 orbit. b) Cooling sequence of Planck, showing the various stages reaching in turn their operational temperature, till the dilution plate actually reached 93 mK on July 3rd. Credit ESA and HFI consortium} \\label{fig:trajcool} \\end{figure}  Once at L2, a calibration and performance verification phase was conducted till mid-august, to insure that all system are working properly and that instrumental parameters are all set at best. From August 13th to 27th, we conducted a ``First Light Survey'' (FLS) in normal operational mode for an ultimate verification of parameters and of the long-term stability of the experiment. We found the data quality to be excellent, and the Data processing Centre pipelines could be operated as hoped to produce the first images. Figure~\\ref{fig:strip} (extracted from the Press release we made on September 17th) illustrates the FLS coverage by showing an image generated from the data acquired from a single 100GHz detector of HFI superimposed to an image of the optical sky by Axel Mellinger. We also released (see the press release in English at ESA's site http://sci.esa.int/science-e/www/object/index.cfm?fobjectid=45543 and in French at http://public.planck.fr/actualiteFLS.php) a comparison of a high latitude field whose emission ought to be dominated by the CMB (shown by a small white square in the figure) as observed by an HFI and LFI detector. They demonstrate an excellent similarity while the two instruments are using quite different technologies. Nine images at all the frequencies covered by Planck of a Galaxy crossing area (indicated by the large white square of the figure) provide visual evidence to the richness of the dataset that Planck shall deliver, allowing very broad scientific studies outreaching its primary cosmological goals. Indeed an important part of Planck long term legacy will be the unique set of  maps of the millimetric and sub-millimetric polarised full sky.  \\begin{figure}\\begin{center} \\psfig{figure=PR1.jpg,width=\\textwidth} \\end{center} \\caption[]{Sky coverage during Planck First Light Survey (FLS). Planck measures a ring of about 1 degree per day (for a full rotation around the sun in a year), yielding this $\\sim 15$ degree strip during the two weeks of the FLS. The detector image shown in Galactic coordinates is superimposed to a background image where the Galactic plane is clearly visible. Credit ESA and HFI \\& LFI consortia. } \\label{fig:strip} \\end{figure}  With the success of the FLS, the normal survey operations have now started. We should therefore be in a position to deliver in December 2010, as planned, an \u201cEarly Release Point Source Catalogue\u201d based on a rapid analysis of the first coverage of the sky which will be issued in time to allow the astrophysical community at large to propose follow-ups  by Herschel during its expected cryogenic life. A first public release based on the data from the nominal 14 months mission is slated for december 2012. The release should contain the clean calibrated time-ordered data of each detector, the nine full sky maps at the six frequencies covered by HFI and the three ones by LFI, possibly supplemented by polarisation maps, as well as maps of identified astrophysical components (CMB, Galactic Emissions, Extragalactic sources catalogue), some ancillary information (\\eg on beams, spectral transmission, etc), accompanied by about 50 scientific papers describing the mission, how the ``products'' were obtained, validated, and the results of a first pass of scientific exploitation by the Planck collaboration itself, encompassing in particular the implications of the measured statistical properties of the CMB. Our anticipation from the measured Helium consumption in flight is that the mission duration will exceed the nominal duration of 14 months, and we plan a further release, about a year later on the basis of the extra data which might allow covering as much as five times in total the entire sky. In addition to an improved sensitivity, this extra duration will foremost allow greater data redundancy and therefore a tighter control of all systematic effect which can be searched for with a longer baseline. This should allow us detecting the gravitational wave stochastic background predicted in one interesting class of inflationary models, providing the long thought after ``smoking gun'' of inflation, or otherwise put meaningful constraints on the viable inflation models remaining.  {\\em In conclusion, \\ps\\ is now in normal operation \\& performances are as expected or better.}\\\\ This gives us confidence that the scientific program of \\ps\\ can be carried through as anticipated. The dataset should in particular allow addressing many key cosmological questions, including the existence of a primordial gravitational wave background, or that of highly revealing deviations from the current minimal model, where the primordial fluctuation can be purely Gaussian, adiabatic, scale-free, in a strictly flat spatial geometry with a dark energy component describable as a pure cosmological constant, and (cosmologically) negligible neutrinos masses.  A rather complete overview of the scientific Program of \\ps\\ can be found in the so-called ``Blue Book'' which was issued in 2004. It can be downloaded from \\\\ {\\em http://www.planck.fr/IMG/pdf/Planck\\_book.pdf}. \\\\ In addition, we submitted a series of pre-launch papers (all with a title starting with ``Planck pre-launch status:'') giving many details of the design and tests of the mission, the instruments, and some of their components. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0911/0911.3271_arXiv.txt": {
        "abstract": "% The habitability of planets is strongly affected by impacts from comets and asteroids. Indications from the ages of Moon rocks suggest that the inner Solar System experienced an increased rate of impacts roughly 3.8 Gya known as the Late Heavy Bombardment (LHB). Here we develop a model of how the Solar System would have appeared to a distant observer during its history based on the Nice model of Gomes et al. (2005). We compare our results with observed debris discs. We show that the Solar System would have been amongst the brightest of these systems before the LHB. Comparison with the statistics of debris disc evolution shows that such heavy bombardment events must be rare occurring around less than 12\\% of Sun-like stars. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0911/0911.1274_arXiv.txt": {
        "abstract": "The prediction of the spin of the black hole resulting from the merger of a generic black-hole binary system is of great importance to study the cosmological evolution of supermassive black holes. Several attempts have been recently made to model the spin via simple expressions exploiting the results of numerical-relativity simulations. Here, I first review the derivation of a formula, proposed in Ref.~\\cite{new}, which accurately predicts the final spin magnitude and direction when applied to binaries with separations of hundred or thousands of gravitational radii.  This makes my formula particularly suitable for cosmological merger-trees and N-body simulations, which provide the spins and angular momentum of the two black holes when their separation is of thousands of gravitational radii. More importantly, I investigate the physical reason behind the good agreement between my formula and numerical relativity simulations, and nail it down to the fact that my formula takes into account the post-Newtonian precession of the spins and angular momentum in a consistent manner. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0911/0911.1056_arXiv.txt": {
        "abstract": "The recent detection of blazar 3C279 by MAGIC has confirmed previous indications by H.E.S.S. that the Universe is more transparent to very-high-energy gamma rays than previously thought. We show that this fact can be reconciled with standard blazar emission models provided photon oscillations into a very light Axion-Like Particle occur in extragalactic magnetic fields. A quantitative estimate of this effect explains the observed spectrum of 3C279. Our prediction can be tested in the near future by the satellite-borne GLAST detector as well as by the ground-based Imaging Atmospheric Cherenkov Telescopes H.E.S.S., MAGIC, CANGAROO III, VERITAS and by the Extensive Air Shower arrays ARGO-YBJ and MILAGRO. ",
        "introduction": "As is well known, in the very-high-energy (VHE) band above $100 \\, {\\rm GeV}$ the horizon of the observable Universe rapidly shrinks as the energy further increases. This comes about because photons from distant sources scatter off background photons permeating the Universe, thereby disappearing into electron-positron pairs~\\cite{stecker1971}. The corresponding cross section $\\sigma (\\gamma \\gamma \\to e^+ e^-)$ peaks where the VHE photon energy $E$ and the background photon energy $\\epsilon$ are related by $\\epsilon \\simeq (500 \\, {\\rm GeV}/E) \\, {\\rm eV}$. Therefore,  for observations performed by Imaging Atmospheric Cherenkov Telescopes (IACTs) -- which probe the energy interval $100 \\, {\\rm GeV} - 100 \\,{\\rm TeV}$ -- the resulting cosmic opacity is dominated by the interaction with ultraviolet/optical/infrared diffuse background photons (frequency band $1.2 \\cdot 10^{3} \\, {\\rm GHz} - 1.2 \\cdot 10^{6} \\, {\\rm GHz}$, corresponding to the wavelength range $0.25 \\, \\mu {\\rm m} - 250 \\, \\mu {\\rm m}$), usually called Extragalactic Background Light (EBL), which is produced by galaxies during the whole history of the Universe. Neglecting evolutionary effects for simplicity, photon propagation is controlled by the photon mean free path ${\\lambda}_{\\gamma}(E)$ for $\\gamma \\gamma \\to e^+ e^-$, and so the observed photon spectrum $\\Phi_{\\rm obs}(E,D)$ is related to the emitted one $\\Phi_{\\rm em}(E)$ by \\begin{equation} \\label{a1} \\Phi_{\\rm obs}(E,D) = e^{- D/{\\lambda}_{\\gamma}(E)} \\ \\Phi_{\\rm em}(E)~. \\end{equation}  Within the energy range in question, ${\\lambda}_{\\gamma}(E)$ decreases like a power law from the Hubble radius $4.2 \\, {\\rm Gpc}$ around $100 \\, {\\rm GeV}$ to $1 \\, {\\rm Mpc}$ around $100 \\, {\\rm TeV}$~\\cite{CoppiAharonian}. Thus, Eq.~(\\ref{a1}) entails that the observed flux is {\\it exponentially} suppressed both at high energy and at large distances, so that sufficiently far-away sources become hardly visible in the VHE range and their observed spectrum should anyway be {\\it much steeper} than the emitted one.  Yet, observations have {\\it not} detected the behaviour predicted by Eq.~(\\ref{a1}). A first indication in this direction was reported by the H.E.S.S. collaboration in connection with the discovery of the two blazars H2356-309 ($z = 0.165$) and 1ES1101-232 ($z = 0.186$) at $E \\sim 1 \\, {\\rm TeV}$~\\cite{aharonian:nature06}. Stronger evidence comes from the observation of blazar 3C279 ($z = 0.536$) at $E \\sim 0.5 \\, {\\rm TeV}$ by the MAGIC collaboration~\\cite{3c}. In particular, the signal from 3C279 collected by MAGIC in the region $E<220$ GeV has more or less the same statistical significance as the one in the range 220 GeV $< E <$ 600 GeV ($6.1 \\sigma$ in the former case, $5.1 \\sigma$ in the latter).  A suggested way out of this difficulty relies upon the modification of the standard Synchro-Self-Compton (SSC) emission mechanism. One option invokes strong relativistic shocks~\\cite{Stecker2007}. Another rests upon photon absorption inside the blazar~\\cite{Aharonian2008}. While successful at substantially hardening the emission spectrum, these attempts fail to explain why {\\it only} for the most distant blazars does such a drastic departure from the SSC emission spectrum show up.  Our proposal -- usually referred to as the DARMA scenario -- is quite different~\\cite{drm}. Implicit in previous considerations is the hypothesis that photons propagate in the standard way throughout cosmological distances. We suppose instead that photons can oscillate into a new very light spin-zero particle -- named Axion-Like Parlicle (ALP) -- and vice-versa in the presence of cosmic magnetic fields, whose existence has definitely  been proved by AUGER observations~\\cite{auger}. Once ALPs are produced close enough to the source, they travel {\\it unimpeded} throughout the Universe and can convert back to photons before reaching the Earth. Since ALPs do not undergo EBL absorption, the  {\\it effective} photon mean free path ${\\lambda}_{\\gamma , {\\rm eff}} (E)$ gets {\\it increased} so that the observed photons cross a distance in excess of ${\\lambda}_{\\gamma}(E)$. Correspondingly, Eq. (\\ref{a1}) becomes \\begin{equation} \\label{a1bis} \\Phi_{\\rm obs}(E,D) = e^{- D/{\\lambda}_{\\gamma , {\\rm eff}}(E)} \\ \\Phi_{\\rm em}(E)~, \\end{equation} from which we see that even a {\\it slight} increase of ${\\lambda}_{\\gamma , {\\rm eff}}(E)$ gives rise to a {\\it huge} enhancement of the observed flux. It turns out that the DARMA mechanism makes ${\\lambda}_{\\gamma , {\\rm eff}}(E)$ shallower than ${\\lambda}_{\\gamma}(E)$ although it remains a decreasing function of $E$. So, the resulting observed spectrum is {\\it much harder} than the one predicted by Eq. (\\ref{a1}), thereby ensuring agreement with observations even for a {\\it standard} SSC  emission spectrum. As a bonus, we get a natural explanation for the fact that only the most distant blazars would demand $\\Phi_{\\rm em}(E)$ to substantially depart from the emission spectrum predicted by the SSC mechanism.  Our aim is to review the main features of our proposal as well as its application to blazar 3C279. ",
        "conclusions": ""
    },
    "0911/0911.0927_arXiv.txt": {
        "abstract": "We demonstrate that stars beyond the virial radii of galaxies may be generated by the gravitational impulse received by a satellite as it passes through the pericenter of its orbit around its parent.  These stars may become energetically unbound (escaped stars), or may travel to further than a few virial radii for longer than a few Gyr, but still remain energetically bound to the system (wandering stars).  Larger satellites (10-100\\% the mass of the parent), and satellites on more radial orbits are responsible for the majority of this ejected population. Wandering stars could be observable on Mpc scales via classical novae, and on 100 Mpc scales via SNIa.  The existence of such stars would imply a corresponding population of barely-bound, old, high velocity stars orbiting the Milky Way, generated by the same physical mechanism during the Galaxy's formation epoch.  Sizes and properties of these combined populations should place some constraints on the orbits and masses of the progenitor objects from which they came, providing insight into the merging histories of galaxies in general and the Milky Way in particular. ",
        "introduction": "One of the triumphs of the last decade is orders of magnitude increase in the numbers of stars catalogued by large sky surveys such as the Sloan Digital Sky Survey and the Two Micron All Sky Survey. These catalogs contain significant samples of stars out to $\\sim$100kpc from the center of the Galaxy, allowing these regions to be probed through number counts to far lower surface brightness than previously possible and resulting in the discovery of numerous new dwarf satellite galaxies \\citep[e.g.][]{2005ApJ...626L..85W,2006ApJ...647L.111B}, as well as tidal tails from dwarf galaxies possibly disrupted long ago \\citep[e.g.][]{2003ApJ...599.1082M, 2005MNRAS.357...17B}.  We move from this decade of large sky surveys into a decade of continual all-sky surveying with the introduction of the Panoramic Survey Telescope \\& Rapid Response System \\citep[Pan-STARRS, see][]{2004SPIE.5489...11K} and the prospect of the Large Synoptic Survey Telescope \\citep[LSST, see ][]{ 2007AAS...210.6605I}. These surveys will be sensitive to ever fainter magnitude objects and hence probe ever lower surface brightness structures in ever increasing volumes of space.  The repeated nature of these surveys also adds the extra dimension of time, enabling large-sample statistical studies of variable phenomena in a way that was not previously possible.  For example, \\cite{2006AJ....131.2980S} points out that LSST will detect many new classical novae (hereafter CN).  While interesting in their own right, the special relationship between the peak brightness and decline time relationship of CN, coupled with their high luminosity, will allow us to map a volume out to 40 Mpc from Earth. This will be the first map of stars between galaxies and outside of clusters or groups of galaxies (i.e. intercluster stars) ever produced.  But should we expect to find anything there?  Red giant stars, planetary nebulae and classical novae  have been detected outside of galaxies but {\\it within clusters} \\citep[i.e. intracluster light - hereafter ICL, for example see][]{2004MNRAS.355..159W,2005ApJ...618..692N}. These populations are thought to begin their lives in galaxies: stars are liberated from their original hosts during galactic collisions, galactic harassment and tidal shredding, which generate tails of debris that can be ripped from the galaxy system by the tidal field of the cluster \\citep{1996Natur.379..613M}. This evolutionary picture is confirmed by the general agreement of observations and simulations on the level of ICL \\citep{2004ApJ...607L..83M},which photometric studies define as $0  -  50\\%$ of the total cluster light \\citep{2004bdmh.confE..26A,2004IAUS..217...86F,2004MNRAS.355..159W}. The same dynamical processes can be appealed to on smaller scales  to explain the existence of ICL around  less massive, looser clusters and galaxy groups \\citep{2004bdmh.confE..26A}, as well as diffuse stellar halos around individual galaxies \\citep{2005ApJ...635..931B}. Thus stellar halos and the ICL can be thought of  as testaments to galactic interactions during the epochs of galaxy and cluster formation.  We explore whether dynamical interactions could plausibly lead to stars {\\it beyond} the virial radii of the parent galaxy, group or cluster, and, if so, what the existence of such a population might tell us about the nature of hierarchical structure formation. \\citet{1992ApJ...397L..75S} have already shown that N-body simulations of mergers of massive galaxies can give rise to unbound particles that form a gaussian tail to the energy distribution, in agreement with a statistical mechanical description of violent relaxation. Here, we develop a simple description of interactions that allows us to illustrate how the size of the escaped population should depend on the mass ratio and orbits of the progenitor systems. In \\S \\ref{Mechanism for Escape} we establish that a tidal impulse at perigalacticon is a viable path for  stars from satellite dwarf galaxies to escape from the combined dark matter potentials of the parent and satellite galaxies. In \\S \\ref{Illustration} we first use restricted three-body simulations to test the mechanism described in \\S \\ref{Mechanism for Escape} and then analyze N-body results to establish trends, and quantitative expectations for the escaping population. In \\S \\ref{Discussion} we discuss the implications of our results for future surveys of giant stars, CN, supernovae and high velocity stars. In \\S \\ref{Conclusions} we summarize our conclusions. ",
        "conclusions": "We've outlined how tidal interactions can eject stars which are loosely bound to a satellite galaxy to beyond the virial radius of the satellite's parent.  We determined that more massive satellites contribute the most to this escaped (entirely unbound from the satellite/parent system) or wandering (bound, but traveling beyond 2 virial radii for more than $\\sim$Gyr) population of stars, as do satellites on highly radial orbits.  The majority of this population of stars originated during the epoch of galaxy formation, when large interacting satellites were more frequently infalling on radial orbits, approximately from a redshift of around 4 to 3.  If this is indeed the case, we expect an old, uniform (to first order) population of escaped and wandering stars to exist beyond a few virial radii of the Milky Way.  Subcategories (e.g. classical novae, supernovae) of this population should be detectable by LSST and Pan-STARRS. We estimate a lower limit of 0.05\\% of stars to be members of this population. An accurate number prediction warrants further investigation using full cosmological simulations which follow interacting subhalos as they form the parent halo.  \"Not all those that wander are lost\"-Tolkien"
    },
    "0911/0911.2780_arXiv.txt": {
        "abstract": " ",
        "introduction": "Primordial density perturbations are traditionally described by a Gaussian distribution, characterised by an almost scale-invariant power spectrum. However the detailed information about the primordial density perturbations over a range of cosmological scales offers the opportunity to test in detail the nature of the primordial perturbations, both their scale-dependence and Gaussianity \\cite{Komatsu:2008}.  The local model for non-Gaussianity has proved a remarkably popular description of possible deviations from a purely Gaussian distribution of primordial perturbations. In the simplest case the primordial Newtonian potential includes a contribution from both the local value of a linear Gaussian field and a quadratic term proportional to the square of the local value of the Gaussian field: \\be \\label{originalPhi} \\Phi(\\vec{x}) = \\varphi_G(\\vec{x}) + \\fNLlocal \\left( \\varphi_G^2(\\vec{x}) - \\langle\\varphi_G^2\\rangle \\right) \\,, \\ee where ${f}^{\\rm local}_{\\rm NL}$ is a dimensionless parameter characterising the deviations from Gaussianity~\\cite{Komatsu:2001rj}. $\\langle\\varphi_G^2\\rangle$ denotes the ensemble average, or equivalently the spatial average in a statistically homogeneous distribution. Note that following Komatsu and Spergel~\\cite{Komatsu:2001rj} we adopt a sign convention for the Newtonian potential $\\Phi$ which is the opposite of that used, e.g., by Mukhanov et al~\\cite{Mukhanov:1990me}.  In Fourier space we define the power spectrum and bispectrum as \\bea \\langle\\Phi_{\\vec{k_1}}\\Phi_{\\vec{k_2}} \\rangle = (2\\pi)^3 P_{\\Phi}(k_1) \\delta^{3}(\\vec{k_1}+\\vec{k_2}) \\,,\\\\ \\langle \\Phi_{\\vec{k_1}} \\Phi_{\\vec{k_2}} \\Phi_{\\vec{k_3}} \\rangle = (2\\pi)^3 B_{\\Phi}(k_1,k_2,k_3) \\delta^{3}(\\vec{k_1}+\\vec{k_2}+\\vec{k_3}) \\,, \\eea and the amplitude of the bispectrum relative to the power spectrum is conventionally given by the non-linearity parameter \\be \\label{deffnl} \\fnl(k_1,k_2,k_3) \\equiv \\frac{B_{\\Phi}(k_1,k_2,k_3)}{2 \\left[ P_{\\Phi}(k_1) P_{\\Phi}(k_2) + P_{\\Phi}(k_2) P_{\\Phi}(k_3) + P_{\\Phi}(k_3) P_{\\Phi}(k_1) \\right]} \\,. \\ee In the special case of the local model (\\ref{originalPhi}) we have $\\fnl=\\fNLlocal$ and it is clear that $\\fnl$ is, by construction, a constant parameter independent of spatial position or scale.  This local model turns out to be a very good description of non-Gaussianity in some simple physical models for the origin of structure. In particular it can describe the primordial density perturbation on super-Hubble scales predicted, up to second-order, using the $\\delta N$-formalism \\cite{starob85,SS,LR} if the local integrated expansion is a function of a Gaussian random field, $N(\\sigma_g)$, at some initial time, $\\sigma_g(\\vec{x})=\\sigma(t_i,\\vec{x})$. We then have \\be \\label{Phizeta} \\Phi = \\frac35 \\zeta = \\frac35 \\left[ N - \\langle N \\rangle \\right] \\,. \\ee A good example is provided by the simplest curvaton scenario~~\\cite{Enqvist:2001zp,Lyth:2001nq,Moroi:2001ct}, where quantum fluctuations of a weakly-coupled field during inflation are well described by a Gaussian random field and the primordial density perturbation is determined by the density of the curvaton field when it decays, which is proportional the the square of the initial local field~\\cite{Linde:1996gt,Lyth:2002my}.   In this paper we consider extensions of the simple local model (\\ref{originalPhi}). In particular we will characterise and quantify the scale-dependence of the parameter $\\fnl$ which arise in realistic inflationary models for the origin of structure. Scale-dependence arises due to two key features. Firstly we examine a multi-variate local model where the local expansion is a quadratic function of more than one Gaussian random field. In this case scale-dependent $\\fnl$ can arise if the Gaussian fields have differing scale-dependence, leading to a change with scale in the correlation of quadratic terms with the linear perturbation. Secondly, we show that scale-dependent $\\fnl$ can arise even when the expansion is a function of a single canonical scalar field. In this case, it is due to the development of intrinsic non-Gaussianity associated with the non-linear evolution of the initially Gaussian fluctuations after Hubble-exit. We will refer to this case as a quasi-local model.  In the case of the so-called equilateral $\\fnl$, which arises for example from DBI inflation (see e.g. \\cite{DBI}), the scale dependence has already been quite well studied both in theoretical models \\cite{Chen:2005fe,Khoury:2008wj,Byrnes:2009qy,Leblond:2008gg} and forecasts have been made for future observational prospects \\cite{LoVerde:2007ri,Sefusatti:2009xu}. For a discussion of the different possible shape dependences of the bispectrum see \\cite{Fergusson:2008ra}. However for local-type models there has been little previous consideration of a scale dependence. The scale dependence of local-type $\\fnl$ was first calculated in Byrnes et al \\cite{Byrnes:2008zy} for a specific model of hybrid inflation with two-fields. It was found that although the scale dependence is slow-roll suppressed, it depends on a particular combination of slow-roll parameters which is not negligible and can easily be larger than the spectral index of the power spectrum. The scale dependence of $\\fnl$ was considered in the case of an ekpyrotic universe in \\cite{Khoury:2008wj}. In the case of an exact solution of two-field inflation, which can give rise to a large non-Gaussianity, it was shown that $\\fnl$ is scale independent \\cite{Byrnes:2009qy}. The observational prospects for local-type models were considered in \\cite{Sefusatti:2009xu} which showed that the CMB data is sensitive to a scale dependence of $\\fnl$. However they used a very simple Ansatz for the scale-dependent $\\fnl$: here we calculate for the first time the full scale dependence of a scale-dependent quasi-local $\\fnl$.  We note that higher-order contributions to the primordial perturbations are expected beyond quadratic order. These only affect the bispectrum at subleading order in the scalar field perturbations, although in some cases they might provide the dominant contribution to observables \\cite{Boubekeur:2005fj}. In the language of \\cite{Byrnes:2007tm} they are loop corrections. The effective scale-dependence of $\\fnl$ due to higher-order terms is examined explicitly in \\cite{Kumar:2009ge} and was considered previously, for example, in \\cite{Suyama:2008nt}.  Notice that the results depend at leading order on the infra-red cut-off, and that no infra-red complete theory is yet known \\cite{Seery:2009hs}. It would be interesting to further develop  this interesting issue  in a case in which the infra-red  effects are well-understood. Note that the loop term is not well described as having constant scale dependence, unlike the cases we consider in this paper. We also note that \\cite{Kumar:2009ge} only consider the scale dependence coming from the logarithmic term which depends on the cut-off, while the other terms which multiply this will also have a scale dependence in general. This scale dependence can be calculated using the methods presented in this paper.    This paper  is organized as follow. In Section \\ref{mvarsec}, we focus on multi-variate local models, which are models that contain contributions to the primordial curvature perturbation from more than one Gaussian field, in which the scale dependence of $\\fnl$ is due to the different scale dependences of the fields that drive local expansion. As an example of this we study a mixed inflaton and curvaton scenario in sec.~\\ref{sec:inflatonandcurvaton}. In Section \\ref{sfsec}, we show that scale dependence of $\\fnl$ can arise also for single field models, due to the non-linear evolution of initially Gaussian fluctuations after Hubble exit. In Section \\ref{multisec} we extend the discussion to systems involving multiple fields, taking into account the  non-linear evolution of fluctuations after horizon crossing in this context. We conclude in sec.~\\ref{concsec}. ",
        "conclusions": "\\label{sec:conclusion}\\label{concsec}  We have studied generalisations of the local model of non-Gaussianity, thereby seeing in which cases the standard assumption that it can be described by a single, scale-independent parameter is valid. We have shown that this is only strictly valid in specific models, where the primordial curvature perturbation is generated by a single scalar field which acts as an isocurvature (``test'') field during inflation and has a quadratic potential. An example of this is the simplest curvaton scenario. However as soon as the scalar field has interactions these generate non-Gaussianity of the field perturbations after Hubble exit, and these give rise to a scale dependent non-Gaussianity which we call quasi-local, meaning that in the limit that the scale dependence of $\\fnl$ goes to zero one recovers the local model.  We have also shown that even in a multi-variate local model where more than one Gaussian field contributes to the primordial curvature perturbation the effective $\\fnl$ has a scale dependence unless all fields have the same scale dependence. An example of this is a mixed inflaton and curvaton scenario or a multi-curvaton model, where the curvatons have quadratic potentials with different masses.   We have developed a formalism, based on the $\\delta N$ approach, which allows us to obtain compact expressions for $\\nnl=\\,d \\ln{|\\fnl|}/d \\ln{k}$, in both the single and multi-field case. We have also discussed more generally the validity of defining a scale dependence of $\\fnl$ with respect to a single scale when the bispectrum is generally a function of three variables, except in the case of an equilateral triangle. We have shown that the scale dependence of $\\fnl$ is independent of the shape of the triangle it describes, provided that we consider scale variations which preserve its shape. It is then appropriate to perform the simpler calculation of $\\fnl$ for an equilateral triangle, and use this to calculate $\\nnl$ by taking its logarithmic scale dependence.   We applied our formalism to discuss several specific models. Our results suggest that the scale dependence is typically first order in slow roll. The precise value however depends quite sensitively on details of the model which, makes it in principle possible to use $\\nnl$ to discriminate observationally between different models.  We have shown that while the simplest realisation of the curvaton model has an exactly scale independent $\\fnl$, almost any extension to a more realistic scenario does give rise to a scale dependence. In the case of an interacting curvaton or a multiple-curvaton scenario the scale dependence is likely to be suppressed compared to the spectral index. This is because $\\nnl$ in these scenarios depends schematically on $\\eta_{{\\rm curvaton}}$ which is normally positive and, in light of the observational preference for a red spectrum, this is likely to be small. A detection of both $\\fnl$ and $\\nnl$ would therefore be a signal of non-trivial dynamics in the curvaton scenario, or that the inflaton perturbations have also made a significant contribution to the observed power spectrum.  In the case of a mixed scenario, in which the primordial curvature perturbation has contributions from both the inflaton and curvaton perturbations it is quite natural for there to be a consistency relation between the scale dependence of the power spectrum and the bispectrum, $\\nnl\\simeq-2(n-1)$, and this could be observable with Planck assuming a large enough fiducial value for $\\fnl$, close to the current observational bounds. In this case the bispectrum would be larger on small scales and this could make a detection of $\\fnl$ using large scale structure data more likely. In any scenario where one includes the Gaussian perturbations of the inflaton field as well as the perturbations of a second field which generates non-Gaussianities one will have a scale dependent $\\fnl$, unless both fields have exactly the same spectral index.  For the case where non-Gaussianity is generated during slow-roll hybrid inflation one generally has a non-negligble $\\nnl$ \\cite{Byrnes:2008zy}. In this case, since the magnitude of $\\fnl$ can only grow during inflation, larger scales, which exit the horizon earlier, will be more non-Gaussian and one necessarily has $\\nnl<0$.  It should be straightforward to extend the formalism we have presented here to the primordial trispectrum (the four-point function) \\cite{Seery:2006js, Byrnes:2006vq} and we expect that the two non-linearity parameters which model it, $\\tau_{NL}$ and $g_{NL}$, would also have a first order in slow roll scale dependence. We have chosen the notation $\\nnl$ in a way that is extendible in an obvious manner to study this. For example in the case that a single field generates the primordial curvature perturbation one has a non-Gaussianity consistency relation, $\\tau_{NL}=(6\\fnl/5)^2$. Then assuming our formalism can be extended in the obvious way it follows that $n_{\\tau_{NL}}=2\\nnl$."
    },
    "0911/0911.5396_arXiv.txt": {
        "abstract": "{ We consider cosmologies in which a dark-energy scalar field interacts with cold dark matter. The growth of perturbations is followed beyond the linear level by means of the time-renormalization-group method, which is extended to describe a multi-component matter sector. Even in the absence of the extra interaction, a scale-dependent bias is generated as a consequence of the different initial conditions for baryons and dark matter after decoupling. The effect is  enhanced significantly by the extra coupling and can be at the 2-3 percent level in the range of scales of baryonic acoustic oscillations. We compare our results with N-body simulations, finding very good agreement. \\\\ PACS numbers: 95.35.+d, 95.36.+x, 98.80.Cq }    \\newpage ",
        "introduction": "\\setcounter{equation}{0}  The  power spectrum of matter perturbations reflects the evolution of the Universe since the time of matter-radiation equality. For given initial conditions, determined by the primordial spectrum (usually assumed to be scale invariant), the growth of perturbations depends on the cosmological scenario. The calculation of the present matter power spectrum can constrain this scenario through the comparison of the deduced spectrum with the observed large-scale structure. A major technical difficulty in the realization of such a program is the failure of linear perturbation theory to describe present-day fluctuations with characteristic length scales below roughly 100 Mpc. At length scales below about 10 Mpc, the evolution is highly non-linear, so that only numerical N-body simulations can capture the dynamics of the formation of galaxies and clusters of galaxies. Fluctuations with length scales of around 100 Mpc fall within the mildly non-linear regime, for which analytical methods have been developed. These scales (corresponding to wavenumbers in the $0.03-0.25\\,h$/Mpc range) are of particular interest, because they correspond to the sound horizon at decoupling, which can be determined by reconstructing the  oscillatory behavior of the matter power spectrum due to baryonic acoustic oscillations (BAO).  The various analytical methods \\cite{CrSc1}--\\cite{matsubara} essentially amount to resummations of subsets of perturbative diagrams of arbitrarily high order, in a way analogous to the renormalization group (RG). In this work we shall follow the approach of  \\cite{Max2}, named time-RG or TRG. In the context of the RG the various observables depend on a characteristic energy scale, and evolve as this scale is varied. The TRG uses time as the flow parameter that describes the evolution of physical quantities, such as the spectrum of perturbations. The method is characterized by conceptual simplicity. It has been applied to ${\\rm \\Lambda CDM}$ and quintessence cosmologies \\cite{Max2}, as well as models with massive neutrinos \\cite{lesgourgues}.  The fundamental equations in the TRG approach are the ``equations of motion'', i.e. the continuity, Euler and Poisson equations. From these, equations can be derived for the time evolution of correlation functions for the density and velocity fields.  The various spectra are obtained through appropriate Fourier transforms of the correlation functions. The method results in a coupled infinite system of integro-differential equations for the time evolution of the spectrum, bispectrum etc. The crucial approximation, that makes a solution possible, is to neglect the effect of the higher-order correlation functions in the evolution equations of the lower-order ones. The calculations performed so far take into account the spectrum and bispectrum and set the higher-level spectra to zero.  The procedure of truncating the system of equations is commonly employed in the applications of the Wilsonian RG to field theory or statistical physics. (For a review, see \\cite{erg}.) The accuracy of the calculation can be determined either by enlarging the truncated system (by including the trispectrum, for example) and examining the stability of the results, or by comparing with alternative methods. The second approach is often followed, because enlarging the truncation can increase the complexity of the calculation considerably. In the case of the TRG, the agreement with results from N-body simulations for ${\\rm \\Lambda CDM}$ has been confirmed \\cite{Max2}. Also, a comparative analysis of several analytical methods, using N-body simulations as a reference, has been carried out in ref. \\cite{carlson}. The study demonstrates that the TRG remains accurate at the 1-2\\% level over the whole BAO range  at all redshifts.  In fig. 5 of ref. \\cite{carlson} the deviation of the TRG prediction from a reference spectrum derived through simulations  for ${\\rm \\Lambda CDM}$ is depicted. At a redshift $z=1$ the deviation is at the 1\\% level, or smaller, over the whole depicted range $0\\leq k \\leq 0.15\\,h$/Mpc. For $z=0$ the deviation may exceed 1\\% for some values of $k$, but stays below 2\\% over the whole depicted range. Based on these findings and the analysis of \\cite{Max2} we estimate an accuracy of 1\\% for $z=1$ and 2\\% for $z=0$ over the range $0\\leq k \\leq 0.2\\,h$/Mpc. The accuracy of 1-loop standard perturbation theory (SPT) can be inferred from fig. 1 of ref. \\cite{carlson}. At $z=1$ the accuracy is at the 1\\% level for $0\\leq k \\leq 0.1\\,h$/Mpc, while at $z=0$ it is at the 2\\% level for $0\\leq k \\leq 0.05\\,h$/Mpc. We depict these ranges in figs. \\ref{ddcc}--\\ref{bias1} of the current paper. The additional approximations that we make in this paper for the study of models with non-zero coupling between dark matter and dark energy induce uncertainties at the sub-percent level. We discuss this issue in detail in subsection 3.1. For this reason the level of accuracy of our results for the power spectra in the coupled case is expected to be similar to that for ${\\rm \\Lambda CDM}$. Its magnitude is set by the truncations in the evolution equations (the omission of the effect of the trispectrum and higher spectra), which are similar in all models.  The purpose of the present work is threefold: \\\\ 1) We extend the formalism to more complicated models. We introduce two modificiations to previous studies: a) Within the matter sector we allow for an arbitrary number of species with independent spectra. These include baryonic matter (BM) and cold (pressureless) dark matter (CDM). One may also consider contributions from massive neutrinos etc. b) We allow for an interaction between CDM and dark energy (DE). We consider a class of quintessence models, in which there is direct coupling between CDM and the quintessence field. The form of the interaction is a generalization of the universal coupling to all species present in scalar-tensor theories in the Einstein frame. It is modelled through the dependence of the mass of the CDM particles on the quintessence field \\cite{wetcosmon}-\\cite{baccigalupi}. \\\\ 2) We test the accuracy of the method in this enlarged framework by comparing with available N-body simulations. We perform our numerical analysis for a model for which results from simulations are given in ref. \\cite{baldi}. In the context of coupled quintessence, the cosmological evolution can be very diverse \\cite{amendola,baccigalupi,perturbations1,perturbations2}. It is very time-consuming to study exhaustively every model through N-body simulations. Our approach provides an alternative method, which can be much faster, while retaining the necessary accuracy. It is also important to note that, while the N-body simulations are highly accurate at the length scales of galaxies and clusters of galaxies, they are less accurate in the BAO range because of the required large volumes. On the other hand, analytical methods, like ours, are more accurate in the quasi-linear regime of large length scales. The two  approaches can be viewed as complementary. \\\\ 3) We provide predictions for observables for which non-linear corrections can be important. As such, we study the bias between dark and baryonic matter in the BAO range for models of coupled quintessence.   The couplings between the matter sector and DE are constrained by observations. For the BM-DE coupling, the bound from the Cassini spacecraft \\cite{Cassini} limits its order of magnitude to be below $10^{-3}$. As such small couplings produce negligible effects on the power spectrum, we assume that the BM-DE coupling is exactly zero. The coupling between CDM and DE is constrained by various considerations, such as the modification of the spectrum of the cosmic microwave background (CMB) or that of the matter distribution. A common feature of the class of models we are considering is that the presence of an additional long-range force between CDM particles, induced by the DE field, modifies their clustering properties. The various observable consequences have been discussed in the literature \\cite{largescale}--\\cite{baldi}. The strength of the CDM-DE interaction is constrained through the comparison with the observed CMB and matter spectra. It has been shown that, for particular models, the CDM-DE coupling must be considerably smaller than the gravitational one \\cite{bean,lavacca}. Such constraints cannot be considered generic, because the evolution of the cosmological background and the perturbations around it varies considerably from model to model. We work within the model of \\cite{baldi}, because our main objective is to compare with the results of N-body simulations presented there. The couplings that we consider are roughly consistent with the bounds deduced in \\cite{bean,lavacca} for a model similar to ours.  In this work we follow a novel approach for the derivation of the fundamental equations. We derive the continuity, Euler and Poisson equations on an expanding background, starting from the conservation of the energy-momentum tensor. In order to cast these equations in a form that generalizes the standard expressions on a static background, we need to impose a certain hierarchy between the density perturbations, the velocity field and the potentials. Our derivation makes it straightforward to generalize these equations in future studies in order to take into account non-zero pressure and higher-order terms. Next, we derive the system of differential equations for the spectrum and bispectrum, within a truncation that neglects higher-level spectra. We integrate these equations numerically in order to produce the non-linear spectra at low redshift and compare with  the results of N-body simulations. We study in detail the difference, usually characterized as bias, between the spectra of dark and baryonic matter. We find that the CDM-DE coupling enhances significantly the bias of the decoupled case.   In the following section we derive the evolution equations for the spectra of dark and baryonic matter. The details of the derivation for the case of one massive component are given in appendix A. The generalization for an arbitrary number of massive components is presented in appendix B. The results of the numerical integration of the evolution equations are presented in section 3. We compare them with the results of N-body simulation. We also discuss in detail the form of the bias in the BAO range. ",
        "conclusions": "\\setcounter{equation}{0}  In this paper we have extended the TRG formalism introduced in ref.~\\cite{Max2} in two respects. Firstly, we described the matter sector of the theory keeping track of the CDM and BM components separately, instead of treating them as a single fluid. Secondly, we introduced a new scalar force that couples differently to CDM and BM. As was discussed recently in \\cite{smith}, an accurate computation of the evolution of the different components is mandatory if one wants to achieve high precision modeling of structure formation. In order to test the accuracy of the TRG method for this more general class of cosmological scenaria, we analyzed the same cosmologies considered in \\cite{baldi}, and compared our results with those of the N-body simulations presented there. The agreement is very good up to $k \\simeq 0.5 \\,\\mathrm{h/Mpc}$ at $z=0$, where the non-linear power spectrum is roughly twice the linear one. These results confirm the reliability of the TRG as a computational tool that can fill the gap between linear and non-linear scales.   Even in the absence of an extra force, or in the case that the extra force couples universally to all matter species, the evolution of BM and CDM differs as a consequence of different initial conditions after decoupling. In linear theory, the initial unbalance between BM and CDM fluctuations is almost completely washed out by the present epoch. However, when non-linearities are taken into account the bias persists. We have seen that the effect is at the percent level in the BAO range at low redshift in the uncoupled case, but it may grow up to the $2.5\\%$ level when a non-zero coupling is turned on with a value compatible with present bounds (obtained within the linear approximation \\cite{bean,lavacca}). As a result, these models will receive significant constraints from future galaxy surveys, which aim to measure the power spectrum within the BAO range with an accuracy at the percent level. A thorough investigation of this model-dependent issue goes beyond the purpose of this paper and is postponed for future work.   The most interesting outcome of our analysis from the point of view of observations concerns the form of the bias between CDM and BM spectra. In linear theory, the bias is {\\it scale independent} in recent epochs \\cite{tocchini}. The non-linear corrections remove this feature, and make the bias {\\it scale dependent}, consistently with the results of \\cite{baldi}. We confirm and extend these results for the BAO region ($0.03\\,h$/Mpc $\\lta k \\lta 0.25\\, h$/Mpc), where we expect the TRG method to be reliable. The effect becomes more pronounced with increasing coupling between CDM and DE (assuming that the respective coupling for BM vanishes). The form of the bias provides an additional observational handle for the differentiation between various models. For example, in models with massive neutrinos the growth of the spectrum induced by the non-linearities at small length scales is compensated by the free-streaming of neutrinos \\cite{lesgourgues}. On the other hand, the bias is expected to remain scale-dependent at the non-linear level if there is a substantial CDM-DE coupling."
    },
    "0911/0911.3390_arXiv.txt": {
        "abstract": "We  have investigated the evolution of a pair of interacting planets embedded in a gaseous \\langeditorchanges{disc,} considering \\langeditorchanges{the} possibility of the resonant capture of a Super-Earth by a Jupiter mass gas giant. First, we have examined the situation where the Super-Earth is on the internal orbit and the gas giant on the external one. It has been found that the terrestrial planet is scattered from the disc or the gas giant captures the Super-Earth into an interior 3:2 or 4:3 \\langeditorchanges{mean-motion} resonance. The stability of \\langeditorchanges{the latter configurations} depends on the initial planet positions and \\langeditorchanges{on} eccentricity evolution. The behaviour of the system is different if the Super-Earth is the external planet. We have found that instead of being captured in the \\langeditorchanges{mean-motion} \\langeditorchanges{resonance,} the terrestrial planet is trapped at the outer edge of the gap opened by the gas giant. \\langeditorchanges{This effect} prevents \\langeditorchanges{the} occurrence of the first order \\langeditorchanges{mean-motion} commensurability. \\langeditorchanges{These} results are particularly interesting in light \\langeditorchanges{of} recent exoplanet discoveries and provide predictions of what will become observationally testable in the near future. ",
        "introduction": "So far we know only \\langeditorchanges{a} few Super-Earths, \\langeditorchanges{i.e.,} planets less massive than $10 M_\\oplus$. However, there is a good chance that in the near future, COROT and \\langeditorchanges{the} Kepler mission will find more terrestrial type planets. An interesting possibility is the existence \\langeditorchanges{of} Super-Earths close \\langeditorchanges{to} Jupiter-like planets. \\langeditorchanges{Different} mass objects embedded in a protoplanetary disc will migrate with different rates. The final configurations will depend on the intricate interplay among many physical processes including planet-planet, disc-planet and planet-star interactions. Here we focus on the possible resonant configurations of a gas giant and a Super-Earth orbiting a Solar-type star.  The early divergent conclusions concerned with the occurrence \\langeditorchanges{of} terrestrial planets \\langeditorchanges{in} hot Jupiter systems (Armitage, \\cite{armitage}, Raymond et al., \\cite{raymond05}) have been clarified by the most recent studies (Fogg \\& Nelson, \\cite{fonel07a}, \\cite{fonel07b}, Raymond et al., \\cite{raymond}), which predict that terrestrial planets can grow and be retained \\langeditorchanges{in} hot-Jupiter \\langeditorchanges{systems,} both interior and exterior to the gas giant. Here we consider the evolution of an already formed Super-Earth and a gas giant, both embedded in a gaseous disc. In our studies we are interested in the possibility of the formation of resonant planetary configurations. First, we have examined the behaviour of the system when the terrestrial planet is on the internal orbit and the gas giant on the external one. Then we have reversed the situation such that the Jupiter is the \\langeditorchanges{internal} planet and the Super-Earth is \\langeditorchanges{the external planet}. %  \\langeditorchanges{Two-dimensional} hydrodynamical simulations have been employed in order to determine the occurrence \\langeditorchanges{of first-order mean-motion} resonances as the outcome \\langeditorchanges{of} convergent orbital migration. Our predictions have \\langeditorchanges{interesting} astrobiological \\langeditorchanges{implications,} which we discuss in \\langeditorchanges{the Conclusions section}. ",
        "conclusions": "A large-scale orbital migration \\langeditorchanges{in} young planetary systems might play an important role in \\langeditorchanges{shaping} their architectures. \\langeditorchanges{Tidal} gravitational forces are able to rearrange the planet positions according to their masses and the disc parameters. The final configuration after \\langeditorchanges{disc} dispersal might be what we actually observe in the extrasolar systems. In particular, \\langeditorchanges{convergent} migration can bring planets into \\langeditorchanges{mean-motion} resonances \\langeditorchanges{as} has been found in Podlewska \\& Szuszkiewicz (\\cite{podszusz}) \\langeditorchanges{in the} case where the terrestrial planet is on the internal orbit.  When the Super-Earth is on the external \\langeditorchanges{orbit,} it is captured at the outer edge of the gap opened by the gas giant (Podlewska \\& Szuszkiewicz, \\cite{paperII}). Despite \\langeditorchanges{the} absence of \\langeditorchanges{mean-motion} resonance, the trapping at the outer edge of the gap can slow down the migration of the \\langeditorchanges{Super-Earth.} Thus, depending on whether the Super-Earth is inside or outside \\langeditorchanges{the} gas giant \\langeditorchanges{orbit,} the most likely planet configurations will be different. We claim that the Super-Earth can survive in close proximity to the gas giant. \\langeditorchanges{The} terrestrial planet is \\langeditorchanges{either} locked in \\langeditorchanges{mean-motion} commensurability or it is captured in the trap which prevents \\langeditorchanges{fast} migration toward the star.  Our results have an interesting implication \\langeditorchanges{for} astrobiological studies. Namely, if we have a gas giant in or near the habitable \\langeditorchanges{zone,} \\langeditorchanges{then a} terrestrial planet locked in the mean motion resonance or captured in the trap at the outer edge of the gap  can be also located in the habitable zone. Some of the known extrasolar gas giants as well as the recently discovered Super-Earth in the Gliese 581 system (Udry et al., \\cite{udry}) are in the habitable \\langeditorchanges{zone.} We can expect that future observations will reveal terrestrial planets close to gas giants."
    },
    "0911/0911.3445_arXiv.txt": {
        "abstract": "We present new high spatial resolution ($\\lesssim0$\\farcs1) 1--5\\mic\\ adaptive optics images, interferometric 1.3\\,mm continuum and $^{12}$CO 2--1 maps, and 350\\mic, 2.8 and 3.3\\,mm fluxes measurements of the {\\hv} system. Our adaptive optics images unambiguously demonstrate that {\\hvab}--C is a common proper motion pair. They further reveal an unusually slow orbital motion within the tight {\\hvab} pair that suggests a highly eccentric orbit and/or a large deprojected physical separation. Scattered light images of the {\\hvc} edge-on protoplanetary disk suggest that the anisotropy of the dust scattering phase function is almost independent of wavelength from 0.8 to 5\\mic, whereas the dust opacity decreases significantly over the same range. The images further reveal a marked lateral asymmetry in the disk that does not vary over a timescale of 2 years. We further detect a radial velocity gradient in the disk in our $^{12}$CO map that lies along the same position angle as the elongation of the continuum emission, which is consistent with Keplerian rotation around an 0.5--1$M_\\odot$ central star, suggesting that it could be the most massive component in the triple system. To obtain a global representation of the {\\hvc} disk, we search for a model that self-consistently reproduces observations of the disk from the visible regime up to millimeter wavelengths. We use a powerful radiative transfer model to compute synthetic disk observations and use a Bayesian inference method to extract constraints on the disk properties. Each individual image, as well as the spectral energy distribution, of {\\hvc} can be well reproduced by our models with fully mixed dust provided grain growth has already produced larger-than-interstellar dust grains. However, no single model can satisfactorily simultaneously account for all observations. We suggest that future attempts to model this source include more complex dust properties and possibly vertical stratification. While both grain growth and stratification have already been suggested in many disks, only a panchromatic analysis, such as presented here, can provide a complete picture of the structure of a disk, a necessary step towards quantitatively testing the predictions of numerical models of disk evolution. ",
        "introduction": "Circumstellar disks are an ubiquitous outcome of the stellar formation process and they are believed to be the birth place of planetary systems. The growth of dust particles towards planetesimal sizes along with their vertical settling due to gas drag are processes that are believed to be the first steps towards planet formation. Hydrodynamical models have shown that these processes can be efficient early in the disk evolution \\citep[e.g.,][]{weidenschilling97, dullemond04, barriere05}. To test these models, it is necessary to obtain an observation-based description of the structure and dust content of protoplanetary disks as a function of the age of the system and other relevant parameters (e.g., stellar mass).  The dust component of protoplanetary disks has long been studied via its thermal emission from near-infrared to millimeter wavelengths which is frequently associated with low-mass pre-main sequence T\\,Tauri stars \\citep{kenyon87, bertout88, strom89, beckwith90}. Both grain growth and dust settling can alter the overall shape of the spectral energy distribution (SED) of a T\\,Tauri star \\citep{dalessio01, dullemond04, dalessio06}. Indeed, several studies that analyzed (elements of) the SED of young stars have concluded that both grain growth and settling is occuring in protoplanetary disks \\citep[e.g.,][and references therein]{beckwith91, mannings94, furlan06, kessler06, rodmann06, natta07}. Unfortunately, such studies suffer from the absence of spatial information inherent to photometric measurements and the high optical thickness of disks in the near- to mid-infrared regime. As a result, comparing an object's SED to radiative transfer models leaves many ambiguities \\citep{chiang01}. For instance, the inferred total dust mass and the maximum size of the dust grains are inversely correlated because of the dependency of dust opacity on grain size. In addition, most of these studies, which focus on a single type of observations (e.g., millimeter fluxes, silicate emission feature), only probe a limited region of the disk and a small fraction of the entire grain size distribution.  To solve for the ambiguities inherent to SED studies, it is critical to obtain spatially-resolved observations. Such observations include thermal emission mapping with (sub)millimeter interferometers \\citep[e.g.,][]{keene90, simon92, lay94, dutrey96, andrews07} and scattered light imaging with optical and near-infrared high-resolution instruments \\citep[e.g.,][]{burrows96, roddier96, stapelfeldt98}. Disks are generally optically thin at long wavevengths, so the former type of observations can probe the entire disk structure. Furthermore, they are very sensitive to the presence of millimeter-sized particles. On the other hand, scattered light images, which only probe dust grains at the disk surface, are very sensitive to the size distribution of micronic grains especially when images at multiple wavelengths are analyzed simultaneously \\citep[][and references therein]{watson07ppv}. Both types of observations have already yielded important pieces of evidence supporting both grain growth and dust settling in disks \\citep[e.g.,][]{duchene03, duchene04, watson04}.  While grain growth and settling appear to occur in protoplanetary disks, detailed {\\it quantitative} tests of hydrodynamical models can only be achieved with a detailed view of the entire structure of a disk. This can only be obtained via a multi-technique, panchromatic approach. Unfortunately, observational and computational limitations have so far limited the number of objects for which such an analysis could be conducted to a handful. The most notable examples are the studies of the ``Butterfly Star'' \\citep{wolf03}, IM\\,Lup \\citep{pinte08} and IRAS\\,04158+2805 \\citep{glauser08}. In the former two cases, these studies have unambiguously shown that the dust population is stratified, possibly indicating that dust settling is already occurring. Increasing the number of disks studied in such detail is necessary to disentangle individual peculiarities from genuine trends associated with disk evolution.  {\\hv} is a triple system located in the Taurus star-forming region. It consists of a 550\\,AU wide pair whose optically brightest component is itself a tight (10\\,AU) visual binary \\citep{simon96}. Spectroscopic and photometric measurements revealed that this subsystem does not currently experience accretion nor does it show infrared excess. They further establish an age of about 2\\,Myr for the system \\citep{white01}. The third component of the system, {\\hvc}, is much fainter yet bluer than {\\hvab}. While \\citet{magazzu94} first thought that this source was an Herbig-Haro object, \\citet{woitas98} later proposed that {\\hvc} is a normal M0 T\\,Tauri star surrounded by an opaque edge-on disk similar to that found in HH\\,30 by \\citet{burrows96}. Subsequent high resolution imaging confirmed this hypothesis \\citep{monin00}. \\citet{stapelfeldt03} produced the first model of high resolution 0.8 and 2.2\\mic\\ scattered light images of {\\hvc}, finding that dust properties similar to those of interstellar dust grains can account for these images. Their 0.8\\mic\\ image also revealed the presence of a roughly spherical envelope, producing a symmetric halo that is more extended than the disk itself, that is likely the remnant of the core from which the system formed.  As a consequence of their particular viewing geometry, edge-on protoplanetary disks offer a unique opportunity to determine their geometry and dust content. As such, they may be the best candidates to study vertical stratification in protoplanetary disks. They also are comparatively easy targets for high-angular resolution instruments since the contrast requirement is strongly relaxed. On the other hand, they are challenging from the modeling point of view because of the difficulty of accurately solving for radiative transfer including anisotropic scattering in high optical depth regions. Previous modeling efforts of edge-on protoplanetary disks have therefore focused on interpreting one type of observation at a time \\citep{burrows96, stapelfeldt98, stapelfeldt03, cotera01, wood02, watson04}. While these studies proved highly valuable to constrain some of the disk properties, no self-consistent model was used to model all data at once, leaving unexplained contradictions.  Our objective in this work is to perform a global analysis of the {\\hvc} disk, combining scattered light images, millimeter interferometric data and the overall SED into a single fit. We present a series of new observations of the system in Section~\\ref{sec:obs} and discuss our empirical results in Section~\\ref{sec:results}. In Section~\\ref{sec:models}, we present radiative transfer models of the {\\hvc} disk and discuss their implications in Section~\\ref{sec:discus}. Section~\\ref{sec:concl} summarizes our results. ",
        "conclusions": "\\label{sec:concl}  We have obtained new 1--5\\mic\\ adaptive optics images of the {\\hv} triple system at the VLT and Keck observatories, including the first 4.8\\mic\\ scattered light image of the edge-on {\\hvc} disk.  All of our images reveal a steady lateral asymmetry in the disk that indicates a departure from pure axisymmetry in the disk. This could be related to the known variability of {\\hvc}. We extact precise relative astrometry from our new images and find that {\\hvab}--C constitutes a common proper motion pair. We also resolved the tight binary system {\\hvab} and found a surprisingly slow orbital motion compared to its projected separation, probably due to orbit eccentricity and/or a large deprojected physical separation. We have also obtained the first spatially resolved 1.3\\,mm continuum and $^{12}$CO 2--1 millimeter maps of the {\\hvc} disk with IRAM's PdBI interferometer. The continuum emission is resolved along the same position angle as the scattered light images, as expected from dust thermal emission. The CO map shows evidence of Keplerian rotation about an 0.5--1$M_\\odot$ central star. We also completed the SED of that component with new {\\it Spitzer}/MIPS 24, 70 and 160\\mic\\ observations as well as observations at 350\\mic, 2.8 and 3.3\\,mm at the CSO, PdBI and CARMA facilities. The 1--3\\,mm spectral index of {\\hvc} is relatively flat ($\\alpha_{\\rm mm}=2.5\\pm0.2$), indicative of the presence of millimeter-sized grains in the disk and/or high optical depth even at millimeter wavelengths.  To interpret these data along with previously published HST images of {\\hvc}, we have computed a grid radiative transfer models that produced synthetic disk observations. While we can reproduce most observational properties of {\\hvc}, a Bayesian analysis of our model grid reveals that the SED and scattered light images of {\\hvc} provide mutually exclusive constraints. As found previously, the scattered light images are best fit with a relatively small total mass and interstellar-like dust properties. Fitting the SED, however, requires almost two orders of magnitude more mass as well as a grain size distribution that extends to millimeter sizes. The mismatch between the two sets of constraints reveals an intrinsic shortcoming of our parametrized model, which assumes perfectly mixed dust. Indeed, the small maximum dust grain size inferred from fitting the scattered light images is impossible to reconcile with the long-wavelength end of the SED. In turn, this suggests that both grain growth and vertical stratification are present in the {\\hvc} disk. While these phenomena have been suggested for other protoplanetary disks in the past, {\\hvc} is one of a handful of objects that have been studied in sufficient details to eventually provide a complete picture of these processes.  Future observations of {\\hvc} can help refine the model proposed here. For instance, high-resolution (0\\farcs2 or better) millimeter mapping of the disk would better resolve it and provide more stringent constraints. In particular, obtaining observations at even longer wavelengths, 7\\,mm or 1.3\\,cm, would probe a regime in which the optical depth through the disk is much smaller and enable a more direct interpretation in terms of disk properties. On the longer run, mapping of the disk with ALMA in the submillimeter regime and in scattered light at 8--10\\mic\\ with the next generation of large ground-based telescopes will help bridge the gap between the scattering and thermal emission regimes at a roughly constant spatial resolution. In particular, very high resolution maps with ALMA may resolve this and other disks along the vertical axis, in order to probe their vertical stratification and, ultimately, to outline an evolutionary sequence."
    },
    "0911/0911.1992_arXiv.txt": {
        "abstract": "We take advantage of good atmospheric transparency and the availability of high quality instrumentation in the $1\\, \\micron$ near-infrared atmospheric window to present a grid of F, G, K, and M spectral standards observed at high spectral resolution ($R\\approx$25,000).  In addition to a spectral atlas, we present a catalog of atomic line absorption features in the $0.95-1.11\\,\\micron$ range. The catalog includes a wide range of line excitation potentials, from 0-13 eV, arising from neutral and singly ionized species, most frequently those of Fe {\\scshape i} and Ti {\\scshape i} at low excitation, Cr\\,{\\scshape i}, Fe {\\scshape i}, and Si {\\scshape i} at moderate excitation, and C {\\scshape i}, S {\\scshape i}, and Si {\\scshape i} having relatively high excitation.  The spectra also include several prominent molecular bands from CN and FeH. For the atomic species, we analyze trends in the excitation potential, line depth, and equivalent width across the grid of spectroscopic standards to identify temperature and surface gravity diagnostics near $1\\, \\micron$. We identify the line ratios that appear especially useful for spectral typing as those involving Ti\\,{\\scshape i} and C {\\scshape i} or S {\\scshape i}, which are temperature sensitive in opposite directions, and Sr\\,{\\scshape ii}, which is gravity sensitive at all spectral types. ASCII versions of all spectra are available in the online version of the journal. ",
        "introduction": "\\label{sec:intro}  Since the advent of infrared arrays in the early 1990s, there have been many low resolution spectroscopic atlases published covering the $J$, $H$, or $K$ bands.  Previous efforts that included the full O--M spectral classes and I--V luminosity classes made use of moderate resolution Fourier transform spectroscopy techniques \\citep[e.\\/g.\\/,][]{Wall00,Mey98,Wall97}. An additional window of atmospheric transparency between these traditional near-infrared bands and the optical window, dubbed the $Y$ band, was highlighted by \\citet{Hill02}. With existing observational capabilities on instruments at the IRTF, Keck, Las Campanas, and UKIRT telescopes, and with upcoming wide-field survey facilities such as Pan-STARRS and VISTA, the $Y$ band is becoming an increasingly appreciated and scientifically important near-infrared bandpass.  Several low resolution ($R\\equiv\\lambda/\\Delta\\lambda\\sim2000$) spectral atlases for M, L, and T dwarfs have been published including $Y$ band data, e.g. \\citet{Leg96}, \\citet{McLe03}, and \\citet{Cush05}.  However, a high-resolution spectroscopic atlas covering earlier spectral types and the full range of luminosity classes has not yet appeared.  To complement existing moderate resolution catalogs in the $J$, $H$, and $K$ bands, we present here a high-resolution infrared spectral atlas for the $Y$ band.  Our data cover the $0.94-1.12\\,\\mu$m region sampled by NIRSPEC \\citep{McLe98} at the W. M. Keck Observatory. We obtained spectra of 20 MK-classified stars ranging in spectral type from F through M and in luminosity class from I--V for use as spectroscopic standards in future work.  The data were processed as briefly described in Section \\ref{sec:data}.  We present a representative grid of spectroscopic standards, a table of identified absorption features, and the resulting atlas of spectral lines in Section \\ref{sec:atlas}.  In Section \\ref{sec:analysis}, we present relevant atomic data and illustrate as well as discuss line sensitivity to stellar effective temperature and surface gravity. ",
        "conclusions": "\\label{sec:summary}  We have presented a spectral atlas for the $Y$ band region of the near infrared portion of the electromagnetic spectrum.  Our data set consists of a grid of MK-classified stars for use as spectroscopic standards spanning the F through M spectral classes and I--V luminosity classes observed at R$\\approx$25,000. The associated catalog of identified lines contains numerous spectral absorption features, both atomic and molecular in origin, that span a wide range of excitation potentials. Many of the atomic features show strong line depth and equivalent width trends across the spectral atlas. This atlas and the identified atomic spectral lines are potentially useful for two-dimensional spectral classification, especially line depth ratios involving Ti {\\scshape i} and C {\\scshape i} or S {\\scshape i}, which are temperature sensitive in opposite directions, and Sr {\\scshape ii}, which is gravity sensitive. ASCII versions of all spectra are available to download with the online version of the journal (an example of the form and content of these spectra can be found in Table \\ref{tab:asciiex})."
    },
    "0911/0911.4862_arXiv.txt": {
        "abstract": "{We observed 51~Peg, the first detected planet-bearing star, in a 55 ks {\\em XMM-Newton} pointing and in 5 ks pointings each with {\\em Chandra} HRC-I and ACIS-S. The star has a very low count rate in the {\\em XMM} observation, but is clearly visible in the {\\em Chandra} images due to the detectors' different sensitivity at low X-ray energies. This allows a temperature estimate for 51~Peg's corona of T$\\lesssim1\\mbox{MK}$; the detected ACIS-S photons can be plausibly explained by emission lines of a very cool plasma near 200~eV. The constantly low X-ray surface flux and the flat-activity profile seen in optical \\ion{Ca}{ii} data suggest that 51~Peg is a Maunder minimum star; an activity enhancement due to a Hot Jupiter, as proposed by recent studies, seems to be absent. The star's X-ray fluxes in different instruments are consistent with the exception of the HRC Imager, which might have a larger effective area below 200~eV than given in the calibration.} ",
        "introduction": "The star 51~Peg (GJ~882, HD~217014) shot to fame in 1995 when \\cite{mayorqueloz1995} detected an exoplanet in its orbit, the planetary parameters being quite unexpected at that time, because 51~Peg~b is a giant planet, located at only 0.05~AU distance. The star itself is a G5V star 15.4~pc away from the Sun. Its properties are quite similar to the Sun's, since 51~Peg is about 4~Gyr old and its mass, radius and effective temperature are comparable to solar values with $R=1.27R_{\\sun}$ \\citep{bainesmcalister2008}, $M=1.11M_{\\sun}$, $T_{eff}\\approx 5790K$ \\citep{fuhrmannpfeiffer1997}. However, 51~Peg is a metal-rich star, for which the metallicities given in the literature vary over a wide range of $+0.05\\leq [Fe/H]\\leq +0.24$, see for example \\cite{valentifischer2005}. Enhanced metallicities are a common feature of stars with giant planets \\citep{gonzales1997, santosisraelianmayor2001}.  The activity profile of 51~Peg turned out to be unspectacular. In the Mount Wilson program \\citep{baliunasdonahuesoon1995}, which monitors the \\ion{Ca}{ii} H and K line fluxes of main sequence stars, the star shows a very low and nearly flat chromospheric activity level from 1977 until 1989 and a slight drop in 1990 and 1991. In the Lowell Observatory program \\citep{halllockwoodskiff2007}, it also shows low activity and little variability in \\ion{Ca}{} fluxes since the beginning of the program in 1994. The star was also observed in a 12.5~ks {\\em ROSAT} PSPC pointing in 1992 and detected as a weak X-ray source.  The coronal activity of 51~Peg is of interest not only because the star is similar to the Sun, but also with regard to recent studies \\citep{kashyapdrakesaar2008}, which claim stars with close-in giant planets to be more X-ray active than stars with far-out ones. ",
        "conclusions": "We have detected X-ray emission from 51~Peg in a 55~ks observation with {\\em XMM-Newton} and 5~ks observations with {\\em Chandra} ACIS-S and HRC-I each. The detection of 51~Peg with a low count rate in the {\\em XMM} pointing and the clear source signal in the much shorter {\\em Chandra} observations can be explained by the different effective response of the detectors at low energies and 51~Peg having an extremely cool corona. Our main results are summarized as follows: \\begin{enumerate} \\item 51~Peg shows weak emission in the \\ion{O}{vii} triplet and emission around 200~eV which can be explained most likely by cool silicon emission lines. \\item A coronal temperature of $\\lesssim 1$~MK is consistent with the detected hardness ratios in different energy bands in both the {\\em XMM} and the {\\em Chandra} pointing as well as in the {\\em ROSAT} observation carried out 16~years earlier. \\item The {\\em Chandra} HRC-I count rate is higher than can be explained by differences in the effective areas of HRC and ACIS-S; HRC's effective area might be larger at low energies than given in the calibration so far. \\item The constant and very low surface X-ray flux level together with a flat-activity behavior in chromospheric \\ion{Ca}{ii}~H and K line fluxes suggests 51~Peg to be a Hot Jupiter-bearing Maunder minimum candidate. \\end{enumerate}"
    },
    "0911/0911.1489_arXiv.txt": {
        "abstract": "We present X-ray  analysis of two Wolf-Rayet (WR) binaries: V444 Cyg and  CD Cru using the data from  observations with  XMM-Newton. The X-ray light curves show the phase locked variability in both binaries, where the flux increased by a factor of $\\sim 2$ in the case of V444 Cyg and $\\sim 1.5$ in the case of CD Cru from  minimum to  maximum. The maximum luminosities in the 0.3--7.5 keV energy band were found to be $5.8\\times10^{32}$ and $2.8\\times10^{32}$ erg s$^{-1}$ for V444 Cyg and CD Cru, respectively. X-ray spectra of these stars confirmed large extinction and revealed hot plasma with prominent emission line features of highly ionized Ne, Mg, Si, S, Ar, Ca and Fe, and are found to be consistent with a two-temperature plasma model.  The cooler plasma at a temperature of $\\sim$ 0.6 keV was found to be constant  at all phases of both binaries, and could be due to a distribution of small-scale shocks in radiation-driven outflows.  The hot components in these binaries were found to be phase dependent. They varied from 1.85 to 9.61 keV for V444 Cyg and from 1.63 to 4.27 keV for CD Cru. The absorption of the hard component varied with orbital phase and found to be maximum  during primary eclipse of V444 Cyg.  The high plasma temperature and variability with orbital phase suggest that the hard-component emission is caused by a colliding wind shock between the binary components. ",
        "introduction": "\\label{sec:reduc} The log of the X-ray observations is given in Table~\\ref{tab:xray_observations}. We made use of the archival data obtained with the XMM-Newton observatory which consists of three co-aligned X-ray telescopes observed simultaneously and covered 30$^\\prime$ $\\times$ 30$^\\prime$ region of the sky. The X-ray photons were recorded with the European Photon Imaging Camera (EPIC), which forms images on three CCD-based detectors: the PN (Str$\\ddot{u}$der et al. 2001), and the twin MOS1 and MOS2 (Turner et al. 2001) with an angular resolution of 6$^{\\prime\\prime}$ (FWHM). During the observations, all the three EPIC detectors were active in full frame mode together with the Thick filter. For V444 Cyg, six separate observations were taken with exposure time ranging from 10 to 20 ks. These observations are spread over the half cycle of binary system. However, for the star CD Cru, three separate observations were taken covering almost full binary cycle.   \\subsection{EPIC data reduction} \\label{sec:src_reduc}  We reduced the X-ray data using standard XMM-Newton Science Analysis System software (SAS version 7.0.0) with updated calibration files (Ehle et al. 2004). Event files for the MOS and the PN detectors were generated  using the tasks {\\sc emchain} and {\\sc epchain}, respectively. These tasks allow calibration of the energy and the astrometry of the events registered in each CCD chip and to combine them in a single data file. We restricted our analysis  to the energy band 0.3--7.5 keV, as the data below 0.3 keV are mostly unrelated to bonafide X-rays, and above 7.5 keV is mostly dominated by the background counts. Event list files were extracted using the SAS task {\\sc evselect}. Data from the three cameras were individually screened for the time intervals with high background when the total count rates (for single events of energy above 10 keV) in the instruments exceed 0.35 and 1.0 $\\rm{counts~s^{-1}}$ for the MOS and PN detectors, respectively. The observations with observation ID 0206240401 for the source V444 Cyg were heavily affected by the high background events, and the data in MOS detectors were not useful. Only half of the observation time ($\\sim$5 ks) of PN was found useful.  Light curves and spectra were extracted using a circular region with the source as the center in the energy range 0.3--7.5 keV of the EPIC detectors. The X-ray sources in the cluster were often found to be largely contaminated due to the emission from the neighboring sources. For this reason, the radii of extraction regions were varied between 20$^{\\prime\\prime}$ and 30$^{\\prime\\prime}$ depending upon the position of the sources in the detector and their angular separation between the neighboring X-ray sources. The background has been estimated from a number of empty regions close to the X-ray source in the same CCD of the detector. X-ray spectra of the sources were generated using SAS task {\\sc especget}, which also computes the photon redistribution as well as the ancillary matrix. For each source, the background spectrum was obtained from regions devoid of any sources chosen according to the source location. Finally, the spectra were re-binned to have at least 20 counts per spectral bin for both the sources in all the observations.   \\subsection{RGS data reduction} \\label{sec:rgs} The Reflection Grating Spectrometers (RGS) are mounted on two of the three XMM-Newton X-ray telescopes and were operated in spectroscopy mode during the observations. We followed the standard procedure as outlined in the XMM-Newton handbook to generate the RGS spectra of the sources. The raw data were processed with the task {\\sc rgsproc} at the position of the sources. The other bright sources, found in the same field observed by RGS, have been excluded from background estimation. In addition, we have filtered the event list for high background level events. For star V444 Cyg, the data obtained in observations with observation ID 0206240401 were contaminated with high background events in RGS detector, therefore the data was not useful. For CD Cru, the two observations ID 0109480101 and 0109480201 were not useful as they are heavily contaminated by a brighter source HD 110432 and affected by high background episodes, respectively.  \\section {Analysis and Results}\\label{sec:results}   \\subsection{X-ray light curves}\\label{sec:WR_lt} \\begin{figure*} \\centering \\subfigure[]{\\includegraphics[width=3.0in]{V444Cyg_PN_JD_total.eps}} \\subfigure[]{\\includegraphics[width=3.0in]{CDCru_JD_PN_total.eps}} \\caption{X-ray light curves of the WR-binaries (a) V444 Cyg and (b) CD Cru in the 0.3--7.5 keV energy band.} \\label{fig:xlc} \\end{figure*}   \\begin{figure*} \\centering \\subfigure[]{\\includegraphics[width=3.0in]{V444Cyg_PN_final.eps}} \\subfigure[]{\\includegraphics[width=3.0in]{V444Cyg_MOS_final.eps}} \\caption{(a) PN and (b) MOS light curves at three energy bands: the total (0.3--7.5 keV), the soft (0.3--2.0 keV) and the hard (2.0--7.5 keV), and hardness ratio  HR curve as a function of orbital phase, where HR = hard/soft of V444 Cyg.} \\label{fig:V444Cygphaselc} \\end{figure*}  \\begin{figure*} \\centering \\subfigure[]{\\includegraphics[width=3.0in]{CDCru_PN_final.eps}} \\subfigure[]{\\includegraphics[width=3.0in]{CDCru_MOS_final.eps}} \\caption{Similar to Fig. \\ref{fig:V444Cygphaselc} but for CD Cru.} \\label{fig:CDCruphaselc} \\end{figure*}  The background subtracted X-ray light curves of the WR-binaries V444 Cyg and CD Cru, as observed with PN detector are shown in Fig. \\ref{fig:xlc}(a) and \\ref{fig:xlc}(b), respectively. The light curves are in the  0.3--7.5 keV (total) energy band show the variability which suggests colliding wind shocks. Further, we performed the $\\rm{\\chi^2}$ test  to measure the significance of the deviations from the mean count rate in order to quantify the constancy of the data over the time-scale of observations. We found the variability in the light curves with a confidence level of greater than 99.999\\% for both binaries. In order to investigate the variability in the different energy bands, the light curve of these binaries obtained with the MOS and PN data are  generated into  two energy bands namely the soft (0.3--2.0 keV) and the hard (2.0--7.5 keV). The hardness ratio (HR) is defined by the ratio of hard to soft band count rates. The total, hard and soft band intensity curves, and the HR curve  as a function of the orbital phase are shown in the subpanels running from top to bottom in  Fig. \\ref{fig:V444Cygphaselc}a  (PN) and \\ref{fig:V444Cygphaselc}b (MOS) for the WR binary V444 Cyg. Similar plots of intensity and HR for the star CD Cru are shown in Fig. \\ref{fig:CDCruphaselc}. The phases of the observations are reckoned using the ephemeris HJD = 2441164.337+4.213E for V444 Cyg (Underhill, Grieve \\& Louth  1990) and  HJD = 2443918.4000+6.2399E  for  CD Cru (Moffat et al. 1990, Niemela, Massey \\& Conti 1980). Here, the phase 0.0 indicates the primary eclipse and the phase 0.5 indicates the secondary eclipse. The light curves in the individual bands show the phase locked variability for both binaries.  For V444 Cyg, the count rates in total energy band were minimum at phase 0.0 and maximum at the phase 0.45. After the phase 0.45, the count rates were decreased upto the phase 0.5. The count rates are increased by a factor of $\\sim 2$ being minimum at phase 0.0\\ in  the total energy band.  In the soft band, the count rates were minimum at phase 0.0 and maximum during the phase 0.45-0.5.  However, in the hard band light curve two minima at phase 0.0 and 0.5 were seen clearly. During the phases from 0.13 to 0.45 (i.e. outside the eclipse)  the count rates in the hard band were constant. The HR curve can reveal the information about the spectral variations.  The HR curve  shows that the emitted X-ray  is harder just after phase 0.0 being maximum at phase 0.13, and afterward decreased till the phase 0.5. In the case of CD Cru, the nature of the variability was found to be similar to that seen in V444 Cyg.  In soft band, the count rates were found to be constant during the phases 0.47 and 0.78 (i.e. outside the eclipse). However,  the count rates are decreased rapidly after the phase 0.47 \\ in the hard band.  A small variation was also observed in the HR curve of the CD Cru. It was maximum at the phase 0.0 and minimum during  the phase 0.47.        \\subsection{X-ray spectra and spectral fits }\\label{sec:WR_spt} \\begin{figure*} \\centerline{\\hbox{\\hspace{0.5in} \\includegraphics[width=60mm, angle=-90]{V444Cyg_obs301_2Tvapec_new_formate.eps} \\hspace{0.15in} \\includegraphics[width=60mm, angle=-90]{V444Cyg_obs201_2Tvapec_new_formate.eps} }} \\vspace{0.15in} \\centerline{\\hbox{\\hspace{0.5in} \\includegraphics[width=60mm, angle=-90]{V444Cyg_obs701_2Tvapec_new_formate.eps} \\hspace{0.15in} \\includegraphics[width=60mm, angle=-90]{V444Cyg_obs501_2Tvapec_new_formate.eps} }} \\vspace{0.15in} \\centerline{\\hbox{\\hspace{0.5in} \\includegraphics[width=60mm, angle=-90]{V444Cyg_obs801_2Tvapec_new_formate.eps} \\hspace{0.15in} \\includegraphics[width=60mm, angle=-90]{V444Cyg_obs401_2Tvapec_new_formate.eps} }} \\vspace{3pt} \\caption{X-ray  spectra of MOS and PN data with the best fit 2T {\\sc vapec} model in upper subpanels of each graph for V444 Cyg. The $\\chi^2$ distribution in terms of ratio are given in lower subpanels of each graph. The last three digits of observation ID and the corresponding phases of the observations are given in each graph. } \\label{fig:spt_V444Cyg} \\end{figure*}   \\begin{figure*} \\centerline{\\hbox{\\hspace{0.5in} \\includegraphics[width=60mm, angle=-90]{CDCRu_obs401_2Tvapec_new_formate.eps} \\hspace{0.15in} \\includegraphics[width=60mm, angle=-90]{CDCru_obs201_2Tvapec_new_formate.eps} }} \\vspace{0.15in} \\centerline{\\hbox{\\hspace{0.5in} \\includegraphics[width=60mm, angle=-90]{CD_Cru_obs101_2Tvapec_new_formate.eps} }} \\vspace{3pt} \\caption{Same as Figure~\\ref{fig:spt_V444Cyg}  but for CD Cru. } \\label{fig:spt_CDCru} \\end{figure*}    The EPIC spectra of WR binaries in different phases are shown in Figs.~\\ref{fig:spt_V444Cyg} and \\ref{fig:spt_CDCru} for V444 Cyg and CD Cru, respectively. Below 1 keV, the spectra were found to be affected by the high extinction as previously observed with ASCA for V444 Cyg (Maeda et al. 1999). Strong emission lines are seen in the MOS and PN spectra of both WR binaries. The most prominent lines found in the spectra along with their laboratory energies are the following: Fe XVII (0.8 keV), Ne X (1.02 keV), Mg XII (1.47 keV), Si XIII (1.853 keV), S XV (2.45 keV), Ar XVII (3.12 keV), Ca XIX+XX (3.9 keV) and Fe XXV (6.63 keV).  In order to trace the spectral parameters at different binary phases, we performed spectral analysis of each data set using simultaneous fitting of EPIC data by two models, (a) plane-parallel shock model ({\\sc vpshock}; Borkowski, Lyerly \\& Reynolds 2001), and (b) models of Astrophysical Plasma Emission Code ({\\sc apec}; Smith et al. 2001), as implemented in the XSPEC version 12.3.0. A $\\chi^2$ -- minimization  gave the best fitted model to the data. The presence of interstellar material along the line-of-sight and the local circumstellar material around the stars can modify the X-ray emission from massive stars. We corrected for the local  absorption in the line-of-sight to the source using the photoelectric absorption cross sections according to Baluci$\\acute{n}$ska-Church \\& McCammon (1992) and modeled as {\\sc phabs} (photoelectric absorption screens) with two absorption components, $\\rm{N_{H}^{ISM}}$ and $\\rm{N_{H}^{local}}$. The $\\rm{N_{H}^{ISM}}$ was estimated using the relation, $\\rm {N_H}$ $\\rm{= 5.0\\times}10^{21}\\times{E(B-V)~cm^{-2}}$ (Vuong et al. 2003), where $\\rm{E(B-V)=A_V/3.1}$, assuming a normal interstellar reddening law towards the direction of the cluster. We used the values of $\\rm{A_V}$ nearly 2.48 mag and 3.56 mag derived for the cluster Berkeley 86 (Massey, Johnson \\& Degioia-Eastwood 1995) and Hogg 15 (Sagar, Munari \\& Boer 2001), respectively. The estimated values of $\\rm{N_{H}^{ISM}}$ towards V444 Cyg and CD Cru are found to be $\\rm{4.0\\times10^{21}~cm^{-2}}$ and $\\rm{6.0\\times10^{21}~cm^{-2}}$, respectively. The $\\rm{N_{H}^{local}}$ was estimated by making a fit to the observed spectra by varying the local environment for the soft ($\\rm{kT_1}$) and the hard ($\\rm{kT_2}$) energy components in terms of $\\rm{N_H^1}$ and $\\rm{N_H^2}$, respectively. Because the WN stars are at  evolved stages, the abundances of He, C, N, O, Ne, Mg, Si, S, Ar, Ca, and Fe were allowed to vary during the fitting procedure to account for the observed line emission. The solar abundances were adopted from Lodders (2003).  First, we fitted {\\sc vpshock} plasma model to derive their spectral features. The constant temperature {\\sc vpshock} plasma model was considered without incorporating mass-loss and orbital parameters of WR stars. However, the model does account for non-equilibrium ionization effects and assumes an equal electron and ion temperature. The best fit {\\sc vpshock} models to the data for the WR stars  are given in Table~\\ref{tab:V444Cyg_fit} and Table~\\ref{tab:CDCru_fit} for V444 Cyg and CD Cru, respectively.   X-ray emitting plasma may not be isothermal and the observed X-ray spectrum may be a superposition of a cool stellar component and a hot colliding wind shock plasma, as a number of emission lines in the spectra can be formed over a range of temperatures. The cooler component is believed to arise from the instabilities in radiation-driven outflows. However, using the X-ray imaging data of XMM-Newton it is not possible to resolve the X-ray emission from colliding wind region and individual stars separately. Therefore, secondly, we fitted two temperature (2T) plasma model \"{\\sc vapec}\" to characterize such components. The form of 2T plasma model was {\\sc phabs(phabs*vapec+phabs*vapec)}. In terms of $\\chi^2$ the 2T plasma model  provides  better goodness-of-fit for both the sources over the {\\sc vpshock} model. The results for the 2T plasma model  along with the data are displayed in Fig.~\\ref{fig:spt_V444Cyg} and Fig.~\\ref{fig:spt_CDCru}  for V444 Cyg and CD Cru, respectively. The best-fit parameters are given in Table~\\ref{tab:V444Cyg_fit} and  Table~\\ref{tab:CDCru_fit} for V444 Cyg and CD Cru, respectively.  \\input{bestfit_WR_V444Cyg.tab} \\input{bestfit_WR_CDCru.tab}  \\clearpage  \\begin{figure} \\centering \\subfigure[]{\\includegraphics[width=3.0in]{analysis_V444Cyg.eps}} \\subfigure[]{\\includegraphics[width=3.0in]{analysis_CDCru.eps}} \\caption{Variation of $\\rm{L_X^S}$, $\\rm{L_X^H}$, $\\rm{N_H^1}$, $\\rm{N_H^2}$, $\\rm{kT_1}$ and $\\rm{kT_2}$ as a function of orbital phase of the the stars (a) V444 Cyg and (b) CD Cru.} \\label{fig:ana_phase} \\end{figure} \\subsection {Evolution of spectral parameters} \\label{sec:spt_evo}  The spectral analysis at different phases of the WR binaries provides the dependence of the best fit values of parameters on orbital phase. The variation of the soft ($\\rm{L_X^S}$) and hard ($\\rm{L_X^H}$) band X-ray luminosities, column densities corresponding to the cool ($\\rm{N_H^1}$) and hot ($\\rm{N_H^2}$) temperature components, and cool(kT$_1$) and hot (kT$_2$) temperatures as a function of orbital phases of V444 Cyg and CD Cru are shown in Fig. \\ref{fig:ana_phase}(a) and \\ref{fig:ana_phase}(b), respectively. In the case of V444 Cyg, the $\\rm{L_X^H}$ was found to be minimum during the  primary  and secondary eclipses and maximum outside the eclipse. It was constant during the phase 0.13 to 0.47. However, the $\\rm{L_X^S}$ was found to be minimum during the primary eclipse only. At the phase of 0.5 the  $\\rm{L_X^S}$ was found to be maximum. Both  $\\rm{L_X^H}$ and $\\rm{L_X^S}$ were found to be minimum during the phase 0.0 and maximum at phase 0.5 in the WR binary CD Cru. The cool component ($\\rm{kT_1}$) was found to be constant with a mean value of 0.61$\\pm$0.05 keV and 0.57$\\pm$0.10 keV for V444 Cyg and CD Cru, respectively. The hard energy component ($\\rm{kT_2}$) was varied from a minimum value of  1.88 keV  to a maximum value of  9.61 keV  for V444 Cyg. It was maximum at primary eclipse and at the phase of 0.29, and minimum during the secondary eclipse and at the phase of 0.13. For CD Cru, $\\rm{kT_2}$ varies from a minimum value 1.63  keV during outside the eclipse to a maximum value of  4.27 keV at primary eclipse. For the star V444 Cyg, $\\rm{N_H^1}$ was maximum at phase 0.29, otherwise it was  constant at all phases. However, $\\rm{N_H^1}$ was  found to be constant throughout the orbital phase with a mean value of 1.14$\\pm$0.11 $\\times \\rm{10^{22}~cm^{-2}}$ for CD Cru. For V444 Cyg, $\\rm{N_H^2}$  was increased from phase 0.0 to its maximum value at phase 0.13. Afterward it was decreased and became almost constant from phase 0.29  to phase 0.47. For CD Cru, $\\rm{N_H^2}$ was minimum at phase 0.0 and maximum at phase 0.78.  \\begin{figure*} \\centering \\includegraphics[width=6.0in, height=5.0in]{RGS_all_500b.eps} \\caption{RGS spectra of the WR binaries at different binary phase for the stars V444 Cyg (upper panel) and CD Cru (lower panel).} \\label{fig:rgs} \\end{figure*} \\subsection{RGS spectra} \\label{sec:an_rgs}   We combined the first and second order spectra of detector RGS1 and RGS2 to inspect the main spectral lines. We obtained the RGS fluxed spectra using task {\\sc rgsfluxer} and is shown in Fig.~\\ref{fig:rgs} for different phases for both the binaries. We identified the prominent emission lines for V444 Cyg at wavelengths Mg XII (6.50 \\AA; 6.74 \\AA), Mg XI (9.23 \\AA), Fe XVIII (11.40 \\AA) and Ca XVI (22.61 \\AA, 21.43 \\AA, 22.30 \\AA). Individual lines show intensity variations from spectrum to spectrum. Unfortunately, our observations are at the limit of detectability with the data available with poor count statistics. Therefore, the signal-to-noise ratio of individual lines are not sufficient to perform a detailed quantitative analysis of the RGS line spectrum for any of the massive stars. ",
        "conclusions": "\\label{sec:discuss} We have analyzed the X-ray emission properties of two WR binaries with strong stellar winds. We discuss below the implications of the X-ray results.  \\subsection{Phase-locked X-ray variability}\\label{sec:discuss_lt}  The X-ray temporal and spectral analyses of WR binaries V444 Cyg and CD Cru showed a phase-locked phenomenon. The phase resolved X-ray spectroscopic observations show that the soft energy component  and the corresponding $\\rm{N_H^1}$ were also found to be nearly constant throughout the  binary phases of both stars. The $\\rm{L_X^S}$ was  found to be minimum during the  primary eclipse only. These results lead to the radiative wind shock origin of soft energy component. In the case of CD Cru, the modulation of $\\rm{L_X^S}$  well matches with the optical light curve reported by Moffat et al. (1990).  Lamontagne et al. (1996) have reported a similar shape of optical photometric light curves for 13 WR binaries including CD Cru and they said that it could be due to the \"atmospheric eclipse\". The X-ray flux in the hard energy band was found to be minimum at both primary and secondary eclipse for both stars.  This implies  that the hard energy component could be due to the presence of wind-wind collision zone which is located somewhere in between the O-type and WR star, and therefore, minimum when either of the star in the binary system is eclipsed. A similar behavior of X-ray light curves was seen in the EINSTEIN, ROSAT and ASCA observations of V444 Cyg (Moffat et al. 1982; Corcoran et al. 1996 ; Maeda et al. 1999).  The  hot temperature ($\\rm{kT_2}$) and the corresponding $\\rm{N_H^2}$ were also varied with orbital  phases of both binaries.  The eccentricity of both binary systems is almost zero, therefore, the variation in the temperature corresponding to the hard energy component could be the result of variation in $\\rm{N_H^2}$, i.e., varying optical depth (see \\S\\ref{sec:intro}).  This further supports the  wind-wind collision phenomenon in these binaries.  \\subsection{X-ray Temperatures of Plasma}\\label{sec:discuss_spt}  The X-ray spectra of WR binaries are well defined by two temperature plasmas. The values of  $\\rm{log(L_X/L_{bol})}$ in the total energy band  are found to be -6.6 and -7.2  in for V444 Cyg and CD Cru, respectively, which are similar to those for other WR (WN) stars observed from {\\sc XMM-Newton} and {\\sc CHANDRA} (G$\\ddot{u}$del \\& Naz$\\acute{e}$ 2009). The temperature of the cool component of WR binaries are comparable to other WR binaries e.g., 0.56--0.67 keV for WR1 (Ignace, Oskinova \\& Brown 2003) and  0.7--0.8 keV for WR 147 (Skinner et al. 2007). However, the temperatures corresponding to hot energy component o f these WR binaries are slightly more than some of the similar binaries. The possible mechanisms for generation of X-rays are discussed below.   \\subsubsection{ Instabilities driven radiative wind shocks } The wind-shock model predicts the intrinsic instability of the line driving mechanism. The standard model estimates the shock velocities by the relation  $\\rm{kT_{sh}=1.95 \\mu v^2_{shock}}$ (Lucy 1982; Luo, McCray \\& Maclow 1990). The temperatures of the cool energy components are found to be almost similar ($\\sim$ 0.6 keV) for both the binary systems (see Table~\\ref{tab:V444Cyg_fit} and Table~\\ref{tab:CDCru_fit}). Adopting the mean particle weight $\\rm{\\mu\\approx}$1.16 (Skinner et al. 2007) for WN stars and $\\rm{\\mu\\approx}$0.62 for O-type stars (Cassinelli et al. 2008), we derived the \"average\" value of shock velocities $\\approx$ 515 $\\rm{km~s^{-1}}$ for WN stars and $\\approx$ 704 $\\rm{km~s^{-1}}$ for O-type stars corresponding to the cool component. These values are about a factor of 2 larger than those predicted by radiative shock  model of Lucy (1982). However, the advance version of the wind shock model by Owocki et al. (1988) predicts  X-ray emission up to 1 keV. Moreover, the temperatures of the cool energy component and $\\rm{N_H^1}$ are not varying with orbital phase. Therefore, it appears that the cool energy component from  V444 Cyg and CD Cru could be generated by either of the binary components and may be explained by instabilities in radiation-driven wind shocks.   \\subsubsection{Magnetically Confined Wind Shock} Babel \\& Montmerle (1997) have suggested that the presence of magnetic fields confine the wind, which may be an important ingredient for the production of X-ray emission (Babel \\& Montmerle 1997). The degree of confinement of wind by the magnetic field is derived in terms of a confinement  parameter $\\rm{ \\Gamma=B^{2}_{0} R^{2}_{\\ast}/\\dot{M}v_{\\infty} }$, where  $\\rm{B_0}$ is the surface equatorial magnetic field strength (ud-Doula \\& Owocki 2002). For a confined wind  model $\\rm{\\Gamma>>1}$. Using the values of stellar radius ($\\rm{R_{\\ast}}$) and stellar mass loss rate ($\\rm{\\dot{M}v}$) as given in  Table~\\ref{tab:par_str} for V444 Cyg binary components, the minimum magnetic field required to confine the wind is derived to be 0.14 kG for the primary and 1.25 kG  for the secondary. Similarly, for CD Cru, we derived minimum magnetic field of 0.16 kG and 1.96 kG for the primary  and the secondary components, respectively. Based on the formula given by  Babel \\& Montmerle  (1997) and using above estimated values of B, the X-ray luminosity are estimated to be $10^{34.66}$~\\egs~and $10^{35.74}$ ~\\egs~for  primary  and the secondary components of V444 Cyg, and $10^{35.06}$~\\egs~and  $10^{36.8}$ ~\\egs~for primary  and the secondary components of CD Cru. Such a high X-ray luminosity has not been observed from any of the WR binaries. However, smaller fields can exist and the confinement may be limited to just above the magnetic equatorial plane.  \\subsubsection{Colliding wind shock model} \\input{model.tab} Assuming that shock velocities have reached to the terminal velocities, the standard model predicts a maximum temperature for the shocked region from winds of massive stars by the relation  $\\rm{kT^{max}_{sh}=1.95 \\mu v^2_{\\infty}}$ (Luo, McCray \\& Maclow 1990). The nominal wind parameters of massive stars (see Table~\\ref{tab:par_str}) and the mean particle weight $\\rm{\\mu}$ for the WN stars and for the O-type stars give the maximum temperature generated by the individual components of binaries. The derived values are 7.80 keV + 7.21 keV for V444 Cyg (O6+WN5) and 10.88 keV + 13.69 keV for CD Cru (O5+WN6). The maximum observed temperature at phase 0.29 corresponding to hot component is  found to be similar to that of maximum possible values in the case of V444 Cyg. However, in CD Cru maximum observed temperature is lesser than that of maximum possible values.  The phase locked variability of  the hard energy component shows that the observed hot X-ray temperatures in massive  binaries could originate from the collision of stellar winds. The distances from the stars where these winds meet are derived using the relations (De Becker 2007; Stevens, Blondin \\& Pollock 1992),  \\begin{equation} \\rm{  r_{OB}=\\frac {1}  {1+\\eta^{1/2}}  D} \\end{equation}  \\noindent where $\\rm{r_{OB}}$ and D are the distance to the collision zone from the primary and the separation between the binary components, respectively. The wind momentum ratio $\\rm{\\eta}$ is expressed as,  \\begin{equation} \\rm{\\eta=\\frac{\\dot{M_2}v_{\\infty,2}}{\\dot{M_1}v_{\\infty,1}}} \\end{equation}  \\noindent where $\\rm{\\dot{M_1}}$ and $\\rm{\\dot{M_2}}$ are mass loss rates of primary and secondary components, respectively, and $\\rm{v_{\\infty,1}}$ and $\\rm{v_{\\infty,2}}$ are terminal velocities of primary and secondary components, respectively. Using the above relations, we derived wind momentum ratios of 7.03 and 24.60 for the stars V444 Cyg and CD Cru , respectively. Distances to the collision zone from the primary are estimated to be 0.27 and 0.17 times of their corresponding binary separation (D) for V444 Cyg and CD Cru, respectively. Therefore, the collision zone exists to be very close to the primary O-type star in case of both the binaries. Further, the collision wind zone will produce a bow shock around the O-type star , the shock volume and the emission measure are dominated by the WR winds, therefore, X-ray emitting plasma also shows non-solar abundances.  The gas in a colliding wind  region could be either adiabatic or radiative and depends upon the value of cooling parameter ($\\chi$), which is defined as the ratio of the cooling time ($\\rm{t_{cool}}$) of the shocked gas to the escape time ($\\rm{t_{esc}}$) from the intershock region (Stevens, Blondin \\& Pollock 1992).  \\begin{equation} \\rm {\\chi = \\frac{t_{cool}}{t_{esc}} = \\frac{v^4_3 d_7} {\\dot{M}_{-7}}} \\end{equation}   \\noindent where $\\rm{v_{3}}$ is the pre-shock velocity in the unit of 10$^3$ $\\rm{km~s^{-1}}$, $\\rm{d_{7}}$ is the distance from stellar center to the shock in the unit of 10$^7$ km and $\\rm{\\dot{M}_{-7}}$ is the mass loss rate in the unit of $\\rm{10^{-7}M_{\\odot} yr^{-1}}$. Using the parameters given in Table~\\ref{tab:par_str} and the values from above estimation, we derived cooling parameters $\\rm{\\chi_1}$ and $\\rm{\\chi_2}$ for primary and secondary components, respectively. The estimated values of $\\rm{\\chi_1}$ and  $\\rm{\\chi_2}$ are given Table~\\ref{tab:model}. The winds from WR stars are found to be clearly radiative ($\\rm{\\chi<1}$) for both the binaries. The intrinsic luminosity can be estimated for each component of the binary using the following relation (Pittard  \\& Stevens 2002) :  \\begin{equation} \\rm {L_X = 0.5~\\Xi~\\dot{M}~v^2 } \\end{equation}  \\noindent where $\\rm{\\Xi}$ accounts for a geometrical and inefficiency factors and v is the wind speed at contact surface (i.e. pre-shock velocities derived from observed temperatures). The values of $\\rm{\\Xi_{1}}$ (0.403 for V444 Cyg ; 0.564 for CD Cru) and $\\rm{\\Xi_{2}}$ (0.033 for V444 Cyg ; 0.0042 for CD Cru) for primary and secondary stars, respectively, are taken from Pittard  \\& Stevens (2002). The estimated values of intrinsic luminosities of each component and the total intrinsic luminosities of binary systems are given in Table~\\ref{tab:model}. These values are nearly 3 orders of magnitude higher than those of the  observed values. A similar results have been found for other WR binaries, e.g.,  HD 159176 (De Becker et al. 2004); WR 147 (Skinner et al. 2007). De Becker et al. (2004) suggested that the disagreement between observed and theoretical predictions could possibly be explained by (a) the kinetic power of the collision should be considered as an upper limit on the X-ray luminosity, (b) higher value of the parameter $\\rm{\\eta}$ should be considered, (c) diffusive mixing between hot and cool material is likely to exist due to the instability of the shock front, and (d) orbital effects should also be included to study such systems. The observed and theoretically predicted values of X-ray luminosities are shown in Fig.~\\ref{fig:model}. This model also predicts the phase-locked variations which can be seen in Fig.~\\ref{fig:ana_phase}. Therefore, it appears that the hard energy component is most likely associated with wind collision zone.  \\begin{figure} \\centering \\includegraphics[width=3.5in, height=3.5in]{model_obs_hard.eps} \\caption{Observed vs Model X-ray luminosity.} \\label{fig:model} \\end{figure}   \\subsubsection{ Non-thermal emission}  There is no strong justification for invoking non-thermal origin of X-ray emission since the observed X-ray emission at higher energies can be fitted satisfactorily with a thermal model (see Fig.~\\ref{fig:spt_V444Cyg} and Fig.~\\ref{fig:spt_CDCru}). However, the present data is limited up to 10 keV and it is difficult to eliminate the contribution of non-thermal components from the spectra which is dominated by the presence of the strong thermal emission, and the non-thermal emission is overwhelmed by the thermal emission (De Becker 2007). Therefore, we can state that the spectra of massive stars in the present sample are fairly well explained by two-temperature plasma models."
    },
    "0911/0911.4924_arXiv.txt": {
        "abstract": "{}{We investigate the dependence of $\\gamma$-ray brightness of blazars on intrinsic properties of their parsec-scale radio jets and the implication for relativistic beaming.}{By combining apparent jet speeds derived from high-resolution VLBA images from the MOJAVE program with millimetre-wavelength flux density monitoring data from Mets\\\"ahovi Radio Observatory, we estimate the jet Doppler factors, Lorentz factors, and viewing angles for a sample of 62 blazars. We study the trends in these quantities between the sources which were detected in $\\gamma$-rays by the $Fermi$ Large Area Telescope (LAT) during its first three months of science operations and those which were not detected.}{The LAT-detected blazars have on average higher Doppler factors than non-LAT-detected blazars, as has been implied indirectly in several earlier studies. We find statistically significant differences in the viewing angle distributions between $\\gamma$-ray bright and weak sources. Most interestingly, $\\gamma$-ray bright blazars have a distribution of comoving frame viewing angles that is significantly narrower than that of $\\gamma$-ray weak blazars and centred roughly perpendicular to the jet axis. The lack of $\\gamma$-ray bright blazars at large comoving frame viewing angles can be explained by relativistic beaming of $\\gamma$-rays, while the apparent lack of $\\gamma$-ray bright blazars at small comoving frame viewing angles, if confirmed with larger samples, may suggest an intrinsic anisotropy or Lorentz factor dependence of the $\\gamma$-ray emission.}{} ",
        "introduction": "One of the most important discoveries of the Energetic Gamma-Ray Experiment Telescope (EGRET) on-board the {\\it Compton Gamma-Ray Observatory} in the 1990s was the detection of over 65 active galactic nuclei (AGN) at photon energies above 100\\,MeV \\citep{har99,mat01}. The detected sources were almost exclusively blazars, a class of highly variable AGN comprised of flat spectrum radio quasars and BL Lac objects. The distinctive characteristic of blazars is a relativistic jet oriented close to our line-of-sight. Synchrotron radiation of energetic electrons in the jet dominates the low energy end of the blazar spectral energy distribution. This emission is strongly beamed due to relativistic effects which increase the observed flux density of a stationary jet by a factor of $\\delta^{2-\\alpha}$ and that of distinct ``blobs'' in the jet by a factor of $\\delta^{3-\\alpha}$.  Here $\\delta$ is the jet Doppler factor and $\\alpha$ is the spectral index defined as $S_\\nu \\propto \\nu^{+\\alpha}$ \\citep{bla79}. The Doppler factor is defined as $\\delta = [\\Gamma (1-\\beta \\cos \\theta)]^{-1}$, where $\\Gamma = (1-\\beta^2)^{-1/2}$ is the bulk Lorentz factor, $\\beta$ is the jet speed divided by the speed of light, and $\\theta$ is the angle between the jet and our line-of-sight. The requirement that the $\\gamma$-ray-bright sources are transparent to $\\gamma\\gamma$ pair production, together with their small sizes deduced from the fast $\\gamma$-ray variability, strongly suggest that the $\\gamma$-ray emission originates in the jet and is also relativistically beamed, in the same manner as the radio emission \\citep{mon95,mat93}. The many correlations found between the $\\gamma$-ray emission detected by EGRET and the radio/mm-wave properties of blazars further support this scenario \\citep{val95, jor01b, jor01a, lah03, kel04, kov05}. Since the $\\gamma$-ray spectrum is typically steeper than the radio spectrum \\citep{abd09b}, it is possible that the $\\gamma$-ray emission is even more enhanced by Doppler boosting than the radio emission.  Although inverse Compton (IC) scattering of soft photons off relativistic electrons in the jet is currently the favoured model for the $\\gamma$-ray emission, there is a substantial controversy within this model about the origin of the target photon field and the location of the emission site. The seed photons could be, for example, synchrotron photons emitted by the same electrons which scatter them later, \\citep[synchrotron self-Compton model, SSC;][]{mar92, blo96} or synchrotron photons emitted by electrons in a different layer of the jet \\citep{ghi05}. The seed photons can also originate in sources external to the jet like the accretion disk, the broad line region clouds or the dust torus \\citep[external Compton model;][]{der92,ghi96,sik94,bla00}. Beside these leptonic models there are also a number of models where $\\gamma$-rays are produced by hadronic processes initiated by relativistic protons co-accelerated with electrons \\citep[e.g.,][]{man93}.  The Large Area Telescope (LAT) on-board the {\\it Fermi Gamma-ray Space Telescope} is a successor to EGRET, with much better sensitivity, larger energy range (up to 300 GeV), better angular resolution, and larger field-of-view \\citep{atw09}. In only its first three months of science operations, the LAT detected 205 bright $\\gamma$-ray sources at $>10\\sigma$ level \\citep{abd09a}, 116 of which are associated with AGN at high galactic latitudes ($|b| \\ge 10^\\circ$) \\citep{abd09b}. Here we refer to these 116 sources as ``LAT-detected sources''. Based on the long-term monitoring of the radio jet motions with the Very Long Baseline Array (VLBA), it was recently shown that the LAT-detected quasars have significantly faster apparent jet speeds and higher core brightness temperatures than non-LAT-detected quasars \\citep{lis09b,kov09}. The LAT $\\gamma$-ray photon flux also correlates with the compact radio flux density and the flares in $\\gamma$-rays and radio seem to happen in the VLBI cores within a typical apparent time separation of up to several months \\citep{kov09}. The $\\gamma$-ray bright blazars also have larger-than-average apparent jet opening angles \\citep{pus09}. These findings indicate that $\\gamma$-ray bright blazars are likely more Doppler boosted than $\\gamma$-ray faint ones, but what remains unknown is a possible dependence of {\\it intrinsic} $\\gamma$-ray luminosity on parsec-scale jet properties such as the bulk Lorentz factor or the viewing angle in the comoving frame of the jet \\citep{lis09b}. In this paper we confirm the connection between the $\\gamma$-ray brightness and the Doppler factor of the parsec-scale jet for a sample of 62 blazars.  We have also combined measurements of the apparent jet speeds and temporal variability at mm-wavelengths to derive jet Lorentz factor and the viewing angle in the observer's frame and in the frame comoving with the jet for 57 blazars in order to study how the $\\gamma$-ray detection probability depends on these intrinsic jet parameters.  \\begin{figure*} \\centering \\includegraphics[width=0.75\\textwidth]{savolainen_v4fg1.eps} \\caption{{\\it Upper left:} Variability Doppler factor ($\\delta_\\mathrm{var}$) distributions for non-LAT-detected (upper sub-panel) and LAT-detected (lower sub-panel) blazars in the Mets\\\"ahovi-MOJAVE sample. Quasars are denoted by shaded bins. {\\it Lower left:} Bulk Lorentz factor ($\\Gamma$) distributions. {\\it Upper right:} Distributions of the angles between the jet and our line-of-sight in the observer's frame ($\\theta$). {\\it Lower right:} Distributions of the viewing angles in the comoving frame of the jet ($\\theta_\\mathrm{src}$). The values were calculated assuming $\\Lambda$CDM cosmology with $H_0 = 71$\\,km\\,s$^{-1}$\\,Mpc$^{-1}$, $\\Omega_\\mathrm{M} = 0.27$, and $\\Omega_\\mathrm{\\Lambda} = 0.73$.} \\label{fig1} \\end{figure*} ",
        "conclusions": "The variability Doppler factors from \\citet{hov09} were determined over a multi-year monitoring program while the $\\gamma$-ray detections used in our analysis are based on just three-months of LAT monitoring.  The sharp distinction seen in the $\\delta_\\mathrm{var}$ values between the LAT-detected and non-LAT-detected sources strongly supports the idea that $\\delta_\\mathrm{var}$ must remain fairly constant with time. Since the LAT-detected and non-LAT-detected sources are treated in the exact same way with respect to the derivation of their $\\delta_\\mathrm{var}$, any temporal variation in the Doppler factor (or in $T_\\mathrm{b,int}$) would thus only serve to destroy the possible correlation and could not create one. Therefore, we consider the result regarding higher $\\delta_\\mathrm{var}$ for the LAT-detected blazars to be very robust.  A similar result regarding the high $\\delta_\\mathrm{var}$ of $\\gamma$-ray bright sources was found earlier using the less uniform EGRET data \\citep{lah03}. The simplest and arguably most likely interpretation is that the $\\gamma$-ray bright blazars are indeed systematically more Doppler boosted than the $\\gamma$-ray weak ones. This interpretation is compatible with the LAT-detected quasars having faster apparent jet speeds \\citep{lis09b}, wider apparent jet opening angles \\citep{pus09}, and higher VLBI brightness temperatures \\citep{kov09}. An alternative interpretation would be that, for some reason, the intrinsic brightness temperature is systematically about a factor of three higher in the LAT-detected blazars than in the non-LAT-detected ones. In the latter case our assumption of a constant limiting brightness temperature would lead to an overestimation of $\\delta_\\mathrm{var}$ in the $\\gamma$-ray bright blazars. This alternative explanation would not, however, explain the faster apparent jet speeds or wider apparent jet opening angles of the LAT-detected sources.  The observed difference in the comoving-frame viewing angle distributions between the $\\gamma$-ray bright and weak blazars is an unanticipated result. The left panel of Fig.~\\ref{fig2} shows that the lack of LAT-detected blazars at large comoving-frame viewing angles can be explained by low Doppler factors of the sources at large $\\theta_\\mathrm{src}$. The beaming model does not, however, explain the lack of LAT-detected sources at {\\it small} values of $\\theta_\\mathrm{src}$. If this lack is real, it may reflect an intrinsic anisotropy of the $\\gamma$-ray emission in the comoving frame of the jet. This would have wide implications for the theoretical models of the high energy emission from blazars, since almost all of these models rely on the assumption that the $\\gamma$-ray emission is (nearly) isotropic in the rest frame of the relativistically moving sub-volume \\citep[e.g.][]{der95}. Possible sources of anisotropy in the $\\gamma$-ray emission include, for example, anisotropic $\\gamma\\gamma$ absorption or anisotropic seed photon field for inverse Compton scattering.  The right hand panel of Fig.~\\ref{fig2} shows that the non-LAT-detected sources at small comoving frame viewing angles have small Lorentz factors, except for \\object{B0804+499}. This may provide an alternative explanation for the apparent lack of $\\gamma$-ray bright sources at small comoving frame viewing angles if the \\textit{intrinsic} $\\gamma$-ray luminosity depends on the bulk Lorentz factor in addition to being relativistically beamed. The fact that \\object{B0804+499} is not a bright $\\gamma$-ray source, despite having a very high Doppler factor and a moderately high Lorentz factor, poses a problem for this explanation, but does not rule it out since the $\\gamma$-ray emission may be intermittent (or there may be something unusual in this particular source)."
    },
    "0911/0911.5113_arXiv.txt": {
        "abstract": "% The SuperAGILE (SA) instrument is a X-ray detector for Astrophysics measurements, part of the Italian AGILE satellite for X-Ray and Gamma-Ray Astronomy launched at 23/04/2007 from India. SuperAGILE is now studying the sky in the 18 - 60 KeV energy band. It is detecting  sources with advanced imaging and timing detection capabilities and good spectral detection capabilities. Several astrophysical sources has been detected and localized, including Crab, Vela and GX 301-2. The instrument has the skill to resolve correctly sources in a field of view of [-40, +40] degrees interval, with the angular resolution of 6 arcmin, and a spectral analysis with the resolution of 8 keV. Transient events are regularly detected by SA with the aid of its temporal resolution (2 microseconds) and using signal coincidence on different portions of the instrument, with confirmation from other observatories. The SA data processing scienti\ufb01c software performing at the AGILE Ground Segment is divided in modules, grouped in a processing pipeline named SASOA. The processing steps can be summarized in \\textit{data reduction}, \\textit{photonlist building}, \\textit{sources extraction} and \\textit{sources analysis}. The software services allow orbital data processing (near real-time), daily data set integration, Temporal Data Set (TDS) processing and TDS processing with source target optimization (TDS\\_SRC). Automatic data processing monitoring and interactive data analysis is possible from an internet connected workstation, with the use of SA data processing Web services. Many solutions were implemented in order to achieve fault tolerance. Archive management and data storage are performed with the help of relational database instruments. ",
        "introduction": "\\begin{deluxetable}{|c|l|l|} \\tablecaption{SASOA software Data Processing Stages \\label{E.16-tbl-1}} \\startdata \\textbf{\\begin{tiny}n.\\end{tiny}} &\\textbf{\\begin{tiny}STAGE\\end{tiny}} & \\textbf{\\begin{tiny}Description\\end{tiny}} \\nl \\hline \\hline \\begin{tiny}1\\end{tiny} & \\begin{tiny}TRIGGER\\end{tiny} & \\begin{tiny}control if new data are transmitted\\end{tiny}  \\nl \\hline \\begin{tiny}2\\end{tiny} & \\begin{tiny}DATA\\_DOWNLOAD\\end{tiny} & \\begin{tiny}if triggered download available data\\end{tiny} \\nl \\hline \\begin{tiny}3\\end{tiny} & \\begin{tiny}COMMANDS\\_PARSING\\end{tiny} & \\begin{tiny}parse run options\\end{tiny} \\nl \\hline \\begin{tiny}4\\end{tiny} & \\begin{tiny}LOCAL\\_SETTINGS\\end{tiny} & \\begin{tiny}set hosts and data coordinates \\end{tiny} \\nl \\hline \\begin{tiny}5\\end{tiny} & \\begin{tiny}DATA\\_FETCH\\end{tiny} & \\begin{tiny}load data input for actual run in local buffer \\end{tiny} \\nl \\hline \\begin{tiny}6\\end{tiny} & \\begin{tiny}FILE\\_NAMING\\end{tiny}& \\begin{tiny}creates file names following the standard\\end{tiny} \\nl \\hline \\begin{tiny}7\\end{tiny} & \\begin{tiny}CORRECTION\\end{tiny}& \\begin{tiny}perform data correction task\\end{tiny} \\nl \\hline \\begin{tiny}8\\end{tiny} & \\begin{tiny}ORBIT\\_REBUILDING\\end{tiny}& \\begin{tiny}rebuild data related on a given orbit\\end{tiny} \\nl \\hline \\begin{tiny}8\\end{tiny} & \\begin{tiny}DATA\\_REDUCTION\\end{tiny}& \\begin{tiny}data separation, reconstruction and equalization\\end{tiny} \\nl \\hline \\begin{tiny}9\\end{tiny} & \\begin{tiny}PHOTONLIST\\_BUILDING\\end{tiny}& \\begin{tiny}analyse data of the reduced photon list\\end{tiny} \\nl \\hline \\begin{tiny}10\\end{tiny} & \\begin{tiny}ATTITUDE\\_CORRECTION\\end{tiny}& \\begin{tiny}correction for the attitude wobbling\\end{tiny} \\nl \\hline \\begin{tiny}11\\end{tiny} & \\begin{tiny}IMAGING\\end{tiny}& \\begin{tiny}image extraction \\end{tiny} \\nl \\hline \\begin{tiny}12\\end{tiny} & \\begin{tiny}FILING\\end{tiny}& \\begin{tiny}archive produced data\\end{tiny} \\enddata  \\end{deluxetable} The SASOA (Super Agile Standard Orbital Analysis) pipeline is composed by a succession of elaboration tasks applied on the SA detected eventlist, the output of a certain task is the input of the following. Macro-tasks launch other subfunctions, some of them are implemented by external scripts and batch programs. In SASOA, a daemon running an endless loop provides to check if data for the incoming contacts are available and when they are found, the trigger for contact data processing is given. Data for a given contact mostly contain the data related to measurements acquired by the AGILE instruments during the last orbit. The failure of a single task of the pipeline may block the production of the final output data for a single run but the failure of a single contact run does not stop the daemon for the next incoming triggers and related data processing run. Different kinds of triggers can be used: the first way used to trigger the sasoa pipeline is to control if a file with extension .ok is generated by the preprocessing system for the incoming contact data, i. e. if a file named \"VC1.002469.009.ok\" is present, data for the contact 2469 are correctly generated and available on the data area of the preprocessing machine. Another trigger modality is to control if a value in a database of the ASDC Data Center is changed. When the system has assured that data are present, then the data are downloaded from the \\textit{gtb} server in Bologna using the \\textit{wget} utility (see M. Trifoglio et Al, \u201c\\textit{Archiving the AGILE Level-1 telemetry data}, Astronomy and Astrophysics, 2008\u201d, ), data from ASDC are downloaded with \\textit{wget}  also, but they have to be downloaded one by one (mirror and -A  option are disabled) to avoid transmission band overloading. Data products (histograms, plots, statistics \\& reports) are generated at different levels, hosts coordinates, data directories paths and other configurations of the processing system are set. The most important input telemetry data used for SA instrument data analysis are eventlist data (SA TM 3905), attitude data (SA TM 3914), and ephemerides data (SA TM 3916). The file names for i/o file of all tasks are generated following the official naming convention, the correction stage performs time synchronization among input data files and other corrections. Then some actions useful to rebuild the orbital organization of the data are performed: if the contact contains more than 1 orbit, then data are divided in the respective parts and the pipeline is launched on the first part of data. After the run end, the system is triggered on the 2nd part. In the Data Reduction stage the eventlist from the previous stage is taken as input (complete orbit or orbital part), with efficiency files and then a reduced output eventlist is generated. In the Sources Extraction stage, the data are corrected removing wobbling effects, then the exposure is evaluated using informations regarding Earth Occultation and SAGA time periods. After these steps the imaging procedure is applied and sources lists are generated, giving position and flux for the detected sources.  \\\\ \\begin{figure}[t] \\epsscale{0.70} \\plotone{E.16_1.eps} \\caption{TDS processing schema} \\label{E.16-fig-1} \\end{figure} \\begin{figure}[t] \\begin{center} \\epsscale{0.70} \\includegraphics[angle=-90, width=8cm]{E.16_2.eps} \\caption{TDS SRC software GUI} \\label{E.16-fig-2} \\end{center} \\end{figure} \\begin{figure}[t] \\epsscale{0.70} \\plotone{E.16_3.eps} \\caption{SA software monitor} \\label{E.16-fig-3} \\end{figure} ",
        "conclusions": "The SA scientific processing system is ready, enabling the access from web solve installation \\& portability problems. Processing stations will be added to fit with the actual number of data analysis software users. All SA software is portable, we have in use running versions of the whole processing system under Linux OpenSUSE (until 10.3 version, for 32 \\& 64bit architecture)."
    },
    "0911/0911.2219_arXiv.txt": {
        "abstract": "While limited to low spatial resolution, the next generation low-frequency radio interferometers that target 21 cm observations during the era of reionization and prior will have instantaneous fields-of-view that are many tens of square degrees on the sky. Predictions related to various statistical measurements of the 21 cm brightness temperature must then be pursued with numerical simulations of reionization with correspondingly large volume box sizes, of order 1000 Mpc on one side. We pursue a semi-numerical scheme to simulate the 21 cm signal during and prior to Reionization by extending a hybrid approach where simulations are performed by first laying down the linear dark matter density field, accounting for the non-linear evolution of the density field based on second-order linear perturbation theory as specified by the Zel'dovich approximation, and then specifying the location and mass of collapsed dark matter halos using the excursion-set formalism. The location of ionizing sources and the time evolving distribution of ionization field is also specified using an excursion-set algorithm. We account for the brightness temperature evolution through the coupling between spin and gas temperature due to collisions, radiative coupling in the presence of Lyman-alpha photons  and heating of the intergalactic medium, such as due to a background of X-ray photons. The hybrid simulation method we present is capable of producing the required large volume simulations with adequate resolution in a reasonable time so a large number of realizations can be obtained with variations in assumptions related to astrophysics and background cosmology that govern the 21 cm signal. ",
        "introduction": "Observations of the 21 cm line of neutral hydrogen are currently considered to be one of the most promising probes of the epoch of reionization (EoR) and possibly even the preceding period, during the so called dark ages \\citep{madau97, loeb04, gnedin04, furlanetto04a, zaldarriaga04,sethi05, bharadwaj05, morales03, loeb04}. Given the line emission, leading to a frequency selection for observations, the 21 cm data provides a tomographic view of the reionization and pre-reionization process \\citep{santos05, furlanetto04c} and can also be an exquisite cosmological probe providing complimentary information when compared to cosmic microwave background anisotropies \\citep{mao08,santos06,mcquinn06,bowman07}. New large area radio interferometers are now being developed, with the promise of measuring this signal, namely, LOFAR\\footnote{http://lofar.org} in the Netherlands, MWA\\footnote{http://web.haystack.mit.edu/arrays/MWA} in Australia and the low-frequency extension of the planned Square Kilometre Array (SKA\\footnote{http://www.skatelescope.org}).  Motivated by the observational possibilities offered by the current and upcoming low frequency radio interferometers a great deal of effort has been underway in order to fully understand and generate the expected 21 cm signal that will be seen by these experiments (see \\citealt{furlanetto06c} for a review). Analytical models can be very useful to quickly generate the signal with high dynamic range, in particular the power spectrum of 21 cm brightness temperature fluctuations \\citep{furlanetto04b,furlanetto06a,sethi05,barkana07} and for testing the ability of a given experiment to constraint cosmological and astrophysical parameters \\citep{santos06,mcquinn06,mao08}. The analytical models have also been useful to understand the possible contributions to the 21 cm signal at high redshifts \\citep{barkana05,pritchard06}, although it is at low redshifts ($z\\lesssim 10$) that they seem to provide a better description of the 21 cm 2-point correlation function \\citep{santos08}. However, these models still have problems dealing with the spatial distribution of the reionization process, such as bubble overlap and ignores complicated astrophysics during reionization.  Numerical simulations, on the other hand, can potentially provide an improved description of the 21 cm brightness temperature signal from first principles. They typically involve a combination of a N-body algorithm for dark matter, coupled to a prescription for baryons and star formation (often complemented by a hydrodynamical simulation), plus a full radiative transfer code to propagate the ionizing radiation \\citep{gnedin00a, razoumov02,sokasian03,ciardi03,kohler05,iliev06,zahn06,mcquinn07,shin07,trac06,baek09}. These simulations have the advantage of properly dealing with the spatial distribution of the fields (such as bubble overlapping or the non-sphericity of the bubbles) and can be used to extract more information than just the 21 cm anisotropy power spectrum. The disadvantage is that they are slow to run, being constrained in dynamical range to sizes typically smaller than 143 Mpc, and producing a large set of simulations for cosmological and astrophysical parameter estimates is unrealistic.  Instead the 21 cm community makes use of hybrid approaches involving simulations and analytical techniques. \\citet{thomas09} generate faster simulations by post-processing a dark matter simulation with a 1-D radiative transfer code to create bubbles around the sources of radiation embedded in the dark matter halos. \\cite{mesinger07} have taken a hybrid approach between the analytical models \\citep{furlanetto04b,furlanetto04c} and the spatial description provided by the full numerical simulations. This is based on 3-D Monte Carlo realizations of the dark matter density field combined with an ``analytical prescription'' in order to generate the halo and ionization fields. \\citet{zahn06} also considered a similar analytical prescription but based directly on the linear density field (see also \\citealt{alvarez09}). Although much faster than previous numerical approaches, the hybrid approach of \\citet{mesinger07} is still constrained to ``small'' sizes depending on the amount of shared memory available in the computer and to low redshifts since it neglects the spin temperature evolution for the 21 cm calculation.  It is useful to note that the proposed low-frequency interferometers that plan to observe the 21 cm signal from reionization have low spatial resolution capabilities but very large fields of view - 5x5 deg$^2$ or more, requiring boxes of at least 1000 Mpc in size if we want to properly simulate a single field of view and pass such a simulation through the observation pipeline in order to understand how instrumental noise and systematic effects impact the observations. At the same time we need to be able to resolve halos of the order of 10$^8 M_\\odot$ corresponding to a cooling temperature of $10^4$ K that is deemed necessary to host the ionizing sources responsible for reionization. This implies a resolution of 0.14 Mpc requiring the use of large boxes with $(7000)^3$ cells, making even a hybrid method not useful. Moreover, at high redshifts we need to take into account detailed physics that determine the 21 cm brightness temperature, such as the heating of the inter galactic medium (IGM) or the Lyman-$\\alpha$ coupling of the spin temperature to the gas temperature, which makes it a challenging exercise to implement a complete radiative transfer/gas dynamics algorithm in a large volume simulation box.  In this paper, we propose a semi-numerical technique, partially following the hybrid approach prescribed in \\citet{mesinger07}, capable of quickly generating an end-to-end simulation of the 21 cm signal even at high redshifts when the effect of the spin temperature is non-negligible. Moreover, this method can be used to simulate very large volumes (e.g. 1000 Mpc), crucial to simulate the field-of-view of next generation of radio telescopes, without sacrificing the speed or requiring unpracticable amounts of computer memory. The code to generate this type of simulation will be provided publicly online\\footnote{\\url{http://www.simfast21.org}} and it will be subject to continuous improvement through calibration against full radiative transfer/hydrodynamic simulations. The large volume simulations created with this method will be part of the SKA Design Studies Simulated Skies ($S^3$) initiative and we hope this approach can prove useful to generate sky models for the future 21 cm experiments, which is crucial to test for calibration issues and foreground removal as well as to study the features of the 21 cm signal that can be probed by a variety of experiments and to check the dependence of various statistics extracted from the observations on the underlying astrophysical and cosmological parameters.  In the next Section we introduce the algorithm to generate the simulation of the 21 cm signal up to high redshifts, while in Section~3 we explain how to extend the simulations to very large volumes. For results presented here, we assumed a flat $\\Lambda$CDM universe with the following cosmological parameters: $\\Omega_m=0.28$, $\\Omega_b=0.046$, $h=0.70$, $n_s=0.96$ and $\\sigma_8=0.817$ based on the results from WMAP5, BAO and SN (see \\citealt{hinshaw09} and references therein). Throughout the paper we quote all quantities in comoving units unless otherwise noted. ",
        "conclusions": "We presented a semi-numerical method capable of quickly generating end-to-end simulations of the 21 cm signal even at the high redshifts where the spin temperature is non-negligible. The algorithm allows to generate brightness temperature boxes with very large volumes, e.g. (1000 Mpc)$^3$, crucial to properly simulate the field of view of the next generation of radio-telescopes, without sacrificing the speed or requiring unfeasible computer resources\\footnote{During the review process of this paper another algorithm was presented by \\citet{mesinger10}.}.  The corresponding code (SimFast21) implements the following prescription: \\begin{itemize} \\item Monte Carlo generation of the dark matter density and velocity fields at z=0 assuming Gaussian distribution functions. \\item Determination of the halo catalog using an excursion-set formalism on the linear density field at any given redshift. For large volumes a biased Poisson sampling is used to introduce small mass halos. \\item Adjustment of the halo and dark matter locations using the Zel'dovich approximation. \\item Simulation of $x_i({\\bf x},z)$ from FFTs of the previous boxes. \\item Extraction of the star formation rate density from the halo mass distribution. \\item Extraction of gas temperature fluctuations $T_K({\\bf x},z)$ . \\item Calculation of the Ly$_\\alpha$ coupling and collisional coupling parameters. \\item Determination of the 21 cm brightness temperature boxes using the above quantities. \\end {itemize} The computational time required for the above steps depends mainly on the volume to be covered (our bubble filtering procedure is now pratically independent of the ionization fraction due to the FFT based algorithm we are now using). This basically affects the halo determination running time, since larger volumes imply a larger number of halos and operations for overlap checking. Nevertheless, the typical time for the computation of the above steps is about 2 hours for the simulation with L=1000 Mpc, which can be considered remarkably fast. We are expecting that further optimization of the code can still reduce this time.  Although much faster than hydrodynamical numerical simulations, our analysis shows that relevant quantities such as the halo mass function, the halo mass power spectrum, the star formation rate or the ionization fraction power spectrum are all consistent with the numerical simulations with which we compared our results to. In particular, the ionization fraction power spectrum is similar to the one obtained with radiative transfer codes (RT) when comparing to the same volume. We still find some differences on the largest scales when comparing to simulations with relatively small volumes (143 Mpc), which we can relate to the production of larger HII bubbles. These differences become much smaller as we increase the simulation volume and the change of size in the largest bubbles becomes less important when compared to the overall size of the simulation.  The dependence on the astrophysical parameters of the simulation was encoded in three functions: the ionization efficiency, $\\zeta$, the Ly$_\\alpha$ spectral distribution function of the sources, $\\epsilon_\\alpha$ and the X-ray spectral distribution function, $\\epsilon_X$. These functions can be easily changed for a model of our choice and the code can then quickly generate new simulations of the signal (even faster if we keep the same cosmology). By combining all the unknowns into a physically meaningful small set of parameters we can easily probe the huge intrinsic parameter space available at high redshifts probed by 21 cm observations. This versatility is included in our algorithm. The possibility to generate reionization simulations from scratch without the need for supercomputers is a great advantage allowing to experiment with different models. Moreover, this allows room for improvement through calibration with full numerical simulations. The code can be a useful tool to generate sky models for future 21 cm experiments, important to test the observation pipeline, study how the foregrounds affect the observations and can be separated given a noise model, and to develop optimal estimators for signal extraction. The code will be released for public use in due course and we welcome the community participation in its updating and upgrading as well as development of additional applications on observational probes of reionization beyond the 21 cm observations."
    },
    "0911/0911.0523.txt": {
        "abstract": "{}{}{}{}{} % 5 {} token are mandatory  \\abstract % context heading (optional) % {} leave it empty if necessary {} % aims heading (mandatory) { We investigate the properties and the environment of radio sources with optical counterpart from the combined VLA-COSMOS and zCOSMOS samples. The advantage of this sample is the availability of optical spectroscopic information,  high quality redshifts, and accurate density determination. } % methods heading (mandatory) {By comparing the star formation rates  estimated from the optical spectral energy distribution with those based on  the radio luminosity, we divide the  radio sources in three families,  passive AGN,  non-passive AGN and  star forming galaxies. These families occupy specific regions of the 8.0-4.5 $\\mu$m infrared color- specific star formation plane, from which we extract the corresponding control samples.} % results heading (mandatory) {Only the passive AGN have a significantly different environment distribution from their control sample. The fraction of  radio-loud passive AGN increases from $\\sim 2\\% $ in underdense regions to $\\sim 15 \\% $ for overdensities (1+$\\delta$) greater than 10. This trend is also present  as a function of richness of the groups hosting the radio sources. Passive AGN in overdensities tend to have higher radio luminosities than those in  lower density environments. Since the black hole mass distribution is similar in both environments, we speculate that, for low radio luminosities, the radio emission is controlled (through fuel disponibility or confinement of radio jet by local gas pressure) by the interstellar medium of the host galaxy, while in other cases it is determined by the structure (group or cluster) in which the galaxy resides.   } % conclusions heading (optional), leave it empty if necessary {} ",
        "introduction": "Radio sources  have been traditionally classified into two broad classes depending upon their emission mechanism: AGN-driven emission has been supposed for early type galaxies, while a star forming origin of the radio power has been invoked for late type galaxies \\citep{Condonreview}. However, not all optical galaxies exhibit radio emission (at least at the depth of  current surveys) and it is important to investigate  the physical mechanism which triggers this emission. For AGN-induced activity, the difference between radio quiet and radio loud galaxies is naturally  connected with the availability of fuel for the central engine, the black hole mass, and its emission efficiency \\citep[see e.g.][and references therein]{Shabala}, while for star forming objects it is closely related to the strength of the star formation episode.  Heating by a central AGN is also thought to be important for the host galaxy evolution, providing the energy to stop the central cooling that drives gas to the central black hole in an intermittent feedback mechanism dubbed \"radio mode\" \\citep[see][]{Croton,Ciotti}. The same intermittent mechanism is indirectly observed  as bubbles in the intra-cluster medium \\citep{Birzan2004} or in the restarted activity of radio sources \\citep{Venturi2004,Bardelli2002}. This heating emission is thought also to be able to suppress the star formation and  to be an important mechanism for creating the galaxy color bimodality \\citep{Croton}. Moreover, it is a fact that many radio AGN reside in early type galaxies \\citep{LO96}, a class of objects which are preferentially found in high density environments.   It has been proposed that cluster/group mergers, galaxy-galaxy mergings or at least tidal disturbance caused by close encounters between galaxies would increase radio activity. For early type galaxies, these phenomena could drive gas more efficently toward the center, while for late type galaxies, these encounters would trigger and/or enhance star formation \\citep[see e.g.][]{Bekki,Vollmer}. The radio luminosity of the AGN also depends on the interaction of the jet coming from the black hole with the surrounding gas of the host galaxy and/or of the hosting cluster \\citep[see e.g.][and references therein]{Shabala}. For these reasons, some degree of  dependence of radio activity on the environment is expected. Although the problem is clear from the qualitative description, controversial results are present in the literature and for mainly two reasons.  First of all, very few papers are primarly based on radio data. Usually, papers in the literature consider AGN selected on the basis of  their optical (mostly estimated from line diagnostics) or X-ray activity \\citep[see][and references therein]{Kauffmann2004}. It seems that the family of optical/X-ray AGN inhabits similar environments to non-active galaxies, with a preference for lower densities at higher stellar masses \\citep{Kauffmann2004,Silverman}. However, \\cite{Best05} claim that the phenomena generating the  emission lines and the radio emission are statistically independent, even if recent results indicate that radio loud emission line AGN favour higher densities than radio quiet ones \\citep{Kauffmann08} \\citep[see also ][for a study in the IR]{Caputi09}. Secondly, only specific environments have been considered, mainly clusters or groups \\citep{Miller,Hill, Giacintucci}, lacking therefore the coverage of a large dynamical range in densities.  The first paper that studied in a complete statistical way, both from the radio and optical side, the environment of radio galaxies was that of \\cite{Bestenv} . This paper uses data from  2dF  Galaxy Redshift Survey \\citep{Colless} and the NRAO VLA Sky Survey \\citep{NVSS} with a well defined estimator of density, i.e. the projected density of the 10$^{th}$ nearest neighbor. Moreover, groups (and clusters) are found with the standard friend-of-friends detection method. The redshift range is $0.02<z<0.1$. The general result is that radio-selected star forming galaxies decrease in number with increasing overdensity, while objects with AGN activity exhibit little dependence upon the local density, except at the very low densities, where the probability to find these objects decreases significantly. \\cite{Bestenv} claims also that the larger scale (i.e. on the scales typical of groups and clusters) is more important than the smaller scale (typical of galaxy pairs or companions) to determine the AGN radio emission. Star forming galaxies do not exhibit dependence on environment at scales larger than one Mpc \\citep{Kauffmann2004}.  In this paper we investigate the environment properties of the radio sources present in the zCOSMOS survey using a set of well defined density estimators in the redshift range [0.1-0.9]. In Section 2 we present the used data and the estimators for various environments; in Section 3 we describe the method adopted to separate galaxies with radio emission induced by AGN from that induced by star formation and  the definition  of the control samples. In Section 4 we discuss the differences between the various samples, while in Section 5 we present the dependences on environment. In Section 6 the luminosity distributions  and radio-optical ratio for a specific class of radio AGN are shown and in Section 7 we discuss the results. The adopted cosmology has $\\Omega_m=0.25$,$\\Omega_{\\Lambda}=0.75$ and $H_0=70$ km s$^{-1}$ Mpc$^{-1}$.   Throughout the paper we adopted the term \"radio galaxy\" for galaxies hosting a radio source, independently of the physical  nature of the emission. ",
        "conclusions": "Studying the environment of radio galaxies requires a division of the radio sources on the basis of their emission mechanism, because AGN and star forming galaxies may be influenced differently by the local density. To do this, we compared the star formation rate  estimated from the spectral energy distribution  with the one derived from the radio luminosity. We found a clear correlation between these two quantities for the star forming galaxies: we defined as AGN those objects for which the radio star formation rate is significantly higher than the optical one. In other words, these objects have a radio luminosity significantly higher than that predicted  by their star formation. We assume a  stringent cut at which the AGN radio emission is more than an order of magnitude higher than the star formation one in the same galaxy. In this way we excluded from our radio-based AGN sample those AGN defined on the basis of their optical emission lines, such as Seyfert 2 and LINERs, for which the ratio of AGN to star forming radio emission  is lower than the adopted threshold.   Since AGN are hosted primarly by early type optical galaxies and star forming objects are basically  spiral galaxies, we expect to recover at least the optical morphology-density relation \\citep{Dressler}, not related to the radio emission. Therefore, to investigate whether there is an environmental effect on the radio emission,  we  defined control samples to be compared with the AGN and star forming radio samples. To achieve this,  we studied  the positions of radio sources and  zCOSMOS galaxies in an infrared color-specific star formation plane. We defined  three populations: passive AGN, non-passive AGN and star forming galaxies. The passive and non-passive division is reminescent of the low and high-excitacion dichotomy described in \\cite{Smolcic09} and probably corresponds to two different processes of accretion (hot and cold, respectively). In fact, the two processes are expected to depend differently on the environment \\citep{Hardcastle}. The radio luminosity functions of the passive and non-passive AGN are similar in shape but not in luminosity range. Non-passive AGN stop at logL$_{radio}$(W Hz$^{-1}$)= 23.4, while the passive AGN reach logL$_{radio}$(W Hz$^{-1}$)= 25.2.  Finally, our radio objects belongs almost all to the FR-I class \\citep[as defined by][]{LO96} and the sample lack of the broad line AGN class. Therefore we focussed mainly to objects triggered by gas accretion and did not considered other mechanism as e.g. mergers.  {\\it Passive AGN:} This population of AGN is hosted by red galaxies. Using the morphological classification derived from HST/ACS images through the ZEST algorithm \\citep{Scarlata,Tasca}, we find that $\\sim 85 \\%$ of both passive AGN and relative control sample are classified as early type (i.e. bulge dominated). The spectra of these galaxies with and without radio emission are indistinguishable and the ratio of their stellar mass functions is rather constant at $\\sim 0.2$ for masses higher than $5 \\times 10^{10} $ M$_{\\odot}$. This value can be interpreted as corresponding  to the average fraction of time during which the central engine is switched on. Note that our ratio is derived comparing the mass functions: if we had used  the observed number of galaxies, we would have found values more similar to those reported in Table 4 of \\cite{Shabala} ( $\\sim 0.5 $ at high masses, $\\sim 0.15$ at intermediate masses and $\\sim 0.01$ at low masses).  This population exhibits a strong dependence on the environment: the fraction of objects with radio emission increases from $\\sim 0.02$ for underdense regions ($1+\\delta<3$) to $\\sim 0.15$ for overdense regions  ($1+\\delta>10$). Similarly, such trend is  also found (but at a low significance level) as a function of  richness of  galaxy groups.  Studying the radio luminosity-stellar mass plot, we found that for luminosities lower than log P(W Hz$^{-1}$)=23.5 there is a correlation between the two quantities, that disappears at higher powers. The correlation between stellar mass and radio luminosity is not unexpected considering the relation between the accretion on to a black hole in the \"radio mode\" to the black hole mass and hot gas fraction \\citep[see equation 10 of][]{Croton}. As stated by \\cite{Croton}, the hot gas fraction is approximately constant for galaxies with $v_{vir}>150$ km s$^{-1}$  and therefore the black hole luminosity scales with the black hole mass, which correlates with the stellar mass. Interestingly, half of the objects in the low radio luminosity regime reside in low densities, while the others are in overdense regions.  The correlation disappears for radio luminosities higher than log L$_{Radio}$(W Hz$^{-1}$)=23.5. Almost all of these radio luminous objects reside in overdense region. On the other hand, no difference is present in the black hole mass between high and low densities and therefore it seems that in richer environments the emission mechanism is more efficient than in low density regions.  Therefore,  we can tentatively say that at low density the emission is determined by the host galaxy stellar mass and at high density by the host structure (group/cluster) in which the galaxy resides. The higher luminosities in denser environments could be caused by a greater fuel supply (due to the central cooling of the gas)  and/or to the fact that also outside the galaxy there is a significant gas density to confine the radio jet. On the other hand, we already know that ellipticals at the center of galaxy groups are richer in hot gas than isolated and non-dominant ones \\citep{Helsdon} and this implies that a mechanism driving gas within dominant galaxies is or has been active.   This can also justify  a longer time spent by the central engine in the \"switched on\" status, as suggested by the increasing ratio of passive  AGN to control sample as a function of  density. It would be interesting to investigate whether the high density AGN have luminosities that correlate with the central cooling flow, following the relation between radio luminosity and accretion mass found by \\cite{Mittal}. Note that our findings are consistent with those of both \\cite{Bestenv}, \\cite{Mandelbaum}.  Finally, the difference in overdensity between radio-loud and control samples remains constant across  the studied redshift range, implying that no significant evolution in this phenomenon is present.   {\\it Non-passive AGN}: This is a less well-defined family. In principle, in this class one can find not only AGN but also objects such as post--starburst galaxies. In  this case, radio emission disappears on a time scale of $10^{8}$ years, which could be longer than the disappearing time of the  star formation signatures in optical. However, we expect that these objects are rare because of the implied short time scales. In fact, only two of the radio detected post--starbursts of the \\cite{Verganipsb} sample and one of the upper limits lie within our AGN region. This means that post-starburst galaxies contaminate our sample by $4 \\%$. The non-passive AGN population occupies a small region of the infrared color-specific star formation rate plot, which is approximately  equivalent to the green valley defined in the optical bands. These galaxies are objects with  spectrophotometric type earlier than Sa-Sb and applying the standard diagnostics based on the emission line ratios we found that $50\\% $ are classified as optical AGN and the remaining are star forming galaxies. The ZEST  morphologies of non-passive AGN are equally divided between early types and spirals (with a  $\\sim 8 \\% $ of irregulars), while the control sample has  percentages of $\\sim 60 \\%$ and $40 \\%$, respectively. The radio emitting objects of this class follow the same environment distribution of the corresponding control sample galaxies and do not show the luminosity-density behaviour.   {\\it Star Forming galaxies} The composite spectra of radio emitting star forming galaxies is redder than the control sample with a smaller [OII] and higher H$\\alpha$ equivalent width.   Considering that the radio detected star forming galaxies are objects where the star formation rate is on average higher than the undetected ones, this implies that the dust content is higher. This increase of dust extinction with increasing star formation is already  known \\citep[see e.g.][]{Pannella} and could justify the relation between the emission lines and the radio SFR estimators shown in the lower panel of Figure \\ref{sfrsfr}. The ZEST morphological classification shows that the radio detected objects tend to be of earlier type than the control sample, as found for the spectrophotometric types.  The radio emitting star forming galaxies show an apparent difference in density distribution with respect to the  control sample, but we have shown that this is due to the different stellar mass distributions of the two samples. Correcting the control sample with the fraction of radio emitting objects, no differences in environment are present between radio and control sample galaxies. This is consistent with the conclusion of \\cite{Kauffmann2004} that the star formation does not  depend on the density for scales larger that the Megaparsec."
    },
    "0911/0911.0670_arXiv.txt": {
        "abstract": "This paper presents a numerical study over a wide parameter space of the likelihood of the dynamical bar-mode instability in differentially rotating magnetized neutron stars. The innovative aspect of this study is the incorporation of magnetic fields in such a context, which have thus far been neglected in the purely hydrodynamical simulations available in the literature. The investigation uses the \\cosmospp code which allows us to perform three-dimensional simulations on a cylindrical grid at high resolution. A sample of Newtonian magneto-hydrodynamical simulations starting from a set of models previously analyzed by other authors without magnetic fields has been performed, providing estimates of the effects of magnetic fields on the dynamical bar-mode deformation of rotating neutron stars. Overall, our results suggest that the effect of magnetic fields is not likely to be very significant in realistic configurations. Only in the most extreme cases are the magnetic fields able to suppress growth of the bar mode. ",
        "introduction": "\\label{sec:intro}  Rotating neutron stars formed following the gravitational collapse of a massive stellar iron core or the accretion-induced collapse of a white dwarf can be subject to various nonaxisymmetric instabilities depending on the amount and degree of differential rotation. The prospects of detection of gravitational radiation from newly born rapidly rotating neutron stars by the current and future generations of kHz-frequency, ground-based gravitational wave interferometers highly motivate the investigation of such instabilities. In particular, if the rotation rate is high enough and shows a high degree of differentiation,  the star is subject to the so-called $m=2$, {\\it dynamical} bar-mode instability driven by hydrodynamics and gravity, with $m$ being the order of the azimuthal nonaxisymmetric fluid mode $e^{\\pm im\\varphi}$ . On the other hand, at lower rotation rates gravitational radiation and viscosity can drive a star {\\it secularly} unstable against bar-mode deformation. These two flavors of the bar-mode instability set in when the ratio $\\beta=T/|W|$ of rotational kinetic energy $T$ to gravitational potential energy $W$ exceeds a critical value $\\beta_{\\rm c}$.   Early studies with incompressible MacLaurin spheroids in Newtonian gravity showed that the onset of the instability arises when $\\beta_c\\sim 0.27$ and $0.14$ for the dynamical and secular cases, respectively~\\citep{chandra69}.  Improvements on these simplified analytic models have been achieved through numerical work. Newtonian and general relativistic analyses of the dynamical bar-mode instability  are available in the literature, using both simplified models based on equilibrium stellar configurations perturbed with suitable eigenfunctions, and more involved models for the core collapse scenario. Newtonian hydrodynamical simulations have shown that the value of $\\beta_{\\rm c}$ is quite independent of the stiffness of the equation of state, provided the star is not strongly differentially rotating (see~\\citet{houser94,new00,liu02} and references therein). On the other hand, relativistic simulations~\\citep{shibata00} have yielded a slightly smaller value of the dynamical instability parameter ($\\beta_{\\rm c}\\sim 0.24-0.25$). They have also shown that the dynamics of the process closely resembles that found in Newtonian theory, that is, unstable models with large enough $\\beta$ develop spiral arms following the formation of bars, ejecting mass and redistributing the angular momentum. Further relativistic simulations~\\citep{baiotti07} have shown the appearance of nonlinear mode-coupling which can limit, and even suppress, the persistence of the bar-mode deformation. It is also worth mentioning that, as the degree of differential rotation becomes higher and more extreme, Newtonian simulations~\\citep{shibata02,shibata03} have also shown that rotating stars are dynamically unstable against bar-mode deformation even for values of $\\beta$ of order 0.01.  Whether the requirements for the development of the instability inferred from numerical simulations are  met by the collapse progenitors remains unclear. Observations of surface velocities imply that a large fraction of progenitor cores are rapidly rotating. However, it has been shown that magnetic torques can spin down the core of the progenitor, leading to slowly rotating neutron stars at birth~\\citep{spruit98}. The most recent computations of the evolution of massive stars, which include angular momentum redistribution by magnetic torques and spin estimates of neutron stars at birth, lead to core collapse progenitors which do not seem to rotate fast enough to guarantee the unambiguous growth of the canonical bar-mode instability~\\citep{heger05,ott06}. These estimates are in agreement with observed periods of young neutron stars. However, rapidly-rotating cores might be produced by an appropriate mixture of high progenitor mass and low metallicity. Recent estimates suggest that about 1\\% of all stars with masses larger than 10 solar masses will produce rapidly rotating cores~\\citep{woosley06}. Recent relativistic simulations of a large sample of rotational core collapse models carried out by~\\citet{dimmelmeier08}, which include a state-of-the-art treatment of the relevant physics of the collapse phase and realistic precollapse rotational profiles~\\citep[see also][]{ott07}, have shown that the critical threshold of the dynamical bar-mode instability is never surpassed, at least early after core bounce. However, a large set of the models investigated by~\\citet{ott07} and~\\citet{dimmelmeier08} do show that models with sufficiently differential and rapid rotation are subject to the low$-\\beta$ instability. In addition it is also worth mentioning the simulations of~\\citet{baiotti08} which show that hypermassive differentially rotating neutron stars form following the merger of low-mass binary neutron stars, when modeled as polytropes. The resulting  hypermassive star undergoes a persistent phase of bar-mode oscillations, emitting large amounts of gravitational radiation prior to its delayed collapse to a black hole.   An important piece of physics that the existing numerical work has so far neglected is the presence of magnetic fields. The amplification and emergence of strong magnetic fields in neutron stars from initially weak magnetic field configurations in the pre-collapse stellar cores are currently under investigation~\\citep[see e.g.,][and references therein]{burrows07,cerda-duran08}. We note, however, that the weakest point of all existing magneto-rotational core collapse simulations to date is the fact that both the strength and distribution of the initial magnetic field in the core are unknown.  In this paper, we show results from a  detailed numerical study  of the effects that magnetic fields may have on the dynamical bar-mode instability in differentially rotating magnetized neutron stars. In particular, we investigate how sensitive the onset and development of the instability is to the presence of magnetic fields, as well as the role played by the magnetorotational instability (MRI) and magnetic braking mechanisms to alter the angular momentum distribution in the star and possibly suppress the bar-mode instability. Our study is motivated by the potential astrophysical implications that the presence of strong magnetic fields may have for post-bounce core collapse dynamics and, in turn, for gravitational wave astronomy.  The uncertainty of the strength and distribution of magnetic fields in collapse progenitors is reflected in our somewhat ad hoc parameterization of the field configuration through a large sample of equilibrium models of rapidly and highly differentially rotating neutron stars. Our sample of Newtonian magnetohydrodynamical (MHD) simulations is based upon the set of purely hydrodynamical models previously analyzed by~\\citet{new00}. The simulations are performed  using the covariant (and adaptive mesh refinement) code \\cosmospp~\\citep{anninos03a,anninos03b,anninos05,fragile05} which allows us to perform three-dimensional simulations on a logarithmically scaled cylindrical grid at high resolution. A number of equilibrium models of rapidly rotating stars with different values of the rotational instability parameter ($\\beta=T/|W|$) and magnetic plasma beta ($\\beta_B=P/P_B$), and different polytropic equations of state are constructed, introducing also different configurations for the magnetic field distribution (of both poloidal and toroidal varieties) and field strengths. The equilibrium models are perturbed by seeding small random perturbations in order to initiate the onset of the bar-mode deformation.  The organization of the paper is as follows. Section~\\ref{sec:methods} discusses our basic formalism, numerical methods, diagnostics, and the construction of initial data; Section~\\ref{sec:results} presents our results in two subsections, one for initially toroidal field configurations and one for poloidal. We conclude with a summary and discussion of our results in Section~\\ref{sec:conclusion}. ",
        "conclusions": "\\label{sec:conclusion}  We have studied the growth of the dynamical bar-mode instability in differentially rotating magnetized neutron stars through a set of numerical Newtonian MHD calculations. The calculations explored both toroidal and poloidal initial field distributions of differing strengths, as well as the role of the equation of state.  For our initially toroidal field configurations, field amplification always saturated at a level insufficient to strongly affect the dynamics of the bar mode. Even the most extreme case (TG53B1, $\\beta_\\mathrm{B,min}=1$), with equal initial magnetic and hydrodynamic pressures within the toroidal loop, gave only a $\\sim30$\\% change in gravitational wave properties. Otherwise evolution proceeded quite similarly to our unmagnetized reference cases.  The effects were larger for the poloidal configurations. In the most extreme case considered (model PG53B10HR with $\\beta_\\mathrm{B,min}=10$), the magnetic field was sufficient to completely suppress the formation of a bar. However, in that case, we started with the magnetic field already within a factor of $\\sim$30 of being dynamically more important than thermal pressure in some portions of the star. It then only took a few dynamical times of field growth for $\\beta_B$ to approach unity, a timescale that is much shorter than the bar deformation time for azimuthal Fourier modes to reach appreciable amplitudes. Also, due to the computationally demanding nature of evolving strong poloidal fields, we cannot verify that this result will not change with increased grid resolution. For less extreme initial conditions, the effects of including poloidal field components were consistently more modest with decreasing initial field amplitude. These lower amplitude calculations are also less demanding computationally and can, like the toroidal cases, be more easily checked for convergence. The principle effects of introducing poloidal fields are in systematically delaying the onset of the bar mode and in suppressing both Fourier mode and gravitational wave amplitudes, all of which become negligible with initial field amplitudes below the threshold of $1/\\beta_\\mathrm{B,min} \\lesssim 10^{-2}$.  Overall, our results suggest that the effect of magnetic fields on the emergence of the bar-mode instability in neutron stars is not likely to be very significant. Particularly considering that collapse progenitor models predict realistic field configurations that are dominantly toroidal in nature with toroidal and poloidal components of order $10^{10}$G and $10^6$G, respectively \\citep{heger05}, substantially below most of the field configurations we have considered. Thus, except in special cases where neutron stars are born very highly magnetized, we might still expect them to be good gravitational wave sources if their rotational kinetic energies exceed the critical bar-mode instability parameter."
    },
    "0911/0911.2079_arXiv.txt": {
        "abstract": " ",
        "introduction": "The INTEGRAL (INTErnational Gamma-Ray Astrophysics Laboratory) satellite produces many observations since its launch in 2002, not only in gamma part of the spectra. The onboard OMC (Optical Monitoring Camera) was designed to obtain the observations in optical $V$ passband. These observations are in fact only a by-product of the mission, but nowadays there are many observations available.  Despite the fact that the database of these measurements is freely available on internet, the analyses are still very rare. The most recent one using the OMC data is that by \\cite{2009IBVS.5881....1J} about a new $\\beta$~Cep star.  This investigation is directly following our previous papers (\\citealt{Zasche2008NewA} and \\citealt{2009NewA...14..129Z}). The selection criteria used here were also the same: maximum number of data points and non-existence of any detailed light-curve analysis of the particular system. There were 33 systems selected for the present paper. ",
        "conclusions": "The light-curve analyses of thirty-three selected systems have been carried out. Using the light curves observed by the Optical Monitoring Camera onboard the INTEGRAL satellite, one can estimate the basic physical parameters of these systems. Despite this fact, the parameters are still only the preliminary ones, affected by relatively large errors and some of the relevant parameters were fixed at their suggested values. The detailed analysis is still needed, especially spectroscopic one, or another more detailed light curve one in different filters. Together with a prospective radial-velocity study, the final picture of these systems could be done. Particularly, the systems V1450~Aql and CY~Lac seem to be the most interesting ones. The first one is massive semi-detached system, which shows total eclipses and the second one due to its eccentric orbit."
    },
    "0911/0911.0760_arXiv.txt": {
        "abstract": "{} {We aim at determining the rotational periods and the starspot properties in very young low-mass stars belonging to the Ori OB1c star forming region, contributing to the study of the angular momentum and magnetic activity evolution in these objects. } {We performed an intensive photometric monitoring of the PMS stars falling in a field of about 10$\\arcmin\\times 10\\arcmin$ in the vicinity of the Orion Nebula Cluster (ONC), also containing the BD eclipsing system 2MASS\\,J05352184-0546085. Photometric data were collected between November 2006 and January 2007 with the REM telescope in the $VRIJHK'$ bands. The largest number of observations is in the $I$ band (about 2700 images) and in $J$ and $H$ bands (about 500 images in each filter). From the observed rotational modulation, induced by the presence of surface inhomogeneities, we derived the rotation periods. The long time-baseline (nearly three months) allowed us to detect rotation periods, also for the slowest rotators, with sufficient accuracy ($\\Delta P/P<2\\%$). The analysis of the spectral energy distributions and, for some stars, of high-resolution spectra provided us with the main stellar parameters (mass, luminosity, effective temperature, mass, age, and $v\\sin i$) which are essential for the discussion of our results. Moreover, the simultaneous observations in six bands, spanning from optical to near-infrared wavelengths, enabled us to derive the starspot properties for these very young low-mass stars.} {In total, we were able to determine the rotation periods for 29 stars, spanning from about 0.6 to 20 days. Thanks to the  relatively long time-baseline of our photometry, we derived periods for 16 stars and improved previous determinations for the other 13. We also report the serendipitous detection of two strong flares in two of these objects. In most cases, the light-curve amplitudes decrease progressively from the $R$ to $H$ band as expected for cool starspots, while in a few cases, they can only be modelled by the presence of hot spots, presumably ascribable to magnetospheric accretion. The application of our own spot model to the simultaneous light curves in different bands allowed us to deduce the spot parameters and particularly to disentangle the spot temperature and size effects on the observed light curves. } {} ",
        "introduction": "\\label{sec:Intro}  The study of the evolution of angular momentum of low-mass stars in the pre-main sequence (PMS) phase has gained great advantage from high-precision photometry of star forming regions (SFR). While during the main sequence (MS) evolutionary phase the total angular momentum of a solar-mass star is lost efficiently via the magnetic-wind breaking process \\citep[e.g.,][]{Chabo95}, in the PMS phase the presence of a circumstellar disc coupled to the star through magnetic fields permits a more efficient angular momentum transfer, thus preventing the spin-up driven by the rapid mass-accretion and contraction of the star \\citep{Shu94}.  An accurate knowledge of the rotation period distribution of the solar and low-mass stars in different SFR environments is of paramount importance for understanding how the disc-locking mechanism is effective in modifying the angular momentum in the PMS stage, and for testing the reliability of the theoretical models of PMS star interiors and their evolutionary tracks.  Measurements of periodic light-variations with typical amplitudes in the range 0.05--0.40 mag in the $V$ and $R$ bands, due to uneven starspot distributions on the stellar photospheres, are a direct way to obtain this physical parameter. There has been a rapid increase in the number of PMS objects with measured rotation periods over the last decade \\citep[see, e.g.,][and references therein]{Mat04,Herbst07,Irwin09}. A broad, bimodal distribution exists at the earliest observable phases (about 1~Myr) for stars more massive than 0.4~M$_{\\sun}$. The fast rotators (50--60\\% of the sample) evolve to the ZAMS with little or no angular momentum loss. The slow rotators continue to lose substantial amounts of angular momentum for up to 5 Myr, creating the even broader bimodal distribution characteristic of 30--120 Myr old clusters. Accretion disc signatures, such as H$\\alpha$ emission and infrared (IR) excess, are observed more frequently among slowly rotating PMS stars, suggesting a connection between accretion and rotation \\citep{rebu06}. According to the present picture, discs appear to influence rotation only for a small percentage of the solar-type stars and only during the first 5 Myr of their life. This time interval is comparable to the maximum life-time of accretion discs derived from NIR studies and may be a useful upper limit to the time available for forming giant planets. For the sparse population of young stars (mainly weak-T Tauri stars, wTTS, with ages between 5 and 30 Myr) in the Orion complex, \\citet{Mari07} found no evidence of bimodality for masses higher than 0.7~M$_{\\sun}$, and a median rotation period of $\\sim$\\,1.5 days, suggesting that the spin-up process at this stage has become dominant over the disc-locking effect. It is then mandatory to explore the rotational properties of as many clusters/associations as possible with ages less than 5 Myr. Indeed, in this age range, the different scenarios for angular momentum evolution proposed by \\citet{lamm} and related to different time-scales of disc-locking and stellar contraction \\citep{Hart02} are expected to produce their stronger effects.  There is growing evidence that the same mechanisms operating in solar-mass stars also regulate the rotation of very-low mass (VLM) stars and brown dwarfs (BDs). Indeed, very recently many BDs showing accretion discs with physical properties similar to their analogs around higher-mass stars have been discovered \\citep[e.g.,][and reference therein]{Luhman05,Alc06,Muz06,Mer07}. However, it seems that both disc interactions and stellar winds are less efficient in braking these objects, although the statistics available for this VLM regime is rather poor.  In some cases, both VLM and BD objects in star forming regions exhibit photometric light curves with high amplitudes and irregular modulations \\citep{Scholz04,Scholz05} that are usually explained by hot spots caused by mass-accretion flow from the disc \\citep{Herbst94}. The same authors remark that wTTS and some classical T Tau stars (cTTS) display instead more regular light curves with smaller amplitudes ascribable to cool photospheric spots that can survive for several stellar rotations.  In this paper, we report the results of an intensive photometric monitoring, in near-infrared (NIR) and optical bands, of a sample of young low-mass stellar and sub-stellar objects (YSOs) in a small area flanking the Orion Nebula Cluster (ONC). The main goal of our observations was the study of the NIR light curve of \\object{2MASS\\,J05352184-0546085}, an eclipsing system composed of two BD components recently discovered by \\citet{stass06}. The long-lasting and intensive simultaneous optical and NIR monitoring also allowed us to study the photometric variability of many YSOs in the field of view around the eclipsing BD. Therefore, in this study we report the characterization of these objects and determine their rotational periods, as well as the properties of their starspots. The present work thus contributes to the study of the early phases of angular momentum evolution and of the behaviour of magnetic activity in young, VLM objects which still needs observational and interpretation efforts.  We derived the rotation periods of our targets by means of time-series analysis of our large photometric data-set gathered in the $VRIJHK'$ bands. The physical characterization (radius, luminosity, effective temperature, mass, and age) of the objects, which is essential for any discussion on angular momentum evolution and stellar activity, is based on our photometric data as well as on literature mid-IR data from the Spitzer space telescope \\citep{rebu06} through an analysis of their spectral energy distribution (SED). For some objects, accurate physical parameters and the projected rotational velocity, $v\\sin i$, were determined by means of high-resolution spectra.  The use of simultaneous multi-band light curves allows an accurate analysis of the spot properties because the different amplitude of the light curves constrains the spot temperatures and filling factors. ",
        "conclusions": "\\label{sec:conclusions}  We presented the results of an intensive photometric monitoring of a 10$\\arcmin\\times 10\\arcmin$ field flanking the Orion Nebula Cluster (ONC) conducted during three consecutive months. The main results of our work can be summarized as follows: \\begin{itemize} \\item We detected rotation periods for 29 stars, spanning from about 0.6 to 20 days, sixteen of which are new periodic variables. Thanks to the relatively long time-baseline we measured the periods with sufficient accuracy ($\\Delta P/P<2\\%$) also for the slowest rotators. We remark that none of the stars with NIR excess has a rotation period shorter than 5\\,days, in agreement with the hypothesis that the angular momentum of the stars is moderated by interaction with the disc.  \\item The analysis of the spectral energy distribution and, for some stars, the high-resolution spectra provided us with $T_{\\rm eff}$ and luminosity, and that allowed us to construct the HR diagram of these stars. We could then assign a photometric membership to the objects by a comparison with PMS evolutionary tracks and derive masses and ages. The majority of the objects with detected period are classified as Orion members and result to be younger than 10\\,Myr.  \\item Our spot modelling code enabled us to derive the starspot properties for five of these star, based on the simultaneous analysis of the light curves in several optical and NIR bands. For one of these, the light curves could only be modelled with hot spots, which are likely related to magnetospheric accretion. For three stars with $T_{\\rm eff}$ in the range 4100--4400\\,K and rotation periods between 4 and 8 days we found cool spots with $\\Delta T=T_{\\rm ph}-T_{\\rm sp}$ in the range 1000--1300\\,K, i.e. larger, on average, than $\\Delta T$ typically found for the sub-giant and giant components of RS~CVn systems. Conversely, the spotted area (8--10\\% of the star surface) is smaller. These differences could be due to the fully convective internal structure of PMS stars which gives rise to a different dynamo action. For the cool ($T_{\\rm eff}\\simeq 3500$\\,K) and ultra-fast ($P\\simeq 0.56$\\,days) star \\#9, a smaller spot contrast ($\\Delta T\\simeq 270$\\,K) is found. This could be due both to the lower effective temperature and the high rotation rate.  For the star with the highest modulation amplitudes (star \\#11), which also displays a remarkable near- and mid-IR excess, the light curves were not consistent neither with hot nor with cool spots. We suggest that the flux modulation could be produced by inhomogeneities of the circumstellar disk.  \\item A very strong flare was detected on star \\#19 (V498\\,Ori) in the $I$ band. The temporal evolution of the event was fully resolved and we evaluated the rise and decay e-folding times as $\\tau_{r}\\simeq 2.5$\\,min and $\\tau_{d}\\simeq 24.9$\\, min, respectively. We estimated an energy released in the $I$ band of nearly $3\\cdot10^{35}$ erg and a 20 times higher energy released in the optical continuum. Another strong flare, which released an energy of about $6\\cdot10^{34}$\\,erg in the $I$ band, was observed on star\\,\\#\\,11 (\\object{NR~Ori}), which is cooler and less massive than star \\#\\,19 and displays the broadest modulation and a strong infrared excess.   \\end{itemize}"
    },
    "0911/0911.5003_arXiv.txt": {
        "abstract": "The final composition of giant planets formed as a result of gravitational instability in the disk gas depends on their ability to capture solid material (planetesimals) during their 'pre-collapse' stage, when they are extended and cold, and contracting quasi-statically. The duration of the pre-collapse stage is inversely proportional  roughly to the square of the  planetary mass, so massive protoplanets have shorter pre-collapse timescales and therefore limited opportunity for planetesimal capture. The available accretion time for protoplanets with masses of 3, 5, 7, and 10 Jupiter masses is found to be 7.82$\\times 10^4$,  2.62$\\times 10^4$, 1.17$\\times 10^4$ and  5.67$\\times 10^3$ years, respectively. The total mass that can be captured by the protoplanets depends on the planetary mass, planetesimal size, the radial distance of the protoplanet from the parent star, and the local solid surface density. We consider three radial distances, 24, 38, and 68 AU, similar to the radial distances of the planets in the system HR 8799, and estimate the mass of heavy elements that can be accreted. We find that for the planetary masses usually adopted for the HR 8799 system, the amount of  heavy elements accreted by the planets is small, leaving them with nearly stellar compositions. ",
        "introduction": "The recent discoveries of wide-orbit massive protoplanets by direct imaging (Kalas et al., 2008; Marois et al., 2008) have presented new type of gaseous planets that giant planet formation theories must be able to explain. Work regarding possible formation scenarios has already been demonstrated by several groups (e.g., Dodson-Robinson et al., 2009b; Nero and Bjorkman, 2009). Systems with massive gaseous planets at large radial distances offer a challenge for theorists mainly because they cannot be explained easily by the core accretion model, the standard model for giant planet formation (e.g., Cameron, 1978; Pollack et al., 1996). The core accretion model fails to form giant planets at large radial distances in situ due to the low surface density and the extremely long accretion timescale (Ida and Lin, 2004; Dodson-Robinson et al., 2009b). As a result, a more detailed investigation of giant planet formation scenarios seems to be required. \\par  In addition, the increasing number of transiting planets provides information regarding the planetary mean densities, and therefore, their compositions. Giant planet formation theories are also required to explain the planets' observed properties, and when possible, provide predictions that can be tested with future observations and upcoming data. Since the heavy element masses of planets can be estimated for transiting planets, further investigation and more detailed understanding of the origin of the heavy elements and the cause for the large variety in heavy elements enrichments observed (e.g., Guillot et al., 2007) is desirable. \\par  While the core accretion model offers an enrichment with heavy elements that is coupled to the formation process (e.g., Pollack et al., 1996; Hubickyj et al., 2005), in the gravitational instability model (Cameron, 1978; Boss, 1997) the planet's enrichment with solids is so far considered as an independent part of the actual formation mechanism. In the disk instability model giant protoplanets are initially formed with stellar compositions, and can be enriched with heavy elements after their formation by accretion of planetesimals (e.g., Helled et al., 2006). The total mass of solids that can be accreted depends on the planetary mass, its radial distance, the available solid material for capture, and the protoplanet's efficiency in accreting the solid planetesimals (Helled and Schubert, 2009). \\par  A gravitationally unstable condensation of a few Jupiter masses in a protoplanetary disk evolves through three phases (DeCampli and Cameron, 1979; Bodenheimer et al., 1980). First, it contracts quasi-statically with cool internal temperatures on the order of a few hundred K, with hydrogen in molecular form,  and with radii a few thousand times the present Jupiter radius. Second, once the central temperature reaches about 2000 K, the H$_2$ dissociates and initiates a dynamical collapse of the entire planet, ending only when the radius has decreased to a few times the present Jupiter radius. Third, it contracts and cools on the long time scale of $\\sim 10^9$ yr. The first phase is the crucial one for the   capture of solids, and the time scale of this phase, which is about $4 \\times 10^5$ yr for 1 Jupiter mass, determines the amount of material that can be captured. As the protoplanets contract quasi-statically during this phase, planetesimals in their feeding zone can be slowed down by gas drag, and get absorbed in the planetary envelopes (e.g., Helled and Schubert, 2009). As a protoplanet contracts, fewer planetesimals pass through its envelope, and instead of being accreted they get ejected from the protoplanet's vicinity. Efficient planetesimal accretion ends once the central temperature reaches $\\sim$ 2000 K and  molecular hydrogen begins to dissociate.  Whether gaseous protoplanets can form as a result of a local gravitational instability in a protoplanetary disk is still a matter of debate. The majority of theoretical and numerical work suggests that the gravitational instability scenario for planet formation is rather unlikely, at least at relatively small radial distances (e.g., Rafikov, 2007; Cai et al., 2009). Other results suggest that protoplanetary disks can break into gaseous protoplanets (e.g., Boss, 1997), although this may be limited to special conditions (Mayer et al., 2007). The extrasolar giant planets observed by direct imaging (Kalas et al., 2008; Marois et al., 2008) have raised the suggestions that such massive gaseous planets at large radial distances were formed as a result of gravitational instabilities (Dodson-Robinson et al., 2009b; Nero and Bjorkman, 2009). The fact that HR 8799 host star is found to be metal-poor ([Fe/H]=-0.47) supports the idea that its planetary system was formed via disk instabilities. More advanced numerical simulations with longer integration times and further theoretical work seem to be required before a robust conclusion regarding this formation mechanism and its limitations can be made.  While recent theoretical work (Boley, 2009; Rafikov, 2009; Nero and Bjorkman, 2009) suggests that gravitational instability can form planets only at distances greater than 50-100 AU, it is possible that these planets can migrate inward as a result of interactions with the disk, after formation. Thus in this paper we use the working hypothesis that massive gaseous planets presently at radial distances $>$ 20 AU could have formed  as a result of gravitational instability. Below we investigate the process of heavy element enrichment via planetesimal capture for the planetary system HR 8799. ",
        "conclusions": "We investigate the possibility of heavy element enrichment for the planetary system HR 8799. We consider planetary masses of 3, 5, 7 and 10 M$_J$, with radial distance of 24, 38, and 68 AU. We find that planetesimal accretion is relatively inefficient for such a system due to main two reasons: First, massive protoplanets have relatively short precollapse stages, and therefore, have limited time for planetesimal capture. Second, the accretion rate is low at large radial distances due to its dependence on the solid surface density and orbital frequency (both decreasing with radial distance). Farther from the star planetesimal velocities are lower, but the accretion time-scale is substantially longer. \\par  Smaller planetesimals can be accreted more efficiently, and if planetesimal sizes are of the order of $\\sim$ 1 km, solid material between a few and tens of M$_{\\oplus}$ can be captured. However, if planetesimals are formed with initial sizes of $\\sim$100 km or larger, the accreted mass is negligible. We therefore conclude that under such conditions wide-orbit gaseous planets, if formed via gravitational instability,  will have nearly stellar compositions. \\par  The capture cross-section used in this work was calculated taking the gravitational enhancement factor $F_g$ which corrects for the three-body effects (see Greenzweig and Lissauer 1990 for details) as unity. As a result, the accreted mass presented here should be taken as a lower bound. The total mass of solids in the protoplanet's feeding zone, which represents the maximum solid mass available for accretion, can be up to 2-3 M$_J$,  increasing with radial distance. However, it is unlikely that this available mass can actually be accreted due to the low accretion rate (see Helled and Schubert, 2009 for details) unless it is found that a significant gravitational focusing occurs. Detailed calculations of the gravitational enhancement factor for extended massive protoplanets with gas drag included seem to be required before a robust estimate for the upper bound of the accreted mass can be made. \\par  Our results suggest that massive protoplanets at wide orbits will have metallicity similar to their parent star. This conclusion does not depend on the stellar metallicity. Naturally, protoplanets around metal-rich stars will contain more high-Z material in their interiors due to the larger fraction of heavy elements (higher dust-to-gas ratio), but as we show here, an increase in the available solid material typically does not lead to further enrichment due to the low accretion rate. Protoplanets around stars with lower metallicity than considered here, will have a smaller fraction of heavy elements due to their initial composition which is metal-poor, and the fact that planetesimal accretion is inefficient as a result of the long accretion timescale and the low solid surface density. Our suggestion that massive gaseous planets at large radial distances will have stellar compositions is therefore independent of stellar metallicity. \\par  Our conclusion that massive protoplanets ($\\ge$ 3 M$_J$) will have relatively low enrichments in heavy elements is valid under the assumption that the high-Z material is captured during the pre-collapse stage in the form of solid planetesimals. It is possible however, that another enrichment mechanism for massive gaseous planets exists. For example, it may be possible that gaseous protoplanets are enriched with heavy elements from 'birth'. Disk fragments form at the corotation of dense spiral waves (Durisen et al. 2008) which can be enhanced with heavier elements (e.g., Haghighipour and Boss 2003, Rice et al. 2004) leading to formation of enriched fragments. In that case, solid concentrations are achieved at the locations where  fragmentation is most likely to occur, and the formed protoplanets will be metal-rich relatively to their host star. In addition, simulation of solids in disks show that regions in the disk can become highly enriched with solid material as a result of vortices in the disk (e.g., Klahr and Bodenheimer, 2006), although such vortices are likely to be destroyed in self-gravitating disks (Mamatsashvili and Rice, 2009). In any case, forming gaseous massive protoplanets which are enriched with heavy elements from birth is possible, and as a result massive gaseous protoplanets at large radial distances could contain a significant amount of heavy elements. It would be desirable to investigate whether enrichment from birth can be distinguished from enrichment at later stages (planetary evolution, internal structure, etc.), and we hope to address this issue in future work. An even more important issue that needs to be addressed to make future progress on this problem is a detailed study of planetesimal formation and collisional evolution in the 50--100 AU region of a gravitationally unstable disk. \\par  Given the assumption that the planets around HR 8799 formed from near-stellar composition, a question arises as to the observable metallicity in the present atmospheres. Even if there is added planetesimal material it tends to be deposited deep inside the planet, and a substantial fraction of it would condense out and form a solid core (Helled et al., 2008). Also, to a lesser extent, the grains in the original composition of the planet would tend to settle out, and at least the rock component could  be added to the core. The heavy element material not in the core would have been mixed by convection before the present time. These considerations suggest a slight depletion in heavy elements in the envelope of the planet as compared with the star. However numerous additional processes, such as cloud formation and possible core erosion, affect the atmosphere between the early phase studied here and the present time. Thus it is difficult to make precise predictions. In the limiting case where the planet has been completely mixed and has uniform composition at the present time, a typical case (38 AU, 10 km planetesimals, 7 M$_J$), the added solid mass is only 2.9 M$_\\oplus$ compared to a total mass of 2226 M$_\\oplus$. Thus even with some enhancement as discussed earlier, the difference between planetary and stellar metallicity would be negligible. Even in the case of the most-enriched planet, the fraction of solid material added is only 4\\% by mass."
    },
    "0911/0911.2623_arXiv.txt": {
        "abstract": "{The observed spectral variation of HD\\,50138 has led different authors to classify it in a very wide range of spectral types and luminosity classes (from B5 to A0 and III to Ia) and at different evolutionary stages as either HAeBe star or classical Be.  } {Based on new high-resolution optical spectroscopic data from 1999 and 2007 associated to a photometric analysis, the aim of this work is to provide a deep spectroscopic description and a new set of parameters for this unclassified southern B[e] star and its interstellar extinction.     } {From our high-resolution optical spectroscopic data separated by 8 years, we perform a detailed spectral description, presenting the variations seen and discussing their possible origin. We derive the interstellar extinction to HD\\,50138 by taking the influences of the circumstellar matter in the form of dust and an ionized disk into account. Based on photometric data from the literature and the new Hipparcos distance, we obtain a revised set of parameters for HD\\,50138.} {Because of the spectral changes, we tentatively suggest that a new shell phase could have taken place prior to our observations in 2007. We find a color excess value of E(B-V) = 0.08\\,mag, and from the photometric analysis, we suggest that HD\\,50138 is a B6-7 III-V star. A discussion of the different evolutionary scenarios is also provided.    } {} ",
        "introduction": "Stars that present the B[e] phenomenon are known to form a heterogeneous group. This group is composed of objects at different evolutionary stages, such as high- and low-mass evolved stars, intermediate-mass pre-main sequence stars, and symbiotic objects (Lamers et al. \\cite{Lamers}). However, for more than 50\\% of the confirmed galactic B[e] stars, the evolutionary stage is still unknown, so that they are gathered in the group of the unclassified B[e] stars. This problem is mainly caused by poor knowledge of their physical parameters, especially their distances.  In this paper, we present our study related to the southern B[e] star \\object{HD\\,50138} (V743 Mon, MWC158, IRAS 06491-0654). The spectrum of this star was discussed for the first time by Merrill et al. (\\cite{Merrill1}). Later, the presence of spectral variability was cited by Merrill (\\cite{Merrill2}) and Merrill \\& Burwell (\\cite{Merrill3}). Since then, numerous papers have mainly considered this star as either a pre-main sequence star, more specifically as a Herbig Ae/Be star, or as a classical Be star. Jaschek et al. (\\cite{Jaschek2}) have even considered this star as a transition object between a classical Be and a B[e] star. Later Lamers et al. (\\cite{Lamers}) and Zorec (\\cite{Zorec}) include this star in the list of unclassified B[e] stars.  The difficulty in obtaining the correct classification of HD\\,50138 is mainly caused by the strong spectral variability as reported in the literature. Pogodin (\\cite{Pogodin}) cites the presence of different scales of spectral variability, from days to months, especially in the line profiles of H$\\alpha$, He{\\sc i} ($\\lambda$ 5876), and Na{\\sc i} lines. Spectral variations have also been identified in the UV, by the analysis of IUE spectra (Hutsem\\'ekers \\cite{Hutsemekers}). These spectral variabilities have been explained by an outburst that possibly happened in 1978-1979 (Hutsem\\'ekers \\cite{Hutsemekers}) and by a shell phase in 1990-1991 (Andrillat \\& Houziaux \\cite{Andrillat}). On the other hand, HD\\,50138 had presented small and non-periodic photometric variations that did not seem to be associated to the spectral changes related to this shell phase (Halbedel \\cite{Halbedel}).  In addition, studies based on polarimetry and spectro-polarimetry have identified an intrinsic polarization that seems to be linked to nonspherical symmetry of matter around this object, probably a circumstellar disk (Vaidya et al. \\cite{Vaidya}; Bjorkman et al. \\cite{Bjorkman}; Oudmaijer \\& Drew \\cite{Oudmaijer}; Harrington \\& Kuhn \\cite{Harrington}). A disk-scattering effect associated to an outflow has also recently been suggested by Harrington \\& Kuhn (\\cite{Harrington2}).  On the other hand, Cidale et al. (\\cite{Cidale}) propose that this object could be a binary system. Based on spectro-astrometry, Baines et al. (\\cite{Baines}) suggest the same possibility, where the companion would be separated by 0.5$\\arcsec$ - 3.0$\\arcsec$. However, up to now, no direct evidence of binarity has been found.  HD\\,50138 was further observed during the Hipparcos mission, and its distance of 290$\\pm$71\\,pc (Perryman et al. \\cite{Perryman}) turned out with 500$\\pm$150\\,pc to be almost doubled according to the new data reduction procedure performed by van Leeuwen (\\cite{vanLeeuwen}). The newly determined distance and the large uncertainties in the previously published stellar classification attempt request and warrant detailed investigation and revision of the stellar parameters of HD\\,50138.  In this study, we present our new optical spectroscopic observations and perform a photometric and spectroscopic analysis of HD\\,50138, aimed on the one hand at describing the observed spectral variations and, on the other, at better constraining the stellar parameters, needed for an improved discussion about the nature of this object.  The paper has the following structure. In Sect.\\,\\ref{obs} we describe our observations. In Sect.\\,\\ref{results}, we present our results. In Sect.\\,\\ref{spectral_descr}, we describe our high-resolution spectra taken on different dates. In Sect.\\,\\ref{SpType}, we derive from the analysis of our spectroscopic data and public photometric measurements, the insterstellar and circumstellar extinction, hence the stellar parameters and the spectral type of this object. In Sect.\\,\\ref{discussion}, we discuss the possible scenarios for explaining the nature of this curious star, and finally in Sect.\\,\\ref{conclusion}, we present our conclusions. ",
        "conclusions": "HD\\,50138 is a very curious star that displays strong spectroscopic variations. Based on analysis of new high-resolution data, we present a detailed description of these variations. Our analysis of the photometric data suggests that HD\\,50138 is a B6-7 III-V star, whose luminosity was tentatively obtained from a careful study of the influence of the possible circumstellar extinction sources and based on the new Hipparcos distance. A new value for the color excess, $E(B-V) = 0.08\\pm 0.01$\\,mag, was derived. In addition, we suggest that a new-shell phase or the formation of one-armed spiral could have taken place before 2007.  Based on our results, a pre-main sequence star or a transition object between Be and B[e] stars, close to or just at the turn-off from the main sequence, or a binary scenario, can be neither confirmed nor discarded; however, an observational campaign, based on photometry, to derive a detailed light-curve and high-resolution spectroscopy associated to a detailed analysis in terms of the line profile appearances and variations have to be performed to confirm the possible shell phases and the slow-down of this star. A careful interferometric analysis, associated to a SED modeling considering different scenarios for the circumstellar dust is in progress (Borges Fernandes et al., in preparation) and will certainly provide better constraints for the circumstellar geometry and the nature of this curious star."
    },
    "0911/0911.4233_arXiv.txt": {
        "abstract": "{ We present the results of soft X-ray studies of the classical nova V2491 Cygni using the Suzaku observatory. On day 29 after outburst, a soft X-ray component with a peak at $\\sim$0.5~keV has appeared, which is tantalising evidence for the beginning of the super-soft X-ray emission phase. We show that an absorbed blackbody model can describe the observed spectra, yielding a temperature of 57~eV, neutral hydrogen column density of 2$\\times$10$^{21}$~cm$^{-2}$, and a bolometric luminosity of $\\sim$10$^{36}$~erg~s$^{-1}$. However, at the same time, we also found a good fit with an absorbed thin-thermal plasma model, yielding a temperature of 0.1~keV, neutral hydrogen column density of 4$\\times$10$^{21}$~cm$^{-2}$, and a volume emission measure of $\\sim$10$^{58}$~cm$^{-3}$. Owing to low spectral resolution and low signal-to-noise ratio below 0.6~keV, the statistical parameter uncertainties are large, but the ambiguity of the two very different models demonstrates that the systematic errors are the main point of concern. The thin-thermal plasma model implies that the soft emission originates from optically thin ejecta, while the blackbody model suggests that we are seeing optically thick emission from the white dwarf. } ",
        "introduction": "Classical novae are a class of cataclysmic variables, which occur in accreting binaries with a white dwarf and a late-type companion. Hydrogen-rich material is accumulated on the white dwarf surface. The surface temperature increases as the material accumulates. When the accreted material reaches a critical mass for nuclear fusion, an thermonuclear runaway process ignites on the white dwarf surface. A review of the evolution of classical novae can be found in e.g., \\cite{starrfield2008}.  After outburst, the system is expected to display X-ray emission via different mechanisms: (1) Non-thermal emission has recently been reported from the classical nova V2491 Cygni by the Suzaku satellite \\citep{takei2009a}. (2) Thin-thermal plasma emission from adiabatic shocks has been reported in some classical novae, mainly in data taken with the Swift satellite (e.g., \\citealt{bode2006,ness2009a}). (3) Photospheric emission from a hot layer of a white dwarf surface, which is characterised by bright, soft blackbody-like continuum emission and absorption lines, similar to the class of Super-Soft X-ray Sources (SSS; \\citealt{kahabka1997}).  Following from various arguments, the time scale of SSS emission is a function of the white dwarf mass and the chemical composition of the post-outburst envelope (e.g., \\citealt{sala2005,hachisu2006}). The optically thick wind theory is an approach to model SSS light curves with some simple assumptions (e.g., \\citealt{hachisu2005,hachisu2006}). In this theory, the onset of the SSS phase is assumed to occur after the wind stops, and the SSS phase fades out once the nuclear fuel on the white dwarf surface is consumed. Meanwhile, recent observations have shown that significant wind velocities can still be detected during the SSS phase (see contribution by Ness in this same journal). It is difficult to estimate the effects of the expansion on model light curves, and, ultimately, the white dwarf mass that would be derived from models that account for the continued expansion during the SSS phase. More advanced models are clearly needed, nevertheless, the duration of the SSS phase is an important observational quantity that will always be needed to determine the white dwarf mass.  In order to determine the duration of the SSS phase, we have to estimate the turn-on and turn-off time of SSS emission. For both quantities, dense observations in time are needed, and intensive monitoring campaigns are currently being performed with Swift (see, e.g., the contributions by Beardmore and by Osborne in this same journal). However, the turn-on time is particularly difficult to be estimated because we need to discriminate SSS emission from other sources of X-ray emission. For example, in RS\\,Oph, significant shock emission was present during the early times and faded on a time scale of $\\sim$30~days. When the SSS phase started around day 30 after outburst, some shock emission was still present that may contaminate the SSS emission (e.g., \\citealt{ness2009b}).  \\cite{rohrbach2009} analysed three Chandra ACIS spectra of V1494 Aql obtained on days 134, 187, and 248 after outburst, and interpreted a soft excess in the last observation as SSS emission. However, in the two earlier ACIS spectra, emission lines from N were present at the same energies. They emphasised that the soft component in all three ACIS spectra could equally well be fitted by a blackbody or a thin-thermal plasma model with high N or O abundance. The conclusion of SSS emission in a CCD-type spectrum has thus to be treated with care, because the energy resolution is not high enough to compare these models. \\cite{rohrbach2009} had additional high-resolution Chandra LETGS spectra at their disposal that clearly showed that the soft component was atmospheric rather than thin-thermal emission. For these purposes, it is thus mandatory to observe with high energy resolution with sufficient sensitivity at soft X-ray energies so that different models can be distinguished.  In this paper, we present the results of soft X-ray studies of the classical nova V2491 Cygni using the Suzaku observatory. Suzaku has taken two observations during the early phase, and the X-ray CCDs aboard Suzaku provided X-ray spectra of this nova \\citep{takei2009a,takei2009b}. A tantalising hint of SSS emission was seen on day 29, although it provides several different interpretations at the same time. We discuss different possibilities of interpreting the soft component. ",
        "conclusions": "The fast nova V2491 Cygni has been observed on days 9 and 29 after outburst with Suzaku. The first observation showed a non-thermal component and an optically thin thermal plasma component with a temperature of $\\sim$2.9~keV \\citep{takei2009a}. Later, on day 29, the non-thermal component has disappeared, and the thin-thermal component has a softer spectrum, yielding a lower temperature \\citep{takei2009b}. In addition, a super-soft component with a peak at 0.5~keV was present. The origin of this component can either be first light of the SSS phase, or it could be the soft extension of the thin-thermal plasma component. With the given combination of limited spectral resolution and soft sensitivity, we can not distinguish between these two possibilities. Technically, both scenarios are possible and plausible. Depending on the correct interpretation, the duration of the SSS phase could be different by about a week. According to \\cite{page2009}, the SSS phase started between days 36 and 42. If the soft component detected by Suzaku is indeed SSS emission, then the start of the SSS phase would be at least one week earlier.  Central to addressing such questions is the spectral resolution. In high resolution, one could clearly distinguish between SSS and a thin-thermal plasma emission. The results of this work are one example that demonstrates the need for high spectral resolution with sufficient sensitivity at soft energies that would have to be implemented in future X-ray observatories like the International X-ray Observatory."
    },
    "0911/0911.4719_arXiv.txt": {
        "abstract": "We present spectroscopic and photometric observations of the Type IIn supernova (SN) 2008iy. \\sn\\ showed an unprecedentedly long rise time of $\\sim$400 days, making it the first known SN to take significantly longer than 100 days to reach peak optical luminosity. The peak absolute magnitude of \\sn\\ was $M_r \\approx -19.1$ mag, and the total radiated energy over the first $\\sim$700 days was $\\sim$2 $\\times 10^{50}$ erg. Spectroscopically, \\sn\\ is very similar to the Type IIn SN~1988Z at late times, and, like SN~1988Z, it is a luminous X-ray source (both supernovae had an X-ray luminosity $L_{\\rm X} > 10^{41}$ erg s$^{-1}$). \\sn\\ has a growing near-infrared excess at late times similar to several other SNe~IIn. The H$\\alpha$ emission-line profile of \\sn\\ shows a narrow P Cygni absorption component, implying a pre-SN wind speed of $\\sim$100 \\kms. We argue that the luminosity of \\sn\\ is powered via the interaction of the SN ejecta with a dense, clumpy circumstellar medium. The $\\sim$400 day rise time can be understood if the number density of clumps increases with distance over a radius $\\sim$1.7 $\\times 10^{16}$ cm from the progenitor. This scenario is possible if the progenitor experienced an episodic phase of enhanced mass loss $<$ 1 century prior to explosion or if the progenitor wind speed increased during the decades before core collapse. We favour the former scenario, which is reminiscent of the eruptive mass-loss episodes observed for luminous blue variable (LBV) stars. The progenitor wind speed and increased mass-loss rates serve as further evidence that at least some, and perhaps all, Type IIn supernovae experience LBV-like eruptions shortly before core collapse. We also discuss the host galaxy of \\sn, a subluminous dwarf galaxy, and offer a few reasons why the recent suggestion that unusual, luminous supernovae preferentially occur in dwarf galaxies may be the result of observational biases. ",
        "introduction": "The recent development of synoptic, wide-field imaging has revealed an unexpected diversity of transient phenomena. One such example is the discovery of a new subclass of very luminous supernovae (VLSNe; e.g., \\citealt{ofek06gy,smith07-2006gy,quimby05ap}). While events of this nature are rare \\citep{miller08es,quimby2009}, each new discovery serves to bracket our understanding of the physical origin of well-established classes of SNe and does, in principle, demand an increased clarity in our understanding of the post-main-sequence evolution of massive stars.  Yet another unusual transient was discovered by the Catalina Real-Time Transient Survey (CRTS; \\citealt{drake2009}), which they announced via an ATel as CSS080928:160837+041626 on 2008 Oct.\\ 07 UT\\footnote{UT dates are used throughout this paper unless otherwise noted.} \\citep{drake08-2008iy}. The transient was classified as a Type IIn SN with a spectrum taken on 2009 Mar.\\ 27 (\\citealt{mahabal2009}; see \\citealt{schlegel96} for a definition of the SN~IIn subclass, and \\citealt{filippenko1997} for a review of the spectral properties of SNe). The SN was later given the IAU designation \\sn\\ \\citep{catelan2009}. \\citet{mahabal2009} noted that the transient was present on CRTS images dating back to 2007 Sep.\\ 13; however, it went undetected by the CRTS automated transient detection software until 2008 because it was blended with a non-saturated, nearby ($\\sim$11\\arcsec\\ separation) star.  Here, we present our observations and analysis of \\sn, which peaked around 2008 Oct.\\ 29 \\citep{catelan2009} and had a rise time of $\\sim$400 days. This implies that \\sn\\ took longer to reach peak optical brightness than any other known SN. Type II SN rise times are typically $\\la$ 1 week (e.g., SNe~2004et and 2006bp, \\citealt{wli-2004et, quimby06bp}; see also \\citealt{patat1993,wli-rates-LF}), and have never previously been observed to rise $\\ga$ 100 days, let alone 400, making \\sn\\ another rare example of the possible outcomes for the end of the stellar life cycle. In addition to an extreme rise time, \\sn\\ is of great interest because the unique circumstellar medium (CSM) in which it exploded may provide a link to very long-lived SNe, such as SN 1988Z, and thus provide clues into the nature of their progenitors.  This paper is organised as follows. Section 2 presents the observations. The data are analysed in Section 3, and the results are discussed in Section 4. We give our conclusions, as well as predictions for the future behaviour of \\sn, in Section 5. ",
        "conclusions": "We have reported on observations of the Type IIn \\sn, which took $\\sim$400 days to reach peak optical output. There are few known SNe with optical rise times $\\ga$ 50 days, and \\sn\\ is the first with a rise time $>$ 1 year. We argue that this long rise to peak is caused by the interaction of the SN ejecta with a dense CSM; radioactivity is unlikely to drive a 400-day rise. Furthermore, the late-time optical decay, $\\sim$0.4 mag (100 days)$^{-1}$, is slower than that of $^{56}$Co, which provides further evidence that radioactive heating is not a dominant energy source for \\sn. Spectroscopically, \\sn\\ is very similar to SN~1988Z at late times. SN~1988Z is understood to have exploded in a dense, clumpy CSM \\citep{chugai1994}, which we argue was also the case for \\sn. We detect \\sn\\ in X-rays with a total luminosity $L_{\\rm X} = (3.7 \\pm 1.2) \\times 10^{41}$ erg s$^{-1}$, which is similar to that of the Type IIn SNe~1988Z and 1995N \\citep{fox2000}. Similar to other SNe~IIn, \\sn\\ has a growing NIR-excess at late times. \\sn\\ had a peak absolute magnitude of $M_r \\approx -19.1$ mag and a total radiated energy of $\\sim 2 \\times 10^{50}$ erg in the optical (assuming no bolometric correction).  The steady increase in optical luminosity over a $\\sim$400-day period means that the wind-density parameter, $w$, increased over a distance of $\\sim$1.7 $\\times 10^{16}$ cm from the SN. We propose two possible scenarios to explain this increase in $w$: (i) the progenitor experienced an episode of LBV-like, eruptive mass-loss $\\sim$55 years prior to the SN, or (ii) the wind speed of the progenitor was increasing during the years leading up to core collapse. We prefer the former scenario, as the later adds unnecessarily complicated pieces to the puzzle without providing a unique solution. Our favoured scenario provides yet another piece of evidence that some SNe~IIn are connected to LBV-like progenitors (see \\citealt{gal-yam2007}, and references therein). We find that the host of \\sn\\ is a subluminous dwarf galaxy, though we caution against premature conclusions that unusual SNe, specifically those that are very luminous or have very long rise times, preferentially occur in low-mass dwarf galaxies.  Finally, we close with some predictions for the late-time behaviour of \\sn. There are examples of SNe~IIn whose luminosity dramatically drops after the SN ejecta overtake the dense CSM (e.g., SN~1994W; \\citealt{sollerman94w}). However, the similarities to SN~1988Z suggest that \\sn\\ could continue interacting, and thus remain luminous, for several years. This would allow long-term monitoring in the radio, X-ray, and optical, like SN~1988Z \\citep{williams2002, schlegel2006, aretxaga88z}. We predict that \\sn\\ is a luminous radio source, like SN~1988Z \\citep{vandyk88z}, though we note that at $z = 0.0411$ deep observations may be necessary for detection. The increasing $K_s$-band luminosity from day 560 to 710 is likely due to the presence of dust, and we predict that \\sn\\ will also be luminous in the mid-IR ($\\sim$3$-$5 $\\mu$m). The data are insufficient to distinguish between newly formed dust or dust that was present prior to the SN, but future medium- to high-resolution spectroscopy could distinguish between these two cases. If new dust is being formed, it should result in a systematic blueshift in the line profiles, as the dust creates an optically thick barrier to radiation from the receding SN ejecta.    \\bigskip  A.A.M. would like to thank D. Poznanski for useful discussions that helped improve this paper and M. Modjaz for discussions concerning the metallicity of the host. We thank Neil Gehrels for approving the {\\it Swift} ToO request for \\sn, which was submitted by Dave Pooley, and the {\\it Swift} team for scheduling and obtaining those observations. We acknowledge the use of public data from the {\\it Swift} archive. Cullen Blake, Dan Starr, and Emilo Falco assisted with the operation of PAIRITEL. We wish to thank the following Nickel observers: P. Thrasher, M. Kislak, J. Rex, J. Choi, I. Kleiser, J. Kong, M. Kandrashoff, and A. Morton. We thank the referee, Nikolai Chugai, for suggestions that helped to improve this paper.  A.A.M. is supported by the NSF Graduate Research Fellowship Program and DOE grant \\#DE-FC02-06ER41453. J.S.B.'s group is partially supported by NASA/{\\it Swift} grant \\#NNX08AN84G.   A.V.F.'s group is grateful for the financial assistance of NSF grant AST--0908886 and the TABASGO Foundation. PAIRITEL is operated by the Smithsonian Astrophysical Observatory (SAO) and was made possible by a grant from the Harvard University Milton Fund, the camera loan from the University of Virginia, and the continued support of the SAO and UC Berkeley. The PAIRITEL project is partially supported by NASA/{\\it Swift} Guest Investigator Grant \\#NNX08AN84G. This publication makes use of data products from the Two Micron All Sky Survey, which is a joint project of the University of Massachusetts and the Infrared Processing and Analysis Center/California Institute of Technology, funded by the National Aeronautics and Space Administration (NASA) and the National Science Foundation (NSF). Some of the data presented herein were obtained at the W. M. Keck Observatory, which is operated as a scientific partnership among the California Institute of Technology, the University of California, and NASA; it was made possible by the generous financial support of the W. M. Keck Foundation. The authors wish to recognise and acknowledge the very significant cultural role and reverence that the summit of Mauna Kea has always had within the indigenous Hawaiian community; we are most fortunate to have the opportunity to conduct observations from this mountain."
    },
    "0911/0911.3763_arXiv.txt": {
        "abstract": "Radial metallicity gradients are observed in the disks of the Milky Way and in several other spiral galaxies. In the case of the Milky Way, many objects can be used to determine the gradients, such as HII regions, B stars, Cepheids, open clusters and planetary nebulae. Several elements can be studied, such as oxygen, sulphur, neon, and argon in photoionized nebulae, and iron and other elements in cepheids, open clusters and stars. As a consequence, the number of observational characteristics inferred from the study of abundance gradients is very large, so that in the past few years they have become one of the main observational constraints of chemical evolution models. In this paper, we present some recent observational evidences of abundance gradients based on several classes of objects.  We will focus on (i) the magnitude of the gradients, (ii) the space variations, and (iii) the evidences of a time variation of the abundance gradients. Some comments on recent theoretical models are also given, in an effort to highlight their predictions concerning  abundance gradients and their variations. ",
        "introduction": "Radial metallicity gradients are observed in the disks of many galaxies, including the Milky Way, galaxies of the Local Group and other objects. Current research topics include the determination of (i) the magnitude of the gradients, (ii) any space variations along the  disk and (iii) possible time variations during the evolution of the host galaxy. In this paper, we present some recent observational evidences of radial abundance gradients. Our main focus is the Milky Way, but it will be shown that the analysis of some objects in the Local Group, particularly M33, is  useful in order to study the main properties of the gradients in our own Galaxy. A brief discussion of some recent theoretical models is also given, in an effort to highlight their predictions concerning the radial abundance gradients and their variations. Some recent reviews and general papers on abundance gradients include: Freeman (\\cite{freeman}), Maciel \\& Costa (\\cite{mc2009}), Rudolph et al. (\\cite{rudolph}), and Stasi\\'nska (\\cite{stasinska04}). Theoretical models are discussed by a number of people, including Fu et al. (\\cite{fu}), Magrini et al. (\\cite{magrini09a}), Cescutti et al. (\\cite{cescutti}), Moll\\'a and D\\'\\i az (\\cite{molla}), Chiappini et al. (\\cite{chiappini03}, \\cite{chiappini01}, \\cite{chiappini97}), Hou et al. (\\cite{hou}), and Gensler (these proceedings). Recent discussions on azimuthal and vertical gradients, not treated here, can be found in Davies et al. (\\cite{davies}) and Ivezi\\'c et al. (\\cite{ivezic}). ",
        "conclusions": "From all objects considered, some tentative conclusions may be drawn: (1) Average abundance gradients are generally between $-0.03$ dex/kpc and $-0.10$ dex/kpc, but a single value for the whole disk may be misleading. (2) Most evidences point to a flattening  out of the gradients  at large galactocentric distances. (3) There are some clear evidences of a flattening  of the gradients  near the galactic bulge. (4) The change of slope in the outer Galaxy occurs in the region around $R \\simeq 10$ kpc. (5) Any further change of the slope needs better data than presently available.  Cepheids may be an exception. (6) There are no evidences of a steepening of the gradients at large galactocentric distances, as suggested  by some theoretical models. (7) Either the gradients do not change appreciably during galactic evolution, or they flatten out at a moderate rate. (8) There are no clear evidences of a steepening of the gradients with time, as suggested by some theoretical models.  \\bigskip\\noindent {\\it Acknowledgements. This work was partially supported by FAPESP and CNPq.}"
    },
    "0911/0911.4972_arXiv.txt": {
        "abstract": "Weak gravitational lensing has been used extensively in the past decade to constrain the masses of galaxy clusters, and is the most promising observational technique for providing the mass calibration necessary for precision cosmology with clusters.  There are several challenges in estimating cluster masses, particularly (a) the sensitivity to astrophysical effects and observational systematics that modify the signal relative to the theoretical expectations, and (b) biases that can arise due to assumptions in the mass estimation method, such as the assumed radial profile of the cluster.  All of these challenges are more problematic in the inner regions of the cluster, suggesting that their influence would ideally be suppressed for the purpose of mass estimation. However, at any given radius the differential surface density measured by lensing is sensitive to all mass within that radius, and the corrupted signal from the inner parts is spread out to all scales.  We develop a new statistic \\ups\\ that is ideal for estimation of cluster masses because it completely eliminates mass contributions below a chosen scale (which we suggest should be about 20 per cent of the virial radius), and thus reduces sensitivity to systematic and astrophysical effects.  We use simulated and analytical profiles including shape noise to quantify systematic biases on the estimated masses for several standard methods of mass estimation, finding that these can lead to significant mass biases that range from ten to over fifty per cent.  The mass uncertainties when using the new statistic \\ups\\ are reduced by up to a factor of ten relative to the standard methods, while only moderately increasing the statistical errors. This new method of mass estimation will enable a higher level of precision in future science work with weak lensing mass estimates for galaxy clusters. % ",
        "introduction": "\\label{S:introduction}  Many scientific applications require robust measurements of the mass in galaxy clusters.  One such application is the use of the dark matter halo mass function to constrain cosmological model parameters, including the amplitude of matter density perturbations, the average matter density, and even the equation of state of dark energy \\citep[e.g., most recently, ][]{2007ApJ...657..183R,2008MNRAS.387.1179M,2009ApJ...692.1060V,2010ApJ...708..645R}. Another example is validation and refinement of models of cluster formation and evolution, which predict relations between the more easily measured optical and X-ray emission, and the underlying dark matter halo \\citep{2006ApJ...650..128K,2007ApJ...668....1N,2008A&A...482..451Z,2009arXiv0906.4370B}. Currently, there are thousands of known clusters selected in various ways that can be used for these applications.  Future surveys such as the Dark Energy Survey (DES)\\footnote{\\texttt{http://www.darkenergysurvey.org/}}, Pan-STARRS\\footnote{\\texttt{http://pan-starrs.ifa.hawaii.edu/public/}}, and the Large Synoptic Survey Telescope (LSST)\\footnote{\\texttt{http://www.lsst.org/lsst}} will provide even larger and deeper samples that can be used for this purpose, requiring greater systematic robustness in the mass measures to complement the smaller statistical errors.  Many different methods have been used to measure the halo profile of clusters and thereby estimate their masses.  Kinematic tracers such as satellite galaxies, in combination with a Jeans analysis or caustics analysis, can give information over a wide range of physical scales and halo masses.  While the issues of relaxation, velocity bias, anisotropy of the orbits and interlopers need to be carefully addressed, recent results suggest a good agreement with theoretical predictions for the form of the density profile \\citep{2003ApJ...585..205B,2004ApJ...600..657K,2003AJ....126.2152R,2005ApJ...628L..97D,2006AJ....132.1275R,2007MNRAS.378...41S}. Hydrostatic analyses of X-ray intensity profiles of clusters % use X-ray intensity and temperature as a function of radius to reconstruct the density profile and estimate a halo mass.  The advantage of thermal gas pressure being isotropic in partially lost due to the possible presence of other sources of pressure support, such as turbulence, cosmic rays or magnetic fields. These extra sources of pressure support cannot be strongly constrained for typical clusters with present X-ray data \\citep{2004A&A...426..387S}, but could modify the hydrostatic equilibrium and affect the conclusions of such analyses. Recent results are encouraging and are in a broad agreement with predictions, although most require concentrations that are higher than those predicted by a concordance cosmology \\citep{2006ApJ...640..691V,2007ApJ...664..123B,2007MNRAS.379..209S}. While the above-mentioned systematic biases cannot be excluded, the small discrepancy could also be due to baryonic effects in the central regions, due to selection of relaxed clusters that may be more concentrated than average \\citep{2006ApJ...640..691V}, or due to the fact that at a given X-ray flux limit, the more concentrated clusters near the limiting mass are more likely to be included in the sample \\citep{2007A&A...473..715F}.  Gravitational lensing is by definition sensitive to the total mass, and is therefore one of the most promising methods to measure the mass profile independent of the dynamical state of the clusters.  Many previous weak lensing analyses have focused on individual clusters (for example, \\citealt{2007MNRAS.379..317H,2007ApJ...667...26P,2009ApJ...702..603A,2009arXiv0903.1103O}). Measuring the matter distribution of individual clusters allows a comparison with the combined baryonic (light and gas) distribution on an individual basis, and so can constrain models that relate the two, such as MOND versus CDM \\citep{2006ApJ...648L.109C}.  However, these measurements can be quite noisy for individual clusters.  Stacking the signal from many clusters can ameliorate this problem, since shape noise and the signal due to correlated structures will be averaged out.  Such a statistical approach is thus advantageous if one is to compare the observations to theoretical predictions, which also average over a large number of halos in simulations.  A final advantage of stacking is that it allows for the lensing measurement of lower-mass halos, where individual detection is impossible due to their lower shears relative to more massive clusters.  Individual high signal-to-noise cluster observations and those based on stacked analysis of many clusters are thus complementary to each other at the high mass end, with the stacked analysis drastically increasing the available baseline in mass.  Extraction of cluster dark matter halo masses from the weak lensing signal is subject to a number of uncertainties, which we discuss in this paper in detail, including the ways that the uncertainties differ for individual versus stacked cluster lensing analyses.  In brief, these uncertainties are: (i) biased calibration of the lensing signal; (ii) modification of the lensing profile in the inner cluster regions due to accidental inclusion of cluster member galaxies in the source sample, intrinsic alignments of those galaxies, non-weak shear, magnification, baryonic effects that modify the initial cluster dark matter halo density profile, and cluster centroiding errors; (iii) contributions to the lensing signal from nonvirialized local structures and large-scale structure (LSS).  Furthermore, parametric modeling of the mass requires the assumption of a form for the dark matter halo profile, which may differ from the intrinsic profile and/or have poorly constrained parameters.  Non-parametric modeling, while not subject to this weakness, results in projected masses that must be converted to three-dimensional enclosed masses to be compared against the theory predictions, all of which are currently phrased in terms of 3d masses.  We quantify the degree to which this conversion depends on assumptions about the density profile.  Generally, we show the effects of many of these uncertainties on the estimated masses from cluster weak lensing analyses, both in the stacked and individual cases, using parametric and non-parametric mass modeling.  Effects that modify the cluster density profile in the inner regions ($\\lesssim 0.5$\\hmpc), are particularly problematic given that the weak lensing signal \\dsr\\ is sensitive to the density profile not just at a projected separation $R$, but also at all smaller separations. We propose a modified statistic, denoted \\ups, that removes the dependence on the projected density between $R=0$ and $R=R_0$, with $R_0$ chosen to avoid scales with systematic uncertainties.  The decrease in systematic errors that results from removing scales below $R_0$ comes at the expense of somewhat increased statistical errors.  We explore the optimal choice of $R_0$, and quantify the degree to which our use of this new statistic to estimate cluster masses lessens systematic biases and increases statistical errors.  Our tools for this investigation include simple, idealised cluster density profiles; more complex and realistic density profiles from $N$-body simulations; and finally, real cluster lensing data from the Sloan Digital Sky Survey (SDSS, \\citealt{2000AJ....120.1579Y}) that was previously analysed by \\cite{2008JCAP...08..006M}.  We begin in Section~\\ref{S:theory} with a discussion of the theoretical aspects of cluster-galaxy weak lensing, including a detailed discussion of the challenges of mass determination, and a summary of typical approaches to parametric and non-parametric mass estimation, with the introduction of a new statistic from which to derive parametric mass estimates.  In Section~\\ref{S:simulations}, we describe the $N$-body simulations that we use to provide sample cluster density profiles. Section~\\ref{S:data} has a description of the SDSS cluster lensing data we use to test for some of the effects that we find using the simulations.  Results for both the theoretical profiles and the real data are presented in Section~\\ref{S:results}.  We conclude with a discussion of our findings and their implications in Section~\\ref{S:conclusions}. ",
        "conclusions": "\\label{S:conclusions}  In this paper, we have assessed the degree to which certain systematic errors in lensing measurement and methods of mass estimation can bias weak lensing cluster mass estimates.  In brief, the challenges we considered included the following. \\begin{itemize} \\item Lensing calibration bias, which leads to changes in the mass $\\propto \\ds^{\\eta}$ for $\\eta$ typically in the range $1.4$--$2$ depending on the radial range and fit method used for the parametric NFW mass fits (lower for \\dsr\\ than for \\ups), or $\\propto \\ds$ for the non-parametric mass estimates within a fixed physical aperture (or a steeper scaling when estimated the mass within some spherical over-density radius) using \\zetac\\ (Section~\\ref{SSS:calibration},~\\ref{SSS:resnonpara}, and~\\ref{SSS:ressimpara}). \\item Offsets of the identified BCG from the minimum of the cluster potential well (Section~\\ref{SSS:offsets}) were incorporated into the lensing profiles using a model from mock catalogues presented in \\cite{2007arXiv0709.1159J}.  This model includes the effects of photometric errors in selecting the wrong BCG, and is therefore an overly conservative estimate in cases where the BCG can be unambiguously identified for all clusters or where X-ray data can precisely locate the cluster centre. \\item The effect of differences between an assumed NFW concentration and the true NFW concentration were studied using pure NFW lensing signals. \\item Differences in the halo profile relative to a pure NFW profile were studied using fits to density profiles from $N$-body simulations. \\item The effects of mass from structures other than the cluster itself on the lensing signal were also studied using the signal from simulations, since we have not included only the mass that is virialized when computing \\ds\\ in the simulations. \\item Contamination of the source sample by cluster member galaxies, intrinsic alignments of those member galaxies, and baryonic effects on the halo density profile were considered to be included among the previous tests, namely changes in NFW concentration (in Section~\\ref{SS:theorresults}), changes in the inner region of the profile using variations of the $N$-body simulation outputs (in Section~\\ref{SS:simresults}), and centroid offsets that modify the signal only in the inner regions of the cluster. \\end{itemize}  When fitting a parametric model (in our case, the NFW profile) to \\dsr, with fixed concentration, we find that the uncertainties due to unknown true concentration plus changes in the lensing profile due to small-scale systematics yield systematic errors that range from a factor of two in mass (when only using small scales in the fits, e.g. $0.1$--$1$\\hmpc) to tens of per cent (when using $R>0.5$\\hmpc) to several per cent (for $R>2$\\hmpc, which yields stable mass estimates but large statistical errors, and which may not be available for individual cluster lensing analyses due to limited telescope FOV). This level of systematic error occurred when allowing a relatively broad variation in concentration ($4<\\cvir <7$), given the disagreement between simulations on the concentration-mass relation at high masses, the large lognormal scatter in this relation, and other systematics such as baryonic effects discussed in Sections~\\ref{SS:challenges-theory}, \\ref{SS:challenges-obs}, and~\\ref{SS:paramodel}. When using a narrower range in concentration, the systematic errors decreased comparably, but are still unacceptably large relative to what is needed for precise cosmological parameter constraints.  The addition of centroiding errors to the list of systematics we considered led to uniform suppression of the mass estimates of order tens of per cent (for $\\rmin=0.1$\\hmpc).  To completely avoid this suppression while fitting to \\dsr\\ and ignoring the possibility of centroiding errors, we found it necessary to restrict the fits to $\\rmin>1$\\hmpc.  Generally, the addition of larger scales, out to $\\sim 2$\\rvir, is useful in minimising the effects of small-scale systematics; going beyond that can lead to excessive contribution from large-scale structure, which will bias the mass estimates if it is not modelled accurately.  Allowing a variation in concentration in the fits is another way to reduce systematic error, at the expense of statistical errors that are increased by 45 per cent, but this scheme is not helpful when dealing with systematics that have a radial profile that does not mimic a change in concentration.  \\ups\\ is still more reliable at removing the impact of small-scale systematics on the lensing signal.  The aperture mass statistic \\zetac\\ led to accurate estimates of projected masses, provided that either (a) the mass in the outer annulus was estimated rather than ignored, or (b) the mass in the outer annulus was ignored, but $R_{o1}\\gg R_1$ (i.e. a large range of transverse separations was included in the first integral in Eq.~\\eqref{E:zetac}).  For many applications, such as the halo mass function, the quantity of interest is the 3d virial mass, for which a density profile must be assumed to do the conversion from the 2d projected mass within $R_1$.  We found that uncertainty in the true density profile led to tens of per cent level biases in the 3d virial masses.  The effect of centroiding errors was to uniformly suppress the aperture masses by $\\sim 10$--$20$ per cent depending on the halo mass, degree of centroiding errors, and transverse separations used for the analysis; these biases were then propagated into the 3d enclosed mass estimates.  The aperture mass-based estimates of the cluster virial mass were substantially noisier than fits to \\dsr\\ using the same range of scales.  Finally, the new statistic we introduce here, \\ups, removes the effect of small scales from the lensing signal, gave superior performance over \\dsr\\ when fitting an NFW profile to the cluster lensing signal. This statement is true not just for the basic tests with pure NFW profiles and profiles from simulations, but also when including the effects of centroiding errors.  The increases in statistical error on the mass can be $\\sim 40$ per cent relative to fitting to \\dsr\\ over the same scales.  % The residual systematic uncertainties after removal of an overall offset in the masses is of order several per cent, when fitting from $0.5<R<4$\\hmpc, as demonstrated using SDSS maxBCG data.  The effects of \\ups\\ in decreasing systematic error are less dramatic when only small scales ($\\le 2$\\hmpc) are used for the mass estimates; however, the residual systematics of order 10 per cent are still at least a factor of two smaller than when fitting to \\dsr.  These conclusions also apply for individual cluster lensing analyses; however, we caution that in that case, we expect additional uncertainties in the true halo profile due to contamination by cluster member galaxies, the lognormal scatter in concentration at fixed mass, mergers, substructure, triaxiality, and projection effects (Section~\\ref{SS:challenges-theory} and~\\ref{SS:challenges-obs}), so the systematic errors will tend to be larger than for stacked analyses using the same mass estimation method.  Next, we will briefly discuss the implications of our findings about mass estimation methods for several previously-published cluster lensing studies.  We begin with \\cite{2009arXiv0903.1103O}, which contains an analysis of circularly-averaged cluster lensing data for thirty individual clusters by comparison with spherical models.  They begin with direct fitting of the tangential shear profile to parametric models, including the NFW profile.  These fits allow all model parameters to vary; for example, the NFW fits have two parameters, a mass and a concentration, unlike the cases we have considered here with a fixed concentration.  Consequently, the estimated masses from the NFW model fits are unlikely to be strongly biased due to modeling assumptions, since the concentration is not fixed.  However, they may still have some systematic bias due to the NFW profile not describing cluster profiles well, due to deviations in individual cluster profiles due to substructure, mergers, and triaxiality, and possibly significant biases due to small-scale systematics such as contamination by cluster member galaxies and centroiding errors.  \\cite{2009arXiv0903.1103O} also use the aperture mass statistic \\zetac\\ to estimate \\mproj, while neglecting the second term in Eq.~\\eqref{E:zetac} and choosing the outer annulus such that it does not contain any significant structures.  As we have seen here, the aperture mass statistic when including both terms properly leads to quite accurate projected mass estimates, or can yield results that are accurate at the several per cent level even without the second term provided that $R_1 \\ll R_{o1}$.  Given the scales that are accessible with the Subaru Suprime-Cam, and the typical cluster redshifts, we should compare against the top portion of Table~\\ref{T:nfwmap}, the rows with $(R_1, R_{o1}, R_{o2})=(0.275, 1.1, 2)$ and $(0.5,1.1,2)$\\hmpc. Those results suggest that for the most massive clusters, neglect of the second term may cause 15--20 per cent suppression of the projected masses.  We find that the suppression is reduced to 5--10 per cent for more typical cluster masses of $10^{14}\\hmsun$. Furthermore, as we have already seen, effects that suppress the signal in the inner cluster regions, such as centroiding errors and contamination by cluster member galaxies, can suppress the aperture masses at the $\\sim 10$ per cent level.  \\cite{2007MNRAS.379..317H} contains an analysis of cluster weak lensing data for twenty individual clusters.  This work utilises parametric mass estimates from the tangential shear distortion averaged in annuli, fitting to an NFW profile with fixed concentration-mass relation from $N$-body simulations using $0.25<R<1.5$\\hmpc.  In this case, we can assess systematic uncertainties as being somewhere between the results for $(\\rmin, \\rmax) = (0.25, 1)$ and $(0.25, 2)$\\hmpc\\ on Fig.~\\ref{F:showsimfit}.  That figure suggests that uncertainties due to differences between the assumed and the true profile lead to $\\sim 50$ per cent variations in the estimated cluster halo masses.  This variation may be manifested as significant noisiness in the mass estimates for a given true mass, as well as a mean bias if the true profiles (with the imposition of systematics such as contamination by cluster member galaxies) differ from the NFW profile with that assumed concentration-mass relation. This problem is in addition to other uncertainties in individual cluster mass estimates noted previously, such as LSS (for which they explicitly increase their error bars) and triaxiality.  \\cite{2007MNRAS.379..317H} also use the aperture mass statistic to estimate projected masses, \\mproj, while estimating the second term in Eq.~\\eqref{E:zetac} due to the outer annulus using the best-fitting NFW model.  In that case, we note that while \\cite{2007MNRAS.379..317H} do not miss mass by excluding the second term in the aperture mass calculation, their conversion from $\\mproj(<R_1)$ to virial radii using spherical overdensities that can define the mass function will be strongly concentration-dependent. While \\cite{2007MNRAS.379..317H} claim that the fact that the masses from the fits to \\dsr\\ and from the aperture mass calculation are consistent shows that their fitting procedure is unbiased, as discussed in Section~\\ref{SSS:resnonpara} this claim is not true.  The fact that \\cite{2006ApJ...640..691V} use the cluster mass estimates from this work to calibrate their mass function constraints is therefore of concern, because of the possible biases due to these systematics in the signal and the large systematic scatter that we have found.  In summary, we believe that weak lensing is the best observational technique to robustly estimate cluster virial masses (regardless of their dynamical state) at the level required for precision cosmology. Given the small statistical errors of recent cluster abundance analyses, the cosmological constraints are already dominated by the systematic precision of the cluster mass determination \\citep{2006ApJ...640..691V}. As we argue in this paper, current methods are inadequate for this purpose because they rely on the information from the inner parts of the cluster, which can be contaminated or modified due to a variety of effects discussed in this paper, and because they do not use numerical $N$-body simulations to calibrate their results.  Our results suggest eliminating lensing information from scales below $R_0$ (for which we suggest the range $0.2<R_0<0.5$\\hmpc\\ or about 15-25 per cent of the virial radius, as determined via an iterative procedure).  Our proposed statistic for parametric estimates of cluster mass, the ADSD \\ups, achieves this by construction, and is consequently more robust to many different systematics and to the details of the model to which the data are fitted, all of which are more problematic in the inner parts of the cluster.  Use of \\ups\\ to estimate cluster masses allows systematic errors to be reduced to the several per cent level, which is up to a factor of 10 smaller than when fitting to the lensing signal \\dsr\\ itself, suggesting that for current and future datasets, \\ups\\ should be the statistic of choice for parametric mass fitting to cluster weak lensing data. While we have focused on clusters in this paper, similar concerns about accurately determining the halo mass would arise also for smaller halos. For these, the stellar component from the galaxy (and possibly the associated redistribution of the dark matter) would modify the mass distribution relative to predictions from pure $N$-body simulations in the inner parts, suggesting that eliminating the inner halo information by using \\ups\\ could lead to more accurate mass determination of group and galaxy type halos as well."
    },
    "0911/0911.0189_arXiv.txt": {
        "abstract": "\\torie is a young, moderate mass binary system, a rarely observed case of spectral-type G-giants of about 3 Solar masses, which are still collapsing towards the main sequence, where they presumably become X-ray faint.  We have obtained high resolution X-ray spectra with \\chandra\\ and find that the system is very active and similar to coronal sources, having emission typical of magnetically confined plasma: a broad temperature distribution with a hot component and significant high energy continuum; narrow emission lines from H- and He-like ions, as well as a range of Fe ions, and relative luminosity, $L_x/L_\\mathrm{bol} = 10\\mthree$, at the saturation limit.  Density, while poorly constrained, is consistent with the low density limits, our upper limits being $n_e<10^{13}\\,\\mathrm{cm\\mthree}$ for \\eli{Mg}{11} and $n_e<10^{12}\\,\\mathrm{cm\\mthree}$ for \\eli{Ne}{9}.  Coronal elemental abundances are sub-Solar, with Ne being the highest at about 0.4 times Solar. We find a possible trend in Trapezium hot plasmas towards low relative abundances of Fe, O, and Ne, which is hard to explain in terms of the dust depletion scenarios of low-mass young stars.  Variability was unusually low during our observations relative to other coronally active stars.  Qualitatively, the emission is similar to post main-sequence G-stars.  Coronal structures could be compact, or comparable to the dimensions of the stellar radii.  From comparison to X-ray emission from similar mass stars at various evolutionary epochs, we conclude that the X-rays in \\tooe are generated by a convective dynamo, present during contraction, but which will vanish during the main-sequence epoch, and possibly to be resurrected during post main-sequence evolution. ",
        "introduction": "} \\note{abstract updated} %  Intermediate mass pre-main sequence (PMS) stars, like their massive cousins, are difficult to study because of their rapid evolutionary time scales. Though not as short as stars above $8\\,\\msun$ which take less than $10^5$ years to reach the main sequence, intermediate mass stars between 2$\\msun$ and $8\\msun$ may only take 10--20 Myr.  In both cases the accretion time scales dominate the evolution time to the zero-age main sequence (ZAMS), in contrast to the low-mass T Tauri stars for which PMS contraction times are longest. Intermediate mass PMS stars are also not easily found and identified, and most existing studies focus on Herbig Ae and Be (HAeBe) stars.  Herbig stars \\citep{Herbig:1960} are recognized as such once they already contracted to high enough photospheric temperatures to be optically identified as A and B stars and are thus already close to the ZAMS. \\Vtworev % Herbig stars mark the transition between formation mechanisms of low-mass and high-mass stars~\\citep{Baines06}.  Herbig Ae stars seem more similar to the low-mass T Tauri stars~\\citep{Waters98, Vink05}. Herbig Be stars are more similar to embedded young massive stars \\citep{Drew97}. Both the Ae and Be stars are all already in fairly late PMS stages.  Most studies of Herbig stars use infrared and optical wavelengths to probe their circumstellar disks and dusty environments.  Recent studies suggest that specifically in HAe stars there is evidence not only for circumstellar disks~\\citep{Mannings97, Grady99} but also indications of dust shadowing and settling indicative of dust grain growth and planetesimal formation \\citep{Acke04, Dullemond04, Grady05}. Some studies also suggest that magnetospheric accretion analogous to classical T Tauri stars is possible ~\\citep{Muzerolle04, Grady04, Guimaraes06}.  Recent modeling of 37 Herbig Ae/Fe stars using UV spectra revealed that all but one show indications of accretion with accretion rates in many cases substantially exceeding $10^{-8}\\, \\msun \\mathrm{yr^{-1}}$ \\citep{Blondel06}.  Binarity also seems to be an important attribute in the formation and evolution of intermediate mass stars.  In a sample of 28 HAeBe stars, \\citet{Baines06} find a binarity fraction of almost 70\\% with a higher binary frequency in HBe stars than in HAe stars.  HAe stars with close companions also seem to lack circumstellar disks~\\citep{Grady05}.  X-ray studies of young intermediate mass stars are still quite rare and to date also focus almost entirely on HAe stars. Systematic studies have shown that these are moderately bright in X-rays ~\\citep{Damiani94, Zinnecker94, Hamaguchi05, Stelzer:al:2006}.  This fact is already quite remarkable since main sequence A-stars lack strong winds or coronae and it suggests that the physical characteristics of HAe stars stars differ from those of main sequence A- and B-stars. Mechanisms suggested range from active accretion to coronal activity to some other form of plasma confinement. \\Vtworev % It is also possible, given the high frequency of binaries among HAeBe stars, that some X-ray sources could be due to late-type companions \\citep{Stelzer:Hu:al:2006, Stelzer:al:2006}. A detailed summary can be found in a recent \\chandra\\ high resolution spectroscopic study of the HAe star HD~104237 in the $\\epsilon$ Chamaeleontis Group \\citep{Testa08}.   \\Vtworev % In this paper we focus on X-ray emission from \\tooe, which was recently determined to be an intermediate mass binary star.  The Orion Trapezium is generally known for its ensemble of the nearest and youngest massive stars \\citep{Schulz:Canizares:al:2001,Schulz:2003,Stelzer:al:2005}. Recent studies now suggest the presence of several intermediate mass stars. \\toriaa\\ harbors the second most massive O-star of the Trapezium, but also two unidentified intermediate mass stars both between $3\\msun$ and $7\\msun$ \\citep{Preibisch99}. The system is particularly interesting in X-rays for its high luminosity and hard spectral properties as well as giant hard X-ray outbursts \\citep{Feigelson02, Schulz06}.  Plasma temperatures during these outbursts exceed $10^8 \\mathrm{K}$ \\citep{Schulz06}.  While the latter authors suggested a possible link of these outbursts to binary interactions involving the closer intermediate mass companion, there is also some evidence that these may be connected to the more distant companion (M.\\ Gagne, private communication).  \\torie\\ is another system now known to contain young intermediate mass stars.  It was long misidentified as B5 to B8 spectral type \\citep{Parenago54, Herbig:1960}. \\citet{Herbig:Griffin:2006} obtained optical spectroscopic radial velocity measurements and identified the system as a binary containing two G~III type stars of masses of about 3--4 \\Msun in a 9.9 day orbit.  Evolutionary tracks constrain the age of the system to 0.5--1.0 Myr making the components of \\torie\\ some of the youngest intermediate mass PMS stars known and far younger than stars in the HAeBe phase. \\torie\\ is not among the optically brightest stars in the Trapezium, but has long been recognized as the second \\Vtworev % brightest Trapezium source in X-rays \\citep{Ku82, Gagne94, Schulz:Canizares:al:2001}.  The \\chandra\\ Orion Ultradeep Project (COUP) observed \\torie\\ (COUP 732) for a total exposure of about 10 days over a time period of 3 weeks and found a low level of variability including one moderate X-ray flare \\citep{Stelzer:al:2005}.  Its luminosity during COUP was determined to be $\\log L_x [\\ergsec] = 32.4$; early \\chandra\\ High Energy Transmission Grating (HETG) spectra indicated plasma temperatures of up to 50 MK \\citep{Schulz:2003}.  The HETG Orion Legacy Project has now accumulated almost 4 days of total exposure of \\torie\\ allowing for an in depth study of its X-ray spectral properties.  The following analysis of \\torie's high resolution X-ray spectrum is aimed to characterize its coronal nature. \\Vtworev % The existence of coronal X-rays confirms predictions that very young intermediate mass stars of less than $4\\msun$ are not fully radiative \\citep{Palla93} and may possess some form of magnetic dynamo. We also compare these properties with the ones observed in \\toriaa, various T \\Vtworev % Tauri stars including the relatively massive T Tauri star, SU~Aur ($2\\msun$), active coronal sources, and post-main sequence evolved G-type giants.  The optical and binary system parameters of \\tooe\\ can be found in \\citet{Herbig:Griffin:2006}. ",
        "conclusions": ""
    },
    "0911/0911.2003_arXiv.txt": {
        "abstract": "We present Advanced Camera for Surveys observations of MACS\\,J1149.5$+$2223, an X-ray luminous galaxy cluster at $z{=}0.544$ discovered by the Massive Cluster Survey.  The data reveal at least seven multiply-imaged galaxies, three of which we have confirmed spectroscopically.  One of these is a spectacular face-on spiral galaxy at $z=1.491$, the four images of which are gravitationally magnified by $8\\ls\\mu\\ls23$.  We identify this as an $L^\\star$ ($M_B\\simeq-20.7$), disk-dominated ($B/T\\ls0.5$) galaxy, forming stars at $\\sim6\\Msolpyr$.  We use a robust sample of multiply-imaged galaxies to constrain a parameterized model of the cluster mass distribution.  In addition to the main cluster dark matter halo and the bright cluster galaxies, our best model includes three galaxy-group-sized halos.  The relative probability of this model is ${\\rm P}(N_{\\rm halo}=4)/{\\rm P}(N_{\\rm halo}<4)\\ge10^{12}$ where $N_{\\rm halo}$ is the number of cluster/group-scale halos.  In terms of sheer number of merging cluster/group-scale components, this is the most complex strong-lensing cluster core studied to date.  The total cluster mass and fraction of that mass associated with substructures within $R\\le500\\kpc$, are measured to be $\\Mtot=(6.7\\pm0.4)\\times10^{14}\\Msol$ and $\\fsub=0.25\\pm0.12$ respectively.  Our model also rules out recent claims of a flat density profile at $\\gs7\\sigma$ confidence, thus highlighting the critical importance of spectroscopic redshifts of multiply-imaged galaxies when modeling strong lensing clusters.  Overall our results attest to the efficiency of X-ray selection in finding the most powerful cluster lenses, including complicated merging systems. ",
        "introduction": "\\label{sec:intro}  Strong gravitational lensing by galaxy clusters is a well-established probe of the extragalactic Universe, offering a magnified view of high redshift galaxies that would otherwise be beyond the reach of present day telescopes \\citep[e.g.][]{Franx97, Ellis01, Kneib04b, Smail07, Swinbank07, Bradley08}.  The detailed cluster mass models required to intepret such observations contain a wealth of information about the mass and structure of cluster cores, against which theoretical predictions can be tested \\citep[e.g.][]{Smith05a,Comerford06,Sand08}.  The majority of spectroscopically confirmed strong lensing clusters are at low redshift, i.e.\\ $z{\\ls}0.3$ \\citep[e.g.][]{Kneib96,Broadhurst05acs,Smith05a,Limousin07a1689,Richard09}. In contrast, only four spectroscopically confirmed strong lensing clusters are known at $z{>}0.5$ \\citep{Gladders02,Inada03,Borys04,Sharon05,Swinbank07,Ofek08}.  The Massive Cluster Survey \\citep[MACS;][]{Ebeling01} offers an unprecedented opportunity to expand the available sample of strong lensing clusters at $0.3\\le z\\ls0.7$.  We present new results from this search: MACS\\,J1149.5$+$2223 (hereafter MACS\\,J1149; 11:49:34.3 $+$22:23:42.5 [J2000]) at $z=0.544$, one of a complete subsample of 12 MACS clusters at $z>0.5$ \\citep{Ebeling07}.  In \\S\\ref{sec:obs} we describe the data; modeling and results are presented in \\S\\ref{sec:results}, and summarized in \\S\\ref{sec:summary}. We assume $\\Ho=70\\kms\\Mpc^{-1}$, $\\oM=0.3$ and $\\oL=0.7$; at $z=0.544$ $1''$ corresponds to $6.35\\kpc$.  All uncertainties and upper/lower limits are stated and/or plotted at $95\\%$ confidence. ",
        "conclusions": "We have presented new \\emph{HST}/ACS and Keck I/LRIS observations of MACS\\,J1149, a massive X-ray selected galaxy cluster at $z=0.544$ discovered in the Massive Cluster Survey.  These data reveal seven robustly identified multiply-imaged galaxies, three of which we have confirmed spectroscopically.  The most spectacular system is a multiply-imaged face-on disk galaxy at $z=1.491$ that we identify as an $L^\\star$ ($M_B\\simeq-20.7$) late-type ($B/T\\ls0.5$) galaxy with an ongoing star formation rate of $\\sim6\\Msolpyr$; the brightest images of this galaxy are magnified by $\\mu=23$.  Future observations using integral field spectrographs should probe its properties in exquisite detail, thanks to the combination of lens magnification and fortuitous viewing angle.  We use the positions and redshifts of robustly identified multiply-imaged galaxies to constrain a detailed model of the mass and structure of the cluster core.  Our fiducial model contains the main cluster halo plus three group-scale halos; the probability of a model this complex, relative to less complex models is ${\\rm P}(N_{\\rm halo}=4)/{\\rm P}(N_{\\rm halo}<4)\\ge10^{12}$ where $N_{\\rm halo}$ is the number of cluster/group-scale halos.  We measure the mass and fraction of mass residing in substructures to be $M(\\le500\\kpc)=6.7\\pm0.4\\times10^{14}\\Msol$ and $\\fsub(\\le500\\kpc)=0.25\\pm0.12$ respectively.  In summary, MACS\\,J1149 is the most complex strong-lensing cluster core studied to date, its relatively dis-assembled nature being qualitatively consistent with the expectation that clusters at high redshifts are on average less mature than those at lower redshifts.  A more complete view will emerge from our analysis of the full sample of MACS clusters at $z>0.5$ (Smith et al., in prep.).  We also obtain a power law density profile slope of $\\gamma=-0.3\\pm0.05$ ($95\\%$ confidence error bars) on scales of $r\\sim3-30\\arcsec$, thereby ruling out density profile slopes as flat as those recently proposed by \\cite{Zitrin09b} at $\\gs7\\sigma$ confidence.  In summary, \\citeauthor{Zitrin09b}'s result can be explained by an absence of multiple-image redshifts of any form in their study, and by them not treating the unknown redshifts as free parameters in their model.  These issues are probably compounded by them mis-identifying some multiple-image systems.  Overall, this underlines the critical importance of measuring spectroscopic redshifts of multiply-imaged galaxies for reliable lens models of strong lensing clusters."
    },
    "0911/0911.3143_arXiv.txt": {
        "abstract": "{The discovery of \\epsba, Bb, a binary brown dwarf system very close to the Sun, makes possible a concerted campaign to characterise the physical parameters of two T dwarfs. Recent observations suggest substellar atmospheric and evolutionary models may be inconsistent with observations, but there have been few conclusive tests to date. We therefore aim to characterise these benchmark brown dwarfs to place constraints on such models. We have obtained high angular resolution optical, near-infrared, and thermal-infrared imaging and medium-resolution (up to R$\\sim$5\\,000) spectroscopy of \\epsba, Bb with the ESO VLT and present $VRIzJHKL'M'$ broad-band photometry and 0.63--5.1\\micron\\ spectroscopy of the individual components. The photometry and spectroscopy of the two partially blended sources were extracted with a custom algorithm. Furthermore, we use deep AO-imaging to place upper limits on the (model-dependent) mass of any further system members. We derive luminosities of log L/L$_{\\sun}$=$-4.699\\pm0.017$ and $-5.232\\pm0.020$ for \\epsba, Bb, respectively, and using the dynamical system mass and COND03 evolutionary models predict a system age of 3.7--4.3\\,Gyr, in excess of previous estimates and recent predictions from observations of these brown dwarfs.  Moreover, the effective temperatures of 1352--1385\\,K and 976--1011\\,K predicted from the COND03 evolutionary models, for \\epsba\\ and Bb respectively, are in disagreement with those derived from the comparison of our data with the BT-Settl atmospheric models where we find effective temperatures of 1300--1340\\,K and 880--940\\,K, for \\epsba\\ and Bb respectively, with surface gravities of log g=5.25 and 5.50. Finally, we show that spectroscopically determined effective temperatures and surface gravities for ultra-cool dwarfs can lead to underestimated masses even where precise luminosity constraints are available.\\thanks{The full resolution spectra of both brown dwarfs are available in electronic form at the CDS via anonymous ftp to cdsarc.u-strasbg.fr}} ",
        "introduction": "\\label{sec:intro}  The characterisation of low-mass stars and brown dwarfs is important for studies of substellar and planetary atmospheres, the reliable application of low-mass evolutionary models, and the derivation of the full initial mass function. With over five hundred L and over one hundred T dwarfs now known\\footnote{http://dwarfarchives.org - the M, L, and T dwarf compendium maintained by Chris Gelino, Davy Kirkpatrick, and Adam Burgasser.}, statistical studies of global properties and detailed studies of the closest objects are now possible.  Binary systems have an important role to play. They allow the determination of dynamical masses, provide a laboratory in which objects with the same age and chemical composition may be compared, and, where they have main-sequence companions, provide external constraints of metallicity and age which isolated objects lack, breaking the substellar mass-luminosity-age degeneracy.   To fully constrain the evolutionary models of substellar objects \\citep[e.g.][]{Burrows:1997, Baraffe:2003, Saumon:2008}, it would be most useful to determine the bolometric luminosity, radius, mass, and age of a range of such objects. Bolometric luminosities can be determined from photometric and spectroscopic observations across a large wavelength range. Masses can be determined in systems where an orbit may be monitored, and finally, the age and metallicity can be inferred from better characterised stars in the same system. To constrain the atmospheric models of brown dwarfs, it is necessary to acquire high signal-to-noise spectra over as wide a wavelength range as possible, allowing robust estimates of the effective temperature and surface gravity to be made.  The discovery of a distant companion (projected separation $\\sim$1500\\,AU) to the high proper-motion ($\\sim$4.7\\,arcsec/yr) K4.5V star, $\\varepsilon$ Indi, was reported by \\citet{Scholz:2003}. One of our nearest neighbours, \\epsa\\ has a well-constrained parallax from HIPPARCOS \\citep{Perryman:1997} as refined by \\citet{vanleeuwen:2007}, putting the system at a distance of 3.6224$\\pm$0.0037\\,pc. This was followed by the discovery of the companion's binary nature \\citep{McCaughrean:2004}. The proximity of \\epsba, Bb to the Earth means Ba is more than a magnitude brighter than any other known T dwarf, and allows unprecedented, detailed spectroscopic studies of these important template objects.   \\epsba, Bb are uniquely suited to provide key insights into the physics, chemistry, and evolution of substellar sources. Although there are a number of other T dwarfs in binary systems, such as the M4/T8.5 binary Wolf\\,940 \\citep{Burningham:2009} and the T5/T5.5 binary 2MASS\\,1534$-$2952 \\citep{Liu:2008}, \\epsba, Bb has a very well-determined distance, a main-sequence primary star with which to constrain age and metallicity, and a short enough orbit \\citep[nominally $\\sim$15\\,years, ][]{McCaughrean:2004} such that the system and individual dynamical masses can soon be determined \\citep[][in prep.]{McCaughrean:2009, Cardoso:2010}. They are also relatively bright, close enough, and sufficiently separated to allow detailed photometric and spectroscopic studies of both components. Importantly, these two objects roughly straddle the L to T transition \\citep[cf.][]{Burgasser:2009} where the atmospheres of substellar objects alter dramatically. The study of these two coeval objects on either side of the transition will help in understanding the processes effecting the change from cloudy to cloud-free atmospheres. Characterisation of this system allows the mass-luminosity-age relation at low masses and intermediate age to be tested, investigation of the atmospheric chemistry, including vertical up-mixing, and detailed investigation of the species in the atmosphere.     To date, spectroscopic observations of T dwarfs have predominantly been either at low-resolution \\citep[e.g., ][]{Burgasser:2002,Chiu:2006}, which allows spectral classification and overall spectral energy distribution modelling to determine luminosities, or high-resolution studies of relatively small wavelength regions to investigate gravity and effective temperature-sensitive features. For example, \\citet{McLean:2003} presented near-IR spectra at a spectral resolution of R$\\sim$2000 of objects spanning spectral types M6 to T8, and discussed broad changes in spectral morphology and dominant absorbers through the spectral sequence. This was complemented by R$\\sim$20\\,000 $J$-band spectra presented in \\citet{McLean:2007} where many H$_2$O and FeH features were identified and the progression of the $J$-band potassium doublet from M to T dwarfs charted.   Previous studies of ultra-cool dwarfs attempting to constrain low-mass evolutionary models have been hampered by ambiguous ages, possible unresolved binarity, and the difficulty associated with acquiring observations of close, faint companions. For example, observations of AB\\,Dor\\,C have roused some controversy over the applicability of current low-mass evolutionary models, with the assumed age of the system being a major source of disagreement. \\citet{Close:2005} determined a dynamical mass for AB\\,Dor\\,C and, using an assumed age of 30--100\\,Myr, argued that evolutionary models predicted a higher luminosity than was observed. However, \\citet{Luhman:2005} countered with an analysis based on an age of 75--150\\,Myr, finding no significant discrepancy between the observations and models. \\citet{Nielsen:2005} further argued that using their slightly revised age of 70$\\pm$30\\,Myr, the models still under-estimated the mass of this object.  Again, this was disputed by \\citet{Luhman:2006} after a re-reduction of the same data used by \\citet{Close:2005}, and then \\citet{Close:2007} concluded that, based on newly acquired spectra, there was no discrepancy between the observations and models. Despite this apparent rapprochement, the situation may nevertheless be further complicated by the suggestion of \\citet{Marois:2005} that AB Dor C may itself be an unresolved binary.  More recently, \\citet{Dupuy:2009} presented a dynamical mass for the binary L dwarf system HD\\,130948BC which, along with an age estimated from the rotation of the main-sequence parent star, suggests that the evolutionary models predict luminosities 2--3 times higher than those observed.  \\citet{Leggett:2008} used 0.8--4.0\\micron\\ spectra at R$\\sim$100--460 and near- to mid-IR photometry of HN Peg B, a T2.5 dwarf companion to a nearby G0V star, to investigate physical properties including dust grain properties and vertical mixing. In the near-IR, the resolution was too low to study spectral lines in detail. However, by fitting the overall spectral morphology and making use of the longer-wavelength data, they were able to place important constraints on vertical mixing and sedimentation. \\citet{Leggett:2009} also reported the physical properties of four T8--9 dwarfs from fitting observed near- to mid-IR spectral energy distributions with the atmospheric and evolutionary models of \\citet{Saumon:2008}. They discussed the effects of vertical transport of CO and N$_2$ and demonstrated the complementary effects of increasing metallicity and surface gravity.   \\citet{Reiners:2007} analysed high-resolution (R$\\sim$33\\,000) optical spectra of three L dwarfs and the combined \\epsba, Bb system, concluding that although some individual features are not well-matched by the model atmospheres and that significant differences remain for some molecular species and alkali metal features, general features are reproduced. \\citet{Smith:2003} also acquired high resolution (R$\\sim$50\\,000) near-IR spectra of (only) \\epsba\\ in the wavelength ranges 1.553--1.559\\micron\\ and 2.308--2.317\\micron. These were fit with the unified cloud models of \\citet{Tsuji:2002} and effective temperatures of 1400\\,K and 1600\\,K were derived for the two spectral regions. Mid-IR spectroscopy of the unresolved \\epsba, Bb system was also acquired by \\citet{Roellig:2004}, who use evolutionary models along with the luminosities of \\citet{McCaughrean:2004} and an assumed age of 0.8--2.0\\,Gyr to derive effective temperatures and surface gravities and then compared composite spectral models to the observed spectrum. Their predictions were revised by \\citet{Mainzer:2007} who derived effective temperatures of 1210--1250\\,K and 840\\,K, for \\epsba\\ and Bb respectively, under the assumption of a system age of $\\sim$1\\,Gyr.  Finally, \\citet{Kasper:2009} presented R$\\sim$400 near-IR NACO/VLT spectroscopy of \\epsba, Bb which were compared to the evolutionary models of \\citet{Burrows:1997} and the atmospheric models of \\citet{Burrows:2006}. They derived effective temperatures of 1250--1300\\,K and 875--925\\,K, and surface gravities of log g (cm\\,s$^{-1}$) 5.2--5.3 and 4.9--5.1, for \\epsba\\ and Bb respectively, by comparing their observed spectra with their spectral models scaled using the distance and a radius predicted by their evolutionary models. We will discuss these results further in Sect.\\,\\ref{sec:model_diffs} in contrast to our new data.  In this paper we present high signal-to-noise photometry from the $V$- to $M'$-band (0.5--4.9\\micron) and medium resolution spectroscopy from 0.6--5.1\\micron\\ of the individual components of the \\epsba, Bb system. In Sect.\\,\\ref{sec:obs}, we describe the observations and data reduction, including the routines employed to extract the partially-blended photometry and spectroscopy. We re-derive the spectral types of both objects according to the updated classification scheme of \\citet{Burgasser:2006} in Sect.\\,\\ref{sec:spec_class} and discuss constraints imposed by the parent main-sequence star in Sect.\\,\\ref{sec:EpsA_constrain}. We derive the luminosities of both sources in Sect.\\,\\ref{sec:lum} and discuss the preliminary dynamical mass measurement of \\citet[][in prep.]{McCaughrean:2009} in Sect.\\,\\ref{sec:mass}. Our observations are compared to evolutionary models in Sect.\\,\\ref{sec:evo_model_comp} and to atmospheric models in Sect.\\,\\ref{sec:atm_model_comp}. We then put limits on the masses of lower-mass companions in Sect.\\,\\ref{sec:Bc_mass_limits}, and finally the predictions of evolutionary and atmospheric models and previous determinations are compared in Sect.\\,\\ref{sec:model_diffs}. ",
        "conclusions": "\\label{sec:conc}  We have presented the results of a comprehensive photometric and spectroscopic study of the individual components of the nearest known binary brown dwarf system, \\epsba, Bb. The relative proximity of these T1 and T6 dwarfs to the Earth resulted in very high quality data, while archival results for the well-studied parent star, \\epsa, provide invaluable additional information. We find the spectra of these brown dwarfs are best matched by the BT-Settl spectral models with T$_{\\rm{eff}}$=1300--1340\\,K and log g=5.50 for \\epsba\\ and 880--940\\,K and 5.25 for \\epsbb, both with a metallicity of [M/H]=$-0.2$.  COND03 evolutionary model predictions for the masses are significantly inconsistent with the measured system mass if the young age range of 0.8--2.0\\,Gyr suggested by \\citet{Lachaume:1999} is used. We find that a system age of 3.7--4.3\\,Gyr is necessary for the COND03 evolutionary models to be consistent with the measured system mass at the observed luminosities, and a review of the literature finds evidence supporting an age of $\\sim$5\\,Gyr for \\eps\\ A\\@. In the age range 3.7--4.3\\,Gyr, the COND03 models predict effective temperatures in the range 1352--1385\\,K and 976--1011\\,K, for Ba and Bb, respectively.  It is clear that there are several areas in which the atmospheric models currently do not reproduce observations and a more detailed analysis of these issues will be the subject of future work. They include the strength and shape of the wide absorption by \\ion{K}{I} and \\ion{Na}{I} in the optical, the possible formation of feldspars in mid-late T dwarfs, and the reaction rates of CO and CH$_{4}$. In addition, the spectral shape in the $L$-band caused by CH$_4$ absorption is poorly reproduced, as is also the case for CH$_4$ absorption at $\\sim$1.6\\micron. The $M$-band spectra, although low resolution, also show that the atmospheric models significantly over-estimate the flux in this region. While the flux levels of the near-IR peaks can be reasonably reproduced, the level of absorption between the peaks tends to be problematic. In particular, we find a feature at 1.35--1.40\\micron\\ in both our object spectra which is not predicted in the atmospheric models.   Neither source has detectable atomic \\ion{Li}{I} absorption at 6708\\,\\AA. The absence of lithium in the more massive component is consistent with the revised, higher age estimates coupled with its probable dynamical mass, while the lack of absorption in the cooler source is expected from its low effective temperature, where lithium is incorporated into molecules.  Although there is significant room for improvement in the atmospheric models, the current match to \\epsba\\ and Bb is nevertheless impressive. When new data on methane opacities become available, we will be able to better reproduce the observed spectra and more reliably compare these spectral models to spectra of objects with less well-constrained physical parameters. Additionally, when these updated atmospheric models are incorporated into the evolutionary models, a fully self-consistent comparison will be possible.  Finally, when the individual dynamical masses become available and if we can obtain a reliable estimate of the age of this system, based on asteroseismological observations of the parent star \\epsa, then \\epsba\\ and Bb will become invaluable benchmark objects with a full set of physical parameters which newer models will have to reproduce, making them more reliable for analysing the properties of isolated ultra-cool field dwarfs.  The predictions of the evolutionary models using luminosity and mass constraints are somewhat different to the derived effective temperature and surface gravity from fitting atmospheric models to observed spectra. These differences may be resolved when the newer atmosphere models are incorporated into the evolutionary models. However, it seems that derivations of the mass of cool brown dwarfs are uncertain even where estimates of the effective temperature, surface gravity, and luminosity exist. We therefore caution against the over-analysis of predicted brown dwarf masses at this time."
    },
    "0911/0911.4887_arXiv.txt": {
        "abstract": "{} {We intended to measure the radial velocity curve of the supergiant companion to the eclipsing high mass X-ray binary pulsar  {EXO~1722--363} and hence determine the stellar masses of the components.} {We used a set of archival K$_{\\rm s}$-band infrared spectra of the counterpart to  {EXO~1722--363} obtained using ISAAC on the VLT, and cross-correlated them in order to measure the radial velocity of the star.} {The resulting radial velocity curve has a semi-amplitude of $24.5 \\pm 5.0$~km~s$^{-1}$. When combined with other measured parameters of the system, this yields masses in the range 1.5 $\\pm$ 0.4 - 1.6 $\\pm$ 0.4 ~M$_{\\odot}$ for the neutron star and 13.6 $\\pm$ 1.6 - 15.2 $\\pm$  1.9 ~M$_{\\odot}$ for the B0--1 Ia supergiant companion. These lower and upper limits were obtained under the assumption that the system is viewed edge-on (i = 90$^\\circ$) for the lower limit and the supergiant fills its Roche lobe ($\\beta = 1$) for the upper limit respectively. The system inclination is constrained to $i>75^{\\circ}$ and the Roche lobe-filling factor of the supergiant is $\\beta>0.9$. Additionally we were able to further constrain our distance determination to be 7.1 $\\le$ d $\\le$ 7.9 kpc for  {EXO~1722--363}. The X-ray luminosity for this distance range is 4.7 $\\times$ 10$^{35}$ $\\le$ L$_{\\rm X}$ $\\le$ 9.2 $\\times$ 10$^{36}$ erg s$^{-1}$.} { {EXO~1722--363} therefore becomes the seventh of the ten known eclipsing X-ray binary pulsars for which a dynamical neutron star mass solution has been determined. Additionally  {EXO~1722--363} is the first such system to have a neutron star mass measurement made utilising near-infrared spectroscopy.} {} ",
        "introduction": "The precise form of the neutron star (NS) equation of state is still unknown. Despite much theoretical work  aimed at determining this fundamental aspect of astrophysics, to eliminate some of the contending theories we must turn to observational data. Presently the only means of determining the mass of neutron stars in accretion driven systems is by observing eclipsing X-ray binary pulsars. Unfortunately only 10 such systems are currently known, and only 6 of these have previously had mass measurements made (e.g. \\cite{ash99}, \\cite{quaintrell03}, \\cite{valbaker05}).  In this paper we present the preliminary results from our on-going work on the mass of the High Mass X-ray Binary (HMXB) accretion driven pulsar  {EXO 1772--363}. The counterpart star within this HMXB is heavily obscured and reddened, necessitating for the first time the utilisation of near-infrared spectroscopy to construct the Radial Velocity (RV) curve in order to obtain an accurate mass solution.  Observations made of  {EXO~1722--363} (alternatively designated IGR J17252--3616), in 1987 by the \\emph{Ginga} X-ray satellite were the first to detect pulsations. These pulsations were found to have a $413.9 \\pm 0.2$s period \\citep{tawara89}. Over an 8 hour period the source appeared to vary substantially in flux, decreasing from 2~mcrab to 0.2 -- 0.3~mcrab within the 6--21~keV band. It was subsequently found that the X-ray flux remained persistent from 20 -- 60~keV, however a cutoff point was found above this flux level in which the source was undetectable.  {EXO~1722--363} at maximum flux and assuming a distance of 10 kpc, had a luminosity calculated as $5 \\times 10^{36}$~erg~s$^{-1}$ \\citep{tawara89}. Later observations by the \\emph{Rossi X-ray Timing Explorer} (RXTE) revealed the eclipsing nature of this system with the eclipse duration determined as $1.7 \\pm 0.1$~days \\citep{corbet05}. Subsequent observations by \\emph{INTEGRAL} followed up by \\emph{XMM-Newton} in 2004 led to a further refinement of the spin and orbital periods to $413.851\\pm0.004$~s and $9.7403\\pm0.0004$~days respectively \\citep{thompson07}.  \\emph{XMM-Newton} observations allowed the source position to be determined more precisely (with an uncertainty of $4^{\\prime\\prime}$) at RA(2000.0) = $17^h 25^m 11.4^s$ and Dec = $-36^\\circ 16 ^\\prime 58.6 ^{\\prime\\prime}$.  {EXO~1722--363} lies within the Galactic plane and as it is heavily reddened, unsurprisingly the counterpart star could not be detected optically. An infrared counterpart was found lying $1^{\\prime\\prime}$ from the X-ray source position \\citep{zurita06} with a corresponding entry in the 2MASS catalogue, 2MASS~J17251139--3616575 (with JHK magnitudes J = 14.2, H = 11.8 and K$_s$ = 10.7). Examination of near infrared K-band spectra obtained with the ESO \\emph{ISAAC} instrument led to our determination of the spectral classification of the mass donor as B0 -- B1 Ia \\citep{mason09}. ",
        "conclusions": "Although the 11 spectra reported here are of relatively low quality, and few in number, they still allow us to make a preliminary determination of the orbit of the supergiant in  {EXO~1722--363} and make a first measurement of the dynamical masses of the stellar components. The results are encouraging for a number of reasons. First, the resulting neutron star mass is consistent with the canonical mass of 1.4~M$_{\\odot}$ measured in most other eclipsing HMXBs, except for that in Vela X-1, \\citep{quaintrell03}. Second, the measured mass and radius of the supergiant, $M \\sim 13 - 15$~M$_{\\odot}$ and $R \\sim 25 - 28$~R$_{\\odot}$, support the B0-1 Ia spectral classification that we have previously determined \\citep{mason09}. This is illustrated by the Hertzsprung-Russell diagram plotted in Fig. \\ref{evo_track}, which shows a close correspondence between the system primary and the properties of other galactic field BSGs \\citep{searle08}. While the similarity in temperature is to be expected - the value for the primary was adopted on the basis of its spectral type, which in turn has been calibrated by the analysis of \\cite{searle08}, the radii for the field BSGs were determined via non-LTE model atmosphere analysis, while that for the primary is instead determined dynamically.  The measurement of the stellar radius and hence bolometric luminosity, has in turn has allowed a more precise determination of the distance to the system by comparison to its observed photometric magnitude and reddening. The refined distance to {EXO~1722--363} of 7.1 - 7.9 kpc results in  an X-ray luminosity ranging from {L$_{\\rm {X_{min}}}$ = 0.47} $\\times$ ~10$^{36}$ to L$_{\\rm {X_{max}}}$ = 9.2 $\\times$ ~10$^{36}$ erg s$^{-1}$. However,  due to the nature of the archive observations used for this work, the uncertainties on the mass and radius parameters are still rather large;  it is our intention in the  near future to propose and obtain more accurate VLT/ISAAC observations to further constrain the orbital solution parameters for this HMXB system.  Finally, comparison to evolutionary tracks in Fig. \\ref{evo_track} might suggest the primary had an initial progenitor mass of $\\sim$ 35 - 40M$_{\\odot}$ and hence the neutron star originated in a more massive star. However we caution that the binary is highly likely to have undergone at least one episode of mass transfer in the past, rendering such a conclusion highly uncertain. As an exemplar we cite  {GX 301-2}, a HMXB composed of a NS and a B hypergiant with a spectroscopic mass of 43$\\pm$10M$_{\\odot}$ \\citep{kaper06}. However \\citep{wellstein99} propose a formation scenario in which two stars of comparable initial masses evolved via quasi conservative mass transfer into the current configuration  post supernova; hence determining progenitor masses for both primary and neutron star based on the current system parameters is non trivial.   \\begin{figure} \\includegraphics[width=9cm]{fig3_res.eps} \\begin{center} \\caption{Position of  {EXO 1722-363} on the Hertzsprung-Russell \\citep{searle08} diagram alongside a sample of O and B supergiants \tfrom differing locations, Galactic B supergiants, \\cite{crowther06}, SMC B supergiants (\\cite{trundle04};          \\cite{trundle05}) and Galactic O stars, \\cite{repolust04}. These are overplotted together with solar metallicity evolutionary tracks from \\cite{meynet00}. Also shown is the lower and upper limits on the luminosity of  {EXO 1772-363}.} \\label{evo_track} \\end{center} \\end{figure}"
    },
    "0911/0911.3966_arXiv.txt": {
        "abstract": "We investigate non-axisymmetric low frequency modes of a rotating and magnetized neutron star, assuming that the star is threaded by a dipole magnetic field whose strength at the stellar surface, $B_0$, is less than $\\sim 10^{12}$G, and whose magnetic axis is aligned with the rotation axis. For modal analysis, we use a neutron star model composed of a fluid ocean, a solid crust, and a fluid core, where we treat the core as being non-magnetic assuming that the magnetic pressure is much smaller than the gas pressure in the core. For non-axisymmetric modes, spheroidal modes and toroidal modes are coupled in the presence of a magnetic field even for a non-rotating star. Here, we are interested in low frequency modes of a rotating and magnetized neutron star whose oscillation frequencies are similar to those of toroidal crust modes of low spherical harmonic degree and low radial order. For a magnetic field of $B_0\\sim 10^7$G, we find Alfv\\'en waves in the ocean have similar frequencies to the toroidal crust modes, and we find no $r$-modes confined in the ocean for this strength of the field. We calculate the toroidal crustal modes, the interfacial modes peaking at the crust/core interface, and the core inertial modes and $r$-modes, and all these modes are found to be insensitive to the magnetic field of strength $B_0\\ltsim10^{12}$G. We find the displacement vector of the core $l^\\prime=|m|$ $r$-modes have large amplitudes around the rotation axis at the stellar surface even in the presence of a surface magnetic field $B_0\\sim10^{10}$G, where $l^\\prime$ and $m$ are the spherical harmonic degree and the azimuthal wave number of the $r$-modes, respectively. We suggest that millisecond X-ray variations of accretion powered X-ray millisecond pulsars can be used as a probe into the core $r$-modes destabilized by gravitational wave radiation. If the $l^\\prime=|m|=2$ $r$-mode is excited, we will have the pulsation of the frequency $\\sim4\\Omega/3$ with $\\Omega$ being the spin frequency of the star. ",
        "introduction": "Oscillation of strongly magnetized neutron stars has attracted an intense attention in recent years, particularly motivated by the discovery of quasi-periodic oscillations (QPOs) of magnetar candidates (e.g., Woods \\& Thompson 2006), which are believed to be one of the observational manifestations of global oscillations of neutron stars that have a strong global magnetic field of order of $B_0\\sim 10^{15}{\\rm G}$ at the stellar surface (e.g., Duncun 1998, Israel et al 2005; Strohmayer \\& Watts 2005, 2006; Watts \\& Strohmayer 2006). Thus, recent theoretical studies of the oscillations of magnetized neutron stars have been mainly concerned with the stars having extremely strong surface magnetic fields $B_0\\gtsim 10^{15}$G, and these studies have suggested that the QPOs observed in the magnetar candidates are attributable to the toroidal crust modes and Alfv\\'en modes of the neutron stars (e.g., Piro 2005; Glampedakis, Samuelsson \\& Andersson 2006; Lee 2007, 2008; Sotani, Kokkotas \\& Stergioulas 2008; Cerd\\'a-Dur\\'an, Stergioulas \\& Font 2009; Colaiuda, Beyer \\& Kokkotas 2009; Bastrukov et al 2009; Sotani \\& Kokkotas 2009). For toroidal modes of strongly magnetized neutron stars, however, it is intriguing, from the theoretical point of view, that Lee (2008) obtained discrete frequency spectra of the magnetic modes, but Sotani, Kokkotas \\& Stergioulas (2008), Cerd\\'a-Dur\\'an, Stergioulas \\& Font (2009), and Colaiuda, Beyer \\& Kokkotas (2009), using a different numerical method from that used by Lee (2008), suggested the existence of continuum frequency spectra of the modes, as originally discussed by Levin (2006, 2007).     The burst oscillation observed in low mass X-ray binaries (LMXBs) can be another example in which a magnetic field plays an important role in the oscillations of neutron stars, although the strength of the field at the surface of the neutron stars in LMXBs is thought to be less than $\\sim10^{10}$G, much weaker than that for the magnetar candidates. For the burst oscillation, the hot spot model (e.g., Strohmayer et al 1997; Cumming \\& Bildsten 2001; Cumming et al 2002) and the Rossby wave model (Heyl 2004; Lee 2004) have been proposed, but it seems none of the models is accepted as the one that fully explains the observational properties of the burst oscillation. In the Rossby wave model, Heyl (2004, 2005) , Lee (2004), and Lee \\& Strohmayer (2005) have examined the possibility that the burst oscillation is produced by low frequency Rossby waves, called $r$-modes in the astrophysical literature, propagating in the surface fluid region (ocean), disregarding the effects of a magnetic field on the low frequency waves. We note that Bildsten \\& Cutler (1995) discussed the modal properties of $g$ modes propagating in the fluid ocean above the solid crust of mass accreting neutron stars as a possible mechanism responsible for the $\\sim 6$Hz QPOs observed in LMXBs. In their paper, on the assumption that the magnetic pressure, $p_B$, is much smaller than the gas pressure, $p$, the critical strength of a magnetic field, $B_c$, below which the magnetic field has no significant effects on the $g$-modes, was estimated to be $B_c\\sim2\\times 10^{10}$G for the accreted envelopes composed of carbon (see also Piro \\& Bildsten 2005). Since their argument is based on a local analysis and on the assumption $p_B\\ll p$, we think it useful to carry out a global modal analysis of low frequency waves propagating in the magnetized fluid ocean of a neutron star.  Accretion powered millisecond X-ray pulsars in LMXBs show small amplitude, almost sinusoidal X-ray time variations, the dominant period of which is thought to be the spin period of the underlying neutron stars. Lamb et al (2009) argued that the millisecond X-ray variations are produced by an X-ray emitting hot spot located at a magnetic pole rotating with the star, and that so long as the symmetry center of the hot spot is only slightly off the rotation axis the X-ray variations produced have small amplitudes and become almost sinusoidal. They also suggested that if the hot spot is located close to the rotation axis, a slight drift of the hot spot away from the rotation axis leads to appreciable changes in the amplitudes and phases of the X-ray variations. Lamb et al (2009) pointed out that a temporal change in mass accretion rates and hence the radius of the magnetosphere, for example, can cause such a drift of the hot spot. We think it is also interesting to consider the possibility that global oscillations of the neutron stars work as a mechanism that perturbs the hot spot periodically.    In this paper, to calculate global oscillations of a rotating and magnetized neutron star, we use the method of series expansion, in which the angular dependence of the perturbations is represented by series expansion in terms of spherical harmonic functions with different spherical harmonic degrees $l$ for a given azimuthal wave number $m$ (e.g., Lee \\& Strohmayer 1996, Lee 2007). We calculate low $m$, low frequency modes of a neutron star composed of a fluid ocean, a solid crust, and a fluid core, where the star is assumed to be threaded by a dipole magnetic field, the strength $B_0$ of which at the surface is smaller than $\\sim10^{12}$G. The method of calculation we employ is presented in \\S 2, and the numerical results are described in \\S 3, and we discuss a local analysis for low frequency modes in the magnetized fluid ocean in \\S 4, and we give discussion and conclusion in \\S 5. ",
        "conclusions": "Lamb et al (2009) proposed that the small amplitude, almost sinusoidal millisecond X-ray pulsation observed in accretion powered millisecond X-ray pulsars may be well explained by the hot spot model, in which the hot spots are assumed to be located at the magnetic poles, which are nearly aligned with the rotation axis. As discussed by Lamb et al (2009), even a small drift of the hot spot could produce appreciable changes in pulsation amplitudes of the X-ray pulsation. If this proposition is correct, it is interesting to pursue a possibility of using the millisecond X-ray pulsation to probe the core $r$-modes excited by gravitational wave radiation (Andersson 1998; Friedman \\& Morsink 1998). In fact, if the core $r$-modes with $|m|\\ge2$ are excited by the emission of gravitational wave, since the $l^\\prime=m=2$ $r$-mode, which is the most strongly destabilized mode among the $r$-modes (e.g., Lockitch \\& Friedman 1999; Yoshida \\& Lee 2000a), produces the surface displacement vector $\\pmb{\\xi}$ whose horizontal and toroidal components at the surface have large amplitudes around the rotation axis as shown by Figure 6, the hot spot could suffer periodic disturbance from the $r$-mode. Note that the $r$-mode induced temperature perturbation, the surface pattern of which may be proportional to $X^r\\left(\\theta\\right)e^{{\\rm i}m\\phi}$, might generate X-ray variations, the amplitudes of which should be very small. If we write the oscillation frequency of $r$-modes as \\be {\\omega/\\Omega}=\\kappa_0+\\kappa_2\\bar\\Omega^2+O\\left(\\bar\\Omega^4\\right), \\ee the coefficient $\\kappa_0$ for the modes is given by \\be \\kappa_0=2m/ \\left[l^\\prime\\left(l^\\prime+1\\right)\\right], \\ee and the coefficient $\\kappa_2$ may depend on the equation of state and the deviation from the isentropic stratification in the core (e.g., Yoshida \\& Lee 2000a,b), where $\\bar\\Omega=\\Omega/\\Omega_0$. Since the neutron star core is nearly isentropic such that $N^2\\sim0$, we only have to consider the $l^\\prime=|m|$ $r$-modes, for which we have $\\omega\\approx\\kappa_0\\Omega=2\\Omega/\\left(|m|+1\\right)$, and we obtain the frequency $\\omega\\approx 2\\Omega/3$ for $m=2$ in the corotating frame of the star and the frequency $\\sigma\\equiv\\omega-m\\Omega\\approx -4\\Omega/3$ in an inertial frame. It may be interesting to point out that if the $r$-mode of $l^\\prime=m=1$ is also excited by some mechanism, this $r$-mode can produce long period variations in an inertial frame since the inertial frame frequency $\\sigma\\sim0$ for this mode. Although no detection of periodicities whose frequency is approximately equal to $4\\Omega/3$ has so far been reported, if we detect the periodicities produced by the core $r$-modes of $l^\\prime=|m|$ in the X-ray millisecond pulsation, we can use the frequency deviation given by $\\Delta\\bar\\omega\\equiv\\bar\\omega-\\kappa_0\\bar\\Omega\\approx\\kappa_2\\bar\\Omega^3$ to derive information about the equations of state and the thermal stratification in the core. The detectability of the $r$-mode pulsation may depend on the thickness of the crust, the property of the fluid ocean, and the strength of the magnetic field and so on.     We have calculated non-axisymmetric low frequency modes of a rotating and magnetized neutron star, where we used a neutron star model composed of a surface fluid ocean, a solid crust, and a fluid core. We have assumed that the star is threaded by a dipole magnetic field but the fluid core can be treated as a non-magnetic region. For this model, we found that for a magnetic field of strength $B_0\\sim10^7$G, Alfv\\'en waves in the surface ocean come in as low frequency modes, which largely modify the gravito-inertial modes in the ocean. However, the oscillation frequencies of the modes in the ocean are dependent on $j_{\\rm max}$, and they do not reach any good convergence even if $j_{\\rm max}$ is increased to $\\sim20$. At present it is not clear that these ocean modes will converge to discrete modes with real frequencies in the limit of $j_{\\rm max}\\rightarrow \\infty$. We also found that no $r$-modes, which are confined to the surface ocean, can be found in the presence of the weak magnetic field. If this is also the case for the accreted envelopes expected in mass accreting neutron stars in binary systems, we need to reconsider the Rossby wave model for the burst oscillation in LMXBs (Heyl 2004, Lee 2004). On the other hand, the toroidal crust modes and the interfacial modes at the core/crust interface, which show good convergence for $j_{\\rm max}\\gtsim10$, are found to be insensitive to magnetic fields of strength $B_0\\ltsim10^{12}$G. We also find that the core $r$-modes and inertial modes are not affected by the magnetic field even if their eigenfunctions extend to the surface through the magnetized crustal and envelope regions. However, this will not be the case for neutron stars having magnetic fields as strong as $B_0\\sim10^{15}$G (e.g., Lee 2008).  In the present paper, we employed for modal analysis a low mass neutron star model having a thick solid crust and a cold and thin surface ocean, for which the toroidal crust modes of low radial order and low spherical harmonic degree are well separated from the $f$- and $p$-modes. If we use more massive neutron stars with a hot accreted fluid envelope and a thin solid crust, the frequencies of the crust modes of low radial order and those of the $f$- and $p$-modes may overlap, leading to more complicated frequency spectra. We think it necessary to conduct similar modal analyses for such neutron star models to clarify the properties of the ocean modes in the presence of a magnetic field and to examine the possibility that the small amplitude millisecond X-ray pulsation can be used as a probe into the core $r$-modes.     \\begin{appendix}"
    },
    "0911/0911.3410.txt": {
        "abstract": "We assess the validity of a single step Godunov scheme for the solution of the magneto-hydrodynamics equations in more than one dimension. The scheme is second-order accurate and the temporal discretization is based on the dimensionally unsplit Corner Transport Upwind (CTU) method of Colella. The proposed scheme employs a cell-centered representation of the primary fluid variables (including magnetic field) and conserves mass, momentum, magnetic induction and energy. A variant of the scheme, which breaks momentum and energy conservation, is also considered. Divergence errors are transported out of the domain and damped using the mixed hyperbolic/parabolic divergence cleaning technique by Dedner et al. (J. Comput. Phys., 175, 2002). The strength and accuracy of the scheme are verified by a direct comparison with the eight-wave formulation (also employing a cell-centered representation) and with the popular constrained transport method, where magnetic field components retain a staggered collocation inside the computational cell. Results obtained from two- and three-dimensional test problems indicate that the newly proposed scheme is robust, accurate and competitive with recent implementations of the constrained transport method while being considerably easier to implement in existing hydro codes.  %This paper aims to compare two different numerical schemes %for computational MHD in multidimensions. %In the constrained transport method, the magnetic field %has a staggered representation whereby field components live %on the face they are normal to while the remaining variables %retain their usual collocation at the cell center. %This provides a framework by which the induction equation %is more naturally updated using Stoke\u0092s theorem and the divergence-free %condition is fulfilled to machine accuracy. ",
        "introduction": "\\label{sec:intro} % % % %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%  A primary aspect in building stable and robust Godunov type schemes for the numerical solution of the compressible magnetohydrodynamics (MHD) equations relies on an accurate way to control the solenoidal property of the magnetic field while preserving the conservation properties of the underlying physical laws. Failure to fulfill either requisite has been reported as a potential hassle leading to unphysical effects such as plasma acceleration in the direction of the field, incorrect jump conditions, wrong propagation speed of discontinuities and odd-even decoupling, see \\cite{Toth00,BK04}. A comprehensive body of literature has been dedicated to this subject and several strategies to enforce the $\\nabla\\cdot\\vec{B}=0$ condition in Godunov-type codes have been proposed, see for example \\cite{ZMC94, RJF95, RMJF98, BS99, Toth00} and, more recently, \\cite{Balsara04, LdZ04, GS05, Rossmanith06, MBMR08}. The robustness of one method over another can be established on a practical base by extensive numerical testing, see \\cite{Toth00, BK04}.  In a first class of schemes, the magnetic field is discretized as a cell-centered quantity and the usual formalism already developed for the Euler equation can be extended in a natural way. Cell-centered methods are appealing since the extensions to adaptive and/or unstructured grids are of straightforward implementation. Moreover, the same interpolation scheme and stencil used for the other hydrodynamic variables can be easily adapted since all quantities are discretized at the same spatial location, thus facilitating the extension to schemes possessing higher than second order accuracy. Unfortunately, numerical methods based on a cell-centered discretization do not naturally preserve Gauss's law of electromagnetism, even if $\\nabla\\cdot\\vec{B}=0$ initially. In the approach suggested by Powell \\cite{Powell94, Powell99}, Gauss's law for magnetism is discarded in the derivation of the MHD equations and the resulting system of hyperbolic laws is no longer conservative by the appearance of a source term proportional to $\\nabla\\cdot\\vec{B}$. Although the source term should be physically zero at the continuous level, Powell showed that its inclusion changes the character of the equations by introducing an additional eighth wave corresponding to the propagation of jumps in the component of magnetic field normal to a given interface. A different approach is followed in the projection scheme \\cite{BB80, ZMC94, RJF95, Crockett05}, where a Helmholtz-Hodge decomposition is applied to resolve $\\vec{B}$ as the sum of an irrotational and a solenoidal part, associated with scalar and vector potentials. A cleaning step allows to recover the divergence-free magnetic field by subtracting the unphysical contribution coming from the irrotational component at the extra cost of solving a Poisson equation. In the approach of Dedner et al. \\cite{Dedner02}, the divergence-free constraint is enforced by solving a modified system of conservation laws where the induction equation is coupled to a generalized Lagrange multiplier. Dedner et al. showed that the choice of mixed hyperbolic/parabolic correction offers both propagation and dissipation of divergence errors with the maximal admissible characteristic speed, independently of the fluid velocity. This approach preserves the full conservation form of the original MHD system at the minimal cost of introducing one additional variable in the system and will be our scheme of choice. Finally, Torrilhon \\cite{Torrilhon05} (see also \\cite{AT08}) showed a general procedure to modify the inter-cell fluxes in the framework of a flux distribution scheme that preserves the value of a certain discrete divergence operator in each control volume.  %In cell-centered methods the divergence-free constraint is preserved %either at the truncation error of the scheme or it is exactly %zero in some particular $\\nabla\\cdot\\vec{B} = 0$ at the %truncation level. Alternative approaches suggested by Toth (Field CD).  A different strategy is followed in the constrained transport (CT) methods, originally devised by \\cite{EH88} and later built into the framework of shock-capturing Godunov methods by a number of investigators, e.g., \\cite{BS99, Balsara04, LdZ04, GS05, GS08}. In this class of schemes, the magnetic field has a staggered representation whereby the different components live on the face they are normal to. Hydrodynamic variables (density, velocity and pressure) retains their usual collocation at the cell center. CT schemes preserve the divergence-free condition to machine accuracy in an integral sense since the magnetic field is treated as a surface averaged quantity and thus more naturally updated using Stokes' theorem. This evolutionary step involves the construction of a line-averaged electric field along the face edges, thereby requiring some sort of reconstruction or averaging of the electromotive force from the face center (where different components are usually available as face centered upwind Godunov fluxes) to the edges. A variety of different strategies have been suggested, including simple arithmetic averaging \\cite{BS99, RMJF98}, solution of 2-D Riemann problems \\cite{LdZ04, FHT06} or other somewhat more empirical approaches \\cite{GS05,GS08,LD09}. The staggered collocation of magnetic and electric field variables in CT schemes makes their extension to adaptive grids rather arduous and costly. Besides, significant efforts have to be spent in order to develop schemes with spatial accuracy of order higher than second. An alternative constrained transport method, based on the direct solution of the magnetic potential equation (thus avoiding staggered grids), has been presented by \\cite{Rossmanith06}.  In the present work we propose a new fully unsplit Godunov scheme for multidimensional MHD, based on a combination of the Corner Transport Upwind of \\cite{Colella90} and the mixed hyperbolic/parabolic divergence cleaning technique of \\cite{Dedner02} (CTU-GLM). The proposed scheme has second order accuracy in both space and time and adopts a cell-centered spatial collocation (no staggered mesh) of all flow variables, including the magnetic field. The scheme is fully conservative in mass, momentum, magnetic induction and energy and the divergence-free constraint is enforced via a mixed hyperbolic/parabolic correction which avoids the computational cost associated with an elliptic cleaning deriving from a Hodge projection. A variant of the scheme, which introduces divergence source terms breaking the conservative properties of some equations, is also presented. We assess the accuracy and robustness of the scheme by a direct quantitative comparison with the 8-wave formulation of \\cite{Powell99} and the recently developed constrained transport method of \\cite{GS05, GS08}. Other similar implementations may be found in \\cite{FHT06,LD09}. The comparison is conveniently handled using the PLUTO code for computational astrophysics \\cite{Mignone07} where both cell-centered and staggered-mesh implementations are available.  Our motivating efforts are driven by issues of simplicity, efficiency and flexibility. In this sense, the benefits offered by a method where all of the primary flow variables are discretized at the same spatial location considerably ease the extension to adaptive grids, to more complex physics and to schemes with higher than second order accuracy. The latter possibility will be explored in a companion paper.  %Besides the ease of implementation, our motivations reside mainly in the convenience %offered by a scheme where all of the primary flow variables are discretized %at the same location, that is, at the cell midpoint. This endeavor finds its %justification in the extension to adaptive grids and to scheme of higher %than second order accuracy which will be presented in a companion paper.  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% ",
        "conclusions": "\\label{sec:conclusions} % % % %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%  A second-order, cell-centered numerical scheme for the solution of the MHD equations in two and three dimensions has been proposed. Fully unsplit integration resorts to the Corner Transport Method of Colella \\cite{Colella90} and the divergence-free condition is controlled by using a constrained formulation of the MHD equations where the induction equation is coupled to a generalized Lagrange multiplier (GLM, \\cite{Dedner02}). The system is hyperbolic, easy to implement and does not require expensive cleaning projection steps associated with the solution of elliptic problems. The GLM scheme is fully conservative in mass, momentum, energy and magnetic induction, although we have also considered a slightly modified variant (EGLM) which infringes momentum and energy conservation.  In order to assess the reliability and accuracy of the schemes we have performed a number of code benchmarks on standard two- and three-dimensional MHD test problems. Results have been compared with two different numerical schemes: a non-conservative cell-centered method based on the 8-wave formulation (8W, \\cite{Powell99}) and the constrained transport (CT) method where the magnetic field has a staggered collocation. Both the GLM and EGLM schemes give excellent results in terms of accuracy and robustness and do not show, in the tests presented here, any evidence for incorrect jump conditions or wrong wave propagation, as found for the eight wave formulation (in agreement with T\\'oth \\cite{Toth00}). This has been verified on problems involving discontinuous waves and holds true for both the conservative GLM formulation \\emph{and} the EGLM variant which breaks momentum and energy conservation. In this perspective, our results seem to indicate that the presence of source terms in the equations does not necessarily lead to erroneous jumps. Instead, we have found the non-conservative formulation to be more robust for problems involving the propagation of oblique strongly magnetized shocks. Although, this behavior may be attributed to discretization, such a study is beyond the scope of the present paper. The comparison has also revealed an excellent quantitative agreement with the CTU-CT scheme (in the version of \\cite{GS05,GS08}) showing errors with comparable magnitude and similar order of convergence while retaining the desired robustness and stability.  For these reasons, we believe that the proposed CTU-GLM and CTU-EGLM schemes provide excellent competitive alternatives to modern staggered-mesh algorithms while being considerably easier and more flexible in their implementations. Owing to the cell-centered collocation of all of the flow fields, the CTU-GLM scheme can be easily generalized to resistive MHD, adaptive and/or unstructured grids and to higher than second-order spatially-accurate numerical schemes. Some of these issues will be presented in forthcoming papers.  %Local divergence errors are accurately monitored and controlled %at a very small computational cost.   %% The Appendices part is started with the command \\appendix; %% appendix sections are then done as normal sections %% \\appendix  %%"
    },
    "0911/0911.1364_arXiv.txt": {
        "abstract": "{In a previous paper, we found that globular clusters in our Galaxy lie close to a line in the ($\\log{R_e}$, ${\\mathit{SB}}_e$, $\\log{\\sigma}$) parameter space, with a moderate degree of scatter and remarkable axi-symmetry. This implies that a purely photometric scaling law exists, that can be obtained by projecting this line onto the ($\\log{R_e}$, ${\\mathit{SB}}_e$) plane. These photometric quantities are readily available for large samples of clusters, as opposed to stellar velocity dispersion data.}{We study a sample of $129$ Galactic and extragalactic clusters on this photometric plane in the V-band. We search for a linear relation between ${\\mathit{SB}}_e$ and $\\log{R_e}$ and study how the scatter around the best-fit relation is influenced by both age and dynamical environment. We interpret our results in terms of testing the evolutionary versus primordial origin of the fundamental line.}{We perform a detailed analysis of surface brightness profiles, which allows us to present a catalogue of structural properties without relying on a given dynamical model.}{We find a linear relation between ${\\mathit{SB}}_e$ and $\\log{R_e}$, in the form ${\\mathit{SB}}_e = (5.25 \\pm 0.44) \\log{R_e} + (15.58 \\pm 0.28)$, where ${\\mathit{SB}}_e$ is measured in mag/arcsec$^2$ and $R_e$ in parsec. Both young and old clusters follow the scaling law, which has a scatter of approximately $1$ mag in ${\\mathit{SB}}_e$. However, young clusters display more of a scatter and a clear trend in this with age, which old clusters do not. This trend becomes tighter if cluster age is measured in units of the cluster half-light relaxation time. Two-body relaxation therefore plays a major role, together with passive stellar population evolution, in shaping the relation between ${\\mathit{SB}}_e$, $\\log{R_e}$, and cluster age. We argue that the $\\log{R_e}$-${\\mathit{SB}}_e$ relation and hence the fundamental line scaling law does not have a primordial origin at cluster formation, but is rather the result of a combination of stellar evolution and collisional dynamical evolution.}{} ",
        "introduction": "\\label{Sect:Introduction}  In a previous paper \\citep[][hereafter Paper I]{MyFirstPaper}, we demonstrated that Galactic globular clusters (Galactic GCs) occupy a narrow region around a line in the ($\\log{R_e}$, ${\\mathit{SB}}_e$, $\\log{\\sigma}$) parameter space. This fundamental line is located within the fundamental plane of Galactic GCs \\citep[see][]{Djorgovski1995} and its existence was noted by \\cite{Bellazzini} based on principal component analysis techniques.  \\cite{Bellazzini} divided the observed Galactic GCs into two sets of inner GCs and outer GCs with respect to the Sun's orbit, and found that outer GCs have a dimensionality of $1$ in the ($\\log{R_c}$,  ${\\mathit{SB}}_0$, $\\log{\\sigma}$) space\\footnote{The photometric quantities considered by \\cite{Bellazzini} are core/central quantities (i.e., core radius $R_c$ and central surface brightness ${\\mathit{SB}}_0$), as opposed to the half-light ones adopted in Paper I and in this paper (i.e., half-light radius $R_e$ and average surface brightness within it ${\\mathit{SB}}_e$).}, while the dimensionality rises to $2$ for inner GCs. He argued that inner GCs, which are closer to the Galactic bulge and interact more frequently with the disk, are more dynamically disturbed than outer GCs. Therefore, he suggested that dynamical interactions with the environment are more likely to disrupt than to preserve a fundamental line, which was then speculated to be primordial in origin.  By projecting the fundamental line onto the ($\\log{R_e}$, ${\\mathit{SB}}_e$) plane, in Paper I we predicted the existence of a purely photometric scaling law in the form of a linear relation between ${\\mathit{SB}}_e$ and $\\log{R_e}$. This relation can be practically studied for a much larger sample than the fundamental plane, because velocity dispersion data is not required. In particular, bright clusters (both young and old) in the LMC, SMC, and Fornax galaxies may be included in the study, together with Galactic GCs. This allows a comparison to be made between the behaviors in the ($\\log{R_e}$, ${\\mathit{SB}}_e$) plane, of clusters that have significantly different ages and are associated with different environments.  The choice of using only photometric variables allows us to consider a relatively large sample of clusters with a substantial age spread. We are then able to attack the problem of the origin and evolution of the ${\\mathit{SB}}_e$-$\\log{R_e}$ scaling law by studying its dependence on cluster age and to look for evidence for or against the picture in which this scaling law originates primordially and diminishes with time because of dynamical evolution.  In the present paper, we follow Paper I in deriving cluster structural parameters in a model-independent way, in particular by avoiding fitting a given dynamical model to observational data. Some minor changes in the model-independent algorithm used are described in Sect.~\\ref{Sect:TheData}, where a detailed account of the data that we collect from the literature is also given. We compile a catalogue of fraction-of-light radii, average surface brightness, and absolute magnitudes, as a machine-readable table. Central cusps in the surface brightness profile have been observed in both old, Galactic GCs \\citep[][]{NoyolaGalacticSlopes} and extragalactic clusters \\citep[][]{NoyolaExtragalacticSlopes} and are predicted to form after core-collapse by N-body simulations \\citep[][]{Tr2009}. Extended cluster envelopes have also been observed in young clusters in external galaxies \\citep[e.g.,][]{LarsenEnvelopes, SchweitzerEnvelopes}. The simple, single-mass, spherical isotropic \\cite{KingModels} models that are usually taken to fit cluster surface brightness profiles have a flat core and sharp tidal cutoff and therefore have difficulties in modeling clusters with these features. The problem of extended envelopes prompted \\cite{CatMcLaughlin} to refit a sample of Milky Way, LMC, SMC, and Fornax clusters with King models, power-law models, and \\cite{WilsonModels} models, finding that, for a sizeable fraction of the old GC population, King models do not perform more effectively than either \\cite{WilsonModels} or power-law models in fitting surface brightness profiles. More generally, the construction of physically-motivated self-consistent models continues to advance, by including features such as tidal flattening \\citep[e.g., see][]{GBVarri}, but model-independent approaches remain a valid complement to model-based analyses \\citep[see][]{KronMayall}. While a certain dependence on specific assumptions, such as the details of the extrapolation of the external part of surface brightness profiles, cannot be avoided, the catalogue of model-independent GC structural parameters presented in this paper naturally emerges as a tool useful for avoiding the potentially limiting assumptions required by dynamical models. The fundamental line and its scatter are presented in Sect.~\\ref{Sect:Results}. Discussion and conclusions are given in Sect.~\\ref{Sect:DiscussionAndConclusions}. ",
        "conclusions": "\\label{Sect:DiscussionAndConclusions} By considering a sample of $129$ Galactic and extragalactic GCs, we have found a linear relation between ${\\mathit{SB}}_e$ and $\\log R_e$, given by Eq.~(\\ref{Eq:SBelogReRelat}). The relevant coefficient $5.25 \\pm 0.44$ is compatible with $5$, i.e., with cluster absolute magnitude not showing a systematic linear trend with $\\log R_e$. Indeed, when examined in the ($\\log R_e$, $M_V$) plane, the cluster absolute magnitude shows no trend with size \\citep[in contrast to the trend exhibited by dwarf spheroidal galaxies; see][]{vandenBerghNearbyDSPH}.  The typical scatter about the relation is $1$ mag in ${\\mathit{SB}}_e$. Young clusters, of age below $4$ Gyr, exhibit a larger scatter of about $1.4$ mag. GC age seems to be an important factor in driving this scatter, because young clusters are found to display a clear trend of residuals to the relation with age, with a correlation coefficient of $r = 0.79$, while old GCs do not exhibit signs of such a correlation. If we assume that all the massive young clusters included in our sample are genuine GC progenitors, our finding suggests that the scaling law being considered shows signs of evolution with age, i.e., that GCs of different ages differ with respect to this scaling law. Moreover, the larger scatter displayed by young clusters points to an evolutionary origin of the relation between ${\\mathit{SB}}_e$ and $\\log R_e$, as opposed to a primordial one. This result is in contrast to the suggestion of \\cite{Bellazzini} that ``at earlier times, globular clusters populated a line in the three-dimensional S-space, i.e., their original dynamical structure was fully determined by a single physical parameter''.  In principle, modeling the processes that generate the relation between ${\\mathit{SB}}_e$ and $\\log R_e$ would require taking into account passive stellar evolution, internal dynamical phenomena (evaporation and core-collapse driven by two-body relaxation), and dynamical interaction with the environment. We find that the trend with age becomes stronger for young clusters if age is measured in units of the half-light relaxation time of each GC, i.e., if \\emph{dynamical age} with respect to two-body relaxation is considered, the correlation coefficient rises to $0.90$. Given this result, we argue that two-body relaxation is the key ingredient in shaping the evolution of the ${\\mathit{SB}}_e$-$\\log R_e$ scaling law, at least for young clusters. This may not be the case for older clusters, which do not show such a trend. The evolution of these clusters in the ($\\log R_e$, ${\\mathit{SB}}_e$) plane is likely to be influenced by their host environment, but the available data is of limited use in resolving this issue. The distance to the host galaxy is only a simple proxy for the degree of dynamical disturbance that the GC suffers from it, while more detailed orbital data is available only for a sample of $48$ old, Galactic GCs \\citep[][]{AMPOrbits48GC}. In any case, we find no trend between the residuals to the ${\\mathit{SB}}_e$-$\\log R_e$ relation and either distance from the host galaxy or pericenter distance and orbital eccentricity, when these data are available. When GC age is measured in units of the bulge-interaction and disk-interaction destruction times based on the orbital data provided by \\cite{AMPOrbits48GC}, weak ($r \\approx 0.5$) trends are found with ${\\mathit{SB}}_e$-$\\log R_e$ relation residuals. While this suggests that dynamical interaction with the galactic environment may exert an influence on GCs with observable effects in the ($\\log R_e$, ${\\mathit{SB}}_e$) plane, to draw conclusive evidence, the correlation with age alone should be removed from the trend and a larger sample of clusters with accurately measured orbits is necessary.  In contrast to what we noted for fundamental plane residuals in Paper I, we find here that central surface brightness profile slopes measured by \\cite{NoyolaGalacticSlopes} and \\cite{NoyolaExtragalacticSlopes} do not correlate with residuals from the ${\\mathit{SB}}_e$-$\\log R_e$ relation. This result suggests that velocity dispersion information is necessary to understand the mechanism producing cuspy GC surface brightness profiles.  \\emph{Acknowledgments.} We wish to thank S. Degl'Innocenti, W. Harris, M. Lombardi, M. Trenti, E. Vesperini, and D. Kruijssen for their comments and helpful suggestions. This research has made use of the SIMBAD data-base, operated at CDS, Strasbourg, France. Part of this work was carried out at the Kavli Institute for Theoretical Physics, Santa Barbara (CA), while we participated in the Program \\emph{Formation and Evolution of Globular Clusters}."
    },
    "0911/0911.3361_arXiv.txt": {
        "abstract": "{ The mass-loss rate is a key parameter of hot, massive stars. Small-scale inhomogeneities (clumping) in the winds of these stars are conventionally included in spectral analyses by assuming optically thin clumps, a void inter-clump medium, and a smooth velocity field.  To reconcile investigations of different diagnostics (in particular, unsaturated UV resonance lines vs. $\\rm H_{\\alpha}$/radio emission) within such models, a highly clumped wind with very low mass-loss rates needs to be invoked, where the resonance lines seem to indicate rates an order of magnitude (or even more) lower than previously accepted values.  If found to be realistic, this would challenge the radiative line-driven wind theory and have dramatic consequences for the evolution of massive stars. } { We investigate basic properties of the formation of resonance lines in small-scale inhomogeneous hot star winds with non-monotonic velocity fields.} { We study inhomogeneous wind structures by means of 2D stochastic and pseudo-2D radiation-hydrodynamic wind models, constructed by assembling 1D snapshots in radially independent slices. A Monte-Carlo radiative transfer code, which treats the resonance line formation in an axially symmetric spherical wind (without resorting to the Sobolev approximation), is presented and used to produce synthetic line spectra.} { The optically thin clumping limit is only valid for very weak lines. The detailed density structure, the inter-clump medium, and the non-monotonic velocity field are all important for the line formation. We confirm previous findings that radiation-hydrodynamic wind models reproduce observed characteristics of strong lines (e.g., the black troughs) without applying the highly supersonic `microturbulence' needed in smooth models.  For intermediate strong lines, the velocity spans of the clumps are of central importance. Current radiation-hydrodynamic models predict spans that are too large to reproduce observed profiles unless a very low mass-loss rate is invoked. By simulating lower spans in 2D stochastic models, the profile strengths become drastically reduced, and are consistent with higher mass-loss rates. To simultaneously meet the constraints from strong lines, the inter-clump medium must be non-void. A first comparison to the observed Phosphorus V doublet in the O6 supergiant $\\lambda$~Cep confirms that line profiles calculated from a stochastic 2D model reproduce observations with a mass-loss rate approximately ten times higher than that derived from the same lines but assuming optically thin clumping.  Tentatively this may resolve discrepancies between theoretical predictions, evolutionary constraints, and recent derived mass-loss rates, and suggests a re-investigation of the clump structure predicted by current radiation-hydrodynamic models. } {} ",
        "introduction": "\\label{Introduction}  Mass loss through supersonic stellar winds is pivotal for the physical understanding of hot, massive stars and their surroundings. A change of only a factor of two in the mass-loss rate has a dramatic effect on massive star evolution \\citep{Meynet94}.  Winds from these stars are described by the line-driven wind theory \\citep{Castor75,Pauldrach86}, which traditionally assumes the wind to be stationary, spherically symmetric, and homogeneous.  Despite this theory's apparent success \\citep[e.g.,][]{Vink00}, evidence for an inhomogeneous and time-dependent wind has over the past years accumulated, recently summarized in the proceedings from the workshop `Clumping in hot star winds' \\citep{Hamann08} and in a general review of mass loss from hot, massive stars \\citep{Puls08}.  That line-driven winds should be intrinsically unstable was already pointed out by \\citet{Lucy70}, and was later confirmed first by linear stability analyses and then by direct, radiation-hydrodynamic modeling of the time-dependent wind \\citep[e.g.,][]{Owocki84,Owocki88,Feldmeier95,Dessart05}, where the line-driven (or line-deshadowing) instability causes a small-scale, inhomogeneous wind in both density and velocity.  \\textit{Direct observational} evidence of a small-scale, clumped stellar wind has, for O-stars, so far only been given for two objects, $\\zeta$ Pup and HD\\,93129A \\citep{Eversberg98,Lepine08}.  Much \\textit{indirect} evidence, however, has arisen from quantitative spectroscopy, where the standard way of deriving mass-loss rates from observations nowadays is via line-blanketed, non-LTE (LTE: local thermodynamic equilibrium) model atmospheres that include a treatment of both the photosphere and the wind. Wind clumping has been included in such codes (e.g., CMFGEN \\citep{Hillier98}, PoWR \\citep{Grafener02}, FASTWIND \\citep{Puls05}) by assuming statistically distributed \\textit{optically thin} density clumps and a void inter-clump medium, while keeping the smooth velocity law. The major result from this methodology is that any mass-loss rate derived from smooth models and density-squared diagnostics ($\\rm H_{\\alpha}$, infra-red and radio emission) needs to be scaled down by the square root of the clumping factor (which describes the over density of the clumps as compared to the mean density, see Sect.~\\ref{wind_stoch}). For example, \\citet{Crowther02}, \\citet{Bouret03}, and \\citet{Bouret05} have concluded that a reduction of `smooth' mass-loss rates by factors $3 \\dots 7$ might be necessary. Furthermore, from a combined optical/IR/radio analysis of a sample of Galactic O-giants/supergiants, \\citet{Puls06} derived upper limits on observed rates that were factors of $2 \\dots 3$ lower than previous $\\rm H_{\\alpha}$ estimates based on a smooth wind.  On the other hand, the strength of UV resonance lines (`P Cygni lines') in hot star winds depends linearly on the density and is therefore not believed to be directly affected by optically thin clumping. By using the Sobolev with exact integration technique (SEI; cf.~\\citeauthor{Lamers87} 1987) on the unsaturated Phosphorus V (PV) lines, \\citet{Fullerton06} for a large number of Galactic O-stars derived rates that were factors of $10 \\dots 100$ lower than corresponding smooth $\\rm H_{\\alpha}$/radio values (provided PV is the dominant ion in spectral classes O4 to O7). Such large revisions would conflict with the radiative line-driven wind theory and have dramatic consequences for the evolution of, and the feedback from, massive stars \\citep[cf.][]{Smith06,Hirschi08}.  Indeed, a puzzling picture has emerged, and it appears necessary to ask whether the present treatment of wind clumping is sufficient. Particularly the assumptions of optically thin clumps, a void inter-clump medium, and a smooth velocity field may not be adequate to infer proper rates under certain conditions.  \\paragraph{Optically thin vs. optically thick clumps.} \\citet{Oskinova07} used a porosity formalism \\citep{Feldmeier03,Owocki04} to scale the opacity from smooth models and investigate impacts from \\textit{optically thick} clumps on the line profiles of $\\zeta$ Pup. Due to a reduction in the effective opacity, the authors were able to reproduce the PV lines without relying on a (very) low mass-loss rate, while simultaneously fitting the optically thin $\\rm H_{\\alpha}$ line. This formalism, however, was criticized by \\citet{Owocki08} who argued that the original porosity concept had been developed for continuum processes, and that line transitions rather should depend on the non-monotonic velocity field seen in hydrodynamic simulations. Proposing a simplified analytic description to account for this velocity-porosity, or `vorosity', he showed how also this effect may reduce the effective opacity.  In this first paper we attempt to clarify the most important concepts by conducting a detailed investigation on the synthesis of UV resonance lines from inhomogeneous two-dimensional (2D) winds. We create both pseudo-2D, radiation-hydrodynamic wind models and 2D, stochastic wind models, and produce synthetic line profiles via Monte-Carlo radiative transfer calculations.  We account for and analyze the effects from a wind clumped in \\textit{both} density and velocity as well as the effects from a non-void inter-clump medium. Especially we focus on lines with intermediate line strengths, comparing the behavior of these lines with the behavior of both optically thin lines and saturated lines.  Follow-up studies will include a treatment of emission lines (e.g., $\\rm H_{\\alpha}$) and an extension to 3D, and the development of simplified approaches to incorporate effects into non-LTE models.  In Sect.~\\ref{wind} we describe the wind models and in Sect.~\\ref{rt} the Monte-Carlo radiative transfer code.  First results from 2D inhomogeneous winds are presented in Sect.~\\ref{2d}, and an extensive parameter study is carried out in Sect.~\\ref{ps}.  We discuss some aspects of the interpretations of these results and perform a first comparison to observations in Sect.~\\ref{Discussion}, and summarize our findings and outline future work in Sect.~\\ref{Conclusions}. ",
        "conclusions": "\\label{Discussion}  \\subsection{The shapes of the intermediate lines} \\label{shapes}  \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics[angle=90]{fig9.ps}} \\caption{Total, absorption part, and re-emission part line profiles for 1D, smooth models with $\\kappa_0=5.0$ (dashed-dotted lines) and $\\kappa_0=5.0/(2\\eta)$ (solid lines, see Sect.~\\ref{shapes}), and for a 2D, stochastic model with density parameters as the default model and a $\\beta=1$ velocity law (dashed lines).} \\label{Fig:1d_fer} \\end{figure}  For intermediate lines, the shape of the absorption part of the default model differs significantly from the shape of a smooth model (see Fig.~\\ref{Fig:stochs}, the middle plot in the lower-left panel). We showed in Sect.~\\ref{eta} that the shapes could be qualitatively understood by the behavior of $\\eta$. This is further demonstrated here by scaling the line strength parameter of a 1D, smooth model, using a parameterization $\\kappa_0 \\propto \\eta^{-1}$ outside the radius $r=1.3$ where clumping is assumed to start. Fig.~\\ref{Fig:1d_fer} displays the line profiles of 1D, smooth models with $\\kappa_0=5.0$ and $\\kappa_0=5.0/(2\\eta)$. These profiles are compared to those calculated from a `real' 2D stochastic model with density-clumping parameters as the default model, but with a $\\beta=1$ velocity field. $\\eta$ was calculated from Eq.~\\ref{Eq:heff}, using the parameters of the default model and a $\\beta=1$ velocity law, and the factor of 2 in the denominator of the scaled $\\kappa_0$ was chosen so that the \\textit{integrated} profile strength of the 2D model was roughly reproduced. From Fig.~\\ref{Fig:1d_fer} it is clear that the 1D model with scaled $\\kappa_0$ well reproduces the 2D results, indicating that indeed $\\eta$ governs the shape of the line profile. We notice also that these profiles display a completely black absorption dip in the outermost wind, as opposed to the default model with a non-monotonic velocity field (see Fig.~\\ref{Fig:stochs}, the middle plot in the lower-left panel).  This is because the $\\beta$ velocity field does not allow for any clumps to overlap in velocity space (see the discussion in Sect.~\\ref{vel_par}), making the mapping of $\\eta$ almost perfect.  Let us also point out that the line shapes can be somewhat altered by using a different velocity law, e.g., $\\beta \\ne 1$. Such a change would affect the distances between clumps as well as the Sobolev length, and thereby the line shapes of both absorption and re-emission profiles. However, in all cases is the shape of the re-emission part similar in the clumped and smooth models.  \\subsection{The onset of clumping and the blue edge absorption dip} \\label{outer}  \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics[angle=90]{fig10.ps}} \\caption{\\textit{Upper panel:} Density structures of one slice in the default stochastic model (upper), in the default stochastic model with a modified $\\delta t$ (middle, see Sect.~\\ref{outer}), and in FPP (lower). \\textit{Lower panel:} Line profiles for the absorption part of an intermediate line for the default model (solid line), for the default model with a modified $\\delta t$ (dashed line), and for the default model with an ionization structure decreasing with increasing velocity (dashed-dotted line, see text).} \\label{Fig:holes} \\end{figure}  We have used $r=1.3$ as the onset of wind clumping in our stochastic models, which roughly corresponds to the radius where significant structure has developed from the line-driven instability in our RH models.  However, \\citet{Bouret03,Bouret05} analyzed O-stars in the Galaxy and the SMC, assuming optically thin clumps, and found that clumping starts deep in the wind, just above the sonic point. Also \\citet{Puls06} used the optically thin clumping approach, on $\\rho^2$-diagnostics, and found similar results, at least for O-stars with dense winds.  With respect to our stochastic models, the qualitative results from Sects.~\\ref{2d}~and~\\ref{ps} remain valid when choosing an earlier onset of clumping. Quantitatively, the integrated absorption in intermediate lines becomes somewhat weaker, because the clumping now starts at lower velocities, and of course the line shapes in this region are affected as well.  The onset of wind clumping will be important when comparing to observations, as discussed in Sect.~\\ref{cmp_obs}.  The stochastic models that de-saturate an intermediate line generally display an absorption dip toward the blue edge (see Figs.~\\ref{Fig:stochs} and \\ref{Fig:1d_fer}), which has been interpreted in terms of low values of $\\eta$ in the outer wind (see Sect.~\\ref{eta}). However, this characteristic feature (not to be confused with the so-called DACs, discrete absorption components) is generally not observed, and one may ask whether it might be an artifact of our modeling technique.  In the following we discuss two possibilities that may cause our models to overestimate the absorption in the outer wind; the ionization fraction and too low clump separations.  Starting with the former, we have so far assumed a constant ionization factor, $q=1$ (cf. Eq.~\\ref{Eq:kappa0}). This is obviously an over-simplification.  For example, an outwards decreasing $q$ would result in less absorption toward the blue edge. Here we merely demonstrate this general effect, parameterizing $q = v_0/v_{\\beta}$ in the stochastic default model (see Table~\\ref{Tab:mod}), with $v_{\\rm 0}=0.1$ the starting point below which $q=1$.  Fig.~\\ref{Fig:holes} (lower panel, dashed-dotted lines) shows how the absorption in the outer wind becomes significantly reduced.  The temperature structure of the wind is obviously important for the ionization balance.  Whereas an isothermal wind is assumed in POF (see Sect.~\\ref{wind_rh}), the FPP model has shocked wind regions with temperatures of several million Kelvin. To roughly map corresponding effects on the line profiles, we re-calculated profiles based on FPP models assuming $q=0$ in all regions with temperatures higher than $T=10^5\\rm\\,K$, and $q=1$ elsewhere.  Since the hot gas resides primarily in the low-density regions, however, the emergent profiles were barely affected, and particularly intermediate lines remained unchanged.  On the other hand, the X-ray emission from hot stars (believed to originate in clump-clump collisions, see FPP) is known to be crucial for the ionization balance of highly ionized species such as C\\,IV, N\\,V, and O\\,VI \\citep[see, e.g., the discussion in][]{Puls08}. X-rays have not been included here, but could in principle have an impact on our line profiles, by illuminating the over-dense regions and thereby changing the ionization balance. \\citet{Krticka09}, however, find that incorporating X-rays does not influence the PV ionization significantly. Finally, non-LTE analyses including feedback from optically thin clumping have shown that this as well has a significant effect on the derived ionization fractions of, e.g., PV \\citep[]{Bouret05,Puls08b}. To summarize, it is clear that a full analysis of ionization fractions must await a future non-LTE application that includes relevant feedback effects from an inhomogeneous wind on the occupation numbers.  In RH models, the average distance between clumps increases in the outer wind, due to clump-clump collisions and velocity stretching \\citep{Feldmeier97,Runacres02}.  Neglecting the former effect, our stochastic models have clumps much more closely spaced in the outer wind\\footnote{The effect is minor in POF, since these RH models only extend to $r \\sim 5$ (see Sect.~\\ref{wind_rh}).}.  We have therefore modified the default model by setting $\\delta t = 3$ outside a radius corresponding to $v_{\\beta}=0.7$. This is illustrated in the upper panel of Fig.~\\ref{Fig:holes}. The mass loss in the new stochastic model is preserved (because the clumps are more extended, see the figure), and this model now better resembles FPP. Recall that differences in the widths of the clumps are expected, since in the default model $f_{\\rm cl} \\approx f_{\\rm v}^{-1}=4$, whereas in FPP $f_{\\rm cl} \\approx 10$. The corresponding line profile shows how the absorption outside $x \\approx 0.7$ has been reduced, as expected from the higher $\\delta t$.  \\subsection{The velocity spans of the clumps} \\label{vgrad}  \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics[angle=90]{fig11.ps}} \\caption{\\textit{Upper:} Velocity spans of density enhancements in the FPP model (squares) and corresponding $\\beta$ intervals (diamonds). \\textit{Lower:} Three density enhancements and corresponding velocity spans in the FPP model, highlighted as in Fig.~\\ref{Fig:contours}.} \\label{Fig:dvdvb} \\end{figure}  In Sect.~\\ref{cmp} it was found that $|\\delta v| > \\delta v_{\\beta}$ in the RH models.  Fig.~\\ref{Fig:dvdvb}, upper panel, shows the velocity spans of density enhancements (identified as having a density higher than the corresponding smooth value) in the FPP model, and demonstrates that, after structure has developed, $|\\delta v|$ is much higher than $\\delta v_{\\beta}$ throughout the whole wind. These high values essentially stem from the location of the starting points of the density enhancements, which generally lie \\textit{before} the velocities have reached their post shock values (see Fig.~\\ref{Fig:dvdvb}, middle and lower panels).  By using a $\\beta$ velocity law (which in principle corresponds to a stochastic velocity law with $v_{\\rm j}=0$ and $\\delta v = \\delta v_{\\beta}$, see Fig.~\\ref{Fig:vj}) together with the density structure from FPP, we simulated a RH wind with low velocity spans. Indeed, for the corresponding intermediate line the equivalent width of the absorption part was $\\sim 35 \\%$ lower than that of the original FPP model. The strong line, on the other hand, remained saturated, because the ICM in FPP is not void. So, again, the RH models would in parallel display de-saturated intermediate lines and saturated strong lines, were it not for the large velocity spans inside the clumps.  We suggest that the large velocity span inside a shell (clump) is primarily of kinematic origin, and reflects the formation history of the shell. The shell propagates outwards through the wind, essentially with a $\\beta=1$ velocity law \\citep{Owocki88}. Fast gas is decelerated in a strong reverse shock at the inner rim of the shell. The shell collects ever faster material on its way out through the wind. This new material collected at higher speeds resides on the star-facing side, i.e.~at smaller radii, of the slower material collected before. Thus, a negative velocity gradient develops inside the shell. The fact that $|\\delta v| \\gg \\delta v_{\\beta}$ in FPP seems to reflect that the shell is formed at small radii, and then advects outwards maintaining its steep interior velocity gradient\\footnote{Actually, the velocity gradient may further steepen during advection, due to faster gas trying to overtake slower gas ahead of it; however, this effect is balanced by pressure forces in the subsonic postshock domain.}. From this formation in the inner, steeply accelerating wind, velocity spans within the shells up to (a few) hundred $\\rm km\\,s^{-1}$, as seen in Fig.~\\ref{Fig:dvdvb}, appear reasonable.  However, the dynamics of shell formation in hot star winds is very complex due to the creation and subsequent merging of subshells, as caused by nonlinear perturbation growth and the related excitation of harmonic overtones of the perturbation period at the wind base \\citep[see][]{Feldmeier95}. Future work is certainly needed to clarify to which extent the large velocity spans inside the shells in RH models are a stable feature (see also Sect.~\\ref{future}).  \\subsection{3D effects} \\label{3d}  A shortcoming of our analysis is the assumed symmetry in $\\Phi$. The 2D rather than 3D treatment has in part been motivated by computational reasons (see Appendix~\\ref{rt_code}).  More importantly though, we do not expect our \\textit{qualitative} results to be strongly affected by an extension to 3D. Within the broken-shell wind model, all wind slices are treated independently, and distances between clumps increase only in the radial direction.  Therefore the expected outcome from extending to 3D is a smoothing effect rather than a reduction or increase in integrated profile strength (similar to the smoothing introduced by $N_{\\Theta}$, see Sect.~\\ref{ang_dep}). Also, we have shown that the main effect from the inhomogeneous winds is on the absorption part of the line profiles (see, e.g., Sect.~\\ref{shapes}). The formation of this part is dominated by radial photons, especially in the outer wind, because of the dependence only on photons released directly from the photosphere.  This implies that most photons stay within their wind slice, restricting the influence from any additional `holes' introduced by a broken symmetry in $\\Phi$ to the inner wind. Of course, these expectations hold only within the broken shell model, because in a real 3D wind the clumps will, for example, have velocity components also in the tangential directions.   \\subsection{Comparison to other studies} \\label{oskow}  To scale the smooth opacity in the formal integral of the non-LTE atmospheric code PoWR, \\citet{Oskinova07} used a porosity formalism in which both $f_{\\rm v}$ and the average distance between clumps enter. Other assumptions were a void ICM, a smooth $\\beta$ velocity field, and a microturbulent velocity $v_{\\rm t} \\approx 50 \\rm\\,km\\,s^{-1}$, the last identified as the velocity dispersion within a clump. However, a direct comparison between their study and ours is hampered by the different formalisms used for the spacing of the clumps. Here we have used the `broken-shell' wind model as a base (see Sect.~\\ref{wind_stoch}), in which each wind slice is treated independently and the distance between clumps increases only in the radial direction (clumps preserve their lateral angles).  This gives a radial number density of clumps, $n_{\\rm cl} \\propto v^{-1}$, the same as used by, e.g., \\citet{Oskinova06}, when synthesizing X-ray emission from hot stars.  In \\citet{Oskinova07}, on the other hand, the distance between clumps increases in \\textit{all} spatial directions. In a spherical expansion, this gives a radial number density of clumps $n_{\\rm cl} \\propto v^{-1}r^{-2}$, i.e., clumps are distributed much more sparsely within this model, especially in the outer wind. Therefore their choice of $L_{\\rm 0}=0.2$ is not directly comparable with $\\delta t =0.2$ in our models. The shapes of the clumps differ between the two models as well; in \\citeauthor{Oskinova07} clumps are assumed to be `cubes', whereas here the exact shapes of the clumps are determined by the values of the clumping parameters. Despite these differences, our findings confirm the qualitative results of \\citeauthor{Oskinova07} that the line profiles become weaker with an increasing distance between clumps as well as with a decreasing $v_{\\rm t}$. These results may be interpreted on the basis of the effective escape ratio, $\\eta$ (see Eq.~\\ref{Eq:heff}).  Both a decrease in $v_{\\rm t}$ and an increase in the distance between clumps mean that the velocity span covered by a resonance zone becomes smaller when compared to the velocity gap between two clumps (see Fig.~\\ref{Fig:fer}, left panel), leading to higher probabilities for line photons to escape their resonance zones without interacting with the wind material.  An important result of this paper is that models that de-saturate intermediate lines require a non-void ICM to saturate strong lines. This is confirmed by the \\citeauthor{Oskinova07} model, in which the ICM is void and strong lines indeed do not saturate \\citep{Hamann09}.  \\citet{Owocki08} proposed a simplified description of the non-monotonic velocity field to account for vorosity, i.e., the velocity gaps between the clumps. Here, the vorosity effect has been discussed using the quantity $\\eta$ (see Sect.~\\ref{eta}), and we have introduced two new parameters to characterize a non-monotonic velocity field, $\\delta v$ and $v_{\\rm j}$. The reason for introducing a new parameterization is that when using a single velocity parameter, we have not been able to simultaneously meet the constraints from strong, intermediate, and weak lines as listed in Sect.~\\ref{2d}. Tests using a `velocity clumping factor' $f_{\\rm vel}=\\delta v / \\Delta v$ as proposed by \\citet{Owocki08}, together with a smooth density structure, have shown that this treatment indeed can reduce the line strengths of intermediate lines, but that the observational constraints from strong lines may not be met.  Still, the basic concept of vorosity holds within our analysis.  For example, one may phrase the high values of $\\delta v$ in the RH models in terms of insufficient vorosity.  \\subsection{Comparison to observations}  We finalize our discussion by performing a first comparison to observations. The two components of the Phosphorus V $\\lambda\\lambda$1118-1128 doublet are rather well separated, and the singlet treatment used here suffices to model the major part of the line complex. Nevertheless, the two components overlap within a certain region (indicated in Fig.~\\ref{Fig:cmp_obs}), so when interpreting the results of this subsection, one should bear in mind that the overlap is not properly accounted for, but treated as a simple multiplication of the two profiles.  We used observed FUSE spectra (kindly provided by A. Fullerton) from HD\\,210839 ($\\lambda$ Cep), a supergiant of spectral type O6\\,I(n)fp.  When computing synthetic spectra, we first assumed optically thin clumping with a constant clumping factor $f_{\\rm cl}=9$ and a smooth $\\beta=1$ velocity field. $f_{\\rm cl}=9$ agrees fairly well with the analysis of \\citet{Puls06}, who derived clumping factors $f_{\\rm cl}=6.5$ for $r \\approx 1.2 \\dots 4.0$ and $f_{\\rm cl}=10$ for $r \\approx 4.0 \\dots 15$, assuming an un-clumped outermost wind.\\footnote{This stratification has been found to be prototypical for O-supergiants and was, together with its well developed PV P Cygni profiles, the major reason for choosing $\\lambda$~Cep as comparison object instead of, e.g., $\\zeta$~Pup, which displays a somewhat unusual run of $f_{\\rm cl}$.} We took the ionization fraction $q=q(r)$ of PV from \\citet{Puls08b}, calculated with the unified non-LTE atmosphere code FASTWIND for an O6 supergiant, using the Phosphorus model atom from \\citet{Pauldrach01}. The feedback from optically {\\it thin} clumping was accounted for and X-rays were neglected. This ionization fraction was then used as input in our MC-1D code when computing the synthetic spectra.  We assigned a thermal plus a highly supersonic `microturbulent' velocity $v_{\\rm t}=0.05$ (corresponding to 110 km\\,s$^{-1}$), as is conventional in this approach.  The mass-loss rate was derived using the well known relation between $\\kappa_0$ and $\\dot{M}$ \\citep[e.g.,][]{Puls08}. For atomic and stellar parameters, we adopted the same values as in \\citet{Fullerton06}.  The dashed line in Fig.~\\ref{Fig:cmp_obs} represents our fit to the observed spectrum, assuming optically thin clumping, resulting in a mass-loss rate $\\dot{M}=0.24$, in units of $10^{-6}\\rm\\,M_{\\odot}\\,yr^{-1}$.  \\citet{Fullerton06} derived $\\langle q \\rangle \\dot{M} = 0.23$ for this star. Because our clumped FASTWIND model predicts an averaged ionization fraction $\\langle q \\rangle \\approx 0.9$ in the velocity regions utilized by \\citeauthor{Fullerton06}, the two rates are in excellent agreement. On the other hand, \\citet{Repolust04} for HD\\,210839 derived $\\dot{M}=6.9$ from $\\rm H_{\\alpha}$ assuming an unclumped wind, yielding $\\dot{M}_{\\rm H_{\\alpha}}=2.3$ when accounting for the reduction implied by our assumed $f_{\\rm cl}=9$ ($\\dot{M}_{\\rm H \\alpha}=\\dot{M}_{\\rm H \\alpha,sm}f_{\\rm cl}^{-1/2}$). This rate is almost ten times higher than that inferred from PV, and thus results in PV line profiles that are much too strong (see Fig.~\\ref{Fig:cmp_obs}, dashed-dotted line). That is, to reconcile the $\\rm H_{\\alpha}$ and PV rates for HD\\,210839 with models that assume optically {\\it thin} clumps also in PV, we would have to raise the clumping factor to $f_{\\rm cl} > 100$.  In addition to this very high clumping factor, the low rate inferred from the PV lines conflicts with the theoretical value $\\dot{M}=3.2$ provided by the mass-loss recipe in \\citet{Vink00} \\citep[using the stellar parameters of][]{Repolust04}, and is also strongly disfavored by current massive star evolutionary models \\citep{Hirschi08}.  Next we modeled the PV lines using our MC-2D code together with a stochastic 2D wind model. The same clumping factor ($f_{\\rm cl}=9$) and ionization fraction (calculated from FASTWIND, see above) were used.  This time, we assigned $v_{\\rm t}=0.005$, i.e., applied no microturbulence. In previous sections, e.g. \\ref{st} and \\ref{shapes}, we showed that stochastic models generally display a line shape different from smooth models, with a characteristic absorption dip at the blue edge as well as a dip close to the line center. Such shapes are not seen in the PV lines in $\\lambda$ Cep. Thus, to better resemble the observed line shapes, we used different values for $\\delta t$ and $\\fic$ in the inner and outer wind (the former modification already discussed in Sect.~\\ref{outer}) and let clumping start close to the wind base. Clumping parameters are given in Table~\\ref{Tab:mod}, model Obs1.  As illustrated in Fig.~\\ref{Fig:cmp_obs}, the synthetic line profiles using $\\dot{M}=2.3$, as inferred from $\\rm H_{\\alpha}$, are now at the observed levels. Because of our insufficient treatment of line overlap, we gave higher weight to the $\\lambda$1118 component when performing the fitting, but the profile-strength ratio between the blue and red component was nevertheless reasonably well reproduced (see also discussion in Sect.~\\ref{dens_par}).  However, though the fit appears quite good, we did not aim for a perfect one, and must remember the deficits of our modeling technique. For example, while the early onset of clumping definitely improved the fit (using our default value, there was a dip close to line center) and might be considered as additional evidence that clumping starts close to the wind base, the same effect could in principle be produced by non-LTE effects close to the photosphere or by varying the underlying $\\beta$ velocity law. Such effects will be thoroughly investigated in a follow-up paper, which will also include a comparison to observations from many more objects.  Clearly, a consistent modeling of resonance lines (at least of intermediate strengths) requires the consideration of a much larger parameter set than if modeling via the standard diagnostics assuming optically thin clumping, and a reasonable fit to a single observed line complex can be obtained using a variety of different parameter combinations.  The analysis of PV lines as done here can therefore, at present, only be considered as a consistency check for mass-loss rates derived from other, independent diagnostics, and not as a tool for directly estimating mass-loss rates. Additional insight might be gained by exploiting more resonance doublets, due to the different reactions of profile strengths and shapes on $\\kappa_0$. The different slopes of the equivalent width as a function of $\\kappa_0$ in smooth and clumped models, especially at intermediate line strengths (Sect.~\\ref{dens_par}), may turn out to be decisive. However, because of, e.g., the additional impact from the ICM density, also this diagnostics requires additional information from saturated lines. Taken together, only a consistent analysis using different diagnostics and wavelength bands, and embedded in a suitable non-LTE environment, will (hopefully) provide a unique view.  \\label{cmp_obs} \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics[angle=90]{fig12.ps}} \\caption{Observed FUSE spectra of the PV doublet $\\lambda\\lambda$1118-1128 for the O6 supergiant $\\lambda$ Cep \\citep{Fullerton06}. The synthetic spectra are calculated for two 1D models assuming optically thin clumping (see Sect.~\\ref{cmp_obs}) and for one 2D stochastic model with parameters as in Table~\\ref{Tab:mod}, model Obs1. The models have mass-loss rates $\\dot{M}\\,[\\rm M_{\\odot}\\,yr^{-1}]$ as given in the figure. The zero point frequency is shifted to the line center of the $\\lambda$1118 component, and the two arrows at the bottom of the figure indicate in which region the two components overlap.} \\label{Fig:cmp_obs} \\end{figure}"
    },
    "0911/0911.1014_arXiv.txt": {
        "abstract": "We present our optical observations of {\\em Swift} GRB 070518 afterglow obtained at the 0.8-m Tsinghua University-National Astronomical Observatory of China telescope (TNT) at Xinglong Observatory. Our follow-up observations were performed from 512 sec after the burst trigger. With the upper limit of redshift$\\sim$0.7, GRB 070518 is found to be an optically dim burst. The spectra indices $\\beta_{ox}$ of optical to X-ray are slightly larger than 0.5, which implies the burst might be a dark burst. The extinction $A_{V}$ of the host galaxy is 3.2 mag inferred from the X-ray hydrogen column density with Galactic extinction law, and 0.3 mag with SMC extinction law. Also, it is similar to three other low-redshift optically dim bursts,which belong to XRR or XRF, and mid-term duration($T_{90}<10$, except for GRB 070419A, $T_{90}$=116s). Moreover, its $R$ band afterglow flux is well fitted by a single power-law with an index of 0.87. The optical afterglow and the X-ray afterglow in the normal segment might have the same mechanism, as they are consistent with the prediction of the classical external shock model. Besides, GRB 070518 agrees with Amati relation under reasonable assumptions. The Ghirlanda relation is also tested with the burst. ",
        "introduction": "Since the first optical counterpart of gamma-ray burst (GRB) was identified on 1997 February 28,  more and more GRBs have been optically detected, especially after {\\em Swift} was launched successfully, which allows early follow-up observations with ground-based optical telescopes. However, about 40 per cent of {\\em Swift} GRBs still failed to be identified with optical or infrared counterparts (Burrows et al. 2008). For the nature of the dark burst, several models have been proposed: extinction from dusty opaque star-forming regions ( e.g. Groot et al. 1998; Reichart \\& Price 2002); high redshift, which diminishes the GRB optical luminosity (Lamb \\& Reichart 2000), as the ly$\\alpha$ forest is shifted into the observational energy range; intrinsically dim bursts (Fynbo et al. 2001; Rol et al. 2005 ); the observational effect, As a result of the late response to the bursts. Jakobsson et al. (2004) have proposed an operational definition of dark burst (i.e. optical-to-X-ray spectral index $\\beta_{\\mathrm{ox}}<0.5$), according to the fireball model.   However, GRBs with an optical afterglow have been studied widely and in depth(e.g. GRB 060218, GRB 080319B). This has resulted in an increasing number of GRBs with known redshift, which has allowed statistical studies to be carried out. The optical luminosity distribution of GRBs varies widely up to several magnitudes (Kann et al. 2006; 2008). After being transformed into a common redshift (e.g., $z=1$), the optical luminosity shows a bi-model phenomenon (Nardini et al. 2006; Liang et al. 2006; Kann et al. 2006; 2008): optically luminous and optically dim. The bi-model distribution is also found to exist in other energy bands (Gendre et al. 2008a,b).  Liang \\& Zhang (2006) found that the apparent bimodality cannot be interpreted as a manifestation of the extinction effect, and they suggested that there might be two types of progenitors or two types of explosion mechanisms in operation. Considering the dark burst is usually a type of GRBs with a faint optical afterglow, optically dim bursts might be related to dark bursts (Nardini et al. 2006). However, the number of well-studied optically dim bursts is very limited. Thus, increasing the number of known optically dim bursts may help us to understand the origin of the bi-model distribution, the nature of optically dim bursts and dark bursts.  One such interesting case is GRB 070518. It was detected by Gamma-ray Burst Alert Telescope (BAT) onboard {\\em Swift} at 14:26:21 (UT) on 2007 May 18. The X-ray telescope (XRT) and ultraviolent/optical telescope (UVOT) on board {\\em Swift} observed the counterpart beginning at 70 and 100 sec after the trigger, respectively. T$_{90}$ (15-350 KeV) is 5.5$\\pm$0.2 s ,T$_{50}$ is 2.9s (Sakamoto et al. 2008a).  The values of duration mean the burst belongs to the short tail of the long group (Kouveliotou et al. 1993). Thus, it might be a classical long-duration GRB or an intermediate-long GRB (Horvath et al. 2008). The ratio of fluence between the 25-50 KeV and 50-100 KeV band (Sakamoto et al. 2008a) is about 1.06, which makes GRB 070518 belong to an X-ray rich burst (XRR), according to the criterion \\footnote{S(25-50KeV)/S(50-100KeV)$=<$0.72 C-GRB; \\\\ 0.72$<$S(25-50KeV)/S(50-100KeV)$=<$1.31 XRR;\\\\ S(25-50KeV)/S(50-100KeV)$>$1.32 XRF.} given by Sakamoto et al. (2008b). The first given magnitude and coordinates of the counterpart are about 18 mag in $white$ band reported by UVOT and RA(J2000) = 16h 56m 47.7s, Dec(J2000) = 55d 17m 42.3s (radius, 90 pen cent containment), respectively (Guidorzi et al. 2007). The afterglow was also observed by other ground-based telescopes later. Besides, the redshift of GRB 070518 was reported to be lower than 0.7 (Cucchiara et al. 2007) based on the photometry of the {\\em Swift} UVOT. No other report about redshift of GRB 070518 has been presented in the literature up to now. In our analysis, the redshift of 0.7 are used to investigate the properties of GRB 070518.  In this paper, we report on optical photometric follow-up observations of GRB 070518 at Xinglong Observatory of National Astronomical Observatories, Chinese Academy of Sciences. Afterglow observations and results are described in $\\S$2. An anaysis and discussion are presented in $\\S$3, and a summary and conclusion are given in $\\S$4.  In our calculations, a flat universe is assumed with matter density $\\Omega_M = 0.3$, cosmological constant $\\Omega_\\Lambda=0.7$, and Hubble constant $H_0=70$ km s$^{-1}$ Mpc$^{-1}$. The formula $f\\propto$$t^{-\\alpha}\\nu^{-\\beta}$ is used for the afterglow decay analysis. ",
        "conclusions": "We have presented optical afterglow observations of GRB 070518 from the TNT and an analysis of the optical and X-ray light curves. The optical afterglow shows a constant power law decay $\\sim$0.87, while the X-ray shows a \"steep-shallow-normal\" decay behavior. Based on the upper limit of redshift of 0.7, the optical luminosity at 12 hours in the rest frame was estimated to be $2\\times10^{28}$ erg s$^{-1}$, which means that the burst is an optically dim burst. The conclusion is also supported by comparing with three other optically dim bursts: GRB 050416A, GRB 060512 and GRB 070419A. The spectral indices of $\\beta_{\\mathrm{ox}}$ at 1000 s, 12 hours,10$^{5}s$ are calculated. All $\\beta_{\\mathrm{ox}}$ were slightly larger than 0.5, indicating that GRB 070518 might be a dark burst or a gray burst, according to the definition of Jakobsson et al. (2004). The optical extinction in its host galaxy inferred from  X-ray hydrogen column density is 3.2 with Galactic extinction law, and 0.3 with SMC extinction law.  Afterglow model has been applied to the light curves. Energy injection is found to occurr for the X-ray afterglow before the start of the normal decay segment. The behavior of the normal segment of the X-ray afterglow and the optical afterglow is consistent with the prediction of the classical external shock model (Urata et al. 2007), indicating that they come from the same origin. As an XRR burst, GRB 070518 agrees with Amati relation. Moreover, it fills the gape of high-luminosity and low-luminosity bursts in the Amati relation, similar to GRB 050416A (Sakamoto et al. 2005). It is also found that GRB 070518 agrees with the Ghalanda relation within reasonable assumptions.  Moreover, a comparison with three optically dim bursts (GRB 050416A, GRB 060512 and GRB 070419A and GRB 070518) reveals that several similar properties were shared: burst duration T$_{90}$$<$10s (expect for GRB 070419A, $T_{90}=116s$), soft spectrum at BAT band $\\Gamma$$>2$, low redshift $z<$1, etc."
    },
    "0911/0911.5401_arXiv.txt": {
        "abstract": "A general equation of state is used to model unified dark matter and dark energy (dark fluid), and it has been proved that this model is equivalent to a single fluid with time-dependent bulk viscosity. In this paper, we investigate scalar perturbation of this viscosity dark fluid model. For particular parameter selection, we find that perturbation quantity can be obtained exactly in the future universe. We numerically solve the perturbation evolution equations, and compare the results with those of $\\Lambda$CDM model. Gravitational potential and the density perturbation of the model studied here have the similar behavior with the standard model, though there exists significant value differences in the late universe. ",
        "introduction": "Astrophysics and cosmology observations in recent years delineate the cosmological picture on its constituents more and more accurately, i.e, the precision cosmology era comes. Except for the long standing puzzling dark matter component, an unknown cosmic matter-energy constituent refereed to the so called dark energy may also exist that accelerates our universe expansion now, which contradicts with our traditional attitude on the behavior of conventional matter but it is eventually confirmed by recent observations like SNe Ia \\cite{sn1} \\cite{sn2} and CMB observations \\cite{CMB}. The cosmological dark sector, often divided as the mysterious DM and DE sectors respectively, takes around $95\\%$ of total energy budget of our universe. The concord $\\Lambda$CDM model could be consistent with most of global astrophysics observational results. But the introduction of the cosmological constant simultaneously results in the yet to answer problems related directly to how to understand the fundamental physics theory, like the fine-tuning and the coincidence problems respectively. At the same time, ``most'' of course does not equal to ``all'', some astrophysics problems still need be clarified and solved in the framework of the $\\Lambda$CDM model \\cite{six}, such as the the core singularities of the cold matter halo profiles. With the aim to understand the cosmic acceleration or dark energy phenomena, many theoretical models have been proposed, like the scalar field models and the modified Einstein gravity models \\cite{r1} \\cite{r2} \\cite{r3}.  Due to the limited scope of our experimental and observational tools, we have not yet been able to understand  the ``dark'' nature and to detect the origin of DM and DE. Purely gravitational probes can not provide enough information to differentiate these two kinds of mysterious constitutions either. Therefore, from the phenomenological and practical point of view, a single (unified DM with DE) fluid description may be more plausible at least in the cosmic evolution description, which utilizes a single equation of state to model the dark matter and dark energy contributions together \\cite{df1} \\cite{df2} \\cite{df3} \\cite{df4} \\cite{df5} \\cite{Liddle}. Generally, such models has a non-constant equation of state (EoS), which reflects both its dynamical and thermodynamics characters. The density dependent equation of state is widely investigated, such as the famous Chaplygin gas and generalized Chaplygin gas models, which assume an EoS form like the $p=-A/\\rho^{\\alpha}$ where the $\\alpha$ is a model parameter. Another practical method to modify the EoS is by the introduction of cosmic viscosity media contribution which replaces the simplest perfect fluid EoS on a more physical and realistic basis. In the homogeneous and isotropic Friedmann-Robertson-Walker frame, only a bulk viscosity term which behaves as an additive pressure contribution can mimic both the two dark components and their coupling effects by playing a main role to influence the cosmic evolution. Different forms of the viscosity coefficients have been proposed like the density $\\rho$ dependent \\cite{d} or the redshift $z$ dependent \\cite{md}.  In this letter, we will proceed our effort on the investigation of the viscosity dark fluid model to study the scalar perturbation of this model, for perturbation analysis can provide us a powerful tool to differentiate and constrain cosmology models finely as the calculation of perturbation quantities links the theoretical models with more plentiful and precise observations, like the cosmic microwave background (CMB) and large scale structure observation. There have been some researches on the perturbation evolution of the viscosity models \\cite{FGR} \\cite{lb} \\cite{z}, by which we know that after corresponding model parameters chosen properly, the Chaplygin gas formulation can be viewed as a special case of the density dependent viscosity model. In the non-perturbative (zero order) level, the Chaplygin gas model can be exactly solved and fit the observational data well. But it has been found that in the perturbation level, there exist some unacceptable behaviors, like the blow up of density perturbation evolution and other peculiar behaviors \\cite{z}  \\cite{ps}. One motivation to build other kinds of the viscosity models is to overcome these difficulties the Chaplygin gas models possess. Here, we will consider a time-dependent viscosity coefficient model, which is equivalent to the introduction of a general Equation of state(EoS) \\cite{RM}. The general EoS is \\begin{equation}\\label{1} p=(\\gamma-1)\\rho+p_{0}+w_{H}H+w_{H2}H^{2}+w_{dH}\\dot H. \\end{equation} In the background, this model can fit the current astrophysics observational datasets consistently. We derive its perturbation equations that govern the evolution of gravitational potential and density perturbation below. We numerically solve the perturbation equation, and compare it with that of conventional $\\Lambda$CDM model and the Chaplygin gas model finding that the dark fluid model behaves well in different scales. Though there exists some value difference between the $\\Lambda$CDM model and the dark fluid model in the late time evolution, their gravitational potential and density contrast shape and evolution behavior are similar by plotting respectively.  This paper is organized as follows: In Sec. \\textbf{II}, we summarize the calculations of scalar perturbation, and give the general evolution equation of the gravitational potential. In Sec. \\textbf{III}, we briefly review the background evolution of the dark fluid model. In Sec. \\textbf{IV}, we discuss the perturbation evolution of the dark fluid model. In Sec. \\textbf{V}, we numerically solve the perturbation equation and compare it with other models. Finally, we present the conclusions in the last section. ",
        "conclusions": "In this paper, we investigate extensively the dark fluid model proposed in \\cite{RM}, equivalently this model can be viewed as a single fluid with time-dependent bulk viscosity. Scale factor and density evolution can be exactly solved in this model. In the background, the dark fluid model can fit the supernova data acceptable. Our main task in this paper is to analysis the behavior of this model in the perturbation level. We derive equations govern the perturbation quantities. For the condition that $T_{1}$ is smaller than $0$, the universe will enter de Sitter phase in the $t\\rightarrow\\infty$ future. We solve exactly the gravitational equation in this condition and obtain the solution for both long and short wave case. Generally, the perturbation evolution equations are solved numerically. When compare the results with those of $\\Lambda$CDM model, we find that \\begin{itemize} \\item In the early time and the large scale, both the gravitational potentials of two models behave as a constant. \\item Though the gravitational potentials of two models have similar behavior and shape, as can be seen form FIG. ~3, there exists about $5\\%-10\\%$ significant value difference in the late time. \\end{itemize}  Perturbation analysis also provides constraint on model parameter. For different selection of parameter $\\tilde\\gamma$, both $\\tilde\\gamma<0$ and $\\tilde\\gamma>0$ can give consistent prediction curve of distance modulus, but numerical results indicate in $\\tilde\\gamma<0$ case, the gravitational potential deviates from $\\Lambda$CDM significantly from the early time. This result strongly constrain the selection region of $\\tilde\\gamma$. We suggest $\\tilde\\gamma>0$, which can also produce positive sound speed naturally."
    },
    "0911/0911.2367_arXiv.txt": {
        "abstract": "The properties of galactic cosmic rays are investigated with the KASCADE-Grande experiment in the energy range between $10^{14}$ and $10^{18}$~eV. Recent results are discussed. They concern mainly the all-particle energy spectrum and the elemental composition of cosmic rays. ",
        "introduction": "To reveal the origin of galactic cosmic rays is the main objective of the KASCADE and KASCADE-Grande experiments.  The results of KASCADE contributed significantly to the understanding of the origin of the knee in the all-particle energy spectrum of cosmic rays at energies around $4\\cdot10^{15}$~eV.  It could be shown that the knee is caused by a fall-off in the flux of light nuclei \\cite{ulrichapp,Apel:2008cd}.  Energy spectra for five dominant elemental groups could be reconstructed (p, He, CNO, Si, and Fe), they exhibit a fall-off behavior approximately proportional to the nuclear charge of the elemental groups.  In the energy region between $10^{17}$ and $10^{18}$~eV a transition from a galactic to an extra-galactic origin of cosmic rays is expected \\cite{behreview}.  Main focus of KASCADE-Grande is to understand the end of the galactic cosmic-ray component through detailed investigations of the energy spectrum and the mass composition in this energy region.   KASCADE-Grande comprises 37 detector stations equipped with plastic scintillators, covering an area of 0.5~km$^2$ to measure the electromagnetic shower component \\cite{grande}. It also includes the original KASCADE experiment, consisting of several detector systems \\cite{kascadenim}. A $200 \\times 200$~m$^2$ array of 252 detector stations, equipped with scintillation counters, measures the electromagnetic and, below a lead/iron shielding, the muonic parts of air showers. A 130~m$^2$ streamer tube detector, shielded by a soil-iron absorber serves to reconstruct the tracks of high-energy ($E_\\mu>0.8$~GeV) muons \\cite{mtdnim}. An iron sampling calorimeter of $16 \\times 20$~m$^2$ area detects hadronic particles \\cite{kalonim}.  It has been calibrated with a test beam at the SPS at CERN up to 350~GeV particle energy \\cite{kalocern}.  \\begin{figure}\\centering \\epsfig{file=dipierro-lat.eps, width=\\columnwidth} \\caption{Measured lateral distribution of a single event for charged particles and muons \\cite{icrc09-dipierro}.} \\label{lat} \\end{figure}  As an example for a measured air shower, the lateral distributions for charged particles and for muons are shown in \\fref{lat} \\cite{icrc09-dipierro}. The charged-particle density is measured with the Grande detectors and the muonic component by the shielded array detectors of KASCADE. ",
        "conclusions": "The recent results from the KASCADE-Grande experiment indicate that we make substantial progress in measuring the energy spectrum and the elemental composition of cosmic rays in the energy region between $10^{16}$ and $10^{18}$~eV. The new data will significantly contribute to the understanding of the end of the galactic component and the transition to an extra-galactic component. The data will be crucial to distinguish between different astrophysical scenarios."
    },
    "0911/0911.3588_arXiv.txt": {
        "abstract": "We present a photometric study of $I$-band variability in the young association Cepheus OB3b. The study is sensitive to periodic variability on timescales of less than a day, to more than 20 days. After rejection of contaminating objects using $V$, $I$, $R$ and narrowband H$\\alpha$ photometry, we find 475 objects with measured rotation periods, which are very likely pre-main-sequence members of the Cep OB3b star forming region.  We revise the distance and age to Cep OB3b, putting it on the self-consistent age and distance ladder of \\cite{2008MNRAS.386..261M}. This yields a distance modulus of 8.8$\\pm$0.2 mags, corresponding to a distance of 580$\\pm$60 pc, and an age of 4-5Myrs.  The rotation period distribution confirms the general picture of rotational evolution in young stars, exhibiting both the correlation between accretion (determined in this case through narrowband H$\\alpha$ photometry) and rotation expected from disc locking, and the dependence of rotation upon mass that is seen in other star forming regions. However, this mass dependence is much weaker in our data than found in other studies. Comparison to the similarly aged NGC 2362 shows that the low-mass stars in Cep OB3b are rotating much more slowly. This points to a possible link between star forming environment and rotation properties. Such a link would call into question models of stellar angular momentum evolution, which assume that the rotational period distributions of young clusters and associations can be assembled into an evolutionary sequence, thus ignoring environmental effects. ",
        "introduction": "\\label{sec:introduction}  There are sound theoretical reasons to expect that accretion processes help determine the angular momentum of T-Tauri stars. Historically, the slow rotation rates of T Tauri stars (relative to their break-up velocity) has been explained by the disc-locking theory \\citep{konigl91,shu94}. In this theory, magnetic field lines connect the star to the disc, enforcing synchronous rotation between the star and the material in the disc at some radius, near where the magnetic field disrupts the disc. The simplistic theory has been expanded in a variety of models where angular momentum is removed from the star by a combination of the disc and an accretion-driven wind \\citep[e.g.][]{fendt07,romanova07,matt08}. Whether the star is spun up or down by the star-disc interaction depends upon the balance between accretion spinning up the star and magnetic (or wind) torques slowing it down. Theoretically, this issue is unresolved; some studies find the star spins up \\citep{bessolaz08}, some find it spins down \\citep[e.g.][]{long05}.  Observationally, the evidence for the influence of accretion disks on rotation is much stronger than it was a few years ago \\citep{herbst07}. Previously, conflicting results had arisen \\citep[e.g.][] {herbst02,stassun99,littlefair05}. These were most likely due to a combination of small sample sizes, and ambiguous diagnostics of the presence of accretion disks. \\cite{rebull06} resolved these issues with a large sample of rotation periods in the ONC, and accretion disk status defined from Spitzer IRAC data. \\cite{rebull06} found a clear correlation between mid-IR excess and rotation, in the sense that stars with mid-IR excess were much more likely to be slow-rotators. An analysis of the slightly older NGC 2264 found the same result \\citep{cieza07}, and confirmed the \\cite{rebull06} result, via a refined analysis of the same data. There is thus strong observational evidence that the star-disc interaction is responsible for extracting angular momentum from young stars. Interestingly, a small population of rapidly rotating stars with mid-IR excesses exists; it is possible that these stars are being spun-up by accretion, hinting that the braking process of young stars is intermittent. Also, in both the ONC and NGC 2264, there exists a significant population of slow rotators with {\\em no} mid-IR excess. These are often interpreted as being recently released from disc-locking, but there are problems with this interpretation \\citep[see][for example]{bouvier07}.  The firm link established between accretion and rotation represents significant progress, but there are still open questions regarding the rotation of young stars. For example, the low-mass stars are rotating more rapidly than the high-mass stars, and appear to spin-up more rapidly as they contract towards the main sequence \\citep{herbst02,irwin07a}. The reason for this is still not known. Also, the angular momentum evolution of young stars is determined by assembling different clusters into an evolutionary sequence, and assuming the period distribution of the older clusters can be modelled using the young clusters as a starting point. In doing so, the possibility of an environmental effect on rotation is ignored. Such an environmental effect may be indicated by the data; the young stars in IC348 rotate much slower than those in the similarly aged NGC 2264 \\citep{littlefair05}.  Here we present a photometric study of the young association Cepheus OB3b (hereafter Cep OB3b). Cep OB3 is a young association covering a region of the sky from approximately 22$^h$46$^m$ to 23$^h$10$^m$ in right ascension and +61$^{\\circ}$ to +64$^{\\circ}$ in declination. The subgroup Cep OB3b lies closest to the molecular cloud, and has a rich pre-main-sequence (PMS) population, confirmed by both spectroscopy \\citep{pozzo03} and X- ray data \\citep{getman06}. In section~\\ref{sec:obs} we present the observations and data reduction techniques applied. Section~\\ref{sec:detect} describes the techniques used to identify periodic variables. Section~\\ref{sec:distance} provides a revision of the age of, and distance to, Cep OB3b. In Sections~\\ref{sec:results} and \\ref{sec:discussion} we present our results and analysis of the data, whilst in section~\\ref{sec:concl} we draw our conclusions. ",
        "conclusions": "\\label{sec:concl}  We present a photometric study of $I$ band variability towards the young association Cepheus OB3b. The study is sensitive to periodic variability on timescales of less than a day, to more than 20 days. The result is a database of 704 periodic variables in the field of Cep OB3b. A random inspection of 200 of these objects suggest that around 97 per cent of these periods are genuine. Colour cuts using $V$, $I$, $R$ and narrowband H$\\alpha$ photometry are used to reject contaminating objects, leaving 475 objects with measured rotation periods, which are very likely pre-main-sequence members of the Cep OB3b star forming region.  We revise the distance and age to Cep OB3b, putting it on the consistent age and distance ladder of \\cite{2008MNRAS.386..261M}. This yields a distance modulus of 8.8$\\pm$0.2 mags, corresponding to a distance of 580$\\pm$60 pc, and an age of 4-5Myrs. For the purposes of this paper we therefore adopt an age of 4.5Myr.  The rotation period distribution confirms the general picture of rotational evolution in young stars, exhibiting both the correlation between accretion and rotation expected from disc locking, and the dependence of rotation upon mass that is seen in other star forming regions. However, this mass dependence is much weaker than seen in other regions. Comparison to the similarly aged NGC 2362 shows that the low-mass stars in Cep OB3b are rotating much more slowly. This points to a possible link between star forming environment and rotation properties. Such a link would call into question models of stellar angular momentum evolution, which assume that associations can be assembled into an evolutionary sequence, thus ignoring environmental effects."
    },
    "0911/0911.1258_arXiv.txt": {
        "abstract": "We present spectra of high-redshift supernovae (SNe) that were taken with the Subaru low resolution optical spectrograph, FOCAS.  These SNe were found in SN surveys with Suprime-Cam on Subaru, the CFH12k camera on the Canada-France-Hawaii Telescope (CFHT), and the Advanced Camera for Surveys (ACS) on the Hubble Space Telescope (HST). These SN surveys specifically targeted $z>1$ Type~Ia supernovae (SNe~Ia). From the spectra of 39 candidates, we obtain redshifts for 32 candidates and spectroscopically identify 7 active candidates as probable SNe~Ia, including one at $z=1.35$, which is the most distant SN~Ia to be spectroscopically confirmed with a ground-based telescope. An additional 4 candidates are identified as likely SNe~Ia from the spectrophotometric properties of their host galaxies. Seven candidates are not SNe~Ia, either being SNe of another type or active galactic nuclei. When SNe~Ia are observed within a week of maximum light, we find that we can spectroscopically identify most of them up to $z=1.1$. Beyond this redshift, very few candidates were spectroscopically identified as SNe~Ia. The current generation of super red-sensitive, fringe-free CCDs will push this redshift limit higher. ",
        "introduction": "\\label{sec:introduction}  Type~Ia supernovae (SNe~Ia) have proven to be very good standard candles for cosmological studies.  They are bright enough to detect at cosmological distances, $z\\sim1.5$, if one uses 8--10~m class ground-based optical telescopes or the Hubble Space Telescope (HST), and their luminosities can be standardised using empirical relations between luminosity and light curve shape (\\cite{phillips1993}) and colour \\citep{tripp1998}.  SNe~Ia have played a leading role in measuring the expansion history of the universe since two independent teams, the Supernova Cosmology Project (SCP) and the High-$z$ Team, discovered the accelerating expansion of the universe (\\cite{perlmutter1999}; \\cite{riess1998}). Since then, many projects to discover and identify SNe~Ia have been organized. For example, the Carnegie Supernova Project (CSP; \\cite{hamuy2006}), the Nearby Supernova Factory (SNfactory; \\cite{aldering2002}), the Harvard-Smithsonian Center for Astrophysics (CfA) supernova survey (\\cite{jha2006}; \\cite{hicken2009}), the Sloan Digital Sky Survey-II Supernova Survey (SDSS-II SN Survey; \\cite{sako2008}; \\cite{frieman2008}), the Supernova Legacy Survey (SNLS; \\cite{astier2006}), the Equation of State: SupErNovae trace Cosmic Expansion (ESSENCE) survey (\\cite{miknaitis2007,woodvasey2007}), and Higher-Z team (\\cite{riess2007}) have detected around 1000 SNe up to redshift $z\\sim 1.5$.  The combination of SN~Ia data with measurements of cosmic microwave background (CMB) fluctuations (\\cite{spergel2003}; \\cite{spergel2007}; \\cite{komatsu2009}), baryon acoustic oscillations (BAO; \\cite{eisenstein2005}) and galaxy cluster number counts (\\cite{vikhlinin2009}) have constrained the dark energy equation of state parameter.  However, the limits are consistent with very different dark energy models, so the fundamental nature of dark energy remains unclear. It is currently one of the biggest mysteries in physics, and combined astronomical observations (SNe~Ia, CMB, BAO and weak lensing) seem to be the only way to constrain its properties.  Since the discovery of the accelerating expansion of the universe, the SCP has been carrying out imaging surveys for SNe~Ia at $z\\gtrsim 1$, an epoch during which the expansion of the universe is expected to be decelerating.  Spectroscopic follow-up observations are an essential part of these surveys, providing spectroscopically determined redshifts and, when necessary and possible, direct confirmation of the SN type. In this paper, we present spectra of SNe and their host galaxies taken with the Faint Object Camera And Spectrograph (FOCAS; \\cite{kashikawa2002}) on the Subaru 8.2-m telescope. Twelve SN candidates shown here were found in ground-based observations targeting blank fields and nearby galaxy clusters with Suprime-Cam (\\cite{miyazaki2002}) on Subaru and the CFH12k camera (\\cite{cuillandre2000}) on the Canada-France-Hawaii telescope (CFHT). The remaining 27 SN candidates were discovered using the Advanced Camera for Surveys (ACS; \\cite{benitez2003}) on HST\\footnote{Based on observations made with the NASA/ESA Hubble Space Telescope and obtained from the data archive at the Space Telescope Institute. STScI is operated by the Association of Universities for Research in Astronomy, Inc. under the NASA contract NAS 5-26555. The observations are associated with program 10496.} targeting high redshift galaxy clusters in a program called the HST Cluster SN Survey (Program number 10496, PI: Perlmutter). These SN searches specifically targeted $z>1$ SNe~Ia, for which there have been relatively few spectroscopically confirmed SNe~Ia (\\cite{aldering1998}; \\cite{coil2000}; \\cite{tonry2003}; \\cite{barris2004}; \\cite{riess2004}; \\cite{lidman2005}; \\cite{riess2007}). In \\S\\ref{sec:observationanddata}, we summarize the SN searches and present data for both imaging and spectroscopy.  Spectroscopic data reductions are described in \\S\\ref{sec:observationanddata}.  SN and host galaxy classifications are shown in \\S\\ref{sec:sntyping} and \\S\\ref{sec:focashostspectraforhstclustersn}, respectively.  In \\S\\ref{sec:typeideff}, we describe factors that influence the classification of SNe.  \\S\\ref{sec:summary} is a summary of the paper. We use the standard $\\Lambda$CDM cosmological parameters of $(H_0, \\Omega_M, \\Omega_\\Lambda)=(70, 0.27, 0.73)$ for calculating the age of the universe at a certain redshift. All magnitudes are measured in the AB system. ",
        "conclusions": "\\label{sec:summary}  We presented spectra of SN candidates obtained with FOCAS on the Subaru 8.2-m telescope. Seven active candidates were identified as SNe Ia, including SCP06G4 at $z_{\\rm{SN}}=1.35$, the most distant SN~Ia to be spectroscopically identified with a ground-based telescope. Redshifts were obtained for all but 7 of the remaining 32 candidates based on the host galaxy spectra. An additional 4 candidates are identified as likely SNe~Ia from the spectrophotometric properties of their hosts.  The spectral properties of these hosts found in the HST Cluster Supernova Survey were examined by comparing their spectra with BC03 model spectra. All of the host galaxy spectra indicate that they are passively evolving galaxies that have quenched their star forming activities.  We also investigated the factors affecting the classification of SNe and found that it is critical to take spectra within a week of maximum light. This requires secure early detection, well-sampled light curves and prompt spectroscopic follow-up.  \\bigskip  We thank the anonymous referee for providing helpful comments and suggestions. TM and YI are financially supported by the Japan Society for the Promotion of Science (JSPS) through the JSPS Research Fellowship. CL acknowledges the financial support from the Oskar Klein Centre at the University of Stockholm. This work was supported in part with scientific research grants (15204012 and 17104002) from the Ministry of Education, Science, Culture, and Sports of Japan, and a JSPS core-to-core program ``International Research Network for Dark Energy''. Financial support for this work was provided in part by NASA through program GO-10496 from the Space Telescope Science Institute, which is operated by AURA, Inc., under NASA contract NAS 5-26555.  This work was also supported in part by the Director, Office of Science, Office of High Energy and Nuclear Physics, of the U.S. Department of Energy under Contract No. AC02-05CH11231.  Part of the Suprime-Cam observations were done during the guaranteed time observation of Suprime-Cam and we thank for the Suprime-Cam instrument team.  We also appreciate much help by the SDF and SXDS project team members.  We also thank for Youichi Ohyama, who helped our observations as a support scientist of FOCAS.  Data analysis were in part carried out on common use data analysis computer system at the Astronomy Data Center, ADC, of the National Astronomical Observatory of Japan."
    },
    "0911/0911.4820_arXiv.txt": {
        "abstract": "The coupled dark energy models, in which the quintessence scalar field nontrivially couples to the cold dark matter, have been proposed to explain the coincidence problem. In this paper we study the perturbations of coupled dark energy models and the effects of this interaction on the current observations. Here, we pay particular attention to its imprint on the late-time Integrated Sachs-Wolfe (ISW) effect. We perform a global analysis of the constraints on this interaction from the current observational data. Considering the typical exponential form as the interaction form, we obtain that the strength of interaction between dark sectors is constrained as $\\beta<0.085$ at $95\\%$ confidence level. Furthermore, we find that future measurements with smaller error bars could improve the constraint on the strength of coupling by a factor two, when compared to the present constraints. ",
        "introduction": "Current cosmological observations, such as the cosmic microwave background (CMB) measurements of temperature anisotropies and polarization \\cite{Komatsu:2008hk} and the redshift-distance measurements of Type Ia Supernovae (SNIa) at $z<2$ \\cite{Kowalski:2008ez}, have demonstrated that the Universe is now undergoing an accelerated phase of expansion and that its total energy budget is dominated by the dark energy component. The nature of dark energy is one of the biggest unsolved problems in modern physics and has been extensively investigated in recent years, both under the theoretical and the observational point of view.  The simplest candidate of dark energy is the cosmological constant, where the equation of state $w$ always remains $-1$. In the standard $\\Lambda$-Cold Dark Matter ($\\Lambda$CDM) cosmology, cold dark matter only interacts with other components by gravity, while dark energy is simply a cosmological constant without any evolution and perturbation. Although this concordance model can fit the current observational data very well \\cite{Komatsu:2008hk,Kowalski:2008ez}, the possibility of dynamical dark energy models cannot be ruled out yet. Furthermore, this model suffers of the fine-tuning and coincidence problems, i.e. why the Universe is dominated by dark energy in late times \\cite{CCproblem1,CCproblem2}.  In order to lift these severe tensions, many alternative dynamical dark energy models, such as quintessence \\cite{Quint1,Quint2,Quint3,Quint4}, phantom \\cite{Phantom}, K-essence \\cite{kessence1,kessence2} and quintom \\cite{Feng:2004ad}, have been proposed recently (for a review see Ref.\\cite{copeland}). For example, the quintessence model, a scalar field $\\phi$ slowly rolling down a potential energy $V(\\phi)$, has the spatial fluctuations and its equation of state $w_{\\phi}$ can evolve with the cosmic time. More importantly, being a dynamical component, the quintessence is naturally expected to interact with the other components, such as the cold dark matter \\cite{Copeland98,Amendola00} or massive neutrinos \\cite{gu03,fardon04}, in the field theory framework. If these interactions really exist, it would open up the possibility of detecting the dark energy non-gravitationally.  In the coupled dark energy models, the quintessence scalar field could nontrivially couple to the cold dark matter component. The presence of the interaction clearly modifies the cosmological background evolutions. The evolution of cold dark matter energy density $\\rho^{}_{c}(a)$ will be different, when compared to that of the minimally coupled models, and dependent on the quintessence field $\\phi\\,$: \\begin{equation} \\rho^{}_{\\rm c}(a)=\\rho^{}_{\\rm c0}a^{-3+\\delta(\\phi)}~,\\label{cdm} \\end{equation} where $a$ is the scale factor, $\\rho^{}_{\\rm c0}$ is the present value of cold dark matter energy density, and $\\delta(\\phi)$ is the modification due to the interactions and could be non-zero during the evolution of Universe. In this case, at early times there will be more (less) cold dark matter energy density, when $\\delta(\\phi)<0$ ($>0$).  On the other hand, the interaction between dark sectors will also affect the evolution of cosmological perturbations, which has been widely investigated recently (see e.g. Refs.\\cite{hwangnoh02,amendola04,koivisto05,lee06, olivares06,brookfield08,bean08,mainini08,valiviita08,he09a,chong09}). Due to the different evolution of cold dark matter energy density, the interaction between dark sectors could shift the matter-radiation equality scale factor $a_{\\rm eq}$, and affect the locations and amplitudes of acoustic peaks of CMB temperature anisotropies and the turnover scales of large scale structure (LSS) matter power spectrum consequently. In addition, the interaction will also affect the late ISW effect \\cite{isw} at large scales which is produced by the CMB photons passing through the time-evolving gravitational potential well, when dark energy or curvature becomes important at later times \\cite{brookfield08,he09b}.  Therefore, with the accumulation of observational data and the improvements of the data quality, it is of great interest to investigate the non-minimally coupled dark energy models from the current observational data. In this paper we study the perturbations of coupled dark energy models and present the constraints on the interactions from the observational data in detail. The structure of the paper is as follows: in Sec.II and Sec.III we show the basic equations of background evolution and linear perturbations of the coupled dark energy model, respectively. In Sec.IV we present the current and future observational datasets we used. Sec.V contains our main global fitting results from the current and future observations, while Sec.VI is dedicated to the summary. ",
        "conclusions": "The coupled dark energy models, in which the quintessence scalar field non-minimally couples to the cold dark matter, could affect the CMB temperature anisotropies and LSS matter power spectrum. In this paper we present constraints on this coupled dark energy model using the latest observations and future measurements. The WMAP5 data alone could only give a weak constraint on the strength of coupling $\\beta$. And then we exploit the capabilities of the late-time ISW effect in constraining the non-minimal coupling, using the cross-correlation signal between the quasar sample of SDSS DR6 \\cite{richards09} and the WMAP5 ILC map \\cite{hinshaw09}. We find that the current ISW data could slightly improve the constraint on $\\beta$. If we add the BAO and SNIa data into our calculations, the constraint on $\\beta$ will improves significantly, namely the $95\\%$ upper limit is $\\beta<0.085$. Finally we simulate the future measurements with smaller error bars and find that the future measurements could improve the constraint on the strength of coupling by a factor of two, when compared to the present constraints."
    },
    "0911/0911.2637_arXiv.txt": {
        "abstract": "{} {We study three Galactic \\HII\\ regions -- RCW~79, RCW~82 and RCW~120 -- where triggered star formation is taking place. Two stellar population are observed: the ionizing stars of each \\HII\\ region and young stellar objects on their borders. Our goal is to show that they represent two distinct populations, as expected from successive star forming events. } {We use near--infrared integral field spectroscopy obtained with SINFONI on the VLT to make a spectral classification. We derive the stellar and wind properties of the ionizing stars using atmosphere models computed with the code CMFGEN. The young stellar objects are classified according to their $K$--band spectra. In combination with published near and mid infrared photometry, we constrain their nature. Linemaps are constructed to study the geometry of their close environment. } {We identify the ionizing stars of each region. RCW~79 is dominated by a cluster of a dozen O stars, identified for the first time by our observations. RCW~82 and RCW~120 are ionized by two and one O star, respectively. All ionizing stars are early to late O stars, close to the main sequence. The cluster ionizing RCW~79 formed 2.3$\\pm$0.5 Myr ago. Similar ages are estimated, albeit with a larger uncertainty, for the ionizing stars of the other two regions. The total mass loss rate and ionizing flux is derived for each regions. In RCW~79, where the richest cluster of ionizing stars is found, the mechanical wind luminosity represents only 0.1\\% of the ionizing luminosity, questioning the influence of stellar winds on the dynamics of these three \\HII\\ regions. The young stellar objects show four main types of spectral features: H$_{2}$ emission, \\brg\\ emission, CO bandheads emission and CO bandheads absorption. These features are typical of young stellar objects surrounded by disks and/or envelopes, confirming that star formation is taking place on the borders of the three \\HII\\ regions. The radial velocities of most YSOs are consistent with that of the ionized gas, firmly establishing that their association with the \\HII\\ regions. Exceptions are found in RCW~120 where differences up to 50 \\kms\\ are observed. Outflows are detected in a few YSOs. All YSOs have moderate to strong near--IR excess. In the [24] versus $K-$[24] diagram, the majority of the sources dominated by H$_{2}$ emission lines stand out as redder and brighter than the rest of the YSOs. The quantitative analysis of their spectra indicates that for most of them the H$_{2}$ emission is essentially thermal and likely produced by shocks. We tentatively propose that they represent an earlier phase of evolution compared to sources dominated by \\brg\\ and CO bandheads. We suggest that they still possess a dense envelope in which jets or winds create shocks. The other YSOs have partly lost their envelopes and show signatures of accretion disks. Overall, the YSOs show distinct spectroscopic signatures compared to the ionizing sources, confirming the presence of two stellar populations. } {} ",
        "introduction": "\\label{s_intro}  Massive stars play a significant role in several fields of astrophysics. They produce the majority of heavy elements and spread them in the interstellar medium, taking an active part in the chemical evolution of galaxies. But they also end their life as supernovae and gamma--ray bursts. Through their strong winds and ionizing fluxes they power \\HII\\ regions and bubbles which are often used to trace metallicity gradients in galaxies. The energy they release in the interstellar medium is thought to trigger second--generation star formation events. Observations of young stellar objects (YSO) in molecular clouds surrounding (clusters of) massive stars lend support to this mechanism \\citep[e.g.][]{walborn02,hatano06}.  A particular case concerns star formation on the borders of \\HII\\ regions. According to the collect and collapse model \\citep{el77}, a dense shell of material is trapped between the shock and ionization front of an expanding \\HII\\ regions. When the amount of collected material is large enough, global shell fragmentation occurs and new stars are formed. The observation of molecular condensations on the borders of several \\HII\\ regions and the subsequent identification of YSOs within these clumps \\citep{de03,de05,za06,za07,de08a,de09,po09} confirms that this mechanism is at work at least in some \\HII\\ regions.  Other mechanisms leading to triggered star formation exist. Some work qualitatively as the collect and collapse model in the sense that the clumps are formed during the \\HII\\ region expansion. For instance, dynamical instabilities of the ionization front \\citep{vishniac83,garcia96} create molecular condensations separated by zones of lower densities. The newly formed clumps grow until they become Jeans unstable and collapse. Alternatively, second generation star formation can happen in pre-existing clumps. If the neutral gas in which the \\HII\\ region expands is not homogeneous, the outer layers of the molecular overdensities are ionized like the borders of a classical \\HII\\ region. A shock front precedes the ionization front inside these clumps, leading to their collapse \\citep{duvert90,ll94}.  A number of questions regarding triggered star formation remain unanswered. The properties of the observed YSOs are poorly known besides a crude classification in class I or class II objects by analogy with low--mass stars. YSOs usually display near--infrared spectra with CO, \\brg\\ and/or H$_{2}$ emission lines \\citep{bik06}. The relation, if any, between objects with different spectroscopic appearance is not clear. Besides, in the regions where the collect and collapse process is at work, the quantitative properties of the ionizing sources of the \\HII\\ regions are not known. In particular, the relative role of ionizing radiation and stellar winds on the dynamics of such regions is debated. The timescales under which material accumulates and fragments, and the properties of the resulting clumps depend on the strength of those two factors \\citep{whit94}. Hence, one might wonder whether the nature of the newly formed objects depends on the properties of the stars powering the \\HII\\ regions.   In the present study, we tackle these questions by investigating the properties of the ionizing stars and YSOs of three Galactic \\HII\\ regions: RCW~79, RCW~82 and RCW~120 \\citep{rod60}. Those regions are known to be the sites of triggered star formation \\citep{za06,za07,po09,kang09}. We have used SINFONI on the VLT to obtain near--infrared spectra of both the ionizing stars and a selection of YSOs in each region. Our main goals were:  \\begin{itemize}  \\item Identify the ionizing stars of each region, and derive their stellar properties using atmosphere models. In particular, we want to determine their ionizing fluxes and mass loss rates in order to better understand the dynamics of the \\HII\\ regions. Equally important is the determination of the age of those stars, since it can be related to the presence of YSOs to quantitatively confirm the existence of triggered star formation.  \\item Constrain the nature of the YSOs on the borders of the \\HII\\ regions. In combination with infrared photometry, spectroscopy can reveal the presence of disks or envelopes. The evolutionary status of those objects can thus be better understood. In particular, it can be clearly shown whether they are stars still in their formation process or objects already on the main sequence.  \\end{itemize}  The paper is organized as follows. In Sect.\\ \\ref{pres_hii} we present the three \\HII\\ regions targeted in this study. Sect.\\ \\ref{s_obs} describes our observations. The analysis of the ionizing stars is presented in Sect.\\ \\ref{s_ex}, while the YSOs are discussed in Sect.\\ \\ref{s_yso}. We discuss our results in Sect.\\ \\ref{s_disc} and summarize our conclusions in Sect.\\ \\ref{s_conc}. ",
        "conclusions": "\\label{s_conc}  We have obtained near--infrared integral field spectroscopy of three \\HII\\ regions (RCW~79, RCW~82, RCW~120) with SINFONI on the VLT. The main results of our study are:  \\begin{itemize}  \\item[$\\bullet$] We have identified the ionizing sources of all three \\HII\\ regions. RCW~79 is powered by a cluster of O stars unraveled here for the first time. The most massive stars have a spectral type O4--6V/III. A number of later type OB stars are also detected. RCW~82 is ionized by two O9--B2V/III stars and a third star, probably of spectral type Be. Finally, a single O6--8V/III star ionizes the \\HII\\ region in RCW~120.  \\item[$\\bullet$] The ionizing stars of each region have been analyzed with atmosphere models computed with the code CMFGEN. The derived stellar properties have been used to build the HR diagram. In the case of RCW~79, the large number of stars allowed a reliable age determination: the ionizing cluster formed 2.3$\\pm$0.5 Myr ago. For the other two \\HII\\ regions, similar ages are suggested but are less well constrained due to the small number of sources. An upper limit of $10^{-7}$ \\myr\\ was derived for the mass loss rate of the ionizing stars of all three regions.  \\item[$\\bullet$] The cumulative mass loss rate due to the O stars at the center of RCW~79 is $< 10^{-6}$ \\myr. The resulting wind mechanical luminosity is about $10^{-3}$ the ionizing luminosity. This is ten times weaker than in the numerical simulations of \\citet{freyer03}, raising the question of the quantitative role of winds on the dynamics of the three \\HII\\ regions studied here. The presence of hot dust emission close to the ionizing stars argues against a strong influence of stellar winds.  \\item[$\\bullet$] Spectroscopy of the infrared sources on the borders of the \\HII\\ regions revealed typical signatures of young stellar objects. Four main categories of YSOs have been identified: (1) sources dominated by H$_{2}$ emission lines, (2) sources showing mainly \\brg\\ emission, (3) sources with CO bandheads in emission and (4) sources with CO bandheads in absorption.  \\item[$\\bullet$] Near and mid--infrared color--color and color--magnitude diagrams indicate that all YSOs have SEDs dominated by non stellar emission. In the Spitzer/IRAC [3.6]$-$[4.5] vs [5.8]$-$[8.0] diagram, essentially all YSOs are grouped at the position between class I and class II objects. Only a couple of H$_{2}$ sources seem to be redder. The trend is more clearly seen in the [24] vs $K$--[24] diagram where the H$_{2}$ sources form clearly a distinct group with larger 24~\\mum\\ emission and redder $K$--[24] colors.  \\item[$\\bullet$] The detailed analysis of the H$_{2}$ emitting sources through line ratios and excitation diagrams indicate that thermal emission is the dominant mechanism in the majority of sources. Slow J--shocks or fast C--shocks are the preferred explanation for this emission. Some sources show a possible weak contribution of fluorescence emission, and three are dominated by this mechanism.  \\item[$\\bullet$] Measurement of radial velocities and line width was performed in a few sources. In general, the derived radial velocities are consistent with that of the ionized gas in the \\HII\\ region, confirming that the YSOs are well associated with the regions. The exception is RCW~120, for which most sources move 20 to 50 \\kms\\ faster than the ionized gas. In sources showing both \\brg\\ and H$_{2}$ lines, the former are systematically wider (150/250 \\kms vs non--resolved/100 \\kms), possibly due to a different spatial origin.  \\item[$\\bullet$] Morphological studies reveal that when both \\brg\\ and \\htwo\\ emission are present, \\brg\\ is centered on the position of the YSO while \\htwo\\ emission is more extended. In some cases, the \\htwo\\ emission peak is offset compared to that of the YSO. RCW120 51 presents a double \\htwo/\\brg\\ cone structure following the ionization front. RCW82 7 (and possibly RCW120 49 and RCW120 C4--67) has a \\htwo\\ velocity structure typical of an outflow.  \\item[$\\bullet$] The spectrophotometric and kinematic properties of YSOs paint the following plausible picture: H$_{2}$ dominated sources are stellar objects surrounded by an envelope in which shocks (due to jets or stellar/disk winds) produce H$_{2}$ emission; other YSOs correspond to more evolved sources in which the envelope partially or totally disappeared, revealing a disk structure from which CO and/or \\brg\\ emission is produced.  \\end{itemize}   Future tailored observations of a few key YSOs will provide more information on the dynamics of the circumstellar material. Additional UV and optical spectroscopy of the ionizing stars of RCW~79 will help to refine the age estimate and to reveal a possible age difference with the stars in the CHII region on the border of RCW~79. This would be a first direct proof of the existence of triggered massive star formation. Better estimates of the mass loss rates will also be possible, helping to quantify the role of stellar winds in the dynamics of \\HII\\ regions."
    },
    "0911/0911.5221_arXiv.txt": {
        "abstract": "A model is presented for generation of fast solar wind in coronal holes, relying on heating that is dominated by turbulent dissipation of MHD fluctuations transported upwards in the solar atmosphere. Scale-separated transport equations include large-scale fields, transverse Alfv\\'enic fluctuations, and a small compressive dissipation due to parallel shears near the transition region. The model accounts for proton temperature, density, wind speed, and fluctuation amplitude as observed in remote sensing and \\emph{in situ} satellite data. ",
        "introduction": "An open question in solar and heliospheric physics is to identify the physical processes responsible for heating the corona and accelerating the fast solar wind streams emanating from coronal holes.    This requires that a fraction of the energy available in photospheric motions be transported through the chromospheric transition region, and dissipated in the corona. The measured speeds of fast solar wind streams require spatially extended heating \\citep{WithbroeNoyes77,HolzerLeer80,Withbroe88}. The physical mechanisms for this transport and dissipation have remained elusive. Some models have resorted to use of a parametrically defined heat deposition (a ``heat function'') that decays exponentially with height, or anomalous heat conduction that redistributes energy along field-lines \\citep{HabbalEA95,McKenzieEA95,BanaszkiewiczEA98}. One-dimensional (1D) models of this kind, extending from the chromosphere to 1\\,AU \\citep{HansteenEA94,HansteenEA97}, have helped in understanding the regulation of the solar wind mass flux and can reproduce fast solar wind streams originating in cool electron coronal holes.  Here we present a model that demonstrates solar wind acceleration due to heating by a quasi-incompressible turbulent cascade triggered by coronal stratification \\citep{MatthaeusEA99}, and supplemented by compressive heating near the base. This model accounts for most presently available coronal and interplanetary observations.  The idea that broadband plasma fluctuations might heat the extended corona and accelerate the solar wind has long been discussed \\citep{Coleman68,BelcherDavis71,Hollweg86,HollwegJohnson88,Velli93a}. However, the mechanisms of transfer of fluctuation energy to small scales---and, in particular, the role of Alfv\\'enic turbulence (as observed in the solar wind) and cascade processes (e.g., phase mixing, ponderomotive driving, shocks, etc.)---have not been described self-consistently. Two recent papers shed light on these relationships \\citep{SuzukiInutsuka05,CranmerEA07}.  \\citet{SuzukiInutsuka05} incorporate 1D compressive nonlinear interactions driven by Alfv\\'en waves and leading to shock heating. This model produces good agreement with solar wind speed profiles. As low-frequency Alfv\\'en waves propagate upward, their wave pressure compresses the plasma. Unable to refract or mode-mode couple into a perpendicular wavenumber cascade, these waves must dissipate in 1D shock fronts. This model provides a valuable demonstration that MHD fluctuations can act as a conduit to transport energy to the requisite altitudes. However the restriction to 1D cascade is at odds with the well-established propensity for an incompressible MHD cascade to proceed mainly through wavevectors perpendicular to a strong mean magnetic field \\citep{RobinsonRusbridge71,ShebalinEA83,% OughtonEA94,BieberEA96,ChoEA02-strongB0}. Furthermore the corona exhibits a clear transverse structuring, and the initial fluctuations must have perpendicular correlation lengths not much larger than a super-granulation scale ($15\\,000$\\,km).  Another recent model \\citep{CranmervanBall05,CranmerEA07} incorporates a low-frequency cascade model \\citep{VerdiniEA06-soho}; however, the treatment of propagation and dissipation differs significantly from the present approach. Their scheme treats nonlinear effects as a perturbation, and it is unclear if it converges for strong turbulent heating. Here we employ a strong turbulence closure. We also do not rely on electron heat conduction to boost radial energy transport. Instead we compute an internal energy associated with the protons only. This approach supports comparisons with results employing improved representation of turbulence, such as shell models \\citep{VerdiniEA09-emp, VerdiniEA09-apj} and (potentially) full MHD simulation.  We find here that reflection of Alfv\\'enic turbulence alone does not lead to a full corona/solar wind stationary state --  a compressible contribution is required. This is supported by recent observations from \\emph{Hinode} \\citep{LangangenEA08,DePontieuEA09}.  % When we include a small component of compressive heating near the coronal base, fast solar wind streams are then accounted for. ",
        "conclusions": "The above model shows that turbulence near the coronal base, originating through chromospheric transmission of fluctuations, can heat the plasma in an expanding coronal hole flux-tube and produce a fast solar wind that matches a number of observational constraints. The turbulence is mainly of the low-frequency Alfv\\'enic type.  A small amount of compressive heating between the transition region and the sonic point appears to be needed to match the observations. This additional heating may be due to type II spicules that supply broadband low-frequency vertical fluctuations at transition region heights, thus launching compressive MHD modes near the coronal base \\citep{DePontieuEA09}.   Most of the fluctuation energy is in low-frequency turbulence, and this sustains a strong anisotropic MHD cascade through reflections from local density gradients. This type of anisotropic cascade is favored in MHD turbulence in the presence of a strong DC magnetic field \\citep{RobinsonRusbridge71,ShebalinEA83,OughtonEA94}. Heat conduction does not enter the present model at all, since it mainly affects electron internal energy, which evolves independently in this approximation. Similar assumptions work well in understanding observations of solar wind turbulence \\citep{BreechEA09,CranmerEA09}. In these ways, the present model differs from other recent models that incorporate turbulent heating \\citep{SuzukiInutsuka05,CranmerEA07}. In particular we believe that this model demonstrates, possibly for the first time, that a model (almost) free of ad hoc heat functions, artificial equations of state, and ad hoc assumptions about heat conduction, can indeed heat the corona and accelerate the solar wind.   We plan further study of the required small amount of compressible heating, attributed here to spicule-driven magnetosonic activity. Another\tuseful extension would be to include separate electron and proton internal energy budgets, which will enable additional observational constraints, and will permit study of the role of kinetic dissipation processes \\citep{BreechEA09,CranmerEA09}."
    },
    "0911/0911.2771_arXiv.txt": {
        "abstract": "We analysed the orientation of galaxy groups in the Local Supercluster (LSC). It is strongly correlated with the distribution of neighbouring groups in the scale till about 20 Mpc.  The group major axis is in alignment with both the line joining the two brightest galaxies  and the direction toward the centre of the LSC, i.e. Virgo cluster. These correlations suggest that two brightest galaxies were formed in filaments of matter directed towards the protosupercluster centre. Afterwards, the hierarchical clustering leads to aggregation of galaxies around these two galaxies. The groups are formed on the same or similarly oriented filaments. This picture is in agreement with the predictions of numerical simulations. ",
        "introduction": "\\citet{b2} was the first who found that major axes of galaxy clusters tend to point towards their neighbours.  Later the existence of this effect was discussed by several authors and usually the significant alignment was reported. The distance between clusters for which the effect was detected changed from $10 h^{-1} Mpc$ till $150 h^{-1} Mpc$ (where $h=H_0/100\\,km\\,s^{-1} Mpc^{-1}$). The strength of the effect decreases with distance \\citep{s2,f3,r1,u1,p2,w4,c1,h3}. These investigations involved both optical and X-ray data, as well as clusters belonging (or not) to superclusters. Nowadays, it is accepted that the effect is not due to selection effects but is real and its distance scale is between $10 - 60 h^{-1} Mpc$. The alignment of galaxy group was studied by \\citet{w3}.  He used CfA group catalog \\citep{g1} and a catalog based on SSRS ( Southern  Sky Redshit Survey) \\citep{m1}. Each group should have at least $4$ objects and less than $100$. Decontamination of the groups by foreground and background objects was performed by removing objects with redshift difference from the group mean over $1000\\, km\\,s^{-1}$. They were $59$ groups and Binggeli effect was observed among galaxy groups till $15- 30 h^{-1} Mpc$. One should note that the investigation of galaxy group orientation is more difficult than in the case of galaxy clusters. The position angle of a group consisting of a few objects is determined with much greater error than for rich clusters. Moreover, statistics with small number are less reliable. The other investigations of galaxy groups were performed by \\citet{p1}. During study the orientation of $92$ out of $100$ compact groups listed in the Catalogue of the Compact Groups \\citep{h5}, they do not find alignment, but location of groups along  long chains is  noted.  The interpretation of the effect has changed with development of theories, but the main idea that this should reflect conditions during the structure formation is still very popular. Numerical simulations gave a better understanding of physical processes leading to structure formation. These simulations were performed in the framework of the cold dark matter (CDM) model presently regarded as the correct description of the large scale structure formation. Using different approaches and codes, these investigations led to the conclusion that the preferred orientation of galaxy clusters in the CDM model is a natural consequence of processes leading to structure formation due to gravitational interaction. \\citet{o1} using large - scale simulation found  that in  $\\Lambda$CDM cosmological model effect reaches the distance up to $30 Mpc$, while in $\\tau$CDM model the range of effects is twice smaller.  Also in SCDM and OCDM models, where smaller scale simulations were performed, some alignment effect could be noted. N - body simulation for standard $\\Lambda$CDM model \\citep{f1} for $3000$ clusters showed the alignment  of neighbouring clusters  in the distance range from $10 - 15 Mpc$, while  the Binggeli effect till about $100 Mpc$ is observed. The strong alignment, decreasing with the increasing distance between clusters, observed till about 100 Mpc was reported \\citep{h6}. The preferred orientation of clusters belonging to superclusters existed also in SHM + N body simulation . Moreover, from some of the numerical simulations, it follows that the structure formation occurred along filamentary structures rather than the walls \\citep{f2,s1,b2,h1,h2,a1,w1,w2}. In order to confirm, or deny, this scheme of structure origin we carry out an analysis of the Local Supercluster (LSC) galaxy groups alignment as well as the distribution of the acute angle between the position angle of the structure and the direction to all remaining clusters. Groups were taken from Nearby Galaxies (NBG) Catalog \\citep{t3}. We investigated also the alignment of the brightest group galaxy and the parent group. We determined the position of the line joining the two brightest galaxies and checked the orientation of this line in respect to parent group and direction towards Virgo cluster. The distributions of these angles, as well as differences between some of these angles were tested for isotropy.  The paper is organised in the following manner. Section 2 describes observational data, section 3 presents  statistical method used in the paper, section 4 presents results and discussions. In section 5 we formulate conclusions. ",
        "conclusions": "We used two samples of data. The first one contains $61$ groups having at least $10$ members, in agreement with West's (1989) wish to have such sample of data. From the second set of data poorer groups were eliminated and $35$ groups having at least $20$ members remained for analysis. All groups are nearby and they are located within the Local Supercluster.  Our main results are: \\begin{enumerate} \\item The group major axis is in alignment with the line joining the two brightest galaxies. \\item The group major axis is in alignment with the direction toward the centre of the LSC. \\item The acute angle between the position angle of the group and direction towards each remaining group is not isotropic. \\item The structures have tendency to point each other when the distance between groups is smaller than $20 Mpc$  \\end{enumerate}  We performed study of the acute angle between the position angle of the group and direction towards each remaining group denoted as $\\phi$. The sample was divided according to the distance between group centres. Each subsample was analysed independently using two coordinate systems. In the equatorial coordinate system the distributions of the $\\phi$ - angle for all subsamples disregarding subsample $D<10 Mpc$ were isotropic. This is not the case of the coordinate system connected with the Local Supercluster. We found that for closer neighbours ($D<10 Mpc$) strong alignment is observed. It is at almost  $5 \\sigma$ level. The Binggeli effect is diminishing with distance increase, vanishing at about $20 Mpc$. We used the K - S test, which is usually applied for alignment investigation. \\citet{c2}  criticised it, because it does not point the place, where the departure of isotropy is observed. They preferred to use the Wilcox test on rank-sum, which, in their opinion, gave a higher confidence signal for alignment. However, \\citet{o1} applied both tests finding a little difference in the obtained statistics. Furthermore,  we used the K - S test only to check  the isotropy of the distribution and not to find the anisotropy location. Our results obtained with the help of K-S test are confirmed by linear regression analysis. The fact that detection of anisotropy is connected with LSC coordinate system supports the point of view that formation of galaxies occurred within protostructures. The analysis of the differences between position angles shows that it is possible that there exists the alignment of the line joining two brightest galaxies with both position angle of the parent group and direction towards Virgo cluster centre. The fact that detection of anisotropy is connected with LSC coordinate system supports the point of view that formation of galaxies occurred within protostructures. The analysis of the differences between position angles shows that it is possible that there exists the alignment of the line joining two brightest galaxies with both position angle of the parent group and direction towards Virgo cluster centre. From the presented analysis of the orientation of galaxy groups in the Local Supercluster the following picture of the structure formation appears. The two brightest galaxies were formed first. They originated in the filamentary structure directed towards the centre of the protocluster. This is the place where the Virgo cluster centre is located now. Due to gravitational clustering, the groups are formed in such a manner that galaxies follow the line determined by the two brightest objects. Therefore, the alignment of structure position angle and line joining two brightest galaxies is observed. The other groups are forming on the same or nearby filament. The flatness of the LSC additionally contributes to the observed alignment of galaxy groups. The majority of the groups lie close to us. Due to completeness of the Catalog, the lack of groups further than the Virgo Cluster centre is observed, but nearby groups are very well selected and they contain only more massive galaxies. This picture is in agreement with predictions of several CDM models, in which structure formation is due to hierarchical clustering. Moreover, the formation is occurring on the filamentary structure. The further investigation considering groups clearly  inside  and outside superclusters  on the greater data set will be very useful to support or reject this picture."
    },
    "0911/0911.5198_arXiv.txt": {
        "abstract": "Uninhibited radiative cooling in clusters of galaxies would lead to excessive mass accretion rates contrary to observations. One of the key proposals to offset radiative energy losses is thermal conduction from outer, hotter layers of cool core clusters to their centers. However, thermal conduction is sensitive to magnetic field topology. In cool-core clusters where temperature decreases inwards, the heat buoyancy instability (HBI) leads to magnetic fields ordered preferentially in the direction perpendicular to that of gravity, which significantly reduces the level of conduction below the classical Spitzer-Braginskii value. However, the cluster cool cores are rarely in perfect hydrostatic equilibrium. Sloshing motions due to minor mergers and stirring motions induced by cluster galaxies or active galactic nuclei (AGN) can significantly perturb the gas. The turbulent cascade can then affect the topology of the magnetic field and the effective level of thermal conduction. We perform three-dimensional adaptive mesh refinement magnetohydrodynamical (MHD) simulations of the effect of turbulence on the properties of the anisotropic thermal conduction in cool core clusters. We show that very weak subsonic motions, well within observational constraints, can randomize the magnetic field and significantly boost effective thermal conduction beyond the saturated values expected in the pure unperturbed HBI case. We find that the turbulent motions can essentially restore the conductive heat flow to the cool core to level comparable to the theoretical maximum of $\\sim 1/3$ Spitzer for a highly tangled field. Runs with radiative cooling show that the cooling catastrophe can be averted and the cluster core stabilized; however, this conclusion may depend on the central gas density. Above a critical Froude number, these same turbulent motions also eliminate the tangential bias in the velocity and magnetic field that is otherwise induced by the trapped $g$-modes. Our results can be tested with future radio polarization measurements, and have implications for efficient metal dispersal in clusters. \\\\ ",
        "introduction": "\\label{section:intro}  X-ray clusters show a strong central surface brightness peak, and many of them have short central cooling times. However, {\\it Chandra} and {\\it XMM-Newton} observations have shown that actual gas cooling rates fall significantly below classical values from a standard cooling flow model; emission lines such as Fe XVII expected from low temperature gas are not observed \\citep{peterson06}. This suggests that besides quasi-hydrostatic equilibrium, clusters might also be in quasi-thermal equilibrium, with some form of heating balancing cooling. Many forms of heating have been postulated. Hybrid models can account for the observations; leading contenders generally involve some combination of AGN heating and thermal conduction. The remarkable discovery by {\\it Chandra} of AGN blown bubbles in clusters has triggered a widespread renaissance in the study of AGN feedback (see \\citet{mcnamara07} for a recent review). It has been suggested that the AGN feedback {\\it alone}, generally fails to distribute heat in the  spatially distributed fashion required to offset cooling (e.g., Vernaleo \\& Reynolds 2006), both in its radial dependence and solid angle coverage. However, recent simulations by Br{\\\"u}ggen \\& Scannapieco (2009) involving subgrid model for Rayleigh-Taylor-driven turbulence show that AGN heating can self-regulate cool cores. In these non-MHD simulations the subgrid turbulence model is crucial to simulate the evolution of the cool core. In a realistic situation, the ability to self-regulate undoubtely depends on a range of the mechanisms used to transmit the energy from the AGN blown bubbles to the ICM (e.g., dissipative sound waves, $pdV$ work, cosmic rays, suppression of both the Rayleigh-Taylor instability and the mixing of the bubble material by (even weak) magnetic fields). These issues are a matter of intense debate. At the same time, it is a remarkable fact that if one simply uses observed temperature profiles in clusters to construct the Spitzer conductive flux from the cluster outskirts, it is very nearly equal to that required to balance the radiative cooling rate as indicated by the observed X-ray surface brightness profile, for some reasonable fraction $f\\sim 0.3$ of the Spitzer value (e.g., Fig. 17 of Peterson \\& Fabian 2006). There is no reason in principle why such close agreement should exist, and it is a tantalizing hint that nature somehow ``knows'' about Spitzer conductivity. Models which invoke both mechanisms are able to reproduce observed temperature and density profiles without any fine-tuning of feedback or conduction parameters \\citep{ruszkowski02,guo08a}, unlike conduction-only models \\citep{bertschinger86,zakamska03}.  Nonetheless, even if no fine-tuning of conductivity is required, it is unsatisfying to invoke an ad-hoc suppression factor $f$; the anisotropic transport of heat along field lines can be calculated from first principles. If the field lines are tangled and chaotic, then $f\\sim 0.2$ might be appropriate \\citep{narayan01}. However, recent work has uncovered instabilities arising in stratified plasmas, which rearrange the field lines on large scales, with strong implications for heat transport. Unlike convective instability, which only arises if entropy declines with radius, these instabilities depend on temperature gradients, and arise irrespective of the sign of the gradient. If $dT/dr < 0$, as in the outer regions of clusters, the resulting magnetothermal instability (MTI; \\citet{balbus00,parrish05,parrish07,parrish08}) drives field lines to become preferentially radial, substantially boosting conductivity to close to the Spitzer value, $f \\sim 1$. On the other hand, if $dT/dr > 0$, as in the cores of cool-core clusters, then the resulting heat buoyancy instability (HBI; \\citet{quataert08, parrish09}, Bogdanovi{\\'c} et al. 2009) will drive field lines to become preferentially tangential, essentially shutting off radial conduction ($ f \\ll 1$), and leading to catastrophic cooling in the cluster core. Thus, it would seem that a more careful calculation vitiates the ad-hoc assumptions of models which invoke conduction.  However, this presumes that there are no competing mechanisms which also affect field line topology. One possibility which has been speculated upon in the literature \\citep{ruszkowski07,guo08b,bogdanovic09} is AGN bubbles, which upon rising would leave radially orientated field lines in their wake. \\citet{guo09} showed that time-variable conduction, perhaps triggered by AGN outbursts, could allow clusters to cycle between the cool core and non-cool core states. Another, perhaps more general mechanism is simply turbulent gas motions, which can be induced by mergers \\citep{ascasibar06}, the motion of galaxies \\citep{kim07,conroy08}, AGN outbursts \\citep{mcnamara07}, or cosmic-ray driven convection \\citep{chandran07,sharma09}\\footnote{Note, however, that since the HBI is a buoyancy instability rather than a convective instability, turbulence cannot be self-consistently generated by the HBI itself, since the induced velocities are small \\citep{bogdanovic09,parrish09}.}. Indeed, \\citet{sharma09a} suggested that the amount of turbulence needed to avoid the stable end-state induced by the HBI could be estimated from the Richardson number \\citep{turner73}, which is simply the ratio of the restoring buoyant force to the inertial $(\\rho {\\bf u}\\cdot \\nabla {\\bf u})$ term. If one defines $Ri \\equiv g r ( d \\, {\\rm ln}T/d \\, {\\rm ln} r)/\\sigma^{2}$ for typical cluster conditions, the critical turbulent velocity is \\citep{sharma09a}:  \\begin{eqnarray} \\sigma \\approx 135 \\, {\\rm km \\, s^{-1}} {g_{-8}}^{1/2} r_{10}^{1/2}\\left( \\frac{d \\, {\\rm ln}T/d \\, {\\rm ln} r}{0.15} \\right)^{1/2} \\left( \\frac{Ri_{c}}{0.25} \\right)^{-1/2}\\\\ \\nonumber \\end{eqnarray}  \\noindent where $g_{-8}$ is the gravitational acceleration in units of $10^{-8} \\, {\\rm cm^{2} s^{-1}}$, $r_{10}$ is a characteristic scale height in units of $10$ kpc, and $Ri_{c}$ is the critical Richardson number. $Ri_{c} \\sim 1/4$ is typical for hydrodynamic flow; order unity differences might exist for the MHD case. Such levels of turbulence are easily seen in cosmological simulations, even in very relaxed clusters \\citep{evrard90,norman99,nagai03,vazza09a, vazza09b}, presumably due to continuing gas accretion and the supersonic motions of galaxies through the ICM. Gas ``sloshing'' in clusters has been invoked to explain observed cold fronts \\citep{ascasibar06}, and biases in X-ray mass profiles, particularly in the cluster core \\citep{lau09}. In the future, direct measurement of turbulent velocities could best be performed by a high-spectral resolution imaging X-ray spectrometer \\citep{rebusco08}, although early upper limits exist from {\\it XMM-Newton} (\\citet{sanders09}; for A1835, non-thermal broadening is less than $\\sim 275 \\, {\\rm km \\, s^{-1}}$). Besides the strong impact on conductive heat transport, the interaction of turbulence with conduction has testable implications for metal dispersal (which is strongly enhanced; \\citep{sharma09}). With an instrument as as the Square Kilometer Array, the field topology itself can potentially be mapped out via the rotation measure of background radio sources.  Our goal in this paper is to quantitatively examine for the first time if indeed turbulent motions can overwhelm the tangential field topology driven by the HBI (we defer the MTI dominated regime to future work), and allow thermal conduction at the level required to avert a cooling catastrophe. We begin with a MHD {\\it FLASH} simulation of a cluster which in its unperturbed state suffers the HBI, consistent with previous work. We then simulate turbulence via a spectral forcing scheme that enables statistically stationary velocity fields (Eswaran \\& Pope 1998). An important feature of our study is that even if the medium is stirred isotropically, the turbulent velocity field itself can be anisotropic, and interact in non-trivial ways with thermal conduction. This is because when the turbulent driving frequency falls below the \\brunt frequency, {\\it g}-modes can be excited; the trapped modes induce tangentially biased vortices, due to the restoring forces which act on vertical motion \\citep{riley00}. Thus, even in the absence of the HBI, the excitation of $g$-modes can tangentially bias the magnetic field. In realistic clusters, $g$-modes also greatly enhance the efficacy of stirring, since trapped modes allow the turbulence to be volume-filling, although our present simulations already have volume-filling turbulence by design. The \\brunt frequency, which determines where $g$-modes can be produced, depends both on the gravitational potential and the gas temperature and density profile. The latter can in turn be affected by heating and thermal conduction. Thus, the cluster potential, gas properties, and thermal conduction interact in often subtle ways. Overall, we find that turbulence can indeed restore conductive heat flow back to the values $f\\sim 0.3$ expected for a highly tangled field, and required to stem a cooling flow.  The organization of this paper is as follows. In \\S\\ref{section:turbulence}, we describe simple scaling relations to guide our intuition; in \\S\\ref{section:methods}, we describe our computational methods; in \\S\\ref{section:results}, our results; we then conclude in \\S\\ref{section:conclusions}. ",
        "conclusions": "\\label{section:conclusions}  We have shown that a very low level of turbulent perturbations, $\\sim 50-150 \\, {\\rm km/s}$ to the ICM --- entirely consistent with the expectations for cosmological infall, galaxy motions, mergers, or AGN activity --- can entirely alter the magnetic field distribution resulting from the HBI instability. Instead of preferentially tangential magnetic fields, the final field configuration can be easily randomized. As expected from simple theoretical considerations, this happens when the typical Froude number (equation \\ref{eq:Fr_MHD}) $Fr^{\\rm MHD} \\gsim \\mathcal{O}(1)$. Note that weak, low frequency perturbing motions lead to trapped $g$-modes and result in preferentially tangential gas motions, due to strong restoring buoyancy forces in the vertical direction; thus, magnetic fields could in principle become tangential even in the absence of thermal conduction and the HBI. However, this tangential velocity bias is also lifted for $Fr_{\\rm MHD} \\gsim \\mathcal{O}(1)$. This has several immediate consequences: \\begin{itemize} \\item{The boost in thermal conduction restores the heat flow from the hot outer layers to the central cool core. This is crucial to averting massive overcooling and heating the cluster cores in a stable fashion. Indeed, our runs with radiative cooling showed that a cooling catastrophe was averted even in the absence of AGN heating. Thus, naive models which assume thermal conductivity with Spitzer fraction $f\\sim 1/3$ may perhaps be reasonably applied to clusters after all. Since conduction is regulated by turbulence, it will in general be time-dependent. A sudden boost in thermal conduction can modulate a transition between a cool core (CC) to a non cool core (NCC) cluster (\\citet{guo09}; we also see some hints of this in Fig \\ref{plot4}). Thus, mergers, while generally insufficient in the pure hydrodynamic case to effect a CC to NCC metamorphosis \\citep{poole08}, may be effective once the boost in thermal conduction due to additional turbulence is taken into account. If AGN are the main source of turbulence, this implies feedback between accretion onto the AGN and thermal conduction. Thus, AGN could exert another level of self-regulation and thermostatic control beyond straight kinetic heating. } \\item{Our predictions might be testable in the future: magnetic topology can be measured by radio polarization measurements (e.g., Pfrommer \\& Dursi 2009), while turbulence can be measured by a high spectral-resolution calorimeter. Alternatively, our predictions can be used to provide indirect constraints on turbulence. For instance, the absence of large-scale ordered fields would require a minimal level of turbulence, and argue against a fully relaxed ICM. The resulting turbulent pressure support affects mass estimates which assume hydrostatic equilibrium \\citep{lau09}, and the use of clusters for cosmology.} \\item{The distribution of metals in clusters is known to be broader than that of the stars. As shown by \\citet{sharma09,sharma09a}, metal mixing in a stratified plasma is much more effective once conduction is at play, because of the weaker nature of restoring buoyancy forces. This allows broad metallicity profiles to be obtained without inversion of the central entropy profile (which is not observed). The restoration of thermal conduction by stirring motions greatly aids this process. In the limit where conduction is so efficient that the cluster becomes isothermal (so that $\\omega_{\\rm BV}^{\\rm HBI} \\rightarrow 0$) or conduction becomes isotropic, fluid elements have the same temperature (and hence density) as their surrounding, and fluid elements become neutrally buoyant. This makes metal mixing maximally efficient. We note that any mechanism which disperses metals over wide distances (e.g., turbulent diffusion, Rebusco et al. 2008) generally excites turbulent velocities sufficient to isotropize the magnetic field. Thus, the relatively broad metallicity profiles in clusters is consistent with non-negligible stirring.} \\end{itemize}  Our stirring methods ensure by construction that the ensuing turbulence is volume-filling. However, this is by no means guaranteed in nature. For instance, galactic wakes will be relatively narrow, with a cross-section of order the galaxy size; similarly, rising AGN bubbles will not induce velocity fluctuations over a large solid angle. A straightforward way to ensure volume-filling turbulence, as is required to stem the HBI, is to excite trapped $g$-modes which are repeatedly reflected and focused within a resonance region where $\\omega < \\omega_{\\rm BV}^{\\rm MHD}$, as discussed in \\S\\ref{section:turbulence}. At the same time, $\\omega > \\omega_{\\rm BV}^{\\rm MHD}$ is required to overwhelm buoyancy forces, an apparently contradictory requirement. Of course, the assumption of a single frequency is too simplistic: stirring motions contain many harmonics and will be scale dependent, generally increasing in frequency as turbulence cascades toward smaller scales. Thus, large scale $g$-modes could fall below the \\brunt frequency, ensuring trapping and amplification, but the resulting small scale turbulence could potentially still have turnover frequencies exceeding $\\omega_{\\rm BV}^{\\rm MHD}$. We will examine this in future work.  Although we have focused on the HBI in this paper, the same Froude number considerations apply with equal force to the MTI unstable outer regions of the cluster, where temperature falls with radius and the MTI causes field lines to become preferentially radial. Here, if $Fr^{\\rm MHD} > \\mathcal{O}(1)$, turbulent motions can also overwhelm buoyant forces and randomize the field. Recently, there has been tantalizing evidence from the polarized emission surrounding the magnetic drapes of galaxies sweeping up magnetic fields in the Virgo cluster, that outside a central region, the magnetic field might be preferentially oriented radially \\citep{pfrommer09}, as one might expect from the MTI. If this result continues to hold up, this would imply the Froude number is below the critical value at these radii (thus providing an indirect constraint on turbulent velocities), or additional physical processes, not modeled here, are at play. \\\\"
    },
    "0911/0911.3820.txt": {
        "abstract": "The present operation of the ground-based network of gravitational-wave laser interferometers in ``enhanced'' configuration and the beginning of the construction of second-generation (or advanced) interferometers with planned observation runs beginning by 2015 bring the search for gravitational waves into a regime where detection is highly plausible. The development of techniques that allow us to discriminate a signal of astrophysical origin from instrumental artefacts in the interferometer data and to extract the full range of information are therefore some of the primary goals of the current work. Here we report the details of a Bayesian approach to the problem of inference for gravitational wave observations using a network (containing an arbitrary number) of instruments, for the computation of the Bayes factor between two hypotheses and the evaluation of the marginalised posterior density functions of the unknown model parameters. The numerical algorithm to tackle the notoriously difficult problem of the evaluation of large multi-dimensional integrals is based on a technique known as Nested Sampling, which provides an attractive (and possibly superior) alternative to more traditional Markov-chain Monte Carlo (MCMC) methods. We discuss the details of the implementation of this algorithm and its performance against a Gaussian model of the background noise, considering the specific case of the signal produced by the in-spiral of binary systems of black holes and/or neutron stars, although the method is completely general and can be applied to other classes of sources. We {also} demonstrate the utility of this approach by introducing a new coherence test to distinguish between the presence of a coherent signal of astrophysical origin in the data of multiple instruments and the presence of incoherent accidental artefacts, and the effects on the estimation of the source parameters as a function of the number of instruments in the network. ",
        "introduction": "\\label{s:intro}  Searches for gravitational waves are entering a crucial stage with the network of ground-based laser interferometers -- LIGO~\\cite{BarishWeiss:1999, lsc-s5-instr}, Virgo~\\cite{virgo} and GEO\\, 600~\\cite{geo} -- now fully operational and engaged in a new year-long data taking period~\\cite{Smith:2009,eligo} at ``enhanced'' sensitivity, which may allow the first direct detection of gravitational radiation. Construction has already begun for the upgrade of the instruments to advanced configuration (second generation interferometers) with installation at the sites that will start at the end of 2011~\\cite{Smith:2009,advligo,advvirgo}. {When} science observations resume at much improved sensitivity {by 2015}, several gravitational-wave events are expected to be observed, opening a new means to explore a variety of astrophysical phenomena (see \\emph{e.g.} Ref.~\\cite{Cutler:2002me,Kokkotas:2008} and references therein).  Coalescing binary systems of compact objects -- black holes and neutron stars -- will be the workhorse source for gravitational wave observations. Ground-based laser interferometers will monitor the last seconds to minutes of the coalescence of these systems. The theoretical modelling of the (in-spiral) waveform is well in hand (see \\emph{e.g.}~\\cite{Blanchet:2006} and references therein), and the search algorithms are well understood~\\cite{lsc-cbc-s5-12-to-18, lsc-cbc-s5y1, lsc-cbc-spins, lsc-cbc-s3s4, lsc-cbc-s2-bbh, lsc-cbc-primordial-bh, lsc-cbc-s2-bns, lsc-cbc-s1}. The detection rate for on-going searches and observations with second generation instruments is estimated to lie in the range $9\\times 10^{-5}\\,\\mathrm{yr}^{-1} - 0.7\\,\\mathrm{yr}^{-1}$ and $0.2\\,\\mathrm{yr}^{-1} - 1000\\,\\mathrm{yr}^{-1}$, respectively, see~\\cite{cbc-rates} for a review. It is likely that in a few years time ground-based laser interferometers will allow us to extract a wealth of new information ranging from the formation and evolution of binary stars, the nature of precursors of (short) gamma-ray bursts, dynamical processes in star clusters, and could yield a new set of standard candles for precise cosmography.  As instruments are beginning to operate at a meaningful sensitivity from an astrophysical, cosmological and fundamental physics point of view, much emphasis is now being placed on the development of methods that offer the maximum discriminating power to separate {disturbances} of instrumental origin from a true astrophysical signal, and to extract the full range of information from the detected signals. Bayesian inference provides a powerful approach to both model selection (or hypothesis testing) and parameter estimation. Despite the conceptual simplicity of the Bayesian framework, there has been only limited use of these methods for ground-based gravitational-wave data analysis due to their computational burden, in this case related to the need to compute large multi-dimensional integrals. Additionally, the Gaussian likelihood functions considered so far do not address the instrumental glitches which are present in data from the current generation of gravitational wave detectors.  Here we present an efficient method to compute concurrently the full set of quantities at the heart of Bayesian inference: the Bayes factor between competing hypotheses and the posterior density functions (PDFs) on the relevant model parameters. The method is based on the \\emph{Nested Sampling} algorithm~\\cite{Skilling:AIP,Skilling:web,Sivia} to perform multi-dimensional integrals, that present the practical and computationally intensive challenge for the implementation of Bayesian methods. We demonstrate the algorithm by considering multi-detector observations of gravitational waves generated during the in-spiral phase of the coalescence of a binary system, modelled using the restricted post$^{2.0}$-Newtonian stationary phase approximation, which is the waveform used so far for searches of non-spinning binary objects~\\footnote{We note that searches for low-mass (that is with total mass $\\le 35\\,M_\\odot$) binary systems in the data now collected by LIGO and Virgo are based on templates computed at the post$^{3.5}$-Newtonian order. All the results presented here apply directly to that case at no additional computational costs.}. Initial results based on this method and applied to simplified gravitational waveforms were reported in~\\cite{VeitchVecchio:2008a,VeitchVecchio:2008b}. An application to the study of different waveform approximants to detect and estimate the parameters of signals generated through the numerical integration of the Einstein's equations for the two body problem in the context of the Numerical INJection Analysis (NINJA) Project was reported in~\\cite{CadonatiElAl:2009,AylottEtAl:2009,AVV:2009}. In this paper we: \\begin{itemize} \\item  Provide for the first time full details about the theoretical and technical issues on which the computation of the evidence is based; \\item  Discuss the errors associated with the computation of the integrals and the associated computational costs; \\item  Show how from the nested samples one can construct at negligible computational cost the {marginalised} posterior PDFs on the source parameters; \\item Demonstrate the performance of this technique {in detecting a binary in-spiral signal against a Gaussian model of background noise in coherent observations using a network of detectors, and} by introducing a new {\\emph{coherence test}} to distinguish between the presence of a coherent signal of astrophysical origin in the data of multiple instruments and the presence of incoherent accidental artefacts, and \\item {Show the effects of the number of instruments in the network on the estimation of the source parameters. } \\end{itemize} The method in its present implementation can be extended in a straightforward way to the full coalescence waveform of binary systems -- in fact the first example applied on the NINJA data set is described in~\\cite{AylottEtAl:2009,AVV:2009} -- and the software is available as part of the LSC Analysis Library Applications (LALapps)~\\cite{lal,lalapps}. Work is already on-going in this direction. Furthermore, the approach can also be extended to other gravitational-wave signals.  The paper is organised as follows: in Section~\\ref{s:modelselection} we review the key concepts of Bayesian inference and the signal generated by in-spiralling gravitational wave signals; in Section~\\ref{s:NSalgo} we describe our implementation of the numerical computation of multi-dimensional integrals for the computation of the marginal likelihood and marginalised posterior density functions based on {nested-sampling} the limited likelihood function; Section~\\ref{s:NSimplementation} contains the implementation details of the algorithm, including a quantification of the errors that affect the results as a function of the choice of the main tuning parameters of the algorithm and the effects on the scaling of the computational costs. Section~\\ref{s:results} contains the results from the applications of this method to several test cases, including a Bayesian coherence test that we introduce here for the first time.  Throughout the paper we use geometric units, in which $G = c = 1$. ",
        "conclusions": "By taking a Bayesian approach to the analysis of data for the detection and characterisation of in-spiral signals, we have been able to implement a conceptually simple yet flexible framework for drawing inference from observations. Although the Bayesian formalism calls for the evaluation and integration of high-dimensional likelihood functions, we have shown that the nested sampling {technique} provides us with a means to both search for and estimate the parameters of a signal. Further work remains to be done in improving the efficiency and reliability of detection at low masses, but the particular implementation that we describe here provides a solid basis for these future improvements.  We have used our implementation to demonstrate the power of Bayesian model selection in classifying putative gravitational wave signals, through the use of the coherence test described in Section \\ref{sec:coherent}. The coherence test goes some way to implementing a robust and Bayesian defence against glitches present in gravitational wave data which do not resemble coherent gravitational waves by providing an internal consistency check within the detector network. Tests of this kind may provide a useful new additional discriminator when analysing candidate gravitational waves, and are only achievable through the treatment of the detector network in a coherent way. Notably, this test is only possible through the use of the Occam factor and the comparison of probability distributions which differ in their dimensionality, and could not be possible with point estimates of maximum likelihood. Indeed, the maximum likelihood of the coherent and incoherent models can be trivially shown to be identical. In addition to such tests, it would also be desirable to model the detector data in a way which is non-stationary or Gaussian, but this is beyond the scope of this work. Furthermore, the use of the Bayesian parameter estimation framework is invaluable in inferring the signal parameters, and produces the full posterior probability distribution, not just maximum likelihood points and estimates of variance. The covariance and interdependence of the parameters does not prevent us from calculating consistent joint PDFs on the full parameter space even when point estimates have little meaning, and this allows us to see the benefit of using a network of detectors in a coherent fashion.  {One of} the main benefits in the use of the nested sampling and Bayesian evidence approach is the ease with which it can be extended or adapted to different signal models with minimal changes needed to the implementation. In other work past and ongoing, we have shown how this method may be applied to the discrimination between candidate waveforms when comparing to a numerical relativity simulation; how one may place bounds on the Compton wavelength of the graviton given a single or multiple observations of an in-spiral signal, and testing the effects of including spins in the detection and estimation of in-spiral signals \\cite{AylottEtAl:2009,WDP:lambdaG}\\,.  Future work on the implementation of the core algorithm will focus on achieving a better reliability and efficiency in the sampling of the parameter space, which should lead to further improvements in the performance in a real world situation, and on testing on the full set of waveform approximants used in present searches for coalescing binary systems. The work presented here forms the basis of the \\texttt{ lalapps\\_inspnest } program, which is can be found in the LALApps software distribution, released under the terms of the GNU General Public Licence \\cite{lalapps}."
    },
    "0911/0911.3176_arXiv.txt": {
        "abstract": "We report on the properties of pre-main-sequence objects in the Taurus molecular clouds as observed in 7 mid- and far-infrared bands with the Spitzer Space Telescope.  There are 215 previously-identified members of the Taurus star-forming region in our $\\sim$44 square degree map; these members exhibit a range of Spitzer colors that we take to define young stars still surrounded by circumstellar dust (noting that $\\sim$20\\% of the bonafide Taurus members exhibit no detectable dust excesses). We looked for new objects in the survey field with similar Spitzer properties, aided by extensive optical, X-ray, and ultraviolet imaging, and found 148 candidate new members of Taurus.  We have obtained follow-up spectroscopy for about half the candidate sample, thus far confirming 34 new members, 3 probable new members, and 10 possible new members, an increase of 15-20\\% in Taurus members. Of the objects for which we have spectroscopy, 7 are now confirmed extragalactic objects, and one is a background Be star. The remaining 93 candidate objects await additional analysis and/or data to be confirmed or rejected as Taurus members.  Most of the new members  are Class II M stars and are located along the same cloud filaments as the previously-identified Taurus members. Among non-members with Spitzer colors similar to young, dusty stars are evolved Be stars, planetary nebulae, carbon stars, galaxies, and AGN. ",
        "introduction": "\\label{sec:intro}  A complete inventory of all the coeval stars in a young stellar association, cluster, or group (hereafter ``association\") enables studies of the initial mass function (IMF), disk fraction, and stellar rotational properties, among other pursuits.  Information from associations with a range of ages enables understanding of the overall formation and evolution of young stars, including the change with time of disk fraction and stellar rotation rate.  However, identifying all of the member stars of a given young association can be quite difficult.  (By ``member,'' we mean objects that are clearly young, close to the same age, and often still associated with their natal cloud.)  Finding all such members requires that one employ multiple observational techniques.  These methods include but are not limited to X-ray surveys (e.g., Alcal\\'a \\etal\\ 1996, Wolk \\etal\\ 2006), H$\\alpha$ surveys (e.g., Ogura \\etal\\ 2002), variability surveys (e.g., Carpenter \\etal\\ 2001, Rebull 2001), ultraviolet (UV) surveys (e.g., Rebull \\etal\\ 2000), and infrared (IR) surveys (e.g., J{\\o}rgensen \\etal\\ 2006, Rebull \\etal\\ 2007). At each wavelength, we can use the fact that stellar youth implies more flux at a given radiometric band (X-rays, H$\\alpha$, UV, IR), or more flux variability, than older stars of comparable mass, allowing separation of association members from field contaminants.  When we combine the information from surveys in  multiple wavelengths, we must remember that  the influence of extinction due to circumstellar matter and/or the molecular cloud is vastly different at different wavelengths; extinction affects UV and optical wavelengths much more strongly than IR. For young associations, many, but not all, legitimate members are identifiable using just one or a few survey methods.  There are advantages and disadvantages to studying nearby young associations. Identification of young members in an association is made easier if the objects are located at smaller heliocentric distances and therefore brighter on the whole.  This is especially the case for low-mass members with very low luminosities; a complete census of these stars and brown dwarfs can be made only for very nearby star-forming regions. On the other hand, nearby star-forming molecular clouds cover large areas of sky and therefore require large investments of observing time. At just 137 pc (with a depth of $\\sim$ 20 pc; Torres \\etal\\ 2007, 2009), the Taurus star-forming region is one of the closest large cloud complexes with hundreds of low-mass stars, ongoing star formation, and objects ranging in age up to $\\sim$5 Myr.  Studies of Taurus objects have significantly influenced our basic understanding of the star-formation process for decades (e.g., Herbig \\& Rao 1972, Kenyon \\etal\\ 2008). However, the Taurus Molecular Cloud is close enough that it subtends more than 100 square degrees of sky; surveying all or even most of it is difficult within typical telescope time allocations.  Infrared surveys led to the discovery that some stars have infrared excesses, interpreted as circumstellar matter (e.g., Aumann \\etal\\ 1984, Beichman \\etal\\ 1986).  Most if not all low-mass stars form with circumstellar accretion disks, resulting in IR excesses for as long as the dusty circumstellar material survives (e.g., Hernandez \\etal\\ 2007). By using IR to survey a star-forming region, the stars with IR excesses are relatively easily distinguished from stars without such excesses. The Spitzer Space Telescope (Werner \\etal\\ 2004) provides an excellent platform for surveying star-forming regions in the mid-IR and far-IR, enabling stars with IR excesses to be identified; the member stars which do not have IR excesses must be recovered using different techniques such as the ones listed above.  Because Spitzer relatively efficiently maps large regions of sky, it is a particularly useful tool for surveying large star forming regions.  We have conducted a large multi-wavelength imaging and spectroscopic survey of the Taurus Molecular Cloud (TMC) in order to test if our inventory of Taurus members with infrared excesses is, in fact, complete.  The Spitzer imaging component is referred to as the Taurus Spitzer Survey, and is described by Padgett \\etal\\ (2008; hereafter P08) and Padgett \\etal\\ (2009; hereafter P09).   It covers $\\sim$44 square degrees from 3.6 to 160 \\mum\\ and is a Spitzer Legacy Project, so enhanced data products have been delivered back to the Spitzer Science Center (SSC), including the catalogs on which this present paper is based.  In addition to the Spitzer component, there are four other major components to our Taurus survey.  XMM-Newton was used by the XMM-Newton Extended Survey of the Taurus Molecular Cloud (XEST) program (e.g., G\\\"udel \\etal\\ 2007 and references therein), which mapped $\\sim$5 square degrees, most of which was also mapped by the Spitzer observations; the XEST data include X-ray imaging but also include ultraviolet data from the XMM-Newton Optical Monitor (Audard \\etal\\ 2007). XEST was deliberately pointed towards aggregates of previously identified Taurus members. In the optical, the Canada-France-Hawaii Telescope (CFHT) survey (Monin \\etal\\ in preparation; G\\\"udel, Padgett, \\& Dougados 2007) mapped $\\sim$28 square degrees (all of which are encompassed by the Spitzer area), and the Sloan Digital Sky Survey (SDSS) (Finkbeiner \\etal\\ 2004; Padmanabhan \\etal\\ 2008) mapped $\\sim$48 square degrees in two perpendicular strips, about half of which overlaps the Spitzer area. Finally, the Five College Radio Astronomy Observatory (FCRAO) millimeter wavelength survey (Goldsmith \\etal\\ 2008) mapped $\\sim$100 square degrees in the CO(1-0) line, covering the Spitzer survey area entirely.  The relative coverages of these surveys is shown in P09. Our extended collaboration has already begun to use this rich dataset to search for new members of Taurus; Scelsi \\etal\\ (2007,2008) identified new candidate members using the XEST data, and Guieu \\etal\\ (2006, 2007) identified new brown dwarf members using the CFHT data to study their disk properties using the Spitzer data. Ongoing investigations include searches for members via emission line spectra (Knapp \\etal, in prep), Herbig-Haro (HH) objects (Stapelfeldt \\etal\\ in prep), and transition disks (McCabe \\etal\\ in prep).  In this paper, we select new candidate Taurus members with infrared excesses using Spitzer and Two-Micron All-Sky Survey (2MASS; Skrutskie \\etal\\ 2006) data.  We construct color-magnitude and color-color diagrams for point sources, then use the locations of previously-identified young stars in these diagrams to select new candidate members with infrared excesses.  To discard obvious extragalactic sources, we examine the source  morphology in all available bands of the multiwavelength Taurus survey. We construct spectral energy distributions (SEDs) from the photometry over all available bands, again discarding objects we believe to be galaxies. Follow-up spectroscopy has been obtained to assess whether or not our new candidate Taurus objects are likely Taurus members.  We also present Spitzer flux densities for the 215 previously-identified members found in the region covered by our Spitzer survey.  Note that (a) the 215 previously-identified objects include those members without infrared excesses identified via other mechanisms; (b) the 215 previously-identified objects are those covered by our Spitzer map -- there are other legitimate Taurus members outside the region we observed, such as in L1551; (c) our new candidate member list is necessarily just those with IR excesses and exclusively within the regions covered by our Spitzer observations. Some objects are resolved in one or more of the Spitzer images, and extended source photometry may be a better representation of the complete flux from the object; many of the extended sources are discussed individually in  other papers (e.g., Tobin \\etal\\ 2008, Stapelfeldt \\etal\\ in prep)   The observations, data reduction, and ancillary data are described in \\S\\ref{sec:obs}.  Section~\\ref{sec:pickysos} describes our young stellar object (YSO) selection process which is based on the colors of previously-identified Taurus members, also presented here.   We describe 34 objects that we have, thus far, identified as new members of Taurus (plus 3 probable new members and 10 possible new members) in \\S\\ref{sec:discussion} and discuss the properties of the new objects in conjunction with (and comparison to) the previously-identified members of Taurus.  Finally, we summarize our main points in \\S\\ref{sec:concl}.  The Appendix contains spectral energy distributions and discussion of some specific objects. ",
        "conclusions": "\\label{sec:concl}  We have presented here Spitzer flux densities for 215 previously-identified members of the Taurus Molecular Cloud young stellar object population.  We constructed Spitzer color-color and color-magnitude diagrams, investigated where the previously identified Taurus members were located, and then used those diagrams to select additional candidate Taurus objects out of a catalog of $\\sim$700,000 objects observed with Spitzer over 44 square degrees.  We used a wealth of supporting data (including ground-based optical imaging) to winnow that list down to 148 candidate new Taurus members.  We obtained follow-up optical spectroscopy for about half the sample, thus far finding 34 new members, 3 probable new members, and 10 possible new members, a potential increase of 15-20\\% of young Taurus members.  Most of the new members are located in close (projected) proximity to the previously-identified Taurus members; most of them are Class II M stars.  In addition to the new members in our sample, there are 60 stars needing additional follow-up observations, and 33 objects pending any follow-up observations at all, so more Taurus members may yet be confirmed out of our list of candidate members.  We also found a background Be star, a new planetary nebula, a new carbon star, many background giants, and 100s of galaxies, just 7 of which made it into our final list of 148 YSO candidates.  As part of this project's classification as a Spitzer Legacy Project, enhanced data products have been delivered back to the Spitzer Science Center (SSC), including the catalogs on which this present paper is based.  This study has demonstrated the unique power of Spitzer to efficiently survey large areas of the sky and provide new information on membership and YSO properties even in nearby star forming regions such as Taurus, which has been extensively studied for decades. Even after many decades of study, our knowledge of membership in Taurus is still incomplete, but definitely improving, thanks to Spitzer.   \\appendix"
    },
    "0911/0911.4560_arXiv.txt": {
        "abstract": "Both analytical and numerical works show that magnetic reconnection must occur in hot accretion flows. This process will effectively heat and accelerate electrons. In this paper we use the numerical hybrid simulation of magnetic reconnection plus test-electron method to investigate the electron acceleration and heating due to magnetic reconnection in hot accretion flows. We consider fiducial values of density, temperature, and magnetic parameter $\\beta_e$ (defined as the ratio of the electron pressure to the magnetic pressure) of the accretion flow as $n_{0} \\sim 10^{6} {\\rm cm^{-3}}$, $T_{e}^0\\sim 2\\times 10^9 {\\rm K}$, and $\\beta_e=1$. We find that electrons are heated to a higher temperature $T_{e}=5\\times 10^9$K, and a fraction $\\eta\\sim 8\\%$ of electrons are accelerated into a broken power-law distribution, $dN(\\gamma)\\propto \\gamma^{-p}$, with $p\\approx 1.5$ and $4$ below and above $\\sim 1$ MeV, respectively. We also investigate the effect of varying $\\beta$ and $n_0$. We find that when $\\beta_e$ is smaller or $n_0$ is larger, i.e, the magnetic field is stronger, $T_e$, $\\eta$, and $p$ all become larger. ",
        "introduction": "Magnetic reconnection serves as a highly efficient engine which converts magnetic energy into plasma thermal and kinetic energy. The first theoretical model for magnetic reconnection was proposed by Sweet and Parker \\citep{Parker 1958}. Ever since then, theoretical studies along this line have been reported in various papers including analytical theories and numerical simulations \\citep{Parker 1963, Drake 2005, Litvinenko 1996}.  Magnetic reconnection has been widely used to explain explosive phenomena in space and laboratory plasmas such as solar flares \\citep{Somov 1997}, the heating of solar corona \\citep{Cargill 1997}, and substorms in the Earth's magnetosphere \\citep{Bieber 1982, Baker 2002}. Magnetic reconnection also plays an important role in high-energy astrophysical environments such as the magnetized loop of the Galactic center \\citep{Heyvaerts 1988}, jets in active galactic nuclei \\citep{Schopper 1998, Larrabee 2003, Lyutikov 2003}, and $\\gamma$ ray bursts \\citep{Michel 1994}. For example, by assuming that particle acceleration is due to direct current (DC) electric field and is balanced by synchrotron radiative losses, Lyutikov (2003) found that the maximum energy of electrons is $\\sim 100$ TeV. Larrabee calculated the self-consistent current density distribution in the regime of collisionless reconnection at an X-type magnetic neutral point in relativistic electron-positron plasma and studied the acceleration of non-relativistic and relativistic electron populations in three-dimensional tearing configurations.  But so far there is very few work on the magnetic reconnection in hot accretion flows. In accretion flows, the gravitational energy is converted into turbulent and magnetic energy due to the magnetorotational instability \\citep[MRI;][]{Balbus 1991}. Both analytical works \\citep[e.g.,][]{Quataert 1999, Goodman 2008, Uzdensky 2008} and three-dimensional magnetohydrodynamic simulations of black hole accretion flow \\citep{Hawley 2002, Machida 2003, Hirose 2004} have shown that magnetic reconnection is unavoidable and plays an important role in converting the magnetic energy into the thermal energy of particles. In fact, turbulent dissipation and magnetic reconnection are believed to be the two main mechanisms of heating and accelerating particles in hot accretion flows. In the case of large $\\beta$ ($\\beta$ is defined as the ratio of the gas pressure to the magnetic pressure), the magnetic reconnection is even likely to the dominant mechanism of electrons heating and acceleration \\citep[][]{Quataert 1999, Bisnovatyi-Kogan 1998}. On the observational side, flares are often observed in black hole systems. One example is the supermassive black hole in our Galactic center, Sgr A*, where infrared and X-ray flares occur almost every day \\citep[e.g.,][]{Baganoff 2001, Hornstein 2007, Dodds-Eden 2009}. To explain the origin of these flares, we often require that some electrons are transiently accelerated into a relativistic power-law distribution by magnetic reconnection within the accretion flow \\citep[e.g.,][]{Baganoff 2001, Yuan 2003, Yuan 2004, Dodds-Eden 2009, Eckart 2009} or in the corona above the accretion flow (e.g., Yuan et al. 2009), and their synchrotron or synchrotron-inverse-Compton emission will produce the flares.  Regarding the study of the electron acceleration by reconnection, recently many works have been done on the relativistic magnetic reconnection with the approach of Particle-in-Cell (PIC) simulation. \\cite{Zenitani 2007} presented the development of relativistic magnetic reconnection, whose outflow speed is on the order of the light speed. It was demonstrated that particles were strongly accelerated in and around the reconnection region and that most of the magnetic energy was converted into a ``non-thermal'' part of plasma kinetic energy. Magnetic reconnection during collisionless, stressed, X-point collapse was studied using kinetic, 2.5D, fully relativistic PIC numerical code in \\cite{Tsiklauri 2007} and they also found high energy electrons. PIC electromagnetic relativistic code was also used in \\cite{Karlicky 2008} to study the acceleration of electrons and positrons. They considered a model with two current sheets and periodic boundary conditions. The electrons and positrons are very effectively accelerated during the tearing and coalescence processes of the reconnection. All these work investigate the circumstance of the Sun.  In this paper we use the numerical hybrid simulation of magnetic reconnection plus test-electron method to investigate the electron heating and acceleration by magnetic reconnection in hot accretion flows. We would like to note that the results should also be applied to corona of accretion flow and jet, if the properties of the plasma in those cases are similar to what we will investigate in the present paper. Self-consistent electromagnetic fields are obtained from the hybrid code, in which ions are treated as discrete particles and electrons are treated as massless fluid. Thereafter, test electrons are placed into the fields to study their acceleration. The input parameters include density, temperature and magnetic field of the accretion flow. The values of density and temperature of astrophysical accretion flows can differ by several orders of magnitude among various objects. Here we take the accretion flow in Sgr A* as reference \\citep[see][for details]{Yuan 2003}, but we also investigate the effect of varying parameters. For the strength of the magnetic field, we consider mainly $\\beta_{e}=1$. Here $\\beta_{e}$ is defined as the ratio of the electron pressure to magnetic pressure, $\\beta_{e}\\equiv P_{e}/P_{\\rm mag}$. Note that what we usually use in hot accretion flows is the ratio of the gas pressure and the magnetic pressure, $\\beta\\equiv P_{gas}/P_{\\rm mag}$. The hot accretion flow is believed to be two-temperature, with $T_{i} \\gg T_{e}$. For a reasonable value of $T_{i}=10T_{e}$ \\citep[ref.][]{Yuan 2003}, the above $\\beta_{e}$ corresponds to the usual $\\beta =10$. This is the typical value in the hot accretion flow as shown by MHD numerical simulations \\citep[e.g., ][]{Hirose 2004}. Given that the value of $\\beta$ is very inhomogenous in the accretion flow \\citep{Hirose 2004}, we also consider $\\beta_{e}=10, 0.1$, and $0.01$.  In Section 2, we use the numerical hybrid simulation of magnetic reconnection to get the structure of self-consistent electric and magnetic fields. In Section 3 we use the obtained electromagnetic fields as the background of test electrons to investigate the electron acceleration. Finally we summarize our results and discuss the application in interpreting the flares in Sgr A* in Section 4. ",
        "conclusions": "In this paper we use hybrid simulation code of magnetic reconnection plus test-electron method to study the electrons heating and acceleration by magnetic reconnection process in hot accretion flows. The self-consistent electromagnetic fields are obtained from the hybrid simulation and they are then used in the test-electron calculation. The fiducial parameters of the background accretion flows we adopt are density $n_{e}=10^6 {\\rm cm}^{-3}$ and electron temperature $T_{e}^0=2\\times 10^{9.3}$K (ions temperature $T_{i}=10 T_{e}^0$). These values are taken from the accretion flow in Sgr A*, the supermassive black hole in our Galactic center \\citep[ref.][]{Yuan 2003}. For the strength of the magnetic field in the accretion flow, we set $\\beta_{e} (\\equiv P_{e}/P_{\\rm mag})=1.0$. This is the typical value according to the MHD numerical simulation of hot accretion flow. We find that the electrons are heated and accelerated by the reconnection. The new distribution can be fitted by a Maxwellian distribution plus a broken power-law. The temperature is increased to $10^{9.65}$K and the power-law is described by $dN(E)\\propto E^{-1.47}$ between 1 and 10 MeV and $dN(E)\\propto E^{-4}$ between 10 MeV and 40 MeV. The fraction of electrons with energy above 1 MeV is $\\sim 8\\%$. Given that the magnetic field and density in accretion flow are inhomogeneous, we also consider $\\beta_{e}=10, 0.1$, and $0.01$ and $n_{e}=10^7 {\\rm cm}^{-3}$. We find that the results are qualitatively similar, namely the distribution of electrons can be fitted by a Maxwellian one with a higher ``heated'' temperature plus a broken power-law. When the magnetic field is stronger ($\\beta_{e}$ is smaller or $n_{e}$ is larger), the heating and acceleration become more efficient. The ``heated'' temperature is higher and more electrons are accelerated. In addition, the ``low-energy'' part of the broken power-law becomes softer (refer to Table 1).  In our work we assume ideal MHD for the hybrid simulation, thus the Joule heating to the accretion flow is neglected. Obviously the electrons will be heated to a higher temperature when this effect is considered."
    },
    "0911/0911.1459_arXiv.txt": {
        "abstract": "Almost 80 years have passed since Trumpler's analysis of the Galactic open cluster system laid one of the main foundations for understanding the nature and structure of the Milky Way. Since then, the open cluster system has been recognised as a key source of information for addressing a wide range of questions about the structure and evolution of our Galaxy. Over the last decade, surveys and individual observations from the ground and space have led to an explosion of astrometric, kinematic and multiwavelength photometric and spectroscopic open cluster data. In addition, a growing fraction of these data is often time-resolved. Together with increasing computing power and developments in classification techniques, the open cluster system reveals an increasingly clearer and more complete picture of our Galaxy. In this contribution, I review the observational properties of the Milky Way's open cluster system. I discuss what they can and cannot teach us now and in the near future about several topics such as the Galaxy's spiral structure and dynamics, chemical evolution, large-scale star formation, stellar populations and more. ",
        "introduction": "In 1930, Robert Trumpler \\nocite{Trumpler1930} published his seminal paper in which he proves the existence of the interstellar medium (ISM).  In that paper, Trumpler showed how distances to open clusters (OCs) derived from the apparent magnitudes of cluster stars of known spectral types are systematically greater than those derived from their apparent sizes, and how the effect increases with distance.  He correctly attributed this phenomenon to the presence of an intervening medium. This was the first paper using OCs as tracers of Galactic structure. ",
        "conclusions": ""
    },
    "0911/0911.2493_arXiv.txt": {
        "abstract": "We describe an exact, flexible, and computationally efficient algorithm for a joint estimation of the large-scale structure and its power-spectrum, building on a Gibbs sampling framework and present its implementation \\textsc{ARES} (Algorithm for REconstruction and Sampling). \\textsc{ARES} is designed to reconstruct the 3D power-spectrum together with the underlying dark matter density field in a Bayesian framework, under the reasonable assumption that the long wavelength Fourier components are Gaussian distributed. As a result \\textsc{ARES} does not only provide a single estimate but samples from the joint posterior of the power-spectrum and density field conditional on a set of observations. This enables us to calculate any desired statistical summary, in particular we are able to provide joint uncertainty estimates. We apply our method to mock catalogs, with highly structured observational masks and selection functions, in order to demonstrate its ability to reconstruct the power-spectrum from real data sets, while fully accounting for any mask induced mode coupling. ",
        "introduction": "\\nextext{Throughout cosmic history a wealth of information on the origin and evolution of our Universe has been imprinted to the large scale structure via the gravitational amplification of primordial density perturbations. Harvesting this information from probes of the large scale structure, such as large galaxy surveys, therefore is an important scientific task to further our knowledge and to establish a conclusive cosmological picture. In recent years great advances have been made, both in retrieving huge amounts of data and increasing sensitivity in galaxy redshift surveys. Especially the recent galaxy surveys, the 2dF Galaxy Redshift Survey \\citep[][]{COLLESS2001} and the Sloan Digital Sky Survey \\citep[][]{SDSS7} provide sufficient redshifts to probe the 3D galaxy distribution on large scales. In particular, the two point statistics of the matter distribution contains valuable information to test standard inflation and cosmological models, which describe the origin and evolution of all observed structure in the Universe. Measuring the power-spectrum from galaxy observations therefore has always attracted great interest. Precise determination of the overall shape of the power-spectrum can for instance place important constraints on neutrino masses, help to identify the primordial power-spectrum, and break degeneracies for cosmological parameter estimation from CMB data \\citep[e.g.][]{hu-98,wmap-spergel,HANNESTAD2003,Efstathiou_2002,PERCIVAL2002,wmap-spergel,VERDE2003}. In addition, several characteristic length scales have been imprinted to the matter distribution throughout cosmic history, which can serve as new standard rulers to measure the Universe. A prominent example of these length scales is the sound horizon, which yields oscillatory features in the power-spectrum, the so called baryon \\nextext{acoustic} oscillations (BAO) \\citep[e.g.][]{SILK1968,PEEBLES1970,SUNYAEV1970}. Since the physics governing these oscillatory features is well understood, precise measurements of the BAO will allow us to establish a new, precise standard ruler to measure the Universe through the distance redshift relation \\citep[][]{BLAKE2003,SEO2003}. Precision analysis of large scale structure data therefore is a crucial step in developing a conclusive cosmological theory.}  \\nextext{Unfortunately, contact between theory and observations cannot be made directly, since observational data is subject to a variety of systematic effects and statistical uncertainties. Such systematics and uncertainties arise either from the observational strategy or are due to intrinsic clustering behavior of the galaxy sample itself \\citep[][]{SANCHEZ2008}. Some of the most prominent uncertainties and systematics are:} \\begin{itemize} \\item survey geometry and selection effects \\item close pair incompleteness due to fiber collisions \\item galaxy biases \\item redshift space distortions \\end{itemize} \\nextext{The details of galaxy clustering, and how galaxies trace the underlying density field are in general very complicated. The bias between galaxies and mass density is most likely non-linear and stochastical, so that the estimated galaxy spectrum is expected to differ from that of the mass \\citep[][]{DEKEL1999}. Even in the limit where a linear bias could be employed, it still differs for different classes of galaxies \\citep[see e.g.][]{COLE2005}. In addition, the apparent density field, obtained from redshift surveys, will generally be distorted along the line-of-sight due to the existence of peculiar velocities.}  \\nextext{ However, the main cause for the systematic uncertainties in large scale power-spectrum estimations is the treatment of the survey geometry \\citep[][]{TEGMARK1995,BALLINGER1995}. Due to the survey geometry the raw power-spectrum yields an expectation value for the power-spectrum, which is the true cosmic power-spectrum convolved with the survey mask \\citep[][]{COLE2005}. This convolution causes an overall distortion of the power-spectrum shape, and drastically decreases the visibility of the baryonic features.}  \\nextext{ The problems, mentioned above, have been discussed extensively in literature, and many different approaches to power-spectrum analysis have been proposed. Some of the main techniques to recover the power-spectrum from galaxy surveys are Fourier transform based, such as the optimal weighting scheme, which assigns a weight to the galaxy fluctuation field, in order to reduce the error in the estimated power \\citep[see e.g.][]{FELDMAN1994,TEGMARK1995,HAMILTON1997A,YAMAMOTO2003,PERCIVAL2004}. Alternative methods rely on Karhunen-Lo\\`{e}ve decompositions \\citep[][]{TEGMARK1997,TEGMARK_2004,POPE2004} or decompositions into spherical harmonics, which is especially suited to address the redshift space distortions problematic \\citep[][]{FISHER1994,HEAVENS1995,TADROS1999,PERCIVAL2004,PERCIVAL2005}. In addition, to these deconvolution methods there exists a variety of likelihood methods to estimate the real space power-spectrum \\citep[][]{BALLINGER1995,HAMILTON1997A,HAMILTON1997B,TADROS1999,PERCIVAL2005}. In order to not just provide the maximum likelihood estimate but also conditional errors, \\cite{PERCIVAL2005} proposed a Markov Chain Monte Carlo method to map out the likelihood surface. }  \\nextext{ Nevertheless, as the precision of large scale structure experiments has improved, the requirement on the control and characterization of systematic effects, as discussed above, also steadily increases. It is of critical importance to propagate properly the uncertainties caused by these effects through to the matter power-spectrum and cosmological parameters estimates, in order to not underestimate the final uncertainties and thereby draw incorrect conclusions on the cosmological model. }  \\nextext{We therefore felt inspired to develope a new Bayesian approach to extract information on the two point statistics from a given large scale structure dataset. We prefer Bayesian methods to conventional likelihood methods, as the yield more general and profound statements about measurements. In example, conventional likelihood methods can only answer questions of the inner form like :\" Given the true value \\(s\\) of a signal, what is the probability distribution of the measured values \\(d\\)?\" A Bayesian method, on the other hand, answers questions of the type :\"Given the observations \\(d\\), what is the probability distribution of the true underlying signal \\(s\\)?\" For this reason, Bayesian statistics answers the underlying question to every measurement problem, of how to estimate the true value of the signal from observations, while conventional likelihood methods do not \\citep[][]{MICHEL1999}.} \\nextext{Since the result of any Bayesian method is a complete probability distribution they permit fully global analyses, taking into account all systematic effects and statistical uncertainties. In particular, here, we aim at evaluating the power-spectrum posterior distribution \\(\\mathcal{P}\\left(\\{P(k_i)\\}|\\{d_i\\}\\right)\\), with \\(P(k_i)\\) being the power-spectrum coefficients of the \\(k_i\\)th mode and \\(d_i=d(\\vec{x}_i)\\) is an observation at position \\(\\vec{x}_i\\) in three dimensional configuration space.} This probability distribution would then contain all information on the two point statistics supported by the data. In order to explore this posterior distribution we employ a Gibbs sampling method, previously applied to CMB data analysis \\citep[see e.g.][]{WANDELT2004,2004ApJS..155..227E,JEWELL2004}.  \\nextext{ Since direct sampling from \\(\\mathcal{P}\\left(\\{P(k_i)\\}|\\{d_i\\}\\right)\\) is impossible or at least difficult, they propose instead to draw samples from the full joint posterior distribution \\(\\mathcal{P}\\left(\\{P(k_i)\\},\\{s_i\\}|\\{d_i\\}\\right)\\) of the power-spectrum coefficients \\(P(k_i)\\) and the 3D matter density contrast amplitudes \\(s_i\\) conditional on a given set of data points \\(\\{d_i\\}\\). This is achieved by iteratively drawing density samples from a Wiener-posterior distribution and power-spectrum samples via an efficient Gibbs sampling scheme (see figure \\ref{fig:flowchart} for an illustration). Here, artificial mode coupling, as introduced by survey geometry and selection function, is resolved by solving the Wiener-filtering equation, which naturally regularizes inversions of the observational response operator in unobserved regions.} In this fashion\\nextext{,} we obtain a set of Monte Carlo samples from the joint posterior, which allows us to compute any desired property of the joint posterior density, with the accuracy only limited by the sample size. In particular\\nextext{,} we obtain \\nextext{the power spectrum posterior} \\(\\mathcal{P}\\left(\\{P(k_i)\\}|\\{d_i\\}\\right)\\) by simply marginalizing \\nextext{the joint posterior} \\(\\mathcal{P}\\left(\\{P(k_i)\\},\\{s_i\\}|\\{d_i\\}\\right)\\) over the auxiliary density amplitudes \\(s_i\\), which is trivially achieved by ignoring the \\(s_i\\) samples.  The Gibbs sampler also offers unique capabilities for propagating systematic uncertainties end-to-end. Any effect, for which we can define a sampling algorithm, either jointly with or conditionally on other quantities, can be propagated seamlessly through to the final posterior.  It is worth noting, that our method differs from traditional methods of analyzing galaxy surveys in a fundamental aspect. Traditional methods consider the analysis task as a set of steps, each of which arrives at intermediate outputs which are then fed as inputs to the next step in the pipeline. Our approach is a truly global analysis, in the sense that the statistics of all science products are computed jointly, respecting and exploiting the full statistical dependence structure between various components.  In this paper we present \\textsc{ARES} (Algorithm for REconstruction and Sampling), a computer algorithm to perform a full Bayesian data analysis of 3D redshift surveys. In section \\ref{LSS_SAMPLER} we give an introduction to the general idea of the large scale structure Gibbs sampling approach, followed by section \\ref{MAP_SAMPLER} and \\ref{PS_SAMPLER}, where we describe and derive in detail the necessary ingredients to sample the 3D density distribution and the power-spectrum respectively. The choice of the prior and the relevance for the cosmic variance are discussed in section \\ref{Prior_and_Variance}. Details concerning the numerical implementation are discussed in section \\ref{NUMERICAL_IMPLEMENTATION}. We then test \\textsc{ARES} thouroughly in section \\ref{Gaussian_Test_Cases}, particularly focussing on the treatment of survey masks and selection functions. In section \\ref{OPERATIONS_GIBBS_SAMPLES} we demonstrate the running median filter, and use it as an example to demonstrate how uncertainties can be propagated to all inferences based on the set of Gibbs samples. Finally we conclude in section \\ref{Conclusion}, by discussing the results of the method and giving an outlook for future extensions and application of our method. ",
        "conclusions": "\\label{Conclusion} In this work, we presented \\textsc{ARES}, a new Bayesian computer algorithm, designed to extract the full information on the two point statistics from any given probe of the three dimensional large scale structure.  The scientific aim of this algorithm is to provide an estimate of the power-spectrum posterior \\(\\mathcal{P}\\left(\\{P(k_i)\\}|\\{d_i\\}\\right)\\), conditional on a data set, while accounting and correcting for all possible sources of uncertainties, such as survey geometry, selection effects and biases. This is achieved by exploring the power-spectrum posterior \\(\\mathcal{P}\\left(\\{P(k_i)\\}|\\{d_i\\}\\right)\\) via a Gibbs sampling approach.  While direct sampling from the power-spectrum posterior is not possible, it is possible to draw samples from the full joint posterior distribution \\(\\mathcal{P}\\left(\\{P(k_i)\\},\\{s_i\\}|\\{d_i\\}\\right)\\) of the power-spectrum coefficients \\(P(k_i)\\) and the three dimensional matter density contrast amplitudes \\(s_i\\) conditional on a given set of data points \\(\\{d_i\\}\\).  The entire Gibbs sampling algorithm therefore consists of two basic sampling steps, in which first a full three dimensional Wiener reconstruction algorithm is applied to the data and then a power-spectrum is drawn from the inverse gamma distribution. In this fashion we obtain a set of power-spectrum and signal amplitude samples, which provide a full representation of the full joint posterior distribution \\(\\mathcal{P}\\left(\\{P(k_i)\\},\\{s_i\\}|\\{d_i\\}\\right)\\). The scientific output of this Bayesian method therefore is not a single estimate but a complete probability distribution, enabling us to calculate any desired statistical summary such as the mean, mode or variance.                      \\begin{figure*} \\centering{\\resizebox{1.\\hsize}{!}{\\includegraphics{./figures/picture13}}} \\caption{Volume rendering of the ensemble variance (upper panels) and the ensemble mean (lower panels) obtained from the mock galaxy catalog analysis.} \\label{fig:DELUCIA_VARIANCE} \\end{figure*}            \\begin{figure*} \\centering{\\resizebox{1.\\hsize}{!}{\\includegraphics{./figures/picture14}}} \\caption{Running median estimates of the power-spectra (upper panels) and the according wiggle functions (lower panels) for the set of Gibbs samples with the Jeffrey's prior (left panels), and the inverse gamma prior (right panels). The black lines represent the ensemble mean of the sample set and the light gray and dark gray shaded regions denote the one and two sigma confidence regions respectively. Additionally we show the according input power-spectra. The blue line shows the cosmological power-spectrum from which the matter field realization was drawn, and the red line is the power-spectrum of this specific matter field realization.} \\label{fig:RUNNING_MEDIAN} \\end{figure*} We also demonstrated, that given a set of Gibbs samples, it is possible to provide an analytic approximation to the power-spectrum posterior \\(\\mathcal{P}\\left(\\{P(k_i)\\}|\\{d_i\\}\\right)\\). This Blackwell-Rao estimator has an analytically appealing form enabling us to calculate any desired moment of the probability distribution in a simple analytic way.  In addition, since the full joint probability distribution is available, it is easy to carefully propagate all uncertainties through to the result of further post-processing analysis steps, such as parameter estimation.  In this work, we focused on thoroughly testing the performance and behavior of our method by applying it to simulated data with controlled properties. These tests were designed to highlight the problematic of survey geometry and selection effects, for the two cases of Gaussian random fields and a mock galaxy catalog based on the Millennium run.  One of the main goals of these tests was to build up intuition on the phenomenological behavior of the Gibbs sampling algorithm, estimating particularly issues, such as the correlation length of the Gibbs chain, burn-in and convergence times. The result of these tests is of special relevance, as it shows how long the Gibbs sampling chain has to run in order to produce a sufficient amount of independent samples.  Through these experiments we found that the longest correlation lengths are dominated by the poorly constrained Nyquist modes of the box, which can be easily alleviated by imposing some prior knowledge on these modes. In doing so we found that the maximal correlation length for the Gibbs chain was on the order of hundred Gibbs samples. Thus, creating a large number of independent samples in a full scale data analysis is numerically very well feasible.  However, the most important result of these tests is, that our method is able to correct for artificial mode coupling due to the survey geometry and selection effects. This was tested by examining the correlation structure of the Gibbs samples, which showed that the maximal residual correlation can be reduced to less than \\(10 \\%\\), demonstrating that this method correctly accounts for geometry effects.  The application of \\textsc{ARES} to a galaxy mock catalog, based on the Millennium run, demonstrated the ability of our method to capture the characteristics of the fully nonlinearly evolved matter field. This is owed to the fact, that the Wiener filter is a linear operation on the data, and as such does not destroy the intrinsic statistical characteristics of the data set.  Nevertheless, a full real data analysis of existing redshift surveys requires the treatment of additional systematic effects such as scale or luminosity dependent biases or redshift space distortion corrections, which we defer to future works. However, the Bayesian framework, as presented here, can take all these effects naturally into account and treats them statistically fully consistent. Beside the possibility to include various kinds of uncertainties, the Gibbs sampling approach also allows for a natural joint analysis of different data sets, taking into account the systematics of each individual data set.  In summary, we showed that \\textsc{ARES} is a highly flexible and adaptive machinery for large scale structure analysis, which is able to account for a large variety of systematic effects and uncertainties. For this reason, \\textsc{ARES} has the potential to contribute to the era of precision cosmology."
    },
    "0911/0911.5566_arXiv.txt": {
        "abstract": "We use high-resolution cosmological hydrodynamical AMR simulations to predict the characteristics of \\lya emission from the cold gas streams that fed galaxies in massive haloes at high redshift. The \\lya luminosity in our simulations is powered by the release of gravitational energy as gas flows from the intergalactic medium into the halo potential wells. The UV background contributes only $<20\\%$ to the gas heating. The \\lya emissivity is due primarily to electron-impact excitation cooling radiation in gas $\\sim 2\\times 10^4$~K. We calculate the \\lya emissivities assuming collisional ionisation equilibrium (CIE) at all gas temperatures. The simulated streams are self-shielded against the UV background, so photoionisation and recombination contribute negligibly to the \\lya line formation. We produce theoretical maps of the \\lya surface brightnesses, assuming that $\\sim 85\\%$ of the \\lya photons are directly observable. We do not consider transfer of the \\lya radiation, nor do we include the possible effects of internal sources of photoionisation such as star-forming regions. Dust absorption is expected to obscure a small fraction of the luminosity in the streams. We find that typical haloes of mass $\\Mv\\sim10^{12-13}\\msun$ at $z\\sim 3$ emit as \\lya blobs (LABs) with luminosities $10^{43-44}\\ergs$. Most of the \\lya comes from the extended ($50-100\\kpc$) narrow, partly clumpy, inflowing, cold streams of $(1-5)\\times 10^4$K that feed the growing galaxies. The predicted LAB morphology is therefore irregular, with dense clumps and elongated extensions. The integrated area contained within surface-brightness isophotes of $2 \\times 10^{-18}\\ergsc$ is $\\sim 2-100\\, {\\rm arcsec}^2$, consistent with observations. The linewidth is expected to range from $10^2$ to more than $10^3 \\kms$ with a large variance. The typical \\lya surface brightness profile is $\\propto r^{-1.2}$ where $r$ is the distance from the halo centre. Our simulated LABs are similar in luminosity, morphology and extent to the observed LABs, with distinct kinematic features. The predicted \\lya luminosity function is consistent with observations, and the predicted areas and linewidths roughly recover the observed scaling relations. This mechanism for producing LABs appears inevitable in many high-$z$ galaxies, though it may work in parallel with other mechanisms. Some of the LABs may thus be regarded as direct detections of the cold streams that drove galaxy evolution at high $z$. ",
        "introduction": "\\label{sec:intro} Hundreds of Lyman-alpha blobs (LAB) have been detected so far in the redshift range $z = 2 - 6.5$, mostly near $z \\sim 3$ \\citep{steidel, matsuda, saito, ouchi, yang}. Their luminosities range from below $10^{43}$ to above $10^{44}\\ergs$, and they extend on the sky to $30-50\\kpc$ and more. The two main open questions are (a) the origin of the extended cold and relatively dense gas capable of emitting \\lya, and (b) the continuous energy source for exciting the gas to emit \\lya.  The emitting hydrogen gas should be at a temperature $T \\gsim 10^4$K, relatively dense and span a much larger area than covered by the stellar component of galaxies. It may arise from outflows or inflows. The energy sources discussed in the literature include photoionisation by obscured AGNs, early starbursts or extended X-ray emission \\citep{haimanb, jimenez, scharf}, as well as compression of ambient gas by superwinds to dense \\lya emitting shells \\citep{mori}, and star formation triggered by relativistic jets from AGNs \\citep{rees}.  Many of the bright LABs are found in the vicinity of massive, star-forming galaxies \\citep{matsudab}. Multi-wavelength observations reveal that a fraction of the LABs are associated with sub-millimetre and infrared sources that indicate very high star-formation rates (SFR) in the range $10^2-10^3 \\sy$ \\citep{chapman, geacha, geachb} or with obscured active galactic nuclei (AGN) \\citep{basu, scarlata}. While stellar feedback and AGNs could in principle provide the energy source for the \\lya luminosity, many LABs are not associated with sources of this sort that are powerful enough to explain the observed \\lya luminosities \\citep{kim, sj}. This indicates that star formation and AGNs are not the sole drivers of LABs, and may not even be the dominant source for LABs.  Indeed, high-redshift galaxies exhibit a generic mechanism that simultaneously provides both the cold gas and the energy source. It is a direct result of the phenomenon robustly established by simulations and theoretical analysis, where high-z massive galaxies are continuously fed by narrow, cold, intense, partly clumpy, gaseous streams that penetrate through the shock-heated halo gas into the inner galaxy, grow a dense, unstable, turbulent disc with a bulge and trigger rapid star formation \\citep{bd03, keresa, db06, ocvirk, keresb, nature, peter, dsc, cd}. Massive clumpy star-forming disks observed at $z\\sim 2$ \\citep{genzel08, genel, foerster2} may have been formed via the smooth and steady accretion provided by cold flows, as opposed to merger events. The streaming of the gas into the dark-matter halo potential well is associated with transfer of gravitational energy to excitations of the hydrogen atoms followed by cooling emission of \\lya \\citep{haiman, fardal, db06, db08, ko08, furlanetto2, furlanetto}.  There were two earlier attempts to compute the \\lya cooling radiation from hydrodynamic SPH cosmological simulations. Based on their analysis, \\citet{fardal} pointed out the potential association of this \\lya cooling emission with the first observed LABs. These early simulations did not allow a proper resolution of the detailed structure of the cold streams as they penetrate through the hot medium. Their shortcomings included intrinsic limitations of the SPH technique, a limited force resolution of $7 \\hkpc$ (comoving), not allowing for radiative cooling below $10^4$K and neglecting the photoionisation by the UV background. \\citet{furlanetto} used SPH simulations with a somewhat higher resolution of $\\sim$1\\,kpc to make more detailed comparisons of the simulated and observed luminosity functions and size distributions of LABs powered by cold accretion. In their pessimistic model, assuming no emissivity from gas that is self-shielded from the ionising background radiation, they end up with low \\lya luminosity. They comment that cooling IGM gas may explain the observations only if one adopts an optimistic scenario, where the self-shielded gas is emitting \\lya at CIE. The latter scenario will also be adopted in this paper. Other sources of ionisation are ignored and radiative transfer is not applied.  \\citet[][hereafter DL09]{dijkstra} have recently worked out an analytic toy model for \\lya cooling radiation from cold streams, based on the general properties of the cold streams as reported from cosmological simulations. They conclude that the streams could in principle provide spatially extended \\lya sources with luminosities, line widths and abundances that are similar to those of observed LABs. They point out that the filamentary structure of cold flows may explain the wide range of observed LAB morphologies. They also highlight the fact that the most luminous cold flows are associated with massive haloes, which preferentially reside in dense large-scale surroundings, in agreement with the observed presence of bright LABs in dense environments. This model presents a successful feasibility test for the role of cold streams in powering the LAB emission, and it provides physical intuition into the way by which this process could be manifested. However, a comparison to our simulations indicates that this simplified model does not capture the detailed hydrodynamic properties of the cold streams. In particular, it significantly over-predicts the cold-gas density, under-predicts the gas temperature, and does not account for the partly clumpy nature of the streams and their characteristic radial distribution in the halo. As a result, the DL09 model underestimates the total \\lya luminosity by a factor of a few, and it therefore has to appeal to the excessive clustering of the LABs in order to boost up the predicted luminosity function for a match with the observations.  In the current paper, we calculate the \\lya emission from the cold gas in high-redshift galactic haloes using state-of-the-art hydrodynamical AMR cosmological simulations. With a maximum resolution better than 70\\,pc in our simulations we can quite accurately map the extended \\lya sources. This allows us to study their individual shapes, morphologies and kinematics. We measure quantities such as the distribution of surface brightness within each halo, the area covered with surface brightness above a given isophotal threshold, the predicted observed linewidth, the typical total \\lya luminosity per halo of a given mass, and the overall \\lya luminosity function of LABs. These predicted properties are compared with the observed LABs.  This paper is organised as follows. In \\se{feasibility} we work out a simple feasibility test where we use a simple toy model to estimate the expected gravitational heating power as a proxy for the \\lya luminosity. In \\se{sim} we introduce the two sets of cosmological simulations used. In \\se{lal} we explain our methodology of computing \\lya emissivity and luminosity as a function of halo mass. In \\se{lyagas} we identify the gas that contributes to the \\lya emission. In \\se{images} we apply our methodology to the simulations and provide predicted images of LABs. In \\se{L(M,z)} we determine the scaling of \\lya luminosity as a function of halo mass and redshift. In \\se{lumfun} we compare our predicted \\lya luminosity function with observational results. In \\se{obs} we measure the predicted surface-density profile, isophotal area and linewidth and compare to observations. In \\se{grav} we show that gravitational heating is the main source of energy driving the \\lya luminosity in our simulations, while the role of photoionisation is minor. In \\se{con} we discuss our analysis and results and draw our conclusions. ",
        "conclusions": "\\label{sec:con}  Hydrodynamical cosmological simulations robustly demonstrate that massive galaxies at high redshifts were fed by cold gas streams, inflowing into dark-matter haloes at high rates along the cosmic web \\citep{keresa, db06, nature}. In this paper we have shown that these streams should be observable as luminous \\lya sources, with elongated irregular structures stretching for distances of over 100 kpc. The release of gravitational potential energy by the instreaming gas as it falls into the halo potential well is the origin of the \\lya luminosity, which ranges from $10^{42}$ to $10^{44}\\ergs$ for haloes in the mass range $10^{11}-10^{13}\\msun$ at $z \\sim 3$. The predicted \\lya emission morphologies and luminosities make such streams likely candidates for the sources of observed high redshift \\lya blobs. Most of the gas in the cold streams is at temperatures in the range $(1-5)\\times 10^4$K, for which the \\lya emissivity is maximised. The hydrogen gas densities in the streams are in the range $0.01-0.1 \\cmpc$, and are higher in clumps that flow with the streams, leading to the high surface brightnesses.  Using state-of-the-art cosmological AMR simulations with 70 pc resolution and cooling to below $10^4$K, and applying a straightforward analysis of \\lya emissivity due to electron impact excitation, we produced maps of \\lya emission from simulated massive galaxies at $z \\sim 3$. We computed the average \\lya luminosity per given halo mass and the predicted luminosity function of these extended \\lya sources, with an uncertainty of $\\pm 0.4$ dex. The properties of the individual images resemble those of the observed LABs in terms of morphology and kinematics. The predicted luminosity function is close to the observed LAB luminosity function. The simulated LABs qualitatively reproduce the observed correlations between the global LAB properties of \\lya luminosity, isophotal area and linewidth.  The LAB properties can be understood using a very simple toy model that accounts for gravitational heating of the inflowing gas. In comparison, the more elaborate toy model of DL09 predicts a steeper power-law relation between luminosity and halo mass (their Fig.~3, with $f_{\\rm g}=0.2$). In the mass range $10^{12} - 10^{13}\\msun$, DL09 predict luminosities in the range $1.5 \\times 10^{42} - 2.5\\times 10^{44} \\ergs$. A comparison to our \\fig{lumiz246} indicates that they underestimate the luminosity at $\\Mv \\sim 10^{12}\\msun$ by a factor of a few. The DL09 toy model matches the observed luminosity function only after the predictions are boosted, trying to account for an overdense region of the universe. The association of the predicted LABs with massive haloes places them preferentially in overdense environments, as observed \\citep{steidel, matsuda, matsudab}, but this is already accounted for by the number density of haloes in a fair sample.  The association of observed LABs with star-forming LBGs \\citep{hayashino}, sub-millimetre galaxies \\citep{chapman, geacha} or active galactic nuclei \\citep{geachc, basu} is not surprising since star formation and AGN activity are triggered by the same process of streaming of cold gas into massive haloes. \\lya emission can be driven by any of these sources of energy. For example, in photoionised gas in star-forming regions the \\lya luminosity is $L_{\\rm L\\alpha} = 1.5 \\times 10^{42}\\, {\\rm SFR}\\, {\\rm erg}\\, {\\rm s}^{-1}$ where SFR is the star-formation rate (in $M_\\odot$ yr$^{-1}$) for continuous star formation for a Kroupa initial mass function \\citep{shp}. Thus the typical LAB luminosity of 10$^{43}$ erg s$^{-1}$ would require SFR $\\sim 10$ M$_\\odot$ yr$^{-1}$. However, most of this radiation is likely absorbed by dust, and confined to the central galaxy.  A limitation of our AMR simulations arises from the artificial pressure floor imposed in order to properly resolve the Jeans mass. This may have an effect on the temperature and density of the cold gas in the streams, with potential implications on the computed \\lya emission. Still, the AMR code is the best available tool for recovering the stream properties. With 70 pc resolution and proper cooling below $10^4$K, the {\\CD} simulations provide the most reliable description of the cold streams so far. The rather small correction that we had to apply to the luminosities extracted from the {\\MN} simulation indicates that the \\MN galaxies can be used to recover the mass and redshift dependence of the global stream properties and the scaling relations between them.  Another source of uncertainty has to do with the simplified way the ionisation by the UV background is handled in the simulations, and with the post-processing calculation of the ionisation state of the cold gas. In the \\CD simulations, the centres of the streams, where the gas density is higher than $0.1 \\cmmc$, were assumed to be self-shielded against the UV background, in agreement with our analytic estimates in \\se{lal}. The ionisation state of the gas, which is a key ingredient in evaluating the \\lya emissivity, was computed in post-processing assuming CIE. Cooling rates were computed for the given gas density, temperature, metallicity, and UV background based on \\textsc{Cloudy} (\\CD) or assuming ionisation equilibrium for H and He, including both collisional- and photo-ionisation (\\MN). Photoionisation is neglected since the streams are sufficiently thick to be self-shielded. An alternative calculation using \\textsc{Cloudy} yielded lower fractions of neutral hydrogen at $n \\leq 0.01 \\cmmc$, and total luminosities that are typically three times smaller than obtained using CIE. This calculation assumed that each cell, with a given hydrogen density $n$, is in the middle of a uniform slab of thickness 1 kpc. Based on our estimates of self-shielding, \\equ{shielding}, we adopt the CIE results as the more reliable approximation.  In our computations so far we did not consider the radiative transfer of \\lya photons through the streams, or through the intervening intergalactic medium. We assume that most of the radiation emitted at $z\\sim 3$ will reach the observer at $z=0$ without undergoing significant attenuation. Resonant line scattering in the optically thick streams will tend to spread out the \\lya emission region, while the repeated Doppler shifts broaden the line profiles, and the increased path length of the random walk amplifies the absorption by dust.  Such effects are expected to modify the images and line profiles, but to have only a small effect on the total \\lya luminosity. An analysis of \\lya emission from our simulated galaxies including the effects of radiative transfer and dust absorption and a more accurate treatment of photoionisation from stars will be presented in a forthcoming paper (Kasen et al, in preparation).  Other hydrogen emission lines, such as L$\\beta$ or H$\\alpha$, are expected to be two orders of magnitude less luminous than \\lya in our model \\citep{baldwin, miller2}. The column densities of $N_{\\rm HI}\\sim 10^{20}\\cmms$ in the cold streams should also be detectable as \\lya absorption in quasar spectra \\citep{prochaska99, wolfe}. Considering all the galaxies that are intersected by a line of sight to a background quasar with an impact parameter $<\\Rv$, we crudely estimate from the simulations that absorption by $N_{20}>1$ hydrogen should be detected in $\\sim 20\\%$ of the galaxies. Alternatively, \\lya photons that are emitted form the central galaxy should be absorbed in the radial streams feeding the same galaxy at $N_{20}>10$, but only in $\\sim 5\\%$ of the galaxies. The streams could also be detectable as low-ionisation metal absorbers \\citep[e.g.][]{gnat} as long as the metallicity in the streams is greater than $\\sim 0.01$ solar (paper in preparation).  Our results support the idea that the observed LABs are direct detections of the cold steams that drive the evolution of massive galaxies at high redshifts. Even though the observed LABs are sometimes associated with central sources that are energetic enough to power the observed \\lya emission, such as starbursts and AGNs, these central sources are very different from each other in the different galaxies, and in many LABs they are absent altogether \\citep{yang, geachb}. The gravitational heating associated with the inflowing cold streams is a natural mechanism for driving the extended \\lya cooling radiation observed as LABs, and this extended \\lya emission is inevitable in most high-redshift galaxies."
    },
    "0911/0911.0669_arXiv.txt": {
        "abstract": "Future galaxy surveys will map the galaxy distribution in the redshift interval $0.5<z<2$ using near-infrared cameras and spectrographs. The primary science goal of such surveys is to constrain the nature of the dark energy by measuring the large-scale structure of the Universe. This requires a tracer of the underlying dark matter which maximizes the useful volume of the survey. We investigate two potential survey selection methods: an emission line sample based on the \\ha\\ line and a sample selected in the H-band. We present predictions for the abundance and clustering of such galaxies, using two published versions of the \\galform\\ galaxy formation model. Our models predict that \\ha\\ selected galaxies tend to avoid massive dark matter haloes and instead trace the surrounding filamentary structure; H-band selected galaxies, on the other hand, are found in the highest mass haloes. This has implications for the measurement of the rate at which fluctuations grow due to gravitational instability. We use mock catalogues to compare the effective volumes sampled by a range of survey configurations. To give just two examples: a redshift survey down to $H_{\\rm AB}=22$ samples an effective volume that is $\\sim 5-10$ times larger than that probed by an \\ha\\ survey with $\\logfha > -15.4$;  a flux limit of at least $\\logfha = -16$ is required for an \\ha\\ sample to become competitive in effective volume. ",
        "introduction": "A number of approaches have been proposed to uncover the nature of the accelerating expansion of the Universe which involve measuring the large scale distribution of galaxies \\citep[e.g ][]{albrecht06,peacock06}. The ability of galaxy surveys to discriminate between competing models depends on their volume. Once the solid angle of a survey has been set, the useful volume can be maximised by choosing a tracer of the large-scale structure of the Universe which can effectively probe the geometrical volume. This depends on how the abundance of tracers drops with increasing redshift, and how much of this decline is offset by an increase in the clustering amplitude of the objects.  Several wide-angle surveys have probed the redshift interval between $0<z<1$ \\citep[e.g ][]{colless03,york00,cannon06,blake09}. The next major step up in volume will be made when the range from $0.5 <z <2$ is opened up with large near-infrared cameras and spectrographs which are mounted on telescopes able to map solid angles running into thousands of square degrees. From the ground, this part of the electromagnetic spectrum is heavily absorbed by water vapour in the Earth's atmosphere and affected by the strong atmospheric OH emission lines. A space mission to construct an all-sky map of galaxies in the redshift range $0.5<z<2$ would have a significant advantage over a ground based survey in that the sky background in the near-infrared (NIR) is around 500 times weaker in space than it is on the ground.  An important issue yet to be resolved for a galaxy survey extending to  $z \\sim 2$ is the construction of the sample and the method by which the redshifts will be measured. One option is to use slitless spectroscopy and target the \\ha\\ emission line.  \\ha\\ is located at a restframe wavelength of $\\lambda = 6563$\\AA{}, which, for galaxies at $z>0.5$, falls into the near-infrared part of the electromagnetic spectrum \\citep{thompson96,mccarthy99,hopkins00,shim09}. \\ha\\ emission is powered by UV ionizing photons from massive young stars. The only source of attenuation is dust, which is less important at the wavelength of \\ha\\ than it is for shorter wavelength lines. This makes \\ha\\ a more direct tracer of galaxies which are actively forming stars than other lines such as \\lya, OII, OIII, H$\\beta$ or H$\\gamma$, which suffer from one or more sources of attenuation (i.e. dust, stellar absorption, resonant scattering) and which are more sensitive to the metallicity and ionisation state of the gas. The second option is to use some form of multi-slit spectrograph to carry out a redshift survey of a magnitude limited sample. The use of a slit means that unwanted background is reduced, allowing fainter galaxies to be targetted. Also, it is easier to identify which spectrum belongs to which galaxy with a slit than it is with slitless spectroscopy. Targets could be selected in the H-band at an effective wavelength of just over  1 micron, which is around the centre of the near infrared wavelength part of the spectrum.   Space missions designed to carry out redshift surveys like the ones outlined above are currently being planned and assessed on both sides of the Atlantic. At the time of writing, the European Space Agency is conducting a Phase A study of a mission proposal called Euclid \\footnote[1]{http://sci.esa.int/science-e/www/object/index.cfm?fobjectid=43226}, one component of which is a galaxy redshift survey. Both of the selection techniques mentioned above are being evaluated as possible spectroscopic solutions. The slit solution for Euclid is based on a novel application of digital micromirror devices (DMDs) to both image the galaxies to build a parent catalogue in the H-band and to measure their redshifts (see Cimatti et~al. 2009 for further details about the Euclid redshift survey). A \\ha\\ mission is also being discussed in the USA \\footnote[2]{http://jdem.gsfc.nasa.gov/}. At this stage, the sensitivity of these missions is uncertain and subject to change. For this reason we consider a range of \\ha\\ flux limits and H-band magnitudes when assessing the performance of the surveys. The specifications and performance currently being discussed for these missions have motivated the range of fluxes that we consider.  A simple first impression of the relative merits of different selections methods can be gained by calculating the effective volume of the resulting survey. This requires knowledge of the survey geometry and redshift coverage, along with the redshift evolution of the number density of sources and their clustering strength. In this paper we use published galaxy formation models to predict the abundance and clustering of different samples of galaxies in order to compute the effective volumes of  a range of \\ha\\ and H-band surveys. Observationally, relatively little is known about the galaxy population selected by \\ha\\ emission or H-band magnitude at $0.5 < z < 2$. Empirically it is possible to estimate the number density of sources from the available luminosity function data and, on adopting a suitable model, to use the limited clustering measurements currently available to infer the evolution of the number density and bias \\citep{shioya08,morioka08,geach08}. Geach et~al. (2009), in a complementary study to this one, make an empirical estimate of the number density of \\ha\\ emitters, and combine this with the predictions of the clustering of these galaxies presented in this paper to estimate the efficiency with which \\ha\\ emitters can measure the large scale structure of the Universe.  The outline of the paper is as follows: in Section~\\ref{sec.models} we give a brief overview of the models. Some general properties of \\ha\\ emitters in the models, such as luminosity functions (LF), equivalent widths (EW) and clustering bias are presented in Section~\\ref{sec.properties} as these have not been published elsewhere. In Section~\\ref{sec.DE} we show how our models can be used to build mock survey catalogues. We analyse the differences in the clustering of \\ha\\ emitters and H-band selected galaxies and present an indication of the efficiency with which different surveys trace large-scale structure (LSS). Finally, we give our conclusions in Section~\\ref{sec.conclusions}. ",
        "conclusions": "\\label{sec.conclusions}  In this paper we have presented the first predictions for clustering measurements expected from future space-based surveys to be conducted with instrumentation sensitive in the near-infrared. We have used published galaxy formation models to predict the abundance and clustering of galaxies selected by either their \\ha\\ line emission or H-band continuum magnitude. The motivation for this exercise is to assess the relative performance of the spectroscopic solutions proposed for galaxy surveys in forthcoming space missions which have the primary aim of constraining the nature of dark energy.  The physical processes behind \\ha\\ and H-band emission are quite different. \\ha\\ emission is sensitive to the instantaneous star formation rate in a galaxy, as the line emission is driven by the number of Lyman continuum photons produced by massive young stars. Emission in the observer frame H-band typically probes the rest frame $R$-band for the proposed magnitude limits and is more sensitive to the stellar mass of the galaxy than to the instantaneous star formation rate.  The {\\tt GALFORM} code predicts the star formation histories of a wide population of galaxies, and so naturally predicts their star formation rates and stellar masses at the time of observation. Variation in galaxy properties is driven by the mass and formation history of the host dark matter halo. This is because the strength of a range of physical effects depend on halo properties such as the depth of the gravitational potential well or the gas cooling time. This point is most striking in our plot of the spatial distribution of \\ha\\ and H-band selected galaxies, Fig.~\\ref{fig.dm}. This figure shows remarkable differences in the way that these galaxies trace the underlying dark matter distribution. \\ha\\ emitters avoid the most massive dark matter haloes and trace out the filamentary structures surrounding them. The H-band emitters, on the other hand, are preferentially found in the most massive haloes. This difference in the spatial distribution of these tracers has important consequences for the redshift space distortion of clustering.  In this paper we have studied two published galaxy formation models, those of \\citet{baugh05} and \\citet{bower06}. The models were originally tuned to reproduce a subset of observations of the local galaxy population and also enjoy notable successes at high redshift. We presented the first comparison of the model predictions for the properties of \\ha\\ emitters, extending the work of \\citet{dell1,dell2} and \\citet{orsi08} who looked at the nature of Lyman-alpha emitters in the models. Observations of \\ha\\ emitters are still in their infancy and the datasets are small. The model predicitions bracket the current observational estimates of the luminosity function of emitters and are in reasonable agreement with the distribution of equivalent widths.  The next step towards making predicitions of the effectiveness of future redshift surveys is to construct mock catalogues from the galaxy formation models \\citep[see ][]{baugh08}. Using the currently available data, we used various approaches to fine tune the models to reproduce the observations as closely as possible. The main technique was to rescale the line and continuum luminosities of model galaxies; another approach was to randomly dilute or sample galaxies from the catalogue. This allowed us to better match the number of observed galaxies. The resulting mocks gave reasonable matches to the available clustering data around $z \\sim 2$. Our goal in this paper was to make faithful mock catalogues. The nature of \\ha\\ emitters in hierarchical models will be pursued in a future paper.  The ability of future surveys to measure the large scale structure of the Universe can be quantified in terms of their effective volumes. The effective volume takes into account the effect of the discreteness of sources on the measurment of galaxy clustering. If the discreteness noise is comparable to the clustering signal, it becomes hard to extract any useful clustering information. Once this point is reached, although the available geometrical volume is increased by going deeper in redshift, in practice there is little point as no further statistical power is being added to the clustering measurments. The error on a power spectrum or correlation function measurement scales as the inverse square root of the effective volume. In the case of flux-limited samples, the number density of sources falls rapidly with increasing redshift beyond the median redshift. Even though the effective bias of these galaxies tends to increase with redshift, it does not do so at a rate sufficient to offset the decline in the number density. The \\galform\\ model naturally predicts the abundance and clustering strength of galaxies needed to compute the effective volume of a galaxy survey.  The differences in the expected performance of \\ha\\ and H-band selected galaxies when measuring the power spectrum is related to the different nature of the galaxies selected by these two methods. \\ha\\ emitters are active star forming galaxies, which makes them have smaller bias compared to H-band selected galaxies. Their redshift distribution is also very sensitive to the details of the physics of star formation: The effect of a top-heavy IMF in bursts in the \\bau\\ model boosts the number density of bright emitters, making the redshift distrubtion of \\ha\\ emitters very flat and slowly decreasing towards high redshifts, in contrast to the predictions of the \\bow\\ model, where a sharp peak at $z\\sim 0.5$ and a rapid decrease for higher redshifts is found. H-band galaxies are less sensitive to this effect, and the redshift distributions are similar in both models. This is why the balance between the power spectrum amplitude (given by the effective bias) and the number density is translated in two different effective volumes for \\ha\\ and H-band selected galaxies.  Although there are differences in detail between the model predictions, they give similar bottom lines for the effective volumes of the survey configurations of each galaxy selection. Comparing the spectroscopic solutions in Table~\\ref{table.veff}, a slit based survey down to $H_{\\rm AB}=22$ would sample 4-10 times the effective volume which could be reached by a slitless survey to $\\logfha = -15.4$, taking into account the likely redshift success rate. To match the performance of the H-band survey, an \\ha\\ survey would need to go much deeper in flux, down to $\\logfha = -16$.  We have also looked at the accuracy with which \\ha\\ emitters and H-band selected galaxies will be able to measure the bulk motions of galaxies and hence the rate at which fluctuations are growing, another key test of gravity and the nature of dark energy. All of the samples we considered showed a small systematic difference between the measured growth rate and the theoretical expectation, at about the $1 \\sigma$ level. The error on the growth rate from an \\ha\\ survey with $\\logfha > -15.4$ was found to be about three times larger than that for a sample with $H_{AB}<22$."
    },
    "0911/0911.1279_arXiv.txt": {
        "abstract": "\\noindent Previous work has shown that Lyman break galaxies (LBGs) display a range in structures (from single and compact to more clumpy and extended) that is different from typical local star-forming galaxies. Recently, we have introduced a sample of rare, nearby ($z<0.3$) starburst galaxies that appear to be good analogs of LBGs. These ``Lyman Break Analogs'' (LBAs) provide an excellent training set for understanding starbursts at different redshifts.  We present an application of this by comparing the rest-frame UV and optical morphologies of 30 LBAs with those of galaxies at $z\\sim2-4$ in the Hubble Ultra Deep Field. We compare LBAs with star-forming s$BzK$ galaxies at $z\\sim2$, and LBGs at $z\\sim3-4$ at the same intrinsic UV luminosity ($L_{UV}\\gtrsim0.3L^*_{z=3}$). The UV/optical colors and sizes of LBAs and LBGs are very similar, while the $BzK$ galaxies are somewhat redder and larger. LBAs lie along a mass-metallicity relation that is offset from that of typical local galaxies, but similar to that seen at $z\\sim2$.  There is significant overlap between the morphologies ($G$, $C$, $A$ and $M_{20}$) of the local and high redshift samples, although the high redshift samples are somewhat less concentrated and clumpier than the LBAs.  Based on their highly asymmetric morphologies, we find that in the majority of LBAs the starbursts appear to be triggered by interactions/mergers. When the images of the LBAs are degraded to the same sensitivity and linear resolution as the images of LBGs and BzK galaxies, we find that these relatively faint asymmetric features are no longer detectable. This effect is particularly severe in the rest-frame ultraviolet. It has been suggested that high redshift galaxies experience intense bursts unlike anything seen in the local universe, possibly due to cold flows and instabilities. In part, this is based on the fact that the majority ($\\sim$70\\%) of LBGs do not show morphological signatures of interactions or mergers. Our results suggest that this evidence is insufficient, since a large fraction of such signatures would likely have been missed in current observations of galaxies at $z\\sim2-4$. This leaves open the possibility that clumpy accretion and mergers remain important in driving the evolution of these starbursts, together with rapid gas accretion through other means. ",
        "introduction": "\\label{sec:intro}  One of the key tasks in galaxy evolution is to understand how the young, forming galaxies observed at high redshift relate to the well-defined Hubble sequence observed at the present epoch. The study of sizes and morphologies of large samples of galaxies as a function of redshift would not have been possible without the {\\it Hubble Space Telescope} (HST). Studies of the rest-frame UV sizes of Lyman Break Galaxies (LBGs) at $z\\sim3-7$ indicate that they are mostly very compact objects with a single core ($r_{1/2}\\sim0.7-1.5$ kpc), and that the size distribution develops a tail of larger sized objects of up to several kpc towards lower redshifts \\citep[e.g.][]{bouwens04,ferguson04,papovich05,oesch09b}. Morphological studies performed by means of (a combination of) visual classifications, quantitative morphological parameters and two-dimensional profile fitting \\citep[e.g.][]{abraham96,lotz04,lotz06, delmegreen05,ravindranath06,conselice09,petty09} have shown that this increase in sizes at $z\\sim2-4$ is related to the accumulation of luminous star-forming clumps or cores within a variety of structures, including spheroid- and disk-like objects and irregular objects. These clumpy systems are expected to coalesce and form a spheroid while a surrounding disk may grow through the continued accretion of gas \\citep{belmegreen08,belmegreen09,dekel09}. Individual clumps could in some cases be the nuclei of star-forming objects that are merging together, or they could be giant starburst regions inside a larger gaseous system induced by merging or a plentiful smooth or ``lumpy'' gas supply. Detailed knowledge on the importance of such processes would, in principle, provide powerful constraints on models of galaxy formation \\citep[e.g.][]{birnboim03,somerville01,somerville08,guo08,guo09,dekel09}, but observationally they are hard to ascertain, especially at high redshift. First, estimates of the (major) merger rate of galaxies by means of galaxy pair counts are difficult and critically depend on the merger time-scale, which may evolve with redshift \\citep{kitzbichler08}. Second, our ability to identify galaxy mergers is a strong function of, e.g., the stage of the merger, the viewing angle, and gas fraction \\citep[e.g.][]{lotz08}. Third, at high redshift it is neither possible to directly measure the (HI) gas fractions of galaxies nor to map the distribution of intergalactic hydrogen gas believed to be the supplying reservoir. However, the relatively wide range in kinematic properties of the emission line gas observed in high redshift galaxies \\citep[e.g.][]{law07,law09,lehnert09,forster09,lowenthal09} may indicate that a variety of the basic mechanisms outlined above could be at play.  One of the most basic tools that can be used to study the origin of these peculiar morphologies at high redshift is to contrast them against local or lower redshift samples of galaxies \\citep[e.g.][]{hibbard97,burgarella06,lotz06,scarlata07,cardamone09,delmegreen09,ostlin09,petty09,rawat09}. However, it is important to keep in mind that the properties of galaxies in the local universe are generally very different from those at high redshift, making it hard to disentangle actual physical differences from observational biases. To better facilitate the straight comparison, the ``Lyman break analogs'' (LBA) project was designed in order to search for local starburst galaxies that share typical characteristics of high redshift LBGs \\citep{heckman05}. In brief, the UV imaging survey performed by the Galaxy Evolution Explorer (GALEX) was used in order to select the most luminous ($L_{FUV}>10^{10.3}$ $L_\\odot$) and most compact ($I_{FUV}>10^{9}$ $L_\\odot$ kpc$^{-2}$) star-forming galaxies at $z<0.3$. Although such galaxies are very rare, they tend to be much more luminous in the UV than typical local starburst galaxies studied previously \\citep{heckman98,meurer97,meurer99} consistent with high SFRs and relatively little dust extinction. The median absolute UV magnitude of the sample is --20.3, corresponding to $\\simeq0.5L_{z=3}^*$, where $L_{z=3}^*$ is the characteristic luminosity of LBGs at $z\\sim3$ \\citep[][i.e., $M_{1700,AB}=-21.07$]{steidel99}. Analysis of their spectra from the Sloan Digital Sky Survey (SDSS) and their spectral energy distributions from SDSS, GALEX, Spitzer and VLA and follow-up spectroscopy subsequently showed that the LBAs are similar to LBGs in their basic global properties, including: stellar mass, metallicity, dust extinction, SFR, and emission line gas properties \\citep{heckman05,hoopes07,basu-zych07,overzier09}.  One of the advantages of LBAs being so bright in the rest-frame UV is that their morphologies can be easily compared with typical galaxies at high redshift without having to artifically brighten them as was done in previous studies \\citep{lotz04,petty09}. In \\citet{overzier08} (Paper I) we analyzed the UV morphologies of a small sample of 8 LBAs observed with HST, finding that most of the UV emission in the LBAs originates in highly compact burst regions in small, clumpy galaxies that appear morphologically similar to LBGs. We argued that if LBGs at high redshift are also small merging galaxies similar to the LBAs, this would be very hard to detect given the much poorer physical resolution and sensitivity. In this paper, we present an analysis of the morphologies of our full data set consisting of rest-frame UV and optical HST images of 30 LBAs and compare with star-forming galaxies at $z\\simeq2,3,4$ in the Hubble Ultra Deep Field (HUDF). The structure of this paper is as follows. In Section 2 we describe our low and high redshift samples, the observations and data reduction methods, and our techniques for producing redshifted simulated images as well as for performing parametrized galaxy morphologies. In Section 3 we compare the rest-frame UV and optical colors, sizes and morphologies of the LBAs, BzKs and LBGs. We discuss our results in Section 4, followed by a summary of the main results. We use the AB magnitude system throughout the paper, and assume a cosmology [$\\Omega_M$,$\\Omega_\\Lambda$,$H_0$] $=$ [0.27,0.73,73.0] (with $H_0$ in km s$^{-1}$ Mpc$^{-1}$) so that the angular scales at $z\\approx0.2$ and 3.0 are about 3 and 8 kpc arcsec$^{-1}$, respectively. ",
        "conclusions": "\\label{sec:disc}  \\subsection{Morphologies of LBAs at different redshifts}  In previous works we have shown that LBAs are clumpy starburst galaxies with peculiar morphologies that are most consistent with merging \\citep{overzier08,overzier09}. Although we currently do not know the relative contribution from major and minor mergers, the diversity in LBA morphologies at least suggests a range in merger conditions and stages albeit that they all have in common that their starbursts are all very young (ages of a few tens of Myr), compact and UV-luminous \\citep{overzier09}. We have shown that, both in the UV and optical, asymmetry appears to be a better indicator for this merger activity than $M_{20}$. The median $A$ is larger than the merger criterion of $A>0.35$  from \\citet{conselice09}, while the median $M_{20}$ is much lower than the merger criterion of $M_{20}\\gtrsim-1.1$ \\citep[e.g.][]{lotz04,conselice09}. We redshifted the LBAs to $z=2-4$ and found a significant drop in the UV asymmetries such that the median $A$ fell well below the merger criterion. In the optical, the drop was smaller and less significant. $M_{20}$ became even smaller as the different star-forming clumps blended into single light concentrations at high redshift. We thus conclude that if galaxies similar to LBAs were observed in the UV at high redshift, all but a few would be mistaken as being relatively smooth and symmetric rather than identified as mergers. In the optical, we estimate that $\\sim$50\\% would be identified as symmetric.  \\citet{lotz08} investigated in detail the complicated relation between the ability to morphologically classify merging galaxies on one hand, and the physical details of the merger on the other. As expected, the merger observability time-scale (i.e., the time during which a merger is identified as such compared to the total merging time) critically depends on, e.g., orbital parameters, gas fraction, dust, SN feedback, and viewing angle. These results explain, for example, why the $G,C,A,M_{20}$ system is much more effective in identifying the roughly two-thirds of local ULIRGs having more than one nuclei that are seen either just prior to final coalescence or in projection at first passage, than single nucleus ULIRGs \\citep{lotz04,lotz08}. In a similar fashion, this ``merger observability'' as defined by \\citet{lotz08} can also explain why a significant fraction of the LBAs do not appear to be particularly disturbed for some or all of the morphological parameters. However, in addition to these effects, our simple simulations clearly show that for LBA-like galaxies observed at high redshift the decreased physical resolution and sensitivity are of equal or even greater importance than the physical details of the merger.  \\subsection{Morphologies of LBAs, BzKs and LBGs}  The median $M_{20}$ of LBGs at $z\\sim3-4$ is larger than that of LBAs, while the $A$ is similar. The fractions of LBGs classified as disturbed based on either one of these parameters in the UV is $\\sim$30\\%, consistent with previous works \\citep{lotz04,lotz06,conselice09}. In the optical, a similar result is obtained based on $A$, while a much smaller fraction is derived based on $M_{20}$. It is interesting to note that the differences observed between $BzK$ galaxies and LBAs are much more significant than those between LBGs and LBAs: in the UV, $BzK$s have the highest $A$ and $M_{20}$ and the lowest $C$ and $G$ of all samples considered here. In the optical, $BzK$s had the largest median half-light radii.  It must be noted, however, that our way of selecting LBAs based on a UV surface brightness criterion was tuned to select objects having similar UV luminosities and sizes as typical LBGs at $z\\sim3$ \\citep{heckman05}. In principle, the LBA selection could be expanded or modified to search for such objects more similar to $BzK$ galaxies, but one must be more careful as such a selection would include a wider variety of galaxies that are not so suitable analogs of high redshift objects compared to LBAs \\citep{hoopes07}. An example of a particularly clumpy LBA is the luminous blue compact dwarf galaxy Haro 11 at $z=0.02$ \\citep{ostlin09}, which lies close to the edge of our LBA selection criteria based on FUV luminosity and surface brightness \\citep{grimes07}. It is similar to the LBAs in most aspects but with a relatively high degree of clumpiness due to three strong light concentrations in an otherwise amorphous, forming galaxy. Its UV morphology at high redshift is similar to that of double nucleated galaxies identified on the basis of a large $M_{20}$ or even a ``by eye''  classification \\citep{overzier08}.  \\subsection{Possible Implications for accretion processes in starbursts at high redshift}  Based on the results presented in the previous sections, we can now make a series of statements purely based on the comparison of the morphologies of LBAs and high redshift starbursts (such as BzKs and LBGs):\\\\  1.  LBAs show clumpy star formation.  Our redshift simulations suggest that real high redshift samples have a similar (LBGs; $z\\sim3-4$) or perhaps a slightly stronger ($BzK$s; $z\\sim2$) degree of clumpy star formation.\\\\  2.  LBAs show luminous UV clumps and faint optical tidal features that, in the majority of cases, are interpreted by us as being due to mergers based on the visual evidence. Our simulations suggest that this information can not be recovered from the rest-frame UV morphological parameters at high redshift, while in the rest-frame optical perhaps 1 out of every 2 objects would be classified as disturbed. This clearly demonstrates that current estimates for the fraction of galaxies at high redshift having disturbed morphologies (as opposed to smooth/symmetric morphologies) is most likely a (weak) lower limit on the true fraction. This statement is independent of the physical mechanism that is the cause of these disturbed morphologies.\\\\  3.  The observations suggest a scenario where either cold gas accretion and internal instabilities can drive clumpy star formation in LBGs at levels higher than any seen in LBAs in the local universe, or where clumpy cold accretion, possibly at the level of major or minor mergers may be responsible for star formation in LBAs and LBGs alike. We will discuss these possibilities in more detail below.\\\\  The latest observational evidence suggests that the fraction of LBGs at $z\\sim3-4$ having disturbed or distorted morphologies is $\\sim$30\\% \\citep[e.g.,][This Paper]{conselice03,lotz06,conselice09}. Although these disturbances do not necessarily need to be explained by galaxy mergers, it does suggest that bulk material is coming in, perhaps in the form of giant gas clouds or minor mergers. \\citet{conselice09} find that the fraction of $z\\sim4$ $B$-dropouts forming a pair with another $B$-dropout is very similar to the fraction of disturbed morphologies: $\\sim$20\\% within a 20 h$^{-1}$ kpc projected radius. Interpreting the similarity in the pair counts and morphologies as evidence for merging, they estimate that a $\\sim10^{10}$ $M_\\odot$ galaxy will undergo a major merger every 1--2 Gyr at $z\\sim3-4$. Similar high merger rates were obtained for $z\\sim2-3$, possibly with an increase towards the most massive and most luminous galaxies \\citep{conselice03,conselice09,bluck09}. With such high inferred merger rates, the mass growth of LBGs due to these mergers could make a significant contribution to the mass growth at $z\\sim3$ compared to that due to star formation fueled by a more continuous gas accretion and sustained over a period of a Gyr at a rate of $\\sim10$ $M_\\odot$ yr$^{-1}$. Perhaps these two channels of formation are not contradictory given that the high density environments of LBGs are likely to be simultaneously associated with both frequent mergers of small galaxies and rapid cold accretion (likely with some level of lumpiness).  \\citet{conselice09} suggest that the remaining 70--80\\% of LBGs, that appear as relatively smooth/symmetric and are forming stars at a similarly high rate as the disturbed objects, could be the result from rapid gas collapse, provided that a sufficiently long amount of time ($\\gtrsim$0.5 Gyr) has passed since their last major merger otherwise this would have been apparent in their morphology given the high (inferred) merger rates. However, our new results based on LBAs present an important caveat to this interpretation: we have shown that, at least for galaxies at high redshift similar to LBAs, in the majority of cases we are not able to detect significant disturbances or asymmetries to their morphologies (mostly because of redshift effects, and not because they are not there). This allows for the possibility that the fraction of galaxies at high redshift that is undergoing interactions could be much higher than currently inferred from the observations. Alternatively, less violent accretion processes coupled with large disk instabilities at high redshift are capable of triggering starbursts at levels only seen in low redshift, merger-induced samples such as the LBAs.  Future deep, high resolution observations of rest-frame optical morphologies and kinematics (see below), together with constraints on the LBG pair fractions and small-scale clustering \\citep{conselice09,cooke09} will perhaps allow us to distinguish between these scenarios.  \\subsection{Relation to Studies of Gas kinematics}  In recent years, the study of high redshift LBGs has advanced considerably beyond studies that are based purely on morphology. Mainly through the use of integral field spectrographs in the near-infrared it has become feasible to determine the basic kinematical properties of the emission line gas as well. Motivated by the high degree of similarity in the properties of our locally selected UV-luminous galaxies and those at high redshift, it is a useful exercise to compare the gas kinematics of LBAs and LBGs. In a first study published by \\citet{basu-zych09}, the bright Pa-$\\alpha$ emission line was used to study the resolved gas dynamics in 3 of our LBAs. In two cases, a mild velocity gradient was found, but in all three cases the kinematics were dominated by the dispersion rather than structured rotation ($v/\\sigma\\sim1-2$). Simulating the data at $z\\sim2$ demonstrated that the (gas) kinematical properties of these objects are similar to the kinematical profiles commonly seen at high redshift \\citep[e.g.][]{law07,law09,genzel08,forster09,lowenthal09}.  \\citet{lehnert09} argued that neither the self-gravity of disks fueled by gas accretion flows nor the internal velocity dispersions of massive star-forming clumps can fully explain the high velocity dispersions observed at high redshift. Furthermore, the most massive of clumps observed possibly would not have been formed at all if the gas turbulence in the disk was not high enough to begin with \\citep{belmegreen09}, while mergers alone may not be sufficient to explain those objects dominated by a large number of massive clumps in so-called ``clump-cluster'' or ``chain galaxy'' configurations at high redshift \\citep{bournaud09}. Instead, \\citet{lehnert09} suggest that the kinematic properties of the ISM in starburst galaxies is affected by the mechanical energy input resulting from massive star formation. In this scenario, the starbursts that are associated with each of the clumps drive blast waves from supernovae (SN) and stellar winds that appear sufficient to give rise to the high gas pressures and high velocity dispersions observed. As shown in \\citet{overzier09}, some LBAs show very high pressures and clear evidence for an ISM that is dominated by starburst- and SN-feedback associated with massive star-forming clumps, consistent with such a scenario. If this is correct, then estimates for the merger rates at high redshift estimated from pair counts or morphologies are perhaps less biased than those obtained from the (ionized) gas kinematics because the latter may not always trace galaxy interactions, if present, very well. In addition, \\citet{robertson08} have shown with simulations that at least some merging systems would still be classified as ``disks'' using the methodology applied to observations at $z\\sim2$ by \\citet{shapiro08}. The simulations of \\citet{robertson08} also show how some of the main structural, spectral and chemical properties of certain $BzK$ galaxies can be explained in the context of gas-rich mergers as well, eventhough it has been claimed that such systems cannot be merger remnants. Similar studies of the gas kinematics in LBAs as initiated by \\citet{basu-zych09} will be very useful, and will allow us to carefully test the methods typically applied to high redshift starbursts in a suitable low redshift comparison sample. Results on the kinematical properties in a much larger sample of LBAs are forthcoming \\citep[][in prep.]{goncalves10}."
    },
    "0911/0911.4176_arXiv.txt": {
        "abstract": "Ultra High Energy Cosmic Rays, UHECR, maybe protons, as most still believe and claim, or nuclei; in particular lightest nuclei as we advocated recently. The two model offer a dramatic different view of UHECR sky because different galactic Lorentz deflection and GZK cut-off. The first (Auger Collaboration) nucleon proposal (2007)\\cite{Auger-Nov07} foresaw  to trace clearly the UHECR GZK Universe reaching far (up to $100$ Mpc) Super-Galactic-Plane, with little angular dispersion. On the contrary  Lightest Nuclei model (2008)\\cite{Fargion2008}, inspired by observed composition and by nearest $Cen_A$  clustering (almost a quarter of the AUGER events) explains (by cut off)  a modest and narrow (few Mpc) Universe view, as well as the puzzling Virgo absence. Lightest nuclei offer a little blurred Astronomy.  Here we address to a part of the remaining scattered  events in the new up-dated Auger map (March 2009-ICRC09). We found within rarest clustering the surprising imprint of a few remarkable galactic sources. In particular we recognize a first trace of Vela, brightest gamma and radio galactic source, and smeared sources along Galactic Plane and Center. We expect in a near future much more clustering along $Cen_A$ and a few more confirm to those galactic sources. The clustering may imply additional tails of fragments (by nuclei photo-dissociation) at half energies ($2-4  \\cdot 10^{19}$eV). The UHECR light-nuclei fragility and opacity may also reflect into a train of secondaries as gamma and neutrinos UHE events at tens-hundred PeVs. These UHE neutrinos might be detectable in a coming future within nearest AUGER and  Array Fluorescence Telescope views,(few $km$ distances) by fast fluorescence   flashing of horizontal up-going $\\tau$ Air-showers. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0911/0911.3426_arXiv.txt": {
        "abstract": "Debris discs -- analogous to the Asteroid and Kuiper-Edgeworth belts in the Solar system -- have so far mostly been identified and studied in thermal emission shortward of 100$\\,\\mu{\\rm m}$. The {\\it Herschel} space observatory and the SCUBA-2 camera on the James Clerk Maxwell Telescope will allow efficient photometric surveying at 70 to 850$\\,\\mu{\\rm m}$, which allow for the detection of cooler discs not yet discovered, and the measurement of disc masses and temperatures when combined with shorter wavelength photometry. The SCUBA-2 Unbiased Nearby Stars (SUNS) survey and the DEBRIS {\\it Herschel} Open Time Key Project are complimentary legacy surveys observing samples of $\\sim\\!500$ nearby stellar systems. To maximise the legacy value of these surveys, great care has gone into the target selection process. This paper describes the target selection process and presents the target lists of these two surveys. ",
        "introduction": "The solar neighbourhood is an ideal testing ground for the study of debris discs and planetary systems. Proximity maximises dust mass sensitivity and can allow systems to be spatially resolved. Systems near the Sun span a wide range of stellar parameters e.g., mass, age, metallicity, multiplicity. Whilst determining these parameters may not be easy, the diversity included in volume limited samples makes them ideal for legacy surveys where one may wish to investigate trends as a function of many system parameters.  This paper presents five all-sky volume limited samples of nearby stellar systems with main-sequence primaries of spectral type A,F,G,K,M. These form the basis of the target lists of two complimentary surveys for debris discs using the SCUBA-2 \\citep[Submillimetre Common User Bolometer Array 2;][]{hol03,aud04} camera on the James Clerk Maxwell Telescope (JCMT), and the {\\it Herschel} space observatory \\citep[][]{pil08}.  The \\SUNSs\\ \\citep[SUNS;][]{mat07} is a large flux-limited survey of 500 systems at $850\\,\\mu{\\rm m}$. The target flux RMS is $0.7\\,{\\rm mJy}/{\\rm beam}$, equal to the extragalactic confusion limit of the JCMT at $850\\,\\mu{\\rm m}$. Shallow $450\\,\\mu{\\rm m}$ images of varying depth will be obtained simultaneously, and deep images at $450\\,\\mu{\\rm m}$ will be proposed to follow-up $850\\,\\mu{\\rm m}$ detections.  The \\DEBRIS\\ (DEBRIS) {\\it Herschel} Open Time Key Program will image 446 systems (356 in common with SUNS) at 110 and $170\\,\\mu{\\rm m}$ using the PACS \\citep[Photodetector Array Camera and Spectrometer;][]{pog08} instrument, with follow-up of around 100 systems at 250, 350 and $500\\,\\mu{\\rm m}$ using the SPIRE \\citep[Spectral and Photometric Imaging Receiver;][]{gri08} instrument. This survey is primarily driven by the $110\\,\\mu{\\rm m}$ band, which has the highest dust mass sensitivity for cold discs such as the Kuiper-Edgeworth belt of our Solar System. The intended flux RMS at $110\\,\\mu{\\rm m}$ is $1.2\\,{\\rm mJy}/{\\rm beam}$, which is twice the predicted extragalactic confusion limit. $170\\,\\mu{\\rm m}$ images are taken simultaneously with a predicted RMS of $1.7\\,{\\rm mJy}/{\\rm beam}$, equal to the predicted extragalactic confusion limit in this band.  The primary goals of these surveys are statistical: In general how do debris disc properties vary with stellar mass, age, metallicity, system morphology (multiplicity, component masses, separations), presence of planets etc. To be able to answer so many questions, and to minimise the risk of unforeseen selection effects, large samples and simple, clearly defined target selection criteria are required. Volume limited samples satisfy these requirements, and as well as maximising the proximity of the targets, the stars nearest the Sun are very widely studied. For example, nearby stars are the main targets of radial velocity, astrometry and direct imaging planet searches. The majority of SUNS and DEBRIS targets also have photometry at 24 and $70\\,\\mu{\\rm m}$ from the MIPS (Multiband Imaging Photometer and Spectrometer) instrument on the {\\it Spitzer} space telescope, which ceased operation at the end of March 2009. This large spectral coverage, from 24 to $850\\,\\mu{\\rm m}$ for over 300 systems will be an incredible resource for detailed spectral energy distribution modelling of systems with debris discs.  Given that we are considering the closest systems to the Sun, substantial effort was required to compile the samples presented here. Late M-type stars within $10\\,{\\rm pc}$ are still being discovered \\citep[e.g.][]{hen06}, and complete homogeneous datasets covering the spectral type and distance ranges we consider do not exist. We have tried to make our sample selection using the most complete and accurate data available at the time of the DEBRIS proposal submission in October 2007. ",
        "conclusions": ""
    },
    "0911/0911.2420_arXiv.txt": {
        "abstract": "{} {A variation of the fundamental constants is expected to affect the thermonuclear rates important for stellar nucleosynthesis. In particular, because of the very small resonant energies of $^8$Be  and $^{12}$C,  the triple $\\alpha$ process is extremely sensitive to any such variations. } {Using a microscopic model for these nuclei, we derive the sensitivity of the Hoyle state to the nucleon-nucleon potential  allowing for a change in the magnitude of the nuclear interaction. We follow the evolution of 15 and 60 $M_{\\sun}$ zero metallicity stellar models, up to the end of core helium burning. These stars are assumed to be representative of the first, Population III stars. } {We derive limits on the variation of the magnitude of the nuclear interaction and model dependent limits on the variation of the fine structure constant based on the calculated oxygen and carbon abundances resulting from helium burning. The requirement that some $^{12}$C and $^{16}$O be present are the end of the helium burning phase allows for permille limits on the change of the nuclear interaction and limits of order $10^{-5}$ on the fine structure constant relevant at a cosmological redshift of $z \\sim 15-20$. } {} ",
        "introduction": "The equivalence principle is a cornerstone of metric theories of gravitation and in particular of general relativity \\citep{willbook}. This principle, including the universality of free fall, the local position and Lorentz invariances, postulates that the local laws of physics, and in particular the values of the dimensionless constants such as the fine structure constant $\\alpha_{em}\\equiv e^2/4\\pi\\varepsilon_0\\hbar c$, must remain fixed, and thus be the same at  any time and in any place. It follows that by testing the constancy of fundamental constants one actually performs a test of General Relativity, that can be extended on astrophysical and cosmological scales~\\citep[for a review, see][]{uzanrevue1,Uzan2009a}  We define a fundamental constant as any free parameter of the fundamental theories at hand~\\citep{weinberg83,duff,tri,barrowbook,jpbook}. These parameters are contingent quantities that can only be measured and are assumed constant since (i) in the theoretical framework in which they appear, there is  no equation of motion for them and they cannot be deduced from other constants and (ii) if the theories in which they appear have been validated experimentally, it means that, at the precision of the experiments, these parameters have indeed been checked to be constant. Hence, by testing for their constancy we extend our knowledge of the domain of validity of the theories in which they appear. In that respect, astrophysics and cosmology allow one to probe larger time-scales, typically of the order of the age of the universe.  One can, however, question the constancy of these dimensionless numbers and the physics which determines their value. This sends us back to the phenomenological argument by~\\citet{dirac37}, known as the `Large Number Hypothesis', according to which the dimensionless ratio $Gm_{\\rm e}m_{\\rm p}/\\hbar c$, or  simply $G$ in atomic units, should decrease as the inverse of the age of the universe, followed by~\\citet{jordan37} who formulated a field theory in which both the fine structure constant and the gravitational constant were replaced by dynamical fields. It was soon pointed out by~\\citet{fierz56} that astronomical observations can set strong constraints on the variations of these constants. This paved the way to two complementary directions of research on the fundamental constants.  On the one hand, from a theoretical perspective, many theories involving ``varying constants'' have been designed. This is in particular the case of theories involving extra-dimensions, such as the Kaluza-Klein mechanism~\\citep{kaluza21,Klein26} and string theory, in which all the constants (including gauge, Yukawa and gravitational couplings) are dynamical quantities \\citep{Wu:1986ac,Wetterich:1987fk, taylor, witten}, or in theories such as scalar-tensor theories of gravity~\\citep{Jordan49,bd61,gefdam} and in many models of quintessence~\\citep{uzan99,dam02, rundil,wetterich,Lee:2004vm, riaz}, aiming at explaining the acceleration of the universe by the dynamics of a scalar field. It is impingent on these models to explain why the constants are so constant today and provide a mechanism for fixing their value~\\citep{dn,dp}. In this respect, testing for the constancy of the fundamental constants is one of the few windows on these theories.  On the other hand, from an experimental and observational perspective, the variation of various constants have been severely constrained. This is the case of the fine structure constant for which the constraint $\\dot\\alpha_{em}/\\alpha_{em}=(-1.6\\pm2.3)\\times10^{-17}\\,{\\rm yr}^{-1}$ at $z=0$ has been obtained from the comparison of aluminium and mercury single-ion optical clocks~\\citep{rosenband}. Over a longer timescale, it was demonstrated that $\\alpha_{em}$ cannot have varied by more than $10^{-7}$ over the last 2 Gyr  from the Oklo phenomenon~\\citep{oklo0,oklo1,Fujii:1998kn,oklo2,oklo3} and over the last 4.5 Gyr from meteorite dating~\\citep{meteo1,dyson,meteo2,meteo3}. At higher redshift, $0.4 < z < 3.5$, there are conflicting reports of a an observed variation of $\\alpha_{em}$ from quasar absorption systems.  Using the many-multiplet method, \\citet{webb01} and \\cite{murph03,murphy07} claim a statistically significant variation $\\Delta\\alpha_{em}/\\alpha_{em} = (-0.54 \\pm 0.12)\\times 10^{-5}$, indicating a smaller value of $\\alpha_{em}$ in the past. More recent observations taken at VLT/UVES using the many multiplet method have not been able to duplicate the previous result \\citep{chand,sri04,quast04,srianand2007}. The use of Fe lines in \\citet{quast04} on a single absorber found $\\Delta \\alpha_{em} / \\alpha_{em} = (-0.05 \\pm 0.17) \\times 10^{-5}$.  However, since the previous result relied on a statistical average of over 100 absorbers, it is not clear that these two results are in contradiction.  In \\citet{chand}, the use of Mg and Fe lines in a set of 23 systems yielded the result $\\Delta \\alpha_{em} / \\alpha_{em} = (0.01 \\pm 0.15) \\times 10^{-5}$ and therefore represents a more significant disagreement and can be used to set very stringent limits on the possible variation in $\\alpha_{em}$. A purely astrophysical explanation for these results is also possible \\citep{amo,amo2}. At larger redshifts, constraints at the percent level have been obtained from the observation of the temperature anisotropies of cosmic microwave background at ($z\\sim10^{3}$) \\citep[\\textsl{e.g. }][]{cmb4,cmb3,cmb1,cmb2} and from big bang  nucleosynthesis (BBN)  ($z\\sim10^{10}$) \\citep[\\textsl{e.g. }][]{kpw,campolive95,Bergstrom,Nollett,Ichikawa02,bbn2,bbn3,Ichikawa04,bbn4,coc07,bbn1}. We refer to \\citet{uzanrevue1,uzanrevue2,uzanrevue3} for recent reviews on this topic. For the time being, there is no constraint on $\\alpha_{em}$ for redshifts ranging from 4 to $10^3$ although it has been proposed that 21~cm observations may allow one  to fill in the range $30<z<100$~\\citep{21cm}.  This article focuses on the effect of the possible variation of the fundamental constants on the stellar evolution of early stars, hence possibly providing constraints in a domain of redshifts where no such constraint is available. A similar issue was actually considered by~\\cite{gamow67} (see also the recent work by~\\citet{adams} who showed that the evolution of the Sun was able to exclude the Dirac model of a varying gravitational constant. In this case, non-gravitational physics is kept unchanged and the evolution of the star is affected only by the modification of gravity. Changing the non-gravitational sector has more drastic implications on stellar physics since the nuclear physics and thus the cross-sections and reaction rates of all the processes should be modified.  \\citet{rozen88} argued that the synthesis of complex elements in stars (mainly the possibility of the $3\\alpha$-reaction as the origin of the production of ${}^{12}{\\rm C}$) sets constraints on the values of the fine structure and strong coupling constants. There have been several studies on the sensitivity of carbon production to the underlying nuclear rates \\citep{livio,Fairbairn,csoto2000,ober2000,ober2002,schlattl2004,tur}. The production of ${}^{12}{\\rm C}$ in stars requires a triple tuning: (i) the decay lifetime of ${}^8{\\rm Be}$, of order $10^{-16}$~s, is four orders of magnitude longer than the time for two $\\alpha$ particles to scatter, (ii) an excited state of the carbon lies just above the energy of ${}^8{\\rm Be}+\\alpha$ and finally (iii) the energy level of ${}^{16}{\\rm O}$ at 7.1197~MeV is non resonant and below the energy of ${}^{12}{\\rm C}+\\alpha$, at 7.1616~MeV, which ensures that most of the carbon synthesized is not destroyed by the capture of an $\\alpha$-particle. The existence of this excited state of $^{12}$C was actually predicted by~\\citet{hoyle} and then observed at the predicted energy by~\\citet{dunbar} as well as its decay~\\citep{cook}. The variation of any constant which would modify the energy of this resonance, known as the Hoyle level, would dramatically affect the production of carbon.  Qualitatively, and perhaps counter-intuitively, if the energy level of the Hoyle level were increased, $^{12}$C would probably be rapidly processed to $^{16}$O since the star would, in fact, need to be hotter for the triple-$\\alpha$ reaction to be triggered. On the other hand, if it is decreased very little oxygen will be produced. From the general expression of the reaction rate (see below for details, definitions of all the quantities entering this expression, and a more accurate computation) $$ \\lambda_{3\\alpha}=3^{3/2}N_\\alpha^3 \\left(\\frac{2\\pi\\hbar^3}{M_\\alpha k_{\\rm B}T}\\right)^{3}\\frac{\\Gamma}{\\hbar} \\exp\\left[-\\frac{Q_{\\alpha\\alpha\\alpha}}{k_{\\rm B}T}\\right], $$ where $Q_{\\alpha\\alpha\\alpha}~\\sim380$~keV is the energy of the resonance, one deduces that the sensitivity of the reaction rate to a variation of $Q_{\\alpha\\alpha\\alpha}$ is $$ s=\\frac{{\\rm d}\\ln\\lambda_{3\\alpha}}{{\\rm d}\\ln Q_{\\alpha\\alpha\\alpha}} =-\\frac{Q_{\\alpha\\alpha\\alpha}}{k_{\\rm B}T}\\sim  \\left(\\frac{-4.4}{T_9}\\right). $$ where $T_9 = T/10^9$K. This effect was investigated by~\\citet{csoto2000} and \\cite{ober2000,ober2002} who related the variation of $Q_{\\alpha\\alpha\\alpha}$ to a variation of the strength of the nucleon-nucleon (N-N) interaction. Focusing on the C/O ratio in red giant stars (1.3, 5 and 20 M$_{\\sun}$ with solar metallicity) up to thermally pulsing asymptotic giant branch stars (TP-AGB)~\\citep{ober2000,ober2002} and in low, intermediate and high mass stars (1.3, 5, 15 and 25 M$_{\\sun}$ with solar metallicity)~\\citep{schlattl2004}, it was estimated that outside a window of 0.5\\% and 4\\% for the values of the strong and electromagnetic forces respectively, the stellar production of carbon or oxygen will be reduced by a factor 30 to 1000 (see also \\citet{pochet}).  Indeed, modifying the energy of the resonance alone is not realistic since all cross-sections, reaction rates and binding energies etc. should be affected by the variation of the constants. One could indeed have started by assuming independent variations of all these quantities but it is more realistic (and hence more model-dependent) to try to deduce their variation from a microscopic model.  Our analysis can then be outlined in three main steps: \\begin{enumerate} \\item Relating the nuclear parameters to fundamental constants such as the Yukawa and gauge couplings, and the Higgs vacuum expectation value. This is a difficult step because of the intricate structure of QCD and its role in low energy nuclear reactions, as in the case of BBN. The nuclear parameters include the set of relevant energy levels (including the ground states), binding energies of each nucleus and the partial width of each nuclear reaction. This involves a nuclear physics model of the relevant nuclei (mainly $^4$He, $^8$Be, $^{12}$C, and $^{16}$O for our study). \\item Relating the reaction rates to the nuclear parameters, which implies an integration over energy of the cross-sections. \\item Deducing the change in the stellar evolution (lifetime of the star, abundance of the nuclei, Hertzprung-Russel (H-R) diagram, etc.). This involves a stellar model. \\end{enumerate}  Let us summarize the main hypothesis of our work for each of these steps.  The {\\bf first step} is probably the most difficult. We shall adopt a phenomenological description of the different nuclei based on a cluster model in which the wave functions of the $^8$Be and $^{12}$C nuclei are approximated by a cluster of respectively two and three $\\alpha$ wave functions. When solving the associated Schr\\\"odinger equation, we will modify the strength of the electromagnetic and nuclear N-N interaction potentials respectively by a factor  $(1+\\delta_\\alpha)$ and $(1+\\delta_{NN})$ where $\\delta_\\alpha$ and $\\delta_{NN}$ are two small dimensionless parameters that encode the variation of the fine structure constant and other fundamental couplings. At this stage, the relation between $\\delta_{NN}$ and the gauge and Yukawa couplings is not known. This will allow us to obtain the energy levels, including the binding energy, of $^2$H, $^4$He, $^8$Be, $^{12}$C and the first  $J^\\pi$ = 0$^+$ $^{12}$C excited energy level. Note that all of the relevant nuclear states are assumed to be interacting alpha clusters. In a first approximation, the variation of the $\\alpha$ particle mass cancels out. The partial widths (and lifetimes) of these states are scaled from their experimental laboratory values, according to their energy dependence. $\\delta_{NN}$ is used as a free parameter. The dependence of the deuterium binding energy on $\\delta_{NN}$ then offers us the possibility of relating this parameter to the gauge and Yukawa couplings if one matches this prediction to a potential model via the $\\sigma$ and $\\omega$ meson masses~\\citep{sig1,sig2,coc07, dd} or the pion mass, as suggested by \\citet{pi1,pi2,pi3}.  The {\\bf second step} requires an integration over energy to deduce the reaction rates as functions of the temperature and of the new parameters $\\delta_\\alpha$ and $\\delta_{NN}$.  The {\\bf third step} involves stellar models and in particular some choices about the masses and initial metallicity of the stars. While theoretically uncertain, it is usually thought that the first stars were massive; however, their mass range is presently unknown, (for a review, see \\cite{ bromm09}). In a hierarchical scenario of structure formation, they were formed a few $\\times~10^8$ years after the big bang, that is at a redshift of $z\\sim10-15$ with zero metallicity (so that we can use the BBN abundances as initial conditions, see Sect.~\\ref{sec_stellar}). We thus focus on Population~III stars (Pop III) with typical masses, 15 and 60 $M_{\\sun}$, assuming no rotation. Our computation is stopped at the end of core helium burning.  The {\\bf final step} would be to use these predictions to set constraints on the fundamental constants, using stellar constraints such the C/O ratio which is in fact observable in very metal poor stars. While this article can be seen as a theoretical investigation that describes the expected effect of a variation of the fundamental constants, it also sheds some interesting light on stellar physics and its sensitivity to fundamental physics.  The article is logically organized as follows. Section 2 recalls the basis of the $3\\alpha$-reaction, Section 3 describes the nuclear physics modeling (first step), Section 4 is devoted to stellar implications and Section 5 to the discussion. Technical details are gathered in the appendices. ",
        "conclusions": "As we have seen in the previous sections, the extreme sensitivity of the triple $\\alpha$ process to the resonant energy of the Hoyle state can lead to very different histories for massive Population III stars.  In particular, we have shown that very slight variations in the nucleon-nucleon potential (of order a few $\\times 10^{-3}$) can lead to very different core compositions at the end of CHeB. We identified two cases (III and IV) corresponding to nearly pure $^{24}$Mg or pure $^{12}$C cores. These cases were present in both the 15 and 60 M$_{\\sun}$ models studied. Below $\\delta_{NN} = -0.005$, the stars end the CHeB phase with a core that is almost completely deprived of carbon, oxygen and neon. This is due to the $^{12}$C production by the 3$\\alpha$ reaction becoming extremely weak compared to the $^{12}$C($\\alpha$,$\\gamma$)$^{16}$O reaction \\citep[for which we have used the rates of][]{kunz02}. As soon as a little amount of $^{12}$C is produced, it is transformed into $^{16}$O, which in turn is transformed into $^{20}$Ne and then $^{24}$Mg due to the high temperature and density at which He burning occurs in these models. Above $\\delta_{NN} = +0.002$, the models end the CHeB phase devoid of $^{16}$O.  We have checked the limiting values for $\\delta_{NN}$ variation with stellar models in two different mass domains. For the 15 $M_{\\sun}$ models, the lower limit is slightly larger than for the 60 $M_{\\sun}$ models. This is due to the fusion phases occurring at higher $T_{\\rm c}$ in the more massive stars, at conditions where the $^{12}$C($\\alpha$,$\\gamma$)$^{16}$O, $^{16}$O($\\alpha$,$\\gamma$)$^{20}$Ne and $^{20}$Ne($\\alpha$,$\\gamma$)$^{24}$Mg reaction rates are largely dominant over the 3$\\alpha$ rate. A weak 3$\\alpha$ reaction is a bigger handicap in the high mass domain. In contrast,  the upper limit is larger for the 60 $M_{\\sun}$ models, because the 3$\\alpha$ reaction rate is a little less extreme at higher $T_{\\rm c}$.  Excluding these cases allows us to set a relatively conservative limit on  $\\delta_{NN}$, \\begin{equation} -0.004 < \\delta_{NN} < +0.002 \\, . \\label{weak} \\end{equation} A more aggressive limit would also exclude case II in which CHeB ends with a  $^{16}$O and $^{20}$Ne core with little or no $^{12}$C. In this case, one could argue \\begin{equation} -0.001 < \\delta_{NN} < +0.002 \\, . \\end{equation} For the remainder of the discussion, we will restrict our attention to the weak limit (\\ref{weak}), as our conclusions can be easily scaled to the stronger limit.  The limit in Eq. (\\ref{weak}) stems directly from the variation in $Q_{\\alpha\\alpha\\alpha}$. Excluding regions III and IV amount to limiting $Q_{\\alpha\\alpha\\alpha}$ to a range 0.3142 -- 0.5100 MeV, or \\begin{equation} -0.17 < \\frac{\\Delta Q_{\\alpha\\alpha\\alpha}}{Q_{\\alpha\\alpha\\alpha}} < 0.34. \\end{equation}  As discussed in \\S \\ref{sens}, a variation in $\\delta_{NN}$ will result in a variation in the deuterium binding energy. Using Eq. (\\ref{BDdNN}), the bound (\\ref{weak}) thus becomes \\begin{equation} -0.023 < {\\Delta B_D \\over B_D}  < +0.011 \\, . \\label{bdbound} \\end{equation} In principle, one would like to next convert the limit on $\\delta_{NN}$  or $B_D$ into a limit on the fundamental constants. Unfortunately, as we have argued earlier, i) the direct limit from the triple $\\alpha$ process based on $\\delta_\\alpha$ is far weaker than that due to $\\delta_{NN}$ and ii) in the absence of some guiding theory of unification, we can not relate the variation of $B_D$ directly to a variation in $\\alpha_{em}$. However, as discussed above and in more detail in \\citet{coc07}, we can use gauge coupling unification to relate a variation in $B_D$ to a variation in $\\alpha_{em}$ through $\\Lambda$.  Ignoring first any variation in the Yukawa couplings and Higgs vev, thus using $\\Delta B_D/ B_D = 18 R \\Delta \\alpha_{em} / \\alpha_{em}$, with $R = 36$ as is expected in the simplest grand unified theories, we obtain \\begin{equation} -3.5 \\times 10^{-5} < {\\Delta \\alpha_{em} \\over \\alpha_{em}}  < +1.8 \\times 10^{-5} \\, . \\label{abound1} \\end{equation} If we further assume the relations between gauge and Yukawa couplings and use Eq. (\\ref{DeltaBd4}), the limit, though more speculative, is actually weakened by a factor of about 2 due to the partial cancellation between the gauge and Yukawa contributions to $B_D$.  The limits on the variation of the fine structure constant derived above corresponds to a variation in $\\alpha_{em}$ between the present time and a period around a redshift $z \\sim 15-20$ where the Population III stars would have been present. These values are compatible with the similar limits (also assuming gauge coupling unification) on the variation of $\\alpha_{em}$ at a redshift of 10$^{10}$ from BBN predictions.  They are larger by a factor of 10 than the values found in the claimed detections \\citep{webb01,murph03,murphy07}  or a factor of 10 weaker than the limits from the non-detection \\citep{chand,sri04,quast04,srianand2007} of a variation in $\\alpha_{em}$ from quasar absorption systems at redshifts $z \\leq 3.5$.  We remind the reader, that in the present work, the variation in $\\delta_{NN}$ is only taken into account for the 3$\\alpha$ reaction. If the other rates are also affected, the limits found here would potentially have to be revised, because they have been determined by anomalies in the evolution that are due to a competition between the efficiency of the various rates. However, being a resonant reaction, the 3$\\alpha$ reaction is expected to be the most sensitive. Following \\citet{ober2000}, the $^{16}$O($\\alpha$,$\\gamma$)$^{20}$Ne reaction is not expected to be sensitive to variations in $\\alpha_{em}$, while the $^{12}$C($\\alpha$,$\\gamma$)$^{16}$O could be more affected by such variations because of subthresholds in the $^{16}$O nucleus. According to the same authors, this last reaction is expected to be strengthened by a weakening of the nucleon-nucleon interaction. In this case, the effects described for cases II and III would be more dramatic than the ones presented here, and so the limits might be tighter.  We conclude by asking: is it reasonable to exclude values of $\\delta_{NN}$  using nucleosynthetic constraints from stellar models? The criteria that we applied assumes the possibility for a ``normal'' succession of burning phases (H $\\rightarrow$ He $\\rightarrow$ C $\\rightarrow$ Ne $\\rightarrow$ O $\\rightarrow$ Si). Although we have not done so here,  a modified code including ``non standard'' fusion phases, would allow us to follow these models further. It is expected that the resulting yields would present large anomalies. Given the current state of abundance determination in extremely metal poor stars, it is highly improbable that the first stars would not produce fair amounts of $^{12}$C and $^{16}$O. The conservative case seem thus to offer a reasonable limit on the variations of the fundamental constants.  \\appendix \\setcounter{equation}{0} \\renewcommand{\\theequation}{A\\arabic{equation}}"
    },
    "0911/0911.4342.txt": {
        "abstract": "We explore the stability of domain wall and bubble solutions in theories with compact extra dimensions. The energy density stored inside of the wall can destabilize the volume modulus of a compactification, leading to solutions containing either a timelike singularity or a region where space decompactifies, depending on the metric ansatz. We determine the structure of such solutions both analytically and using numerical simulations, and analyze how they arise in compactifications of Einstein--Maxwell theory and Type IIB string theory. The existence of instabilities has important implications for the formation of networks of topological defects and the population of vacua during eternal inflation. ",
        "introduction": "The introduction of extra dimensions has long been one of the most useful ideas in the theorist's tool box, providing insight into important theoretical puzzles ranging from the hierarchy problem to the unification of forces and a theory of quantum gravity. However, it has been appreciated since the early days of general relativity that there will always be a proliferation of low-energy effective theories resulting from compactification of extra dimensions. This idea has gained particular force in the context of string theory, where extra dimensions are required for theoretical consistency, and the program of moduli stabilization yields a vast set of low-energy effective theories, referred to as the string theory landscape (or simply ``the landscape\").  Classically, vacua in the landscape are stable, and drive inflation if they have positive energy density (as is the case in our universe if the observed dark energy is a cosmological constant). Quantum mechanically, they can decay via the nucleation of bubbles containing some new phase, a process described in many cases by the Coleman--de Luccia (CDL) instanton~\\cite{Coleman:1980aw,Coleman:1977py}. Because the rate of bubble formation is exponentially suppressed, there is typically less than one bubble formed per Hubble volume per Hubble time, and therefore the inflating vacuum is never completely eaten up. This phenomenon is known as eternal inflation, and is a mechanism by which many different vacuum phases can be populated in different spatial regions.  Given that the postulated existence of compactified extra dimensions is the primary motivation for considering landscapes of four-dimensional vacua, it is important to determine what role extra dimensions might play in the dynamics of eternal inflation. In the four-dimensional effective theory, the overall volume of a compactification appears as a dynamical scalar field that inherits dilatonic couplings to the other low-energy fields. Therefore, the minimal model of a landscape obtained from a theory with extra dimensions contains at least two fields: one with properties determining the four-dimensional vacua (for example a scalar or gauge field), and the field related to the volume, which we will refer to as the volume modulus.  The dilatonic coupling to the volume modulus can cause a general obstruction to finding non-singular static domain walls and CDL bubbles connecting four dimensional vacua. Because of the coupling, domain walls can perturb the volume modulus out of its vacuum, in some cases making it impossible to satisfy the appropriate boundary conditions. This was first observed by Cvetic and Soleng~\\cite{Cvetic:1994ya,Cvetic:1995rp} and was applied to the landscape of Type IIB string theory by two of the authors~\\cite{Johnson:2008vn}. More recently, Yang has studied these effects in 6D Einstein--Maxwell theory~\\cite{Yang:2009wz}. A similar obstruction in more general multifield models can prevent the existence of other topological defects, such as monopoles~\\cite{Kumar:2009pr} and strings~\\cite{Achucarro:2001ii,Achucarro:1998er,Penin:1996si,Yajnik:1986tg,Yajnik:1986wq}.  The static domain wall and CDL bubble solutions result from imposing a rather restrictive ansatz on the metric and fields, which if relaxed, may allow for additional solutions that interpolate between different four-dimensional vacua. To assess this possibility, we revisit and extend our previous work~\\cite{Johnson:2008vn} to domain walls and bubbles assuming only planar, spherical, or hyperbolic symmetry. Using a toy model, we find that whenever static solutions do not exist, planar walls decay into a region where the extra dimensions decompactify. Bubble walls can also seed an expanding region where the extra dimensions decompactify. When the phase inside of the bubble has positive or zero energy density (and in some cases when the interior has negative energy density), unstable walls cause the entire bubble interior to decompactify unless the initial radius is larger than the true vacuum Hubble constant. In this case, even though the initial data contains regions of two different four-dimensional vacua, at late times the configuration contains only the original four-dimensional vacuum and a region of the decompactified phase. We apply some simple bounds to assess the stability of domain walls and bubbles in model landscapes arising from Type IIB string theory and D-dimensional Einstein--Maxwell theory, extending the analysis in Refs.~\\cite{Johnson:2008vn,Yang:2009wz}.  Even when static solutions do exist, they are not always stable to perturbations. Cosmologically, this is important for the formation of networks of domain walls or other topological defects in the early universe. Another type of perturbation that is particularly relevant for eternal inflation is the collision between bubbles. We find that in many cases, decompactification can result to the future of bubble collisions, which has important implications for the observability of collisions in a realistic landscape scenario (see Ref.~\\cite{Aguirre:2009ug} for a review on the observability of bubble collisions).  The instability of domain walls restricts the types of four-dimensional vacua that can be populated by eternal inflation. Although the CDL instanton cannot describe the creation of a region of some new vacuum when the bubble wall is unstable, we expect other types of transitions to mediate vacuum decay. If any such process has a classical description at late times, then even though these transitions might initially seed a region of some other four-dimensional vacuum, the unstable wall will quickly eradicate it, leaving behind a bubble containing only the decompactified phase. This might break the landscape up into sets of disconnected islands~\\cite{Clifton:2007en}, with the members of each island populated only in separate  realizations of an eternally inflating universe. There may even be many eternally inflating vacua that cannot populate any other four-dimensional vacua in the landscape. These considerations make it harder to ignore the question of initial conditions, a problem that eternal inflation is often purported to eliminate. The restriction on purely four-dimensional transitions also motivates a more general and comprehensive study of the possible dynamics during the early universe in theories with extra dimensions. For example, the dynamical compactification mechanism of Ref.~\\cite{Carroll:2009dn} can populate four-dimensional vacua from a higher dimensional space, leading to a multiverse where vacua of different dimensionality coexist.  The plan of the paper is as follows. In Sec.~\\ref{sec:toymodel}, we introduce our two-field toy model. We review the criteria for the existence of static domain walls in the toy model, and perform a numerical analysis of non-static solutions in Sec.~\\ref{sec:domainwalls}. Bubble solutions are analyzed analytically and numerically in Sec.~\\ref{sec:bubbles}. We then briefly describe the implications of extra dimensions for the stability of bubble collisions in Sec.~\\ref{sec:collisions}. The stability of domain walls and bubbles in the Type IIB and Einstein--Maxwell landscapes are described in Sec.~\\ref{sec:EMandtypeIIB}, and we conclude in Sec.~\\ref{sec:conclusions}.  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% ",
        "conclusions": "\\label{sec:conclusions}  The dynamics of the volume modulus can have important implications for the picture of eternal inflation in the string theory landscape. In this paper, we have shown that only a rather restricted set of configurations of the internal manifold and sources can be connected across stable planar or spherical domain walls.  The dilatonic coupling between the volume modulus and domain walls separating different four-dimensional vacua can completely overwhelm the stabilizing potential in the neighborhood of the domain wall. Using analytic methods, we have shown that non-singular static planar domain wall and SO(3,1) bubble solutions exist only when the energy scale associated with the tension of the domain wall is smaller than the energy scale associated with the potential maximum for the volume modulus. Rephrasing this condition, solutions do not exist when the tension of domain walls separating compactified vacua is larger than the tension of domain walls interpolating between the compactified and decompactified phases, indicating that it is energetically favorable for such domain walls to decompactify.  Using a set of numerical simulations, we find that this is indeed the case. When non-singular solutions do not exist, domain wall and bubble solutions seed a growing region of the decompactified phase. In the case of bubble solutions, when the true vacuum has zero or positive energy, the entire bubble interior goes over to the decompactified phase. This completely removes the spacetime region containing the true vacuum. Regardless of the transition mechanism that produces the original region of true vacuum, the true vacuum phase will not exist at late times, and therefore cannot be populated by any such transition.  In a spacetime without a cosmological constant, the presence of an unstable domain wall separating different four-dimensional vacua causes all of space to decompactify. If there is a cosmological constant, roughly a comoving Hubble volume of the background space will be converted to the decompactified phase at late times. In neither case does the spacetime evolve to a solution where different four-dimensional vacua coexist. This conclusion is independent of the creation mechanism, be it the formation of domain walls via the Kibble mechanism, or the nucleation of bubbles of true vacuum during eternal inflation.  On the other hand, our analysis shows that domain walls separating AdS vacua are not troubled by this instability. Although we have not used it in our analysis, it is interesting to note that such domain walls can be BPS, in which case their stability is guaranteed by supersymmetry. Our analysis is sufficient to show that such domain walls are not unstable to decompactification, but there may exist instabilities in other moduli fields if they couple to the domain walls.  The dilatonic coupling of the volume modulus can also play an important role in determining the outcome of a collision between two bubbles. Such collisions might leave detectable experimental signatures of eternal inflation, but only if our observed cosmology can exist to their future. For stable bubble walls, the energy released in a collision has the ability to seed a region where space decompactifies. We have performed a number of numerical simulations, which indicate that this process is in fact quite efficient. Depending on the kinematics and properties of the potential, the decompactified region can accelerate into the bubble interior, rendering such collisions unobservable. Further analysis of these preliminary results will undoubtedly be an important element in determining the observability of bubble collisions arising in context of eternal inflation driven by the string theory landscape.  We have investigated compactifications of Einstein--Maxwell and type IIB string theory, which both give rise to landscapes of vacua, domain walls, and possibly eternal inflation. Fundamental branes act as domain walls separating four-dimensional vacua in these theories, and possess a  dilatonic coupling to the volume modulus. We find that domain walls between dS or Minkowski vacua can be unstable to decompactification, as described above.  The dilatonic coupling becomes especially problematic when the potential fixing the volume modulus is small compared to the tension of the walls. This is a common feature for both KKLT and BBCQ compactifications of type IIB string theory since the volume modulus is unfixed at leading order. In this case, we find that stable domain walls and bubbles arise only when the brane wraps a strongly warped cycle, leading to a reduced tension. The instabilities provide further motivation for finding dS vacua in compactifications where all geometric moduli are fixed at the same scale, as is the case in e.g. generic type IIA compactifications. These compactifications bear some resemblance to the Einstein--Maxwell model, where we find that stability depends on the dimensionality of the cycles that a fundamental brane wraps in the internal manifold.  The existence of unstable domain walls separating different four-dimensional vacua leads us to conclude that only a restricted set of configurations of the internal manifold and sources can be populated by eternal inflation. However, our results do not exclude the existence of solutions where vacua of different dimensionality, separated by domain walls, coexist. Such a solution could contain four-dimensional vacua, separated by intervening regions where a different number of dimensions is kept large, as suggested in Ref.~\\cite{Carroll:2009dn}. Further exploration of these and other ideas will be necessary to determine the full dynamical role of extra dimensions in populating landscapes of vacua."
    },
    "0911/0911.5170_arXiv.txt": {
        "abstract": "% The search for life outside of the Solar System should not be restricted to exclusively planetary bodies; large moons of extrasolar planets may also be common habitable environments throughout the Galaxy.  Extrasolar moons, or exomoons, may be detected through transit timing effects induced onto the host planet as a result of mutual gravitational interaction.  In particular, transit timing variations (TTV) and transit duration variations (TDV) are predicted to produce a unique exomoon signature, which is not only easily distinguished from other gravitational perturbations, but also provides both the period and mass of an exomoon.  Using these timing effects, photometry greater or equal to that of the \\emph{Kepler Mission} is readily able to detect habitable-zone exomoons down to 0.2 $M_{\\oplus}$ and could survey up to 25,000 stars for Earth-mass satellites.  We discuss future possibilities for spectral retrieval of such bodies and show that transmission spectroscopy with JWST should be able to detect molecular species with $\\sim 30$ transit events, in the best cases. ",
        "introduction": "%  In this conference, named `The Pathways Towards Habitable Planets', we are trying to develop a strategy towards detecting life on another planet.  However, by only considering \\emph{planets} we may be unnecessarily restrictive by excluding another potentially common environment for life to thrive - moons.  The possibility of life-bearing moons was addressed by Williams et al. (1997) who noted that a moon of around one third of an Earth mass or larger, residing around a gas giant planet in the habitable zone, would satisfy many of the criteria traditionally placed on habitable planets.  For M-dwarf systems, habitable-zone moons have the added bonus that any tidal locking of the planet or the moon would not result in one side of the moon being locked in perpetual darkness.  Whether or not such moons could be dynamically stable remained a relativity open question until Barnes \\& O'Brien (2002) showed that a habitable-zone Jupiter-like planet could hold onto a 1$M_{\\oplus}$ moon in systems with a host star as small as $M_* \\sim 0.4 M_{\\odot}$, and remain dynamically stable for 5 Gyr.  In light of these two studies, the potential for habitable moons must be taken as being at least theoretically possible.  The formation of a $\\sim 1 M_{\\oplus}$ exomoon around a Jupiter-mass planet is not supported by the current models of the formation of the regular satellites, e.g. see Canup \\& Ward (2006).  However, the possibility of a Jupiter-like planet capturing a terrestrial planet as a moon (an irregular satellite) is a topic on which we could find no detailed studies. Given that such moons would be dynamically stable, and the lessons learnt from the failings of trying to predict exoplanetary systems based on Solar System models prior to the first actual detection, a search for extrasolar moons is not only historically justified but scientifically imperative. ",
        "conclusions": "We have shown that exomoons, in particular habitable exomoons, are detectable by coupling transit timing variations (TTV) and transit duration variations (TDV) together down to 0.2 $M_{\\oplus}$ for \\emph{Kepler}-class photometry.  Around 25,000 stars could be surveyed within \\emph{Kepler}'s field-of-view for Earth-mass habitable exomoons.  Indeed, we propose here that it may even be possible for gas giants of $\\sim 10 M_J$ to harbour a super-Earth mass exomoon, or a `super exomoon'. Subsequent characterisation of exomoons could be achieved by using the planet-moon ephemeris derived from the transit timing signal.  For the case of a system within 10pc and a 0.2 $M_{\\oplus}$ habitable exomoon, we predict that molecular species could be found using transmission spectroscopy with JWST after the binning of $\\sim$30 transit events."
    },
    "0911/0911.5346_arXiv.txt": {
        "abstract": "The properties of galaxies with the reddest observed $R-K$ colors (Extremely Red Objects, EROs), including their apparent division into passive and obscured active objects with roughly similar number densities, are a known challenge for models of galaxy formation.  In this paper we produce mock catalogues generated by interfacing the predictions of the semi-analytical {\\mor} model for the evolution of galaxies in a $\\Lambda$CDM cosmology with the spectro-photometric + radiative transfer code {\\gs} and Infrared (IR) template library to show that the model correctly reproduces number counts, redshift distributions and active fractions of $R-K>5$ sources. We test the robustness of our results against different dust attenuations and, most importantly, against the inclusion of TP-AGB stars in Simple Stellar Populations (SSPs) used to generate galaxy spectra, and find that the inclusion of TP-AGBs has a relevant effect, in that it allows to increase by a large factor the number of very red active objects at all color cuts. We find that though the most passive and the most obscured active galaxies have a higher probability of being selected as EROs, many EROs have intermediate properties and the population does not show bimodality in specific star formation rate (SSFR).  We predict that deep observations in the Far-IR, from 100 to 500 $\\mu$m, are the most efficient way to constrain the SSFR of these objects; we give predictions for future {\\it Herschel} observations, and show that a few objects will be detected in deep fields at best.  Finally, we test whether a simple evolutionary sequence for the formation of $z=0$ massive galaxies, going through a sub-mm-bright phase and then a ERO phase, are typical in this galaxy formation model. We find that this sequence holds for $\\sim25$ per cent of $z = 0$ massive galaxies, while the model typically shows a more complex connection between sub-mm, ERO and massive galaxies. ",
        "introduction": "The wealth of data coming from recent multiwavelength surveys has made it possible to study in detail distant galaxy populations from their UV-optical-Infrared broad-band spectral energy distributions (SED). Among the various categories that have been observationally defined by means of color selection techniques, the Extremely Red Objects (EROs) have attracted wide interest.  Defined on the basis of of their very red $R-K$ color (in the Vega system $R-K>5$ at least), they were first addressed in the context of deep near-infrared surveys \\citep{Elston88, McCarthy92}. Such red colors are expected to arise as a combination of an intrinsically red SED and a large K-correction at $1<z<2$. Then, EROS were regarded as natural counterparts of local elliptical galaxies \\citep[e.g.][]{Cimatti02a}, Moreover, their relatively large $K$-band fluxes suggested high stellar masses, comparable with local ellipticals, leading to the interpretation of EROs as already assembled elliptical galaxies and the redshift range $z>2$ as their formation epoch. The correct prediction of such an early and efficient formation of massive galaxies at high redshift is a well known challenge for models of galaxy formation based on the CDM cosmology: in fact early implementations predicted lower space densities of EROs than observed \\citep[e.g.,][]{Firth02,Smith02}.  On the other hand, it soon became evident that the same red $R-K$ colors might be linked to a completely different class of objects such as the active (starburst) galaxies embedded in dust shells that cause high extinction of the SEDs \\citep{Cimatti02a,PozzettiMannucci00,Smail02}.  Interestingly enough, the relative contribution to the total space density by the two sub-populations is similar, with a passive fraction of $48$ per cent at $K<19.2$ (\\citealt{Kong09}, but a ratio between 30 and 70 per cent has been reported also by \\citealt{Cimatti02a} and \\citealt{Smail02}), and $58$ per cent at $K<20.3$ \\citep{Miyazaki03}.  However, splitting the observed ERO population in two is a difficult task.  The cleanest way relies on high signal-to-noise spectroscopy, able to reveal either emission lines linked to ongoing star formation or the spectral breaks typical of aged stellar populations (see e.g., \\citealt{Cimatti02a}). Alternatively, the presence of a starburst may be revealed by additional imaging in the far-infrared, where the absorbed light is reprocessed \\citep{Kong09}, or indirectly through the morphology of these objects \\citep{StiavelliTreu01,Cimatti03}. Presently, all these techniques can be successfully applied only to a small fraction of the overall population, typically limited to the brightest sources.  As mentioned above, the redshift distribution and the expected physical properties of the two ERO sub-populations have been a long-standing problem for theoretical models of galaxy formation \\citep{Cimatti02a,Somerville04b}.  More recent models, based on improved treatment of baryonic physics \\citep{DeLucia06,Bower06,Menci06,Monaco07,Lagos08}, are able to give a better representation of the formation and assembly of massive galaxies at $z<3$. Only a few of these models have been directly tested by comparing the predicted properties of EROs with data: in most cases a reasonable success in reproducing the abundance of passive EROs has been claimed, giving support to the interpretation of these objects as passive massive galaxies at $z\\sim1-2$ that are the progenitors of local elliptical galaxy. At the same time, the same models do not produce a sufficient number of active EROs.  More specifically, \\citet{Nagamine05} implemented sub-grid baryonic physics in both a Lagrangian and Eulerian code; they found the correct abundance of massive galaxies at $z\\sim 2$, but could find almost no passive ERO, and in order to produce enough red galaxies they were forced to assume a dust attenuation as high as an equivalent $E(B-V)=0.4$ for the whole mock catalogue. \\citet{Kang06} correctly reproduced the space density of massive galaxies at $z\\sim 2$, but underpredicted the number of very red galaxies, in particular of the active class. \\citet{Menci06} proposed a more successful model able to reproduce the color distribution of galaxies at $z<1$ and the space density of massive galaxies at $1<z<2$ at the same time.  This model is able to reproduce the ERO number counts and redshift distribution, but it predicts a relatively high fraction of passive objects ($\\sim$80 per cent at $K<22$).  More recently, \\citet{GonzalezPerez09} compared different implementations of the {\\sc galform} model \\citep{Cole00} to observational data of the ERO population.  Their best fit model \\citep[i.e.,][]{Bower06} produced the correct number of EROs, even when a very red cut (up to $R-K>7$) was used, but again the predicted ERO population was dominated by passive objects, with only a small fraction of active galaxies.  On the other hand the model by \\citet{Baugh06}, which gives a good fit of the sub-mm galaxy population, was found to significantly underestimate ERO number counts.  A remarkable attempt to put quasars, bright sub-mm galaxies, EROs, and low-$z$ ellipticals into a unified scenario was proposed by \\citet{Granato04}. In their model a typical low-z elliptical forms the bulk of its mass in an intense obscured starburst phase at $z\\sim2$, when it is visible in the sub-mm band.  In the galaxy nucleus, star formation triggers accretion onto a seed black hole; the end of the star formation phase is caused by feedback from the resulting quasar, able to wipe away the interstellar medium from the galaxy and quench star formation. After $\\sim1$ Gyr the galaxy gets very red and becomes an ERO. It then evolves passively into a typical local elliptical galaxy.  The above mentioned papers have clarified the role of the ERO population as a powerful test for theoretical models of galaxy formation. These should reproduce both their global number and the diversity of their physical properties.  A successful model would be helpful in investigating the relation of EROs with other observationally selected galaxy populations, in particular with the sub-mm galaxies.  Apparently, the most challenging requirement for published model is to produce a high fraction of active EROs, especially when deep samples are concerned.  In principle, active EROs should be easy to produce in a hierarchical model of galaxy formation, because the redshift range where we observe them corresponds to the peak of the cosmic star formation rate density. \\citet{Kong06} and \\citet{Nagamine05} interpret the tension between their models and data as a sign of over-simplified recipes in treating baryonic physics or dust attenuations. Regarding the latter point, \\citet{Fontanot09a} showed that currently used recipes for dust attenuation, calibrated using low-$z$ samples, provide very scattered extinction values at $z\\sim 2$ and, more interestingly, do not match the predictions of the {\\gs} code \\citep{Silva98} that explicitly solves the equations of radiative transfer. Another important point is the contribution of the TP-AGB stars to the synthetic SEDs of model galaxies (see e.g. \\citealt{Maraston05}), which is not included in the above mentioned models. It has been shown by \\citet{Tonini08, Tonini09} that the inclusion of this extreme phase of stellar evolution has a strong impact on the expected photometry and colors of simulated galaxies at $1 \\la z \\la 4$, where intermediate-age stellar populations dominate the restframe $K$-band starlight. In the case of EROs, an increase of their number density is expected when TP-AGBs are included due to the reddening of rest-frame $V-K$. All these elements critically affect the modeling of galactic SEDs: they are therefore a primary concern for the interpretation of the physical ingredients of models.  In this paper we use the {\\mor} model (\\citealt{Monaco07}, as updated in \\citealt{LoFaro09}), interfaced with the {\\gs} spectro-photometric + radiative transfer code \\citet{Silva98} and to Infrared (IR) template library \\citep{Rieke09}, to produce mock samples of EROs. We both use the {\\it Padova} \\citep{Bertelli94} and the \\citet[][M05 hereafter]{Maraston05} library of SSPs. We find reasonable agreement with data in terms of number counts, redshift distribution and fraction of active objects, especially when M05 library is used (details on the effect of this modification on the whole galaxy population will be the subject of a forthcoming paper). This prompts us to use the model to investigate the physical properties of these object.  We address whether there is a true active/passive bimodality of the ERO population and what is the connection between EROs, sub-mm galaxies and massive galaxies at $z=0$. We finally give predictions of ERO properties for future surveys with the {\\it Herschel} satellite.  This paper is organized as follows. In sec.~\\ref{sec:model} we recall the main feature of the theoretical models {\\mor} and {\\gs}. In sec.~\\ref{sec:results} we present our results: in sec.~\\ref{sec:ero_prop} we analyze the properties of the model ERO sample and in sec.~\\ref{sec:sequence} we discuss the connection between the ERO population, the sub-mm bright galaxies and the massive galaxies at $z=0$.  Finally in sec.~\\ref{sec:final} we give our conclusions.  Throughout the paper we assume magnitudes are in the Vega system (unless otherwise stated), and a cosmological model consistent with WMAP3 results ($\\Omega_0=0.24$, $\\Omega_\\Lambda=0.76$, $h=0.72$, $\\sigma_8=0.8$, $n_{\\rm sp}=0.96$, \\citealt{Komatsu09}). ",
        "conclusions": "\\label{sec:final} \\begin{figure} \\centerline{ \\includegraphics[width=9cm]{flux_evo.ps} } \\caption{The redshift evolution of the total mass ({\\it panel a}), SSFR ({\\it panel b}), observer frame R-K color ({\\it Panel c}), and observer frame $850 \\mu$m flux ({\\it Panel d}) for 5 representative galaxies. In each panel the same greyscale corresponds to the same object. Empty symbols correspond to the $z=0$ mass of the 5 galaxies.} \\label{fig:flux_evo} \\end{figure}  In this paper we have discussed the statistical properties of the ERO population as predicted by the {\\mor} galaxy formation model, in terms of number counts, redshift distributions and active fractions of $R-K>5$ sources. We generated mock photometric catalogues by interfacing {\\mor} with the spectro-photometric + radiative transfer code {\\gs}, and created large samples of model galaxies (so as to have good statistics on the rare EROs). We used both the {\\it Padova} and the M05 SSP libraries; the latters include the contribution of TP-AGB stars. We combined the resulting UV-to-Near-IR synthetic SEDs, computed without the expensive Far-IR emission, with the \\cite{Rieke09} IR template library to estimate IR fluxes longward $3 \\mu m$.  We then selected samples of EROs (and sub-mm galaxies) by applying the corresponding color and flux cuts.  We found that our standard model with {\\it Padova} SSP library is able to reproduce the overall number counts and redshift distribution of the $R-K>5$ sources; the agreement worsens when redder cuts are considered. Other discrepancies, like the excess of faint EROs at deep $K$ fluxes and the underestimate of the number of distant passive EROs, are consequences of well-known points of tension of the model with data, like the dearth of massive $z>2$ galaxies \\citep{Fontanot07b,Cirasuolo08,Marchesini09}, or the excess of $z\\sim1$ small galaxies \\citep{Fontanot09b}; this analysis does not reveal new discrepancies.  Interestingly, {\\mor} is able to roughly reproduce the substantial contribution (of the order of $\\sim 50\\%$, see e.g. \\citealt{Cimatti02a,Miyazaki03}) of active galaxies to the global ERO population; here we define active galaxies as those with ${\\rm SSFR} > 10^{-11}$ Gyr. We ascribe this success to the higher SSFR predicted by the model at higher redshifts \\citep{Fontanot07b,Santini09}.  We showed that the ERO number counts increase considerably when using the M05 SSP library.  The addition of TB-AGBs gives a strong boost to the restframe SEDs of active galaxies longward $6000$ \\AA, reddening them and thus increasing the number of active EROs.  This leads to a good level of agreement even at very red cuts ($R-K>6$), though it does not completely solve the lack of very red passive galaxies at higher redshift. This highlights the importance of an accurate modeling of galactic synthetic SEDs when comparing the prediction of theoretical models of galaxy formation to observations, and confirms that TP-AGBs are a relevant ingredient for the correct reproduction of the high-redshift Universe \\citep{Maraston06}. The details of the modeling of dust attenuation also play a role: we checked that both increasing the fraction of cold gas locked into molecular clouds (as suggested by \\citealt{LoFaro09} to reproduce the properties of Lyman break Galaxies at $z \\ga 4$) or using the analytical approach proposed by \\citet{DeLucia07b} based on an universal attenuation law slightly reduces the number of active EROs ($\\lesssim 0.3$ dex at $K \\sim 20$).  We then considered the physical properties of the predicted ERO population with respect to a control sample of non color-selected galaxies with $K<22$ and $1<z<2$.  Overall the physical properties of EROs do not differ significantly from those of a typical galaxy in the same redshift range, apart that active galaxies with low dust attenuation and low metallicities are excluded.  We do not see any signature for a significant bimodality in the ERO population: the probability of a galaxy to be selected peaks in two regions of the SSFR--$E(B-V)$ space, corresponding to moderately attenuated passive galaxies and strongly attenuated ($E(B-V)>0.2$) active galaxies, but this trend is not visible in the overall ERO distribution.  This result is in agreement with \\citet{GonzalezPerez09}. Deep fields in the Far-IR, from 100 to 500 $\\mu$m, are the best way to constrain the SSFR of these objects, but presently planned surveys with {\\it Herschel} will detect a small number of EROs at best.  We also tested the possible relation between EROs and both the sub-mm ($f_{850 \\mu m}>0.5$ mJy) galaxies and the local massive ($M_\\ast > 10^{11}$ M$_\\odot$) galaxies. This was done to test the suggestion that these different populations reflect well-defined stages of a typical evolutionary sequence of massive galaxies, like sub-mm galaxies evolving into EROs and later into massive passive galaxies. We found that most massive EROs have a sub-mm bright main progenitor and evolve into a massive galaxy at $z=0$, and that many $z=0$ massive galaxies have an ERO or a sub-mm main progenitor, but only a minor fraction, $\\sim25$ per cent, follow the sub-mm -- ERO -- massive galaxy sequence.  Moreover, massive galaxies may easily have more than one sub-mm and/or ERO progenitor.  In fact, star formation histories predicted by {\\mor} are far more complex and exhibit a great degree of diversity, with only a small sample following the well defined sequence described above.  Our work shows that EROs can be used as a powerful constraint for theories of galaxy formation and evolution, but a better and deeper understanding of the distribution of EROs in a configuration space defined by the physical quantities, like star formation activity, stellar mass, metallicity, stellar population ages and dust obscuration, is needed.  Unfortunately, planned surveys with {\\it Herschel} may not reach the combination of depth and area required for detecting and analysing statistically significant samples of EROs. On the other hand, the same instruments may provide the required spectroscopic and photometric follow-up in order to better understand the real nature of already detected EROs."
    },
    "0911/0911.5352_arXiv.txt": {
        "abstract": "We present MMT/Megacam imaging of the Leo~IV dwarf galaxy in order to investigate its structure and star formation history, and to search for signs of association with the recently discovered Leo~V satellite. Based on parameterized fits, we find that Leo~IV is round, with $\\epsilon < 0.23$ (at the 68\\% confidence limit) and a half-light radius of $r_{h} \\simeq 130$ pc.  Additionally, we perform a thorough search for extended structures in the plane of the sky and along the line of sight. We derive our surface brightness detection limit by implanting fake structures into our catalog with stellar populations identical to that of Leo~IV.  We show that we are sensitive to stream-like structures with surface brightness $\\mu_{r}\\lesssim29.6$ mag arcsec$^{-2}$, and at this limit, we find no stellar bridge between Leo IV (out to a radius of $\\sim$0.5 kpc) and the recently discovered, nearby satellite Leo V.  Using the color magnitude fitting package StarFISH, we determine that Leo~IV is consistent with a single age ($\\sim$14 Gyr), single metallicity ($[Fe/H]\\sim-2.3$) stellar population, although we can not rule out a significant spread in these value.  We derive a luminosity of $M_{V}=-5.5\\pm0.3$.  Studying both the spatial distribution and frequency of Leo~IV's 'blue plume' stars reveals evidence for a young ($\\sim$2 Gyr) stellar population which makes up $\\sim$2\\% of its stellar mass.  This sprinkling of star formation, only detectable in this deep study, highlights the need for further imaging of the new Milky Way satellites along with theoretical work on the expected, detailed properties of these possible 'reionization fossils'. ",
        "introduction": "Since 2005, 14 satellite companions to the Milky Way have been discovered (see \\citet{willman09} and references therein).  Despite the fact that many of these objects are less luminous than a typical globular cluster ($-1.5 < M_V < -8.6$), these 14 objects have a range of properties that encompass the most extreme of any galaxies, including: the highest inferred dark matter content \\citep[e.g.][]{simongeha,geha09}, the lowest [Fe/H] content \\citep{Kirby08}, unusually elliptical morphologies \\citep[e.g. Hercules; ][]{Coleman07,sandherc}, and in some cases evidence for severe tidal disturbance \\citep[e.g. Ursa Major II; ][]{Zucker06,munoz09}.  The varied properties of these lowest luminosity galaxies are valuable probes for understanding the physics of dark matter and galaxy formation on the smallest scales.    Of the newly discovered Milky Way (MW) satellites, Leo~IV ($M_V = -5.5$, $r_h \\sim 130$ pc) is among the least studied, despite several signs that it is an intriguing object.  Leo IV appears to be dominated by an old and metal-pool stellar population \\citep{sdsssfh}.  However it also has an apparently complex color magnitude diagram (CMD), with a 'thick' red giant branch, possibly caused by either multiple stellar populations and/or depth along the line of sight \\citep{Belokurov07}. Leo IV also may have a very extended stellar distribution, despite its apparently round ($\\epsilon=0.22^{+0.18}_{-0.22}$) and compact \\citep[$r_{h}=2.5^{+0.5}_{-0.7}$ arcmin;][]{sdssstruct} morphology.  A search for variable stars was recently performed by \\citet{Moretti09}, who used the average magnitude of three RR Lyrae stars to find a distance modulus of $(m-M)=20.94\\pm0.07$ mag, corresponding to $154\\pm5$ kpc.  Interestingly, one of the three RR~Lyrae variables lies at a projected radius of $\\sim$10 arcmin, roughly three times the half light radius, leading to the suggestion that Leo~IV may actually possess a 'deformed morphology'.  Based on Keck/DEIMOS spectroscopy of 18 member stars, Leo~IV has one of the smallest velocity dispersions of any of the new MW satellites, with $\\sigma=3.3\\pm1.7$ km/s \\citep{simongeha}.  A metallicity study of 12 of these spectra \\citep{Kirby08} showed Leo~IV to be extremely metal poor, with $\\langle [Fe/H] \\rangle=-2.58$ with an intrinsic scatter of $\\sigma_{[Fe/H]}=0.75$ -- the highest dispersion among the new dwarfs.  Recently, the MW satellite Leo~V ($M_V \\sim -4.3 \\pm 0.3$) has been discovered, separated by only $\\sim$2.8 degrees on the sky and $\\sim$40 km/s from Leo~IV \\citep{leov}. With a Leo~V distance of $\\sim$180 kpc, this close separation in phase space led \\citet{leov} to suggest that the Leo~IV/Leo~V system may be physically associated. This argument was bolstered by \\citet{leovspec}, who spectroscopically identified two possible Leo~V members 13 arcminutes from the satellite's center (Leo~V's $r_{h}$ is $\\sim$0.8 arcminutes) along the line connecting Leo~IV and Leo~V, suggesting that Leo~V is losing mass.  A recent analysis by \\citet{leoivleov} of two 1 square degree fields situated between Leo~IV and Leo~V shows tentative evidence for a stellar 'bridge' between the two systems with a surface brightness of $\\sim$32 mag arcsec$^{-2}$.   Motivated by all the above, we obtained deep photometry of Leo~IV with Megacam on the MMT.  In this paper, we use these data to present a detailed analysis of both the structure and SFH of Leo~IV.  We also search for any signs of disturbance in Leo~IV which may hint at a past interaction with the recently discovered, nearby Leo~V.  In \\S~\\ref{sec:observations} we describe the observations, data reduction and photometry.  We also present our final catalog of Leo~IV stars. In \\S~\\ref{sec:struct} we derive the basic structural properties of Leo~IV, and search for signs of extended structure.  We quantitatively assess the stellar population of Leo~IV in \\S~\\ref{sec:starform} using both CMD-fitting software and an analysis of its blue plume population.  We discuss and conclude in \\S~\\ref{sec:discuss}. ",
        "conclusions": "\\label{sec:discuss}  In this work we have presented deep imaging of the Leo~IV MW satellite with Megacam on the MMT and study this galaxy's structure and SFH.  In particular, we assess reports in the literature concerning both its stellar population and its possible association with the nearby satellite, Leo~V.   Leo~IV's SFH is dominated by an old ($>12$ Gyr), metal poor ($[Fe/H]\\lesssim-2.0$) stellar population, although we uncover evidence for a young sprinkling of star formation 1-2 Gyrs ago.  Our best-fit StarFISH results indicate that a single metal poor population dominates, although the data is also compatible with a spread in metallicities.  The old population is consistent with the emerging picture that the faintest MW satellites are 'reionization fossils' \\citep[e.g.~][]{Ricotti05,Gnedin06}, who formed their stars before reionization and then lost most of their baryons due to photoevaporation.  The apparent sprinkling of young stars begs the question of what has enabled Leo~IV to continue forming stars at a low level.  There is no sign of HI in Leo~IV, with an upper limit of 609 $M_{\\odot}$ \\citep{Grcevich09}, although we note that this limit is still a factor of $\\sim$2 larger than the stellar mass associated with the young stellar population studied in \\S~\\ref{sec:BP}.  One possible mechanism for late gas accretion, and subsequent star formation, among the faint MW satellites was recently discussed by \\citet{Ricotti09} to help explain the complex SFH and gas content of Leo~T \\citep{leot,leotgas}.  In this scenario, the smallest halos stop accreting gas after reionization as expected, but as their temperature decreases and dark matter concentration increases with decreasing redshift they are again able to accrete gas from the intergalactic medium at late times, assuming they themselves are not accreted by their parent halo until $z\\lesssim1-2$.  This can lead to a bimodality in the SFH of the satellite, with both a $>12$ Gyr population and one that is $<10$ Gyr, as we see in Leo~IV.  One stringent requirement of this model is that the satellite can not have been accreted by its host halo until $z\\lesssim1-2$ (and thus not exposed to tidal stirring and ram pressure stripping until late times, allowing the satellite to retain its newly accreted gas).  Future proper motion studies will be able to test if Leo~IV is compatible with this late gas accretion model.  Another prediction of this model is the possible existence of gas-rich minihaloes that never formed stars, but could serve as fuel for star formation if they encountered one of the luminous dwarfs. More detailed study of this late gas accretion mechanism will be necessary to understand the possible diversity of SFHs in the faint MW satellites.  Additionally, we note that if the apparent segregation of young stars in Leo~IV is real, then it is not an isolated case among the MW satellites.  As has been mentioned in \\S~\\ref{sec:BP}, \\citet{Martin08} noted that Canes~Venatici I has a compact star forming region clearly offset from the galaxy as a whole, with an age of 1-2 Gyrs, similar to Leo~IV.  Additionally, Fornax has several compact clumps and shells that house young stellar populations roughly $\\sim$1.4 Gyrs old \\citep{coleman04,Coleman05,Olszewski06}.  It has been suggested that these could have been the result of a collision between Fornax and a low-mass halo, which was possibly gas-rich. \\citet{pennarubia09} investigated the disruption of star clusters in triaxial, dwarf-sized halos and found that segregated structures can persist depending on the orbital properties of the cluster, providing yet another viable mechanism.  More work is needed to distinguish between all of the above scenarios and to properly model the emerging diversity of SFHs among the new, faint MW satellites.   Structurally, Leo~IV appears to be very round, with $\\epsilon \\lesssim 0.23$ (at the 68\\% confidence limit) and a half light radius ($\\sim 130$ pc) which is typical of the new MW satellites.  An exhaustive search for signs of extended structure in the plane of the sky has ruled out any associated streams with surface brightnesses of $\\mu_{r}\\lesssim29.6$.  The extent of Leo~IV along the line of sight is less than $\\sim$15 kpc, a limit that will be difficult to improve upon given the inherent limitations of using the spread in BHB magnitudes to measure depth.  We find no evidence for structural anomalies or tidal disruption in Leo~IV.  We do not have the combination of image depth and area necessary to confirm the stellar bridge, with a surface brightness of 32 mag arcsec$^{-2}$, recently reported in between Leo~IV and Leo~V \\citep{leoivleov}.  Indeed, Leo~V is almost certainly disrupting, as discussed by \\citet{leovspec}, due to the presence of two member stars $\\sim$13 arcminutes ($>13 r_{h}$) from Leo~V's center along the line connecting the putative Leo~IV/Leo~V system.  The nature of Leo~V is still very ambiguous, with the kinematic data being consistent with it being dark matter free -- suggesting that perhaps Leo~V is an evaporating star cluster \\citep{leovspec}.  It is thus critical to obtain yet deeper data on these two systems and the region separating them to uncover their true nature.   The probable detection of a small population of young stars illustrates once again that it is crucial to obtain deep and wide field follow up for all of the newly detected MW satellites.  Every new object has a surprise or two in store upon closer inspection."
    },
    "0911/0911.5164_arXiv.txt": {
        "abstract": "New astrophysical instruments such as skA (square kilometer Array) and IXO (formerly Constellation X) promise the discovery of tens of thousands of new isolated rotating neutron stars (pulsars), neutron stars in low-mass X-ray binaries (LMXBs), anomalous X-ray pulsars (AXPs), and soft gamma repeaters (SGRs).  Many of these neutron stars will experience dramatic density changes over their active lifetimes, driven by either stellar spin-up or spin-down, which may trigger phase transitions in their dense baryonic cores. More than that, accretion of matter onto neutron stars in LMXBs is believed to cause pycno-nuclear fusion reactions in the inner crusts of neutron stars. The associated reaction rates may be drastically altered if strange quark matter would be absolutely stable. This paper outlines the investigative steps that need to be performed in order to explore the thermal response of neutron stars to rotationally-driven phase transitions in their cores as well as to nuclear burning scenarios in their crusts. Such research complements the exploration of the phase diagram of dense baryonic matter through particle collider experiments, as performed at RHIC in the USA and as planned at the future Facility for Antiproton and Ion Research (FAIR) in Darmstadt, Germany. ",
        "introduction": "On the Earth, particle collider experiments enable physicists to cast a brief glance at the properties of ultra-dense and hot baryonic matter (Fig. \\ref{fig:FAIR}). On the other hand, it is estimated that galaxies like our Milky Way contain between $10^8$ and $10^{10}$ collapsed stars known as neutron stars, which harbor ultra-dense matter permanently in their cores. This key feature together with the unprecedented progress in observational astrophysics, which is expected to be excelled by future observatories such as the square kilometer Array (skA) and Constellation-X, make neutron stars superb astrophysical laboratories for a wide range of physical studies. These studies concern nuclear fusion processes on the stellar surface, pycnonuclear reactions in electron degenerate matter at sub-nuclear densities, and the possible formation of boson condensates and other novel states of baryonic matter--like color superconducting quark matter--at super-nuclear densities.  (For overviews, see, for instance \\cite{glen97:book,weber99:book,blaschke01:trento,lattimer01:a,rajagopal01:a,% weber05:a,klahn06:a_short,page06:a,sedrakian07:a,alford08:a}.) More than that, there is the very intriguing theoretical suggestion that strange quark matter could be more stable than atomic nuclei, known as the strange quark matter hypothesis, in which case neutron stars should be largely composed of pure strange quark matter \\cite{alcock86:a,alcock88:a,madsen98:b}. If quark matter exists in neutron stars it ought to be a color superconductor \\cite{rajagopal01:a,alford01:a,alford98:a,rapp98+99:a}. Other testable implications of the strange quark matter hypothesis concern the \\begin{figure}[tb] \\begin{center} \\includegraphics*[width=0.60\\textwidth,angle=0]{FAIR_FWeber3.eps} \\caption[]{Schematic phase diagram of strongly interacting matter \\cite{FAIR09:phasediagram}. The dashed circle indicates the portion of the phase diagram that can be explored by studying either hot and newly formed proto neutron stars (arrow pointing downwards), or cold rotating neutron stars (vertical double-headed arrow) such as millisecond pulsars and neutron stars in low-mass X-ray binaries (LMXBs).} \\label{fig:FAIR} \\end{center} \\end{figure} possible existence of a new class of white dwarfs, known as strange white dwarfs \\cite{glen92:crust,glen94:a}, and the drastic alteration of heavy-ion reaction rates in the deep crustal layers of neutron stars \\cite{golf09:a}, if strange quark matter nuggets should be present in these layers. ",
        "conclusions": ""
    },
    "0911/0911.2328_arXiv.txt": {
        "abstract": "Gamma-ray bursts (GRBs) are the brightest events in the universe. They have been used in the last five years to  study  the cosmic chemical evolution, from the local universe to the first stars.  The sample  size is still relatively small when compared to field galaxy  surveys. However, GRBs show a universe that is surprising. At $z>2$, the cold interstellar medium in galaxies is chemically evolved, with a mean metallicity of about 1/10 solar. At lower redshift ($z<1$), metallicities of the ionized gas are relatively low, on average 1/6 solar. Not only is there no evidence of redshift evolution in the interval $0<z<6.3$, but also the dispersion in the $\\sim30$ objects is large. This suggests that the metallicity of host galaxies is not the physical quantity triggering GRB events. From the investigation of other galaxy parameters, it emerges that active star-formation might be a stronger requirement to produce a GRB.  Several recent striking results strongly support the idea that GRB studies open a new view on our understanding of galaxy formation and evolution, back to the very primordial universe at $z\\sim8$. ",
        "introduction": "During the last decade, the chemical evolution of the universe has been investigated using a new class of objects: gamma-ray bursts (GRBs). GRBs are the brightest sources in the universe, but were first detected only in 1967 by a US military satellite (Klebesadel et al.\\ 1973), because their emission does not last long. For this reason, their cosmological origin was demonstrated only in 1997, when the first redshift was measured (Metzger et al.\\ 1998). Today, after more than twelve years, the number of events with spectroscopic redshift is still relatively low, about 200. Nevertheless, on April 23 2009 the highest spectroscopic redshift ever was measured, and this happened to be a GRB, GRB~090423, at $z=8.2$ (\\cite[Salvaterra et al.\\ 2009]{salvaterra09}; Tanvir et al.\\ 2009). This is not only a very exciting success for GRB science, but it also demonstrates that the field is still potentially and effectively crucial for the understanding of our universe.  GRBs are very luminous, but do not shine for very long (it cannot be any different, otherwise we would not be here to tell). Their $\\gamma$-ray emission lasts at most a few minutes, during which they radiate the same energy emitted by the Sun over its entire life, 10 Gyr. Long-duration GRBs (more than a few seconds, the majority of those detected) originate from the final core collapse of a massive star, a supernova (Woosley, 1993). Short-duration GRBs (shorter than a few seconds) have  likely a different progenitor (Katz  \\& Canel 1996): the coalescence of two compact objects (neutron stars or black holes). In both classes, rotation is the key ingredient producing the collimated emission. The GRB rate is of the order of one event every $10^5$ years in a galaxy. This means that, integrating over the entire universe and considering the collimated emission, few events are detectable from $\\gamma$-ray satellites.  During the last two years, a number of particularly interesting discoveries have shown that the universe probed by GRBs is surprisingly exciting. Apart from the already mentioned GRB~090423, in March 2008 the brightest source ever was recorded. This was GRB~080319 at $z=1.9$ (7.5 Gyrs after the Big Bang), nicknamed the ``naked eye''  GRB because it had an optical magnitude $m=5.6$ at its maximum (Bloom et al.\\ 2009). In September 2008, the at-the-time second most distant object ever was detected, GRB~080913B at $z=6.7$ (Greiner et al.\\ 2009).  Thanks to this rich phenomenology and the large redshift range spanned, there is no doubt that GRBs are very effectively probing, among other things, the chemical evolution of the universe, all the way from the local universe to the epoch of first stars, more than 13 Gyr ago. In this paper, we will summarize the results obtained in the last five years.   \\begin{figure}[b] \\begin{center} \\includegraphics[width=4.5in]{fig1.eps} \\caption{Fraction of GRB-DLAs (filled histograms) and QSO-DLAs (empty histograms) per HI and ZnII column-density bin (left- and right-hand side panels, respectively). The QSO-DLA histograms are complete for $\\log N_{\\rm HI} > 20.2$ and $\\log N_{\\rm ZnII} > 12.4$. The completeness level for GRB-DLAs is not well determined. It is apparent that column densities in GRB-DLAs are generally higher than in QSO-DLAs. This can either indicate that GRB-DLAs originate in bigger galaxies, or that the volume density of the gas is higher (e.g., the GRB sightline is crossing a region closer to the galaxy center) than QSO-DLAs, or both.} \\label{f1} \\end{center} \\end{figure}  \\begin{figure}[b] \\begin{center} \\includegraphics[width=4in]{fig2.ps} \\caption{Redshift evolution of the metallicity relative to solar values, for 17 GRB-DLAs at $z>2$, 16 GRB hosts at $z<1$  and $\\sim250$ QSO-DLAs in the interval $0<z<4.4$. Error bars are not available for all GRB-DLAs. Errors for GRB hosts are not estimated. Errors for QSO-DLAs are generally smaller than 0.2 dex. The dashed line is the best-fit linear correlation for QSO-DLAs. The solid line is the mean metallicity predicted by semi-analytic models for galaxy formation (Somerville et al.\\ 2001). The GRB-DLAs metallicity in $2<z<4.5$ is on average 2.5 times higher than the average value in QSO-DLAs in the same redshift interval.} \\label{f2} \\end{center} \\end{figure} ",
        "conclusions": "It is well known that GRBs shine through a universe that is hard to see in other ways. They are incredibly bright and last only a short time. These two properties make them very different from QSOs, which are not as bright, and do not fade away, making the investigation of nearby galaxies much more complicated. From luminous GRB afterglows, it is possible to measure the redshift and localize faint galaxies.  Most GRBs are associated with the death of a massive star, thus with a star-forming region. It is well known that the SFR of the universe was much higher in the past than it is today (Hopkins \\& Beacom 2006), therefore GRBs might be the most efficient way of identifying the evolution of the SFR density. We also know that the SFRD is dominated by small star-forming galaxies, which are probably the most common galaxies in the distant universe (Pozzetti et al.\\ 2009).  Identification of distant galaxies with GRBs is affected by a different bias than traditional galaxy surveys, because GRBs are not detected through optical instruments, but with $\\gamma$-ray and X-ray satellites. Some of these galaxies can be faint, because dust extinguished, or because too far, or because intrinsically faint. With GRBs it is possible to explore extreme regimes of galaxy parameters, thus they are important to understand galaxy formation and evolution. GRB hosts identified in the optical and NIR at  $z<2$ are generally small (on average the stellar mass of the Large Magellanic Cloud) and star-forming galaxies, although SFRs span a large interval.  GRBs are probes of the state of the chemical enrichment of the universe, from the local universe, back to the time of the formation of first stars. Metallicities of the cold ISM in host galaxies at $z > 2$ is not low. The measured average value is 1/10 solar, and the dispersion is large, about a factor of five. Relatively high metallicity is confirmed also for the highest redshift detections (Savaglio 2006; Totani et al.\\ 2006; Price et al.\\ 2007), which means that there is no indication of redshift evolution. Low metallicities of the GRB progenitor are theoretically predicted (no mass loss) in order to keep a high angular momentum, and have a highly collimated jet.  Relatively high metallicities, in this case from ionized gas of the host galaxies, are confirmed also at $z<1$. The average value is 1/6 solar, with a dispersion of a factor of two, indicating that the metal content in host galaxies is not evolving so fast.  The sample is not very large and systematic uncertainties are still not totally under control, therefore more observations and detections are very important. Nevertheless, the lack of evidence of redshift evolution and the observed large dispersion suggest that GRBs do not happen necessarily in metal poor galaxies. Star formation, on the other hand, might be a more important physical trigger.  GRBs are extremely important for our understanding of the primordial universe and the formation and evolution of heavy elements. The enlightening discoveries of the last few years are a clear indication that the investigation is affected by our technical capabilities which have dramatically improved recently. Dedicated instruments and observational programs have opened a new window in the hidden universe and show that this is more surprising and fascinating than expected."
    },
    "0911/0911.2434_arXiv.txt": {
        "abstract": "The observed super-massive black hole (SMBH) mass -- galaxy velocity dispersion ($M_{\\rm bh} - \\sigma$) correlation may be established when winds/outflows from the SMBH drive gas out of the potential wells of classical bulges. Here we present numerical simulations of this process in a static isothermal potential. Simple spherically symmetric models of SMBH feedback at the Eddington luminosity can successfully explain the $M_{\\rm bh} - \\sigma$ and nuclear cluster mass $M_{\\rm NC}-\\sigma$ correlations, as well as why larger bulges host SMBHs while smaller ones host nuclear star clusters. However these models do not specify how SMBHs feed on infalling gas whilst simultaneously producing feedback that drives gas out of the galaxy.   More complex models with rotation and/or anisotropic feedback allow SMBHs to feed via a disc or regions not exposed to SMBH winds, but in these more realistic cases it is not clear why a robust $M_{\\rm bh} - \\sigma$ relation should be established.  In fact, some of the model predictions contradict observations. For example, an isotropic SMBH wind impacting on a disc (rather than a shell) of aspect ratio $H/R \\ll 1$ requires the SMBH mass to be larger by a factor $\\sim R/H$, which is opposite to what is observed. We conclude that understanding how a SMBH feeds is as important a piece of the puzzle as understanding how its feedback affects its host galaxy.  Finally, we note that in aspherical cases the SMBH outflows induce differential motions in the bulge. This may pump turbulence that is known to hinder star formation in star forming regions. SMBH feedback thus may not only drive gas out of the bulge but also reduce the fraction of gas turned into stars. ",
        "introduction": "It is believed that the centres of most galaxies contain super-massive black holes (SMBHs) whose mass $\\mbh$ correlates with the velocity dispersion $\\sigma$ of the host galaxy \\citep{Ferrarese00,Gebhardt00,Tremaine02}. Similarly, there is a correlation between $\\mbh$ and the mass of the bulge, $M_{\\rm bulge}$ for large SMBH masses \\citep{Magorrian98,Haering04,GultekinEtal09}. Observations also suggest that the masses of nuclear star clusters (NC) ($10^5 \\msun \\simlt M_{\\rm NC} \\simlt 10^8 \\msun$) correlate with the properties of their host dwarf ellipticals \\citep{FerrareseEtal06,Wehner06} in a manner that is analogous to the one between SMBHs and their host ellipticals.  These empirical relations can be explained in a very natural way if the growth of host galaxies and their central SMBHs or NCs are linked by feedback. This was first pointed out by \\cite{SilkRees98}, who highlighted the potential importance of SMBH heating and outflows before any robust observational evidence for such a link had been found. Subsequently, \\cite{Fabian99} argued that radiation pressure acting on cold gaseous clouds in the bulge could give rise to the observed correlations. \\cite{KP03} argued for the existence sub-relativistic outflows from the very central regions of AGN, which prompted \\cite{King03,King05} to study the motion of a shell of gas swept up by a wind/outflow from a central black hole in a galactic isothermal dark matter potential. \\cite{King05} demonstrated that the shell will be expelled from the potential provided the black hole mass exceeds a critical value that, as a function of $\\sigma$, turns out to be close to the observed $M_{\\rm BH}$--$\\sigma$ relation.  The main result of \\cite{King03,King05} can be deduced using a simpler order of magnitude ``weight argument''. According to this argument, the SMBH luminosity is assumed to be limited by the Eddington value. Radiation pressure drives a wind. \\cite{KP03} argue that wind velocity is comparable to the escape velocity from the inner accretion disc, e.g., $v \\sim 0.1 c$. The momentum outflow rate is assumed to be \\begin{equation} \\dot P_{\\rm SMBH} \\approx {L_{\\rm Edd}\\over c} = \\frac{ 4 \\pi G M_{\\rm BH}}{\\kappa}\\;; \\label{pismbh} \\end{equation} \\noindent here $\\kappa$ is the electron scattering opacity and $M_{\\rm BH}$ is the SMBH mass. This result is natural to the order of magnitude: $L_{\\rm Edd}/ c$ is the radiation momentum flux which is presumably passed to the wind as radiation accelerates the outflow \\citep[cf.][]{KP03}.  The argument also assumes that the black hole wind is optically thin at the point of interaction with the ambient gas. Because the cooling time of the shocked gas is short on scales appropriate for observed bulges \\citep{King03,King05}, the bulk energy of the outflow is thermalised and quickly radiated away. It is then only the momentum push (equation \\ref{pismbh}) of the outflow on the ambient gas that is important since it is this that produces the outward force on the gas\\citep[as in the earlier model by][]{Fabian99}. The weight of the gas is $W(R) = GM(R)[M_{\\rm total}(R)]/ R^2$, where $M_{\\rm gas}(R)$ is the enclosed gas mass at radius $R$ and $M_{\\rm total}(R)$ is the total enclosed mass including dark matter. For an isothermal potential, $M_{\\rm gas}(R)$ and $M_{\\rm total}(R)$ are proportional to $R$, so the result is \\begin{equation} W = {4f_g\\sigma^4\\over G}\\;. \\label{w} \\end{equation} Here $f_g$ is the baryonic fraction and $\\sigma^2 = GM_{\\rm total}(R)/2R$ is the velocity dispersion in the bulge. By requiring that momentum output produced by the black hole just balances the weight of the gas, it follows that \\begin{equation} M_{\\sigma} = {f_g\\kappa\\over \\pi G^2}\\sigma^4, \\label{msigma} \\end{equation} \\noindent which is consistent with the observed the $M_{\\rm BH}$--$\\sigma$ relation.  To order of magnitude, relation \\ref{w} should hold for any potential at the virial radius. Therefore the model by \\cite{King03,King05} appears to be a promising explanation of the observed correlations between SMBHs and their host galaxies.  However, additional complications, not considered in \\cite{King03,King05}, may be important. Amongst these are (1) the self-gravity of the gas, because gas that accumulates at the centre of the potential may begin to dominate it and therefore alter the result; (2) finite angular momentum of the gas may lead to formation of a disc; (3) collimated and variable outflows. \\\\  The goal of this paper is to verify that the predictions of the analytical model of \\cite{King03,King05} hold and to explore more realistic settings, of the kind just described. Note that we concentrate on the early stages of galaxy evolution, when we would expect galaxies to be gas-rich and the baryonic mass within dark matter haloes is dominated by gas rather than stars. At later times, when the SMBH is likely to be less luminous and the galaxy is less gas-rich, we would expect different forms of feedback to become important, such as relativistic jets and the associated radio bubbles \\citep[e.g.,][]{ChurazovEtal02} or pre-heating due to the inverse Compton effect \\citep[e.g.,][]{SazonovEtal04}. However, we do not include these forms of feedback in our current simulations. ",
        "conclusions": "Our main conclusions are \\begin{itemize}  \\item Predictions of spherically symmetric models with black hole feedback tied to Eddington limit luminosity as in the models of \\cite{King03,King05} are confirmed numerically. Self-gravity of the gas complicates evolution of the system, and a self-consistent treatment of star formation and its feedback is necessary for further progress.  \\item As suggested earlier based on analytical arguments, it is the dynamical time in the bulge, $R_{\\rm b}/\\sigma$, that determines whether or not the SMBH can reach its limiting $M_\\sigma$ value. Central regions of smaller bulges, where the dynamical time is observed to be shorter than the Salpeter time, could be smothered with infalling gas despite the SMBH feedback. This process may be the origin of the nuclear star clusters in ``smaller'' galaxies.  \\item A net angular momentum in the shell is essential in determining the fate of the shell and the SMBH feeding. There is no well defined $M_\\sigma$ mass in that case since the momentum thrust required is different in different directions. Gas near the symmetry axis is blown out easier than in the spherically symmetric case, whereas gas settled into a disc requires $\\sim R/H$ more thrust than in the latter case.  \\item If the SMBH is fed through a several kpc scale disc, the SMBH mass would have to be $\\sim R/H$ larger than in the spherical case to expel the feedback-resistant self-shielding disc. This directly contradicts the recent observations of \\cite{Hu09} that show that pseudo-bulges have lighter SMBHs than their classical counter-parts. This may imply that black holes are not fed by large (kpc or larger) discs but rather by flows with a small specific angular momentum.  \\item We also noted that SMBH outflows in a realistic, i.e., aspherical, situation may pump turbulence (differential motions) in the bulge. Turbulence is known to hinder star formation. Therefore we believe that AGN feedback not only expels the gas from the galaxy but it may also reduce the amount of mass turned into stars during bulge formation.  \\end{itemize}"
    },
    "0911/0911.3118_arXiv.txt": {
        "abstract": "We propose a mechanism for the creation of cosmic string loops with dynamically stabilised windings in the internal space. Assuming a velocity correlations regime in the post-inflationary epoch, such windings are seen to arise naturally in string networks prior to loop formation. The angular momentum of the string in the compact space may then be sufficient to ensure that the windings remain stable \\emph{after} the loop chops off from the network, even if the internal manifold is simply connected. For concreteness we embed our model in the Klebanov-Strassler geometry, which provides a natural mechanism for brane inflation, as well a being one of the best understood compactification schemes in type IIB string theory. We see that the interaction of angular momentum with the string tension causes the loop to oscillate between phases of expansion and contraction. This, in principle, should give rise to a distinct gravitational wave signature, the future detection of which could provide indirect evidence for the existence of extra dimensions. ",
        "introduction": "The existence of string loops with dynamically stabilised winding in a compact space was first demonstrated by Iglesias and Blanco-Pillado \\cite{BlancoPillado:2005dx}, using the Klebanov-Strassler geometry \\cite{Klebanov:2000hb}. They considered strings at the tip of the throat, with geodesic wrappings in the $S^3$ which regularises the conifold singularity. Although they derived a lower bound for the angular momentum of a loop with a given number of windings - below which the windings became unstable - the result must remain of purely theoretical interest to cosmology as long as a specific mechanism for winding formation (and hence for the formation of the string angular momentum) is not considered. The purpose of this paper is to propose such a mechanism, which leads to the formation of string configurations such as those investigated in \\cite{BlancoPillado:2005dx} and to investigate the resulting string dynamics with specific reference to their cosmologically observable consequences.  \\n Assuming that a velocity correlations regime in the post-inflationary epoch as in \\cite{Avgoustidis:2005vm} leads naturally to the formation of geodesic windings we show that the winding number ($n$), total energy ($E$) and angular momentum ($l$) of the string are specified precisely by the model parameters. That is, by the parameters which define the Klebanov-Strassler geometry (in this case specifically by the value of the warp factor $a_0$) and those which determine the scale of the string network ($\\alpha$, $t_i$). Substituting for $l$ and $n$ in the bound referred to above then demonstrates the stability of these windings, at least under the assumption that $l$ remains approximately constant over small time scales after the moment of initial loop formation. By assuming also that the total energy of the string remains approximately constant (i.e. by neglecting the loss of $E$ and $l$ via gravitational wave emission over small time scales), we then determine the equation of motion for the four dimensional string radius $r(t)$, and solve to find a (generically) oscillating solution. Crucially we observe that the qualitative behaviour of the loop depends on the value of the warp factor $a_0$ with $a_0^2<1/2$ leading to an initial phase of expansion and $a_0^2>1/2$ leading to an initial phase of contraction. The fixed point solution $a_0^2 = 1/2$ is a static, non-oscillatory solution. In both oscillatory modes (initially contracting \\emph{and} initially expanding) we find that the period of the oscillation is inversely proportional to $a_0^2$, and proportional to the initial size of the loop $(\\alpha t_i)$. Following \\cite{Avgoustidis:2005vm} we refer to these objects as non-topological cycloops \\footnote{The term 'cycloop' was first coined in \\cite{Avgoustidis:2005vm} to refer to cosmic string loops with \\emph{smooth} windings wrapping cycles in the internal space - as opposed, for example, to non-smooth, step-like windings which give rise to necklace configurations from a four dimensional perspective \\cite{Matsuda:2006ju, Matsuda:2005ez, Lake:2009nq}. However in their original conception Avgoustidis and Shellard used it only to refer to windings which are topologically trapped. We therefore propose the term 'non-topological cycloops' to refer to string loops with dynamically stabilised smooth windings (in this case geodesics) around a simply connected compact manifold.}.  \\n The layout of the paper is then as follows: In Section 2 we briefly review the relevant Klebanov-Strassler background, focusing on the geometry of the conifold tip. In Section 3 we show how the assumption of a velocity correlations regime yields a dynamical model of winding formation after the end of inflation. Section 4 recaps the generic results of \\cite{BlancoPillado:2005dx} which we then combine with the results of the previous section to give explicit expressions for the winding number, energy and angular momentum of the string. The equation of motion for the loop radius is then derived and solved in Section 5 and a brief summary of the main results together with a discussion of their cosmological implications and possible consequences for experimental observations is presented in section 6. Finally the two appendices at the end of this work deal with issues which arise within the text: Appendix I outlines the method of Eulerian substitution of the third kind which is used to integrate the differential equation involving $\\dot{r}(t)$ and $r(t)$ derived in Section 5. Appendix II gives a detailed description of the Hopf fibration of the 3-sphere, which is introduced briefly in section 2 and used throughout the following analysis. ",
        "conclusions": "We have argued that a velocity correlations regime in the post-inflationary epoch leads naturally to the formation of string loops with geodesic windings in the compact space.  For strings at the tip of the conifold throat of the Klebanov-Strassler geometry we were able to show that the quantities which determine the dynamical evolution of the loop (i.e. the initial winding number $n(t_i)$, energy $E(t_i)$ and string velocity/angular momentum in the compact space, $\\dot{s}(t_i)$/ $l(t_i)$) are \\emph{uniquely} determined by the parameters $a_0^2$, $\\alpha$ and $t_i$. Crucially these windings were found to have sufficient angular momentum in the compact directions to remain stable, even after the string chops off from the network to form a loop.  \\n The interaction between the tension and the angular momentum in the compact space was found to play a significant role in the dynamical evolution of the string, including - perhaps surprisingly - the evolution in four dimensions. By assuming energy loss (and angular momentum) via gravitational radiation to be negligible, we determined equations of motion for the four-dimensional radius $r(t)$ and the string velocity $v(t) \\sim \\dot{s}(t)R$, which we believe to be valid over small time scales after the moment of loop formation, $t=t_i$.  \\n We found that the qualitative behaviour of the string depends crucially on the square of the warp factor, $0<a_0^2<1$, with $a_0^2<1/2$ leading to an oscillatory solution characterised by an initial phase of expansion, whilst $a_0^2>1/2$ leads to oscillations with an initial contracting phase. In each case the string was seen to oscillate between it's initial radius $r(t_i)=(\\alpha t_i)$ and a secondary critical value defined by $ r_{c2}=\\frac{(1-a_0^2)}{a_0^2}(\\alpha t_i)$, with the period of oscillation given by $T = \\frac{(\\alpha t_i)} {a_0^2}$.  \\n As noted above, in the two oscillatory regimes it is the interaction of the angular momentum with the string tension which \"drives\" the dynamical evolution converting kinetic energy into rest mass during expansion (with the inverse process occurring in the contracting phase). The string is seen to evolve \\emph{towards} a static, minimum energy configuration where $l=2\\pi a_0^2 T_1$ and $\\dot{s}^2 = a_0^2/R^2$, but is unable to satisfy these two conditions simultaneously at a point where $\\dot{r}(t)=0$. By contrast, in the $a_0^2=1/2$ scenario we find that the energy minimising conditions \\emph{are} satisfied simultaneously at the moment of loop formation. In this case the tension exactly offsets the effect of angular momentum so that the string remains static at it's original radius $r(t_i)=(\\alpha t_i)$ and is non-oscillatory.  \\n The meaning of the term \"small\" above is somewhat ambiguous, but it seems reasonable to assume that our solution will provide a valid approximation over at least one full oscillatory cycle of the loop; that is, over a time period $\\Delta t \\sim T = \\frac{(\\alpha t_i)}{a_0^2}$. For time periods $\\Delta t >> T$ our initial analysis must be extended to include the effects of gravitational wave emission on $E(t)$ and $l(t)$ (or equivalently on $r(t)$ and $\\dot{s}(t)$). We may expect the qualitative effect of energy and angular momentum loss to be relatively simple, as long as the string retains sufficient angular momentum for the extra-dimensional windings to remain stable. In this case it is likely that the loss of $E$ and $l$ due to gravitational wave emission will simply act to damp the oscillations of $r(t)$ and $\\dot{s}(t)$. What is unclear however is whether or not the damping coefficient itself is likely to be time-dependent; For example, it is possible that smaller oscillations lead to greater rates of emission per unit length (as the rate of acceleration $\\ddot{r}(t)$ is higher in this case) so that the damping itself increases with time (for an individual loop).  \\n In the case that the windings eventually become unstable however (as $l(t)$ drops below the threshold for ensuring their stability) the string dynamics are likely to become complicated, and it is not clear whether the process of winding contraction (i.e. of windings \"falling off\" the $S^3$) may even be accommodated within an analysis which uses an ansatz of the form (\\ref{eq:ansatz}) to describe the string configuration. This is because the coordinates $r(t)sin\\sigma$, $r(t)cos\\sigma$, $\\psi(\\sigma,t)$, $\\theta(\\sigma,t)$ and $\\phi(\\sigma,t)$ are treated as \\emph{independent} variables with respect to the determination of the equations of motion. We are therefore unable to take account of the continuously connected nature of the string when string sections \"move\" from one direction to another (e.g. in \"falling off\" the $S^3$ to form part of the four dimensional rest mass of the loop).  \\n The present analysis could of course still be improved by accounting for the effects of gravitational wave emission under the assumption that that loops retain their windings, which would indeed be valid up to the point where $l(t)$ drops below the critical value (which may also be calculated). Such an improved analysis could proceed as follows: one could compute the stress energy tensor for an oscillating loop and look for solutions with this as a source to the Einstein equations. This should allow us to estimate the rate of loss of $E$ and $l$ via gravitational wave emission and, as stated, we expect this to produce a damping term in the equations for $r(t)$ and $\\dot{s}(t)$.  \\n Calculating the emission spectrum for an oscillating loop would also be of immense practical interest. The gravitational wave signature of such a loop, whose self-oscillation is caused by the presence of angular momentum in the compact space, may differ significantly from that of a string whose oscillations, though superficially similar, do not result from self-interaction. Gravitational wave emission from loops oscillating with period $\\omega \\sim L^{-1}$ (where $L$ is the loop size) have been intensively studied in four dimensions \\cite{Vaschpati+Vilenkin:1985, Burden:1985, Garfinkle+Vaschpati:1987, Allen+Shellard:1992}. However in such cases the loop is not undergoing \\emph{genuine} phases of expansion and contraction, but rather experiencing \"wiggles\" of a size comparable to it's own length. Although Weinberg \\cite{Weinberg:1972} has has shown that - in an FRW universe - the power of a weak, isolated, periodic source (to lowest order in $G$) may be given by a single formula \\emph{regardless} of the exact nature of the source, it is not immediately clear that this should hold in extra-dimensional scenarios. Additionally such sources (i.e. loops) have no angular momentum to shed in the process of emission. In fact even if we were to study loops whose self-oscillation was due to their \"rotational\" motion in Minkowski space (c.f. \\cite{Durrer:1989, Vilenkin+Shellard:2000}), the angular momentum which would be carried away by gravitational radiation would be very different from that lost via emission from oscillating cycloops \\footnote{We note also that one possible further extension of the current analysis would be to consider a wound string with rotational motion in both the compact \\emph{and} Minkowski directions.}. The detection of such a signature by the future generation gravitational wave detectors could then provide indirect evidence of the presence of compact dimensions. We hope therefore to be able to provide such an analysis in the near future.  \\n Although the full string theory compactification favours exponentially small values of $a_0^2$, it is also interesting to note that the possibility of obtaining evidence from the gravitational wave signature of wound strings exists even in the case of static loops (i.e. the $a_0^2 = 1/2$ case). These loops, which form automatically with a configuration which meets the energy minimising conditions, look just like ordinary string loops from a four-dimensional perspective. However they contain an \"unseen\" angular momentum which is not directly manifested in their dynamical evolution in Minkowski space. As stated above, we expect that even \"standard\" four dimensional string loops may undergo periodic oscillations with $\\omega \\sim L^{-1}$ creating ripples in space-time and giving rise to gravitational waves. Such fluctuations \\emph{along} the length of the string - though not in the total four dimensional length itself - would likely still occur in this case (though it is possible that the existence of the angular momentum term may lend a certain 'rigidity' to the circular string configuration, making it more resistant to deformation) but the string must also now shed it's angular momentum. The gravitational wave signature of even a non-self-oscillating loop in the extra-dimensional scenario is therefore also likely to differ significantly from the standard case of an un-wound string, though perhaps to a less significant degree than in the self-oscillating case. However the possibility of the indirect detection of compact dimensions from cosmic strings therefore remains, even in the absence of self-oscillating loops. We hope also to be able to offer an analysis of this interesting possibility in a future letter.  \\n It remains for us to outline some of the possible limitations of the analysis we have attempted: We have assumed throughout the present work that the velocity correlations regime leads naturally to a) geodesic windings and b) movement of the string parallel to itself (i.e. along the geodesics). However these assumptions may be questioned. The rationale for adopting such an approach (which simplified the resulting analysis considerably) was simple - we expect velocity correlations to impart a constant angular momentum density to each point along the string. Therefore we expect each \"point\" (or infinitesimal string segment) to travel along a geodesic curve, both before and after loop formation. However it has not been proved that each point along the string traversing a separate geodesic, leads to a winding configuration that is itself geodesic. Likewise it does not necessarily follow that the resulting motion of the string as a whole is parallel to itself. In general, we may expect that the geometry of the internal space plays a role in determining the exact nature of the winding configuration and the resulting string motion. Moreover we have not included the contribution from the Ramond-Ramond (RR) sector, which in principle will couple to the string. This charge term is ultimately crucial for distinguishing between cosmic strings and cosmic super-strings.  \\n Whilst, in principle, motion of the string perpendicular to its length may easily be accounted for (see end of Appendix II) the most significant problem in the analysis of string loops with non-geodesic windings arises from the resulting $\\sigma$-dependence of the expressions for $E$ and $l$. We believe that the present analysis may therefore be improved by a more thorough investigation of the winding process itself, and though we expect our expressions for $n(t_i)$ (\\ref{eq:n2}) and $v(t_i) \\sim \\dot{s}(t_i)R$ (\\ref{eq:omega_l_v}) to remain valid (in the KS case), the appropriate string ansatz may well be more complicated. Although it seems unlikely that the qualitative behaviour of the string will differ significantly than described in the scenarios above, finding analogous (and \\emph{quantitatively} different) results may be extremely difficult. But ultimately exact \\emph{quantitative} predictions will be needed for any comparison with future experimental data, and we believe such a project to be worthwhile. \\begin{center} {\\bf Acknowledgments} \\end{center} M.L is supported by an STFC studentship and the Queen Mary EPSTAR consortium. This work was supported in part by NSERC of Canada. \\renewcommand{\\theequation}{A-\\arabic{equation}} \\setcounter{equation}{0}  %"
    },
    "0911/0911.0722.txt": {
        "abstract": "The Extreme ultraviolet SpectroPhotometer (ESP) is one of five channels of the Extreme ultraviolet Variability Experiment (EVE) onboard the NASA {\\it Solar Dynamics Observatory} (SDO). The ESP channel design is based on a highly stable diffraction transmission grating and is an advanced version of the Solar Extreme ultraviolet Monitor (SEM), which has been successfully observing solar irradiance onboard the {\\it Solar and Heliospheric Observatory} (SOHO) since December 1995. ESP is designed to measure solar Extreme UltraViolet (EUV) irradiance in four first order bands of the diffraction grating centered around 19 nm, 25 nm, 30 nm, and 36 nm, and in a soft X-ray band from 0.1 to 7.0~nm in the zeroth order of the grating. Each band's detector system converts the photo-current into a count rate (frequency). The count rates are integrated over 0.25~sec increments and transmitted to the EVE Science and Operations Center for data processing. An algorithm for converting the measured count rates into solar irradiance and the ESP calibration parameters are described. The ESP pre-flight calibration was performed at the Synchrotron Ultraviolet Radiation Facility of the National Institute of Standards and Technology. Calibration parameters were used to calculate absolute solar irradiance from the Sounding Rocket flight measurements on 14 April 2008. These irradiances for the ESP bands closely match the irradiance determined for two other EUV channels flown simultaneously, EVE's Multiple Euv Grating Spectrograph (MEGS) and SOHO's Charge, Element and Isotope Analysis System / Solar EUV Monitor (CELIAS/SEM). ",
        "introduction": "\\label{S-Introduction} The {\\it Solar Dynamics Observatory} (SDO) is the first NASA Living With a Star mission with its launch planned for February 2010. It will provide accurate measurements of the solar atmosphere characteristics with high spatial and temporal resolution at many wavelengths simultaneously. These measurements will help us understand the solar activity cycle, the dynamics of energy transport from magnetic fields to the solar atmosphere, and the influences of this energy transport on the Earth's atmosphere and the heliosphere. Dynamic changes of the solar radiation in the extreme ultraviolet (EUV) and X-ray regions of the solar spectrum are efficient drivers of disturbances in the Earth's space weather environment. The Extreme ultraviolet Variability Experiment (EVE) is one of three instrument suites on SDO. EVE measures the solar EUV irradiance with unprecedented spectral resolution, temporal cadence, accuracy, and precision. Furthermore, the EVE program will incorporate physics-based models of the solar EUV irradiance to advance the understanding of solar dynamics based on short- and long-term activity of the solar magnetic features \\cite{Woods06, Woods09}.  ESP is one of five channels \\cite{Woods06} in the EVE suite. It is an advanced version of the SOHO/CELIAS Solar Extreme ultraviolet Monitor (SEM) \\cite{Hovestadt95, Judge98}. SEM measures EUV solar irradiance in the zeroth diffraction order (0.1 to 50.0~nm bandpass) and two (plus and minus) first order diffraction bands (26 to 34~nm bandpasses) centered at the strong He II 30.4~nm spectral line. More than 13 years of EUV measurements have shown that the SEM is a highly accurate and stable EUV spectrometer \\cite{Judge08} that has suffered only minor degradation, mainly related to deposition of carbon on the SEM aluminum filters.  The ESP design is based on a highly stable diffraction transmitting grating \\cite{Schattenburg90, Scime95}, very similar to the one used in SEM. ESP has filters, photodiodes, and electronics with characteristics (transmission, sensitivity, shunt resistance, {\\it etc.}) that are susceptible to some change over the course of a long mission. Because of such possible degradation ESP shall be periodically calibrated throughout the mission. The calibration program includes a pre-flight calibration followed by a number of sounding rocket under-flights with a nominally identical prototype of the flight instrument that is typically calibrated shortly before and shortly after each under-flight. Calibration of the SDO/EVE and EVE/ESP soft X-ray and EUV flight instruments was performed at the National Institute of Standards and Technology (NIST) at the Synchrotron Ultraviolet Radiation Facility (SURF-III). NIST SURF-III has a number of Beamlines (BL) with different specifications and support equipment ({\\it i.e.} monochromators, translation stages and calibration standards). The beamline used for a given instrument depends on the instrument's characteristics and calibration requirements such as spectral ranges, entrance aperture, alignment, {\\it etc.} Before the flight ESP was integrated into the EVE package, it was first calibrated on SURF BL-9. BL-9 is equipped with a monochromator capable of scanning through a wide EUV spectral band which allows the instrument efficiency profile for ESP to be measured in increments of 1.0~nm over a spectral window of about 15 to 49~nm. After ESP was mounted onto EVE, the second ESP calibration was performed on SURF BL-2 \\cite{Furst93}. BL-2 illuminates the instrument with the whole synchrotron irradiance spectrum, which can be shifted in its spectral range by changing the energy of the electrons circulating in the electron storage ring producing the synchrotron radiation. The intensity of the EUV beam is controlled by the current (and therefore number) of circulating electrons. Because the EUV beam consists of a wide spectrum of photon energies, this type of calibration is a radiometric calibration. A major part of this paper is related to the results of the ESP BL-9 and BL-2 calibrations.  The scientific objectives of ESP are given in Section 2. Section 3 presents a short overview of the ESP channel. Section 4 describes an algorithm to calculate solar irradiance from count-rates measured on each of ESP's photometer bands, measurement conditions ({\\it i.e.} ESP temperature, dark current, band response, {\\it etc}.), ESP calibration data and orbit parameters. Section 5 shows results from ESP ground tests. Section 6 summarizes ESP pre-flight calibration results. A comparison of ESP solar irradiance measurements from the sounding rocket flight of 14 April 2008 with data from other EUV channels (EVE/MEGS sounding rocket instrument and SOHO/SEM) are given in Section 7. Concluding remarks are given in Section 8. ",
        "conclusions": "\\label{S-Conclusions} As part of the SDO/EVE suite of channels, ESP will provide highly stable and accurate measurements of absolute solar irradiance in five EUV wavelength bands. Its high temporal cadence, low latency, and spectral overlap with the other SDO instruments, will provide important details of the dynamics of rapidly changing irradiance related to impulsive phases of solar flares. An algorithm to convert measured count rates and other ESP data into solar irradiance was described. ESP has many significant design improvements over its SOHO/SEM predecessor, including the possibility of measuring on orbit changes of dark currents, electronics gain changes, contamination degradation and pin-holes of thin-film filters, visible light scatter, and energetic particles background. All optical components of ESP (diodes, filters, grating) were also separately tested and calibrated prior to an end-to-end calibration. The ESP was calibrated at SURF BL-9 using quasi-monochromatic wavelengths, and at BL-2 to obtain a radiometric calibration. Efficiency profiles determined during BL-9 calibration were used as reference data for calculation of irradiance in each band during BL-2 calibration. Measurements of solar irradiance from ESPR during the EVE sounding rocket flight of 14 April 2008 were compared with corresponding EVE/MEGS spectra and SOHO/SEM irradiance measurements. These comparisons show good agreement. Both the ESP flight and rocket instruments are fully calibrated and ready for solar measurements. %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% \\begin{acks} The authors thank Frank Eparvier, Mike Anfinson, Rick Kohnert, Greg Ucker, Don Woodraska, Karen Bryant, Gail Tate, Matt Triplett, and the rest of the LASP EVE Team at the University of Colorado at Boulder for their many contributions during ground tests of the ESP. We also thank Don McMullin of the Space Systems Research Corporation for his discussions and long-term support of this mission, and the  NIST SURF Team: Rob Vest, Mitch Furst, Alex Farrell, and Charles Tarrio for support of ESP calibration at BL-9 and BL-2 and diffraction grating transmission measurements at the SURF reflectometer facility. This work was supported by the University of Colorado award 153-5979. \\end{acks}   %%% BIBLIOGRAPHY %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%"
    },
    "0911/0911.4448_arXiv.txt": {
        "abstract": "A digital frequency multiplexing (DfMUX) system has been developed and used to tune large arrays of transition edge sensor (TES) bolometers read out with SQUID arrays for mm-wavelength cosmology telescopes. The DfMUX system multiplexes the input bias voltages and output currents for several bolometers on a single set of cryogenic wires. Multiplexing reduces the heat load on the camera's sub-Kelvin cryogenic detector stage. In this paper we describe the algorithms and software used to set up and optimize the operation of the bolometric camera. The algorithms are implemented on soft processors embedded within FPGA devices operating on each backend readout board. The result is a fully parallelized implementation for which the setup time is independent of the array size. ",
        "introduction": "\\label{sec-introduction} A new generation of mm-wavelength experiments use (or will use) arrays of hundreds or thousands of transition edge sensor (TES) bolometers operating at sub-Kelvin temperatures to image the cosmic microwave background (CMB) to unprecedented precision in the search for new galaxy clusters or the polarization imprint left by inflation. One key element in the cryogenic design for these experiments is the reduction of readout wiring, which is accomplished for APEX-SZ~\\cite{2008JLTP..151..697M}, the South Pole Telescope (SPT)~\\cite{2004astro.ph.11122R}, EBEX \\cite{2008SPIE.7020E..68G} and Polarbear~\\cite{2008AIPC.1040...66L} by multiplexing the bolometer sky-signals in the frequency domain.  In this paper we describe tuning algorithms and software developed for the new digital backend electronics~\\cite{2008ITNS...55...21D} in use for EBEX and POLARBEAR. This system is the evolution of the analog backend~\\cite{2005ApPhL..86k2511L} deployed on APEX-SZ and the SPT. The new system provides an order of magnitude reduction in power consumption and size, improvements in low-frequency noise performance, and a faster, more robust setup. As most of the system complexity is encapsulated in reprogrammable firmware, it is easily adapted to the requirements of specific experiments.  \\begin{figure}[htbp] \\includegraphics[width=7cm]{schematic} \\caption{Schematic diagram showing $N$ bolometers multiplexed through a single set of wires and SQUID in the Digital frequency domain multiplexer system.} \\label{fig-schematic} \\end{figure}  Figure \\ref{fig-schematic} shows a circuit with $N$-bolometers multiplexed on a single set of wires and read out through a SQUID operating in shunt feedback~\\cite{2008ITNS...55...21D, 2005ApPhL..86k2511L}. The box on the left labelled `DfMUX' (for digital frequency multiplexer') represents firmware implemented on a Field Programmable Gate Array (FPGA) that performs the signal processing for the system digitally. A digital multifrequency synthesizer (DMFS) produces a sine-wave bias carrier for each channel and sums them to form a ``comb'' in frequency space. Each bolometer in the multiplexer module is in series with an inductor-capacitor (LC) filter that selects a single bias that is uniquely positioned in frequency-space.  Intensity variations from the sky-signal change the bolometer resistance, amplitude modulating the bolometer current, and appearing as side-bands adjacent to the carriers. These currents are summed at the input of the SQUID, which operates as a transimpedance amplifier. To avoid flux-burdening the SQUID with the large bias carriers, an inverted copy of the carrier ``comb,'' called the nuller, is injected at the SQUID input from a second DMFS. The sky-signal modulated comb is digitized after the SQUID circuit and each bolometer signal is separately demodulated in the digital multifrequency demodulator (DMFD), implemented in firmware.  Each detector channel has its own demodulator, which provides only the in-phase ``I'' component. Including a second demodulator for each channel doubles the usage of FPGA fabric, increasing the power consumption, and does not provide any useful additional information. To align the demodulator phase with the carrier, accounting for phase shifts in the cryostat, each module has an extra $N^{th}+1$ demodulator that can be used to momentarily measure the out-of-phase ``Q'' component of each detector. In addition to aligning the demodulator with the carrier, this ``helper'' demodulator is employed to provide extra information in the algorithms described below.  Frequency domain multiplexing permits the reduction of cryogenic wires from two per bolometer to two per $N$ bolometers, where $N$ is the multiplexing factor. The system is currently implemented with $N=8$ and $N=16$.  Extension to much higher multiplexing factors is feasible but requires a substantial improvement in SQUID bandwidth, using methods such as those demonstrated in Ref.~\\cite{2009arXiv0901.1919L}. ",
        "conclusions": "The algorithms used to setup and tune large arrays of TES bolometers for observations of the Cosmic Microwave Background radiation with mm-wavelength experiments using a digital frequency domain SQUID multiplexer system are described. The system provides substantial performance and practical improvements over previous generation analog systems.   This work is supported by the Natural Sciences and Engineering Research Council of Canada, Canadian Institute for Advanced Research, Canada Research Chairs program, a NASA GSRP grant, a Minnesota Graduate School Dissertation Fellowship, and the Canadian Foundation for Innovation. We thank Xilinx Canada for their contribution of the FPGA devices."
    },
    "0911/0911.4481_arXiv.txt": {
        "abstract": "Massive black hole (BH) mergers can result in the merger remnant receiving a ``kick,'' of order 200 km~s$^{-1}$ or more, which will cause the remnant to oscillate about the galaxy centre. Here we analyze the case where the BH oscillates through the galaxy centre   perpendicular or parallel to the plane of the galaxy  for a model galaxy consisting of an exponential disk, a Plummer model bulge, and an isothermal dark matter halo. For the perpendicular motion we find that there is a strong resonant  forcing of the disk radial motion near but somewhat less than the ``resonant radii'' $r_R$ where the BH oscillation frequency is equal one-half, one-fourth, ($1/6$, etc.) of the radial epicyclic frequency in the plane of the disk. Near the resonant radii there can be a strong enhancement of the radial flow and disk density which can lead to shock formation. In turn the shock may trigger the formation of a ring of stars near $r_R$. As an example, for a BH mass of $10^8~M_\\odot$ and a kick velocity of $150$ km s$^{-1}$, we find that the resonant radii lie between $0.2$ and $1$ kpc. For BH motion parallel to the plane of the galaxy we find that the BH leaves behind it a supersonic wake where star formation may be triggered.   The shape  of the wake is calculated as well as the slow-down time of the BH. The differential rotation of the disk stretches the wake into ring-like segments. ",
        "introduction": "Recent breakthroughs in numerical General Relativity have led to predictions of large recoil velocities of merged binary supermassive black-holes (BH) as the binary radiates away linear momentum as gravitational waves during the final stages of merger (Gonz\\'alez et al. 2007;  Campanelli et al. 2007;  Lousto, Campanelli, \\& Zlochower 2009). Typical kick velocities are of order $\\sim 200$~km~s$^{-1}$ (Bogdanovi\\'c, Reynolds, \\& Coleman 2007). One possible result of these mergers is to cause the resulting remnant black hole (BH) to oscillate through the stellar disk on timescales on the order of Gyr (Kornreich \\& Lovelace 2008, hereafter KL08; Blecha \\& Loeb 2008; Fujita 2009). If the merged BH receives an impulse from the merger, some of that motion will be transmitted to the galaxy by dynamical friction. It is clearly of interest to  know whether the motion of the BH through the disk generates observable changes in morphology and dynamics in the galaxy. Binary black-holes have already been observed as a double nucleus in a quasar (e.g., Decarli et al. 2009). Further, Comerford et al. (2009) show that BH binaries resulting from recent mergers are observable in the spectra of as many as $40\\%$ of Seyfert 2 galaxies. The possible observable effects of free and ``wandering''  BH merger remnants on disk galaxies, however, have not yet been fully analyzed. de la Fuente Marcos \\& de la Fuente Marcos (2008) discuss the formation of stars in the wakes of runaway black holes ejected from the host galaxy. Here we consider more slowly moving BHs which remain bound to the host galaxy.  This work first analyzes the case where the ejected BH oscillates with angular frequency $\\Omega_{bhz}$ through the galaxy center in a direction normal to the disk. Later, we discuss the case where the BH is ejected parallel to the plane of the galaxy. Sufficiently large gas accretion by the binary BH system is predicted to drive the orbital and BH spins into alignment normal to the plane of the galaxy (Bogdanovi\\'c et al. 2007). However, the amount is uncertain due to the uncertainty in  the ratio of the two viscosity coefficients for the $(r,z)$ and $(r,\\phi)$ motion in the disk (e.g., Natarajan \\& Armitage 1999). This alignment  favors BH ejection in the plane of the galaxy (Campanelli et al. 2007). The merging process preceding the ejection is assumed to be sufficiently slow that the  host galaxy has returned to an axisymmetric equilibrium state. The BH oscillations are observed in simulations to persist for many periods (KL08; Blecha \\& Loeb 2008; Fujita 2009). Under these conditions we find that there is a strong resonant  forcing of the disk radial motion near radii $r_R$ where $2n\\Omega_{bhz}= \\Omega_r(r_R)$, with $n=1,~2,..$, where $\\Omega_r$ is the radial epicyclic frequency. Near $r_R$ there can be a strong enhancement of the density in  a ring of the disk and shock formation which may trigger the formation of a ring of stars. We note that models for the formation of observed ``ring galaxies'' assume the (single) passage of one galaxy nearly through the center of another (Theys \\& Spiegel 1976, 1977; Lynds \\& Toomre 1976). The gravitational interaction of the two galaxies causes the formation of a pronounced ring(s) of enhanced density  (in one or both of the galaxies) where star formation is expected to be enhanced.  For the case of BH ejection in the plane of the galaxy we find that the BH leaves behind it a supersonic wake where star formation may be triggered.  The shape of this wake is calculated and the slow-down time of the BH is estimated. The differential rotation of the disk stretches the wake into ring-like segments.    Section 2.1 describes the gravitational potential of the equilibrium disk galaxy at the time of the BH merger.    This includes an exponential disk of stars and gas,  a Plummer model the bulge, and an isothermal dark matter halo. Section 2.2 derives the dynamical equations of the system as driven by the radial force due to a BH oscillating perpendicular to the disk. The frequencies and amplitudes of the BH oscillations are taken from the   $N$--body simulations of KL08, Section 2.3 discusses the forced oscillations of the gas disk, and \\S 2.4 the possible parametric instability of the disk. Section 3 discusses  the response of a galaxy disk to a BH ejected in the plane of the galaxy. Section 4 gives the conclusions of this work. ",
        "conclusions": "Here, we first analyzed the linear axisymmetric perturbations of a gas disk driven by a black hole oscillating vertically along the axis of symmetry. We find that there is a strong resonant  forcing of the disk radial motion near ``resonant radii'' $r_R$ where the BH oscillation frequency is equal one-half, one-fourth, ($1/6$, etc.) of the radial epicyclic frequency in the plane of the disk. Near the resonant radii there can be a strong enhancement of the radial flow velocity and disk density  which can lead to shock formation. The shock formation occurs during one period of oscillation of the BH which is assumed longer than the BH damping time due to dynamical friction. This shock may trigger the formation of a ring of stars near $r_R$. As an example, for a BH mass of $10^8~M_\\odot$ and a kick velocity of $150$ km s$^{-1}$, we find that the resonant radii lie between $0.2$ and $1$ kpc. The magnitude of the disk response is  proportional to the BH mass, inversely proportional to the Toomre~$Q$ of the gas disk, and decreases rapidly as $r_R$ increases. The resonant radii increase as the initial BH kick velocity increases (KL08).   For BH motion parallel to the plane of the galaxy we find that the BH leaves behind it a supersonic wake which over time gets contorted into a complicated shape by the galaxy's differential rotation.   The shape  of the wake is calculated for an illustrative case as well as the slow-down time of the BH. The differential rotation of the disk stretches the wake into ring-like segments of radii approximately equal to the maximum excursion of the BH, $x_m$, where the BH most slowly.  Many other processes may be involved in the formation and evolution of nuclear rings in galaxies (e.g.,  van de Ven \\& Chang 2009)."
    },
    "0911/0911.1092_arXiv.txt": {
        "abstract": "Using the nuclear symmetry energy that has been recently constrained by the isospin diffusion data in intermediate-energy heavy ion collisions, we have studied the transition density and pressure at the inner edge of neutron star crusts, and they are found to be $0.040$ fm$^{-3}$ $\\leq \\rho _{t}\\leq 0.065$ fm$^{-3}$ and $0.01$ MeV/fm$^{3}$ $\\leq P_{t}\\leq 0.26$ MeV/fm$^{3}$, respectively, in both the dynamical and thermodynamical approaches. We have also found that the widely used parabolic approximation to the equation of state of asymmetric nuclear matter gives significantly higher values of core-crust transition density and pressure, especially for stiff symmetry energies. With these newly determined transition density and pressure, we have obtained an improved relation between the mass and radius of neutron stars. ",
        "introduction": "\\label{introduction}  Studying the properties of  neutron stars, which are among the most mysterious objects in the universe, allows us to test our knowledge of matter under extreme conditions. Theoretical studies have shown that neutron stars are expected to have a liquid core surrounded by a solid inner crust~\\cite{Cha08}, which extends outward to the neutron drip-out region. While the neutron drip-out density $\\rho _{\\rm out}$ is relatively well determined to be about $4\\times 10^{11}$ g/cm$^{3}$ or $0.00024~{\\rm fm}^{-3}$~\\cite{Rus06}, the transition density $\\rho _{t}$ at the inner edge is still largely uncertain mainly because of our very limited knowledge on the nuclear equation of state (EOS), especially the density dependence of the symmetry energy ($E_{\\rm sym}(\\rho)$) of neutron-rich nuclear matter~\\cite{Lat00,Lat07}. These uncertainties have hampered our understanding of many important properties of neutron stars~\\cite{Lat00,Lat07,Lat04} and related astrophysical observations~\\cite{Lat00,Lat07,BPS71,BBP71,Pet95a,Pet95b,Ste05,Lin99,Hor04,Bur06,Owe05}. Recently, significant progress has been made in constraining the EOS of neutron-rich nuclear matter using terrestrial laboratory experiments (See Ref.~\\cite{LCK08} for the most recent review). In particular, the analysis of the isospin-diffusion data \\cite{Tsa04,Che05a,LiBA05c} in heavy-ion collisions has constrained tightly the $E_{\\rm sym}(\\rho)$ in exactly the same sub-saturation density region around the expected inner edge of neutron star crust. The extracted slope parameter $L=3\\rho_0\\frac{\\partial E_{\\rm sym}(\\rho)}{\\partial\\rho}|_{\\rho=\\rho_0}$ in the density dependence of the nuclear symmetry energy was found to be $86\\pm25$ MeV, which has further been confirmed by a more recent analysis~\\cite{Tsa09}. With this constrained nuclear symmetry energy, we have obtained an improved determination of the values for the transition density and pressure at the inner edge of neutron star crusts. This has led us to obtain significantly different values for the radius of the Vela pulsar from those estimated previously. Also, we have found that the widely used parabolic approximation (PA) to the EOS of asymmetric nuclear matter gives much larger values for the transition density and pressure, especially for stiff symmetry energies. ",
        "conclusions": ""
    },
    "0911/0911.4354_arXiv.txt": {
        "abstract": "Using the results of extensive Monte Carlo simulations we discuss corrections to the linear mixing rule in strongly coupled binary ionic mixtures.  We analyze the plasma screening function at zero separation, $H_\\mathit{jk}(0)$, for two ions (of types $j=1,2$ and $k$=1,2) in a strongly coupled binary mixture. The function $H_\\mathit{jk}(0)$ is estimated by two methods: (1) from the difference of Helmholtz Coulomb free energies at large and zero separations; (2) by fitting the Widom expansion of $H_\\mathit{jk}(x)$ in powers of interionic distance $x$ to Monte Carlo data on the radial pair distribution function $g_{jk}(x)$. These methods are shown to be in good agreement. For illustration, we analyze the plasma screening enhancement of nuclear burning rates in dense stellar matter. ",
        "introduction": "More than 30 years ago \\cite{HV76} the linear mixing rule for multicomponent strongly coupled mixtures was shown to be highly accurate. However, only recent studies \\cite{PCR09,Mixt_New} have achieved enough accuracy to describe the corrections to the linear mixing rule for a wide range of plasma parameters; previous attempts, e.g.\\ \\cite{DWSC96,DWS03}, were restricted at least by a limited number of data points. We discuss the corrections to the linear mixing rule in application to the plasma screening of nuclear reactions in strongly coupled mixtures. Following Ref.\\ \\cite{DWS99} we apply two approaches to calculate the screening enhancement: one is based on the thermodynamic relations and the other on fitting the mean-field potentials. The main advance of the present work is in using a much wider set of numerical data and most precise thermodynamic results. ",
        "conclusions": "\\label{Sec:approx}   We suggest to use the following approximation for the enhancement factor for all $\\Gamma$ and mixture composition \\begin{equation} h^0_{jk}=h^\\mathrm{lin}_{jk}/ \\left[1 +C_{jk}\\left(1-C_{jk}\\right)\\, \\left(h^\\mathrm{lin}_{jk}/h^\\mathrm{DH}_{jk} \\right)^2 \\right]. \\label{appr} \\end{equation} Here, $h^\\mathrm{lin}_{jk}$ is given by (\\ref{h0lin}), $ h^\\mathrm{DH}_{jk}=3^{1/2} Z_j\\,Z_k\\ZZmid^{1/2}\\Gamma_\\mathrm{e}^{3/2}/\\Zmid^{1/2}$ is the well known Debye-H\\\"{u}ckel enhancement parameter, and   % $C_{jk}= 3Z_j\\,Z_k\\ZZmid^{1/2}\\Zmid^{-1/2}/\\left[\\left(Z_j+Z_k\\right)^{5/2}-Z_j^{5/2}-Z_k^{5/2}\\right] $. Eq.\\ (\\ref{appr}) reproduces the Debye-H\\\"{u}ckel asymptote at low $\\Gamma$ and the linear mixing at strong coupling.  \\begin{figure} \\begin{center} \\leavevmode \\epsfxsize=150mm \\epsfbox{ChugunovAI_fig2.eps} \\end{center} \\caption{(Color online) Enhancement factors $h_{jk}^0/\\Gamma_{11}^{3/2}$ vs $\\Gamma_{11}$ for three binary ionic mixtures. } \\label{Fig_hvsG} \\end{figure}  In Fig.\\ \\ref{Fig_hvsG} we show the dependence of the approximated enhancement factors $h_{jk}^0/\\Gamma_{11}^{3/2}$ on $\\Gamma_{11}$. The figure contains three panels; each for a specific binary ionic mixture. Each panel shows three groups of four lines. They are (from top to bottom) $h^0_{22}/\\Gamma_1^{3/2}$, $h^0_{12}/\\Gamma_1^{3/2}$ and $h_{11}^0/\\Gamma_1^{3/2}$. Two of any four lines (solid and thick dashed lines) are almost the same in the majority of cases. This couple represents the approximation (\\ref{appr}) and the thermodynamic enhancement factor (\\ref{h0}), respectively. The dotted horizontal lines refer to the Debye-H\\\"{u}ckel model and the dash-dot lines are the linear mixing results. One can see that our approximation is in a good agreement with thermodynamic results for most of cases, especially in panel (a) (for all mixtures with not too large $Z_2/Z_1$). If $Z_2/Z_1$ becomes too large [panel (c)], the thermodynamic model of $h_{11}^0/\\Gamma_1$, calculated in accordance with \\cite{Mixt_New}, has a specific feature ($h_{11}/\\Gamma_{11}^{3/2}$ increases at $\\Gamma_{11}\\sim10^{-2}$), while our approximation has not. We expect that this feature is not real, but results from not too accurate extractions of the enhancement factors from thermodynamic data. The free energy is almost fully determined by larger charges $Z_2$ which also dominate by number ($99\\%$) in panel (c). Using Eq.\\ \\ref{h0} to get $h_{11}^0$, one should differentiate the free energy with respect to $N_1$, which provides vanishing contribution to the free energy. Hence this procedure is very delicate and can strongly amplify the errors of original thermodynamic approximation. We expect that our approximation can be more accurate than the original thermodynamic result. Another, less probable option is that we still have not enough data to prove the presence of the feature of $h^0_{11}/\\Gamma_{11}^{3/2}$.    To conclude, we have calculated the enhancement factors of nuclear reactions in binary ionic mixtures by two methods and showed good agreement of the results. We have proposed a simple approximation of the enhancement factors valid for any Coulomb coupling. This approximation is almost the same as thermodynamic ones for not too specific mixtures. It does not confirm some questionable features of the enhancement factors for mixtures with large $Z_2/Z_1$ and small $x_1$.   \\begin{acknowledgement} We are grateful to D.G.~Yakovlev and A.Y.~Potekhin for useful remarks. Work of AIC was partly supported by the Russian Foundation for Basic Research (grant 08-02-00837), and by the State Program ``Leading Scientific Schools of Russian Federation'' (grant NSh 2600.2008.2). Work of HED was performed under the auspices of the US Department of Energy by the Lawrence Livermore National Laboratory under contract number W-7405-ENG-48. \\end{acknowledgement}"
    },
    "0911/0911.0161_arXiv.txt": {
        "abstract": "Assuming that density waves trigger star formation, and that young stars preserve the velocity components of the molecular gas where they are born, we analyze the effects that non-circular gas orbits have on color gradients across spiral arms. We try two approaches, one involving semi-analytical solutions for spiral shocks, and another with magnetohydrodynamic (MHD) numerical simulation data. We find that, if non-circular motions are ignored, the comparison between observed color gradients and stellar population synthesis models would in principle yield pattern speed values that are systematically too high for regions inside corotation, with the difference between the real and the measured pattern speeds increasing with decreasing radius. On the other hand, image processing and pixel averaging result in systematically lower measured spiral pattern speed values, regardless of the kinematics of stellar orbits. The net effect is that roughly the correct pattern speeds are recovered, although the trend of higher measured $\\Omega_p$ at lower radii (as expected when non-circular motions exist but are neglected) should still be observed. We examine the \\citet{mar09} photometric data and confirm that this is indeed the case. The comparison of the size of the systematic pattern speed offset in the data with the predictions of the semi-analytical and MHD models corroborates that spirals are more likely to end at Outer Lindblad Resonance, as these authors had already found. ",
        "introduction": "Until the detection of an azimuthal color gradient across one of the arms of the SAc galaxy M~99 \\citep[][GG96 hereafter]{gon96}, only sparse evidence of star formation triggered by spiral density waves had been found in stellar counts in the Milky Way\\citep{sit89,sit91,ave89} and M~31\\citep{efr80a,efr80b,efr85}. More recently, applying the same method defined by GG96 and described below, \\citet[][MG09 hereafter]{mar09} examined a sample of 13 spiral galaxies of types A and AB, and found color gradients consistent with theoretical expectations in 10 of their objects.\\footnote{See also \\citet{gro09} for an infrared study of young stellar complexes in NGC~2997.} Although they did not compute stellar orbits, both GG96 and MG09 implicitly assumed that all stars, including those recently born near spiral arms, move in circular orbits. MG09 did investigate the effects of variable circular speeds and variable densities on observed color gradients and found them to be negligible. However, complicated non-circular motions have been reported in studies about the migration of young stars following star formation triggered by spiral shocks \\citep{yua69,wie79,fer08}. \\citet{bash81} had also noticed that the ballistic orbits calculated from galactic HII region complexes follow non-circular trajectories that initially move along (and not across) the spiral arms.  In the present work, we examine how the presence of non-circular motions would modify the pattern speeds derived from color gradients under the assumption of circular orbits.  We will explore two different approaches. The first one (see \\S~\\ref{appr1}) involves the semi-anaylitical solutions obtained for spiral shocks \\citep{rob69,shu73,gitt04}; we assume that newly born stars preserve the orbital motion of the shocked gas where they form. The second approach (see \\S~\\ref{appr2}) is based on data from MHD simulations; we follow the gas flow vectors near spiral arms. We always assume that stars are triggered all along the studied regions of spiral arms in the respective model. Other methods, not discussed here, may involve orbit calculations for young stellar groups. ",
        "conclusions": "Under the assumption that the orbits of young stars preserve the velocity components of the parent molecular clouds where they form, we have analyzed the effects that non-circular motions would have on azimuthal color gradients. Semi-analytical calculations and MHD simulations show that the spiral pattern speeds derived from the comparison between color gradients and stellar population synthesis models, assuming purely circular motions, would have values systematically higher than the real ones for regions within corotation. Also inside corotation, the effect would decrease with galactocentric radius. Non-circular motions, however, would not prevent the detection of legitimate azimuthal color gradients in real galaxies.  Using a synthetic image, we have also analyzed the effects of image processing and pixel averaging on the method applied in GG06 and MG09 to detect color gradients and derive pattern speeds. We have found that pixel averaging (due to image processing) systematically decreases the derived $\\Omega_p$, such that it nearly compensates for the systematic effect introduced when neglecting non-circular motions in the analysis. The net result is that the correct spiral pattern speeds and resonance locations can be obtained. Nevertheless, a residual trend of slightly higher pattern speeds at lower radii can still be discerned ({\\it solid} triangles in Fig.~\\ref{myOMEGAS} and Fig.~\\ref{diffOM_OLR}), so that it is possible to detect the presence of non-circular motions and confirm the link between star formation and disk dynamics.  We have re-examined the results obtained by MG09. When normalizing the mean radii where gradients were found by the end point radii, in order to treat the whole sample as a single galaxy, we were able to reproduce the trend of $\\delta \\Omega_{p}$ with radius expected if non-circular motions are ignored (as these authors did). The size of the observed $\\delta \\Omega_{p}$ in the data only matches the theoretical expectations if spiral patterns end at the OLR, and not at the 4:1 resonance (cf.\\ Fig.~\\ref{myOMEGAS} with both Fig.~\\ref{diffOM_4_1} and Fig.~\\ref{diffOM_OLR}).  Future studies of azimuthal color gradients in disk galaxies by orbit calculations should include the effects of disk heating \\citep{asi99}, and a correct modeling of the IMF as a discrete distribution. These factors may be important, since high mass stars may take different trajectories when compared to low mass stars, due to the effects of dynamical friction~\\citep{chandra43}."
    },
    "0911/0911.5091_arXiv.txt": {
        "abstract": "{The dynamics of prominence fine structures  is a challenge to understand the formation of cool plasma prominence embedded in the hot corona.} {Recent observations from  the high resolution \\emph{Hinode}/SOT telescope allow us to compute  velocities perpendicularly to the line-of-sight or transverse velocities. Combining simultaneous observations obtained in H$\\alpha$ with \\emph{Hinode}/SOT and the MSDP spectrograph operating in the Meudon solar tower we derive the velocity vectors of a quiescent prominence.} {The velocities perpendicular to the line-of-sight  are measured by time slice technique, the Dopplershifts by the bisector method.} {The Dopplershifts of bright threads derived from the MSDP reach 15 km s$^{-1}$ at the edges of the prominence and are between $\\pm$ 5 km s$^{-1}$ in the center of the prominence. Even though they are minimum values due to seeing effect, they are of the same order as the transverse velocities.} {These measurements are  very important because they suggest that the vertical structures shown in SOT may not be real vertical magnetic structures in the sky plane. The vertical structures could be a pile up of dips in more or less horizontal magnetic field lines in a 3D perspective, as it was proposed by many MHD modelers. In our analysis we also calibrate the \\emph{Hinode} H$\\alpha$ data using MSDP observations obtained simultaneously.} ",
        "introduction": "The existence of cool structures in so called prominences and filaments during a few solar rotations embedded in the hot corona has been a mystery since the beginning of their spectrographic observations \\citep{Azambuja48}. Many reviews concerned the study of quiescent prominences \\citep{Schmieder89,Tandberg94,Labrosse09,Mackay09}. Since that time period, it is a challenge to derive what is the best mechanism to succeed to maintain cool plasma in the corona. A very popular idea is that the plasma is frozen in magnetic field lines and  stays cool due to the  low transverse thermal conduction \\citep{Demoulin89}. Many magnetic and static models have been developed on this idea \\citep{Kuperus74,Kippenhahn57,Aulanier98,Dudik08}. But a big question remains: how to get cool plasma inside the field lines into the corona and how to keep it there?  It is recognized that this material should come from the chromosphere by levitation or by injection \\citep{Saito73}. Sufficient mass should be extracted from the chromosphere by magnetic forces which inject or lift the plasma or by pressure forces which evaporate the plasma and then cool it to prominence temperatures. Many models have been developed in this sense, i.e. thermal non-equilibrium models \\citep{Mariska85,Karpen03,Karpen05}. Levitation models are proposed through possible magnetic reconnection \\citep{Ballegooijen89}. Injection models could be due to injection through the reconnection of magnetic field during canceling flux. These models indicate that the plasma in prominences should have a large dynamics and static models  may be obsolete. Recent observations at the Swedish Solar Telescope (SST) show highly dynamic plasma in filament threads \\citep{Lin03,Lin05}. It was also a first attempt to compute the velocity vectors of the filament threads. They conclude that the threads inclination from the horizontal were around 16 degrees with a net flow in both directions of 8 km s$^{-1}$.  Fine counterstreaming flow is often observed along horizontal threads or in the barbs \\citep{Zirker98,Schmieder91,Schmieder08}. \\emph{Hinode}/SOT movies (available with the electronic version of the paper) reveal strong dynamics in the prominence fine structures. The spicules close to the barbs could allow to inject plasma inside the fine threads. Is that sufficient to feed continuously the main core of the prominence?  A mass budget should be done. Using \\emph{Hinode} observations  at the limb \\citet{Berger08} and \\citet{Chae08} tried to answer  these questions. They reported different velocity measurements. \\citet{Berger08} found upflows of  dark bubbles around 20 km s$^{-1}$ and down flows of bright knots less than 10 km s$^{-1}$ looking at Ca II H images in the line center. \\citet{Chae08} analysed an hedgerow prominence and found horizontal displacements before observing downflows of bright knots suggesting the existence of  vortex motions.  \\begin{figure} \\centering \\hspace*{-1.5cm} \\includegraphics[width=0.60\\textwidth,clip=]{fig1.ps} \\caption{(a)  EIT 304 \\AA\\ image observed on April 20, 2007 at 01:00 UT, the field of view is 1100$\\times$750 arcsec, (b) Filament observed in H$\\alpha$ at Meudon on April 22, 2007 at 15:28 UT, (c) MDI longitudinal magnetic field on April 21 at 08:00 UT. The magnetic inversion line is represented by the dashed line in the middle of the EUV filament channel and in MDI image. The letters F1 and F2 indicate approximately the location of the two filament fragments visible in H$\\alpha$. }\\label{meudon} \\end{figure}   Prominence dynamics at the limb look quite different from filament dynamics on the disk. The integration  along the line of sight complicates the interpretation. \\citet{Mein91} show that more than 15 threads may be integrated along the line of sight and the resulting velocities should depend on two different Gaussian distributions. At the edges of prominences the velocity values are higher because fewer threads are integrated and velocity cells are larger than intensity knots revealing that {\\bf bunches} of threads may move with the same velocity plasma (Dopplershifts). Similar results have been  found from completely different approaches. Therefore to reproduce Lyman line profiles observed by SOHO/SUMER, \\citet{Gunar07} introduce 10 threads perpendicularly to the line-of-sight with a random velocity distribution in a 2D non LTE radiative transfer code.  We proposed in this study to study both the velocity perpendicular to the line of sight using one hour of observations of \\emph{Hinode}/SOT in H$\\alpha$ combined with Dopplershifts observed  also in H$\\alpha$ with the Multichannel Subtractive Double Pass spectrograph (MSDP) operating in the Meudon solar tower. These observations were obtained simultaneously during a coordinated observing program (JOP178). It is the first time that using \\emph{Hinode} data such fine structures are resolved in prominences and that oscillations and transverse velocities can be derived \\citep{Okamoto07,Berger08,Chae08}. \\emph{Hinode} has a much better spatial resolution than MSDP by a factor of 5, but the MSDP observations are very useful to calculate the Dopplershifts and to calibrate the intensity of  \\emph{Hinode}/SOT observations. ",
        "conclusions": "For the first time an H$\\alpha$ hedgerow prominence has been observed simultaneously by a high spatial resolution telescope (\\emph{Hinode}/SOT) and by a spectrograph (MSDP) operating in  the Meudon solar tower on April 25 2007. \\emph{Hinode}/SOT allows us to get the velocities perpendicular to the line-of-sight V(x,z). The second -- the line-of-sight velocities  V(y) derived from Dopplershifts. The prominence shows strong dynamics in the SOT movie with dark cavities rising from the limb with an upward velocity reaching 24 km s$^{-1}$ and downflowing vertical-like bright threads. These threads are moving horizontally to avoid the dark cavities. During the rise of cavities, ahead of them,  are observed bright curved fine structures from time to time with high velocities similar to the speed rise of the bubble. The \\emph{Hinode}/SOT observations have been calibrated by using the MSDP data. The integrated H$\\alpha$ intensity of the threads reaches 1.5 $\\times 10^5$ erg/s/sr/cm$^2$. The contrast of the dark cavities is between 70 and 90$\\%$.  \\begin{figure} \\vspace*{-1cm} \\centering \\hspace*{-2cm} \\includegraphics[width=0.60\\textwidth,clip=]{fig7.ps} \\vspace*{-1cm}  \\caption{Transverse velocities in SOT bright structures using time slice technique (axis x unit is time, axis y unit is arc sec along the slide). Top/medium/bottom frame corresponds to slice A/ D/ E, drawn in Fig. \\ref{doppler}. They show the velocities measured respectively in points A1, A2, D1, D2, E1, E2 (Table 1). Positive/negative velocities correspond to up/down flows. The large value +24  km s$^{-1}$ corresponds to the speed of a rising bubble from the limb or the flow  speed of its bright edge. Fine threads close to the limb are spicules. Wave  pattern corresponds to oscillations with 15 to 20 min of period. Adjacent pixels in a slice have coherent velocities. } \\label{sot_velocity} \\end{figure}  The transverse velocities V(x,z)  of the bright threads are computed by time slice technique and these values are of the order of  a few km s$^{-1}$ to 6 km s$^{-1}$ reaching 11 km s$^{-1}$ for individual fine threads. The pattern of the Dopplershift map show elongated cells nearly perpendicular to the limb. They are {\\bf wider than the MSDP spatial resolution}. The time slice maps exhibit several pixels close to each other along the slice having a similar velocity trend. This means that fine threads close to each other have a coherent displacement. According to the observations of the prominence three days before when it is still on the disk as a filament, it appears that only the feet or barbs are enough dense to be observed. The prominence would represent the barb threads integrated along the line-of-sight as the filament is crossing the limb. The structures are not vertical in the sky plane (x,z) as suggested by the movie. Dopplershifts and transverse velocities are of same order of magnitude (less than 6 km s$^{-1}$). In Table 1  we have selected individual threads with the largest transverse velocities in region with higher Dopplershifts. The other parts of the prominence exhibit coherent velocity much smaller (1 to 2 km s$^{-1}$) difficult to measure. The narrow profiles of H$\\alpha$ lines of the prominence indicate  that the different threads integrated along the line of sight have similar velocities. The dispersion of the velocities along the line-of-sight is small.  The longitudinal magnetic field observed (by the SOHO/MDI instrument) in the filament channel on the disk and on both edges of the inversion line is weak. The strength of the small polarities are less than 10 Gauss. The prominence lies in a quiet region and corresponds to a quiescent filament. The small polarities can change rapidly and this would explain the fast dynamics of the structures.  In  a flux tube model \\citep{Aulanier98,Dudik08} the H$\\alpha$ filament  is considered as cool material trapped in shallow dips along long magnetic  field lines. The feet are extension of the flux tube disturbed  laterally by parasitic polarities. The barbs are piled up dips touching the  photosphere. When a parasitic polarity cancels or moves,  the feet move and can even disappear \\citep{Aulanier98b,Schmieder06,Gosain09}. {\\bf In this 3D perspective the prominence material would be  trapped in inclined field lines   and the downflow motion would occur along the shallow dips.  The brightness would result of the integration of the threads along the line-of-sight \\citep[see Fig. 3e in ][]{Dudik08}.} \\citet{Aulanier98b} explain very well through their magnetic extrapolation the relationship between parasitic polarities and the flux tube itself and their evolution. The flux of parasitic polarity overcomes that of the twisted flux tube and destroys the twisted configuration. The bubbles would be structures more magnetized that the surrounding and represented  by the separatrices. A small increase of  magnetic pressure in the bubble would lead to the rise of it in the atmosphere. Strong currents can be created in the quasi separatrice layers (QSL) around the separatrices  by photospheric displacements of the parasitic polarities \\citep{Demoulin96}. Energy release is expected. This could correspond to the brightening {\\bf rims} associated with filaments \\citep{Heinzel95}. They are not systematically visible due to the dense plasma of filaments. In SOT observations it could correspond to the bright top edge of the cavities where reconnection could occur and expel plasma. This would explain the fast velocity material in brighter threads surrounding the dark cavities. On the next days, such bubbles are not observed in the prominence because the feet and parasitic polarities related to it would be on the back side of the disk. {\\bf In the arcade model with dips proposed by \\citet{Heinzel01} the prominence observed on April 25 may would consist of vertical threads trapped in dips and piled up giving the impression of vertical continuous threads. The downflows of 1 to 5 km s$^{-1}$ would be due to shrinkage or successive reconnections of field lines.}  Another explanation for the buoyancy of the dark cavities or bubbles could be adiabatic expansion of a heated volume of plasma \\citep{Berger08}. This is not exclusive of the magnetic pressure increase scenario and both magnetic and thermal buoyancy may play a role in the formation of these structures. We would like to measure the magnetic field in the bubbles and in the prominence. An interesting aspect would be also to analyse the EIS and SUMER data to see if the dark bubbles are filled with hot material. Such measurements are needed to know which mechanism is valid for the formation of these dark low cavities."
    },
    "0911/0911.0999_arXiv.txt": {
        "abstract": "We have used archival RXTE PCA data to investigate timing and spectral characteristics of the transient XTE J1817-330. The data pertains to 160 PCA pointed observations made during the outburst period 2006, January 27 to August 2. A detailed analysis of Quasi-Periodic Oscillations (QPOs) in this black hole X-ray binary is carried out. Power density spectra were obtained using the light curves of the source. QPOs have been detected in the 2-8 keV band in 10 of the observations. In 8 of these observations, QPOs are present in the 8-14 keV and in 5 observations in the 15-25 keV band. XTE J1817-330 is the third black hole source from which the low frequency QPOs are clearly detected in hard X-rays. The QPO frequency lies in $\\approx$ 4-9 Hz and the rms amplitude in 1.7-13.3\\% range, the amplitude being higher at higher energies. We have fitted the PDS of the observations with Lorentzian and power law models. Energy spectra are derived for those observations in which the QPOs are detected to investigate any dependence of the QPO characteristic on the spectral parameters. These spectra are well fitted with a two component model that includes the disk black body component and a power law component. The QPO characteristics and their variations are discussed and its implication on the origin of the QPOs are examined. ",
        "introduction": "\\label{sec:intro}  Accretion powered X-ray binaries are the brightest X-ray sources in our Galaxy. These binaries contain either an accreting neutron star or an accreting black hole as the X-ray source. Based on estimates of the mass of the accreting object and its X-ray characteristics, about 40 X-ray sources have been classified as black hole binaries (Remillard and McClintock 2006a). Of these 20 have reliable mass estimates and are, therefore, regarded as confirmed black holes while the remaining 20 are considered to be black hole candidates (Remillard and McClintock 2006a). A majority of the black hole binaries (BHBs) are transients and most of them have a low mass optical companion.\\\\  The black hole X-ray binaries have several distinctive X-ray characteristics. During the outburst there is a strong soft component that originates in the inner region of the hot accretion disk. The presence of a hard X-ray component in energy spectra is another common feature of black hole binaries. The hard X-rays arise through Compton scattering of low energy photons from the accretion disk in a hot optically thin plasma. The resulting hard X-rays are referred to as the thermal Compton component (McClintock and Remillard 2006).\\\\  Low frequency quasi-periodic oscillations (QPOs) in $\\approx$ 0.1-40 Hz range and high frequency QPOs $\\approx$ 50-450 Hz also occur in many BHBs. Presence of several distinct spectral states and transition from one state to another at irregular intervals, is another distinguishing feature of BHBs. Spectral and temporal characteristics of the sources vary from one state to another. Remillard and McClintok (2006a) have broadly classified the spectral state for BHB as (a) High or Soft state (HS) dominated by the thermal component, (b) Low or Hard State (LH) marked by the hard X-ray power law and (c) Steep power law (SPL) or very high state characterized by steep slope power law component. For a more complete description of the spectral states McClintock and Remillard (2006) included two more states namely an Intermediate state (IM) that occurs when the source moves from LH to HS state and an extreme LH state which they called as quiescent state. Gierlinski, Done and Page (2008) have included an additional state termed as Ultra soft state (US) which is an extreme case of high/soft state found in this source. The US state is characterized by a very weak high energy tail with a low disk temperature and very low hardness ratio $\\leq$ 0.1 (Gierlinski, Done and Page 2008). The low frequency QPOs are usually detected in the SPL and LH states and have been observed in 14 BHBs so far (Remillard and McClintock 2006a). Detection of LFQPO in XTE J1817-330 indicates that LFQPOs can also be found sometime in the HS state.\\\\  An X-ray transient known as XTE J1817-330 was discovered by Remillard et al. (2006b) on 2006 January 26, with the All Sky Monitor (ASM) on Rossi Timing X-ray Explorer (RXTE) (Levine et al. 1996). At the time of its detection the 2-12 keV flux was  0.93 Crab. The intensity then rose to a peak value of 1.9 Crab on January 28, and then declined to 1.2 Crab by January 30. Subsequently it decayed exponentially with a decay time of 27 days (Sala et al. 2007). From its high/soft state at the time of the outburst, the source declined to a low/hard state characterized by kT = 0.2 keV. A change in the intensity state of the source occurred around $\\sim$ February 9 (Shaw et al. 2006). Hard X-rays (20-60 keV) flux increased from 46 $\\pm$ 2 mCrab on February 9  to 79 $\\pm$ 2 mCrab on February 14 accompanied by a decrease of soft X-ray (2-10 keV) flux from 840 $\\pm$ 4 mCrab to 670 $\\pm$ 4 mCrab during this period indicating the onset of the hard state (Kuulkers et al. 2006).\\\\  Its energy spectrum was studied with the RXTE, XMM-Newton, Integral and Swift instruments. Sala et al. (2007) measured the spectra using data from the XMM-Newton and Integral instrument when the source was in a high/soft state during 2006 February-March. Their spectral results indicated that its energy spectrum was typical of a BHB source with a dominant thermal disk component well described by kT $\\sim$ 0.7-0.9 keV and a thermal Compton power law component with a photon index of $\\sim$ 2-3 (Sala et al. 2007). Its spectral characteristics were observed with X-ray telescope (XRT) on the Swift satellite covering different stages of the outburst over 160 days. During this period the source made transition from the high/soft state in the initial outburst to a low hard state near the end of the outburst. The XRT spectra in the high soft state in 0.6-10 keV are well described by a two component model consisting of a thermal component from a optically thick and geometrically thin accretion disk and a hard power law component. The temperature of the inner disk producing the soft component declined from $\\sim$ 0.8 keV during the initial outburst phase to $\\sim$ 0.2 keV near the end of the outburst (Rykoff et al. 2007).    A detailed study of the spectral evolution of XTE J1817-330 at different phases of the outburst was carried out by Gierlinski, Done and Page (2008) using the RXTE and the Swift data. The spectra of 150 PCU2 (RXTE) observations were modeled with the two component spectral model consisting of a disk component and a hard thermal Compton power law component. Using the same XRT data as used by Rykoff et al. (2007), they also found that the inner disc temperature declined from $\\sim$ 0.9 keV to $\\sim$ 0.2 keV as the source intensity declined. Based on this they claimed that the accretion disk recedes when the source transits from the high/soft state to the low/hard state. It may also be noted that apart from XTE J1118+480, this black hole binary has the lowest absorption along the line of sight among all the bright black hole candidates as obtained from the Chandra and Swift spectral data (Miller et al. 2006a;b).  \\\\  Power density spectra of XTE J1817-330 obtained from the first two X-ray observations during 06:03-17:04 UTC on 2006 February 24, revealed strong QPOs at $\\sim$8.5 Hz (Homan, Miller and Wijnands 2006). Third observation (13:20-13:43 UTC) on the same day showed only a weak QPO at around 6.4 Hz. The last 2 observations (14:55-17:04 UTC) again indicated the presence of strong QPOs at 5.0 Hz.\\\\  Rupen, Dhawan and Mioduszewski (2006a), detected with VLA a radio object in the error box of the X-ray source having a flux density of 2.1 mJy at 1.4 GHz on 2006 January 31. The radio  source faded away by 2006 February 2 (Rupen, Dhawan and Mioduszewski 2006b). A bright optical counterpart of the X-ray source was found at the time of the outburst with V = 11.3 magnitude and its brightness decreased to V = 15.5 by February 10 (Torres et al. 2006). The optical star was also detected in the near-infrared with a K magnitude of 15.0 on 2006 February 7 (D'Avanzo et al. 2006). Near-UV observations with the Optical Monitor on the XMM-Newton showed variations in the UV flux correlated to the hard X-ray flux variations (Sala et al. 2007).  \\\\  All these characteristics strongly suggest that XTE J1817-330 is most likely a black hole binary. This transient was repeatedly observed with the Proportional Counter Array (PCA) on RXTE in the pointed mode during the period 2006 January 27 - August 2. We have carried out detailed timing analysis of the PCA data to study the properties of the QPOs and in this paper we present the results of this analysis. ",
        "conclusions": "\\label{Discussion section} The LFQPOs occur most frequently when the power law flux is the dominating component in the energy spectrum. Some times they are also present in the high luminosity state with the presence of a hard component. From table \\ref{table1.tab} and \\ref{table2.tab} it will be noticed that in all the cases of the detection of LFQPOs from XTE J1817-330, except the observations of MJD 53764 and 53775 in which no QPOs are detected, the ratio of the power law flux to the thermal disk flux lies in 0.20 to 1.13 range consistent with its occurrence only in the states with a significant power law component. Also note that the LFQPOs have significant coherence (Q $=$ $\\nu$/$\\delta\\nu $) with the Q in range of 3-11 and their rms amplitude vary from a few percent to as high as 13\\%.\\\\  Variation of the QPO frequency with the source intensity is another feature detected in some BHBs. The fundamental QPO frequency in XTE J1817-330 varies in a narrow band of 4.4-5.9 Hz. We have investigated the variation of the QPO frequency in the 2-8 keV band with the thermal disk component flux ( d$_{bb}$ ). This is shown in Fig \\ref{fig:fig9} and a trend similar to that of Fig \\ref{fig:fig7}(a) is seen here indicating that the frequency is correlated with the thermal disk component. As expected a clear 1:2 relationship of the QPO fundamental frequency and that of the first harmonic is seen.  All the characteristics of the LFQPOs reported by us in this paper from XTE J1817-330 are similar to those seen in the other black hole binaries and further strengthen the black hole nature of this source (Remillard et al. 2003).\\\\  Correlation of the properties of LFQPOs with the spectral parameters of the BHBs has been studied in detail for several sources (Muno, Remillard and Morgan 2001; Tomsick and Kaaret 2001; Remillard et al. 2003; Belloni, Psaltis and van der Klis 2002; Vignarca et al. 2003; Rossi, Homan and Belloni 2004). These studies show that the LFQPO characteristics are generally well correlated with the thermal disk and the power law components of the energy spectra. While the QPO frequency is closely correlated with the disk flux, the amplitude of the QPOs for the fundamental frequency is found to track the flux of the power law component (Remillard et al. 2003). In general the QPO amplitude is higher for the higher energy X-rays up to about 20 keV and then it tends to decrease at the higher energy. Most of the LFQPOs detected in the black hole binaries occur below 10 keV. In two of the BHBs namely GRS 1915+105 and XTE J1550-564 the QPOs have been reported above 20 keV (Trudolyubov, Churazov \\& Gilfanov 1999; Remillard et al. 2003). In some cases the QPO detection is claimed in a broad spectral band of 2-60 keV but since no breakdown of QPO properties is given in the different energy intervals say below 10 keV and above 10 keV, it is not obvious whether the QPOs have indeed been detected in hard X-rays. The enigmatic source GRS 1915+105 is the only black hole binary in which the QPOs in 0.8-3.0 Hz have been reported at an energy up to 124 keV (Tomsick and Kaaret 2001). It is found that in this object, the amplitude of the fundamental frequency QPO increases with energy up to 29 keV and then decreases in 30-60 keV and 60-124 keV bands. Remillard et al. (2003) investigated the QPO characteristics in the transient XTE J1550-564 from the RXTE - PCA observations and detected LFQPOs in 2-13 keV and 13-30 keV energy channels making it only the second black hole binary in which the LFQPOs have been detected up to 30 keV. We have detected the QPOs from XTE J1817-330 in five of the observations up to an energy of $\\sim$ 25 keV, making it only the third BHB showing unambiguous presence of the LFQPOs at higher energy. From Table \\ref{table1.tab} it may be noticed that the amplitude of the fundamental QPOs is always higher in the 8-14 keV channel, being in 3.2-13.3 \\% range, compared to the values of 1.7-7.0 \\%  in the 2-8 keV channel. At still higher energy 15-25 keV, the QPO amplitude is comparable to or slightly higher than that in the 8-14 keV indicating that it has reached a plateau level.\\\\  A detailed study of the QPO centroid frequency, its coherence, amplitude, phase lag and their dependence on the photon energy, is of vital importance to understand the origin of the QPOs and pin point the emission process. The QPOs are believed to originate in the innermost region of the accretion disk and the most common models explain their generation to the modulation of the disk flux by the Keplerian motion of the localized hot regions termed as 'blobs'.  Lehr, Wagoner and Wilms (2000) have developed a model to compute by Monte Carlo simulations, the energy dependence of the QPO amplitude to probe the site of their origin in the accretion disk. They use two components with repeated Compton scattering to produce the high energy X-rays and assume a radial dependence of the disk temperature. They computed the energy dependence of the QPO amplitude for GRS 1915 + 105 and found it to be in agreement with the observation of Morgan and Remillard (1997). Thus they are able to localize the QPO origin in the inner disk. The amplitude of the QPOs will either increase or decrease with energy depending on the region of the disk in which the QPOs are produced and the temperature of the corona and its gradient. It is reasonable to assume that the QPO frequency is related to the dynamical time scale of the blobs and therefore, the LFQPOs observed by us in the 4.4-5.9 Hz from XTE J1817-330 will originate farther out in the accretion disk. In the outer region, the Compton scattering corona will be relatively cooler and the QPO amplitude will decrease with increasing energy. In our case we have detected a marginal increase in the amplitude of the QPOs at the higher energy in the two observations, a decrease in the amplitude in the other two observations and no change in amplitude in one case. This suggests that the site of QPO generation is itself dynamically varying in XTE J1817-330 due to variation in the X-ray luminosity which in turn depends on the accretion rate.\\\\  We have detected the QPOs in the 8-14 keV band but not in the 2-8 keV for the observations of MJD 53790.2 and 53790.3. This is similar to the detection of the QPOs at the high energy and its absence at the low energy in some observations from GRS 1915+105 (Chakraborti and Manickam 2000). This behavior of GRS 1915+105 was explained by Chakrabarti and Manickam (2000) on the basis of \"on\" and \"off\" (burst and quiescent) state of the source with the shock oscillation model. Further detailed studies of the LFQPOs at the higher energy are required to test the validity of the model. \\\\"
    },
    "0911/0911.3448_arXiv.txt": {
        "abstract": "Weak gravitational lensing provides a sensitive probe of cosmology by measuring the mass distribution and the geometry of the low redshift universe. We show how an all-sky weak lensing tomographic survey can jointly constrain different sets of cosmological parameters describing dark energy, massive neutrinos (hot dark matter), and the primordial power spectrum. In order to put all sectors on an equal footing, we introduce a new parameter $\\beta$, the second order running spectral index. Using the Fisher matrix formalism with and without CMB priors, we examine how the constraints vary as the parameter set is enlarged. We find that weak lensing with CMB priors provides robust constraints on dark energy parameters and can simultaneously provide strong constraints on all three sectors. We find that the dark energy sector is largely insensitive to the inclusion of the other cosmological sectors. Implications for  the planning of future surveys are discussed. ",
        "introduction": "In the last few decades, a wealth of cosmological data (from large scale structure \\citep[][2dF]{Peacock:2005}; the cosmic microwave background \\citep[][WMAP5]{Komatsu:2009}; supernovae (\\citealt[][SNLS]{Astier:2006}; \\citealt[][ESSENCE]{Miknaitis:2007}; \\citealt{Kowalski:2008a}); weak lensing \\citep{Schrabback:2009}) has revolutionised our vision of the Universe.  In this concordance cosmology, initial quantum fluctuations are believed to have seeded dark matter perturbations in which the Large Scale Structure we observe today has formed. Within this concordance model the Universe is composed only of a small proportion of baryons (4\\%), the rest being dark matter (25\\%, which can be hot or cold) and dark energy.  One of the main challenges today is to understand the nature of the mysterious dark energy which causes cosmic acceleration and constitutes 75\\% of the Universe's energy density \\citep{DETF, Peacock:2006}.  There exists a wealth of potential models for dark energy. To distinguish these models the determination of the dark energy equation of state $w$ has gained importance since some models can result in very different expansion histories. Current data can constrain the dark energy equation of state $w$ to 10\\%, with the assumption of flatness, but a percentage level sensitivity as well as redshift evolution information are required in order to understand the nature of dark energy.  Future cosmic shear surveys show exceptional potential for constraining the dark energy equation of state $w(z)$  \\citep{DETF, Peacock:2006} and have the advantage of directly tracing the dark matter distribution (see \\citealt{Hoekstra:2008} for a review).  In fact, cosmic shear surveys have the potential to constrain all sectors of our cosmological model.  As shear measurements depend on the initial seeds of structure, it can be used to probe the slope and running of the initial power spectrum \\citep[see e.g.][]{Liu:2009} which is central to our understanding of the inflationary model \\citep[see e.g.][]{Hamann:2007}. Shear measurements have also been used to complement neutrino constraints from particle physics  (\\citealt{Tereno:2008}; \\citealt*{Ichiki:2008}) and galaxy surveys \\citep*{Takada:2006}, and future weak lensing surveys will provide bounds on the sum of the neutrino masses, the number of massive neutrinos and the hierachy (\\citealt*{Hannestad:2006, Kit2008b}; \\citealt{debernardisdraft2009}).  The parameters that describe each sector are degenerate in the cosmic shear power spectrum, which means fixing parameters in one sector may results in anomalies being detected in another.  We think all sectors should therefore be considered simultaneously. As each sector provides information about a different area of physics, we also argue they should also be considered on equal footing, i.e. have roughly the same number of parameters describing them, so that one sector is not favoured. Indeed, evidence for a departure from $\\Lambda$CDM may come from any sector.  In section 2, we describe the cosmological parameter set we consider, which includes the three sectors of dark energy, initial conditions and neutrinos. We introduce a new parameter to include the second order running of the initial power spectrum. We also present the weak lensing tomography and the Fisher matrix forecast methods, and briefly discuss systematic effects.  In section 3, we present the weak lensing constraints we expect from future space based surveys with and without \\textit{Planck} priors and investigate the stability of the results as new parameters are added to the analysis. In section 4, we optimise such a survey using the Figure of Merit as well as constraints on all sectors.  In section 5, we present our conclusions.  \\section[]{Method}  \\subsection{Cosmology} \\label{Cosmology}  We start by describing the 12-parameter cosmological model which includes dark energy, dark matter (hot and cold) and initial conditions sectors. Throughout this paper, we work within a Friedmann-Robertson-Walker cosmology. Our cosmological model contains baryonic matter, cold dark matter (CDM) and dark energy, to which we add massive neutrinos (i.e. hot dark matter -- HDM). We also consider different parametrizations of the primordial power spectrum (defined in section \\ref{Power_spectrum}.  We allow for a non-flat geometry by including a dark energy density parameter $\\Omega_\\mr{DE}$ together with the total matter density $\\Omega_m$, such that in general $\\Omega_m+\\Omega_\\mr{DE}\\neq 1$. The dynamical dark energy equation of state parameter, $w=p/\\rho$, is expressed as function of redshift and is parametrized by  a first-order Taylor expansion in the scale factor $a$ \\citep{ChevPol2001, Linder:2003}: \\bq w(a)=w_n + (a_n -a)w_a,\\eq where $a=(1+z)^{-1}$. The pivot redshift corresponding to $a_n$ is the point at which $w_a$ and $w_n$ are uncorrelated.   Our most general parameter space consists of: \\begin{enumerate} \\item Total matter density -- $\\Omega_m$ (which includes baryonic matter, HDM and CDM) \\item Baryonic matter density -- $\\Omega_b$ \\item Neutrinos (HDM) -- $m_\\nu$ (total mass), $N_\\nu$ (number of  massive neutrino species) \\item Dark energy parameters -- $\\Omega_{\\mathrm{DE}},\\; w_0,\\; w_a$ \\item Hubble parameter -- $h$ \\item Primordial power spectrum parameters -- $\\sigma_8$ (amplitude), $n_s$ (scalar spectral index), $\\alpha$ (running scalar spectral index), $\\beta$ (defined in section \\ref{Power_spectrum}) \\end{enumerate} We shall refer to this fiducial cosmology as `$\\nu\\mr{QCDM+\\alpha+\\beta}$'. We choose fiducial parameter values based on the five-year WMAP results \\citep{WMAP5} similar to those used in \\citet{Kit2008b}. The values are given in Table \\ref{Cosmo_models}.  \\setcounter{table}{0} \\begin{table*} \\caption{Cosmological parameter sets used in our calculations. For each parameter set, the ticks ($\\checkmark$) and crosses ($\\times$) indicate whether a parameter is allowed to vary or not, respectively.} \\label{Cosmo_models} \\begin{tabular}{lcccccccccccc} Parameters & $w_0$\t&$w_a$&$\\Omega_\\mr{DE}$&$\\Omega_m$&$ \\Omega_b$&$h$&$\\sigma_8$&$n_s$&$\\alpha$&$\\beta$&$m_\\nu$&$N_\\nu$ \\\\ \\hline Fiducial values & $-0.95$&$0$&$0.7$&$0.3$&$ 0.045$&$0.7$&$0.8$&$1$&$0$\\quad&$0$\\quad&$0.66$&$3$\\\\ \\hline QCDM & $\\checkmark$ & $\\checkmark$ & $\\checkmark$ & $\\checkmark$ & $\\checkmark$ & $\\checkmark$ & $\\checkmark$ & $\\checkmark$ & $\\times$ & $\\times$ & $\\times$ & $\\times$ \\\\ $\\mr{QCDM}+\\alpha$ & $\\checkmark$ & $\\checkmark$ & $\\checkmark$ & $\\checkmark$ & $\\checkmark$ & $\\checkmark$ & $\\checkmark$ & $\\checkmark$ & $\\checkmark$ & $\\times$ & $\\times$ & $\\times$ \\\\ $\\mr{QCDM}+\\alpha+\\beta$ & $\\checkmark$ & $\\checkmark$ & $\\checkmark$ & $\\checkmark$ & $\\checkmark$ & $\\checkmark$ & $\\checkmark$ & $\\checkmark$ & $\\checkmark$ & $\\checkmark$ & $\\times$ & $\\times$ \\\\ $\\nu$QCDM & $\\checkmark$ & $\\checkmark$ & $\\checkmark$ & $\\checkmark$ & $\\checkmark$ & $\\checkmark$ & $\\checkmark$ & $\\checkmark$ & $\\times$ & $\\times$ & $\\checkmark$ & $\\checkmark$ \\\\ $\\nu\\mr{QCDM}+\\alpha$ & $\\checkmark$ & $\\checkmark$ & $\\checkmark$ & $\\checkmark$ & $\\checkmark$ & $\\checkmark$ & $\\checkmark$ & $\\checkmark$ & $\\checkmark$ & $\\times$ & $\\checkmark$ & $\\checkmark$ \\\\ $\\nu\\mr{QCDM}+\\alpha+\\beta$ & $\\checkmark$ & $\\checkmark$ & $\\checkmark$ & $\\checkmark$ & $\\checkmark$ & $\\checkmark$ & $\\checkmark$ & $\\checkmark$ & $\\checkmark$ & $\\checkmark$ & $\\checkmark$ & $\\checkmark$ \\\\ \\end{tabular} \\end{table*}   \\subsection{Matter power spectrum} \\label{Power_spectrum}  The matter power spectrum is defined as: \\bq \\langle \\delta(\\mathbf{k})\\delta^\\ast(\\mathbf{k}')\\rangle= {(2\\pi)}^3\\delta_D^3(\\mathbf{k}-\\mathbf{k}') P(k) \\eq and can be modelled by: \\bq P(k,z)=\\frac{2\\pi^2}{k^3}A_sk^{n_s(k)+3}{T^2(k,z)}\\left(\\frac{D(z)}{D(0)}\\right)^2,\\eq where $A_s$ is the normalisation parameter, $T(k,z)$ is the transfer function and $D(z)$ is the growth function. The primordial spectral index is denoted by $n_s(k)$, and can depend on the scale $k$.  In our cosmological model, the shape of the primordial power spectrum is of particular interest, since it may mimic some of the small-scale power damping effect of massive neutrinos. In the concordance model, the primordial power spectrum is generally parametrized by a power-law \\citep[see e.g.][]{Kosowsky:1995, Bridle:2003} \\bq\\mathcal{P}_\\chi(k)=A_s\\left(\\frac{k}{k_{s0}}\\right)^{n_s-1}.\\eq  We parametrize the running of the spectral index by using a second-order Taylor expansion of $\\mathcal{P}_\\chi$ in log-log space, defining the running as $\\alpha = \\mr{d} n_s/\\mr{d}\\ln k |_{k_0}$, so that the primordial power spectrum is now scale-dependent, with the scalar spectral index defined by \\citep{Spergel:2003, Hannestad:2002} \\bq n_s(k)=n_s(k_0)+\\frac{1}{2}\\frac{ \\mr{d} n_s}{ \\mr{d} \\ln k}\\bigg|_{k_0} \\ln\\left(\\frac{k}{k_0}\\right),\\eq where $k_0$ is the pivot scale.  We use a fiducial value of ${k_0=0.05 \\mr{Mpc^{-1}}}$ for the primordial power spectrum pivot scale.  Although it is motivated by simplicity and standard slow-roll inflation theory, the second-order truncated Taylor expansion is limited and may lead to incorrect parameter estimation \\citep*[see][]{Abazajian:2005, Leach:2003}. In order to test this, we allow an extra degree of freedom in the primordial power spectrum by adding a third-order term in the Taylor expansion, which we call $\\beta:$ \\bq n_s(k)=n_s(k_0)+\\frac{1}{2!}\\alpha\\ln\\left(\\frac{k}{k_0}\\right)+\\frac{1}{3!}\\beta \\ln\\left(\\frac{k}{k_0}\\right)^2,\\eq where $\\beta=\\mr{d}^2 n_s/\\mr{d} \\ln k^2 |_{k_0}$.  We use the \\cite{EH97} analytical fitting formula for the time-dependent transfer function to calculate the linear power spectrum, which includes the contribution of baryonic matter, cold dark matter, dark energy and massive neutrinos, with the modification in the transfer function suggested by \\citet*{Kiakotou:2008jk}. We use the \\citet{Smithetal} correction to calculate the non-linear power spectrum. The matter power spectrum is normalised using $\\sigma_8$, the root mean square amplitude of the density contrast inside an $8\\,h^{-1}\\mr{Mpc}$ sphere.  Following \\citet{EH97}, we assume $N_\\nu$, the number of massive (non-relativistic) neutrino species, to be a continuous variable, as opposed to an integer.  Neutrino oscillation experiments do not, at present, determine absolute neutrino mass scales, since they only measure the difference in the squares of the masses between neutrino mass eigenstates \\citep{Quigg:2008}. Cosmological observations, on the other hand, can constrain the neutrino mass fraction, and can distinguish between different mass hierarchies (see \\citealt*{Elgaroy:2005} for a review of the methods). The \\citeauthor{EH97} transfer function assumes a total of three neutrino species (i.e. $N_\\mr{massless}+N_\\nu=3$), with degenerate masses for the most massive eigenstates, i.e. if $m_\\nu$ is the total neutrino mass, then \\bq m_\\nu=\\sum^{N_\\nu}_{i=0}m_i=N_\\nu m_i, \\eq where $m_i$ is the same for all eigenstates. Thus, $N_\\nu=2$ for the normal mass hierarchy, and $N_\\nu=1$ for the inverted mass hierarchy, while $N_\\nu=3$ corresponds to the case where all three neutrino species have the same mass \\citep[see][]{Quigg:2008}. The temperature of the relativistic neutrinos is assumed to be equal to $(4/11)^{1/3}$ of the photon temperature \\citep{KolbTurner1990}.  Dark energy affects the matter power spectrum in three ways. Its density $\\Omega_\\mr{DE}$ changes the normalisation and $k_{eq}$, the point at which the power spectrum turns over. $\\Omega_\\mr{DE}$ and the dark energy equation of state parameter $w$ change the growth factor at late times by changing the Hubble rate. In addition to this, for departures from a cosmological constant the shape of the matter power spectrum on large scales is affected through dark energy perturbations. In all our calculations, we only consider small deviations from $w_0=-1$, and so neglect dark energy perturbations, only considering the first two mechanisms.  In comparing parameter constraints in different parameter spaces, we shall use six parameter sets. We start with the simplest set (QCDM) to which we add neutrino and additional primordial power spectrum parameters. In all cases the central fiducial model (given in Table \\ref{Cosmo_models}) is the same, but the number of parameters marginalized over varies.  \\subsection{Weak lensing tomography}  The cosmological probes considered in this paper are tomographic cosmic shear and the CMB. In weak lensing surveys, the observable is the convergence power spectrum. In our analysis, we calculate this quantity from the matter power spectrum via the lensing efficiency function. Our convergence power spectrum therefore depends on the survey geometry and on the matter power spectrum. We use the power spectrum tomography formalism by \\citet{Hu:2004}, with the background lensed galaxies divided into 10 redshift bins. Cosmological models are then constrained by the power spectrum corresponding to the cross-correlations of shears within and between bins. The 3D power spectrum is projected onto a 2D lensing correlation function using the \\citet{Limber:1953} equation: \\bq C_\\ell^{ij}=\\int \\mr{d}z \\frac{H}{D^2_A} W_i(z)W_j(z)P(k=\\ell/D_A,z),\\eq where $i$, $j$ denote different redshift bins. The weighting function $W_i(z)$ is defined by the lensing efficiency: \\bq W_i(z)=\\frac{3}{2}\\Omega_m \\frac{H_0}{H}\\frac{H_0 D_{OL}}{a}\\int_z^\\infty \\mr{d} z^{\\prime}\\frac{D_{LS}}{D_{OS}}P(z^{\\prime}), \\eq where the angular diameter distance to the lens is $D_{OL}$, the distance to the source is $D_{OS}$, and the distance between the source and the lens is $D_{LS}$ (see \\citealt{Hu:2004} for details). Our multipole range is $10<\\ell<5000$.  The galaxies are assumed to be distributed according to the following probability distribution function \\citep*{Smail1994}: \\bq P(z)=z^a \\exp\\left[-\\left(\\frac{z}{z_0}\\right)^b\\right],\\eq where $a=2$ and $b=1.5$, and $z_0$ is determined by the median redshift of the survey $z_m$ \\citep[see e.g][]{AR2007}.  Our survey geometry follows the parameters for a `wide' all-sky survey with $A_s=20 000\\,\\mr{sq\\; degrees}$. The survey parameters are shown in Table \\ref{survey_params}. The median redshift of the density distribution of  galaxies is $z_\\mr{median}$ and the observed number density of galaxies is $n_g$. We include photometric redshift errors $\\sigma_z(z)$ and intrinsic noise in the observed ellipticity of galaxies $\\sigma_\\epsilon$. We follow the definition $\\sigma_\\gamma^2=\\sigma_\\epsilon^2$, where $\\sigma_\\gamma$ is the variance in the shear per galaxy \\citep[see][]{Bartelmann:2001}.  \\setcounter{table}{1} \\begin{table} \\begin{center} \\caption{Fiducial parameters for the all-sky weak lensing survey considered.} \\label{survey_params} \\begin{tabular}{@{}lr@{}} \\hline $A_s$/sq degree &20 000\\\\ $z_\\mr{median}$&0.9\\\\ $n_g/\\mr{arcmin}^{2}$&35\\\\ $\\sigma_z(z)/(1+z)$&0.025\\\\ $\\sigma_\\epsilon$&0.25\\\\ \\hline \\end{tabular} \\end{center} \\end{table}   \\subsection{Error forecast}  The predictions for cosmological parameter errors presented in this paper use the Fisher matrix formalism. The Fisher matrix gives us the lower bound on the accuracy with which we can estimate model parameters from a given data set \\citep*{Fisher1935, Tegmark:1997, Kitching:2009}. In calculating forecast survey errors, we are implicitly making assumptions about the parameter set \\citep*[see the discussion of nested models in][]{Heavens:2007}.  We want to know whether our constraints are robust against variations in the parameterisation of the cosmological model. This is of particular importance when dark energy constraints are considered, because of the degeneracies with other parameters (see Fig. \\ref{pk_deg}).  The forecast parameter precision is improved by combining independent experiments. Using this technique, joint constraints or error forecasts can be obtained by combining weak lensing with other observational techniques.  The Fisher matrix for the shear power spectrum is given by \\citep{Hu:2004}: \\bq\\label{eq_10} F_{\\alpha\\beta} = f_\\mr{sky}\\sum_\\ell{\\frac{(2\\ell+1)\\Delta \\ell}{2}}\\mr{Tr}\\left[D_{\\ell\\alpha}\\widetilde{C}_\\ell^{-1}D_{\\ell\\beta}\\widetilde{C}_\\ell^{-1}\\right],\\eq where the sum is over bands of multipole $\\ell$ of width $\\Delta \\ell$, $\\mr{Tr}$ is the trace, and $f_\\mr{sky}$ is the fraction of sky covered by the survey. Equation \\ref{eq_10} assumes the likelihod obeys a Gaussian distribution with zero mean.  The observed power spectra for each pair $i,j$ of redshift bins are written as the sum of the lensing and noise spectra: \\bq \\widetilde{C}_\\ell^{ij}=C_\\ell^{ij}+N_\\ell^{ij}.\\eq The derivative matrices are given by \\bq [D_{\\alpha}]^{ij}=\\frac{\\partial C_\\ell^{ij}}{\\partial p_\\alpha} ,\\eq where $p_\\alpha$ is the vector of parameters in the theoretical model.  In order to quantify the potential for a survey to constrain dark energy parameters, we use the Figure of Merit, as defined by the Dark Energy Task Force \\citep{DETF}: \\bq \\mr{FoM}=\\frac{1}{\\Delta w_n\\Delta w_a}.\\eq  \\subsection{\\textit{Planck} priors}  In this article, together with lensing-only constraints, we also include joint lensing and CMB constraints. To combine constraints from different probes we add the respective Fisher matrices: $\\bf{F}_\\mr{joint}=\\bf{F}_\\mr{lensing}+\\bf{F}_\\mr{CMB}$. For our CMB priors, we use the forthcoming \\textit{Planck} mission as our survey. The \\textit{Planck} Fisher matrix is calculated following \\citet{Rassat:2008}, which estimates errors using information from the temperature and E mode polarisation (i.e. TT, EE, and TE), with the help of the publicly available \\textsc{camb} code \\citep*{Lewis:2000}. We conservatively do not use information from B modes, and only use the 143~GHz channel, assuming other frequencies will be used for the foreground removal. This is conservative compared to other \\textit{Planck} priors in the literature. More details are given in \\citet[][Appendix B]{Rassat:2008}. The full parameter set for the \\textit{Planck} calculation is: $\\{\\Omega_\\mr{DE},\\,w_0,\\,w_a,\\,\\Omega_m,\\,\\Omega_b,\\,m_\\nu,\\, N_\\nu,\\,h,\\,\\sigma_8,\\,n_s,\\,\\alpha, \\,\\tau\\}$. We use the same central values as for our weak lensing calculations, as described above, with a fiducial value for the reionisation optical depth $\\tau=0.09$ and subsequently marginalize over $\\tau$. We consider neutrino parameters as $\\Omega_\\nu h^2$ and $N_\\nu$ and use a Jacobian to translate this into constraints on $m_\\nu$ and $N_\\nu$.    \\subsection{Systematic effects} Weak lensing measurements are affected by systematic effects which reduce the precision on cosmological parameters and introduce bias (see \\citealt{Refregier:2003a} and \\citealt{Schneider:2006} for a review). In this section we discuss some of these effects.  Intrinsic correlations can contaminate the lensing signal. Solutions include using tomography or 3D lensing, which decouple the long-distance line-of-sight effects from the the physical proximity of the galaxies \\citep[see e.g.][]{King:2003, Bridle:2007, Kit2008b}. Using this approach, a nulling technique for shear-intrinsic ellipticity has been proposed by  \\citet{Joachimi:2008, Joachimi:2009}.  Measurement systematics are due to PSF effects \\citep[see][]{KSB1995}. The results presented in \\citet{Bridle:2009} indicate that the required  accuracy will be reached for the next generation of all-sky weak lensing surveys.  Redshift distribution systematics, which lead to an uncertainty in the median galaxy redshift, cause an uncertainty in the amplitude of the matter power spectrum \\citep{Hu:1999b}. The problem of photometric redshift systematics can be met given a number of galaxies in the spectroscopic calibration sample of $10^4-10^5$ \\citep{Ma:2006, AR2007}.  Finally there are theoretical uncertainties on the shape of the matter power spectrum. The existing non-linear corrections to the matter power spectrum \\citep{Peacock:1996ys, Ma:1998} are only accurate to about $10\\%$ and disagree with one another to this level in the non-linear r\\'egime \\citep[see][]{Huterer:2001}. Newer prescriptions such as those by \\citet{Smithetal} offer more accurate predictions particularly for non-$\\Lambda\\mr{CDM}$ cosmological models. The error in the non-linear part may still be significant  if effect of massive neutrinos is included \\citep*[see e.g.][]{Saito:2008}. Current semi-analytical models need to be improved to match the degree of statistical accuracy expected for future weak lensing surveys. The solution is to run a suite of $N$-body ray-tracing simulations \\citep[see e.g.][]{White:2004, Huterer:2005, Hilbert:2009, Sato:2009, Teyssier:2009}. ",
        "conclusions": "The main aim of this paper was to put three different sectors of the cosmological model: dark energy, dark matter (cold and neutrinos) and initial conditions on equal footing while forecasting constraints for a future weak lensing survey. To do this we introduce a new parameter $\\beta$, which models the second order running spectral index. We have forecast errors for an all-sky tomographic weak lensing survey, and for weak lensing+\\textit{Planck},  using different cosmological parameter sets, studying the effect of the addition of parameters in the model. We have shown that error forecasts for some parameters are stable against changes in the parameter set (Table \\ref{table_lensing}), and that degeneracies between the dark energy parameters $w_0$ and $w_a$ are not significantly affected by the addition of parameters (Fig. \\ref{ellipses_article_joint}). We have investigated the shift in the optimal pivot scale for primordial power spectrum constraints with the addition of extra parameters (Fig. \\ref{Prim_v_k0}).  We have shown that parameter constraints are nevertheless dependent on the parameterisation of the primordial power spectrum. With the addition of CMB priors, we have shown that we can obtain improved constraints in this sector, which reduces the dependency on constraints on the parameter set.  We have obtained predicted constraints for both the total neutrino mass and the number of massive neutrino species. This article has a resonance with \\citet{debernardisdraft2009}, in which it is shown that the neutrino mass hierarchy can be constrained using cosmic shear, using a more general parameterisation of the neutrino mass splitting. We find that for the parameters that are common between the two articles there is an agreement between the predicted errors, despite the slightly different parameter sets and assumptions. The results obtained here also mirror those in \\citet{Zunckel:2007}, in which it is found that the neutrino mass constraints are degraded when the hypothesis space is enlarged. This is due to degeneracies between the neutrino mass and other parameters. Adding priors only tightens the constraints when the additional information comes from independent experiments,  which reduces the freedom in the degenerate parameters. We have used information from the CMB, which constrains the primordial power spectrum particularly well, resulting in improved constraints on our neutrino parameters.  In the neutrino sector, parameter constraints would be improved by a hierarchical parameterisation in which different neutrino species have non-degenerate masses. This would model more accurately the process whereby each massive species becomes non-relativistic at a different redshift. While the different transition redshifts have only a very small effect on the CMB anisotropy power spectrum, the effect is non-negligible in future cosmic shear experiments which measure the matter power spectrum to a sufficient accuracy to discriminate between different mass hierarchies.  This article concludes that a future all-sky weak lensing survey with CMB priors provides robust constraints on dark energy parameters and can simultaneously provide strong constraints on all parameters. The results presented here show that error forecasts from our weak lensing survey are stable against the addition of parameters to the fiducial model, and that this stability is improved by adding CMB priors."
    },
    "0911/0911.3397_arXiv.txt": {
        "abstract": "We report on a new class of fast-roll inflationary models.  In a huge part of its parameter space, inflationary perturbations exhibit quite unusual phenomena such as scalar and tensor modes freezing out at widely different times, as well as scalar modes reentering the horizon during inflation.  One specific point in parameter space is characterized by extraordinary behavior of the scalar perturbations. Freeze-out of scalar perturbations as well as particle production at horizon crossing are absent. Also the behavior of the perturbations around this quasi-de Sitter background is dual to a quantum field theory in flat space-time. Finally, the form of the primordial power spectrum is determined by the interaction between different modes of scalar perturbations. ",
        "introduction": "Inflation has become the standard paradigm for the description of the early universe, solving the monopole, horizon and flatness problems as well as providing for a mechanism to seed structure in the universe. Various explicit models for this period of exponential expansion, most relying on the dynamics of real scalar ``inflaton'' fields, mostly agree with current observations of perturbations of an isotropic background, such as cosmic microwave background (CMB) fluctuations and large scale structure (LSS) (although see \\cite{Copi:2008hw} and references therein). When analyzing these models, we mostly rely on the assumption of ``slow-roll'' \\cite{Stewart:1993bc, Sasaki:1995aw, Martin:2002vn} -- the slow evolution of the inflaton field -- although fast-roll models have also been considered \\cite{Linde:2001ae}. Slow roll is an approximation made to solve the perturbation equations, but not an absolute requirement for an extended period of accelerated expansion. Even though the scalar spectral index has arguably been measured by the WMAP satellite\\cite{Dunkley:2008ie} to be smaller than unity, $n_s\\approx0.96$, this constrains the slow-roll parameters \\begin{eqnarray} \\epsilon\\equiv-\\frac{\\partial_t{H}}{H^2}\\,, \\eta\\equiv2\\frac{\\partial_\\phi^2H}{H}\\,, \\zeta^2=\\frac{1}{2}\\frac{\\partial_\\phi H \\partial_\\phi^3H}{H^2}\\,, \\end{eqnarray} all in reduced Planck units $M_p^2=\\frac{1}{8\\pi G}=1$, only if they are small. This was exactly the assumption made when solving the perturbation equations, be it using Hankel functions \\cite{Stewart:1993bc}, the $\\Delta N$ formalism \\cite{Sasaki:1995aw, Starobinsky:1986fxa}, or the WKB approximation \\cite{Starobinsky:1982ee, Martin:2002vn}. If the slow-roll parameters are large, then it is impossible to make any general statement about their values from the measurement of $n_s$.  While it is true that $0<\\epsilon<1$ for successful inflation, higher order slow-roll parameters may be large, contrary to intuition. Besides exploring unknown corners of inflationary model space, there are many additional motivations to study models with large $\\eta$. For example, string theory motivated models of inflation are typically characterized by large values of $\\eta$ \\cite{Kachru:2003sx}. Also, inflationary non-Gaussianities are suppressed by powers of the slow-roll parameters \\cite{Maldacena:2002vr}, so that one might expect large non-Gaussianities in fast-roll inflationary models.  As the problem of analyzing the behavior of fast-roll inflationary models is technically very challenging, so far those models received only limited attention. Essentially, there two ways to study this largely unexplored realm. One of them is finding exact solutions to a coupled system of non-linear differential equations. Another approach is to solve these equations numerically. In our paper, we adopt both approaches, providing an exact solution for a subset of initial conditions as well as numerically charting the phase space of a new class of inflationary models with large acceleration, i.e.  large $\\eta$. This allows us to study the behavior of both background and perturbations in new and interesting regimes, where we find various unexpected and even downright bizarre phenomena.  While our model can be made to agree with observations for a small range of parameters, it possesses several most unusual features for parameter ranges that are not compatible with observations. Some of these features might seem counter-intuitive.  For example, according to general lore, perturbation modes that exit the Hubble horizon, freeze out and remain classical until the present epoch. We find that is not true for a particular model of this family. The modes remain quantum mechanical at all times and wavelengths, even long after they cross the Hubble scale. Also, it is generally believed that the interactions among the different modes of scalar perturbations are suppressed by powers of the slow roll parameters, so that only insignificant amounts of non-Gaussianity are created during single-field inflation. Instead we show that it is the interaction between different modes that defines the very form of the scalar power spectrum. Other unusual features include the fact that tensor and scalar modes can freeze out at different times and even unfreeze during inflation. Finally, it turns out that for a specific member of this family, the theory of scalar perturbations is described by a quantum field theory in flat space time.  \\begin{figure*}[t!] (a)\\includegraphics[width=0.45\\linewidth]{V_of_phi_p0_04} (b)\\includegraphics[width=0.45\\linewidth]{V_of_phi_p-2} \\caption{(a)$V(\\phi)$ for $p=0.04$. Inflation takes place close to the top of the potential, rolling down from $\\phi\\approx0$ towards larger values of $|\\phi|$. The potential is very shallow and turns negative only for $\\phi>\\frac{\\sqrt{24}}{0.04}\\approx122$, i.~e.~about $30$ e-folds after the end of inflation. For this model, the power spectrum of scalar perturbations has a spectral index that is compatible with observations, $n_s=1-p=0.96$. (b)$V(\\phi)$ for $p=-2$. For negative $p$, the potential goes to $-\\infty$ for $\\phi\\rightarrow\\infty$. Inflation takes place rolling from small values of $|\\phi|$ towards $\\phi=0$ and ends through tunneling. The scalar power spectrum possess some interesting features, see text.} \\label{fig:potentialp2} \\end{figure*}  Our paper is organized as follows. In section \\ref{sec:most_curious_model} we present an explicit one-parameter family of fast-roll models whose background equation is exactly solvable for certain initial conditions. In section \\ref{sec:linear_perturbations} we discuss the spectrum of scalar and tensor perturbations, which can be reliably computed despite the fact that the slow-roll parameter $\\eta$ may be larger than unity. In section \\ref{sec:p-2} we present an in-depth analysis of the perturbations for a special member of the family of models, where the scalar perturbations do not follow the usual behavior of oscillating, being stretched to the horizon, and freeze-out, but instead obey an exactly massless harmonic oscillator type equation. In section \\ref{sec:conclusions} we conclude and offer some ideas for future work on this and similar classes of models. ",
        "conclusions": "\\label{sec:conclusions}  In this paper we discuss the behavior of the family of inflationary models \\begin{eqnarray} \\label{eq:Vconclusion} V&=&H_0^2M_P^2 e^{-\\frac{p}{4}\\frac{\\phi^2}{M_P^2}}\\left(3-\\frac{p^2}{8}\\frac{\\phi^2}{M_P^2}\\right)\\,, \\end{eqnarray} with arbitrary $p$ and $H_0<M_{\\mathrm{pl}}$.  Exploring the phase space of this family of models for large $|p|$ reveals the existence of a non-slow-roll attractor which at late times can be approximated by \\begin{eqnarray} \\dot{\\phi}&=&\\frac{p}{2}H_0\\phi\\,. \\label{eq:nonslowrollattrlate} \\end{eqnarray} For small values of $|p|$ this attractor coincides with the slow-roll attractor, see equation (\\ref{eq:attractors}).  For all values of $|p|$, this family of models is observationally ruled out. However, they turn out to exhibit many fascinating features, some quite counterintuitive. The reason why the standard lore based on the slow roll approximation fails in this case is that the value of $p$ essentially determines the value of the slow roll parameter $\\eta=2\\frac{\\partial_\\phi^2 H}{H}$, so that larger values of $|p|$ correspond to larger values of $\\eta$, see equation (\\ref{eq:eta_p}) .  Table \\ref{tab:p_overview} show the physical properties of our family of models for different values of the parameter $p$. \\begin{table} \\begin{tabular}{|c|c|} \\hline {\\bf value of} $\\mathbf{p}$ & {\\bf model features}\\\\ \\hline $-\\infty < p < -6$ & widely separated freeze-out\\\\ & $\\infty$ long inflation for one initial condition\\\\ & $|\\eta|>1$ \\\\ \\hline $-6<p<-4$ & widely separated freeze-out\\\\ & $|\\eta|>1$\\\\ & $\\infty$ long classical inflation\\\\ \\hline $-4<p<-2$ & no freeze out\\\\ & $|\\eta|>1$\\\\ & $\\infty$ long classical inflation\\\\ \\hline $p=-2$ &  no freeze out\\\\ & no gravitational particle production\\\\ & no stochastic eternal inflation\\\\ & non-suppressed mode interactions\\\\ & (pre)thermalization\\\\ & dual to a QFT in flat space\\\\ & $|\\eta|\\approx 1$\\\\ & $\\infty$ long classical inflation\\\\ \\hline $-2<p<-2+\\sqrt{3}$ & modes reenter the horizon\\\\ &$\\infty$ long classical inflation\\\\ \\hline $-2+\\sqrt{3}<p<0$& $\\infty$ long classical inflation\\\\ \\hline $p=0$ & exact de Sitter space-time\\\\ \\hline $0<p<0.01$ & finite duration of inflation\\\\ \\hline $0.01<p<0.07$ & compatible with observations\\\\ \\hline $0.07<p<\\infty$ & finite duration of inflation\\\\ \\hline \\end{tabular} \\caption{Behaviour of our model $V=H_0^2 e^{-\\frac{p}{4}\\phi^2}(3-\\frac{p^2}{8}\\phi^2)$ for different values of $p$.} \\label{tab:p_overview} \\end{table} The most interesting physics takes place in the model with $p=-2$ (see the Fig. 2b). First of all, to leading order in a $1/a$ expansion, the scalar perturbations in this model turn out to be described by a quantum field theory in flat spacetime (\\ref{eq:Sadm}) . The linear scalar perturbations are free massless harmonic oscillators and thus never freeze out\\footnote{In fact, models with $-4<p<-2$ are also characterized by the absence of freeze out. Here, the scalar perturbations are harmonic oscillators with time-varying but strictly positive masses. We leave a thorough examination of this issue to future work.}. Gravitational particle production, stochastic kicks uplifting the background value of the inflaton field and a regime of stochastic eternal inflation are completely absent.  If the initial state for the scalar fluctuations $u_k=a\\delta\\phi_k$ has a non-zero occupation number for a range of momenta, one finds that the corresponding occupation numbers $n_k$ typically approach a spectrum of form (\\ref{eq:WeakTurbSpectrum}) due to interactions between the modes $u_k$, with a spectral index $n_s=\\frac{3}{2}$. Interestingly, the effects of the interaction between the different modes $u_k$ do not get redshifted away.  As a curious aside, we notice that the shape of the potential (\\ref{eq:potential}) is strongly reminiscent of the SUGRA form \\begin{eqnarray} V_{\\mathrm{SUGRA}}&=&e^K\\left(K^{ij} D_iW D_{\\bar{j}}\\bar{W} - 3|W|^2\\right)\\, , \\end{eqnarray} where the K\\\"ahler potential $K$ is a real function of the complex fields $\\phi^i, \\phi^{\\bar{i}}$, the superpotential $W$ is a holomorphic function of the complex fields $\\phi^i$, the covariant derivative $D_iW=\\partial_{\\phi^i} W + \\partial_{\\phi^i}K W$ and $K^{i\\bar{j}}=(\\partial_{\\phi^{i}\\phi^{\\bar{j}}} K)^{-1}$. Assuming the K\\\"ahler potential $K=-\\frac{q}{4}\\phi\\bar{\\phi}$ and the superpotential $W=H_0=\\mathrm{constant}$, the potential can be computed to \\begin{eqnarray} V_{\\mathrm{SUGRA}}&=&-H_0^2 e^{-\\frac{q}{4}|\\phi|^2}\\left(3+\\frac{q}{4}|\\phi|^2\\right)\\, . \\end{eqnarray} For $q=-2$, this agrees with (\\ref{eq:Vconclusion}) for $p=-2$ up to an overall sign \\begin{eqnarray} V_{\\mathrm{SUGRA}}^{q=-2}&=&-V^{p=-2}\\, , \\end{eqnarray} possibly indicating that this model may be embedded into string theory.  Finally, we remark that it might be possible to have both modes without freeze-out and a scalar power spectrum with $n_s\\approx{}0.96$ for a wider class of models, where the attractor solution has $z=\\mathrm{constant}$ instead of $z^\\prime=\\mathrm{constant}$. In this case, it is immediately clear that the equations of motion for the perturbations $u_k$ again simplify to those of a free harmonic oscillator. What is not clear is whether in these models inflation actually takes place: the inflationary condition from (\\ref{eq:inflation_condition}) corresponds to a highly non-trivial differential equation for the Hubble parameter $H$. We leave a thorough examination of this class of models for future work."
    },
    "0911/0911.3732_arXiv.txt": {
        "abstract": "{} { We probe the dependence of the electron-to-proton mass ratio, $\\mu = m_{\\rm e}/m_{\\rm p}$, on the ambient matter density by means of radio astronomical observations. } {The ammonia method, which has been proposed to explore the electron-to-proton mass ratio, is applied to nearby dark clouds in the Milky Way. This ratio, which is measured in different physical environments of high (terrestrial) and low (interstellar) densities of baryonic matter is supposed to vary in chameleon-like scalar field models, which predict strong dependences of both masses and coupling constant on the local matter density. High resolution spectral observations of molecular cores in lines of NH$_3$ $(J,K) = (1,1)$, HC$_3$N $J = 2-1$, and N$_2$H$^+$ $J = 1-0$ were performed at three radio telescopes to measure the radial velocity offsets, $\\Delta V \\equiv V_{\\rm rot} - V_{\\rm inv}$, between the inversion transition of NH$_3$ (1,1) and the rotational transitions of other molecules with different sensitivities to the parameter \\dmm $\\equiv (\\mu_{\\rm obs} - \\mu_{\\rm lab})/\\mu_{\\rm lab}$. } {The measured values of $\\Delta V$ exhibit a statistically significant velocity offset of $23\\pm4_{\\rm stat}\\pm3_{\\rm sys}$ \\ms. When interpreted in terms of the electron-to-proton mass ratio variation, this infers that \\dmm = $(2.2\\pm0.4_{\\rm stat}\\pm0.3_{\\rm sys})\\times10^{-8}$. If only a conservative upper bound is considered, then the maximum offset between ammonia and the other molecules is $|\\Delta V| \\leq 30$ \\ms. This provides the most  accurate reference point at $z = 0$ for \\dmm\\ of $|\\Delta \\mu/\\mu| \\leq 3\\times10^{-8}$. } {} ",
        "introduction": "\\label{sect-1}  Testing the variability of dimensionless physical constants is an important topic in contemporary laboratory and astrophysical experiments. The physical constants are not supposed to vary in the Standard Model (SM) of particle physics but do vary quite naturally in grand unification theories, multidimensional theories, and whenever there is a coupling between light scalar fields and baryonic matter. In particular, light scalar fields have been widely discussed in the context of dark energy, since they provide negative pressure that may be responsible for the cosmic acceleration detected at low redshifts, $z<1$ (Caldwell \\etal\\, 1998; Peebles \\&  Rata 2003). If this scalar field does exist, then a question arises: why has it not been detected in local tests of the equivalence principle or fifth force searches? A solution was suggested using the so-called `chameleon' models (Khoury \\& Weltman 2004; Brax \\etal\\ 2004; Mota \\& Shaw 2007). These models assume that a light scalar field acquires both an effective potential and effective mass because of its coupling to matter that strongly depends on the ambient matter density. In this way, this scalar field may evade local tests of the equivalence principle and fifth force experiments since the range of the scalar-mediated fifth force for the terrestrial matter densities is too small to be detected. This is not the case for space-based tests, where the matter density is considerably lower, an effective mass of the scalar field is negligible, and an effective range for the scalar-mediated force is large. The present paper deals with one such astronomical test based on spectral observations of molecular clouds in the Milky Way disk. Additional tests employing polarization of the light from the stars and a modification of the Sunyaev-Zel'dovich effect in the cosmic microwave background due to a coupling between a chameleon-like scalar field and photons are described, respectively, in Burrage \\etal\\ (2009) and Davis \\etal\\ (2009).  Astronomical spectroscopy can probe the physical constants that describe atomic and molecular discrete spectra: the fine-structure constant, $\\alpha = e^2/(\\hbar c)$, and the electron-to-proton mass ratio, $\\mu = m_{\\rm e}/m_{\\rm p}$. In GUTs, these constants mediate the strength of fundamental forces: $\\alpha$ is the coupling constant of the electromagnetic interaction, $m_{\\rm e}$ is related to the vacuum expectation value of the Higgs field, namely to the scale of the weak nuclear force, and $m_{\\rm p}$ is proportional to the quantum chromodynamics scale. We note that the predicted variabilities of the fine-structure constant and the electron-to-proton mass ratio are not independent and that the variations in $\\mu$ may exceed those of $\\alpha$ (Calmet \\& Fritzsch 2002; Langacker \\etal\\ 2002; Dine \\etal\\ 2003; Flambaum \\etal\\ 2004). Thus, a hypothetical variation in $\\mu$ is expected to be easier to detect than that in $\\alpha$.  Despite many efforts, the variability in $\\alpha$ and $\\mu$ has never been detected. Data obtained from high precision frequency measurements in laboratory experiments with atomic clocks and from  astronomical observations provide only upper limits. For example, laboratory experiments have delivered the following results: fractional temporal variations in $\\mu$ are restricted to a level of $\\dot{\\mu}/\\mu = (3.8\\pm5.6)\\times10^{-14}$ yr$^{-1}$ (Shelkovnikov \\etal\\, 2008), and $\\dot{\\mu}/\\mu = (1.6\\pm1.7)\\times10^{-15}$ yr$^{-1}$ (Blatt \\etal\\, 2008), whereas the current level for $\\alpha$ is $\\dot{\\alpha}/\\alpha = (-1.6\\pm2.3)\\times10^{-17}$ yr$^{-1}$ (Rosenband \\etal\\, 2008). Here $\\Delta \\mu$ and $\\Delta \\alpha$ are the relative changes between the values measured at two different epochs.  The most stringent upper limits to \\daa\\, and \\dmm\\, obtained from astronomical observations of extragalactic objects are restricted by a few units of ppm (1ppm = $10^{-6}$). Here $\\Delta \\alpha/\\alpha = (\\alpha' - \\alpha)/\\alpha$, where $\\alpha, \\alpha'$ denote the values of the fine-structure constant in the laboratory and the specific absorption/emission line system of a Galactic or extragalactic object (the same definition is applied to \\dmm). However, there are claims of a variability in both $\\alpha$ and $\\mu$ at the 5$\\sigma$ and 4$\\sigma$ confidence levels, respectively, although they are in contrast to other null results and the whole issue is highly controversial. The current observational status is the following.  From the measurements of the relative radial velocity shifts of different absorption lines (e.g., \\ion{Mg}{ii}, \\ion{Si}{ii}, \\ion{Fe}{ii}, \\ion{Zn}{ii}) of 143 QSO absorption systems obtained with the Keck/HIRES in the redshift range $0.2 < z < 4.2$, Murphy \\etal\\ (2004) claimed that $\\Delta\\alpha/\\alpha = -5.7\\pm1.1$ ppm. On the other hand, Chand \\etal\\ (2004) using VLT/UVES spectra of bright QSOs analyzed 23 absorption systems, which are not in common with those studied by the Keck telescope, and failed to reproduce Murphy \\etal's result. At first, they claimed a very stringent limit of $\\Delta\\alpha/\\alpha = -0.6\\pm0.6$ ppm, but Murphy \\etal\\ (2008a) demonstrated convincingly that the error estimate of Chand \\etal\\ (2004) is underestimated. When these errors are properly accounted for, the new weighted mean becomes $\\Delta\\alpha/\\alpha = -6.4\\pm3.6$ ppm. The new value has a scatter that is larger than the quoted errors implying that there are systematic errors that are comparable or even larger than the statistical ones. However, Srianand \\etal\\ (2008) showed that excluding two systems that deviate from the mean at the 3$\\sigma$ level when reanalyzing VLT/UVES systems leads to $\\Delta\\alpha/\\alpha = 0.1\\pm1.5$ ppm. We note that the only claimed non-zero measurement of \\daa\\ corresponds to an uncertainty in the wavelength scale calibration smaller 100 \\ms. This precise calibration of the Keck/HIRES spectra is difficult to achieve as shown by Griest \\etal\\ (2010). Using iodine exposures to calibrate the normal Th-Ar Keck data pipeline output, they found absolute wavelength offsets of 500 \\ms to 1000 \\ms with drifts of more than 500 \\ms over a single night, and drifts of nearly 2000 \\ms over several nights. Their conclusion is that these systematic uncertainties make it difficult to use Keck/HIRES QSO spectra to constrain the change in the fine structure constant at the $10^{-6}$ level. The most stringent constraint on $|\\Delta\\alpha/\\alpha| < 2$ ppm was found from the \\ion{Fe}{ii} lines at $z = 1.15$ towards the bright quasar HE~0515--4414 (Quast \\etal\\ 2004; Levshakov \\etal\\ 2005; Molaro \\etal\\ 2008).  Bounds on  $\\mu$ variations are most effectively obtained from observations of the Werner and Lyman series of the molecular hydrogen H$_2$ in damped Ly-$\\alpha$ systems (DLA). The electron-vibro-rotational transitions have different dependences on the reduced mass and can be used to constrain the variability of $\\mu$ (Thompson 1975; Varshalovich \\& Levshakov 1993). The measurements of $\\mu$ rely on the same H$_2$ systems observed  with VLT/UVES and claim a variability of $\\Delta\\mu/\\mu = 24\\pm6$ ppm (Reinhold \\etal\\ 2006) or no variability with $-15\\ {\\rm ppm} < \\Delta\\mu/\\mu < 57$ ppm (Levshakov \\etal\\ 2002), $|\\Delta\\mu/\\mu| \\leq 49$ ppm (Wendt \\& Reimers 2008), $\\Delta\\mu/\\mu = 2.6\\pm3.0$ ppm (King \\etal\\ 2008), and $\\Delta\\mu/\\mu = -7\\pm8$ ppm (Thompson \\etal\\ 2009). Radio astronomical observations place constraints on $\\mu$-variations of $|\\Delta \\mu/\\mu| < 1.8$ ppm at redshift $z = 0.68$ (Murphy \\etal\\ 2008b), and $|\\Delta \\mu/\\mu| < 0.6$ ppm at $z = 0.89$ (Henkel \\etal\\ 2009).  One should keep in mind, however, that laboratory experiments and spectra of extragalactic objects probe very different time-scales and different regions of the universe, and the connection between them is quite model dependent (Mota \\& Barrow 2004).  Astronomical measurements of the dimensionless constants are based on the comparison of the line centers in the absorption/emission spectra of astronomical objects and the corresponding laboratory values. To distinguish the line shifts due to the radial motion of the object from those caused by the variability of constants, lines with different sensitivity coefficients, $Q$,  to the variations in $\\mu$ and/or $\\alpha$ should be employed. Unfortunately, the observable optical lines have very close sensitivity coefficients with differences of $|\\Delta Q|$ that do not exceed 0.05. Combined with the resolving power of spectrographs at modern optical telescopes, this small difference leads to an upper limit on \\daa\\, and \\dmm\\, of $\\sim$1 ppm, the optimal value achievable in observations of extragalactic objects with existing facilities. However, one can probe the values of fundamental constants at a considerably more accurate level if, firstly, nearby objects in the Milky Way are observed and, secondly, spectral lines from other frequency ranges, infrared and radio, are employed. This statement is based on the following considerations.  As mentioned above, the scalar field models suggest a coupling between the scalar fields and baryonic matter. This coupling results in a functional dependence of $\\mu$ and $\\alpha$ on $\\rho$, the local matter density (Olive \\& Pospelov 2008). This coupling alone may prevent a positive detection of the variations in $\\mu$ and other dimensionless constants in laboratory studies, since they are performed under the same terrestrial conditions and the effective range of the scalar-field-mediated force is smaller than 1 mm at terrestrial matter densities (Olive \\& Pospelov 2008). On the other hand, the density in cold molecular clouds is only  $10^3 - 10^5$ particles per cm$^3$, i.e., $\\sim 10^{19}$ times lower than in terrestrial environments. Because of this extremely large difference between the matter densities non-zero values of \\dmm\\ and \\daa\\ are predicted. It is very important to compare the above ratio of $\\sim10^{19}$ with the differences in $\\rho$ between molecular clouds themselves, to find that they are negligible, i.e., at most one or two orders of magnitude for a given tracer. In the context of \\dmm\\ (or \\daa) this means that measurements of all dense molecular clouds are identical irrespective of their location in space and time (redshift). Thus, nearby Galactic molecular clouds can be observed to ensure a strong signal, i.e., that the line profiles can be centered with high accuracy.  Molecular lines originating in these clouds are mainly observed in cm and mm radio bands. An advantage of radio observations is that very narrow spectral lines (of line widths $\\sim$100 \\ms) arising in cold molecular cores can be observed with extremely high spectral resolution, FWHM $\\sim25$ \\ms. In optical observations of extragalactic objects, we have in general broader sources, $b \\ga 1$ \\kms, and use lower resolution spectrographs, FWHM $\\ga 4$ \\kms\\ (Levshakov \\etal\\ 2007). Taking into account that the uncertainty in the line position is roughly 1/10th of the pixel size, this infers a gain of about two orders of magnitude in the accuracy of the measurements of \\daa\\ and/or \\dmm\\ if radio spectra of the local interstellar objects are used.  Complementary to this, differences in the sensitivity coefficients for lines from the  microwave, far IR, and radio ranges are much larger than those from optical and UV ranges. For example, the difference between the sensitivity coefficient of the $(J,K)=(1,1)$ inversion transition of ammonia (\\amm) and that of any rotational transition in another molecule is $\\Delta Q = 4.46 - 1 = 3.46$ (Flambaum \\& Kozlov 2007). When compared with UV transitions, this infers a gain of about 70 times in the sensitivity of the line positions to the change in $\\mu$.  Our preliminary study of cold molecular clouds (Levshakov \\etal\\  2008a, hereafter LMK) was based on published results of high quality spectral radio observations obtained with the 100-m Green Bank telescope (Rosolowsky \\etal\\ 2008; Rathborne \\etal\\ 2008) and the 45-m Nobeyama radio telescope (Sakai \\etal\\ 2008). Comparison of the relative radial velocities of the \\amm\\ $(J,K) = (1,1)$ inversion transition and CCS $J_N = 2_1 - 1_0$ rotational transition measured a systematic shift of $\\Delta V \\equiv V_{lsr}({\\rm CCS}) - V_{lsr}({\\rm NH}_3) \\sim 60$ \\ms, which when interpreted in terms of the electron-to-proton mass ratio variation infers that \\dmm $\\sim 6\\times10^{-8}$. We also noted that a similar offset between \\nnhp\\ and \\amm\\ measured by Pagani \\etal\\ (2009) in the cold dark cloud L183 could have a similar physical origin (Molaro \\etal\\ 2009).  In this paper, we present results of our own spectral observations of cold and compact molecular cores in the Taurus giant molecular complex obtained with the Medicina 32-m telescope, the Effelsberg 100-m telescope, and the Nobeyama 45-m telescope. ",
        "conclusions": "\\label{sect-7}  Using the 32-m Medicina, 100-m Effelsberg, and 45-m Nobeyama radio telescopes, we have performed precise spectral measurements of the relative radial velocity differences between the rotational transitions in \\hcccn\\ $(J=2-1)$ and \\nnhp\\ ($J=1-0$) and the inversion transition in \\amm\\ $(J,K)=(1,1)$. We detect a velocity offset of $\\Delta V = 23\\pm4_{\\rm stat}\\pm3_{\\rm sys}$ \\ms\\ between the rotational and inversion lines observed in cold molecular cores with dominating thermal motion. We do not find any plausible systematic effects that could mimic an offset of about 20 \\ms\\ between rotational and inversion transitions and would be regularly reproduced in observations of different cold molecular cores with different facilities. The measured positive offset is qualitatively consistent with our preliminary result, $\\Delta V \\simeq 60$ \\ms, obtained on the basis of the GBT observations of Rosolowsky \\etal\\ (2008) and Rathborne \\etal\\ (2008) of the NH$_3$/CCS pairs in molecular cores in the Perseus molecular cloud and the Pipe Nebula (LMK). However, the rest-frame frequency of the CCS $J_N = 2_1-1_0$ transition is not well known, and before deciding whether the difference between the current and preliminary results is statistically significant or not, new high precision laboratory measurements of the CCS frequency should be carried out.  If we assume that the measured velocity offset is caused by the electron-to-proton mass ratio variation, then \\dmm = $22\\pm4_{\\rm stat}\\pm3_{\\rm sys}$ ppb. To account for the conditions of the terrestrial laboratory experiments, the non-zero $\\Delta\\mu$ would require chameleon-like scalar field models that predict a strong dependence of mass and coupling constant on the ambient matter density.  In any case, our estimate defines a strongest conservative upper limit to $|\\Delta\\mu/\\mu| \\leq 3\\times10^{-8}$, which can be considered as a reference point at $z = 0$. We note that extragalactic molecular clouds have gas densities similar to those in the interstellar clouds of the Milky Way. Thus, the value of \\dmm\\ in high-$z$  molecular systems is expected to be at the same level as in the interstellar clouds, i.e., \\dmm $\\sim 10^{-8}$, provided that no temporal dependence of the electron-to-proton mass ratio is present.  To be completely confident that the derived velocity shift is not caused by kinematic effects in the clouds but reflects the density-modulated  variation of \\dmm, new high precision radio-astronomical observations are needed for a wider range of objects. These observations should also target different molecules (i.e., not NH$_3$) with tunneling transitions sensitive to the changes in $\\mu$ and molecules with $\\Lambda$-doublet lines that also exhibit enhanced sensitivity to variations in $\\mu$ and $\\alpha$ (Kozlov  2009).  For the molecular transitions in question, it is very important to measure their rest-frame frequencies with an accuracy of about 1 \\ms\\ in laboratory experiments. In some cases the present uncertainties in the rest-frame frequencies are larger than the errors of radio-astronomical measurements, thus preventing unambiguous conclusions.  In addition, searches for variations in the fine-structure constant $\\alpha$ within the Milky Way disk using mid- and far-infrared fine-structure transitions in atoms and ions (Kozlov \\etal\\ 2008), or searches for variations in the combination of $\\alpha^2/\\mu$ using the [\\ion{C}{ii}] $\\lambda158$ $\\mu$m line and CO rotational lines (Levshakov \\etal\\ 2008b), or the [\\ion{C}{i}] $\\lambda609$ $\\mu$m line and low-lying rotational lines of $^{13}$CO (Levshakov \\etal\\ 2009) would be of great importance for cross-checking these results.     \\begin{table*}[t!] \\centering \\caption{Target list. } \\label{tbl-1} \\begin{tabular}{rlcccllc} \\hline \\hline \\noalign{\\smallskip} & \\multicolumn{1}{c}{Source} & R.A. & Dec.  & $V_{lsr}$ & \\multicolumn{1}{c}{Obs.$^a$} & \\multicolumn{1}{c}{Molecules$^b$} & References$^c$\\\\ \\multicolumn{1}{c}{\\#} & & \\multicolumn{1}{c}{(J2000.0)} & \\multicolumn{1}{c}{(J2000.0)} & (\\kms) \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} 1 &L1355 & 02:53:12.1 & $+$68:55:52 & $-3.7$ & $m, e$   & \\amm\\  & 14\\\\ 2 &NH3SRC8 & 03:25:26.3 & $+$30:45:05 & $+4.1$ & $m, e, n $ & \\amm, \\nnhp  & 5\\\\ 3 &NH3SRC28 & 03:27:26.4 & $+$29:51:08 & $+5.6$ & $e$   & \\amm\\  & 5\\\\ 4 &NH3SRC36 & 03:27:55.9 & $+$30:06:18 & $+4.7$ & $m$   & \\amm\\ & 5\\\\ 5 &NH3SRC91 & 03:30:15.2 & $+$30:23:39 & $+5.9$ & $m$ &\\amm & 5\\\\ 6 &NH3SRC95 & 03:30:32.1 & $+$30:26:19 & $+6.1$ & $m$   & \\amm, \\hcccn & 5\\\\ 7 &NH3SRC97 & 03:30:50.5 & $+$30:49:17 & $+7.7$ & $e$   & \\amm & 5\\\\ 8 &L1489    & 04:04:49.0 & $+$26:18:42 & $+6.7$ & $e, n$ & \\amm,  \\hcccn, \\nnhp & 1, 4\\\\ 9 &L1498    & 04:10:51.4 & $+$25:09:58 & $+7.8$ & $e, n$ & \\amm, \\hcccn, \\nnhp & 1, 3, 4, 6\\\\ 10 &L1495    & 04:14:08.2 & $+$28:08:16 & $+6.8$ & $e, n$ & \\amm, \\hcccn, \\nnhp & 1, 6\\\\ 11 &L1521F   & 04:28:39.8 & $+$26:51:35 & $+6.5$ & $e$ & \\amm, \\hcccn & 4, 6, 10\\\\ 12 &L1524    & 04:29:22.7 & $+$24:34:58 & $+6.5$ & $e, n$ & \\amm, \\hcccn, \\nnhp & 4, 8, 10\\\\ 13 &B217S-W   & 04:27:46.4 & $+$26:17:52 & $+7.0$ & $e$ & \\amm, \\hcccn & 1, 4, 12\\\\ 14 &L1400K    & 04:30:52.0 & $+$54:51:55 & $+3.3$ & $e, n$ & \\amm, \\hcccn & 1, 3, 4, 6\\\\ 15 &TMC-2     & 04:32:48.6 & $+$24:25:12 & $+6.2$ & $e, n$ & \\amm, \\hcccn, \\nnhp & 1\\\\ 16 &L1536     & 04:33:23.5 & $+$22:42:46 & $+5.6$ & $e, n$ & \\amm, \\hcccn, \\nnhp & 1, 4\\\\ 17 &CB22      & 04:40:33.1 & $+$29:55:04 & $+6.0$ & $e$ & \\amm, \\hcccn & 4\\\\ 18 &TMC-1C    & 04:41:38.8 & $+$26:00:42 & $+5.2$ & $m,e,n$ & \\amm, \\hcccn, \\nnhp & 1, 2, 3, 4\\\\ 19 &TMC-1/HC3N& 04:41:41.8 & $+$25:41:42 & $+5.9$ & $m$ & \\amm, \\hcccn & 9, 11\\\\ 20 &CB23      & 04:43:27.7 & $+$29:39:11 & $+6.0$ & $e, n$ & \\amm, \\hcccn, \\nnhp & 4, 6\\\\ 21 &L1517B    & 04:55:18.8 & $+$30:38:04 & $+5.8$ & $e, n$ & \\amm, \\hcccn, \\nnhp & 2, 4, 6\\\\ 22 &L1512     & 05:04:09.6 & $+$32:43:09 & $+7.1$ & $m, e, n$ & \\amm, \\hcccn, \\nnhp & 1, 2, 4, 6, 8\\\\ 23 &L1544     & 05:04:16.5 & $+$25:10:48 & $+7.2$ & $e, n$ & \\amm, \\hcccn & 3, 4, 6, 8, 9\\\\ 24 &L134A     & 15:53:36.2 & $-$04:35:26 & $+2.7$ & $e$ & \\amm & 4, 6, 8\\\\ 25 &L183      & 15:54:06.4 & $-$02:52:23 & $+2.4$ & $m, e, n$ & \\amm, \\hcccn, \\nnhp & 4, 6, 7, 8\\\\ 26 &L43       & 16:34:35.0 & $-$15:46:36 & $+0.7$ & $e, n$ & \\amm, \\hcccn, \\nnhp & 4 \\\\ 27 &L260-NH3  & 16:47:06.6 & $-09$:35:21 & $+3.5$ & $e, n$ & \\amm, \\hcccn, \\nnhp & 1, 8, 12\\\\ 28 &L234A     & 16:48:06.8 & $-$10:51:48 & $+2.9$ & $e$ & \\amm, \\hcccn & 4, 8\\\\ 29 &L234S-E   & 16:48:08.6 & $-$10:57:25 & $+3.2$ & $e$ & \\amm & 4\\\\ 30 &L63       & 16:50:15.4 & $-$18:06:06 & $+5.7$ & $e$ & \\amm, \\hcccn & 4, 8\\\\ 31 &CB68      & 16:57:19.4 & $-$16:09:25 & $+5.2$ & $e$ & \\amm, \\hcccn & 4 \\\\ 32 &L492      & 18:15:49.6 & $-$03:46:14 & $+7.7$ & $e$ & \\amm, \\hcccn & 6\\\\ 33 &L778      & 19:26:32.5 & $+$23:58:42 & $+9.9$ & $m$ & \\amm & 1, 4\\\\ 34 &B335      & 19:37:01.3 & $+$07:34:07 & $+8.4$ & $e$ & \\amm, \\hcccn & 4, 8\\\\ 35 &L694-2    & 19:41:04.4 & $+$10:57:02 & $+9.5$ & $m, e$ & \\amm & 6, 13\\\\ 36 &L1049-2   & 20:41:30.5 & $+$57:33:44 & $+0.1$ & $e$ & \\amm & 15 \\\\ 37 &L1251T3   & 22:29:50.3 & $+$75:13:25 & $-3.8$ & $e, n$ & \\amm, \\hcccn, \\nnhp & 4 \\\\ 38 &L1251T2   & 22:30:31.2 & $+$75:13:38 & $-4.0$ & $e, n$ & \\amm, \\hcccn, \\nnhp & 4 \\\\ 39 &L1251T1   & 22:31:02.3 & $+$75:13:39 & $-4.2$ & $e, n$ & \\amm, \\hcccn, \\nnhp & 4 \\\\ 40 &L1251C    & 22:35:53.6 & $+$75:18:55 & $-4.7$ & $e, n$ & \\amm, \\hcccn, \\nnhp & 16,17 \\\\ 41 &L1197     & 22:37:02.1 & $+$58:57:22 & $-3.1$ & $e, n$ & \\amm, \\hcccn, \\nnhp & 6,13 \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} \\multicolumn{8}{l}{$^a$ Observations: $m$~-- Medicina 32-m, $e$~-- Effelsberg 100-m, and $n$~-- Nobeyama 45-m radio telescopes. }\\\\ \\multicolumn{8}{l}{$^b$ Detected transitions: NH$_3$ $(J,K) = (1,1)$,  HC$_3$N $J = 2-1$, N$_2$H$^+$ $J = 1-0$.  }\\\\ \\multicolumn{8}{l}{$^c$ References: (1) Benson \\& Myers 1989; (2) Myers \\etal\\ 1983; (3) Howe \\etal\\ 1994; (4) Jijina \\etal\\ 1999;  }\\\\ \\multicolumn{8}{l}{\\,\\, (5) Rosolowsky \\etal\\ 2008; (6) Crapsi \\etal\\ 2005; (7) Pagani \\etal\\ 2009; (8) Fuller \\& Myers 1993;  }\\\\ \\multicolumn{8}{l}{\\,\\, (9) Churchwell \\etal\\ 1984; (10) Codella \\etal\\ 1997; (11) Benson \\& Myers 1983; (12) Motte \\& Andre 2001;}\\\\ \\multicolumn{8}{l}{\\,\\, (13) Lee \\etal\\ 2007; (14) Park \\etal\\ 2004; (15) Lee \\etal\\ 1999; (16) Caselli  \\etal\\ 2002; (17) Walsh \\etal\\ 2004.}\\\\   \\end{tabular} \\end{table*}    \\begin{figure*}[t] \\vspace{0.0cm} \\hspace{-0.2cm}\\psfig{figure=13007f1.eps,height=12cm,width=10cm} \\vspace{-1.0cm} \\caption[]{Sketch of the hyperfine transitions for the NH$_3$ (1,1), HC$_3$N $(2-1)$, and N$_2$H$^+$ $(1-0)$ states. Line identification numbers are given in each panel (see Tables \\ref{tbl-2}~--~\\ref{tbl-4}). } \\label{fg1} \\end{figure*}   \\begin{figure*}[t] \\vspace{0.0cm} \\hspace{-1.2cm}\\psfig{figure=13007f2.eps,height=18cm,width=20cm} \\vspace{-7.0cm} \\caption[]{Spectra of NH$_3$ (1,1) and HC$_3$N ($2-1$) toward the cores L1512 and L183 obtained at the Medicina 32-m radio telescope. The histogram shows the data, the solid curve shows the fit, and the residual is plotted below each profile. The horizontal thick bars mark spectral windows used in the fitting procedure. The data are the arithmetic means of all the observations. The size of the resolution element (pixel) is 62 \\ms\\ for \\amm, and 80 \\ms\\ for \\hcccn. For each spectrum, the signal-to-noise ratio (S/N) per pixel at the maximum intensity peak is depicted. } \\label{fg2} \\end{figure*}  \\begin{figure*}[t] \\vspace{0.0cm} \\hspace{-1.2cm}\\psfig{figure=13007f3.eps,height=18cm,width=20cm} \\vspace{-1.5cm} \\caption[]{Spectra of NH$_3$ (1,1) and HC$_3$N ($2-1$) toward the cores L1498, L1495, and L1400K obtained at the Effelsberg 100-m radio telescope. The histogram shows the data, the solid curve shows the fit, and the residual is plotted below each profile. The horizontal thick bars mark spectral windows used in the fitting procedure. The data are the arithmetic means of all the observations. The size of the resolution element (pixel) is 15 \\ms\\ for \\amm, and 20 \\ms\\ for \\hcccn. For each spectrum, the signal-to-noise ratio (S/N) per pixel at the maximum intensity peak is depicted. } \\label{fg3} \\end{figure*}  \\begin{figure*}[t] \\vspace{0.0cm} \\hspace{-1.2cm}\\psfig{figure=13007f4.eps,height=18cm,width=20cm} \\vspace{-1.5cm} \\caption[]{Same as Fig.~\\ref{fg3} but for the cores CB22, TMC-1C (offset $\\Delta\\alpha,\\Delta\\delta = 0,40$ arcsec),  and L1517B. } \\label{fg4} \\end{figure*}   \\begin{figure*}[t] \\vspace{0.0cm} \\hspace{-1.2cm}\\psfig{figure=13007f5.eps,height=18cm,width=20cm} \\vspace{-1.5cm} \\caption[]{Same as Fig.~\\ref{fg3} but for the cores L1512, L183, and L260-\\amm. } \\label{fg5} \\end{figure*}   \\begin{figure*}[t] \\vspace{0.0cm} \\hspace{-1.2cm}\\psfig{figure=13007f6.eps,height=18cm,width=20cm} \\vspace{-1.5cm} \\caption[]{Same as Fig.~\\ref{fg3} but for the cores L234A, B335, and L1251C. } \\label{fg6} \\end{figure*}   \\begin{figure*}[t] \\vspace{0.0cm} \\hspace{-1.2cm}\\psfig{figure=13007f7.eps,height=18cm,width=20cm} \\vspace{-1.5cm} \\caption[]{Spectra of NH$_3$ (1,1) and N$_2$H$^+$ ($1-0$) toward the cores L1498, L1536, and CB23 obtained at the Nobeyama 45-m radio telescope. The histogram shows the data, the solid curve shows the fit, and the residual is plotted below each profile. The horizontal thick bars mark spectral windows used in the fitting procedure. The data are the arithmetic means of all the observations. The size of the resolution element (pixel) is 49 \\ms\\ for \\amm, and 25 \\ms\\ for \\nnhp. For each spectrum, the signal-to-noise ratio (S/N) per pixel at the maximum intensity peak is depicted. } \\label{fg7} \\end{figure*}   \\begin{figure*}[t] \\vspace{0.0cm} \\hspace{-1.2cm}\\psfig{figure=13007f8.eps,height=18cm,width=20cm} \\vspace{-1.5cm} \\caption[]{Same as Fig.~\\ref{fg7} but for the cores L1517B, L1512, and L183. } \\label{fg8} \\end{figure*}   \\begin{figure*}[t] \\vspace{0.0cm} \\hspace{-1.2cm}\\psfig{figure=13007f9.eps,height=18cm,width=20cm} \\vspace{-1.5cm} \\caption[]{Same as Fig.~\\ref{fg7} but for the cores L183 (offset $\\Delta\\alpha,\\Delta\\delta = -31.5,25$ arcsec), L43, and L260-\\amm. } \\label{fg9} \\end{figure*}   \\begin{figure*}[t] \\vspace{0.0cm} \\hspace{0.0cm}\\psfig{figure=13007f10.eps,height=15cm,width=15cm} \\vspace{-0.1cm} \\caption[]{ Doppler velocity differences, $\\Delta V$, between the \\hcccn\\ $J=2-1$ and \\amm\\ $(J,K) = (1,1)$ transitions for the data obtained at the 32-m (filled squares) and the 100-m (filled circles) telescopes, and between the \\nnhp\\ $J=1-0$ and \\amm\\ $(J,K) = (1,1)$ transitions for the 45-m telescope data (open circles). The 1$\\sigma$ statistical errors are indicated. Given in parentheses are the coordinate offsets in arcsec. } \\label{fg10} \\end{figure*}   \\begin{figure*}[t] \\vspace{0.0cm} \\hspace{0.0cm}\\psfig{figure=13007f11.eps,height=15cm,width=15cm} \\vspace{-0.8cm} \\caption[]{ Doppler velocity differences, $\\Delta V$, between the \\hcccn\\ $J=2-1$ and \\amm\\ $(J,K) = (1,1)$ transitions for the data obtained at the 32-m and 100-m telescopes, and between the \\nnhp\\ $J=1-0$ and \\amm\\ $(J,K) = (1,1)$ transitions for the 45-m telescope data (1$\\sigma$ statistical errors are indicated). The parameter $\\beta$ is the ratio of the Doppler $b$-parameters: $\\beta = b$(\\amm)$/b$(\\hcccn), or $\\beta = b$(\\amm)$/b$(\\nnhp). Given in parentheses are $1\\sigma$ errors. The filled circles mark sources with thermally dominated motions. } \\label{fg11} \\end{figure*}   \\begin{table}[t!] \\centering \\caption{Hyperfine components of the NH$_3$ $(J,K) = (1,1)$ transition. The velocity offsets, $v_i$, are relative to 23.694495487 GHz. Data from Kukolich (1967) and Ho (1977). Shown in parentheses are $1\\sigma$ errors. } \\label{tbl-2} \\begin{tabular}{ccccccr@{.}l} \\hline \\hline \\noalign{\\smallskip} \\# & $F'_1$ & $F'$ & $F_1$ & $F$ & Weight & \\multicolumn{2}{c}{$v_i$, \\ms} \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\medskip} 1& 1 & 1.5 & 0 & 0.5 & 0.317 & $-19549$&98(48)  \\\\[-2pt] & 1 & 0.5 & 0 & 0.5 & 0.159 & $-19411$&73(61)  \\\\[2pt] 2& 1 & 1.5 & 2 & 2.5 & 0.357 & $-7815$&25(63)  \\\\[-2pt] & 1 & 1.5 & 2 & 1.5 & 0.040 & $-7372$&80(75)  \\\\[-2pt] & 1 & 0.5 & 2 & 1.5 & 0.198 & $-7233$&48(58)  \\\\[2pt] 3 & 2 & 1.5 & 2 & 2.5 & 0.071 & $-250$&92(60)  \\\\[-2pt] & 1 & 1.5 & 1 & 1.5 & 0.198 & $-213$&00(47)  \\\\[-2pt] & 2 & 2.5 & 2 & 2.5 & 1.000 & $-132$&38(47)  \\\\[-2pt] & 1 & 0.5 & 1 & 1.5 & 0.040 & $-75$&17(75)  \\\\[-2pt] & 2 & 1.5 & 2 & 1.5 & 0.643 & $192$&27(56)  \\\\[-2pt] & 2 & 2.5 & 2 & 1.5 & 0.071 & $311$&03(80)  \\\\[-2pt] & 1 & 1.5 & 1 & 0.5 & 0.040 & $322$&04(35)  \\\\[-2pt] & 1 & 0.5 & 1 & 0.5 & 0.079 & $460$&41(46)  \\\\[2pt] 4 & 2 & 1.5 & 1 & 1.5 & 0.040 & $7351$&32(49)  \\\\[-2pt] & 2 & 2.5 & 1 & 1.5 & 0.357 & $7469$&67(73)  \\\\[-2pt] & 2 & 1.5 & 1 & 0.5 & 0.198 & $7886$&69(72)  \\\\[2pt] 5 & 0 & 0.5 & 1 & 1.5 & 0.317 & $19315$&90(71)  \\\\[-2pt] & 0 & 0.5 & 1 & 0.5 & 0.159 & $19851$&27(35)$^\\dagger$  \\\\[-2pt] \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} \\multicolumn{8}{l}{$^\\dagger$The error of 0.35 \\ms\\ corresponds to the difference}\\\\[-2pt] \\multicolumn{8}{l}{between theoretically predicted (Ho 1977) and astrono-}\\\\[-2pt] \\multicolumn{8}{l}{mically determined (Rydbeck \\etal\\ 1977) frequencies.} \\end{tabular} \\end{table}    \\begin{table}[t!] \\centering \\caption{Hyperfine components of the HC$_3$N $J = 2-1$ transition. The velocity offsets, $v_i$, are relative to 18.19621694  GHz. Data from Lapinov (2008, private comm.) and M\\\"uller \\etal\\ (2005). } \\label{tbl-3} \\begin{tabular}{ccccr} \\hline \\hline \\noalign{\\smallskip} \\# & $F'$ & $F$ & Weight & \\multicolumn{1}{c}{$v_i$, \\ms} \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\medskip} 1 & 1 & 1 &  0.333 &   $-35548$.77(2.82)\\\\[-2pt] 2 & 1 & 2 &  0.022  &  $-14167$.97(2.83)\\\\[-2pt] 3 & 3 & 2 &  1.867  &  $-1540$.96(2.82)\\\\[-2pt] 4 & 2 & 1 &  1.000 &      0.00(2.82)\\\\[-2pt] 5 & 1 & 0 &  0.444  &   17806.60(2.82)\\\\[-2pt] 6 & 2 & 2 &  0.333  &   21380.64(2.82)\\\\ \\noalign{\\smallskip} \\hline \\end{tabular} \\end{table}   \\begin{table}[t!] \\centering \\caption{Hyperfine components of the N$_2$H$^+$ $J = 1-0$ transition. The velocity offsets, $v_i$, are relative to 93.1737722 GHz. Data from M\\\"uller \\etal\\ (2005). } \\label{tbl-4} \\begin{tabular}{ccccccr} \\hline \\hline \\noalign{\\smallskip} \\# & $F'_1$ & $F'$ & $F_1$ & $F$ & Weight & \\multicolumn{1}{c}{$v_i$, \\ms} \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\medskip} 1 &0&1 &1&0& 0.076 & $-8008.2(13.8)$  \\\\[-2pt] &0&1 &1&1& 0.100 & $-8008.2(13.8)$  \\\\[-2pt] &0&1 &1&2& 0.252 & $-8008.2(13.8)$  \\\\[2pt] 2&2&1 &1&2& 0.017 & $-622.6(13.8)$  \\\\[-2pt] &2&1 &1&0& 0.134 & $-622.6(13.8)$  \\\\[-2pt] &2&1 &1&1& 0.278 & $-622.6(13.8)$  \\\\[2pt] 3&2&3 &1&2& 1.000 & 0.00(13.2)  \\\\[2pt] 4&2&2 &1&2& 0.111 & 956.9(13.5)  \\\\[-2pt] &2&2 &1&1& 0.604 & 956.9(13.5)  \\\\[2pt] 5&1&1 &1&1& 0.051 & 5541.6(13.5)  \\\\[-2pt] &1&1 &1&2& 0.159 & 5541.6(13.5)  \\\\[-2pt] &1&1 &1&0& 0.218 & 5541.6(13.5)  \\\\[2pt] 6&1&2 &1&1& 0.111 & 5982.4(13.5) \\\\[-2pt] &1&2 &1&2& 0.604 & 5982.4(13.5)  \\\\[2pt] 7&1&0 &1&1& 0.143 & 6934.2(13.2)  \\\\ \\noalign{\\smallskip} \\hline \\end{tabular} \\end{table}   \\begin{table*}[t!] \\centering \\caption{Radial velocities, $V_{lsr}$, Doppler parameters, $b$, and corresponding $\\chi^2_\\nu$ values normalized per degree of freedom. Data from the Medicina 32-m and Effelsberg 100-m radio telescopes. The numbers in parentheses correspond to $1\\sigma$ statistical errors. } \\label{tbl-5} \\begin{tabular}{l r@{.}l r@{.}l  r@{.}l r@{.}l r@{.}l r@{.}l r@{.}l  c} \\hline \\hline \\noalign{\\smallskip} & \\multicolumn{14}{c}{ $V_{lsr}$, \\kms\\,/\\,b, \\kms\\,/\\,$\\chi^2_\\nu$ } \\\\[2pt] & \\multicolumn{8}{c}{ NH$_3$ (1,1) } & \\multicolumn{6}{c}{ HC$_3$N $(2-1)$ } & $\\Delta V$  \\\\[2pt] \\multicolumn{1}{c}{Object} &  \\multicolumn{2}{c}{outer} & \\multicolumn{2}{c}{inner} & \\multicolumn{2}{c}{central} & \\multicolumn{2}{c}{total} & \\multicolumn{2}{c}{low} & \\multicolumn{2}{c}{high} & \\multicolumn{2}{c}{total} & \\multicolumn{1}{c}{(\\ms)} \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip}  \\multicolumn{16}{c}{\\it Medicina 32-m radio telescope} \\\\ \\noalign{\\smallskip}  L1512 &  7&0800(60) & 7&0795(47) & 7&0760(41) & 7&0790(28) & \\multicolumn{2}{c}{ } & 7&0965(30) & \\multicolumn{2}{c}{ } & +17.5(4.1) \\\\[-2pt] &  0&1299(93) & 0&116(18) & 0&1138(51) & 0&1189(47) & \\multicolumn{2}{c}{ } &  0&1037(96) & \\multicolumn{2}{c}{ } \\\\[-2pt] &  1&04 & 0&89 & 0&38 & 0&89  & \\multicolumn{2}{c}{ } & 0&78 & \\multicolumn{2}{c}{ } \\\\ L183 &  2&4000(63) & 2&3970(49) & 2&4140(60) & 2&4075(32) & \\multicolumn{2}{c}{ } & 2&404(15) & \\multicolumn{2}{c}{ } & $-3.5(15.3)$ \\\\[-2pt] &  0&129(18) & 0&148(13) & 0&1297(55) & 0&1363(51) & \\multicolumn{2}{c}{ } &  0&134(61) & \\multicolumn{2}{c}{ } \\\\[-2pt] &  1&16 & 0&88 & 0&74 & 1&30  & \\multicolumn{2}{c}{ } & 0&97 & \\multicolumn{2}{c}{ } \\\\  \\multicolumn{16}{c}{\\it Effelsberg 100-m radio telescope} \\\\ \\noalign{\\smallskip}  L1498 &  7&8025(16) & 7&8080(13) & 7&8065(13) & 7&8060(8) & 7&8290(38) & 7&8280(15) & 7&8280(14) & +22.0(1.6) \\\\[-2pt] &  0&1125(31) & 0&1138(64) & 0&1227(14) & 0&1177(13)  & 0&105(13) & 0&0904(52) & 0&0910(48)  \\\\[-2pt] &  1&64 & 0&87 & 0&80 & 1&22  & 1&19 & 0&55 & 0&99 \\\\ L1495 &  6&7690(30) & 6&7830(27) & 6&7895(19) & 6&7840(14) & \\multicolumn{2}{c}{ } & 6&8175(38) & 6&8175(37) & +33.5(4.0) \\\\[-2pt] &  0&1421(57) & 0&147(17) & 0&1425(20) & 0&1446(19)  & \\multicolumn{2}{c}{ } & 0&113(14) & 0&115(12)  \\\\[-2pt] &  1&36 & 0&69 & 0&87 & 1&09  & \\multicolumn{2}{c}{ } & 0&86 & 0&89 \\\\ L1400K &  3&2580(28) & 3&2555(21) & 3&2650(19) & 3&2600(13) & 3&2690(50) & 3&2780(20) & 3&2770(18) & +17.0(2.2) \\\\[-2pt] &  0&1277(50) & 0&1111(59) & 0&1224(24) & 0&1165(20)  & 0&107(16) & 0&0986(96) & 0&0961(94)  \\\\[-2pt] &  1&40 & 1&00 & 0&88 & 1&18  & 0&52 & 0&70 & 1&0 \\\\ CB22 &  5&9665(26) & 5&9675(20) & 5&9740(17) & 5&9705(11) & \\multicolumn{2}{c}{ } & 5&9945(30) & 5&9935(28) & +23.0(3.0) \\\\[-2pt] &  0&1256(60) & 0&125(11) & 0&1357(19) & 0&1305(18)  & \\multicolumn{2}{c}{ } & 0&088(17) & 0&086(14)  \\\\[-2pt] &  1&10 & 0&89 & 0&75 & 1&02  & \\multicolumn{2}{c}{ } & 0&44 & 1&02 \\\\ TMC-1C$^a$ &  5&2405(36) & 5&2570(36) & 5&2575(28) & 5&2540(19) & 5&2815(56) & 5&2880(22) & 5&2875(20) & +33.5(2.8) \\\\[-2pt] &  0&0997(70) & 0&126(10) & 0&1135(32) & 0&1213(29)  & 0&086(24) & 0&1055(70) & 0&1032(66)  \\\\[-2pt] &  1&17 & 0&85 & 0&43 & 0&94  & 0&75 & 1&11 & 0&92 \\\\ L1517B &  5&7800(20) & 5&7830(17) & 5&7855(18) & 5&7835(10) & 5&8200(76) & 5&8250(23) & 5&8245(20) & +39.0(1.5) \\\\[-2pt] &  0&1131(40) & 0&1229(61) & 0&1201(17) & 0&1216(16)  & 0&126(60) & 0&1073(69) & 0&1100(68)  \\\\[-2pt] &  1&18 & 0&79 & 0&55 & 0&96  & 0&91 & 0&70 & 0&82 \\\\ L1512 &  7&1165(24) & 7&1120(13) & 7&1130(13) & 7&1130(8) & \\multicolumn{2}{c}{ } & 7&1345(9) & 7&1335(9) & +20.5(1.2) \\\\[-2pt] &  0&1244(35) & 0&1166(62) & 0&1135(15) & 0&1132(13)  & \\multicolumn{2}{c}{ }  & 0&0784(43) & 0&0786(41)  \\\\[-2pt] &  1&30 & 0&76 & 0&77 & 1&03  & \\multicolumn{2}{c}{ } & 0&94 & 1&04 \\\\ L183 &  2&4200(13) & 2&4170(11) & 2&4220(12) & 2&4195(7) & \\multicolumn{2}{c}{ } & 2&4335(32) & 2&4320(30) & +12.5(3.1) \\\\[-2pt] &  0&1062(25) & 0&1149(40) & 0&1166(13) & 0&1118(11)  & \\multicolumn{2}{c}{ }  & 0&097(18) & 0&104(11)  \\\\[-2pt] &  1&14 & 0&75 & 0&54 & 0&98  & \\multicolumn{2}{c}{ } & 0&71 & 1&00 \\\\ L260-NH$_3$ &  3&4605(15) & 3&4635(11) & 3&4580(15) & 3&4615(8) & 3&4565(50) & 3&4645(21) & 3&4630(20) & +1.5(2.2) \\\\[-2pt] &  0&1089(30) & 0&1082(39) & 0&1126(15) & 0&1084(13)  & 0&084(14)  & 0&0973(72) & 0&0951(61)  \\\\[-2pt] &  1&58 & 0&93 & 1&16 & 1&39  & 1&26 & 0&84 & 1&15 \\\\ L234A &  2&8935(50) & 2&9065(57) & 2&9030(35) & 2&9030(25) & \\multicolumn{2}{c}{ } & 2&9165(40) & 2&9185(45) & +15.5(5.1) \\\\[-2pt] &  0&1337(93) & 0&170(45) & 0&1601(33) & 0&1596(32)  & \\multicolumn{2}{c}{ }  & 0&136(14) & 0&135(15)  \\\\[-2pt] &  0&98 & 0&73 & 0&83 & 0&90  & \\multicolumn{2}{c}{ } & 0&92 & 0&73 \\\\ B335 &  8&3350(25) & 8&3430(27) & 8&3250(20) & 8&3365(14) & \\multicolumn{2}{c}{ } & 8&3545(51) & 8&3585(50) & +22.0(5.2) \\\\[-2pt] &  0&2312(55) & 0&2326(36) & 0&2138(22) & 0&2294(19)  & \\multicolumn{2}{c}{ }  & 0&196(12) & 0&195(13)  \\\\[-2pt] &  0&84 & 1&29 & 0&89 & 1&14  & \\multicolumn{2}{c}{ } & 0&61 & 0&99 \\\\ L1251C &  $-4$&7390(22) & $-4$&7360(21) & $-4$&7200(21) & $-4$&7320(12) & \\multicolumn{2}{c}{ } & $-4$&7005(70) & $-4$&6975(70) & +34.5(7.1) \\\\[-2pt] &  0&1396(55) & 0&1614(64) & 0&1577(21) & 0&1588(18)  & \\multicolumn{2}{c}{ }  & 0&146(18) & 0&141(19)  \\\\[-2pt] &  1&77 & 1&18 & 0&67 & 1&48  & \\multicolumn{2}{c}{ } & 0&65 & 0&74 \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} \\multicolumn{16}{l}{$^a$ Offset $(\\Delta\\alpha,\\Delta\\delta) = (0'',40'')$. } \\end{tabular} \\end{table*}   \\begin{table*}[t!] \\centering \\caption{Radial velocities, $V_{lsr}$, Doppler parameters, $b$, and corresponding $\\chi^2_\\nu$ values normalized per degree of freedom. Data from the Nobeyama 45-m radio telescope. The numbers in parentheses correspond to $1\\sigma$ statistical errors. } \\label{tbl-6} \\begin{tabular}{l r@{.}l r@{.}l  r@{.}l r@{.}l r@{.}l r@{.}l r@{.}l  c} \\hline \\hline \\noalign{\\smallskip} & \\multicolumn{14}{c}{ $V_{lsr}$, \\kms\\,/\\,b, \\kms\\,/\\,$\\chi^2_\\nu$ } \\\\[2pt] & \\multicolumn{8}{c}{ NH$_3$ (1,1) } & \\multicolumn{6}{c}{ N$_2$H$^+$ $(1-0)$ } & $\\Delta V$  \\\\[2pt] \\multicolumn{1}{c}{Object} &  \\multicolumn{2}{c}{outer} & \\multicolumn{2}{c}{inner} & \\multicolumn{2}{c}{central} & \\multicolumn{2}{c}{total} & \\multicolumn{2}{c}{low} & \\multicolumn{2}{c}{high} & \\multicolumn{2}{c}{total} & \\multicolumn{1}{c}{(\\ms)} \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} L1498 &  7&8215(30) & 7&8210(25) & 7&8180(25) & 7&8200(15) & 7&8450(42) & 7&8400(43) & 7&8430(30) & +23.0(3.4) \\\\[-2pt] &  0&1080(54) & 0&116(11) & 0&1129(30) & 0&1124(26) & 0&092(15) &  0&0987(40) & 0&1025(68) \\\\[-2pt] &  1&37 & 0&65 & 1&07 & 1&00  & 0&84 & 1&16 & 1&08 \\\\ L1536 &  5&5900(28) & 5&5935(24) & 5&6030(27) & 5&5955(15) & 5&629(11) & 5&6295(92) & 5&6290(80) & +33.5(8.1) \\\\[-2pt] &  0&1284(66) & 0&1306(94) & 0&1391(27) & 0&1315(26) & 0&121(24) &  0&117(40) & 0&118(20) \\\\[-2pt] &  0&90 & 1&16 & 1&07 & 1&15  & 0&76 & 0&77 & 0&75  \\\\ CB23 &  6&0345(30) & 6&0365(22) & 6&0325(22) & 6&0345(14) & 6&0515(46) & 6&0385(45) & 6&0460(33) & +11.5(3.6) \\\\[-2pt] &  0&1027(65) & 0&106(11) & 0&1071(28) & 0&1064(25) & 0&097(13) &  0&119(12) & 0&1043(92) \\\\[-2pt] &  1&47 & 0&84 & 0&86 & 1&08  & 0&64 & 1&40 & 1&00 \\\\ L1517B &  5&8015(62) & 5&7980(60) & 5&7965(58) & 5&7985(34) & 5&8190(50) & 5&8215(53) & 5&8195(38) & +21.0(5.1) \\\\[-2pt] &  0&107(12) & 0&119(28) & 0&1228(64) & 0&1171(57) & 0&109(21) &  0&138(16) & 0&127(14) \\\\[-2pt] &  1&13 & 0&45 & 1&15 & 0&84  & 1&00 & 1&03 & 1&03 \\\\ L1512 &  7&0925(59) & 7&1030(38) & 7&1080(42) & 7&1025(25) & 7&1290(45) & 7&1260(41) & 7&1270(35) & +24.5(4.3) \\\\[-2pt] &  0&127(14) & 0&112(21) & 0&1210(48) & 0&1158(43) & 0&0847(92) &  0&085(11) & 0&0850(90) \\\\[-2pt] &  0&59 & 0&93 & 1&00 & 0&94  & 0&89 & 0&89 & 0&86 \\\\ L183 &  2&4195(27) & 2&4165(26) & 2&4275(33) & 2&4225(16) & 2&4395(50) & 2&4485(45) & 2&4445(32) & +22.0(3.6) \\\\[-2pt] &  0&1490(78) & 0&1563(65) & 0&1536(29) & 0&1563(27) & 0&149(17) &  0&1484(61) & 0&154(56) \\\\[-2pt] &  1&55 & 1&01 & 1&75 & 1&64  & 0&91 & 0&82 & 0&91 \\\\ L183$^a$ &  2&4245(19) & 2&4265(18) & 2&4375(27) & 2&4265(11) & 2&4190(26) & 2&4135(26) & 2&4165(18) & $-10.0(2.1)$ \\\\[-2pt] &  0&1579(50) & 0&1681(45) & 0&1816(22) & 0&1610(20) & 0&1325(64) &  0&1624(69) & 0&1423(54) \\\\[-2pt] &  1&18 & 0&82 & 1&32 & 1&43  & 1&18 & 1&01 & 1&30 \\\\ L43 & 0&7020(25) & 0&7040(27) & 0&7150(27) & 0&7050(16) & 0&7485(31) & 0&7470(25) & 0&7480(19) & +43.0(2.5) \\\\[-2pt] &  0&1523(63) & 0&1748(64) & 0&1859(29) & 0&1677(25) & 0&1368(88) &  0&173(17) & 0&1512(93) \\\\[-2pt] &  1&46 & 1&37 & 0&84 & 1&43  & 1&81 & 0&90 & 1&60 \\\\ L260-NH$_3$ & 3&4815(25) & 3&4800(21) & 3&4750(25) & 3&4790(14) & 3&4700(59) & 3&4710(52) & 3&4690(37) & $-10.0(4.0)$ \\\\[-2pt] &  0&1053(55) & 0&119(10) & 0&1144(26) & 0&1131(23) & 0&0997(82) &  0&111(26) & 0&1061(79) \\\\[-2pt] &  1&26 & 0&83 & 0&53 & 1&06  & 1&34 & 1&12 & 1&24 \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} \\multicolumn{16}{l}{$^a$ Offset $(\\Delta\\alpha,\\Delta\\delta) = (-31.5'',25'')$. } \\end{tabular} \\end{table*}"
    },
    "0911/0911.5578_arXiv.txt": {
        "abstract": "We estimate the detection efficiency of binary gravitational lensing events through the channel of high-magnification events.  From this estimation, we find that binaries in the separations ranges of $0.1 \\lesssim s\\lesssim 10$, $0.2\\lesssim s\\lesssim 5$, and  $0.3\\lesssim s\\lesssim 3$ can be detected with $\\sim 100\\%$ efficiency for events with magnifications higher than  $A=100$, 50, and 10, respectively, where $s$ represents the projected separation between the lens components normalized by the Einstein radius.  We also find that the range of high efficiency covers nearly the whole mass-ratio range of stellar companions.  Due to the high efficiency in wide ranges of parameter space, we point out that majority of binary-lens events will be detected through the high-magnification channel in lensing surveys that focus on high-magnification events for efficient detections of microlensing planets.  In addition to the high efficiency, the simplicity of the efficiency estimation makes the sample of these binaries useful in the statistical studies of the distributions of binary companions as functions of mass ratio and separation.  We also discuss other importance of these events. ",
        "introduction": "If a star is gravitationally microlensed by a lens system composed of two masses, the resulting light curve can dramatically deviate from the smooth and symmetric one of a single-lens event.  This deviation is caused by the formation of caustics for binary-lens systems. The caustics represent the source positions at which the lensing magnification of a point source becomes infinite. The set of caustics forms one, two, or three close curves each of which is composed of concave curves that meet at points.  By analyzing the light curve of a binary-lens event, it is possible to obtain information about the lens system because the structure of the caustic system and the resulting light curve vary depending on the mass ratio and the projected separation between the components of binary lenses.    Since the pioneering work by \\citet{chang80, chang84}, binary lensing has been a subject of intense theoretical studies.  \\citet{schneider86} made a comprehensive study of binary lenses in order to learn about caustics in quasar macrolensing.  \\citet{witt90} developed a simple algorithm for finding caustics of binary lenses.  \\citet{witt95} studied lensing magnification inside caustic and found that the minimum magnification when the source is inside a caustic is greater than 3.  \\citet{rhie97} found that the maximum number of images for multiple-lens systems.  With the beginning of microlensing surveys, theoretical studies became even more active.  \\citet{gaudi97} pointed out that microlensing is an efficient method to detect close binaries.  \\citet{distefano97} mentioned various channels of detecting binaries including repeating events. \\citet{dominik99b} studied the lensing behavior in the extreme cases of binary separations and mass ratios. \\citet{dominik99a} and \\citet{albrow99b} mentioned possible degeneracies in modeling light curves of binary-lens events. \\citet{han99} and \\citet{han01} studied the astrometric behavior of binary-lens events.  \\citet{bozza00,bozza01} derived analytic expressions for the location of caustics and studied the motion of images of microlensed stars. \\citet{gaudi02a, gaudi02b} investigated the photometric and astrometric behaviors in the region very close to caustics.  \\citet{graff02} devised a method to measure the mass of the binary-lens system from the analysis of light curves of caustic-crossing events.  In addition to the theoretical studies, binary-lens events were actually detected from various surveys \\citep{udalski94, udalski98, alcock99, alard95, afonso00, albrow99a, albrow00, albrow01, an02, smith02, albrow02, abe03, kubas05, jaroszynski04, cassan04, jaroszynski05, jaroszynski06}.   With the active researches in both theoretical and observational fields, binary microlensing has developed into a useful tool to study stellar astrophysics.  The most active field of application is the stellar atmosphere for which microlensing is used to probe detailed structures on the surface of source stars by using the high resolution of caustic-crossing events \\citep{albrow99a, albrow01, abe03}.   Microlensing can also be used to probe the distributions of binary companions of Galactic stars as functions of mass ratio and separation.  These binary distributions provide important observational constraints on theories of star formation.  Since microlensing is sensitive to low-mass companions that are difficult to be detected by other methods, it is in principle possible to make complete distribution down to the lower mass limit of binary companions.  Despite the importance, the progress of this application of binary lensing has been stagnant. There are two main reasons for this.  The first reason is caused by the difficulty in estimating the detection efficiency of binary-lens events.  In previous lensing surveys, most binary-lens events were discovered through the channel of caustic-crossing events, in which the caustic crossings were accidently discovered from the sudden rise of the source star flux.  Due to the haphazard nature of caustic crossings, it was difficult to estimate the detection efficiency that is essential for the statistical studies of binary companions.  The second reason is that microlensing is mainly sensitive to binaries over a narrow range of projected separations.  This limits especially the study of the distribution of binary separations.  %   When the first microlensing surveys \\citep{alcock93, aubourg93} were started, the prime scientific goal was constraining the Galactic dark matter in the form of massive compact halo objects.  However, the most important science of the current lensing surveys \\citep{udalski03, bond02, beaulieu98, yoo04} is detections of extrasolar planets. With the change of the goal, the observational strategy of the surveys also changed.  These changes include the employment of early warning system and the operation of follow-up observations.  Another important change is the adoption of an observational strategy focusing on high-magnification events for which the probability of planet detection is high \\citep{griest98}.   In this paper, we emphasize the importance of binary-lens events to be detected through the high-magnification channel, especially for the statistical studies of Galactic binaries. We demonstrate the high efficiency of the high-magnification channel to binaries in a wide range of parameter space and the simplicity of the efficiency estimation.  We also discuss other importance of these events. ",
        "conclusions": "In this paper, we emphasized the importance of binary-lens events to be discovered through the channel of high-magnification events. Due to the high efficiency in wide ranges of parameter space, we pointed out that majority of binary-lens events will be detected through the high-magnification channel in current planetary microlensing surveys, In addition to the high efficiency, the simplicity of the efficiency estimation makes the sample of these binaries useful in the statistical studies of the distributions of binary companions as functions of mass ratio and separation. We also discussed other importance of these events."
    },
    "0911/0911.2358_arXiv.txt": {
        "abstract": "Based on the joint-observations of the radio broadband spectral emissions of solar eclipse on August 1, 2008 at Jiuquan (total eclipse) and Huairou (partial eclipse) at the frequencies of 2.00 -- 5.60 GHz (Jiuquan), 2.60 -- 3.80 GHZ (Chinese solar broadband radiospectrometer, SBRS/Huairou), and 5.20 -- 7.60 GHz (SBRS/Huairou), the authors assemble a successive series of broadband spectrum with a frequency of 2.60 -- 7.60 GHz to observe the solar eclipse synchronously. This is the first attempt to analyze the solar eclipse radio emission under the two telescopes located at different places with broadband frequencies in the periods of total and partial eclipse. With these analyses, the authors made a semiempirical model of the coronal plasma density of the quiet Sun, which can be expressed as $n_{e}\\simeq 1.42\\times10^{9}(r^{-2}+1.93r^{-5}),(cm^{-3})$ in the space range of $r=1.039 - 1.212 R_{\\odot}$, and made a comparison with the classic model. ",
        "introduction": "Without the high-resolution solar radioheliograph, the solar radio eclipse observations may be the best way to investigate the spatial structure of the solar radio emission. During the minimum phase of solar cycle, the influence of the active regions is very weak, we may study the quiet solar coronal atmospheric model by means of the radio eclipse observations$^{[1]}$.  On August 1, 2008, a grand solar total eclipse could occurred on the regions of northern Europe and northwestern China. In Jiuquan of Gansu province, people could observe the whole course of eclipse happening. In Beijing, people could also watch the partial solar eclipse clearly. Around this event, the Sun was in a quiet phase between solar cycle 23 and 24. There was no solar flare, no corona mass ejection (CME), no sunspot, and no any active regions. The solar disk looked like a fairly clean plate. Such total eclipse might provide a god-given opportunity to test the quiet Sun atmospheric models by means of radio observations. During this solar eclipse, we made an extensive broadband spectral radio joint-observations at Jiuquan and Beijing Huairou. Section 2 introduces the observed instruments and the related data processing. Section 3 is a deduction of the quiet Sun atmospheric model from the above data. And some discussions and conclusions make up the final part of the paper. ",
        "conclusions": "In this study work, we analyze the joint-observations of radio broadband spectral emissions during the solar eclipse on August 1, 2008 at Jiuquan (total eclipse, observed in the frequency of 2.00 -- 5.60 GHz) and Huairou (partial eclipse, observed in the frequency of 2.60 -- 3.80 GHz and 5.20 -- 7.60 GHz). Using these telescopes, we can assemble a successive series of broadband spectrum with frequency of 2.60 -- 7.60 GHz to observe the solar eclipse synchronously. With these analyses, we obtain a semiempirical quiet Sun coronal model (equation (7)) which indicates that the distribution of solar coronal plasma density in respect to the height is not so steep as the classic model, corresponding to the height of $2.7\\times10^{4}$ -- $1.5\\times10^{5}$ km ($r=1.039 - 1.212R_{\\odot}$) above the solar photosphere, and the bright temperature is in the range of $9.57\\times10^{4} - 2.04\\times10^{5}$ K. This is consistent with the result of Ref [14].  However, our results are obtained from the range of 2.60 -- 7.60 GHz, and only valid in the limited space of $r=1.039 - 1.212R_{\\odot}$. It is necessary to enlarge the observing frequency range in the future observations. Additionally, the height angle of the eclipse of August 1, 2008 was very small. When the eclipse occurred, the line-of-sight was fairly close to the ground, with many disturbances coming from the subaerial factors. The telescope of EcBS/Jiuquan, a new instrument, needs to be improved in many aspects. Fortunately, a more significant solar total eclipse will occur on July 22, 2009 in a more extended area along the Changjiang River, during 8-11 o'clock in the morning, and the height angle is very high. The observational condition is very good. We will make more detailed observations to get the further supports for our semiempirical model."
    },
    "0911/0911.1267_arXiv.txt": {
        "abstract": "In this paper we study various particle physics effects of a light, scalar dark energy field with chameleon-like couplings to matter. We show that a chameleon model with only matter couplings will induce a coupling to photons. In doing so, we derive the first microphysical realization of a chameleonic dark energy model coupled to the electromagnetic field strength. This analysis provides additional motivation for current and near-future tests of axion-like and chameleon particles. We find a new bound on the coupling strength of chameleons in uniformly coupled models. We also study the effect of chameleon fields on Higgs production, which is relevant for hadron colliders. These are expected to manufacture Higgs particles through weak boson fusion, or associated production with a $Z$ or $W^\\pm$. We show that, like the Tevatron, the LHC will not be able to rule out or observe chameleons through this mechanism, because gauge invariance of the low energy Lagrangian suppresses the corrections that may arise. \\vspace{3mm} \\begin{flushleft} \\textsf{% \\textbf{Keywords}: Dark energy theory, Weak interactions beyond the Standard Model, Cosmology of theories beyond the Standard Model \\\\ \\textbf{PACS codes}: 95.36.+x; 12.15.-y, 14.80.Va, 98.80.-k} \\end{flushleft} ",
        "introduction": "\\label{sec:introduction}  We now have compelling evidence that the rate of expansion of the universe is accelerating, requiring the universe to be dominated by an unknown form of matter, known as `dark energy,' characterized by the equation of state $p \\approx - \\rho$. This dark energy could be a cosmological constant unnaturally tuned at the level of 1 part in $10^{123}$ \\cite{Weinberg:1988cp}, or it may be associated with the potential energy of a scalar field \\cite{Ratra:1987rm,Wetterich:1994bg,Zlatev:1998tr,Carroll:1998zi}. If a scalar field drives the accelerated expansion then it must be extremely light, with mass of order $H_0\\sim 10^{-33}\\mbox{ eV}$. The existence of light scalar fields results in new, long range ``fifth forces,'' which are tightly constrained by experiment \\cite{Will:2001mx}.\tTo avoid these constraints the energy scale controlling the coupling of such scalar fields to matter must be many orders of magnitude larger than the Planck scale \\cite{Carroll:1998zi}. Explaining why this coupling is so weak is a major problem for dynamical dark energy models.  These problems can be circumvented by allowing the scalar field mass mass to be determined by a \\emph{chameleon} mechanism \\cite{Khoury:2003aq,Khoury:2003rn,Mota:2003tm,Clifton:2004st, Mota:2006fz}. This makes the field heavy in a dense environment, but light in vacuum.  Such scalar fields hide from experimental searches for fifth forces \\cite{Khoury:2003aq,Khoury:2003rn} and modifications of gravity \\cite{Hu:2007nk,Brax:2008hh} in a novel way: in the interior of a massive object the chameleon is very heavy, and has a correspondingly short interaction length. An observer outside the body feels a scalar force sourced only by a thin shell of matter at the surface of the object. This suppresses unwanted fifth forces \\cite{Khoury:2003aq,Khoury:2003rn}. Observations impose few constraints on the strength of the coupling of the scalar field to the Standard Model, allowing the energy scale of the coupling to be many orders of magnitude below the Planck scale \\cite{Mota:2006fz}.  Chameleonic scalar fields have been shown to have a successful classical phenomenology.\tTheir equation of state depends on the energy density of the universe. In the early universe the chameleon behaves as an additional matter component, but at late times it has an equation of state with roughly $w \\approx -1$. This enables it to drive an era of accelerated expansion \\cite{Brax:2004qh}. Weak constraints on the initial conditions ensure that the chameleon does not disrupt the dynamics of the early universe \\cite{Brax:2004qh,Mota:2006fz}. In this paper we discuss these theories in the framework of quantum mechanics. This is of interest because the cosmological constant problem is essentially quantum mechanical. There is no profit in replacing the cosmological constant by some other theory which is equally unnatural once quantum corrections are taken into account.  The chameleon mechanism makes it mandatory for the dark energy scalar field to interact with conventional matter, and these interactions are potentially very strong. They may help us unravel the microphysics of dark energy if their indirect consequences could be measured \\cite{Brax:2007tk,Weltman:2008fp,Weltman:2008ng,Burrage2009190, Brax:2009bk,Chou:2008gr,Upadhye:2009iv}. If the interactions are sufficiently strong, however, then an unambiguous \\emph{direct} consequence may be observable: new, light scalar quanta must be present in the beam pipe of any particle accelerator, raising the prospect of observing dark energy in the laboratory. Fully consistent quantum mechanical realizations of the chameleon have not yet been constructed, and may suffer from the same naturalness and tuning difficulties as quintessence \\cite{Kolda:1998wq,Peccei:2000rz,Doran:2002bc}. In this paper we do not attempt to address naturalness concerns, or the quantum-mechanical construction of the model. Instead, supposing such models to exist, we determine constraints which are imposed by experiment.  Real or virtual chameleon-like particles will certainly be produced in particle collisions if dark energy couples to the electroweak gauge bosons. In this paper, we will provide a microscopic derivation of this coupling for any chameleon-like model. The models which furnish chameleon-like scalar particles at low energy are effective theories, valid below some energy scale $M_c$. In Ref.~\\cite{Brax:2009aw}, corrections to electroweak precision observables from particles with chameleonic or axionic couplings were analysed, without making a commitment to any specific choice of physics at energies above $M_c$. It was shown that very weak constraints on the energy scale of the chameleon coupling to matter ensure that the chameleon does not have an observable effect on measurements of the $Z$-width. For processes involving fermions and $\\sugroup(2) \\times \\ugroup(1)$ gauge bosons, it was shown that large effects could always be absorbed into renormalizations of the Fermi constant, $G_F$, and the gauge boson masses. When Higgs processes are included it is no longer clear that large effects can be hidden in this way, potentially allowing Higgs production to function as a diagnostic of dark energy physics and its coupling to the Standard Model.  In this paper, we focus on the consequences of such dark energy quanta for Higgs production. This is a key target for both the \\emph{Large Hadron Collider} (LHC) and the \\emph{Tevatron}. The details of Higgs production depend on which couplings occur in the theory. To achieve a successful chameleon phenomenology we assume conformal couplings to matter.% \\footnote{All chameleon models developed to date have conformal couplings to matter.  However there exist related models of scalar fields with environment dependent properties which do not require conformal couplings \\cite{Luty:2003vm,Nicolis:2008in}.} This has implications for interactions with gauge boson kinetic terms, which will be described in more detail in~\\S\\ref{sec:scalar-corrections} below. The Standard Model gauge bosons have conformally invariant kinetic terms and develop no couplings to a scalar field of this type. In this paper we show that violations of conformal symmetry, arising from couplings to matter species, can be communicated to the kinetic term by quantum corrections. Therefore, a coupling \\emph{is} generated at energies below $M_c$ via loops of charged heavy particles. This coupling has important phenomenological consequences which will be briefly recalled. In particular we will show that this coupling induces a coupling to the electromagnetic field and in so doing provides a motivation for the effects probed in quantum laser experiments \\cite{Brax:2007hi,Brax:2007ak,Mirizzi:2005ng,Gies:2007su, Ahlers:2007st,Chou:2008gr,Upadhye:2009iv}. Our results apply for any light scalars with the requisite chameleon-like couplings, whatever their origin.  What scale, $M_c$, should be associated with these new degrees of freedom? A conservative choice would be the GUT scale, $M_c \\sim 10^{16}$ GeV, which may be connected with an early inflationary stage, or a seesaw explanation of neutrino mass \\cite{Minkowski:1977sc,GellMann:1980vs,Yanagida:1979as}. If so, heavy neutrinos of mass $M_c$ might exist but would be sterile, having no $\\sugroup(2) \\times \\ugroup(1)$ quantum numbers. This would not lead to the required coupling. Alternatively, if supersymmetry is realized in nature then $M_c$ could be associated with the scale at which it is spontaneously broken, perhaps of order $1$--$10$ TeV. In any supersymmetric completion of the Standard Model there exist fermionic partners of $\\gamma$, $W^\\pm$, $Z$ and the Higgs, known as gauginos and Higgsinos. (For a review see \\cite{Binetruy:2006ad} and references therein.) The mass eigenstates of these particles (``charginos'') would naturally be of order $M_c$. In this article we remain agnostic about the nature\tof whatever particles circulate within loops. We give our calculation in a form which can be specialized immediately to the minimal supersymmetric Standard Model (MSSM), but our conclusions are general and do not depend on the details of a specific implementation. In any case, the calculation we describe can be adapted easily to any heavy particle carrying the requisite quantum numbers.  The outline of this paper is as follows.  In \\S\\ref{sec:axionic} we introduce our model and show that a coupling between gauge bosons and dark energy is generated at low energy by integrating out heavy particles. Some of the phenomenology associated with this coupling is presented. In \\S\\ref{sec:higgs} we briefly review Higgs production at a hadron collider, emphasizing the role of vertices between the electroweak gauge bosons and the Higgs. In \\S\\ref{sec:scalar-corrections} we compute corrections to the Higgs production rate arising from the low energy theory written down in \\S\\ref{sec:axionic}. We show that the effective Lagrangian includes new contact interactions between the gauge bosons and the Higgs. At scales smaller than $1/M_c$ these contact interactions resolve into heavy particle loops. We compute the correction to Higgs production from both effects. In \\S\\ref{sec:conclusions} we conclude by arguing that the effect of chameleon-like particles on Higgs production is expected to be rather small.  Throughout this paper, we adopt units in which $c = \\hbar = 1$. We set the reduced Planck mass, $\\Mp \\equiv (8 \\pi G)^{-1/2}$, to unity. Our metric convention is $(-,+,+,+)$. ",
        "conclusions": "\\label{sec:conclusions}  In this paper we have derived the low energy effective theory which governs interactions between the gauge bosons of the electroweak sector and a dark energy scalar field. The dark energy is taken to have conformal couplings to the matter species of the Standard Model. It is possible that couplings of this type allow so-called ``chameleon'' behaviour, in which the field dynamically adjusts its mass to be large in regions of high average density, and small elsewhere. If such theories exist then they would lead to an unambiguous prediction of light dark energy quanta interacting in the beam pipe of any particle accelerator. Our low energy theory applies strictly in any model containing heavy charged fermions, but a very similar effective Lagrangian would apply for a model containing heavy charged particles of any spin. As a specific example, any supersymmetric Standard Model must contain charginos and neutralinos. These carry the quantum numbers of the $\\sugroup(2) \\times \\ugroup(1)$ gauge group and have masses of order the supersymmetry-breaking scale. However, our calculation is not restricted to the supersymmetric case.  One might worry that the presence of chameleonic quanta would change our predictions for the outcome of particle physics experiments. In Ref.~\\cite{Brax:2009aw} we argued this did not happen for any process without Higgs quanta. In this paper we have extended our argument to include Higgs production. In particular, using the low energy theory we have computed the oblique corrections to each of the Higgs, $Z$ and $W^\\pm$ propagators at energies below the mass, $M$, of the heavy charged fermions. Such corrections modify the rates $\\Gamma(ZZ \\rightarrow h)$ and $\\Gamma(W^+ W^- \\rightarrow h)$, where $h$ is the lightest neutral Higgs, and in principle could change the rate of production of this particle at a hadron collider. We find that the corrections diverge quadratically with the scale chosen as the cut off for the effective theory. This does not predict a large enhancement to the production of Higgs bosons from interactions with dark energy, instead indicating that the interactions with dark energy make this process sensitive to the UV physics.  Other contributions exist, generated by processes taking place at high energy, which are integrated out of the low-energy description. We determine these ``threshold corrections'' by integrating out heavy fermion loops which mediate interactions between the lightest neutral Higgs and the gauge bosons.  The scale of these corrections can then be reinterpreted as\t a cutoff in an oblique calculation, of order $\\Lambda \\sim \\beta_H^{-1} \\egauge^2 \\coupling^2 / M^2 G_F^2$, and therefore leads to at most small effects. When $M$ is much larger than the Standard Model scale $G_F^{-1/2}$ it is entirely negligible. In an unconstrained theory, we might have anticipated a cutoff of order $\\Lambda \\sim \\beta_H^{-1}$, because at this scale the effective dark energy theory becomes invalid. If this were true, it would be possible to contemplate corrections to the Higgs production rate of $\\Or(1)$ or larger, which would lie within the discovery reach of the LHC.  Unfortunately, the corrections are much smaller. Large terms could only arise from the relevant operator $H Z^a Z_a$, but its appearance is forbidden by gauge invariance above the scale of electroweak symmetry breaking. Therefore, we expect the coefficient of this term to be at most $G_F^{-1/2}$, rather than $\\beta_H^{-1}$. Instead, the leading correction comes from the operators $H^\\dag H \\Tr F^{ab} F_{ab}$, $|D^2 H|^2$ and $(D_a D_b H)^\\dag (D^a D^b H)$. These are dimension-six operators, because of the $\\sugroup(2)$ nature of the Higgs doublet. Accordingly the cutoff is suppressed by $(M^2 G_F)^{-1}$, making it small. This must be typical of any UV correction because there is no relevant operator we can write down which will couple $\\chi$ to the gauge fields.\tTo get a larger effect, it appears to be necessary to break the gauge invariance of the theory. This does not rule out the possibility that dark energy could be responsible for an enhancement in the Higgs production rate, but such a scenario would apparently require exotic physics.  In the context of particle colliders where strong magnetic fields are present, the coupling derived in \\S\\ref{sec:axionic} implies a coupling of the chameleon to synchrotron radiation. This would lead to emergence of chameleon-like particles due to the Primakov effect. For most chameleon theories, the large mass assumed by chameleon-like particles in a dense environment implies that the beam pipe acts as a reflecting wall, preventing dark energy particles from leaking out. Hence collider experiments would actually take place in a dark energy bath. The analysis of this phenomenon is left for future work.  \\ack We would like to thank Ben Allanach, Neil Barnaby, Cliff Burgess, Douglas Shaw and James Stirling for helpful discussions. CB is supported by the German Science Foundation (DFG) under the Collaborative Research Centre (SFB) 676. ACD and DS are supported by STFC and the Cambridge Centre for Theoretical Cosmology. AW is supported by the Cambridge Centre for Theoretical Cosmology."
    },
    "0911/0911.1860_arXiv.txt": {
        "abstract": "{Hundreds of radar-dark patches interpreted as lakes have been discovered in the north and south polar regions of Titan. We have estimated the composition of these lakes by using the direct abundance measurements from the Gas Chromatograph Mass Spectrometer (GCMS) aboard the Huygens probe and recent photochemical models based on the vertical temperature profile derived by the Huygens Atmospheric Structure Instrument (HASI). Thermodynamic equilibrium is assumed between the atmosphere and the lakes, which are also considered as nonideal solutions. We find that the main constituents of the lakes are ethane (C$_2$H$_6$) ($\\sim$ 76--79\\%), propane (C$_3$H$_8$) ($\\sim$ 7--8\\%), methane (CH$_4$) ($\\sim$ 5--10\\%), hydrogen cyanide (HCN) ($\\sim$ 2--3\\%), butene (C$_4$H$_8$) ($\\sim$ 1\\%), butane (C$_4$H$_{10}$) ($\\sim$ 1\\%) and acetylene (C$_2$H$_2$) ($\\sim$ 1\\%). The calculated composition of lakes is then substantially different from what has been expected from models elaborated prior to the exploration of Titan by the Cassini-Huygens spacecraft.} ",
        "introduction": "\\label{sec:intro}  The surface of Saturn's haze-shrouded moon Titan had long been proposed to have oceans or seas, on the basis of the stability of liquid methane and ethane at the ground level (Flasar 1983; Lunine et al. 1983; Lorenz et al. 2003). Ground-based radar observations  ruled out the presence of a global ocean in the 1990s (Muhleman et al. 1995), but the presence of isolated lakes was not precluded (Campbell et al. 2003). A large, dark, lake-like feature subsequently named Ontario Lacus was detected at Titan's south polar region by the Cassini ISS system in 2005 (McEwen et al. 2005) and hundreds of radar dark features with a variety of properties consistent with liquid-filled lakes were found in the northern hemisphere by the Cassini RADAR system (Stofan et al. 2007).  The chemical composition of the lakes of Titan is still not well determined. Good quality spectral data of the Ontario Lacus have been obtained by the Visual and Infrared Mapping Spectrometer (VIMS) aboard Cassini but the only species that seems firmly identified is C$_2$H$_6$ (Brown et al. 2008); the atmosphere contains so much CH$_4$ that it is very difficult to detect the surface liquid phase of this molecule even if it is dominant in the lakes. Because the detection of other compounds in the lakes of Titan remains challenging in the absence of in situ measurements, the only way to get a good estimate of the chemical composition of these lakes is to elaborate a thermodynamic model based on theoretical calculations and laboratory data. Several models, that investigate the influence of photochemistry and the atmospheric composition on the chemical composition of liquids formed on the surface of Titan, have been elaborated in the pre-Cassini years (Lunine et al. 1983; Dubouloz et al. 1989; McKay et al. 1993; Tokano 2005). Based on atmospheric observations these models assumed surface bodies of liquid on Titan to contain a mixture of C$_2$H$_6$, CH$_4$ and N$_2$ and a large number of dissolved minor species.  \\begin{table}[h] \\label{atmoscompo} \\caption[]{Assumed composition of Titan's atmosphere at the ground level.} \\begin{center} \\begin{tabular}{lcc} \\hline \\hline \\noalign{\\smallskip} Atmosphere\t\t\t& Mole fraction\t\t\t    \t\t& Determination         \\\\ \\hline H$_2$\t\t\t\t& $9.8 \\times 10^{-4}$\t\t\t& Huygens GCMS$^{(a)}$\t\\\\ N$_2$\t\t\t\t& 0.95                          \t\t\t& This work\t\t\t\t\\\\ CH$_4$\t\t\t\t& 0.0492 \t\t\t\t\t\t& Huygens GCMS$^{(b)}$\t\\\\ CO\t\t\t\t\t& $4.70 \\times 10^{-5}$\t\t\t& Cassini CIRS$^{(c)}$  \\\\ $^{40}$Ar\t\t\t\t& $4.32 \\times 10^{-5}$ \t    \t\t& Huygens GCMS$^{(b)}$\t\\\\ C$_2$H$_6$          \t\t& $1.49 \\times 10^{-5}$         \t\t& This work\t\t\t\t\\\\ \\hline \\end{tabular} \\tablecomments{$^{(a)}$Owen \\& Niemann 2009; $^{(b)}$Niemann et al. 2005; $^{(c)}$De Kok et al. 2007. N$_2$ and C$_2$H$_6$ abundances are determined from our model (see text).} \\end{center} \\end{table}  However,  Cassini-Huygens measurements have improved our knowledge of the structure and composition of Titan's atmosphere, requiring the solubilities to be recomputed under actual Titan conditions. In particular, the Gas Chromatograph Mass Spectrometer (GCMS) aboard Huygens and the Cassini Composite Infrared Spectrometer (CIRS) provided new atmospheric mole fraction data (see Table 1 and Niemann et al. 2005). Moreover, near-surface brightness temperatures at the high latitudes where the lakes exist have now been determined (Jennings et al. 2009).  Here, we propose a model that takes into account these recent advances and thus provides the most up-to-date chemical composition of Titan's lakes as a function of their location on the satellite's surface. Our model considers the same assumptions as those made by Dubouloz et al. (1989) (hereafter D89) when they calculated the composition of the hypothetical ocean proposed to exist on Titan in the years prior the Cassini-Huygens exploration. The lakes are then considered as nonideal solutions in thermodynamic equilibrium with the atmosphere. This assumption is supported by recent calculations who showed that raindrops could reach the ground in compositional equilibrium with the atmosphere (Graves et al. 2008). ",
        "conclusions": "\\label{sec:discussion}  \\begin{figure} \\begin{center} \\includegraphics[angle=0,width=8.5cm]{canevas_Ych4_1-mole.eps} \\caption{\\label{influ_CH4}(a)--(h): composition of lakes as a function of the methane atmospheric mole fraction, assuming a surface temperature of $93.65$ K. The vertical dashed line corresponds to the methane atmospheric mole fraction measured by Huygens at the ground level.} \\end{center} \\end{figure}  The use of more up-to-date thermodynamic data and the recent Cassini-Huygens measurements, results in our calculated composition of the lakes differing from that of the hypothetical global ocean determined by D89. Indeed, considering the D89 case most closely corresponding to actual Titan conditions ($T= 92.5$ K, $Y_{\\rm Ar}=0$ and $Y_{\\rm CH_{4}}= 0.0492$; see their Fig.~1), they obtained mole fractions of C$_2$H$_6$, CH$_4$ and N$_2$ of $\\sim$ 0.35, 0.60 and 0.05, respectively. The most striking difference in the comparison is that the mole fraction of N$_2$ is more than 10 times lower in our results than in theirs, and some compounds considered as minor in the ocean of D89 are not negligible in the composition of our lakes. However, our model reproduces the results of D89 when we adopt precisely the same thermodynamic parameters. The fact that relatively small differences (order of a few percent) between our thermodynamic data extracted from the NIST database and those of D89 lead to significantly different results illustrates the non-linearity of the system of equations determining the composition of the liquid.  It is difficult to quantify the errors on the predicted mole fractions of the different species because our model is based on thermodynamic data much of which is provided either without uncertainties or is derived from extrapolation of data to low temperature. However, numerical tests aiming to investigate the sensitivity of our model to the photochemical models we employ indicate that a $\\pm$ 30\\% variation of the C$_2$H$_6$ precipitation rate induces similar variations for the C$_3$H$_8$ to C$_6$H$_6$ lake mole fractions but poorly alters the N$_2$ to C$_2$H$_6$ lake mole fractions given in Table 3. Moreover, as shown by Fig. 2, a variation of the methane atmospheric abundance hardly affects the lake mole fractions of the precipitates. Pure numerical errors due to the algorithm have been estimated absolutely negligible.  We also outline that the assumption of thermodynamic equilibrium between the lakes and the atmosphere is a crude approximation since in fact the humidity and temperature of the atmosphere in contact with the lakes is determined by dynamic processes such as convection and wind advection (Mitri et al.  2007).  Our solubility calculations imply that a number of species produced by methane photolysis and energetic particle chemistry in Titan's upper atmosphere should be readily detectable with a mass spectrometer carried to the surface of a liquid-filled lake by a Huygens-like entry probe (Coustenis et al. 2009). The measured abundances of multiple minor constituents in the lake, coupled to measurements and models of stratospheric abundances and production rates, and direct temperature measurements of the lake surface, will constrain lake properties that are of interest in understanding the methane hydrologic cycle. For example, at the winter pole a seasonally deposited upper-layer of liquid methane might exist on top of a longer-lived ethane-methane liquid reservoir by virtue of methane's lower density and limited vertical mixing in the cold lakes (Stevenson and Potter 1986). Such a transient layer would be bereft of minor components compared with our values thanks to the slow sedimentation rate of the high altitude aerosols compared to the seasonal (meteorological) methane deposition rate; our solubility values provide a means of calculating the extent to which the longer-lived liquid reservoir below has mixed into the methane meteorological layer.  (The extreme cold of the tropopause of Titan prevents the hydrocarbon constituents other than methane and possibly ethane from passing directly to the lower atmosphere in the gas phase; thus the lakes must be seeded by stratospheric aerosol sedimentation).  \\begin{table}[h] \\caption[]{Chemical composition of lakes at the poles and the equator.} \\begin{center} \\begin{tabular}{lcc} \\hline \\hline \\noalign{\\smallskip} &  Equator (93.65 K)\t\t\t\t\t& Poles (90 K)\t\t\t\\\\ \\hline \\multicolumn{3}{l}{Main composition (lake mole fraction)}\t\t\t\t\t\t\t\t\\\\ N$_2$     \t\t\t\t&  $ 2.95\\times 10^{-3}$  \t\t\t& $4.90\\times 10^{-3}$ \t\t\\\\ CH$_4$   \t\t\t\t&  $5.55\\times 10^{-2}$  \t\t\t& $9.69\\times 10^{-2}$  \t\t\\\\ Ar     \t\t\t\t\t&  $2.88\\times 10^{-6}$ \t\t\t& $5.01\\times 10^{-6}$ \t\t\\\\ CO     \t\t\t\t    &  $2.05\\times 10^{-7}$  \t\t\t& $4.21\\times 10^{-7}$  \t\t\\\\ C$_2$H$_6$   \t\t\t&  $7.95\\times 10^{-1}$  \t\t\t& $7.64\\times 10^{-1}$  \t\t\\\\ C$_3$H$_8$   \t\t\t&  $7.71\\times 10^{-2}$  \t\t\t& $7.42\\times 10^{-2} $  \t\t\\\\ C$_4$H$_8$   \t\t\t&  $1.45\\times 10^{-2}$  \t\t\t& $1.39\\times 10^{-2}$  \t\t\\\\ H$_2$\t\t\t\t    &  $5.09\\times 10^{-11}$\t\t\t& $3.99\\times 10^{-11}$\t\t\\\\ \\hline \\multicolumn{3}{l}{ Solutes (lake mole fraction)}\t\t\t\t\t\t\t\t\\\\ HCN    \t\t\t\t& $2.89\\times 10^{-2}$\t(s)   \t\t& $2.09\\times 10^{-2}$\t(s) \t\\\\ C$_4$H$_{10}$ \t\t& $1.26\\times 10^{-2}$ (ns)  \t\t& $1.21\\times 10^{-2}$\t(ns)\t\\\\ C$_2$H$_2$   \t\t& $1.19\\times 10^{-2}$ (ns) \t\t& $1.15\\times 10^{-2}$\t(ns) \t\\\\ C$_6$H$_6$   \t\t& $2.34\\times 10^{-4}$ (ns)  \t    & $2.25\\times 10^{-4}$\t(ns)    \\\\ CH$_3$CN  \t\t\t& $1.03\\times 10^{-3}$ (ns)  \t\t& $9.89\\times 10^{-4}$\t(ns) \t\\\\ CO$_2$    \t\t\t& $3.04\\times 10^{-4}$ (ns)  \t\t& $2.92\\times 10^{-4}$\t(ns) \t\\\\ \\hline \\end{tabular} \\tablecomments{(s): saturated; (ns) non saturated.} \\end{center} \\label{nomcompo} \\end{table}  A longer-lived ethane-methane lake sampled at the summer pole might still have undergone pole-to-pole transport on timescales of tens of millenia thanks to the precession of perihelion of Saturn's orbit around the Sun (Aharonson et al. 2009). The abundances of minor constituents compared to our values which assume accumulation over geologic time, coupled with the aerosol sedimentation rate, could  be used to ``date'' the liquid reservoir and hence test whether they have been cycled on the ``Milankovitch\" timescale. For species that are chemically inactive and occur as gas only in the atmosphere, such as the noble gases, the lakes provide a second reservoir other than the atmosphere to measure abundances. Variations from the atmospheric value of, for example, $^{40}$Ar/$^{36}$Ar, might hint at contact between the lakes and a much deeper crustal reservoir of liquid methane and ethane not in contact with the atmosphere (Hayes et al. 2008).  Finally, our results provide the chemical data needed to compute the amount of deposition of various hydrocarbons and nitriles in fluvial valleys in the Titan's midlatitudes, as a function of the flow of methane runoff from convective storms, allowing potential tests of models of fluvial erosion (Perron et al. 2006)."
    },
    "0911/0911.5358_arXiv.txt": {
        "abstract": "In the tidal disruption of a star by a black hole, roughly half of the stellar mass becomes bound and falls into the black hole, while the other half is ejected at high velocity.  Several previous studies have considered the emission resulting from the accretion of bound material; we consider the possibility that the unbound debris may also radiate once it has expanded and become transparent.  We show that the gradual energy input from hydrogen recombination compensates for adiabatic loses over significant expansion factors.  The opacity also drops dramatically with recombination, and the internal energy can be radiated by means of a cooling-transparency wave propagating from the surface layers inward.  The result is a brief optical transient occurring $\\sim 1$ week after disruption and lasting 3-5 days with peak luminosities of $10^{40}-10^{42}$~\\ergss, depending on the mass of the disrupted star.  These recombination powered transients should accompany the x-ray/ultraviolet flare from the accretion of bound material, and so may be a useful signature for discriminating tidal disruption events, especially for lower and intermediate mass black holes. ",
        "introduction": "Stars unfortunate enough to pass too close to a massive black hole are torn apart by tidal gravity forces.  In the process, roughly half of the stellar material becomes bound, circularizes, and may eventually accrete onto the black-hole.  The other half of the star is ejected from the system at high differential velocities.  Tidal disruption events offers one means of probing a central black hole in quiescent galaxies.  The dominant observational signature should be a bright x-ray flare powered by accretion of bound material onto the black hole \\citep{Rees_1988,Ulmer_1999,Ramirez_dis09}.  Candidate disruption flares have been detected in x-ray/ultraviolet surveys \\citep{Donley_Flare, Gezari_2006, Esquej_Flare} in some case with counterpart emission in the optical \\citep{Gezari_2009}.  The light curves decline as power laws with luminosities comparable to the Eddington luminosity for $10^6-10^7$~\\msun\\ black holes. Nevertheless, it can be difficult to uniquely identify these flares as tidal disruption events as opposed to other forms of nuclear activity.  It is therefore useful to determine additional discriminating signatures of tidal disruption events. A few possibilities have been suggested.  For deeply penetrating orbits, strong tidal compression perpendicular to the orbital plane leads to the formation of shocks, which breakout out in short ($\\sim 100$~sec) but bright ($L \\sim 10^{43}$~\\ergss) x-ray bursts \\citep{Kobayashi_TD, Brassart_2008, Guillochon_TD}. Later on, streams of material falling onto the black hole may undergo collisions producing short UV flares of luminosity $10^{40}-10^{41}$~\\ergss\\ \\citep{Cannizzo_1990,Kim_TD}.  In addition, the x-ray/UV luminosity from accretion may be absorbed and reprocessed by the unbound debris into optical emission, leading to optical luminosities of $\\sim 10^{40}$~\\ergss\\ \\citep{Bogdi_TD, Strubbe_TD}. If a substantial fraction of the accreting material is blown in a wind, the optical radiation escaping from this outflow may be much brighter, $10^{42}-10^{43}$~\\ergss\\ \\citep{Strubbe_TD}.  Optical emission is of particular interest, as observational surveys are typically most sensitive at these wavelength.  Little attention has been given to  emission from the unbound debris of disruption.  This material, which expands differentially with velocities $\\sim 10^4$~\\kms\\ in some ways resembles a supernova remnant, though usually lacking any radioactive isotopes to power the light curve.  While the initial internal energy of the star can be fairly significant ($\\Eint \\sim 10^{48}$~ergs), most of it will be lost to adiabatic expansion before the ejecta becomes sufficiently transparent to radiate.  For an ideal gas with adiabatic index $\\gamma = 5/3$, the energy and temperature evolve adiabatically as $E \\propto \\rho^{(\\gamma -1)} \\propto \\rho^{2/3}$.  A decline in density by at least a factor of $\\sim 10^9$ is required before the debris might become translucent, at which point the internal energy and temperature will have decreased by a factor $\\sim 10^6$.  Hence the apparent lack of interest in diffusive thermal emission from the unbound remnants of tidal disruption.  There is, however, another important energy reservoir of the debris -- ionization energy, which amounts to $\\sim 10^{46}$~ergs per solar mass of ionized hydrogen and is not subject to adiabatic degradation.  When recombination eventually sets in at temperatures $T_i \\sim 10^4$~K, the energy released per atom $\\epsilon_i$ is an order of magnitude greater than the mean thermal energy, $\\epsilon_i/kT_i \\approx 15$. Recombination is therefore a gradual process, steadily inputting energy to maintain the gas temperature near $T_i$. The thermodynamic evolution is described by $E \\propto \\rho^{(\\gamma_3 -1)}$ where $\\gamma_3$ is the third generalized adiabatic index, now a function of density and temperature.  In the region of partial ionization $\\gamma_3 - 1 \\approx 2 kT/\\epsilon_i \\approx 0.1$ (see \\S\\ref{sec:thermo}) so that the energy decline is indeed gradual.  When the effects of ionization are included, the debris can expand by a much larger factor, and reach much lower densities, before the internal energy is depleted.  When near neutrality is reached, the opacity also drops sharply due to the elimination of electron scattering.  It may be possible for radiation to escape at this time, taking the form of a cooling-transparency wave that propagates from the surface inwards.  The escape of photons will be aided by the aspherical geometry of the debris, which has been stretched by tidal forces into a thin stream, with lower optical depth along the direction perpendicular to the orbital plane  It is therefore plausible that the unbound debris of tidal disruptions gives rise to a brief optical transient powered primarily by the energy tapped from recombination.  We call such an event a recombination transient (RT). The light curves, even if they are dim, might be relevant for upcoming synoptic optical surveys which will probe for the first time extra-galactic transients with luminosities between classical novae ($M_R > -10$~mag) and supernovae ($M_R < -14$~mag).  Here we elaborate on the basic physical ideas and make approximate predictions of the light curves. ",
        "conclusions": "The purpose of this paper is has been to explore whether the unbound debris of tidal disruption might give rise to luminous optical emission, and to determine under what conditions this emission may be observationally relevant.  The analytical results predict a brief (3-5 day) optical transient with peak luminosities in the range $2\\times 10^{40}(M_\\ast/M_\\odot)^{5/2}$~\\ergss.  The energy for the emission is stored mostly in the ionization energy of hydrogen, and is released when the ejecta becomes neutral and transparent at temperatures $T \\approx 5000$~K.  This physics will be relevant not only for the tidal disruption of stars by massive black holes, but to the weak ejection of hydrogen envelopes in general by, e.g., stellar collisions or pulsations.  The predicted bolometric luminosities of the recombination transients fall into the interesting -- that is to say, relatively unexplored -- range between novae and supernovae.  Upcoming synoptic surveys will begin to probe these sorts of intermediate events.  The RT of 1~\\msun\\ stars are relatively dim events which would only be visible in the nearby universe ($\\sim 100$~Mpc).  The disruption of more massive stars ($M \\ga 5~\\msun$), on the other hand, can be quite bright.  Such events may not be as rare as one might first expect. There is strong evidence that the initial mass function of stars within the Galactic center is top heavy \\citep{alexander05}.  Perhaps massive stars are more common in the cores of other galaxies as well.  For massive black holes, $\\Mbh > 10^6~\\msun$, the optical luminosity from the accretion of bound material likely surpasses the emission from the unbound debris.  Nevertheless, the RT may still be discernible as a precursor, and may provide a important complementary signature. While the properties of the accretion flare depends most sensitively on the black hole mass, the luminosity of the RT depends primarily on the properties of the disrupted star, with little dependence on \\Mbh.  In principle observations of both components would allow one to solve for the parameters of the event, providing insight into both the black hole mass and stellar population in the centers of distant galaxies.  More basically, detection of a RT would provide confirmation that the event witnessed was in fact a legitimate tidal disruption,  Recombination transients appear to be of greatest interest for probing lower mass black holes ($M_{\\rm h} < 10^6M_\\odot$). In these systems, the optical luminosity of the transient is likely to exceed that of accretion, and should outshine the nucleus of the galaxy. This offers one means of searching for low mass black holes in dwarf galaxies, or for intermediate mass black hole in globular clusters.  If black holes do lurk in these numerous locations, the recombination transient might be the most prominent optical signature of a stellar disruption."
    },
    "0911/0911.0782_arXiv.txt": {
        "abstract": "Teutsch\\,145 and Teutsch\\,146 are shown to be open clusters (OCs) orbiting well inside the Solar circle, a region where several dynamical processes combine to disrupt most OCs on a time-scale of a few $10^8$\\,yrs. BVI photometry from the GALILEO telescope is used to investigate the nature and derive the fundamental and structural parameters of the optically faint and poorly-known OCs Teutsch\\,145 and 146. These parameters are computed by means of field-star decontaminated colour-magnitude diagrams (CMDs) and stellar radial density profiles (RDPs). Cluster mass estimates are made based on the intrinsic mass functions (MFs). We derive the ages $200^{+100}_{-50}$\\,Myr and $400\\pm100$\\,Myr, and the distances from the Sun $\\ds=2.7\\pm0.3$\\,kpc and $\\ds=3.8\\pm0.2$\\,kpc, respectively for Teutsch\\,145 and 146. Their integrated apparent and absolute magnitudes are $m_V\\approx12.4$, $m_V\\approx13.3$, $M_V\\approx-5.6$ and $M_V\\approx-5.3$. The MFs (detected for stars with $m\\ga1\\,\\ms$) have slopes similar to Salpeter's IMF. Extrapolated to the H-burning limit, the MFs would produce total stellar masses of $\\sim1400\\,\\ms$, typical of relatively massive OCs. Both OCs are located deep into the inner Galaxy and close to the Crux-Scutum arm. Since cluster-disruption processes are important, their primordial masses must have been higher than the present-day values. The conspicuous stellar density excess observed in the innermost bin of both RDPs might reflect the dynamical effects induced by a few $10^8$\\,yrs of external tidal stress. ",
        "introduction": "\\label{Intro}  Regions interior to the Solar circle represent a harsh environment to the long-term survival of open clusters (OCs). The low-mass ones in particular, dissolve into the field in less than $\\approx1$\\,Gyr (e.g. \\citealt{Friel95}; \\citealt{OldOCs}).  Theoretical and N-body predictions (e.g. \\citealt{Spitzer58}; \\citealt{LG06}; \\citealt{BM03}; \\citealt{GoBa06}; \\citealt{Khalisi07}), coupled to observational evidence (e.g. \\citealt{vdB57}; \\citealt{Oort58}; \\citealt{vHoerner58}; \\citealt{Piskunov07}) consistently indicate that the disruption-time scale (\\tdis) near the Solar circle is shorter than $\\sim1$\\,Gyr and depends on cluster mass as $\\tdis\\sim M^{0.62}$ (\\citealt{LG06}). Thus, $75\\la\\tdis(Myr)\\la300$ should be expected for clusters with an initial mass within $10^2 - 10^3\\ms$. Besides, disruption processes are more effective for the more centrally located and lower-mass OCs (see \\citealt{OldOCs} for a review on these effects).  It is in this context that the discovery and characterisation of new OCs towards the inner Galactic regions play an important role. Here we establish the nature and derive astrophysical parameters of the poorly-studied, faint OCs Teutsch\\,145 and 146. Both clusters were discovered by Phillip Teutsch in a systematic survey of several Milky Way fields near the Galactic plane using red, blue, and infrared First and Second Generation DSS images downloaded from the ESO Online Digitized Sky Survey (\\citealt{KTA06}). We are dealing with $1^{st}$ Galactic Quadrant clusters, which makes them particularly suitable to explore dynamical cluster properties within the Solar circle.  Algorithms designed to deal with field-star contamination in densely populated fields, together with other tools to study colour-magnitude diagrams (CMDs) and stellar radial density profiles (RDPs) have been developed by our group in previous 2MASS\\footnote{The Two Micron All Sky Survey, All Sky data release (\\citealt{2mass1997}) - {\\em http://www.ipac.caltech.edu/2mass/releases/allsky/}} studies (e.g. \\citealt{ProbFSR}; \\citealt{AntiC}; \\citealt{LKstuff}). In the near future, VISTA\\footnote{http://www.vista.ac.uk/ } together with other surveys with large telescopes, will deepen by about 4 magnitudes the presently-available near infrared photometry in a large area throughout the Galactic plane. The new - and deeper - photometry will probably require specific algorithms to be analysed. Thus, besides the more direct goal of deriving parameters of two interesting objects, in the present work we apply our analytical tools to the optical data of the two faint open clusters, Teutsch\\,145 and 146, obtained with the 3.58m GALILEO telescope (TNG)\\footnote{http://www.tng.iac.es/ }.  Since they are located in the $1^{st}$ Quadrant (with the associated enhanced disruption rates), the heavy field contamination should be properly taken into account for the intrinsic properties to be assessed. In this context, our main goal in this work is to determine whether such clusters can be characterised as typical OCs or if they present signs of dissolution. In addition, we will derive their fundamental and structural parameters, most of these for the first time.  \\begin{table} \\caption[]{Position and angular size} \\label{tab1} \\renewcommand{\\tabcolsep}{2.2mm} \\renewcommand{\\arraystretch}{1.25} \\begin{tabular}{cccccc} \\hline\\hline Cluster&$\\alpha(2000)$&$\\delta(2000)$&$\\ell$&$b$&D\\\\ &(hms)&($\\degr\\,\\arcmin\\,\\arcsec$)&(\\degr)&(\\degr)&(\\arcmin)\\\\ (1)&(2)&(3)&(4)&(5)&(6)\\\\ \\hline Teutsch\\,145&18:42:29&$-$05:15:12&27.24&$-$0.41&1.9\\\\ Teutsch\\,146&18:51:34&$+$00:11:10&33.11&$+$0.06&1.6\\\\ \\hline \\end{tabular} \\begin{list}{Table Notes.} \\item Col.~6: optical diameter measured in the DSS images. \\end{list} \\end{table}  This paper is organised as follows. In Sect.~\\ref{RecAdd} we provide details on the observations, photometric calibrations and reductions. In Sect.~\\ref{CMDs} we build the colour-magnitude diagrams, discuss the field decontamination and derive the fundamental parameters. In Sect.~\\ref{struc} we derive structural parameters. In Sect.~\\ref{MF} we estimate cluster mass and build the mass functions. In Sect.~\\ref{Discus} we discuss the parameters of both OCs. Concluding remarks are given in Sect.~\\ref{Conclu}. ",
        "conclusions": "\\label{Conclu}  We use BVI photometry obtained with the TNG (3.58m) telescope to derive astrophysical parameters and investigate the nature of the two optically faint and poorly known OCs Teutsch\\,145 ($m_V=12.4$) and Teutsch\\,146 ($m_V=13.3$). Located in the $1^{st}$ Quadrant, both OCs present heavily field-contaminated CMDs, which makes their nature and properties difficult to establish from optical studies in smaller telescopes.  Decontaminated CMDs show that the two clusters exhibit similar properties, basically a well-populated MS together with a few giants. From these we derive ages of $200^{+100}_{-50}$\\,Myr and $400\\pm100$\\,Myr, and distances from the Sun $\\ds=2.7\\pm0.3$\\,kpc and $\\ds=3.8\\pm0.2$\\,kpc, respectively for Teutsch\\,145 and 146. Their mass functions, detected for stars more massive than $\\approx1\\,\\ms$, present slopes similar to Salpeter's IMF. Extrapolated to the H-burning limit, both cluster masses are of the order of $1400\\,\\ms$, which would characterise them as relatively massive OCs. Intrinsically, they are bright OCs, with integrated $M_V\\approx-5.6$ and $M_V\\approx-5.3$, respectively. However, given the limited spatial range of the observations, the present-day mass values may be somewhat higher.  Teutsch\\,145 and 146 are located in the inner Galaxy (more than 2\\,kpc inside the Solar circle), a region where cluster-disruption processes are important. Besides, they are close to the Crux-Scutum arm. Thus, given the ages, their primordial masses must have been higher than the present-day values. With respect to the radial density distribution of stars, they both present a cusp in the innermost region, which might reflect dynamical effects induced by the important external tidal stresses acting along a few $10^8$\\,yrs. The present analysis may shed light on issues such as cluster stability, tidal disruption rates and the future cluster evolution in such harsh environment."
    },
    "0911/0911.0261_arXiv.txt": {
        "abstract": "Faraday Rotation Measure (RM) Synthesis, as a method for analyzing multi-channel observations of polarized radio emission to investigate galactic magnetic fields structures, requires the definition of complex polarized intensity in the wavelength range $-\\infty <\\lambda^2 < \\infty$. The problem is that the measurements at negative $\\lambda^2$ are not possible. We introduce a simple method for continuation of the observed complex polarized intensity $P(\\lambda^2)$ into the domain  $\\lambda^2<0$ using symmetry arguments. The method is suggested in context of magnetic field recognition in galactic disks where the magnetic field is supposed to have a maximum in the equatorial plane. The  method is quite simple when applied to a single Faraday-rotating structure on the line of sight. Recognition of several structures on the same line of sight requires a more sophisticated technique.  We also introduce a wavelet-based algorithm which allows us to consider a set of isolated structures in the ($\\phi,\\lambda^2$) plane (where $\\phi$ is the Faraday depth). The method essentially improves the possibilities for reconstruction of complicated Faraday structures using the capabilities of modern radio telescopes. ",
        "introduction": "Observations of polarized radio emission are the main sources of information on magnetic fields of galaxies. The basic idea of magnetic field analysis from polarized radio emission data originates in the classical paper of \\cite{Burn1966MNRAS.133...67B} (for a later development see \\cite{1998MNRAS.299..189S}). In particular, \\cite{Burn1966MNRAS.133...67B} noted that the complex polarized intensity $P$ obtained from a radio source is related to the Faraday dispersion function $F(\\phi)$  as \\begin{equation} \\label{p_to_f} P(\\lambda^2) =  \\int_{-\\infty}^{\\infty} F(\\phi) e^{2i\\phi \\lambda^2}  d \\phi. \\end{equation} $F(\\phi)$ is the fraction of radiation with the Faraday depth $\\phi$ multiplied by intrinsic complex polarization and it is an important emission characteristic of interest. Here the Faraday depth $\\phi$ is defined by \\begin{equation} \\phi(z) = -0.81\\int_{z}^{0} \\Brpa n_e dz', \\label{fardep} \\end{equation} where $\\Brpa$ is the line-of-sight magnetic field component measured in $\\mu$G, $n_e$ is the thermal electron density measured in cm$^{-3}$ and the integral is taken from the observer at $z=0$ over the region which contains both, magnetic fields and free electrons, and $z$ is measured in parsecs. Following Eq.~(\\ref{p_to_f}) $P$ is the inverse Fourier transform of $F$. Correspondingly, the Faraday dispersion function $F$ is the Fourier transform of the complex polarized intensity: \\begin{equation} \\label{f_to_p1} F(\\phi) = {{1} \\over{\\pi}} \\hat P(k), \\label{Burn} \\end{equation} where $k=2\\phi$, and the Fourier transform is defined as \\begin{equation} \\label{four} f\\left( x \\right) = {{1}\\over{2\\pi}} \\int_{-\\infty}^{\\infty} {\\hat {f}\\left( k \\right)e^{i k x}d k}, \\quad \\hat {f}\\left( k \\right) = \\int_{-\\infty}^{\\infty} {f\\left( x \\right)e^{ - i k x} dx}. \\end{equation}   \\begin{figure}\\centering{ \\includegraphics[width=0.35\\textwidth]{fig1a_Q.eps} \\includegraphics[width=0.35\\textwidth]{fig1b_Q.eps} \\includegraphics[width=0.35\\textwidth]{fig1c_Q.eps} \\includegraphics[width=0.35\\textwidth]{fig1d_Q.eps}} \\caption {RM Synthesis reconstruction of a standard example from \\protect\\cite{2005A&A...441.1217B}: (a) initial $F(\\phi)$ which is chosen purely real; (b) amplitude of $PI(\\lambda^{2})$; (c) $F(\\phi)$  reconstructed with whole domain $\\lambda^{2}>0$: real part - thin solid, imaginary part - dashed, amplitude - thick solid; (d) $F(\\phi)$ reconstructed from the data of spectral band $ 0.6 < \\lambda < 0.78$\\,m. The spectral window of observations is indicated in panel (b) by horizontal dashed line.} \\label{classex} \\end{figure}  Implementation of multichannel spectro-polarimetry on modern radio telescopes provided observations of $P$ over a wide range of $\\lambda$ \\citep[e.g.][]{2000A&A...356L..13H} which made the use of Eq.~(\\ref{Burn}) possible. This is the idea of Faraday Rotation Measure Synthesis (RM Synthesis) \\citep{2005A&A...441.1217B} which opened new perspectives in investigations of magnetic field of galaxies and clusters of galaxies \\citep{2003A&A...403.1031H,2005A&A...441..931D,2008arXiv0804.4594B,2009arXiv0905.3995H}.  A key problem of RM Synthesis application is that $P$ is defined only for $\\lambda^2>0$ and in practice can be observed only in a finite spectral band. Moreover, the maximum of $P$ in practice can be located outside the available spectral band { (see e.g. Fig.~\\ref{classex}b)}. Development of robust methods for the reconstruction of $F$ from $P$ in a given spectral range becomes crucial for the practical implementation of RM Synthesis.  Fig.~\\ref{classex} shows results of RM Synthesis applied to a standard test as exploited by \\cite{2005A&A...441.1217B}. Panel (a) shows the function $F$, which includes three {\\it real-valued} box-like structures, panel (b) - the corresponding polarized intensity $P$ (the dashed horizontal line shows the spectral window $ 0.6 < \\lambda < 0.78$\\,m). We used a channel spacing of $\\delta\\lambda=0.4$\\,cm. Hereafter, $F$ and $P$ are numerically evaluated in arbitrary but mutually consistent units. Note that $F$ is in general a complex-valued function. Its modulus defines the emission and its phase defines the intrinsic position angle. Panel (c) shows the result of the straightforward application of the RM Synthesis algorithm to the physical range $\\lambda^2 >0$, while $P(\\lambda^2)$ is set to zero for all negative $\\lambda^2$. We see that the real part of the reconstructed signal is the same as the initial one (except that it has a twice lower amplitude), however, the reconstructed signal obtains a substantial imaginary part with a shape which is quite remote from the real part. This leads to a change of the emission distribution and a loss of any information concerning the position angle (apart from the central point of the emission region, where the position angle correctly is zero). In the context of chaotic magnetic fields in galaxy clusters this loss is less important \\citep{2005A&A...441..931D}, but in galactic magnetic field studies it becomes crucial because the intrinsic position angle determines the orientation of the regular magnetic field component perpendicular to the line of sight. Fig.~\\ref{classex}d shows that the reconstruction becomes much more difficult if we restrict the data to a relatively narrow spectral band $ 0.6 < \\lambda < 0.78$\\,m. We see that even the sign of the reconstructed real part can be wrong. In that case the algorithm for finite spectral band introduced by \\cite{2005A&A...441.1217B} was used.  A general message obtained from Fig.~\\ref{classex} is that in order to envisage possible ways to get a practical implementation of RM Synthesis one has to include some additional information based on the nature of the physical phenomena which provide the Faraday rotation. Here we concentrate our efforts on the problems associated with missing $P(\\lambda^2)$ for $\\lambda^2<0$. ",
        "conclusions": "The development of multi-channel observations of polarized radio emission opens promising perspectives in the understanding of cosmic magnetic fields on galactic and intergalactic scales. The first fruitful applications of RM Synthesis suggested in this context include the recognition of local structures in the Milky Way \\citep{2003A&A...403.1031H}, clusters of galaxies \\citep{2005A&A...441..931D} and spiral galaxies \\citep{2009arXiv0905.3995H}. However, in general the RM Synthesis algorithm contains a fundamental problem emerging from the fact that the reconstruction formula requires the definition of complex polarized intensity in the range $-\\infty <\\lambda^2 < \\infty$. In this paper we introduce a simple method for continuation of observed complex polarized intensity $P(\\lambda^2)$ into the domain of negative $\\lambda^2<0$. The method is suggested in context of magnetic field recognition in galactic disks, for which the magnetic field strength is supposed to have a maximum in the equatorial plane.  The suggested method is quite simple when applied to a single structure on the line of sight. Recognition of several structures on the same line of sight requires a more sophisticated technique. The problem of structure separation is resolved using the wavelet decomposition. A simple test example demonstrates the applicability of this method. { The polarization angle reconstruction is significantly improved over the standard technique.} The wavelets can be useful to also overcome some other problems of RM Synthesis, related to the multi-band structure of the observational domain in $\\lambda$-space, noise filtration, etc \\citep[e.g.][]{1997ApJ...483..426F,2001MNRAS.327.1145F}. The method essentially improves the possibilities for reconstruction of complicated Faraday structures using the capabilities of modern radio telescopes.  Finally note that our simple examples illustrate that the extension of the observational band into the long-wavelength domain is helpful for the recognition of structures with sharp boundaries, while the short-wavelength domain is crucial for the reconstruction of smooth structures.  \\label{lastpage}"
    },
    "0911/0911.0796_arXiv.txt": {
        "abstract": "{\\bf galaxies: star clusters; globular clusters: general} An overview of our current understanding of the formation and evolution of star clusters is given, with main emphasis on high-mass clusters. Clusters form deeply embedded within dense clouds of molecular gas. Left-over gas is cleared within a few million years and, depending on the efficiency of star formation, the clusters may disperse almost immediately or remain gravitationally bound. Current evidence suggests that a few percent of star formation occurs in clusters that remain bound, although it is not yet clear if this fraction is truly universal. Internal two-body relaxation and external shocks will lead to further, gradual dissolution on timescales of up to a few hundred million years for low-mass open clusters in the Milky Way, while the most massive clusters ($>10^5$ M$_\\odot$) have lifetimes comparable to or exceeding the age of the Universe. The low-mass end of the initial cluster mass function is well approximated by a power-law distribution, $\\rd N/\\rd M \\propto M^{-2}$, but there is mounting evidence that quiescent spiral discs form relatively few clusters with masses $M>2\\times10^5$ M$_\\odot$. In starburst galaxies and old globular cluster systems, this limit appears to be higher, at least several $\\times 10^6$ M$_\\odot$. The difference is likely related to the higher gas densities and pressures in starburst galaxies, which allow denser, more massive giant molecular clouds to form. Low-mass clusters may thus trace star formation quite universally, while the more long-lived, massive clusters appear to form preferentially in the context of violent star formation. ",
        "introduction": "The appeal of star clusters as tools to study (extra)galactic star-formation histories is at least two-fold: the most massive clusters tend to be long-lived, and therefore potentially carry information about the entire star-formation histories of their host galaxies. Furthermore, they are bright and can be observed at much greater distances than individual stars. Young, massive, compact clusters have now been observed in many external galaxies, and it is commonly assumed that at least some of these are young counterparts of the ancient \\emph{globular} clusters (GCs) which are ubiquitous in all major galaxies. Thus, the view that GC formation required unique physical conditions in the early Universe (e.g., Peebles \\& Dicke 1968) has largely been abandoned. However, it remains much less clear how \\emph{direct} the link is between star formation and the formation of (massive) clusters.  In general terms, we might ask the question, \\emph{What are the conditions for the formation and survival of massive clusters?} Clearly, finding an answer is essential if we wish to use such clusters effectively as probes of galaxy formation and evolution. This contribution begins with a broad (and, by necessity, highly incomplete) overview of what is known about the overall properties of star cluster systems. It should then become clear that it is difficult to answer the question as stated above, since there is no clear-cut physical criterion that allows us to decide when a cluster should be classified as `massive'. It is therefore useful to rephrase the problem in a way that makes it more tractable. The remainder of the current review will thus focus on three main themes: (i) the general problem of cluster formation, (ii) dynamical evolution and (iii) the shape of the initial cluster mass function (ICMF). ",
        "conclusions": "The use of star clusters as tracers of extragalactic stellar populations relies on a close link between star formation in general and cluster formation. While several recent studies have converged on a fraction of $\\sim10$\\% of stars forming in clusters that remain bound for at least a few $\\times10^7$ years, the ratio of clusters to field stars varies enormously in old GC systems, both from galaxy to galaxy and within galaxies (Harris 1991; Harris 2010). It is hard to rule out that cluster dissolution is partly responsible for these differences, particularly if it is mass independent (Whitmore \\ea 2007), but so far there is no robust way to predict what fraction of GCs may have been lost over a Hubble time in this way. While the \\emph{timing} of, e.g., a major burst of star formation may be inferred from a peak in the cluster age distribution, the \\emph{strength} of such a burst remains much more poorly constrained.  Apart from differences in the formation efficiency of clusters, the mass spectrum may also vary with environment. There are hints that the ICMF in quiescent discs may be less top heavy than in violent starbursts, so that \\emph{massive} clusters predominantly trace the latter. Although GC formation is no longer viewed as `special', this suggests that GCs formed under conditions that are more similar to those in present-day starbursts than in discs. Low-mass clusters, formed under quiescent conditions long ago, may have dissolved by now.  In the coming years, observations of molecular gas in external galaxies with the {\\sl Atacama Large Millimeter Array} are likely to provide a tremendous boost in our understanding of cluster formation under conditions that differ from those in local star-forming regions. The newly refurbished {\\sl HST} is more capable than ever, and will allow more detailed constraints on the age and mass distributions of clusters in different galaxies, so that the role of environment in determining the ICMF, disruption and the cluster-formation efficiency can be better quantified. On the theoretical front, cosmological simulations will provide a more detailed picture of galaxy formation and evolution, down to scales where individual clusters can be followed (e.g., Bournaud \\ea 2008; Prieto \\& Gnedin 2008)."
    },
    "0911/0911.0333_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\label{sec:1} Galactic H$\\alpha$ emitting objects are tracers of pre and post main-sequence stars as well as of nebulae, cataclysmic variables, Be stars and other more exotic objects like Luminous Blue Variables (LBV) and Wolf-Rayet stars. IPHAS (\\cite{iphas1}, \\cite{iphas2}, \\cite{iphas3}, \\cite{iphas4}) the most complete survey of the galactic plane carried out so far, is complete in the magnitude range $r'=13$ to $20$ for $-5<|b|<5$. The classical surveys, mostly based on objective-prism photographic observations, are complete up to magnitude 9 (see for example \\cite{old1} and \\cite{old2}). In this context, the main aim of our project is to carry out a photometric survey covering the existing gap down to $r'\\approx$14 with observations and data reduction automatically performed. ",
        "conclusions": "\\label{sec:5} We described the software for the robotization of TROBAR, with the aim of performing a survey of the Galactic Plane, looking for H$\\alpha$ emitting objects. The system is able to allocate and observe the targets in a MySQL archive when they are best visible and dependinfg on observing constraints. The code is now almost complete and we foresee to start robotic operations by the end of the present year. \\\\ \\newline"
    },
    "0911/0911.2100_arXiv.txt": {
        "abstract": "We study some of the cosmological imprints of pre-inflationary particles. We show that each such particle provides a seed for a spherically symmetric cosmic defect. The profile of this cosmic defect is fixed and its magnitude is linear in a single parameter that is determined by the mass of the pre-inflationary particle. We study the CMB and peculiar velocity imprints of this cosmic defect and suggest that it could explain some of the large scale cosmological anomalies. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0911/0911.1874_arXiv.txt": {
        "abstract": "{} {The blue straggler phenomenon is not yet well explained by current theory; however, evolutionary models of star clusters call for a good knowledge of it. Here we try to understand the possible formation scenario of HD~73666, a blue straggler member of the Praesepe cluster.} {We compile the known physical properties of HD~73666 found in the literature, focusing in particular on possible binarity and the abundance pattern.} {HD~73666 appears to be slowly rotating, has no detectable magnetic field, and has normal abundances, thereby excluding close binary evolution and mass transfer processes. There is no evidence of a hot radiation source.} {With the use of theoretical results on blue straggler formation present in literature, we are able to conclude that HD~73666 was probably formed by physical collision involving at least one binary system, between 5 and 350\\,Myr (50\\,Myr if the star is an intrinsic slow rotator) ago.} ",
        "introduction": "Blue stragglers are stars that lay on the extension of the main sequence and are bluer and brighter compared to the main sequence turn-off stars. These objects are found in star clusters, dwarf galaxies and in the field. The existence of blue stragglers probably can be explained only by an interaction between two or more stars, and so to understand this phenomenon we study the interaction of stars in stellar systems \\citep[see][]{leonard,bailyn95,sandquist97,sills97,chen2004,ahumada2007}.  Recently, \\citet{ahumada2007} listed the most frequently cited theories to explain the formation of blue stragglers. They could be (1) horizontal-branch stars that appear above the main sequence turn-off point, (2) stars of second or third generation, (3) stars that have extended their main-sequence life due to some internal mixing (this would generate a chemically peculiar blue straggler), (4) stars formed by collision of two single stars, (5) the result of mass transfer in a close binary system, (6) produced by merger of the components of a binary system or (7) the result of collision between two or more binary systems. Of those seven theories they considered the last four as major channels of formation.  Recently, comprehensive catalogues of blue stragglers in open and globular clusters have been published. \\citet{ahumada1995} created  the first consistent catalogue of blue stragglers in open clusters, which was then expanded by \\citet{demarchi} and finally superseded by \\citet{ahumada2007}. These catalogues make it possible to analyse blue stragglers on a solid statistical base, leading, for example, to the conclusion that the number of blue stragglers increases with cluster age \\citep[see Fig.~5 by ][]{ahumada2007}. In the light of this result, evolutionary models of open clusters need to consider this not well understood phenomenon. For this reason it is important to find clues which allow us to distinguish between different blue straggler formation channels.  The A1V star HD~73666 (40 Cnc, $V$ = 6.61) is an extreme blue straggler \\citep{ahumada2007} in the nearby, well-studied cluster Praesepe (NGC~2632). This paper will show that it is a particularly important example, and provides an excellent test for theories of blue straggler formation. This paper can also be seen as the continuation of \\citet{conti1974} in the light of more than thirty years of new astronomical knowledge and instrumental development.  In Sect.~\\ref{sec_membership} we discuss the membership of HD~73666 in the Praesepe cluster. Fundamental parameters and other physical properties are given in Sect.~\\ref{sec_parameters}, where we show that the star is a blue straggler. In Sect.~\\ref{formation} we describe the possible formation scenario. Conclusions are gathered in Sect.~\\ref{sec_discussion}. ",
        "conclusions": "\\label{sec_discussion} We next consider whether the surface CNO abundances could provide information about a stellar collision. We consider the predictions of \\citet{sills2005} on the surface helium and CNO abundance of stragglers formed after the collision of two low mass stars (0.6\\,\\M). They found that the He and CNO abundance change from the original abundance of the two colliding stars but not fast enough to be visible within 350\\,Myr, that is the estimated maximum age of HD~73666 since becoming a blue straggler (see Sect.~\\ref{minmax}).  If instead we consider a collision between a 2~M$_{\\odot}$ star (operating via CNO cycle) and a 0.5~M$_{\\odot}$ star, we would expect that the CNO present in the core would stay in the core of the remnant star, as He does in low mass stars. Note however that the collision of a high mass and a low mass star, as forming mechanism for HD~73666, is somewhat less probable than the collision of two stars of masses near 1\\,$M_\\odot$, considering that in the center of the cluster the mean stellar mass is about 0.8\\,M$_{\\odot}$ \\citep{adams}. It is also important to mention that the models proposed by \\citet{lombardi1996} and \\citet{lombardi2002} show a small mass-loss, during the collision (between 1 and 10\\% of the total mass of the colliding stars).  We conclude that the abundances of CNO in HD~73666, which are similar to those of the turn-off stars, are consistent with collisional formation.  Since we did not find any evidence to contradict the merging and collisional formation scenario, it seems probable that HD~73666 was formed through a collisional mergers of two stars or a merger of the two components of a close binary system or a collisional mergers between binary systems.  It is important to stress that we are not able to establish which of these two formation mechanisms is the most probable for HD~73666. \\subsection{Minimum and Maximum Times Since Formation}\\label{minmax} A {\\it minimum age} can be deduced if the blue straggler was formed by a physical stellar collision. \\citet{sills97,sills01,sills02} showed that the remnants of physical stellar collisions, in particular of off-axis collisions, should be very fast rotating objects with typical rotational velocities similar or greater than the break-up velocity. This brought to the conclusion that if blue stragglers are formed also by stellar collisions a mechanism to reduce the angular momentum must exist. \\citet{sills2005} showed that such spin-down mechanism can be effective through mass loss and that the star loses about 80\\% of its angular momentum within 5\\,Myr, while a blue straggler like HD~73666 need about 1.4\\,Myr to reach the main sequence after the collision. HD~73666 shows an unusual slow rotation for an A1V star, so that if the star is actually slowly rotating the minimum age of the star as a blue straggler could be set close to 5\\,Myr.  \\citet{luca1} derived a non-zero microturbulence velocity: \\vmic\\ = 1.9 $\\pm$ 0.2\\,\\kms, but this is a fairly normal value for a star of this mass, and appears to indicate the disappearing of large scale convection zones, near the surface of the star. This is in agreement with several recent modelling of collisionally formed blue stragglers \\citep[e.g. ][]{sills97,GP08a}.  The {\\it maximum age} of HD~73666 since becoming a blue straggler is given by the age of the isochrone on which it lies in Fig.~\\ref{hr}. This is about 350\\,Myr. If HD~73666 is an intrinsically slow rotator \\citep[\\vsini\\ $\\leq$ 90\\,\\kms;][]{charbonneau}, a { \\it further maximum age limit} is provided by the time it takes for a slowly-rotating A1V star to develop chemical abundance anomalies, such as Am characteristics, by diffusion since the end of convection in its outer envelope. \\citet{talon06} computed evolutionary models with diffusion for stars of 1.7 to 2.5\\,M$_{\\odot}$. Their most relevant model, labeled 2.5P2, for a star of mass 2.5\\,M$_{\\odot}$ and rotation speed 15\\,\\kms, shows clear abundance anomalies before 50\\,Myr, which we can conservatively set as the maximum age of HD~73666. This is, of course, valid if HD~73666 is now an intrinsically slow rotator. The fact that HD~73666 appears so far off the ZAMS does not contradict the given maximum age of 50\\,Myr since collisionally formed blue stragglers appear on the main sequence not on the ZAMS, but at an already evolved stage. \\subsection{Conclusion} We conclude that the Praesepe blue straggler, HD~73666, was likely formed by physical stellar collision and merger of two low-mass stars, between 5 and 350\\,Myr ago (50\\,Myr if the star is an intrinsic slow rotator) if current models are correct.  On the basis of our knowledge of HD~73666 it is not possible to distinguish between a direct stellar collision and binary coalescence as the formation mechanism for HD73666.  HD~73666 could be a perfect object to test current models of collisionally formed blue stragglers. The wide and detailed knowledge available about this star and the environment in which this star is present would allow to test the reliability of current models and to give important constraints for their future development."
    },
    "0911/0911.3089_arXiv.txt": {
        "abstract": "We introduce a new form of coupling between dark energy and dark matter that is quadratic in their energy densities. Then we investigate the background dynamics when dark energy is in the form of exponential quintessence. The three types of quadratic coupling all admit late-time accelerating critical points, but these are not scaling solutions. We also show that two types of coupling allow for a suitable matter era at early times and acceleration at late times, while the third type of coupling does not admit a suitable matter era. ",
        "introduction": "Cosmological observations strongly suggest that the expansion rate of the universe is accelerating and that matter in the universe is dominated by non-baryonic cold dark matter (see e.g.~\\cite{Dunkley:2008ie}). However, what exactly causes this acceleration is not well understood, and one of the main challenges of modern cosmology is to understand the nature of this mysterious dark energy. The existence of some form of dark matter is long known, as implied by the flattened galactic rotation curves observed by Zwicky as early as 1933. Several experiments have been carried out (see, for example~\\cite{Collaboration:2009nf}) in search of candidate dark matter particles. The fact that dark matter only interacts weakly with standard matter means that it is difficult to detect such particles directly. Neither dark energy nor dark matter have been detected directly. Only the total dark sector energy-momentum tensor is known from its combined gravitational effect. In order to separate the two components, we have to assume a model for them. It is possible that these components interact with each other, while not being coupled to standard model particles. Such a possibility can lead to new approaches to the coincidence problem (``how do dark matter and dark energy attain the same order-of-magnitude value at the right time to allow for the observed large-scale structure?\"). It can also produce interesting new features in large-scale structure, such as a large-scale gravitational bias~\\cite{Amendola:2001rc} and a violation of the weak equivalence principle by dark matter on cosmological scales~\\cite{Koyama:2009gd}.  In this paper we study a class of cosmological models with interactions in the dark sector. Various models of the coupling between dark energy and dark matter have been proposed and investigated (see, e.g.~\\cite{Wetterich:1994bg,Amendola:1999qq,Billyard:2000bh,Zimdahl:2001ar, Farrar:2003uw,Chimento:2003iea,Olivares:2005tb, Sadjadi:2006qp,Guo:2007zk,Kim:2007dp,Boehmer:2008av,He:2008tn, Chen:2008ca,Quartin:2008px, Pereira:2008at,Quercellini:2008vh,Valiviita:2008iv}). We consider only the background dynamics -- for cosmological perturbations of coupled dark energy models, see e.g.~\\cite{Amendola:2002bs,Koivisto:2005nr,Olivares:2006jr,Mainini:2007ft, Bean:2007ny,Valiviita:2008iv,Vergani:2008jv,Pettorino:2008ez, Schaefer:2008qs,Schaefer:2008ku,LaVacca:2008kq,He:2008si,Bean:2008ac, Corasaniti:2008kx,Chongchitnan:2008ry,Jackson:2009mz,Gavela:2009cy, LaVacca:2009yp,He:2009mz,CalderaCabral:2009ja,He:2009pd, Koyama:2009gd,Valiviita:2009nu,Majerotto:2009np}.  There is no fundamental theory that selects a specific coupling in the dark sector, and therefore any coupling model will necessarily be phenomenological, although some models will have more physical justification than others. Here we analyse the background dynamics for a new model of coupling. This model improves the one previously introduced in~\\cite{Boehmer:2008av,Valiviita:2008iv}, which was motivated by simple models of inflaton decay during reheating and of curvaton decay to radiation.  The background description of a coupled model with quintessence dark energy density $\\rho_{\\varphi}$ and dark matter density $\\rho_c$ is given by the energy balance equations \\begin{align} \\dot \\rho_c  &= - 3H\\rho_c+Q_c\\,, \\label{cc}\\\\ \\dot \\rho_{\\varphi} &= - 3H(1+w_{\\varphi})\\rho_{\\varphi}+ Q_\\varphi\\,, \\qquad Q_\\varphi=-Q_c := Q \\,. \\label{kg1} \\end{align} Here $Q_A$ is the rate of energy transfer to species $A$. It follows that \\begin{equation} \\label{sq} Q~\\left\\{ \\begin{array}{l} >0\\\\ <0 \\end{array} \\right.\\quad \\Rightarrow \\quad \\mbox{energy transfer is}~\\left\\{ \\begin{array}{l} \\mbox{dark matter $\\to$ dark energy}\\\\ \\mbox{dark energy $\\to$ dark matter} \\end{array} \\right. \\end{equation} The dark energy equation of state parameter is \\begin{equation} w_\\varphi:=\\frac{p_\\varphi}{\\rho_\\varphi}=\\frac{\\frac{1}{2} \\dot\\varphi^2-V(\\varphi)} {\\frac{1}{2}\\dot\\varphi^2+V(\\varphi)}\\,.\\label{wp} \\end{equation} The modified Klein-Gordon equation follows from Eq.~(\\ref{kg1}) as \\begin{equation} \\ddot \\varphi +3H \\dot \\varphi + \\frac{dV}{d\\varphi}= \\frac{Q}{\\dot\\varphi}\\,. \\label{kg} \\end{equation} For quintessence with an exponential potential, \\begin{equation} V(\\varphi) = V_0 \\exp \\left(-\\kappa \\lambda \\varphi\\right)\\,,\\qquad \\kappa^2:=8\\pi G\\,, \\end{equation} where $\\lambda$ is a dimensionless parameter and $V_0>0$. We neglect the radiation and therefore the evolution equations are \\begin{align} \\dot \\rho_b &= - 3H\\rho_b \\,,\\\\ \\dot H &= -\\frac{\\kappa^2}{2}\\left[\\rho_c+\\rho_b+\\dot\\varphi^2 \\right]\\,, \\label{ray} \\end{align} where baryons $\\rho_b$ are not coupled to the dark sector. The Friedman constraint is \\begin{align} \\label{fried} H^2= \\frac{\\kappa^2}{3}\\left(\\rho_c+\\rho_b+\\rho_{\\varphi}\\right). \\end{align}  We define effective equation of state parameters for the dark components which describe the equivalent uncoupled model in the background: $\\dot \\rho_c +3H(1+ \\wc)\\rho_c=0 $, $\\dot \\rho_\\varphi +3H(1+ \\wv)\\rho_\\varphi=0 $. By Eqs.~(\\ref{cc}) and~(\\ref{kg1}), \\begin{equation} \\label{weff} \\wc = \\frac{Q}{3H\\rho_c}\\,,\\qquad \\wv= w_\\varphi - \\frac{Q}{3H\\rho_\\varphi}\\,. \\end{equation} ",
        "conclusions": "We have made a comprehensive analysis of the background dynamics for a new class of models with quadratic dark sector coupling, which are a simple physically-motivated generalization of the linear coupling couplings with constant rates of energy transfer, given in~\\cite{Boehmer:2008av,Valiviita:2008iv,CalderaCabral:2008bx}. The two species which interact are dark matter and quintessence with an exponential self-interaction potential.  Of the three different terms in the general quadratic coupling, we found that the term $Q= \\mathcal B \\rho_{c}^2$ leads to a universe without a standard matter era, whereas the other two terms, $Q= \\mathcal A \\rho_{\\varphi}^2$ and $Q= \\mathcal C \\rho_c\\rho_{\\varphi}^2$, do admit a standard matter era and an evolution that connects this to a late-time attractor. This attractor is accelerated provide the potential is flat enough. The models we have analysed provide us with the following partial answers. These features are valid for both signs of $\\mathcal A$ or $\\mathcal C$, i.e the evolution is not affected by the direction of the energy transfer. But in the $\\mathcal{C}>0$ case the instability of the matter era is more generic; so there is in a way more room for a transition from the matter era to the accelerated attractor. In other words, in theses case there are less restrictions on the initial conditions for this desired transition between asymptotic states to occur.  The critical point for late-time acceleration, G, is not a scaling solution, since \\begin{equation} \\Omega_{c*}=0\\,, \\qquad \\Omega_{\\varphi *}=1\\,. \\end{equation} This is similar to the asymptotic behaviour of the standard $\\Lambda$CDM model. Thus the quadratic models do not produce a constant non-zero and finite ratio $\\Omega_{c*}/\\Omega_{\\phi *}$, and therefore do not address the coincidence problem in this sense. The linear coupling $Q=\\Gamma \\rho_c$ leads the same critical point G; this accelerated solution is an attractor when $\\Gamma>0$, i.e when dark matter is decaying into dark energy.  The quadratic models which admit a viable background evolution can be compared to observations in order to constrain the parameters $\\alpha$ and $\\gamma$. This will require an investigation of the cosmological perturbations in these models."
    },
    "0911/0911.1790_arXiv.txt": {
        "abstract": "We present the results of a search for transient radio bursts of between 0.125 and 32 millisecond duration in two archival pulsar surveys of intermediate galactic latitudes with the Parkes multibeam receiver. Fourteen new neutron stars have been discovered, seven of which belong to the recently identified ``rotating radio transients'' (RRATs) class. Here we describe our search methodology, and discuss the new detections in terms of how the RRAT population relates to the general population of pulsars. The new detections indicate (1) that the galactic z-distribution of RRATs in the surveys closely resembles the distribution of pulsars, with objects up to 0.86~kpc from the galactic plane; (2) where measurable, the RRAT pulse widths are similar to that of individual pulses from pulsars of similar period, implying a similar beaming fraction; and (3) our new detections span a variety of nulling fractions, and thus we postulate that the RRATs may simply be nulling pulsars that are only ``on\" for less than a pulse period. Finally, the newly discovered object PSR J0941--39 may represent a link between pulsars and RRATs. This bizarre object was discovered as an RRAT, but in follow-up observations often appeared as a bright ($\\sim$10 mJy) pulsar with a low nulling fraction. It is obvious therefore that a neutron star can oscillate between being an RRAT and a pulsar. Crucially, the sites of the RRAT pulses are coincident with the pulsar's emission, implying that the two emission mechanisms are linked, and that RRATs are not just pulsars observed from different orientations. ",
        "introduction": "\\label{sec:intro} The transient radio sky has revealed many classes of exotic objects since the discovery of pulsars in late 1967 \\citep{hewishetal68}. While the first $\\sim$50 pulsars were discovered by searching for their single-pulse emission, the advent of filterbanks and digital signal processing % searching in the sub-second radio sky to use more efficient and successful allowed the use of Fourier search techniques, which are more sensitive to periodic objects than single pulse searches and have since permitted the discovery of more than 1500 pulsars. However, the recent discovery of sparsely bursting rotating radio transients (RRATs, \\citealt{mclaughlinetal06}), and a one-off burst tentatively identified as having an extragalactic origin \\citep[][]{ERB} has spurred renewed interest in the search for single-pulse emission after 30 years of Fourier transform searches dominating surveys for pulsars.  The more recent history of single-burst searches commenced with the search of \\citet{shaun} for impulsive signals between 0.001 and 800~ms using a transient event monitoring system at 843~MHz at the Molonglo cylindrical telescope. While they found only pulsars and ambiguous events, they were able to put an upper limit on the event rate of $1-25$~ms duration events of 0.017~events\\,\\,s$^{-1}$\\,sr$^{-1}$ for pulses with energy density $> 10^{-28}$~J\\,m$^{-2}$\\,Hz$^{-1}$. \\citet{nice99} also recognised that some pulsars might be undetectable in periodicity searches and subsequently re-searched data from a 68 deg$^2$ area of sky along the galactic plane for single pulses with duration between $0.5-8$~ms, revealing one new, slow (2.3~s) pulsar. The survey which discovered RRATs \\citep[][hereafter M06]{mclaughlinetal06} was the first large-scale examination of a significant data set for impulsive bursts since the early pulsar surveys, while the search of the Magellanic Cloud surveys \\citep{ERB} and a transient search with Arecibo telescope \\citep[][hereafter D09]{palfa} and a deeper search of the M06 data \\citep[][hereafter K09]{evanrrats} followed using similar techniques.  There are several known sources of fast transient radio emission which we expect to detect as single pulses on sub-second time scales. Periodic rotators which have irregular emission properties include the aforementioned RRATs, which are a form of rotating neutron star, emitting pulses very sparsely at several to many times the rotation period of the star. At present it is unclear from the few observed RRATs how they relate to the neutron star population as a whole, or how they may relate to other sub-populations of neutron stars such as anomalous X-ray pulsars or magnetars. It is possible that they represent only a phase through which all pulsars undergo when they reach a sufficient age \\citep[e.g.][]{mauranewpaper,evan}, or that are simply distant, low-luminosity pulsars with extreme pulse amplitude modulation \\citep{patrick06}. Such distinctions are important for reconciling the so-called ``birthrate problem''; it has been recently noted that the net neutron star birthrate calculated from apparently disparate neutron star populations (X-ray dim isolated neutron stars (XDINS), magnetars, RRATs, pulsars) exceeds the rate of the core-collapse supernovae that are believed to produce them \\citep{popov06,evan}. Recent works are endeavouring to perform analyses that make the distinction clearer from pulse timing (e.g. \\citealt{mauranewpaper}).  High pulse-to-pulse modulation and periodic nulling represent types of signal variance observed in the single pulse emission of some pulsars \\citep[e.g.][]{herfindalrankin07,sometimesapulsar}. Fourier transform searches become less sensitive to such sources due to the objects' erratic emission and the often short duty cycles and long rotation periods of the stars. Such neutron star-related phenomena promise insights into the localisation and power source of radio pulsar emission. Single pulse searches have the ability to recognise these sources, and may serve to reveal a bridging population or define a distinction between regular pulsars that emit steadily and the seemingly more rare ``part-time'' pulsars that only emit detectable emission some fraction of the time, revealing whether RRATs are (in physical mechanism) just extreme examples of the nulling and modulation seen in some pulsars. %  Theoretical arguments predict fast transients from a variety of non-neutron star origins, such as coalescing systems of relativistic objects \\citep{hansenlyutikov01}, evaporating black holes \\citep{rees77}, or supernova events \\citep[e.g.][]{colgatenoerdlinger71}. The aforementioned detection of a single, apparently extragalactic, 5~ms and 30~Jy burst detected by \\citet{ERB} indicated that the detection of such sources in other galaxies might be possible. While the Lorimer et al. detection was the primary motivator for performing our search, fourteen other discoveries associated with intermittent neutron star phenomena resulted, and these are presented and analysed below. No further detections were made of bursts attributable to an extragalactic origin. Fifteen seemingly non-astrophysical, frequency-swept bursts were discovered with a similar frequency-dependent sweep rate (or dispersion measure) to the Lorimer et al. detection (except that they occurred in all 13 instead of just 3 beams of the Parkes multibeam receiver), however a deeper analysis and the interpretation of those signals will be presented elsewhere (see Burke-Spolaor et al., in prep).   This paper presents the detections from a blind search for single pulses in more than 900 hours of archival data from the pulsar surveys of \\citet{ED} and \\citet{BJ}. The data is described in section \\ref{sec:sample}, and the search methodology in section \\ref{sec:methods}. The fourteen new objects and other results uncovered by this search are presented in section \\ref{sec:results}. The implications of these detections are discussed in section \\ref{sec:discussion}.   \\begin{figure*} { \\includegraphics[trim=0mm 5mm 15mm 0mm,clip,width=0.75\\textwidth]{Figs/viewsusdemo.ps} } \\caption{A detection plot for a burst from J1825--33. {\\sc Upper left:} Time series from each of the 13 beams dedispersed at the DM of brightest detection; each field plots the power versus time, normalised by the RMS of the time series. All fields plot the same abscissa and ordinate range, and the beam (and sample in the time series) of brightest detection is indicated by a red star. In the other beams, the sample of peak detection is marked by a vertical dotted line, and the detected width of the pulse is marked by a horizontal bar. {\\sc Lower left:} Time series for the maximum-SNR beam dedispersed at the DM of brightest detection, stacked with the time series from one DM trial higher (``1up''), one trial lower (``1dn''), and the DM~$=0$ time series. As in the upper left panel, the ordinate gives normalised power (SNR), while the abscissa gives time since the beginning of the file (in seconds). {\\sc Bottom right:} A false colour image of the received power as a function of frequency (in MHz) versus time since the beginning of the file (in seconds), revealing the frequency-dependent structure of the detected pulse in the beam of brightest detection. {\\sc Upper right:} Signal-to-noise ratio as a function of DM trial and boxcar filter size, and best-fit information about the detection (DM, SNR, pointing position of the beam of brightest detection, and boxcar width ``tscr''). Pulses with resolved complex structure can show multiple peaks and other structure in their SNR vs. DM curve, though genuinely dispersed pulses (and the curve above at the correct DM$\\pm5$) generally obey the function given by \\citet{cordesmclaughlin}.} \\label{fig:viewsus} \\end{figure*} ",
        "conclusions": "\\label{sec:discussion} \\subsection{Pulsar transient varieties and the nature of RRAT emission}  Regular pulsars are themselves intermittent radio sources, showing detectable emission only for the fraction of time that their beams sweep Earth. The individual pulses, however, are observed to show intrinsic pulse-to-pulse flux and shape differences, the origin of which is yet unknown (though various theories exist). The variations observed in pulsars can be broken down into two main types, being: {\\it pulse-to-pulse modulation}, in which the object emits continuously, however at an intrinsically modulated amplitude, and {\\it nulling}, during which the pulsar beam is somehow quenched, emitting no detectable radiation for a period of time. In many pulsars, nulling appears to be a real phenomenon (with no underlying emission when in null states), with nulls observed to occur from zero to 95\\% of a pulsar's observed rotations \\citep{wang}. Some pulsars appear to null for days-weeks, with a period derivative that is dependent upon the emission state \\citep{sometimesapulsar}.  The range of nulling fractions observed in pulsars is easily described by a windowing function, in which the number of sequential pulses observed at a given time is defined by an on-window of some length. In this description, regularly emitting pulsars have an indefinite on-window, while pulsars with increasing percentages of nulling have on-windows of increasingly short timescale, with extreme nullers (nulling fraction $\\gtrsim75$\\%) becoming more likely to be detected in single pulse searches than periodicity searches. Various degrees of pulsar windowing are illustrated in Figure \\ref{fig:windows}. It is clear that one could incorporate the RRATs into the general pulsar population in this way, by just extending the non-nulling window to a time of length $t<P$, causing only one pulse to be observed at a time. It is a possibility that in this model, pulsars start off as continuous emitters, and gradually increase their nulling fraction until they become RRATs, which ultimately evolve further to emit no radio emission.  The alternative view of RRAT emission, related to pulse-to-pulse modulation, was suggested by \\citet{patrick06}, who indicated that if PSR B0656+14 were observed at several times its distance from Earth, the highly variable single pulses (emitting at up to 420 times the average pulse flux) could emulate the single pulse behaviour observed by M06's distant discoveries. The pulse-energy-distribution analysis of K09, however, makes an argument against this hypothesis (although the analysis was not done on ``classic RRATs'' as we define them below). It appears obvious that there will ultimately be some fraction of RRATs that belong to this class, but deep integrations would be expected to show underlying emission; the real issue, therefore, is in what fraction of RRATs belong to this class, and if some do not, what are the properties of those sources that do not?  Whether an effect of modulation or nulling, the current competing theories of RRAT emission can be broken into three possibilities, in which the RRATs represent: \\begin{itemize} \\item[(1)] a distinct class (intrinsic life-long single-pulsing) % \\item[(2)] faint pulsars with extreme pulse-to-pulse modulation \\item[(3)] an evolutionary phase of pulsars, e.g. where the magnetic field degrades or the spin period increases, leading to increased nulling fractions \\end{itemize} Scenarios (1) and (3) intertwine similarly with questions about pulsars observed to null (relating to whether nulling is an inherent phase of all pulsars or an effect of the local pulsar environment). The clarification of RRATs' nature not only has innate value, but the three scenarios above have quite different implications for how the RRAT birthrate integrates into that of pulsars, and how they contribute to the imbalance of core-collapse supernova rates and neutron star birthrates noted by \\citet{evan}. For instance, in scenario (1) the latent RRAT population may actually exceed the pulsar population. Scenarios (2) and (3), on the other hand, are in part solutions to the birthrate problem, in that the birthrates contributed by RRATs are effectively absorbed into those of pulsars. If RRATs are simply distant, highly modulated pulsars, the sources detected as RRATs are simply abnormalities among the distant, otherwise undetectable pulsars sitting below the average luminosity limit of periodicity surveys; these distant and low-luminosity pulsars are accounted for in pulsar birthrate analyses \\citep[e.g.][]{lorimerandfriends} with statistical estimates, using luminosity-dependent scaling factors.   Although the original RRATs of M06 have a number of apparent properties which set them apart from regular pulsars, as detailed in Sec. \\ref{sec:explore}, we define the {\\it classic RRAT} as an object which emits only non-sequential single bursts with no otherwise detectable emission at the rotation period. When using the term {\\it nulling pulsar}, we refer to an object which emits for some fraction of its rotations, however the emission coming in bursts of several sequential pulses. Finally, we refer to all objects which emit pulses for less than $\\sim$25\\% of their rotations as {\\it extreme pulsar transients}.  Below we discuss the properties of our new transients and compare them with the RRATs of M06 and with the properties of the general pulsar population.  \\subsection{A breakdown of our discoveries}\\label{sec:varieties} What is immediately striking about the 14 objects detected in this search is the variety of dormancy timescales, covering a broad range of the nulling and modulation phenomena related above. For some objects, such as PSR J1850+15, the lack of previous detection by Fourier search techniques can be explained by the presence in the data of intense periodic RFI, which lowers sensitivity to fainter periodic sources. The pulsars PSR J0735--62 and PSR J1850+15 appear to simply be regular pulsars which were likely missed in previous searches due to the modulation and long period of PSR J0735--62, and the presence of interference in the discovery pointing of PSR J1850+15.  The ``classic RRATs'' in our search include PSRs J1129--53, J1226--32, J1654--23, J1753--38, and J1753--12. We also place the objects from which only one pulse was observed in this category; although we have not detected further events from the sources, the pulse widths are typical of pulsars and the pulsation rates implied by our observations are consistent with the range of pulsation rates observed in the RRATs published by M06, which ranged from one per four minutes to one per three hours. We note that the bright single pulses emitted by PSRs J1226--32 and J1753--38 are too bright and have a pulsation rate too high to represent a long tail of a log-normal pulse energy distribution, immediately indicating that scenario (2) above (extreme modulation) is not the case for all RRATs.  This is less clear in the case of PSR J1753--12, and we note that our sources may represent a mix of highly modulated low-luminosity pulsars, and true RRATs.  Apart from PSR J0941--39, the remaining objects appear to be nulling pulsars with high nulling fractions. Only two of these objects were in an ``on'' state for a sufficient fraction of observing time to be Fourier-detectable. Additionally, it is difficult to say whether the source PSR J1825--33 is truly a nulling pulsar, or whether this source had an event which briefly activated its outer gap, as we have since not detected any further emission from the object.  PSR J0941--39 does not fit cleanly into any currently identified pulsar category, sometimes acting like a single-pulse-emitting RRAT, and sometimes like a nulling (though not extreme-nulling!) pulsar. This object may play an important role in understanding RRAT emission.  \\subsection{Exploration of RRATs and the properties of our new discoveries}\\label{sec:explore} This search represents an important step in understanding extreme pulsar intermittents in several ways. First, we have surveyed the largest area to date for such sources in a region off the galactic plane, and added the discovery of 12 highly intermittent pulsars to a small known population of such objects. Given that neutron stars are mainly born from massive stars in the plane of the galaxy, the high galactic latitudes predispose us to sample a population which has had a sufficiently high kick velocity and sufficient time to reach its current position. Also, the automatic RFI rejection and additional candidate visualisation techniques appear to have given us a higher sensitivity to low DM single pulses, as seen by the fact that M06 discovered no sources at \\mbox{DM $<88$~pc\\,cm$^{-3}$} while we have discovered sources at a broad range of DM, with twelve of the fourteen detections at less than \\mbox{88\\,pc\\,cm$^{-3}$}, and with previously known pulsars positively identified down to \\mbox{${\\rm DM}=3$\\,pc\\,cm$^{-3}$.} Additionally, the latitude ranges of past RRAT searches (K06, D09, K09) did not exceed 5\\,degrees. This increased DM sensitivity and high, complementary latitude range makes our discoveries salient objects for assessing the inherent properties of nearby, RRAT-like bursting sources and how neutron stars with extreme intermittence relate to the Galactic population of pulsars.  Throughout this discussion we analyse all of our sources which are sufficiently intermittent to not be detectable in a Fourier search---that is, the ten sources not plotted in Figure \\ref{fig:stacks}. Where appropriate, we distinguish the analysis of ``classic RRATs'' and that of ``nulling pulsars.'' For reference we also include the M06 RRATs and the four D09 objects which, as noted by their study, did not show detectable underlying regular emission. Additionally, we include the five objects which were not noted by K09 to exhibit Fourier-detectable emission or clear nulling behaviour.  The sparse emission and dissimilarity of magnetic field strength to pulsars which emit giant pulses on $\\mu$s timescales \\citep[][]{cognard} are what most definitively led McLaughlin et al. to the conclusion that their sources were a new neutron star population, however several other of the observed properties of the original 11 sources also distinguish them from the general pulsar population. This includes their low duty cycles, high single-pulse luminosities, and their tendency to be distant ($>2$~kpc, which may have partially been due to a selection effect; this particular RRAT feature gives more weight to the claims of \\citealt{patrick06} that RRATs may simply be distant, highly modulated pulsars). To further our understanding of the properties of RRATs, here we explore their various distinguishing features in comparison to the extreme intermittents detected in our search.  The galactic distribution of our new objects (of the classic RRATs and of our 14 discoveries in general) agree well with the distribution of pulsars, concentrated more highly towards the Galactic plane and the Galactic centre. Though not statistically rigourous with a sample of fourteen objects, the galactic $z$-height distribution ($z=d\\cdot{\\rm sin}(l)$) of our discoveries above $z=200$~pc is well-fit with an exponential and a scale height of 300~pc, in good agreement with the galactic pulsar distribution. This agreement and the discovery of RRATs in the high latitude surveys indicate that the RRATs are not a young phenomenon confined to the galactic plane like anomalous X-ray pulsars \\citep[e.g.][]{deadant}. The object of second largest $z$-height in our sample is the RRAT J1610--17 with a scale height of $z=0.86$~kpc. With a birth at $z=0$~pc and a proper motion outward from the galactic plane at 100~km\\,s$^{-1}$, the implied age of this object is $\\gtrsim10^7$ years.   \\begin{figure} { \\includegraphics[angle=270,trim=5mm 15mm 0mm 1mm,clip,width=0.475\\textwidth]{Figs/DUTYvsPER.ps} }  \\caption{The duty cycle versus period for various detections (as detailed in \\S\\ref{sec:varieties}). Among the sources presented in this paper, pulsars are highlighted in blue, RRATs in black, nulling pulsars in grey, and PSR J0941--39 is represented by the red symbols. The duty cycles of our detections are shown in two ways: using the integrated pulse width (measured after all available on-pulses were stacked at the pulsar period) and using the width of the brightest single pulse. M06, D09, and K09 reported only single-pulse widths. The dotted lines represent traces of $w=0.125$~ms and 32~ms, beyond which our search sensitivity to individual pulses drops significantly.} \\label{fig:dutycycles} \\end{figure}   Figure \\ref{fig:dutycycles} plots the duty cycles of known pulsars and the objects analysed here. As expected, the single pulse widths, $w_{\\rm s}$, gleaned from our new objects lie primarily within the optimal sensitivity range set by the box-car matched filters in our search, as described in section \\ref{sec:methods}.  The most notable feature of this diagram is that upon plotting our integrated pulse width measurements (filled circles) with the single pulse width measurements (open circles), it becomes apparent that when using duty cycles measured from individual pulses (as were used in the calculations of M06, D09, and K09), the values measured for our RRAT-like objects become systematically lower by a factor of $\\sim$3, and up to a factor of more than 10 for the extreme case of J0941--39. There may be several causes of this; for pulsars with nulling, pulse drifts and mode changes are often observed, which can work to broaden the integrated pulsar profile (illustrated well by PSR J2033+00's pulses in Fig. \\ref{fig:stacks20} above). Multi-component integrated profiles, made up of a superposition of otherwise narrow pulses emitted from a range in latitudes, can also broaden an object's duty cycle between single and integrated pulse measurements (as has occurred in the extreme case of PSR J0941--39 with its triple-peaked profile). It is clear that were an integrated pulse width measured for the sources of M06, D09, and K09, the duty cycle distribution of RRATs would more closely resemble that of the general pulsar population, instead of appearing to sit characteristically lower than pulsars (the composite profiles for two RRATs shown in \\citet{mauranewpaper} verify this suggestion). As the RRAT and pulsar duty cycle distributions are similar, it is probable that the beaming fraction of such objects will likewise be akin to that of pulsars.   \\begin{figure} \\begin{centering} \\includegraphics[angle=270,trim=3mm 7mm 0mm 7mm,clip,width=0.45\\textwidth]{Figs/LvsD.ps} \\caption{Dependence of luminosity on the distance of the sources. The lower curve represents our single pulse flux sensitivity limit for a pulse width of 32~ms, while the upper curve shows a representative limit above which a highly modulated pulsar would become detectable in a Fourier search, as detailed in \\S\\ref{sec:explore}} \\label{fig:lvd} \\end{centering} \\end{figure}  Next we examine the distribution of the single-pulse luminosities, as shown in Figure \\ref{fig:lvd}. The points cover a broad range in both distance and luminosity, appearing to show a downturn with proximity to the observer. The lower single-pulse sensitivity limit of our search obeying $L\\propto d^2$ is plotted (at 66~mJy for our 32~ms filter). In light of the possibility that scenario (2) (modulation) is true, we consider the following: Fourier transform searches for SNR~$\\geq 8$ pulses as done by \\citet{ED} over the 4.4 minute survey pointings have a sensitivity limit to pulsars (given a typical duty cycle of 5\\%) with an average single-pulse peak flux of $\\langle S_{\\rm peak}\\rangle \\geq 3.3$~mJy. The prominence and statistics of extremely modulated pulsars (like PSR B0656+14) are not yet well characterised, however the study of \\citet{patrick06a} of 187 pulsars at 1~GHz revealed that typical modulation rarely exceeds ten times the average pulse energy (with pulses exceeding this range typically lasting nanoseconds, and called ``giant pulses'', \\citealt[e.g.][]{cairns04}). Their follow-up study of PSR B0656+14, as previously noted, revealed infrequent single pulse emission peaking up to 420 times the average single pulse intensity. Here we consider that such extreme sources will be relatively uncommon, and take a population of hypothetical sources bursting at half the maximum $S_{\\rm max}/\\langle S_{\\rm peak}\\rangle$ ratio of B0656+14, that is, a bursting population with typical $S_{\\rm max} = 210\\langle S_{\\rm peak}\\rangle$. If such a population is contained in our data, sources above the limit of $\\langle S_{\\rm peak}\\rangle \\geq 3.35$~mJy will have bursts at $S_{\\rm max} \\simeq 700$~mJy. Sources brighter than this limit will be detectable in Fourier searches and be tagged as pulsars, however below this limit the periodic emission will have been missed, with only the brightest bursts detectable in single pulse searches. This limit is plotted in Fig. \\ref{fig:lvd}, and interestingly, appears to correspond to the upper envelope of the maximum source luminosities found by our single pulse search.  The natural interpretation of this is that the nearby, highest-luminosity sources exhibiting erratic emission had bright enough low-level emission that instead of being detected as nulling pulsars or RRATs, they were discovered in periodicity searches and taken into the population as pulsars (albeit modulated or nulling ones). When a similar test is done for the RRATs found by M06 (whose longer, 35.5 minute pointings have higher sensitivity to periodic signals and a higher likelihood of containing the rarer brightest outbursts of a highly modulated pulsar; a periodicity search of that survey is described in \\citealt{pksmb}), two sources, PSRs J1317--5759 and J1819--1458, still lie significantly above the limit of the Fourier analysis. In the framework of our interpretation, these two sources are therefore either examples of extreme sources among modulated pulsars (however nearby equivalents among the known pulsar population, apart from B0656+14, have yet to be recognised), or alternatively they are simply examples of genuine RRATs, where the intrinsic pulsar emission is quite strong, however usually invisible because of a very high nulling fraction.  \\begin{figure} \\begin{centering}  \\includegraphics[angle=270,trim=5mm 4mm 0mm 5mm,clip,width=0.45\\textwidth]{Figs/PHIST.ps} \\caption{Normalised period distributions for pulsars and the single-burst emitting RRATs of M06, D09, K09, and this survey.} \\label{fig:periods} \\end{centering} \\end{figure}  Among the four of our RRATs for which we could measure periods, although very statistically poor, the previously noted RRAT feature---that they have a tendency to longer periods than pulsars---is observed in our new objects, where 50\\% have periods longer than one second. Plotted in Figure \\ref{fig:periods} are the normalised distributions for our and the published ``classic RRAT'' sources as compared with the known pulsar population for pulsars with $P>50$~ms (to restrict the analysis to non-recycled pulsars). When a Komolgorov-Smirnov test is run on the un-normalised distributions, it results in a high probability, $P=0.999985$ ($>4\\sigma$), that the distributions are not drawn from the same parent population. Few pulsars have periods greater than $P = 3$~s and because sources such as anomalous X-ray pulsars/magnetars/soft gamma-ray repeaters all have periods greater than a few seconds, it is tempting to draw an association between such sources and RRATs; the recent release of timing data on six RRATs by \\citet{mauranewpaper} show a further similarity of RRATs to these other neutron star types in their placement in the period-period derivative diagram, and a similar result in periodicity distribution differences. However, the broad distribution of source periods and the presence of 7 RRATs with periods below one second suggests that a more basic factor in the population, for instance age, which can create the observed tendency to longer periods. For the RRATs which have had timing analysis done, all appear to have ages of a few times $10^{5-6}$ years. However, the statistical analysis of \\citet{mauranewpaper}, which compared various derived RRAT and pulsar quantities, does not indicate a strong difference in the distribution of RRAT and pulsar ages, though does indicate a strong probability that the distribution of RRAT and pulsar magnetic field strengths are significantly different. There appears to be no relation between RRAT period and distance from the galactic plane (a proxy, albeit a loose one, for pulsar age).  The source PSR J0941--39 contributes a peculiar voice to this discussion, appearing for hours, to possibly weeks at a time like a classic RRAT---erratically emitting single, non-sequential pulses at a relatively constant rate---while at other times appearing as an intense ($\\sim 10$ mJy) pulsar with a $<10$\\% nulling fraction, extraordinary sub-pulse drift, and latitude-dependent modulation. The object appears superficially to tie a close link between nulling phenomena and RRATs, possibly signifying a pulsar which is in some way transitioning from a pulsar mode to an RRAT mode. If this is the case, the source provides a link both between RRATs and pulsars as well as a link to older, non-radio-emitting neutron stars---i.e. that the RRAT phase represents a degradation of the radio emission mechanism before advancing beyond the radio pulsar death line. As a timing solution is acquired for this pulsar, it will be interesting to note whether it is of an advanced age. It is rather surprising how luminous this pulsar is, and what a small nulling fraction it possesses when in its ``pulsar'' state. It is important to note that the site of the RRAT emission is coincident with the three peaks of the profile when the pulsar is on. Thus, the RRAT emission is from the same site as the normal pulsar emission, again implying similar beaming fractions for the two states.   The properties of the new transients found in this search, and particularly of the five new RRATs, give various indicators that RRAT behaviour is well described by either or both of scenarios (2) or (3) above. The idea that observed RRAT behaviour may be an observed consequence of extreme ($>99$\\%) nulling is supported by the high nulling fraction of the objects tagged as ``nulling pulsars'' above; while previously the ``record'' nulling fraction was held by PSR B1713--40 at 95\\% \\citep{wang}, PSRs J1825--33, J1647--36, and J1534--46 all have implied nulling fractions of more than $97$\\%, telling us that these objects may imply a  ``bridging population'' of high-nulling-fraction pulsars, while objects with higher nulling fractions (smaller and less frequent on-windows) exhibit RRAT behaviour. A full statistical nulling fraction analysis of the \\citet{ED} and \\citet{BJ} pulsars may indicate that the nulling fraction of the RRATs discovered in this search fall naturally in the tail of pulsars with significant nulling fractions shown to exist by \\citet{wang} in other surveys. This proposition, along with the lack of RRAT features that are physically distinct from pulsars (except their period distribution)---including their galactic and duty cycle distribution---point to a strong relationship to the ``normal'' radio pulsar population. The story is complicated and strengthened by the apparent mode-switching of PSR J0941--39, which lends significant support for scenario (3) in that it is clear that the object in its RRAT mode is not simply a modulated pulsar, and it is likewise clear that the object is both RRAT and pulsar in nature, signifying a transition between the two populations and more speculatively, implying that nulling pulsars may undergo stepped increases in their nulling fraction with age. The analyses of \\citet{wang} indicate no correlation between nulling fraction and characteristic pulsar age, however the number of pulsars exhibiting nulls appears to increase with proximity to the radio death line (see their Fig. 8). If all RRATs are relics of nulling pulsars, the broadening of integrated pulse measurements supports the fact that despite their sparse emission, the emitted pulses may still be exhibiting drifting or mode-changing phenomena.     We have re-examined the intermediate galactic latitude surveys of \\citet{ED} and \\citet{BJ} for single pulses of radio emission of duration $0.125<w<32$~milliseconds. This search revealed 14 intermittent sources which appear to be neutron stars. Periods ranging 0.2 to 6.2~s were measured for 12 of these sources. While only four of the objects could be detected retrospectively by Fourier search techniques, the remaining 10 exhibited intermittent emission behaviour akin to both extreme nulling pulsars and the classic RRATs discovered by \\citet{mclaughlinetal06}. The sensitivity of our search techniques to bursts of lower dispersion measure allowed us to find objects which were closer than those detected by previous searches.  The general properties of the new sources appear to paint details of the extreme intermittent and RRAT population that concur with having similar emission modes and galactic distribution to pulsars, despite being significantly less emissive than the general population. It was found that the duty cycle distribution of our detections mimics that of pulsars. These similarities, and the discovery in this search of pulsars with extremely high nulling fractions, lead us to suggest that RRATs may represent an extreme type of nulling pulsar whose window of on-activity is shorter than the rotation period of the object, and that the sources observed to exhibit RRAT behaviour are a natural extension of the tail of pulsars observed with high nulling fractions. Furthermore, PSR J0941--39 appears to oscillate between a clear ``nulling pulsar'' state and a more intermittent ``RRAT'' mode. This source may represent the first detection of an object in transition from a bright pulsar with an average nulling fraction to a single-bursting source, closely tying nulling pulsars to RRATs and indicating a possible evolutionary progression. Further studies and timing are warranted and upcoming on this object.  We conclude by noting that identifying RRATs is not a simple task, as they may represent a mixed population of intrinsically single-pulsing pulsars and highly modulated pulsars whose brightest single pulses singly exceed the system noise. Modulated pulsars appearing as RRATs are expected to be few, as among the local population of pulsars, PSR B0656+14 is currently the only pulsar known to be as modulated as the population required for RRAT emission. Nearby pulsars which are part of such a population may in the future be identified by pulse energy distribution studies of nearby highly modulated pulsars, using long observations which have a chance of detecting the rarer brightest outbursts from such objects. Statistical studies of these pulsars may help separate ``true'' RRATs from simply highly modulated pulsars.  To discern the nature of RRATs, other extreme pulsar intermittents, and highly modulated objects, it will be critical in coming years to both identify more such sources, and to continue timing analyses on the known sources to determine their relationship to other neutron star populations. Of additional interest are the statistics of intermittent phenomena in extreme nulling/modulated pulsars among the known pulsar population, the analysis of which may reveal further insight into the relationship between regular pulsars and the vast variety of intermittent behaviours observed in these equivocal neutron stars."
    },
    "0911/0911.3795_arXiv.txt": {
        "abstract": "We compute one--loop corrections to the annihilation of non--relativistic particles $\\chi$ due to the exchange of a (gauge or Higgs) boson $\\varphi$ with mass $\\mu$ in the initial state. In the limit $m_\\chi \\gg \\mu$ this leads to the ``Sommerfeld enhancement'' of the annihilation cross section. However, here we are interested in the case $\\mu \\lsim m_\\chi$, where the one--loop corrections are well--behaved, but can still be sizable. We find simple and accurate expressions for annihilation from both $S-$ and $P-$wave initial states; they differ from each other if $\\mu \\neq 0$. In order to apply our results to the calculation of the relic density of Weakly Interacting Massive Particles (WIMPs), we describe how to compute the thermal average of the corrected cross sections. We apply this formalism to scalar and Dirac fermion singlet WIMPs, and show that the corrections are always very small in the former case, but can be very large in the latter. Moreover, in the context of the Minimal Supersymmetric Standard Model, these corrections can decrease the relic density of neutralinos by more than 1\\%, if the lightest neutralino is a strongly mixed state. ",
        "introduction": "The existence of non--baryonic Dark Matter (DM) in the universe is by now well established \\cite{Silk_rev}. Recent observations determine the universal average DM density quite accurately. The exact value and its uncertainty depend somewhat on the assumptions made in the fit (``priors''). For example, within a minimal $\\Lambda$CDM model, and combining data on the cosmic microwave background (CMB) anisotropies with observations of supernovae of type 1a and with analyses of baryon acoustic oscillations, one finds \\cite{Komatsu2009}: \\begin{equation} \\label{range} \\Omega_{\\rm CDM}h^2 = 0.1131\\pm 0.0034 \\,. \\end{equation} Here $\\Omega_{\\rm CDM}$ is the energy density of cold dark matter in units of the critical density, and $h$ is the scaled Hubble parameter such that $H_0=100 \\, h \\,\\mathrm{km}\\, \\mathrm{sec^{-1}}\\, \\mathrm{Mpc^{-1}}$ where $H_0$ is the current Hubble parameter. Introducing additional parameters in the fit can increase the allowed range; for example, significantly larger values of $\\Omega_{\\rm CDM} h^2$ are allowed if the number of particles that were relativistic when the CMB decoupled is kept free \\cite{Komatsu2009}. However, the minimal model describes the data well. Moreover, quite soon data from the Planck satellite are expected to reduce the error on $\\Omega_{\\rm CDM} h^2$ to the level of 1.5\\% using CMB measurements alone \\cite{planck}.  From the particle physics point of view, Weakly Interacting Massive Particle (WIMPs) are among the most attractive DM candidates. In standard cosmology their thermal relic density is naturally of the right order of magnitude. Owing to their weak interactions, they can be probed through both direct and indirect detection experiments \\cite{Silk_rev}, although no convincing signal has yet been found. Moreover, the existence of WIMPs is {\\em independently} motivated in many extensions of the Standard Model (SM) of particle physics, the most prominent example being extensions based on supersymmetry (SUSY).  In order to fully exploit the precision of present cosmological data, the theoretical error on the prediction for the WIMP relic density should not exceed the error inferred from observations. For a given cosmological model, the WIMP relic density is determined uniquely by its annihilation cross section into SM particles, if all WIMPs were produced thermally. This requires that the post--inflationary Universe was sufficiently hot, with temperature exceeding about 5\\% of the WIMP mass. Moreover, one has to assume that no other, heavier particles decayed into WIMPs after WIMP decoupling. Under these conditions, which are satisfied for standard cosmology, the uncertainty of the current WIMP relic density is essentially given by the uncertainty of the WIMP annihilation cross section. In order to calculate this cross section to percent level accuracy, at least leading radiative corrections will have to be included.\\footnote{These corrections also contribute to the cross sections for WIMP annihilation into SM particles in our galaxy at present times, which affect the size of indirect WIMP detection signals. However, currently there are very large uncertainties in the backgrounds to these signals; in case of charged particles, propagation through the galaxy adds additional uncertainty. Percent level correction to signals for indirect WIMP detection are therefore not significant.}  In this paper we calculate one class of potentially sizable radiative corrections, which are due to the exchange of a boson between the WIMPs prior to their annihilation. If the mass $\\mu$ of the exchanged boson vanishes, the one--loop expression diverges in the limit of vanishing relative velocity $v$ between the annihilating WIMPs. These large ``Sommerfeld'' corrections then have to be re--summed to all orders in the relevant coupling \\cite{summing, Iengo2009, Cassel2009} (for earlier, related work see \\cite{hisano}). However, since WIMPs couple neither to photons nor to gluons, a new light boson with sizable coupling to the WIMPs, but not to SM particles, has to be introduced. The required hierarchy $\\mu \\ll m_\\chi$, where $m_\\chi$ is the mass of the WIMP, then raises naturalness issues.  Here we instead study the case $\\mu \\lsim m_\\chi$. Examples are the exchange of the light Higgs boson in the minimal supersymmetric extension of the SM (MSSM), and the exchange of the $Z$ boson in most ``generic'' WIMP models. If $\\mu \\lsim m_\\chi$ and $v \\ll 1$ these corrections can be significantly bigger than the ``generic'' estimate $\\alpha/\\pi \\sim$ (a few) $\\times 10^{-3}$, while remaining small enough to permit simple one--loop calculations (and without threatening unitarity \\cite{ralston}). We calculate these corrections in the non--relativistic limit, following Iengo \\cite{Iengo2009}. Since WIMPs decouple at temperature $T \\simeq m_\\chi / 20$, an expansion in the relative velocity usually (but not always \\cite{Griest1991}) converges quite fast, and can therefore be used in the calculation of the one--loop correction. We also show how to compute the relevant thermal average over the corrected cross sections. We first apply our formalism to simple models with scalar or fermionic SM singlet WIMPs; the corrections are always small in the former case, but can be sizable in the latter scenario. We then analyze neutralino annihilation in the MSSM, and show that these corrections can reduce the relic density of strongly mixed neutralinos by more than 1\\%, comparable to the projected uncertainty of the value to be inferred from Planck data.  The rest of this paper is organized as follows. In Sec.~\\ref{sec:corr_A} we introduce the formalism, which was initially proposed to treat non--perturbative Sommerfeld enhancement \\cite{Iengo2009}. In particular, we factorize the correction to the amplitude due to the boson exchange between two incoming particles. In Sec.~\\ref{sec:relicdensity} we briefly review the standard calculation of dark matter relic density, and describe the thermal averaging of the corrected cross section. Numerical results for three WIMP models will be presented in Sec.~\\ref{sec:numerical_calculation}. Finally we will summarize. Some details of our numerical procedure are described in Appendix A, while Appendix B shows that our approximate treatment indeed reproduces the leading terms of an exact calculation of radiative corrections associated with the initial state in a simple scalar model. ",
        "conclusions": "We have calculated one--loop corrections to the WIMP annihilation cross section, and the corresponding corrections to the relic density, due to the exchange of a relatively light boson between the WIMPs.  The formalism for calculating the corrections to the cross section is described in Sec.~2. We used a non--relativistic formalism \\cite{itz,Iengo2009}, i.e. we only included small loop momenta. The correction can then also be understood as re--scattering of the WIMPs prior to their annihilation. The motivation for this is that these configurations can give rise to sizable corrections if the exchanged boson is lighter than the WIMP and the relevant coupling is not too small. Note that these corrections are universal, i.e. they are the same for all final states. However, we saw that they do depend on the partial wave of the initial state, being smaller for annihilation from the $P-$wave if the exchanged boson is massive. The main result of this Section is that these corrections can be described by loop functions which only depend on the ratio of the mass of the exchanged boson and the cms three--momentum of the annihilating WIMPs; simple yet accurate parameterizations for these loop functions are given in Eqs.(\\ref{approximation_I}). We also checked explicitly that the correction is independent of the form of the coupling of the exchanged boson, except for the case of axial vector exchange in the $S-$wave, where an additional factor of $-3$ is required. Moreover, the corrections can be used for scalar as well as fermionic WIMPs.  In Sec.~3 we computed the resulting correction to the relic density. We saw that the thermal averaging and the calculation of the ``annihilation integral'' further reduce the relative importance of the corrections in case of annihilation from a $P-$wave initial state. The main model--independent result of our paper is shown in Fig.~\\ref{fig:delj}, which shows the relative size of the corrections to the annihilation integral -- or, almost equivalently, to the relic density -- as a function of the ratio of the masses of the exchanged boson and the WIMP. We saw that for ${\\cal O}(1)$ coupling between boson and WIMP, the corrections are significant unless the boson is heavier than the WIMP.  Finally, in Sec.~4 we applied this formalism to several WIMP models. We showed that these corrections are very small for a scalar singlet WIMP, but can be very large for a Dirac fermion singlet WIMP annihilating into a light scalar singlet. Finally, we analyzed the ``well--tempered'' neutralino, which is a mixture of $U(1)_Y$ gaugino and higgsinos, and found corrections comparable to the anticipated post--PLANCK precision of the observationally determined WIMP relic density only for a narrow range of neutralino masses near 350 GeV. The main reason is that the relevant coupling is always below 0.2 in this case. In a supersymmetric scenario, bigger couplings of WIMPs to Higgs bosons are e.g. possible for ``singlino'' Dark Matter in scenarios where the MSSM is extended by an additional Higgs singlet superfield \\cite{singlino}.  In the MSSM diagonal couplings of the lightest neutralino to gauge or Higgs bosons are always suppressed by mixing angles. Off--diagonal couplings may be large, however. These can give rise to sizable corrections if these couplings involve particles that are only slightly heavier than the lightest neutralino. Such states, with sizable couplings to the lightest neutralino, often exist in regions of parameter space where co--annihilation is important. In order to treat this, one has to extend the formalism presented here to scenarios where the particles in the loop have slightly different masses than the annihilating external particles. Moreover, if co--annihilation with a sfermion is important, there are vertex corrections involving the exchange of an SM fermion, rather than a boson, between the co--annihilating particles. This offers another avenue for future work.   \\subsubsection*{Acknowledgment}  We thank A. Pukhov for his help with using micrOMEGAs. KIN is supported in part by Grants-in-Aid for Scientific Research from the Ministry of Education, Culture, Sports, Science, and Technology of Japan, and by the Grant-in-Aid for Nagoya University Global COE Program, ``Quest for Fundamental Principles in the Universe: from Particles to the Solar System and the Cosmos'', from the Ministry of Education, Culture, Sports, Science and Technology of Japan. MD and JMK are partially supported by the Marie Curie Training Research Networks ``UniverseNet'' under contract no. MRTN-CT-2006-035863, and ``UniLHC'' under contract no. PITN-GA-2009-237920."
    },
    "0911/0911.0663_arXiv.txt": {
        "abstract": "{The presence of vortices in accretion discs has been debated for more than a decade. Baroclinic instabilities might be a way to generate these vortices in the presence of a radial entropy gradient. However, the nature of these instabilities is still unclear and 3D parametric instabilities can lead to the rapid destruction of these vortices.} {We present new results exhibiting a subcritical baroclinic instability (SBI) in local shearing box models. We describe the 2D and 3D behaviour of this instability using numerical simulations and we present a simple analytical model describing the underlying physical process. } {We investigate the SBI in local shearing boxes, using either the incompressible Boussinesq approximation or a fully compressible model. We explore the parameter space varying several local dimensionless parameters and we isolate the regime relevant for the SBI. 3D shearing boxes are also investigated using high resolution spectral methods to resolve both the SBI and 3D parametric instabilities. } {A subcritical baroclinic instability is observed in flows stable for the Solberg-Ho\\\"iland criterion using local simulations. This instability is found to be a nonlinear (or subcritical) instability, which cannot be described by ordinary linear approaches. It requires a radial entropy gradient weakly unstable for the Schwartzchild criterion and a strong thermal diffusivity (or equivalently a short cooling time). In compressible simulations, the instability produces density waves which transport angular momentum outward with typically $\\alpha\\lesssim 3\\times 10^{-3}$, the exact value depending on the background temperature profile. Finally, the instability survives in 3D, vortex cores becoming turbulent due to parametric instabilities. } {The subcritical baroclinic instability is a robust phenomenon, which can be captured using local simulations. The instability survives in 3D thanks to a balance between the 2D SBI and 3D parametric instabilities. Finally, this instability can lead to a \\emph{weak} outward transport of angular momentum, due to the generation of density waves by the vortices.} ",
        "introduction": "The existence of long-lived vortices in accretion discs was first proposed by \\cite{W44} in a model of planet formation. This idea was reintroduced by \\cite{BS95} to accelerate the formation of planetesimals through a dust trapping process. Moreover, large scale vortices can lead to a significant transport of angular momentum thanks to the production of density waves \\citep{JG05b}. Many physical processes have been introduced in the literature to justify the existence of these vortices such as Rossby wave instabilities \\citep{LLC99}, planetary gap instabilities \\citep{VE06,VAAP07}, 3D circulation models \\citep{BM05}, MHD turbulence \\citep{FN05} and the global baroclinic instability \\citep{KB03}.  Baroclinic instabilities in accretion discs has regained interest in the past few years. A first version of these instabilities (the global baroclinic instability or GBI) was mentioned by \\cite{KB03} using a purely numerical approach and considering a disc with a radial entropy gradient. However, the presence of this instability was affected by changing the numerical scheme used in the simulations, casting strong doubts on this instability as a real physical processes. The linear properties of the GBI were then investigated by \\cite{K04} and \\cite{JG05}. However, only transient growths were found in this context. \\cite{K04} speculated that these transient growths could lead to a nonlinear instability explaining the result of \\cite{KB03}, but it is still unclear whether transient growths are relevant for nonlinear instabilities \\citep[see e.g.][and reference therein for a complete discussion of this point]{LL05}. The nonlinear problem was then studied in local shearing boxes by \\cite{JG06} using fully compressible numerical methods. These authors found no instability in the Keplerian disc regime and concluded that the GBI was ``either global or nonexistent''.  Nevertheless, baroclinic processes were studied again by \\cite{PJS07} using anelastic global simulations and including new physical processes. As in \\cite{KB03}, a weak radial entropy gradient was imposed in the simulation. However, a cooling function was also included in their model, in order to force the system to relax to the imposed radial temperature profile. In these simulations, \\cite{PJS07} observed spontaneous formation of vortices with radial entropy gradients compatible with accretion disc thermodynamics. In a subsequent paper \\citep{PSJ07}, these vortices were found to survive for several hundreds orbits. According to the authors, their disagreement with \\cite{JG06} was due to a larger Reynolds number and the use of spectral methods, a possibility already mentioned by \\cite{JG06}.  In this paper, we revisit baroclinic instabilities in a local setup (namely the shearing box) both in the incompressible-Boussinesq approximation and in fully compressible simulations. The aim of this paper is to clarify several of the points which have been discussed in the literature for the past 6 years about the existence, the nature and the properties of baroclinic instabilities. We show that a subcritical baroclinic instability (or shortly SBI) \\emph{does exist} in shearing boxes. This instability is \\emph{nonlinear} (or subcritical) and \\emph{strongly linked to thermal diffusion}, a point already mentioned by \\cite{PSJ07}. We also present results concerning the behaviour of the SBI in 3 dimensions, as vortices are known to be unstable because of parametric instabilities \\citep{LP09}. We begin in \\S\\ref{equations} by presenting the equations, introducing the dimensionless numbers and describing the numerical methods used in the rest of the paper. \\S\\ref{2Dincomp} is dedicated to 2D incompressible-Boussinesq simulations and a qualitative understanding of the instability. We present in \\S\\ref{compressibility} our results in compressible shearing boxes and in \\S\\ref{3Dincomp} the 3D behaviour of the instability. Conclusions and implications of our findings are discussed in \\S\\ref{conclusions}. ",
        "conclusions": "} In this paper, we have shown that a subcritical baroclinic instability was found in shearing boxes. Our simulations produce vortices with a dynamic similar to the description of \\cite{PJS07} and \\cite{PSJ07}. We find that the conditions required for the SBI are (1) a negative Richardson number (or equivalently a disc unstable for the radial Schwarzschild criterion) (2) a non negligible thermal diffusion or a cooling function and (3) a finite amplitude initial perturbation, as the instability is subcritical\\footnote{Note that point (1) and (2) were already mentioned by \\cite{PSJ07}, although in a somewhat different form.}.  When computed in compressible shearing boxes, the vortices sustained by the SBI produce density waves which could lead to a \\emph{weak} outward angular momentum transport ($\\alpha\\sim 10^{-3}$), a process suggested by \\cite{JG05b}. However, we would like to stress that this transport cannot be approximated by a turbulent viscosity, as the transport due to these waves is certainly a non local process. Several uncertainties remain in the compressible stratified shearing box model. In particular, the imposed temperature profile that is subsequently maintained by a heating source is somewhat artificial. This occurs while angular momentum transport causes a significant density redistribution causing the local model to deviate from the original form. Clearly one  should therefore consider global compressible simulations with a realistic thermodynamic treatment to draw any firm conclusion on the long term evolution and the angular momentum transport aspect of the SBI.  Although the vortices produced by the SBI are anticyclones, the precise pressure structure inside these vortices is not as simple as the description given by \\cite{BS95}. In particular, since the vorticity is of the order of the local rotation rate, vortices are not necessarily in geostrophic equilibrium and pressure minima can be found in the centre of strong enough vortices. Moreover, these vortices interact with each other, which leads to complex streamline structures and vortex merging episodes. This leads us to conclude that the particle concentration mechanism proposed by \\cite{BS95} might not work in its present form for SBI vortices.  In 3D, vortices are found to be unstable to parametric instabilities. These instabilities generates random gas motions (``turbulence'') in vortex cores, but they generally don't lead to a proper destruction of the vortex structure itself. More importantly, the presence of a weak hydrodynamic turbulence in vortex cores will lead to a diffusion of dust and boulders, balancing the concentration effect described above. In the end, dust concentration inside these vortices might not be large enough to trigger a gravitational instability and collapse to form planetesimals.  Given the instability criterion detailed above, one can tentatively explain why \\cite{JG06} failed to find the SBI in their simulation. First, we would like to stress that our compressible simulations are not more resolved than theirs. Since these simulations use similar numerical schemes, the argument of a too small Reynolds number proposed by \\cite{PSJ07} seems dubious. We note however that \\cite{JG06} did not include two other important ingredients: a thermal diffusivity and finite amplitude initial perturbations. As shown in this paper, if one of these ingredient is missing, the flow never exhibits the SBI and it gets back to a laminar state rapidly. Note that even if this argument explains the negative result of \\cite{JG06}, it also indicates that the SBI should be absent from the simulations presented by \\cite{KB03}, as no thermal diffusion nor cooling function was used (the initial conditions were not explicitly mentioned). This suggests that either the global baroclinic instability described by \\cite{KB03} and the SBI are two separate instabilities, or strong numerical artefacts have produced an artificially large thermal diffusion in \\cite{KB03} simulations.  In the limit of a very small kinematic viscosity, the scale at which the instability appears has to be related to the diffusion length scale $(\\mu/S)^{1/2}$. In a disc, one would therefore expect the instability around that scale, provided that the entropy profile has the right slope (condition 1). However, as the instability produces enstrophy, vortices are expected to grow, until they reach a size of the order of the disc thickness where compressible effects balance the SBI source term. We conclude from this argument that even if the SBI itself is found at small scale (e.g. because of a small thermal diffusivity), the end products of this instability are large scale vortices, with a size of the order of a few disc scale heights.  We would like to stress that the present work constitutes only a \\emph{proof of concept} for the subcritical baroclinic instability. To claim that this instability actually exists in discs, one has to study the disc thermodynamic in a self consistent manner, a task which is well beyond the scope of this paper. Moreover, it is not possible at this stage to state that the SBI is a solution to the angular momentum transport problem in discs. Although this instability creates some transport through the generation of waves, this transport is weak and large uncertainties remain on its exact value. Finally, the coexistence of the SBI and the MRI \\citep{BH91a} is for the moment unclear. In the presence of magnetic fields, vortices produced by the SBI might become strongly unstable because of magnetoelliptic instabilities \\citep{MB09}. Although the nonlinear outcome of these instabilities is unknown, one can already suspect that vortices produced by the SBI will be weaken in the presence of magnetic fields."
    },
    "0911/0911.2516_arXiv.txt": {
        "abstract": "{Eight Accreting Millisecond X--ray Pulsars (AMXPs) are known to date. Although these systems are well studied at high energies, very little information is available for their optical/NIR counterparts. Up to now, only two of them, SAX J1808.4$-$3658 and IGR J00291+5934, have a secure multi-band detection of their optical counterparts in quiescence.} {All these systems are transient Low-Mass X--ray Binaries. Optical and NIR observations carried out during quiescence give a unique opportunity to constrain the nature of the donor star and to investigate the origin of the observed quiescent luminosity at long wavelengths. In addition, optical observations can be fundamental as they ultimately allow us to estimate the compact object mass through mass function measurements.} {Using data obtained with the ESO-Very Large Telescope, we performed a deep optical and NIR photometric study of the fields of XTE J1814$-$338 and of the ultracompact systems XTE J0929$-$314 and XTE J1807$-$294 during quiescence in order to look for the presence of a variable counterpart. If suitable candidates were found, we also carried out optical spectroscopy.} {We present here the first multi-band ($VR$) detection of the optical counterpart of XTE J1814$-$338 in quiescence together with its optical spectrum. The optical light curve shows variability in both bands consistent with a sinusoidal modulation at the known 4.3 hr orbital period and presents a puzzling decrease of the $V-$band flux around superior conjunction that may be interpreted as a partial eclipse. The marginal detection of the very faint counterpart of XTE J0929$-$314 and deep upper limits for the optical/NIR counterpart of XTE J1807$-$294 are also reported. We also briefly discuss the results reported in the literature for the optical/NIR counterpart of XTE J1751$-$305.} {Our findings are consistent with AMXPs being systems containing an old, weakly magnetized neutron star, reactivated as a millisecond radio pulsar during quiescence which irradiates the low-mass companion star. The absence of type I X--ray bursts and of hydrogen and helium lines in outburst spectra of ultracompact ($P_{\\rm orb} <1$ hr) AMXPs suggests that the companion stars are likely evolved dwarf stars.} ",
        "introduction": "Some Low Mass X--ray Binaries (LMXBs) exhibit sporadic outburst activity while for most of their time they remain in a state of low-level activity; these systems are commonly referred to as X-ray transients \\citep{White84}. Historically, the classification of X-ray transients is made on the basis of their spectral hardness. The outbursts of Soft X-ray Transients (SXTs), characterized by equivalent bremsstrahlung temperatures $\\leq 15$ keV, are often accompanied by a pronounced increase in the luminosity of their (faint) optical counterparts and by the onset of type I X-ray burst activity \\citep{Maraschi77,Woosley76}. During quiescence, these transient systems are very faint in X-rays ($10^{32} - 10^{33}$ erg s$^{-1}$) and their optical luminosity drops by as much as 6-7 mag, giving a unique opportunity for the study of their companion stars. These properties clearly associate SXTs with LMXBs containing an old, weakly magnetized, neutron star, (for a review see, e.g. \\citealt{Campana98a}). In the following we will refer to these systems as Low-Mass X--ray Transients (LMXTs).  It has long been suspected that millisecond radio pulsars are the spun-up products of sustained mass transfer onto neutron stars in LMXBs. According to this scenario, the companions in the LMXBs evolve and transfer matter onto the neutron star on a long time-scale, spinning it up to periods as short as a few ms. This model has gained strong support about ten years ago thanks to the launch of the {\\it Rossi X-ray Timing Explorer} ({\\it RXTE}) satellite that discovered kilohertz quasi-periodic oscillations (kHz QPOs) as well as coherent oscillations during type-I X-ray bursts in a number of LMXBs (e.g. \\citealt{Strohmayer96}). These periodicities are interpreted as the millisecond spin periods of weakly magnetic neutron stars. In April 1998 the first Accreting Millisecond X-ray Pulsar (AMXP) was discovered, with a 401 Hz (2.5 ms) X-ray pulsations (SAX J1808.4-3658; \\citealt{Wijnands98}). This conclusion was further strengthened by the discovery of additional systems in the following years. Eight such systems are now known: SAX J1808.4$-$3658 \\citep{Wijnands98, Chakrabarty98b}; XTE J1751$-$305 \\citep{Markwardt02}; XTE J0929$-$314 \\citep{Galloway02}; XTE J1807$-$294 \\citep{Markwardt03a}; XTE J1814$-$338 \\citep{Markwardt03b, Strohmayer03}; IGR J00291+5934 \\citep{Galloway05}; HETE J1900.1-2455 \\citep{Morgan05} and SWIFT J1756.9-2508 \\citep{Krimm07}\\footnote{Very recently, episodes of coherent ms X-ray pulsations were discovered in archival data of the LMXTs Aql X-1 \\citep{Casella08} and SAX J1748.9-2021 \\citep{Altamirano08}. We will not consider these sources in the present paper.}. These findings directly confirmed evolutionary models that link the neutron stars of LMXBs to those of millisecond radio pulsars, the former being the progenitors of the latter. All these systems are LMXTs, have orbital periods in the range between 40 min and 4.3 hr and spin frequencies from 1.7 to 5.4 ms. These eight accreting millisecond pulsars are well studied at high energies, especially in X$-$rays, both during outburst and quiescence (see \\citealt{Wijnands05b} for a review). On the other hand, with the significant exception of SAX J1808.4$-$3658 and IGR J00291+5934, their optical/NIR quiescent counterparts are only poorly known. The optical light curve of SAX J1808.4$-$3658 in outburst and quiescence shows variability modulated at the orbital period, in antiphase with the X$-$ray light curve \\citep{Giles99, Homer01, Campana04b}. This is unlike other quiescent transients that normally show a double-humped morphology, due to an ellipsoidal modulation, and indicates that the companion star is exposed to some irradiation. The same behaviour is seen in the quiescent optical light curve of IGR J00291+5934 (D'Avanzo et al. 2007)\\footnote{A recent paper of \\citet{Jonker08} pointed out evidence for a change in the quiescent level of IGR J00291+5934 and for strong flaring episodes. A slight phase difference in the folded light curve is also found with respect to that reported in D'Avanzo et al. (2007)}. In a theoretical study on SAX J1808.4$-$3658 \\citet{Burderi03} proposed, on the base of previous works (\\citealt{Stella94}; Campana et al. 1998), that the irradiation is due to the release of rotational energy by the fast spinning neutron star, switched on, as a radio pulsar, during quiescence. Following this idea, Campana et al. (2004) and D'Avanzo et al. (2007) measured the required irradiating luminosity needed to match the optical flux of SAX J1808.4$-$3658 and IGR J00291+5934, respectively, and found that it is a factor of about 100 larger than the observed quiescent X$-$ray luminosity for both systems. Neither accretion$-$driven X$-$rays nor the intrinsic luminosity of the secondary star or the disk can account for it. So, these authors conclude that the only source of energy available within these systems is the rotational energy of the neutron star, reactivated as a millisecond radio pulsar.  Among the eight accretion-powered millisecond X--ray pulsars known to date, four of them are in ultracompact systems, with orbital periods shorter than 60 minutes. Three of these ultracompact systems, XTE J1751$-$305, XTE J0929$-$314 and XTE J1807$-$294 are remarkably homogeneous, with measured orbital periods of 42.4, 43.6 and 40.1 minutes respectively, well below the minimum period for a system with a donor composed primarily of hydrogen ($P_{\\rm orb} \\leq 80$ min; \\citealt{Rappaport82}).  Optical and NIR observations performed in the past by different groups only led to deep upper limits for the counterparts of XTE J1751$-$305 \\citep{Jonker03} and XTE J1814$-$338 \\citep{Krauss05} or to the detection of very faint candidates, if any (\\citealt{Monelli05} for XTE J0929$-$314). No observations of XTE J1807$-$294 during quiescence have been reported to date, while HETE J1900.1-2455 is still in outburst and SWIFT J1756.9-2508 has been discovered very recently (June 2007). The intrinsic faintness of the targets during quiescence, in combination with the large interstellar absorption and high stellar crowding of the relevant fields clearly explain the lack of detections at optical wavelengths. ",
        "conclusions": "We have presented an extensive study of the optical counterparts of Accreting Millisecond X--ray Pulsars during quiescence. The results of our observational campaigns (also presented in Campana et al. 2004 and D'Avanzo et al. 2007), together with data collected from public archives, provided the first comprehensive optical study of the AMXPs class during quiescence. As we discussed in Sec. 1, the eight accreting millisecond X--ray pulsars known to date can be divided into two subclasses using the orbital period as a discriminant parameter: compact ($P_{\\rm orb} > 1$ hr) and ultracompact ($P_{\\rm orb} < 1$ hr) systems. The compactness of these systems implies very low-mass companion stars as inferred from the measured X--ray mass functions. So, it is reasonable to foresee a very low optical luminosity during quiescence, when the main contribution to the total optical flux is expected to come from the companion star. On the other hand, in such narrow systems, irradiation from the compact object might play a key role in the observed luminosity, by heating the inner face of the companion star (this energy is then redistributed through the star; \\citealt{Podsiadlowski91}). We reported the first detection of the quiescent optical/NIR counterparts of XTE J1814$-$338 (Sec. 3.2). Phase resolved photometry of this source revealed a highly sinusoidal optical light curve, modulated at the orbital period with a minimum when the neutron star is seen behind the companion, with respect to the observer. This behaviour, reminiscent of what observed for SAX J1808.4$-$3658 (Homer et al. 2001; Campana et al. 2004) and IGR J00291+5934 (D'Avanzo et al. 2007), suggests a companion star heated by the compact object. This is at variance with the ``classical'' ellipsoidal modulation observed in wider LMXTs. The irradiating luminosity required to account for the observed optical flux is in all cases at least one order of magnitude greater than the observed quiescent X--ray luminosity. An explanation for this observed ``optical excess'' could be found if we assume that these systems host an active millisecond radio pulsar. According to the standard model of formation of millisecond pulsars these neutron stars have been recycled in interacting low-mass X-ray binaries, i.e. have been spun up by accretion after substantial decay of their surface magnetic fields. During quiescence (i.e. when the accretion rate onto the neutron star drops by orders of magnitude or even stops) the radio pulsar may turn on and start to lose rotational energy in the form of a relativistic particle wind that can irradiate the companion star. This wind eventually can be stopped by the pressure of the material flowing from the companion star generating a shock front. The presence of such an extended feature may eclipse the light coming from a residual accretion disk (if present during quiescence), providing an explanation for the puzzling behaviour of the $V-$band light curve of XTE J1814$-$338 we presented in Fig.~\\ref{fig:lc} and discussed in Sec.~3.5.  As for the compact AMXP systems, the detection of the observed optical flux during quiescence of the ultracompact system XTE J0929$-$314 can be accounted for only by invoking an irradiating luminosity about one order of magnitude greater than the quiescent X--ray luminosity (Fig. 8). Taking our estimates of the irradiating luminosity as a lower limit for the spin-down luminosity of a classical rotating magnetic dipole, we could estimate the neutron star's magnetic field for both systems. The derived limits are in agreement with the standard scenario of weakly magnetized ($10^7-10^8$ G) neutron stars to be the progenitors of millisecond radio pulsars (Sec. 3.5 and 4.1.1). We note that no radio pulsations were detected from XTE J0929$-$314 during quiescence. Beaming factor and intrinsic low radio luminosity were suggested as the most likely explanations for this non-detection \\citep{Iacolina09}. XTE J1751$-$305 and XTE J1807$-$294 lie both on the Galactic plane and are therefore heavily absorbed (we estimated $A_V \\sim 5.2$ mag and $A_V \\sim 2.4$ mag, respectively). Even assuming a high spin-down luminosity ($L_{irr} = 1 \\times 10^{34}$ erg s$^{-1}$), the upper limits of the photometry are consistent with a non detection of the quiescent optical counterpart.  As we discussed in Sec. 4, there are two main basic and competing scenarios for the formation of ultracompact binary systems. The first involves an evolved main-sequence or a He-burning companion star (both with a residual presence of hydrogen) while the second predicts for the presence of a He or C/O white dwarf. We note that according to both scenarios the companion star of these ultracompact systems should be a very low-mass ($M_c \\sim 0.01 \\, M_{\\odot}$) low-luminosity star. As shown in Figures 7, 10 and 11, with our data we cannot constrain any of the evolutionary models for ultracompact binary systems. However, as we reported in Sec. 4, the detection of hydrogen from one of these systems would be a key discriminant in determining the validity of the first scenario (which involves an evolved main-sequence or a He-burning star) with respect to the second one (white dwarf companion). Optical spectroscopy of XTE J0929$-$314 carried out during outburst did not reveal any clear evidence of hydrogen or helium lines (Nelemans, Jonker and Steeghs 2006). In addition, we note that, at variance with what is observed for all the compact systems, none of the ultracompact AMXPs showed Type-I X--ray bursts. These events originate from explosive nuclear ignition on the surface of neutron stars. At a critical point, degenerate hydrogen and/or helium burning ignites explosively, suddenly heating up the entire surface, enough to emit strong X-rays \\citep{Maraschi77,Woosley76}. The absence of Type-I bursts among ultracompact AMXPs finds a natural explanation if we assume that there is no hydrogen and/or He left in the donor star, as expected in the case of a C/O white dwarf companion star, suggesting that the other evolutionary channel for these systens is not populated.  With quiescent unabsorbed magnitude in the range $R \\sim 22-23$, compact AMXPs are well within the spectroscopic capabilities of telescopes of the 8-meter class. An estimate of the radial velocity of the companion star can lead to a determination of the optical mass function and, consequently, to place constraints on the mass of the two stars. Firm upper limits on the neutron star mass could be fundamental to discriminate between different equations of state. Phase-resolved spectroscopy of these sources would be the natural continuation of the work presented in this paper."
    },
    "0911/0911.1723_arXiv.txt": {
        "abstract": "We present the results from multi-frequency VLBA observations from 327 MHz to 8.4 GHz of the gigahertz-peaked spectrum radio source PKS\\,1518+047 (4C\\,04.51) aimed at studying the spectral index distribution across the source. Further multi-frequency archival VLA data were analysed to constrain the spectral shape of the whole source. The pc-scale resolution provided by the VLBA data allows us to resolve the source structure in several sub-components. The analysis of their synchrotron spectra showed that the source components have steep spectral indices, suggesting that no supply/re-acceleration of fresh particles is currently taking place in any region of the source. By assuming the equipartition magnetic field of 4 mG, we found that only electrons with $\\gamma \\leq 600$, are still contributing to the radio spectrum, while electrons with higher energies have been almost completed depleted. The source radiative lifetime we derived is 2700$\\pm$600 years. Considering the best fit to the overall spectrum, we find that the time in which the nucleus has not been active represents almost 20\\% of the whole source lifetime, indicating that the source was 2150$\\pm$500 years old when the radio emission switched off. ",
        "introduction": "The radio emission of extragalactic powerful radio sources is due to synchrotron radiation from relativistic particles with a power-law energy distribution. They are produced in the ``central engine'', namely the active galactic nucleus (AGN), and reaccelerated in the hot spot, that is the region where the particles, channelled through the jets, interact with the external medium. The energy distribution of the plasma is described by a power-law: $N(E) = N_o E^{-\\delta}$ that results in a power-law radio spectrum $S_{\\nu}\\propto  \\nu^{- \\alpha}$, with a moderate steepening $S_{\\nu}\\propto \\nu^{- \\alpha + 0.5}$ caused by energy losses in the form of radio emission. The spectral steepening occurs at the break frequency, $\\nu_{\\rm b}$, which is related to the age of relativistic electrons \\citep[see e.g.][]{pacho70}. This model, known as the {\\it continuous injection} (CI) model, requires a continuous injection of power-law distributed electrons into a volume permeated by a constant magnetic field. The model has been applied to the interpretation of the total spectra of the radio sources \\citep[see e.g.][]{mm99}, assuming that the emission of the lobes (supposed to have a constant and uniform magnetic field) dominates over that of core, jets and/or hot spots.\\\\ It is nowadays clear that powerful radio sources represent a small fraction of the population of the active galactic nuclei associated with elliptical galaxies, implying that the radio activity is a transient phase in the life of these systems. The onset of radio activity is currently thought to be related to merger or accretion events occurring in the host galaxy. \\\\ The evolutionary stage of each powerful radio source is indicated by its linear size. Intrinsically compact radio sources with linear size LS $\\leq$ 1 kpc, are therefore interpreted as young radio objects in which the radio emission originated a few thousand years ago \\cite[e.g.][]{cf95, sn00}. These sources are characterised by a rising synchrotron spectrum peaking at frequencies around a few gigahertz and are known as gigahertz-peaked spectrum (GPS) objects. When imaged with high resolution, they often display a two-sided morphology which looks like a scaled-down version of the classical and large (up to a few Mpc) double radio sources \\citep{pm82}. For these sources, kinematic \\citep{pc03} and radiative \\citep{mm03} studies provided ages of the order of 10$^{3}$ and 10$^{4}$ years, supporting the idea that they are young radio sources whose fate is probably to become classical doubles. However, it has been claimed \\citep{alexander00, marecki06} that a fraction of young radio sources would die in an early stage, never becoming the classical extended radio galaxies with linear sizes of a few hundred kpc and ages of the order of 10$^7$ -- 10$^8$ yr. Support for this scenario comes from a statistical study of GPS sources by \\citet{gugliu05}: they showed that the age distribution of the small radio sources they considered has a peak around 500 years. \\\\ So far, only a few objects have been recognised as dying radio sources. They are difficult to find owing to their extremely steep radio spectrum which makes them under-represented in flux-limited catalogues. The objects 0809+404 \\citep{kunert06} and 1542+323 \\citep{kunert05} are possible examples of young radio sources that are fading. These sources, with an estimated age of 10$^{4}$ -- 10$^{5}$ years, have been suggested as ``faders'' due to the lack of active structures, like cores and hot-spots, although their optically-thin radio spectra do not display the typical form expected for a fader.\\\\  In this paper, we present results from multi-frequency VLBA and VLA data of the GPS radio source PKS\\,1518+047 (RA=15$^{h}$ 21$^{m}$ 14$^{s}$ and DEC= 04$^{\\circ}$ 30$^{'}$ 21$^{''}$, J2000), identified with a quasar at redshift $z = 1.296$ \\citep{sk96}. It is a powerful radio source (Log $P_{\\rm 1.4\\, GHz}\\,\\, {\\rm (W/Hz)} = 28.53$), with a linear size of 1.28 kpc. This source was selected from the GPS sample of \\citet{cstan98} on the basis of its steep spectrum $\\alpha_{\\rm 1.4 GHz}^{\\rm 8.4 GHz} = 1.2$, uncommon in young radio sources. The goal of this paper is to determine by means of the analysis of the radio spectrum and its slope in the optically thin region, whether this source is in a phase where its radio activity has, perhaps temporarily, switched off.\\\\  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. At the redshift of the target 1$^{''}$ = 8.436 kpc. The spectral index is defined as $S {\\rm (\\nu)} \\propto \\nu^{- \\alpha}$.\\\\ ",
        "conclusions": "We presented results from multi-frequency VLBA and VLA observations of the GPS radio source PKS\\,1518+047. The analysis of the spectral index distribution across the whole source structure showed that all the source components have very steep optically-thin synchrotron spectra. The radio spectra are well explained by energy losses of relativistic particles after the cessation of the injection of new plasma in the radio lobes. This result, together with the lack of the source core and active non-steep spectrum components like hot spots and knots in the jets, suggests that no injection/acceleration of fresh particle is currently occurring in any region of the source. For this reason, PKS\\,1518+047 can be considered a fading young radio source, in which the radio emission switched off shortly after its onset. When the supply of energy is switched off, given the high magnetic fields in the plasma, the spectral turnover moves rapidly towards low frequencies, making the source undetectable at the frequencies commonly used for radio surveys. However, in PKS\\,1518+047 the time elapsed since the last particle acceleration is of the order of a few hundred years, suggesting that PKS\\,1518+047 still has a GPS spectrum because adiabatic losses have not had enough time to affect the source spectrum, shifting the peak far from the GHz regime. If the interruption of the radio activity is a temporary phase and the radio emission from the central engine will re-start soon, it is possible that the source will appear again as a GPS without the severe steepening at high frequencies. If this does not happen, the fate of this radio source is to emit at lower and lower frequencies, until it disappears at frequencies well below the MHz regime.\\\\"
    },
    "0911/0911.4276_arXiv.txt": {
        "abstract": "{For almost one year the Large Area Telescope on board the \\textit{Fermi} observatory has been surveying high-energy phenomena in our Universe. We will present an overview of the status of the mission and of some results from the first year of observations, focusing on the topics of particular interest for the high-energy Physics community: detection of high-energy $\\gamma$-ray bursts, the discovery of new populations of $\\gamma$-ray sources, non-confirmation of the excess of diffuse GeV $\\gamma$-ray emission seen by EGRET and, in greater detail, the recent measurement of the cosmic-ray electron spectrum from 20 GeV to 1 TeV.}  \\FullConference{The 2009 Europhysics Conference on High Energy Physics,\\\\ July 16 - 22 2009\\\\ Krakow, Poland}  \\begin{document}  Our Universe is home to violent phenomena, which can reach energies much higher than the collisions we produce in our particle accelerators. Therefore, the observations of high-energy astrophysical phenomena can provide deep insights on fundamental Physics, as well as information for understanding how our Universe works. In the last decades we expanded our investigation to a wide range of messengers, from charged cosmic rays to gravitational waves, from neutrinos to $\\gamma$-rays. High-energy $\\gamma$-rays play a key role, because they are relatively easy to detect and they have low interaction probabilities during their propagation. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0911/0911.4873.txt": {
        "abstract": " ",
        "introduction": "\\label{chap2} Actuellement, la physique des acc\u00e9l\u00e9rateurs nous permet d'atteindre des \u00e9nergies de l'ordre de 14 TeV dans le centre de masse. Ce seuil est associ\u00e9 aux particules supersym\u00e9triques ou \u00e0 des constituants des quarks et des leptons; ceci pourra \u00eatre v\u00e9rifi\u00e9 dans l'avenir avec des acc\u00e9l\u00e9rateurs puissants tel que le grand collisionneur de hadrons LHC (Large Hadron Collider) au CERN. Un deuxi\u00e8me seuil pourrait \u00eatre associ\u00e9 \u00e0 la grande unification qui se propose d'interpr\u00e9ter les trois interactions \u00e9lectromagn\u00e9tique, nucl\u00e9aire forte et faible et pouvant se manifester \u00e0 des \u00e9nergies de l'ordre de $10^{14}eV$. Parmi les cons\u00e9quences de ces th\u00e9ories, la d\u00e9sint\u00e9gration du proton et l'existence de certaines particules exotiques tel que les monop\u00f4les magn\u00e9tiques, les nucl\u00e9arites, etc.... L'observation de telles particules ne peut se faire aupr\u00e8s des acc\u00e9l\u00e9rateurs. Dans ce cas, les exp\u00e9riences sans acc\u00e9l\u00e9rateurs, en particulier les exp\u00e9riences souterraines, jouent un r\u00f4le tr\u00e8s important pour la recherche des particules traces et la violation de certaines lois de conservation.\\\\ Les exp\u00e9riences souterraines sont souvent install\u00e9es dans des sites miniers ou dans des tunnels, la roche couvrant l'exp\u00e9rience joue le r\u00f4le de filtre qui absorbe les particules charg\u00e9es des RC dont le nombre est \u00e9lev\u00e9 au niveau du sol et peuvent rendre impossible la d\u00e9tection des \u00e9v\u00e8nements peu fr\u00e9quents. Les exp\u00e9riences souterraines sont nombreuses et sont destin\u00e9es \u00e0 diff\u00e9rents sujets de recherche en physique fondamentale. N\u00e9anmoins on peut les diviser en deux types, d'une part celles qui se basent sur l'\u00e9tude des rayons cosmiques et d'autre part celles qui ne d\u00e9pendent pas de ces derniers.\\par Les exp\u00e9riences souterraines qui ne sont pas destin\u00e9es \u00e0 l'\u00e9tude des rayons cosmiques s'int\u00e9ressent g\u00e9n\u00e9ralement aux probl\u00e8mes de violations des nombres leptonique ou baryonique en particulier la d\u00e9sint\u00e9gration du proton \\cite{MAC1}.\\par Les exp\u00e9riences souterraines bas\u00e9es sur les informations fournies par l'\u00e9tude des rayons cosmiques ont pour objectifs de r\u00e9soudre les probl\u00e8mes relatifs : \\begin{enumerate} \\item \u00e0 l'astrophysique tel que: \\begin{itemize} \\item  l'\u00e9tude de la composition et du spectre \u00e9nerg\u00e9tique des RC \\item l'\u00e9tude de la variation temporelle et spatiale des RC, recherche des sources ponctuelles et d\u00e9tection des neutrinos de l'effondrement gravitationnel. \\item l'\u00e9tude des neutrinos atmosph\u00e9riques et leur oscillation. \\end{itemize} \\item aux particules \u00e9l\u00e9mentaires et \u00e0 la recherche des particules exotiques telles que les \t\t\t\t\tmonop\u00f4les magn\u00e9tiques, les nucl\u00e9arites, les Q-balles, les WIMPs, les particules supersym\u00e9triques, etc... \\end{enumerate} En g\u00e9n\u00e9ral le nombre d'\u00e9v\u00e8nements attendus est tr\u00e8s petit ce qui n\u00e9cessite des d\u00e9tecteurs assez massifs et volumineux.\\par Les d\u00e9tecteurs souterrains utilis\u00e9s sont : les scintillateurs liquides, les d\u00e9tecteurs Cherenkov \u00e0 eau , les calorim\u00e8tres avec reconstruction de traces, les tubes \u00e0 streamer et les d\u00e9tecteurs plastiques. Pour l'\u00e9tude des RC de haute \u00e9nergie (astronomie des muons et des neutrinos), on choisit g\u00e9n\u00e9ralement un des trois premiers types de d\u00e9tecteurs ou la combinaison de deux d'entre eux, afin d'augmenter la surface de d\u00e9tection et d'avoir une bonne r\u00e9solution spatiale. Dans le cas de l'\u00e9tude de la d\u00e9sint\u00e9gration du proton, la comp\u00e9tition se fait entre les d\u00e9tecteurs Cherenkov \u00e0 eau et les calorim\u00e8tres.\\par Parmi les laboratoires souterrains nous citons: Le laboratoire Baksan en Russie, l'exp\u00e9rience Homestake, l'exp\u00e9rience IMB et Soudan 1-2 aux Etats Unis, l'exp\u00e9rience NUSEX sous le Mont Blac en France, l'exp\u00e9rience Kolar Gold Field en Inde, l'exp\u00e9rience Kamiokande et super-Kamiokande au japan enfin le laboratoire de Gran Sasso en italie.  \\hskip 12pt Il est g\u00e9n\u00e9ralement admis que la distribution du temps d'arriv\u00e9e des rayons cosmiques galactiques est un ph\u00e9nom\u00e8ne al\u00e9atoire. Cependant il existe quelques m\u00e9canismes introduisant des modulations dans cette distribution. La d\u00e9tection des rayons gamma de haute \u00e9nergie ($10^{14}-10^{16}\\;eV$) provenant des sources galactiques ou extragalactiques a conduit aux d\u00e9veloppement de diff\u00e9rents mod\u00e8les qui d\u00e9crivent ces sources dans l'espace environnant (pulsars, syst\u00e8mes binaires, etc...).\\\\ Weekes \\cite{dtr1} a sugg\u00e9r\u00e9 que l'analyse des rayons cosmiques (RC) peut r\u00e9v\u00e9ler la pr\u00e9sence de composantes dont l'origine est un pulsar ou un objet c\u00e9leste similaire qui n'est pas entour\u00e9 de n\u00e9bulosit\u00e9, ayant une courte p\u00e9riode et qui n'est pas situ\u00e9 \u00e0 de grandes distances. De ce fait, une analyse des distributions temporelles des rayons cosmiques est importante. Dans le cas des rayons cosmiques primaires charg\u00e9s, les modulations des RC sont r\u00e9duites ou \u00e9limin\u00e9es par l'effet des champs magn\u00e9tiques galactiques et extragalactiques pr\u00e9sents entre la source et la terre. Ceci provoque la perte de la p\u00e9riodicit\u00e9 du signal provenant des sources situ\u00e9es \u00e0 une distance inf\u00e9rieure \u00e0 150~pc. Les m\u00eames effets sont observ\u00e9s dans le cas des RC arrivant d'une supernovae. Les rayons cosmiques de haute \u00e9nergie gardent la p\u00e9riodicit\u00e9 du signal, ils peuvent pr\u00e9server la modulation et donc on ne peut exclure la pr\u00e9sence de structures dans la distribution temporelle. Il est donc int\u00e9ressant d'examiner ces distributions ind\u00e9pendamment de toutes les hypoth\u00e8ses.\\par Plusieurs recherches utilisant des d\u00e9tecteurs souterrains ou en surface exploitant \"les gerbes \u00e9lectromagn\u00e9tiques dans l'air (EAS)\", n'ont pas report\u00e9 de modulation dans la distribution temporelle des RC de haute \u00e9nergie et par cons\u00e9quent confirment l'arriv\u00e9e al\u00e9atoire des rayons cosmiques. Parmi ces \u00e9tudes nous citons: \\begin{itemize} \\item Morello et al. \\cite{dtr4}, \u00e0 l'aide des mesures EAS avec deux diff\u00e9rents seuils d'\u00e9nergie $E>50\\;GeV$ et $E>2\\;10^{4}\\;GeV$. \\item La collaboration NUSEX \\cite{dtr5}, est une exp\u00e9rience souterraine sous le Mont blanc utilisant un d\u00e9tecteur compos\u00e9 de plans de Fer s\u00e9par\u00e9s par des plans de tube \u00e0 streamer. Cette exp\u00e9rience a mesur\u00e9 le flux des muons \u00e0 une profondeur de $400\\;m.w.e.$ \\footnote{$\\ 1\\  m.w.e (1 meter water equivalent)= 1hg/cm^2$} \\item La collaboration MACRO \\cite{MAC01}, a mesur\u00e9 le flux des muons singuliers et multiples \u00e0 l'aide du d\u00e9tecteur MACRO fonctionnant avec 2 et 4 supermodules seulement et en utilisant uniquement la partie inf\u00e9rieure du d\u00e9tecteur. \\end{itemize} \\par Deux exp\u00e9riences ont confirm\u00e9 la pr\u00e9sence d'une composante non al\u00e9atoire dans la distribution du temps d'arriv\u00e9e des rayons cosmiques de haute \u00e9nergie : \\begin{enumerate} \\item Bhat et al. \\cite{dtr7} ont analys\u00e9 les impulsions de la lumi\u00e8re Cherenkov produites dans l'atmosph\u00e8re par des RC d'\u00e9nergie primaire $E_{p}\\geq 100\\;TeV$. Ils ont observ\u00e9 que la corr\u00e9lation se produit au temps $t<40s$. L'exp\u00e9rience a utilis\u00e9 deux photomultiplicateurs avec un grand angle d'observation \u00e0 2743 m  d'altitude.\\\\ La figure (\\ref{fig:dt1}) donne la distribution des temps d'arriv\u00e9e de deux \u00e9v\u00e8nements successifs et montre une d\u00e9viation sup\u00e9rieure \u00e0 $5\\;\\sigma$ pour $\\Delta t<40s$. \\begin{figure} \\centering \\includegraphics{figures/bhat.eps} \\caption{\\em \\textbf{R\u00e9sultat de l'exp\u00e9rience de Bhat et al. \\cite{dtr7}. L'histogramme montre la distribution observ\u00e9e des temps d'arriv\u00e9e des \u00e9v\u00e8nements dans un bin de 10s. La courbe repr\u00e9sente l'ajustement par la fonction exponentiel th\u00e9orique (en consid\u00e9rant que l'arriv\u00e9e des \u00e9v\u00e9nements est compl\u00e8tement al\u00e9atoire). Une d\u00e9viation de cette distribution a \u00e9t\u00e9 observ\u00e9e pour $t<40s$.}} \\label{fig:dt1} \\end{figure} \\item Badino et al. \\cite{dtr8}, en utilisant un d\u00e9tecteur souterrain dans le tunnel du Mont blanc \u00e0 une profondeur de $5000 hg/cm^2$ de roche, ont observ\u00e9 un exc\u00e8s de muons en \u00e9tudiant la distribution temporelle des muons. Les muons d\u00e9tect\u00e9s ont une \u00e9nergie sup\u00e9rieure \u00e0 $E_\\mu=380GeV$ et sont originaires des RC primaires d'\u00e9nergie $E_p\\approx\\;50TeV$. L'exc\u00e8s a \u00e9t\u00e9 observ\u00e9 pour un intervalle de temps s\u00e9parant deux \u00e9v\u00e8nements successifs d'environ $38s$. Cet exc\u00e8s de muons (1\\%) mis en \u00e9vidence dans un intervalle de temps de quelques dizaines de secondes, r\u00e9v\u00e8le la pr\u00e9sence d'une composante modul\u00e9e superpos\u00e9e \u00e0 la distribution temporelle poissonnienne. \\end{enumerate} La suite de ce chapitre, est consacr\u00e9 \u00e0 l'\u00e9tude des distributions de temps d'arriv\u00e9e des muons cosmiques, r\u00e9alis\u00e9e avec les donn\u00e9es du d\u00e9tecteur MACRO fonctionnant avec l'ensemble de ces supermodules (six). Notre objectif est de conclure sur la nature de ces distributions et ainsi se situer par rapport aux r\u00e9sultats obtenus par les exp\u00e9riences cit\u00e9es ci-dessus.\\par Un grand \u00e9chantillon de muons a \u00e9t\u00e9 utilis\u00e9 compar\u00e9 \u00e0 celui utilis\u00e9 dans les travaux ult\u00e9rieurs de MACRO \\cite{MAC1}. La collecte de tel \u00e9chantillon \u00e9tait possible en utilisant la p\u00e9riode durant laquelle le d\u00e9tecteur fonctionnait avec les six supermodules y compris la partie sup\u00e9rieure du d\u00e9tecteur \"attico\". Il faut noter que cette exp\u00e9rience est en mesure de d\u00e9tecter environ $6\\;10^6$ muons/ans gr\u00e2ce \u00e0 sa grande acceptance. De ce fait, nous nous proposons d'\u00e9tudier le temps d'arriv\u00e9e des muons de diff\u00e9rentes \u00e9nergies collect\u00e9s durant la p\u00e9riode allant de 1995 jusqu'\u00e0 juin 2000. Cette analyse porte sur trois types d'\u00e9v\u00e8nements : \\begin{itemize} \\item \u00e9v\u00e8nements singuliers (multiplicit\u00e9 \u00e9gale \u00e0 1). Un tel \u00e9v\u00e8nement correspond \u00e0 un $\\mu$ simple qui est reconstruit aussi bien sur les fils que sur les strips. Ces muons sont g\u00e9n\u00e9r\u00e9s par des rayons cosmiques primaires ayant des \u00e9nergies de l'ordre de 20 TeV. Avec une telle \u00e9nergie, la particule charg\u00e9e des RC n'est pas influenc\u00e9e par les modulations solaires. Ces muons nous permettent d'obtenir des informations sur les RC primaires ayant une \u00e9nergie dans le domaine $10^{13}-10^{14}$ eV. \\item \u00e9v\u00e8nements doubles (multiplicit\u00e9 \u00e9gale \u00e0 2). Un tel \u00e9v\u00e8nement correspond \u00e0 deux $\\mu$ parall\u00e8les et reconstruits sur les fils et sur les strips. De tels muons sont originaires surtout des primaires charg\u00e9s ayant une \u00e9nergie de l'ordre de 200 TeV ou des $\\gamma$ de haute \u00e9nergie mais avec une faible section efficace. Cette zone d'\u00e9nergie est int\u00e9ressante puisque diff\u00e9rents types de sources galactiques peuvent produire des rayons gamma de telles \u00e9nergies. \\item \u00e9v\u00e8nements multiples (multiplicit\u00e9 $\\geq$3). C'est un groupe d'au moins 3 $\\mu$ reconstruits (sur les fils et les strips). Ils correspondent aux primaires ayant une \u00e9nergie $>$ 200 TeV. \\end{itemize} \\par Dans notre \u00e9tude nous proposons de comparer la distribution temporelle des muons \u00e0 celle obtenue d'un processus poissonnien.  \\label{chapnuc} Propos\u00e9e par Witten \\cite{nuc1}, la mati\u00e8re nucl\u00e9aire \u00e9trange \"Strange Quark Matter\" (SQM) pr\u00e9sente un \u00e9tat de la chromodynamique quantique. C'est un agr\u00e9gat de quarks u, d et s en nombre \u00e9gale, entour\u00e9 par un nuage d'\u00e9lectrons assurant la stabilit\u00e9 \u00e9lectrique et formant ainsi la dite \"nucl\u00e9arite\". \\par Un grand int\u00e9r\u00eat a \u00e9t\u00e9 apport\u00e9 \u00e0 l'\u00e9tude des nucl\u00e9arites, Bodmer \\cite{nuc2} a sugg\u00e9r\u00e9 que les nucl\u00e9arites peuvent \u00eatre plus stables que la mati\u00e8re nucl\u00e9aire ordinaire. Farhi et Jaffe \\cite{nuc3} ont montr\u00e9 que les nucl\u00e9arites peuvent \u00eatre stables sur un large domaine de masses, allant de celles des noyaux les plus l\u00e9gers (quelques quarks) jusqu'\u00e0 celles des \u00e9toiles \u00e0 neutrons ($\\sim 10^{57}$ quarks), \u00e0 conditions qu'ils se composent de quarks u, d et s en quantit\u00e9s approximativement \u00e9gales. En effet, la taille des nucl\u00e9arites peut d\u00e9passer largement celle des noyaux les plus massifs, car contrairement\t\u00e0 ces derniers, il n'y pas de limitation due \u00e0 la r\u00e9pulsion coulombienne.\\par En partant de ces conditions, la mati\u00e8re \u00e9trange doit \u00eatre extr\u00eamement stable. Elle peut former des syst\u00e8mes li\u00e9s avec n'importe quel nombre de quarks, et peut \u00eatre consid\u00e9r\u00e9e comme un candidat \u00e0 la mati\u00e8re obscure de l'univers \\cite{nuc4}. Il est suppos\u00e9 que les nucl\u00e9arites ont \u00e9t\u00e9 produits aux premiers instants de l'univers, ($T > 200 MeV$ et $t < 10^6 s$) \\cite{nuc5}. La production des nucl\u00e9arites, peut avoir lieu aussi dans les ph\u00e9nom\u00e8nes astrophysiques violents se produisant au sein de l'univers dont on peut citer les collisions entre \u00e9toiles \u00e0 neutrons et les explosions des \u00e9toiles \\cite{nuc6}.\\par La possibilit\u00e9 que les nucl\u00e9arites bombardent la terre a \u00e9t\u00e9 introduite par de R\u00f9jula et Glashow \\cite{nuc7}.\\par Dans ce qui suit, nous discuterons la perte d'\u00e9nergie de ces particules et leur possibilit\u00e9 pour arriver aux diff\u00e9rents niveaux de d\u00e9tection. ",
        "conclusions": "Nous avons pr\u00e9sent\u00e9 les r\u00e9sultats de l'\u00e9tude des distributions des muons cosmiques avec une \u00e9nergie plus grande que 1.3 TeV au sommet de la montagne de Gran Sasso. Les donn\u00e9es analys\u00e9es ont \u00e9t\u00e9 collect\u00e9es \u00e0 l'aide du syst\u00e8me des tubes \u00e0 streamer du d\u00e9tecteur MACRO dans sa configuration compl\u00e8te. \\\\ Les muons singuliers, doubles et multiples arrivant de toutes les directions ainsi que ceux venant des zones s\u00e9lectionn\u00e9es formant des c\u00f4nes ont \u00e9t\u00e9 consid\u00e9r\u00e9s \\cite{dtr17}.\\\\ Les r\u00e9sultats de notre \u00e9tude concernant les diff\u00e9rentes distributions temporelles montrent que les donn\u00e9es exp\u00e9rimentales sont compatibles avec la fonction Gamma d'ordre M. Ce qui nous a permis d'affirmer que nos r\u00e9sultats sont en accord avec l'arriv\u00e9e al\u00e9atoire des muons dans notre d\u00e9tecteur et permet ainsi d'exclure la pr\u00e9sence de toute composante, qui ne soit pas al\u00e9atoire dans les temps d'arriv\u00e9e des $\\mu$, similaire \u00e0 celle observ\u00e9e par Bath et al. \\cite{dtr7}, ou par Badino et al. \\cite{dtr8}, et ceci pour des \u00e9nergies de l'ordre de 20 TeV ($\\mu$ singuliers) et sup\u00e9rieures \u00e0 20TeV ($\\mu$ doubles et multiples). L'\u00e9tude des distributions temporelles des $\\mu$ arrivant des c\u00f4nes, a abouti aux m\u00eames conclusions.\\\\ Le test de Kolmogorov-Smirnov, nous a permis de mesurer la d\u00e9viation entre la distribution th\u00e9orique et celle exp\u00e9rimentale. Aucune d\u00e9viation n'a \u00e9t\u00e9 mise en \u00e9vidence dans notre \u00e9chantillon de donn\u00e9es, sauf dans le cas d\u00fb c\u00f4ne 3 et ceci est d\u00fb \u00e0 la faible statistique.\\\\ Dans cette \u00e9tude nous n'avons mis en \u00e9vidence l'existence d'aucun signal qui peut \u00eatre la signature d'une source discr\u00e8te tels que les pulsars.  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% CHAPITRE IV %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% \\chapter{Recherche des variations du flux de muons} \\label{chapss} %\\hskip 12pt \\subsection{Introduction} Durant presque dix ans de prise de donn\u00e9es \u00e0 une profondeur de 3800 m.w.e, MACRO a rassembl\u00e9 un des grands \u00e9chantillons de muons collect\u00e9s par les exp\u00e9riences souterraines de ce type. Le flux de muons peut pr\u00e9senter des variations d'origine galactique, solaire ou m\u00eame terrestre. Lorsque ces variations ne sont pas al\u00e9atoires, elles peuvent r\u00e9v\u00e9ler la pr\u00e9sence des modulations au niveau du signal d\u00e9voilant ainsi l'existence des objets galactiques \u00e9metteurs. Les clusters de muons souterrains peuvent \u00eatre produites par des \u00e9v\u00e9nements violents comme les \"Gamma Ray Burst\" sur l'\u00e9chelle de temps de quelques secondes, ou par des variations m\u00e9teorologiques soudaines sur l'\u00e9chelle de temps de plusieurs heures. Les modulations p\u00e9riodiques sont li\u00e9s \u00e0 la variation de la temp\u00e9rature de la haute atmosph\u00e8re, qui diminue durant la nuit (hiver) produisant les variations journali\u00e8res (saisonni\u00e8res) de la densit\u00e9 de l'atmosph\u00e8re et cependant du flux de muons.\\\\ Le temps d'arriv\u00e9e des muons cosmiques de haute \u00e9nergie a \u00e9t\u00e9 analys\u00e9 pr\u00e9c\u00e9demment, o\u00f9 nous avons pu confirmer l'absence de modulations dans le flux de muons. \\\\ Cette \u00e9tude va \u00eatre argument\u00e9e dans ce chapitre par la recherche de possibles variations de flux. Ainsi deux m\u00e9thodes sont utilis\u00e9es : \\begin{itemize} \\item la recherche d'\u00e9ventuel cluster d'\u00e9v\u00e8nements (groupement d'\u00e9v\u00e8nements) dans le flux de muons. \\item la recherche de variations p\u00e9riodiques. \\end{itemize}  Pour l'analyse des s\u00e9ries de temps des muons, d\u00e9tect\u00e9s par MACRO, deux approches compl\u00e9mentaires ont \u00e9t\u00e9 consid\u00e9r\u00e9es : la recherche de p\u00e9riodicit\u00e9 en utilisant la m\u00e9thode Lomb-Scargle et la recherche de cluster d'\u00e9v\u00e8nement en utilisant la m\u00e9thode scan statistics. Les deux m\u00e9thodes compl\u00e8tent l'analyse de la distribution des temps d'arriv\u00e9e des muons de haute \u00e9nergie men\u00e9e dans le chapitre \\ref{chapdt}. Pour les deux approches aucune d\u00e9viation de \"l'hypoth\u00e8se nulle\" n'a \u00e9t\u00e9 mise en \u00e9vidence. Ce r\u00e9sultat confirme les conclusions du chapitre \\ref{chapdt} et permet \u00e0 nouveau de conclure sur la nature al\u00e9atoire des temps d'arriv\u00e9e des muons.  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% CHAPITRE V %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% \\chapter{\u00c9tude de la perte d'\u00e9nergie des nucl\u00e9arites}  En se basant sur la formule de De R\u00fbjula (\\ref{derujula}) relative \u00e0 la perte d'\u00e9nergie des nucl\u00e9arites et en utilisant la mod\u00e9lisation de Shibata pour le calcul de la densit\u00e9 de l'atmosph\u00e8re, nous avons r\u00e9ussi \u00e0 \u00e9valuer la perte d'\u00e9nergie des nucl\u00e9arites le long de leur parcours dans l'atmosph\u00e8re. Pour les diff\u00e9rentes exp\u00e9riences r\u00e9alis\u00e9es \u00e0 diff\u00e9rentes altitudes dans l'atmosph\u00e8re (CAKE (40km), exp\u00e9riences port\u00e9es sur des avions civil (11km), SLIM (5.29km), exp\u00e9riences situ\u00e9es au niveau de la mer) nous avons d\u00e9termin\u00e9 la perte d'\u00e9nergie dans le plan des param\u00e8tres (masse,$\\beta$). En comparant le r\u00e9sultat obtenu dans la figure (\\ref{acc}) avec celui du \\cite{nuc17}, on remarque une diff\u00e9rence de masse d'un facteur de $\\approx 3$. Ceci est d\u00fb \u00e0 la fa\u00e7on avec laquelle a \u00e9t\u00e9 calcul\u00e9e la densit\u00e9 de l'atmosph\u00e8re.\\\\ Les nucl\u00e9arites ayant une masse $>7\\:10^{12} GeV/c^2$ ne sentent pas la pr\u00e9sence de l'atmosph\u00e8re. Au dessous de cette masse, et pour une exp\u00e9rience port\u00e9e \u00e0 une altitude de 5.29 Km (l'exp\u00e9rience SLIM ), le minimum de masse d\u00e9tectable d\u00e9cro\u00eet d'un facteur 2 par rapport aux exp\u00e9riences install\u00e9es au niveau de la mer. \\\\ Pour l'exp\u00e9rience souterraine MACRO (3800 m.w.e), les nucl\u00e9arites de masse sup\u00e9rieure \u00e0 $1.53\\;10^{16} GeV/c^2$ (figure \\ref{acc}) ne subissent aucune perte d'\u00e9nergie, leur vitesse $\\beta$ est intacte. Par contre le param\u00e8tre $\\beta$ d\u00e9cro\u00eet exponentiellement pour des masses dans l'intervalle $[1.6\\;10^{13},6\\;10^{15}]GeV/c^2$. Le CR39 du d\u00e9tecteur MACRO est sensible aux nucl\u00e9arites de masses $>2\\;10^{13}GeV/c^2$ pour lesquelles $3.37\\;10^{-4}<\\beta<0.9\\;10^{-3}$. Dans le cas du MAKROFOL, MACRO est sensible aux nucl\u00e9arites de masses $>3\\;10^{13}GeV/c^2$ pour lesquelles $0.12\\;10^{-3}<\\beta<0.9\\;10^{-3}$).\\\\ La contribution de l'atmosph\u00e8re dans le calcul de la perte d'\u00e9nergie devient consid\u00e9rable pour les nucl\u00e9arites avec des masses $< 7\\:10^{12}GeV/c^2$ et ayant une vitesse $\\beta = 10^{-3}$ en haut de l'atmosph\u00e8re.  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% CONCLUSION %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% \\addcontentsline{toc}{chapter}{Conclusion} \\chapter*{Conclusion} Le travail que nous avons men\u00e9 au pr\u00e8s de l'exp\u00e9rience MACRO s'est articul\u00e9 autour de la recherche des corr\u00e9lations possibles dans le temps d'arriv\u00e9e des muons cosmiques ainsi que sur la recherche des variations dans le flux de muons. \\par Nous disposons d'un lot de donn\u00e9es d'environ $38 \\;millions$ de muons d'\u00e9nergie sup\u00e9rieure \u00e0 1.3 TeV au sommet de la montagne du Gran Sasso, collect\u00e9 \u00e0 l'aide des tubes \u00e0 streamer du d\u00e9tecteur MACRO dans sa configuration compl\u00e8te.\\par L'analyse des distributions des temps d'arriv\u00e9es des muons singuliers, doubles et multiples arrivant de toutes les directions du ciel ainsi que ceux venant des zones s\u00e9lectionn\u00e9s formant des c\u00f4nes, montre qu'elles sont compatibles avec la fonction gamma. Ce qui confirme l'hypoth\u00e8se d'arriv\u00e9e al\u00e9atoire des RC. Ce r\u00e9sultat a \u00e9t\u00e9 appuy\u00e9 par le test de Kolmogorov-Smirnov, dont aucune d\u00e9viation entre la distribution th\u00e9orique et celle exp\u00e9rimentale n'a \u00e9t\u00e9 mise en \u00e9vidence.\\\\ Nos r\u00e9sultats ne sont pas ainsi analogue \u00e0 ceux obtenus par Bath et al. \\cite{dtr7}, ou par Badino et al. \\cite{dtr8}, et ceci pour des \u00e9nergies de l'ordre de 20 TeV ($\\mu$ singuliers) et sup\u00e9rieures \u00e0 20TeV ($\\mu$ doubles et multiples). \\par La recherche de cluster d'\u00e9v\u00e9nements, dans le flux de muons, par la m\u00e9thode scan statistics n'a pr\u00e9sent\u00e9 aucune d\u00e9viation significative de l'hypoth\u00e8se nulle. Ce qui nous a permis de confirmer l'absence d'\u00e9ventuel cluster d'\u00e9v\u00e9nements dans le flux de muons et donc aucune variation du flux n'a \u00e9t\u00e9 signal\u00e9. \\\\ En analysant les s\u00e9ries de temps des muons par la recherche des p\u00e9riodicit\u00e9s avec la m\u00e9thode Lomb-Scargle, nous avons observ\u00e9 la modulation saisonni\u00e8re et des signaux \u00e0 la position des ondes journali\u00e8re solaire et sid\u00e9rale. Aucune autres d\u00e9viation n'a \u00e9t\u00e9 signal\u00e9. En se basant sur la formule de De R\u00fbjula (\\ref{derujula}) relative au calcul de la perte d'\u00e9nergie des nucl\u00e9arites et en utilisant la mod\u00e9lisation de Shibata pour le calcul de la densit\u00e9 de l'atmosph\u00e8re, la contribution de la perte d'\u00e9nergie dans l'atmosph\u00e8re a \u00e9t\u00e9 introduite afin de d\u00e9terminer la vitesse minimale avec laquelle les nucl\u00e9arites peuvent atteindre diff\u00e9rents niveaux o\u00f9 sont install\u00e9s les d\u00e9tecteurs. \\\\ La contribution de l'atmosph\u00e8re, pour les nucl\u00e9arites ayant une vitesse $\\beta =10^{-3}$ en haut de l'atmosph\u00e8re, devient consid\u00e9rable pour les masses $<7\\:10^{12} GeV/c^2$. Au dessous de cette masse, et pour une exp\u00e9rience port\u00e9e \u00e0 une altitude de 5.29 Km (l'exp\u00e9rience SLIM ), la masse minimale d\u00e9tectable d\u00e9cro\u00eet d'un facteur 2 par rapport aux exp\u00e9riences install\u00e9es au niveau de la mer. Pour les exp\u00e9riences souterraines, les masses sup\u00e9rieures \u00e0 $1.53\\;10^{16} GeV/c^2$ ne subissent aucune perte d'\u00e9nergie.  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% BIBLIO %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% \\evenheading{\\hfill Bibliographie}%"
    },
    "0911/0911.4531_arXiv.txt": {
        "abstract": "{We present results of an ongoing systematic study of the large-scale properties of neutral hydrogen (\\HI) gas in nearby radio galaxies. Our main goal is to investigate the importance of gas-rich galaxy mergers and interactions among radio-loud AGN. From an \\HI\\ study of a complete sample of classical low-power radio galaxies we find that the host galaxies of extended Fanaroff $\\&$ Riley type I radio sources are generally \\HI\\ poor ($\\lesssim 10^8 M_{\\odot}$) and show no indications for gas-rich galaxy mergers or ongoing gas-rich interactions. In contrast, the host galaxies of a significant fraction of low-power compact radio sources contain enormous discs/rings of \\HI\\ gas (with sizes up to 190 kpc and masses up to $2 \\times 10^{10} M_{\\odot}$). This segregation in \\HI\\ mass with radio source size likely indicates that these compact radio sources are either confined by large amounts of gas in the central region, or that their fueling is inefficient and different from the fueling process of classical \\FRI\\ radio sources. To a first order, the overall \\HI\\ properties of our complete sample (detection rate, mass and morphology) appear similar to those of radio-quiet early-type galaxies. If confirmed by better statistics, this would imply that low-power radio-AGN activity may be a short phase that occurs at some point during the lifetime of many early-type galaxies. We discuss how upcoming \\HI\\ surveys (e.g. with ASKAP and Apertif) are essential for studying in a statistical way the the connection between the presence and morphology of a radio-loud AGN and the properties of the cold \\HI\\ gas associated with its host galaxy.}  \\FullConference{Panoramic Radio Astronomy: Wide-field 1-2 GHz research on galaxy evolution - PRA2009\\\\ June 02 - 05 2009\\\\ Groningen, the Netherlands}  \\begin{document} ",
        "introduction": "Active Galactic Nuclei (AGN) are believed to be triggered when gas and matter is deposited onto a super-massive black hole in the center of the host galaxy. For this to happen, the gas needs to lose sufficient angular momentum to be transported deep in the potential well of the galaxy, until it eventually fuels the AGN. Galaxy mergers and interactions have often been invoked as a possible triggering mechanism for AGN activity. During a merger/interaction, perturbations in the galactic potential may stir the gas in a galaxy and enhance cloud-cloud collisions . This may result in the fueling of the central black hole, the exact timing of which could be related to secondary processes, such as bar formation or the merging of individual black holes of the progenitor galaxies.  We are in the process of systematically investigating the merger/interaction history of well-defined samples of various types of nearby radio galaxies through observations of the large-scale neutral hydrogen (\\HI) gas. While in a galaxy merger or interaction part of the gas is driven into the central region, another part may be expelled in large structures \\citep[tidal-tails, bridges, shells, etc., e.g.][]{bar02}. If the environment is not too hostile, part of these gaseous structures (which often have a too low surface density for violent star formation to occur) remain bound to the host galaxy as relic signs of the galaxies' violent past. \\HI\\ observations, in particular when combined with optical information about the stellar content, are therefore an excellent tool to trace and date galaxy mergers or interactions over relatively long time-scales. \\HI\\ observations of nearby radio galaxies have the additional advantage that absorption studies can be done against the radio continuum, allowing in many cases to detect \\HI\\ to much lower column densities than what can be traced in emission. Besides mergers/interactions, other AGN feeding mechanisms (such as the accretion of hot gas from the circum-galactic medium) are often ascribed to certain types of radio galaxies \\citep[e.g.][]{all06,balma08}. Indications for such processes may also be reflected in the in the \\HI\\ content (or lack thereof) of nearby radio galaxies.  In this paper we present the \\HI\\ results of a complete sample of nearby low-power radio galaxies. This complete sample consists of 22 radio galaxies from the B2-catalogue ($F_{\\rm 408MHz} \\gtrsim 0.2$ Jy) up to a redshift of {\\sl z} = 0.041, plus radio galaxy NGC~3894 (which has properties similar to the B2 sources, but falls just outside the declination completeness limit of the B2 sample). Sources in dense clusters and BL Lac objects are excluded from our sample. Our complete sample contains both classical \\FRI\\ as well as compact radio sources, all with a radio power 22.0 $<$ Log ($P_{\\rm 1.4\\ GHz}$) $<$ 24.8 (and with no bias in $P_{\\rm 1.4\\ GHz}$ between the compact and extended sources). Observations were done with the Very Large Array in C-configuration and the Westerbork Synthesis Radio Telescope.  The next Section gives an overview of the results of this \\HI\\ study of low-power radio galaxies \\citep[to be published in an upcoming paper; ][]{emo10} as well as a preview of our ongoing work on large-scale \\HI\\ in more powerful radio galaxies. In Sect. \\ref{sec:askap} we discuss how upcoming \\HI\\ all sky-surveys with ASKAP and Apertif form a crucial next step in this research. ",
        "conclusions": ""
    },
    "0911/0911.1778_arXiv.txt": {
        "abstract": "{} {We report on a detection of an early rising phase of optical afterglow (OA) of a long  GRB\\,090726. We resolve a complicated profile of the optical light curve. We also investigate the relation of the  optical and X-ray emission of this event. } {We make  use of  the optical photometry  of this  OA  obtained  by the 0.5~m telescope of AI~AS~CR, supplemented by the data obtained by  other observers, and the X-ray {\\it Swift}/XRT data.   } {The  optical  emission  peaked  at  $\\sim17.5$~mag($R$)  at $t-T_{0} \\approx ~500$~s. We  find  a complex  profile  of  the  light  curve during the early phase of  this  OA: an  approximately power-law rise, a rapid transition to a plateau, a weak  flare superimposed  on the  center of  this  plateau,  and a slowly  steepening early decline followed  by a power-law  decay. We  discuss several  possibilities  to explain  the short  flare on  the  flat top of the optical light  curve at $t-T_{0}\\approx500$~s; activity of the central engine is favored although reverse shock cannot be ruled out. We show that power-law outflow with $\\Theta_{\\rm obs}/\\Theta_{\\rm c}>2.5$ is the best case for OA of GRB\\,090726. The initial Lorentz factor is $\\Gamma_{0}\\approx230-530$ in case of propagation  of the  blast wave in a homogeneous medium, while propagation of this wave in  a wind  environment  gives  $\\Gamma_{0} \\approx 80-300$. The value  of $\\Gamma_{0}$ in  GRB\\,090726  thus falls into the lower half of the range observed in GRBs and it may even lie  on the  lower  end. We  also show that both the optical and X-ray  emission decayed simultaneously and that the spectral  profile  from  X-ray to the optical band did not vary. This is true for both the time period before and after the break in the X-ray light curve. This  break  can be  regarded  as achromatic. The  available  data  show that neither  the  dust  nor  the  gaseous  component  of  the  circumburst medium underwent any evolution during the decay of this OA, that is after $t-T_{0} < 3000$~s. We  also show that this OA belongs to the least luminous ones in the phase of its power-law decay  corresponding to that observed for the ensemble of OAs of long GRBs.  } {} ",
        "introduction": "}   GRB\\,090726 was a long gamma-ray burst (GRB) localized by {\\it Swift}/BAT on the  26$^\\mathrm{th}$~July  2009  at 22:42:27.8\\,UT. It displayed a single peak  profile,  with  $T_{90}$  (15-350~keV)  $= 67 \\pm 15$~s. Fluence in the 15--150~keV  band  was  (8.6$\\pm$1.0)$\\times 10^{-7}$~erg~cm$^{-2}$~s$^{-1}$. The  1-s  peak  flux  was  0.7$\\pm$0.2 photon cm$^{-2}$ s$^{-1}$ (Page et al. 2009).  The GRB  coordinates  were  available  via  GCN  within  65\\,s, although because of  an  Earth-limb  constraint  {\\it Swift} could not follow-up until $\\sim$3.6\\,ks  after  the  trigger. Redshift  of  this  GRB is $z = 2.71$, as determined  from the lines  in optical  spectrum of the OA (Fatkhullin et al. 2009).  In this letter, we  report  the early  observations  of optical afterlow (OA) of this  GRB  performed  by  the  robotic  0.5\\,m telescope (D50) of the Astronomical Institute  of  AS~CR  in Ond\\v{r}ejov, Czech Republic. The first image of the  D50  started at  22:45:41.0\\,UT, i.e. 194~s  after the trigger. The telescope  was taking images until 23:35:35\\,UT (for 50 minutes), until a pointing limit was reached. On all of the 107 unfiltered, 20\\,s exposures the OA can  be  detected, although, to  improve  the  signal  to  noise ratio, we co-added some of the images. ",
        "conclusions": ""
    },
    "0911/0911.0474_arXiv.txt": {
        "abstract": "We develop a new approach to the well-studied anti-correlation between the optical-to-X-ray spectral index, $\\alpha_{ox}$, and the monochromatic optical luminosity, $l_{opt}$.  By cross-correlating the SDSS DR5 quasar catalog with the XMM-Newton archive, we create a sample of 327 quasars with X-ray S/N $>$ 6, where both optical and X-ray spectra are available. This allows $\\alpha_{ox}$ to be defined at arbitrary frequencies, rather than the standard 2500 $\\mbox{\\AA}$ and 2 keV.  We find that while the choice of optical wavelength does not strongly influence the $\\alpha_{ox}- l_{opt}$ relation, the slope of the relation does depend on the choice of X-ray energy.  The slope of the relation becomes steeper when $\\alpha_{ox}$ is defined at low ($\\sim$ 1 keV) X-ray energies.  This change is significant when compared to the slope predicted by a decrease in the baseline over which $\\alpha_{ox}$ is defined.  The slopes are also marginally flatter than predicted at high ($\\sim$ 10 keV) X-ray energies. Partial correlation tests show that while the primary driver of $\\alpha_{ox}$ is $l_{opt}$, the Eddington ratio correlates strongly with $\\alpha_{ox}$ when $l_{opt}$ is taken into account, so accretion rate may help explain these results. We combine the $\\alpha_{ox}-l_{opt}$ and $\\Gamma$-L$_{bol}$/L$_{Edd}$ relations to naturally explain two results: 1) the existence of the $\\Gamma-l_x$ relation as reported in \\citet{Young09} and 2) the lack of a $\\Gamma-l_{opt}$ relation.  The consistency of the optical/X-ray correlations establishes a more complete framework for understanding the relation between quasar emission mechanisms. We also discuss two correlations with the hard X-ray bolometric correction, which we show correlates with both $\\alpha_{ox}$ and Eddington ratio.  This confirms that an increase in accretion rate correlates with a decrease in the fraction of up-scattered disk photons. ",
        "introduction": "Optical/UV and X-ray continua in quasars are thought to originate in two physically distinct, but nevertheless related regions.  The optical/UV spectrum is dominated by a \u2018Big Blue Bump\u2019 continuum feature \\citep[e.g., ][]{Elvis94}, most likely from thermal emission arising from an accretion disk \\citep{Shields78, MS82, Ward87}.  A hot gas lying above the disk in some unknown geometry then up-scatters these disk photons to form a hard power-law ($<\\alpha_x> \\sim$ -1) \\citep[e.g.,][]{Mateos05, Mainieri07} in the X-ray spectrum \\citep{HM91, Zdziarski00, Kawaguchi01}.  While the shapes of the optical and X-ray continua are well known, the connecting continuum in the extreme ultraviolet (EUV) is not, because of absorption from our own galaxy.  The EUV and soft X-rays form the ionizing continuum (h$\\nu >$ 13.6 eV) needed to ionize a quasar's characteristic broad lines.  The ionizing continuum is also responsible for any line-driven outflows from the accretion disk \\citep{Proga07}.  Since direct observation of this continuum is not possible, it is often parameterized by an imaginary power-law with spectral index $\\alpha_{ox}$ running from 2500 $\\mbox{\\AA}$ to 2 keV in the rest frame.  The importance of the ionizing continuum motivates the decades long study of the anti-correlation between $\\alpha_{ox}$ and L$_{2500}$ \\citep{AT82, KC85, Tananbaum86, AM87, Wilkes94, PIF94, AWM95}; however, controversy continues.  While some studies have found that $\\alpha_{ox}$ depends primarily on redshift rather than optical luminosity \\citep{Bechtold03, Kelly07}, others have suggested that $\\alpha_{ox}$ may be entirely independent of redshift and optical luminosity \\citep{Yuan98, Tang07}, arguing that if the dispersion in L$_{2500}$ is much greater than the dispersion in L$_{2keV}$, an apparent but artificial correlation will result.  A narrow range in optical luminosity will enhance this effect.  In recent years, studies of the $\\alpha_{ox}-L_{2500}$ relation have become more comprehensive, with higher X-ray detection rates and largely complete samples that span a wide range of luminosities \\citep{Vignali03, Strateva05, Steffen06, Just07}.  For example, \\citet{Steffen06} work spans five and four decades in UV and X-ray luminosity, respectively, which helps minimize the effect of an artificial $\\alpha_{ox}-L_{2500}$ correlation due to selection effects.  In addition, \\citet{Strateva05} estimated the dispersion in L$_{2500}$ and L$_{2keV}$, finding that the difference in dispersions was not enough to artificially reproduce a correlation.  These studies used partial correlation analysis to conclude that the $\\alpha_{ox}-$z correlation is insignificant, while the $\\alpha_{ox}-L_{2500}$ anti-correlation is highly significant.  However, while the significance of the relation continues to be better understood, little is known of the underlying physics.  Current disk-corona models \\citep{Beloborodov99, Zdziarski99, Nayakshin00, Malzac01, Sobolewska04a, Sobolewska04b} do not directly predict the $\\alpha_{ox}-L_{2500}$ relation.  Previous studies of the $\\alpha_{ox}-L_{2500}$ relation have largely lacked spectral slopes in the X-rays, and so used the traditional endpoints of 2500~$\\mbox{\\AA}$ and 2 keV, where an X-ray spectral slope ($\\Gamma \\sim$ 2, where $\\Gamma$ = -$\\alpha$ + 1 for F$_{\\nu} \\propto \\nu^{\\alpha}$) is assumed in order to obtain the X-ray flux at 2 keV.  A systematic study with both optical and X-ray spectra enables an investigation of the relation at different frequencies than those traditionally used, revealing clues about the relation's physical underpinnings.  The sample selection and measurement of $\\alpha_{ox}$ are described in \\S2, and results of the data analysis are given in \\S3.  Section 4 discusses the results, specifically their connection to other optical/X-ray relations found for quasars.  We assume a standard cosmology throughout the paper, where H$_0$ = 70 km s$^{-1}$ Mpc$^{-1}$, $\\Omega_M = 0.3$, and $\\Omega_\\Lambda = 0.7$ \\citep{Spergel03}. ",
        "conclusions": "We find that the slope of the $\\alpha_{ox}-l_{opt}$ relation depends on the X-ray energy at which $\\alpha_{ox}$ is defined, though there is no significant dependence on optical wavelength.  While a change in slope is predicted when the baseline over which $\\alpha_{ox}$ is defined increases or decreases, the slopes at the lowest X-ray energy are significantly steeper than predicted, while the slopes at high X-ray energies are marginally flatter than predicted.  This suggests that the efficiency of low energy X-ray photon production depends more strongly on the optical/UV photon supply, whereas the efficiency of high energy X-ray photon production remains relatively constant with respect to the seed photon supply.  Partial correlation tests show that Eddington ratio is highly correlated with $\\alpha_{ox}$ when controlling for $l_{opt}$, and this partial correlation increases in significance with X-ray energy, suggesting that accretion rate likely plays a role in the results. Nevertheless, $l_{opt}$ remains the primary driver of $\\alpha_{ox}$.  To check the results, we redo the calculations for a sample including low S/N sources (CENSORED) and for a sample including only high S/N sources (HIGHSNR).  The CENSORED sample shows the same trends as in the SPECTRA sample, though the results are less significant.  This is likely due to the diluting effect of absorption in the low S/N sources, which emphasizes the importance of X-ray spectra in eliminating sources with absorption.  The HIGHSNR sample reproduces the trend in the slope of the $\\alpha_{ox}-l_{opt}$ relation found for the SPECTRA sample, indicating that the results are not affected by undetected absorption.  The dispersion of the relation does decrease significantly for the HIGHSNR sample, suggesting that undetected absorption does introduce some error into the determination of X-ray fluxes for sources with moderate S/N.  By combining the $\\alpha_{ox}-l_{opt}$ relation with the $\\Gamma$-L$_{bol}$/L$_{Edd}$ relation, we can naturally explain two results: 1) the existence of the $\\Gamma-l_x$ relation as reported in \\citet{Young09} and 2) the lack of a $\\Gamma-l_{opt}$ relation.  Given the measured dispersions and correlation coefficients, it is likely that the $\\alpha_{ox}-l_{opt}$ and $\\Gamma$- L$_{bol}$/L$_{Edd}$ relations are the primary relations, while the $\\Gamma-l_x$ relation is secondary.  Understanding the relations in this order will give a more complete framework for understanding the relation between optical and X-ray emission in quasars.  In order to fully understand the physical implications of the results shown here, a study of the selection, orientation and variability effects is needed.  Such a study is now in progress.  Finally, for the purposes of aiding studies of the hard X-ray bolometric correction, we include a description of the $\\kappa_{2-10keV}$ - $\\alpha_{ox}$ correlation and dispersion.  We find a significant partial correlation between $\\kappa_{2-10keV}$ and L$_{bol}$/L$_{Edd}$, when taking bolometric luminosity into account as the third variable.  This is consistent with a decreasing fraction of up-scattered X-ray photons as accretion rate increases."
    },
    "0911/0911.5584_arXiv.txt": {
        "abstract": "We study the impact of cosmic inhomogeneities on the interpretation of  SNe observations.  We build an inhomogeneous universe model that can confront supernova data and yet is reasonably well compatible with the Copernican Principle. Our model combines a relatively small local void, that gives apparent acceleration at low redshifts, with a meatball model that gives sizeable lensing (dimming) at high redshifts. Together these two elements, which focus on different effects of voids on the data, allow the model to mimic the concordance model. ",
        "introduction": "Until the nature of the dark energy is completely understood, the ``safe'' consequence of the success of the concordance model is that the flat isotropic and homogeneous $\\Lambda$CDM model is a good phenomenological fit to the real inhomogeneous universe and not that the concordance model is the actual model of the universe. It is therefore useful to look for alternative models that fit the data. Here we will discuss the possibility that the late time evidence for dark energy can be explained by the late time formation of nonlinear large-scale inhomogeneities, so that  the so-called coincidence problem becomes a hint of new physics.   In the following, we will assume that the spacetime of the inhomogeneous universe is accurately described by small perturbations around the Friedmann-Lema\\^{i}tre-Robertson-Walker (FLRW) solution whose energy content and spatial curvature are defined as Hubble-volume spatial averages over the inhomogeneous universe. Following Ref.~\\cite{Kolb:2009rp} we will call this Hubble-volume average the Global Background Solution (GBS). The cosmological model obtained through observations, on the other hand, will be called the Phenomenological Background Solution (PBS). The idea is to associate the concordance model with the PBS, while the actual GBS is the Einstein-de Sitter (EdS) model. ",
        "conclusions": ""
    },
    "0911/0911.2471_arXiv.txt": {
        "abstract": "We present the current photometric dataset for the Sloan Lens ACS (SLACS) Survey, including HST photometry from ACS, WFPC2, and NICMOS. These data have enabled the confirmation of an additional 15 grade `A' (certain) lens systems, bringing the number of SLACS grade `A' lenses to 85; including 13 grade `B' (likely) systems, SLACS has identified nearly 100 lenses and lens candidates. Approximately 80\\% of the grade `A' systems have elliptical morphologies while $\\sim $10\\% show spiral structure; the remaining lenses have lenticular morphologies. Spectroscopic redshifts for the lens and source are available for every system, making SLACS the largest homogeneous dataset of galaxy-scale lenses to date. We have created lens models using singular isothermal ellipsoid mass distributions for the 11 new systems that are dominated by a single mass component and where the multiple images are detected with sufficient signal-to-noise; these models give a high precision measurement of the mass within the Einstein radius of each lens. We have developed a novel Bayesian stellar population analysis code to determine robust stellar masses with accurate error estimates. We apply this code to deep, high-resolution HST imaging and determine stellar masses with typical statistical errors of 0.1 dex; we find that these stellar masses are unbiased compared to estimates obtained using SDSS photometry, provided that informative priors are used. The stellar masses range from $10^{10.5}$ to $10^{11.8}$ M$_\\odot$ and the typical stellar mass fraction within the Einstein radius is 0.4, assuming a Chabrier IMF. The ensemble properties of the SLACS lens galaxies, e.g. stellar masses and projected ellipticities, appear to be indistinguishable from other SDSS galaxies with similar stellar velocity dispersions. This further supports that SLACS lenses are representative of the overall population of massive early-type galaxies with M$_* \\gtrsim 10^{11} {\\rm M_\\odot}$, and are therefore an ideal dataset to investigate the kpc-scale distribution of luminous and dark matter in galaxies out to $z \\sim 0.5$. ",
        "introduction": "Early-type galaxies have been found to be remarkably homogeneous, as evidenced by the tight scaling relationships of the Fundamental Plane \\citep{fp,dressler} and the small scatter of their nearly-isothermal central density profiles \\citep[e.g.,][]{rk,rkk,paper3,koopmans09}. However, there are many details of this homogeneity that remain unresolved, including an explanation for the tilt of the Fundamental Plane with respect to the virial plane and for the scatter, or finite thickness, orthogonal to the plane. \\citet{paper7} have used lensing and stellar velocity dispersions to show that the dynamical mass scales with the lensing (i.e., total) mass of early-type galaxies, suggesting that non-homology is unlikely to cause the observed tilt. Joint dynamical and stellar population modeling has suggested that more massive early-type galaxies have higher dark matter fractions \\citep[e.g.,][]{padmanabhan} that could explain most of the tilt \\citep{tortora}. These models depend on assumptions about the initial mass function and do not readily explain how the luminous bulge and a dark matter halo might `conspire' to produce an isothermal total mass density profile; a well-defined set of galaxies with accurate photometric data is required in order to perform the detailed modeling necessary to accurately separate the dark and luminous mass components.  A particularly powerful tool is the combination of precise strong gravitational lensing mass measurements, stellar velocity dispersions, and accurate stellar mass estimates. The advent of the Sloan Lens ACS Survey \\citep[SLACS;][papers I-VIII, respectively]{paper1,paper2,paper3,paper4,paper5,paper6,paper7,paper8} has made this type of modeling feasible for the first time for a large and uniformly selected dataset. One of the key features of SLACS is that each lens contains all of the information necessary to perform a robust decomposition of the matter content in the central regions of galaxies. For example, the SDSS spectroscopy provides a precise measurement of the stellar velocity dispersion and the lens and source redshifts, while high-resolution Hubble Space Telescope (HST) imaging yields a detailed view of the lensed background source and the surface brightness profile of the lensing galaxy.  In this paper we present the definitive SLACS sample and an essential component necessary for disentangling the dark and luminous matter in early-type galaxies: multi-band optical and near-infrared imaging from HST. The photometric colors provided by HST allow us to accurately constrain the stellar mass of the lensing galaxy, thereby determining the normalization of the luminous component of the mass distribution. Although other authors have investigated the stellar populations of SLACS lenses using SDSS photometry \\citep[][hereafter G09]{grillo,grillo09}, we find that the HST imaging is necessary to adequately perform model photometry and yield unbiased stellar mass estimates, to the level of accuracy that is comparable with that of the other observables.  The paper is organized as follows.  We begin by describing our new HST observations in Section 2. We then provide a detailed summary of the SLACS dataset in Section 3, including the introduction of 15 newly confirmed strong lenses. In Section 4 we discuss our model photometry, and our stellar mass estimates are detailed in Section 5. A discussion of the ensemble properties of SLACS lenses, including comparison with `twin' galaxies from the SDSS, is given in Section 6 before concluding with a summary of the final SLACS dataset in Section 7. All magnitudes are on the AB scale \\citep{oke}, a standard concordance cosmology with $\\Omega_{m}=0.3$ $\\Omega_{\\Lambda}=0.7$ $h=0.7$ is assumed, and base-10 logarithms are denoted by `log' while natural logs are written as `ln.' ",
        "conclusions": ""
    },
    "0911/0911.0368_arXiv.txt": {
        "abstract": "We examine the relationship between galaxies and the intergalactic medium at $z < 1$ using a group of three closely spaced background QSOs with $z_{\\rm em} \\approx 1$ observed with the \\textsl{Hubble Space Telescope}.  Using a new grouping algorithm, we identify groups of galaxies and absorbers across the three QSO sightlines that may be physically linked. There is an excess number of such groups compared to the number we expect from a random distribution of absorbers at a confidence level of $99.9$\\%.  The same search is performed with mock spectra generated using a hydrodynamic simulation, and we find the vast majority of such groups arise in dense regions of the simulation. We find that at $z<0.5$, groups in the simulation generally trace the large-scale filamentary structure as seen in the projected 2-d distribution of the \\HI\\ column density in a $\\sim 30h^{-1}$~Mpc region.  We discover a probable sub-damped Lyman-$\\alpha$ system at $z=0.557$ showing strong, low-ionisation metal absorption lines. Previous analyses of absorption across the three sightlines attributed these metal lines to \\HI. We show that even when the new line identifications are taken into account, evidence remains for planar structures with scales of $\\sim 1$~Mpc absorbing across the three sightlines.  We identify a galaxy at $z=0.2272$ with associated metal absorption in two sightlines, each 200~kpc away. By constraining the star formation history of the galaxy, we show the gas causing this metal absorption may have been enriched and ejected by the galaxy during a burst of star formation 2 Gyr ago. ",
        "introduction": "\\label{sec:introduction}  Ever since QSO absorption lines were found to be produced by intergalactic gas unassociated with the background QSO, researchers have been speculating on their connection to galaxies \\citep[e.g.][]{bahcall_absorption_1969}. By measuring galaxy positions close to QSO sightlines, many groups have attempted to identify galaxies responsible for absorption lines, with varying degrees of success. It is particularly timely to study this relationship given the importance that gaseous inflows and outflows are now believed to have in regulating galaxy star formation rates. The relationship between QSO absorbers and galaxies can reveal important clues as to how the intergalactic medium (IGM) and galaxies interact.  It is well established that strong metal line absorbers seen in QSO sightlines are closely associated with galaxies. This connection was first seen between a \\CaII\\ doublet and galaxy at $z=0.53$ \\citep{boksenberg_existence_1978}, and subsequently confirmed for a number of galaxies using \\MgII\\ absorption \\citep{Bergeron91}. \\MgII\\ has a rest wavelength of $\\sim2800$~\\AA\\ and an easily identified doublet, thus it is straightforward to detect such absorbers associated with low-redshift galaxies using optical spectra. At $z<1$, strong \\MgII\\ systems (with rest equivalent width $>0.3$~\\AA) are generally accompanied by a galaxy within a few tens of physical kpc of the QSO sightline \\citep{Bergeron91, Steidel94b, Churchill00b, Kacprzak08}.  Studies of galaxies near these absorbers have shown that the covering factor of the halos causing the absorption must be less than unity, and there may be a different mechanisms associated with very high equivalent width \\MgII\\ systems (galactic winds?) and mid- to low-equivalent width \\MgII\\ systems (infall?).  With UV spectrographs available on the \\textsl{Hubble Space Telescope} (\\textsl{HST}), shorter rest wavelength metal lines and the \\HI\\ \\lya\\ transition can also be observed at low redshift. At $z<1$, strong \\CIV\\ doublet absorption is generally found within $50-100$ physical kpc of a galaxy \\citep{Chen01}, and its clustering strength with galaxies is similar to that of galaxy-galaxy clustering \\citep{Morris06}. \\citet{Adelberger05} used a large sample of Lyman break selected galaxies at $z \\approx 2.5$ around high-redshift QSOs to show that this relationship between \\CIV\\ and galaxies also holds at higher redshifts. Damped \\lya\\ absorbers (DLAs), predominantly neutral \\HI\\ absorbers with column densities \\NHI~$> 10^{20.3}$~absorbers~\\cmm, are also found to be closely associated with a wide range of galaxy types \\citep{Rao03, Chen05}.  These results have led to the picture of `halos' of absorbing gas around galaxies that cause the observed metal lines. A halo of gas around each galaxy gives rise to \\MgII\\ absorption up to a radius of a few tens of kpc, and \\CIV\\ absorption up to larger radii. Evidence for this picture has been assembled on a statistical basis from many single QSO sightlines close to galaxies, and there are no direct detections of a halo of metal-enriched gas around a galaxy using multiple nearby sightlines. The detailed geometry of such gas will depend on how it was placed there -- perhaps by winds induced by star-formation in the galaxy \\citep[e.g.][]{Theuns02}.  The relationship between galaxies and the more tenuous IGM as represented by the \\lya\\ forest (\\HI\\ absorption with \\NHI~$ < 10^{16}$ \\cmm) is even less clear. The first \\textsl{Hubble Space Telescope} observations of the low-redshift \\lya\\ forest showed many more lines than expected from an extrapolation of the high-redshift number counts \\citep{Morris91, Bahcall91}. This trend was subsequently confirmed by \\citet{Weymann98} using a larger number of \\textsl{HST} UV spectra.  A combination of the reduced ionizing background radiation at low redshifts and thus an increase in the fraction of neutral hydrogen \\citep{Theuns_loz_evol_98, Dave99}, and structure evolution \\citep{Scott02}, is believed to explain the unexpectedly high \\HI\\ line density.  This large line density opened the possibility of comparing the low redshift IGM and galaxy distributions.  The first results towards QSO sightline 3C~273 showed little correlation between galaxy and \\HI\\ except for the strongest \\NHI\\ absorbers \\citep{Morris93}, a conclusion that has been supported in more recent studies \\citep{Penton02, Chen05corr, Morris06, Wilman07}. An anti-correlation between absorber rest equivalent width (EW) and the impact parameter of nearby galaxies was reported in \\citet{Lanzetta95}, \\citet{Chen98} and \\citet{Chen01b} for lines with rest EW $> 0.3$~\\AA.  This was initially interpreted as evidence that a significant fraction of \\lya\\ absorbers could be attributed to gaseous halos associated with galaxies. However, other similar studies did not find such a strong relationship \\citep{Stocke95, LeBrun96, Tripp98}, and concluded that an anti-correlation was only significant for the \\HI\\ absorbers with large rest EWs ($>0.24$~\\AA).  The current accepted picture of the low redshift \\lya\\ forest has been heavily influenced by the results of N-body smoothed particle hydrodynamics (SPH) simulations \\citep[e.g.][]{Theuns_sph_sim_98, Dave99}. These predict dense, condensed gas very close to galaxies, shocked gas further out, and diffuse gas far from galaxies. The gas density increases with proximity to galaxies, explaining the correlation of equivalent width with impact parameter. There is no clean differentiation between gas that is associated with galaxies and the diffuse IGM.  Closely spaced QSO sightlines, either due to a fortunate asterism of background QSOs or a single lensed QSO, can be used to search for correlations in \\HI\\ and metal absorption across the sightlines \\citep[e.g.][]{Bechtold94a, Dinshaw95}.  Here `closely spaced' means separations on the expected size scale of \\lya\\ absorbers; up to a few 100 kpc based on SPH simulations and theoretical arguments \\citep{Schaye01}. Lensed QSO sightlines show that spectra of the \\lya\\ forest are very similar over 10s of kpc at $z=2.3$ \\citep{Smette95} and non-lensed close sightlines show correlations exist even over 100s of kpc. SPH simulations of the \\lya\\ forest are consistent with such correlation lengths, although they are interpreted as a coherence length in the IGM rather than the size of individual `clouds'. Indeed, \\citet{Rauch95} have shown there are not enough available baryons in the universe to explain the observed incidence of \\lya\\ absorbers if they are spherical structures with such large radii.  Here we analyse the distribution of galaxies in a field containing three closely spaced QSO sightlines - towards LBQS~0107-024A, LBQS~0107-024B and LBQS~0107-0232. This QSO group has been studied extensively, and showed some of the first evidence for very large scale correlations in \\lya\\ forest absorption \\citep{Dinshaw95}. Using multiple QSO sightlines opens the possibility of directly constraining the geometry of the absorbing gas, both around individual galaxies close to the sightlines and in large scale structures traced by galaxies.  At what distances and velocity scales do we expect to see associations between gas and galaxies? Absorbers could plausibly be associated with galaxies on the scale of a galactic halo, galaxy group, galaxy cluster, or large scale structure.  Velocity dispersions for early type galaxies are typically $\\lesssim 400$~\\kms\\ \\citep{Sheth03}, and their rotational velocities are $\\lesssim 150$~\\kms\\ \\citep{Franx91}. Rotational velocities in spirals are usually less than $\\sim 400$~\\kms\\ \\citep{Ramella89}. \\citet{Eke04} find groups of galaxies in the 2df galaxy redshift survey have median velocity dispersions of $\\sim260$~\\kms. Large galaxy clusters at $z<0.15$ show velocity dispersions of $\\sim800$~\\kms\\ \\citep{Fadda96}. While the intra-cluster medium is expected to be too hot for any \\HI\\ gas to exist, galaxies with associated cold gas and large relative velocity dispersions may be present in the cluster. \\citet{Colberg05} find dark matter filaments strung between galaxy cluster `knots' in $\\Lambda$CDM simulations that have a typical radius of 2~Mpc and length of $5-10$~Mpc.  Winds that expel gas from galaxies could introduce a further velocity offset between galaxies and any associated absorbers. At $z\\sim3$, the velocity of \\lya\\ emission seen in Lyman-break selected galaxies has a large offset ($\\sim600$~\\kms) with respect to the position of nebular lines, interpreted as being partially due to winds \\citep[e.g.][]{Adelberger05}. Low redshift starburst galaxies can show similarly large wind velocities \\citep{Heckman90}. In our analysis we generally consider three velocity cutoffs for association between absorbers or between absorbers and galaxies: 200, 500 and 1000~\\kms.  We intend to analyse galaxies around multiple QSO sightlines to address the questions: Are absorbers more likely to be found near galaxies?  Are groups of galaxies more or less likely to be associated with absorbers than single galaxies?  Do large scale structures seen in absorption across multiple sightlines coincide with large scale structures seen in the galaxy distribution? Can we put any constraints on the geometry of the absorbing gas associated with galaxies using the multiple sightlines?  Do we see evidence that galaxies are linked with metal-line systems in one or more sightlines?  In this paper we identify candidate structures at $z=0.501$ and $z=0.535$ comprised of both galaxies and \\HI\\ absorption spanning the three sightlines. Similar structures we identify in an SPH simulation are likely to arise inside filamentary large-scale structure in high density regions of the universe.  We also analyse a bright galaxy at $z=0.2272$ with associated metal absorption in two sightlines.  We use a cosmology with $\\Omega_{\\rm m}=0.3$, $\\Omega_{\\Lambda}=0.7$ and $H_0=70$~\\kmsmpc, where $\\Omega_{\\rm m}$ and $\\Omega_{\\Lambda}$ are the ratios of the matter density and cosmological constant energy density to the critical density, respectively. We give distances scaled by the parameter $h_{70}$, where $H_0 \\equiv 70h_{70}$~\\kmsmpc. All the distances given are physical (proper) distances unless stated otherwise. We convert between velocity, redshift and wavelength differences using: \\begin{equation} \\frac{\\Delta v}{c} = \\frac{\\Delta \\lambda}{\\lambda} = \\frac{\\Delta z}{1+z}\\, , \\end{equation} where $c$ is the speed of light, and $\\lambda$ and $z$ are the mean wavelength and redshift where the difference is being measured. These relations assume $v \\ll c$, which is true for the velocity differences of $\\lesssim1500$~\\kms\\ we consider.  We calculate the impact parameter of a galaxy from a QSO sightline by first calculating the comoving line of sight distance $D_{\\rm C}$ for a galaxy with redshift $z$ using \\begin{equation} D_{\\rm C} = \\frac{c}{H_0} \\int^z_0 {dz' \\over \\sqrt{\\Omega_m (1+z')^3 + \\Omega_\\Lambda}}\\, , \\end{equation} where $c$ is the speed of light. The physical (proper) impact parameter, $\\rho$, is then given by: \\begin{equation} \\rho = \\frac{\\Delta\\theta}{1+z} D_{\\rm C} \\, , \\end{equation} where $\\Delta\\theta$ is the angular separation in radians between the QSO and the galaxy.  To convert the observed total flux $F$ (\\ergscmm) of an emission line from a galaxy at redshift $z$ to a luminosity $L$~(\\ergs) we use the relationship: \\begin{equation} L = 4 \\pi D_{\\rm L}(z)^2 \\, F \\, , \\end{equation} where $D_{\\rm L}(z) = (1+z)D_{\\rm C}(z)$ is the luminosity distance.  All column density measurements have units of absorbers per cm$^2$ where they are not stated explicitly, and all logarithms are to the base 10.  The paper is structured as follows: Section 2 describes the galaxy and QSO absorption line samples. Section 3 presents our analysis of \\HI\\ absorption associations across the three QSO sightlines. Section 4 presents our analysis of galaxy-absorber associations across the three sightlines. Section 5 describes the simulated galaxies and mock sightlines with which we compare our observations. Sections 6 and 7 discuss and summarise our results. ",
        "conclusions": "The simulations in Figure~\\ref{fig:sim_nhi} suggest two things: (1) the filaments visible over this 2-d projection maintain their geometry in comoving space through redshift. Since the comoving separation between the sightlines changes with redshift, the triple samples large scale structure on different characteristic scales at different redshifts. (2) The column density that traces the filamentary structures falls with decreasing redshift.  At $z=1$, filamentary structures are traced by gas with log\\NHI\\ of 14.3, where as at $z=0$ it is traced by gas with log\\NHI\\ of about 13.5, and higher column densities now traced more collapsed, spherical structures. This implies that triple absorbers identified above a single minimum detectable equivalent width (which can be converted to a column density) will trace different kinds of structures at different redshifts -- filaments and even `void' regions at higher redshifts, and groups of halos at low redshift. Thus triples will tend to correspond to different kinds environments at different redshifts, depending on the combination of comoving separation and the column density sensitivity.  The excess of groups with absorbers in one or more sightlines compared to a distribution of random absorbers tells us the same thing as the nearest-neighbour test -- that \\lya\\ absorbers are correlated with galaxies on these scales. The significance of this excess appears to decrease as we accept absorbers with larger velocity separations from a galaxy into a group. This suggests that single absorbers are only correlated with any single galaxy on scales of $\\lesssim 500$~\\kms. This this is consistent with the results of \\citet{Morris06}.  As we move to groups with absorbers in two or three sightlines, the number of groups becomes smaller. Thus the statistics are more uncertain and our interpretation is more tentative. However, we see a clear excess of galaxies with nearby absorbers in two and three sightlines at all velocity differences compared to a random absorber distribution. Interestingly, the excess of triple LOS-galaxy groups is very significant ($>99$\\%) for velocity separations of $750$ and $1000$~\\kms. This suggests that the observed triple LOS-galaxy groups trace Mpc-scale structures that appear in both \\HI\\ gas and galaxies. We emphasize that while a $1000$~\\kms\\ separation within a single real structure appears unrealistically large if we interpret it as entirely due to the Hubble flow (a Mpc corresponds to $\\sim 70$~\\kms\\ at $z=0.5$), if some of the absorbers arise in winds of the magnitude described in the introduction, then such large velocity differences can still be present in a single structure less than a couple of Mpc across.  The galaxies and absorbers that make up the seven triple LOS-galaxy groups identified in the real absorbers are shown in the Figure~\\ref{fig:galspec} $\\langle z\\rangle = 0.501$ and $\\langle z\\rangle = 0.535$ panels. For both structures, there is no obvious link between any one of the galaxies and its associated absorber, in the sense that none of the absorbers are close enough to the galaxies to be caused by a `halo' of gas associated with that galaxy. A more appealing explanation is that they are caused by a group of galaxies at a knot or filament of the cosmic web and its associated absorption. This absorption may be due to large diffuse IGM gas that spans the three sightlines, or more compact gas clouds closely associated with individual galaxies in the structure that have not been targeted in our spectroscopy.  In future papers we intend to present GMOS, VIMOS and DEIMOS data providing a much larger sample of galaxy redshifts in this field with a better-known selection function, along with \\textsl{HST} Cosmic Origins Spectrograph (COS) spectra of the three QSOs. The COS spectra will resolve \\HI\\ \\lya\\ forest lines, detect \\OVI\\ and other metal line absorbers, and allow us to measure the temperature, column density and ionisation state of the IGM."
    },
    "0911/0911.0829_arXiv.txt": {
        "abstract": "{ Shell-type supernova remnants (SNRs) exhibit correlations between radio surface brightness, SNR diameter, and ambient medium density, that between the first two quantities being the well known $\\Sigma$--$D$ relation. } { We investigate these correlations, to extract useful information about the typical evolutionary stage of radio SNRs, as well as to obtain insight into the origin of the relativistic electrons and magnetic fields responsible for the synchrotron emission observed in radio. } { We propose a scenario, according to which the observed correlations are the combined effect of SNRs evolving in a wide range of ambient conditions, rather than the evolutionary track of a ``typical'' SNR. We then develop a parametric approach to interpret the statistical data, and apply it to the data sample previously published by Berkhuijsen, as well as to a sample of SNRs in the galaxy M~33. } { We find that SNRs cease to emit effectively in radio at a stage near the end of their Sedov evolution, and that models of synchrotron emission with constant efficiencies in particle acceleration and magnetic field amplification do not provide a close match to the data. We discuss the problem of the cumulative distribution in size, showing that the slope of this distribution does not relate to the expansion law of SNRs, as usually assumed, but only to the ambient density distribution. This solves a long-standing paradox: the almost linear cumulative distribution of SNRs led several authors to conclude that these SNRs are still in free expansion, which also implies very low ambient densities. Within this framework, we discuss the case of the starburst galaxy M82. } { Statistical properties of SNR samples may be used to shed light on both the physics of electron acceleration and the evolution of SNRs. More precise results could be obtained by combining data of several surveys of SNRs in nearby galaxies. } ",
        "introduction": "\\label{sect:Introd}  \\indent Radio emission is a quite common property of shell-type supernova remnants (SNRs). The intensity of the (synchrotron) radio emission is related to the magnetic field strength and the amount of accelerated electrons. However, the mechanisms leading to both the magnetic field amplification and the electron injection at the SNR shock and their respective efficiencies remain poorly constrained. To investigate these processes observationally, there have been detailed studies, mostly in X-rays, of some selected SNRs (see e.g., Cassam-Chena\\\"\\i\\ et al.\\ \\cite{cea07}, Cassam-Chena\\\"\\i\\ et al.\\ \\cite{cea08}, and references therein). Relevant, complementary information should also be extracted from a statistical analysis of SNR data samples.  On the observational side, a well-known (even though not widely accepted) statistical relation is the so-called ``\\SgD\\ relation'', namely the empirical correlation discovered between the SNR size ($D$) and its radio surface brightness ($\\Sg$). Various authors (e.g., Clark \\& Caswell \\cite{cc76}, Milne \\cite{m79}, Caswell \\& Lerche \\cite{cl79}, Case \\& Bhattacharya \\cite{cb98}, Uro\\v sevi\\'c et al.\\ \\cite{uea05}) have investigated this correlation. Originally applied as a tool for estimating SNR distances, it has been found to be unreliable for this purpose. Some authors (e.g., Green \\cite{g05}, Uro\\v sevi\\'c et al.\\ \\cite{uea05}) have also argued that this correlation may be affected by selection effects. For Galactic SNRs, Green (\\cite{g05}) highlighted the selection effects: (1) in the surface brightness, with a completeness limit of about $10^{-20}\\U{\\surfbrig}$, while SNRs below this limit are predominantly in regions where the Galactic background is low; (2) in the angular size, so that young but distant SNRs may be missed. For SNRs in other galaxies, Uro\\v sevi\\'c et al.\\ (\\cite{uea05}) also showed that there may be a selection effect in the integrated flux (valid for unresolved or mildly resolved SNRs), with the effect of leading to an observed slope shallower than the intrinsic one.  Even though selection effects may affect this correlation to some level, we believe that the relation itself has a physical origin, and that one can therefore extract from it information about the processes involved in the SNR radio emission: therefore, understanding its origin could eventually contribute to constraining the efficiency of these processes. There have been several attempts (e.g., Shklovsky \\cite{s60}, Van der Laan \\cite{vdl62}, Poveda \\& Woltjer \\cite{pw68}, Kesteven \\cite{k68}, or more recently Duric \\& Seaquist \\cite{ds86}, Berezhko \\& V\\\"olk \\cite{bv04}) to explain the \\SgD\\ relation as the average evolutionary track of a ``typical'' SNR. In all of these cases, the slope of the correlation is assumed to correspond to that of the SNR evolutionary track on the \\SgD\\ parameter plane.  However, there is clear evidence that this assumption is incorrect. Berkhuijsen (\\cite{b86}) has found tight correlations of both $\\Sg$ and $D$ with the ambient density ($\\nz$), in the sense that smaller (and brighter) SNRs are typically located in a denser medium (indeed, the correlation between $\\nz$ and $D$ had already been known for several years; see e.g., Fig.~4 in McKee \\& Ostriker \\cite{mo77}). The best-fit results given by Berkhuijsen are \\begin{eqnarray} \\label{eq:DBerk} D&\\simeq&15\\,\\nz^{-0.39\\pm0.04}\\U{pc}\t\t\t\t\t\\\\ \\label{eq:SgBerk} \\Sg&\\simeq&6\\E{-20}\\nz^{1.37\\pm0.21}\\U{\\surfbrig} \\end{eqnarray} Berkhuijsen concluded that the \\SgD\\ relation ($\\Sg\\propto D^\\xi$) is just a secondary effect, while the two primary relations are those of $D$ and $\\Sigma$ with $\\nz$.  Although this conclusion may appear rather extreme, it is quite obvious that the \\SgD\\ relation contains SNRs that evolve in very different ambient conditions. We then assume that the correlations between $\\Sg$, $D$, and $\\nz$ do not directly reflect the evolution of a ``typical'' object, but are rather the combined effect of the evolution of different SNRs expanding in different ambient conditions. That the correlations with $\\nz$ are statistically so tight also suggests that $D$ and $\\Sg$ are more sensitive to the ambient conditions than to the SNR evolution. The combined effect of these correlations is that a clear trend of $\\nz$ across the \\SgD\\ relation is observed (see Fig.~\\ref{Fig1}), namely that smaller SNRs are preferentially located in higher-density environments.  \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics{12244f01.eps}} \\caption{Distribution of SNRs from Berkhuijsen (\\cite{b86}) sample in the $\\lg D$--$\\lg\\Sg$ parameter plane ($D$ is measured in pc, while $\\Sg$ is in $\\U{\\surfbrig}$). The dot sizes are proportional to $\\lg\\nz$ values, the legend showing in order sizes corresponding to $\\lg\\nz$ from --1.5 to 1.0, in steps of 0.25 ($\\nz$ being measured in $\\U{cm^{-3}}$). It is apparent from this figure that smaller (and brighter) SNRs are typically located in a denser medium.} \\label{Fig1} \\end{figure}  In principle, correlations of $\\Sg$ and $D$ with $\\nz$, accounting for at least qualitatively the trend of $\\nz$ across \\SgD\\ relation, could also appear in the case of the evolution of an individual object, provided that it expands in a medium with a highly structured, fractal, density distribution. This would reflect that, during its life, a SNR always preferentially expands towards the direction in which the ambient density is lower. Therefore, at any time, the ``effective'' ambient density would be close to the ``lowest'' ambient density in the volume occupied by the SNR. In this way, one could explain why large SNRs typically seem to expand in a low-density medium.  However, it is hard to justify, in this scenario, the absence of small SNRs in low-density media (which would be the case, when the supernova itself is located in a low density region), as well as that low ambient density values are measured for all extended SNRs (which requires that such low density regions are ubiquitous in the Galaxy, on scales of tens of parsecs or even less). In addition, if the effects of the fractal interstellar medium were dominant, virtually all SNRs should have a much brighter limb on one side, which is not observed. Therefore, for all these reasons, we conclude that the ``fractal ambient density'' hypothesis is implausible, and we do not consider it any longer in this paper.  One of our goals is to show that the best-fit line usually referred to as the ``\\SgD\\ relation'' provides only a minor part of the information present in the data, while additional information could be extracted by analyzing in detail the distribution of points in the $\\Sg$--$D$--$\\nz$ parameter space. For this reason we propose a rather general (parametric) scenario, with the aim of constraining the physics of the electron injection and magnetic field behavior in SNR shocks and/or the SNR evolutionary phase in which they are most likely to be observed in radio.  The plan of the paper is the following. In Sect.~2, we present the basic ideas, assumptions and formulae that will be used in the rest of the paper. Section 3 is devoted to a re-analysis of the data of Berkhuijsen (\\cite{b86}), using the criteria introduced in Sect.\\ 2, and extracting constraints on the dependence of the particle and magnetic field efficiencies on the basic shock parameters. For the sake of comparison, we also apply the same technique to a sample of SNRs in galaxy M~33. Section 4 shows how the cumulative distribution of SNRs with diameter may be interpreted within our framework, and how we can explain the puzzling linear trend of this distribution without requiring, as usually done, that SNRs expand linearly to large sizes. Section 5 presents our conclusions. ",
        "conclusions": "\\label{sect:Conclus}  We have shown that studies of the statistical properties of SNR samples may provide insight into the physics of electron acceleration and the time evolution of SNRs. We have proposed a new scenario, along the lines of previous work by Berkhuijsen (\\cite{b86}), which interprets in a natural way the observed correlations between radio surface brightness $\\Sg$, size $D$, and ambient density $\\nz$ in a sample of SNRs. The main parameter of SNR evolution that enters into these correlations is the time at which a SNR ceases to behave as a radio source, and we find that this endpoint is located close to the end of the Sedov phase. We otherwise find that the observed correlations mostly reflect that the sample consists of SNRs located in very different ambient conditions; while the evolution of individual SNRs plays a secondary role, and cannot be extracted by simply studying correlations between pair of quantities.  Within this framework, we present a new approach to analyzing the statistical data, based on a 2-dimensional fit to $\\Sg$ as a function of $D$ and $\\nz$. We show that the slope of $\\Sg(D)$ at constant $\\nz$ should represent more closely the true evolutionary track of an individual SNR than the well known ``\\SgD\\ relation'', which is obtained without including information about $\\nz$.  In this work, we have used data published by Berkhuijsen (\\cite{b86}). Although this data sample is rather limited, our method of analysis applied to these data is already capable of discriminating to some level between different theoretical models. For instance, models prescribing constant efficiencies for both magnetic field (turbulent) amplification and electron acceleration (e.g., Berezhko \\& V\\\"olk \\cite{bv04}) are well outside the parameters region allowed by the data. On the other hand, models assuming a constant acceleration efficiency but a constant post-shock magnetic field are marginally (about 2-$\\sg$) consistent with the data. For the sake of comparison, we have applied the same technique to a sample of SNRs in M~33. Although this sample could be affected by selection effects, it is completely independent from the sample of Berkhuijsen, and the parameters that we derive from the two samples are in close agreement, within the statistical errors. This may indicate again that our technique is robust andthe assumptions at the basis of our statistical analysis are correct. With the use of larger and more accurate data samples the confidence interval will become narrower, and the statistical results will be capable of constraining models more tightly.  In nearby galaxies, deep and detailed surveys of SNRs have become available in both radio and X-rays. The number of SNRs in other galaxies has increased considerably, with reasonably good radio and X-ray flux measurements and in various cases also information about their angular size. Although SNRs in external galaxies are more difficult objects to observe than Galactic SNRs, their distances are well known (since they correspond to the distance of the parent galaxy); in addition, selection effects can be modeled in a more accurate way for a sample of SNRs that are all at the same distance. Therefore, a natural extension of the present work is to elaborate our diagnostic tools along the guidelines shown in this paper, and to use them to investigate SNRs samples in various nearby galaxies. These results will be presented in a forthcoming paper."
    },
    "0911/0911.0532_arXiv.txt": {
        "abstract": "We present spectroscopic observations of a sample of 15 embedded young stellar objects (YSOs) in the Large Magellanic Cloud (LMC). These observations were obtained with the {\\it Spitzer Infrared Spectrograph} (IRS) as part of the \\spec\\ Legacy program. We analyze the two prominent ice bands in the IRS spectral range: the bending mode of CO$_2$ ice at 15.2\\,$\\mu$m and the ice band between 5 and 7\\,$\\mu$m that includes contributions from the bending mode of water ice at 6\\,$\\mu$m amongst other ice species. The 5$-$7\\,$\\mu$m band is difficult to identify in our LMC sample due to the conspicuous presence of PAH emission superimposed onto the ice spectra. We identify water ice in the spectra of two sources; the spectrum of one of those sources also exhibits the 6.8-$\\mu$m ice feature attributed in the literature to ammonium and methanol. We model the CO$_2$ band in detail, using the combination of laboratory ice profiles available in the literature. We find that a significant fraction ($\\ga$\\,50\\%) of CO$_2$ ice is locked in a water-rich component, consistent with what is observed for Galactic sources. The majority of the sources in the LMC also require a pure-CO$_2$ contribution to the ice profile, evidence of thermal processing. There is a suggestion that CO$_2$ production might be enhanced in the LMC, but the size of the available sample precludes firmer conclusions. We place our results in the context of the star formation environment in the LMC. ",
        "introduction": "The formation of stars in the high-redshift Universe occurred in a metal-poor environment. Much of what we know about the process of star formation is deduced from observations in nearby, Galactic star-forming regions with essentially solar metallicity. However, many of the physical processes that affect star formation and their observational diagnostics are expected to depend on metallicity, for instance cooling and gas-phase and grain surface chemistry. The nearest templates for the detailed study of star formation under metal-poor conditions are the Magellanic Clouds (MCs). The Interstellar Media (ISM) of the Large and Small Magellanic Cloud (LMC and SMC) have significantly lower metallicities ($Z_{\\rm LMC}\\sim 0.4 $\\,Z$_\\odot$ and $Z_{\\rm SMC}\\sim 0.2 $\\,Z$_\\odot$), than the solar neighborhood ISM ($Z_{\\rm Gal}\\sim $ Z$_\\odot$). Only relatively recently are we able to study the details of the star formation process in the MCs \\citep[for a review of recent results see][] {oliveira09}.  The star-formation process is a complex interplay of various chemo-physical processes. During the onset of gravitational collapse of a molecular cloud, sufficiently dense cores can only develop if the heat produced during the contraction can be dissipated. The most efficient cooling mechanisms are via radiation through fine-structure lines of carbon and oxygen (e.g., via the emission lines of [C\\,{\\sc ii}] at 158\\,$\\mu$m and [O\\,{\\sc i}] at 63 and 146\\,$\\mu$m), and rotational transitions in abundant molecules such as CO, O$_{2}$ and water \\citep{goldsmith78}. These cooling agents all contain at least one heavy atom --- H$_{\\rm 2}$ has no permanent dipole moment, and emission through its rotational transitions is very unlikely, making it an inefficient coolant. Besides also contributing to the thermal balance in the molecular cloud, dust grains are crucial in driving cloud chemistry, as dust opacity shields molecules from radiation and grain surfaces enable chemical reactions to occur that would not happen in the gas phase. Surface chemistry is also thought to play an important role in the formation of H$_{2}$ \\citep[e.g.,][]{williams03} and H$_{2}$O \\citep[e.g.,][]{hollenbach09}. Simple theoretical arguments suggest that, even over the range covered by the Milky Way and the MCs, cooling and chemistry rates and timescales are affected by metallicity \\citep{banerji09}. Not only is it crucial to understand the ISM chemistry at low metallicity, it is also important to investigate the chemical evolution during the star formation process.  In cold molecular clouds, layers of ice form on the surface of dust grains. In particular, when densities increase dramatically during the cold collapse phase, most molecules accrete onto dust grain surfaces. In fact, as water, CO, and also CO$_{2}$ are completely frozen out of the gas phase in these environments, ices are substantial reservoirs of the available oxygen and carbon \\citep[see review by][]{bergin07}. Even if CO seems to form exclusively in the gas-phase, there is laboratory evidence that water ice is predominantly formed by hydrogenation on icy grains \\citep{ioppolo08}. Therefore, understanding ice chemistry is crucial to understanding the gas-phase chemistry and the physical conditions within molecular clouds. Water is by far the most abundant ice (typically $10^{-5}-10^{-4}$ with respect to H$_{2}$), followed by CO$_{2}$ and CO, with a combined abundance of 10$-$30\\% with respect to water ice \\citep{gibb04}. Other ice species have also been identified (methanol, ammonia, ammonium, methane and formaldehyde), albeit in much smaller concentrations \\citep{vandishoeck04}. While water and CO can be observed from the ground, CO$_2$ can only be observed from space and its ubiquitous presence in a variety of environments was one of the major discoveries of the {\\it Infrared Space Observatory ({\\it ISO})} \\citep[see][ and references therein]{ehrenfreund03}.  It is still unknown whether metallicity affects ice chemistry. The most obvious effect of a lower metal content is a reduced availability of carbon and oxygen, the cornerstones of interstellar chemistry. Indeed the CO abundance in the diffuse ISM in the LMC is lower than that measured in Galactic environments \\citep{fukui01}. In metal-poor clouds, it could be expected that ice abundances are generally depressed (i.e. all ices could be less abundant as a result of lower abundances of C and O), but it is possible that the abundance balance of ice species is also reshaped. The lower dust:gas ratio should make cloud cores more transparent to ambient UV light, possibly exposing ices to more radiative processing, and making complex species more abundant than in the Galaxy.  Interstellar ices have been detected in the envelopes of heavily embedded young stellar objects (YSOs) in the MCs. The first evidence of ices in the envelope of a massive YSO in any extra-galactic environment was a serendipitous discovery by \\citet{vanloon05a}. The spectrum of IRAS\\,05328$-$6827 in the LMC, obtained with the Infrared Spectrometer \\citep[IRS,][]{houck04} onboard the {\\it Spitzer Space Telescope} \\citep[\\spitzer,][]{werner04} and the Infrared Spectrometer And Array Camera at the Very Large Telescope (ISAAC/VLT), shows clear absorption signatures of water ice at 3.1\\,$\\mu$m and CO$_{2}$ ice at 15.2\\,$\\mu$m. \\citet{oliveira06} and \\citet{shimonishi08} identified more embedded YSOs with ice signatures in the LMC. Three embedded YSOs have been recently identified in the SMC \\citep{vanloon08}. There is a suggestion that the stronger UV radiation field \\citep{welty06} and/or higher dust temperature \\citep{aguirre03} could be responsible for a higher CO$_{2}$-to-water column density ratio observed towards a few YSOs in the LMC when compared to Galactic YSOs \\citep{vanloon05a,shimonishi08}. These studies are however hampered by small sample size and limited spectral resolution and thus not yet conclusive.  CO$_2$ ice is found to be abundant and ubiquitous in a variety of molecular environments in the Galaxy, from quiescent clouds \\citep{whittet07} to the envelopes of high- and low-luminosity YSOs \\citep{gerakines99,pontoppidan08}. This molecule is generally believed to form exclusively via surface reactions, even though its formation mechanism is not fully understood \\citep[for a recent discussion of this issue see][]{pontoppidan08}. The most prominent CO$_2$ solid state features are the 4.3-$\\mu$m stretching and 15.2-$\\mu$m bending modes \\citep[e.g.,][]{ehrenfreund97}. Comparisons of observed 15.2-$\\mu$m profiles with laboratory spectra suggest that CO$_2$ ices are dominated by components that are either water- or CO-rich, rather than pure CO$_2$ \\citep{gerakines99,pontoppidan08,whittet09}. For some YSOs, narrow substructure in the observed profiles provides evidence of thermal processing \\citep[e.g.,][]{gerakines99,white09}. Therefore the CO$_2$ ice profile is a sensitive probe of its molecular environment.  Another ice band in the IRS range (5$-$7\\,$\\mu$m) is less well understood, with contributions not only from water ice at 6\\,$\\mu$m but also ammonia (NH$_3$), methanol (CH$_3$OH), formic acid (HCOOH) and formaldehyde (CH$_2$O) ices \\citep{boogert08}. This superposition of multiple ice components is possibly responsible for the discrepancy between the column densities determined from the 3-$\\mu$m (O$-$H stretching mode) and the 6-$\\mu$m (O$-$H bending mode) water ice features, with the 6-$\\mu$m column density systematically higher \\citep{gibb04,boogert08}. There is an ice feature at 6.8\\,$\\mu$m attributed to ammonium (NH$_4^+$) \\citep{boogert08}. This band also falls in a wavelength range strongly affected by emission features generally attributed to polycyclic aromatic hydrocarbon (PAH) molecules. The libration mode of water ice is a very broad feature that peaks at 13\\,$\\mu$m and extends between 10$-$30\\,$\\mu$m; in practice this feature is very difficult to isolate both because of its width and the fact that it is hard to disentangle from the more prominent silicate dust features.  In a series of papers we will investigate ice chemistry in the envelopes of massive embedded YSOs in the low metallicity environment of the MCs. The \\spitzer\\ Legacy survey of the LMC \\citep[SAGE,][]{meixner06} identified large numbers of YSO candidates for the first time in any extra-galactic environment \\citep{whitney08}. The follow-up Legacy program with the \\spitzer-IRS includes a variety of point and extended sources among which YSOs feature prominently \\citep[\\spec,][]{kemper09}. We selected the LMC sample described here from both the \\spec\\ program and the \\spitzer\\ archive, as those objects that have spectral energy distributions (SEDs) consistent with early embedded YSOs and exhibiting in their spectrum the absorption feature associated with CO$_{2}$ ice, the most prominent ice band in the \\spitzer-IRS spectral range.  In this paper we present the first detailed analysis of CO$_2$ ice profiles for massive YSOs in the LMC. The paper is organized as follows. First we introduce the sample and discuss the issues of continuum determination and constraining the PAH spectrum. We then describe the observed CO$_2$ profiles in terms of laboratory components and measure column densities. We also investigate features in the 5$-$7\\,$\\mu$m range. Finally we interpret our results in terms of possible effects of low metallicity. ",
        "conclusions": "In this section we compare in more detail the properties of ices in the LMC YSOs to Galactic samples and interpret our results in terms of the star formation environment in the LMC.  \\subsection{Observed properties of the CO$_2$ ice profiles}  As we have seen in Section\\,4.4, a large fraction of CO$_2$ ice is locked in a water-rich ice matrix, with a possible contribution from CO$_2$ ice in a CO-rich layer. Figure\\,\\ref{water_rich} compares the fraction of CO$_2$ ice in the water-rich component for the LMC sample and Galactic samples from the literature: high-luminosity YSOs and background sources\\footnotemark \\citep{gerakines99,whittet09} and low-luminosity YSOs \\citep{pontoppidan08}. \\footnotetext{Background objects are bright sources that coincidentally sit behind the molecular cloud but are not associated with it in any way. They represent quiescent sightlines.} The column densities for this component are derived using the polar\\,+\\,apolar modelling approach (Section\\,4.3.2) for all samples. In the Galaxy, luminous YSOs and background sources display a larger fraction of CO$_2$ ice in the water-rich component ($\\sim$\\,85\\%) when compared to low-luminosity sources ($\\sim$\\,70\\%). The typical fraction for the LMC objects is $\\sim$\\,60\\%. This suggests that, even though both in the LMC and in the Galaxy the water-rich component is the major contributor to CO$_2$ ice, its fraction might be smaller for the LMC sources. However, we note that the LMC sample is small, and systematic uncertainties could bring the samples into better agreement. Even though a similar profile decomposition was used to isolate the water-rich contribution, there are differences in the modelling approach (e.g., number and properties of the components considered).  A contribution from pure-CO$_2$ ice is also needed to fit the majority of the profiles. Whether this contribution arises from inclusions in annealed ices or a pure-CO$_2$ layer resulting from CO desorption is less obvious. If CO$_2$ forms in either water- and/or CO-rich ice mixtures as has been widely suggested, then the presence of pure CO$_2$ diagnoses thermal processing. Therefore, our analysis of the morphology (i.e. shape and composition) of the 15.2-$\\mu$m CO$_2$ ice profiles of a sample of embedded YSOs in the LMC reveals only marginal differences when compared to Galactic YSOs. In the next section we discuss the concentration of CO$_2$ ice in relation to water ice.  \\subsection{CO$_2$ ice abundance in LMC}  Using compiled column-density measurements (Figure\\,\\ref{ratio}), we recompute the Galactic $N({\\rm CO}_2)/N({\\rm H_2O})$ ratio. Filled and open circles are for Galactic YSOs and background sources, respectively. Measurements for LMC YSOs are shown as star symbols. The YSOs in the \\citet{gerakines99} sample are luminous, while those in the \\citet{nummelin01}, \\citet{pontoppidan08} and \\citet{zasowski09} samples are low-luminosity objects. The \\citet{knez05} and \\citet{whittet07,whittet09} samples are exclusively of background sources.  Measurements of the column density of CO$_2$ ice are for either the 15.2-$\\mu$m bending mode \\citep{gerakines99,knez05,pontoppidan08,zasowski09,whittet09} or the 4.3-$\\mu$m stretching mode \\citep{nummelin01,shimonishi08}. Measurements of H$_2$O ice are for the 3-$\\mu$m stretching mode, except for \\citet{zasowski09}, who used the 6-$\\mu$m bending mode. They did not account for the contribution of other ice species to the 6-$\\mu$m optical depth. Therefore, they have probably overestimated the column density of water ice. To compare them with the other samples, we need to scale them accordingly. \\citet{boogert08} compared the column densities of water ice at 3 and 6\\,$\\mu$m, and they found that the 6-$\\mu$m value is on average higher than the 3-$\\mu$m value by a factor of 1.74 (the median of 49 sources; the standard deviation is 0.62). The adjusted column densities of \\citet{zasowski09} appear in Figure\\,\\ref{ratio}. We have also converted their column densities for CO$_2$ ice (computed using A\\,=\\,1.5\\,$\\times$\\,10$^{-17}$\\,cm\\,molecule$^{-1}$ ) to use the same band strength as the other measurements discussed here (A\\,=\\,1.1\\,$\\times$\\,10$^{-17}$\\,cm\\,molecule$^{-1}$).  Figure\\,\\ref{ratio} shows the strong correlation between $N({\\rm CO}_2)$ and $N({\\rm H_2O})$. The median value for the number density ratio is $N({\\rm CO}_2)/N({\\rm H_2O})$\\,$\\sim$\\,0.2 for luminous Galactic YSOs and background sources \\citep{gerakines99,whittet07,knez05,whittet09}. For low-luminosity Galactic YSOs, that ratio is $N({\\rm CO}_2)/N({\\rm H_2O})$\\,$\\sim$\\,0.3 \\citep{nummelin01,boogert04a,pontoppidan08} with a very large spread. \\citet{zasowski09} reported a ratio of 0.16; however if the correction factor derived in the previous paragraph is applied to the measured water column densities then that ratio becomes 0.28, fully consistent with the previous value. In fact, the right-hand panels (b and d) in Figure\\,\\ref{ratio} show that, with this correction, the low-luminosity measurements are essentially indistinguishable for the different samples. Note that there is a larger scatter for the measurements of low-luminosity objects when compared to the high-luminosity and background sample. Therefore, Figure\\,\\ref{ratio} suggests a clear difference between the high-luminosity YSO and background samples, and the low-luminosity YSO samples \\citep[see also][]{pontoppidan08}.  \\citet{shimonishi08} estimated a ratio $N({\\rm CO}_2)/N({\\rm H_2O})$\\,$\\sim$\\,0.45\\,$\\pm$\\,0.17 for five LMC YSOs. As mentioned previously, we revised some of Shimonishi's estimates downwards. For the seven LMC YSOs with both column-density measurements we estimate a median ratio $N({\\rm CO}_2)/N({\\rm H_2O})$\\,=\\,0.32\\,$\\pm$\\,0.15 (1\\,$\\sigma$). A single Shimonishi's measurement drives this large spread, it would otherwise be $N({\\rm CO}_2)/N({\\rm H_2O})$\\,=\\,0.32\\,$\\pm$\\,0.05 (1\\,$\\sigma$). Estimated total luminosities in our sample are 5\\,$-$50\\,$\\times$\\,10$^{3}$\\,L$_{\\odot}$, comparable to the massive YSO sample in the Galaxy. Figure\\,\\ref{ratio} (left panels) shows that if the high-luminosity Galactic and LMC samples are compared the $N({\\rm CO}_2)/N({\\rm H_2O})$ ratio is higher in the LMC, by a factor 1.6. It is curious, however, that the ratio measured for the LMC objects is consistent with the {\\em low-luminosity} Galactic sample. Note that column-density measurements do not depend on any modelling details but are prone to systematic uncertainties, namely in the continuum determination, that are generally difficult to quantify.  One should investigate the homogeneity of the samples discussed here. The LMC sample is restricted to very luminous YSOs, and it should, in terms of luminosity at least, be comparable to the sample of luminous Galactic YSOs observed with {\\it ISO}. \\spitzer\\ studies of star formation concentrated on nearby star-forming regions, molecular clouds and isolated globules that form predominantly intermediate- and low-luminosity objects. These low-luminosity Galactic samples include a variety of evolutionary stages from more embedded YSOs to Herbig AeBe, T Tauri and zero-age main-sequence stars. The different {\\em timescales} involved in the formation of objects with different luminosity may leave an imprint on the observed ice chemistry. Furthermore, YSOs in the LMC should in general appear more evolved due to the smaller dust column density in LMC objects when compared to Galactic objects.  Galactic and LMC measurements sample different {\\em spatial scales} (Section\\,2), from a few parsecs in the LMC to milliparsecs in the Galaxy. If there is a gradient of ice abundances between the more immediate YSO cocoon and the wider molecular environment in which it is immersed \\citep{pontoppidan06}, then the fact that ice properties are integrated over different spatial scales could affect global abundance ratios.  Such issues should be considered, since models of molecular clouds and their chemical evolution during the star formation process suggest that both temporal and radial variations of molecular concentrations are important \\citep[e.g.,][]{lee04,rodgers03,garrod08,vanweeren09,hollenbach09}. \\citet{lee04} singles out CO abundances (and the balance of its gas and solid phases) as particularly critical in determining the chemistry of star-forming cores. Even though they did not explicitly discuss CO$_2$ formation, the evolution of the two molecules is obviously closely linked. The balance of water in the gas phase and in ice form is also controlled by time- and space-dependent effects \\citep{hollenbach09}.  \\subsection{The star formation environment in the LMC}  The focus of this project is how the metal-poor star-formation environment in the LMC influences the chemistry of the envelopes that surround YSOs. Thus we now look at what is known about the environment in molecular clouds in the LMC.  The structure of Giant Molecular Clouds (GMCs) in metal-poor galaxies including the LMC is different from that in the Galaxy. Observations suggest that in metal-poor galaxies there are large reservoirs of H$_2$ that are not traced by CO \\citep[e.g.,][]{israel97}. The magnitude of this H$_2$-to-CO excess was found to depend strongly on both metallicity and strength of the radiation field. This has been interpreted as the result of selective photodissociation of CO in the outer parts of metal-poor GMCs \\citep[e.g.,][]{israel97,bolatto99}: H$_2$ readily self-shields while CO is shielded from the photodissociating radiation mostly by dust, which is less abundant at low metallicities. In this case, CO emission would trace only the inner parts of metal-poor GMCs. If CO cores are smaller in the LMC, due to the harsher UV radiation field \\citep[e.g.,][] {welty06} and lower dust abundance \\citep[e.g.,][]{bernard08} then we could naively expect CO$_2$ production to be reduced. A lower production rate of CO$_2$ could just result from the lower abundance of gas-phase CO observed in the LMC \\citep[e.g.,][]{bel86,fukui01,pineda09}. It could also result from a smaller density of dust grains that act as seeds for the surface chemistry. Both effects would then tend to reduce the value of $N({\\rm CO}_2)/N({\\rm H_2O})$. However as described above this is not what we observed in the LMC.  The formation of CO$_2$ ice in a wide variety of environments (e.g., both quiescent sightlines and low- and high-luminosity YSOs) argues against the need for enhanced radiative processing. Still, laboratory experiments do show that UV irradiation enhances CO$_2$ production \\citep[e.g.,][]{dhendecourt86}, and models of diffusive surface chemistry can produce high CO$_2$ densities at slightly elevated temperatures \\citep[e.g.,][]{ruffle01}. The LMC has a stronger ambient UV radiation field when compared to the Galaxy \\citep[by a factor of 1\\,$-$10,][] {bel86,welty06}. As a result of the stronger radiation field, the average ISM dust temperature is also generally higher in the LMC \\citep[12\\,$\\la$\\,$T_{\\rm d}$\\,$\\la$\\,35\\,K,][]{bernard08} than in the Galaxy \\citep[19\\,$\\la$\\,$T_{\\rm d}$\\,$\\la$\\,25\\,K,][]{reach96}  --- dust temperatures in cloud cores are lower than in the diffuse ISM. Both these effects could increase CO$_2$ production and increase $N({\\rm CO}_2)/N({\\rm H_2O})$, consistent with what is observed.  We should also consider the possibility that the formation of water ice is somewhat inhibited in the LMC. The largest column density of water ice measured so far in the LMC is $\\sim$\\,5\\,$\\times$\\,10$^{18}$\\,cm$^{-2}$ \\citep{shimonishi08}. This is below the highest column densities measured towards Galactic YSOs, 10\\,$-$30\\,$\\times$\\,10$^{18}$\\,cm$^{-2}$ \\citep{gibb04,pontoppidan08}. Due to the small number of LMC objects investigated to date and possible selection effects this may not be a significant difference. Water ice seems to form rather easily in molecular clouds. However the A$_{\\rm V}$-threshold for water ice formation ($A_{\\rm V}$\\,$\\sim$\\,3\\,mag for typical Galactic molecular clouds) is affected by the strength of the incident radiation field and dust temperature \\citep{hollenbach09}. Still, it is difficult to see how water ice formation could be inhibited when CO$_2$ ice formation is not."
    },
    "0911/0911.2701_arXiv.txt": {
        "abstract": " ",
        "introduction": "One of the most important open questions for cosmology is whether dark energy is a dynamical component of the universe or a cosmological constant. A plethora of experiments and future observations are currently planned with the aim of improving our understanding of this question (see for instance \\cite{Albrecht:2006um} for a review). One of the most popular models of dynamical dark energy is quintessence, where the acceleration of the universe is driven by a scalar field with negative pressure. Standard quintessence is described by a minimally coupled canonical scalar field \\cite{Zlatev:1998tr}. In this case, scalar fluctuations propagate at the speed of light and sound waves maintain quintessence homogeneous on scales smaller than the horizon scale \\cite{Ferreira:1997au}. Quintessence clustering takes place only on scales of order the Hubble radius,  so that its effect is strongly limited by cosmic variance.  However, quintessence modifies the growth evolution of dark matter through its different expansion history. Thus, future galaxy catalogs and weak lensing surveys will have a great potential in constraining the dark energy properties, in particular its equation of state, through a detailed study of the evolution of dark matter.   The study of quintessence perturbations -- where by quintessence we indicate a dark energy sector described by a single scalar degree of freedom -- has been recently revived in \\cite{Creminelli:2008wc}. Here the most general theory of quintessence perturbations around a given background was derived using the tools developed in \\cite{Creminelli:2006xe,Cheung:2007st}, formulated in the context of an effective field theory. An important conclusion is that quintessence with an equation of state $w < -1$ can be free from ghosts and gradient instabilities \\cite{Creminelli:2006xe,Creminelli:2008wc}.  In this regime, stability can be guaranteed by the presence of higher derivative operators \\cite{ArkaniHamed:2003uy,Creminelli:2006xe} with the requirement that the speed of sound of propagation is extremely small, $|c_s| \\lesssim 10^{-15} $. On cosmological scales, higher derivative operators are phenomenologically irrelevant and quintessence simply behaves as a perfect fluid with negative pressure but practically zero speed of sound \\cite{Creminelli:2008wc}. Interestingly, this description applies also in the non-linear regime,  i.e.~when perturbations in the quintessence energy density become non-linear,  as long as the effective theory remains valid.    This study motivates the possibility that quintessence has a practically zero speed of sound. Apart from these theoretical considerations, the fact that the speed of sound of quintessence may vanish opens up new observational consequences. Indeed, the absence of quintessence pressure gradients allows instabilities to develop on all scales, also on scales where dark matter perturbations become non-linear. Thus, we expect quintessence to modify the growth history of dark matter not only through its different background evolution but also by actively participating to the structure formation mechanism, in the linear and non-linear regime, and by contributing to the total mass of virialized halos.   In the linear regime, a series of articles have investigated the observational consequences of a clustering quintessence. In particular, they have studied the different impact of quintessence with $c_s=1$ or $c_s=0$ on the cosmic microwave background \\cite{DeDeo:2003te,Weller:2003hw,BeanDore,Hannestad:2005ak}, galaxy redshift surveys \\cite{Takada:2006xs}, large neutral hydrogen surveys \\cite{TorresRodriguez:2007mk}, or on the cross-correlation of the integrated Sachs-Wolfe effect in the cosmic microwave background with large scale structures \\cite{Hu:2004yd,Corasaniti:2005pq}.  On non-linear scales, the dependence of the dark matter clustering on the equation of state of a homogeneous quintessence, i.e.~with $c_s=1$, has been investigated using N-body simulations in a number of articles (see for instance \\cite{Alimi:2009zk} and references therein for a recent account).  A popular analytical approach to study non-linear clustering of dark matter without recurring to N-body simulations is the spherical collapse model \\cite{Gunn:1972sv}. In this approach, one studies the collapse of a spherical overdensity and determines its critical overdensity for collapse as a function of redshift. Combining this information with the extended Press-Schechter theory \\cite{Press:1973iz,Bond:1990iw} one can provide a statistical model for the formation of structures which allows to predict the abundance of virialized objects as a function of their mass. Although it fails to match the details of N-body simulations, this simple model works surprisingly well and can give useful insigths into the physics of structure formation. Improved models accounting for the complexity of the collapse exist in the literature and offer a better fit to numerical simulations. For instance, it was shown in \\cite{Sheth:1999mn} that a significant improvement can be obtained by considering an ellipsoidal collapse model. See also \\cite{Maggiore:2009rv,Maggiore:2009rw} for recent theoretical developments and new improvements in the excursion set theory.   The spherical collapse can be generalized to include a cosmological constant (see for instance \\cite{Lahav:1991wc}) and quintessence with  $c_s=1$ \\cite{Wang:1998gt} (see also \\cite{Weinberg:2002rd,Battye:2003bm} for subsequent applications). If quintessence propagates at the speed of light it does not cluster with dark matter but remains homogeneous. Indeed, pressure gradients contribute to maintain the same energy density of quintessence between the inner and outer part of the spherical overdensity. A study of the spherical collapse model with different quintessence potentials was performed in \\cite{Mota:2004pa}. For a nice review on structure formation with homogeneous dark energy see also \\cite{Percival:2005vm}.   In this paper we study the spherical collapse model in the case of quintessence with zero speed of sound. This represents the natural counterpart of the opposite case $c_s=1$. Indeed, in both cases there are no characteristic length scales associated to the quintessence clustering\\footnote{The characteristic length scale associated to the quintessence clustering is the sound horizon scale, i.e., $L_s \\equiv a \\int c_s dt/a$. As mentioned above, this vanishes for $c_s=0$ so that clustering takes place on all scales. For $c_s=1$ we have $L_s = 2 H_0^{-1}$, which is much larger than the scales associated to the spherical collapse.}  and the spherical collapse remains independent of the size of the object. For this study we describe quintessence using the model developed in \\cite{Creminelli:2008wc}. As explained, this description remains valid also when perturbations become non-linear and can thus be applied to the spherical collapse model.  As the spherical collapse occurs on length scales much smaller than the Hubble radius, we will describe it using a convenient coordinate system where the effect of the Hubble expansion can be treated as a small perturbation to the flat spacetime. In these ``local'' coordinates the description of the spherical collapse becomes extremely simple and the well-known cases can be easily extended to quintessence with $c_s=0$. In this case pressure gradients are absent and quintessence follows dark matter during the collapse. Thus, in contrast with the non-clustering case $c_s=1$ where quintessence and dark matter are not comoving, for $c_s=0$ the collapsing region is described by an exact Friedmann-Lema\\^{i}tre-Robertson-Walker (FLRW) universe. Note that even though the energy density of quintessence develops inhomogeneities as long as the collapse proceeds, the pressure inside and outside the overdense region remains the same. Thus, as explained below, our model does not give the same description of clustering quintessence as that proposed by \\cite{Mota:2004pa} and studied, for instance, in \\cite{Maor:2005hq,Manera:2005ct,Nunes:2004wn}.  We will see that, besides quantitative differences with respect to the $c_s=1$ case -- a different threshold for collapse and a different dark matter growth function -- the $c_s=0$ case has a remarkable qualitatively new feature. Quintessence clusters together with dark matter and participates in the total mass of the virialized object, contributing to their gravitational potential.   The plan of the paper is the following. In section \\ref{sec:model} we describe quintessence models with $c_s=0$. In section \\ref{sec:local_coords} we study spherical collapse solutions first in known cases (dark matter only, $\\Lambda$CDM and $c_s=1$ quintessence) and then in the case of a clustering dark energy,  $c_s=0$. It turns out to be much simpler to describe these solutions in coordinates for which the metric is close to Minkowski around a point in space. The equation for the evolution of the spherical collapse are solved in section \\ref{sec:solving} and the threshold for collapse is calculated in the various cases. This leads to the calculation of the dark matter mass function in section \\ref{sec:massfunction}. In section \\ref{sec:quintmass} we study the accretion of quintessence to the dark matter haloes and its effect on the total mass function. The contribution of quintessence to the mass may be distinguished from the dark matter and baryon component in cluster measurements, as we briefly discuss in section \\ref{sec:3masses}. Conclusions and future directions are discussed in section \\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}      Using the spherical collapse model, we have studied how the clustering of quintessence with negligible speed of sound can affect the prediction for the mass function of dark halos. As quintessence does not develop pressure gradients, it follows geodesic motion remaining comoving with the dark matter. In contrast to the case where quintessence remains smooth, spherical collapsing regions behave as exact closed FLRW solutions with the quintessence contributing to the overdensity. To study the spherical collapse we found it useful to use Fermi coordinates where the effect of the expansion can be treated as a perturbation around Minkowski spacetime locally around a spatial point. In contrast to comoving coordinates, a unique coordinate system can be employed to describe both the interior and exterior of the collapsing region.    Quintessence with zero speed of sound modifies dark matter clustering with respect to the smooth quintessence case through the linear growth function and the linear threshold for collapse. Besides these conventional effects there is a more important and qualitatively new phenomenon: quintessence mass adds to the one of dark matter, contributing to the halo mass by a fraction of order $(1+w) \\Omega_Q/\\Omega_m$. This effect is quite difficult to model and in this paper we have adopted a simplified treatment which gives an accurate estimate of the high mass tail of the distribution where the effect is more relevant.   As dark energy plays an active role in the formation of structures, the distinction between what we call dark matter and dark energy becomes fuzzy. The distinction between the two components can be probed by the different redshift dependence. It is quite remarkable that our knowledge of structure formation still allows a completely inhomogeneous dark energy component on short scales.     In this paper we did not attempt to study experimental constraints and forecasts. However, we have seen that for values of $w$ which are not too close to the cosmological constant one, let's say $|1+w| \\gtrsim 0.1$, the predictions for the high mass tail of the mass function in the $c_s=0$ and the $c_s=1$ cases are quite distinctive. The effect of clustering quintessence gives order one changes in the expected number of clusters. Whether the effect will be measurable or not will depend on the possibility of breaking the degeneracy with cosmological parameters, most notably  $\\sigma_8$ and $\\Omega_m$, and on the good recostruction of the limiting mass of the survey. Similar concerns have been addressed in the context of the effects of primordial non-Gaussianity on the mass function in \\cite{LoVerde:2007ri}.  Our work can be continued in various directions.  A better theoretical treatment of quintessence accretion would strengthen our predictions for the mass function, although at a certain point only a (challenging) numerical simulation can make a fully reliable prediction. On the data side it would be interesting to understand what is the minimum value of $|1+w|$ that allows a distinction between $c_s=0$ and $c_s=1$ using the forthcoming cluster mass function measurements. One should also explore whether other probes, lensing for example, are better suited to study the non-linear clustering of quintessence. It would also be interesting to explore the effect of quintessence on the halo dynamics, but this probably requires a way to deal with the formation of caustics. In a more model-independent way, one could try to use data on clusters to constrain the presence of extra gravitational mass besides the dark matter mass. We leave all this for future work."
    },
    "0911/0911.1158_arXiv.txt": {
        "abstract": "We present near--IR images, obtained with the {\\it Hubble Space Telescope} ({\\it HST}) and the WFC3/IR camera, of six passive and massive galaxies at redshift $1.3<z<2.4$ (SSFR$<10^{-2}$ Gyr$^{-1}$; stellar mass $M\\sim10^{11}$ M$_{\\odot}$), selected from the Great Observatories Origins Deep Survey (GOODS). These images, which have a spatial resolution of $\\sim1.5$ kpc, provide the deepest view of the optical rest--frame morphology of such systems to date. We find that the light profile of these galaxies is regular and well described by a S\\'ersic model with index typical of today's spheroids. Their size, however, is generally much smaller than today's early types of similar stellar mass, with four out of six galaxies having $r_e\\sim1$ kpc or less, in quantitative agreement with previous similar measures made at rest--frame UV wavelengths. The images reach limiting surface brightness $\\mu\\sim$26.5 mag arcsec$^{-2}$ in the F160W bandpass; yet, there is no evidence of a faint halo in the galaxies of our sample, even in their stacked image.  We also find that these galaxies have very weak ``morphological {\\cal k}--correction'' between the rest--frame UV (from the ACS $z$--band), and the rest--frame optical (WFC3 $H$--band): the S\\'ersic index, physical size and overall morphology are independent or only mildly dependent on the wavelength, within the errors. ",
        "introduction": "Most of the stars observed in today's early--type galaxies have formed at high redshifts ($z>2$, see Renzini 2006), but the observations have not yet constrained how that stellar mass has assembled in the presently observed systems.  Several groups have reported that at least in the range $M>M\\sim10^{11}$ M$_{\\odot}$, the stellar mass function of these high--redshift ``elliptical galaxies'', appears to be rapidly increasing from $z\\sim3$ to $z\\sim1$, pointing to this epoch as one of major assembly for modern bright galaxies (e.g. Bundy~et~al.~2006; Fontana~et~al.~2006; Arnouts~et~al.~2007; Scarlata~et~al.~2007; Ilbert~et~al.~2009). A number of mechanisms have been proposed as drivers of this evolution, from merging to gradual in--situ accretion (Naab et al. 2009; van Dokkum et al. 2009; Hopkins et al. 2010).  Elliptical galaxies at $z>1.5$ are often observed to be smaller, by factors $\\sim3-5\\times$, than their local counterparts of similar stellar mass, and thus to have $\\approx 30-100\\times$ higher stellar density (Daddi~et~al.~2005; Trujillo~et~al.~2006; Trujillo~et~al.~2007; Toft~et~al.~2007; Zirm~et~al.~2007; Longhetti~et~al.~2007; Cimatti~et~al.~2008, Van~Dokkum~et.~al.~2008, Buitrago~et~al.~2008, Toft~et~al.~2009), a fact that still lacks an explanation in terms of an evolutionary mechanism.  Some have suggested that current observations at high redshift could have missed low surface-brightness halos surrounding the galaxies, due to the relatively limited sensitivity combined with the $(1+z)^4$ cosmological dimming (Hopkins~et~al.~2009b; Mancini~et~al.~2009). Others have proposed that surveys of the local universe (e.g. the SDSS, Stoughton~et~al.~2002) could have missed a relatively large fraction of very compact, yet massive early--type galaxies due to the limited ground--based angular resolution (1'' corresponds to a physical scale of $\\sim1$ kpc at $z=0.05$). While recent works (Valentinuzzi~et~al.~2009; Trujillo~et~al.~2009) reported the identification of dense and massive galaxies in the local universe, their spatial abundance appears much smaller than the $z\\sim2$ galaxies.  Furthermore, while small high--redshift samples have been imaged with {\\it HST} and NICMOS at rest--frame optical wavelengths (Buitrago~et~al.~2008; Trujillo~et~al.~2007; Toft~et~al.~2007; Zirm~et~al.~2007; Longhetti~et~al.~2007; Van~Dokkum~et.~al.~2008), most of the galaxies have only been observed in the rest--frame UV with ACS, and the dependence of the morphology of these galaxies on the wavelength has not been characterized (Cimatti~et~al.~2008; Daddi~et~al.~2005).  In this letter we use the unprecedented sensitivity and angular resolution of WFC3/IR images recently acquired in the HUDF to study the optical rest--frame morphology of a small but well defined sample of low SSFR galaxies at z$\\sim$2. AB magnitudes are used throughout this paper, and, when needed, a $\\Lambda$-CDM world model with $H_0$=70 km~s$^{-1}$~Mpc$^{-1}$, $\\Omega_M=0.3$ and $\\Omega_{\\Lambda}=0.7$. ",
        "conclusions": "Massive galaxies ($M\\sim10^{11}$ M$_{\\odot}$) with early spectral type at $z\\sim2$ appear to often have a relatively smooth and regular morphology, nearly independent of the wavelength, and with light profile which is well described by a S\\'ersic model, i.e. their morphology can be classified as ``elliptical'', although systems with an exponential component are also observed. This suggests that a tight correlation between spectral and morphological properties similar to what is observed in the local universe, i.e. the back bone of the Hubble sequence, is already in place among bright galaxies at $\\sim25$\\% of the cosmic time. The S\\'ersic analysis has also quantitatively shown that the morphology of these galaxies depends very mildly on the wavelength, between the rest--frame UV and optical windows.  One of our six galaxies is almost certainly an AGN, while another with a spatially unresolved nuclear component is a candidate one. Such kind of contamination should not be surprising given the similarity between the SED of obscured AGN and quiescent galaxies. By the same token, however, it is also possible that both types of sources contribute to the luminous output of the same galaxy, a fact that we have not directly investigated in this work. The apparently high level of contamination (33\\% of our sample) observed at $z\\sim2$, the cosmic era when both AGN and star--formation activity both reach their maximum (Ueda~et~al.~2003), raises the intriguing question of whether this is direct evidence of a link between the AGN and the mechanisms responsible for the cessation of star formation in the galaxies. We plan to report on this issue using larger samples from imminent new observations with WFC3 of the GOODS fields. The high AGN contamination is also likely to bias statistical conclusions made from UV/Optically--selected samples of passive galaxies at $z\\sim2$, unless sensitive X--ray and mid--IR photometry (Daddi~et~al. 2005, 2007), or at least multi--wavelength high--resolution {\\it HST} imaging, can be used (like in the GOODS) to cull these sources.  In surveys where such information is not available, like some recent ground--based ones (e.g. van Dokkum et al. 2009), the AGN so frequently encountered among the galaxies cannot be recognized.  Four out of six of our $z\\sim2$ ellipticals lie below the local mass--size relation for early--type galaxies, in agreement with previous findings that they can be up to $\\sim3$--5 times more compact at given mass, and thus $\\sim30$--100 times denser, than their local counterparts.  Of the remaining two galaxies, however, one (\\# 19389), which has an unresolved central component, is at the boundary of the local relation, while the other (\\# 24626) is well within it. We have not found any evidence of a low--surface brightness halo surrounding these galaxies, either in the individual images, or in their stacked one. At most, only $<2$\\% of the light (and hence of the mass, assuming a constant M/L ratio) can be segregated in such a halo. Galaxy \\# 24626, requires a two--component light profile, one S\\'ersic model with high $n$ (spheroid) embedded in a disk, to provide a good fit to the observations, but both have the same, relatively large, effective radius, $r_e\\sim3$ kpc. Thus, what is emerging here is that these galaxies are characterized by a dispersion of properties that also includes systems similar to the local ones, as some have already suggested (Cimatti~et~al.~2008; Mancini~et~al.~2009), which however still needs to be studied and characterized. Until this work is done, whether and/or how the ultra--compact $z\\sim2$ ellipticals evolve into the local ones or whether ultra--compact systems are also present in the local universe but have so far eluded detection will remain open questions."
    },
    "0911/0911.3180_arXiv.txt": {
        "abstract": " ",
        "introduction": "An issue arising frequently in the literature on NGC 5128 is a lack of consensus about the best distance to adopt for this important nearby elliptical galaxy. As with any other galaxy, the assumed distance directly affects the astrophysics of all its components through the calculation of every intrinsic scale length and every luminosity, at every wavelength.  Some researchers use a distance of 3.9 Mpc from a weighted average of several methods which measure the properties of old resolved stars \\citep{rej04}. Others quote a value of d = 3.4 Mpc which is based only on Cepheids \\citep{fer07}. This uncomfortable uncertainty of almost 20\\% is much larger than one would like for a galaxy we all agree is a keystone object regardless of the wavelength region in which we work. Clearly the need exists for some kind of ``best estimate'' distance.  Fortunately there are now several distance indicators available to us, and the time is right to construct a useful mean distance measurement. None of these methods on its own can be considered by definition to give the {\\it correct} value, because each is limited by its own set of systematic (and random) uncertainties. The strategic advantage of combining several methods is that it will allow us to use many different stellar components of the galaxy, each with its distinctive merits and each of which is at least partly independent of the others. ",
        "conclusions": "The extragalactic distance scale is a classic astronomical subject whose roots extend far into the past. It will already be evident from the discussion in preceding sections that the calibrations and zeropoints of standard candles are a constant source of concern, revision, and debate in the distance scale literature. With the appearance of new distance calibrators, several specialist conferences have been dedicated to this topic in the last two decades. Among the newcomers in the group of well established distance candles over the last $\\sim 20$ years are TRGB, PNLF, and SBF, three of the five distance indicators used to determine distance to NGC~5128.  \\begin{table} \\begin{center} \\caption{Summary of Distance Calibrations}\\label{table} \\begin{tabular}{lc} \\hline Method & $(m-M)_0$ \\\\ \\hline \\\\ Cepheids         &  $27.67 \\pm 0.20$ \\\\ TRGB (I), (JHK)  &  $27.91 \\pm 0.08$ \\\\ PNLF             &  $27.92 \\pm 0.10$ \\\\ LPV (Miras)      &  $27.96 \\pm 0.11$ \\\\ SBF              &  $27.74 \\pm 0.14$ \\\\ \\\\ $\\langle m-M \\rangle_0$ & $27.91 \\pm 0.05^a$ \\\\ \\\\ $\\langle d \\rangle$ (Mpc) & $3.82 \\pm 0.09^a$ \\\\ \\hline \\end{tabular} \\medskip\\\\ $^a$Average of the first four methods listed (see text).\\\\ \\end{center} \\end{table}  Although we have stressed the advantages of combining independent approaches, none of our four key resolved-star standard candles can ultimately be said to be totally independent of all the others.  The distance scale for the nearby universe requires the careful assembly and intercomparison of many techniques, with checks and balances at every outward step. Convergence may at times seem remarkably slow and difficult. Nevertheless, at this stage in the development of the subject, wide agreement to within $\\pm 0.1$ mag has been achieved for galaxies within the Local Group. We use this ``near-field'' Local Group region as our starting point.  Brief comparisons of the basis of each method show how they are interrelated at their starting points and how none are truly ``independent'' of the others.  But These same points of overlap provide several strong consistency checks:  \\noindent (i) The fundamental calibration of the Cepheid PL relation relies on trigonometric parallaxes of nearby Cepheids and main-sequence fitting to young Milky Way star clusters containing Cepheids, often supplemented by Baade-Wesselink method parallaxes and the PL slope from the LMC \\citep[for only a recent sampling of the vast literature, see][]{van07,an07,gro08}.  \\noindent (ii) The TRGB luminosity calibration \\citep[well reviewed by][]{rizzi+07} relies on observations of the tip luminosity in Milky Way globular clusters (particularly $\\omega$ Centauri) and the old-halo components of Local Group dwarf galaxies, all of whose distances rely in turn on the luminosities of the old RR Lyrae and horizontal-branch stars.  \\noindent (iii) The calibration of the Mira PL relation rests on local Milky Way long-period variables and importantly on the LMC LPVs; the LMC fiducial distance relies in turn on a mixture of indicators including RR Lyraes, Population I and II Cepheids, and the expansion-shell parallax to SN1987a.  \\noindent (iv) Finally, the PNLF calibration that we adopt here depends strongly on the distance to M31, which in turn is calibrated from Cepheids, Miras, RR Lyraes, and TRGB.  As soon as we step beyond the Local Group, measurements of distances to individual galaxies are occasionally affected by larger disagreements among the methods that are sometimes still hard to understand (see the papers above for several examples).  For NGC 5128, however, the various methods have finally begun to converge, showing encouraging agreement to within their internal uncertainties.  It is worth noting that the true uncertainty applicable to each method is currently dominated by the {\\sl external} accuracy of the calibration, which is $\\pm 0.1$ mag or greater. The {\\sl internal} measurement uncertainty -- e.g. the apparent magnitude of the RGB tip, or $M*$ for the PNLF -- is now below 0.1 mag thanks to large observational samples and rigorous numerical analysis methods.  Table 1 summarizes our results.  The uncertainty quoted for each one is the combination of internal and external errors. A simple weighted mean of all five methods gives $(m-M)_0 = 27.89 \\pm 0.04$ or $d=3.77 \\pm 0.08$ Mpc. Given the discussion above, however, we recommend a final average based on the four primary, resolved-star methods (Cepheids, TRGB, PNLF, Miras). This average gives a slightly larger distance of $(m-M)_0 =27.91 \\pm 0.08$ or $d=3.8 \\pm 0.1$ Mpc.  The mutual agreement among these methods is about as good as we have for any galaxy beyond the Local Group. It appears that, to within $\\pm 0.1$ Mpc, the recommended distance of 3.8 Mpc for NGC 5128 is well supported by the evidence at hand."
    },
    "0911/0911.2715_arXiv.txt": {
        "abstract": "We present an analysis of the components of solar wind proton temperature perpendicular and parallel to the local magnetic field as a function of proximity to plasma instability thresholds. We find that $T_{\\perp p}$ is enhanced near the mirror instability threshold and $T_{\\parallel p}$ is enhanced near the firehose instability threshold. The increase in $T_{\\perp p}$ is consistent with cyclotron-resonant heating, but no similar explanation for hot plasma near the firehose limit is known.  One possible explanation is that the firehose instability acts to convert bulk energy into thermal energy in the expanding solar wind, a result with significant implications for magnetized astrophysical plasma in general. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0911/0911.2838_arXiv.txt": {
        "abstract": "The statefinder diagnosic is a useful method for distinguishing different dark energy models. In this paper, we investigate the new agegraphic dark energy model with interaction between dark energy and matter component by using statefinder parameter pair  $\\{r, s\\}$ and study its cosmological evolution. We plot the trajectories of the new agegraphic dark energy model for different interaction cases in the statefinder plane. As a result, the influence of the interaction on the evolution of the universe is shown in the statefinder diagrams. ",
        "introduction": "\\label{sec:introduction}  Recent astronomical observations of type Ia supernovae (SNIa) indicate that the universe is undergoing an accelerating expansion~ \\cite{Riess98,Perl99}. This cosmic acceleration has also been confirmed by other observations, such as observations of large scale structure (LSS)~\\cite{Tegmark04,Abazajian04} and measurements of the cosmic microwave background (CMB) anisotropy~\\cite{Spergel03,Bennett03}. Nowadays the most well-accepted idea is that a mysterious dominant component --- dark energy --- with large enough negative pressure is responsible for this the cosmic acceleration. Despite the fact that the cosmological origin of dark energy remains enigmatic at present, physicists have to face the intriguing physical problem and try to understand the ultimate nature of dark energy. Among all theoretical models, the preferred candidate of dark energy is the Einstein's cosmological constant $\\Lambda$. The simplest cosmological model is the so-called $\\Lambda$CDM (or LCDM) model, which consists of a mixture of the cosmological constant $\\Lambda$ and the cold dark matter (CDM). The LCDM model provides an excellent explanation for the acceleration of the universe and the existing observational data. However, the cosmological constant faces two difficulties, namely, the ``fine-tuning\" problem and the ``cosmic coincidence\" problem. The former asks why the cosmological constant observed today is so much smaller than the Plank scale, while the latter asks why the energy densities of dark energy and matter are on the same order today. Theorists have made lots of efforts to try to resolve the cosmological constant problem but all these efforts have turned out to be unsuccessful.  In order to alleviate or even solve these two problems, many dynamical dark energy models have been proposed, whose equation of state is no longer a constant but slightly evolves with time. The dynamical dark energy scenario is often realized by some scalar-field mechanism which suggests that the energy with negative pressure is provided by a scalar field evolving down a proper potential. A lot of scalar field dark energy models have been studied, such as quintessence, \\textit{K}-essence, tachyon, phantom and quintom etc.. Besides, some interacting models have been discussed in many works to help understand or alleviate the coincidence problem by considering the possible interaction between dark energy and dark matter owing to their unknown nature. For reviews of dark energy, see, e.g., Ref.~\\cite{dark energy}.  On the other hand, in various dark energy models, the property of dark energy is strongly model-dependent. In order to be capable of differentiating those competing cosmological dark energy scenarios, a sensitive diagnostic for the many dark energy models is a must. For characterizing the expansion history of the universe, one defines the geometric parameters $H=\\dot{a}/a$ and $q=-\\ddot{a}/aH^2$, namely, the Hubble parameter and the deceleration parameter; here $a(t)$ is the scale factor of the Friedmann-Robertson-Walker (FRW) universe. It is clear that $\\dot{a}>0$ means that the universe is undergoing an expansion and $\\ddot{a}>0$ means the universe is experiencing an accelerated expansion. The cosmic acceleration indicates that $q$ should be less than zero. However the deceleration parameter on its own does not characterize the current acceleration phase uniquely. The presence of a fairly large degeneracy in $q$ is reflected in the fact that rival dark energy models can give rise to the same value of $q_0$ at the present time. Under such circumstances, a robust diagnostic of dark energy, statefinder parameter pair $\\{r(z), s(z)\\}$, was introduced by Sahni et al.~\\cite{Sahni:2002fz} and Alam et al.~\\cite{Alam:2003sc}. In addition, more recently, two new diagnostics of dark energy, $Om$ and acceleration probe $\\bar{q}$, were introduced by Sahni, Shafieloo and Starobinsky~\\cite{Sahni:2008xx}.  The statefinder probes the expansion dynamics of the universe through high derivatives of the scale factor $\\ddot a$ and $\\dddot a$ and is a natural next step beyond the Hubble parameter $H$ depending on $\\dot a$ and the deceleration parameter $q$ depending on $\\ddot a$. The statefinder pair $\\{r, s\\}$ is defined as \\begin{equation} r\\equiv\\frac{\\dddot a}{aH^3},\\label{defr} \\end{equation} \\begin{equation} s\\equiv\\frac{r-1}{3(q-\\frac{1}{2})} .\\label{defs} \\end{equation} It is a ``geometrical'' diagnostic, in the sense that it is constructed from a space-time metric directly, and it is more universal than ``physical'' variables, which depend upon properties of physical fields describing dark energy. So, in order to see the qualitatively different cosmological evolution behaviors of dark energy models in degeneracy of $H_0$ and $q_0$, we can plot statefinder parameter diagrams corresponding to these dark energy models by theoretical calculation. As a reference the spatially flat LCDM scenario corresponds to a fixed point $\\{ r, s \\} = \\{1,0\\} $ in this diagram. Departure of a given dark energy model from this fixed point provides a good way of establishing the ``distance'' of this model from the LCDM~\\cite{Alam:2003sc}. On the other hand, the statefinder can also be extracted from data coming from SuperNova Acceleration Probe (SNAP) type experiments~\\cite{Sahni:2002fz,Alam:2003sc}. Therefore, the statefinder diagnostic combined with future SNAP observations may possibly be used to discriminate between different dark energy models. In this paper, we just apply the statefinder diagnostic to the new agegraphic dark energy (NADE) model.  We will investigate the features of the NADE model with interaction with matter component from the statefinder view point. In Sec.~\\ref{sec:agegrapgic}, we will briefly review the NADE model and introduce an interacting model of NADE. In Sec.~\\ref{sec:evolution}, we will study the cosmological evolution of the interacting NADE model. In Sec.~\\ref{sec:statefinder}, we will apply the statefinder diagnostic to the interacting NADE model. In the last section we will give conclusions. ",
        "conclusions": "\\label{sec:conclusion}  In summary, we have studied the interacting NADE model from the statefinder viewpoint in this paper. Since the accelerated expansion of the universe was discovered by astronomical observations, many cosmological models have been proposed to interpret this cosmic acceleration. This leads to a problem of how to discriminate between these various contenders. The statefinder diagnosis is a useful tool for distinguishing different cosmological models by constructing the parameters $\\{r, s\\}$ using the higher derivative of the scale factor. Moreover, the value of $\\{r, s\\}$ of today can be viewed as a discriminator for testing various cosmological models if it can be extracted from precise observational data in a model-independent way. On the other hand, although we are lacking an underlying theory of dark energy, we still can make some efforts to probe the properties of dark energy according to some principle of quantum gravity. The NADE model, constructed in light of the Karolyhazy relation and corresponding energy fluctuations of space-time, is seen to possess some features of quantum gravity theory and provides us with an attempt to explore the essence of dark energy. In addition, some physicists believe that the involvement of interaction between dark energy and dark matter leads to some alleviation and more understanding to the coincidence problem. Thus, it is worthwhile to investigate the interacting NADE model. We do this by applying the statefinder parameters as a diagnostic tool and plot the statefinder trajectories in the $s-r$ plane. We learn that the interaction between dark energy and dark matter can significantly affect the evolution of the universe, and the contribution of the interaction can be diagnosed out explicitly in this method. In addition, we show cosmological evolution of $E(z)$. For this interacting dark energy model, the parameters $n$ and $\\alpha$ both play important roles and thus affect the cosmological evolution. But, to determine $n$ and $\\alpha$, we need more precise data provided by future experiments. We hope that the future high-precision observations can offer more and more accurate data to determine these parameters precisely and consequently shed light on the essence of dark energy."
    },
    "0911/0911.2465_arXiv.txt": {
        "abstract": "{Direct detection of exoplanets requires high dynamic range imaging. Coronagraphs could be the solution, but their performance in space is limited by wavefront errors~(manufacturing errors on optics, temperature variations, etc.), which create quasi-static stellar speckles in the final image.} {Several solutions have been suggested for tackling this speckle noise. Differential imaging techniques substract a reference image to the coronagraphic residue in a post-processing imaging. Other techniques attempt to actively correct wavefront errors using a deformable mirror. In that case, wavefront aberrations have to be measured in the science image to extremely high accuracy.} {We propose the self-coherent camera sequentially used as a focal-plane wavefront sensor for active correction and differential imaging. For both uses, stellar speckles are spatially encoded in the science image so that differential aberrations are strongly minimized. The encoding is based on the principle of light incoherence between the hosting star and its environment.} {In this paper, we first discuss one intrinsic limitation of deformable mirrors. Then, several parameters of the self-coherent camera are studied in detail. We also propose an easy and robust design to associate the self-coherent camera with a coronagraph that uses a Lyot stop. Finally, we discuss the case of the association with a four-quadrant phase mask and numerically demonstrate that such a device enables detection of Earth-like planets under realistic conditions.} {The parametric study of the technique lets us believe it can be implemented quite easily in future instruments dedicated to direct imaging of exoplanets.} ",
        "introduction": "Exoplanets are typically~$10^7$ to~$10^{10}$ fainter than their host and are often located within a fraction of an arcsecond from their star. Numerous coronagraphs have been proposed to reduce the overwhelming light of a star to obtain a direct imaging of extrasolar planets~\\citep{Rouan00,Mawet05,Guyon05}. Several of them provide observations~\\citep{SchneiderG98,Boccaletti04}. But performance is limited by wavefront errors in the upstream beam for all these coronagraphs and the final focal plane image shows stellar speckles. The effect of most of these aberrations can be corrected by adaptive optics~(AO) or eXtreme AO~\\citep[XAO,][]{Verinaud08} but the uncorrected part generates quasi-static speckles, which limit the image contrast\\,\\citep{Cavarroc06,macintosh05}. To reduce this speckle noise, differential imaging techniques attempt to subtract a reference image of the stellar speckles from the science image~(star plus companion).  Several ways are used to measure that reference: spectral characteristics~\\citep{racine99,marois00,marois04}, polarization states~\\citep{Baba03,Stam04}, differential rotation in image~\\citep{SchneiderG98b,marois06}, or incoherence between stellar and companion lights~\\citep{Guyon04}. The self-coherent camera with which we work uses the last property~\\citep{Baudoz06,Galicher07}. But before using one of these {\\it a posteriori} techniques, we may actively correct quasi-static wavefront errors so that a first speckle reduction is achieved and differential imaging techniques have less to do.  Because of the low level of aberrations that must be achieved~(a few nanometers), the wavefront sensor of that loop has to measure for phase and amplitude errors in the final science image to avoid differential errors introduced by classical wavefront sensor~\\citep[Shack-Hartmann for example,][]{shack71}. \\citet{Codona04} suggest using the incoherence between companion and stellar lights and use a modified Mach-Zender interferometer to encoded the stellar speckles.  The instrument we propose, the self coherent camera~(SCC), is based on the same property and uses Fizeau interferences. We insist on how~SCC can be used both as a wavefront sensor for active correction~(called step~A in this paper) and as a differential imaging technique~(step~B). In~\\citet{Galicher08}, we numerically demonstrated that, applying step~B after step~A, a self-coherent camera associated with a~$32$x$32$ deformable mirror and a perfect coronagraph detects earths~(contrast of~$2\\,10^{-10}$) from space in a few hours under realistic assumptions~(zodiacal light, photon noise, read-out noise, phase errors of~$20$\\,nm rms, $20\\%$~bandwidth and $1\\%$~effective bandwidth,~$8$\\,m telescope with a~$25\\%$ throughput and $\\lambda_0=800$\\,nm). Here, we report results from a parametric study of the~SCC, and propose an easy and robust design to associate it with coronagraphs that use a Lyot stop. We explain the performance in the case of a four-quadrant phase mask coronagraph~\\citep[FQPM,][]{Rouan00}. Section\\,\\ref{sec : sccprin} recalls the principle of the technique and presents the estimators of the pupil complex amplitude~(phase and amplitude errors, step~A) and companion images~(step~B). Section~\\ref{sec : assump} provides the assumptions and criterions used for the parametric studies. Section~\\ref{sec : DM} is a general study~(no SCC) of one intrinsic limitation for deformable mirrors and presents the best contrast they can provide. The signal-to-noise ratios on both~SCC estimators~(wavefront and companion) are developed in Sect.~\\ref{sec : snr}. Section~\\ref{sec : reference} estimates the required stability of the reference beam. We report the effects of amplitude aberrations on the~SCC performance in Sect.~\\ref{sec : amplitude}. The last two sections are the most important in the paper. The first one is the study of the chromatism on the SCC~performance when a perfect coronagraph is used~(Sect.~\\ref{sec : chrom}). In the second one~(Sect.~\\ref{sec : sccrealcoro}), we propose a device to associate the SCC and any coronagraph using a Lyot stop. We discuss the case of a~FQPM coronagraph: earths are detected using such a coronagraph with the~SCC just by adding a small hole to the Lyot stop. ",
        "conclusions": "In Sect.~\\ref{sec : DM}, we provided the intrinsic limitation for deformable mirrors controlled by the algorithm of~\\citet{Borde06}, under realistic yet optimistic assumptions~(no dead actuators, continuous face sheet). It is important to keep in mind that this limitation does not depend on the technique used to estimate for wavefront errors since we assumed a perfect estimation. One way to improve the~DM best contrast could be an apodization of the pupil so that the uncorrected speckles would diffract their light in a more restricted area. However, all the techniques proposed to apodize a pupil~\\citep{vanderbei03,kasdin05,Guyon05,Pluzhnik06} come with throughput problems or manufacturing limitations.  In Sects.~\\ref{sec : snr} to~\\ref{sec : chrom}, we gave the results of the parametric study of a self-coherent camera~(SCC) used as a focal plane wavefront sensor and associated with a perfect coronagraph and a deformable mirror. Several points do not seem to be critical for the technique: reference flux~(Sect.~\\ref{sec : snr}), error on the exact position of the reference image, and diameter of the reference pupil~(Sect.~\\ref{subsec : gamma}). On the contrary, two points are more critical: \\begin{itemize} \\item {\\bf optical path difference between the reference and the image channels.} We have to know and control this~OPD with an accuracy of about~$\\lambda_0/6$. If we associate the SCC with a coronagraph using a Lyot stop as described in Sect.~\\ref{sec : sccrealcoro}, the hardwar optical path difference is always zero and we only have to control the end of the setup. Only common optics are used in this setup. We could put the device in a closed box to avoid differential air variations. We plan to check for the level of the variations in the optical path difference in the device of~Fig.~\\ref{fig : sccfqpm} in a laboratory experiment. \\item {\\bf chromatism.} As shown in Sect.~\\ref{sec : chrom}, chromatism is the most critical point in the technique. The main consequence is the reduction of the corrected area, in other words, the field of view of the image as shown in~Fig.~\\ref{fig : fieldlimit}. The uncorrected speckles spread light and limit the contrast of the detection. We are studying software and hardware solutions to minimize that effect. For example, for the former we could develop a more sophisticated wavefront estimator using a~$\\chi^2$ minimization with regularization terms to account for the spectral dispersion of the speckles~($\\lambda_0/D$) and of the fringes~($\\xi_0/D$). Hardware solutions are certainly more appropriate. We presented the Wynne compensator in Sect.~\\ref{subsec : wynne}. According to numerical simulations, it would enable working with a classical bandpass~($\\sim15\\%$, $R_\\lambda\\simeq6$) in the visible light with a~$64$x$64$~DM. We will test such a Wynne compensator in a laboratory experiment. A second hardware solution to overcome the chromatism limitation could be the association of the~SCC with an integral field spectrometer at modest spectral resolution~($R_\\lambda\\simeq100-150$). This solution is very attractive because it would directly provide a companion spectra, but no work has been done on it yet. \\end{itemize}  We showed in Sect.~\\ref{sec : amplitude} that the self-coherent camera can estimate for both phase and amplitude aberrations. The impact of the amplitude aberrations is quite critical: their level must be smaller than~$1/1000$ to reach a~$5\\,\\sigma$ detection of~$10^{-9}$ at~$5\\,\\lambda_0/D$ with an~$f^{-3}$ power spectral density. That limitation is not intrinsic to the~SCC and would limit any high-contrast imaging system.  Section~\\ref{sec : sccrealcoro} presented a very simple and robust design that associates the self-coherent camera with any coronagraph which uses a Lyot stop. The most interesting point is that the optical path difference between the two channels is constant per construction. In Sect.~\\ref{subsec : sccfqpm}, we studied in detail the case of the association of the~SCC with a~FQPM coronagraph~(SCC-FQPM) and we showed in Sect.~\\ref{subsubsec : sccfqpm_speed} that the performance is very attractive and comparable to the case of a perfect coronagraph that is presented in~\\citet{Galicher08}. Detections of~$2\\,10^{-10}$ earths, $8\\,10^{-10}$ super-earth, and~$10^{-9}$ Jupiter under realistic assumptions are numerically demonstrated for an~SCC-FQPM in polychromatic light~($R_\\lambda=5$) using a Wynne compensator~(reducing the effective bandwidth to~$R_{\\lambda\\,\\mathrm{eff}}=150$) in $\\sim7$\\,h$25$\\,min with a~$4$\\,m space telescope.  The next steps are laboratory demonstrations of both~SCC capabilities: focal-plane wavefront estimation~(step~A, Sect.~\\ref{subsec : stepa}) and companion estimation by differential imaging~(step~B, Sect.~\\ref{subsec : stepb}). We will also attempt to overcome the poor estimate of the astigmatism in the~FQPM transition direction. New algorithms have already been developed for using a~DM interaction matrix including the impact of the whole instrument: DM, coronagraph, and~SCC. It will be tested in a laboratory experiment very soon. We also study the~SCC association with other coronagraphs like an annular groove phase mask~\\citep{Mawet05}.    \\vspace{1cm} We thank R\\'{e}mi~Soummer for private communications about his paper ``Fast computation of Lyot-style coronagraph propagation''\\citep{Soummer07b}, which was very useful for simulating the polychromatic images and the different values of the~$\\gamma$ parameter."
    },
    "0911/0911.3844_arXiv.txt": {
        "abstract": "By comparing near-infrared spectra with atmospheric models, we infer the effective temperature, surface gravity, projected rotational velocity, and radial velocity for 21 very-low-mass stars and brown dwarfs. The unique sample consists of two sequences in spectral type from M6--M9, one of 5--10~Myr objects and one of $>$1~Gyr field objects. A third sequence is comprised of only $\\sim$M6 objects with ages ranging from $<$1~Myr to $>$1~Gyr. Spectra were obtained in the $J$ band at medium (R$\\sim$2,000) and high (R$\\sim$20,000) resolutions with NIRSPEC on the Keck~II telescope. Synthetic spectra were generated from atmospheric structures calculated with the {\\tt PHOENIX} model atmosphere code. Using multi-dimensional least-squares fitting and Monte Carlo routines we determine the best-fit model parameters for each observed spectrum and note which spectral regions provide consistent results. We identify successes in the reproduction of observed features by atmospheric models, including pressure-broadened K~{\\sc i} lines, and investigate deficiencies in the models, particularly missing FeH opacity, that will need to be addressed in order to extend our analysis to cooler objects. The precision that can be obtained for each parameter using medium- and high- resolution near-infrared spectra is estimated and the implications for future studies of very low mass stars and brown dwarfs are discussed. ",
        "introduction": "Determining the physical properties of brown dwarfs and very low-mass stars is important for our understanding of a broad range of topics including star and planet formation, circumstellar disks, dust formation in cool atmospheres, and the initial mass function. The direct measurement of mass and/or radius for a brown dwarf or very low-mass stars is possible only for certain binary systems, which are rare{\\footnote{An archive of very-low-mass binaries is maintained by N.~Siegler (\\url{http://vlmbinaries.org/); only a small fraction of these are double-lined spectroscopic binaries or visual binaries suitable for astrometric monitoring.}}. Determining other physical properties like effective temperature and surface gravity from bolometric luminosity estimates requires a precise distance measurement and assumptions about age (to determine effective temperature) plus mass and radius (to infer surface gravity from evolutionary models). Although radius is not expected to change very much for objects older than a few 100~Myr, reliance on evolutionary models is problematic as they are poorly calibrated at young ages. Commonly used conversions from spectral type to effective temperature (e.g. \\citealt{Golimowski04} for field objects and \\citealt{Luhman03} for young M dwarfs) rely on model isochrones and monotonic relationship between spectral type and effective temperature, which might not accurately represent the complicated formation and evolution of very-low-mass objects \\citep[e.g.][]{Stassun06}. Evolutionary models also require atmosphere models to provide boundary conditions. Thus, synthetic spectra from such atmosphere models potentially offer a more direct method of inferring physical properties by comparison with observed spectra.  The NIRSPEC Brown Dwarf Spectroscopic Survey (BDSS, \\citealt{McLean03,McLean07}) is a large sample of high-quality near-infrared spectra obtained with the Keck~II 10-meter telescope. One of the goals of the BDSS is to provide as direct a measurement of effective temperature, surface gravity, and metallicity as is possible for non-eclipsing objects. In principle this can be accomplished by comparing observed infrared spectra with synthetic spectra from atmosphere models, without relying on spectral type, distance measurements, or estimates of radius, mass, and age. Other authors have applied this promising but as yet imperfect technique to brown dwarfs. The primary difficulty of applying this method to low-mass stars and brown dwarfs is the complexity of the spectra to be modeled, leading to imposed limitations on resolution, wavelength range, and/or sample size. For example, \\citet{Mohanty04a,Mohanty04b} used narrow wavelength ranges at high resolution and \\citet{Cushing08} used a broad wavelength range at low resolution. \\citet{Reiners07} combined high-resolution spectra and broad wavelength coverage for a small sample of objects.  Many uncertainties remain in the determination of physical properties from the comparison of observed and synthetic spectra so it is important to establish the utility of this method using a larger sample of objects and a broader wavelength range. We accomplish this by drawing on medium- (R$\\sim$2,000) and high-resolution (R$\\sim$20,000) $J$-band spectra from the extensive BDSS data base{\\footnote{A public archive is maintained by I.~S.~McLean and collaborators (\\url{http://www.astro.ucla.edu/~mclean/BDSSarchive/)}}. Our goal is to show the extent to which synthetic spectra can be used to derive physical properties (effective temperature [T$_{eff}$], surface gravity [log(g)], projected rotational velocity [$v$sin$i$], and radial velocity [RV)]) rather than the measurement and analysis of trends in the observational data. This paper focuses on field ($\\ge$~1~Gyr) and young (1 to $<$100~Myr) M dwarfs, the majority of which are predicted to be substellar. Future papers will deal with a larger sample of young M dwarfs, young and field L dwarfs, and finally T dwarfs.  Section~2 describes the target selection, NIRSPEC observations, and data reduction. Section~3 outlines the calculation of atmospheric structures and synthetic spectra with the {\\tt PHOENIX} model atmosphere code and our novel spectral fitting method. Section~4 presents the analysis of the spectral fits, including an overview of the results for each resolution and wavelength range and for the inferred physical properties. Our results are discussed in \\S~5. Section~6 summarizes our conclusions and the implications of this work for future studies of brown dwarfs. ",
        "conclusions": "We have estimated the effective temperature, surface gravity, absolute radial velocity and projected rotational velocity for a sample of young brown dwarfs and field M dwarfs. Comparison of high-resolution synthetic and observed spectra of young objects provide more accurate determination of surface gravity than from lower resolution observations. However, the small wavelength range and lack of fidelity of FeH line strength in the atmosphere models result in anomalously high temperatures for late-M type objects. Despite the mismatch in line strengths, radial velocity is well-determined for all but the coolest field objects considered in this analysis. Additionally, the richness of lines at high resolution breaks the degeneracy between sources of line-broadening so that $v$sin$i$ is also well-determined for most objects. The principal results of the paper are as follows:  (1) Adopted (weighted-mean) effective temperatures from high-resolution spectral fits for objects M7 and earlier fall within or very close to previously published estimates and results from medium-resolution spectral fits. For objects M8 and later, temperatures from medium-resolution fits are similar to previously published values and temperatures from high-resolution fits are systematically too hot.  (2) Adopted (weighted-mean) surface gravities from high-resolution fits neatly separate objects with T$_{eff}$ from 1800--3000~K according to age, with 1--10~Myr objects lying broadly between log(g)=3.5 and 4.3, while objects from 0.5--5~Gyr are in the log(g)=5.2--5.4 range. Surface gravities from medium-resolution fits are generally similar to the adopted values except for M8--M9 young objects, which are systematically low, and field objects, which are systematically high.  (3) The age sequence of M6 objects yields an average temperature of 2880~$\\pm$~60~K.  (4) For the Upper Scorpius objects the mean surface gravity is log(g)=3.87, which is very close to the value of log(g)=3.91 predicted by the DUSTY00 evolutionary model for 5~Myr objects between 2500 and 3000~K.  (5) FeH features are not yet properly reproduced by atmosphere models, and this impacts fitting spectra to high-resolution observations. The recent \\citet{Dulick03} line lists provide better correspondence for individual features, but the lines are not strong enough in the models. Improving FeH in the atmosphere models is urgently needed.  (6) Physical properties of brown dwarfs are likely best measured by comparing observed and synthetic spectra using a combination of medium- and high-resolution spectra simultaneously.   This work represents the first in a series of analysis of medium- and high-resolution observed and synthetic spectra of brown dwarfs. Future work will extend the study to later spectral types, cooler atmosphere models (where T$_{eff}$ and log(g) might become more intertwined, e.g. \\citealt{Burgasser06}), and different dust treatments (i.e. the {\\tt PHOENIX}-{\\it cond} models). Further analysis will also explore effects of metallicity on atmospheric structure and resultant spectra. While the focus thus far is on the large database of high-resolution spectra provided by the BDSS, some combination of high resolution and broad spectral coverage will likely prove ideal for accurately inferring the physical properties of brown dwarfs using atmosphere models. Therefore, we are developing a procedure to combine all available observed spectral data for simultaneous model fits."
    },
    "0911/0911.4463_arXiv.txt": {
        "abstract": "We propose a model of asymmetric dark matter (DM) where the dark sector is an identical copy of both forces and matter of the standard model (SM) as in the mirror universe models discussed in literature. In addition to being connected by gravity, the SM and DM sectors are also connected at high temperature by a common set of heavy right-handed Majorana neutrinos via their Yukawa couplings to leptons and Higgs bosons. The lightest nucleon in the dark (mirror) sector is a candidate for dark matter. The out of equilibrium decay of right-handed neutrino produces equal lepton asymmetry in both sectors via resonant leptogenesis which then get converted to baryonic and dark baryonic matter. The dark baryon asymmetry due to higher dark nucleon masses leads to higher dark matter density compared to the familiar baryon density that is observed. The standard model neutrinos in this case acquire masses from the inverse seesaw mechanism. A kinetic mixing between the $U(1)$ gauge fields of the two sectors is introduced to guarantee the success of Big-Bang Nucleosynthesis. ",
        "introduction": "It now appears established that dark matter accounts for about one quarter of the energy density, $\\Omega$, of the universe and plays an essential role in the formation of large scale structure in it. The identity of dark matter, however, remains unknown since all the particles in the successful standard model can be ruled out as candidates. What the dark matter particles are, how they interact with visible matter and how their relic abundance originates, constitute some of the fundamental mysteries of particle physics and cosmology today. Adding to this puzzle is the observation that baryon contribution to $\\Omega$ is only about a fifth (i.e., of the same order) of the dark matter contribution. This raises the question: could the two have a common origin?  The most popular class of candidates for dark matter are the stable weakly interacting massive particles (WIMPs), which arise in many well-motivated TeV scale extensions of the standard model. Their stability is guaranteed by some symmetry. While the WIMP dark matter does not decay due to the symmetry, pairs of them can annihilate and their relic density is determined by the freeze-out of the annihilation from equilibrium to the SM particles. The fact that their observed relic density can be naturally explained by the weak scale annihilation cross section makes these models quite appealing. However in most models, the matter and dark matter contributions to $\\Omega$ are unrelated.  Coming to the question of the symmetry that ensures the stability of WIMPs, in most models one uses a $Z_2$ symmetry (e.g., R-parity in supersymmetry or KK-parity in the case of extra dimension models). On the other hand, one could quite easily imagine that the stability of the WIMP is guaranteed by a continuous $U(1)$ symmetry similar to baryon number (call it ``dark baryon number''). In such a case, observed dark matter density would represent an asymmetry between dark matter and anti-dark matter densities exactly as the case for the observed asymmetry between familiar matter and anti-matter. If these two asymmetries could arise from a common mechanism, it would be a major step towards understanding why their contributions to $\\Omega$ are of the same order. It is the goal of this paper to propose such a scenario.  The baryon-anti-baryon asymmetry in the universe (BAU) is of order $10^{-10}$ and can arise from the laws of microphysics if the three conditions proposed by Sakharov~\\cite{Sakharov:1967dj} are satisfied, namely baryon number violation, CP violation and out-of-equilibrium in the early universe. The SM, however, fails to realize the third condition since it requires that the Higgs mass must be less than 40 GeV; even if this condition was satisfied, with three generations of fermions, it cannot explain why the asymmetry is as large as the one observed. An elegant way to generate the BAU is through leptogenesis~\\cite{Fukugita:1986hr} in the framework of seesaw mechanism~\\cite{seesaw} which naturally explains the smallness of the observed neutrino masses. In these models, lepton asymmetry is generated through the out-of-equilibrium decay of the very heavy right-handed (RH)  neutrinos which then get converted to baryon asymmetry through the non-perturbative $B+L$ violating electroweak sphaleron process~\\cite{Kuzmin:1985mm}.  The appealing mechanism of baryogenesis via leptogenesis combined with the fact that the relic abundances of baryons and dark matter are of the same order of magnitude inspire us to think whether genesis of dark matter could also have its origin in a manner similar to leptogenesis. The mirror universe models discussed in the literature~\\cite{mirror} appear to be a natural setting for this. In the mirror models, the universe has two kinds of matter and forces: one consisting of a standard model sector with forces and matter that we are familiar with, such as quarks, leptons, W, Z, etc., and a parallel sector which is an exact replica (e.g., consisting of mirror duplicates of quarks and leptons, W, Z, $\\gamma$, gluons and the Higgs boson) called the mirror sector~\\cite{mirror}. Forces (except for gravity) in one sector do not affect the matter in the other sector. In what follows we will denote the mirror particles by a prime on a symbol. The two sectors communicate with each other by gravity and possibly some SM singlet interactions which are very weak at the current age of the Universe. There is a dark baryon and lepton number in the mirror sector which is the exact analog of the familiar baryon and lepton number. The main hypothesis of this paper is that the same leptogenesis mechanism that could be producing matter-anti-matter asymmetry, is also producing asymmetry of dark matter-anti-dark-matter. This then links the dark matter energy density of the Universe to that contributed by matter making them of the same order, providing a resolution of the puzzle stated in the beginning~\\cite{adm}\\footnote[1]{For earlier suggestion that mirror baryons constitute dark matter of the Universe, see~\\cite{mdm}. Our model is, however, very different from these models in many respects}.  A key ingredient in our attempt to connect the matter asymmetry to dark matter asymmetry is the assumption that the visible and the mirror sectors talk to each other not only through gravity but also through a common set of three right-handed neutrinos coupled to leptons and Higgs fields in each sector through Yukawa couplings\\cite{zurab}, as shown in Fig.~\\ref{po}. Since the RH neutrinos are standard model singlets, this is consistent with gauge invariance. Also mirror symmetry makes the $N\\ell H$ couplings on both sides equal. The out-of-equilibrium decays of right-handed neutrinos in the early universe can then produce lepton number asymmetries in both sectors, which are then transferred to baryon and dark baryon numbers through the sphaleron processes in each sector. If one imposes exact mirror symmetry on the theory, the primordial lepton asymmetries generated in each sector are equal and after sphaleron interaction produce the same number density for baryons and dark baryons in the early universe. Since we expect the symmetry breaking pattern in both sectors to be different for the model to be consistent with cosmology (see below), the resulting energy density contributions can be different and in the ratio $\\Omega_B : \\Omega_{DM} \\approx 1:5$ if we require that mass of the dark baryons is five times the mass of the familiar SM baryons. This mass difference can arise from the difference in the scales of two $SU(3)_c$ strong interactions ($\\Lambda_{QCD}, \\Lambda^\\prime_{QCD}$). It turns out that this difference depends on the ratio of two electroweak scales $v_{\\rm wk}$ and $v'_{\\rm wk}$, with $v_{\\rm wk}:v'_{\\rm wk}\\sim 1:10^{3}$ giving the required difference between $\\Omega_{B}$ and $\\Omega_{DM}$.  The spectra of the SM neutrinos $\\nu_\\alpha$ and the dark neutrinos $\\nu^\\prime_{\\alpha}$ are determined by the inverse seesaw~\\cite{inverse} and type-I~\\cite{seesaw} seesaw mechanisms, respectively, with the mixing between the $\\nu_\\alpha$ and $\\nu'_\\alpha$ given by the ratio $v_{\\rm wk}/v'_{\\rm wk}$. The dark neutrinos can decay into the SM particles due to this mixing.  In addition to the common set right-handed neutrinos, the SM sector and the mirror sector can also be connected through Higgs interaction and the kinetic mixing between the $U(1)$ gauge bosons consistent with gauge invariance (as shown in Fig.~\\ref{po}). The photon sector mixing is necessary for the model to be consistent with Big Bang Nucleosynthesis (BBN). In this work we assume the photon in the mirror sector acquires a mass around 50 MeV through the spontaneous symmetry breaking so that the mirror-electrically charged particles which are heavier, pair annihilates into it before the BBN, and the mirror photon itself decays to the electron-positron pair through the kinetic mixing. To generate this small mass two Higgs doublets are needed in the mirror sector. The kinetic mixing between the photon and mirror photon determines the lifetime of the mirror photon, so that success of BBN puts a lower bound on this mixing. Furthermore, it also determines the interactions between the dark matter particles and the nucleons; therefore, using the current constraint from the direct detection experiments produces an upper bound on this mixing.  The common right-handed neutrinos generate a mixing between SM neutrinos and mirror neutrinos. The exchange of mirror neutrinos contribute to the neutrinoless double beta decay process ($0\\nu\\beta\\beta$), which also puts some mild constraints on the model.  \\begin{figure}[htb] \\begin{center} \\includegraphics[width=10cm]{portal.eps} \\caption{The connection between the standard model and the dark sector.}\\label{po} \\end{center} \\end{figure} ",
        "conclusions": "In this work, we have proposed that the dark matter of the Universe be identified with the lightest baryon of a possible mirror duplicate of the standard model with the only difference between the two sectors being in the symmetry breaking patterns. Prior to spontaneous symmetry breaking, this model has no free parameters due to mirror symmetry. The lightest dark nucleon is stable due to the mirror analog of baryon number and becomes the dark matter. It is an asymmetric dark matter with its anti-mirror baryon part suppressed in a manner analogous to the matter-anti-matter asymmetry in the standard model sector. The introduction of a common set of right-handed neutrinos connecting the two sectors allows a common mechanism for the genesis of the matter-anti-matter asymmetry in both sectors thereby helping us to understand why the dark matter and normal baryon contribution to the energy density of the Universe are not too different from each other. One only has to make the assumption that the dark nucleon is five times heavier than the familiar nucleon -- an assumption that is easily understood if the electroweak scales in the two sectors are different. We show that this model can be consistent with the constraints of BBN and neutrinoless double beta decay. The mirror photon in our model is massive but mixes with the normal photon to avoid the BBN constraints."
    },
    "0911/0911.1888_arXiv.txt": {
        "abstract": "Recent analysis of the Milky Way's satellite galaxies reveals that these objects share a common central mass density, even though their luminosities range over  five orders of magnitude. This observation can be understood in the context of galaxy formation theory by quantifying the factors which restrict the central mass density to a small range. One limit is set by the maximum mass that can collapse into a given region by the hierarchical growth of structure in the standard cold dark matter cosmology.  Another limit comes from the natural thresholds which exist for gas to be able to cool and form a galaxy. The wide range of luminosities in these satellites reflect the effects of supernova feedback on the fraction of cooled baryons which are retained. \\vspace{0.5cm} ",
        "introduction": "The satellite galaxies of the Milky Way are providing a valuable guide to our understanding of cosmology. Though they constitute only a single statistical sample,  the standard cold dark matter cosmological model (\\lcdm) must at least be able to account for their abundance and properties. Attaining such consistency is not proving to be straightforward.  Initial friction between observation and theory concerned the number of observed satellites, which is greatly exceeded by the number of subhalos that \\lcdm~predicts in a host halo of the Milky Way's size \\cite[and over 1000 citing articles]{Klypin99,Moore99}. A proposed resolution to this was that photoionisation has suppressed the formation of small galaxies in these subhalos \\cite{Bullock01,Benson02,Somerville02}.  Observations of satellite galaxies also indicate that these systems boast extremely high mass to light ratios, much greater than are deduced for more luminous galaxies. Mass does not appear to follow light, as might be naively expected. This has been recently quantified by \\scite{Strigari08}, leading to a particularly intriguing result.  The central mass of each satellite, as contained within a certain chosen radius (300 pc), appears to be the same for all the satellites, remaining approximately constant over seven orders of magnitude in luminosity.  A suggested explanation for this result, by the authors themselves, is that this common central density may be the hallmark of a common formation time, a suggestion which appeals to the results of N-body simulations. Navarro et al. (1996,1997) found that the central density of halos was proportional to the cosmic background density at the very earliest stages in their assembly,  so halos carry a permanent imprint of their formation epoch. \\pagebreak  Confirmation and extension of these ideas has come from the combination of N-body simulations with semi-analytic models to predict the properties of galaxy populations \\cite{Cooper09,LiHelmi09,Maccio09,Okamoto09}. All these calculations identify the early reionisation of gas and minimum temperature for atomic cooling ($10^4$K) as the features which lead to common central mass in satellites.  From N-body simulations of  three different Milky-Way sized host halos, \\scite{Kravtsov09} deduces the central masses of subhalos from their accretion epoch and the detailed results of higher-resolution simulations. Central mass is found to be a power-law in subhalo virial mass, $M_{\\rm v}$. The empirical result of constant central mass thus corresponds, in turn, to a power-law relationship between virial mass and luminosity.  These connections between this singular observational result and the physics of galaxy formation can be understood with particular clarity by investigating the expected distribution of halos in the central mass -- total mass plane. This distribution is investigated in \\S\\ref{Distribution}, paying particular attention to the physical explanation behind each feature and generalising the argument to different measurement radii and host halo masses.  Section \\ref{Luminosities} takes the argument further to understand how the underlying subhalo masses might map on to the luminosities of the galaxies contained. Rather than demonstrating consistency between these particular results and a particular galaxy formation model, the relative influence of infall, cooling and supernova feedback are explained at a more general level. It is then shown how these processes can lead to the wide range of luminosities observed. ",
        "conclusions": "We have considered the issue of the present-day central density of satellite halos in the context of cold dark matter cosmology theory. By generating satellite populations for a range of host halo masses, the pattern of their distribution in the central mass - total mass plane was understood in terms of physical thresholds.  The central density must be high enough for the satellite galaxy to survive the effects of reionisation, and/or for the temperature to be high enough to allow atomic cooling. It must also be lower than the natural maximum density which can arise through gravitational collapse from \\lcdm~ initial conditions.  These boundaries restrict the population of satellite hosts to a very small range in central mass, but {\\em not} in total mass. The explanation for this effect leads to a simple analytic approximation for the central mass (\\ref{peak}) given in terms of the general values of host halo mass and enclosure radius. The merger tree calculations support of this conclusion over about an order of magnitude in both factors.  While half the satellites occupy just a tiny portion of the total range in central mass, this same half spans three orders of magnitude in pre-infall virial mass. Further analysis then helps understand how these virial masses might map onto stellar masses, which will explain the range in luminosities observed by \\scite{Strigari08}.  By reference to the merger tree formation histories (Fig.\\ref{formation}) it can be seen that halos may only be able to cool effectively for {\\em part} of their histories. Due to relatively long free-fall times, only a fraction of the baryons that are initially available in these halos will infall into the galaxies. This is one reason why mass-to-light ratios might be high for smaller systems.  A second reason is reheating. In the case of satellites, where hot halo gas can be stripped upon accretion, the traditional argument (that a set fraction of supernova energy is carried out of the galaxy in the reheated gas) is shown to lead to a simple relationship between final stellar mass and halo mass.  This makes it clear that even quite modest energy conversion will dramatically decrease the fraction of the baryons in satellite halos which can be retained in their galaxies.  For high, but reasonable values of the conversion efficiency, the galaxies which occupy a range in central mass of just  $0.1 \\log_{10}(M_r/\\Msun)$ are distributed across six orders of magnitude in $M_{\\rm baryons}/\\Msun$. This is in keeping with the wide range of luminosities observed for the Milky Way's satellites."
    },
    "0911/0911.3419_arXiv.txt": {
        "abstract": "Long-period transiting exoplanets provide an opportunity to study the mass-radius relation and internal structure of extrasolar planets. Their studies grant insights into planetary evolution akin to the Solar System planets, which, in contrast to hot Jupiters, are not constantly exposed to the intense radiation of their parent stars. Observations of secondary eclipses allow investigations of exoplanet temperatures and large-scale exo-atmospheric properties. In this short paper, we elaborate on, and calculate, probabilities of secondary eclipses for given orbital parameters, both in the presence and absence of detected primary transits, and tabulate these values for the forty planets with the highest primary transit probabilities. ",
        "introduction": "\\label{introduction}  Secondary eclipses of exoplanets provide unique insight into their astrophysical properties such as surface temperatures, atmospheric properties, and efficiency of energy redistribution. In \\citet{kvb08, kvb09}, we demonstrate that the probability of detecting transits or eclipses among known radial-velocity (RV) planets is sensitively dependent on the values of orbital eccentricity $e$ and argument of periastron $\\omega$. with some combinations' of $e$ and $\\omega$ making transit/eclipse searches among long-period planets viable. Though it is feasible to detect planetary eclipses from space \\citep[e.g., ][]{ldl09} and even from the ground \\citep[e.g., ][]{slm09}, the difference in signal-to-noise ratio between transits and eclipses makes detections of the former much more straightforward. In this paper, we calculate the probability of planetary eclipses with or without the knowledge of the existence of a primary transit. ",
        "conclusions": "\\label{discussion}  Fig. \\ref{newgep} plots the {\\it a priori} values (open circles) and conditional lower limits (crosses) of $P_e$ for 203 planets as functions of $e$ and $\\omega$ values tabulated in \\citet{b06}. The {\\it left} panel clearly shows that, for low values of $e$, the existence of a transit practically guarantees an observable eclipse, whereas for higher eccentricities, this is not the case due to the correspondingly weaker constraint on inclination angle imposed by an existing transit \\citep[cf. eq. 8 in ][]{kvb09}. The {\\it right} panel demonstrates that, for $\\omega \\sim 3\\pi / 2$, a detected transit ensures the existence of an observable eclipse , whereas for $\\omega \\sim \\pi / 2$, this constraint is much weaker \\citep[cf. fig. 1 in ][]{kvb09}. Thus, the presence of a planetary transit greatly affects the likelihood of existence of a secondary eclipse. Both the {\\it a priori} value and the conditional lower limit of $P_e$ remain sensitively dependent on the combination of $e$ and $\\omega$.  Table \\ref{table1} shows the forty extrasolar planets with the (currently) highest transit probabilities, their orbital elements $P$, $e$, and $\\omega$, and the explicit values for eclipse probabilities ($P_e$: {\\it a priori} value; $P_e^{\\prime}$: conditional lower limit). For instance, HD 118203 b has $P_e = 7.05\\%$, but {\\it if a primary transit is observed} ($P_t = 9.11\\%$), then the existence of a secondary eclipse is almost certain ($P^\\prime_e\\geq94.65\\%$). See table 1 in \\citet{kvb09} for the equivalent inverse case of $P_t = f(P_e)$.   \\begin{figure*} \\begin{center} \\begin{tabular}{cc} \\includegraphics[angle=270,width=6.6cm]{o_vonBraun_f1.ps} & \\includegraphics[angle=270,width=6.6cm]{o_vonBraun_f2.ps} \\\\ \\end{tabular} \\end{center} \\caption{{\\small The {\\it a priori} values of $P_e$ (Eq. \\ref{eclprob}; open circles) and conditional lower limits of $P_e$ (Eq. \\ref{newprob}; crosses) for 203 planets from the \\citet{b06} catalog, plotted as a function of $e$ and $\\omega$. For purpose of comparison, we assume solar and Jupiter masses and radii for all systems. } } \\label{newgep} \\end{figure*}    \\begin{table}[!ht] \\caption{$P_t$, {\\it a priori} $P_e$, and conditional (revised) $P^\\prime_e$ for the 40 exoplanets from \\citet{b06} with highest $P_t$ values. See \\S\\ref{discussion} \\label{table1}} \\smallskip \\begin{center} {\\small \\begin{tabular}{lcccccc} \\tableline \\noalign{\\smallskip} Planet & $P$ (days) & $e$ & $\\omega (\\deg)$ & $P_t$ (\\%) & $P_e$ (\\%) & $P^{\\prime}_e$ (\\%)\\\\ \\noalign{\\smallskip} \\tableline \\noalign{\\smallskip} HD 41004 B b &      1.33 &  0.08 &  178.50 &   25.81 &   25.70 &  100.00 \\\\ HD 86081 b &      2.14 &  0.01 &  251.00 &   16.93 &   17.19 &  100.00 \\\\ GJ 436 b &      2.64 &  0.16 &  339.00 &   16.50 &   18.50 &  100.00 \\\\ 55 Cnc e &      2.80 &  0.26 &  157.00 &   15.17 &   12.33 &   99.36 \\\\ GJ 674 b &      4.69 &  0.20 &  143.00 &   14.93 &   11.72 &   95.96 \\\\ HD 46375 b &      3.02 &  0.06 &  114.00 &   13.58 &   12.11 &  100.00 \\\\ HD 187123 b &      3.10 &  0.01 &    5.03 &   12.81 &   12.78 &  100.00 \\\\ HD 83443 b &      2.99 &  0.01 &  345.00 &   12.77 &   12.82 &  100.00 \\\\ HD 179949 b &      3.09 &  0.02 &  192.00 &   12.74 &   12.86 &  100.00 \\\\ HD 73256 b &      2.55 &  0.03 &  337.00 &   12.66 &   12.95 &  100.00 \\\\ HD 102195 b &      4.12 &  0.06 &  109.90 &   10.85 &    9.69 &  100.00 \\\\ HD 76700 b &      3.97 &  0.09 &   29.90 &   10.82 &    9.84 &  100.00 \\\\ HD 75289 b &      3.51 &  0.03 &  141.00 &   10.47 &   10.03 &  100.00 \\\\ 51 Peg b &      4.23 &  0.01 &   58.00 &   10.35 &   10.12 &  100.00 \\\\ $\\tau$ Boo b &      3.31 &  0.02 &  188.00 &   10.21 &   10.27 &  100.00 \\\\ HD 88133 b &      3.42 &  0.13 &  349.00 &   10.16 &   10.68 &  100.00 \\\\ BD -10 3166 b &      3.49 &  0.02 &  334.00 &   10.15 &   10.32 &  100.00 \\\\ HAT-P-2 b &      5.63 &  0.51 &  184.60 &    9.44 &   10.24 &  100.00 \\\\ HD 17156 b &     21.20 &  0.67 &  121.00 &    9.14 &    2.47 &   33.05 \\\\ HD 118203 b &      6.13 &  0.31 &  155.70 &    9.11 &    7.05 &   94.65 \\\\ $\\upsilon$ And d &      4.62 &  0.02 &   57.60 &    8.69 &    8.38 &  100.00 \\\\ HD 68988 b &      6.28 &  0.15 &   40.00 &    8.20 &    6.76 &  100.00 \\\\ HIP 14810 b &      6.67 &  0.15 &  160.00 &    7.86 &    7.10 &  100.00 \\\\ HD 162020 b &      8.43 &  0.28 &   28.40 &    7.84 &    6.02 &   93.77 \\\\ HD 217107 b &      7.13 &  0.13 &   20.00 &    7.76 &    7.11 &  100.00 \\\\ HD 168746 b &      6.40 &  0.11 &   17.40 &    7.63 &    7.16 &  100.00 \\\\ HD 185269 b &      6.84 &  0.30 &  172.00 &    7.30 &    6.72 &  100.00 \\\\ HD 49674 b &      4.94 &  0.29 &  283.00 &    6.68 &   11.94 &  100.00 \\\\ HD 69830 b &      8.67 &  0.10 &  340.00 &    6.24 &    6.68 &  100.00 \\\\ HD 130322 b &     10.71 &  0.03 &  149.00 &    5.76 &    5.62 &  100.00 \\\\ HD 33283 b &     18.18 &  0.48 &  155.80 &    5.31 &    3.56 &   82.03 \\\\ HD 38529 b &     14.31 &  0.25 &  100.00 &    5.21 &    3.17 &   74.40 \\\\ HD 55 Cnc b &     14.65 &  0.02 &  164.00 &    4.67 &    4.63 &  100.00 \\\\ HD 13445 b &     15.76 &  0.04 &  269.00 &    4.47 &    4.85 &  100.00 \\\\ HD 27894 b &     17.99 &  0.05 &  132.90 &    4.43 &    4.12 &  100.00 \\\\ HD 108147 b &     10.90 &  0.53 &  308.00 &    4.14 &   10.09 &  100.00 \\\\ HD 6434 b &     22.00 &  0.17 &  156.00 &    4.02 &    3.50 &  100.00 \\\\ HD 190360 c &     17.11 &  0.00 &  168.00 &    3.94 &    3.93 &  100.00 \\\\ HD 20782 b &    585.86 &  0.93 &  147.00 &    3.92 &    1.29 &   40.33 \\\\ GJ 876 c &     30.34 &  0.22 &  198.30 &    3.85 &    4.44 &  100.00 \\\\ HD 99492 b &     17.04 &  0.25 &  219.00 &    3.83 &    5.29 &  100.00 \\\\ \\noalign{\\smallskip} \\tableline \\end{tabular} } \\end{center} \\end{table}"
    },
    "0911/0911.1136_arXiv.txt": {
        "abstract": "Motivated by detections of hypervelocity stars that may originate from the Galactic Center, we revist the problem of a binary disruption by a passage near a much more massive point mass. The six order of magnitude mass ratio between the Galactic Center black hole and the binary stars allows us to formulate the problem in the restricted parabolic three-body approximation. In this framework, results can be simply rescaled in terms of binary masses, its initial separation and binary-to-black hole mass ratio. Consequently, an advantage over the full three-body calculation is that a much smaller set of simulations is needed to explore the relevant parameter space. Contrary to previous claims, we show that, upon binary disruption, the lighter star does not remain preferentially bound to the black hole. In fact, it is ejected exactly in 50\\% of the cases.  Nonetheless, lighter objects have higher ejection velocities, since the energy distribution is independent of mass. Focusing on the planar case, we provide the probability distributions for disruption of circular binaries and for the ejection energy. We show that even binaries that penetrate deeply into the tidal sphere of the black hole are not doomed to disruption, but survive in $20\\%$ of the cases. Nor do these deep encounters produce the highest ejection energies, which are instead obtained for binaries arriving to $0.1-0.5$ of the tidal radius in a prograde orbit. Interestingly, such deep-reaching binaries separate widely after penetrating the tidal radius, but always approach each other again on their way out from the black hole. Finally, our analytic method allows us to account for a finite size of the stars and recast the ejection energy in terms of a minimal possible separation.  We find that, for a given minimal separation, the ejection energy is relatively insensitive to the initial binary separation. ",
        "introduction": "Hypervelocity Stars (HVSs) are stars with a velocity exceeding the escape velocity of the Galaxy.  Currently, $16$ such stars have been observed, 15 of which are thought to originate from our Galactic Centre \\citep{BGK+05,BGK+06,BGK+07,BGK+09,HHO+05} and one from the Large Magellanic Cloud \\citep{ENH+05}.  All have been observed with radial velocities between 300~km/s  and 800~km/s and almost all are located over 50~kpc away. For stars whose likely origin is within the Galaxy and taking the galactic potential into account, this translates to velocities of over 1000~km/s from the bulge.  The observational strategy is such that most of the discovered HVSs are faint B stars.  They have escape velocities from the surface of the order of 600~km/s, well below the ejection velocity of 1000~km/s. Thus, the standard mechanisms for producing high velocity runaway stars, such as star-scattering and explosion as a supernova of one component of a binary cannot work. A dynamical interaction with a massive compact object is likely involved.  In this paper, we adopt one of the leading models for the formation of HVSs: the breakup of a binary as it approaches the black hole in the Galactic Center \\citep{Hil88}. Simple analytical arguments can be made to show the potential of this model to explain HVSs.  If the binary of total mass $m$ has separation $a$, then tidal forces from the black hole overcome the binary's mutual gravitational forces at the tidal radius $r_t=a(M/m)^{1/3}$, where $M$ is the mass of the black hole. The relative velocity of the binary components is of order $v_0=(Gm/a)^{1/2}$.  If the binary approached the black hole with negligible energy, its center of mass moves at the tidal radius with velocity of order $v_{\\rm BH}=(GM/r_t)^{1/2}=v_0 (M/m)^{1/3}$ relative to the black hole. There are three ways to estimate the energies of the individual components of the binary, assuming that they arrived with negligible total energy. It is instructive to consider all three: \\begin{itemize} \\item {\\bf Kinetic energy:} Adding or subtracting the relative velocity of the components, $v_0$, to the velocity around the black hole, results in an additional kinetic energy of order $v_0v_{\\rm BH} \\sim v_0^2 (M/m)^{1/3}$. \\item {\\bf Gravitational potential energy:} The displacement of order $a$ in the position of each component of the binary, at a distance of about $r_t$ from the black hole, results in a change in gravitational energy of $GMa/r_t^2 \\sim v_0^2 (M/m)^{1/3}$. \\item {\\bf Work:} The energy of each of the component in the black hole frame, is changing only due to mutual forces between the binary components. The force is of order $Gm/a^2$ and the length, in the black hole frame, over which it acts is $r_{\\rm t}$. Therefore the work is $Gmr_t/a^2 \\sim v_0^2(M/m)^{1/3}$. \\end{itemize} All these estimates provide an energy of order $v_0^2 (M/m)^{1/3}$. If the binary dissolves, one component of the binary stays bound to the black hole, the other escapes, with velocities at infinity\\footnote{We ignore the galactic potential in this paper.  By ``the velocity at infinity\" we mean the velocity of the object once it escaped the gravitational potential of the black hole, but did not yet climb out of the potential of the rest of the galaxy.} of $v_0(M/m)^{1/6}$. The encounter with the black hole therefore allowed for a larger velocity by a factor of $(M/m)^{1/6}$ than the orbital velocity of the binary.  For the parameters of our Galaxy and stars, $(M/m)^{1/6}\\sim 10$, allowing ejections with thousands of kilometers per second.  This is the theoretical framework in which \\cite{Hil88} predicted the existence of HVSs. Later, it was discussed by \\cite{YuT03} and \\cite{GQ03}. After the observational discovery, many more papers on the subject appeared \\citep{GPS05,GL06,GiL06,BKG+06,SHM07,PHA07,KBG+08}, aiming to predict the properties of the ejected and/or captured stars. Other papers, \\cite[e.g.][]{GL07,A09}, focused instead on the fate of binaries that are not dissolved, but that in fact coalesce. The investigations so far have used three-body simulations or analytic methods that relied on results from three-body simulations. Consequently, only a limited set of parameters (e.g. for the binary mass ratio) have been explored.  In this paper, we show  that this problem can be investigated with methods related to those used in the study of asteroids in the Solar System, which exploit the enormous disparity in mass between the bodies involved. Specifically, we can formulate it in terms of a restricted 3 body problem, i.e. the motion of a single massless particle under the influence of external time dependent forces. Our treatment is valid as long as the binary components are closer to each other than to the black hole. Since $r_{\\rm t}=a(M/m)^{1/3}$ our approximation requires $(M/m)^{1/3} \\gg 1$, which is fully satisfied in the case of the black hole at the Galactic Center and a binary of B stars, $(M/m)^{1/3} \\cong 100$. With the advantages of this analytic method, we can reach general conclusions that do not depend on the physical properties of the system, such as masses and binary semi-major axis. Moreover, the orbit integration is faster and more stable, allowing us to handle more easily cases of close encounters between the bodies.   In \\S \\ref{sec:parabolic} we outline the formulation of the 3-body problem in terms of a restricted parabolic problem.  In \\S \\ref{sec:radial} we use the restricted radial problem to describes binaries that penetrate deep into the tidal radius.  The radial problem has a singularity at the time that the binary encounters the black hole, and we use the results of the parabolic problem to pass the singularity and continue to describe the evolution of the binary on its trajectory away from the black hole.  In \\S \\ref{sec:verification} we compare our results with 3-body numerical integrations, and find excellent agreement.  We then use the parabolic and radial formalism to investigate the probability of dissolving a circular, planar binary and to obtain quantitative estimates of the ejection velocities. We outline these results in \\S \\ref{sec:results}. ",
        "conclusions": "The ultimate goal of our work is to statistically caracterize the population of stars originated from tidally disrupted binaries, and compare it with observations of hyper velocity stars.  To this purpose, we derive in this paper the equations of motion and energy for a member of a close binary, which suffers an encounter with a third far more massive body. For the general case, we assume a parabolic orbit for an effective center of mass of the binary (\\S\\ref{sec:parabolic}). This is in contrast with previous works, \\citep[e.g.][]{BKG+06} who considered binaries on hyperbolic or elliptical orbits. However, our assumption is justified since the orbits of binaries that are candidates  to produce hyper velocity stars are very eccentric. The periapsis passage of a very eccentric orbit could be modeled by a parabolic orbit to a good approximation. This is in agreement with the finding of \\cite{BKG+06} that the initial binary velocity  is of little influence on the final outcome. Nevertheless, a binary on a hyperbolic orbit has a total positive energy, and allows for a new disruption channel, in which both stars are ejected. This channel is relatively rare, since the energy of the binary is typically small compared to the typical ejection energies. Moreover, such double ejections could not lead to hypervelocity stars, since the energy is limited to the original-- small-- positive energy of the binary. Our formalism with zero total energy (except the small negative binary binding energy), does not allow for double ejections.  A parabolic trajectory for the binary approaches a radial one for a very close encounter with the massive body. This observation leads us to explicitly adopt a radial orbit with zero energy, in order to follow the limiting case of a deep penetrator (\\S\\ref{sec:radial}). This simpler set of equations allow us to easily trace a close encounter, that otherwise would require high accuracy when calculated with a full three-body code.   Our formalism can be applied quite generally to explore the fate of a binary with arbitrary orbital parameters. However, in this paper, we only focused on results for circular coplanar binaries. The inclination of the binary, as well as its eccentricity, is expected to affect our results quantitatively. We reserve such a study to a forthcoming paper. Nevertheless, we can already reach some conclusions and note quantitative differences with results found in literature.  The main feature of HVSs is of course their unusually high radial velocity. For equal mass stars, $m_1=m_2=m_{1,2}$, the expression $Gm_1m_2/R_{\\rm min}$ is simply $m_{1,2}v_{\\rm esc}^2/4$, where $v_{\\rm esc}$ is the escape velocity from the surface of the stars. Since the maximal ejection energy is $1.6 ~Gm_1m_2/R_{\\rm min}(M/m)^{1/3}$ (Fig.~\\ref{fig:emin_emax}), we derive for the Galactic black hole a corresponding velocity of $0.9 ~v_{\\rm esc}(M/m)^{1/6}\\cong 9 \\,v_{\\rm esc}$. The escape velocity for main sequence stars is about 600-800~km/s in the mass range of $1-10 \\,M_\\odot$. Therefore velocities can be as large as $\\sim 7000$~km/s even for binaries limited to the plane. Scaling for the different masses (and thus $v_{\\rm esc}$) assumed by \\cite{Hil88} and by \\cite{BKG+06}, we find that a velocity of $0.9 ~v_{\\rm esc}(M/m)^{1/6}$ is respectively a factor of $\\sim 1.2$ and $\\sim 1.8$ higher than their maximum value of $4000$ Km/s. Even higher velocities can be acheived when the binary mass ratio is large and the {\\it lighter} star is ejected. Then, the ejected star travels with maximal velocity of $1.3 \\,v_{\\rm esc}(M/m)^{1/6}$, which is around $\\sim 10,000$~km/s for our Galactic center black hole where $M/m \\sim 10^6$.  For a comparison with observations, it is also important to determine the ejection probability. \\cite{BKG+06} find that the probability for disruption as a function of the penetration factor goes roughly as $P_{\\rm ej} \\approx 1-D/2.2$ . The maximum $D$ for disruption is therefore very similar to ours, $D_{\\rm max} \\sim 2.1$ (see Fig.~\\ref{fig:vmean_comp}). In contrast, we remark a quantitatively different behaviour of the probability function for binaries that approach closely the black hole, $D\\rightarrow 0$. Specifically, their $P_{\\rm ej} \\rightarrow 1$ implies that all binaries in this limit are disrupted. This may look intuitively sound since binaries that penetrate so deeply experience very strong tides. However, as we stressed in \\S\\ref{sec:radial}, the stars in such binaries separate, but approach each other again on the way back from the black hole. In the planar case we find that $\\sim 20\\%$ of the deeply penetrating binaries survive (Fig.~\\ref{fig:dis_f}). Even when taking a uniform distribution for the inclination of the binary plane into account, the percentage of survival is still $10\\%$  for close encounters. In the planar case, Figure~\\ref{fig:dis_f} indicates that the disruption probability is, in fact, not monotonic in $D$: for prograde orbits, it peaks at $D=0.15$  and it is almost unity ($98.0-98.8\\%$) for $0.06 < D < 0.3$.  For large mass ratios $m_1 / m_2 > 10$, \\cite{BKG+06} find that the heavier has consistently more chance to be ejected. We showed here that the probability is 50\\%. This result is not limited to zero inclination of eccentricity. This fact together with the rarity of a very massive star somewhat weakens the claim that the star SO-2 was created by a disruption of a binary, in which a $60\\, M_\\odot$ companion was ejected \\citep{GQ03}. This conclusion was based on the observed orbital parameters of SO-2. However, the timescale to significantly change its shor periapsis distance may quite short, and a careful study of the dynamical processes in the inner regions of the Galactic Center is needed to asses it."
    },
    "0911/0911.3902_arXiv.txt": {
        "abstract": "We investigate the potential role of string and monopole-type junctions in the frustration of domain wall networks using a velocity-dependent one-scale model for the characteristic velocity, $v$, and the characteristic length, $L$, of the network. We show that, except for very special network configurations, $v^2 \\lsim (HL)^2 \\lsim (\\rho_\\sigma + \\rho_\\mu)/\\rho_m$ where $H$ is the Hubble parameter and $\\rho_\\sigma$, $\\rho_\\mu$ and $\\rho_m$ are the average density of domain walls, strings and monopole-type junctions. We further show that if domain walls are to provide a significant contribution to the dark energy without generating exceedingly large CMB temperature fluctuations then, at the present time, the network must have a characteristic length $ L_0 \\lsim 10  \\,\\Omega_{\\sigma 0}^{-2/3} \\, {\\rm kpc}$ and a characteristic velocity $v_0 \\lsim 10^{-5} \\, \\Omega_{\\sigma 0}^{-2/3}$ where $\\Omega_{\\sigma 0}=\\rho_{\\sigma 0}/\\rho_{c 0}$ and $\\rho_c$ is the critical density. In order to satisfy these constraints with $\\Omega_{\\sigma 0} \\sim 1$, $\\rho_{m 0}$ would have to be at least $10$ orders of magnitude larger than $\\rho_{\\sigma 0}$, which would be in complete disagreement with observations. This result provides very strong additional support for the conjecture that no natural frustration mechanism, which could lead to a significant contribution of domain walls to the dark energy budget, exists. ",
        "introduction": "There is now very strong observational evidence that our Universe is currently undergoing a period of accelerated expansion \\cite{Komatsu:2008hk,Frieman:2008sn}. In the framework of General Relativity, the observed acceleration of the Universe ought to be explained by the existence of an exotic dark energy component which violates the strong energy condition. In  \\cite{Bucher:1998mh} it was suggested, for the first time, that a frozen domain wall network could be responsible for such acceleration and several other authors have subsequently advocated this possibility \\cite{Carter:2004dk,Battye:2005hw,Battye:2005ik,Carter:2006cf}.  The evolution of cosmological domain wall networks has been studied in detail using both high-resolution numerical simulations and a semi-analitical velocity-dependent one-scale (VOS) model \\cite{PinaAvelino:2006ia,Avelino:2006xy,Avelino:2006xf,Battye:2006pf,Avelino:2008ve,Avelino:2009tk}. Strong analytical and numerical evidence was provided that a frustrated domain wall network, accounting for a significant fraction of the energy density of the Universe at the present time, will never emerge from realistic phase transitions. These results appear to rule out a significant contribution of domain walls to the dark energy budget.  In all these studies the energy of the domain walls was assumed to be very small. Still, domain wall junctions can be responsible for freezing a domain wall network if they are heavy enough. In that case their contribution to the energy budget could not be neglected but it has been argued in \\cite{Avelino:2006xf,Avelino:2008ve} that it would spoil the dark energy properties associated with a frustrated domain wall network. More recently, a mechanism for dynamical frustration of domain wall networks involving kinky vortons was suggested in  \\cite{Battye:2009ac}.  In this paper we investigate the impact of string and monopole-type junctions on the dynamics of domain wall networks. Accounting for the detailed contribution of the junctions on domain wall network simulations is not a trivial task. In fact, the standard Press-Ryden-Spergel (PRS) algorithm \\cite{Press} (implemented in all field theory numerical simulations in order to ensure fixed comoving resolution) increases artificially the impact of the junctions on the overall network dynamics during the course of the simulations. This effect is not very important for the light junctions which are usually considered in such simulations but could be relevant in the case of heavy junctions.  In order to overcome this potential problem we develop a semi-analytical VOS model which incorporates  the contribution of string and monopole-type junctions to the overall dynamics of domain wall networks. We use this model to constrain the characteristic length and velocity of domain wall networks and discuss the corresponding cosmological implications, in particular for dark energy. ",
        "conclusions": "In this paper we have imposed strong constraints on the characteristic length and velocity of domain wall networks with string and monopole-type junctions using a semi-analitical VOS model. We have shown that a successful domain wall scenario for dark energy would require that $ L_0 \\lsim 10 \\, {\\rm kpc}$ and $v_0 \\lsim 10^{-5}$. We have demonstrated that such small values of $L_0$ and $v_0$ could only be achieved if the contribution of the monopole-type junctions to the total density of the universe was several orders of magnitude larger than that of domain walls and strings (assuming $\\Omega_{\\sigma 0} \\sim 1$), in complete disagreement with observations.  These results highlight the main difficulty with alternative mechanisms for the frustration of domain wall networks. The inclusion of additional degrees of freedom such as heavy junctions and friction may slightly reduce the characteristic length and velocity of the domain walls but is insufficient to lead to frustration, due to the limited amount of matter with which domain walls can interact while conserving energy and momentum at the present time (the mechanism for the frustration of domain wall networks proposed in \\cite{Battye:2009ac} is expected to face similar problems). Our results provide further evidence for the absence of a successful mechanism which can lead to the frustration of cosmologically relevant domain wall networks."
    },
    "0911/0911.2629_arXiv.txt": {
        "abstract": "FEARLESS (Fluid mEchanics with Adaptively Refined Large Eddy SimulationS) is a new numerical scheme arising from the combined use of subgrid scale (SGS) model for turbulence at the unresolved length scales and adaptive mesh refinement (AMR) for resolving the large scales. This tool is especially suitable for the study of turbulent flows in strongly clumped media. In this contribution, the main features of FEARLESS are briefly outlined. We then summarize the main results of FEARLESS cosmological simulations of galaxy cluster evolution. In clusters, the production of turbulence is closely correlated with merger events; for minor mergers, we find that turbulent dissipation affects the cluster energy budget only locally. The level of entropy in the cluster core is enhanced in FEARLESS simulations, in accord with a better modeling of the unresolved flow, and with its feedback on the resolved mixing in the ICM. ",
        "introduction": "The formation and evolution of the large scale structure, which proceeds through the process of hierarchical clustering, plays a key role for the conversion of potential gravitational energy into kinetic and internal energy in galaxy clusters and groups. In this process, cluster mergers induce bulk motions in the intra-cluster medium (henceforth ICM), with velocities of the order of $1000\\ \\mathrm{km\\ s^{-1}}$. The shearing instabilities associated with the merger events and the ubiquitous shock waves in the ICM are thus expected to stir the gas and make the flow turbulent.  Cluster formation is a typical example of turbulence generation in a strongly clumped medium. As the Eulerian approach is generally considered to be superior to the Lagrangian one in modeling hydrodynamical instabilities \\cite{ams07,tbm08,mmb09}), grid-based numerical codes are a suitable tool to address this problem, especially when the Adaptive Mesh Refinement (henceforth AMR, \\cite{Berger1984,bc89,n05}) technique is used. On the other hand, as in most of the astrophysical fluids, the Reynolds number in the ICM is estimated to be $Re \\gg 1$ (although with many issues on the value of the kinematic viscosity; \\cite{nm01,snb03,fsc03,rmf05,j07}), implying a number of degrees of freedom of the dynamical system \\cite{ll87,snh06} which is not manageable even in state-of-the-art numerical simulations\\footnote{In case of cosmological simulations, the problem is even more severe because the evolution of galaxy clusters has to be followed in the cosmological context of their collapse, i.e. in simulated computational domains much larger than the cluster scale itself.}. Therefore, even with AMR, it is usually not possible to resolve all the turbulent cascade, down to the dissipative length scale \\cite{snh06}.  In many fields of computational fluid dynamics, the influence of unresolved turbulence on the resolved scales is modeled by means of heuristic subgrid-scale (SGS) models, coupled to the large scales of the system, for which the fluid equations are solved (Large Eddy Simulations, LES; \\cite{Lesieur1996}). LES are customary, for example, in simulations of Type Ia supernovae \\cite{nh95,rhn02,rh05,snhr06,rhs07}. Recently, \\cite{sb08} developed a similar tool designed for studying Rayleigh-Taylor driven turbulence in the interaction between AGN outflows and the ICM \\cite{bsh09,bs09}.  In this contribution, we describe a novel numerical tool that unites LES and AMR, called FEARLESS (Fluid mEchanics with Adaptively Refined Large Eddy SimulationS). FEARLESS combines the adaptive refinement of the regions where turbulent flows develop with a consistent modeling of the SGS turbulence, and is argued to be a suitable tool in simulations of turbulent clumped flows. The numerical formalism and its first application to the physics of galaxy clusters is presented in \\cite{mis09}. The main features and results are collected here. ",
        "conclusions": ""
    },
    "0911/0911.5523_arXiv.txt": {
        "abstract": "We present the results from the first large ($>100$ source) 3.5\\,mm polarimetric survey of radio loud active galactic nuclei (AGN). This wavelength is favorable within the radio--mm range for measuring the intrinsic linearly polarized emission from AGN, since in general it is only marginally affected by Faraday rotation of the electric vector position angle, and depolarization. The  $I$, $Q$, $U$, and $V$ Stokes parameter observations were performed with the XPOL polarimeter at the IRAM 30\\,m Telescope on different observing epochs from July 2005 (when most of the measurements were made) to October 2009. Our sample consists of 145 flat-radio-spectrum AGN with declination $>-30^{\\circ}$ (J2000.0) and flux density $\\apgt1$\\,Jy at $\\sim86$\\,GHz, as measured at the IRAM 30\\,m Telescope from 1978 to 1994. This constraint on the radio spectrum causes our sample to be dominated by blazars, which allows us to conduct new statistical studies on this class of high-luminosity, relativistically-beamed emitters. We detect linear and circular polarization (above minimum $3\\sigma$ levels of $\\sim$1.5\\,\\%, and $\\sim$0.3\\,\\%) for 76\\,\\%, and 6\\,\\% of the sample, respectively. We find a clear excess in degree of linear polarization detected at 86\\,GHz with regard to that at 15\\,GHz by a factor of $\\sim2$. Over our entire source sample, the luminosity of the jets is anti--correlated with the degree of linear polarization. Consistent with previous findings claiming larger Doppler factors for brighter $\\gamma$-ray blazars, quasars listed in our sample, and in the \\emph{Fermi} Large Area Telescope Bright Source Catalog (LBAS), show larger luminosities than non--LBAS ones, but our data do not allow us to confirm the same for BL~Lac objects. We do not find a clear relation between the linear polarization angle and the jet structural position angle for any source class in our sample. We interpret this as the consequence of a markedly non-axisymmetric character of the 3\\,mm emitting region in the jets. We find that intrinsic circular polarization is the most likely mechanism for generation of the circular polarization detected in our observations. Our new data  can be used to estimate the 3.5\\,mm AGN contribution to measurements of the linear polarization of the cosmic microwave background, such as those performed by the Planck satellite. ",
        "introduction": "\\label{Intr}  Radio loud active galactic nuclei (AGN) are known to produce powerful pairs of highly collimated relativistic jets of magnetized plasma, which are able to extend far beyond the boundaries of the host galaxy. Their relativistic nature, their magnetic fields, and their non-isotropic geometry determine their most relevant observational properties: superluminal motions \\cite[e.g.,][]{Gomez:2001p201, Jorstad:2005p264}, Doppler boosted (decreased) emission of the jet pointing at a small (large) angle to the observer \\citep[e.g.,][]{Kadler:2004p5629}, rapid intrinsic emission variability \\cite[e.g.,][]{Agudo:2006p203, Fuhrmann:2008p267, 2008Natur.452..966M} and changes of the jet structure \\cite[e.g.,][]{2007AJ.134.799J, Agudo:2007p132}, and intense polarized synchrotron and inverse-Compton emission along all the electromagnetic spectrum \\cite[e.g.,][]{Ostorero:2006p406, Abdo:2010, Marscher:2010p11374}.  Polarimetric very long baseline interferometry (VLBI) observations at mm wavelengths, combined with single dish radio, mm, and optical linear polarimetric observing campaigns, allow one to (i) connect the location of the emitting regions at different observing bands and (ii) infer properties about the nature of the innermost moving and stationary knots of emission in the jets and the magnetic field in such regions, which usually cannot be resolved by VLBI \\cite[e.g.,][]{2007AJ.134.799J, 2008Natur.452..966M, Marscher:2010p11374}.  VLBI polarimetric surveys of radio loud AGN also provide relevant information about their jets; see \\citet{2003ApJ...589..733P} for a survey on a sample of 177 sources in the Calteck-Jodrell Bank Flat-Spectrum survey observed with the Very Long Baseline Array (VLBA) at 5\\,GHz, and \\citet{Lister:2005p261, Homan:2006p238} for the first results from the MOJAVE survey on 133 sources observed with the VLBA at 15\\,GHz. \\citet{Lister:2005p261} and \\citet{2003ApJ...589..733P} found the source cores (i.e., the innermost visible jet region at a given observing frequency with an instrument capable of resolving such a jet) to be weakly linearly polarized (\\aplt5\\,\\%), but with significantly larger fractional polarization for the jet regions downstream. Indeed, \\citet{Lister:2005p261} reported a general increase of linear polarization degree with increasing distance outward from the core in quasars and BL~Lacertae (BL~Lac) objets, but with BL~Lac jets more polarized in general. They also found the cores and jets of radio galaxies to be much weaker or not linearly polarized compared to those in quasars and BL~Lac objects. \\citet{Homan:2006p238} found ``strong'' circular polarization ($\\ge$0.3\\%, as defined by them), usually in the cores of the sources (but not always), within $\\sim$15\\% of their sample. They did not find significant correlations between the degree of circular polarization in the core and other relevant source properties.  No polarization survey of a large number of AGN has been performed at mm wavelengths thus far. It is known that AGN jets and their cores are affected by non-negligible Faraday rotation measures (RM) typically ranging from $\\sim10^2$\\,rad\\,m$^{-2}$ \\citep[e.g.,][]{2003ApJ...589..126Z, Zavala:2004p138, Gabuzda:2001p378, 2004MNRAS.351L..89G, 2008ApJ...675...79A} to $\\sim10^4$\\,rad\\,m$^{-2}$ \\citep[e.g.,][]{2002ApJ...566L...9Z, Attridge:2005p220, Gomez:2008p30}. These large RM values $\\sim10^4$\\,rad\\,m$^{-2}$ have only been found for a few sources, typically in the jets of radio galaxies  \\citep{Zavala:2004p138}. Since the observed electric vector linear polarization angle of a source $\\chi_{\\rm{obs}}=\\chi_{\\rm{int}}+\\rm{RM}~\\lambda^{2}$ (where $\\rm{RM}~\\lambda^{2}$ is the amount of Faraday rotation and $\\chi_{\\rm{int}}$ is the intrinsic polarization angle), a $\\lambda=$3.5\\,mm (86\\,GHz) survey would be much less affected by Faraday rotation than are cm wavelength observations. At this wavelength, $|\\rm{RM}|\\approx7000$\\,rad\\,m$^{-2}$ would be required to produce Faraday rotation by $\\approx5$\\,$^{\\circ}$, the median linear polarization angle uncertainty in the measurements presented in this paper. In contrast, such a value of $|\\rm{RM}|$ translates into rotations of $\\approx160^{\\circ}$ and $\\approx1400^{\\circ}$ at 2\\,cm (15\\,GHz) and 6\\,cm (5\\,GHz), respectively. If AGN jet RM in the radio and mm emitting regions are similar in general -- which is still to be confirmed observationally-- the reduced effect of Faraday rotation in mm wavelength observations also makes them much less sensitive to Faraday depolarization, hence reflecting the intrinsic linear polarization properties of the sources with better fidelity than radio observations.  The mm wavelength emission of radio loud AGN is typically dominated by compact regions in their jets; either in their innermost cores of emission, or in their pc scale superluminal knots \\citep[e.g.,][]{2007AJ.134.799J, 2008Natur.452..966M}. Indeed, the results from the 86\\,GHz VLBI survey of \\citet{Lee:2008p301} have shown that, in general, $\\sim50$\\,\\% of the total flux density of their imaged sources is concentrated in regions not further than $\\sim1.5$\\,milli-arcseconds from the mm core. In contrast, because of the longer lifetime of synchrotron radiation at cm wavelengths \\citep{Rybicki:1979p6159}, such emission is more spread out along the pc scale of the jets up to distances $>5$\\,milli-arcseconds \\citep{2003ApJ...589..733P, Lister:2005p261}. Moreover, owing to synchrotron opacity, the bright cm wavelength emission from compact cores in AGN lies in regions farther downstream in the jet than at mm wavelengths \\citep[e.g.,][]{Lobanov:1998p5354, Marscher:2006p362, 2007AJ.134.799J}.  In the absence of a large mm-VLBI polarimetric survey, the statistical analysis of a single-dish polarization-mm survey can yield relevant insights into the intrinsic polarization properties of the innermost compact regions of relativistic jets in AGN. Here we present the results from the first 3.5\\,mm polarimetric survey over a large ($>100$ source) sample of radio loud AGN.  Our results are also a valuable resource for surveys dedicated to studies of the polarization properties of the cosmic microwave background (CMB), specially that conducted by the ESA's Planck mission (see the mission \\emph{Bluebook}\\footnote{\\tt www.rssd.esa.int/SA/PLANCK/docs/Bluebook-ESA-SCI(2005)1\\_V2.pdf}). For such studies, radio loud AGN are considered the most important CMB-foreground contributors away from the Galactic plane at angular scales $\\aplt0.5^{\\circ}$ up to $\\sim100$\\,GHz observing frequency \\citep{Tucci:2005p8610}. Hence, \\emph{a priori} knowledge of their polarization is important to account for their contribution to the polarization of the CMB data \\citep[e.g.,][]{RubinoMartin:2008p8597,LopezCaniego:2009p8509}. ",
        "conclusions": "\\label{Conc}  We have performed the first large 3.5\\,mm polarization survey of radio loud AGN over a sample of 145 bright and flat radio spectrum sources, dominated by blazars, observed with the IRAM 30\\,m Telescope. At such wavelengths, Faraday rotation and depolarization are expected to have a weak effect on the observed linear polarization, which thus reflects the intrinsic linear polarization properties of the sources with high fidelity.  We detect linear and circular polarization above $3\\sigma$ levels for 76\\,\\% and 6\\,\\% of the sample, respectively. We have shown that the fractional linear polarization at 86\\,GHz clearly exceeds that at 15\\,GHz by a factor of $\\sim2$ for all classes of sources considered here. This implies both a larger degree of magnetic field order in the region where the bulk of the 86\\,GHz emission is produced, and considerable Faraday depolarization at 15\\,GHz, although we can not quantify the relative magnitude of these effects with the present data.  Consistent with previous work suggesting that the class of  $\\gamma$-ray bright blazars possess, in general, larger Doppler factors than weaker  $\\gamma$-ray bright blazars, we have found that LBAS quasars from our sample are more luminous than non-LBAS quasars. However, our data do not allow us to claim differences in the luminosity distributions of LBAS and non-LBAS BL~Lac objects.  Our entire source sample shows a trend of lower total flux luminosity for larger degrees of linear polarization, indicating that the level of magnetic field order in the innermost jet emitting regions is anti-correlated with jet luminosity.  Unlike other authors, who observed at radio frequencies (hence affected by Faraday rotation), we do not find a clear relation between the linear polarization angle and the jet structural position angle of BL~Lacs. This is also true for the remainder of the source subsamples considered in this work, despite the tendency of the electric-vector linear polarization angle to align with the jet position angle found at radio frequencies on quasars. This implies that the sources in our sample do not contain the conditions required for their linear polarization angles to lie either parallel or perpendicular to the jet axis (in the observer's frame), i.e., that they have a markedly non axisymmetric character. Indeed, the 3-dimensional character of a non-negligible population of jets in AGN at different scales is known to be complex \\cite[e.g.,][]{Fomalont:2000p4880,2003ApJ...589..733P,Jorstad:2005p264,Lister:2005p261,Agudo:2006p331, Agudo:2007p132}.  In the case of purely axisymmetric jets, cancellation of linearly polarized emission parallel and perpendicular to the jet axis forces $\\chi$ to align only either parallel, or perpendicular to the axis \\citep[e.g,][]{Lyutikov:2005p321,Cawthorne:2006p409}. Our results contradict this phenomenology, which suggests that the theoretical requirements for relativistic jets to show such behavior are not satisfied when observed at 3\\,mm. This implies that the 3\\,mm emitting region in relativistic jets in blazars is primarily non axially-symmetric at this wavelength. Hence, either the magnetic field or the emitting particle distributions (or both) responsible for the synchrotron radiation at 3\\,mm in a large fraction of blazars must have a markedly 3 dimensional (non-axisimmetric) character in order to account for the $|\\chi-\\phi_{\\rm{jet}}|$ distributions of our observations.  We have shown that the circularly polarized emission detected from our observations is most likely generated by intrinsic synchrotron emission. This has an important consequence. Intrinsic circular polarization cannot be produced in pure pair (electron-positron, $e^{+}e^{-}$) plasma jets. Thus, sources in which such mechanisms govern circular polarization production must be composed, at least in part, by an electron-proton ($p^{+}e^{-}$) plasma. This may be taken as a tool for AGN jet composition diagnostics through future mm wavelength polarimetric surveys. If a large fraction of  sources show $m_{\\rm{C}}$ with moduli as large as those detected at cm wavelengths, this may imply that their mm jet emitting regions are primarily composed of $p^{+}e^{-}$ plasma."
    },
    "0911/0911.0350_arXiv.txt": {
        "abstract": "These pages provide you with an example of the layout and style for 100\\% reproduction which we wish you to adopt during the preparation of your paper. This is the output from the \\LaTeX{} document class you requested. \\vspace{1pc} ",
        "introduction": "Text should be produced within the dimensions shown on these pages: each column 7.5 cm wide with 1 cm middle margin, total width of 16 cm and a maximum length of 19.5 cm on first pages and 21 cm on second and following pages. The \\LaTeX{} document class uses the maximum stipulated length apart from the following two exceptions (i) \\LaTeX{} does not begin a new section directly at the bottom of a page, but transfers the heading to the top of the next page; (ii) \\LaTeX{} never (well, hardly ever) exceeds the length of the text area in order to complete a section of text or a paragraph. Here are some references: \\cite{Scho70,Mazu84}.  \\subsection{Spacing}  We normally recommend the use of 1.0 (single) line spacing. However, when typing complicated mathematical text \\LaTeX{} automatically increases the space between text lines in order to prevent sub- and superscript fonts overlapping one another and making your printed matter illegible.  \\subsection{Fonts}  These instructions have been produced using a 10 point Computer Modern Roman. Other recommended fonts are 10 point Times Roman, New Century Schoolbook, Bookman Light and Palatino. ",
        "conclusions": ""
    },
    "0911/0911.1128.txt": {
        "abstract": "We track the coevolution of supermassive black holes (SMBHs) and their host galaxies through cosmic time. The calculation is embedded in the \\texttt{GALFORM} semi-analytic model which simulates the formation and evolution of galaxies in a cold dark matter (CDM) universe. The black hole (BH) and galaxy formation models are coupled: during the evolution of the host galaxy, hot and cold gas are added to the SMBH by flows triggered by halo gas cooling, disc instabilities and galaxy mergers. This builds up the mass and spin of the BH, and the resulting accretion power regulates gas cooling and subsequent star formation. The accretion flow is assumed to form a geometrically thin cool disc when the accretion rate exceeds $0.01\\dot{M}_{\\mathrm{Edd}}$, and a geometrically thick, radiatively inefficient hot flow when the accretion rate falls below this value. The resulting quasar optical luminosity function matches observations well, and the mass of the SMBH correlates with the mass of the galaxy bulge as in the observed $\\MBH-M_{\\mathrm{bulge}}$ relation. The BH spin distribution depends strongly on whether we assume that the gas in any given accretion episode remains in the same plane or it fragments into multiple, randomly aligned accretion episodes due to its self-gravity. We refer to these cases as the ``prolonged\" and ``chaotic\" accretion modes respectively. In the chaotic accretion model there is a clear correlation of spin with SMBH mass (and hence host galaxy bulge mass). Massive BHs ($M>5\\times10^8~\\Msun$) are hosted by giant elliptical galaxies and are rapidly spinning, while lower mass BHs are hosted in spiral galaxies and have much lower spin. Using the Blandford--Znajek mechanism for jet production to calculate the jet power, our model reproduces the radio loudness of radio galaxies, LINERS and Seyferts, suggesting that the jet properties of active galaxy nuclei (AGN) are a natural consequence of {\\em both} the accretion rate onto {\\em and} the spin of the central SMBH. This is the first confirmation that a CDM galaxy formation model can reproduce the observed radio phenomenology of AGN. ",
        "introduction": "\\label{sec:introduction}  Active galaxies can be classified according to the importance of the radio emission from their nucleus. Objects showing strong emission at radio wavelengths belong to the ``radio-loud\" class, whereas those with negligible emission belong to the ``radio-quiet\" class. Radio-loud AGN are associated with large-scale radio-emitting jets while radio-quiet AGN show very little or negligible jet activity. Sikora \\etal (2007) found that AGN have a bimodal distribution on the radio-optical plane, with radio-loud objects being about $10^{3}$ times brighter in radio than radio-quiet objects (see also Kellermann \\etal 1989; Xu \\etal 1999). The origin of this dichotomy remains unknown. However, it has been proposed in many studies that a first step towards explaining why some AGN launch prominent jets while others do not is to understand the nature of the central engine, the accreting BH.  AGN jets are believed to extract rotational energy from the BH and accretion disc through magnetic fields. Analytical studies have contributed to an understanding of the nature of the jets (Blandford \\& Znajek 1977; Macdonald \\& Thorne 1982; Blandford \\& Payne 1982; Begelman, Blandford \\& Rees 1984), but a breakthrough has come from magneto-hydrodynamical (MHD) simulations of the accretion flow. These self-consistently calculate the turbulent magnetic field dynamo which is the physical origin of the stresses which transport angular momentum outwards, allowing material to accrete onto the BH (Balbus \\& Hawley 1998). Embedding such calculations in a fully general relativistic framework produces relativistic jets from the accretion flow, and the collimation and jet power depend on the BH spin (McKinney \\& Gammie 2004; De Villiers \\etal 2005; Hawley \\& Krolik 2006).  Thus, it now seems clear that the jet power depends on the vertical (poloidal) magnetic field strength close to the BH, but the simulations are still not yet at the level where they can predict this {\\it ab. initio} (e.g. Beckwith \\etal 2008), so analytic models are still required. The most popular of these, the Blandford--Znajek (hereafter BZ) mechanism, uses the magnetic field as a means to tap the spin of the BH. This gives a strong spin dependence on the jet power.  Indeed, many authors have proposed spin as the physical parameter that determines the radio loudness of an AGN (see Wilson \\& Colbert 1995; Hughes \\& Blandford 2003). The spin paradigm, as this idea is called, offers a plausible theoretical explanation for the wide range of jet luminosities in AGN and could be the basis for understanding the observed dichotomy between radio-loud and radio-quiet AGN.  Alternatively, the radio--loud, radio--quiet switch may just be determined by the physical state of the accretion flow. Stellar--mass BH binary systems in our galaxy show a clear spectral transition at about 1\\% of the Eddington luminosity, $L\\sim 0.01L_{\\mathrm{Edd}}$, from a hot, optically thin accretion flow at low luminosities (e.g. an advection-dominated accretion flow or ADAF; Narayan \\& Yi 1994) to a cool, geometrically thin disc (Shakura \\& Sunyaev 1973) at higher luminosities (Esin \\etal 1997, see e.g. the review by Done, Gierlinski \\& Kubota 2007). The collapse of the radio jet observed across this transition clearly relates to the large drop in pressure (and hence scale height, $H$) between the hot and cool flow (e.g. Fender, Belloni \\& Gallo 2004). Even without BH spin, this produces a clear dichotomy of radio properties (e.g. Jester 2005).  It seems most plausible that {\\em both} these mechanisms, along with the mass accretion rate, affect jet power.  The BZ mechanism has a ``hidden\" dependence on mass accretion rate and the scale height of the flow because the magnetic field strength close to the BH saturates, due to the dynamo, into rough equipartition with the pressure in the flow. This depends on the mass accretion rate for either the disc or the hot flow, but the pressure in the hot flow is much larger than that in the cool disc so the field strength and hence jet power abruptly drop at this transition (Meier 2001; 2002)  Sikora \\etal (2007) used these ideas to interpret observations of radio galaxies and concluded that the data could be explained if the jet depends on mass accretion rate, accretion model and spin. However, they also required a mass-spin correlation, such that the most massive BHs, which, according to the $\\MBH-M_{\\mathrm{bulge}}$ relationship (Magorrian \\etal 1998; McLure \\& Dunlop 2002; Marconi \\& Hunt 2003; Haring \\& Rix 2004), reside in massive elliptical galaxies, have higher spin than lower mass BHs, which reside in spirals. They speculated that this could occur through the hierarchical growth of structure, where the major mergers which give rise to giant elliptical galaxies trigger large amounts of gas accretion with a given angular momentum direction onto the BH, spinning it up to the maximal value. By contrast, spiral galaxies have not experienced a recent major merger, and grow mainly through smooth accretion and multiple minor mergers with random angular momentum, resulting in low spin and hence weak jet luminosities.  Some of these ideas have recently been explored using simplified models of galaxy formation by Berti \\& Volonteri (2008) and Lagos \\etal (2009). In this paper, we calculate the growth of BHs by accretion and mergers, their acquisition of spin and their accretion rates within a specific theory of galaxy formation in a CDM model which has been extensively tested against a large range of observations of the galaxy population. Not only is our calculation of the joint evolution of BHs and galaxies self-consistent, but we also incorporate several accretion and jet-launching models from the literature which allows us to perform a quantitative comparison with observations of AGN.  The paper is organised as follows. In Section~2, we explain how we calculate the growth of the mass and spin of SMBHs through the accretion of hot gas from the halo (radio mode) and cold gas from the galactic disc and mergers (quasar mode). Most BHs grow primarily by gas accretion except the most massive ones which form late and increase their mass substantially by merging with other SMBHs. A BH is spun up to approximately the maximal value when it accretes its own mass from a flow at constant angular momentum.  In Section~3, we present the resulting distributions of BH mass, gas accretion rate and spin. We consider two possible modes of accretion.  In the ``prolonged accretion\" case, the gas is accreted in a single episode. Even minor mergers trigger gas flows onto the nucleus that often deposit a mass greater than the mass of the recipient BH. Thus, most BHs are typically spun up to the maximal value. This model predicts radio properties that do not reproduce the observed luminosity function. We then consider a ``chaotic accretion'' case in which the accretion episodes are limited by the self-gravity of the disc. Gas flows onto the nucleus give rise to a series of accretion episodes each typically augmenting the BH mass by only a small factor. Successive accretion events are uncorrelated, resulting in low spins for the relatively low mass BHs which grow primarily by accretion. On the other hand, the most massive holes which build up substantial mass through BH-BH mergers are spun up to high (but not maximal) spin values. In Section~4, we show that adopting this accretion model, the optical luminosity from disc-accreting objects matches the quasar luminosity function reasonably well.  In Section~5, we incorporate an explicit BZ model for the jet power (Meier 2002) and use this to predict the radio luminosity function which, in the chaotic accretion case, agrees well with observations. On the basis of our derived optical and radio luminosity distributions, we propose, in Section~6, a ``grand unification of AGN activity\", in which the accretion flow and jet luminosity are related to AGN type through their mass, spin and mass accretion rate. Fundamentally, our calculation shows, for the first time, how the coeval growth of BHs and their host galaxies result in optical and radio properties that explain the AGN activity seen in the local Universe. ",
        "conclusions": "We have presented a model of AGN activity in which the evolution of the galaxy is calculated in its full cosmological context and the luminosity of AGN at radio and optical wavelengths is calculated using a model for gas accretion and the generation of jets. We first considered the evolution of SMBH spin. We have found that the different astrophysical processes that influence the growth of SMBHs have a significant effect on the global spin distribution. For example, if accretion of gas of constant angular momentum dominates the growth of SMBHs, the associated holes will be rapidly rotating. However, if the gas that is fed into the SMBH has random angular momentum, a bimodal spin distribution results. In this case, high spin values occur only for the most massive BHs ($\\gtrsim10^8~\\Msun$), and these are mainly due to gas poor major mergers, where the BH growth is dominated by the BH-BH merger.  We have coupled this mass, spin and mass accretion rate evolution to a model for the accretion flow and jet. The accretion flow is assumed to form a geometrically thick, hot, radiatively inefficient flow (ADAF) for $\\dot{m}<0.01$, and to collapse to a thin disc at higher mass accretion rates. The jet power couples strongly to the accretion mode since it depends on the vertical (poloidal) magnetic field component close to the BH horizon. The collapse by two orders of magnitude in the scale height of the flow results in a similar drop in radio power. This already produces a dichotomy in radio properties which explains the distinction between radio-loud and radio quiet objects.  However, we also find that for our model to match the slope of the observed radio luminosity function, it is necessary for low mass BHs to generate relatively less jet power than high mass BHs. This is readily achieved in models where the jet couples {\\em strongly} to spin, such as in the classic BZ jet mechanism, {\\em and} where lower mass BHs have lower spin than the most massive BHs, as in our chaotic accretion model.  Coupling the chaotic accretion model to the BZ jet mechanism results in an AGN population which reproduces the diversity of nuclear activity seen in the local Universe. In particular, the model accounts for the radio and optical luminosities of the FR-I, BLRG, Seyfert and LINER galaxy populations. This is the first consistent demonstration that a great part of the phenomenology of AGN can be naturally explained by the coeval evolution of galaxies and BHs, coupled by AGN feedback, in a cold dark matter universe.  In future work, we will address the evolution of AGN across cosmic time, as a crucial test of galaxy formation models."
    },
    "0911/0911.2739_arXiv.txt": {
        "abstract": "The perturbed Friedmann-Lema\\^{\\i}tre-Robertson-Walker models allow many different possibilities for the 3-manifold of the comoving spatial section of the Universe. It used to be thought that global properties of the spatial section, including the topology of the space, have no feedback effect on dynamics. However, an elementary, weak-limit calculation shows that in the presence of a density perturbation, a gravitational feedback effect that is algebraically similar to dark energy does exist. Moreover, the effect differs between different 3-spaces. Among the well-proportioned spaces, the effect disappears down to third order in several cases, and down to fifth order for the Poincar\\'e dodecahedral space $S^3/I^*$. The Poincar\\'e space, that which also is preferred in many observational analyses, is better-balanced than the other spaces. ",
        "introduction": "What is the nature of the dark energy (for simplicity, let us use the term ``dark energy'' to generically signify either a form of dark energy or a cosmological constant, i.e. the parameter $\\Omega_\\Lambda$) that has been detected empirically to very high significance in observations of the Universe with several different telescopes and different types of surveys? This is one of the two key questions of this conference. However, the ``detection'' of dark energy would be meaningless without the underlying general-relativistic models of the Universe. These models are the Friedmann-Lema\\^{\\i}tre-Robertson-Walker models.  What did de~Sitter \\citep{deSitt17}, Friedmann \\citep{Fried23,Fried24}, Lema\\^{\\i}tre \\citep{Lemaitre31ell} and Robertson \\citep{Rob35} state were the two properties of the constant curvature 3-manifold of spatial sections of the Universe that would need to be determined by observations? They referred to both curvature and topology. It is through the link between curvature and the various density parameters (e.g. the total density $\\Omtot$) that we infer from observations that dark energy must (according to the exact-FLRW model) exist. What is the role of topology? We now know that theoretically, there are many different possibilities for the 3-manifold of the comoving spatial section of the Universe for any of the three curvatures \\citep{LaLu95,Lum98,Stark98,LR99,RG04} (for a shorter introduction, see \\cite{Rouk00BASI}). For many years, it was thought that topology has no effect on the Friedmann equation of an FLRW model, since the Einstein equations are, by definition, local, while topology (used in this context to mean the $\\pi_1$ homotopy group of a 3-manifold) is a global property of a 3-manifold. Indeed, for a perfectly homogeneous FLRW model, this argument appears to be correct.  However, this parallel session is about the fact that the Universe is not perfectly homogeneous. The planet Earth, the Solar System, the Galaxy, clusters of galaxies and the cosmic web of large-scale structure exist and constitute violations of perfect homogeneity. Many of the contributions in this session constitute work that may lead to (or already has led to) claims that the present ``dark energy'' is an approximation error rather than a physical phenomenon. That is, the Einstein equations should in principle be solved by an inhomgeneous solution rather than by an homogeneous solution with perturbations added afterwards.  Is there a topological feedback effect on the Friedmann equation in an almost-FLRW Universe, i.e. in one that is close to homogeneous but contains perturbations? An elementary, weak-limit calculation for a flat universe with one simply-connected direction shows that in the presence of a density perturbation, a gravitational feedback effect does exist \\citep{RBBSJ06}. Moreover, it is algebraically similar to dark energy. In more general 3-manifolds,  several further interesting characteristics are found \\citep{RBBSJ06,RR09}. ",
        "conclusions": "Throughout the history of cosmology since ancient times, preferences of the model of space have mostly shifted between finite spaces with a boundary and infinite unbounded spaces. Riemannian geometry made it possible to have flat (or hyperbolic) finite 3-spaces of constant curvature without boundaries. Pseudo-Riemannian geometry extends these to the FLRW space-time models of the Universe. Infinity is not a real number, and for the Universe to be a physical object, it would be most reasonable for the comoving volume, and hence total mass-energy, of the Universe to be finite.  \\begin{table} \\caption{Properties of empirically favoured ($\\mathbb{R}^3, T^3, S^3/I^*$) spatial 3-sections and some related 3-spaces  \\label{t-summary}} $ \\begin{array}{c c c c } \\hline \\mathrm{comoving} \\rule[-1.5ex]{0ex}{3.5ex} & {\\mathrm{finite?}} & {\\mathrm{missing}} & \\ddot{x}_{\\mathrm{resid}} \\propto \\\\ \\mbox{3-space} % & & {\\mathrm{fluctuations?}} & \\\\ & & {\\mathrm{(WMAP)}} & \\\\ \\hline \\rule{0ex}{2.5ex}  % \\mbox{infinite flat space} \\;\\;  { { \\mathbb{R}^3 }} & {\\mathrm{ no} } & {\\mathrm{ no} } & {\\mathrm{N/A} } \\\\ \\mbox{1-torus} \\;\\;  {{ T^1 := {S^1}\\times \\mathbb{R}^2 }} & {\\mathrm{ no} } & {\\mathrm{ no} } & { (x/L)^1} \\\\ \\mbox{3-torus} \\;\\;  {{ {T}^3 }} & {\\mathrm{ yes} } & {\\mathrm{ yes} } & { (x/L)^3} \\\\ \\mbox{octahedral space} \\;\\;  {{ S^3/T^* }} & {\\mathrm{ yes} } & {\\mathrm{ no} } & { (x/\\rC)^3} \\\\ \\mbox{truncated cube space} \\;\\;  {{ S^3/O^* }} & {\\mathrm{ yes} } & {\\mathrm{ no} } & { (x/\\rC)^3} \\\\ \\mbox{Poincar\\'e dodecahedral space} \\;\\;  {{ S^3/I^* }} &{\\mathrm{ yes} } & {\\mathrm{ yes} } & { (x/\\rC)^5 } \\\\ \\hline \\end{array}$ \\end{table}  The preferred models of comoving space, given the most recent cosmological observational catalogues, especially including the WMAP data, are summarised in Table~\\ref{t-summary}. The other models discussed above are included as well. The displacement is written as a scalar $x$ in all three cases for simplicity. Full expressions of the dominant terms for $T^3$ and $S^3/I^*$ are given in Eq.~(\\ref{e-test-object-T3-3rdorder}) and Eq.~(\\ref{e-residgrav-Poincare-exact}) respectively. While $\\mathbb{R}^3$ is frequently considered to be the implicit spatial section of the Concordance model, this is rarely stated explicitly, since it is not usually meant to be interpreted literally as a global model. It has the physically undesirable characteristic of giving the Universe an infinite amount of mass-energy, and has difficulty in explaining the lack of fluctuations above 10{\\hGpc} (in projection, above 60 degrees on the sky) in the WMAP cosmic microwave background maps.  Among the finite spaces, the well-proportioned spaces are, in general, good candidates for explaining the lack of $> 10${\\hGpc} fluctuations, but constraints on the curvature radius $\\rC$ make the fundamental domains of the octahedral and truncated-cube spaces a little too large compared to the surface of last scattering. For this reason, ``no'' is written for these spaces in Table~\\ref{t-summary}. The best candidates would appear to be the regular 3-torus $T^3$ and the Poincar\\'e dodecahedral space $S^3/I^*$. The latter appears to constitute an extremum among the class of possible 3-manifolds, in that the residual gravity effect is exceptionally well balanced, down to fifth order. We could say that ``some spaces are more equal than others'' and that the Poincar\\'e space is the space that is ``the most equal'' \\citep{Orwell45}.  This characteristic is a gravitational, geometrical, topological property of the Poincar\\'e space, provided that an inhomogeneity exists in the space, and it was clearly unknown to the group that first proposed the Poincar\\'e space as the best fit to the WMAP data \\citep{LumNat03}.  Is this just a coincidence? Or is it possible that gravity, geometry, topology and inhomogeneity together determine the most likely space to have been realised by early universe physical processes, in the way derived above, and that this most likely space is the one first proposed in 2003 in order to fit the WMAP data \\citep{LumNat03}?"
    },
    "0911/0911.2025_arXiv.txt": {
        "abstract": "The \\textit{Fermi Gamma-ray Space Telescope} observed the bright and long GRB090902B, lying at a redshift of $z = 1.822$. Together the Large Area Telescope (LAT) and the Gamma-ray Burst Monitor (GBM) cover the spectral range from 8 keV to $>$300 GeV.  Here we show that the prompt burst spectrum is consistent with emission from the jet photosphere combined with nonthermal emission described by a single powerlaw  with photon index -1.9. The photosphere gives rise to a strong quasi-blackbody spectrum which is somewhat broader than a single Planck function and has a characteristic temperature of $\\sim 290$ keV.  We model the photospheric emission with a multicolor blackbody and its shape indicates that the photospheric radius increases at higher latitudes. We derive the averaged photospheric radius $R_{\\rm ph}= (1.1 \\pm 0.3) \\times 10^{12} \\, Y^{1/4} \\mathrm{cm}$ and the bulk Lorentz factor of the flow, which is found to vary by a factor of two and has a maximal value of $\\Gamma = 750 \\, Y^{1/4}$. Here $Y$ is the ratio between the total fireball energy and the energy emitted in the gamma-rays. We find that during the first quarter of the prompt phase the photospheric emission dominates, which explains the delayed onset of the observed flux in the LAT compared to the GBM.  We interpret the broad band emission as synchrotron emission  at $R \\sim 4 \\times 10^{15}$ cm. Our analysis emphasize the  importance of having high temporal resolution when performing spectral analysis on GRBs, since there is strong spectral evolution. ",
        "introduction": "The mechanism giving rise to prompt emission in gamma-ray bursts (GRBs) has long been a puzzle. The emission typically peaks in the 100-1000 keV range and is modeled with the empirical Band function, which consists of two exponentially joined power-laws \\citep{band}.  In general the narrow spectral ranges available during observations of GRBs have made it difficult to unambiguously determine the emission mechanism. This has been remedied by the launch of the {\\it Fermi Gamma-ray Space Telescope} which regularly observes GRBs with its two instruments, the Large Area Telescope \\citep[LAT, nominal energy range 20\\,MeV--$>$300\\,GeV;][]{lat} and the Gamma-ray Burst Monitor \\citep[GBM, 8\\,keV--40\\,MeV;][]{gbm}, thus covering an unprecedented spectral range.  The bright, long GRB090902B was detected by {\\it Fermi} and was one of the brightest observed to date by the LAT \\citep{chi09}. Over 200 photons were detected at energies above 100\\,MeV, including 39 photons with energy above 1\\,GeV \\citep{deP09}. The LAT also detected a photon with energy $33.4^{+2.7}_{-3.5}$\\,GeV, the highest seen in any GRB so far. The burst lies at a redshift of $z=1.822$ \\citep{cuc09}, which yields an isotropic energy of $E_{\\rm iso}= (3.63 \\pm 0.05) \\times 10^{54}$ ergs, and the afterglow emission was seen at X-ray, optical, NIR and radio wavelengths.  The prompt spectrum of GRB090902B over the  energy range 8 keV-33 GeV shows clear deviation from the generally expected Band function \\citep{chi09}. The time-integrated spectrum is best fitted by the addition of a separate power-law component (photon index of $\\sim -1.9$) to the Band function. The power-law is detected at both lower and higher energies around the Band component. During the first half of the prompt emission phase, the Band power-law index at energies below the spectral peak significantly violates the optically thin synchrotron limit; photon index of $\\alpha = -2/3$ \\citep{preece98,preece02}.  \\cite{ryde05} suggested that GRB spectra are a superposition of two spectral components: photospheric blackbody emission and an accompanying nonthermal component. Over the limited energy range of 20-2000 keV (using data from BATSE on {\\it CGRO}) studied, an adequate model consists of a single Planck function and a power-law. However, the nonthermal component, which is modelled by the power-law, is not well characterized, since the burst spectral shape is in general dominated by the thermal component (see further \\citet{ryde08,rydepeer09}).  The \\textit{Fermi} data now allow an improved characterization of the two components due to the increased energy range. In this paper, we investigate in more detail whether the observed two components in the broadband spectrum of GRB090902B could be attributed to such a photospheric model. ",
        "conclusions": "The presence of a strong photospheric component in GRB spectra has been discussed by several authors \\citep[e.g. ][]{good86,MRRZ,rm05}\\footnote{see also e.g. \\citet{pac86,mes00,LU,DM,DS,R,GS,peer06,ruffini}}. Here we have identified a strong photospheric component in GRB090902B, which is modelled by a multicolor blackbody.   It dominates over  the nonthermal component, being nearly 100\\% at early times.  Varying energy injection at the central engine, or variations caused by interaction with the progenitor material  can cause the observed variability in the light curve. The flow is advected through the photosphere, where the thermal emission escapes. A fraction of the kinetic energy stored in the flow is later dissipated and emitted as nonthermal emission, e.g. synchrotron radiation. It is therefore expected that both the thermal and nonthermal emission reflect the properties and original energy injection at the central engine and thus are correlated with each other. This is consistent with the high correlation observed,  with a time lag of approximately 0.5 s. Such a lag corresponds to a shock radius $R_{\\rm sh} - R_{\\rm ph} = 2 \\, c \\Gamma^2 t_{\\rm lag} (1+z)^{-1} \\sim 4 \\times 10^{15} (\\Gamma/580)^2   \\mathrm{cm} \\sim R_{\\rm sh}$ at which the optical depth $\\tau \\sim (R_{\\rm sh}/R_{\\rm ph})^{-2} \\sim 10^{-7}$ ensuring an optically thin-emission site.  \\cite{chi09} estimated a lower limit on the bulk flow Lorentz factor $\\Gamma_{\\rm min}  \\sim 1000$  by combining the energy of the 11.16 GeV photon, observed at 11 s after the GBM trigger, with the variability timescale  in the LAT data which was determined to be $t_{\\rm v} \\sim 0.1$ s: The gamma-ray opacity for pair production should be less than unity in order to allow the photons to escape without attenuation from the flow. This  Lorentz factor  is slightly higher than the value we determined above for the flow advected through the photosphere. However, the estimations of $\\Gamma$ based on opacity arguments are sensitive to uncertainties in determining the variability time  \\citep{zhangpeer09}.  Indeed, assuming that the variability timescale in the nonthermal component is set by the angular timescale at the  shock radius, we have $t_{\\rm v} = t_{\\rm ang}^{sh} =   (1+z)^{} \\,\\,R_{\\rm sh}/ (2 \\, c \\Gamma^2)= 0.5$ s.  This value is consistent with the  fact that there is not much variability strength on timescales shorter than 1 s, as measured by the power density spectrum of the LAT data.   A variability timescale of  0.5 s yields  $\\Gamma_{\\rm min} \\sim 750$, which is close the value we estimate. Finally, we note that the observed variability timescale in the GBM data should  reflect the angular timescale of the photosphere. The determined value $t_v^{\\rm GBM} \\sim 50$ ms \\citep{chi09} is larger than $(1+z)^{} \\,\\, R_{\\rm ph}/ (2 \\, c \\Gamma^2)\\sim  2 \\times 10^{-4}$ s, which is to be expected as shown by Pe'er (2008).  As shown above, the photospheric emission allowed us to estimate $R_0$ assuming nondissipative acceleration: $ R_{\\rm 0} \\sim10^{9}\\, \\xi^{-4} Y^{-3/2} \\mathrm{cm} $.   This is larger than the physical size  around a solar mass black hole, typically assumed to be at the centre of the GRB engine, with a Schwarzschild radius of $R_{\\rm sch} = M/c^2 \\sim  10^7 \\mathrm{cm}$ \\citep{pac86}. This might be an indication of that the fireball dynamics includes significant dissipation during the acceleration phase and/or subphotospheric heating, which would lower the estimated value of  $ R_{\\rm 0} $. Alternatively, the larger size we find might correspond to the radius of the stellar core, at which the jet is launched. \\cite{thom07} argue that internal dissipation within the star prevents the Lorentz factor from significantly rising until the jet escapes the core of the progenitor star, such as a Wolf-Rayet star. However,  \\cite{lazz09}) showed that the dissipative shocks in the jet continue far beyond the stellar radius, thereby  influencing the jet dynamics.  Some bursts have no strong evidence of a narrow thermal component. For cases like GRB080916C which has a standard Band spectrum covering 6-7 orders of magnitude, the lack of a thermal component can be taken as an argument for a Poynting flux dominated flow \\citep{zhangpeer09}. For other cases that show a standard Band spectrum in a narrower energy range, like during the second half of the prompt phase in 090902b, it is possible that a modified photosphere spectrum with subphotospheric heating can account for the data \\citep{peer07,bel09,lazz2009}. Broadband observations of more GRB prompt emission spectra are desirable to more definitely diagnose the composition of GRB jets."
    },
    "0911/0911.1817_arXiv.txt": {
        "abstract": "We present a comprehensive multiband spectral and polarimetric study of the jet of 3C 264 (NGC 3862). Included in this study are three {\\it HST} optical and ultraviolet polarimetry data sets, along with new and archival {\\it VLA} radio imaging and polarimetry, a re-analysis of numerous {\\it HST} broadband data sets from the near infrared to the far ultraviolet, and a {\\it Chandra} ACIS-S observation.  We investigate similarities and differences between optical and radio polarimetry, in both degree of polarization and projected magnetic field direction. We also examine the broadband spectral energy distribution of both the nucleus and jet of 3C 264, from the radio through the X-rays. From this we place constraints on the physics of the 3C 264 system, the jet and its dynamics.  We find significant curvature of the spectrum from the near-IR to ultraviolet, and synchrotron breaks steeper than 0.5, a situation also encountered in the jet of M87.  This likely indicates velocity and/or magnetic field gradients and more efficient particle acceleration localized in the faster/higher magnetic field parts of the flow. The magnetic field structure of the 3C 264 jet is remarkably smooth; however, we do find complex magnetic field structure that is correlated with changes in the optical spectrum.  We find that the X-ray emission is due to the synchrotron process; we model the jet spectrum and discuss mechanisms for accelerating particles to the needed energies, together with implications for the orientation of the jet under a possible spine-sheath model. ",
        "introduction": "The characterization of jet polarization properties provides a powerful diagnostic of jet physics, particularly with respect to magnetic field configuration and particle acceleration.  Extragalactic jets generally emit a continuum of radiation from radio through optical, and often into the X-ray regime. Through matched resolution comparisons of flux density measurements at various frequencies, we can glean morphological information  about particle acceleration regions and jet energetic structure.   The closest kiloparsec-scale radio-optical jet is that of M87 at a distance of $16{\\rm ~Mpc}$ (Tonry 1991).  By using matched resolution polarimetry at different wavelengths and combining this with multi-wavelength imaging and X-ray imaging and spectroscopy, much information about jet energetics and magnetic fields, three dimensional structure and particle acceleration can be constrained (Perlman \\& Wilson 2005 and references therein). There are relatively few jets for which there exists HST optical polarimetry compared to the number that have been detected.  As of this writing, a total of ten optical jets have HST polarimetry observations (M87 - e.g. Perlman et al. 1999; 3C 273 - Thomson et al. 1993; 3C 293 - Floyd et al.  2006; 3C 15 - Perlman et al. 2006 and Dulwich et al. 2007; 3C 346 - Perlman et al. 2006 and Dulwich et al. 2009; 3C 66B, 3C 78, 3C 264 and 3C 371 - Perlman et al. 2006;  PKS 1136--135 - Cara et al., in prep.).   This is in comparison to the $\\sim 34$ detected optical extragalactic jets {\\footnote{see http://home.fnal.gov/\\textasciitilde jester/optjets/ maintained by S. Jester}}. Given this dearth of polarimetry observations, it is not surprising that there are  few constraints on the configuration of the magnetic fields in optically emitting regions of extragalactic jets. Here we present a comprehensive study of the jet of  3C 264 that includes radio and optical polarimetry, as well as X-ray observations of the jet.  At a redshift of 0.0217 (Baum et al. 1990), and hence a distance of 94 Mpc,  3C 264 is among the closest known bright radio galaxies with an optical jet. Also detected at X-ray energies, its relative proximity makes the 3C 264 jet a prime candidate for deep optical, radio and X-ray studies, as we present here.  It is hosted by NGC~3862, a large elliptical galaxy offset to the south-east from the center of the cluster Abell 1367, and is classified as a Fanaroff-Riley type I radio source (Fanaroff \\& Riley, 1974). On large scales it exhibits a twin-tail radio structure extending to the north-east (Bridle \\& Vallee 1981). On arcsec scales it consists of a compact core, with a one-sided, nearly knot-free jet extending also to the north-east and a weak counter jet (Lara et al. 1997). The optical jet counterpart extends for only roughly $2''$ beyond the galaxy core, the latter arcsecond of which is considerably fainter than the inner jet (Perlman et al. 2006). This short length combined with the ``smoothness'' of the jet makes studying this object difficult.  In particular, in the optical and near-IR  useful observations are really only possible with HST, as most adaptive-optics systems require a nearby, bright point source.  3C 264 was part of the original Third Cambridge catalog with $S_{\\rm 159~MHz}= 37{\\rm ~Jy}$ (Edge et al. 1959) and the optical jet counterpart was discovered by Crane et al (1993) with the pre-COSTAR FOC instrument on HST. Since this initial optical discovery, over $50{\\rm ~ks}$ of HST observing time has been dedicated to this source alone (e.g. Sparks et al. 1994; Baum et al. 1997, 1998; Hutchings et al. 1998; Martel et al. 1999, 2000; Noel-Storr et al. 2003). More recent observations have concentrated on deep optical polarimetry with WFPC2 (Perlman et al. 2006) and optical and UV polarimetry with ACS (Capetti et al. 2007). It was first detected in X-rays with {\\it Einstein} (Elvis et al. 1981), and more recently observed by {\\it Chandra} and {\\it XMM-Newton}, with the nuclear emissions studied by Donato et al. (2004) and Evans et al. (2005), and the hot atmosphere by Sun et al. (2007). Also, several radio observations have been performed with the VLA (Lara et al. 1999; Lara et al. 2004; this paper), the global VLBI network (Lara et al. 1997) and the EVN and MERLIN arrays (Baum et al. 1997; Lara et al. 1999; Lara et al. 2004).  In this paper we attempt to bring all of this observational information into a coherent picture of the physics of the 3C 264 jet. In \\S 2 we detail the observations and data reduction techniques used. In \\S 3 we present comparisons between optical and radio polarimetry, optical and radio spectral index maps, and a comparison of the resolved and unresolved jet emissions from {\\it Chandra}. Lastly, in \\S 4 we discuss the implications of these results, and present a simple model that attempts to explain them. Throughout the paper we assume a flat Friedmann cosmology with $H_0 = 72{\\rm ~km ~s^{-1} ~Mpc^{-1}}$, (corresponding to the latest WMAP cosmology, see Dunkley et al. 2009), giving a projected linear scale of $0.42{\\rm ~kpc/''}$. ",
        "conclusions": "In the preceding sections, we have discussed in detail optical and radio polarization images of the jet of 3C 264, as well as its optical and broadband spectrum and the results of a deep {\\it Chandra} X-ray observation.  As already discussed, this jet displays a smooth morphology in both optical and radio. The differences we see in polarization morphology, while less drastic than those seen in e.g., M87 (Perlman et al. 1999), are nevertheless significant. We also find significant variations in the optical and radio-optical spectrum of the jet.  X-ray emission is also seen from the jet on scales between $0.7-1.8''$ from the nucleus. In this section, we discuss the inter-relationship of these properties and their relevance to the overall physics of the 3C 264 jet.  We will also discuss other jets where radio and optical polarimetry has recently been done, in particular M87 (Perlman et al. 1999, 2001, Perlman \\& Wilson 2005), 3C 15 (Dulwich et al. 2007), and 3C 346 (Worrall \\& Birkinshaw 2005, Dulwich et al. 2009), and their relationship to this source.  \\subsection{The Physics of the 3C 264 System}  3C 264 is a highly interesting radio galaxy that contains many elements that are important to understanding both the AGN phenomenon as well as jet physics. The data discussed herein allow us to place significant new constraints on the physics of the nucleus and jet, and also to comment on the nature of the host galaxy environment.  The latter is quite interesting in and of itself, as the host galaxy to 3C 264, namely NGC 3862, lies within -- but not at the center of -- a rich cluster that has complex substructure. Cortese et al. (2004) have found that Abell 1367, the surrounding cluster, lies at the intersection of two filaments within the Great Wall, and is likely in the midst of its formation process.  The cluster has a large number of subclumps and grouping, as revealed by multi-object spectroscopy and X-ray studies (Cortese et al. 2004; Sun et al. 2005, 2007).  NGC 3862 itself lies at the center of a compact group within the cluster that consists of 18 galaxies (Lipovka et al. 2005).  Our {\\it Chandra X-ray} data reveal that the complicated temperature structure of the A1367 system continues down to scales of kpc, with evidence of a cooling-flow  like structure within the subgroup.  NGC 3862 was dubbed ``optically dull'' by Elvis et al. (1981), referring to the relative strength of nuclear X-ray to optical emission.  Further investigation of galaxies with similar characteristics (sometimes called X-ray bright optically normal galaxies, or XBONGs) have normally found that they can be fitted into the scheme of conventional accretion models if there is a moderate amount of nuclear gas ($N_{\\rm H} \\sim 10^{21}$ cm$^{-2}$) and dust (E(B-V) $\\sim 0.5-0.8$) covering large solid angles (e.g., Civano et al. 2007). While our HST images do show a dusty disk a few hundred parsecs from the AGN, this disk does not appear to continue all the way to the center, and indeed, we find no evidence of reddening or absorption of the nucleus, based both on our HST and {\\it Chandra} data.  A similar conclusion was reached by Hutchings et al. (1998) based on analysis of the nuclear spectrum. Thus this distinct lack of nuclear gas makes 3C 264 rather different from most XBONGs.  However, in 3C\\,264 there is a bright radio nucleus, and the radio and X-ray emissions could be enhanced relative to the optical continuum if the continuum nuclear emission\\footnote{We note the presence of narrow forbidden lines in the optical spectrum (Hutchings et al. 1998) } is primarily of a non-thermal origin.  Our data make it clear that the primary radiation mechanism of the 3C 264 jet and nucleus in the radio through ultraviolet is synchrotron emission. A broken-power-law synchrotron model (left panel of Fig.~\\ref{fig:radiotoxray}) describes the overall characteristics of the radio-to-optical spectrum of the nucleus rather well, especially given the non-contemporaneous nature of the data.  We have assumed a spherical blob of radius 2~pc, in which an electron spectrum of minimum Lorentz factor $\\gamma_{\\rm min} = 300$ produces radiation of spectral index $\\alpha = 0.5$ in a minimum-energy magnetic field of $\\sim 9$~mG, breaking by $\\Delta \\alpha = 0.7$ at $\\gamma_{\\rm break} = 7 \\times 10^4$ (see e.g., Worrall \\& Birkinshaw 2006 for relevant formulae). The dominant X-ray inverse Compton mechanism is synchrotron self-Compton, but this gives a negligible contribution to the observed X-rays.  It is interesting that the synchrotron lifetime of electrons producing emission close to the break frequency is $\\sim 5$ yr, of the same order of magnitude as the light-travel time across the source. However, $\\Delta \\alpha$ is greater than the value of 0.5 expected for a simple energy-loss model, and is in the range $0.6-0.9$ commonly found for the resolved jets of low-power radio galaxies, and not yet readily explained by acceleration models (see Worrall 2009 for a review). We return to this latter issue in section 4.2. While a synchrotron model can describe the overall features of the nuclear continuum of 3C\\,264, we note that the more extreme case of a radio and X-ray bright (and unabsorbed) radio galaxy with an optically dull continuum, PKS~J2310-437 (Tananbaum et al. 1997, Worrall et al. 1999), has proved to be more problematic to fit to a simple synchrotron model (Bliss et al. 2009).  Also supporting a synchrotron nature for the nuclear emissions in the optical are the polarized emissions found by our HST observations.  A similar conclusion was reached by Capetti et al. (2007), who note an optical polarization $\\sim 8\\%$ for the core and a position angle for the electric field vectors $\\sim 5^\\circ$.  This fractional polarization is too high to be reached with scattering (see Capetti et al. for discussion) and is most consistent with synchrotron radiation.  It  should be noted, however, that the magnetic field orientation seen by Capetti et al. (2007) is somewhat different (by about 25 degrees)  from that seen either in the K-band, A-array data we (or Lara et al. 2004) present, and is also different from that seen in most of  the VLBI jet (Kharb et al. 2009).  Unfortunately we cannot use our {\\it WFPC2}  polarimetry data to comment on this latter issue, as the core is saturated.  The optical emissions from the jet have been known to be due to synchrotron radiation for more than a decade (e.g., Baum et al. 1997).  Our data allow us to extend this conclusion considerably.  Not only do they allow us to place much better constraints on the shape of the jet's optical SED, but they also show that synchrotron radiation is the dominant jet emission mechanism up to X-ray energies. While the X-ray spectrum of the resolved jet is not measured, the emission is undoubtedly of synchrotron origin.  X-rays are measured only from the outer jet where contamination from core emission is least, but they arise from a relatively large region.  We have modelled this section of jet with a cylinder of radius 20~pc \\citep{Lara04} and length 550~pc, and used $\\gamma_{\\rm min} = 300$. Again inverse Compton scattering, including on the cosmic microwave background radiation, gives a negligible contribution to the observed X-ray emission.  Again a synchrotron model of a broken power law fits well the data (right panel of Fig. 13).  To find a minimum value of $\\Delta \\alpha$ we allow the break to occur at as low a frequency as possible, but as for the core the value of $\\Delta \\alpha$ (0.82) is greater than 0.5, and in the range $0.6-0.9$ commonly found for the resolved jets of low-power radio galaxies for which the X-ray spectrum is measured.  The jet spectral index below the break is $\\alpha=0.7$, the minimum-energy magnetic field is $\\sim 330 \\mu$G, and the Lorentz factor of electrons emitting radiation close to the break is $\\gamma_{\\rm break} = 1.6 \\times 10^6$.  The synchrotron lifetime of electrons producing radiation at the break frequency ($\\sim 140$ yr) is again of the same order of magnitude as the light-travel time across the source.  As noted above, the jet of 3C 264 shows rather subtle magnetic field direction changes, both in the optical/UV and radio.  The first region where changes are seen to occur are the knot A/B region, which displays a low polarization region in the interknot region between them, that extends northwest  along the northern  side of knot A as well as northeast along the eastern side of knot B.   Twists in the 30-90$^\\circ$ range are seen to the side of this feature. The low polarization at the knot maximum is likely to be the result of field superposition, and the polarization angle changes may be due to the occurrence of perpendicular and oblique shocks at the location of the knot maxima. A similar pattern of changing position angles in knot C of the 3C 15 jet has, alternatively, been modelled as the superposition of a helical strand of magnetic field on a more ordered field structure. We do not have the angular resolution in our X-ray data to see anything that might be related to the latter feature (as we do in 3C 15), however, a comparison of the two sources and their associated maps reveals some similarity.   A second region is seen with significant rotations in the orientation of the magnetic field, namely the knot D region.  This region shows low polarization in the jet  interior and particularly at the knot maximum, with magnetic field vectors that  follow the jet's flux contours downstream.  Such a configuration is very similar to  some of the knots in the M87 jet, particularly knots D and F.  Perlman et al. (1999) modeled those features as shocks that occur within the jet interior, in a  stratified jet where the optical/UV  emitting particles are also concentrated in the  jet interior, while lower energy, radio emitting particles are more evenly distributed through the jet.  Such a model can be applied  for the most part to these knots as well, with the possible exception that here, we do not have sufficient evidence in our spectral index maps to say that the knots are much more optically bright than the edges, although the $\\alpha_{ro}$ map does show that they do have a slightly  flatter radio-optical spectrum.   Thus, as in the M87 system,  the knots in the 3C 264 jet regions are likely to be  sites of enhanced {\\it in situ} particle acceleration.   The nature of the dust ring in 3C 264 and its relationship to the jet is an interesting topic.  We see some apparent steepening in the jet's optical spectrum (from $\\alpha_O \\approx 0.75$ to $\\alpha_O \\approx 0.9$) near this location, as well as significant changes in polarization in both the optical and radio. These changes include minima in the radio polarization and changes by 20-40$^\\circ$ in the direction of the magnetic field vectors in the optical.  The former does not  necessarily indicate an interaction between the jet and the dust ring, but rather is more consistent with absorption of jet radiation   by the dust within the ring.  Moreover, a closer look at Figure 4 reveals that the spectral and polarization changes begin about $0.1''$ closer to the nucleus than the inner edge of the ring.   In addition, the changes in the direction of the optical polarization vectors are much stronger near the  northern edge and continue in the region beyond the ring.  This indicates that the spectral and polarization changes are primarily intrinsic to the jet rather than caused by any screening by or interaction with the ring.  Thus, we believe that the disk and jet most likely occupy physically distinct regions of the galaxy.  A similar conclusion was reached by Martel et al. (2000) on the basis of the continuous distribution of line emission at the location of the disk, with the galaxy subtracted long-slit spectrum of Hutchings et al. (1998) showing no evidence of line emission from the jet itself.  Beyond $1''$, the optical polarization is much harder to measure, due to the much lower optical flux.  Interestingly, however, where it is measurable, we see an increased  optical polarization in this region, as high as 40\\%, with an orientation slightly different than that seen in the rest of the inner jet, perhaps reflecting the turn towards a more northerly orientation that is seen in the radio flux contours at  slightly greater distances (Figures 2, 3, 8 and Lara et  al. 2004).  The overall morphology of the radio through X-ray spectrum of the core+jet (Figure 7, {\\it bottom right panel}), combined with the power-law slope of the optical-to-X-ray spectrum, is most consistent with synchrotron radiation. Interestingly, the  spectral index observed for the core + inner jet in the X-rays is reasonably similar (within $2 \\sigma$) to  that observed in our data in the ultraviolet for the entire jet, as well as that of the inner jet (the outer jet is too faint to see in the UV data).  This implies little curvature of the spectrum between the UV and  X-rays, and also implies that the $\\nu F_\\nu$ peak of the jet's spectrum is located at a few  $\\times 10^{14}$ Hz, over a broad range of jet components.  The spectral break required for the inner 3C 264 jet  is $\\Delta \\alpha = 0.6$; slightly but not significantly smaller than seen in M87 ($\\Delta \\alpha = 0.7$, Perlman \\& Wilson 2005).   Also in accord with the  synchrotron radiation model, we see that the width of the jet is very similar in the radio, optical and UV, for as long as the data allow us to track it (Figure 14). Interestingly, the jet FWHM increases smoothly with distance from the nucleus in  all three bands, with the width $W \\propto r^{1.1}$ approximately.  \\begin{figure} \\label{width} \\centerline{\\includegraphics*[width=3.5in]{3c264_Opt_Jet_Width.ps}} \\caption{The width of the 3C 264 jet in the UV, optical and radio.  Three datasets were used: The VLA K-band, A array data (crosses), Our Stokes I, F555W data (diamonds), and the NUVMAMA image (asterisks).  Note the constant rate of increase in jet width, as observed in all three bands out to about $1.2''$ from the core.} \\end{figure}    \\subsection{Relationship to Other Jets}  It is useful to discuss these results in the broader context of the class of optical and X-ray jets.  Ten such jets have now been observed  polarimetrically with HST:  M 87 (Capetti et al. 1997, Perlman et al. 1999), 3C 273 (Thomson et al. 1993), 3C 293 (Floyd et al. 2006), 3C 15, 3C 66B, 3C 78, 3C 264, 3C 346, 3C 371 (Perlman et al. 2006; see also Dulwich et al. 2007, 2009)  and most recently PKS 1136-135 (Cara et al., in prep.).  Of these the vast majority (all but 3C 273 and PKS 1136-135) are relatively low-power sources, albeit with a variety of optical and radio morphologies. These objects are all at a wide range of distances, ranging from just 16 Mpc in the case of M 87 to $z=0.534$ in the case of PKS 1136-135.  Of these ten, just three (M 87, Perlman et al. 1999; 3C 15 and 3C 346; Dulwich et al. 2007, 2009)  have previously had deep analyses of their radio and optical total flux and polarization morphology, as well as broadband spectra and X-ray emission published (importantly, note that the existing  HST polarimetry of 3C 273 is not deep enough for these purposes).   Compared to the other objects, 3C 264's  optical and X-ray jet is the most compact:  at $2''$ its angular length is less than any other optical/X-ray  jet for which polarimetry has been done, except for 3C 78, which is very similar in length, although its physical length of about 850 pc (projected) is not that much shorter than the 1.5 kpc projected extent of the M87 jet. This limits our ability to extend the detail we can obtain in the  radio and optical imagery to the X-rays, as the jet of 3C 264 is only barely  resolved with {\\it Chandra}.  Thus with these observations we cannot comment in detail on the correlation between changes in optical polarization and the  location of X-ray components, previously noted by Perlman \\& Wilson (2005),  Perlman et al. (2006) and Dulwich et al. (2007, 2009).  The dominant polarization configuration observed in the jet of 3C 264, in both the radio and optical, is one of a high polarization (40-50\\%) along the jet edges, and somewhat lower polarization in the jet center, with magnetic field vectors oriented primarily along the direction of the jet.  Such  a configuration is seen in most of the knot A-B complex of the M 87 jet, albeit with  considerably more detail because of the greater proximity of that source.    What is missing in 3C 264, however (as compared to M87), are the presence of  strong \"shock-like\" components, which are seen not only in the A-B complex of the M 87 jet, but also commonly in the other two jets as well as other regions of the M 87 jet.  By contrast, the jet of 3C 264 has a very smooth polarization morphology, unlike any of the other three jets analyzed deeply so far and similar only to 3C 78 in detail. Perhaps related to this, we note that 3C 264 and 3C 78 represent the lowest power jets of those that have been observed polarimetrically with HST. It is therefore possible that the overall character of the polarization morphology is partly a function of the  jet power.  We will test this in our future studies of the remaining jets in the Perlman et al. (2006) sample (3C 78, 3C 371 and 3C 66B,  for which future papers are all under preparation).  We turn now to an effect observed in all four jets: the spectral energy distribution (SED) peaks at $\\sim $ optical-UV energies, and is followed by a power-law extension, with a break steeper than the canonical $\\Delta\\alpha=1/2$. In M87,  the spectrum appears to continue to steepen at X-ray energies; however this is not the case in 3C 264.    Although a possible formal  explanation for this steeper break is that electrons do not diffuse in pitch angle as they radiate (Kardashev 1962, Pacholczyk 1970), a more plausible explanation is that the magnetic field decreases away from the sites of particle acceleration. This naturally  results in spectral breaks greater than $1/2$, assuming that electrons, after they are accelerated propagate toward  regions with lower magnetic field (Wilson 1977). Recent work by Reynolds (2009) has expanded upon this aspect, associating steeper breaks with inhomogeneities and/or differential losses within the source. Another possibility that can work in parallel with the magnetic field gradients is that of velocity gradients in the flow.  This idea was first posited by Laing and Bridle (Laing 1996, Bridle 1996) to deal with observed details of radio polarization maps of jets, which suggested a spine-sheath configuration for the jet, with a faster flow in the central spine, similar to a firehose. Later work invoked lateral velocity gradients to model the sub-pc scale jets of TeV emitting BL Lacertae objects (Georganopoulos \\& Kazanas 2003), and can be also applied to transverse velocity gradients that may characterize the kpc scale jets of M87 and 3C 264.  A velocity gradient across the jet would effectively reduce and/or inhibit a field component in the direction perpendicular to the velocity gradient, and thus would allow for a different polarization morphology in the jet center, as observed here.  In this scenario,  the high energy emission would come primarily from the fast moving spine, while the low energy emission would be dominated by emission from the slower sheath (Ghisellini et al. 2005). For small angles $\\theta<1/\\Gamma_{sp}$, where $\\Gamma_{sp}$ is the bulk Lorentz factor of the spine, the observed cooling break is the canonical $1/2$ break predicted by simple theory. At larger angles, however, the high energy emission of the spine is beamed out of our line of sight more than the sheath emission, and this results in spectral breaks steeper that $1/2$ (see Figure 2 of Georganopoulos \\& Kazanas 2003).   It is important to point out that these two possible explanations are not mutually exclusive;  indeed, this can be seen by the breadth of work that has developed on the SED of the M87 jet, where the spectrum continues to steepen through the X-rays.  This caused Perlman \\& Wilson (2005) to invoke particle acceleration in an energy-dependent part of the cross-section of the jet to explain this, while other authors (Fleishman 2006, Liu \\& Shen 2007, Sahayanathan 2008) have  expanded on this by invoking two-zone, and in one case a diffusive, shock models.  Work on other jet sources has expanded upon this by invoking velocity shear and deceleration (NGC 315: Worrall et al. 2007, 3C 31:  Laing et al. 2008;  see also Wang et al. 2009). In these sources, both toroidal and longitudinal fields are needed to account for details of the polarization morphology at the jet edges.  These last two sources are rather similar in flux and polarization morphology to 3C 264, albeit longer and without optical emission from their jets.  This is particularly true when the details of the F330W polarization maps (Figure 9) are taken into account. Note, however, that while the work to date suggest that particle acceleration is enhanced in shock-like knots (e.g., Perlman \\& Wilson 2005), a continuous mode of particle acceleration also appears required for most X-ray jets (see e.g., Perlman \\& Wilson 2005; Jester et al. 2001, 2005, 2006; Hardcastle et al. 2003, 2006, 2007).  Another issue is orientation.  For M 87, this is reasonably well constrained now, as the superluminal motions (Biretta et al. 1995, 1999, Cheung et al. 2007) seen in its jet restrict the angle to our line of sight to values $\\leq 18^\\circ$ (although n.b., this is in conflict  with estimates derived from the observed geometry of the disk and knot A, e.g., Ford et al. 1994).  A similarly small viewing angle is seen in the 3C 346 jet, which Dulwich et al. (2009) has constrained to be at an orientation of 14$^{+8}_{-7}$ degrees, before bending 22 degrees at an oblique shock several arcseconds down the jet.  This value is consistent with earlier VLBI analysis as well (Cotton et al. 1995).  By comparison,      for 3C 264, modeling  suggests an orientation angle as large as $\\theta\\sim 50^\\circ$ (Baum et al 1997, Lara et al. 1997), based on the jet to counterjet ratio as well as VLBI and ROSAT data.  No firm constraints exist on the orientation of the 3C 15 jet, the only other one for which a detailed comparison of radio  and optical polarimetry has been published to date (Dulwich et al. 2007). As was argued by Perlman \\& Wilson (2005) in the case of M87 most of the synchrotron optical and X-ray emission comes from the spine of the jet. If we assume that the high energy emitting spine is also faster than the mostly radio emitting sheath, a result supported by numerical simulations of relativistic jets (Aloy et al 2003), then for sources at greater angles to the line of sight, as possibly is 3C 264, the high energy emission from the fast spine will be more de-amplified  than the slower sheath emission. In this scenario, most of the radio emission comes from the slower sheath, independent of orientation, and this may explain the similar radio polarization of 3C 264 and M 87 as shear due to transverse velocity gradients in the plasma is likely to cause the magnetic field to align with the jet.  The difference between highly aligned (e.g., M87) and less aligned (e.g., 3C 264) sources appears at higher energies.  In the case of highly aligned jets, one expects that the optical and X-ray emission we observe is dominated by the fast spine, where a significant part of the high energy particle acceleration takes place, most probably in shocks manifested through magnetic field orientations perpendicular to the jet axis -- as indeed is seen particularly in the polarization data for the inner knots, e.g., Perlman et al. (1999), Perlman \\& Wilson (2005). However, in 3C 264, which is seen at a considerably larger angle, any high-energy emission coming from the fast spine will be de-amplified and thus the X-ray emission we see is most likely originating in the slower sheath.  This may explain the fact the optical polarization of 3C 264 is in general parallel to the jet axis, following the radio jet polarization pattern.  It also suggests that 3C 264 requires a mechanism for distributed particle acceleration, as suggested also for the much more powerful jet of 3C 273 by Jester et al. (2007). This scenario needs testing with optical and radio polarimetry of additional jets.   \\subsection {Summary}  We have presented new VLA, {\\it Chandra} and {\\it HST} imaging and polarimetric data on the jet of 3C 264 that allow us to make a comprehensive, multiband study of the energetics and structure of this jet, and the relationship of these factors to the high-energy emission.  We present polarization maps in the radio and optical, as well as total flux and spectral index maps in both bands.  We also present the first X-ray image of the jet and discuss its relationship to the structure seen in other bands.  The dominant emission mechanism of the jet in the X-rays is found to be synchrotron emission. We find significant curvature of the jet spectrum both between the radio and optical and also within the optical-UV band. We present a model of the jet emission that includes a smooth extrapolation of the optical-UV spectrum into the X-ray with a  break at particle Lorentz factors $\\gamma = 1.6 \\times 10^6$ that is steeper than the canonical 0.5.   We discuss mechanisms for producing the observed high-energy emission and the anomalously large spectral breaks observed in 3C 264 and M87, and we suggest that it can be due to differential losses, velocity gradients and/or magnetic field gradients coupled with more effective particle acceleration taking place in the faster and/or higher magnetic field parts of the flow.  It is, however, important to note that a full test of these mechanisms would require higher-angular resolution data than currently available, as {\\it Chandra} barely resolves the jet emission in the X-rays.   In the polarization data,  we find a smooth yet complex magnetic field structure in the jet which, while without exhibiting sharp, 'knotty' structures (as seen in some other high-resolution polarization maps, e.g., M87), displays a structure that is correlated with changes in the optical spectrum.    While such a configuration is seen also in other jets, such as M87, here the orientation is rather different, with 3C 264 being at a rather larger angle to our line of sight.  Thus in 3C 264 the spine-sheath configuration manifests itself in the form of more subtle differences between optical and radio polarization.   These include significantly higher polarization at the jet edges, as well as twists in the magnetic field orientation near the positions of two knots and in the interknot regions, seen particularly in the optical polarization data in the knot A/B region and near knot D.  Such a configuration is also consistent with the observed spectral energy distribution of 3C 264, but requires a distributed particle acceleration mechanism."
    },
    "0911/0911.1210_arXiv.txt": {
        "abstract": "{We report an inter-comparison of some popular algorithms within the artificial neural network domain ( viz., Local search algorithms, global search algorithms, higher order algorithms and the hybrid algorithms) by applying them  to  the standard benchmarking problems  like the IRIS data, XOR/N-Bit parity and  Two Spiral. Apart from giving  a brief  description  of these algorithms, the results obtained   for  the above  benchmark  problems  are presented in the  paper. The results  suggest  that  while  Levenberg-Marquardt algorithm yields  the lowest RMS error for the   N-bit Parity and the Two Spiral problems, Higher Order Neurons algorithm gives the best results  for the IRIS data problem. The best results for the XOR problem are  obtained  with the  Neuro Fuzzy  algorithm. The  above  algorithms were also applied for  solving  several  regression problems such  as cos(x)  and  a few  special  functions like the Gamma function, the complimentary Error  function  and  the   upper tail cumulative  $\\chi^2$-distribution function. The results of these  regression problems   indicate that, among  all the ANN algorithms used in the present study,  Levenberg-Marquardt algorithm  yields  the best  results. Keeping in view  the highly non-linear behaviour and  the wide  dynamic range of these functions, it is suggested that these functions  can be also considered as standard benchmark  problems for function approximation using artificial neural networks. } ",
        "introduction": "\\label{1} An artificial neural network (ANN) is an interconnected group of artificial neurons that use a mathematical model for information processing  to accomplish variety of tasks [1]. They  can be configured in various arrangements to perform a range of tasks including pattern recognition and classification [2].  While the theory and the implementation of ANN has been around for more than 50 years, it is only recently that  they have found wide spread practical applications. This is primarily  due to the advent of high speed, low cost computers that can support the rather computationally intensive requirement of an ANN of any real complexity. \\par Artificial neural networks  have been  used  successfully in various fields, like, pattern recognition, financial analysis, biology, engineering and so on,  because  of  their   merits  such as self-learning, self-adapting, good robustness and  capability of dealing with non-linear problems. They  have  also been  employed  extensively  in several branches  of  astronomy  for  automated  data   analysis  and  other  applications  like   star/galaxy classification,  time series  analysis (e.g  prediction of solar activity) , determination of  photometric  redshifts,   characterization  of peculiar   objects  such as  QSO's,  ultraluminous  IR galaxies[3][4].   With the increase of quantity and the distributing complexity of astronomical data,  its scientific exploitation requires a variety of automated tools, which are capable   of  performing  variety of tasks , such as data preprocessing, feature selection, data reduction, data mining  and data analysis[5].   In some recent applications the IUCAA  group (and their  collaborators) have used ANNs with   remarkable   success  for   problems   like  star/galaxy classification,  stellar spectra classification  etc.  Employing   principal  component  analysis (PCA)  for  reducing  the dimensionality of the data   and   a  multilayer  backpropagation  network based  ANN scheme, a fast  and robust method  has  been  developed  in [6]  for  classifying  a library of  optical stellar spectra for  O to M type stars.  It has been demonstrated in their work that the   PCA   when  combined  with ANN   reduces  the network  configuration  (and  hence computational time)   drastically   without  compromising on the classification accuracy.  An automated  classification  scheme  based  on the   above idea  has been also used   for  classifying    1273  stars   in the  CFLIB  data base [7]   with  an  added  advantage  that   by employing    a   generalized  PCA  technique,  the  authors  were  able    to restore   missing  data  for   300 stars.   A supervised  back-propagation  algorithm  was  used  to classify 2000 bright sources  from the   Calgary database  of IRAS (Infrared Astronomical  Satellite) spectra  into  17 predefined  classes  and   a success rate of   80$\\%$ has  been  reported  by  the authors  [8]. Stellar  spectra classification   using   the probabilistic  Neural Network (PNN)  for  automated  classification of  about  5000   Solan Digital Sky  Survey (SDSS)  spectra  into  158 spectral  classes  has   also   been  performed  in [9]   with some encouraging  results.   The use of  ANN-based   technique  to develop  a  pipeline for  automated   segregation of  star/galaxies   to be observed  by the  Tel-Aviv University Ultra-Violet Experiment  (TAUVEX)  is also validated  by using  synthetic  spectra  in the  UV  region  as the  training  set  and  International Ultraviolet  Explorer (IUE) low  resolution  spectra  as the test  set [10][11]. \\par An  important  research activity in the field of  neural networks is  to compare the performance of  different ANN algorithms  on benchmark problems and to develop more efficient algorithms for solving real world problems  with noisy  and scarce data. It has been  also noticed  by  several workers,  that neural network algorithms are often benchmarked rather poorly [12].  More importantly,  it is also found in the literature   that   performance  of any algorithm is  only compared  to  the standard backpropagation algorithm [13]  even though  there  are  several powerful and  widely  used  algorithms   readily  available now.   Keeping this in mind, we  present  a detailed  study in  this paper  where  the   performance   of three generations of neural network algorithms i,e Ist order algorithms (Standard Backpropagation and Resilient Backpropagation), IInd order algorithms (Conjugate Gradient, Levenberg-Marquardt, Radial Basis Function, Simulated Annealing),  the Hybrid models like the Higher Order Neuron model and  Neuro-Fuzzy system, is examined   by applying  them to  standard  benchmarking problems  like IRIS data, XOR/N-Bit parity and  Two Spiral data.  In addition to benchmark problems discussed above,  we have also applied  the above mentioned  neural  network algorithms  for   solving  several  regression problems such  as cos(x)  and   a few  special  functions like the Gamma function, the complimentary Error  function  and  the upper tail cumulative  $\\chi^2$-distribution function.  A  short  introduction to  ANN methodology and a brief  description of the ANN algorithms used  in the present  work   has also been  presented  in the paper so that  the manuscript can be  easily followed  by  researchers who are not experts in the field of  neural networks. ",
        "conclusions": "Artificial  neural network   algorithms  have been applied to a variety of problems in various  diverse areas of physics, biology, medicine, computer research etc. The main aim of most of these studies has been to use ANN-based algorithms (generally standard backpropagation) as an  alternative method  to conventional analysis  for achieving  better results.  While  comparative performance of  some  ANN algorithms  like standard backpropagation, fuzzy logic, genetic algorithms, fractals etc., has been studied  for various  applications, a rigorous  intercomparison of  some of powerful algorithms ( e.g. the ones studied in this work)  is still  missing from the literature.  The primary aim of this  work  has been to provide a rigorous comparative study of various  powerful algorithms, by first applying them to standard benchmark problems  and  then apply them  as a regression tool  for  approximating  functions like cos(x) and  a few special functions. Our  results  suggest  that  while  Levenberg-Marquardt algorithm yields  the lowest RMS  error for the   N-bit Parity and the Two Spiral problems, Higher Order Neurons algorithm gives the best results  for the IRIS data problem. The best results for the XOR problem are  obtained  with the  Neuro Fuzzy  algorithm.  It is worth mentioning  here  that  benchmark  problems ( IRIS, XOR/N-Bit Parity and 2-Spiral)  have been also studied by numerous other workers.  For example, using a 2:2:1 configuration for the XOR problem,  Wang [28] has reported that one can achieve an accuracy of  $\\sim$80$\\%$ with $\\sim$5000 epochs of training. Other benchmark problems like Parity and cos(x)  has been also studied by the same author, but comparison is done only for the backpropagtion and simulated annealing techniques. Likewise, the 2-Spiral problem has been studied by several workers  using different algorithms like  vanilla backpropagaion [36] with configuration of 2:5:5:1, generalised regression model [37], vector quantization method [38], input coding scheme [39].  Complicated ANN configurations like 2:20:20:1 for the 2-Spiral problem with 50000 training epochs   and 4:4:2:1 for the IRIS problem  with 30000 training epochs has also been attempted  by   Lee [40] for achieving  reasonably accurate results  for these benchmark problems. \\par Regarding  application of  neural  network algorithms  for  solving   regression problems,  such as  evaluation of special  functions like the Gamma function, the complimentary Error  function  and  the   upper tail cumulative  $\\chi^2$-distribution function, we believe that such  an  attempt  has been made  in this work for the first time. The results obtained  in this work  indicate that, among  all the ANN algorithms used in the present  study,  Levenberg-Marquardt algorithm  yields  the  best  results. Conventionally,  two groups of approximations are found in the literature which are used for  calculating these special functions.   One group consists of standard numerical algorithms  which, at least in theory, allow computation of the above integral to arbitrarily high  precision. However, computation using  these numerical  algorithms requires massive computation.  The  second group consists of so called \"ad hoc  approximations\"  which require  only a few carefully chosen numeric constants. Importantly, a serious limitation of most of the approximations in both the groups is that they are designed to work  in a predefined range of input parameter  values and the accuracy of the approximation rapidly   deteriorates   when  the input parameters  takes a  value outside  the predefined  range. \\par In order to appreciate the complexity of evaluating the special functions studied in this paper, it is worth discussing  here  some of important approximations used for these functions.  A well known  method  for calculating the gamma function is the so called  Lanczos approximation [41] which computes the value of $\\Gamma(z)$  any positive real argument(z)  with a high level of accuracy. Likewise, a compilation of useful approximations used for evaluating the upper tail integrals for the Gaussian and $\\chi^2$ distributions  can be found in [42] and [43], respectively. \\par Although the comparative performance  of different ANN algorithms is in general problem dependent, we feel that the  study undertaken in this paper  does  give an insight into the power of various  powerful ANN algorithms.  Since for real world problems it is not an easy task to  identify the most  suitable  ANN algorithm by just having a look at the  problem, our results  suggest  that  while  investigating the comparative performance  of  other ANN algorithm,  the  Levenberg-Marquardt algorithm  deserves  a  serious   consideration  and  cannot  be  rejected  outright  because  of  its  training  time  overheads."
    },
    "0911/0911.1998_arXiv.txt": {
        "abstract": "We present new Keck/DEIMOS spectroscopic observations of hundreds of individual stars along the sightline to Andromeda's first three discovered dwarf spheroidal galaxies (dSphs) -- And~I, II, and III, and leverage recent observations by our team of three additional dSphs, And~VII, X, and XIV, as a part of the SPLASH Survey (Spectroscopic and Photometric Landscape of Andromeda's Stellar Halo).  Member stars of each dSph are isolated from foreground Milky Way dwarf and M31 field contamination using a variety of photometric and spectroscopic diagnostics.  Our final spectroscopic sample of member stars in each dSph, for which we measure accurate radial velocities with a median uncertainty (random plus systematic errors) of 4 -- 5~km~s$^{-1}$, includes 80 red giants in And~I, 95 in And~II, 43 in And~III, 18 in And~VII, 22 in And~X, and 38 in And~XIV.  The sample of confirmed members in the six dSphs are used to derive each system's mean radial velocity, intrinsic central velocity dispersion, mean abundance, abundance spread, and dynamical mass.  This combined data set presents us with a unique opportunity to perform the first systematic comparison of the global properties (e.g., metallicities, sizes, and dark matter masses) of one-third of Andromeda's total known dSph population with Milky Way counterparts of the same luminosity.  Our overall comparisons indicate that the family of dSphs in these two hosts have both similarities and differences.  For example, we find that the luminosity -- metallicity relation is very similar between $L$ $\\sim$ 10$^{5}$ -- 10$^{7}$ $L_\\odot$, suggesting that the chemical evolution histories of each group of dSphs is similar.  The lowest luminosity M31 dSphs appear to deviate from the relation, possibly suggesting tidal stripping.  Previous observations have noted that the sizes of M31's brightest dSphs are systematically larger than Milky Way satellites of similar luminosity.  At lower luminosities between $L$ = 10$^{4}$ -- 10$^{6}$~$L_\\odot$, we find that the sizes of dSphs in the two hosts significantly overlap and that four of the faintest M31 dSphs are {\\it smaller} than Milky Way counterparts.  The first dynamical mass measurements of six M31 dSphs over a large range in luminosity indicates similar mass-to-light ratios compared to Milky Way dSphs among the brighter satellites, and smaller mass-to-light ratios among the fainter satellites.  Combined with their similar or larger sizes at these luminosities, these results hint that the M31 dSphs are systematically less dense than Milky Way dSphs.  The implications of these similarities and differences for general understanding of galaxy formation and evolution is summarized. ",
        "introduction": "\\label{intro}  Dwarf spheroidal (dSph) galaxies represent important laboratories for constraining feedback processes in galaxies and for testing cosmological models on the smallest scales (Klypin et~al.\\ 1999; Moore et~al.\\ 1999; Bullock, Kravtsov, \\& Weinberg 2000).  These low luminosity systems ($-1.5 \\lesssim M_V \\lesssim -13$) show no evidence of ongoing star formation, contain little or no interstellar matter, and are almost always distributed around massive hosts, such as the Milky Way.  Within the framework of the $\\Lambda$CDM paradigm, theories suggest that the observed dSph satellites are the present-day counterparts to previously accreted objects that were tidally destroyed to build the spheroids and stellar halos of massive galaxies (Bullock et~al.\\ 2001; Bullock \\& Johnston 2005).  This process has been directly verified in the outskirts of giant galaxies with the discovery of luminous substructure in the form of tidal streams.  For example, the accretion mechanism, where these dwarf galaxies are shredded and eventually melded into diffuse stellar components, is abundant in the Milky Way (e.g., the Sagittarius, Orphan, and Cetus Polar Streams -- Ibata, Gilmore, \\& Irwin 1994; Belokurov et~al.\\ 2007; Newberg, Yanny, \\& Willet 2009), in the Andromeda (M31) spiral galaxy (e.g., the Giant Southern Stream -- Ibata et~al.\\ 2001), as well as in several other systems such as NGC~5907 (Martinez-Delgado et~al.\\ 2008) and NGC~4013 (Martinez-Delgado et~al.\\ 2009).  Despite representing the most common morphological class among all galaxies in the Universe, the overall count of known dSph satellites around the Milky Way and M31 are significantly lower than first-order expectations.  Simulations predict large galaxy halos to host {\\em thousands} of dark matter substructures with masses similar to those of dSph galaxies (Klypin et~al.\\ 1999; Diemand, Kuhlen, \\& Madau 2007; Springel et~al.\\ 2008).  While models exist for explaining the mismatch, they can only be tested with precise mass measurements of the observed dwarfs. Recently, large observational efforts have successfully uncovered many new dSphs around both the Milky Way (e.g., Willman et~al.\\ 2005; Zucker et~al.\\ 2006a; 2006b; Sakamoto \\& Hasegawa 2006; Belokurov et~al.\\ 2006; Belokurov et~al.\\ 2007; Irwin et~al.\\ 2007; Walsh, Jerjen, \\& Willman 2007; Liu et~al.\\ 2008; Belokurov et~al.\\ 2008; Belokurov et~al.\\ 2009) and M31 (Martin et~al.\\ 2006; Zucker et~al.\\ 2007; Ibata et~al.\\ 2007; Majewski et~al.\\ 2007; Irwin et~al.\\ 2008; McConnachie et~al.\\ 2008; Martin et~al.\\ 2009), effectively doubling the total sample size in the past few years.  Follow up spectroscopic observations of individual stars in the new Milky Way satellites (e.g., by Kleyna et~al.\\ 2005; Munoz et~al.\\ 2006; Simon \\& Geha 2007, Martin et~al.\\ 2007; Koch et~al.\\ 2009; Geha et~al.\\ 2009) have established that these systems are the most dark matter dominated objects known, and also provided a means to test proposed solutions of the ``missing satellites problem'' (see also Strigari et~al.\\ 2008a; 2008b).  Unfortunately, it is difficult to draw general conclusions based on the Milky Way system alone.   \\begin{table*} \\begin{center} \\caption{} \\begin{tabular}{llccccc} \\hline \\hline \\multicolumn{1}{c}{Date} & \\multicolumn{1}{c}{Mask} & \\multicolumn{2}{c}{Pointing center:} & \\multicolumn{1}{c}{Field PA} & \\multicolumn{1}{c}{Exp.} & \\multicolumn{1}{c}{No. Sci.}  \\\\ & \\multicolumn{1}{c}{} &\\multicolumn{1}{c}{$\\alpha_{\\rm J2000}$} & \\multicolumn{1}{c}{$\\delta_{\\rm J2000}$} &\\multicolumn{1}{c}{($^\\circ$E of N)} & \\multicolumn{1}{c}{Time} & \\multicolumn{1}{c}{Targets\\tablenotemark{1}} \\\\ & \\multicolumn{1}{c}{} & \\multicolumn{1}{c}{($\\rm^h$:$\\rm^m$:$\\rm^s$)} & \\multicolumn{1}{c}{($^\\circ$:$'$:$''$)} & & \\multicolumn{1}{c}{(s)} & \\\\ \\hline 2005 Nov 05  & d1\\_1 (And~I)    & 00:45:48.61  &  +38:05:46.8 & $+0.0$  & 3 x 1,200 & 153   \\\\ 2006 Sep 16  & d1\\_2 (And~I)    & 00:46:13.95  &  +38:00:27.0 & $+90.0$ & 3 x 1,200 & 150   \\\\ 2005 Sep 06   & d2\\_1 (And~II)  & 01:17:07.46  &  +33:29:25.1 & $-90.0$ & 3 x 1,200 & 139  \\\\ 2005 Sep 06   & d2\\_2 (And~II)  & 01:16:43.29  &  +33:34:25.8 & $+0.0$  & 3 x 1,200 & 139  \\\\ 2005 Sep 08   & d3\\_1 (And~III) & 00:36:03.83  &  +36:27:27.4 & $+90.0$ & 3 x 1,200 & 128  \\\\ 2005 Sep 08   & d3\\_2 (And~III) & 00:35:39.61  &  +36:21:41.8 & $+0.0$  & 3 x 1,200 & 126  \\\\ \\hline \\end{tabular} \\tablenotetext{1}{10 of the stars in And I, 17 in And II, and 15 in And III were observed on both masks.} \\label{table:masks1} \\end{center} \\end{table*}  Recently, a number of large surveys have targeted the outskirts of M31 with wide field imagers on 4-meter class telescopes (e.g., Ferguson et~al.\\ 2002; Ostheimer 2003; Ibata et~al.\\ 2007; McConnachie et~al.\\ 2009).  These studies, combined with spectroscopic follow up such as the SPLASH Survey (Spectroscopic and Photometric Landscape of Andromeda's Stellar Halo), have resulted in ground-breaking discoveries related to the properties of M31's stellar halo, which itself was discovered in the surveys \\citep{guh05,irwin05}.  For example, the field halo of M31 exhibits a power-law surface brightness profile similar to the Milky Way \\citep{ostheimer03,guh05,irwin05}, is filled with substructure \\citep{ibata07,gilbert07,gilbert09}, and is metal-poor in its outskirts \\citep{kalirai06a,chapman06,koch08}.  These recent surveys have also led to the discovery of new dSph galaxies in M31, over 35 years after \\cite{vandenbergh72} found the first such systems. Similar to the discovery rate in the Milky Way, the past five years have led to thirteen new galaxies, And~IX \\citep{zucker04}, And~X \\citep{zucker07}, And~XI--XIII \\citep{martin06}, And~XIV \\citep{majewski07}, And~XV--XVI \\citep{ibata07}, And~XVII \\citep{irwin08}, And~XVIII--XX \\citep{mcconnachie08}, and And~XXI--XXII \\citep{martin09}.\\footnote{And~XVIII is a distant background Local Group dSph.}  The total population of known dSphs is therefore similar in M31 and the Milky Way, and considering the spatial coverage of the surveys in both galaxies (e.g., the PAandAS Survey -- Ibata et~al.\\ 2007; McConnachie et~al.\\ 2008; McConnachie et~al.\\ 2009), the overall census is incomplete by at least a factor of a few.  Given the distance to M31, 780~kpc, most of our information on the properties of its dSph population has come about from studies of the bright red giant branch (RGB) stars in these systems.  For example, \\cite{mcconnachie06} present a detailed analysis of the structural properties of six of these galaxies and find both similarities and differences with the Galactic population.  For example, similar to the Galactic system, M31 dSphs with higher central surface brightnesses  are found further from their host.  A key difference however is that the scale radii of the M31 dSphs is approximately twice as large as the Milky Way dSphs at the same luminosity. \\cite{mcconnachie06} also show that the central surface brightness of the M31 satellites is smaller than that of comparable Milky Way systems.  The only Hubble Space Telescope {\\it HST} observations of the M31 satellites are the WFPC2 studies of Da Costa et~al.\\ (1996; 2000; 2002), who obtained optical photometry of stars in each of And~I, II, and III down to the level of the horizontal branch.  Their study provides the first constraints on the age and abundance distributions of stars in these three systems.  Indeed, Da Costa et~al.\\ find that, similar to the Milky Way dSphs, the M31 satellites have also had diverse evolutionary histories.  Unlike for the Milky Way satellites, there exist very few {\\it spectroscopic} studies of the M31 dSph system, most of which were only able to characterize a few stars in each galaxy.  \\cite{cote99b} obtained the first such spectra using HIRES on Keck.  They measured velocities for seven RGB stars in And~II and estimated both the radial velocity of the galaxy and the central velocity dispersion from this small sample.  \\cite{guh00} also obtained spectra of a few tens of individual RGB stars in each of And I, III, V, and VI.  Their study led to a dynamical mass estimate of M31 based on the mean radial velocity of each dSph \\citep{evans00}. Unfortunately, the data set could not be used to probe the internal kinematics of the satellites given low numbers of stars in each galaxy, and large uncertainties in the velocity measurements.  \\cite{guh00} also looked at several stars in And~VII with HIRES on Keck and measured both the radial velocity and the total velocity dispersion of the galaxy (e.g., the combined true intrinsic dispersion plus the dispersion due to velocity errors).  Recently, \\cite{chapman05} and \\cite{chapman07} have used Keck/DEIMOS to spectroscopically confirm six to eight members of And~IX and And~XII, however these data do not constrain the internal kinematics of the dSph.  For example, the inclusion of one marginal member in the outer radial bin of And~IX inflates the velocity dispersion by a factor of two.  A re-analysis of these data by \\cite{collins09} establishes upper limits on the velocity dispersions of both galaxies\\footnote{At the time of the writing of this paper, the \\cite{collins09} study appears as a preprint.  The 1-$\\sigma$ error bars on the velocity dispersions are consistent with $\\sigma_v$ = 0~km~s$^{-1}$ for both dSphs.}, and they also do not resolve the dispersions of And~XI and XIII. Finally, \\cite{letarte09} have presented a new analysis of seven and eight spectroscopically confirmed stars in And~XV and And~XVI.  The radial velocity uncertainties in these measurements are large, ranging from $\\sigma_v$ = 6 -- 25~km~s$^{-1}$, with one star as high as $\\sigma_v$ = 52~km~s$^{-1}$.  Summarizing, of the entire 18 galaxies comprising the known M31 dSph population, And~II and And~VII are the only systems with secure velocity dispersion measurements, based on high resolution studies of 7 and $\\sim$20 stars each.  In this paper, we present the first results from a large project aimed at spectroscopically observing hundreds of stars in the M31 dSph system, as a part of the overall SPLASH Survey.  These data will allow the first detailed comparison of the properties of M31 satellites to the Milky Way dSph system.  For example, we will accurately measure velocities of (a minimum of) dozens of individual stars in each galaxy, thereby resolving their intrinsic central velocity dispersions and leading to estimates of the total system masses.  For the best data sets ($>$500 velocities in a single dSph), we will also probe for any radial changes in the velocity dispersion, and test for rotation, thereby constraining dynamical models for dSphs (e.g., Kroupa 1997).  Such high quality data sets can also be used to better understand how luminosity maps to dark matter subhalo mass.  In addition to the kinematical analysis, spectroscopic observations of large numbers of stars in each dSph can provide a detailed analysis of radial trends in the chemical abundances of stars, and perhaps even constrain the $\\alpha$-abundance of these satellites.  Therefore, these observations directly relate to a better understanding of both population gradients and the evolutionary history of dSphs.  Comparisons of the results from the survey with the well studied dSphs of the Milky Way will aid in our understanding of whether the observed differences in the satellite systems (see \\S\\,\\ref{globalproperties}) are caused by in situ processes within the galaxies, or whether they are shaped by different interaction histories of the dSphs with their massive hosts.  In this first paper, we address the internal kinematics, chemical compositions, sizes, and total masses of And I, II, III, VII, X, and XIV, based on a large sample size of individual radial velocities.  Future papers will address other M31 dSphs, for which we have already collected a substantial amount of Keck/DEIMOS observations.  The layout of this paper is as follows.  We describe the photometric and spectroscopic observations of these dSphs as well as the data reduction in the next section.  This includes a detailed discussion of the uncertainties in the velocity measurements, which are critical to understand in order to resolve the intrinsic dispersions of the (kinematically cold) satellites.  Membership of stars in the satellites is established and velocity histograms of the confirmed red giants are presented in \\S\\,\\ref{membership}.  The kinematics are analyzed to yield mass-to-light ratios for the galaxies in \\S\\,\\ref{meanvelocity}. Next, we measure the chemical abundance of each dSph in \\S\\,\\ref{chemabund}.  Unlike previous studies, we measure the abundances of each confirmed dSph star using two independent methods, photometrically and spectroscopically.  The results from this overall study are discussed in \\S\\,\\ref{globalproperties}, where we compare the properties of all six M31 dSphs (e.g., their sizes, metallicities, and masses) with measured trends for Milky Way satellites.  The main results from the paper are summarized in \\S\\,\\ref{conclusion}.   \\begin{figure*} \\begin{center} \\leavevmode \\includegraphics[width=5.4cm,angle=270]{radecall.eps} \\end{center} \\caption{The distribution of stars from the KPNO wide-field Mosaic observations of each dSph are shown as small points.  The selection function used to isolate these stars involved a simple cut on the color-magnitude diagram (CMD) to enhance the dSph RGB with respect to the field population, which is still abundantly present in the maps.  Superimposed on the photometric data are the objects targeted with Keck/DEIMOS spectroscopy (darker points).} \\label{fig:radecall} \\end{figure*} ",
        "conclusions": "\\label{conclusion}  We present the first systematic large scale spectroscopic survey of M31's dSphs as a part of the SPLASH Survey, a detailed study of And~I, II, III, VII, X, and XIV.  The fundamental properties of these galaxies are established from a sample of confirmed member stars, including the mean radial velocity, velocity dispersion, mean abundance, abundance spread, and total (stellar + dark matter) mass.  The kinematics indicate that the radial velocity and intrinsic velocity dispersion of the six satellites are $v_{\\rm rad}=-375.8\\pm1.4$~km~s$^{-1}$ and $\\sigma_{v}=10.6\\pm1.1$~km~s$^{-1}$ for And~I, $v_{\\rm rad}=-193.6\\pm1.0$~km~s$^{-1}$ and $\\sigma_{v}=7.3\\pm0.8$~km~s$^{-1}$ for And~II, $v_{\\rm rad}=-345.6\\pm1.8$~km~s$^{-1}$ and $\\sigma_{v}=4.7\\pm1.8$~km~s$^{-1}$ for And~III, $v_{\\rm rad}=-309.4\\pm2.3$~km~s$^{-1}$ and $\\sigma_{v}=9.7\\pm1.6$~km~s$^{-1}$ for And~VII, $v_{\\rm rad}=-168.3\\pm1.2$~km~s$^{-1}$ and $\\sigma_{v}=3.9\\pm1.2$~km~s$^{-1}$ for And~X, and $v_{\\rm rad}=-481.0\\pm1.2$~km~s$^{-1}$ and $\\sigma_{v}=5.4\\pm1.3$~km~s$^{-1}$ for And~XIV.  Both photometric and spectroscopic metallicities are measured for all confirmed member stars in each galaxy, and the two independent measurements agree nicely with one another.  We find that the six dSphs are metal-poor systems with a  mean iron abundance of [Fe/H]$_{\\rm phot}$ = $-1.45\\pm0.04$ ($\\sigma_{\\rm [Fe/H]}$ = 0.37) for And~I, [Fe/H]$_{\\rm phot}$ = $-1.64\\pm0.04$ ($\\sigma_{\\rm [Fe/H]}$ = 0.34) for And~II, [Fe/H]$_{\\rm phot}$ = $-1.78\\pm0.04$ ($\\sigma_{\\rm [Fe/H]}$ = 0.27) for And~III, [Fe/H]$_{\\rm phot}$ = $-1.40\\pm0.30$ for And~VII, [Fe/H]$_{\\rm phot}$ = $-1.93\\pm0.11$ ($\\sigma_{\\rm [Fe/H]}$ = 0.248) for And~X, and [Fe/H]$_{\\rm phot}$ = $-2.26\\pm0.05$ for And~XIV (where two of the satellites lack a metallicity dispersion measurement).  The dynamical mass of each satellite is estimated to first order using the \\cite{illingworth76} formalism.  More accurate mass measurements at the half-light radius of each galaxy are also calculated using the methods described in \\cite{wolf09}.  As summarized in Table~2 and 3, these yield half-light radius mass-to-light ratios of $M_{1/2}$/$L_{1/2}$ = 31 $\\pm$ 6 for And~I, 13 $\\pm$ 3 for And~II, 19 $\\pm$ 12 for And~III, 7.7 $\\pm$ 2.8 for And~VII, 63 $\\pm$ 40 for And~X, and 102 $\\pm$ 71 for And~XIV.  These calculations are found to agree nicely with masses measured through full modeling of the individual radial velocities assuming the dark matter follows a five parameter density profile.  The comparison of the overall properties of M31's dSphs with Milky Way dSphs of similar luminosity indicates some interesting similarities and differences.  For example, we find that the luminosity -- metallicity relation is very similar between $L$ $\\sim$ 10$^{5}$ -- 10$^{7}$ $L_\\odot$, suggesting that the chemical evolution histories of each group of dSphs is similar.  The lowest luminosity M31 dSphs appear to deviate from the relation, possibly suggesting tidal stripping.  The previous finding of McConnachie \\& Irwin (2006) indicating that the brightest M31 dSphs are systematically larger than Milky Way counterparts of the same luminosity does not persist cleanly to lower luminosities.  We find that the sizes of dSphs in the two hosts significantly overlap for $L$ = 10$^{4}$ -- 10$^{6}$~$L_\\odot$, and that four of the faintest M31 dSphs are {\\it smaller} than Milky Way satellites of similar luminosity.  The first dynamical mass measurement comparisons between the M31 and Milky Way dSphs hints that the M31 satellites are systematically less dense than Milky Way counterparts."
    },
    "0911/0911.1483_arXiv.txt": {
        "abstract": "\\label{abstract} We study the dynamical evolution of the young star cluster Arches and its dependence on the assumed initial stellar mass function (IMF). We perform many direct $N$-body simulations with various initial conditions and two different choices of IMFs. One is a standard Kroupa IMF without any mass segregation. The other is a radially dependent IMF, as presently observed in the Arches. We find that it is unlikely for the Arches to have attained the observed degree of mass segregation at its current age starting from a standard non-segregated Kroupa IMF.  We also study the possibility of a collisional runaway developing in the first $\\sim 2-3\\,\\rm{Myr}$ of dynamical evolution.  We find that the evolution of this cluster is dramatically different depending on the choice of IMF: if a primordially mass segregated IMF is chosen, a collisional runaway should always occur between $2-3\\,\\rm{Myr}$ for a broad range of initial concentrations. In contrast, for a standard Kroupa IMF no collisional runaway is predicted. We argue that if Arches was created with a mass segregated IMF similar to what is observed today then at the current cluster age a very unusual, high-mass star should be created.  However, whether a collisional runaway leads to the formation of an intermediate-mass black hole (IMBH) depends strongly on the mass loss rate via winds from massive stars.  Growth of stellar mass through collisions can be quenched by strong wind mass loss.  In that case, the inter-cluster as well as intra-cluster medium are expected to have a significant Helium enrichment which may be observed via Helium recombination lines.  The excess amount of gas lost in winds may also be observed via X-ray observations as diffused X-ray sources. ",
        "introduction": "\\label{intro} One of the biggest uncertainties in star cluster evolution studies lies in determining the true initial stellar mass function (IMF).  Traditionally all studies focusing on the dynamical evolution of star clusters assume an initially fully formed cluster with a given number of single and binary stars, all at zero age main sequence (ZAMS) at $t=0$ and no gas.  It is also often assumed that the IMF is a standard power-law (or power-laws with different indices in different mass ranges) with no primordial radial variation in the cluster \\citep[e.g., ][]{1959ApJ...129..608S,1979ApJS...41..513M,2001MNRAS.322..231K}.   However, star formation simulations indicate that neither of the above simplifications may be valid.  For example, star formation takes place over a time-period and the formation timescale depends on many different physical processes including the Jeans mass of the protostellar clouds, radiative feedback and turbulence \\citep[e.g., ][]{2006ApJ...641L.121T}.  Moreover, depending on these detailed physical processes stars of different mass ranges may form over different timescales \\citep[e.g., ][]{2007MNRAS.381L..40D,2009ApJ...699..850K}.  In addition, deviations from the standard MFs are observed in many clusters both at the high mass and the low mass ends \\citep[e.g., ][]{2004MNRAS.354..367E}.  In particular at the high mass end observed MFs are generally top-heavy compared to standard power-laws in dense cluster cores like the Arches cluster and starburst regions \\citep{2002A&A...394..459S,2006ApJ...653L.113K,2009MNRAS.394.1529D}. Mass segregation is also observed in old globular clusters \\citep[e.g., ][]{1993A&A...273..100S,1998ApJ...492..540H,2008ApJ...685..247B}. Whether the observed mass segregation is imprinted from the star formation epoch or it is attained via dynamical evolution of the cluster at a later time is still an open question. Although a conclusive general answer to this question does not exist, it has been shown in the case of the very young Trapezium cluster in the Orion nebula that the observed degree of mass segregation cannot be explained simply via dynamical mass segregation and a preferential formation of high-mass stars near the center is necessary to explain the observations \\citep{1998MNRAS.295..691B}.  In order to bypass the above mentioned assumptions a detailed model for the star formation processes during the formation of the cluster is needed. Interstellar clouds contract and fragment due to density perturbations to form pre-stellar cores (PSC) that collapse on timescales that depend on their Jeans masses to form stars \\citep[e.g., ][]{1966PThPh..36..515N}.  Depending on the compactness of the assortment of these PSCs they can merge with other PSCs and grow before they collapse to form stars \\citep[e.g., ][]{1979ApJ...229..242S,2003MNRAS.338..817E,2007ApJ...661..262D}.  Using this simple model \\citet{2007MNRAS.381L..40D} have attempted to eliminate the above-mentioned assumptions regarding cluster IMFs by building an IMF from a distribution of PSCs in a molecular cloud.  In their simplified model they assume that each PSC collapses to form a single star of mass proportional to the parent PSC. Comparing the collision timescale between the PSCs with the collapse timescale of the PSCs they show that near the central regions of a dense cluster the PSCs collide and grow in mass before they can collapse to form a star. Thus, near the core where the number density of the PSCs is significantly higher than that in the halo, it is more likely to form stars that are heavier than stars formed from PSCs in the halo.  Using the Arches cluster as a particular example they show that their semi-analytical model naturally results in a mass segregated cluster similar to the Arches cluster and the observed top-heaviness near the center is thus explained as imprinted from the star formation epoch.  The effects of primordial mass segregation on the global dynamical evolution of clusters have been explored previously in a parametric way showing interesting differences in the global properties of those clusters and their dynamical evolution \\citep[e.g.,][]{2008ApJ...685..247B,2004ApJ...604..632G}. The observed top-heavy MF in the central regions of the Arches cluster \\citep[e.g.,][]{2006ApJ...653L.113K,2007JKAS...40..153K} thus makes it an interesting candidate for such studies.  If the observed radially dependent top-heaviness in Arches MF is indeed primordial, then this is a direct measurement of the degree of mass segregation at formation of this cluster, thus it eliminates the necessity for parameterization.  The Arches cluster is also unique for the following reasons. The cluster has a projected distance of $\\sim 30\\,\\rm{pc}$ from the Galactic center \\citep{1995AJ....109.1676N,2002ApJ...565..265P,2006ApJ...653L.113K,2007JKAS...40..153K}.  Estimates from detailed proper motion observations with a baseline of $\\sim 4\\,\\rm{yr}$ indicate that the real Galactocentric distance of the Arches cluster is $\\approx 27\\,\\rm{pc}$ \\citep{2008ApJ...675.1278S}.  Thus the Arches is the closest massive cluster from the Galactic center. The proximity of the cluster from the Galactic center necessitates this cluster to be compact to be saved in the face of possible early tidal disruption. \\citet{1998Natur.394..448S} a decade ago found $\\gtrsim 100$ O stars within $0.6\\,\\rm{pc}$ of the cluster, indicating Arches to be the densest young cluster in the Galaxy.  The high estimated total cluster mass ($\\sim 10^4-10^5\\,\\rm{M_\\odot}$), and the high luminosity $10^8\\,\\rm{L_\\odot}$ thus make this young cluster a power house cluster for Galactic standards \\citep[e.g.,][]{1998Natur.394..448S}. The central stellar density of the Arches cluster is $\\sim 10^5\\,\\rm{M_\\odot/pc^3}$ \\citep[e.g.,][]{1998Natur.394..448S,2007MNRAS.378L..29P}, which is typical only for some old Galactic globular clusters.    The core radius ($r_c$) of the Arches cluster is $\\lesssim 0.2\\,\\rm{pc}$ \\citep[e.g.,][]{1999ApJ...525..750F,2002A&A...394..459S,2006JPhCS..54..217S}. The high central density and the high mass thus makes the Arches cluster a very good candidate to study for possibility of collisional runaway.  The Arches cluster is also very young, the cluster age ($t_{cl}$) is only $2\\pm1\\,\\rm{Myr}$ \\citep[e.g.,][]{2008ApJ...675.1278S}.  The timescale for a successful collisional runaway is only $\\sim 3\\,\\rm{Myr}$ since past that age, the cluster is depleted of its high-mass stars due to Supernova explosions \\citep[e.g.,][]{2006MNRAS.368..141F,2007ASPC..367..707F}. Hence, if conducive for a collisional runaway, the Arches cluster may be undergoing that process at the present time.  Two main scenarios to create the observed top-heavy radially dependent MF have been suggested in previous studies: 1. The observed MF is primordial, as mentioned earlier \\citep[e.g., ][]{2007MNRAS.381L..40D}.  2. The observed MF is a result of dynamical evolution starting from a standard IMF with no primordial radial variation \\citep[e.g.,][]{2007MNRAS.378L..29P}.  Using numerous numerical simulations we study the dynamical evolution of the Arches cluster assuming the observed MF was indeed primordial.  We compare the results from these simulations with the results from another set of simulations where the initial cluster properties are kept unchanged but a standard IMF with no radial variation is used instead.  We investigate whether it is possible to attain the observed degree of mass segregation starting from an unsegregated cluster. We verify that the MFs and cluster parameters of the simulated clusters at $t_{cl}\\approx2\\,\\rm{Myr}$ are consistent with observations. We then investigate the differences in the dynamical evolution of the simulated Arches-like clusters with the two above-mentioned choices of IMFs focusing on the possibilities of a collisional runaway.  Furthermore, we study the final fate and discuss the possible observational signatures of a collisional runaway depending on the stellar wind prescription.  The arrangement of this paper is as follows. In \\S\\ref{numerical} we describe our simulations.  In \\S\\ref{results} we summarize our key results.  This is followed by \\S\\ref{effects} where we discuss the key observable signatures of a collisional runaway.  We summarize and conclude in \\S\\ref{conclusion}. ",
        "conclusions": "\\label{conclusion} We carried out many $N$-body simulations with initial conditions chosen such as to resemble the Arches cluster. Using a standard \\citet{2001MNRAS.322..231K} IMF with no radial variation ({\\tt NSeg}) and a radially dependent IMF that corresponds to the level of mass segregation inferred from observations \\citep[{Seg}; e.g.][]{2007MNRAS.381L..40D}, we study possible consequences and discuss potential observable signatures to distinguish between the two cases.  We show that the degree of mass segregation via dynamical evolution of the simulated clusters during the first $2\\,\\rm{Myr}$ is minor both for the {\\tt Seg} and the {\\tt NSeg} IMF (Figure\\ \\ref{plot:MF}).  Thus, while runs with the {\\tt Seg} IMF reproduce well all parts of the currently observed slopes in the MF at $\\approx 2\\,\\rm{Myr}$, those with the {\\tt NSeg} IMF do not produce enough mass segregation within the Arches lifetime to match the observed MF as well (Figure\\ \\ref{plot:MF}).  Therefore, our results indicate that the present day observed degree of mass segregation in the Arches cluster cannot be a result of dynamics only and at least some degree of primordial mass segregation is required.  We further show that the choice of IMF changes the overall dynamical evolution significantly for Arches like clusters over a range of initial $W_0$.  For the {\\tt Seg} IMF there is clear evidence of collisional runaway, producing stars up to $\\sim 10^3\\ \\rm{M_\\odot}$ over a range of initial concentrations (Table\\ \\ref{tab:runs}, Figures\\ \\ref{plot:tvsm_w8-1}, \\ref{plot:tvsm_w7-1}, \\ref{plot:wvsm}).  In contrast, with the more standard {\\tt NSeg} IMF no collisional runaway is observed with the same initial conditions. We find that the sequence of collisions in the simulations with successful collisional runaway starts at around $2\\,\\rm{Myr}$.  Note that the current estimated age of the Arches is $2 \\pm 1\\,\\rm{Myr}$ \\citep[e.g., ][]{2002A&A...394..459S}. Hence, if the observed top-heavy MF within $r \\sim 2r_c$ is indeed primordial, then it is possible that a collisional runaway has started or is bound to happen.  Unfortunately, the uncertainty in the age estimation prevents us from predicting whether a collisional runaway should already have taken place in the Arches.  An interesting aspect of the collisional runaways observed in this study is the general tendency of a double runaway in cases where the collisional runaway is strong (e.g., runs {\\tt SW8-1,2} and runs {\\tt SW7-1,2}, see Table\\ \\ref{tab:runs}). Double collisional runaways have been observed and studied in numerical simulations before \\citep{2006ApJ...640L..39G}.  However, those simulations concluded that primordial binaries are a necessity for a double collisional runaway.  It is very interesting that with the primordially mass-segregated clusters even without any primordial binaries double collisional runaways can take place.  Due to the flatness in the high mass end of the IMF in this case, there is a larger reservoir of high-mass stars. Thus the probability to start a second runaway is higher. The ultimate fate of the smaller runaway is dependent on the statistical fluctuations of the simulation (see \\S\\ref{results}).  However, it is possible that the VMSs do not collide prior to compact object formation and both stay bound to the cluster after their respective supernova explosions (e.g., run {\\tt SW8-1}). In such a case they may form IMBH binaries or coalesce at the center of the cluster. If possible to form, this can create a very strong gravity wave signal for LISA \\citep{2006ApJ...640L..39G}.  If the mass loss from stellar evolution driven winds is too high then the collisional mass growth is quenched \\citep{2009A&A...497..255G}.  By evolving the collision products separately following each successive collision in a sequence of runaway collisions according to the prescription in \\citet{2009A&A...497..255G} we find that the VMSs are not produced as a result of a collisional runaway.  Instead, a Helium star of a few tens of a Solar mass is created (Figure\\ \\ref{plot:wind}).  The mass lost via these winds is extremely Helium rich, with $X_{He, S}$ reaching beyond $0.9$.  In this case the cluster as well as its surrounding medium will be enriched by this Helium rich gas, which could be observed through recombination at distances beyond $3-7\\,\\rm{pc}$ from the Arches cluster.  Further indications of a collisional runaway might be provided by the runaway object itself, as such an object has a rather unusual radius and mass loss rate compared to ordinary massive stars. Constraints on the radius, mass loss rate, and the wind speeds may be obtained using the Baldwin effect \\citep[e.g.,][]{1977ApJ...214..679B,2001A&A...372..963V}.  Furthermore, during the evolution of the runaway star, a Pistol like stellar object may also be produced (\\S\\ref{pistol}).  However, since the duration for which the star attains Pistol like high luminosities is short, it is not possible to predict whether a Pistol like star should already have been created in Arches given the uncertainties in the age estimation and the statistical fluctuations in the exact onset and the details of the simulated collisional runaways.  Nevertheless, we should point out that the above recipe to follow the runaway star for these high stellar winds is not completely self-consistent.  Each collision in a cluster is dependent on the masses and radii of the collision progenitors.  Since, severe wind mass loss from the stars changes the masses of the stars significantly, the chain of collisions in the runaway may also be different. However, this model still provides us with an estimate of the total mass loss from the VMS via stellar winds, its stellar properties, as well as the composition of the ejected material if a runaway did indeed happen.  We thank Farhad Yusef-Zadeh for helpful suggestions.  This research was supported by NASA Grant NNX08AG66G and NSF Grant AST-0607498 at Northwestern University.  This research was partly done at KITP while the authors participated in the spring 2009 program on ``Formation and Evolution of Globular Clusters\", and was supported in part by the NSF Grant PHY05-51164."
    },
    "0911/0911.1356_arXiv.txt": {
        "abstract": "We investigate the Spitzer/IRAC properties of 36 $z\\sim7$ $z_{850}-$dropout galaxies and 3 $z\\sim8$ $Y_{098}$ galaxies derived from deep/wide-area WFC3/IR data of the Early Release Science, the ultradeep HUDF09, and wide-area NICMOS data. We fit stellar population synthesis models to the SEDs to derive mean redshifts, stellar masses, and ages. The $z\\sim7$ galaxies are best characterized by substantial ages ($>100$~Myr) and $M/L_V\\approx0.2$. The main trend with decreasing luminosity is that of bluing of the far$-UV$ slope from $\\beta\\sim-2.0$ to $\\beta\\sim-3.0$. This can be explained by decreasing metallicity, except for the lowest luminosity galaxies ($0.1~L^*_{z=3}$), where low metallicity and smooth SFHs fail to match the blue far$-UV$ and moderately red $H-[3.6]$ color. Such colors may require episodic SFHs with short periods of activity and quiescence (``on-off'' cycles) and/or a contribution from emission lines. The stellar mass of our sample of $z\\sim7$ star forming galaxies correlates with SFR according to $\\log M^{*}=8.70(\\pm0.09)+1.06(\\pm0.10)\\log SFR$, implying star formation may have commenced at $z>10$. No galaxies are found with SFRs much higher or lower than the past averaged SFR suggesting that the typical star formation timescales are probably a substantial fraction of the Hubble time. We report the first IRAC detection of $Y_{098}-$dropout galaxies at $z\\sim8$. The average rest-frame $U-V\\approx0.3$ (AB) of the 3 galaxies are similar to faint $z\\sim7$ galaxies, implying similar $M/L$. The stellar mass density to $M_{UV,AB}<-18$ is $\\rho^*(z=8)=1.8^{+0.7}_{-1.0}\\times10^6~M_\\sun~$Mpc$^{-3}$, following $\\log\\rho^*(z)=10.6(\\pm0.6) - 4.4(\\pm0.7) \\log(1+z)$ $[M_\\sun~$Mpc$^{-3}]$ over $3<z< 8$. ",
        "introduction": "Until recently, only a modest number of relatively bright $z\\gtrsim7$ galaxies were known, mostly from wide-area NICMOS \\citep[][R.~J.~Bouwens et al. in preparation]{Bo08,Oe09a} and ground-based searches\\citep{Ou09,Ca09,Hi09}. The arrival of WFC3/IR aboard HST has dramatically improved the situation by identifying large numbers of $z\\gtrsim7$ galaxies by their redshifted $UV$ light. Here we report on $z\\gtrsim7$ galaxies selected from the WFC3/IR Early Release Science (ERS) observations over the GOODS-South field (\\citealt{Wil09}, R.~J.~Bouwens et al. in preparation), complemented with candidates from the recent ultradeep survey with WFC3/IR over the HUDF09 field (\\citealt{Oe09b,Bo09a}; see also \\citealt{Mc09,Bu09,Ya09}).  Little is known about the stellar masses, metal production, and the contribution of star formation to reionization in these galaxies. Mid-infrared observations with the InfraRed Array Camera (IRAC; \\citealt{Fa04}) on {\\it Spitzer} % can been used to constrain the stellar masses and ages, which has lead to the surprising discovery of quite massive $\\sim10^{10}M_\\sun$ galaxies at $z\\gtrsim6$ \\citep{Ey05,Ya06,St09} and appreciable ages ($200-300$Myr) and $M/L$s as early as $z\\sim7$ \\citep{Eg05,La06,Go09}. The overall results suggest the galaxies formed substantial amounts of stars at even earlier times, well into the epoch of reionization \\citep{St07, Ya06, La09}.   In this Letter, we study the stellar populations of the largest sample of $z\\gtrsim7$ galaxies with IRAC measurements to date, focusing on correlations with luminosity and stellar mass, and implications for the mass density and $z\\sim8$. We adopt an $\\Omega_M=0.3,\\Omega_\\Lambda=0.7$ cosmology with $H_0=70$~km~s$^{-1}$Mpc$^{-1}$. Magnitudes are in the AB photometric system \\citep{Ok83}. ",
        "conclusions": "Using a large sample of 36 $z_{850}-$dropout galaxies based on deep/wide-area WFC3/IR data from the Early Release Science, ultradeep data from the HUDF09, and wide-area NICMOS programs, we investigate the stellar population properties at extreme redshifts $z\\gtrsim7$. The main results are:  \\begin{itemize} \\item The average rest-frame far$-UV$ slope at $z\\sim7$ becomes bluer with decreasing luminosity, from $\\beta\\sim-2.0$ ($L^*_{z=3}$) to $\\beta\\sim-3.0$ ($0.1~L^*_{z=3}$), as reported by \\cite{Bo09b}. The rest-frame $U-V$ becomes bluer as well, but is still moderately red $U-V\\approx0.3$ at $0.1~L^*_{z=3}$, apparently excluding extremely young ages $<100$Myr. If the ages inferred from the simple model fits are correct, galaxies started forming stars very early-on, perhaps as high as $z>10$. The blue far$-UV$ slope and red  $U-V$ colors remain a challenge to fit, however, even for sub-solar metallicity models. Episodic SFHs with periods of activity and quiescence and/or a ($\\sim0.2$ mag) contribution of emission lines to the $[3.6]-$band may be required to resolve the mismatch. \\item The derived stellar masses correlate with the SFRs at $z\\sim7$ according to $\\log M^*=8.70(\\pm0.09)+1.06(\\pm0.10)\\log~SFR$, with relatively low scatter $\\sim0.25~$dex. Emission line contributions of  $\\approx0.2$ mag to both $[3.6]$ and $[4.5]$ would shift the relation by $\\approx-0.2$ dex in mass. The absence of galaxies with SFRs much lower or higher than the past averaged SFR (i.e., strongly bursting or suppressed) suggests that that the typical star formation timescales are probably a substantial fraction of the Hubble time. Note that instantaneously quenched galaxies may fade too quickly to be selected as dropout galaxies. \\item The first Spitzer/IRAC detection of $z=8$ galaxies and their red average $H_{160}-[3.6]\\approx0.55$ suggest that luminous early galaxies may have substantial $M/L_V\\approx0.15$, similar to $z=7$ galaxies. The derived stellar mass density then increases gradually with time following $\\log\\rho^*(z)=10.6(\\pm0.6) - 4.4(\\pm0.7) \\log(1+z)$ $[M_\\sun~$Mpc$^{-3}]$ over $3<z< 8$. \\end{itemize}       Deeper IRAC data on $z>7$ galaxies are needed to bolster these results. More nebular emission line measurements of $z>2$ galaxies would help to understand the possible contribution to the broadband fluxes. Additional modeling of the distribution of SFR versus $M^*$ is needed to decipher the SFHs of $z\\sim7$ galaxies. Larger samples would enable a more secure assessment of the mass density evolution beyond $z=8$."
    },
    "0911/0911.0570_arXiv.txt": {
        "abstract": "In this review, we present a brief description of observational efforts to understand the Galactic thick disk and its relation to the other Galactic components. This review primarily focused on elemental abundance patterns of the thick disk population to pin down the process or processes that were responsible for its existence and evolution. Kinematic and chemical properties of disk stars establish that the thick disk is a distinct component in the Milky Way. The chemical enrichment and star formation histories hold clues to the bigger picture of understanding the Galaxy formation. ",
        "introduction": "Deciphering the history of the birth and the growth of the Milky Way Galaxy is one of the major outstanding astrophysical problems. Detailed studies of our Galaxy would help to gain insights into the formation and evolution of other galaxies, and thereby large scale structure formation in the universe. As inhabitants, we have much better access to our Galaxy to resolve  its building blocks: the stars. The properties of stars that constitute the Galaxy are clues to the processes involved in the making of the Milky Way Galaxy we see today. The observational data on stellar motions and photospheric abundances played decisive roles in the progression of our understanding of the Galaxy. The major, perhaps the first, quantitative study of stellar kinematic and chemical properties of a large sample of nearby dwarfs by Eggen, Lynden-Bell and Sandage (1962), known as ELS, laid the foundation for the Galaxy formation models.  In this seminal paper ELS showed that the eccentricity of stellar orbits and the photometrically derived UV excess ($\\delta$ (U-B)), an indicator of metallicity, are correlated: stars with the largest excess (metal-poor)  are found to be moving in highly elliptical orbits, where as the stars with little or no excess (solar metallicity stars) are found to have circular orbits. ELS attributed this to the rapid collapse of a proto-galactic cloud forming the halo quickly, and the rotationally supported disk component thereafter. Further studies, based on improved observations of globular clusters in the Galaxy, led to the fact that mergers play significant role in the evolution of the Galaxy \\cite[(Searle $\\&$ Zinn 1978)]{SearleZinn1978}. Thanks to large scale surveys of both the kinematic and photometric properties of stars, we now know that the Galaxy consists of four major components: the halo, the thick disk, the thin disk and the bulge. There are excellent reviews which trace the progressive understanding of the subject \\cite[(e.g; Gilmore 1989; Majewski 1993; Freeman $\\&$ Bland-Hawthorn2002)]{majewski93}. In this review, we primarily concentrate on the thick disk component and its relation to the Milky Way galaxy. ",
        "conclusions": "Significant progress has been made, over the last two decades, towards establishing various observational properties. However, we still lack information on the metal-poor side of the thick disk. This has implications to favour a particular merging model. The issue of understanding the MWTD is related to accurate astrometry for a large number of stars at relatively farther distances. Hopefully, this would be resolved with the proposed GAIA space mission data expected to be available in 2015/16. Systematic studies of thick disk in large number of spiral galaxies with varying mass would help to constrain the mass of the merging satellites in the simulations of a particular galaxy with a particular thick disk to thin disk mass ratio. As far as decoding the disk components in other galaxies, we require more powerful tools like next generation large aperture telescopes which are at the horizon. As shown in Figure~2, thick disk structure is much more complex and observations point to multiple process for this structure: mergers, accretion of stars from the debris, heating of the thin disk through scattering, and radial mixing.  Hopefully, in the near future, an unified model will be developed that would match all the well defined observed parameters."
    },
    "0911/0911.0093_arXiv.txt": {
        "abstract": "Using \\xmm\\ and \\chandra, we achieved phase-connected timing of the 105~ms X-ray pulsar \\psr\\ that provides the first measurement of the spin-down rate of a member of the class of Central Compact Objects (CCOs) in supernova remnants.  We measure $\\dot P = (8.68 \\pm 0.09) \\times 10^{-18}$, and find no evidence for timing noise or variations in X-ray flux over 4.8~yr.  In the dipole spin-down formalism, this implies a surface magnetic field strength $B_s = 3.1 \\times 10^{10}$~G, the smallest ever measured for a young neutron star, and consistent with being a fossil field. In combination with upper limits on $B_s$ from other CCO pulsars, this is strong evidence in favor of the ``anti-magnetar'' explanation for their low luminosity and lack of magnetospheric activity or synchrotron nebulae. While this dipole field is small, it can prevent accretion of sufficient fall-back material so that the observed X-ray luminosity of $L_x = 5.3\\times 10^{33}(d/7.1\\ {\\rm kpc})^2$~erg~s$^{-1}$ must instead be residual cooling. The spin-down luminosity of \\psr, $\\dot E =  3.0 \\times 10^{32}$~erg~s$^{-1}$, is an order-of-magnitude smaller than $L_x$. Fitting of the X-ray spectrum to two blackbodies finds small emitting radii, $R_1 = 1.9$~km and $R_2 = 0.45$~km, for components of $kT_1 = 0.30$~keV and $kT_2 = 0.52$~keV, respectively.  Such small, hot regions are ubiquitous among CCOs, and are not yet understood in the context of the anti-magnetar picture because anisotropic surface temperature is usually attributed to the effects of strong magnetic fields. ",
        "introduction": "After the first four timing observations of \\psr\\ in 2004 and 2006 (Paper~2), it was evident that the frequency derivative $\\dot f$ was too small to measure without obtaining a phase-coherent series of observations spanning several years. The most economical way to measure $\\dot f$ is to begin with a logarithmically spaced sequence of observations that maintains the absolute cycle count $\\phi(t)/2\\pi$ and measures the frequency $f$ with increasing accuracy, until the small quadratic term in the phase ephemeris \\begin{displaymath} {\\phi(t) \\over 2\\pi} = f(t-t_0) + {1 \\over 2}\\dot f (t-t_0)^2 + . . . \\end{displaymath} begins to make a significant contribution to the phase. In this case, it was also deemed possible (Paper~2) that accretion from fall-back material could contribute timing irregularities known as torque noise, of a magnitude comparable to the existing upper limits on dipole spin-down.  This made it all the more important to obtain a well-sampled ephemeris that could test for such effects.  \\begin{deluxetable*}{llrccccc}[t] \\tabletypesize{\\scriptsize} \\tablewidth{0pt} \\tablecaption{Log of X-ray Timing Observations of \\psr } \\tablehead{ \\colhead{Mission} & \\colhead{Instr/Mode} & \\colhead{ObsID/Seq\\#} & \\colhead{Date} & \\colhead{Exposure} & \\colhead{Start Epoch} & \\colhead{Period\\tablenotemark{a}} & \\colhead{$Z^2_3$} \\\\ \\colhead{} & \\colhead{} & \\colhead{} & \\colhead{(UT)} & \\colhead{(ks)} & \\colhead{(MJD)} & \\colhead{(ms)} & \\colhead{} } \\startdata {\\it XMM} & EPIC-pn/SW  & 0204970201   & 2004 Oct 18 & 30.6 & 53296.001 & 104.912638(39) & 121.7 \\\\ {\\it XMM} & EPIC-pn/SW  & 0204970301   & 2004 Oct 23 & 30.5 & 53301.984 & 104.912612(55) & \\phantom{1}77.2 \\\\ {\\it XMM} & EPIC-pn/SW  & 0400390201   & 2006 Oct 08 & 29.7 & 54016.245 & 104.912610(47) & \\phantom{1}92.4 \\\\ \\chandra\\ & ACIS-S3/CC  & 6676/500630  & 2006 Nov 23 & 32.2 & 54062.256 & 104.912592(40) & \\phantom{1}94.0 \\\\ {\\it XMM} & EPIC-pn/SW  & 0400390301   & 2007 Mar 20 & 30.5 & 54179.878 & 104.912600(43) & 107.5 \\\\ \\chandra\\ & ACIS-S3/CC  & 9101/500964  & 2007 Nov 12 & 33.1 & 54426.674 & 104.912615(40) & \\phantom{1}95.8 \\\\ \\chandra\\ & ACIS-S3/CC  & 9102/500965  & 2008 Jun 16 & 31.2 & 54628.080 & 104.912593(40) & 115.5 \\\\ {\\it XMM} & EPIC-pn/SW  & 0550670201   & 2008 Sep 19 & 21.2 & 54728.758 & 104.912563(84) & \\phantom{1}61.7 \\\\ {\\it XMM} & EPIC-pn/SW  & 0550670301   & 2008 Sep 21 & 31.0 & 54730.066 & 104.912594(46) & \\phantom{1}84.2 \\\\ {\\it XMM} & EPIC-pn/SW  & 0550670401   & 2008 Sep 23 & 34.8 & 54732.079 & 104.912573(42) & \\phantom{1}80.3 \\\\ {\\it XMM} & EPIC-pn/SW  & 0550670501   & 2008 Sep 29 & 33.0 & 54738.016 & 104.912596(40) & \\phantom{1}97.4 \\\\ {\\it XMM} & EPIC-pn/SW  & 0550670601   & 2008 Oct 10 & 36.0 & 54750.006 & 104.912641(40) & \\phantom{1}90.0 \\\\ \\chandra\\ & ACIS-S3/CC  & 9823/501015  & 2008 Nov 21 & 30.1 & 54791.803 & 104.912645(55) & \\phantom{1}69.8 \\\\ \\chandra\\ & ACIS-S3/CC  & 9824/501016  & 2009 Feb 20 & 29.6 & 54882.483 & 104.912591(53) & \\phantom{1}80.4 \\\\ {\\it XMM} & EPIC-pn/SW  & 0550671001   & 2009 Mar 16 & 27.0 & 54906.255 & 104.912610(48) & \\phantom{1}97.9 \\\\ {\\it XMM} & EPIC-pn/SW  & 0550670901   & 2009 Mar 17 & 26.0 & 54907.609 & 104.912592(60) & \\phantom{1}82.9 \\\\ {\\it XMM} & EPIC-pn/SW  & 0550671201   & 2009 Mar 23 & 27.3 & 54913.582 & 104.912616(46) & 104.8 \\\\ {\\it XMM} & EPIC-pn/SW  & 0550671101   & 2009 Mar 25 & 19.9 & 54915.654 & 104.912544(85) & \\phantom{1}75.8 \\\\ {\\it XMM} & EPIC-pn/SW  & 0550671301   & 2009 Apr 04 & 26.0 & 54925.539 & 104.912590(52) & \\phantom{1}97.9 \\\\ {\\it XMM} & EPIC-pn/SW  & 0550671901   & 2009 Apr 10 & 30.5 & 54931.535 & 104.912620(48) & \\phantom{1}82.9 \\\\ {\\it XMM} & EPIC-pn/SW  & 0550671801   & 2009 Apr 22 & 28.0 & 54943.897 & 104.912621(69) & \\phantom{1}67.4 \\\\ \\chandra\\ & ACIS-S3/CC  & 10128/501055 & 2009 Jun 02 & 33.2 & 54984.888 & 104.912596(37) & 117.9 \\\\ \\chandra\\ & ACIS-S3/CC  & 10129/501056 & 2009 Jul 29 & 32.2 & 55041.231 & 104.912633(40) & 105.6 \\enddata \\tablenotetext{a}{Barycentric period derived from a $Z^2_3$ test. The \\cite{lea83} uncertainty on the last digits is in parentheses.} \\label{logtable} \\end{deluxetable*}  Unfortunately, maintaining cycle count would require more observations classified as time-constrained than the \\chandra\\ {\\it X-ray Observatory} allocates to any one project, while for \\xmm, semiannual visibility windows for this source are only 40 days long, not wide enough to securely bridge over the intervening 5 month gaps. It took until 2008 to implement a strategy that uses the two satellites in a coordinated fashion, filling two \\xmm\\ visibility windows with six observations each of variable spacing, while requesting pairs of \\chandra\\ observations to bridge the gaps between and after the two \\xmm\\ windows.  In this manner, a phase-coherent timing solution spanning 1.7~yr could be achieved, which could also be extrapolated backward to incorporate the earlier observations. All but one of the approved observations in 2008--2009 were performed as planned.  Due to a loss of contact with the \\xmm\\ spacecraft, the last observation of 2008 was rescheduled to 2009, but this did not compromise the success of the program.  We were thus able to obtain a fully coherent timing solution incorporating all of the observations listed in Table~\\ref{logtable}, including the earliest ones, spanning 4.8~yr in total.  All of the \\xmm\\ observations used the pn detector of the European Photon Imaging Camera (EPIC-pn) in ``small window'' (SW) mode to achieve 5.7~ms time resolution, and an absolute uncertainty of $\\approx 3$~ms on the arrival time of any photon. We processed all EPIC data using the emchain and epchain scripts under Science Analysis System (SAS) version xmmsas\\_20060628\\_1801-7.0.0.  The leap second at the end of 2008 was inserted manually. Simultaneous data were acquired with the EPIC~MOS cameras, operated in ``full frame'' mode.  Although not useful for timing purposes because of the 2.7~s readout, the location of the source at the center of the on-axis MOS CCDs allows a better background measurement to test for flux variability, an important indicator of accretion, than the EPIC-pn SW mode.  The \\chandra\\ observations used the Advanced Camera for Imaging and Spectroscopy (ACIS) in continuous-clocking (CC) mode to provide time resolution of 2.85~ms. This study uses data processed by the pipeline software revisions v7.6.9--v8.0.  Reduction and analysis used the standard software package CIAO (v3.4) and CALDB (v3.4.2).  The photon arrival times in CC mode are adjusted in the standard processing to account for the known position of the pulsar, spacecraft dither, and SIM offset. Absolute accuracy of the time assignment in \\chandra\\ CC-mode is limited by the uncertainty in the position of the pulsar, which was determined in Paper~2 from an ACIS image. The typical accuracy of $\\approx 1$ pixel then corresponds to an uncertainty of $\\approx 3$~ms, which is similar to the \\xmm\\ accuracy, and is 0.03 rotations in the case of \\psr.  We will show that the measured dispersion in pulse arrival times is comparable to this uncertainty.  \\subsection{Results of Timing Analysis}  For each observation in Table~\\ref{logtable} we transformed the photon arrival times to Barycentric Dynamical Time (TDB) using the pulsar coordinates determined in Paper~2 and reproduced in Table~\\ref{ephemeris}. Diffuse SNR emission is a significant source of background. To maximize the pulsar signal-to-noise ratio, we used a source extraction radius of $12^{\\prime\\prime}$ for the \\xmm\\ observations, and five columns ($2.\\!^{\\prime\\prime}4$) for the \\chandra\\ CC-mode data. An energy cut of $1-5$~keV was found to maximize pulsed power. The pulse profile and value of the period in each observation was derived from a $Z^2_3$ periodogram \\citep{buc83}, a choice of harmonics that was found to minimize the uncertainties. The resulting profiles were cross-correlated, shifted, and summed to create a master pulse profile template.  Individual profiles were then cross correlated with the template to determine the time of arrival (TOA) at each epoch.  \\begin{figure}[b] \\centerline{ \\hfill \\includegraphics[scale=0.36,angle=270]{f1.eps} \\hfill } \\caption{ Pulse phase residuals for the 23 timing observations of \\psr\\ from the best fitting ephemeris of Table~\\ref{ephemeris}. Error bars are 1 $\\sigma$. } \\label{phaseplot} \\end{figure}  Starting with the dense series of \\xmm\\ observations in 2008 September, the TOAs were iteratively fitted using the TEMPO\\footnote{http://www.atnf.csiro.au/research/pulsar/tempo} software.  We fitted the three observations from September 19--23 to a linear ephemeris, and added TOAs one at a time using a quadratic ephemeris, finding that the new TOA would match to $<0.1$ cycles the predicted phase derived from the previous set.  After adding the final observation to the ephemeris, we then worked backward in time from 2008 September to 2004 October until all 23 observations had been included. The quadratic term contributes $-8.7$ cycles of rotation over the 4.8~yr span of the ephemeris, which yields the small uncertainty in $\\dot P$ discussed below. The phases and errors were determined by cross-correlating with a final iterated template of co-added profiles produced from the best fitting ephemeris given in Table~\\ref{ephemeris}. Figure~\\ref{phaseplot} shows the residuals of the individual observations from the best fit, which demonstrates the validity of the solution.  The weighted rms of the phase residuals is 3.4~ms, or 0.032 pulse cycles, which is comparable to the individual measurement errors. There is no evidence of timing noise or higher derivatives in the residuals.  The variation of $Z^2_3$ (pulsed power) among the observations listed in Table~\\ref{logtable} is also consistent with statistical expectations. Figure~\\ref{lightcurve} shows the summed pulse profile from the 16 \\xmm\\ observations, for which it is possible to measure a reasonably accurate background (unlike the \\chandra\\ CC-mode data).  The sensitivity of the results to $\\dot P $ and $B_s$ can be estimated analytically. If a pulsar is spinning down smoothly, then a coherent timing solution spanning time $T$ will have an uncertainty in frequency of $\\delta f = \\delta P/P^2 \\approx 0.1/T$, and will be sensitive to a period derivative of $\\dot P_{\\rm \\min} = \\delta \\dot P \\approx 0.2(P/T)^2$. This in turn will measure $B_s = 3.2 \\times 10^{19}\\sqrt{P\\dot P}$ with fractional uncertainty \\begin{displaymath} {\\delta B_s \\over B_s} \\ \\approx \\ 1.0 \\times 10^{38}\\ {P^3 \\over B_s^2\\ T^2} \\end{displaymath} \\begin{displaymath} \\ = \\ 0.05 \\left(B_s \\over 10^{10}\\,{\\rm G}\\right)^{-2} \\left(T \\over 4.8\\ {\\rm yr}\\right)^{-2} \\left(P \\over 0.105\\ {\\rm s}\\right)^3. \\end{displaymath} A coherent ephemeris for \\psr\\ spanning 4.8~yr has an uncertainty on $\\dot P$ of $\\delta \\dot P \\approx 1 \\times 10^{-19}$, corresponding to a $5\\sigma$ detection limit of $B_s = 5 \\times 10^9$~G. In comparison, the measured values $P=0.104912611147(4)$~s and $\\dot P = 8.68(9) \\times 10^{-18}$ have fitted uncertainties consistent with these analytic estimates, and imply in the dipole spindown formalism a surface magnetic field strength $B_s = 3.1 \\times 10^{10}$~G, a spin-down luminosity $\\dot E = -I\\Omega\\dot\\Omega = 4\\pi^2I\\dot P/P^3 = 3.0 \\times 10^{32}$~erg~s$^{-1}$, and characteristic age $\\tau_c \\equiv P/2\\dot P = 192$~Myr, all with negligible statistical error.  \\begin{deluxetable}{ll} \\tablewidth{0pt} \\tablecaption{Spin Parameters of \\psr } \\tablehead{ \\colhead{Parameter} & \\colhead{Value} } \\startdata Right ascension, R.A. (J2000)\\tablenotemark{a} & $18^{\\rm h}52^{\\rm m}38^{\\rm s}\\!.57$ \\\\ Declination, Decl. (J2000)\\tablenotemark{a}    & $+00^{\\circ}40^{\\prime}19^{\\prime\\prime}\\!.8$ \\\\ Epoch (MJD TDB)\\tablenotemark{b}               & 54597.00000046 \\\\ Spin period, $P$                               & 0.104912611147(4)~s \\\\ Period derivative, $\\dot P$                    & $(8.68 \\pm 0.09) \\times 10^{-18}$ \\\\ Valid range of dates (MJD)                     & 53296 -- 55041 \\\\ Surface dipole magnetic field, $B_s$           & $3.1 \\times 10^{10}$~G \\\\ Spin-down luminosity, $\\dot E$                 & $3.0 \\times 10^{32}$~erg~s$^{-1}$ \\\\ Characteristic age, $\\tau_c$                   & 192~Myr \\enddata \\tablenotetext{a}{Measured from \\chandra\\ ACIS-I ObsID 1982 (Paper~1). Typical \\chandra\\ ACIS coordinate uncertainty is $0\\farcs6$.} \\tablenotetext{b}{Epoch of ephemeris corresponds to phase zero in Figure~\\ref{lightcurve}.} \\label{ephemeris} \\end{deluxetable}  \\begin{figure} \\centerline{ \\hfill \\includegraphics[scale=0.35,angle=270]{f2.eps} \\hfill } \\caption{ Summed pulse profile from the 16 \\xmm\\ observations of \\psr\\ listed in Table~\\ref{logtable}, folded using the ephemeris of Table~\\ref{ephemeris}. Phase zero corresponds to the TDB epoch in Table~\\ref{ephemeris}.  The pulsed fraction after correcting for background is $f_p = 64\\pm2\\%$. } \\vspace{-0.1in} \\label{lightcurve} \\end{figure}  \\subsection{Flux and Spectral Analysis}  \\begin{deluxetable*}{lcccc}[b] \\tablewidth{0pt} \\tablecaption{\\xmm\\ Spectral Fits to \\psr } \\tablehead{ \\colhead{Parameter}  & \\colhead{BB} & \\colhead{BB+BB} & \\colhead{NSA} & \\colhead{Comp BB} } \\startdata $N_{\\rm H}$~($10^{22}$ cm$^{-2}$)         & $1.32 \\pm 0.05$        & $1.82^{+0.23}_{-0.18}$  & $1.67^{+0.06}_{-0.07}$ & $1.53^{+0.09}_{-0.08}$ \\\\ $kT_1$ (keV)                              & $0.460 \\pm 0.007$      & $0.30\\pm0.05$ & $0.287 \\pm 0.007$  & $0.396 \\pm 0.014$    \\\\ $R_1$ (km)                                & $0.72\\,d_{7.1}$        & $1.9\\,d_{7.1}$  & 13.06\\tablenotemark{a}  &  $0.93\\,d_{7.1}$  \\\\ $kT_2$ (keV)                              & . . .                  & $0.52\\pm0.03$   & . . . & . . .        \\\\ $R_2$ (km)                                & . . .                  & $0.45\\,d_{7.1}$  & . . . & . . .       \\\\ $\\Gamma$    & . . . & . . .  & . . .   & $4.80^{+0.34}_{-0.25}$ \\\\ $d$ (kpc)                                 & . . . & . . . &  $23.7^{+4.0}_{-2.6}$ & . . .  \\\\ $F_x(0.5-10\\ {\\rm keV}$)\\tablenotemark{b} & $2.0 \\times 10^{-13}$  & $2.0 \\times 10^{-13}$ & $2.0 \\times 10^{-13}$  & $2.0 \\times 10^{-13}$  \\\\ $L_{\\rm bol}$ (erg s$^{-1})$\\tablenotemark{c} & $3.0 \\times 10^{33}\\,d_{7.1}^2$ & $5.3 \\times 10^{33}\\,d_{7.1}^2$ & . . .  & $3.6 \\times 10^{33}\\,d_{7.1}^2$ \\\\ $\\chi^2_{\\nu}(\\nu)$                       & 1.38(204)  & 1.07(201) & 1.10(204) & 1.12(203) \\enddata \\tablecomments{Spectrum fitted over $0.7-7$~keV. Errors are 90\\% confidence for one interesting parameter ($\\chi^2 = \\chi^2_{\\rm min} + 2.7$).} \\tablenotetext{a}{Fixed parameter $R^{\\infty}=13.06$~km corresponding to $R=10$~km with $M = 1.4\\,M_{\\odot}$.} \\tablenotetext{b}{Absorbed flux in units of erg cm$^{-2}$ s$^{-1}$. Average of EPIC pn and MOS.} \\tablenotetext{c}{Unabsorbed, bolometric luminosity assuming $d = 7.1$~kpc. Average of EPIC pn and MOS.} \\label{spectable} \\end{deluxetable*}  In Papers 1 and 2 we presented several observations of \\psr\\ that were consistent with steady flux and a blackbody spectrum of $kT \\approx 0.46$~keV.  The luminosity of $3.4 \\times 10^{33}\\,d^2_{7.1}$ erg~s$^{-1}$ corresponded to a blackbody radius of only $\\approx 0.8\\,d_{7.1}$~km.  This is typical result for a CCO, and indicates a concentrated hot spot that remains to be understood.  Using the methods described in Paper 1, we now combine all 16 \\xmm\\ observations of \\psr\\ into one spectrum in order to search for deviations from a blackbody that might indicate temperature variations on the surface, or cyclotron lines.  EPIC pn and MOS spectra were fitted jointly. The accumulated exposure allows us to fit the spectrum over $0.7-7$~keV, an increase in coverage over the $1-5$~keV range used in the previous papers. While a single blackbody can still be fitted with a temperature of 0.46~keV, the high signal-to-noise ratio and expanded energy range of the summed spectrum reveal systematic deviations from the fit; the reduced $\\chi^2_{\\nu} = 1.38$ for 240 degrees of freedom is unacceptable (see Figure~\\ref{spectrum} and Table~\\ref{spectable}). Similar to the results from other CCOs, we find that a fit to two blackbodies is significantly improved ($\\chi^2_{\\nu} = 1.07$), with temperatures of $kT_1 = 0.30$~keV and $kT_2 = 0.52$~keV, while the corresponding radii $R_1 = 1.9\\,d_{7.1}$~km and $R_2 = 0.45\\,d_{7.1}$~km still cover only a small fraction of the surface. Here we define radius using the phase-averaged luminosity $L_{1,2} = 4\\pi R_{1,2}^2\\sigma T_{1,2}^4$ regardless of the unknown emission geometry, which could be, for example, a concentric annulus. Adoption of the two-blackbody fit results in a higher bolometric luminosity than the single blackbody, $5.3 \\times 10^{33}\\,d_{7.1}^2$ erg~s$^{-1}$ vs. $3.0 \\times 10^{33}\\,d_{7.1}^2$ erg~s$^{-1}$. We don't combine the \\chandra\\ spectra here for an independent fit because of the increased background and other uncertainties involved in analyzing CC-mode data for this faint source.  \\begin{figure}[b] \\centerline{ \\includegraphics[scale=0.37,angle=270]{f4a.eps} } \\centerline{ \\includegraphics[scale=0.37,angle=270]{f4b.eps} } \\caption{ Spectrum of the 16 summed \\xmm\\ observations of \\psr. The EPIC pn spectrum is black, while the average of the MOS1 and MOS2 spectra is red. Top: Fit to a single blackbody model. Bottom: Fit to a double blackbody model. The parameters of the fits are given in Table~\\ref{spectable}. } \\label{spectrum} \\end{figure}  \\begin{figure}[b] \\centerline{ \\includegraphics[scale=0.37,angle=270]{f5a.eps} } \\centerline{ \\includegraphics[scale=0.37,angle=270]{f5b.eps} } \\caption{ Spectrum of the 16 summed \\xmm\\ observations of \\psr. The EPIC pn spectrum is black, while the average of the MOS1 and MOS2 spectra is red. Top: Fit to a nonmagnetic neutron star hydrogen atmosphere model. Bottom: Fit to a Comptonized blackbody model. The parameters of the fits are given in Table~\\ref{spectable}. } \\label{spectrum2} \\end{figure}  Additional spectral models that were explored include the nonmagnetic neutron star hydrogen atmosphere of \\citet{zav96}, and a simplified Comptonized blackbody, as described by \\citet{hal08} for application to anomalous X-ray pulsars. Each of these models fits nearly as well as the two-blackbody model because they can eliminate the residual excess at high energy that is left by a single blackbody fit (see Figure~\\ref{spectrum2} and Table~\\ref{spectable}). For the neutron star atmosphere (NSA), we fixed the parameters $M=1.4\\,M_{\\odot}$ and $R^{\\infty}=13.06$~km (corresponding to $R=10$~km), treating the effective temperature and distance as free parameters. As expected, the fitted temperature is smaller than those of the blackbody models, but the fitted distance of 23.7~kpc is a factor of 3.3 higher than the known value.  This indicates that in the atmosphere model only $\\sim 10\\%$ of the surface is emitting.  We are not motivated to search for an atmosphere model that fits to the full surface area, as was done for the Cas~A CCO by \\citet{ho09}, because the larged pulsed fraction of \\psr\\ clearly requires a small hot spot. The same interpretation attaches to the Comptonized blackbody model, in which the inferred blackbody radius is only $0.93\\,d_{7.1}$~km.  \\begin{figure} \\centerline{ \\hfill \\includegraphics[scale=0.47,angle=0]{f3.eps} \\hfill } \\caption{ Energy-resolved pulse profiles from the summed 16 \\xmm\\ observations of \\psr\\ of Figure~\\ref{lightcurve} decomposed into four energy bands containing equal numbers of photons. Background has been subtracted. The pulse shape does not differ significantly as a function of energy. } \\label{fourfold} \\end{figure}  The large intrinsic pulsed fraction, $f_p = 64\\pm2\\%$ (defined as the fraction of counts above the minimum in the light-curve, corrected for background; see Figure~\\ref{lightcurve}) confirms that the emitting area must be small and far from the rotation axis, while the line of sight also makes a large angle to the rotation axis. The pulse shape does not appear to be a function of energy (Figure~\\ref{fourfold}), which also suggests that the emitting area is small. The fitted column density from any model in Table~\\ref{spectable} agrees reasonably well with the $N_{\\rm H}$ derived from fits to the SNR spectrum. \\citet{sun04} find $N_{\\rm H} = (1.54 -  1.78) \\times 10^{22}$~cm$^{-2}$ from fitting a variety of equilibrium and non-equilibrium ionization models to {\\it ASCA\\/} and \\chandra\\ spectra of the SNR, while \\citet{gia09} find $N_{\\rm H} = (1.52 \\pm 0.02) \\times 10^{22}$~cm$^{-2}$ from fitting a non-equilibrium ionization model to the \\xmm\\ observations of 2004.  We then searched for any indication of variability by extracting and fitting the spectrum of each individual observation, fixing the spectral parameters (except for the total flux) to those of the summed spectrum.  In order to minimize statistical and systematic errors in this comparison, we restricted the fitted energy range to $1-5$~keV for the individual spectra. The resulting individual fluxes are shown in Figure~\\ref{fluxes}, where they are seen to be constant within errors. The mean $1-5$~keV flux is $1.9 \\times 10^{-13}$ erg~cm$^{-2}$ s$^{-1}$. The typical $1 \\sigma$ uncertainty in each observation is $\\approx 10\\%$, and the rms dispersion among the 23 observations is $10\\%$.  Unlike the unique case of \\one, there is no evidence for cyclotron absorption lines in the spectrum of \\psr\\ or in any other CCO.  \\one\\ has strong absorption features at 0.7 and 1.4~keV \\citep{san02,mer02a} that can be attributed to the cyclotron fundamental energy $E_c$ and its first harmonic in a field of $\\approx 8 \\times 10^{10}$~G \\citep{big03}, according to the relation $E_c = 1.16(B/10^{11}\\,{\\rm G})/(1+z)$~keV, where $z \\approx 0.3$ is the gravitational redshift. The spectrum of \\psr\\ can now be understood in terms of its weaker surface $B$-field. For $B_s = 3.1 \\times 10^{10}$~G, the cyclotron fundamental and first harmonic are at 0.27~keV and 0.54~keV, below the region accessible to X-ray spectroscopy (Figure~\\ref{spectrum}), especially because of the larger intervening column density to \\psr.  The same explanation applied to other well-observed CCOs suggests that they also have weaker surface $B$-fields than \\one.  \\begin{figure}[b] \\centerline{ \\hfill \\includegraphics[scale=0.36,angle=270]{f6.eps} \\hfill } \\caption{ Fluxes in the $1-5$ keV band for the 23 individual observations of \\psr, fitted to the two blackbody model given in Table~\\ref{spectable}. Errors bars are 1 $\\sigma$. The weighted mean flux is indicated by the dot-dashed line. } \\label{fluxes} \\end{figure} ",
        "conclusions": "Measurement of the spin-down rate of the 105~ms \\psr\\ in \\snr\\ was achieved using X-ray observations coordinated between \\xmm\\ and \\chandra.  The resulting phase-connected ephemeris spanning 4.8~yr requires only the first derivatives of the spin frequency, with no measurable timing noise superposed. The straightforward interpretation of this result is dipole spin-down due to a weak surface magnetic field of $B_s = 3.1 \\times 10^{10}$~G, the smallest measured for any young neutron star. While this is the first measurement of the spin-down rate of a CCO, upper limits on the period derivatives of \\one\\ and \\puppsr, as well as spectral features in those two, also indicate weak $B$-fields.  The properties of the CCO class were loosely defined in the past, but it is now evident that a weak $B$-field is the physical reason that a large fraction of the neutron stars were so named.  The body of evidence supports the ``anti-magnetar'' model for the origin of such CCOs: neutron stars born spinning ``slowly'' may as a result be endowed with only weak magnetic fields.  \\psr\\ falls in a region of ($B,P$) space that overlaps with moderately recycled pulsars.  Single radio pulsars with similar spin parameters may therefore be former CCOs rather than recycled, and they may be much younger than their characteristic ages. X-ray observations of pulsars in this region may find evidence of their relative youth via surface thermal emission.  Ongoing timing studies of the CCOs \\one\\ and \\puppsr\\ will able to measure their spin-down rates and can refine the correspondence between spectral features interpreted as electron cyclotron lines and surface $B$-field.  Those CCOs such as \\psr, with no apparent absorption lines, may simply have weaker $B$-fields than the others by a factor $2-3$, so that the cyclotron resonance is below the soft X-ray band.  As for CCOs that have not yet been seen to pulse, there is no reason to suppose that they are of a different physical class.  Rather, their surface temperature distributions and/or viewing geometries may deliver only weak modulation that happens to fall below the sensitivity of existing data, while deeper observations may succeed in discovering their spins. It would be of great interest to discover pulsations from the CCO in Cas~A, the youngest known neutron star, which in all respects is the prototype of the CCO class.  At an age of 330~yr, it is an order of magnitude younger than the others, while having similar spectral properties \\citep{pav09}.  In our model, its spin parameters were fixed at birth, and should conform closely to those of the older CCOs.  A recent \\chandra\\ observation of the Cas~A CCO did not detect pulsations (see Appendix), but the High Resolution Camera (HRC) that was employed lacks energy resolution, a capability that may be crucial in the same way that it was for the \\xmm\\ discovery of the pulsar in Puppis~A \\citep{got09}.  A remaining theoretical puzzle about CCOs is the origin of their surface temperature anisotropies, in particular, the one or two warm/hot regions that are smaller than the full neutron star surface.  The high signal-to-noise spectrum accumulated from \\psr\\ reveals similar temperature structure as the other CCOs. The measured spin-down power of \\psr\\ is too small by an order of magnitude to contribute to these excess emissions. We considered whether accretion of fall-back material can be responsible for this effect, but find it unlikely in the case of \\psr\\ because of its slow, steady spin-down.  If other CCOs are found to have similar spin-down properties, the same arguments would apply to them.  Therefore, it is important to measure the spin-down rate of \\puppsr, the CCO in Puppis~A, and investigate further its apparent phase-dependent emission line at 0.8~keV \\citep{got09}, which may be indicative of accretion because it is in emission.  The small regions of high surface temperature are properties that CCOs share with magnetars, although most magnetars are hotter and more luminous. The total X-ray luminosities of CCOs are consistent with slow cooling scenarios, and there would be no need to hypothesize a mechanism of magnetic field amplification beyond simple flux freezing if it were not for their strongly anisotropic surface temperature distributions.  Paradoxically, the explanations that have been considered for hot spots invoke strong magnetic fields just below the surface, which is exactly the basis of the magnetar model.  In this sense, CCOs are potentially magnetars waiting to emerge, or maybe ordinary neutron stars in which the original field was submerged by especially massive fall-back of supernova debris. This is an unsatisfying picture, if only because the contrast between field strength in different regions must be a factor of $10^4$ or larger to affect the surface temperature distribution while maintaining the slow spin-down and allowing soft X-ray cyclotron lines. Resolving this essential mystery of the CCOs will undoubtedly generate new insight into the physics of neutron stars."
    },
    "0911/0911.0436_arXiv.txt": {
        "abstract": "We present  a new algorithm for identifying  the substructure within simulated dark  matter haloes.  The  method is an extension  of that proposed by \\citet{tmy04}  and \\citet{giocoli08}, which identifies a subhalo  as a  group of  self-bound  particles that  prior to  being accreted by the  {\\it main} progenitor of the  host halo belonged to one and the same  progenitor halo (hereafter `satellite').  However, this definition does  not account for the fact  that these satellite haloes themselves may also  have substructure, which thus gives rise to sub-subhaloes,  etc.  Our new  algorithm identifies substructures at all levels of this hierarchy, and we use it to determine the mass function  of all  substructure (counting  sub-haloes, sub-subhaloes, etc.).  On average, haloes which formed more recently tend to have a larger  mass fraction in  substructure and  to be  less concentrated than average haloes of the  same mass.  We provide quantitative fits to these correlations.  Even  though our algorithm is very different from  that of  \\citet{getal04}, we  too find  that the  subhalo mass function  per  unit  mass  at  redshift $z=0$  is  universal.   This universality extends  to any redshift  only if one accounts  for the fact  that host  haloes of  a given  mass are  less  concentrated at higher redshifts,  and concentration and  substructure abundance are anti-correlated.  This universality  allows a simple parametrization of the subhalo  mass function integrated over all  host halo masses, at any given time.  We  provide analytic fits to this function which should be useful  in halo model analyses which  equate galaxies with halo substructure when interpreting clustering in large sky surveys. Finally,  we  discuss systematic  differences  in  the subhalo  mass function that arise from different definitions of (host) halo mass. ",
        "introduction": "In  the  standard  scenario   of  structure  formation,  galaxies  are surrounded by extended dark matter haloes, which form by gravitational instability   seeded  by  some   initial  density   fluctuation  field $\\delta(\\vec{x},z)$.  According to the simple spherical collapse model \\citep[e.g.][]{pee80}, a halo of mass  $M$ collapses at a redshift $z$ when the  linear density  fluctuation field -  smoothed on a  scale $R \\sim    M^{1/3}$   -   first    exceeds   some    critical   threshold \\citep{ps74,bond91,lc93,st02}.   Small   systems  collapse  at  higher redshifts when the universe  was denser, and then merge hierarchically to  form  larger haloes.   Galaxy  clusters are  at  the  top of  this hierarchy: they  represent the  most massive virialized  structures in the present-day universe, and  may host thousands of galaxies.  Recent studies,     using    high     resolution     numerical    simulations \\citep[e.g.][]{m98,tds,swtk01}, have shown that the cores of subhaloes accreted  along  the  host  merging  history  may  survive  until  the present-time to form the so-called substructure population.  Different studies of galaxy  formation and evolution have attempted to correlate such substructures with satellite galaxies, with conflicting results.  $N$-Body  simulations of  galaxy-size haloes seem  to predict more  substructures   than  observed  \\citep{mooreetal99,setal03}.   A number of different solutions to this `substructure problem' have been suggested, among  others early cosmic  reionization \\citep{susa}, mass loss and  gas stripping \\citep{ketal04,mac09},  and a downturn  in the power spectrum at small scales \\citep{kam00}.  From  the observational point  of view,  recent kinematic  analyses of Milky-Way satellites  \\citep{simonmw,walk} have shown  that, if galaxy formation in low-mass dark  matter haloes is strongly suppressed after reionization,  the   circular  velocity  function   of  simulated  CDM subhaloes can be brought  into approximate agreement with the observed circular  velocity  function  of  the  Milky  Way  satellite  galaxies \\citep[e.g.,][]{koposov09}.    \\citet{Ishiy08,Ishiy09}  claim  instead that  there is  no missing  satellites  problem if  one considers  the scatter  in  the  substructure  mass  function due  to  the  different formation history of the haloes.   The idea here is that the Milky-Way and  Andromeda haloes  reside  within an  underdense  region, so  they accreted their  substructures at above average  redshifts (compared to other objects  of the same  mass today), and  so they are  expected to possess fewer satellites \\citep[see also][]{vtg05}.  Knowledge of  the substructure population  is important, not  only for galaxy  formation  studies,  but  also  for  estimating  the  expected observable $\\gamma$-ray  flux from dark  matter particle annihilations within  the Milky  Way.  The  expected flux  depends  strongly on  the substructure population, its spatial distribution within the halo, and on the DM particle cross section \\citep{pbb,diemand07,gioc08}.  The observed thickness of stellar  disks in spiral galaxies is another imprint    of    the    substructure    population   of    the    host haloes. Semi-analytical models \\citep{benson} and numerical simulations \\citep{kaz1,kaz2} show  that merging events  in the central  region of the halo are responsible for disk thickening.  Hence, the substructure mass function, its redshift evolution and the satellite accretion rate all represent  key ingredients for a  more complete model  of the disk structure of spiral  galaxies, not to mention flares,  bars, and faint filamentary structures above the disk plane.  Different algorithms  have been proposed to  identify substructures in simulated dark matter haloes.  All these algorithms yield subhalo mass functions  that  resemble a  power-law  at the  low  mass  end with  a logarithmic    slope   in    the   range    from   $-0.8$    to   $-1$ \\citep{gill1,getal04,delucia04,shaw06,giocoli08,mad08,wetzel09},   both at redshift $z=0$ and at  higher redshifts.  In this work we highlight some differences  among four  such methods and  present a  new subhalo finder which we believe is  useful for studies of galaxy formation and evolution in that it is directly linked to the merger tree of the host halo.  In this paper we focus on the study of the ``dark'' subhalo population and on its redshift evolution, to help clarify the role of dark matter and   its   dynamics  in   the   structure   formation  history.    In Section~\\ref{bkgnd}  we describe  several  substructure identification algorithms already  in the  literature, including the  one we  will be modifying.   Section~\\ref{sim}  describes  the  cosmological  $N$-Body simulation  we  use,  and   the  required  post-processing.   Our  new algorithm  to  identify substructures,  and  an  analytic  fit to  the resulting  subhalo  mass function  is  presented in  Section~\\ref{r2}. Section~\\ref{r3} discusses  how the  scatter in the  substructure mass fraction at fixed mass and  at different redshifts correlates with the formation  history, and other  structural parameters  of the  halo.  A summary and conclusions are given in Section~\\ref{sc}. In Appendix \\ref{previous} we compare the results of this paper with two previous ones: \\citet{giocoli08} and \\cite{getal04} that studied the substructures population on the same cosmological simulation but with different algorithms. In Appendix  \\ref{massris} we propose some fitting functions useful to rescale the virial radius and mass when different definitions of the enclosed virial density are adopted. ",
        "conclusions": "\\label{sc}  We estimated the  mass function of substructures by  following all the branches (rather than just the main branch) of the merger history tree of  a halo.   This  distribution, $\\mathrm{d}N/\\mathrm{d}m_{\\rm  sb}$, depends on  the mass of the  host halo, when the  halo was identified, and its formation time.  However  when the number of substructures per host  halo   mass  is  considered,   the  distribution  turns   to  be ``universal'' at  a given redshift,  but simulations with  higher mass resolution are  required to  determine if this  is true even  when the host halo mass is smaller than $10^{11}h^{-1}M_\\odot$.  This universal shape is characterized  by a power law of slope  $-0.9$ at low masses, times  an exponential  cut-off when  the substructure  mass approaches that of the host halo.   This distribution can be integrated to obtain the substructure virial mass fraction and the mean number of clumps at a given host halo mass.  At fixed  host halo mass, the  haloes which formed  at higher redshift tend to have less substructure.  In addition, more massive host haloes have  more substructures  than  less massive  ones.   The first  trend arises because satellites which spend  more time in the potential well of their host lose a larger  fraction of their initial mass (or can be totally disrupted).  Massive haloes formed more recently, so this same mechanism explains  the second trend.   Similar trends are seen  if we replace  halo  formation  time  (which  is not  observable)  with  the concentration parameter of the halo density profile (which is).  We also  showed that the  Poisson model, with  a mean that  depends on host halo mass, is a reasonable description of the substructure counts within a host  halo.  However, the second and  third factorial moments of  the substructure counts  are $\\sim  5$ percent  larger than  for a Poisson  model with the  same mean  counts.  As  observations improve, this  may become  important in  Halo Model  parametrization  of large scale structure which relate halo substructure to galaxies.  At fixed  mass, the  mean number of  substructures above  some minimum mass is larger  at high redshift.  Although this  is not unexpected -- the  higher  redshift  haloes  have  had less  time  to  destroy  their substructures  -- we  argued  that the  redshift  dependence could  be well-approximated  by  assuming   that  the  anti-correlation  between substructure abundance and host halo concentration we measure at $z=0$ does  not evolve.   At  fixed  mass, higher  redshift  haloes are  less concentrated --  in effect,  our assumption is  that the  same physics which affects the concentrations  also affects the substructure.  This allows us to provide a  simple prescription for the redshift evolution of substructures, which we  use to provide analytic approximations for the result of integrating the substructure mass function over all host halo   masses.   These  results   should  find   use  in   halo  model interpretations of the (evolution  of the) galaxy clustering signal in large scale sky surveys \\cite[e.g.,][]{vdb07,2dfsdss}."
    },
    "0911/0911.4954_arXiv.txt": {
        "abstract": "{Recent data from cosmic ray experiments such as PAMELA, Fermi, ATIC and PPB-BETS all suggest the need for a new primary source of electrons and positrons at high ($\\gsim 100 \\gev$) energies. Many proposals have been put forth to explain these data, usually relying on a single particle to annihilate or decay to produce \\epp. In this paper, we consider models with multiple species of WIMPs with significantly different masses. We show if such dark matter candidates $\\chi_i$ annihilate into light bosons, they naturally produce equal annihilation rates, even as the available numbers of pairs for annihilation $n_{\\chi_i}^2$ differ by orders of magnitude. We argue that a consequence of these models can be to add additional signal naturally at lower ($\\sim 100 \\gev$) versus higher ($\\sim \\tev$) energies, changing the expected spectrum and even adding bumps at lower energies, which may alleviate some of the tension in the required annihilation rates between PAMELA and Fermi. These spectral changes may yield observable consequences in the microwave Haze signal observed at the upcoming Planck satellite. Such a model can connect to other observable signals such as DAMA and INTEGRAL by having the lighter (heavier) state be a pseudo-Dirac fermion with splitting 100 keV (1 MeV). We show that variations in the halo velocity dispersion can alleviate constraints from final state radiation in the galactic center and galactic ridge. If the lighter WIMP has a large self-interaction cross section, the light-WIMP halo might collapse, dramatically altering expectations for direct and indirect detection signatures.} ",
        "introduction": "\\label{sec:intro} Signals from a wide range of sources have prompted a rethinking of the nature of dark matter. The positron fraction of cosmic rays has shown a steep rise \\cite{Adriani:2008zr}, implying the presence of a new primary source of antimatter. The Fermi \\cite{Abdo:2009zk}, ATIC \\cite{:2008zzr} and PPB-BETS \\cite{Torii:2008xu} experiments all show a harder spectrum of \\epp at high ($\\gsim 100 \\gev$) energies than had been predicted from secondary production, with a break at $\\sim$ 1 TeV seen at Fermi and HESS  \\cite{Collaboration:2008aaa} \\cite{Aharonian:2009ah}. Data from INTEGRAL suggest a new source of low energy positrons, while the DAMA modulation signal has extended to 8.2$\\sigma$.  Remarkably, there is a simple ``unified'' framework that explains all of these anomalies \\cite{ArkaniHamed:2008qn}. One of the principle changes is that rather than annihilating to standard model channels,  the WIMP can annihilate dominantly into a new light, unstable boson, which in turn maintains equilibrium with the standard model and decays through a small mixing parameter \\cite{Finkbeiner:2007kk}. The light mass kinematically prevents production of antiprotons \\cite{Cholis:2008vb}, allowing consistency with the PAMELA antiproton measurements \\cite{Adriani:2008zq}. Such scenarios have been widely discussed in the context of present dark matter anomalies \\cite{Finkbeiner:2007kk,Pospelov:2007mp,Cholis:2008vb,ArkaniHamed:2008qn,Pospelov:2008jd,Nelson:2008hj,Cholis:2008qq,Nomura:2008ru}.  In SUSY models, the LSP is stable by R-parity, and thus there is a unique WIMP candidate. In contrast, in these models, the symmetry stabilizing the dark matter against decay is not R-parity, and thus there is no compelling reason that there should be a single stable neutral particle. Indeed, it is very possible - perhaps even likely - that there are a variety of parities in the theory, and that many different particles together comprise the dark matter. We shall see that in models where dark matter freezes out by annihilation into a light boson, the annihilation rates of different components of dark matter {\\em are naturally equal}, even though $n_\\chi^2$ can differ between the species by orders of magnitude. As a consequence, the spectra of positrons and electrons can have interesting structures, including crests and troughs. This phenomenon arises naturally, without any tuning of the parameters. Such a model can alleviate some tension between PAMELA, which seems to require somewhat larger cross sections at lower energy, and Fermi, which requires lower cross sections but extending to higher energy.  In such models, one can explain the INTEGRAL signal with a heavy ($\\sim \\rm TeV$) mass state exciting into an MeV excited state as in \\cite{Finkbeiner:2007kk}, while explaining DAMA through a lighter ($\\sim 100 \\gev$) state exciting into a $\\sim 100 \\kev$ excited state, as in \\cite{TuckerSmith:2001hy,Chang:2008gd}. This scenario, to employ multiple states for inelastic and exciting dark matter (which we refer to as MiXDM) was proposed in \\cite{Katz:2009qq}. Although DAMA and INTEGRAL both play roles in motivating the spectrum considered, we shall not focus on these aspects, focusing instead on the higher energy cosmic ray signals (and their synchrotron counterparts). Because of the motivation in \\cite{Katz:2009qq},  we shall generically refer to these models, in which multiple states annihilate comparably into light bosons, as ``MiXDM'' models.   In this paper, we shall explore the possibility of testing such a scenario. In section \\ref{sec:multicompmodel}, we will review the setup and in section \\ref{sec:mixmodels} the model building issues, including models where different annihilation channels arise. In section \\ref{sec:multicompcr}, we discuss the cosmic ray signatures from such a scenario, and discuss the interesting spectral features that can arise. In \\ref{sec:Haze}, we go beyond the $e^+e^-$ signals, we shall argue that signatures in the microwave haze of the inner galaxy may help to show if the dark sector has multiple components. In section \\ref{sec:3_masses} we consider further generalizations, with more states, while in section \\ref{sec:scatter}, we note the important changes to direct detection experiments that can arise. Finally, we conclude in section \\ref{sec:Conclusions}. ",
        "conclusions": "\\label{sec:Conclusions}  Cosmic ray signals from PAMELA and Fermi among other have prompted a reexamination of our assumptions of the nature of dark matter. Two of the most central are the assumptions that annihilations are into standard model states, and that there is a single parity, namely R-parity, that stabilizes the WIMP.  Relaxing the latter opens the possibility that the dark matter is composed of multiple particles, with different masses and relic abundances. Such as scenario has been previously motivated to explain both DAMA and INTEGRAL. We have argued that such a ``MiXDM'' scenario, in which DM annihilates as in standard XDM into light force carriers can naturally lead to similar annihilation rates, even if the different WIMP relic densities are different by an order of magnitude. The consequence is to modify possible particle physics interpretations of the spectra. Softer spectra can arise naturally from the presence of the second, lighter WIMP state, while even bumps and dramatic spectral features can arise in the positron and electron spectra from the presence of these lighter states.  We find that these models are in good agreement with the Fermi $e^{+}+e^{-}$ flux and the PAMELA positron fraction, with the heaviest mass restricted by the HESS upper bounds. These cross sections are very similar to what is needed to explain microwave and gamma-ray excesses in the inner galaxy region,  while still below bounds from the CMB. Limits from the Galactic Ridge, while possibly constraining, are weakened when velocity-dependent effects are taken into account.  The presence of the lighter state is possibly detectable in a variety of ways. Beyond the possible features in the $e^+$ and $e^{\\pm}$ spectra, the changes in the synchrotron emission spectrum can also help distinguish it. In particular, synchrotron radiation from $e^{\\pm}$ of DM origin at frequencies up to $\\sim 500$ GHz, probed by Planck, are potentially different for these models. Unfortunately, the uncertainties in the magnetic field, ISRF and electron diffusion, as well as the great number of possible particle physics models, currently it is not possible to discriminate from synchrotron emission alone between a one DM and a two DM species scenarios that both fit the PAMELA and Fermi data. Yet, when combined with a better understanding of all the factors involved in the calculations of the microwave Haze and the synchrotron radiation(B-field, ISRF, diffusion models, dust emission, Planck instrumental errors), such a tool may offer an important handle, especially in combination with other indirect and direct DM searches.   Interesting additional possibilities exist as well. One can consider a third, heavier WIMP, although we find that its effects are highly constrained by the HESS electron limits, and that it is difficult to generate any features at high energies consistent with data. A possibility exists for subdominant components to have very large scattering rates, potentially collapsing the lighter WIMP halo, and radically altering the predictions for direct detection scattering rates.     \\vskip 0.2in \\noindent {\\bf Acknowledgments} The authors would like to thank Nima Arkani-Hamed, Gregory Dobler, Douglas Finkbeiner, Joseph David Gelfand, Lisa Goodenough, Ronnie Jansson, Dmitry Malyshev, Tracy Slatyer and Itay Yavin  for fruitful discussions. NW is supported by DOE OJI grant \\# DE-FG02-06ER41417 and NSF CAREER grant PHY-0449818. IC is supported by DOE OJI grant \\# DE-FG02-06ER41417 and also by the Mark Leslie Graduate Assistantship.   \\begin{appendix}"
    },
    "0911/0911.3059_arXiv.txt": {
        "abstract": "Sensitive and subarcsecond resolution ($\\sim$ 0.7\\arcsec) CH$_3$OH(7$_{-2,6}$ $\\rightarrow$ 6$_{-2,5}$) line and 890 $\\mu$m continuum observations made with the Submillimeter Array (SMA) towards the hot molecular circumbinary ring associated with the young multiple star {\\it Ori 139-409} are presented.  The CH$_3$OH(7$_{-2,6}$ - 6$_{-2,5}$) emission from the ring is well resolved at this angular resolution revealing an inner cavity with a size of about 140 AU. A LTE model of a Keplerian disk with an inner cavity of the same size confirms the presence of this cavity. Additionally, this model suggests that the circumbinary ring is contracting with a velocity of V$_{inf}$ $\\sim$ 1.5 km s$^{-1}$ toward the binary central compact circumstellar disks reported at a wavelength of 7 mm. The inner central cavity seems to be formed by the tidal effects of the young stars in the middle of the ring. The ring appears to be not a stationary object. Furthermore, the infall velocity we determine is about a factor of 3 slower than the free-fall velocity corresponding to the dynamical mass. This would correspond to a mass accretion rate of about 10$^{-5}$ M$_\\odot$/yr. We found that the dust emission associated with Ori 139-409 appears to be arising from the circumstellar disks with no strong contribution from the molecular gas ring. Furthermore, a simple comparison with other classical molecular dusty rings ({\\it e.g.} GG Tau, UZ Tau, and UY Aur) suggests that {\\it Ori 139-409} could be one of the youngest circumbinary rings reported up to date. Finally, our results confirm that the circumbinary rings are actively funneling fresh gas material to the central compact binary circumstellar disks, {\\it i.e.} to the protostars in the very early phases of their evolution. ",
        "introduction": "Most of the main sequence low-mass stars in our Galaxy are binaries or multiple systems, however its formation processes are still not well understood. Our poor understanding about the formation of the multiple stars has been mainly attributed to the modest sample of relatively well studied protobinary or multiple systems at millimeter and infrared wavelengths \\citep{mo2007,Lau2004}. Surveys have been made at millimeter wavelengths in an attempt to find new very young multiple systems, however these have been strongly limited by sensitivity \\citep{dutrey2001}. Now, Submillimeter interferometers begin to open a new ``{\\it radio window}'', providing highly-excited molecular and continuum observations with a better signal-to-noise ratio (in comparison to the millimeter observations) and a sub-arcsecond resolution which may be able to trace the innermost parts of close-by molecular and dusty circumbinary or circum-multiple disks.  {\\it Ori 139-409} is part of a multiple system of compact millimeter and centimeter continuum sources located in southern most part of Orion South \\citep{Zapataetal2004a,Zapataetal2005,Zapataetal2007, eis2006,eis2008,Zap2004b} and at a distance of 414 pc \\citep{Mentenetal2007}. This source is associated with a hot (T $\\geq$ 100 K) and dense ($\\rho$ $\\simeq$ 10$^8$ cm$^3$) molecular circumbinary ring surrounding two compact dusty disks (with truncated sizes of less than 50 AU) which are associated with intermediate mass protostars \\citep{Zapataetal2007}. The circumbinary molecular gas ring has a deconvolved size of $0\\hbox{$.\\!\\!^{\\prime\\prime}$}64 \\pm 0\\hbox{$.\\!\\!^{\\prime\\prime}$}03 \\times 0\\hbox{$.\\!\\!^{\\prime\\prime}$} 45 \\pm 0\\hbox{$.\\!\\!^{\\prime\\prime}$}04$ (or $\\rm 294~AU \\pm 14~AU \\times 207~AU \\pm 18$ AU) with a PA of 87$^\\circ\\pm 7^\\circ$.  \\citet{zapa2009} reported a highly collimated SO(6$_5$$\\rightarrow$5$_4$) molecular and monopolar outflow with a northeast-southwest orientation ejected from {\\it Ori 139-409} and that forms part of a more extended (2$'$) monopolar CO(2$\\rightarrow$1) outflow that emanates from this region \\citep{Schmid-Burgketal1990}.  In this paper, we present submillimeter and subarcsecond line and continuum observations made with the Submillimeter Array toward the protobinary system {\\it Ori 139-409}. We find that the molecular emission is well resolved at this angular resolution, revealing an inner cavity at the center of the molecular structure with a size of about 140 AU. A LTE model of a Keplerian disk with a central hole agrees well with the line observations and further reveals that the circumbinary ring is infalling towards the compact circumstellar disks. ",
        "conclusions": "Our Submillimeter interferometric observations together with the LTE model revealed that the circumbinary ring surrounding the protobinary {\\it Ori 139-409} is infalling toward the central compact circumstellar disks.  Moreover, the channel-velocity maps of the CH$_3$OH(7$_{-2,6}$ - 6$_{-2,5}$) line emission show the presence of an inner cavity centered where the circumstellar dusty disks are located. Our model of an infalling and Keplerian disk with a central cavity seems to confirm this hypothesis (see Figure 1). The size of this cavity is estimated to be of about 140 AU $\\pm$ 20 AU.  There are a few mechanisms that would have formed the inner cavity observed in the circumbinary ring. One of them would be due to the dynamic interaction of the circumbinary disk and the binary circumstellar compact disks located in its center. When a multiple system is formed the tidal effects of the stars creates gaps or holes in the center of the circumbinary disks \\citep{mo2007}. This is thought to be the case for the circumbinary rings: GG Tau, SR 24 and UY Aur, see \\citet[][]{Gui1999,Duv1998,And2005}. On the other hand, photo-dissociation or opacity effects of this molecule (CH$_3$OH) close to the central binary protostars could be other possibilities. The photo-evaporation in flattened molecular gas disks is expected to be very important principally in the formation of massive stars because of their strong UV fields. However, as we have a protobinary system with an enclosed dynamical mass of 9 M$_\\odot$ (or say  about 4 M$_\\odot$ for each star), this seems not to be an likely alternative. Moreover, the SED for Ori 139-409 at centimeter wavelengths does not show any contribution of free-free emission (Figure 3).  It is interesting to note that dust emission arising from {\\it Ori 139-409} is well fitted by a single power-law with $\\alpha$ = 2.44 as obtained in our Figure 3. The continuum emission at 1.3 cm and 7 mm arises completely from the circumstellar disks as reported by \\citet[][]{Zapataetal2007} and suggested by \\citet[][]{Zapataetal2004a}. We additionally tapered the 7 mm continuum data from \\citet[][]{Zapataetal2007} at an angular resolution similar to these observations and did not detected any emission from the ring. If such contribution was stronger one then might be expect a power-law with two components (one associated with the circumstellar disks obtained with a high angular resolution, and the other component with both contributions from the circumbinary ring and the compact disks), but this is not the case for {\\it Ori 139-409}. Furthermore, the deconvolved size at 850 $\\mu$m of the circumbinary ring is smaller as compared with the one of the CH$_3$OH(7$_{-2,6}$-6$_{-2,5}$) line emission (see Table 1). Assuming that the emission at 850 $\\mu$m is optically thin, isothermal (100 K) and a gas-to-dust ratio of 100, we estimated an enclosed mass of 0.03 $M_\\odot$ for the source detected at this wavelength. This value is similar to the masses of the circumstellar disks reported at 7 mm, suggesting that we are seeing thermal emission from the disks.  These arguments thus suggest that the continuum emission seems to be only originate from the dusty circumstellar disks with not strong contribution from the circumbinary molecular gas ring. This physical phenomenon might be also the case for the circumbinary disk in the multiple star system SR 24N, where the circumbinary disk is only detected in the CO line emission and not in the millimeter continuum \\citep[][]{And2005}. The dust emission from the circumstellar disks also appears to be very faint. In the case of GG tau, the mm. dust emission is detected in the ring as well as in the position of the circumstellar disks, see \\citet[][]{Gui1999}.  Both position-velocity diagrams made at different angles show a good agreement (Figure 4). In the top image of Figure 4 it is also clear that the kinematics of the molecular gas in the ring could be also fitted by a rigid body law, as suggested in Zapata et al. (2007).  If one do a simple comparison with other well known protobinary systems as GG Tau, UZ Tau, and UY Aur, one notes that the protobinary system {\\it Ori 139-409} could be one of the youngest protobinary system reported up to now. A clear example is that all the former multiple systems have counterparts at infrared wavelengths associated with the central protobinary system and/or the circumbinary disks \\citep[][]{Ro1996,Ghez1994,Hio2007} suggesting that they are likely class I or later type objects. This is contrary to what is observed in {\\it Ori 139-409} where the infrared emission (2.2, 8.8 and 11.7 $\\mu$m ) is not detected at all \\citep[][]{Gau1998,smi2004}. {\\it Ori 139-409} is still very embedded in its parent molecular cloud and therefore is related with a true class 0 object. Moreover, there are a few pronounced differences between the circumbinary disk observed in GG Tau, UZ Tau, and UY Aur and that one observed in the {\\it Ori 139-409}. The millimeter continuum emission from those circumbinary disks in Taurus and Auriga complexes are very strong compared to that observed in {\\it Ori 139-409}, while the molecular emission seems to go in the other sense, with the molecular emission from {\\it Ori 139-409} stronger.  The size of the cavity found in the Ori 139-409 system is in good agreement with the size of the cavity suggested for the L1551 binary system (of about 180 AU for $q=0.4$) estimated from theoretical models by \\citep{Picha2005}. The L1551 binary system appears to be a very similar young stellar object to Ori 139-409. Both binary systems have de-projected separations of 50--80 AU, semi-major disk axes of about 10-40 AU, and masses of the components of approximately 0.05 M$_\\odot$. A possible reason of why the circumbinary cavity in Ori 139-409 is something smaller than the one in L1551 could due to the former is very young and the stars in the center have not time to made the final size of the cavity."
    },
    "0911/0911.1789_arXiv.txt": {
        "abstract": "As part of the WIYN High Image Quality Indiana Irvine (WHIQII) survey, we present 123 spectra of faint emission-line galaxies, selected to focus on intermediate redshift ($.4 \\lesssim z \\lesssim  .8$) galaxies with blue colors that appear physically compact on the sky. The sample includes 15 true Luminous Compact Blue Galaxies (LCBGs) and an additional 27 slightly less extreme emission-line systems.  These galaxies represent a highly evolving class that may play an important role in the decline of star formation since $z \\sim 1$, but their exact nature and evolutionary pathways remain a mystery.  Here, we use emission lines to determine metallicities and ionization parameters, constraining their intrinsic properties and state of star formation. Some LCBG metallicities are consistent with a ``bursting dwarf'' scenario, while a substantial fraction of others are not, further confirming that LCBGs are a highly heterogeneous population but are broadly consistent with the intermediate redshift field. In agreement with previous studies, we observe overall evolution in the luminosity-metallicity relation at intermediate redshift.  Our sample, and particularly the LCBGs, occupy a region in the empirical R$_{23}$-O$_{32}$ plane that differs from luminous local galaxies and is more consistent with dwarf Irregulars at the present epoch, suggesting that cosmic ``downsizing'' is observable in even the most fundamental parameters that describe star formation.  These properties for our sample are also generally consistent with lying between local galaxies and those at high redshift, as expected by this scenario.  Surprisingly, our sample exhibits no detectable correlation between compactness and metallicity, strongly suggesting that at these epochs of rapid star formation, the morphology of compact star-forming galaxies is largely transient. ",
        "introduction": "\\label{sec:intro} Faint blue galaxies have long been a subject of study \\citep[e.g.,][and references therein]{koo94,ellis97}. The most extreme class, Compact Narrow Emission Line Galaxies (CNELGs) were first noted by \\citet{koo94}. Generally speaking, they are luminous ($\\lesssim B_*$) and starbursting as indicated by blue optical colors (i.e. B-V $\\lesssim 0.6$).  They have bright, narrow emission lines, are compact (i.e. $r_h \\lesssim 3$ kpc), and are predominantly found at intermediate redshift ($0.4 < z < 1.2$). Later studies of these objects \\citep[e.g.,][]{guz98,bvz06cnelg} and similar populations such as luminous compact galaxies \\citep[e.g.,][]{ham01} and Luminous Compact Blue Galaxies \\citep[LCBGs, e.g.,][]{noeske06}, have revealed that these galaxies have very high star formation rates, perhaps -- depending on the definitions used -- representing up to $\\sim 45\\%$ of total star formation at intermediate redshift, and 20\\% of the field number density \\citep{phillips97,guz97,noeske06}.  At the present epoch, however, these galaxies are far less common \\citep{koo94,phillips97,guz97,werk04,noeske06}, suggesting that these objects may play an important role in the global decline of star formation \\citep{madau96} and the evolution in the blue luminosity function \\citep{lilly95,ellis97}.  Despite this, it remains unknown what their end state is in the local universe.  While there are subtle differences in the terminology used in the literature, we follow the definition of \\citet{noeske06}, and use the ``LCBG'' term, as essentially all intermediate redshift objects meeting LCBG criteria also show narrow emission lines.  Two primary scenarios have been suggested for the evolutionary scenario of LCBGs.  The ``bursting dwarf'' hypothesis \\citep[e.g.,][]{koo95,guz98,hoyos05} suggests that LCBGs are low mass galaxies undergoing a single major starburst that will exhaust their gas.  They then fade after the starburst by several magnitudes into present day dwarf spheroidal or elliptical (dSph/dE) galaxies.  While it is possible that underlying pre-starburst stellar populations might exist that prevent the galaxies from fading this much, the scenario still predicts low metallicities, consistent with dwarf galaxies rather than bright galaxies.  The alternative interpretation is that the LCBGs are {\\it in situ} bulge-formation of lower luminosity spirals \\citep[e.g.,][]{ham01,bvz06cnelg}.  The bulges-in-formation scenario predicts higher metallicities than those likely to be observed in local dwarf galaxies \\citep{KZ99}. The bulge and dwarf scenarios are not mutually exclusive -- the LCBGs may be a heterogeneous population, or, i.e. evolve to become mostly lower luminosity spirals now \\citep{bvz01,ham01,noeske06,bvz06cnelg}.   The WIYN High Image Quality Indiana Irvine (WHIQII) survey aims to address the question of the evolutionary paths of LCBGs by examining the metallicities of a large sample derived using the strong $\\oii$, $\\oiiia$, $\\oiiib$, and $H\\beta$ lines; primarily through the use of the $R_{23} = (\\oii+\\oiii)/ H\\beta$ and $O_{32} = \\oiii / \\oii$ indicators (see \\S \\ref{sec:met}).     This paper is organized as follows: in \\S \\ref{sec:isamp} we describe the imaging and initial sample selection.  In \\S \\ref{sec:dat}, we discuss observation and reductions of of the spectroscopic survey and selection of the sample used for analysis.  In \\S \\ref{sec:numden} we briefly address the number density of LCBGs determined from WHIQII. In \\S \\ref{sec:met} we discuss the metallicities of the sample and how they are derived as well as the luminosity-metallicity (LZ) relation.   \\S \\ref{sec:disc} discusses trends and local comparisons, and in \\S \\ref{sec:conc} we present our conclusions.  Where relevant, we assume a WMAP5 \\citep{kom09WMAP} $\\Lambda$CDM cosmology. ",
        "conclusions": "\\label{sec:conc} In this paper, we have analyzed the LCBGs found in the WHIQII survey to constrain their evolutionary path to the current epoch.  \\begin{enumerate} \\item LCBGs are likely a heterogeneous population.  The range of metallicities is too large to be consistent with an all-dwarf or all-spiral scenario, and some of the galaxies show metallicities simply too high to evolve into a typical present day dwarf.  Combined with other results like those of \\citet{ham01, Garland04, bvz06cnelg, noeske06}, this study suggests that they are a mix of true dwarfs and in situ bulges forming in lower-luminosity spirals.  The WHIQII data effectively rule out either an all-dwarf or all-bulge descendant population. \\item WHIQII objects, including LCBGs, show signs of redshift evolution in the luminosity-metallicity relation.  They have lower metallicities at fixed stellar mass or luminosity, consistent with the results of \\citet{KK04}. \\item Compact star forming galaxies have no strong correlation between size, $r_h$, and oxygen abundance or $M_*$.  That is, apparent size does not appear to correlate with the likely end product. Thus, apparently compact objects at intermediate redshift are not necessarily intrinsically compact or low mass galaxies. \\item LCBGs appear to have high ionization parameters relative to local galaxies, consistent with the intermediate redshift field \\citep{KK04}.  While this result may be physically related to their compactness and intense star formation, it is also true of other field galaxies at intermediate redshift.  Hence, the observed evolution of LCBGs may be generic ``downsizing'' enhanced by selection effects. \\end{enumerate}    If LCBGs are indeed a heterogeneous population, it may be that unusually blue galaxies seem uncommon simply because the \\emph{global} star formation rate declines, more or less irrespective of star-forming galaxy type \\citep{noe07,cw08}.  Hence, the conditions in their star forming regions would be intermediate between local galaxies and high-redshift objects such as Lyman Break Galaxies \\citep{steidel98lbgs,shap01lbgs,shap03lbgs}.  This is one of the many faces of galaxy ``downsizing,'' and the WHIQII data appear to be consistent with this scenario.  Hence, while LCBGs are a straightforward category to observationally define, the scatter in their properties means that they must be individually examined to understand their place in the wider context of galaxy formation.  \\vspace{10mm} We wish to thank Lisa Kewley for extensive and insightful comments; we also thank David Koo, Janice Lee, Alice Shapley, Marla Geha, and the anonymous referee for conversation and helpful suggestions, as well as J.D. Smith for assistance with the imaging observations. This research was supported by NSF grant AST05-06054 and the Center for Cosmology at UC Irvine.  The authors wish to recognize and acknowledge the very significant cultural role and reverence that the summit of Mauna Kea has always had within the indigenous Hawaiian community.  We are most fortunate to have the opportunity to conduct observations from this mountain.  {\\it Facilities:} \\facility{WIYN (MiniMo)}, \\facility{Keck:I (LRIS)}"
    },
    "0911/0911.3745_arXiv.txt": {
        "abstract": "This brief article sums up the possible imprints of loop quantum gravity effects on the cosmological microwave background. We focus on semi-classical terms and show that \"Big Bounce\" corrections, together with the \"pre Big Bounce\" state, could modify the observed spectrum. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0911/0911.3798_arXiv.txt": {
        "abstract": "We analyse the influence of rotation on shapes of pulse profiles of fast-rotating (millisecond) pulsars. Corotation has two opposing effects: 1) the caustic enhancement of the trailing side (TS) by aberration and retardation (AR), which squeezes the emission into a narrower phase interval; 2) the weakening of the TS caused by the asymmetry of curvature radiation about the dipole axis. Analysis of the radii of curvature of electron trajectories in the inertial observer's frame (IOF) enables these two effects to be considered together. We demonstrate that for dipolar magnetic field lines on the TS there exists a `caustic phase' beyond which no emission can be observed. This phase corresponds to the zero (or minimum) curvature of the IOF trajectories and \\emph{maximum} bunching of the emission. The maximum gradient of polarisation angle (PA) in the S-shaped PA curve is also associated with the curvature minimum and occurs at exactly the same phase. The asymmetry of trajectory curvature with respect to the dipole axis affects the curvature emissivity and the efficiency of pair production, suggesting a \\emph{minimum} at the caustic phase. Emission over a fixed range of altitudes, as expected in millisecond pulsars, leads to broad leading profiles and sharp peaks with a cutoff phase on the TS. We apply our results to the main pulse of the 5 ms pulsar J1012+5307. ",
        "introduction": "We analyse two specific issues that are important for the appearance of pulse-averaged profiles of pulsars: 1) the radial extent of a region emitting at a given radio frequency, and 2) the asymmetry about the magnetic axis of the curvature of electrons' trajectories measured in the observer frame. The influence of both of these factors on profile shape is closely associated with the rotation of the magnetosphere, so they should be most noticeable for rapidly rotating objects such as the millisecond pulsars (MSPs). Unlike normal pulsars, MSPs do not exhibit radius-to-frequency mapping (RFM, see figs.~16 and 17 in Kramer et al.~1999). This might possibly be interpreted as near-surface emission with negligible radial extent (McConnell et al.~1996). However, Dyks, Rudak \\& Demorest (2009) have shown that at least some components in radio pulse profiles are produced when our sightline cuts through the curvature emission from radially-extended plasma streams. In this stream-cut scenario, the observed location of a component % is determined by the frequency-independent moment at which our line of sight crosses the plane of a thin stream. When this happens, the component becomes simultaneously recorded at all frequencies that are emitted at the altitude of the cut. There is no RFM, because positions of components are fixed by the geometry of the cut.  This interpretation requires that a broad range of radio frequencies is emitted by the stream at a fixed altitude. Such a possibility arises naturally in the stream-cut scenario, because density of plasma in a stream decreases in the transverse direction from the stream axis. Therefore, regardless of the altitude that is considered, a broad range of plasma densities (and plasma frequencies) is available at different transverse distances from the stream axis. It is then possible to have a radially-extended emission region (a stream) which provides a range of plasma frequencies at each altitude (the `radius-to-frequency' mapping is replaced with the `stream-diameter-to-frequency' mapping). Thus, the radially-extended emission at a fixed radio frequency is possible without violating the foundation of the RFM idea, that relates the observed frequency to the local plasma frequency at the emission point.\\footnote{Needless to say, even in the stream the density must also decrease with altitude, so the RFM is present, but becomes a second-order effect.}   The effect of rotation on a radially-extended source can be expected to make intrinsically symmetric pulse profiles appear asymmetric with respect to their centre, changing the phases of their components, their intensities and their widths. In the following section (\\ref{widths}) we introduce a simple (small-angle) analytical model to determine the likely observable effects in the profiles of fast-rotating pulsars with emitting regions of finite height. These caustic effects are well studied {\\it numerically} and are routinely employed to interpret high-energy pulse profiles of pulsars (Morini 1983; Smith 1986; Romani \\& Yadigaroglu 1995; Cheng et al.~2000, Dyks \\& Rudak 2003). Here, however, we focus on the {\\it analytical} study of the phenomenon, which is feasible for low emission altitudes ($r \\la 0.1\\rlc$) or small distances from the dipole axis. This limitation therefore allows possible applications to the radio profiles as well as to specific gamma-ray peaks, eg.~those generated by the trailing caustics of the two-pole caustic model.  The results are used to create artificial but instructive pulsar profiles in Section (\\ref{models}). In Section (\\ref{intensity}) we analyse the curvature of the particle trajectories and examine the impact of considering curvature radiation on model profiles. Then, in Section (\\ref{MSP}), we see how far our results can model the main pulse profile of J1012+5307. Finally, conclusions are drawn in Section (\\ref{conc}). ",
        "conclusions": "\\label{conc}   We find that the shape of a profile is affected by two opposing effects of rotation: 1) the caustic enhancement of the TS, and 2) the increase of $\\rhiof$ on the TS which decreases the centripetal acceleration and can thereby weaken the emissivity. The first can be viewed as a direct result of the straightening of an electron trajectory in the IOF, because the radiation emitted within $\\Delta r$ is received within $\\dphi \\simeq \\Delta r/\\rhiof$. The radiation that constitutes the caustic peaks is emitted from the inflection points of IOF trajectories (or from regions of minimum IOF curvature in a non-orthogonal case). Formally, the caustic peaks are expected to occur at the phase of the center of the relativistically-shifted S-curve of the polarisation angle. The caustic phase, if calculated for the maximum $s=1$, becomes the `maximum possible detection' phase for tangent-to-$\\vec B$ radiation within the open region. Fast-rotating and highly-inclined pulsars are then expected to have pulse profiles with sharply defined (distinctively marked) trailing sides. On the LS, AR causes radiation from the outflowing electrons to be detected at a phase that monotonically increases with the emission altitude. The LS of  the profile is thus expected to grow only gradually, creating a leading wing (the extent of which in practice is likely determined by the cooling timescale of the electrons).  Whether the leading or trailing side of the profile is the more enhanced by rotation is likely to be determined by the intrinsic dependence of the coherent emission mechanism on $\\rhiof$. In the case of the CR the emissivity is entirely driven by the macroscopic centripetal acceleration and we may expect a minimum in the zero curvature region. Otherwise, a caustic peak will dominate the profile (cf.~Fig.~\\ref{profis}b with \\ref{dsn}c). Statistical analysis of millisecond profiles, could therefore provide limits on the possible value of the exponent $a$ in the relation $F_{coh}\\propto \\rhiof^a$. Inspection of published profiles suggests that $a\\la 0$ (as opposed to $a=-2/3$), although the number of good-quality millisecond profiles is very limited, and the interpretation of profiles in terms of a simple sum of emission components ignores possible absorption effects.  According to popular electrodynamic assumptions, the accelerating electric field is symmetrical with respect to the main meridian. However, the IOF curvature radii are intrinsically symmetric about the caustic phase at a given height (and asymmetric with respect to the dipole axis), so circular cones centered on this phase are theoretically likely. Because of the IOF-curvature asymmetry within the open region of fast rotating MSPs, the traditional, dipole-axis-centered nest of cones may be considered less natural than the decentered geometry of Fig.~\\ref{caps}. In the extreme case of the fastest known MSPs ($P\\sim 1.5$ ms), the caustic phase occurs at the outermost trailing edge of the polar cap, so it becomes difficult to generate a conal structure at all (possibly yielding the simplicity of the MP of B1937$+$21).  The main pulse of \\jtt exhibits asymmetry of flux and width which can be interpreted in terms of the corotational effects within a radially-extended emission region. The observed width of the main pulse suggests that the outermost components have $s \\simeq 1$. This is roughly consistent with the widths of the outermost components if $r_1 \\sim \\rns$ and $\\Delta r \\sim \\rns$. However, based solely on the symmetry of $\\rhiof$, the profile should then be symmetrical with respect to the caustic phase located somewhere within the main pulse window, which is not the case. This implies that either the profile is dominated by unknown $\\rhiof$-independent effects, or the emission region is located fully on the LS of the caustic coordinates. In the uniform-altitude scenario (same $r_1$ and $r_2$ for all components) this requires the rings/cones to be located asymmetrically with respect to the main meridian plane, but the location is not uniquely constrained because of the degeneracy of parameters. In the alternative cone-altitude scenario, the asymmetry with respect to $\\scst$ can still be achieved by a region that is symmetrical with respect to the main meridian in th CF. However, it is hard physically to justify the cone-altitude geometry. Neither of the scenarios is naturally consistent with the observed width of the main pulse, which is the same as the opening angle of the near-surface polar beam.  A possible solution to this problem is a model in which the polar tube at low altitudes is fully filled with the radio emission, and the minima observed in the MP result from absorption effects. The absorption could be provided by dense, opaque, radially-extended plasma columns in the polar tube. In this model the MP is not regarded as a sum of emission components. Rather, it is a single, box-like feature, with the minima of different width indented by the plasma columns. We expect that the minima will have essentially the same widths as the emission components in the emissive version of the model presented in Section 2. That is, the minima on the LS/TS will correspondingly be broad/narrow. Then, if the absorption occured (on average) above $0.14\\rlc$, it would be possible to avoid the problem of profile symmetry around $\\phcst$. Such an absorption-based scenario is supported by the box-like look of the MP of J1012$+$5307, and presents a promising subject for future study."
    },
    "0911/0911.4492_arXiv.txt": {
        "abstract": "The eclipsing low-mass X-ray binary 4U 1822--371 is the prototypical accretion disk corona (ADC) system. We have obtained new time-resolved UV spectroscopy of 4U~1822--371 with the Advanced Camera for Surveys/Solar Blind Channel on the {\\it Hubble Space Telescope} and new $V$- and $J$-band photometry with the 1.3-m SMARTS telescope at Cerro Tololo Inter-American Observatory. We use the new data to construct its UV/optical spectral energy distribution and its orbital light curve in the UV, $V$, and $J$ bands. We derive an improved ephemeris for the optical eclipses and confirm that the orbital period is changing rapidly, indicating extremely high rates of mass flow in the system; and we show that the accretion disk in the system has a strong wind with projected velocities up to 4000~km~s$^{-1}$. We show that the disk has a vertically-extended, optically-thick component at optical wavelengths. This component extends almost to the edge of the disk and has a height equal to $\\sim$0.5 of the disk radius. As it has a low brightness temperature, we identify it as the optically-thick base of a disk wind, not as the optical counterpart of the ADC. Like previous models of 4U~1822--371, ours needs a tall obscuring wall near the edge of the accretion disk, but we interpret the wall as a layer of cooler material at the base of the disk wind, not as a tall, luminous disk rim. ",
        "introduction": "4U~1822--371 is a low-mass X-ray binary star (LXMB) with an orbital period of $P_{orb} = 5.57$~hr \\citep{mas80}. The compact star in the system is an X-ray pulsar with spin period $P_{pulse} = 0.593$~s \\citep{jon01}  and is, therefore, a neutron star. The UV/optical spectrum of 4U~1822--371 is rich with emission and absorption lines \\citep{cha80, mas82a, mas82c, har97, cow82, cow03, jon03, hut05}. Doppler tomography of the \\ion{He}{2}~$\\lambda4686$ and \\ion{O}{6}~$\\lambda3811$ emission lines confirms expectations that there is an accretion disk around the neutron star fed by a stream of gas from the secondary star \\citep{cas03}.  4U~1822--371 is one of the rare eclipsing LXMBs \\citep{mas80, whi81}. Many analyses of the eclipse have been published \\citep{whi81, whi82, mas82a, mas82b, hel89, bap02, cow03} and, although the detailed results of the various analyses often differ greatly, there is substantial agreement on several points. The eclipse is a transit of the secondary star across the accretion disk and, since the eclipse is very broad at optical wavelengths, the accretion disk is large, although exactly how large is uncertain: \\citet{mas82b} find $r_{\\rm disk}/a = 0.28$ to 0.38 and \\citet{hel89} find $r_{\\rm disk} / a = 0.58\\pm 0.08$, where $r_{\\rm disk}$ is the radius of the disk and $a$ is the separation of the stars. The X-ray eclipse is also broad. Its flux at minimum is at about 50\\% of the uneclipsed flux level and it has no sharp features, showing that the X-ray emission comes from a partially-eclipsed extended cloud around the neutron star - an accretion disk corona (ADC) \\citep{whi81,mas82a}. The corona extends to $\\sim 0.5 r_{\\rm disk}$ in the plane of the disk and is vertically extended. Estimates of the ratio of its height to its radius in the plane of the disk range from 0.3 to 1 \\citep{whi82, mas82a, hel89}. The orbital inclination lies in the range $78^\\circ < i < 84^\\circ$.  The ratio of the X-ray luminosity of 4U1822--371 to its optical luminosity, $L_X/L_{opt}$, lies between $\\sim 15$ and $\\sim 65$, much less than that for typical LXMBs, indicating that the ADC is optically thick and blocks most of the X-ray emission from the inner disk and neutron star \\citep{gri78,mas80}. Fits to the X-ray spectrum using realistic Comptonization models yield electron temperatures near $T_e \\approx 5 \\times 10^7$~K and electron scattering optical depths in the range $\\tau_{e}=13$--26 depending on the geometry of the ADC and the Comptonization model \\citep{par00, iar01}. The mean unabsorbed apparent luminosity of 4U~1822--371 in the 0.1--100 keV energy range is $1.15 \\times 10^{36}\\ \\textrm{erg s}^{-1}$ for a distance of 2.5~kpc \\citep{iar01}. As most of the X-ray flux escapes roughly perpendicularly to the orbital plane and is beamed away from the Earth, the true luminosity is much greater than the apparent luminosity. If one corrects for the ADC obscuration by simply multiplying $L_X$ by a factor to raise $L_X/L_{opt}$ to the typical value for LXMBs ($\\sim 500$), the true X-ray luminosity is $\\sim 10^{37}\\ \\textrm{erg s}^{-1}$. Even this should be taken as a lower limit because the disk is observed edge-on, reducing $L_{opt}$. A true luminosity near the Eddington luminosity is not impossible.  All solutions for the X-ray eclipse have invoked a vertically-extended wall of optically-thick material whose height varies with angle around the neutron star \\citep{whi81}. The wall generally has been identified as a vertically extended disk rim, and, indeed, \\citet{hel89} showed the wall must be near the edge of the disk.  The rim invoked by \\citet{hel89} has $0.08 < h_{rim}/r_{\\rm disk} < 0.22$, while the maximum height of the rim invoked by \\citet{whi82} approached $h_{rim}/r_{\\rm disk} \\approx 0.5$. \\citet{iar01} also needed a thick belt of absorbing material around the ADC to match the observed X-ray spectral energy distribution (SED). Placing the belt at the edge of the disk, they found the angle subtended by the absorbing region is $16^\\circ$, or $h_{rim}/r_{\\rm disk} \\approx 0.29$. These rims are an order of magnitude taller than the expected thickness of the disk proper and cannot be pressure supported. \\citet{whi82} invoked turbulence to support the rim but the turbulence would have to be highly supersonic, needing Mach numbers between 7 and 15. While it is not inconceivable that supersonic turbulence could be generated by the impact of the accretion stream on the disk, it is unclear how supersonic turbulence could be maintained around the entire circumference of the disk.  It has been universally assumed that the only contributors to the UV/optical continuum are 1) the accretion disk, 2) its extended rim, which may be irradiated and have a higher temperature on its surface towards the neutron star, and 3) the irradiated secondary star \\citep{mas82b,hel89,puc95}. It also has been assumed that the ADC is optically thin at UV/optical wavelengths, even though it is optically thick to electron scattering at X-ray wavelengths. These contradictory preclude a UV/optical contribution from the ADC or from vertical structures other than the rim.  The ADC and the eclipses lend particular importance to 4U~1822--371. ADC sources are more numerous than suggested by the small number that have been identified. The luminosity of an ADC is typically $\\sim100$ times less than the luminosity of the inner disk and neutron star, so ADCs are visible only in those systems with orbital inclinations near enough to $90^\\circ$ that the outer accretion disk or the ADC itself blocks most of the flux from the more luminous components of the binary \\citep{whi82}. Many ADC sources must be lurking among the low-inclination LXMBs, perhaps 5 -- 10 times more than are currently known. 4U~1822--371 is the prototype for these systems. Models for accretion onto neutron stars and black holes in X-ray binaries or onto supermassive black holes in AGN need much more than flat, unadorned accretion disks. They must also include high-temperature flows, vertical extension, winds, and irradiation, all of which can vary with time \\citep{nar98,bla99,wij99,nar00,las01}. ADCs are examples of high-temperature accretion flows and vertically-extended structures. The eclipses of 4U~1822--371 offer the opportunity to map these structures in detail.  We have obtained new {\\it HST} objective-prism ultraviolet spectroscopy and $V$ and $J$ photometry of 4U~1822--371, from which we construct its SED and its orbital light curves in the UV, $V$, and $J$ bands (Section 2). In Section 3 we derive an improved ephemeris for the optical eclipses and confirm that the orbital period is changing rapidly, indicating extremely high rates of mass flow in the system. In Section 4 we show that the accretion disk has a strong wind with projected velocities up to 4000~km~s$^{-1}$. The core of this paper is Section 5, in which we model the orbital light curves with our light curve synthesis program. We show that much of the disk is vertically extended and optically thick at UV/optical wavelengths. This vertical structure extends almost to the edge of the disk and has a height equal to $\\sim$0.5 of the disk radius. As this structure has a brightness temperature near $3 \\times 10^4$~K, we identify it as the optically-thick base of the disk wind, \\emph{not} as the optical counterpart of the ADC. We, too, need a tall obscuring wall near the edge of the accretion disk, but we interpret the wall as layer of cooler material at the base of the disk wind, not as a tall disk rim. Our results are summarized in Section 6. ",
        "conclusions": "Combining our new times of eclipses with the previously published times, we have derived an improved ephemeris for the eclipses in the UV/optical light curves of 4U~1822--371 (equation~\\ref{Ephemeris}). The quadratic term in the ephemeris yields a timescale for a change in orbital period of ${{P} / {\\dot P}} = (3.0 \\pm 0.3) \\times 10^{6}\\ \\textrm{yr}$. The deduced rate of mass transfer is large, probably greater than $3 \\times 10^{-8}\\ M_\\odot \\ \\textrm{yr}^{-1}$. Mass cannot be accreting onto the neutron star at this rate without violating the Eddington limit, so much of the transferred mass must be lost from the system.  We have inferred a strong disk wind from the \\ion{C}{4} emission line in the UV spectrum. The wind is axially symmetric and has projected outflow velocities up to 4000~km~s$^{-1}$. The flux in the \\ion{C}{4} emission lines does not decrease during eclipses. Since regions less than $\\sim 0.5r_{\\rm disk}$ above the disk are at least partially eclipsed, the \\ion{C}{4} emission arises from regions yet further above the disk. The outflowing material closer to the disk either does not produce \\ion{C}{4} emission or is hidden by other optically-thick structures.  Much of this paper was devoted to developing a model for the accretion disk in 4U~1822--371 from fits to the UV/optical orbital light curves. The following properties of the accretion disk appear to be robustly determined. The disk is large; the fitted radius is $r_{\\rm disk}/a = 0.40 \\pm 0.01$, close to the tidal truncation radius. To avoid an eclipse that is deeper and narrower than the observed eclipse, the neutron star and the central regions of the accretion disk must be obscured. The dual requirements that the models must produce synthetic light curves that fit the observed eclipse and must obscure the neutron star and central disk forces most of the obscuring material to be at intermediate radii in the disk, not at the edge of the disk. We modeled the obscuring material as an optically thick torus. The torus extends more than 3/4 of the way to the edge of the disk at UV wavelengths and nearly all the way to the edge in the $V$- and $J$-bands; and it extends to a height of $\\sim 0.5r_{\\rm disk}$ above the disk at all wavelengths.  The surface brightness of the outer (visible) wall of the torus is only $\\sim 26,000$~K. While this estimate is highly uncertain, the true brightness temperature is unlikely to be more than a factor of 2 -- 3 higher and is far less than coronal temperatures. We have, therefore, interpreted the torus as the optically-thick base of the disk wind, not as the optical counterpart of the X-ray ADC.  We modeled the large hump in the orbital light curve of 4U~1822--371 with a combination of heated face of the secondary star and a variable-height, optically-thick wall located at the edge of the disk. All previous models for the orbital light curve have needed a similar wall, but we now suggest that the wall is actually slightly closer to the neutron star than the edge of the disk and is part of the disk wind, not the disk rim. The wall is a cooler, darker layer at the base of the disk wind just above the disk.  Our results thus yield a more complex view of the structure of the disk in 4U~1822--371. The previous X-ray observations have shown that there is a vertically-extended, optically-thick ADC stretching roughly half way out to the edge of the disk. Our UV/optical observations show that most of the outer disk emits a cooler disk wind. The wind is is optically thick close to the accretion disk but becomes optically thin at a height of $\\sim 0.5r_{\\rm disk}$ above the disk. Close to the disk the wind is yet cooler and darker, forming a relatively dark wall around disk. The disk itself appears to be nearly entirely buried within and obscured by these vertically-extended structures."
    },
    "0911/0911.1612_arXiv.txt": {
        "abstract": "s{ Nancay radiotelescope is involved in high precision timing since 20 years. Since 2004, a coherent dedispersion instrumentation enables numerous routine observations on more than 200 pulsars using half of the time if this 100-meters class radiotelescope. Two main programs are currently conducted. A large set of young and old pulsars is timed for a multi-wavelength approach, complementary to the very successful high energy observations of pulsars done by FERMI. A set of highly stable millisecond pulsars is monitored as our contribution to the European Pulsar Timing Array in order to probe the cosmological Gravitational Wave Background. } ",
        "introduction": "Pulsars are highly magnetized neutron stars from which we receive collimated radio beams every rotation. The most stable pulsars are used as clock for many different kind of studies. Such pulsars located in a relativistic binary are used to put constrains on the Gravitation theories \\cite{2006Sci...314...97K}. An array of highly stable pulsars distributed over the sky is used to search for a cosmological gravitational waves background \\cite{2006ApJ...653.1571J}. An effort towards an international Pulsar Timing Array is taking place in every large telescope in the world in order to share all the observations and set the best limit on a gravitational waves background coming from the early Universe.  With a collecting area corresponding to a 100m dish, the Nan\\c cay radiotelescope is a centimetric telescope among the largest in the  world. Based on a Kraus design, the telescope has two receivers providing a continuous coverage from 1.1 to 3.5GHz. The 1.4 to 2GHz range is a good choice to observe pulsars between embarrasing interstellar effects at the lower frequencies and faint emission received from pulsars at higher ones.  After a description of the Nan\\c cay instrumentation, I will present few results in the context of this search for a gravitational waves background.  \\begin{figure}[t] \\begin{center} \\epsfig{figure=1909tpl.ps,angle=270,width=2.8in} \\epsfig{figure=1909-20090623.ps,angle=270,width=2.8in} \\caption{(left) Integrated profile obtained for the millisecond pulsar PSR J1909-3744 at Nan\\c cay used as a template to determine Times Of Arrival. (right) A typical daily profile for the same pulsar, obtained on June 26, 2009.\\label{fig:1909}} \\end{center} \\end{figure} ",
        "conclusions": ""
    },
    "0911/0911.1424_arXiv.txt": {
        "abstract": "We performed an analysis of the mid-infared properties of the 12\\,$\\mu$m Seyfert sample, a complete unbiased 12\\,$\\mu$m flux limited sample of local Seyfert galaxies selected from the {\\em IRAS} Faint Source Catalog, based on low resolution spectra obtained with the Infrared Spectrograph (IRS) on-board {\\it Spitzer} Space Telescope. A detailed presentation of this analysis is dicussed in Wu et al. (2009). \\\\ We find that on average, the 15-30\\,$\\mu$m slope of the continuum is $<$$\\alpha_{15-30}$$>$=-0.85$\\pm$0.61 for Seyfert 1s and -1.53$\\pm$0.84 for Seyfert 2s, and there is substantial scatter in each type. Moreover, nearly 32\\% of Seyfert 1s, and 9\\% of Seyfert 2s, display a peak in the mid-infrared spectrum at 20\\,$\\mu$m, which is attributed to an additional hot dust component.  The Polycyclic Aromatic Hydrocarbon (PAH) equivalent width decreases with increasing dust temperature, as indicated by the global infrared color of the host galaxies. However, no statistical difference in PAH equivalent width is detected between the two Seyfert types, 1 and 2, of the same bolometric luminosity.  Finally, we propose a new method to estimate the AGN contribution to the integrated 12\\,$\\mu$m galaxy emission, by subtracting the \"star formation\" component in the Seyfert galaxies, making use of the tight correlation between PAH 11.2\\,$\\mu$m luminosity and 12\\,$\\mu$m luminosity for star forming galaxies. ",
        "introduction": "Active galaxies emit radiation from their nucleus which is due to accretion onto a super-massive black hole (SMBH) located at the center. The fraction of the energy emitted from these active nuclei (AGN), compared with the total bolometric emission of the host can range from a few percent in moderated luminosity systems (L$_{\\rm bol} <10^{11}$L$_{\\odot}$), to more than 90\\% in quasars (L$_{\\rm bol} >10^{12}$L$_{\\odot}$) \\citep[see][ and references therein]{Ho08}. As a subclass, Seyfert galaxies are the nearest and brightest AGNs, with 2-10keV X-ray luminosities less than $\\sim10^{44}$ergs$^{-1}$ and their observed spectral line emission originates principally from highly ionized gas. Seyferts have been studied at many wavelengths and classified as Seyfert 1s (Sy 1s) and Seyfert 2s (Sy 2s), with the type 1s displaying features of both broad (FWHM$>$2000 km s$^{-1}$) and narrow emission lines, while the type 2s only narrow-line emission.   The differences between the two Seyfert types have been an intense field of study for many years. Are they due to intrinsic differences in their physical properties, or are they simply a result of dust obscuration that hides the broad-line region in Sy 2s? A so-called unified model has been proposed \\citep[see][]{Antonucci93, Urry95}, suggesting that Sy 1s and Sy 2s are essentially the same objects viewed at different angles. A dust torus surrounding the central engine blocks the optical light when viewed edge on (Sy 2s) and allows the nucleus to be seen when viewed face on (Sy 1s). Optical spectra in polarized light \\citep{Antonucci85} have indeed demonstrated for several Sy 2s the presence of broad lines, confirming for these objects the validity of the unified model.  However, the exact nature of this orientation-dependent obscuration is not clear yet. Recently, more elaborate models, notably the ones of \\citet{Elitzur08}, \\citet{Nenkova08}, and \\citet{Thompson09} suggest that the same observational constraints can also be explained with discrete dense molecular clouds, without the need of a torus geometry.  Mid-IR spectroscopy is a powerful tool to examine the nature of the emission from AGNs, as well as the nuclear star-formation activity. Since IR observations are much less affected by dust extinction than those at shorter wavelengths, they have been instrumental in the study of obscured emission from optically thick regions in AGNs. This is crucial to understand the physical process of galaxy evolution. With the advent of the {\\em Infrared Space Observatory (ISO)}, local Seyferts have been studied by several groups \\citep[see][ for a review]{Verma05}. The recent launch of the {\\em Spitzer} Space Telescope \\citep{Werner04} has enabled the study of large samples of AGN with substantially better sensitivity and spatial resolution, in an effort to quantify their mid-IR properties \\citep{Buchanan06, Sturm06, Deo07, Gorjian07, Hao07, Tommasin08}.  We have used archival Spitzer/IRS \\citep{Houck04} observation and have embarked in a detailed study Seyferts of the extended 12\\,$\\mu$m galaxy sample.  This sample of 116 Seyfers appears to be the best for studying in an unbiased manner the affects of an AGN to physical properties of a galaxy, since all galaxies emit a nearly constant fraction ($\\sim$7\\%) of their bolometric luminosity at 12\\,$\\mu$m \\citep{Spinoglio89}. A total of 103 Seyferts have been observed by various Spitzer programs using the low-resolution (R$\\sim$64-128) modules of IRS. Among these objects, 47 are optically classified as Sy 1s and 56 as Sy 2s \\citep{Rush93}. The details of this work are summarized here and can be found in \\citet{Wu09}. High-resolution (R$\\sim$600) IRS spectroscopy for 110 objects are also available and their analysis is presented in \\citet{Tommasin08, Tommasin09}.  \\vspace*{-0.2cm} ",
        "conclusions": "Based on the analysis of Spitzer/IRS low resolution spectra for a complete unbiased sample of Seyfert galaxies selected from the {\\em IRAS} Faint Source Catalog based on their 12\\,$\\mu$m fluxes we find that:  1. The 12\\,$\\mu$m Seyferts display a variety of mid-IR spectral shapes. The mid-IR continuum slopes of Sy 1s and Sy 2s are on average $<\\alpha_{15-30}>$=-0.85$\\pm$0.61 and -1.53$\\pm$0.84 respectively, though there is substantial scatter for both types. We identify a group of objects with a local maximum in their mid-IR continuum at $\\sim$20\\,$\\mu$m, which is likely due to the presence of a warm $\\sim$150 K dust component and 18\\,$\\mu$m emission from astronomical silicates. Emission lines, such as the [NeV]\\,14.3\\,$\\mu$m/24.3\\,$\\mu$m and [OIV]\\,25.9\\,$\\mu$m lines, known to be a signature of an AGN are stronger in the average spectra of Sy 1s than those of Sy 2s.  2. PAH emission is detected in both Sy 1s and Sy 2s, with no statistical difference in the relative strength of PAHs between the two types. This suggests that the volume responsible for the bulk of their emission is likely optically thin at $\\sim$12\\,$\\mu$m.  3. The 11.2\\,$\\mu$m PAH EW of the 12\\,$\\mu$m Seyfert sample correlates well with the IRAS color of the galaxies as indicated by the flux ratio of F$_{25}$/F$_{60}$. PAH emission is more suppressed in warmer galaxies, in which the strong AGN activity may destroy the PAH molecules.  4. The FIR luminosities of the 12\\,$\\mu$m Seyferts are dominated by star-formation. Their mid-IR luminosity increases by the additional AGN contribution. A method to estimate the AGN contribution to the 12\\,$\\mu$m luminosity, in a statistical sense, has been proposed and applied to the sample.   \\vspace*{-0.2cm}"
    },
    "0911/0911.0422_arXiv.txt": {
        "abstract": "\\PRE{\\vspace*{.3in}} Dark matter with Sommerfeld-enhanced annihilation has been proposed to explain observed cosmic ray positron excesses in the 10 GeV to TeV energy range. We show that the required enhancement implies thermal relic densities that are too small to be all of dark matter. We also show that the dark matter is sufficiently self-interacting that observations of elliptical galactic dark matter halos exclude large Sommerfeld enhancement for light force carriers. Resonant Sommerfeld enhancement does not modify these conclusions, and the astrophysical boosts required to resolve these discrepancies are disfavored, especially when significant self-interactions suppress halo substructure. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0911/0911.2427_arXiv.txt": {
        "abstract": "Constructing dynamical maps from the filtered output of numerical integrations, we analyze the structure of the $\\nu_\\odot$ secular resonance for fictitious irregular satellites in retrograde orbits. This commensurability is associated to the secular angle $\\theta = \\varpi - \\varpi_\\odot$, where $\\varpi$ is the longitude of pericenter of the satellite and $\\varpi_\\odot$ corresponds to the (fixed) planetocentric orbit of the Sun. Our study is performed in the restricted three-body problem, where the satellites are considered as massless particles around a massive planet and perturbed by the Sun.  Depending on the initial conditions, the resonance presents a diversity of possible resonant modes, including librations of $\\theta$ around zero (as found for Sinope and Pasiphae) or $180$ degrees, as well as asymmetric librations (e.g. Narvi). Symmetric modes are present in all giant planets, although each regime appears restricted to certain values of the satellite inclination. Asymmetric solutions, on the other hand, seem absent around Neptune due to its almost circular heliocentric orbit.  Simulating the effects of a smooth orbital migration on the satellite, we find that the resonance lock is preserved as long as the induced change in semimajor axis is much slower compared to the period of the resonant angle (adiabatic limit). However, the librational mode may vary during the process, switching between symmetric and asymmetric oscillations.  Finally, we present a simple scaling transformation that allows to estimate the resonant structure around any giant planet from the results calculated around a single primary mass. ",
        "introduction": "Both Jupiter and Saturn are rich in irregular satellites. Jupiter has at least 54 moons in this group, while Saturn contains more than 30 members. A smaller population is also observed around Uranus and Neptune, although it is not clear whether this is intrinsic or due to observational bias.  Irregular satellites suffer strong perturbations from the Sun which significantly affect their orbits evolution. Whipple \\& Shelus (1993) numerically integrated the orbit of the Jovian moon Pasiphae for $10^5$ years under the effects of all planets, and found that its longitude of pericenter $\\varpi$ is locked in a secular resonance where $\\varpi - \\varpi_{\\rm Jup}$ librates around $180^\\circ$. Here $\\varpi_{\\rm Jup}$ is the longitude of pericenter of the heliocentric orbit of Jupiter. Since $\\varpi_{\\rm Jup} = \\varpi_\\odot + \\pi$, where $\\varpi_\\odot$ is the Jupiter-centric longitude of the pericenter of the Sun, this implies that $\\varpi - \\varpi_{\\odot}$ oscillates around zero. By analogy from asteroidal motion, we will call this the $\\nu_5$ secular resonance for satellite orbits.  In the same year, Saha \\& Tremaine (1993) found that the Jovian moon Sinope also displays a similar resonance lock in the $\\nu_5$ resonance, altough in this case the motion alternates between libration and circulation on timescales of the order of $10^5-10^6$ years. Extending the orbital evolution for $5 \\times 10^7$ years, Nesvorn\\'y et al. (2003) showed that Pasiphae's libration is also temporary, and will switch to a circulation of the resonant angle in approximately $3 \\times 10^7$ years.  \\'Cuk \\& Burns (2004) performed a detailed search for other irregular moons in secular resonances, finding that the prograde Saturn moon Siarnaq appears to the trapped in a pericenter secular resonance with Saturn (i.e. $\\nu_6$), also displaying intermittent libration, where the resonant angle oscillates around $180^\\circ$. While the Uranus irregular satellite Stephano is not in libration, the angle $\\varpi - \\varpi_{\\odot}$ shows a quasi-resonant behavior with a very long-period circulation.  Finally, Beaug\\'e \\& Nesvorn\\'y (2007) found that although the retrograde Saturn moon Narvi is not currently in resonant motion, long-term simulations show a future libration in the $\\nu_6$ secular resonance. Contrary to previous cases, here the resonant angle displays libration around values different from zero or $180^\\circ$, in what appears to be an asymmetric libration point. Since all these resonant configurations, independently of the planetary mass, involve the relative behavior of the longitude of pericenter of the satellite and that of the planetocentric orbit of the Sun, in all future references we will denote this commensurability as $\\nu_\\odot$.  At present it is not clear whether the observed resonant population in the outer planets is evidence of a past smooth orbital migration of the satellites, or whether it is simply due to chance. Beaug\\'e \\& Nesvorn\\'y (2007) found that stability criteria alone yield satellite populations close to $\\nu_\\odot$, although the proximity of some individual moons (particularly Pasiphae and Sinope) appear sufficiently detached from a random distribution to be statistically significant.  One of the problems in understanding a possible relationship between the $\\nu_\\odot$ resonance and the past evolution of the irregular satellites, is the lack of a detailed analysis of the resonant structure. As will be shown in Section 2, analytical models, even of high-order, are not sufficiently precise and are unable to reproduce the main characteristics of the secular commensurability. Semi-analytical models, where the secular perturbations are evaluated by numerical averaging of the exact Hamiltonian, also suffer the same limitations since they are equivalent to a first-order theory. For this reason, so far most dynamical studies have been restricted to numerical integrations of individual orbits.  In this paper we present a numerical study of the structure of the $\\nu_\\odot$ resonance for retrograde orbits. The resonant behavior is obtained applying two filters on the numerical output: the first is constructed to eliminate the short-period variations (frequencies comparable with the mean motions). The resulting data is then filtered once again to reduce the effects caused by the non-resonant secular angular variable. The results are displayed as dynamical maps showing the location, type and extension of the librational domains for several values of the integrals of motion.  The details of the numerical method are outlined in Section 3, which includes an application to the region of the phase space in the vicinity of the Jovian moon Sinope. Section 4 shows the structure of the $\\nu_\\odot$ resonance in the Saturn system. In both cases the maps are constructed for a single value of the proper semimajor axis. Section 5 discussed the effects of an ad-hoc migration acting of the satellites orbits, and its effects on the resonant motion. Next, in Section 6 we present a simple scaling law that allows us to relate the resonant structure for any planetary mass. Finally, the application the real satellites are briefly discussed in Section 7, while concluding remarks close the paper in Section 8. ",
        "conclusions": "We have presented a numerical study of the structure of the $\\nu_\\odot$ secular resonance for retrograde satellites around the giant planets. This commensurability is associated to a libration of the secular angle $\\theta=\\varpi-\\varpi_\\odot$, where $\\varpi = \\Omega-\\omega$ is the longitude of pericenter of the irregular moon and $\\varpi_\\odot$ the corresponding angle for the planetocentric orbit of the Sun. We have shown that the resonant domain presents a diversity of libration modes, including both symmetric and asymmetric librations. Each mode appears restricted to certain intervals of the mean-mean inclination $i^*$ of the satellite. Librations around $\\theta=0$ occur for large values of $i^*$ (typically $i^* \\gtrsim 150^\\circ$), while oscillations around $\\theta=180^\\circ$ seem to dominate for $i^* \\lesssim 140^\\circ$. Asymmetric solutions appear at intermediate inclinations, although planets with very small eccentricities (such as Neptune) may not contain asymmetric solutions at all. Calculating the mean-mean inclination for real irregular satellites shows a good agreement with this predictions. Pasiphae ($i^*=148^\\circ$), Sinope ($i^*=157^\\circ$) both display librations around $\\theta=0$, while the resonant angle of Narvi ($i^*=141^\\circ$) is in an asymmetric mode.  Simulating a smooth orbital migration of satellite orbits have shown that the resonant lock is maintained as long as the adiabatic condition is met; in other words, the characteristic timescale of the orbital variation is much longer than the libration period. We have used these results to map the location and type of libration centers as a function of the proper semimajor axis, thus giving a more global view of the region of the phase space dominated by the $\\nu_\\odot$ commensurability.  Finally, we presented a simple scaling law that can be used to relate the resonant structure in any the satellite system of planetary mass. Since Jupiter, Saturn and Uranus have similar orbital eccentricities, it is expected that the structure of the secular resonance should also be very similar, except for a scaling in the semimajor axis of the irregular moons. Neptune, however, has an almost circular orbit, which inhibits the formation of asymmetric libration points. For satellites around this planet, it is expected that symmetric solutions should dominate.  From the scaling law, it can be shown that a change in the semimajor axis of the satellite is equivalent to a similar one on the planet. In other words, the same results would obtained if the heliocentric semimajor axis of the planet suffered a slow migration while the planetocentric orbit of the moon was kept fixed. Although we have only done a few simulations with a fast (non-adiabatic) migration, it appears that the resonant configuration is not so easily preserved, and any originally librating body was usually expelled from the resonance. However, additional simulations are required to confirm is this behavior is general.  In this paper we have considered only the restricted three-body problem and, consequently, neglected the gravitational effects of the other giant planets. This is indeed an approximation, since these additional perturbations may cause significant orbital variations (e.g. Carruba et al. 2004, \\'Cuk \\& Burns 2004). In particular, the temporary nature of the $\\nu_\\odot$ resonance lock of real satellites is due to the gravitational effects of the other planetary bodies. In consequence, the results presented here must be viewed only as a starting ground for the complete secular behavior for retrograde irregular moons. Surely the full picture will be far more complex."
    },
    "0911/0911.5341_arXiv.txt": {
        "abstract": "We present a set of ultra-large particle-mesh simulations of the \\lya forest targeted at understanding the imprint of baryon acoustic oscillations (BAO) in the inter-galactic medium. We use 9 dark matter only simulations which can, for the first time, simultaneously resolve the Jeans scale of the intergalactic gas while covering the large volumes required to adequately sample the acoustic feature.  Mock absorption spectra are generated using the fluctuating Gunn-Peterson approximation which have approximately correct flux probability density functions (PDFs) and small-scale power spectra. On larger scales there is clear evidence in the redshift space correlation function for an acoustic feature, which matches a linear theory template with constant bias. These spectra, which we make publicly available, can be used to test pipelines, plan future experiments and model various physical effects. As an illustration we discuss the basic properties of the acoustic signal in the forest, the scaling of errors with noise and source number density, modified statistics to treat mean flux evolution and mis-estimation, and non-gravitational sources such as fluctuations in the photo-ionizing background and temperature fluctuations due to HeII reionization. ",
        "introduction": "Oscillations of the baryon-photon plasma in the early universe, also known as Baryon Acoustic Oscillations (BAO), imprint a distinct signature on the clustering of matter \\citep{PeeYu70,SunZel70} which provides a ``standard ruler'' by which we can measure the expansion history of the Universe (see \\citealt{EisHu98,MeiWhiPea99} for a detailed description of the physics in modern cosmologies and \\citealt{ESW07} for a comparison of Fourier and configuration space pictures).  These oscillations have been traditionally measured in the Cosmic Microwave Background \\citep[CMB; see][for the latest results]{WMAP10} but with the advent of new, large-volume galaxy redshift surveys BAO have been detected in galaxy clustering at low $z$ as well \\citep{Eis05,Cole05,Hut06,Bla07,Pad07,Per07,Per09}.  In principle, the BAO technique becomes even more powerful as one moves to higher redshift, where there is much more volume available to be surveyed and where the acoustic scale is more deeply in the linear regime. This latter fact helps in two ways. First, the power spectrum or correlation function can be computed quite accurately with only linear perturbation theory once one specifies the baryon-to-photon ratio and matter-radiation ratio.  These are both measured accurately from CMB acoustic peaks (see \\citealt{WhiCoh02} for a review) so we have a template with which to fit the data. Second, the acoustic oscillations are less damped by non-linear evolution providing more modes which contain measurable signal.  Tracing the enormous volumes required with galaxies requires a heavy investment in telescope time. However, in principle any tracer of the mass field will do, including the neutral hydrogen in the inter-galactic medium \\citep{Whi03} or galaxies \\citep{Cha08}.  Tracing neutral hydrogen in galaxies via its redshifted $21\\,$cm emission is a key goal for proposed future radio telescopes.  However, even with current technology it is relatively straightforward to obtain a low resolution spectrum of distant quasars and study the \\lya forest of absorption lines which map the neutral hydrogen along the line-of-sight. At $z\\simeq 2-3$ the gas making up the inter-galactic medium (IGM) is thought to be in photo-ionization equilibrium, which results in a tight density-temperature relation for the absorbing material with the neutral hydrogen density proportional to a power of the baryon density \\citep{HuiGne97,Mei09}. Since pressure forces are sub-dominant, the neutral hydrogen density closely traces the total matter density on large scales. The structure in QSO absorption thus traces, in a calculable way, slight fluctuations in the matter density of the universe back along the line-of-sight to the QSO, with most of the \\lya forest arising from over-densities of a few times the mean density.  Motivated by the upcoming Baryon Oscillation Spectroscopic Survey \\citep[BOSS;][a part of SDSS-III]{BOSS}, which will deliver an unprecedented number of quasar spectra probing the \\lya forest at $z\\sim 2-3$, and access to the world's fastest supercomputer, ``Roadrunner'', we have produced a set of ultra-large particle-mesh simulations of cosmological structure formation to further investigate BAO-\\lya science. In this paper we describe our simulations, which simultaneously resolve the Jeans scale and the acoustic scale, and the construction of mock quasar spectra with properties close to those observed at $z\\sim 2-3$ (Sec.~\\ref{sec:sim}). We present several examples to illustrate the use of these spectra in understanding BAO-\\lya science, describing the 2-point correlation function of the Lyman-$\\alpha$ forest flux in Sec.~\\ref{sec:twopt}, and the impact of two non-gravitational signals (fluctuations in the ionizing background and HeII reionization) in Sec.~\\ref{sec:nongrav}. We conclude in Sec.~\\ref{sec:conclude}. In the hope that these spectra can be more generally useful, we have made them publicly available\\footnote{{\\tt http://mwhite.berkeley.edu}}. ",
        "conclusions": "\\label{sec:conclude}  Baryon acoustic oscillations have become one of the premier methods for determining cosmological distances and hence the expansion history of the Universe.  The structure in the spectrum of distant quasars, which is thought to trace the structure of the IGM at near mean density, has been suggested as a relatively cheap method for measuring BAO at high redshift. Motivated by the upcoming BOSS experiment, and the availability of the world's fastest supercomputer, ``Roadrunner'', we have performed a set of ultra-large particle-mesh simulations in order to study the BAO signature in the inter-galactic medium.  {}From the density and velocity fields we have used the ``fluctuating Gunn-Peterson approximation'' (FGPA) to produce mock \\lya spectra. These simulations are the first to be able to cover the large volume necessary to constrain the acoustic scale while simultaneously resolving the Jeans scale of the gas, and the small-scale power and pixel distribution in our mock spectra approximately match those seen in observations at $z\\simeq 2.5$.  We anticipate that these mock spectra will be very useful in testing pipelines, calibrating analysis tools and planning future projects so we have made them publicly available\\footnote{{\\tt http://mwhite.berkeley.edu}}.  In the simulations the \\lya flux follows the mass fluctuations on large scales, with negligible scale-dependent bias \\citep{SHWL09}. This is in some sense to be expected.  We have assumed the FGPA in producing our mock spectra and within this approximation the flux is a highly non-trivial but deterministic function of the underlying density field.  Since the small-scale statistics in our simulations approximately match observations, we are able to make a reasonable estimate of the error in the flux auto-correlation function.  We discuss the scaling of the error with the number of quasar sightlines and the noise in each spectrum. In the limit of low noise and many sightlines the error on the correlation function is well approximated by the Gaussian expression.  We have presented preliminary investigations of an evolving mean flux, fluctuations in the photo-ionization rate and HeII reionization which generate ``extra'' power on the acoustic scale and reduce the contrast of the acoustic peak. Gravitational instability produces a well-defined pattern of higher order correlations which is not obeyed by non-gravitational contributions such as the above, allowing (in principle) a diagnostic of non-gravitational physics in the forest. As an example, we demonstrate that the 3-point cross-correlation function in models with HeII reionization has a different scale dependence than the 3-point function in gravity-only simulations, regardless of the equation-of-state assumed in the latter.  \\vspace*{1cm}"
    },
    "0911/0911.4798_arXiv.txt": {
        "abstract": "The most massive globular cluster in the Milky Way, $\\omega$ Centauri, is thought to be the remaining core of a disrupted dwarf galaxy\\cite{lee99,bf03}, as expected within the model of hierarchical merging\\cite{freeman,diemand}. It contains several stellar populations having different heavy elemental abundances supplied by supernovae\\cite{johnson} --- a process known as metal enrichment. Although M22 appears to be similar to $\\omega$  Cen\\cite{marino}, other peculiar globular clusters do not\\cite{carretta,georgiev}. Therefore $\\omega$  Cen and M22 are viewed as exceptional, and the presence of chemical inhomogeneities in other clusters is seen as `pollution' from the intermediate-mass asymptotic-giant-branch stars expected in normal globular clusters\\cite{ventura}. Here we report Ca abundances for seven globular clusters and compare them to $\\omega$ Cen. Calcium and other heavy elements can only be supplied through numerous supernovae explosions of massive stars in these stellar systems\\cite{timmes}, but the gravitational potentials of the present-day clusters cannot preserve most of the ejecta from such explosions\\cite{baum}. We conclude that these globular clusters, like $\\omega$  Cen, are most probably the relics of more massive primeval dwarf galaxies that merged and disrupted to form the proto-Galaxy. ",
        "introduction": "\\subsection{The Filter Transmission} The Ca filter system was designed to include Ca II H and K lines at $\\lambda$ 3968 and 3933 \\AA, respectively, with a full-width half maximum (FWHM) of approximately 90 \\AA. The CN band absorption strengths at $\\approx$ $\\lambda$ 3885 \\AA~ are often very strong  in stellar spectra and the lower limit of the Ca filter is set to avoid  contamination by the CN band$^{12}$. The Ca filter used at Cerro Tololo Inter-American Observatory (CTIO) has a similar FWHM, approximately 90 \\AA, but its passband is shifted approximately 15 \\AA~ to the longer wavelength (Alistair Walker, private communication) compared to that in Anthony-Twarog \\emph{et al.}$^{12}$ In Supplementary Figure \\ref{sifig:filter}, we show the transmission functions of the both Ca filters. In the figure, we also show synthetic spectra for the CN normal and the CN strong RGB stars in typical intermediate metallicity globular clusters (GCs) as an illustration. The effect of the CN band on the $hk$ index is negligible as will be discussed below.  \\subsection{The Central Wavelength Drift of the Ca Filter} The CTIO Ca filter transmission function shown in Supplementary Figure \\ref{sifig:filter} is that measured with a collimated beam. It is known that the passband of the narrow band interference filter depends on the angle of the incidence beam following, \\begin{equation} \\lambda = \\lambda_0\\left( 1-\\frac{\\sin^{2}\\beta}{n^{*2}}\\right) ^{1/2}, \\end{equation} where $\\lambda_0$ is the wavelength of peak transmittance at normal incidence, $\\beta$ is the angle of incidence of the collimated beam on the filter and $n^*$ is the effective refractive index of the filter\\cite{clarke}. Therefore, when the Ca filter is used with a fast telescope, the filter passband can be significantly different from that shown in Supplementary Figure \\ref{sifig:filter}. The CTIO 1-m telescope used for our survey is a slow telescope with $f$/10.5 and the effect resulted from the angular dependency of a converging beam is expected to be very small. Assuming $\\beta$ $\\approx$ 1/21 radians for the converging beam at the CTIO 1-m telescope and $n^*$ $\\approx$ 1.4 for the CTIO Ca filter, the peak wavelength of the Ca filter will be shifted by 2.3 \\AA~ to the shorter wavelength. Given the much larger FWHM of the CTIO Ca filter, this may contribute small effect. We investigate contributions to the $hk$ index resulted from the shifted Ca passband using synthetic spectra for typical intermediate metallicity RGB stars in our GCs. Our calculations integrating over the filter transmission curve show that this effect contributes no larger than 0.011 mag to our $hk$ measurements, in the sense that the shifted Ca passband to the shorter wavelength gives slightly larger $hk$ values. We emphasize that, since our results are based on a single instrument setup (the same telescope, filters and the CCD camera) during the observations of our science targets and the photometric standards, this effect is expected to be cancelled out during our photometric calibrations. Also importantly, our main results presented here rely on the split or the spread in the $hk$ index of RGB stars of an individual GC. Therefore, the shifted Ca passband affects similar degree to the $hk$ index of the RGB stars in a GC and does not contribute to the apparent RGB split or the spread in the $hk$ index of an individual GC.  \\subsection{Effect of radial motions and internal velocity dispersions of GCs} The mean radial motion of 139 GCs in our Milky Way Galaxy\\cite{harris} is $|v_r|$ = 110 km/s, equivalent to the wavelength shift by $|\\Delta\\lambda|$ $\\approx$ 1.4 \\AA~ at $\\lambda$ 3950\\AA. We calculate the contribution due to the mean radial motion of GCs to the $hk$ index using the shifted CTIO Ca passband and synthetic spectra for typical intermediate metallicity GC RGB stars. We find that the net effect is negligibly small, $|\\Delta hk|$ $<$ 0.006 mag. Among our eight GCs, NGC1851 has the largest radial velocity, $v_r$ = 321 km/s, equivalent to the wavelength shift by $\\Delta \\lambda$ = 4.2 \\AA~ to the longer wavelength at $\\lambda$ 3950\\AA. We calculate the contribution due to the radial motion of NGC1851 using the shifted CTIO Ca passband and the red-shifted synthetic spectra with a fixed CN abundance for the cluster. We obtain $\\Delta hk$ $\\approx$ 0.015 mag, in the sense that the red-shift is resulted in a slightly larger $hk$ index. As we discussed above, the difference in the $hk$ index due to the high radial motion of NGC1851 does not affect our results presented here, since the $hk$ indices of RGB stars in NGC1851 will be affected by similar degree and the mean radial motion of the cluster does not produce an apparent split or a spread in the $hk$ index. Perhaps, this effect may become important in the inter-cluster comparisons, which is beyond the scope of our study.  What concerns us most about the high radial velocities of some GCs, in particular for red-shift, is the potential contamination by the strong CN band absorption features at $\\lambda$ 3885 \\AA~ as shown in Supplementary Figure \\ref{sifig:filter}. For example, NGC1851 has a bimodal CN distribution and some RGB stars show very strong CN band absorption strengths\\cite{hesser}. Due to its high radial velocity away from us (i.e. red-shifted), the CN band absorption features in the CN-strong RGB stars may affect the $hk$ index and, subsequently, may produce an apparent RGB split of the cluster as shown in Figure 2 or Supplementary Figure \\ref{sifig:n1851dist}. We calculate the CN band contributions using the shifted CTIO Ca passband and the red-shifted synthetic spectra for the CN-normal and the CN-strong RGB stars (see discussion below). Our calculations suggest that the net effect is negligibly small, $\\Delta hk$ $\\leq$ 0.003 mag, and the high radial velocity of NGC1851 combined with a bimodal CN distributions does not produce the RGB split in the $hk$ index.  We also investigate the effect of the internal velocity dispersion of an individual GC. Assuming $\\sigma_{LOS}$  = 15 km/s, equivalent to $\\Delta\\lambda$ $\\leq$ 0.2 \\AA~ at $\\lambda$ 3950\\AA, we obtain $\\Delta hk$ $\\leq$ 0.001 mag following the same method described above, and the effect from the internal velocity dispersions of GCs does not affect our results.  \\subsection{Summary of uncertainties on the $hk$ index} Supplementary Table \\ref{tab:exterr} summarizes the uncertainties in our $hk$ index measurements relevant to the CTIO Ca passband. (The variations in the $hk$ index due to differences in elemental abundances of GC RGB stars will be discussed below.) As discussed above, the effects due to the shifted CTIO Ca passband and the radial motions of GCs do not affect our results, since both effects contribute similar degree to the $hk$ indices among RGB stars in a GC. (i.e. They only affect the zero point of the $hk$ index and they do not affect the $\\Delta hk$ distributions). In addition to our photometric measurement errors which will be discussed below, the effects due to the internal velocity dispersions of GCs and the differential interstellar Ca II absorption (see discussion below) can affect our $hk$ index measurements. However, their contributions to our $hk$ measurements are no larger than 0.022 mag and they do not affect our main conclusion presented here. Therefore, our results strongly suggest that the split or the spread in the $hk$ index of RGB stars in GCs are related to the variations in elemental, in particular calcium, abundances among RGB stars in a GC, which will be discussed below. ",
        "conclusions": ""
    },
    "0911/0911.3105.txt": {
        "abstract": "*{I will introduce some statistical tools that any cosmologist should know %about in order to be able to understand  recently published results from the %analysis of  cosmological data sets. I will not present a complete and rigorous %introduction to statistics, but I will touch upon useful tools....} \\abstract{ The advent of large data-set in cosmology has meant that in the past 10 or 20  years our knowledge  and understanding of the Universe has changed not only quantitatively but also, and most importantly, qualitatively. Cosmologists are interested in studying the origin and evolution of the physical Universe. They rely on data where a host of useful  information is enclosed, but is encoded in a  non-trivial way. The challenges in extracting this information    must be overcome to  make the most of the  large experimental effort. Even after having analyzed a decade or more of data and having  converged to a standard cosmological model (the so-called and highly successful LCDM model) we should keep in mind that this model is   described by 10 or more physical parameters and if we want to study deviations from the standard model the number of parameters is even larger. Dealing with  such a high dimensional parameter space and finding parameters constraints  is a challenge on itself.  In addition, as gathering data is such an expensive and difficult process,  cosmologists  want to be able to compare and combine different data sets both for testing for possible disagreements (which could indicate new physics)  and for improving  parameter determinations.  Finally, always because experiments are so expansive, cosmologists in many cases want to find out {\\it a priori}, before actually doing the experiment, how much one would be able to learn from it. For all these reasons,  more and more sophisiticated statistical techniques are being employed in cosmology, and it has become crucial  to know some statistical background to understand recent literature in the field. % Here,   I will introduce some statistical tools that any cosmologist should know about in order to be able to understand  recently published results from the analysis of  cosmological data sets. I will not present a complete and rigorous introduction to statistics as there are several good books which are reported in the references. The reader should refer to those. I will take a practical approach and  I will touch upon useful tools such as  statistical inference,  Bayesians vs Frequentist approach, chisquare and goodness of fit, confidence regions, likelihood, Fisher matrix approach,   Monte Carlo methods and  a brief introduction to model testing. Throughout, I will use practical examples often taken from recent literature to illustrate the use of such tools. Of course this will not be an exhaustive guide: it should be interpreted as a ``starting kit\", and the reader is warmly encouraged to read the references to find out more.} ",
        "introduction": "\\label{sec:1}  As cosmology has made the transition from a data-starved science to a data-driven science, the use of increasingly sophisticated statistical tools  has increased. As  explained in detail below, cosmology is intrinsically related to statistics, as theories of the origin and evolution of the Universe  do not predict, for example, that a particular galaxy will form at a specific point in space and time or that  a specific patch of the cosmic microwave background will have a given temperature: any theory will predict average statistical properties of our Universe, and we can only observe a particular realization of that.  It is often said that cosmology has entered the precision era:  ``precision\"  requires a good knowledge of  the error-bars  and thus confidence intervals  of a measurement. This is an inherently statistical statement. We should try however to go even further, and achieve also ``accuracy\" (although cosmology    does not have a particularly stellar track record in this regard). This requires quantifying systematic errors (beyond the statistical ones) and it is also requires statistical tools. For all these reasons,  knowledge of basic statistical tools has become indispensable to understand the recent cosmological literature.  Examples of applications where  probability and statistics are crucial in Cosmology  are:  {\\it i)} Is the universe homogenous and isotropic on large scales? {\\it ii)} are the initial conditions consistent with being Gaussian? {\\it iii)} is there a detection of non-zero tensor modes? {\\it iv)} what is the value of the density parameter of the Universe $\\Omega_m$ given the WMAP data for a LCDM model? {\\it v)} what are the allowed values at a given confidence level for the primordial power spectrum spectral slope $n$?  {\\it vi)}  what is the best fit  value of the dark energy equation of state parameter  $w$?  {\\it vii)} Is a model with equation of state parameter different from -1 a better fit to the data than a model with non-zero curvature?  {\\it viii)}  what will be the constraint on the parameter $w$ for a survey with given characteristic?  The first three questions address the hypothesis testing issue. You have an hypothesis and you want to check wether the data are consistent with it. Sometimes, especially for addressing issues of ``detection\"  you can  test the null hypothesis: assume the quantity is zero and test wether the data are consistent with it.  The next three questions  are ``parameter estimation\" problems: we have a model, in this example the LCDM model,  which is characterized by some free parameters  which we would like to measure.  The next question, {\\it vii)}, belong to ``model testing\": we have two models and ask which one is a better fit to the data. Model testing comes in several different flavors: the two models to be considered may have different number of parameters or equal number of parameters, may be have some parameters in common or not etc.  Finally question {\\it viii)} is on ``forecasting\", which is particularly  useful for or quickly forecasting the performance of future experiments and for  experimental design.    Here we will mostly concentrate in the issue of parameter estimation but also touch upon the other applications. ",
        "conclusions": "\\label{sec:conlcusions} I have given a brief overview of statistical techniques that are frequently used in the cosmological literature. I have presented several examples often from the literature to put these techniques into context. This is not an exhaustive list nor a rigorous treatment, but a starter kit  to ``get you started\". As more and more sophisticated  statistical techniques are used to make the most of  the data, one should always remember that they need to be implemented and used correctly: \\begin{itemize} \\item data gathering is an expensive and hard task: statistical techniques make possible  to make the most of the data \\item always beware of systematic effects \\item  an incorrect treatment of the data will give non-sensical results \\item   there will always be  things that are beyond the statistical power of a given  data set \\end{itemize} Remember: ``Treat your data with respect!\""
    },
    "0911/0911.0943_arXiv.txt": {
        "abstract": "Pulsars are amongst the most stable rotators known in the Universe.  Over many years some millisecond pulsars rival the stability of atomic clocks.  Comparing observations of many such stable pulsars may allow the first direct detection of gravitational waves, improve the Solar System planetary ephemeris and provide a means to study irregularities in terrestrial time scales.  Here we review the goals and status of current and future pulsar timing array projects. ",
        "introduction": "The pulsar timing technique (e.g. Lorimer \\& Kramer 2005 for an overview and Edwards, Hobbs \\& Manchester 2006 for a detailed description) has enabled some of the most exciting recent results in astrophysics.  For instance, the first extra-Solar planets were discovered orbiting the pulsar PSR~B1257$+$12 (Wolszczan \\& Frail 1992), the first evidence for gravitational waves was provided by the binary system PSR~B1913$+$16 (Hulse \\& Taylor 1974) and the most stringent tests of general relativity in the strong-field regime have been carried out using the double pulsar system (Kramer et al. 2006).   Even though the PSR~B1913$+$16 system has provided evidence for the existence of gravitational waves, most researchers regard this as indirect evidence and not a direct detection of such waves.  Sazhin (1978) and Detweiler (1979) showed that low-frequency gravitational waves will induce pulsar timing residuals\\footnote{Pulse times of arrival for a given pulsar are compared with a physical model of the pulsar.  The differences between the measured arrival times and predicted arrival times using the model are known as the pulsar timing residuals and denote physical effects that are not included in the model.} that could be detectable with high-precision observations of the pulse arrival times.  However, with a single pulsar it is difficult, and perhaps impossible, to distinguish between the effects of gravitational waves and many other effects such as the irregular rotation of the neutron star, interstellar medium propagation effects or inaccuracies in the planetary ephemeris.  Hellings \\& Downs (1983) showed that such effects could be distinguished by looking for correlated behaviour in the timing residuals of many pulsars.  In brief, the timing residuals caused by irregular rotation or propagation effects should be uncorrelated between different pulsars.  Irregularities in the terrestrial atomic time scale would produce completely correlated timing residuals.  The angular correlation function expected for an isotropic, stochastic gravitational wave background (Hellings \\& Down 1983) is shown in Figure~\\ref{fg:corr} (using simulated data described in Hobbs et al. 2009).  The first attempt to create a ``pulsar timing array'' in which enough pulsars are observed with sufficient timing precision to search for correlated signals was reported by Foster \\& Backer (1990).  More recently Jenet et al. (2005) calculated the number of pulsars, the timing precision and data span required in order to make a significant detection of the expected stochastic background of gravitational waves.   In summary, at least 20 millisecond pulsars are required, observed for five years with timing precision around 100\\,ns. In this review paper we discuss the current status of various existing pulsar timing array projects that aim to achieve this goal.  In the future, telescopes such as the Square Kilometre Array may make gravitational wave detection using pulsars commonplace.  The current status of this project is described in \\S\\ref{sec:future}.  \\begin{figure*} \\begin{center} \\includegraphics[width=5cm,angle=-90]{hdCurve3.ps} \\end{center} \\caption{The expected correlation in the timing residuals of pairs of pulsars as a function of angular separation for an isotropic gravitational wave background.  The solid line gives the theoretical curve.  The dots are calculated by correlating simulated data sets for 20 pulsars in the presence of a gravitational wave background that produces white timing residuals.}\\label{fg:corr} \\end{figure*} ",
        "conclusions": "Many pulsars are currently being observed as part of pulsar timing array projects.  Correlated timing residuals allow 1) irregularities in terrestrial time standards to be identified, 2) inaccurate planetary masses in the Solar System ephemeris to be corrected and 3) the detection of gravitational wave signals.  Combining observations from Northern and Southern hemisphere telescopes should allow the detection of gravitational waves on a time scale of 5-10 years."
    },
    "0911/0911.3546_arXiv.txt": {
        "abstract": "We present a study, based on  archival \\xmmn observations, of the extended X-ray emission associated  with the inner  disk of M33.  After the exclusion  of point sources with \\lx  $> 2 \\times 10^{35} \\ergsec$ (0.3-6  keV), we investigate both the  morphology and  spectrum of  the  residual X-ray  emission, comprising  the integrated signal from unresolved  discrete sources and diffuse components. This residual emission  has a  soft X-ray  spectrum which  can be fitted  with a two-temperature thermal model,  with $\\rm kT \\approx 0.2$ keV  and $\\approx 0.6$ keV, the  cooler  component accounting  for  the bulk  of  the  luminosity. There  is evidence that the X-ray emitting plasma has a subsolar metal abundance. The soft  X-ray surface brightness distribution shows a strong  correlation with FUV emission and since the latter serves as a tracer of the inner spiral arms of M33, this is indicative of a close connection  between recent star-formation activity and the production of soft X-rays. Within 3.5 kpc  of the nucleus of M33, the soft X-ray and   FUV  surface   brightness  distributions   exhibit  similar  radial profiles.  The  implication is  that  the ratio  of  the  soft X-ray  luminosity (0.3-2.0 keV) per unit disk area to  the star formation rate (SFR) per unit disk area remains  fairly constant within this  inner disk region. We  derive a value for   this  ratio   of  $1-1.5   \\times  10^{39}   \\ergsec(\\Msun  {\\rm yr}^{-1})^{-1}$, which is towards the top  of the range of similar estimates for several other nearby face-on spiral galaxies (\\eg M51, M83).   In the same region, the ratio of soft  X-ray  luminosity  to  stellar  mass  (the  latter  derived  from  K-band photometry) is  $4 \\times 10^{28} \\ergsec  \\Msun^{-1}$, a factor  of 5-10 higher than is  typical of  dwarf elliptical galaxies (\\eg M32, NGC3379),  suggesting that 10-20\\%  of the unresolved emission seen in M33 may originate in its old stellar population. The remainder of the soft X-ray emission is found to be equally split between two spatial components, one which closely traces the spiral arms of the galaxy and the other more smoothly distributed across the inner disk of M33. The former must represent a highly clumped low-filling factor component linked to sites of recent or ongoing star formation, encompassing \\hii regions, X-ray source complexes and X-ray superbubbles, whereas the distribution of the latter gives few clues as to its exact origin. ",
        "introduction": "\\label{sec:m33intro}  On the basis of current observations,  the X-ray emission of spiral galaxies may be separated into several components.  A very significant contribution generally arises  from a  set  of resolved  point  sources which  correspond  to the  most luminous  examples of the  galactic X-ray  binary population,  encompassing both high-mass  (HMXB)  and low-mass  (LMXB)  systems.  To  this  must  be added  the integrated emission  from large  numbers of lower-luminosity  sources, including supernova  remnants (SNRs), more  quiescent forms  of X-ray  binary, cataclysmic variables, and coronally  active stars.  Finally we need  to include the thermal X-ray  signal  emanating from  concentrations  of  truly  diffuse million-K  gas resident both in  the galactic disk and potentially  extending into the galactic halo.  The extent to which these components can be distinguished from each other depends on the orientation and distance  of the target galaxy and, of course, on the spatial resolution and threshold sensitivity afforded by the observation.  Early X-ray  studies based on  \\einstein observations focussed on  the brightest point  sources  seen  in  a  number  of nearby  galaxies,  whilst  also  finding underlying extended  emission in complex  structures (\\citealt{fabbiano89}). The superior spatial resolution and soft response of {\\it Rosat} further progressed the field, for example allowing the extent of the apparently diffuse emission to be traced  both in  and out of  the plane  of the disk  for face-on  and edge-on systems  respectively (\\eg  \\citealt{read97}; \\citealt{dahlem98}).   More recent investigations, which  exploit the enhanced collecting  area, spatial resolution and  spectral  sensitivity   of  \\xmmn  and  {\\it  Chandra},   have  probed  the point-source populations in galaxies  to much fainter thresholds than previously attainable        (\\eg       \\citealt{strickland04a};       \\citealt{colbert04}). The X-ray luminosity functions (XLF) typical of the HMXB and   LMXB   populations   seen   in   spiral   galaxies   have   thereby   been determined (\\citealt{grimm05}; \\citealt{kilgard05}; \\citealt{fabbiano06}). By extrapolation of the  XLF to fluxes below the threshold at which individual sources  can be resolved, the total  aggregated luminosity of a particular class of source may also be inferred (\\citealt{sazonov06})  In  two previous papers  (\\citealt{warwick07}, hereafter  W07; \\citealt{owen09}, hereafter OW09),  we studied  the morphology of  the underlying  unresolved soft X-ray emission  in a  sample of six  nearby face-on  spiral galaxies and  made a comparison with other wavelength, particularly FUV, measurements.  These studies confirmed the close connection  between recent star-formation and the production of  soft  X-rays.   We  also  investigated  how the  soft  X-ray  luminosity  to star-formation rate  (SFR) ratio varied radially within  the individual galaxies as well as  from galaxy  to  galaxy,  and found  some  evidence  for an  enhanced efficiency in the production of soft  X-rays in regions of high SFR density. The link between soft X-rays and star  formation can be further explored by studying galaxies closer to us,  in which the majority of HMXBs and  LMXBs can be readily resolved, so as to provide a relatively clear perspective of the underlying soft X-ray emission. It is in this context that we focus in the present paper on M33.  M33   is   an  Sc   spiral   galaxy   with   an  inclination   of   $56^{\\circ}$ (\\citealt{zaritsky89}) at  a distance of  795 kpc (\\citealt{vdburgh91}).  As the the  third largest  galaxy in  the Local  Group, the  relative proximity  of M33 permits detailed study  of its discrete X-ray source  population and enables the separation  of  these sources  from  any  residual emission  in  the  disk to  a relatively  faint threshold.  The relatively  low Galactic  foreground  \\nh~ ($7.5 \\times 10^{20}  \\cm2$; \\citealt{kalberla05})  in the direction  of M33, the  low to moderate inclination of the galaxy, and its relatively high SFR in comparison to other nearby systems (0.3-0.7\\Msun yr$^{-1}$,  \\citealt{hippelein03}) make M33 an ideal candidate for the present study.  Early  \\einstein observations  of  M33 revealed  a  total of  17 point  sources, including  the  bright   ULX  M33  X-8  (\\citealt{long81};  \\citealt{markert83}; \\citealt{trinchieri88}).  \\citet{trinchieri88}  also   found  evidence  for  the presence of a  soft diffuse component in the plane of  the galactic disk. \\rosat observations   (\\citealt{schulman95};  \\citealt{long96})   expanded   the  known population of  X-ray point sources  in the  direction of M33  to a total  of 184 (\\citealt{haberl01}). More recently, using \\xmmn  data, a total of 350-400 X-ray point  sources to a  limiting X-ray  luminosity of  $10^{35} \\ergsec$  have been identified    and   categorized    (\\citealt{foschini04};   \\citealt{pietsch04}; \\citealt{grimm05};  \\citealt{misanovic06};  \\citealt{grimm07}).   These  results have   been   complemented   by   recent   studies   of   M33   using   \\chandra (\\citealt{plucinsky08};  \\citealt{williams08}) which  have increased  the source statistics particularly in the confused inner regions of the galaxy. Most of the emphasis  to  date has  been  on  the X-ray  emission  from  point sources,  but \\citet{haberl01},  \\citet{pietsch04}   and  \\citet{tullman08}  have   noted  the presence  of seemingly  diffuse structures  along the  spiral arms.  It  is this latter component which we examine in this paper.  Here, we  focus on the spectral  and spatial properties of  the unresolved X-ray emission   from  the   inner  disk   of   M33,  deduced   from  archival   \\xmmn observations. In \\S\\ref{sec:m33obs}, we describe the set of observations used to construct soft-band X-ray images and  outline the methods used in the subsequent data analysis.  In \\S\\ref{sec:m33lx}, we  briefly examine the properties  of the resolved discrete  point source population  and the spatial distribution  of the residual  soft X-ray  emission, which  remains after  the contribution  of these sources is excluded.  We go on  to compare the soft X-ray morphology with \\galex FUV  measurements and 2MASS K band data. We follow  this with  spectral analysis of the  bright point source population and unresolved    residual       emission (\\S\\ref{sec:m33spectrum}). We examine the  spatial extent of the X-ray emission in  comparison  with star  formation  data and the underlying  mass distribution (\\S\\ref{sec:m33star}) for M33, and compare these relationships to those found in other galaxies. Finally   we  discuss   the  implication   of   our  results (\\S\\ref{sec:m33disc})        and         summarize        our        conclusions (\\S\\ref{sec:m33conc}).   \\begin{table*} \\caption{Details of the \\xmmn observations of M33.} \\small \\centering \\begin{tabular}{lcccccccccc} \\hline Galaxy    & Observation ID & Start Date  & Filter~$^{a}$   & & \\multicolumn{2}{c}{Target co-ordinates}  & & \\multicolumn{2}{c}{Useful exposure (ks)} \\\\ &                & (yyyy-mm-dd) & pn/MOS1/MOS2    & & RA (J2000)       & Dec (J2000)             & & pn            & MOS 1+2 \\\\ \\hline M33  &  0102640101$^{c}$  &  2000-08-04  &  M/-/-   & & $01^h33^m51.0^s$ & $+30\\deg39\\arcm37\\arcs$~$^{b}$ & & 7.1  &  -     \\\\ & 0102640201  &  2000-08-04   &  M/M/M  & &  $01^h34^m40.0^s$ & $+30\\deg57\\arcm48\\arcs$   & & 11.8   & 31.5   \\\\ &  0102640301  &  2000-08-07  &  M/M/Tn  & &  $01^h33^m32.0^s$ & $+30\\deg52\\arcm13\\arcs$   & &  3.6   & 9.5    \\\\ &  0102640401  &  2000-08-02  &  Tk/Tk/Tk  & & $01^h32^m51.0^s$ & $+30\\deg36\\arcm49\\arcs$  & &  9.1  &  23.1  \\\\ &  0102640501  &  2001-07-05  &  M/M/M    & &  $01^h33^m02.0^s$ & $+30\\deg21\\arcm24\\arcs$  & &  9.2  &  23.1  \\\\ &  0102640601  &  2001-07-05  &  M/M/M    & &  $01^h34^m08.0^s$ & $+30\\deg46\\arcm06\\arcs$  & &  4.5  &  11.9  \\\\ &  0102640701  &  2001-07-05  &  M/M/M    & &  $01^h34^m10.0^s$ & $+30\\deg27\\arcm00\\arcs$  & &  6.9  &  21.9  \\\\ &  0102640801  &  2001-07-07  &  -/M/M    & &  $01^h34^m51.0^s$ & $+30\\deg42\\arcm22\\arcs$  & &  -  &  3.2  \\\\ &  0102640901  &  2001-07-08  &  M/M/M    & &  $01^h34^m04.0^s$ & $+30\\deg57\\arcm25\\arcs$  & &  3.9  &  11.2  \\\\ &  0102641001  &  2001-07-08  &  M/M/M    & &  $01^h33^m07.0^s$ & $+30\\deg45\\arcm02\\arcs$  & &  1.5  &  16.3  \\\\ &  0102641101  &  2001-07-08  &  M/M/M    & &  $01^h32^m46.0^s$ & $+30\\deg28\\arcm19\\arcs$  & &  8.0  &  21.0  \\\\ &  0102641201  &  2000-08-02  &  Tk/Tk/Tk  & & $01^h33^m38.0^s$ & $+30\\deg21\\arcm49\\arcs$  & &  12.0  &  7.2  \\\\ &  0102642001  &  2001-08-15  &  M/M/M     & & $01^h34^m51.0^s$ & $+30\\deg42\\arcm22\\arcs$  & &  8.8  &  22.3  \\\\ &  0102642101  &  2002-01-25  &  M/M/M     & & $01^h34^m34.0^s$ & $+30\\deg34\\arcm11\\arcs$  & &  10.0  &  24.3  \\\\ &  0102642201  &  2002-01-25  &  M/M/M     & & $01^h34^m56.0^s$ & $+30\\deg50\\arcm52\\arcs$  & &  11.5  &  27.3  \\\\ &  0102642301$^{c}$  &  2002-01-27  &  M/M/M     & & $01^h33^m33.0^s$ & $+30\\deg33\\arcm07\\arcs$  & &  9.9  &  24.1  \\\\ &  0141980101  &  2003-07-11  &  M/M/M     & & $01^h33^m07.0^s$ & $+30\\deg45\\arcm02\\arcs$  & &  6.2  &  13.2  \\\\ &  0141980301  &  2003-07-25  &  -/M/M     & & $01^h34^m08.0^s$ & $+30\\deg46\\arcm06\\arcs$  & &  -   &  1.2  \\\\ &  0141980501  &  2003-01-22  &  M/M/M     & & $01^h33^m51.0^s$ & $+30\\deg39\\arcm37\\arcs$  & &  1.9  &  17.0  \\\\ &  0141980601  &  2003-01-23  &  M/M/M     & & $01^h32^m51.0^s$ & $+30\\deg36\\arcm49\\arcs$  & &  11.0  &  25.9  \\\\ &  0141980701  &  2003-01-24  &  M/M/M     & & $01^h33^m38.0^s$ & $+30\\deg21\\arcm49\\arcs$  & &  4.4  &  11.4  \\\\ &  0141980801$^{c}$  &  2003-02-12  &  M/M/M     & & $01^h33^m51.0^s$ & $+30\\deg39\\arcm37\\arcs$  & &  7.8  &  19.8  \\\\    \\hline \\end{tabular} \\\\ \\raggedright $^{a}$ - Tn = thin filter, M = medium filter, Tk = thick filter \\\\ $^{b}$ - Assumed position of the galactic nucleus.\\\\ $^{c}$ - Observations used for spectral analysis.\\\\ \\label{table:m33obs} \\end{table*}     \\begin{figure*} \\centering \\rotatebox{270}{\\scalebox{0.45}{\\includegraphics{m33methodxsoft.ps}}} \\rotatebox{270}{\\scalebox{0.45}{\\includegraphics{m33methodxmed.ps}}} \\rotatebox{270}{\\scalebox{0.45}{\\includegraphics{m33methodpsf.ps}}} \\rotatebox{270}{\\scalebox{0.45}{\\includegraphics{m33methodsub.ps}}} \\rotatebox{270}{\\scalebox{0.45}{\\includegraphics{m33finalfuvlin.ps}}} \\rotatebox{270}{\\scalebox{0.45}{\\includegraphics{m33finalklin.ps}}} \\caption{{\\it Top-left panel:}  Adaptively smoothed version of the \\xmmn pn+MOS image of M33 in the soft (0.3-1 keV) band. {\\it Top-right panel:} The same image in the medium (1--2 keV) band. {\\it Middle-left panel:} Simulated image of the bright sources in M33, with the source mask contours overlaid. The bright central region is dominated by the ULX M33 X-8. {\\it Middle-right panel:} The residual emission in M33 in the 0.3-1 keV band obtained by subtracting the bright source model and applying the spatial mask. {\\it Bottom-left  panel:} The {\\it GALEX} FUV  ($\\lambda_{eff}\\approx1528$\\AA) image on the same spatial scale as the X-ray  data. {\\it Bottom-right panel:} 2MASS K band image on the same spatial scale as  the X-ray data. The ellipse shown in each image represents the area of the galaxy used for X-ray analysis (see text for details). All the images are displayed with logarithmic amplitude scaling and are 40\\arcm on a side. } \\label{fig:m33xuv1} \\end{figure*} ",
        "conclusions": "\\label{sec:m33conc}  We have used archival \\xmmn  observations to examine the residual X-ray emission observed from M33 after the exclusion of the bright point source population to a limit of \\lx $~> 2  \\times 10^{35} \\ergsec$.  Using the  same methodology as in  OW09, we have investigated the spectral and spatial  properties of the X-ray emission within an inner disk region extending up to 3.5 kpc from the nucleus of the galaxy.  The observed X-ray spectrum can be modelled as thermal emission with cool and hot components of $\\approx  0.2$ keV  and $\\approx  0.6$ keV respectively, with the cooler emission providing the dominant contribution to the luminosity. There is some evidence for a subsolar metallicity consistent with other indicators of low-metallicity in M33.  The  strong correlation established between  X-ray and  FUV  morphologies confirms the close linkage between X-ray emission and recent star formation. Detailed comparison of  soft X-ray and FUV radial profiles in  the inner disk of M33 reveals  the ratio of X-ray emission  to SFR to be  $1-1.5 \\times 10^{39} \\ergsec$ (\\Msun~yr$^{-1}$)$^{-1}$ (in  the 0.3-2 keV band).  This matches the  predictions of \\citet{mas08} for an extended burst of star  formation occurring 10 Myr ago, with  an efficiency of mechanical energy conversion to X-rays of $\\sim1\\%$.  The soft X-ray emission to mass ratio found for  M33 is $4 \\times 10^{28} \\ergsec\\Msun^{-1}$, a  factor of  5 higher than  the corresponding  value for dwarf elliptical galaxies and spirals with low SFR. This implies that up to $\\sim20\\%$ of the observed soft X-ray emission originates in the old stellar source population, in source types such as cataclysmic variables and active binaries.  With the contribution of the old stellar population subtracted, the soft X-ray emission is found to be equally split between two spatial components, one which closely traces the spiral arms of the galaxy and the other more smoothly distributed across the inner disk region. The constraints on the cooling timescale implied by the presence of spiral features in the soft X-ray images suggest the presence of a highly clumped component, encompassing sites of on-going starformation, \\hii regions and hot gas bubbles. The nature of the smoothly distributed component is much less certain. Plausibly it may represent the integrated emission from a whole range of processes and source types including supernovae occurring in the interarm regions, individual sources and source complexes with luminosity not far below the applied luminosity threshold and an accumulated distribution of hot gas which has managed to leak away from the site of its original production."
    },
    "0911/0911.3400_arXiv.txt": {
        "abstract": "We have measured the 2D 2-point correlation function, \\twopcf, between low column density \\lya\\ absorbers and galaxies at a redshift $\\rm{z}\\sim1$. We measured \\lya\\ absorbers between redshifts ${\\rm z}=0.68\\rightarrow1.51$ over a total redshift path length of $\\Delta{\\rm z}=1.08$ from HST STIS E230M absorption spectra towards the quasars HE 1122-1648 ($\\rm{z}=2.4$) and PKS 1127-145 ($\\rm{z}=1.187$).  The column density of the \\lya\\ absorbers ranged from $13.2\\leq$\\logcoldens$\\leq17.4$, with a median column density of \\logcoldens$=14.0$.  A total of 193 galaxy redshifts within the surrounding 6.8\\arcmin$\\;\\times\\;$5.7\\arcmin\\ field of view of both quasars were identified in a ${\\rm R}$ magnitude limited survey ($21.5\\leq \\rm{R}_{\\rm{Vega}}\\leq24.5$) using the FORS2 spectrograph at the VLT, of which 95 were higher than the minimum redshift of $\\rm{z}=0.68$ to be used in the correlation function. A $3\\sigma$ upper-limit of \\twopcf$=2.8$ was found when 145 \\lya\\ absorber-galaxy pairs were binned in redshift space, in a bin of size $\\Delta\\sigma=1.0$, $\\Delta\\pi=2.0$ \\hMpcsev\\ along the projected separation and line of sight distances respectively.  The upper-limit in the cross-correlation was found to be $5.4\\sigma$ lower than the central peak in the galaxy auto-correlation within the same redshift range, \\acf, which was in our data equal to $10.7\\pm1.4$. Thus we have shown for the first time that the clustering between low column density absorbers and galaxies at a redshift of 1 is weaker than that between galaxies at the same redshift. ",
        "introduction": "\\label{sec-intro} This paper is the third in a series that investigates the extent to which \\lya\\ absorbers in the inter-galactic medium (IGM) are associated with galaxies.  At the current epoch only a minor fraction of the baryons in the universe exist within galaxies. Stars, galaxies and other stellar remnants make up 6\\% of the baryonic matter density of the universe, $\\Omega_{\\rm b}$, while within the galaxies cold ($\\rm{T}\\lesssim10^4$ K), condensed gas contributes a further $\\sim1.7\\%$ \\citep{Bregman2007}. Therefore at least 90\\% of baryons in the universe lie beyond the galaxies, in the IGM.  Part of this remaining fraction, which is only $10-100$ times the mean cosmic density, is more susceptible to shock heating. Shock waves that are caused by gravitational collapse travel along the filamentary structure heating the gas to $\\rm{T}=10^5-10^7$ K. These baryons form the diffuse warm-hot intergalactic medium (WHIM) \\citep{CenOstriker1999,Fukugita1998}. The detection and study of the WHIM using metals detected in ultra-violet \\citep{Verner1994,Tripp2000a,Danforth2005,Thom2008} and x-ray spectroscopy \\citep{Nicastro2002,Fang2006} is a major science motivation for the future space instruments XEUS (the X-Ray Evolving Universe Spectrometer) and the new COS (Cosmic Origins Spectrograph) on the Hubble-Space Telescope (HST).  A third of the gas that remains is probably located in the diffuse IGM at temperatures $\\rm{T}\\sim10^4-10^5$ K \\citep{Sembach2004,Lehner2007,Danforth2008}.  The neutral fraction of this ionised plasma is observed in quasar absorption spectra. Under the influence of dark matter this gas is thought to have collapsed to form clouds that are part of a filamentary network that traces the layout of the galaxies.  Two original opposing models questioned whether this gas is correlated with galaxies but is part of the giant cosmic web \\citep{Morris1993}, or whether the clouds are discrete and isolated with the gas contained within galaxy haloes \\citep{Lanzetta1995}.  Even though opinion is siding with the former there is probably no definitive answer to this question, as the typical location of these clouds is expected to be a function of their column density. At \\logcoldens$\\sim21$, the density of gas in an absorber would be comparable to that in the outer regions of a galaxy. Therefore at a small separation it would be pointless to draw a line where a galaxy terminates and the absorber begins.  In the low \\citep{Rao2003,Zwaan2005} and intermediate redshift universe \\citep{Brun1997,Chen2003} clouds with a column density \\logcoldens$\\gtrsim20.3$, (the damped \\lya\\ absorbers systems) have been been shown in imaging and spectroscopy to be associated with galaxies.  Meanwhile those clouds with a column density several decades less dense, \\logcoldens$\\leq14$, are more likely to make up the filaments or lie in the voids of the IGM. The absorber-galaxy association is also dependent on the epoch, for example because of the decrease in the number density of \\lya\\ lines since redshift $\\rm{z}\\sim1.7$ \\citep{Sargent1980,Bahcall1991,Janknecht2006}.  \\subsection{The Issues to Resolve and Investigations so far} The 2D 2-point cross correlation between \\lya\\ absorber and galaxy positions is a powerful way to quantitatively measure the relationship between galaxies and the IGM. It is hoped that these measurements will provide an insight into the degree of association between galaxies and the surrounding neutral gas. The degree of correlation is thought to depend on the redshift, the column density of the absorbing cloud, the galaxy magnitude, the presence of a galactic wind, and the line-of-sight distance and projected separation (distances $\\pi$ and $\\sigma$) between the galaxy and the absorber.  A common way to analyse the \\lya\\ absorber-galaxy cross-correlation is to contrast the results with the galaxy auto-correlation. If these values are comparable or the cross-correlation is greater at small separation, then the \\lya\\ absorbers could be claimed to be more of a feature of the galactic halo rather than the IGM.  We now give a short overview of such measurements that have been made over different redshifts.  \\subsubsection{Observations at Low Redshift} \\citet{Weber2006} (hereafter RW06) studied the HI-galaxy correlation at redshifts $\\rm{z}\\sim0$.  5317 gas rich galaxies that were detected in the HIPASS galaxy survey \\citep{Zwaan2003} were correlated with 129 absorbers from 27 lines of sight.  These spectra were gathered from STIS and GHRS data and all absorbers had \\logcoldens$=12.41-14.81$.  It was found that at small separations there is a strong correlation between absorbers and gas rich galaxies that is similar in strength to the galaxy auto-correlation function.  This supports the idea of matter clustered within the vicinity of galaxies (within the nodes of the cosmic web) be they absorbing clouds or other galaxies. However, unlike the galaxy auto-correlation that decays radially along the $\\sigma$ and $\\pi$ direction, \\twopcf\\ remained high out to 10 \\hMpc\\ or 1000 \\kms\\ along the line of sight. The proposed origin of this was the accretion of gas. In this scenario the gas falls in from the more diffuse IGM, drains along filaments and falls into a galactic halo causing an elongation in the redshift space correlation along the line of sight.  The gas rich galaxies that were detected using 21 cm radiation also tended to have a low mass ($\\rm{M}\\sim10^{8}$ \\Msolar), and it is galaxies such as these that \\citet{Keres2005} claims are dominated by this method of accretion.  \\subsubsection{A simulation of \\twopcf\\ at low redshift} The results of RW06 were compared with simulations in \\citet{Pierleoni2008}. Three simulations, each with a different galaxy wind model were run to a redshift of zero. Either no winds, strong winds or extremely strong winds were present.  Spectra were generated from 999 sightlines that sampled the simulation. The \\lya\\ absorbers that had a column density in the range \\logcoldens$=12.41-14.81$ were then correlated with $\\sim5000$ galaxies that resided in haloes of mass $\\rm{M}=8\\times10^{10}-10^{13.5}$ \\Msolar. However, contrary to RW06, \\citet{Pierleoni2008} found the galaxy auto-correlation was predicted to be stronger than the absorber-galaxy cross-correlation.  The correlation function was recalculated using only 27 sightlines to match the sample from RW06, and the sample noise increased. Bootstrap errors showed a difference between \\twopcf\\ and \\acf\\ at only a $1\\sigma$ significance level.  \\citet{Pierleoni2008} suggest that it is this sparse sampling that make the cross and auto-correlation functions indistinguishable. They also attribute the large ``finger-of-god'' observed in Figure 3 of RW06 to this effect.  \\subsubsection{Observations at low to intermediate redshift} \\begin{enumerate} \\item{{\\it\\citet{ChenMulchaey2009,Chen2005}}\\\\ \\\\ \\citet{Chen2005} studied the sightline towards PKS 0405-123, and compared the positions of 112 \\lya\\ absorbers from STIS echelle spectra with 482 galaxies within the redshift range $\\rm{z}=0.1\\rightarrow0.5$. These were collected using a magnitude limited survey down to $R\\leq20$.  The most striking result found by \\citet{Chen2005} was a difference in the level of correlation that depended on whether the galaxy in question was dominated by emission or absorption spectral lines.  The correlation function between galaxies that were dominated by emission, and \\lya\\ absorbers that had \\logcoldens$\\geq14.0$, was found to be similar to the galaxy auto-correlation. No significant cross-correlation was seen for those galaxies dominated by absorption.  \\citet{ChenMulchaey2009} then expanded on this result with two further lines of sight towards quasars HE 0226-4110 and PG 1216+069. It was found that within a projected separation of 1 $\\rm{h_{100}^{-1}\\;Mpc}$ \\lya\\ absorbers with \\logcoldens$\\ge14$ had a similar cross-correlation with galaxies as the auto-correlation of emission line dominated galaxies. This was $\\sim6$ times weaker than the self-clustering of absorption line dominated galaxies. They suggest that the cross-correlation was similar between emission dominated galaxies and \\lya\\ absorbers because both these occupy the same halo system.  The cross-correlation between any type of galaxy and \\lya\\ absorber with \\logcoldens$\\le13.5$ was observed to be very weak.} \\\\ \\item{{\\it{Papers I and II}}\\\\ \\\\ This paper is an extension to the investigation first conducted in paper I, \\citet{M_J_2006} (hereafter MJ06). Here the location of 636 galaxies, with redshifts determined using CFHT spectroscopy, were compared with 381 \\lya\\ absorbers found in the UV spectra of quasars taken with the Faint Object Spectrograph during the Hubble Key Project.  It was found that, excluding the high column density systems (\\logcoldens$\\geq17$), absorbers are correlated with galaxies at redshifts of $\\rm{z}\\leq0.8$ and that this correlation is weaker than the corresponding galaxy auto-correlation.  Paper II \\citep{Wilman2007} (hereafter W07) using the MJ06 dataset computed the 2-point correlation function along the line-of-sight ($\\pi$) and projected separation ($\\sigma$). A conclusion drawn from this work suggested that at small co-moving separations of $\\sigma< 0.4$, $\\pi<2\\;\\rm{h}^{-1}$Mpc, the correlation function increases marginally as the column density increases from \\logcoldens$=13-16$, then there is suggestion of a sharp increase in correlation above \\logcoldens$\\geq17$.  This was proposed as tentative evidence for the column density at which absorbers become a part of the galactic halo.  It was argued in W07 that a ``finger-of-god'' may also be caused by a large anisotropy, for example a quasar line of sight that runs parallel and close to to a large filament. However \\citet{Pierleoni2008} ruled out a ``finger-of-god'' being caused by this geometric effect in their simulation by showing that the \\twopcf\\ result remained the same no matter which orientation the sightlines had been directed. Instead a ``finger-of-god'' was caused if a sightline passed through a virialised system of gas and a paired galaxy.} \\end{enumerate}  \\subsubsection{Observations at high redshift} In order to observe the IGM-galaxy relationship at high redshift (where the \\lya\\ absorption has now reached optical wavelengths) \\citet{Adelberger2003} correlated the \\lya\\ lines from 8 quasar HIRES or ESI spectra with 431 Lyman-break-galaxies.  It was expected that greater concentrations of gas would be located near to the galaxies. This was found to be so as gas with a greater optical depth than the mean density was found to lie within $\\sim10$ \\hMpc\\ of the nearest galaxy.  However, evidence arose of an anti-correlation within $\\sim0.5$ \\hMpc.  It was hypothesised that this was caused by super-winds that, travelling at $\\sim600$ \\kms\\ \\citep{Pettini2001}, clear the immediate vicinity of the galaxies and heat the surroundings, decreasing the neutral fraction of HI and so increasing the transmission of QSO flux.  With a larger data set of 23 sightlines and 1044 UV selected galaxies \\citet{Adelberger2005} found that evidence for an anti-correlation at small separation significantly decreased. Only in a third of the cases was weaker absorption now detected around galaxies.  \\subsection{The motivation for this work} No results have been obtained for the redshift window $\\rm{z}\\sim1$. The main reason for this is the difficulty of acquiring a sufficient number of galaxy redshifts at this epoch that are within the same fields of view as quasars for which there is high resolution UV echelle spectral data. The aim of this investigation was to start to bridge this gap.  Thus we hope to shed light on the \\lya\\ absorber-galaxy correlation at redshifts $\\rm{z}\\sim1$. Then by comparison with the other works mentioned above determine whether the correlation evolves with redshift.  The structure of this paper will be as follows. The data reduction for the FORS2 galaxy survey and absorbers from the HST STIS E230M spectra will be described in Sections 2 and 3 respectively. In Section 4 the basics of the 2D 2-point cross correlation function used will be outlined.  Results are given in Section 5 and we present our discussion and conclusions in Section 6.  Throughout the paper we have adopted the cosmological parameters $\\Omega_{\\rm{M}}=0.3$, $\\Omega_{\\rm{\\Lambda}}=0.7$, $\\rm{h_{70}}=0.7$ and ${\\rm H}_0=100\\rm{h}_{70}$ km~s$^{-1}$~Mpc$^{-1}$. ",
        "conclusions": "The two straw-man models mentioned in the introduction were that the gas follows the galaxy distribution and collects in galaxy haloes, or it forms part of a diffuse network, located near but not necessarily causally linked with galaxies.  Figure \\ref{fig:absgalran} showed that we did not detect significant correlation between galaxies  and \\lya\\ absorbers that have $13\\leq$\\logcoldens$\\leq17.4$.  As only 145 absorber-galaxy pairs existed within $\\Delta\\sigma=0-5$, $\\Delta\\pi=0-20$ \\hMpcsev, in Figure \\ref{fig:absgalran}(a), no cuts in the column density were made as this would further increase an already large uncertainty.  An overwhelming majority of the \\lya\\ absorbers from the STIS spectra however had \\logcoldens$\\leq15$.  We calculated a $3\\sigma$ upper-limit of \\twopcf$=2.8$. The decrease in the $3\\sigma$ upperlimit to 1.9 in Figure \\ref{fig:abspert22}, caused by the increase in binsize to $\\Delta\\sigma=\\Delta\\pi=2$ \\hMpcsev\\ is further evidence that \\twopcf\\ is near to 0 when averaged over these distances.  This result was supported by the 1D plots of $\\xi_{\\rm AG}(\\pi)$ and $\\Xi_{\\rm AG}(\\sigma)$ along the line of sight and projected separation in Figure \\ref{fig:1Dabspert}.  The value in the central bin for the galaxy auto-correlation function was \\acf$=10.7\\pm1.4$. This was a $5.9\\sigma$ increase on the value in the equivalent bin of the absorber-galaxy cross-correlation that had a value of \\twopcf$=0.6\\pm0.9$.  Even when the value for the auto-correlation was reduced by the larger binning in Figure \\ref{fig:galpert22}, the difference between the 3-$\\sigma$ upper-limit in the cross correlation and the central bin in the auto-correlation was still $5.5\\sigma$.  This evidence implies that at a redshift of $\\rm{z}\\sim1$ galaxies that inhabit the nodes and filaments of the dark matter cosmic web are not strongly correlated with the low column density \\lya\\ absorbers (\\logcoldens$\\leq17$) that loosely trace this filamentary structure.  With a catalogue of 95 galaxies within a redshift range of $\\rm{z}=0.68-1.51$ that produced 138 galaxy pairs, it is unlikely that we would have been able to reproduce the result and distribution of pairs of the galaxy auto-correlation function found by \\citet{Pollo2005} and \\citet{Guzzo2008}. However our results were similar within the margin of error. When these studies are compared with the results at $\\rm{z}\\leq0.2$ there is a noticeable decrease in the auto-correlation.  When our results are compared with the values of the cross-correlation at redshifts $\\lesssim0.5$ the data suggest that the cross-correlation between low column density \\lya\\ absorbers and galaxies does not have a strong redshift dependence. This may be because of the flat evolution in the \\lya\\ absorber density distribution at redshifts $\\lesssim1.7$ \\citep{Janknecht2006}.  This conclusion is a stark contrast to that which was found by RW06. In that investigation \\twopcf\\ had been found to be comparable with \\acf\\ with values $\\approx11$ in the closest bin. There are three possible reasons for this difference. First of all the dataset of RW06 was at low redshift, where clustering is thought to have increased because of gravitational collapse. However an increase of this magnitude over a change in redshift of $\\rm{z}\\sim1\\rightarrow0$ does not seem plausible, particularly when results gathered by \\citet{Wilman2007}, \\citet{Chen2005,ChenMulchaey2009} and \\citet{Pierleoni2008} cite lower values for \\twopcf\\ at redshifts of $\\rm{z}\\lesssim0.5$ and 0.  Another reason may be because of the galaxy population. RW06 used the HIPASS galaxies that were detected using 21 cm emission and targeted in a blind survey. Therefore as well as being gas rich it included galaxies of any mass. Hence a strong correlation and velocity dispersion along the line-of-sight were expected as gas falls through the potential of the filamentary structure and is accreted directly by the low-mass galaxies. This is the method of `cold-accretion' proposed in \\citet{Keres2005}.  Our magnitude limited galaxies (M$_B\\sim-17\\rightarrow-21$) were collected in an incomplete survey at higher redshifts than that of the HIPASS catalogue.  Consequently the galaxies for which we were able to get spectroscopic data would have a higher mass than most of the HIPASS candidates. \\citet{Keres2005} claim that the dominant mode of accretion for galaxies with a dark halo mass $\\rm{M}\\gtrsim10^{11.4}$ \\Msolar\\ is the more conventional method; where gas is shock heated to the virial temperature of the halo before cooling and collapsing onto the galaxy. This could mean less neutral gas is expected in the immediate vicinity of the galaxies in our sample when compared to those from RW06.  Another reason why the cross-correlation may seem larger in \\citet{Weber2006} is because of the smaller bin size of $\\Delta\\sigma=\\Delta\\pi=0.1$ ${\\rm h_{100}^{-1}}$Mpc that was used. We could argue that the absorbers that are at $\\sigma<0.1$ ${\\rm h_{100}^{-1}}$Mpc are contained within the galaxy halo and may even be virialised, so \\twopcf\\ per bin is likely to substantially increase.  Alternatively these discrepancies could be due to the problem suggested by \\citet{Pierleoni2008}. \\citet{Weber2006,Wilman2007,Chen2005,ChenMulchaey2009} and to a greater extent our dataset all either contain few pairs and/or few lines-of-sight. The appearance or absence of artifacts such as a strong ``finger-of-god'' in the correlation pattern could be due to the small size of the sample.  Unfortunately because we only had 13 galaxies in the catalogue that were dominated by absorption features we were unable to test the hypothesis of \\citet{Chen2005} and \\citet{ChenMulchaey2009} that \\twopcf\\ is stronger for those galaxies that are dominated by emission and comparable to the emission dominated galaxy auto-correlation function. However we still see a low correlation in the smallest bin, despite the abundance of emission dominated galaxies in our sample.  We are hoping to improve on these findings with more data collected from three further STIS E230M lines-of-sight and galaxies from the VIMOS instrument. This will enable us to correlate galaxies and \\lya\\ absorbers out to a projected separation of $\\sigma\\sim7$ \\hMpcsev\\ at a redshift of $\\rm{z}\\sim1$.  We also plan to investigate the 2D 2-point correlation function using the Galaxies-Intergalactic Medium Interaction Calculation (GIMIC)  SPH simulation \\citep{Crain2009}."
    },
    "0911/0911.4220_arXiv.txt": {
        "abstract": "We describe a 4D Eigenvector 1 (4DE1) space that serves as a surrogate H-R diagram for quasars. It provides a context for describing and unifying differences between all broad line AGN. Quasar spectra can be averaged in a  non-random way using 4DE1 just as stellar spectra can be averaged non-randomly within the OBAFGKM classification sequence. We find that quasars with FWHM H$\\beta$ less than (Population A) and greater than (Population B) 4000 km s$^{-1}$ show many significant differences that may point to an actual dichotomy. Broad line profile measures and fits reenforce the idea of a dichotomy because they are fundamentally different: Pop.A - Lorentzian-like and Pop.B - double Gaussian. The differences have implications both for BH mass estimation and for inferences about source structure and kinematics. ",
        "introduction": "\\label{intro}  Studies of quasar spectra have fallen out of fashion over the past ten years. This impression is supported by the small number of citations involving optical/UV spectroscopic work found from a random NED search of citations for even the brightest (e.g., PG) sources. One can identify at least two reasons for this paucity of studies: 1) a widespread belief that all broad-line spectra are essentially the same and/or 2) the impression that a deeper understanding of broad-line phenomenology and physics is unobtainable. The one exception, reflecting reason 1) involves computation of mean/median composite spectra from large samples (e.g. 2dF, SDSS). Of course if reason 1) is not correct then these results will reenforce reason 2). Just as  the indiscriminate averaging of OBAFGKM  stellar spectra would yield bizarre conclusions.  The purpose of this review is to convince the reader that reason 1) is incorrect and that understanding quasar spectral diversity can stimulate new ideas that will remove the sense of hopelessness reflected in reason (2).  The question is how best to represent quasar phenomenological diversity and unite it into a coherent picture. The history of stellar studies immediately comes to mind although quasars are almost certainly more complicated sources. At the very least source orientation to our line-of sight likely adds a serious complication. We have been working on a spectroscopic unification for quasars that, in fact, embraces all broad line emitting AGN. Gradually the hope emerged of finding a diagnostic space capable of serving as a surrogate H-R diagram for quasars. A 2D H-R diagram works quite well because, among other things, stellar orientation plays no major role in spectroscopic studies. The full power of the stellar H-R diagram requires exploitation of more parameter dimensions. Certainly an equivalent spectroscopic diagnostic space for quasars will require more than two dimensions if only to remove the orientation-physics degeneracy. In 2000 we  proposed a 4D Eigenvector 1 parameters space (hereafter 4DE1; \\citealt{Sulentic00a}). 4DE1 remains the most promising way to emphasize the spectroscopic diversity while also contextualizing the diverse types of broad-line emitting sources.  4DE1 space emphasizes observables with the largest intrinsic dispersions including the most statistically significant line profile differences. 4DE1 has roots in the Principal Component Analysis (PCA) of the Bright (PG) Quasar Sample (\\citealt{BG92}; BG92) as well as in correlations that emerged from ROSAT \\citep{Wang96}. 4DE1 in its simplest form involves two BG92 measures: 1) full width half maximum of broad H$\\beta$ (FWHM H$\\beta$) and 2) equivalent width ratio of optical FeII and broad H$\\beta$: R$_{FeII}$ = W(FeII $\\lambda$4750 blend)/W(H$\\beta$). We added measures of 3) the soft X-ray photon index ($\\Gamma$$_{soft}$) and 4) CIV$\\lambda$1549 broad-line profile displacement at half maximum (c(1/2)). Other points of departure from BG92 involve comparison of radio-quiet (RQ) sources with a large radio-loud (RL) sample. We also subordinate [OIII]$\\lambda\\lambda$4959,5007 measures (although see: \\citealt{Zamanov02,Marziani03a,Marziani09}). 4DE1 tells us that all quasars are not similar whether we observe them at optical, UV or X-ray wavelengths. ",
        "conclusions": ""
    },
    "0911/0911.2232.txt": {
        "abstract": "In this paper, we study the distribution functions that arise naturally during self-similar radial infall of collisionless matter. Such matter may be thought of either as stars or as dark matter particles. If a rigorous steady state is assumed, then the system is infinite and is described by a universal distribution function given the self-similar index. The steady logarithmic potential case is exceptional and yields the familiar Gaussian for an infinite system with an inverse-square density profile. We show subsequently that for time-dependent radial self-similar infall, the logarithmic case is accurately described by the Fridmann and Polyachenko distribution function. The system in this case is finite but growing. We are able to embed a central mass in the universal steady distribution only by iteration\\textbf{,} except in the case of massless particles. The iteration yields logarithmic corrections to the massless particle case and requires a `renormalization' of the central mass. A central spherical mass may be accurately embedded in the Fridmann and Polyachenko growing distribution however. Some speculation is given concerning the importance of radial collisionless infall in actual galaxy formation. ",
        "introduction": "The relation between the formation of black holes and of galaxies has developed into a key astrophysical question, from the early papers by Kormendy and Richstone \\cite{KR1995}, to the more recent discoveries by Magorrian et al. \\cite{Ma98}, Ferrarese and Merritt \\cite{FM2000}, and Gebhardt et al.\\cite{Geb2000}. Recent papers \\cite{Graham04,KB2009} also establish a correlation between black hole mass and bulge luminosity deficit. These papers establish a strong correlation between what is essentially the black hole mass and the surrounding stellar bulge mass (or velocity dispersion). The origin of this proportionality, which extends well beyond the gravitational dominance of the black hole, remains uncertain. But it is generally taken to imply a coeval growth of the black hole and bulge.  Various proposals have been offered to explain the black hole mass-bulge mass proportionality as a consequence of the AGN phase. There is as yet no generally accepted scenario although a kind of `auto-levitation', based on radiative feed-back from the accreting black hole to the star-forming gas that in turns limit accretion, is plausible. In any event there remains the question of the origin of the black hole seed masses. In some galaxies at very high red shift the inferred black hole masses are already of order $10^{9}~M_{\\odot}$ \\cite[e.g.]{Kurk2007} after about one Ga of cosmic time. This may require frequent, extremely luminous early events \\cite[e.g.]{W09}, or it may suggest an alternate growth mechanism.  The latter possibility is reinforced by the detection of a change in the normalization of the black hole mass-bulge mass proportionality in the sense of relatively larger black holes at high red shift \\cite[e.g.]{Mai2007}. As suggested in that paper it seems that the black holes may grow first, independently of the bulge. The collisionless matter that we invoke might be stars or it might be the dark matter itself.  Recently \\cite{PFP2008} have studied the possible size of the dark matter component in black hole masses. By assuming that the dimensional or 'pseudo-phase-space density' \\cite[ i.e. Henriksen 2006b, e.g.]{H2006} is strictly constant they convert the relativistic accretion of the dark matter into an adiabatic Bondi flow problem and obtain the resulting accretion rate. Then by adopting the mass proportionality between bulge and black hole and fitting boundary conditions from cosmological halo simulations, they deduce that between 1\\% and 10\\% of the black hole mass could be due to dark matter.  If we accept this result at face value, a seed mass of say $10^{6}M_{\\odot}$ could have grown from dark matter. It would now be part of a super massive black hole that subsequently grew in the AGN phase. Some seeds may be primordial. As early as 1978, through fully general relativistic numerical collapse calculations, \\cite{BH1979} predicted primordial black hole masses in the range $10^{2}$ to $10^{6}M_{\\odot}$.  % \\begin{comment} In this and subsequent papers we will attempt to embed a black hole (or at least a central mass) in a distribution of particles that arises naturally during the formation of the galactic core. We will predict the consequent density cusp profile and that of the velocity dispersion variation in the cusp. It is perhaps significant in the light of the extensive orbital study of (\\cite{V3CdZ08}) that we are most successful in a time-dependent case. For these authors suggest that the central mass may render the orbits chaotic and non-stationary. We find in the steady systems that a series of iterations diverges at the centre. \\end{comment} {}  Stellar density cusps surrounding black holes have been studied extensively previously. Classic studies by \\cite{P1972} and by \\cite{BW76} dealt with the problem of feeding the black hole from a filled loss cone (nearly radial orbits). In addition to these diffusion studies, Young \\cite{Y1980} explored the cusps produced by the adiabatic growth of a black hole in a pre-existing isothermal stellar environment. This was extended by \\cite{Q1995} and by \\cite{MH2002} to more general environments. In a cosmological context, Bertschinger \\cite{Bertschinger} also studied the growth of a central black hole by radial infall. The conclusions were that the black hole induced cusps were never flatter than $r^{-1.5}$ (the isothermal and cosmological case) and that no black hole mass-bulge mass correlation was established \\cite{MH2002}. The latter conclusion has spurred the investigation of coeval dynamical growth of the black hole and bulge in contrast to adiabatic growth \\cite{MH2003}. In this paper we discuss possible coeval growth due to radial infall of stars and dark matter that both feed the black hole and establish a collisionless cusp with a self-similar distribution function.  % \\begin{comment} Such concerns have passed from the abstract to the practical with the detection of stellar cusps around galactic nuclear black holes \\cite{G2009}. This latter paper reports stellar cusp power laws for the central Milky Way in the range $-1.1\\pm0.3$, significantly shallower than the adiabatic $-1.5$ or the zero flux \\textbf{\\cite{BW76} }$-1.75$. We seek to find under what general conditions such power laws may arise during collisionless infall. In the present paper that is restricted to radial infall, we do not find cusps flatter than $r^{-2}$.  An effective alternate method of evolving flat cusps is to invoke tight binary black hole systems produced by mergers that 'scour' the stellar environment \\cite[for e.g.]{MS2006,NM99}. Depending on the power law assumed for the initial stellar environment, Merritt and Szell simulate scoured power laws that can be as flat as $-1$ in an initially $-1.5$ cusp, and as flat as $-0.5$ in an initially $-1$ stellar cusp. Flatter values are reduced to essentially constant density cores. The $-1$ slope as noted above is a reasonable fit to the Milky Way (ibid).  Subsequently, but taking at least the central relaxation time, the cusps should regenerate to the zero flux condition ($r^{-1.75}$) according to the simulations in \\cite{MS2006}. However this will extend only out to about $0.2$ of the gravitational influence radius of the black hole. This regeneration is not thought to be relevant to the Milky Way central stellar cusp, but may be present on small scales elsewhere.  Such a picture is seductive, especially given the recent detection of a strong correlation between the nuclear black hole mass and the central luminosity deficit \\cite{KB2009}. However the correlation in itself only implicates the influence of the black hole. It does not necessarily require the merger history, which in any case is unlikely to be the same for different galaxies. Consequently we explore in this series of papers \\cite{HLeDMcM09a,HLeDMcM09b,HLeDMcM09c} whether cusps as flat as those resulting from scouring might also be produced during the dynamical formation of the black hole or mass concentration. In this first paper we present a summary of the situation for radial orbits since the arguments are typical but easily carried out in that case. \\end{comment} {}  In this and subsequent papers \\cite{HLeDMcM09a,HLeDMcM09b,HLeDMcM09c} we will embed a black hole (or at least a central mass) in a distribution of particles that arises naturally during the formation of the galactic core. The `natural evolution' is taken to be that of self-similarity since power-law behaviour is observed both in reality and in simulations. In this fashion we introduce some uniqueness into the form of the resulting distribution function that describes the collisionless matter that surrounds the central mass. We will predict the consequent density cusp profile and that of the velocity dispersion variation in the cusp.  In this first paper we confine ourselves to radial infall. The first general result on radial infall was in fact given in \\cite{FG84} and the Bertschinger \\cite{Bertschinger} result follows by putting their parameter $\\epsilon_{FG}=1/3$. However neither this work nor that of Bertschinger attempted to infer closed forms for the equilibrium distribution function. This was begun by Henriksen \\& Widrow (1995) (hereinafter \\cite{HW95}). In \\cite{H2006} it is pointed out that\\textbf{ }$a=9/8$ ($\\epsilon=3$) yields the Bertschinger solution in its entirety, including the recently heralded power law of the proxy for phase space density. The parameter $\\epsilon=3\\epsilon_{FG}$ and it is related to the index called `$a$' below, both by simulations and theory.  We deduce in this paper two principal distribution functions that arise naturally. One is steady and infinite and can not contain a central mass exactly. However we can treat the central mass as a perturbation to the gravitational field and iterate on the Boltzmann and Poisson equations to find the flattest possible cusp in a steady infinite system surrounding a central mass. There is a special logarithmic case in this category that yields a Gaussian distribution function and an inverse square law density cusp. Our derivation is new but the result agrees with the form deduced in \\cite{LB1967} based on statistical mechanics.  Our second result is most important in that it represents a growing cusp wherein a central mass may be embedded exactly. It is the Fridmann and Polyachenko distribution function, which has not previously been identified in this physical context. We derive it theoretically and verify it with simulations. It also corresponds to a logarithmic potential for which the self-similar index $a$ (see below) is unity. The compatibility with a central mass allows us to give a growth rate for that mass. It is perhaps significant in light of the extensive orbital study of \\cite{V3CdZ08} that we are most successful with our embedding in a time dependent case. For these authors suggest that the central mass may render the orbits chaotic and nonstationary.  We are aware that such a radial system is unstable to the radial orbit instability (ROI) on small scales \\cite[hereinafter for MacMillan et al. 2006]{MWH2006}. However even fully cosmological simulations show an outer envelope wherein the orbits are trending to be radial. Moreover isolated halos, which also show statistical relaxation that is begging to be understood, show quite radial orbits in the envelope \\cite{MWH2006}. In such cases the `central mass' may be a central `bulge' of dark matter that forms rapidly. We speculate below that the growth of this bulge in such an envelope may be described in terms of radial infall, and that this infall may continue hierarchically to smaller scales after relaxation by the ROI and clump-clump interactions \\cite{MH2003}. Such interactions would remove angular momentum from some particles in favour of others and so create the continuing radial infall.  We begin the next section with the general formulation in spherical symmetry. Subsequently in section \\ref{sec:Radial-Orbit-Steady} we discuss the various possible rigorously steady distribution functions (DF from now on) for radial orbits. We show in section \\ref{sec:The-Logarithmic-Case} that the DF of Fridman and Polyachenko \\cite[hereafter called FPDF]{FP1984} describes a system of radial orbits that is growing self-similarly. This is contrasted in the same section with an infinite steady system for which the DF is Gaussian. After some discussion, we\\textbf{ }give our conclusions. ",
        "conclusions": "In this paper we have presented unique distribution functions based on the self-similar evolution of a gravitating system of radial orbits. The rigorously steady state corresponds to an infinite system while the time dependent system is finite and growing. Both of these systems have been confirmed by numerical simulations, with the confirmation of the growing mode being most dramatic.  Both systems have been derived coherently from the basic equations by using a convenient formulation of self-similarity. We have considered how these distributions are affected by the presence of a central spherical mass, which may not necessarily be point-like. The steady system is perturbed outside the central mass to an inverse cube law with logarithmic corrections as we discuss further below. The growing time-dependent case forms a perfect inverse-square density cusp that can contain a central mass and is growing proportionally to the time since formation. The material with which we are concerned is collisionless and so may refer to dark matter in an early stage of formation or to stars at a later stage. The restriction to radial orbits in this paper is removed in subsequent papers, although both in theory and in the numerical simulations a bias to radial orbits remains on the larger scales.  For this latter reason we think that the results of this paper may be relevant in reality for two reasons. In the first instance the distribution functions may develop outside a rapidly formed large central bulge wherein angular momentum has been important. This may correspond to the NFW core radius, and our detailed results are compatible with the density profile of the simulations outside this radius. The second plausible scenario is that, on much smaller scales surrounding an actual black hole, various instabilities may allow radial infall of stars and dark matter onto the centre. This fits the time-dependent infall model and allows the black hole to grow proportionally to the time from the onset of the radial infall. We detail these results in what follows.  In section (\\ref{sec:Radial-Orbit-Steady}) on the steady state we deduced the steady, spherically symmetric DF with radial orbits that yields infinite systems with power law profiles. This was found previously but we rederive it here as a limit of the time-dependent equations. We have perturbed this DF by embedding a central mass. The DF is universal for these systems ($f=K|E_{o}-E|^{1/2}$) but the potential-density pair depends on the self-similar index $a$ which in turn depends on dominant constants or boundary conditions. If it is a memory of a cosmological fluctuation profile, then $a=3\\epsilon/(2(\\epsilon+1))$.  This steady DF can not contain a central mass without being perturbed, except for the Keplerian limit wherein $a=3/2$ and there is a mass-less cusp with an inverse cube density profile. Iterating the equations to find the perturbation produced by an embedded central mass predicts a transition region in the halo of the mass that is an $r^{-3}$ profile modified only by logarithmic corrections. It does not correspond to the power law density of stars found close to the black hole in the galactic centre\\cite{G2009}, but may apply outside a bulge mass confined to the NFW scale radius\\textbf{.}  Since the density is linear in the potential we were also able to find a solution \\emph{that broke the self-similarity} by seeking non-power-law solutions of Poisson's equation (pure power laws only exist in the limits) with the density (\\ref{eq:lindens}). The density profile of the cusp can not be flatter than $r^{-2.5}$ near a central point mass, but it may be as flat as $r^{-2}$ at large scales in the low density limit. Such behaviour is too steep to explain the cusp observed around the Sagittarius A{*} black hole. All possible descriptions of a steady system of radial orbits are clearly excluded by this observation. These descriptions are not excluded outside a central bulge however.  In the section concerning the logarithmic potential exception we showed that it corresponded to continuing time-dependent infall. We found both theoretically and numerically that the Fridmann and Polyachenko DF is the distribution established by continuing radial infall. Remarkably, it allows a growing central point mass (or indeed a bulge mass on a larger scale) to be embedded in the infall self-consistently. % \\begin{comment} We were able to contrast this result with a rigorously steady, infinite system of radial orbits. This is described by a Gaussian DF that has also been inferred in the past for fully relaxed, infinite systems \\textbf{\\cite{LB1967}}. \\end{comment} {} In both the steady and the time-dependent cases the density profile is $r^{-2}$% \\begin{comment} , but a black hole is not readily incorporated into the steady case even by iteration \\end{comment} {}. The time-dependent logarithmic case allows a consistent calculation of the central mass growth rate according to $M_{*}\\propto t$. Such a growth rate from a surrounding envelope was found in \\cite{MH2003}. This simulation used a true N-body code that treated particle-particle scattering, which scattering led eventually to radial infall of some of the particles.  The growth rate of a central mass (i.e. a collisionless concentration, not a true black hole) from a reservoir of radial orbits is zero if there is a rigorous steady state. % \\begin{comment} However in a state of self-similar virialisation we can expect them to settle into a bulge as they become trapped by the increasing mass. \\end{comment} {} The growth rate is not zero if the central mass is a black hole, since then the outward bound radial orbits are suppressed. In that case however the steady state is only an approximation except in an infinite system. For a finite system the true timescale would be the free-fall time of the bulk of the accreting mass.  One is inclined not to take these growth estimates seriously, since they simply assume an endless supply of radial particles, and hence an arbitrarily filled loss cone. In fact the radial alignment required to hit a growing black hole from a few hundred parsecs is at least one part in $10^{8}$ to $10^{10}$ depending on the mass of the black hole! This suggests that instead \\cite[see e.g.]{MH2003} the actual growth involving radial orbits may be by way of a multi-stage process. In the first stage, radial orbits accrete from the galactic halo to form a bound spherical bulge of intermediate size, due to finite angular momentum about the centre \\cite{MH2003}. They are trapped there either by the usual mechanism of self-similar infall as the potential increases in time with increasing internal mass, or by dissipative interactions. If there is substructure in the collisionless matter (e.g. stars and dark matter clumps), then these are able to produce dissipational collisions. Ultimately these collisions and tidal interactions can lead to a more gradual growth of a more central mass \\cite[e.g., in the Carnegie meeting]{MH2003}.  Recently a high resolution study \\cite{Stetal2009} of sub-structure in an isolated halo revealed that this sub-structure disappears in the inner few parsecs. this coincides with the region where the halo is becoming spherical and where the density power-law is flattening to less that $1$. The interactions leading to the sub-structure disappearance may well lead to relaxation.  In addition the radial orbit instability can lead to the development of a bar \\cite{MWH2006}. This bar can then transport angular momentum away from the bulge by the ejection of particles. Such `interrupted accretion' may repeat several times on the way to the actual central object. The rapid accretion of a bulge is in fact the way in which dark matter halos are thought to grow \\cite{Zhao2003,Lu2006} initially. This is then followed by a slower growth phase. The DF (\\ref{eq:FPDF}) can be used to describe the environment of the central mass on each scale of the interrupted cascade.  In this connection we refer to the work of Mutka \\cite{Mutka09} on gravitationally lensed galaxies with double images. He concludes that there are two classes of density cusps with the larger sample (about 80\\%) showing a logarithmic density slope of $\\approx-1.95$ well inside the NFW scale radius. The other 20\\% show this slope as $\\approx-1.45$. These may be unresolved triple image lens and ,if so, the measured value should be rejected.  Mutka's result is a measure of the total mass distribution rather than just the dark matter. Perhaps we are seeing enhanced relaxation in the mixture of stars and dark matter, that leads towards an isothermal cusp, rather than the shallower cusps of the dark matter simulations. It is significant that this inverse square slope is also frequently found by direct dynamical modeling of galaxies \\cite{vanderMarel09}.  However an inverse square slope is not restricted to a system of purely radial orbits as the isotropic isothermal distribution shows. In a subsequent paper we survey anisotropic distribution functions in spherical spatial symmetry that also have a self-similar memory. Some of these also provide an inverse square density profile.  The significance of the hierarchy of co-evolving structures is that there will always be a mass correlation between them. Thus if the mass concentration derives its ultimate mass $M_{\\bullet}$ from a halo of radius $r_{h}$, while $r_{s}$ encloses the mass that forms the ultimate bulge mass $M_{s}$ then \\begin{equation} \\frac{M_{\\bullet}}{M_{s}}=\\frac{r_{h}}{r_{s}}.\\end{equation} This assumes the pure inverse square density law, which might in fact have a logarithmic correction. In the subsequent paper, we shall find a slightly more general correlation that involves the self-similar memory. Taken at face value this simple relation gives $r_{h}/r_{s}\\approx100$.  This paper comprises a systematic derivation of the distribution functions that arise through self-similar evolution of gravitating systems of radial orbits. In addition we have studied the effects of a central mass on these distributions. We do not find density profile as flat as those proposed in \\cite{MS2006} and \\cite{NM99} as a result of density scouring by merging black holes. Thus our calculations probably do not describe the cusp around Sagittarius A{*}, although this does not exclude other applications. Subsequent papers in this series address this problem by allowing anisotropy in velocity space.% \\begin{comment}"
    },
    "0911/0911.2776_arXiv.txt": {
        "abstract": "We present evidence for very high gas fractions and extended molecular gas reservoirs in normal, near-infrared selected (BzK) galaxies at z$\\sim$1.5. Our results are based on multi-configuration CO[2-1] observations obtained at the IRAM Plateau de Bure Interferometer. All six star forming galaxies observed were detected at high significance. High spatial resolution observations resolve the CO emission in four of them, implying sizes of the gas reservoirs  of order of 6--11~kpc and suggesting the presence of ordered rotation. The galaxies have UV morphologies consistent with clumpy, unstable disks, and UV sizes that are consistent with those measured in CO. The star formation efficiencies are homogeneously low within the sample and similar to those of local spirals -- the resulting gas depletion times are $\\sim0.5$~Gyr, much higher than what is seen in high-z submm galaxies and quasars. The CO luminosities can be predicted to within 0.15~dex from the observed star formation rates and stellar masses, implying a tight correlation of the gas mass with these quantities. We use new dynamical models of clumpy disk galaxies to derive dynamical masses for our sample. These models are able to reproduce the peculiar spectral line shapes of the CO emission. After accounting for the stellar and dark matter masses we derive molecular gas reservoirs with masses of 0.4--1.2$\\times10^{11}M_\\odot$. The implied conversion (CO luminosity-to-gas mass) factor is very high: $\\alpha_{\\rm CO}=3.6\\pm0.8$, consistent with a Galactic conversion factor but four times higher than that of local ultra-luminous IR galaxies that is typically used for high-redshift objects. The gas mass in these galaxies is comparable to or larger than the stellar mass, and the gas accounts for an impressive 50--65\\% of the baryons within the galaxies' half light radii. We are thus witnessing truly gas-dominated galaxies at $z\\sim1.5$, a finding that explains the high specific SFRs observed for $z>1$ galaxies. The BzK galaxies can be viewed as scaled-up versions of local disk galaxies, with low efficiency star formation taking place inside extended, low excitation gas disks. These galaxies are markedly different than local ULIRGs and high-z submm galaxies and quasars, where  higher excitation and more compact gas is found. ",
        "introduction": "Over the last decade, deep and wide multiwavelength galaxy surveys have been key in addressing critical issues regarding galaxy evolution.  By using a variety of color selection techniques, and deriving photometric and/or spectroscopic redshifts, different galaxy populations have now been probed to unprecedented detail up to at least $z\\sim3$. Also, the history of cosmic star formation as well as the corresponding build--up of stellar mass has now been constrained using different observational approaches.  Less well understood are the physical processes that regulate the star formation activity in these young systems and the question of if and how galaxy star formation at high redshift differs from what is seen in the local Universe. While it has been established that star formation rates (SFRs) in galaxies were on average higher in the past and that the cosmic SFR density peaked at $z>1$ (see, e.g., LeBorgne et al. 2009), it is still a matter of debate how much of this increase is due to more frequent merging/interactions between galaxies, or other processes. Analysis of the ultraviolet (UV) and optical rest frame morphology and H$\\alpha$ velocity fields of high redshift galaxies suggest that the galaxies responsible for the bulk of the SFR density at $z\\sim1$~to~3 are disks (e.g., Bell et al. 2005; Elbaz et al. 2007; Genzel et al. 2008). These findings confirmed early hints for the presence of extended disks at high redshift, based on damped Ly$\\alpha$ studies (e.g., Wolfe et al. 1986). Some authors, however, advocate a predominant role of mergers based on, e.g., irregularities in high signal to noise (S/N) 3D observations of emission lines (e.g., Flores et al. 2006). On the other hand, it is also found that distant star forming galaxies show irregular morphologies with high clumpiness, especially in the rest-frame UV (e.g., Cowie et al. 1996; Chapman et al. 2003; Daddi et al. 2004; Elmegreen \\& Elmegreen 2005). In this context it is interesing to note that the existence of such bright clumps can explain the observed kinematic irregularities in high--z galaxies (i.e., without the need to require the presence of mergers, Immeli et al. 2004; Bournaud et al. 2008). More direct evidence for an important role of in situ, quiescent star formation in distant galaxies came from the estimate that the duty cycle of star formation is high in massive $z\\sim2$ galaxies (Daddi et al. 2005; 2007a; see also Caputi et al. 2008), with typical durations of $\\approx0.5$--1~Gyr. This is much longer than the typical duration expected for starbursts ($\\simlt100$~Myr)  triggered by mergers, based on simulations (e.g., Mihos \\& Hernquist 1996; Di~Matteo et al. 2008) and also on the observed gas consumption timescales for IR luminous merging galaxies (e.g., Downes \\& Solomon 1998).  A recent major step forward in characterizing the nature of star formation in distant galaxies was based on the discovery that star forming galaxies define a narrow locus in the stellar mass-SFR plane (in this paper we will refer to the ratio of SFR to stellar mass as the specific SFR, SSFR).  This correlation has been estimated to have a slope of 0.7--1.0 in log space and a dispersion of a factor of only $\\sim$2 from redshifts $z\\sim0.1$ to $z\\sim3.0$ (Noeske et al. 2007; Elbaz et al. 2007; Daddi et al. 2007a; Magdis et al. 2009).  The normalization in this correlation, i.e. the SSFR at a given stellar mass, is higher by a factor of 30 at $z=2$--3 compared to the local Universe. At even higher redshifts the situation is unclear as the SFRs are typically less well constrained (as they are primarily based on UV luminosities). However, the available data is still consistent with the picture that the correlation extends to earlier epochs ($3.5<z<7$) with a normalization similar to that seen at $z=2$ (Daddi et al. 2009; Stark et al. 2009, Gonzales et al. 2009). The existence of this correlation has major implications for our understanding of galaxy formation. For example, the cosmic evolution of the normalization of the relation (Daddi et al. 2009; Pannella et al. 2009; Gonzales et al. 2009) is similar to the cosmic evolution of the SFRD (Hopkins \\& Beacom 2006; 2008; Le Borgne et al. 2009, etc). It is possible in fact to describe most of the SFRD evolution as the combined result of the evolution of the stellar mass-SFR correlation and that of the space densities (e.g., the mass/luminosity functions) of galaxies. The possible peak at $z\\sim2.5$--3  (Pannella et al. 2009; Magdis et al. 2009; Daddi et al. 2009a) of the SSFRs is thus likely more fundamentally related to the physics underlying galaxy formation than the peak at $z>1$ of the SFRD evolution, as it is hinting at a change in the physical processes that regulate star formation.   The remarkable tightness of this relation implies a very high degree of homogeneity in these galaxies. In other words, despite the physical complexity of the process by which stars are formed, for a normal, near-IR selected star forming galaxy, the star formation rate is ultimately linked to its stellar mass (which in itself is presumably related to the mass of the underlying dark matter halo).   \\begin{figure*}[!ht] \\centering \\includegraphics[width=13.0cm,angle=0]{f1.eps} \\caption{Velocity averaged maps of the CO[2-1]  emission in our sample BzK galaxies. The (cleaned) images were obtained using natural weighting and are based on data from all observations (as indicated on top of the individual panels). Contour levels start at $\\pm2\\sigma$ and are in steps of $1\\sigma$ (see Tab.~\\ref{tab:1}). The size and orientation of the beam is indicated in the bottom--left corner.  The panels in the top two rows are 25$''$ in size, those in the bottom row are 45$''$ in size. The cross in each panel corresponds to the VLA 1.4~GHz radio continuum position (see Tab.~\\ref{tab:1}) and are $\\pm2''$ ($\\pm 17$~kpc for $z=1.5$) in size. } \\label{fig:2D}% \\end{figure*}     \\begin{figure*}[!ht] \\centering \\includegraphics[width=13.0cm,angle=0]{f2.eps} \\caption{ACS images in the F775W band (i-band) of the target BzK galaxies. CO[2-1]  contours are the same as shown in Fig.\\ref{fig:2D}. The panels in the top two rows are $7.5''$ in size. The panels in the bottom are 15$\"$  in size. For reference, 1$''$ subtends 8.5~kpc at $z=1.5$. } \\label{fig:acs} \\end{figure*}   \\begin{figure*} \\centering \\includegraphics[width=13.0cm,angle=0]{f3.eps} \\caption{Spectra of CO[2-1] emission binned in steps of 25~km~s$^{-1}$. The yellow color indicates the regions where positive emission is detected. These regions have been used to derive total integrated fluxes and the velocity averaged maps in Fig.~\\ref{fig:2D}. The red lines show best fitting double Gaussian profiles to the spectra (see Tab.~\\ref{tab:2} for fit values). In the fits, we forced the FWHM of the two Gaussian functions to be the same within a given galaxy. For the four galaxies with the highest quality spectra we also show single Gaussian fits. Those yield significantly worse fits to the wings of the CO profiles, except for the face-on galaxy BzK-16000 that is consistent with a single Gaussian. Velocities are computed with respect to the tuning frequencies of the observations as listed in Tab.~\\ref{tab:1}. } \\label{fig:1D} \\end{figure*}    \\begin{figure*} \\centering \\includegraphics[width=17cm,angle=0]{f4.eps} \\caption{{Left:} uv-coverage of the four target BzK galaxies that were observed also with the higher resolution B-configuration. {Right:} Signal amplitude versus baseline length in the uv plane. The solid curves show circular Gaussian models with FWHM as labeled. As a reminder, for a point/unresolved source the visibility amplitudes would be constant with uv distance. } \\label{fig:uvcov} \\end{figure*}   In order to make further progress in understanding how galaxy formation proceeds in distant galaxies, and to shed light on the origin of the stellar mass-SFR correlations, we need to start probing the properties of molecular gas reservoirs in these galaxies, since it is the molecular gas out of which stars form. If the assembly of galaxies is mostly driven by mergers, one would expect most star formation in merger-driven starbursts with concentrated gas cores (like in submm galaxies) and very efficient star formation, and much less molecular gas and star formation in regular objects. If, on the other hand, most star formation is due to smooth gas flows, then star formation would predominantly take place in more quiescent disk galaxies, accreting gas from their surrounding medium. In addition, one could speculate that the tightness in the stellar mass-SFR correlation might be pointing to homogeneous properties of the molecular gas reservoirs and star formation modes inside the galaxy populations.  The increase in the normalization of this correlation with cosmic lookback time could then possibly be simply linked to an increase of the average gas fractions in galaxies, as well as in the ratio of the molecular (H$_2$) over the atomic (HI) hydrogen fraction (see, e.g., Obreschkow \\& Rawlings 2009).  The presence of massive gas reservoirs seem to be a major requirement for justifying high SFRs, with the lack thereof possibly supporting instead the need for changes in the initial stellar mass function (IMF).  Recently, we reported the detection of large molecular gas reservoirs in two near-IR selected massive galaxies at $z\\sim1.5$ (Daddi et al.\\ 2008). Based on the knowledge available at the time, i.e., the previously known empirical  correlation between CO line luminosity ($L'_{\\rm CO}$) and total IR luminosity ($L_{\\rm IR}$) those galaxies should not have been detected in CO. However, observations were still attempted based on the results of Daddi et al. (2005; 2007) that suggested a high duty cycle and long duration for star formation in massive high-$z$ galaxies. The subsequent detections of CO[2-1] emission indeed implied low star formation efficiencies ($SFE$, defined as the ratio $L_{\\rm IR}/ L'_{\\rm CO}$). These SFEs are similar to what is seen in local spirals and much lower than previously seen in high redshift galaxies. Subsequent studies revealed low-excitation CO [3-2] to [2-1] line ratios in these galaxies (Dannerbauer et al. 2009) that are more typical of local spirals.  Given that both galaxies were detected in CO, this early result suggested commonplace large molecular gas reservoirs in distant massive galaxies (Daddi et al. 2008). However, improved statistics would clearly be needed in order to formally confirm the result and to estimate the variations in gas properties among the high-$z$ galaxy populations. Higher spatial resolution observations are further required to constrain the sizes of the molecular emission; i.e., given the low SFEs one would expect extended reservoirs rather than the compact emission lines seen in submm selected galaxies (SMGs) for which Tacconi et al. (2006; 2008) find typically half light radii of order 2~kpc.  In this paper we report on low-resolution (D-configuration) observations with the IRAM Plateau de Bure Interferometer (PdBI) of CO[2-1] (redshifted to 3mm) in four additional BzKs (near-IR selected galaxies at $z\\sim1.5$) thus tripling the number of sources in our sample. In addition, we report on higher resolution ($\\sim1$--1.5$''$; B-configuration) observations of four galaxies. We also discuss the properties of these six galaxies with respect to the `parent' BzK sample. We use the information gathered on these six galaxies in order to independently constrain the molecular gas mass and, conversely, the infamous CO luminosity to gas mass conversion factor $\\alpha_{CO}$.  The paper is organized as follows. In section~\\ref{sec:data} we present the new PdBI observations and their analysis. In Sect.~\\ref{sec:sample} we discuss the relation of the six galaxies to the parent sample and the derivation of accurate SFRs and stellar masses. The SFEs are characterized in Sect.~\\ref{sec:LCO} and compared to other galaxy populations. We introduce numerical models of clumpy disks in Sect.~\\ref{sec:models} and use the models to calibrate dynamical masses for our sample in Sect.~\\ref{sec:Mdyn}. This section also discusses implications for the implied gas masses and the conversion factors.  A discussion and summary of our results is presented in Sect.~\\ref{sec:discussion}~and~\\ref{sec:summary}. We apply a concordance WMAP3 cosmology throughout the paper. ",
        "conclusions": "\\subsection{Massive, extended reservoirs of low-excitation molecular gas}  CO[2-1] detections of two near-IR selected (BzK) galaxies at $z=1.5$ were presented in Daddi et al. (2008). We have now observed a total of six galaxies in CO[2-1] and detected all of them. This clearly supports the ubiquity of massive molecular gas reservoirs in these galaxies, a result confirmed by Tacconi et al. (2010).  In our observations we were able to spatially resolve robustly the CO[2-1] emission in this kind of galaxies and derive typical CO emission FWHM in the range of 6--11~kpc. This is a factor of 2--3 larger than what is found in SMGs (Tacconi et al. 2006; 2008; Bouch\\'e et al. 2008). The presence of extended gas disks is quite consistent with our previous finding based on the low CO excitation in these objects (Dannerbauer et al. 2009). These findings are indicative of colder and/or lower density gas than what is found in SMGs, a situation more similar to what is found in local spiral galaxies.  \\subsection{Normal star-forming galaxies at high redshift as extended, clumpy rotating disks}  Our sample is based on the BzK-selection technique and we have shown it to be representative of ``normal'', typical star-forming galaxies at high redshift. We now summarize the many pieces of evidence that these CO-detected BzK galaxies are large disk galaxies rather than on-going violent mergers.  First, we clearly detect double peaked profiles and the velocity width in the blue- and redshifted components is very similar. This would be a highly contrived coincidence if this was due to a merger: the components would need to have equal mass and be observed at similar inclinations.  Second, the UV rest-frame morphology disfavors on-going major mergers: in no case do we see evidence for a double-galaxy encounter.  The irregular and clumpy morphology could lead one to speculate that this is due to the presence of multiple minor merging events, but in all likelihood we are looking at clumps in single massive disks (Elmegreen et al. 2007, 2009a, Bournaud et al. 2008, Genzel et al. 2008).  Other evidence for rotating disks include the position of these sources in the velocity-size plane. Bouch\\'e et al. (2007) used H$\\alpha$ resolved spectroscopy of optical and near-IR selected galaxies and identified a dichotomy between SMGs and extended disks in such velocity-size diagrams: SMGs have much higher velocities and smaller sizes. Using CO velocity widths and CO sizes for both BzK galaxies and SMGs, we confirm a dichotomy between the two classes: in the notation of Bouch\\'e et al., our CO-detected galaxies have 'maximum rotational velocities' $v_d = v_{FWHM}/2.35 \\simlt200$~km~s$^{-1}$ and exponential disk scale lengths $r_d>3$~kpc, while the SMGs in Bouch\\'e et al. have $v_d>200$~km~s$^{-1}$ and $r_d<3$~kpc. The latter are better understood in terms of merging systems, while the former appear consistent with being non-interacting disk galaxies (or interacting only with low-mass satellites).   Further evidence that we are looking at disk galaxies comes from the CO properties. We have already discussed the large spatial sizes of the CO reservoirs, and the low gas excitation. Furthermore, the SFE in these galaxies is considerably lower than that of both SMGs and local ULIRGs, suggesting a different star formation {\\em mode}.  The gas depletion timescales are of order of 0.5~Gyr and imply a high duty-cycle (see also Daddi et al.\\ 2005; 2007; 2008) -- this is quite different from the timescales expected for a merger-induced event that would result in a $\\simlt100$Myr starburst (Mihos \\& Hernquist 1996; Greve et al. 2005; Di~Matteo et al. 2008). As can be seen from Fig.~\\ref{fig:both_plot_6CO}, SMGs typically also have depletion timescales of $<$100\\,Myr.  \\subsection{Origin of massive disks at high redshift}  The fact that we are presumably looking at non-interacting rotating disks instead of on-going mergers does not in itself imply that these massive disks were not assembled by past mergers. Remnants of major mergers can have massive rotating gas disks, especially when the gas fraction is high (e.g., Robertson \\& Bullock 2008). Numerical models predict that, under some physical assumptions, merger remnants can contain extended gas disks (e.g. Hopkins et al. 2009). However, Bournaud \\& Elmegreen (2009) argued that the overall spatial distribution of the mass in high-redshift disks is inconsistent with the majority of them being merger remnants, suggesting that their mass assembly was dominated by relatively smooth mass infall along gas streams (see also Dekel et al. 2009b). This is because  the morphology of these disk galaxies, in particular their supermassive star-forming clumps, cannot be accounted for if the main assembly channel is by mergers. Our CO observations add further support to this picture. Indeed, we find extended CO reservoirs with spatial sizes consistent with what is measured from their stellar light. The end product of galaxy mergers are characterized by highly concentrated molecular gas distributions, even when some molecular gas is found in the outer regions. This holds both for early-stage interactions (pre-merger galaxies like Arp~105 -- Duc et al. 1997), on-going mergers (like NGC~520, Yun \\& Hibbard 2001), and merger remnants (like NGC~7252, Hibbard et al. 1994). This dynamical effect results from the gravitational torquing of the gas, which is well reproduced by numerical simulations (Barnes \\& Hernquist 1991; Di~Matteo et al. 2008) and is independent of the gas fraction. If major mergers played a role in the formation of $z=1.5$ BzK galaxies, their large molecular disks must have been replenished by the smooth infall of baryons at a rate higher than the merger-driven mass assembly in order to keep extended gas disks. Cold flow-driven galaxy formation, indeed, does predict a large population of massive, extended star-forming disks at high redshift with high gas fractions (e.g., Dekel et al. 2009b).   \\subsection{Global properties of star formation at high redshift}  In a simple cartoon, the $z=1.5$ BzK galaxies appear to be disks like local spiral galaxies, but containing a substantially higher gas content and much larger star-forming clumps. A similar comparison holds for SMGs vs. local ULIRGs, with the former being the scaled--up version of the latter but having overall higher gas content (Tacconi et al. 2006). Most of the observed properties of the two classes can be understood within this scheme: there is a population of disk galaxies having low excitation CO, spatially extended gas reservoirs, low SFE (and hence long gas consumption timescales), low SSFR and optically thin UV emission. These are the familiar spirals in the local Universe, and likely most of the near-IR (or optical) selected disk galaxies in the distant Universe. Conversely, there are galaxy populations with opposite properties: high CO excitation, compact gas reservoirs, high SFE, high SSFR and with optically thick UV emission. These are the local ULIRGs and high redshift SMGs, both classes likely corresponding to major (wet) mergers of massive star forming galaxies.    \\subsection{New constraints for galaxy formation models and the IMF}  Based on our results of a remarkably low scatter in observed SFE and gas fractions for the CO-detected BzK galaxies, we suggest that the molecular gas content plays an important role in driving the tight stellar mass-SFR correlation in high redshift sources. A simple scenario may be that the gas properties are tightly connected to that of the underlying dark matter halo, possibly via the regulation of cold gas accretion rates through cosmic times (e.g., Keres et al. 2005; see also Bouch\\'e et al. 2010). This in turn could regulate the amount of star formation in each halo and its time integral, i.e. the stellar mass, which, as a result, are tightly connected. Ultimately, there might be a stellar mass-SFR correlation because the gas content of galaxies tightly correlates with the masses of hosting dark matter halos over cosmological timescales.  Galaxy formation models based on numerical hydrodynamic codes (e.g., Finlator et al. 2006) and semianalytical renditions (e.g., Kitzbichler \\& White 2007; see Fig.~17~and~18 in Daddi et al. 2007) had indeed predicted the existence of tight correlations between stellar mass and SFR before they were recovered by observations, with roughly a correct slope (or possibly slightly shallower, see, e.g., Daddi et al. 2007a). This is probably one of the rare cases where models of galaxy formation have preceded observations, showing truly predictive power.  Also, such models had qualitatively predicted an increase in the normalization of the correlation toward higher redshifts.  However, existing models fail in reproducing the actual magnitude of the normalization increase to $z\\sim2$ (falling short by factors of 3--5 compared to observational estimates; see e.g. Daddi et al. 2007a; Dav\\'e 2008). The SFRs estimated for distant galaxies are generally too high compared to the prediction of models, a fact that has been long known for SMGs (Bough et al.\\ 2005). Dav\\'e (2008) suggest that this discrepancy between models and observations can be solved by adopting substantial evolution in the stellar initial mass function (IMF). This conclusion was also reached by Baugh et al.\\ (2005) in their modeling of SMGs (i.e., for galaxies with extreme SFRs of typically 1000~$M_\\odot$~yr$^{-1}$ that have space densities at least an order of magnitude lower than that of near-IR selected massive galaxies at $z\\sim2$). However, new theoretical frameworks have recently been proposed in which feeding of star formation through cold flow filaments (Dekel et al. 2009a; see also Keres et al. 2005) can more easily justify high star formation rates in distant galaxies. In addition, we have also found large gas reservoirs, reaching $10^{11}M_\\odot$ in some cases, i.e. the gas accounts for a major fraction of the baryons in these sources. We are thus witnessing a phase in galaxy formation in which massive galaxies were truly gas--dominated systems.  These high gas masses can explain the high observed SFRs and one does not need to advocate non standard IMFs to reduce the SFR, at least for the normal, near-IR selected (BzK) galaxies. We note that quantitative predictions of the models by Oppenheimer \\& Dave (2006; 2008) on the gas fractions in galaxies at $z\\sim2$ (with similar stellar masses to our CO--detected galaxies) result in expected gas mass fractions of order of 15\\%, i.e., much smaller than what we derived. Therefore, our study suggests that the molecular gas fraction in galaxy simulations should be increased in order to reproduce the observed IR luminosities, without necessarily introducing a top heavy IMF.      We have presented results from a CO[2-1] survey of normal, BzK-selected star-forming galaxies at $z \\simeq 1.5$. The analysis of the integrated spectra and the spatially-resolved CO observations was made in combination with high-resolution optical imaging and multi-wavelength data. We also used appropriate numerical simulations of clumpy disk galaxies (BEE07) to derive dynamical masses and compare them with the observed CO luminosities. The main findings and implications can be summarized as follows:  \\begin{itemize}  \\item All observed galaxies were detected with high $S/N$ ratios. We detected extended CO reservoirs with typical sizes of $\\sim 6$--11~kpc, and presented evidence for rotation from spatially-resolved CO observations. This is indicative of the molecular gas being situated in disks that are similar in size and shape to the UV rest-frame morphology.  \\item We have provided direct evidence for a high $\\alpha_{\\rm CO}$ conversion factor (relating CO luminosity to gas mass), based on the dynamical masses of the systems. Our results show that the conversion factor of  BzK galaxies is high, similar to that of local spirals. Our derived $\\alpha_{\\rm CO}$ is a factor of 4--5 larger than what has been derived for local ULIRGs and typically used for high redshift SMGs.  \\item Our observations imply high gas fractions, $\\simeq 50$--65\\% of the baryons in our $z \\simeq 1.5$ galaxy sample, and quiescent star formation with relatively low SFE and long gas consumption timescales ($\\sim$ 0.5~Gyr).  \\item Our results also point to a strong correlation between the molecular gas content and other physical properties of these sources. Once the stellar mass and SFR of a given galaxy are known, the CO luminosities can be predicted to within 40\\% rms. This and the overall similar SFEs seen in these galaxies and local spirals leads us to propose a relation between CO luminosity and bolometric luminosity that could allow one to systematically estimate the CO luminosity of massive disk galaxies, presumably at any redshift $z\\simlt2$ as the relation can fit local and $z=1.5$ observations with an overall scatter of only 0.25~dex.   \\item Our sample consists of normal star-forming galaxies at high redshift, and we have shown that it is representative of the majority of $z>1$ near-IR selected star-forming galaxies. Our results shed more light on the nature of these high-redshift disk galaxies. In particular, the large spatial size of the molecular gas reservoirs suggest that they did not assemble mostly through past episodes of violent mergers. Instead, our observations are more consistent with smooth mass infall along cosmic flows and subsequent internal evolution as the main drivers of disk galaxy formation and star formation at high-redshift.  \\end{itemize}  Based on observations with the IRAM Plateau de Bure Interferometer. IRAM is supported by INSU/CNRS (France), MPG (Germany) and IGN (Spain). We acknowledge the use of GILDAS software (http://www.iram.fr/IRAMFR/GILDAS).  We are grateful to Len Cowie for help in cross-checking the spectroscopic redshift of BzK-12591 from Keck observations, prior to our CO follow-up.  We thank Claudia Maraston for providing Chabrier IMF models from her library of synthetic stellar populations templates. We thank Padeli Papadopoulos for comments and the anonymous referee for a constructive report. This research was supported by the ERC-StG grant UPGAL 240039.  We acknowledge the funding support of French ANR under contracts ANR-07-BLAN-0228, ANR-08-JCJC-0008 and ANR-08-BLAN-0274. Support for this work was provided by NASA, Contract Number 1224666 issued by the JPL, Caltech, under NASA contract 1407. DR acknowledges support from from NASA through Hubble Fellowship grant HST-HF-51235.01 awarded by the Space Telescope Science Institute, which is operated by the Association of Universities for Research in Astronomy, Inc., for NASA, under contract NAS 5-26555. The work of DS was carried out at Jet Propulsion Laboratory, California Institute of Technology, under a contract with NASA. The simulations were performed using HPC resources from GENCI-CCRT (Grant 2009-042192)."
    },
    "0911/0911.0897_arXiv.txt": {
        "abstract": "We present the analysis of an observation by \\xmm\\ that exhibits strongly variable, low-energy diffuse X-ray line emission. We reason that this emission is due to localised solar wind charge exchange (SWCX), originating from a passing cloud of plasma associated with a Coronal Mass Ejection (CME) interacting with neutrals in the Earth's exosphere. This case of SWCX exhibits a much richer emission line spectrum in comparison with previous examples of geocoronal SWCX or in interplanetary space. We show that emission from O\\,VIII is very prominent in the SWCX spectrum. The observed flux from oxygen ions of $\\rm 18.9\\,keV\\,cm^{-2}\\,s^{-1}\\,sr^{-1}$ is consistent with SWCX resulting from a passing CME. Highly ionised silicon is also observed in the spectrum, and the presence of highly charged iron is an additional spectral indicator that we are observing emission from a CME. We argue that this is the same event detected by the solar wind monitors \\ace\\ and \\wind\\ which measured an intense increase in the solar wind flux due to a CME that had been released from the Sun two days previous to the \\xmm\\ observation. ",
        "introduction": "We report on a case of extreme solar wind charge exchange (SWCX) detected by X-ray signatures within an observation taken with the \\xmm\\ telescope \\citep{jansen}. SWCX occurs when an ion in the solar wind interacts with a neutral atom and the ion gains an electron into an excited state. If the ion is in a sufficiently high charge state, single or multiple X-ray or UV photons are emitted in the subsequent relaxation of the electron. In the case of geocoronal SWCX, i.e. in the vicinity of the Earth, the neutrals involved in the charge exchange interaction are exospheric hydrogen. A SWCX spectrum is characterised by emission lines from ions such as C, N, O, Mg and Fe, without a continuum component.  For the general user of the \\xmm\\ observatory, the short-term variations from geocoronal SWCX \\citep{snowden2004, carter2008}, or the longer term variations that occur from SWCX interactions in interplanetary space \\citep{smith2005} or at the heliosheath \\citep{cravens2000, koutroumpa2007}, can produce a significant diffuse signal below $\\sim 2$\\,keV. This signal may constitute an additional source of background when studying astrophysical sources, which must be characterised before being removed from an astrophysical data set. However, geocoronal SWCX provides important information concerning the constituents of the solar wind, supplementing or contributing additional information to that received by upstream solar wind monitors. Solar wind compositional signatures such as abundances and abundance ratios are indicative of the conditions prevailing near the Sun during the acceleration of the solar wind, a process that is not fully understood \\citep[and references therein]{richardson2004}. Geocoronal SWCX has been proposed as a mechanism to image the magnetosheath around the Earth \\citep{collier2005, collierlxo}, leading to an increase in understanding of transport processes within the plasma in and near the bow shock. This case of SWCX therefore holds interest for both the astronomical and solar-terrestrial communities.  The \\xmm\\ observation under study in this paper was made on the $21^{st}$ of October, 2001. We present the analysis of this observation and reason that the unusual X-ray signatures seen are due to a Coronal Mass Ejection (CME) that was recorded on $19^{th}$ of October 2001 by the Solar and Heliospheric Observatory (SOHO) \\citep{domingo1995} and which subsequently passed by the Earth. CMEs involve an ejection of high density plasma with characteristics different to that of the ambient solar wind; for example unusually high Fe charge states or enhanced alpha particle to proton ratios \\citep{zurbuchen}. CMEs may pass by the Earth, depending on their location of origin in the solar corona and passage through interplanetary space. The absolute frequency of CMEs increases around solar maximum, although at solar minimum, CMEs occur approximately weekly. The event under analysis in this paper occurred close to solar maximum in 2001. We use additional data from both the Advanced Composition Explorer (\\ace, \\citet{stone1998}) and the \\wind\\ \\citep{acuna1995} spacecraft to support our argument. We also analyse \\xmm\\ observations before and after our case observation and the nearest \\chandra\\ observation in time to this period.  The $21^{st}$ of October 2001 event was first identified as a particularly noteworthy observation during a systematic search for SWCX within the \\xmm\\ public archive as described in \\citet{carter2008}. That study performed an analysis of 187 observations of the EPIC-MOS instruments \\citep{turnershort} taken in full-frame imaging mode to search for variable diffuse emission in an energy band concentrating on the SWCX indicators of O\\,VII and O\\,VIII emission between 0.5\\,keV and 0.7\\,keV.  This observation proved to be the most spectrally rich example of SWCX found within the sample, and its possible association with a known CME warranted a more detailed study. In fact, as we shall show, the dominant diffuse component in the entire observation was due to X-rays from SWCX. In this paper we perform a more detailed spectral analysis than reported in the survey paper of \\citet{carter2008}. We extend our analysis to include the EPIC-pn which was also in full frame mode for this observation.  The paper is organised as follows. In section \\ref{secmulti} we describe several relevant \\xmm\\ observations of the same target as the $21^{st}$ of October event and a \\chandra\\ observation taken around the same time. Section \\ref{secnewdata} contains a spectral analysis of the \\xmm\\ data and in Section \\ref{secview} we discuss the viewing geometry and orientation of \\xmm\\ during the observation. We end the paper with Section \\ref{secdiscuss} containing our discussions and conclusions. ",
        "conclusions": "\\label{secdiscuss} We consider the possibility that the SWCX enhancement of \\ourc\\ is linked to the CME event of the 19th October 2001 \\citep{wang2005}. The delay between the occurrence of the CME at the solar corona and its arrival near Earth would be approximately two and a half days. Increased solar proton fluxes were registered by both \\ace\\ and \\wind\\ and therefore this plasma cloud would have passed in the immediate vicinity of the Earth. It is not always the case that enhancements in solar proton fluxes, and any accompanying highly charged ions, are registered by increased incidents of soft proton flaring or SWCX enhancements by \\xmm. However, the arrival of the peak of the low energy enhancement as seen by \\xmm\\ is consistent with the delay expected as the feature passes in sequence from \\ace\\ to \\wind, on to a region intersected by the line of sight of \\xmm.  We have shown that line emission from O\\,VIII is very prominent and dominates that of O\\,VII, contrary to signatures of heliospheric SWCX \\citep{koutroumpa2009a}. We have also shown in our spectral analysis, that the observed flux from SWCX emission is much greater than that from a simple estimate of the expected emission, based on the abundances of a slow solar wind. Also, mid-energy emission lines in the regime 0.70\\,keV to 2.00\\,keV infer the presence of highly charged states of iron, as is often seen in a CME \\citep{zurbuchen, zhao}. We see no significant compositional changes in the line emission over the duration of the \\xmm\\ observation. In addition, we have observed emission at 2.00\\,keV from highly charged states of silicon, implying a very high temperature plasma. A CME, rather than a steady state solar wind, would explain the large enhancements, flux observed and the richness of the spectrum as seen by \\xmm. This case is the richest spectrally of those examined by \\citet{carter2008}.  CMEs have been used to explain the results of other X-ray observations in the literature pertaining to the diffuse X-ray background. \\citet{henley2008} invoked a CME to explain differences between results obtained from \\xmm\\ and \\suzaku, when determining Halo and Local Bubble X-ray spectra. They also observed emission from Mg\\,XI and Ne\\,IX, although emission lines from oxygen were less significant. They attribute this emission to a possible localised enhancement in solar wind density crossing the line of sight of \\xmm. \\citet{smith2005} attributed the anomalously high level of O\\,VIII seen in their observation of a nearby molecular cloud to SWCX, and noted this was unlikely to be due to SWCX from a steady state solar wind. Instead they conclude that their enhancement was due to charge exchange from a CME and the interstellar medium, probably at a distance of a few AU from the Sun, due to the depletion of neutral gas available for charge exchange near the Sun. We eliminate the possibility that the emission seen in \\ourc\\ is due to SWCX occurring at the heliospheric boundary or at a large distance from Earth. Short-term variations can occur for heliospheric SWCX, especially if observing along the helium focusing cone \\citep{robertson2003a, robertson2003b}, but the pointing of \\xmm\\, which does not insect the region of peak emission from this area, argues against this case. In addition, the abundant emission line spectrum and the variations in the fluxes of the major ions in the spectrum which reflect the variations in solar proton flux support a geocoronal occurrence of SWCX. We conclude that the SWCX interaction we have observed occurs between ions from a CME and neutrals in the exosphere of the Earth, at a relatively close distance to the Earth, but not confined to the magnetosheath within the bow shock.  Although data regarding the ion states of the solar wind for the period of \\ourc\\ from the solar wind monitors \\ace\\ and \\wind\\ are sparse, we have been able to identify ions from a rich set of emission lines from a passing CME. Not all CMEs detected by \\ace\\ will be detected by \\wind, or indeed intersect the line of sight of \\xmm. \\xmm\\ was not optimised to study the magnetosheath or near Earth regions. However, we have shown that \\xmm\\ can be used to provide additional compositional information of the solar wind plasma, especially for the highest charge state ions, to that obtained by upstream solar wind monitors, providing the observing geometries and inclinations of the incoming wave fronts are favourable."
    },
    "0911/0911.3150_arXiv.txt": {
        "abstract": "We present broadband (gamma-ray, X-ray, near-infrared, optical, and radio) observations of the gamma-ray burst (GRB) 090709A and its afterglow in an effort to ascertain the origin of this high-energy transient.  Previous analyses suggested that GRB\\,090709A exhibited quasi-periodic oscillations with a period of 8.06\\,s, a trait unknown in long-duration GRBs but typical of flares from soft gamma-ray repeaters.  When properly accounting for the underlying shape of the power-density spectrum of GRB\\,090709A, we find no conclusive ($> 3\\sigma$) evidence for the reported periodicity.  In conjunction with the location of the transient (far from the Galactic plane and absent any nearby host galaxy in the local universe) and the evidence for extinction in excess of the Galactic value, we consider a magnetar origin relatively unlikely.  A long-duration GRB, however, can account for the majority of the observed properties of this source.  GRB\\,090709A is distinguished from other long-duration GRBs primarily by the large amount of obscuration from its host galaxy ($A_{K,\\mathrm{obs}} \\gtrsim 2$\\,mag). ",
        "introduction": "\\label{sec:intro} A variety of astrophysical sources are capable of producing short, intense flashes of high-energy emission detectable by the current generation of gamma-ray satellites.  These sources span an incredible range of the observable universe, from electrical discharges associated with thunderstorms on Earth \\citep{fbm+94} to the deaths of the earliest known stars in the universe \\citep{tfl+09,sdc+09}.  Gamma-ray bursts (GRBs) are the most luminous class of these high-energy transients ($L_{\\mathrm{iso}} \\approx 10^{50}$--$10^{52}$\\,erg\\,s$^{-1}$). At least two distinct progenitor systems are thought to produce GRBs \\citep{kmf+93}.  It is now widely accepted that most GRBs with duration $t_{90} \\gtrsim 2$\\,s arise as a byproduct of the core collapse of massive stars (hereafter referred to as ``long-duration'' GRBs; e.g., \\citealt{wb06} and references therein). The origin of ``short-duration'' GRBs is still a hotly debated topic. They likely arise from an older stellar population (e.g., \\citealt{gso+05,hsg+05,bcb+05,cdg+05,bpc+05,bpp+06,gcg+06}), possibly due to the merger of a neutron star-neutron star (NS-NS) or black hole-neutron star (BH-NS) binary \\citep{elp+89,npp92}.  Soft gamma-ray repeaters (SGRs), on the other hand, are distinguished from GRBs by repeated outbursts with isotropic energy release of $E_{\\gamma,\\mathrm{iso}} \\lesssim 10^{46}$\\,erg (e.g., \\citealt{wt06,o07,m08}). The discovery of periodic oscillations from bright SGR flares \\citep{mgi+79}, along with the measurement of a spindown in their periods \\citep{kds+98}, allows for an estimation of the magnetic field strength.  Unlike typical rotation-powered radio pulsars, SGRs (as well as their counterparts discovered in quiescence, the anomalous X-ray pulsars, or AXPs) are likely powered by their intense magnetic fields ($B \\gtrsim 10^{14}$\\,G; \\citealt{dt92}).  Together, SGRs and AXPs are now thought to comprise a single class of young, highly magnetized neutron stars (``magnetars'').  At least three discoveries have in recent years challenged this simple classification picture.  First, both GRB\\,060614 and GRB\\,060505 had high-energy durations in excess of a few seconds. Yet despite being quite nearby ($z \\approx 0.1$), neither exhibited evidence for an associated supernova to quite deep limits \\citep{gnb+06,gfp+06,dcp+06,fwt+06,ocg+07}.  Second, many GRBs that would have been classified as having short duration ($t_{90} \\lesssim 2$\\,s) by the less sensitive \\textit{BATSE} satellite exhibit soft (and oftentimes faint) X-ray tails extending hundreds of seconds in time when observed by \\swift\\ (e.g., \\citealt{bcb+05,pmg+09}) -- perceived duration is, after all, at once a redshift-, sensitivity-, and bandpass-dependent quantity.  Finally, the X-ray and optical afterglow light curves of GRB\\,070610 ($t_{90} = 5$\\,s) displayed dramatic flares on time scales as short as several seconds \\citep{kck+08,sks+08,cdg+08}. Though almost certainly of Galactic origin, the source of this variability is still poorly understood.  Here we present observations of GRB\\,090709A, a high-energy transient whose classification as a traditional long-duration GRB has been called into question.  \\citet{GCN.9645} reported the detection of quasi-periodic oscillations (period $P = 8.06$\\,s) in the high-energy light curve, typical of the observed properties of SGRs and AXPs.  However, unlike all currently known magnetars, the localization of GRB\\,090709A is inconsistent with both the Galactic plane and any nearby galaxy.  Through a detailed analysis of the high-energy light curve, the broadband afterglow, and the immediate environment, we attempt to uncover the origin of GRB\\,090709A.  Throughout this work we adopt a standard $\\Lambda$CDM cosmology ($\\Omega_{\\Lambda} = 0.73$; $\\Omega_{m} = 0.27$; $H_{0} = 71$\\,km\\,s$^{-1}$\\,Mpc$^{-1}$). All quoted uncertainties are 1$\\sigma$ (i.e., 68\\%) confidence intervals unless otherwise stated.  Spectral and temporal power-law indices are provided using the convention $f_{\\nu} \\propto t^{-\\alpha} \\nu^{-\\beta}$ \\citep{spn98}.  UT dates are used throughout this work.  In the final stages of preparing this manuscript, \\citet{dei+09} posted an analysis of X-ray (\\textit{XMM-Newton} and \\textit{Swift}-XRT) and gamma-ray (\\textit{Swift}-BAT and \\textit{Integral}-SPI/ACS) observations of GRB\\,090709A.  These authors reach largely similar conclusions regarding the origin of GRB\\,090709A, although they favor a somewhat larger distance for the event.  We attempt to highlight both the differences and similarities between our work in what follows. ",
        "conclusions": ""
    },
    "0911/0911.1363_arXiv.txt": {
        "abstract": "We report on nine wide common proper motion systems containing late-type M, L, or T companions.  We confirm six previously reported companions, and identify three new systems.  The ages of these systems are determined using diagnostics for both stellar primaries and low--mass secondaries and masses for the secondaries are inferred using evolutionary models.  Of our three new discoveries, the M3+T6.5 pair G 204-39 and SDSS J1758+4633 has an age constrained to 0.5-1.5 Gyr making the secondary a potentially useful brown dwarf benchmark. The G5+L4 pair G 200-28 and SDSS J1416+5006 has a projected separation of $\\sim$25,000 AU making it one of the widest and lowest binding energy systems known to date.  The system containing NLTT 2274 and SDSS J0041+1341 is an older M4+L0 ($>$4.5 Gyr) pair which shows H$\\alpha$ activity in the secondary but not the primary making it a useful tracer of age/mass/activity trends. Two of the nine systems have discrepant component ages which emerge from stellar or ultracool diagnostics indicating possible shortcomings in our understanding of the age diagnostics of stars and brown dwarfs. We find a resolved binary frequency for widely-separated ($>$ 100 AU) low--mass companions (i.e. at least a triple system)  which is at least twice the frequency found for the field ultracool dwarf population.  The ratio of triples to binaries and quadruples to binaries is also high for this sample: 3:5 and 1:4, respectively, compared to 8-parsec sample values of 1:4 and 1:26.  The additional components in these wide companion systems indicates a formation mechanism that requires a third or fourth component to maintain gravitational stability or facilitate the exchange of angular momentum.  The binding energies for the nine multiples discussed in this text are among the lowest known for wide low-mass systems, suggesting that weakly bound, low--to--intermediate mass (0.2M$_{\\sun}$ $<$ M$_{tot}$ $<$1.0M$_{\\sun}$) multiples can form and survive to exist in the field (1-8 Gyr). ",
        "introduction": "Ultracool dwarfs (UCDs) comprise the late-type M, L, and T dwarf spectral classifications  \\citep[e.g.,][and references therein]{2005ARA&A..43..195K} and include brown dwarfs--objects that do not support stable hydrogen fusion  (\\citealt{1962AJ.....67S.579K}; \\citealt{1963PThPh..30..460H}).  UCDs sample the low-mass extremum of star formation processes and are abundant in nearly every Galactic environment.  The low temperatures and high pressures in the photospheres of UCDs give rise to abundant molecular species, whose complex chemistry and opacities result in highly-structured spectral energy distributions (SEDs).  Disentangling the physical characteristics---mass, age, surface gravity, metallicity, atmospheric properties, etc.--- that modulate these spectral features is a critical step for testing theoretical models.  However, individual characterization of Galactic brown dwarfs is challenging because their thermal evolution leads to a degeneracy between mass, age, and physical properties derived from observables such as luminosity and effective temperature (T$_{eff}$).  While spectral analyses can constrain physical properties for some systems (e.g. \\citealt{2006ApJ...639.1095B}; \\citealt{2007ApJ...656.1136S}; \\citealt{2007MNRAS.381.1400W}; \\citealt{2008ApJ...678.1372C}), calibration of these techniques requires detailed studies of well-understood benchmark systems.  One useful group of UCD benchmarks are those which are resolved companions to nearby, well-characterized stars.  Assuming coevality, the physical properties of the primary, such as metallicity and age---which are extremely difficult to measure for low-mass stars---can be applied to the companion.  In particular, independent age determinations are critical to break the mass/age/observable degeneracy for the brown dwarf companion.  Despite the apparent scarcity of wide UCD companions to nearby stars  ($\\sim$2-3\\%; \\citealt{2001ApJ...551L.163G}, \\citealt{2008ApJ...683..844L}), several have been identified and used to calibrate spectral analysis techniques (e.g., \\citealt{2006ApJ...639.1095B}; \\citealt{2007ApJ...656.1136S}), as well as to critically test atmospheric (e.g., \\citealt{2008ApJ...682.1256L}) and structure/evolutionary models (e.g.,  \\citealt{2004ApJ...609..885M};  \\citealt{2008arXiv0807.2450D}). The frequency and characteristics of widely-separated stellar-UCD pairs also puts important constraints on the star formation processes and the subsequent dynamical evolution of stellar systems (e.g., \\citealt{2003ApJ...586..512B}; \\citealt{2003ApJ...587..407C,2007ApJ...660.1492C}; \\citealt{2007ApJ...668..492A}; \\citealt{2009ApJ...691.1265L}).  However, the known population of UCD companions remains small and does not yet fully sample the range of ages, masses and metallicities found among unassociated field sources.   In the past decade, multiplicity surveys focused on the field UCD population have distinguished two classes: \\begin{itemize} \\item Roughly 10-20\\% of the field UCDs are found to be closely-separated ($\\rho$ $<$ 20 AU), near-equal mass, small total mass (M$_{tot}<$0.2) UCD-UCD multiples (e.g. \\citealt{2003AJ....126.1526B}, \\citealt{2003ApJ...587..407C}, \\citealt{2003ApJ...586..512B}, \\citealt{2007ApJ...671.2074A}, \\citealt{2008AJ....135..580R}) \\item A smaller fraction are found to be widely-separated ($\\rho$ $>$ 100AU) from a much more massive stellar companion (e.g. \\citealt{2001AJ....121.3235K}, \\citealt{2001AJ....122.1989W}, \\citealt{2005ApJ...625..385A}) \\end{itemize}  In the first case, the typically tight separations for UCD binaries is well-established (e.g. Allen et. al\\ 2007), and early studies by  \\citet{2003ApJ...586..512B}, and \\citet{2003ApJ...587..407C} identified a maximum separation limit/minimum binding energy for field UCD-UCD pairs of E$_{b} \\sim$ 2$\\times$10$^{42}$ erg. However, the recent discovery of a number of young UCD systems (ages $<$10 Myr) and a handful of field systems that are more widely-separated ($\\rho$ $>$ 100 AU) and more weakly bound (E$_{b}$ $<<$ 10$^{42}$ erg), questions whether separation limits can be considered constraints for formation models or if wide UCD binaries are a normal (albeit rare) mechanism of UCD formation (\\citealt{2005ApJ...633..452K, 2006ApJ...649..306K}, \\citealt{2007ApJ...663..394K}, \\citealt{2009ApJ...691.1265L}, and \\citealt{2009arXiv0902.4742A}, \\citealt{2007ApJ...659L..49A}, \\citealt{2005A&A...440L..55B}, \\citealt{2005A&A...439L..19P}, \\citealt{2007ApJ...667..520C}, \\citealt{2009arXiv0903.3251R}).  In contrast, systems in the second category ($\\rho$ $>>$ 100 AU) have binding energies that are several orders of magnitude smaller than the minimum set for UCD-UCD pairs. \\citet{2005AJ....129.2849B} noted a higher binary frequency among UCDs that are widely-separated from a stellar primary, suggesting the need for higher masses or an angular momentum sink in multibody interactions to form these systems. Recent work by \\citet{2006A&A...458..817W} suggests that for a low-mass primary fragment formed in the cooler outer parts of a circumstellar disk ($\\rho$ $>$ 100 AU), and spinning at a fast enough rate, H$_{2}$ dissociation is likely to trigger a Secondary Fragmentation phase, thereby potentially giving rise to a closely-separated (a$\\sim$5 AU (m$_{system}$/0.1M$_{\\sun}$) UCD binary.  Current observational evidence suggests that widely-separated stellar companions exist out to distances of $\\sim$0.1pc (\\citealt{1984ApJ...281L..41L}; \\citealt{1987ApJ...312..367W}).  Beyond this separation, perturbations from passing stars and giant molecular clouds will likely disrupt the companions over the lifetime of the Galaxy.   Separations of stellar-UCD and, especially, UCD-UCD multiples appear to fall well below the perturbation limit, suggesting dynamical sculpting occurs only in the natal environment (\\citealt{2003ApJ...586..512B}; \\citealt{2007ApJ...660.1492C}). However, the current sample of such systems is far from complete.  In large part this is due to the challenge of covering a large area of the sky, and ascertaining evidence for companionship between two objects.  For stars, common proper motions have been the standard characteristic for identifying co-moving objects at large angular separations (\\cite{1961AJ.....66..528V,1944AJ.....51...61V}; \\citealt{1979lccs.book.....L}; \\citealt{2002AJ....124.1190L}).  Historically, optical proper motion catalogs lacked the depth to detect late-type M, L, and T dwarfs. In addition, the recent discovery of UCDs has largely precluded astrometric measurements due to short temporal baselines, making an extensive common proper motion search difficult. In the past few years, large UCD proper motion samples (e.g., \\citealt{2008MNRAS.384.1399J}; \\citealt{2008MNRAS.390.1517C}; \\citealt{2009AJ....137....1F}) and near-IR proper motion surveys have become available (e.g., \\citealt{2005A&A...435..363D,2007A&A...468..163D,2009MNRAS.394..857D}), making it possible to perform a search in the reverse direction: using the UCD proper motion to find a stellar companion.  In this study, we used a proper motion catalog of UCDs from \\citet{2009AJ....137....1F} (hereafter, the BDKP catalog) to conduct a common proper motion search for main sequence companions to Hipparcos  (\\citealt{1997A&A...323L..49P}) or LSPM-N (\\citealt{2002AJ....124.1190L}) catalog stars.  We have uncovered nine systems, six of which have been briefly noted in the literature and three of which are presented here for the first time.   In section 2 we discuss our target list, the criteria for companionship and the reliability of our matches.  In section 3 we discuss follow-up photometry as well as optical and near-IR spectroscopy of our candidate systems.  In section 4 we apply age diagnostic tests to the primaries and secondaries and calculate masses of the UCD secondaries. In section 5 we explore the stability of the nine systems as well as multiplicity and formation mechanisms for a large sample of UCD field companions.  We summarize our results in section 6. ",
        "conclusions": "We have provided a detailed analysis of nine wide companion systems containing UCDs. Seven of the systems have parallax measurements, six of which are precise Hipparcos measurements.  Combining catalog information with new spectroscopic observations of the primary and secondary components,  a best age range for each system was determined.  Assuming co-evality with the secondaries and combining best age ranges with bolometric luminosity ranges, masses were estimated from the \\citet{1997ASPC..119....9B} evolutionary models.  Seven of the UCDs were determined to be very low--mass stars, one was determined to be substellar, and one has a questionable age therefore undetermined mass.  Two of the nine systems, G171-58 with 2MASS J0025+4759 and G 63-23 with 2MASS J1320+0957,  have significant differences in the component system ages indicating possible shortcomings in our understanding of the age diagnostics of stars and ultracool dwarfs.  Using a compiled list of known wide companion systems containing a UCD, we find that the frequency of tight resolved binaries is at least twice as high for wide companion UCDs as for isolated field equivalents.   The ratio of triples to binaries is 3:5 and quadruples to binaries is 1:4 for wide companion systems with resolved UCD secondaries versus 1:4 and 1:26 for the 8-parsec sample.  The higher frequency of higher order multiples suggests that a third or fourth component may be required to maintain gravitational stability or to facilitate the exchange of angular momentum in these loosely bound systems.  The Jeans criterion was investigated against a large sample of companion systems and we conclude that using the Jeans length to set the separation scale is sufficient for constraining the lowest binding energy UCD companion systems down to M$_{tot}\\sim$0.2M$_{\\sun}$.  However, the tight separation of the closely bound, near equal-mass UCD systems is not explained by the allowed envelope set by the Jeans length. The distinguishing characteristics of objects now known at varying mass ratios, total masses, separations, and ages suggests that more specific predictions from relevant theories will help distinguish the dominant formation mechanism for the UCD population."
    },
    "0911/0911.2795.txt": {
        "abstract": "The effect  of a magnetic field on the linear phase of the advective-acoustic instability is investigated, as a first step toward a magnetohydrodynamic (MHD) theory of the stationary accretion shock instability taking place during stellar core collapse. We study a toy model where the flow behind a planar stationary accretion shock is adiabatically decelerated by an external potential. Two magnetic field geometries are considered: parallel or perpendicular to the shock. The entropy\u00d0vorticity wave, which is simply advected in the unmagnetized limit, separates into five different waves: the entropy perturbations are advected, while the vorticity can propagate along the field lines through two Alfv\\'en waves and two slow magnetosonic waves. The two cycles existing in the unmagnetized limit, advective\u00d0acoustic and purely acoustic, are replaced by up to six distinct MHD cycles. The phase differences among the cycles play an important role in determining the total cycle efficiency and hence the growth rate. Oscillations in the growth rate as a function of the magnetic field strength are due to this varying phase shift. A vertical magnetic field hardly affects the cycle efficiency in the regime of super-Alfv\\'enic accretion that is considered. In contrast, we find that a horizontal magnetic field strongly increases the efficiencies of the vorticity cycles that bend the field lines, resulting in a significant increase of the growth rate if the different cycles are in phase. These magnetic effects are significant for large-scale modes if the Alfv\\'en velocity is a sizable fraction of the flow velocity. ",
        "introduction": "The delayed energy deposition by neutrinos may be a viable explosion mechanism of massive stars \\citep{bethe85} if multidimensional hydrodynamical instabilities are taken into account during the first second after shock bounce, in the inner 200km of the massive star \\citep{buras06b,marek09}. By capturing some gas in large scale convective cells, these instabilities increase the exposure time of this gas to the heating by neutrinos \\citep{marek09,murphy08b,fernandez09b}. The $l=1,2$ deformation of the shock seems to be mainly due to the Standing Accretion Shock Instability (SASI), discovered by \\cite{blondin03}. This asymmetric instability may influence the kick and spin of the residual pulsar \\citep{scheck06,blondin07a}, trigger the excitation of g-modes \\citep{burrows06} and the emission of gravitational waves \\citep{marek09b,ott09}. These potential consequences of SASI are based on numerical simulations where the magnetic field is neglected, because the magnetic pressure is thought to be small compared to the thermal pressure in a typical core collapse. Other models of core collapse supernovae, in which magnetic forces play a dynamical role, usually rely on a fast rotation rate from which the magnetic energy is extracted by field winding and local instabilities \\citep{akiyama03,moiseenko06,shibata06,burrows07b,obergaulinger09}. Surprisingly, \\cite{endeve08} argued that the turbulent flow associated with SASI may be able to amplify an initially weak magnetic field to a dynamically significant strength of $ 10^{15} $G, even if the progenitor is not rotating. This latter result, if confirmed, shows the importance of understanding SASI in the context of MHD even if the progenitor is weakly magnetized.  As a first step in this direction, we investigate the linear phase of SASI in a moderate magnetic field. The linear mechanism of SASI has been identified as an advective-acoustic cycle by \\cite{blondin03,ohnishi06,foglizzo07,scheck08,fernandez09a}. The interplay of acoustic waves and advected perturbations of entropy/vorticity has been illustrated using a simple toy model by \\cite{foglizzo09} and \\cite{sato09}, in which the flow is planar and adiabatic. A similar cycle is expected to exist in a MHD flow, but should be modified by the magnetic field depending on its strength and direction, and could potentially involve entropy waves, fast and slow magnetosonic waves, and Alfv\\'en waves. The local stability of an isolated fast MHD shock has been established by \\cite{gardner64}, and its interaction with MHD waves has been characterized by \\cite{mckenzie70}. However, the global stability of the cycles resulting from the linear coupling of these MHD waves through a region of deceleration below the shock has never been studied. For this purpose we extend the toy model of \\cite{foglizzo09} to a magnetized fluid, and investigate two possible orientations of the magnetic field.  It should be noted that the existence of SASI type instabilities in a magnetized flow has been recently proposed by \\cite{camus09} to explain the oscillations of the termination shock in their MHD simulations of the relativistic pulsar wind in the Crab nebula. The present study can thus be considered more generally as a first step toward understanding the MHD extension of the classical advective-acoustic cycle.  In Sect.~2 we describe the stationary flow of our MHD toy model and define the method of our linear analysis. The efficiencies of all existing cycles are calculated in Sect.~3, and the growth rate of the resulting eigenmode is computed in Sect.~4. The results are summarized and discussed in Sect.~5. Most algebraic derivations have been separated from the main text for the sake of clarity, and can be found in the Appendices. ",
        "conclusions": "The main results of this study can be summarized as follows: \\begin{enumerate} \\item In the presence of a magnetic field, the advective-acoustic cycle splits in up to 5 different cycles. The acoustic wave becomes a fast magnetosonic wave, while the entropy-vorticity splits into an entropy wave, 2 Alfv\\'en waves and 2 slow magnetosonic waves. \\item The acoustic cycle becomes a fast magnetosonic cycle. Its efficiency is not significantly affected by the presence of a magnetic field in any of the configurations considered in this paper. \\item The propagation of the vorticity through slow and Alfv\\'en waves leads to a phase difference between the different cycles, which interact either constructively or destructively depending on the mode considered. The consequence of these interferences is a more irregular eigenspectrum. \\item In the superAlfv\\'enic regime that we investigate, a vertical magnetic field hardly changes the coupling efficiencies in the compact toy model. It only affects the cutoff associated with the size $H_{\\nabla}$ of the deceleration region by changing the vertical structure of the slow waves. \\item The vortical cycles are strongly amplified by a horizontal magnetic field if the field lines are bent ($ \\bf{k}.\\bf{B} \\neq 0 $). Both the Alfv\\'en cycles and the slow magnetosonic cycles are amplified. This amplification of the vortical cycles leads to faster growth rates of the modes with $n_x \\neq 0$. \\end{enumerate}  These results confirm that the mechanism of the advective-acoustic cycle can be generalized to a magnetized environment, such as core collapse supernovae or possibly the termination shock of pulsar wind nebulae. In order to estimate the significance of these results for SASI during core collapse, we can estimate the magnetic field strength needed to affect the advective-acoustic cycle. The desynchronization effect studied in Sects.~\\ref{Bvert_mode} and \\ref{Bhor_mode} requires $v_{\\rm Ash} \\sim v_{\\rm sh}$ for low frequency, low $n_x$ modes. Eq.~(\\ref{Bhor_vAmin}) can be extrapolated to a spherical core collapse as follows: \\begin{equation} \\frac{v_{\\rm A}}{v} \\sim  \\frac{R_{\\rm sh}}{2 \\sqrt{l(l+1)} H}. \\label{sphere_vAmin} \\end{equation} For the dominant $l=1-2$ modes, with $R_{\\rm sh} = 110{\\rm km}$ and $H=50{\\rm km}$, the desynchronisation of the cycles is expected for $v_{\\rm A}/v \\sim 0.4-0.8$ .  The magnetic field strength at which the amplification of the vortical cycles becomes significant is typically $v_{\\rm Ash} \\sim (0.5-1)\\times v_{\\rm sh}$ in our toy model. A more detailed study should determine whether this amplification takes place at the shock or in the region of deceleration.  Interestingly the magnetic effects described in this paper seem to be significant when the Alfv\\'en speed is comparable with the advection speed rather than when the magnetic pressure is comparable with the thermal pressure. These two criteria are related through the Mach number $\\M$: \\begin{equation} \\frac{P_{\\rm mag}}{P_{\\rm th}} = \\frac{\\gamma}{2}\\M^{2}\\left( \\frac{v_{\\rm A}}{v}\\right)^{2}. \\end{equation} In the subsonic flow the pressure ratio is thus smaller than $v_{\\rm A}/v$, by a factor $\\sim 0.06$ at the shock if $\\M_{\\rm sh} \\sim 0.3$, and as low as $10^{-3}$ if $\\M\\sim0.05$ at the coupling radius. This raises the possibility that the magnetic field may affect the growth of SASI through the magnetic tension, even if the field is so weak that the magnetic pressure is negligible.  The strength of the magnetic field present in the iron core of a star before the collapse is very uncertain. The best estimate to date is: $B_{\\phi}\\sim 5.10^9{\\rm G}$ and $B_{r}\\sim 10^6{\\rm G}$ \\citep{heger05}, which is too weak to have any effect on SASI: indeed, flux conservation during the collapse ($B/\\rho^{2/3} = \\rm cste$, e.g. \\cite{shibata06}) would lead to $v_{\\rm Ash}/v_{\\rm sh} \\sim 0.001$ at the shock ($r_{\\rm sh} \\sim150{\\rm km}$, $\\rho_{\\rm sh}\\sim 10^9{\\rm g.cm^{-3}}$), and $v_{\\rm A}/v \\sim 0.025$ at the proto-neutron star surface ($r_{\\rm PNS} \\sim50{\\rm km}$, $\\rho_{\\rm PNS}\\sim10^{11}{\\rm g.cm^{-3}}$).  A reference Alfven speed near the surface of a strongly magnetized proto-neutron star can be estimated by assuming that the fossil magnetic field is strong enough to give birth to a magnetar ($B\\sim 5.10^{15}{\\rm G}$) by simple conservation of the magnetic flux : \\begin{eqnarray} \\frac{v_{\\rm A}}{v} & \\sim &0.93 \\times \\frac{B_{\\rm NS}}{5.10^{15}{\\rm G}}\\left(\\frac{4.10^{14}{\\rm g.cm^{-3}}}{\\rho_{\\rm NS}}\\right)^{2/3}\\left(\\frac{r}{50{\\rm km}}\\right)^{2} \\nonumber \\\\ & & \\left(\\frac{\\rho}{10^{11}{\\rm g.cm^{-3}}}\\right)^{7/6}\\frac{0.3M_{\\odot}.s^{-1}}{\\dot M}. \\end{eqnarray} Such a field could be enough to affect SASI significantly if the amplification of the vortical cycle takes place at the coupling radius, which is slightly above the proto-neutron star surface \\citep{scheck08}. Conversely, if the amplification takes place at the shock, even a fossil magnetic field corresponding to magnetar strength would have a negligible effect on SASI ($v_{\\rm Ash}/v_{\\rm sh} \\sim 0.027$). Note however that even stronger magnetic fields are sometimes considered in the literature, with Alfv\\'en waves propagating \\emph{above} the shock \\citep{suzuki08}.  Given the simplicity of our toy model, we cannot expect it to capture more than the first-order effects of the magnetic field on the advective-acoustic cycle. More accurate estimates should consider a more realistic setup that includes radial convergence and non-adiabatic effects, which are expected to affect the relative importance of the entropy and vorticity cycles. These effects could be estimated in the cylindrical geometry used by \\cite{yamasaki08}.  The role of the Alfv\\'en surface has not been considered in the present study, because the region of acoustic feedback was assumed to be superAlfv\\'enic for the sake of simplicity. If the magnetic field were vertical and strong enough, the magnetic extension of the advective-acoustic cycle to MHD cycles would have to take into account the possible coupling processes taking place at the Alfv\\'en surface. The physics of transAlfv\\'enic accretion flows will be considered in a forthcoming study (Guilet, Fromang \\& Foglizzo, in prep.).  The amplification of the magnetic field observed in the numerical simulations by \\cite{endeve08} seems to take place in the nonlinear regime of the MHD-SASI instability, and is thus beyond the scope of the present linear study. The topology of the magnetic field considered by \\cite{endeve08} is initially vertical (a split monopole), for which our analysis did not reveal any significant magnetic destabilization. Based on the current understanding of the nonlinear saturation process of SASI by parasitic instabilities \\citep{guilet09}, we would rather anticipate a larger saturation amplitude of SASI if the field were horizontal near the shock, due to i) the higher growth rate of the cycles involving the bending of field lines by vorticity perturbations, ii) the resistance of magnetic tension to the development of parasitic instabilities such as Rayleigh-Taylor and Kelvin-Helmholtz instabilities. As noted in Guilet et al. (2009) however, the Rayleigh-Taylor instability may still be able to develop as a parasitic instability in the direction perpendicular to the magnetic field, without bending the field lines. Besides, the presence of other magnetically induced parasitic instabilities cannot be ruled out. Altogether, these qualitative statements are insufficient to explain the growth of the magnetic energy observed by \\cite{endeve08}."
    },
    "0911/0911.2226_arXiv.txt": {
        "abstract": "We have recently carried out the first wide-field hydrogen Paschen-$\\alpha$ line imaging survey of the Galactic Center (GC), using the NICMOS instrument aboard the Hubble Space Telescope. The survey maps out a region of $2253 {\\rm~pc^2}$ ($416 {\\rm~arcmin^2}$) around the central supermassive black hole (Sgr A*) in the 1.87 and 1.90 \\micron\\ narrow bands with a spatial resolution of $\\sim 0.01$ pc (0\\farcs2 FWHM) at a distance of 8 kpc. Here we present an overview of the observations, data reduction, preliminary results, and potential scientific implications, as well as a description of the rationale and design of the survey. We have produced mosaic maps of the Paschen-$\\alpha$ line and continuum emission, giving an unprecedentedly high resolution and high sensitivity panoramic view of stars and photo-ionized gas in the nuclear environment of the Galaxy. We detect a significant number of previously undetected stars with Paschen-$\\alpha$ in emission. They are most likely massive stars with strong winds, as confirmed by our initial follow-up spectroscopic observations. About half of the newly detected massive stars are found outside the known clusters (Arches, Quintuplet, and Central). Many previously known diffuse thermal features are now resolved into arrays of intriguingly fine linear filaments indicating a profound role of magnetic fields in sculpting the gas. The bright spiral-like Paschen-$\\alpha$ emission around Sgr A* is seen to be well confined within the known dusty torus. In the directions roughly perpendicular to it, we further detect faint, diffuse Paschen-$\\alpha$ emission features, which, like earlier radio images, suggest an outflow from the structure. In addition, we detect various compact Paschen-$\\alpha$ nebulae, probably tracing the accretion and/or ejection of stars at various evolutionary stages. Multi-wavelength comparisons together with follow-up observations are helping us to address such questions as where and how massive stars form, how stellar clusters are disrupted, how massive stars shape and heat the surrounding medium, how various phases of this medium are interspersed, and how the supermassive black hole interacts with its environment. ",
        "introduction": "\\label{s:intro}  The Galactic center (GC) is a unique laboratory for a detailed study of star formation and its impact on the nuclear environment of galaxies. Because of its proximity, the GC can be imaged at resolutions more than 100 times better than the nearest galaxies like our own (e.g., the Andromeda galaxy). This unmatched high-resolution capability is particularly important to the study of the formation and destruction processes of stellar clusters as well as the overall population of massive stars in such an environment. Indeed, the GC is known to host three very massive clusters~\\citep[the Arches and Quintuplet clusters, as well as the central parsec cluster surrounding Sgr A$^*$; ages $\\sim 2-7 \\times 10^6$ yrs, masses $\\sim 3\\times10^{4} M_\\odot$ each; e.g.,][]{cot96,ser98,fig99,fig02,gen03}, which are responsible for about 5\\% of the total Lyman continuum flux for the Galaxy, in less than 0.01\\% of its volume. Massive stars are known to exist outside the clusters ~\\citep{cot06}, however, locating these important stars has been sporadic and incomplete~\\citep[e.g. ][]{mun06a,mau07,mau09}. These isolated massive stars may have formed individually or in small groups, although the extreme conditions in the GC make this mode of star formation unusually difficult.   They may also have been spun out of the clusters due to their internal dynamics. This process is expected to be accelerated by the strong external tidal force in the GC~\\citep{kim99}.  \\begin{figure*} % \\centerline{ \\includegraphics[width=1.1\\textwidth,angle=0]{fig1.jpg} } \\caption{% A multi-wavelength montage of the GC: the VLA 20 cm continuum~\\citep[red; ][]{yus84}, \\spitzer\\ 8 $\\mu$m~\\citep[green; ][]{stolo06,are08}, and \\chandra\\ ACIS-I 1-9 keV~\\citep[blue; ][]{wan02,mun09}. The three known massive young clusters are shown, as well as the outer border of the \\pa\\ survey. } \\label{f:fig1} \\end{figure*} \\bigskip  The star formation process depends critically on the interplay of massive stars with the interstellar medium (ISM). Massive stars release large amounts of ultraviolet radiation. They also produce strong stellar winds throughout their short lives and die in supernova explosions. Such energetic stellar output dramatically shapes the morphology of the ISM, as is revealed in existing radio, mid-infrared, and X-ray maps ~\\citep[Fig.~\\ref{f:fig1}; e.g., ][]{yus84,yus02,wan02,wan06,mun09}. Perhaps the best-known example of this is the Radio Arc region which surrounds the Arches and Quintuplet clusters: a collection of arc-like thermal filaments, synchrotron-bright non-thermal linear filaments, and diffuse, hot gas. Some of the structures within this region, the Sickle HII region ionization front and photoevaporated pillars in particular, suggest that gas is being collected and compressed~\\citep[][Cotera et al. in preparation]{cot06}, and is morphologically resembles  known regions outside the GC where such processes are believed to have led to the formation of a new generation of stars (e.g. M16; \\citet{smi05}). The energy injection from massive stars and possibly from the central supermassive black hole (SMBH), that is coincident with the radio source Sgr A*, may also strongly affects the thermal and/or dynamical properties of the ISM. Molecular gas in the GC, for example, has unusually high turbulent velocities and temperatures compared to clouds in other parts of the Galaxy~\\citep[e.g., ][]{mor96}. This may bias the stellar initial mass function toward heavier stars and favor the production of massive star clusters~\\citep[][ and references therein]{mor93,wan06}.  To constrain the dynamical process of the known clusters, the overall population of massive stars, and their formation modes, we sought to obtain a uniform survey of massive stars and their interplay with the ISM across the central 90 parsecs. Ionized warm gas provides the ideal tracer of massive stars and their interplay with their surroundings. Although previous radio continuum observations are free of extinction and trace both thermal and non-thermal emission, the scale-dependence of radio interferometers makes it difficult to accurately measure flux densities on arcsecond scales. In addition, fine features (e.g., ionization fronts) are often found in the midst of diffuse non-thermal emission (e.g., the Radio Arc), making decomposition of thermal and non-thermal components challenging.  Radio recombination line observations typically have a substantially lower signal-to-noise and spatial resolution than radio continuum data. In the near-IR, the 2.16\\micron\\ Br$\\gamma$ line is accessible from the ground (no suitable filter is available for existing and planned space-based platforms), but has also not proved effective for large scale mapping of warm ionized gas. To our knowledge, no large-scale Br$\\gamma$ line survey of the GC has ever been published. The difficulty lies in the lack of the combination of high spatial resolution and stable point spread function (PSF) in a ground-based survey. Observations with adaptive optics (AO) can have superb resolution, but only over small fields of view (typically $15\\arcsec - 40\\arcsec$). The AO PSF also changes strongly, both spatially and with time, and has extended PSF wings, making the PSF subtraction extremely difficult.  In addition, photometric conditions rarely, if ever, persist for the time required for such a survey, hampering the study of large-scale features.  We have carried out the first large-scale, high-resolution, near-IR survey of the GC, using {\\it HST} NICMOS. The primary objective of the survey is to map out the hydrogen Paschen-$\\alpha$ (P$\\alpha$ for short) line (wavelength 1.876 $\\micron$) in a field of $\\sim 39^\\prime \\times 15^\\prime$ around Sgr A$^*$~\\citep[corresponding to 90 pc $\\times$ 35 pc at the GC distance of 8~kpc; ][ Fig.~\\ref{f:fig1}]{ghe08}. This field is known to be rich in clustered massive star formation (i.e., including the three known clusters). NICMOS observations of the P$\\alpha$ emission were taken previously only for a few discrete regions, centered on the three known massive clusters (see \\S~\\ref{ss:comparion} for further discussion). Our survey field also includes much unexplored territory. In particular, the region between Sgr A and Sgr C along the Galactic plane (to the right in Fig.~\\ref{f:fig1}) shows a number of compact {\\it Spitzer} IRAC sources (several with radio counterparts) likely related to some of the earliest stages of massive star activities. This field selection also facilitates an unbiased study of the apparent asymmetry between the positive and negative Galactic longitude sides relative to Sgr A$^*$.   \\begin{figure*} % \\centerline{ \\includegraphics[width=1.15\\textwidth,angle=0]{fig2.jpg} } \\caption{% NICMOS maps from our survey: (a) calibrated F187N mosaic, (b) continuum-subtracted F187N map, and (c) continuum-subtracted and point source removed F187N  (diffuse \\pa) map. These maps are presented with a bin size of 0\\farcs4. The intensity is logarithmically scaled from of 0.01 ${\\rm~mJy~arcsec^{-2}}$ to 10 ${\\rm~mJy~arcsec^{-2}}$ in (a) and to 3 ${\\rm~mJy~arcsec^{-2}}$ in (b) and (c). } \\label{f:fig2} \\end{figure*} \\bigskip  Our \\hst\\ NICMOS observations overcome all of the potential problems of a ground-based survey. The wide NIC3 camera on this instrument provides a resolution of $\\sim0\\farcs2$ with a stable PSF. This capability is important for detecting and cleanly removing point-like sources as well as resolving fine structures of extended features. Both the stable PSF and the absence of atmosphere for our \\hst\\ survey means that the entire survey is photometrically accurate and consistent. The extra magnitude of extinction along the sight-line to the GC at 1.87\\micron\\ (as compared to 2.16 \\micron) is more than compensated for by an intrinsic P$\\alpha$ line intensity 12 times greater than the Br$\\gamma$ line~\\citep{hum87}. In addition, the background at P$\\alpha$ with NICMOS is a factor of $\\sim 800$ lower than the sky background at Br$\\gamma$ as observed from the ground. This low background of \\hst\\ NICMOS also allows us to detect faint point-like sources and extended features, particularly over our large mosaicked field. ",
        "conclusions": "\\label{s:dis}  While the data analysis, multi-wavelength comparison, follow-up observations, and  modeling are still ongoing and will be presented in later papers, we discuss here the implications of the survey for addressing several key issues related to star formation and its impact on the GC environment:  \\subsection{How do massive stars form in the GC?}  Two large questions remain unanswered about massive star formation in the GC: (1) Do massive stars form exclusively in the clusters, or in looser associations throughout the field?  and (2) Is star formation continuous, or does it happen in discrete bursts? Until recently, massive (e.g., emission-line) stars have mostly be found within the 3 GC clusters. A number of searches for young massive stars at infrared and X-ray wavelengths have been carried out ~\\citep[e.g., ][]{cot99,mun06a,mau07,mau09}. These studies, each with diverse search criteria, have yielded a rather incomplete picture of the young massive star population in this GC.  The P$\\alpha$ emission is a sensitive tracer of massive stars and their environs. The emission traces warm ionized gas, produced and maintained primarily by Lyman continuum radiation from massive stars. Such gas can be associated with individual stars (e.g., massive stellar winds), their immediate vicinities (circumstellar material), and surrounding ISM structures. The fact that the majority of our new sources are actually located outside the three known clusters has strong implications for understanding the dynamics of massive stars with ages of only a few $ \\times 10^{6}$ yr. The exact time scales for clusters to be disrupted and the place where the stars end up provide information about the internal dynamics of the clusters as well as the large-scale tidal force of the GC~\\citep[][]{kim99,kim00,kim04,kim03,por02,por03}. The spatial distribution of the massive stars detected in our survey can be used to better constrain these processes throughout the field. The distributions of the  stars and their types as well as follow-up work on both radial velocities and proper motions~\\citep[][]{sto08} will allow us to constrain how quickly the massive stellar clusters are currently dissolving, how many star clusters have been disrupted, and whether or not the stars are formed in massive clusters which have been subsequently dispersed, or formed in relative isolation.   \\subsection{What are the state and dynamics of the ISM?} In the high pressure and high density environment of the GC, the diffuse P$\\alpha$-emitting gas tends to be found at ionization fronts bordering molecular clouds. The gas density and pressure are expected to change drastically at such fronts~\\citep[e.g., ][]{sco98}. The high resolution P$\\alpha$ map resolves various ionization fronts, which helps to determine the local neutral gas density and the thermal pressure of diffuse warm ionized gas. The comparison between Figs.~\\ref{f:fig1} and \\ref{f:fig2} demonstrates the complicated relationship among various stellar and ISM components. Within the Sickle region (Fig.~\\ref{f:fig4}b), for instance, we see two distinctive ionization front morphologies.  We see columns of photo-evaporating gas similar to the elephant trunks or pillars seen in regions such as M16 in one part, but  also a remarkably smooth ionization front located on another side of the large HII region.  Although very different, both are morphologically perpendicular to the brightest non-thermal radio filaments, which are uniquely found in the GC field and have been a long-standing mystery. They are known to trace enhanced magnetic field and/or cosmic rays. But the underlying physical processes remain unclear. The observed relative morphological configuration then suggests that the large-scale, inter-cloud, vertical (poloidal) magnetic field is combing through the dense gas and gets enhanced in the process. The expansion of the magnetized inter-cloud material, probably accompanied with cosmic ray acceleration, is likely driven by the strong mechanical energy output from the Quintuplet and Arches clusters, consistent with the barrel-shaped morphology of the nonthermal radio filaments (Fig.~\\ref{f:fig1}). Therefore, the feedback from the clusters likely plays an important role in shaping the ISM and the eco-system of the GC in general.  \\subsection{Comparison with other observations} \\label{ss:comparion}  A few P$\\alpha$ observations were made previously in the GC region. The Sgr A West region, covering the inner $\\sim$ 4 pc of the Galaxy, was observed using both NIC3 and NIC2~\\citep[e.g., ][]{sco03}. Several other programs concentrated only on the central parsec with NIC1. A NIC2 P$\\alpha$ observation was also made for the core of the Arches cluster~\\citep{fig02}. The P$\\alpha$ emission from the Pistol Nebula (Fig.~\\ref{f:fig4}b) was covered with a $2 \\times 2$ mosaic of NIC2 observations, which had a smaller field of view~\\citep[19\\farcs2 on a side; ][]{fig99}. Our map shows considerably more structure in the diffuse \\pa\\ emission of the Nebula than the image shown in~\\citet{fig99}, apparently due to the higher sensitivity of our NIC3 observations. In addition, our new data provide additional epochs for variability analysis in these regions with the existing observations ~\\citep[e.g., ][]{blu01}.  Our study of the GC region complements other efforts to understand Galactic star forming regions, such as Orion~\\citep[e.g., ][]{fei05}, the Carina nebula~\\citep[e.g., ][]{tap06}, NGC 3603~\\citep[][]{sto04,sto06}, and Westerlund 1~\\citep[e.g., ][]{mun06b}. It is crucial to compare the time scales and efficiency with which stars form in the GC to those in other regions of the Galactic disk in order to gain a detailed understanding of the interplay between massive stars and the ISM.  Ultimately, understanding the gas removal, via the formation of massive stars and their impact on the ISM or other processes~\\citep[e.g., Type Ia supernova-driven Galactic bulge wind; ][]{tan09,li09},  will help to determine why some host SMBHs accrete at rates near their Eddington limits, while others, including Sgr A*, are at rates orders of magnitude lower."
    },
    "0911/0911.0686_arXiv.txt": {
        "abstract": "Future galaxy redshift surveys aim to measure cosmological quantities from the galaxy power spectrum. A prime example is the detection of baryonic acoustic oscillations (BAOs), providing a standard ruler to measure the dark energy equation of state, $w(z)$, to high precision. The strongest practical limitation for these experiments is how quickly accurate redshifts can be measured for sufficient galaxies to map the large-scale structure. A promising strategy is to target emission-line (i.e.\\ star-forming) galaxies at high-redshift ($z\\sim0.5$--$2$); not only is the space density of this population increasing out to $z\\sim2$, but also emission-lines provide an efficient method of redshift determination. Motivated by the prospect of future dark energy surveys targeting H$\\alpha$ emitters at near-infrared wavelengths (i.e.\\ $z>0.5$), we use the latest empirical data to model the evolution of the H$\\alpha$ luminosity function out to $z\\sim2$, and thus provide predictions for the abundance of H$\\alpha$ emitters for practical limiting fluxes. We caution that the estimates presented in this work must be tempered by an efficiency factor, $\\epsilon$, giving the redshift success rate from these potential targets. For a range of practical efficiencies and limiting fluxes, we provide an estimate of $\\bar{n}P_{0.2}$, where $\\bar{n}$ is the 3D galaxy number density and $P_{0.2}$ is the galaxy power spectrum evaluated at $k=0.2\\,h\\,{\\rm Mpc}^{-1}$. Ideal surveys must provide $\\bar{n}P_{0.2}>1$ in order to balance shot-noise and cosmic variance errors. We show that a realistic emission-line survey ($\\epsilon=0.5$) could achieve $\\bar{n}P_{0.2}=1$ out to $z\\sim1.5$ with a limiting flux of $10^{-16}$\\,erg\\,s$^{-1}$\\,cm$^{-2}$. If the limiting flux is a factor 5 brighter, then this goal can only be achieved out to $z\\sim0.5$, highlighting the importance of survey depth and efficiency in cosmological redshift surveys. ",
        "introduction": "One of the greatest challenges the current generation of cosmologists faces is to understand the physics underlying the apparent acceleration of the expansion of the Universe (e.g.\\ Riess et al.\\ 1998; Perlmutter et al.\\ 1999). Contemporary models favour the influence of a dark energy that has come to dominate the energy density of the universe during the last 8\\,billion years. Unfortunately dark energy is outside the realm of the standard model, and requires new physics to explain. Nevertheless, many mechanisms have been proposed, and the potential for establishing which (if any) is correct experimentally, has caused great fervour amongst the astronomical community over the past decade.  The reward for investing a large amount of effort into determining the physics of dark energy is of course a profound advancement of our understanding of the fundamental nature of the universe.  A range of dark energy models exist (see Peebles \\& Ratra\\ 2003 and Copeland et al.\\ 2009 for reviews), however the two most prominent scenarios attribute the accelerating expansion to (a) a `cosmological constant' ($\\Lambda$) analogous to a non-zero quantum mechanical vacuum energy that has now come to dominate the overall energy density of the universe (but 120 orders of magnitude smaller than the value predicted by quantum physics); or (b) a dynamic scalar field (`quintessence') which varies with both time and space. Both models require general relativity to hold on cosmological scales. A third alternative to explain the acceleration is the failure of general relativity on large scales, such that the gravity theory itself needs to be modified (e.g.\\ Dvali et al.\\ 2000).  One way of distinguishing $\\Lambda$ from quintessence is to measure the evolution of the expansion of the universe, which is controlled by the dark energy equation of state, $w(z)$; the ratio between dark energy pressure $P$ and density, $\\rho$. For $\\Lambda$ models, $\\rho=-P/c^2$ for all time, such that $w(z)=-1$. Detecting a varying $w(z)$ would be a possible indication for quintessence.  However, if gravity is modified, such behaviour could be just an indirect effect of the failure of general relativity.  Such degeneracy between dark energy and modified gravity can be lifted only by measuring the growth rate of cosmic structure $f(z)$ (or its integral $G(z)$), which is governed by the interplay between the strength of gravity and the expansion rate of the Universe.  Thus, measuring $w(z)$ and $f(z)$ as a function of redshift represents the most powerful combination of observational probes to distinguish among competing models (Guzzo et al.\\ 2008, Wang et al.\\ 2008). This is the primary goal of current and future dark energy surveys (Albrecht et al.\\ 2009).  Although the accelerated expansion was discovered using observations of Type Ia supernovae (Riess et al.\\ 1998; Perlmutter et al.\\ 1999), results from these observations are now dominated by systematic errors (e.g. Hicken et al.\\ 2009). Future studies of dark energy therefore aim to exploit different observations and of those proposed, the use of Baryon Acoustic Oscillations (BAO) as standard rulers appears to have the lowest level of systematic uncertainty (Albrect et al.\\ 2006). BAO are a series of peaks and troughs in the power spectrum, which quantifies the clustering strength of matter as a function of scale. They occur because primordial cosmological perturbations excite sound waves in the relativistic plasma of the early universe: when the plasma breaks down at recombination, the radiation can be observed as the Cosmic Microwave Background (CMB), while the fluctuations in the baryonic material give rise to BAO (Silk\\ 1968, Peebles \\& Yu\\ 1970, Sunyaev \\& Zel'dovich\\ 1970, Bond \\& Efstathiou\\ 1984, 1987, Holtzman\\ 1989). The BAO signal is on large-scales, which are predominantly in the linear regime today. It is therefore expected that BAO should also be seen in the galaxy distribution (Goldberg \\& Strauss\\ 1998, Meiksin, White \\& Peacock\\ 1999, Springel et al.\\ 2005, Seo \\& Eisenstein\\ 2005, White\\ 2005, Eisenstein, Seo \\& White\\ 2007, Kazin et al.\\ 2009), and can be used as a standard ruler, leading to measurements of the angular diameter distance $D_A(z)$ and the Hubble expansion rate $H(z)$, and therefore $w(z)$ (Seo \\& Eisenstein\\ 2003, Blake \\& Glazebrook\\ 2003, Hu \\& Haiman\\ 2003, Wang\\ 2006).  The acoustic signature has now been convincingly detected at low redshift (Percival et al.\\ 2001, Cole et al.\\ 2005, Eisenstein et al.\\ 2005, Huetsi\\ 2006) using the 2dF Galaxy Redshift Survey (2dFGRS; Colless et al.\\ 2003) and the Sloan Digital Sky Survey (SDSS; York et al.\\ 2000). Further analyses of the SDSS have led to competitive constraints on cosmological models (Percival et al.\\ 2007, Gaztanaga et al.\\ 2009, Percival et al.\\ 2009). Ongoing spectroscopic surveys aiming to use BAO to analyse dark energy include the Baryon Oscillation Spectroscopic Survey (BOSS; Schlegel, White \\& Eisenstein\\ 2009), the Hobby-Eberly Dark Energy Experiment (HETDEX; Hill et al.\\ 2008) and the WiggleZ survey (Glazebrook et al.\\ 2007). Ongoing photometric surveys such as the Dark Energy Survey (DES: {\\tt http://www.darkenergysurvey.org}), the Panoramic Survey Telescope \\& Rapid Response System (Pan-STARRS: {\\tt http://pan-starrs.ifa.hawaii.edu}) aim to find BAO using photometric redshifts.  The power spectrum (or correlation function) of the galaxy distribution also contains key information on the growth rate of structure $f(z)$ (Kaiser\\ 1987). This produces large-scale motions towards density maxima, that contribute a peculiar velocity component to the measured galaxy redshifts used to reconstruct cosmic structure in 3D.  The net effect is to produce an anisotropy in the power spectrum that can be measured to extract an estimate of the growth rate $f(z)$, modulo the bias factor of the galaxies being observed. The importance of this well-known effect in the context of dark energy has become evident only in recent times, when redshift surveys of sufficient size at $z\\sim1$ have started to become available (Guzzo et al.\\ 2008).  Thus, a redshift survey of galaxies provides us with the ability to obtain an estimate of both key probes of cosmic acceleration, the expansion rate and the growth rate.  In order to reduce shot-noise and cosmic variance in `precision' measurements of BAO and redshift distortions, the ultimate observational challenge is to accurately measure a large number (tens or hundreds of millions) of redshifts for galaxies spread over a significant interval of cosmic time, spanning the transition from matter domination to dark energy domination in the universe, and covering the majority of the extragalactic ($|b|>20^\\circ$) sky,\\ $\\sim$$2\\times10^4$ square degrees. Such a survey can be conducted using a dedicated survey telescopes from a space platform, as proposed by the Joint Dark Energy Mission (JDEM: {\\tt http://jdem.gsfc.nasa.gov}) and European Space Agency's Euclid and SPACE satellite mission concepts ({\\tt http://sci.esa.int/euclid}; Cimatti et al.\\ 2009). Some of the ongoing and planned future BAO surveys, such as WiggleZ, will target emission-line galaxies -- i.e.\\ generally star-forming galaxies with easily identifiable redshifts. The goal of this work is to make a prediction for the abundance of H$\\alpha$ emitting galaxies that these dark energy surveys can expect using the existing empirical evidence of past and recent H$\\alpha$ surveys out to $z\\sim2$.   In this work we use empirical data to build a simple phenomenological model of the evolution of the H$\\alpha$ luminosity function (LF) since $z\\sim2$, and therefore predict the number counts of H$\\alpha$ emitters in redshift ranges pertinent to future dark energy surveys (the empirical model can also be used as a fiducial point for semi-analytic predictions for the abundance of star forming galaxies, e.g. Baugh et al.\\ 2005; Bower et al.\\ 2006; Orsi et al.\\ 2009 in prep). In Section 2 we describe the model, list the principal predictions and draw the reader's attention to some important caveats. In Section 3 we discuss the implications of the number count predictions on planned dark energy surveys, and in Section 4 we comment on the relevance of cosmological surveys in the near-IR from a terrestrial base. For luminosity estimates, throughout we assume a fiducial cosmological model of $H_0 = 70$\\,km\\,s$^{-1}$\\,Mpc$^{-1}$, $\\Omega_{\\rm m} = 0.3$ and $\\Omega_\\Lambda = 0.7$.  \\begin{figure} \\includegraphics[width=0.49\\textwidth]{f1.ps} \\caption{Evolution of the H$\\alpha$ luminosity function, assuming our simple model of $L^\\star\\propto(1+z)^Q$ out to $z=1.3$ and no evolution to $z<2.2$. The panels show the LF at $z\\sim0$, 0.2, 0.9 \\& 2.2, with observational data overlaid (all data has been corrected to the same fiducial cosmology used throughout this work and {\\it not} corrected for extinction). Note that not all of the observational data shown here was used to construct the model (see \\S2.1), however the model is a good representation of the observed LFs out to $z\\sim2$. The largest discrepancy occurs at $z\\sim1$, where there is some scatter between different surveys. However, in part, this is due to the mixture of survey strategies, and cosmic variance in the small fields observed.  The model LFs have been truncated at the luminosity limit corresponding to a flux of $10^{-16}$\\,erg\\,s$^{-1}$\\,cm$^{-2}$ at each epoch (vertical dotted lines). Although there are hints that the faint-end slope is steepening out to $z\\sim1$, in the flux regime of practical interest this does not have a significant impact on our counts (also see \\S2.2.1 \\& Figure\\ 3). } \\end{figure}   \\begin{figure} \\includegraphics[width=0.49\\textwidth]{f2.ps} \\caption{A comparison of the predicted number counts of H$\\alpha$ emitters from the simple model to observed counts integrated over the redshift range $0.75 < z < 1.90$. The shaded region indicates the 1$\\sigma$ uncertainty on the model counts. We compare to the observational data of the (slitless spectroscopic) surveys of McCarthy et al.\\ (1999), Hopkins et al.\\ (2000) and Shim et al.\\ (2009), where the integrated counts have been calculated from the respective luminosity functions, uncorrected for dust extinction (so the counts include incompleteness corrections specific to each survey). Note that all-sky redshift surveys are unlikely to probe below flux limits of $\\sim$$10^{-16}$\\,erg\\,s$^{-1}$\\,cm$^{-2}$, where uncertainties due to the poorly constrained faint-end slope become more important to the count predictions (see \\S2.2.1 for more details).} \\end{figure}    \\begin{figure} \\centerline{\\includegraphics[width=0.47\\textwidth]{f3.ps}} \\caption{Predicted redshift distribution ${\\rm d}N/{\\rm d}z$ of H$\\alpha$ emitters for limiting fluxes of $1$--$5\\times 10^{-16}$\\,erg\\,s$^{-1}$\\,cm$^{-2}$ (thick to thin lines). Note that the transition between $L^\\star$ evolution and non-evolution at $z=1.3$ introduces the sharp fall-off in counts towards high-$z$. For comparison, we also show the redshift distribution for the same $L^\\star$ evolution and fixed $\\phi^\\star$, but allowing the faint end slope to steepen monotonically from $-1.35$ at $z=0$ to $-1.6$ at $z=2$. The impact this change has on the predicted counts in the flux limits of practical interest is negligible, and (as expected) more pronounced at fainter limits. } \\end{figure} ",
        "conclusions": "We have presented a simple prescription for the prediction of the abundance of H$\\alpha$ emitters over $0 < z < 2$, based on empirical data. The model is simplistic, due to limited available data; it assumes a fixed space density, fixed faint end slope, and only $L^\\star$ evolution out to $z=1.3$. There is no luminosity evolution to higher redshifts, consistent with current H$\\alpha$ observations at this redshift. Despite its simplicity, the model adequately mimics the observed luminosity functions of a range of H$\\alpha$ surveys (including a mixture of spectroscopic, grism and narrowband strategies). Using the luminosity function model as a basis, we predict the redshift distribution of H$\\alpha$ emitters corresponding to a spectral coverage that extends to 2$\\mu$m.  Our results have particular relevance to dark energy experiments attempting to measure cosmological information from the power spectrum of galaxies detected in all-sky H$\\alpha$ surveys in the near-IR. We use the parameter $\\bar{n}P_{0.2}$ as a measure of the effectiveness of a redshift survey, and make predictions for this value for a range of redshift, limiting flux and success rate (i.e.\\ efficiency). To achieve $\\bar{n}P_{0.2}=1$ out to $z=2$, emission-line surveys should be aiming for limiting fluxes of $\\sim$10$^{-16}$\\,erg\\,s$^{-1}$\\,cm$^{-2}$. However, this estimate is reliant on a high success rate of the sampling of the H$\\alpha$ population: redshift surveys need to aim for high-efficiencies, since any decline in redshift yield (i.e.\\ failing to obtain redshifts for detections) has the same effect on $\\bar{n}P_{0.2}$ as increasing the flux limit (illustrated in Figure\\ 5 of this work). Assuming a more likely situation of 50\\% efficiency, a realistic target for proposed surveys is $\\bar{n}P_{0.2}=1$ at $z=1.5$. At higher redshifts the sharply declining number counts have a severe effect on one's ability to measure $w(z)$ at the desired precision."
    },
    "0911/0911.0365_arXiv.txt": {
        "abstract": "{}{The methodology for modeling grain-surface chemistry has been greatly improved by taking into account the grain size and fluctuation effects. However, the reaction rate coefficients currently used in all practical models of gas-grain chemistry are inaccurate by a significant amount. We provide expressions for these crucial rate coefficients that are both accurate and easy to incorporate into gas-grain models. } {We use exact results obtained in earlier work, where the reaction rate coefficient was defined by a first-passage problem, which was solved using random walk theory. } {The approximate reaction rate coefficient presented here is easy to include in all models of interstellar gas-grain chemistry. In contrast to the commonly used expression, the results that it provides are in perfect agreement with detailed kinetic Monte Carlo simulations. We also show the rate coefficient for reactions involving multiple species. } {} ",
        "introduction": "The chemistry in interstellar clouds of gas and dust is very complex, consisting of reactions taking place in the gas phase as well as on the surfaces of dust grains \\citep{herbst95}. Of the latter reactions, the formation of molecular hydrogen \\citep{hollenbach71b} is of particular importance, but many more processes have been identified \\citep{hasegawa92}. Modeling of these complex gas-grain chemical networks is important to the understanding of the observed chemical complexity in interstellar clouds and the effects of those molecules on gravitational collapse and star formation \\citep{herbst95,hartquist95,tielens05}.  While for gas-phase reactions, rate equations are widely used and fully appropriate, for grain surface chemistry they were found to be inaccurate for a wide range of astrophysically relevant conditions \\citep{tielens95,caselli98,biham02,biham05,lohmar06,lederhendler08}. To address this problem, modified rate equations were proposed \\citep{caselli98,stantcheva01,garrod08}. They were found to provide improved results, but involve ad-hoc and uncontrolled approximations. The master equation, describing the time evolution in the probability distribution of the number of reactants on the grain, provides a complete and accurate description of the reactions on grain surfaces \\citep{biham01,green01,biham02}.  The surface chemistry was also studied using stochastic simulations \\citep{tielens82,charnley01}. Kinetic Monte Carlo simulations (KMC) have been developed for different surface morphologies \\citep{chang05,cuppen05}.  An important advantage of approaches based on differential equations is that they can easily be coupled to the rate equations of gas-phase chemistry. The difficulty with the master equation is that the number of equations quickly proliferates as the number of reactive species on the grains increases, thus approximations are needed \\citep{stantcheva02,stantcheva03}. From the master equation, it is possible to construct a set of moment equations that account correctly for the reaction rates on grain surfaces \\citep{lipshtat03,barzel07a,barzel07b}.  Here we consider the form of the rate coefficients that enter the models. Apart from KMC simulations, all approaches to grain-surface chemistry feature the \\emph{sweeping rate} $A$ as a crucial parameter. This rate coefficient governs the reaction rate on the surface, where there is competition between diffusion-mediated reaction and desorption. Without a precise definition of $A$, the expression $A=a/S$ is widely used, where $a$ is the hopping rate between adjacent adsorption sites and $S$ is the number of these sites on the grain. This expression does not account for the competition between reaction and desorption. It also does not account for the peculiarities of two-dimensional diffusion.  \\citet{lohmar06} introduced a proper definition of the rate coefficient $A$ in terms of the two-particle \\emph{encounter probability}. These authors evaluated it exactly for a single-species reaction on a closed two-dimensional surface where the adsorption sites form a regular lattice \\citep{lohmar06,lohmar08}. The result was compared to the conventional approximation for $A$, which was found to deviate significantly. However, the exact expression contains a large sum of terms, which is costly to evaluate in practice.  In this Letter, we present an expression for the reaction rate coefficient that is both accurate and easy to evaluate. We compare the reaction rates obtained using this expression to those obtained from KMC simulations and find perfect agreement. We also generalize the expression to the case of reactions involving multiple species. ",
        "conclusions": "We have used exact results for the sweeping rate in diffusion-mediated reactions on dust grain surfaces to provide an approximate formula for the crucial reaction rate coefficient.  This expression is numerically accurate, and easy and efficient to implement even in complex gas-grain reaction networks.  For single-species reactions, the expression was employed in the exact analytical master equation framework, and comparison with KMC simulations shows excellent agreement throughout.  The results were also generalized to reactions of multiple species.  We strongly advocate the incorporation of the reaction rate coefficients presented here into \\emph{all} models of interstellar gas-grain chemistry.  This could be easily achieved and, in stark contrast to the commonly used approximation, it provides accurate results at virtually no additional computational cost.  This is an important step towards an accurate modeling of the fascinating chemistry of interstellar clouds."
    },
    "0911/0911.4441.txt": {
        "abstract": "Young isolated radio-quiet neutron stars are still hot enough to be detectable at X-ray and optical wavelengths due to their thermal emission and can hence probe cooling curves. An identification of their birth sites can constrain their age.\\\\ For that reason we try to identify the parent associations for four of the so-called Magnificent Seven neutron stars for which proper motion and distance estimates are available. We are tracing back in time each neutron star and possible birth association centre to find close encounters. The associated time of the encounter expresses the kinematic age of the neutron star which can be compared to its characteristic spin-down age. Owing to observational uncertainties in the input data, we use Monte-Carlo simulations and evaluate the outcome of our calculations statistically. \\\\ \\rxja{} most probably originated from the Upper Scorpius association about $\\unit[0{.}3]{Myr}$ ago. \\rxjn{} was either born in the young local association TW Hydrae about $\\unit[0{.}4]{Myr}$ ago or in Trumpler 10 $\\unit[0{.}5]{Myr}$ in the past. Also \\rxjs{} and RBS 1223 seem to come from a nearby young association such as the Scorpius-Centraurus complex or the extended Corona-Australis association. For RBS 1223 also a birth in Scutum OB2 is possible.\\\\ We also give constraints on the observables as well as on the radial velocity of the neutron star. Given the birth association, its age and the flight time of the neutron star, we estimate the mass of the progenitor star.\\\\ Some of the potential supernovae were located very nearby ($<\\unit[100]{pc}$) and thus should have contributed to the $^{10}$Be and $^{60}$Fe material found in the Earth's crust.\\\\ In addition we reinvestigate the previously suggested neutron star/ runaway pair PSR B1929+10/ $\\zeta$ Ophiuchi and conclude that it is very likely that both objects were ejected during the same supernova event. ",
        "introduction": "\\label{sec:intro}  Neutron stars have very large proper motions which, with known distances, indicate high space velocities \\citep[e.g.][thereafter Ho05]{1994Natur.369..127L,1997MNRAS.289..592L,1997MNRAS.291..569H,1998ApJ...505..315C,2002ApJ...568..289A,2005MNRAS.360..974H}. Those high velocities that are usually larger than those of the progenitor stars are the result of asymmetric supernova explosions assigning the new-born neutron star a kick velocity for that a number of mechanisms have been suggested \\citep[e.g.][]{1996PhRvL..76..352B,1996A&A...306..167J,2005ASPC..332..363J,2006ApJ...639.1007W,2009arXiv0906.2802K}.\\\\ Another fact is that about $\\unit[46]{\\%}$ of O and $\\unit[10]{\\%}$ of B stars also show high space velocities (see \\citealt{1991AJ....102..333S} for a discussion). Two scenarios \\citep[thereafter H01]{2001A&A...365...49H} are accepted to produce those so-called ``runaway stars'' \\citep{1961BAN....15..265B}. The binary-supernova scenario is related to the formation of the high velocity neutron stars: The runaway and neutron star are the products of a supernova within a binary system. The velocity of the former secondary is comparable to its original orbital velocity. The second scenario, which will not be further discussed within this paper, is the dynamical ejection due to gravitational interactions between massive stars in dense clusters. H01 proposed examples for each of the scenarios. In \\autoref{sec:BSS} of this paper we will review and further investigate the binary supernova scenario for PSR B1929+10 and the runaway O9.5V star $\\zeta$ Oph. Instead of restricting the unknown pulsar radial velocity to a certain interval, we will use a velocity distribution to cover the whole spectrum of possible radial velocities. We then use more recent data for the pulsar to strengthen the hypothesis.\\\\ Beforehand, we introduce our sample of OB associations and clusters in \\autoref{sec:assocsample} and the procedure utilised in \\autoref{sec:procedure}.\\\\ In \\autoref{sec:idparassoc} we investigate the origin of four isolated radio-quiet X-ray emitting neutron stars. So far, seven such sources have been confirmed. They are not only bright X-ray sources, but are also detected in the optical due to a hot cooling surface. In two cases a parallax is obtained. Brightness, parallax and temperature yield their radii. From spectra, one can in principle determine their mass and composition, which eventually may lead to constraints on the equation of state. However, only seven such sources have been identified up to now for which they were named ``The Magnificent Seven'' (M7) (\\citealt{2001ASPC..234..225T}; for a recent review see \\citealt{2008AIPC..968..129K}; two new candidates were recently published by \\citealt{2009A&A...498..233P} and \\citealt{2008ApJ...672.1137R}). With known luminosity, one only needs to estimate the age, in order to probe cooling curves. For young neutron stars, the characteristic spin-down age only gives an upper limit. We will therefore try to estimate their age from kinematics.\\\\ Our aim in this paper is to identify the parent associations of four members of the M7 and find constraints on their observables and radial velocity.\\\\ We give a summary of our results and draw our conclusions in \\autoref{sec:conclusions}.  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% ",
        "conclusions": "\\label{sec:conclusions}  \\begin{table*} \\centering \\caption{Summary of the results discussed in \\autoref{sec:idparassoc}. Given here are those associations that we conclude to be the most probable birth places of the four M7 members. Columns 3 to 5 give the position of the potential supernova (SN) as equatorial coordinates and distance to the Sun, column 6 gives the approximate time in the past at which the SN may have occurred and columns 7 to 10 give constraints on the neutron star parameters as they would be needed to reach the respective association at the given time.}\\label{tab:summary} \\footnotesize \\setlength\\extrarowheight{2pt} \\begin{tabular}{l | l | >{$}c<{$} >{$}c<{$} >{$}c<{$} | c | >{$}c<{$} >{$}c<{$} o{3.2} o{3.2}} \\toprule RX J & Association & \\multicolumn{3}{c|}{Position of SN} & Time& \\multicolumn{4}{c}{Present-Day Parameters}\\\\ &\t\t\t\t\t\t & \\multicolumn{1}{c}{$\\unit[\\alpha]{[^\\circ]}$} & \\multicolumn{1}{c}{$\\unit[\\delta]{[^\\circ]}$} & \\multicolumn{1}{c|}{$\\unit[d_{\\odot}]{[pc]}$} & of SN & \\multicolumn{1}{c}{$\\unit[\\pi]{[mas]}$} & \\multicolumn{1}{c}{$\\unit[v_r]{[km/s]}$} &  \\multicolumn{1}{c}{$\\unit[\\mu_{\\alpha}^*]{[mas/yr]}$} & \\multicolumn{1}{c}{$\\unit[\\mu_{\\delta}]{[mas/yr]}$}\\\\\\midrule 1856.5-3754 & US\t\t\t\t\t\t & 243{.}21^{+0{.}58}_{-0{.}52} & -23{.}66^{+0{.}31}_{-0{.}29} & 140\\ldots163 & $\\unit[0{.}3]{Myr}$ & 5{.}41^{+0{.}33}_{-0{.}28} & 193^{+45}_{-32} & 326{.}7+0{.}8 & -59{.}1+0{.}7\\\\\\hline 0720.4-3125 & TWA\t       & 187{.}10^{+3{.}44}_{-2{.}91} & -37{.}98^{+1{.}53}_{-1{.}38} & 48\\ldots67 & $\\unit[0{.}4]{Myr}$ & 4{.}06^{+0{.}47}_{-0{.}49} & 502^{+111}_{-88} & -93{.}9+2{.}2 & 52{.}8+2{.}3\\\\ & Tr 10\t       & 135\\ldots140 & -40{.}19^{+0{.}58}_{-0{.}59} & 335\\ldots375 & $\\unit[0{.}5]{Myr}$ & 41{.}93^{+0{.}24}_{-0{.}20} & 290^{+143}_{-110} & -93{.}8+2{.}2\t& 53{.}1+2{.}2\\\\\\hline 1605.3+3249  & Sco-Cen\t\t\t\t & 255{.}09^{+0{.}57}_{-0{.}66} & -18{.}44^{+1{.}00}_{-1{.}06} & 129\\ldots171 & $\\unit[0{.}5]{Myr}$ & 3{.}38^{+0{.}47}_{-0{.}41} & 343^{+86}_{-80} & -43{.}7+1{.}6 & 148{.}8+2{.}6\\\\ & Ext. R CrA& 261{.}99^{+1{.}12}_{-0{.}99} & -39{.}84^{+2{.}10}_{-2{.}00} & 70\\ldots115 & $\\unit[0{.}5]{Myr}$ & 4{.}41^{+0{.}90}_{-0{.}85} & 321^{+109}_{-58} & -43{.}7+1{.}7 & 148{.}8+2{.}6\\\\\\hline RBS 1223\t\t\t  & Sco-Cen\t\t & 254{.}09^{+2{.}11}_{-2{.}34} & -6{.}29^{+3{.}23}_{-5{.}09} & 109^{+36}_{-19} & $\\unit[0{.}5]{Myr}$ & 4{.}63^{+0{.}84}_{-1{.}27} & 351^{+100}_{-104} & -210+20 & 83+18\\\\ & Ext. R CrA\t\t & 294{.}82^{+7{.}24}_{-7{.}92} & -22{.}94^{+4{.}33}_{-6{.}29} & 80^{+22}_{-10} & $\\unit[0{.}5]{Myr}$ & 5{.}78^{+1{.}33}_{-1{.}21} & 345^{+114}_{-90} & -209+20 & 85+18\\\\ & Sct OB2A\t\t & 277{.}15^{+2{.}47}_{-2{.}60} & -12{.}04^{+2{.}13}_{-1{.}83} & 517^{+18}_{-19} & $\\unit[1]{Myr}$ & 2{.}70^{+0{.}96}_{-0{.}44} & 75\\ldots310 & -216+19 & 53+10\\\\\\bottomrule \\end{tabular} \\end{table*}  We reviewed the prominent neutron star/ runaway case PSR B1929+10/ $\\zeta$ Oph (H01) and reinvestigated it with more recent data for the pulsar. Although we did not put constraints on the radial velocity, we could confirm previous results from \\citet{2008AstL...34..686B} and eventually derive a radial velocity of $\\approx\\unit[250]{km/s}$ of the pulsar which is needed to find the two objects at nearly the same position in US at the same time in the past. However, we should bear in mind that we had to increase the errors of the parallax and proper motion components. If they are truly as small as published by \\citet{2004ApJ...604..339C} the scenario is less likely.\\\\  We investigated the motion of four members of the M7 and for each neutron star simultaneously the motion of 140 OB associations and clusters to confine potential parent associations of the neutron stars. A summary of our results is given in \\autoref{tab:summary} (for individual discussions please see previous sections).\\\\ We find that \\rxja{} most probably originated from US about $\\unit[0{.}3]{Myr}$ ago. \\rxjn{} was very likely born in TWA about $\\unit[0{.}4]{Myr}$ ago or may come from Tr 10 where the supernova then would have occured $\\unit[0{.}5]{Myr}$ in the past. \\rxjs{} and RBS 1223 also were seemingly born from a close young association such as the Scorpius-Centraurus complex or Ext. R CrA. For RBS 1223 also a birth in Scutum OB2 is possible.\\\\  Thus, we find that the most probable birth places were located close to the Sun. The supernova distances are consistent with the Local Bubble and hence may have contributed to its formation or reheating (\\citealt{2002A&A...390..299B}; F08). There is no doubt that the Sco OB2 association to which US belongs experienced some supernovae in the past (\\citealt{2001ApJ...560L..83M,2002A&A...390..299B}; F08).  There have been other investigations showing that many neutron stars may have been born within the Gould Belt \\citep{2003A&A...406..111P}, a torus-like structure around the Sun with a radius of $\\approx\\unit[600]{pc}$ \\citep{2000A&A...359...82T}, to which most of the associations just mentioned also belong.\\\\ Some of the four investigated M7 neutron stars may have formed within $\\approx\\unit[250]{pc}$ within the last $\\unit[0{.}5]{Myr}$, so that they could have contributed to the $^{10}$Be and $^{60}$Fe found in the Earth's crust.  Adopting a birth scenario of each neutron star, we can derive estimations of the mass of the progenitor star assuming contemporary star formation in the association as well as the kinematic neutron star age which we here always find to be lower than the characteristic spin-down age (\\autoref{tab:agemass}) that is an upper limit to the true age. From cooling models we know that, given the surface temperature, for the neutron stars investigated here, the characteristic age is too high. This is also plausible since those objects are still very young. Thus, our kinematic ages better fit with cooling models (see \\autoref{fig:cooling}).\\\\ Since \\rxja{} is the coolest of the four M7 members investigated (\\rxjn{}, \\rxjs{} and RBS 1223 show similar temperatures), it should be the oldest. However, it seems, that \\rxja{} is slightly younger than the other three (or has a similar kinematic age if we adopt the parallax of \\citealt{2000IAUS..195..437W}). This might be owing to different parameters at birth or different cooling. Also the kinematic age of \\rxja{} is of about a factor of $13$ smaller than its characteristic age which is a large difference. This may indicate further doubts on the characteristic age. \\citet{2003ApJ...592L..75M} propose that \\rxja{} is an isolated magnetar that has spun down by the propeller effect. In this case, dipole braking, which is assumed to calculate characteristic ages, is no longer valid. Moreover, \\rxja{} does only show a very small pulse fraction making it difficult to derive $\\dot{P}$ \\citep{2008ApJ...673L.163V}.\\\\ In addition, effective temperatures may be influenced by hot spots and do not reflect the real surface temperature. Infact, this is the case for \\rxjn{} \\citep{2006A&A...451L..17H,2009A&A...498..811H} and RBS 1223 \\citep{2005A&A...441..597S,2007Ap&SS.308..181H} which appear hotter, thus must be actually shifted somewhat to smaller temperatures (even more consistent with cooling curves).\\\\  \\stepcounter{footnote} \\begin{table} \\centering \\caption{Kinematic ages $\\tau_{kin}$ of the four M7 members compared with their characteristic spin-down ages $\\tau_{char}$ (taken from the ATNF pulsar database${}^{\\thefootnote}$, \\citealt{2005AJ....129.1993M}). The fifth column gives the effective temperature as listed in Table 1 of \\citet{2005fysx.conf...39H} (see references therein). Column 6 gives the median masses (for TWA a mass consistent with its mass function) of the progenitor stars predicted using evolutionary models from T80, MM89 and K97 (see text) given the difference between the ages of the associations and the time since the predicted supernovae.% and column 7 gives the respective spectral types on the main sequence according to \\citet{Schmidt-Kaler1982}.} }\\label{tab:agemass} \\footnotesize \\setlength\\extrarowheight{2pt} \\begin{tabular}{l l d{2.1} d{1.1} d{4.0} d{3.0} }% c} \\toprule RX J\t&\tAssociation\t& \\multicolumn{1}{c}{$\\tau_{kin}$} & \\multicolumn{1}{c}{$\\tau_{char}$}  &\t\\multicolumn{1}{c}{$T_{eff,\\infty}$} & \\multicolumn{1}{c}{$M_{prog}$} \\\\ &\t\t\t\t\t\t\t&\t\\multicolumn{1}{c}{[Myr]}\t\t\t\t & \\multicolumn{1}{c}{[Myr]}\t\t\t\t\t& \\multicolumn{1}{c}{[$\\unit[10^3]{K}$]} & \\multicolumn{1}{c}{[$\\mathrm{M_{\\odot}}$]} \\\\\\midrule 1856.5-3754 & US \t\t\t\t\t&  \\approx0,3\t\t& 3,8  \t\t\t\t& 696\t\t  \t&  \\approx29  \\\\\\hline 0720.4-3125 & TWA   \t\t\t&  \\approx0,4\t\t& 1,9\t \t\t\t\t& 1044  \t\t&  \\approx10  \\\\ & Tr 10 \t\t\t&  \\approx0,5\t\t& \t \t\t\t\t\t&    \t\t\t\t&  \\approx10  \\\\\\hline 1605.3+3249 & Sco-Cen\t\t\t&  \\approx0,5\t\t& \\multicolumn{1}{c}{--}\t&  1113   & \\approx29 \\\\ & Ext. R CrA\t&  \\approx0,5\t\t& \t\t\t\t\t\t&   \t\t\t  & \\approx15   \\\\\\hline RBS 1223\t\t&\tSco-Cen\t\t\t&  \\approx0,5\t\t& 1,5  \t\t\t\t& 998    \t\t&  \\approx29  \\\\ & Ext. R CrA  &  \\approx0,5\t\t& \t\t\t\t\t\t&\t\t\t\t\t\t&  \\approx15  \\\\ & Sct OB2\t\t\t&  \\multicolumn{1}{c}{$\\approx1$}\t&\t\t&\t\t&  \\approx38  \\\\ \\bottomrule \\end{tabular} \\end{table} \\footnotetext{\\href{http://www.atnf.csiro.au/research/pulsar/psrcat/}{http://www.atnf.csiro.au/research/pulsar/psrcat/}}  \\begin{figure} %\\centering \\includegraphics[width=0.45\\textwidth, viewport= 85 265 480 590]{Fig12.pdf} \\caption{The four M7 members inserted into a cooling diagram. Filled stars mark the characteristic spin-down age (see \\autoref{tab:agemass}) whereas horizontal lines characterise an area of the kinematic age (lower and upper values from associations in tables of \\autoref{sec:idparassoc}). Open diamonds show the kinematic age for the associations summarised in \\autoref{tab:summary}. Effective temperatures can be found in \\autoref{tab:agemass}.\\newline The purple set of cooling curves was adopted from \\citet{2006PhRvC..74b5803P} (solid lines, for masses of $1{.}05$, $1{.}13$, $1{.}22$, $1{.}28$, $1{.}35$, $1{.}45$, $1{.}55$, $1{.}65$ and $\\unit[1{.}75]{M_{\\odot}}$ from top to bottom; Model from \\citealt{2005PhRvC..71d5801G}), the green set has been kindly provided by A. D. Kaminker (dashed lines, includes superconductive protons and neutrons, for masses from $1{.}1$ to $\\unit[1{.}9]{M_{\\odot}}$ from top to bottom; see also \\citealt{2005MNRAS.363..555G}).} \\label{fig:cooling} \\end{figure}"
    },
    "0911/0911.5554_arXiv.txt": {
        "abstract": "We investigate a series of steady-state models of galaxy clusters, in which the hot intracluster gas is efficiently heated by active galactic nucleus (AGN) feedback and thermal conduction, and in which the mass accretion rates are highly reduced compared to those predicted by the standard cooling flow models. We perform a global Lagrangian stability analysis. We show for the first time that the global radial instability in cool core clusters can be suppressed by the AGN feedback mechanism, provided that the feedback efficiency exceeds a critical lower limit. Furthermore, our analysis naturally shows that the clusters can exist in two distinct forms. Globally stable clusters are expected to have either: 1) cool cores stabilized by both AGN feedback and conduction, or 2) non-cool cores stabilized primarily by conduction. Intermediate central temperatures typically lead to globally unstable solutions. This bimodality is consistent with the recently observed anticorrelation between the flatness of the temperature profiles and the AGN activity \\citep{dunn08} % and the observation by Rafferty et al. (2008) \\citep{rafferty08} that the shorter central cooling times tend to correspond to significantly younger AGN X-ray cavities. ",
        "introduction": "Recent high-resolution Chandra and XMM-Newton observations of cluster cool cores show a remarkable lack of emission lines from the gas at temperatures below about ~ 1/3 of the ambient cluster temperature, which is suggestive of one or more heating mechanisms maintaining the hot gas at keV temperatures for a period at least comparable to the lifetime of galaxy clusters. Since global thermal instability may result in a cooling catastrophe and a strong cooling flow, a successful quasi-equilibrium model for the ICM must be globally stable, or at least only have instabilities which grow on extremely long timescales. In \\citet{Guo08b}, we performed a formal global Lagrangian stability analysis (in the spirit of that performed by \\citep{kim03}) to investigate thermal instability in quasi-equilibrium galaxy clusters with thermal conduction and AGN feedback heating. ",
        "conclusions": ""
    },
    "0911/0911.2468_arXiv.txt": {
        "abstract": "% We analyse the potential migration of massive planets forming far away from an inner planetary system. For this, we follow the dynamical evolution of the orbital elements of a massive planet undergoing a dissipative process with a gas disc centred around the central sun.  We use a new method for post-Newtonian, high-precision integration of planetary systems containing a central sun by splitting the forces on a particle between a dominant central force and additional perturbations. In this treatment, which allows us to integrate with a very high-accuracy close encounters, all gravitational forces are integrated directly, without resorting to any simplifying approach. After traversing the disc a number of times, the planet is finally trapped into the disc with a non-negligible eccentricity ",
        "introduction": "It has been suggested that the most massive planets in a planetary system can be formed by a process of gas collapse, independently of metallicity, whilst the lighter components would have formed via core accretion. This can lead to a situation in which massive planets that have originated relatively far away from the inner system, migrate inwards into it \\citep[see][and contribution by Font-Ribera et al. in this same volume]{Font-RiberaEtAl09}.  In this process, these more massive planets with larger semi-major axis, cross the gas disc centred around the central sun. When going through this dissipative process, the planets lose kinetic energy because of the friction with the gas.  As a consequence, the inner disc is heated up and the semi-major axis of the planet shrinks. After some passages, the planet is trapped in the disc with a residual eccentricity. We propose this scenario as a plausible way of explaining the existence of massive planets distributed around a sun with non-zero eccentricities. We give results based on high-accurate dynamical simulations about the distribution of the orbital elements of the trapped objects in the disc.  We find that the massive planets are typically captured after some $\\sim 10^{5}$ yrs and the final eccentricity is non-negligible ($e \\sim 0.1$). ",
        "conclusions": "We have recently developed an integrator {\\tt bhint} specialised for dynamical processes in the vicinity of a very massive particle, which relies in the assumption that the very massive particle dominates the motion of the smaller ones \\citep{LoeckmannBaumgardt08}. For this new mechanism, we retain the Hermite scheme as a basis. The bottleneck is, of course, the number of particles to be used. Nonetheless, we resort to special-purpose hardware, the GRAPE, a card specially developed to integrate the calculation of Newtonian gravitational forces.  The peak performance of one of these cards is of 130 Gflop, roughly equivalent to 100 single PCs, which makes possible long simulations with a realistic particle number.  \\parpic[r][r]{\\mbox{\\resizebox{0.45\\hsize}{!} {\\includegraphics{p_AmaroSeoane_1_f1.eps}}}}  We set initially a disc made out of $10^3$ small particles which is ``hosting'' a sun in the centre and follows a simple $1/r$ profile. The integrated mass is of some 5 Jupiters and the radius of some $30$ AU. The thickness of the disc is of about the diameter of the central sun and has a gap around the central sun which extends $0.1$ AU.  The mass of the central sun is $1\\,M_{\\odot}$. The mass forming the disc are all single-mass. The massive planet, a massive particle of 5 Jupiters is set in an orbit such that the initial eccentricity with the sun is of $e = 0.95$.  Initially, the particle is 100 AU away from the sun and the inclination angle is $i = 90$ degrees.  The system (disc plus interloper) is integrated until the interloper is trapped by the disc. In the figure we show the evolution of the orbital parameters. Whilst we cannot discuss them in detail because of the publication limits, we note that after some $4 \\times 10^5$ yrs the inclination has almost not changed as compared to its initial value. Then, after a short time of $75 \\times 10^3$ yrs elapses, it abruptly decays from almost 80 degrees to a very small number, to be finally trapped in the disc after $5 \\times 10^5$ yrs. The energy, whilst it decays from the initial high value of 0.95, is of $e \\sim 0.1$ when the massive planet is totally trapped in the disc, within a final semi-major axis which is well within the range of expectation.  An extended and detailed scrutiny of the parameter space of this capture process we propose will be soon published elsewhere \\citep{ASEtAl09}.  The adressing of this scenario has direct bearing on our understanding of planetary dynamics and migration mechanisms."
    },
    "0911/0911.2697_arXiv.txt": {
        "abstract": "{} {Physical and evolutionary properties of the sub-mJy radio population are not entirely known. The radio/optical analysis of the ATESP 5~GHz sample has revealed a significant class of compact flat/inverted radio-spectrum sources associated to early-type galaxies up to redshift $2$. Such sources are most plausibly triggered by an AGN, but their observational properties are not entirely consistent with those of standard radio galaxy populations. In the present work we aim at a better understanding of the radio spectra of such sources and ultimately of the nature of AGNs at sub-mJy flux levels. In particular we are interested in assessing whether the AGN component of the sub--mJy population is more related to efficiently accreting systems - like radio-intermediate/quiet quasars - or to systems with low accretion rates - like e.g. FRI radio galaxies - or to low radiative efficiency accretion flows - like e.g. ADAF. } {We used the ATCA to get multi-frequency (4.8, 8.6 and 19 GHz) quasi-simultaneous observations for a representative sub-sample of ATESP radio sources associated with early-type galaxies (26 objects with $S>0.6$ mJy). This can give us insight into the accretion/radiative mechanism that is at work, since different regimes display different spectral signatures in the radio domain. } {From the analysis of the radio spectra, we find that our sources are most probably jet-dominated systems. ADAF models are ruled out by the high frequency data, while ADAF+jet scenarios are still consistent with flat/moderately inverted-spectrum sources, but are not required to explain the data. We compared our sample with high ($\\gsimeq 20$ GHz) frequency selected surveys, finding spectral properties very similar to the ones of much brighter ($S>500$ mJy) radio galaxies extracted from the Massardi \\etal (2008) sample. Linear sizes of ATESP 5~GHz sources associated with early type galaxies are also often consistent with the ones of brighter B2 and 3C radio galaxies, with possibly a very compact component that could be associated at least in part to (obscured) radio-quiet quasar-like objects and/or low power BL Lacs.} {} ",
        "introduction": "\\label{sec:introduction} After many years of studies, AGNs are now recognised to contribute significantly at radio fluxes below 1 milliJy (mJy). Multi-wavelength studies of deep radio fields show that star-forming galaxies dominate at microJy ($\\mu$Jy) levels \\citep[Richards~\\etal][]{Ric99}, while radio sources associated with early--type galaxies and plausibly triggered by AGNs, are the most significant source component ($\\sim 60-70\\%$ of the total) at higher flux densities ($>100-200$ $\\mu$Jy) with a further 10\\% contribution from broad-/narrow-line AGNs (see e.g. Gruppioni~\\etal~\\citealp{Gru99}; Georgakakis~\\etal~\\citealp{Geo99}; Magliocchetti~\\etal~\\citealp{Mag00}; Prandoni~\\etal~\\citealp{Pra01b}; Muxlow~\\etal~\\citealp{Mux05}; Afonso~\\etal~\\citealp{Afo06}; Mignano ~\\etal~\\citealp{Mig08}; Smolcic~\\etal~\\citealp{Smo08}). This somehow unexpected presence of large numbers of AGN--type sources at the sub-mJy level has given a new and interesting scientific perspective to the study of deep radio fields, since a better understanding of the physical and evolutionary properties of such low/intermediate power AGNs may have important implications for the determination of the black-hole-accretion history of the Universe as derived from radio-selected samples. Of particular interest is the possibility of assessing whether the AGN component of the sub--mJy population is more related to efficiently accreting systems - like radio-intermediate/quiet quasars - or to systems with very low accretion rates - like e.g. FRI radio galaxies \\cite[Fanaroff \\& Riley][]{Fan74}. The latter scenario (radio mode) is supported by the presence of many optically inactive early type galaxies among the sub--mJy radio sources. The quasar mode scenario, on the other hand, may be supported by the large number of so-called radio-intermediate quasars observed at mJy levels \\cite[see e.g. Lacy \\etal][]{Lac01} and by the modelling work of Jarvis \\& Rawlings \\cite{Jar04}, who predict a significant contribution of radio-quiet quasars at sub-mJy levels. Another issue is represented by the possible role played at low radio fluxes by low radiation efficiency accretion mechanisms, associated to optically thin discs, such as the so-called {\\it advection dominated accretion flows} (ADAF) and modifications (ADIOS, CDAF, etc; see Narayan \\& Yi~\\citealp{Nar94}; Quataert \\& Narayan~\\citealp{Qua99}; Abramowitcz~\\etal~\\citealp{Abr02})  One especially suited radio sample to study the phenomenon of low-luminosity nuclear activity, possibly related to low radiation/accretion processes and/or radio-intermediate/quiet QSOs, is the so-called ATESP 5~GHz survey, carried out with the Australia Telescope Compact Array (ATCA) by Prandoni~\\etal~\\cite{Pra06}, which covers a $2\\times 0.5$ sq.~degr. region of the wider original 1.4 GHz ATESP survey \\citep[Prandoni~\\etal][]{Pra00a,Pra00b,Pra01a}. The ATESP 5~GHz survey consists of 131 radio sources with flux densities larger than $S\\sim 0.4$ mJy detected either at 1.4~GHz or 5~GHz or both. Interestingly, the analysis of the 1.4 and 5~GHz ATESP surveys has revealed a significant flattening of the source radio spectra going from mJy to sub-mJy flux levels (Prandoni~\\etal\\citealp{Pra06}). Such flattening is mostly associated to early-type galaxies as demonstrated by deep ($R<25$) UBVRIJK multi-colour imaging. At the flux limit of the sample star-forming galaxies are starting to appear, but are not yet the dominant population: $\\sim 14\\%$ of the sources are identified with broad/narrow-line AGNs and another $\\sim 64\\%$ is identified with early-type galaxies, most probably triggered by an AGN \\cite[see Mignano~\\etal][]{Mig08}. The absence of emission lines in the latter together with low radio luminosities (typically $10^{22-25}$ W/Hz) suggests that these are FRI radio galaxies. However, $>60\\% $ of them show flat ($\\alpha >-0.5$, where $S\\sim \\nu^{\\alpha}$) and/or inverted ($\\alpha >0$) spectra and rather compact linear sizes ($d<10-30$ kpc, \\citealp[see Mignano~\\etal][]{Mig08}). Both compactness and spectral shape suggest a core emission with strong synchrotron or free-free self-absorption. Such sources are known to exist among FRI radio galaxies, but they are relatively rare. It is therefore important to confirm the existence of such a large class of flat/inverted spectrum objects through simultaneous radio observations at different radio wavelengths.  If real, these sources may represent a composite class of objects very similar to the so-called low power ($P_{408 MHz}<10^{25.5}$ W/Hz) compact ($<10$ kpc) - LPC - radio sources studied by Giroletti~\\etal~\\citep{Gir05}. LPC host galaxies do not show signatures of strong nuclear activity in the optical (and X-ray) bands, and preliminary results indicate that multiple causes can produce LPC sources: geometrical-relativistic effects (low power BL-Lacertae objects), youth (GPS-like sources), instabilities in the jets, frustration by a denser than average ISM and a premature end of nuclear activity (sources characterised by low accretion/radiative efficiency, i.e. ADAF/ADIOS systems). Some difficulties remain however. First, if the objects are young GPS sources, much higher ($>10^{25}$ W/Hz) radio luminosities are expected; second, if they are old ADAF/ADIOS sources, the expected luminosities are much lower ($<10 ^{21}$ W/Hz, Doi \\etal \\citealp{Doi05}). In the latter case there may yet be different solutions, as for example if the ADAF source coexists with a radio jet, as was suggested by Doi \\etal \\citep{Doi05} in the case of a few low luminosity AGNs \\cite[see also the ADAF-jet model of Falcke \\& Biermann][]{Fal99}.  Another possibility is that we are dealing with radio-quiet/intermediate QSOs, where the activity in the optical band is obscured by dust in the galaxy. This scenario may be supported by the results reported by Kl\\\"ockner~\\etal~\\citep{Klo09}, who have recently observed with the EVN 11 $z>2$ radio-intermediate obscured quasars (several with flat-spectrum), detecting seven of them. The detected radio emission accounts for 30-100\\% of the entire source flux density, and the physical extent of this emission is $\\lsimeq 150$ pc. The missing flux implies the likely existence of radio jets of physical sizes between $\\gsimeq 150$ pc and $\\lsimeq 40$ kpc.  The aim of the present paper is a better understanding of the physical properties of the AGNs associated to early-type galaxies in the ATESP 5~GHz sample. The study of the radio spectra, based on multi-wavelength (simultaneous) observations, is essential in this respect, since different source types ({\\it e.g.} GPS vs BL Lac) and/or accretion scenarios ({\\it e.g.} jet-dominated vs ADAF/ADIOS-dominated) lead to different shapes of the radio spectrum. Particularly useful is the availability of high radio frequency information, since, for example the location of the frequency at which the radio spectrum peaks may be an indication of the presence or absence of outflows. A pure ADAF model without outflow is expected to peak in the sub-mm (the \"sub-mm bump\") and has an inverted ($\\alpha>0$) spectrum in the cm range \\cite[see e.g. Nagar~\\etal][]{Nag01}. As outflows become important, the peak shifts to longer wavelengths (above 1 cm), and depending on which frequencies were observed, the spectra might be classified as flat ($\\alpha>-0.5$) or steep ($\\alpha < -0.5$).  In this paper we report on quasi-simultaneous observations at 4.8, 8.6 and 19 GHz of a subset of ATESP 5~GHz sources associated to early-type galaxies. Our results will be compared to the ones obtained by new high frequency ($>10$ GHz) radio surveys. Such surveys have started only recently, notably with the ATCA at 20~GHz (Ricci~\\etal~\\citealp{Ric04}; Sadler~\\etal~\\citealp{Sad06}; Massardi~\\etal~\\citealp{Mas08}), and detailed spectral studies based on high radio frequency observations are therefore still scarce. Other deeper recent high frequency surveys are the one by Tucci \\etal \\cite{Tuc08}, who observed sources with $S > 20$ mJy at 33 GHz and compared the composition of the source counts with the high frequency model counts of  De Zotti \\etal \\cite{Dez05}, and the 9C survey at 15 GHz (a 124 source sample complete to 25 mJy and a 70 source sample complete to 60 mJy, Bolton \\etal \\citealp{Bol04}). In the latter case the observed source counts were discussed in Waldram \\etal \\cite{Wal07}. The data presented in this paper represent the first high frequency ($>10$ GHz) systematic follow-up of a sample of sub-mJy radio sources.  The paper is organized as follows. We discuss the selection of the sample, the follow-up observations at 4.8, 8.6 and 19 GHz and the data reduction in Sects.~\\ref{sec:sample} and ~\\ref{sec:radiodata}; source variability at 4.8 GHz (where data at two epochs are available) is discussed in Sect.~\\ref{sec:var}; a complete discussion of the radio spectra, in comparison with high-frequency selected samples, is given in Sect.~\\ref{sec:results}. A further comparison with (FRI and FRII) radio galaxies and radio-quiet QSOs is presented in Sect.~\\ref{sec:pdim}, where we discuss source linear sizes and radio powers. A brief summary of the main results of this work is given in Sect.~\\ref{sec:sum}. Throughout this paper we use the $\\Lambda$CDM model, with $H_0=70$, $\\Omega_m=0.3$ and $\\Omega_{\\Lambda}=0.7$. ",
        "conclusions": "\\label{sec:results}  \\begin{figure*}[t] \\includegraphics[scale=0.4]{alphaplot.zlow.ps}\\includegraphics[scale=0.4]{alphaplot.zhigh.ps}  \\caption{ High-- against low--frequency spectral index. Red symbols indicate ATESP sources: filled symbols are used for $\\alpha_H$ derived between 8.6 and 19 GHz and empty symbols are used for $\\alpha_H$ derived between 4.8 and 19 GHz; triangles indicate upper/lower limits in $\\alpha_H/\\alpha_L$ respectively. Blue diamonds indicate sources associated with early type optical spectra from  the Massardi \\etal \\cite{Mas08} sample; blue dots indicate sources associated with quasars or quasar candidates in the Massardi \\etal \\cite{Mas08} sample. The typical error box for the Massardi \\etal sources is indicated in the upper left corner. {\\it Left panel:} objects with redshift $< 0.5$. {\\it Right panel:} objects with either redshift $>0.5$, or unknown redshift. The dashed line represents the location of sources with power-law spectra (i.e. with $\\alpha_H=\\alpha_L$). Horizontal and vertical dotted lines respectively indicate the $\\alpha_H= 0$ and $\\alpha_L= 0$ lines defining the four quadrants discussed in the text.} \\label{fig:alfapanel} \\end{figure*}  Figure~\\ref{fig:fluxpanel} shows the overall radio spectra of the 26 ATESP 5~GHz sources associated with early-type galaxies discussed in this work. Red circles represent non-simultaneous ATESP 1.4 and 5 GHz observations, while blue diamonds represent the new quasi-simultaneous observations at 4.8, 8.6 and 19 GHz. We notice that low frequency spectral behaviour (steep or flat) derived from 1.4 and 5 GHz ATESP data is generally confirmed by high frequency simultaneous observations. This can be considered as a first qualitative confirmation of the fact that the sub-mJy flat-/inverted-spectrum population associated with early-type galaxies revealed by the ATESP 1.4 and 5~GHz surveys (Prandoni~\\etal\\citealp{Pra06}; Mignano~\\etal\\citealp{Mig08}) and discussed in Sect.~\\ref{sec:introduction} is genuine. The most notable exception is represented by the {\\it variable} source J225344-401928 (see Sect.~\\ref{sec:var}), which is remarkably characterised by a very inverted spectrum between 1.4 and 5 GHz ($\\alpha_L=1.4$), not confirmed as real by high frequency simultaneous observations ($\\alpha_M\\sim \\alpha_H\\sim 0$).  For a more detailed analysis of the radio spectra we will focus below on two spectral indices, $\\alpha_L$ and $\\alpha_H$ (listed in Table~\\ref{tab:atesp fluxes}). These two spectral indices can be used in a sort of \"two-colour\" diagram as was done in previuos works (see Sadler \\etal \\citealp{Sad06}; Massardi et al \\citealp{Mas08}; Tucci \\etal \\citealp{Tuc08}). Note that the analysis of the two--colour diagram is not exhaustive, and the spectral index between 4.8 and 8.6 GHz (e.g. $\\alpha_M$) may give further independent information on the spectrum, as previously noted in Sect.~\\ref{sec:var}.  We will also compare our data with recent high frequency surveys (Sadler \\etal \\citealp{Sad06}; Massardi \\etal \\citealp{Mas08}; Tucci \\etal \\citealp{Tuc08}) and will try to determine how and to what extent frequency selection and flux limits influence the results. The ATESP 5~GHz survey is the only one selected at low frequency, while the other three are based on $20-30$ GHz observations. Also there is quite a difference in the flux limit, which reaches from $\\sim 0.6$ mJy (our sample) via 20 mJy at 33 GHz (Tucci et al. \\citealp{Tuc08}), 100 mJy at 20 GHz (Sadler \\etal \\citealp{Sad06}) to 500 mJy at 20 GHz (Massardi \\etal~\\citealp{Mas08}). One extra selection criterion used for our sample was that we followed-up at high frequencies only sources associated with galaxies having early-type spectra; for that reason we also defined a sub-sample of 36 sources associated with galaxies from the Bright Source Sample discussed by Massardi \\etal \\cite{Mas08}. Since the latter sources are all bright in the radio, these are presumably all classical radio galaxies; their typical radio powers put most of them effectively in the FRII range.  Figure \\ref{fig:alfapanel} shows two--colour diagrams of our (red symbols) and the Massardi~\\etal galaxy (blue diamonds) samples. Also plotted are all sources from the Massardi~\\etal sample that are associatd with quasars or quasar candidates (blue dots). We divided the sources in a nearby ($z <0.5$, left panel) and a far ($z > 0.5$, right panel) sample, in order to check if distance may have an effect upon the spectral indices. In the far sample we included quasar candidates with unknown redshift, the large majority of which (presumably) are at redshifts $>0.5$ on the basis of their optical magnitude distribution.  For a physical interpretation of the two-colour diagram we divided it in four quadrants, as done in previous works: 1) sources with $\\alpha_L<0, \\; \\alpha_H<0$, conventionally named {\\it steep}\\footnote{In fact in this quadrant fall both steep ($\\alpha<-0.5$) and flat ($-0.5<\\alpha<0$) sources}; 2) sources with $\\alpha_L<0, \\; \\alpha_H>0$, named {\\it upturn}; 3) sources with $\\alpha_L>0, \\; \\alpha_H>0$, named {\\it inverted}; 4) sources with $\\alpha_L>0, \\; \\alpha_H<0$, named {\\it peaked}. We point out that the radio spectral index of the overall synchrotron emission from a jet-dominated radio source is expected to be from steep to moderately inverted ($-0.7<\\alpha<0.2$) at the frequencies of this work, depending on the relative contributions of extended (optically-thin) and base (self-absorbed) jet components. Jet-dominated sources are therefore expected to mostly fall into the {\\it steep} quadrant. Instead, for an ADAF, one expects $0.2 < \\alpha < 1.1$ up to mm wavelengths (see Nagar~\\etal\\citealp{Nag01}), with $\\alpha$ varying according to the accretion rate (f.i. $\\alpha = 0.4$ when $L\\sim 10^{-4}L_{\\rm{Edd}}$ and $\\alpha \\sim 1$ when $L\\leq 10^{-7}L_{\\rm{Edd}}$; Mahadevan~\\citealp{Mah97}). ADAF systems are therefore expected to fall into the {\\it inverted} quadrant. Co-existence of moderate outflows can flatten the spectrum (with $\\alpha$ remaining positive), whereas strong outflows can shift the peak of the radio emission from mm to cm wavelengths (Quataert \\& Narayan~\\citealp{Qua99}).  As can be seen in Fig.~\\ref{fig:alfapanel} our ATESP sources, as well as the sources from the much brighter Massardi~\\etal galaxy sample, are mostly located in  the lower part of the plots: 23/24 and 33/36 sources for the ATESP and the Massardi sample respectively have $\\alpha_H <0.2$. Although these numbers are significant only at about the $3\\sigma$ level due to the relatively poor statistics involved, there can be little doubt that both samples consistently contain more sources with $\\alpha_H < 0.2$. This is of course hardly surprising for the ATESP sample, which was originally selected at 5~GHz, but the high frequency selection of the Massardi sample does not make much difference, at least in this respect. Less obvious a priori is the general steepening ($\\alpha_L>\\alpha_H$) of the spectra from low to high frequency. This is evident from the large number of sources lying below the diagonal dashed line, corresponding to power-law spectra (19/24 and 31/36 for ATESP and Massardi). These numbers are significant only at the $2\\sigma$ level, but again both samples behave consistently and actually appear to cover identical regions in the two-colour diagram. Both facts favour a jet-dominated scenario, with steep-spectrum ($\\alpha_H<-0.5$) sources dominated by synchrotron optically-thin emission and flat-spectrum ($\\alpha_H>-0.5$) sources dominated by synchrotron optically-thick emission coming from the base of the jet. Nevertheless jet+ADAF models can also be consistent with values $0\\lsimeq\\alpha_H \\lsimeq 0.2$. Pure ADAF models ($\\alpha_{L}$ and $\\alpha_{H}\\ge 0.2$) are ruled out by the present data (only one object in the Massardi galaxy sample can be considered as a pure ADAF candidate), indicating that such radiatively inefficient accretion regimes must be very rare also at the sub-mJy levels probed by the ATESP 5~GHz sample.  \\begin{table}[t] \\centering \\caption{Mean values and standard deviations of the low and high frequency spectral indices. } \\label{tab:means} \\begin{tabular}{lrrrr} \\hline \\hline Sample$^{(a)}$ & $\\alpha_L$ & $\\sigma$ & $\\alpha_H$ & $\\sigma$   \\\\ \\hline ATESP          &  $-0.18$ & (0.11)  & $-0.79$ & (0.16) \\\\ Massardi G  &  $-0.01$ & (0.08)  & $-0.35$ & (0.09) \\\\ ATESP + Massardi G far    &  $-0.08$  & (0.07)  & $-0.50$ & (0.08) \\\\ Massardi Q far     &  $0.21$  & (0.04)  & $-0.16$ & (0.03) \\\\ \\hline \\multicolumn{5}{l}{\\footnotesize $^{(a)}$ 'Massardi G' and 'Massardi Q' respectively refer to galaxies and}\\\\ \\multicolumn{5}{l}{\\footnotesize \\hspace{0.3cm} stellar identifications of the Massardi~\\etal sample. 'Far' refers to }\\\\ \\multicolumn{5}{l}{\\footnotesize \\hspace{0.3cm} $z>0.5$ sources (or unknown redshift for quasar candidates)}\\\\ \\end{tabular} \\end{table}  For a more quantitative analysis of the two-colour diagram we statistically compared both the low-- and high--frequency spectral index distribution of the ATESP and the Massardi \\etal galaxy samples by implementing the survival analysis methods described in Feigelson \\& Nelson \\cite{Fei85} and Isobe \\etal \\cite{Iso86}. This allows us to properly treat upper/lower limits in $\\alpha_H/\\alpha_L$ respectively. Both $\\alpha_L$ and $\\alpha_H$ spectral index distributions are consistent with being drawn from the same parent population. In fact this hypothesis is rejected at only a $<2\\sigma$ confidence level (Peto \\& Prentice generalized Wilcoxon test, Isobe \\& Feigelson \\citealp{Iso90}). This is especially interesting when we consider that the ATESP and  the Massardi~\\etal samples cover very different ranges in flux density (by two or three orders of magnitude).  A similar comparison can be made between galaxies and quasars. To improve our statistics we consider the ATESP and the Massardi \\etal galaxies as a single population, and we limit our analysis to $z>0.5$ sources, where most quasars (or quasar candidates) are located (see right panel of Fig.~\\ref{fig:alfapanel}). In this case we find that the hypothesis that $z>0.5$ galaxies and quasars are extracted from the same parent population is rejected at a $>4.5\\sigma$ level (Peto \\& Prentice generalized Wilcoxon test, Isobe \\& Feigelson \\citealp{Iso90}). Table~\\ref{tab:means} summarises the mean values for the samples discussed above (Kaplan-Meier estimator, Isobe \\& Feigelson \\citealp{Iso90}). We notice that the Massardi sub-sample of $z>0.5$ quasars and quasar candidates (Massardi Q far) tend to be flatter or even more inverted than sources associated with galaxies, indicating that quasar-like objects are more commonly associated with compact sources and less with the classical extended radio lobes in which the optically thin synchrotron radiation with steep radio spectra is dominant. On the other hand, no difference is seen in spectral steepening: if we compare the curvature $C=\\alpha_H - \\alpha_L$ (see e.g. Gregorini \\etal \\citealp{Gre84}), we find $C=-0.42 \\pm 0.11$ (ATESP + Massardi G far), and $C=-0.37 \\pm 0.05$ (Massardi Q far).  Following Sadler \\etal \\cite{Sad06} we consider in more detail the distribution of sources in the four quadrants of the two-colour diagram, which represent different types of radio sources (defined earlier in this section). The distribution of objects over the four quadrants is shown in Table~\\ref{tab:twocolour} (where U stands for {\\it Upturn}, I for {\\it Inverted}, S for  {\\it Steep} , P for {\\it Peaked}). Most ATESP sources are in the lower left quadrant (61\\%, Table~\\ref{tab:twocolour}), as for classical FRI and FRII radio galaxies. This may be due to the early type galaxy optical pre-selection, and in fact the entries of the Massardi et al. galaxy sample (M-G in Table~\\ref{tab:twocolour}) are, considering the errors, practically the same as those of the ATESP sample, despite the very different flux range covered by the two samples.  In Table~\\ref{tab:twocolour} we also compare the distribution of objects in three samples, all selected at high ($\\geq 20$ GHz) frequencies, where no optical pre-selection was applied (Tucci~\\etal~\\citealp{Tuc08}; Sadler~\\etal~\\citealp{Sad06}; Massardi~\\etal~\\citealp{Mas08}). There appear to be some trends with flux density, which we can attribute partly to selection effects and partly to the different composition of the overall samples, in which quasar spectra may be more or less dominant, depending on the limiting flux density. The {\\it upturn} and {\\it steep} spectra (i.e. the left hand side of the two-colour diagram) tend to become more numerous at lower flux densities, while the inverse happens with {\\it inverted} and {\\it peaked} spectra (right hand side). These trends are individually significant only at the $2\\sigma$ level, again due to the small numbers involved, but appear to be quite consistent in all four quadrants. If we sum the $U$ and $S$ sources (those with $\\alpha_L < 0$) and the $I$ and $P$ sources ($\\alpha_L > 0$), the dependence on the flux density becomes quite obvious: whereas about 18 \\% of the sources have a positive $\\alpha_L$ ($I$ and $P$) at the flux density level of the Tucci \\etal (\\citealp{Tuc08}) sample, at the much higher flux levels of the Massardi \\etal sample this strongly increases to about 60 \\%; this difference is definitely significant at about 4$\\sigma$. The fact that the number of rising spectra increases with flux limit was already noticed by Bolton \\etal \\cite{Bol04}. No doubt these trends may be due to selection effects; however, the possible increase of {\\it peaked} sources at higher flux densities is harder to explain by selection effects only.  \\begin{table}[t] \\centering \\caption{Spectral type statistics in five samples$^{(a)}$. } \\label{tab:twocolour} \\begin{tabular}{l|rrrr|rrrrrr} \\hline \\hline & \\multicolumn{2}{c}{ATESP} & \\multicolumn{2}{c|}{M-G} & \\multicolumn{2}{c}{Tuc} & \\multicolumn{2}{c}{Sad} & \\multicolumn{2}{c}{M}    \\\\ & \\multicolumn{2}{c}{\\scriptsize $>0.6$ mJy} & \\multicolumn{2}{c|}{\\scriptsize $>500$ mJy} & \\multicolumn{2}{c}{\\scriptsize $>20$ mJy} & \\multicolumn{2}{c}{\\scriptsize $>100$ mJy} & \\multicolumn{2}{c}{\\scriptsize $>500$ mJy}  \\\\ & \\multicolumn{4}{c|}{(percent)} & \\multicolumn{6}{c}{(percent)} \\\\ \\hline U &   9 & $\\pm$5  & 8 & $\\pm$5 & 29 & $\\pm$6 & 22 & $\\pm$5 &  10 & $\\pm$1 \\\\ I  & 13 &   7 & 12 &   6 &   7 &  3 & 20 &   4 & 28 &   5  \\\\ S & 61 & 14 & 47 & 11 & 53 &  7 & 34 &   6 & 30 &   5 \\\\ P & 17 &   8 & 33 & 11 & 11 &  3 & 24 &   5 & 32 &   5 \\\\ \\hline \\multicolumn{11}{l}{$^{(a)}$ ATESP 5~GHz sample (26 sources);  Massardi~\\etal galaxy}\\\\ \\multicolumn{11}{l}{\\hspace{0.4truecm} sample (M-G, 36 sources);  Tucci~\\etal\\cite{Tuc08} sample (Tuc,}\\\\ \\multicolumn{11}{l}{\\hspace{0.4truecm} 102 sources);  Sadler~\\etal\\cite{Sad06} sample (Sad, 101 sources)}\\\\ \\multicolumn{11}{l}{\\hspace{0.4truecm} Massardi~\\etal\\cite{Mas08} overall sample (M, 218 sources)}\\\\ \\end{tabular} \\end{table}   \\label{sec:sum}  We presented follow up quasi-simultaneous observations at 4.8, 8.6 and 19 GHz for a representative sub-sample of ATESP sources associated to early type galaxies (26 sources with $S>0.6$ mJy). The new data were discussed together with the original 1.4 and 5 GHz ATESP data (Prandoni \\etal \\citealp{Pra00b,Pra06}), with the aim of better understanding the radio spectral properties of such sources and ultimately the nature of AGNs at sub-mJy flux levels. In particular we were interested in understanding whether the AGN component of the sub--mJy population is more related to efficiently accreting systems - like radio-intermediate/quiet quasars - or to systems with low accretion rates - like e.g. FRI radio galaxies - or to low radiative efficiency accretion flows - like e.g. ADAF.  From the analysis of the radio spectra, we found that our sources are most probably jet-dominated systems, with steeper sources dominated by optically-thin synchrotron emission (typical of FRI and FRII radio galaxies) and flatter or moderately-inverted sources dominated by optically-thick emission coming from the base of the jet. Pure ADAF models seem to be ruled out by the high frequency data, while ADAF+jet scenarios could be consistent with flat/moderately inverted sources, but are not required to explain the data.  By comparing our sample with some recent high frequency ($\\gsimeq 20$ GHz) surveys we found strong spectral similarities between our and the much brighter Massardi~\\etal galaxy sample, possibly suggesting that our radio sources are mostly lower luminosity counterparts of the Massardi~\\etal (FRII) radio galaxies. In addition {\\it peaked} GPS-like sources seem to be preferentially found at higher flux densities than the ones probed by the ATESP sample, and quasar-like {\\it upturn} sources are rare ($9\\pm 5\\%$). This latter result is partly due to the ATESP early type galaxy pre-selection, but argues against an obscured radio-quiet quasar population hidden among the ATESP early type galaxies. Nevertheless compact quasar-like objects may be present among flat/moderately inverted jet-dominated sources.  To further investigate such a hypothesis we compared radio powers and linear sizes of the ATESP 5~GHz sample with the ones of two well-known brighter radio source catalogues (B2 and 3CR) and with the sample of radio-quiet quasars studied by Leipski \\etal\\citep{Lei06}. Unfortunately such a comparison suffers from the large number of unresolved ATESP sources (or the large number of size upper limits). A significant fraction of ATESP sources have sizes consistent with the ones of B2 and 3CR radio galaxies. Such sources seem to confirm and extend to lower radio luminosity the size-power relation found for B2 radio galaxies (de Ruiter \\etal\\citealp{Rui90}). On the other hand, a component of very compact $d<10$ kpc radio sources associated to early type galaxies seem to have sizes more consistent with the ones of radio-quiet quasars. Among them we may have a few low power BL Lac. For a more conclusive analysis we need higher resolution (possibly VLBI) data for the ATESP sample."
    },
    "0911/0911.5248_arXiv.txt": {
        "abstract": "Gamma-ray bursts (GRBs) are the most luminous astrophysical events observed so far. They are conventionally classified into long and short ones depending on their time duration, $T_{90}$. Because of the advantage their high redshifts offer, many efforts have been made to apply GRBs to cosmology. The key to this is to find correlations between some measurable properties of GRBs and the energy or the luminosity of GRBs. These correlations are usually referred to as luminosity relations and are helpful in understanding the GRBs themselves. In this paper, we explored such correlations in the X-ray emission of GRBs. The X-ray emission of GRBs observed by \\emph{Swift} has the exponential functional form in the prompt phase and relaxes to a power-law decay at time $T_p$. We have assumed a linear relation between $\\log L_{X, p}$ (with $L_{X, p}$ being the X-ray luminosity at $T_p$) and $\\log [T_p/(1+z)]$, but there is some evidence for curvature in the data and the true relationship between $L_{X, p}$ and $T_p/(1+z)$ may be a broken power law. The limited GRB sample used in our analysis is still not sufficient for us to conclude whether the break is real or just an illusion caused by outliers. We considered both cases in our analysis and discussed the implications of the luminosity relation, especially on the time duration of GRBs and their classification. ",
        "introduction": "Gamma-ray bursts (GRBs), which can last from milliseconds to nearly an hour, are the most luminous astrophysical events observed so far. The parameter of $T_{90}$, which is defined as the time interval during which the background-subtracted cumulative counts increase from $5\\%$ to $95\\%$, is usually used to denote the time duration of GRBs. Those with $T_{90} > 2 \\, \\mathrm{s}$ are conventionally described as long/soft GRBs and those with $T_{90} < 2 \\, \\mathrm{s}$ as short/hard GRBs~\\citep{Kouveliotou:1993yx}.  In~\\citet{Willingale:2006zh}, it was demonstrated that the X-ray decay curves of GRBs measured by \\emph{Swift} can be fitted using one or two components---the prompt component and the optional afterglow component---both of which have exactly the same functional form comprised of an early falling exponential phase and a following power-law decay. The prompt component contains the prompt gamma-ray emission and the initial X-ray decay. The transition time $T_p$ from the exponential phase to the power-law decay in the prompt component defines an alternative estimate of the GRB duration, which is comparable with $T_{90}$ for most GRBs~\\citep{O'Brien:2006sa, O'Brien:2006jz}. \\citet{O'Brien:2007apss} proposed the classification of GRBs into long and short ones at $T_p = 10 \\, \\mathrm{s}$ instead.  GRBs can be observed at very high redshifts due to their high luminosities. For example, the recently observed GRB 090423 has a redshift of $z \\approx 8.2$~\\citep{Tanvir:2009ex, Salvaterra:2009ey}. It may be possible to calibrate GRBs as standard candles~\\citep[see, for example,][etc.]{Dai:2004tq, Ghirlanda:2004fs, Firmani:2005gs, Lamb:2005cw, Liang:2005xb, Xu:2005uv, Wang:2005ic, Ghirlanda:2006ax, Schaefer:2006pa, Wang:2007rz, Li:2007re, Amati:2008hq, Basilakos:2008tp, Qi:2008zk, Qi:2008ag, Kodama:2008dq, Liang:2008kx, Wang:2008vja, Qi:2009yr}. The key to the calibration of GRBs is to find correlations between some measurable properties (the luminosity indicators) of GRBs and the energy (the isotropic energy $E_{\\gamma, \\mathrm{iso}}$ or the collimation-corrected energy $E_{\\gamma}$) or the luminosity (e.g., the peak luminosity $L$) of GRBs. These correlations are usually referred to as luminosity relations, which are useful both in applying GRBs to cosmology and understanding GRBs themselves. The GRB luminosity relations used in cosmological studies in the literature include the relations of $\\tau_{\\mathrm{lag}}$ (the spectral lag, i.e., the time shift between the hard and soft light curves)--$L$~\\citep{Norris:1999ni}, $V$ (the variability, a quantitative measurement on the spikiness of the light curve; there exist several definitions of $V$, mainly depending on the smoothing time intervals the reference curve is built upon, and on the normalization as well)--$L$~\\citep{Fenimore:2000vs, Reichart:2000kq}, $E_{\\mathrm{peak}}$ (peak energy of the spectrum)--$E_{\\gamma, \\mathrm{iso}}$~\\citep{Amati:2002ny}, $E_{\\mathrm{peak}}$--$E_{\\gamma}$~\\citep{Ghirlanda:2004me}, $E_{\\mathrm{peak}}$--$L$~\\citep{Schaefer:2002tf}, and $\\tau_{\\mathrm{RT}}$ (minimum rise time of the light curve)--$L$~\\citep{Schaefer:2006pa}. For each of the above luminosity relations, there is only one luminosity indicator involved. More complicated luminosity relations, which include two luminosity indicators are also discussed in the literature; see, for example, \\citet{Yu:2009xd} and references therein.  In this paper, we explore the correlation between $T_p$ and the X-ray luminosity of GRBs at $T_p$ and discuss its implications, especially on the time duration of GRBs and their classification. ",
        "conclusions": "We plot the logarithm of $L_{X, p}$ versus the logarithm of $T_p/(1+z)$ in Figure~\\ref{fig:Lp_vs_Tpz}, which include $47$ GRBs. We can see that most of the GRBs ($34$ GRBs) lie in the range of $2 \\, \\mathrm{s} < T_p/(1+z) < 100 \\, \\mathrm{s}$. There is obviously a correlation between $L_{X, p}$ and $T_p/(1+z)$. However, when fitting the GRBs to the relations of Eq.~(\\ref{eq:luminosity_relation}), we have different options depending on how we view the three GRBs with the largest $T_p/(1+z)$, i.e., those with $T_p/(1+z) > 100 \\, \\mathrm{s}$. For the first choice, we can simply include all the data points in the fit (see the top panel of Figure~\\ref{fig:fit_result}), which leads to a result that the GRB with the largest $T_p/(1+z)$ lies outside the $2 \\sigma$ confidence region of the fit. Alternatively, it is also possible that the GRBs with the largest $T_p/(1+z)$, instead of just being outliers, may indeed reveal some trend of the luminosity relation at large $T_p/(1+z)$. In this case, we cannot simply ignore the GRBs with large $T_p/(1+z)$ just because their quantity is small and, if they are taken more seriously, it seems that the samples are split into two groups at some value of $T_p/(1+z)$ based on the slope of the luminosity relation of Eq.~(\\ref{eq:luminosity_relation}). To show this, we perform a fit using only the GRBs with $T_p/(1+z) > 2 \\, \\mathrm{s}$ (see the bottom panel of Figure~\\ref{fig:fit_result}), which gives a result quite different from the prior fit. From the bottom panel of Figure~\\ref{fig:fit_result}, we can see that if the best fit line is extended to the range of $T_p/(1+z) < 2 \\, \\mathrm{s}$, all the GRBs with $T_p/(1+z) < 2 \\, \\mathrm{s}$ lie below the line. As a comparison, we also fit the GRBs with $2 \\, \\mathrm{s} < T_p/(1+z) < 100 \\, \\mathrm{s}$ (see Figure~\\ref{fig:fit_result_subl}), which show that the difference is indeed introduced by the three GRBs with the largest $T_p/(1+z)$ when compared with the fit of GRBs with $T_p/(1+z) > 2 \\, \\mathrm{s}$. We tabulate the fit results in Table~\\ref{tab:fit_result}. From the table, we can see that the values of $b$ for the first two cases are quite different. In addition, it is also interesting to note that, for the first case, the slope $b$ is close to the slope ($-0.74_{-0.19}^{+0.20}$) of a similar luminosity relation about $T_a$ and luminosity at $T_a$ presented in~\\citet{Dainotti:2008vw}. For the second case, the slope $b$ is close to the index ($-1.5 \\pm 0.16$) of the power-law declination of the average luminosity of X-ray flares as a function of time presented in~\\citet{Lazzati:2008da}. These coincidences may be worth noting in future studies with a bigger sample of GRBs.  \\begin{table}[htbp] \\centering {\\tiny \\begin{tabular}{cccc} \\hline \\hline GRB set & $a$ & $b$ & $\\sigma_{\\mathrm{int}}$ \\\\ \\hline All available GRBs & $50.91 \\pm 0.23$ & $-0.89 \\pm 0.19$ & $1.06 \\pm 0.13$ \\\\ $\\frac{T_p}{1+z} > 2 \\, \\mathrm{s}$ & $51.96 \\pm 0.32$ & $-1.73 \\pm 0.25$ & $0.78 \\pm 0.11$ \\\\ $2 \\, \\mathrm{s} < \\frac{T_p}{1+z} < 100 \\, \\mathrm{s}$ & $51.09 \\pm 0.32$ & $-0.74 \\pm 0.30$ & $0.63 \\pm 0.09$ \\\\ \\hline \\end{tabular} } \\caption { Results of the fit to the luminosity relation of Eq.~(\\ref{eq:luminosity_relation}). } \\label{tab:fit_result} \\end{table}  \\begin{figure}[htbp] \\centering \\includegraphics[width = 0.45 \\textwidth]{Lp_vs_Tpz.eps} \\caption { $L_{X, p}$ (in $\\mathrm{erg} \\, \\mathrm{s}^{-1}$) versus $T_p/(1+z)$ (in seconds). $47$ GRBs are included with $37$ GRBs that have $T_p/(1+z)$ greater than $2$ seconds and $3$ of which have $T_p/(1+z)$ greater than $100$ seconds. The red ones are conventional short GRBs ($T_{90} < 2 \\, \\mathrm{s}$). } \\label{fig:Lp_vs_Tpz} \\end{figure}  \\begin{figure}[htbp] \\centering \\includegraphics[width = 0.45 \\textwidth]{fit_all.eps} \\\\ \\includegraphics[width = 0.45 \\textwidth]{fit_long.eps} \\caption { Best fit of the GRBs to the luminosity relation of Eq.~(\\ref{eq:luminosity_relation}) and the corresponding $2 \\sigma$ confidence region. Top: all GRBs are included in the fit. Bottom: only those with $T_p/(1+z) > 2 \\, \\mathrm{s}$ are included in the fit. } \\label{fig:fit_result} \\end{figure}  \\begin{figure}[htbp] \\centering \\includegraphics[width = 0.45 \\textwidth]{fit_subl.eps} \\caption { Best fit of the GRBs to the luminosity relation of Eq.~(\\ref{eq:luminosity_relation}) and the corresponding $2 \\sigma$ confidence region. Same as Figure~\\ref{fig:fit_result} except that only GRBs with $2 \\, \\mathrm{s} < T_p/(1+z) < 100 \\, \\mathrm{s}$ are included in the fit. } \\label{fig:fit_result_subl} \\end{figure}  Generally speaking, in statistical analysis, except for the different measurement precisions, we should treat all the data points equally and should not give more attention to some data points over the others. However, since our sample of GRBs here is very limited, and there are only a few GRBs with very large $T_p/(1+z)$, some selection rules may have already been imposed implicitly on the sample itself. Considering this, it may be unfair to the GRBs with large $T_p/(1+z)$ to be treated as outliers just because their quantity is small, especially when they may reveal some trend in the luminosity relation. This is why we consider both cases in the analysis, i.e., whether the few GRBs with large $T_p/(1+z)$ should be treated as just outliers or taken more seriously. Correspondingly, the relation between $L_{X, p}$ and $T_p/(1+z)$ could be a simple power law or a broken power law with a change in the slope of Eq.~(\\ref{eq:luminosity_relation}) at some characteristic value of $T_p/(1+z)$. For the present sample of GRBs used in our analysis, it is not sufficient for us to conclude which case is real and, for the later one, to determine the exact value of $T_p/(1+z)$ where the slope of the luminosity relation changes. Here, we leave it open to future studies with more GRBs and discuss in the following the implications of the luminosity relation, especially in the situation where there is a change in the slope at some value of $T_p/(1+z)$.  First of all, we emphasize that the luminosity relation is between the luminosity and $T_p$, though $T_p$ and $T_{90}$ can both act as an estimate of the GRB duration and are comparable to each other for most GRBs, as can be seen from Figure~\\ref{fig:Tp_vs_T90}. A similar relation seems not to exist between the luminosity and the $T_{90}$. See Figure~\\ref{fig:Lp_vs_T90z}; the corresponding data points turn out to be very dispersive. Despite the fact that most GRBs in Figure~\\ref{fig:Tp_vs_T90} are distributed around the line on which they are equal, there are some of them that have been considerably different from $T_p$ and $T_{90}$. In fact, the two quantities differ from each other significantly from their derivation. $T_{90}$ is calculated directly by using the BAT data in the $15$--$150$ keV band, while for the calculation of $T_p$, the BAT data are first extrapolated to the XRT band of $0.3$--$10$ keV in order to be combined with the XRT data. For $T_{90}$, the emphasis is on the percent of the fluence of a burst, while for $T_p$, the emphasis is on the transition in a GRB light curve from the exponential decay in the prompt phase to the initial power-law decay. In the time interval of $T_{90}$, $90\\%$ of the total fluence is observed, while the ratio between the observed fluence in the time interval of $T_p$ and the total fluence depends not only on the temporal decay index of the initial power-law decay duration in the prompt component, but also on the shape of the light curve in the afterglow component. In addition, we must remember that large flares have been masked out from the light curves before performing the fit that allows us to derive $T_p$~\\citep{Willingale:2006zh}.  \\begin{figure}[htbp] \\centering \\includegraphics[width = 0.45 \\textwidth]{Tp_vs_T90.eps} \\caption { $T_p$ versus $T_{90}$ for the $107$ GRBs. Times are in seconds. On the cyan dash-dotted line the two quantities have equal values and the red dashed lines correspond to $T_p$ and $T_{90}$ equal to $2$ seconds respectively. } \\label{fig:Tp_vs_T90} \\end{figure}  \\begin{figure}[htbp] \\centering \\includegraphics[width = 0.45 \\textwidth]{Lp_vs_T90z.eps} \\caption { $L_{X, p}$ (in $\\mathrm{erg} \\, \\mathrm{s}^{-1}$) versus $T_{90}/(1+z)$ (in seconds). The set of GRBs is the same as that in Figure~\\ref{fig:Lp_vs_Tpz}. The red ones are conventional short GRBs ($T_{90} < 2 \\, \\mathrm{s}$). } \\label{fig:Lp_vs_T90z} \\end{figure}  As stated above, because the data are limited, there are two possible models for the luminosity relation of Eq.~(\\ref{eq:luminosity_relation}), i.e., the luminosity relation, using one set of values for the parameters (the intercept $a$ and the slope $b$), may be applicable to all the GRBs except for some outliers. or there may be a change in the slope $b$ of the luminosity relation at some value of $T_p/(1+z)$. If there is a change in slope this may suggest that GRBs could be classified into two groups based on their values of $T_p/(1+z)$. Since $T_p/(1+z)$ is an estimate of the GRB duration, this is in fact an indication of how we should classify GRBs into long and short ones and is actually the same as the proposal by~\\citet{O'Brien:2007apss} to use $T_p$ as the criterion for the classification of long and short GRBs. In principle, we should use the quantity in the burst frame ($T_p/(1+z)$) instead of that in the observer frame ($T_p$). However, for a large portion of the observed GRBs, the redshifts are not available. Generally speaking, due to the diversity of GRBs, the classification of a GRB is unlikely to be completely determined by only its time duration, not to mention that the time duration of long and short GRBs overlaps near the demarcation point. Let us assume the demarcation point in the burst frame for the long and short GRBs to be \\begin{equation} \\label{eq:dempoint_burst} T_p/(1+z) = T_{l/s} . \\end{equation} Then, as an approximate method, an effective redshift $z_{\\mathrm{eff}}$ can be defined, such that GRBs can be classified into long and short ones in the observer frame at \\begin{equation} \\label{eq:dempoint_observer} T_p = (1 + z_{\\mathrm{eff}}) T_{l/s} . \\end{equation} In addition, since a similar relation does not hold if we replace $T_p$ with $T_{90}$ as mentioned previously, the change in the slope $b$, which may be a reflection of different mechanisms, if confirmed, would favor $T_p$ over $T_{90}$ as a criterion for the classification of long and short GRBs."
    },
    "0911/0911.0001_arXiv.txt": {
        "abstract": "Interpretations of indirect searches for dark matter (DM) require theoretical predictions for the annihilation or decay rates of DM into stable particles of the standard model. These predictions include usually only final states accessible as  lowest order tree-level processes, with  electromagnetic bremsstrahlung and the loop-suppressed two gamma-ray line as exceptions. We show that this restriction may lead to severely biased results for DM tailored to produce only leptons in final states and with mass in the TeV range. For such models, unavoidable electroweak bremsstrahlung of $Z$ and $W$-bosons has a significant influence both on the branching ratio and the spectral shape of the final state particles. We work out the consequences for two situations: Firstly, the idealized case where DM  annihilates at tree level with 100\\% branching ratio into neutrinos. For a given cross section, this leads eventually to ``minimal yields'' of photons, electrons, positrons and antiprotons. Secondly, the case where the only allowed two-body final states are electrons. The latter case is typical of models aimed at fitting cosmic ray $e^{-}$ and $e^{+}$ data. We find that the multimessenger signatures of such models can be significantly modified with respect to results presented in the literature. ",
        "introduction": "Despite the numerous cosmological and astrophysical indications for the presence of nonbaryonic dark matter (DM), the particle nature of DM has yet to be identified. One strategy towards DM ``detection\" is to search for its self-annihilation (or decay) products in our Galaxy, provided that the annihilation cross section $\\sigv$ (or the decay rate) is large enough. Restricting the space of candidates to weakly interacting massive thermal relics, $\\sigv$ (at freeze-out) is fixed by the DM abundance while the DM mass $m_X$  lies in this class of models typically within one or two orders of magnitude off the weak scale $m_W$. Nevertheless, considerable model-dependence remains because of our ignorance of the final states produced in the annihilation process, which vary from model to model.  Under the sole hypothesis that massive dark matter annihilates into standard model particles, a few years ago an interesting conservative upper bound on $\\sigv$ was derived using data on the least detectable final states, namely neutrinos~\\cite{Beacom:2006tt,Yuksel:2007ac}. (Similar considerations apply to decaying particles~\\cite{PalomaresRuiz:2007ry}, although we will not mention this further.) It is natural to ask if these conservative bounds can be improved or consolidated further. It was shown previously in Ref.~\\cite{Berezinsky:2002hq} that electroweak bremsstrahlung  leads to a break-down of perturbation theory and a non-negligible branching ratio in electromagnetic channels for $m_X\\gg m_W$, even if at tree level DM couples only neutrinos. This argument was used then in Refs.~\\cite{Kachelriess:2007aj,Bell:2008ey} to derive  constraints on the DM annihilation cross sections from diffuse gamma-ray observations, that turned out to be similarly restrictive as the original one from Ref.~\\cite{Beacom:2006tt,Yuksel:2007ac}.  More recently, the PAMELA collaboration published data showing an ``anomalous growth'' of the cosmic ray positron fraction~\\cite{Adriani:2008zr}, while the measured antiproton fraction agrees with expectations from simples models~\\cite{Adriani:2008zq}. Additionally, new data on the overall electron plus positron spectrum were presented, most notably from the Fermi space telescope~\\cite{Abdo:2009zk,Grasso:2009ma}. These data have prompted a plethora of models trying to explain the data by engineering relatively heavy ($m_X\\agt 1\\,$TeV) DM candidates with large annihilation cross sections (or decay rates) and small or vanishing branching ratios (br's) into hadronic final states. In the analysis of these models, the role of radiated $W,Z$ bosons and their phenomenological impact has been generally ignored.  The purpose of this article is to revisit the topic of  electroweak bremsstrahlung effects with several goals in mind. First of all, albeit the qualitative conclusions of Refs.~\\cite{Kachelriess:2007aj} and \\cite{Bell:2008ey} agree, their results differ quantitatively. Here, we repeat this calculations analytically and cross-check them against numerical results using Madgraph. Second and more important for its recent phenomenological appeal, similar effects also arise when the tree-level final state is a charged lepton $\\ell^{+} \\ell^{-}$, and therefore we present an analogous calculation for an $e^{+} e^{-}$ pair as final state. Third, besides photons, the fragmentation of the emitted $W$  and $Z$ bosons produces also electrons and positrons, protons and antiprotons,  as well as neutrinos and antineutrinos (hereafter, simply ``neutrinos''). Hence we can use observations of different cosmic species to derive complementary constraints on DM annihilations. Finally, we include in the present analysis new data on the antiproton~\\cite{Adriani:2008zq} and positron fraction~\\cite{Adriani:2008zr} from the PAMELA satellite as well as on diffuse gamma-rays from the Fermi Gamma-ray Space Telescope~\\cite{MA}.  This article is organized as follows: In Sec.~ \\ref{analytical} we review first some general considerations about the favoured DM annihilation modes and then our analytical calculation of the branching ratio for electroweak bremsstrahlung. In Sec.~\\ref{numerical} we describe the spectra of secondaries found numerically, while Sec.~\\ref{constraints} is devoted to the constraints from present observations. We conclude in Sec.~\\ref{conclusions}. ",
        "conclusions": "We discussed the role of electroweak bremsstrahlung for indirect dark matter signatures, which has been typically neglected in phenomenological studies heretofore. Our approach was  to calculate the branching ratios and the fragmentation spectra $\\d N_i/\\d x$ of secondaries considering electroweak radiation  only from the final state. Therefore, our results are applicable mainly to models tailored to produce only leptons as final state and with DM mass in the TeV range. In particular, our analysis applies to several models trying to match features in $e^+e^-$ CR data with TeV-scale DM, but e.g.\\ not to the benchmark case of neutralino annihilations in the MSSM (where final states containing hadrons or gauge bosons are anyway allowed at tree-level).  An  important phenomenological consequence of electroweak bremsstrahlung is its impact on the predicted photon and electron/positron spectra. Secondaries from $W$ and $Z$ decays can dominate in a certain $x$ range the spectra, leading to changes in the normalization and the shape of the secondary spectra. Thus it is mandatory to assess the importance of these ``corrections'' in a specific model under consideration, before one attempts to fit cosmic ray or photon data. Models viable at tree level may be ruled out by the more realistic spectrum or, vice versa, large boost factors required might be moderately loosened.  On more general grounds, at least in cases where no new light particles are introduced in the spectrum, it appears that DM candidates in the  few TeV mass range contributing a significant fraction of the CR electron/positron flux should lead also to a non-negligible antiproton flux, close to the current bound. Additionally, for the same class of candidates we expect  measurable signatures in both the gamma-ray and the neutrino flux. These ``multimessenger'' signatures are not surprising and have often been implicitly assumed in past phenomenological works on indirect DM detection. However, it is interesting that the conclusion holds virtually unchanged also for heavy DM candidates engineered to produce only leptons as tree level final states."
    },
    "0911/0911.5624_arXiv.txt": {
        "abstract": "Nuclear disks and rings are frequent galaxy substructures, for a wide range of morphological types \\citep[from S0 to Sc, see e.g.,][]{2010AIPC....M,2010AIPC....S}. We have investigated the possible minor-merger origin of inner disks and rings in spiral galaxies through collisionless N-body simulations. The models confirm that minor mergers can drive the formation of thin, kinematically-cold structures in the center of galaxies out of satellite material, without requiring the previous formation of a bar. Satellite core particles tend to be deposited in circular orbits in the central potential, due to the strong circularization experienced by the satellite orbit through dynamical friction. The material of the satellite core reaches the remnant center if satellites are dense or massive, building up a thin inner disk; whereas it is fully disrupted before reaching the center in the case of low-mass satellites, creating an inner ring instead. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0911/0911.4622_arXiv.txt": {
        "abstract": "The Feynman, Metropolis and Teller treatment of compressed atoms is extended to the relativistic regimes. Each atomic configuration is confined by a Wigner-Seitz cell and is characterized by a positive electron Fermi energy. The non-relativistic treatment assumes a point-like nucleus and infinite values of the electron Fermi energy can be attained. In the relativistic treatment there exists a limiting configuration, reached when the Wigner-Seitz cell radius equals the radius of the nucleus, with a maximum value of the electron Fermi energy $(E_e^F)_{max}$, here expressed analytically in the ultra-relativistic approximation. The corrections given by the relativistic Thomas-Fermi-Dirac exchange term are also evaluated and shown to be generally small and negligible in the relativistic high density regime.  The  dependence of the relativistic electron Fermi energies by compression for selected nuclei are  compared and contrasted to the non-relativistic ones and to the ones obtained in the uniform approximation. The relativistic Feynman, Metropolis, Teller approach here presented overcomes some difficulties   in the Salpeter approximation generally adopted  for compressed matter in physics and astrophysics. The treatment is then extrapolated to compressed nuclear matter cores of stellar dimensions with $A\\simeq (m_{\\rm Planck}/m_n)^3 \\sim 10^{57}$ or $M_{core}\\sim M_{\\odot}$. A new family of equilibrium configurations exists for selected values of the electron Fermi energy  varying in the range $0 < E_e^F \\leq (E_e^F)_{max}$. Such configurations fulfill global but not local charge neutrality. They have electric fields on the core surface, increasing for decreasing values of the electron Fermi energy reaching values much larger than the critical value $E_c = m_e^2c^3/(e\\hbar)$, for $E_e^F=0$.  We  compare and contrast our results with the ones of Thomas-Fermi model in strange stars. ",
        "introduction": "\\label{sec:1}  In a classic article Baym, Bethe and Pethick \\cite{baym} presented the problem of matching, in a neutron star,  a liquid  core, composed of $N_n$ neutrons, $N_p$ protons and $N_e$ electrons, to the crust taking into account the electrodynamical and surface tension effects. After discussing the different aspects of the problem  they concluded: {\\it The details of this picture requires further elaboration; this is a situation for which the Thomas-Fermi method is useful}. This statement, in first instance, may appear surprising: the Thomas-Fermi model has been  extensively applied in atomic physics (see e.g Gomb\\' as \\cite{gombas}, March \\cite{march}, Lundqvist and March \\cite{lundqvist}), also has been applied extensively in atomic physics in its relativistic form (see e.g  Ferreirinho, Ruffini and Stella \\cite{ruffini}, Ruffini and Stella \\cite{ruffinistella81}) as well as in the study of atoms with  heavy nuclei  in the classic works of Migdal, Popov and Voskresenskii  \\cite{migdal,migdal77}. Similarly there have been considerations of relativistic Thomas-Fermi model for quark stars  pointing out the existence of critical electric fields on their surfaces (see e.g. Alcock, Farhi, Olinto \\cite{alcock},  Usov \\cite{usov}). Similar results have also been obtained by Alford et al. \\cite{alford} in the transition at very high densities, from the normal nuclear matter phase in the core   to the color-flavor-locked phase  of quark matter in the inner core of hybrid stars. No example  exists to the application of the electromagnetic Thomas-Fermi model for neutron stars. This problem can indeed  be approached with merit by studying the simplified but rigorous concept of a nuclear matter core of stellar dimensions which  fulfills the relativistic Thomas-Fermi equation as discussed in  \\cite{rrx200613, prl}. As we will see this work leads to the prediction of the existence of a critical electric field at the interface between the core and the crust of a neutron star.  In \\cite{rrx200613, prl} we have first generalized the treatment of heavy nuclei by enforcing self-consistently the condition of beta equilibrium in the relativistic Thomas-Fermi equation. Using then the existence of scaling laws we have extended the results from heavy nuclei to the case of nuclear matter cores of stellar dimensions. In both these treatments we had there assumed the Fermi energy of the electrons $E_e^F=0$. The aim of this article is to proceed with this dual approach and to consider first the case of compressed atoms and then, using the existence of scaling laws, the compressed nuclear matter cores of stellar dimensions with a positive value of their electron Fermi energies.  It is well known that Salpeter has been among the first to study  the behavior of matter under extremely high pressures by considering a Wigner-Seitz cell of radius $R_{WS}$ \\cite{salpeter}. Salpeter assumed  as a starting point the nucleus point-like and a uniform distribution of electrons within a Wigner-Seitz cell. He then considered corrections to the above model due to the inhomogeneity of electron distribution. The first correction corresponds to the inclusion of the lattice energy $E_C=-(9N_p^2 \\alpha)/(10 R_{WS})$, which results from the point-like nucleus-electron interaction and, from the electron-electron interaction inside the cell of radius $R_{WS}$. The second correction is given by a series-expansion of the electron Fermi energy about the average electron density $n_e$ given by the uniform approximation. The electron density is then assumed equals to $n_e [1+\\epsilon(r)]$ with $\\epsilon(r)$ considered as infinitesimal. The Coulomb potential energy is assumed to be the one of the point-like nucleus with the uniform distribution of electrons of density $n_e$  thus the correction given by $\\epsilon(r)$ is neglected on the Coulomb potential. The electron distribution is then calculated at first-order by expanding the relativistic electron kinetic energy about its value given by the uniform approximation considering as infinitesimal the ratio $eV/E^F_e$ between the Coulomb potential energy $eV$ and the electron Fermi energy $E^F_e = \\sqrt{[c P^F_e(r)]^2 + m^2_e c^4}-m_e c^2 - e V$. The inclusion of each additional Coulomb correction results in a decreasing of the pressure of the cell $P_{S}$ by comparison to the uniform one (see \\cite{inpreparation}  for details).  It is quite difficult to assess the self-consistency of all the recalled different approximations adopted by Salpeter. In order to validate  and also to see the possible limits of the Salpeter approach,  we consider the relativistic generalization of the Feynman, Metropolis, Teller treatment \\cite{fmt} which takes automatically and globally into account all electromagnetic and special relativistic contributions. We show explicitly how this new treatment leads in the case of atoms to electron distributions markedly different from the ones often adopted in the literature of constant electron density distributions. At the same time it allows to overcome some of the difficulties in current treatments.  Similarly the point-like description of the nucleus often adopted in literature is confirmed to be unacceptable in the framework of a relativistic treatment.   In Sec.~\\ref{sec:2} we first recall the non-relativistic treatment of the compressed atom by Feynman, Metropolis and Teller. In Sec.~\\ref{sec:3} we generalize that treatment to the relativistic regime by integrating the relativistic Thomas-Fermi equation, imposing also the condition of  beta equilibrium. In Sec.~\\ref{sec:4} we first compare the new treatment with the one corresponding to a uniform electron distribution often used in the literature and to the Salpeter treatment. We also compare and contrast the results of the relativistic and the non-relativistic treatment.   We then proceed to analyze the case of compressed nuclear matter cores of stellar dimensions.  In Sec. \\ref{sec:5}, using the same scaling laws adopted in  \\cite{rrx200613, prl} we turn to the case of nuclear matter cores of stellar dimensions with mass numbers $A\\approx(m_{\\rm Planck}/m_n)^3 \\sim 10^{57}$ or $M_{core}\\sim M_{\\odot}$  where  $m_n$  is the neutron  mass and $m_{\\rm Planck}=(\\hbar c/G)^{1/2}$ is the Planck mass. Such a configuration present global but not local charge neutrality. Analytic solutions for the ultra-relativistic limit are obtained. In particular we find:  1) \\quad explicit analytic expressions for the electrostatic field and the Coulomb potential energy,  2) \\quad an entire range of possible Fermi energy for the electrons between zero and a maximum value $(E_e^F)_{max}$, reached when $R_{WS}=R_c$, which can be expressed analytically,  3) \\quad the explicit analytic expression of the ratio between the  proton number $N_p$ and the mass number $A$  when $R_{WS}=R_c$.  We turn then in Sec. \\ref{sec:6} to the study of the compressional energy of the nuclear matter  cores of stellar dimensions for selected values of the electron Fermi energy. We show that the solution with $E_e^F=0$  presents the largest value of the electrodynamical structure.  We finally summarize the conclusions in Sec. \\ref{sec:7}. ",
        "conclusions": "\\label{sec:7}  We have generalized to the relativistic regime the classic work of Feynman, Metropolis and Teller, solving a compressed atom by the Thomas-Fermi equation in a Wigner-Seitz cell.  In the relativistic generalization the equation to be integrated is the relativistic Thomas-Fermi equation, also called the Vallarta-Rosen equation \\cite{vallarta}. The integration of this equation does not admit any regular solution for a point-like nucleus and both the nuclear radius and the nuclear composition have necessarily  to be taken into account \\cite{ruffini, ruffinistella81}. This introduces a fundamental difference from the non-relativistic Thomas-Fermi model where a point-like nucleus was adopted.  As in  previous works \\cite{ruffini, ruffinistella81, rrx200613}  the protons in the nuclei have been assumed to be at constant density, the electron distribution has been derived by the Thomas-Fermi relativistic equation and  the neutron component has been derived by   the  beta equilibrium between neutrons, protons and electrons.  We have also examined for completeness the relativistic generalization of the Thomas-Fermi-Dirac equation by taking into due account the exchange terms \\cite{dirac30}, adopting  the general approach of Migdal, Popov and Voskresenskii \\cite{migdal77}, and shown that these effects, generally small, can be neglected in the relativistic treatment.  There are marked differences between the relativistic and the non-relativistic treatments.  The first and the most general one is that the existence of a finite size nucleus introduces necessarily a limit to the compressibility: the dimension of the Wigner-Seitz cell can never be smaller then the nuclear size. Consequently the electron Fermi energy which in the non-relativistic approach can reach arbitrarily large values, reaches in the present case a perfectly finite value whose expression has been given in analytic form. There are in the literature many papers adopting a relativistic treatment for the electrons together with a point-like approximation for the nucleus, which is clearly inconsistent (see e.g. \\cite{gilles1, gilles2}).  The second is the clear difference of the electron distribution as a function of the radius and of the nuclear composition as contrasted to the uniform approximation often adopted in the literature (see e.g.\\cite{mishustin}) which we have explicitly shown in the Fig.~\\ref{fmtw} of Sec. \\ref{sec:4}. Inferences based on the uniform approximation are not appropriate both in the relativistic and in the non-relativistic regime.  The third, one of the most relevant, is that the relativistic Feynman-Metropolis-Teller treatment allows to treat globally and in generality the electrodynamical interaction within the atom and the relativistic corrections leading to a softening of the dependence of the electron Fermi energy on the compression factor, as well as a gradual decrease of the exchange terms in proceeding from the non-relativistic to the fully relativistic regimes. It is then possible to derive, as shown in Table \\ref{xx1} of Sec.~\\ref{sec:4} a consistent equation of state for compressed matter which overcomes some of the difficulties of existing treatments describing the electrodynamical effect by a sequence of approximations which have lead to the occurrence of unphysical regimes e.g. the existence of negative pressure as in the Salpeter approach. As a direct application of this treatment we have reconsidered the study of white dwarfs within the relativistic Feynman, Metropolis, Teller approach and evaluate their effects on the value of the radii, of the masses of the equilibrium configurations as well as on the numerical value of the critical mass \\cite{inpreparation}. We have there compared and contrasted the results obtained by Chandrasekhar with a uniform approximation with the ones obtained by the equation of state of Salpeter and the ones following from the treatment presented in this article.   We have then extrapolated these results to the case of nuclear matter cores of stellar dimensions for $A\\approx(m_{\\rm Planck}/m_n)^3 \\sim 10^{57}$ or $M_{core}\\sim M_{\\odot}$. The aim here is to explore the possibility of obtaining for these systems a self-consistent solution presenting global and not local charge neutrality. The results generalize the considerations presented in the previous article   corresponding to a nuclear matter core of stellar dimensions with null Fermi energy of the electrons \\cite{prl}. The ultra-relativistic approximation allows to obtain analytic expressions for the fields. The exchange terms can in this approximation be safely neglected. An entire family of configurations exist with values of the Fermi energy of the electrons ranging from zero to a maximum value $(E_e^F)_{max}$ which is reached when the Wigner- Seitz cell coincides with the core radius. The configuration with $E_e^F=(E_e^F)_{max}$ corresponds to the configuration with $N_p=N_e$ and $n_p=n_e$. For this limiting value of the Fermi energy  the system fulfills both the global and the local charge neutrality and correspondingly no electrodynamical structure is present in the core. All the other configurations presents overcritical  electric fields close to their surface. The configuration with $E_e^F=0$ has the maximum value of the electric field at the core surface, well above the critical value $E_c$ (see Fig. \\ref{efieldf}, Fig.~\\ref{efieldf1} and Fig.~\\ref{ANp} of Section \\ref{sec:5}). All these cores with overcritical electric fields are  stable against the vacuum polarization process due to the Pauli blocking by the degenerate electrons \\cite{physrep}. We have also compared and contrasted our treatment of the relativistic Thomas-Fermi solutions to the corresponding one addressed in the framework of strange stars \\cite{alcock} pointing out in these treatments some inconsistency in the definition of the Coulomb potential.  We have finally compared the compressional energy of configurations with selected values of the electron Fermi energy. In both systems of the compressed atoms and of the nuclear matter cores of stellar dimensions a maximum value of the Fermi energy has been reached corresponding to the case of Wigner-Seitz cell radius $R_{WS}$ coincident with the core radius $R_c$.  Both  problems  considered, the one of a compressed atom and the one of compressed nuclear density core of stellar dimensions, have been treated by the solution of the relativistic Thomas-Fermi equation and by enforcing the condition of beta equilibrium. They are theoretically well defined and, in our opinion, a necessary step in order to approach the more complex problem of a neutron star core and its interface with the neutron star crust.  Neutron stars are composed of two sharply different components: the liquid core at nuclear and/or supra-nuclear density consisting of neutrons, protons and electrons and a crust of degenerate electrons in a lattice of nuclei \\cite{baym, harrisonweel} and possibly of free neutrons due to neutron drip when this process occurs (see e.g. \\cite{baym}). Consequently, the boundary conditions for the electrons at the surface of the neutron star core will have generally a positive value of the electron Fermi energy in order to take into account the compressional effects of the neutron star crust on the core \\cite{letter3}. The case of zero electron Fermi energy corresponds to the limiting case of absence of the crust.   In a set of interesting papers \\cite{glendenning92, glendenning95, glendenning97, glendenning99, glendenning00, glendenning01} Glendenning has relaxed the local charge neutrality condition for the description of the mixed phases in hybrid stars. In such configurations the global charge neutrality condition, as opposed to the local one, is applied to the limited regions where mixed phases occur while in the pure phases the local charge neutrality condition still holds. We here generalize Glendenning's considerations by looking to a violation of the local charge neutrality condition on the entire configuration, still keeping its overall charge neutrality. This effect does not occur in processes occurring just locally, but it requires the global description of the equilibrium configuration. As exemplified in \\cite{ inpreparationa, inpreparationb}, where gravitational effects are taken into account both in the $T=0$ and $T\\neq0$ cases, the local properties need the previous knowledge of the entire equilibrium configuration and therefore the description is necessarily global.   In conclusion the analysis of compressed atoms following the relativistic Feynman, Metropolis, Teller treatment presented in the first part of this article has important consequences in the determination of the mass-radius relation of white dwarfs \\cite{inpreparation} leading to the possibility of a direct confrontation of these results with observations, in view of the current great interest for the cosmological implications of the type Ia supernovae \\cite{phillips,riess98,perlmutter, riess04}. The results presented in the second part of this article on nuclear matter cores of stellar dimensions evidence the possibility of having the existence of critical electromagnetic fields in the interface of the core and the neutron star crust. The results here obtained in a simplified but rigorous approach of the application of the relativistic Feynman, Metropolis, Teller treatment to the constant density cores in beta equilibrium  conform  to the prediction of Baym, Bethe and Pethick \\cite{baym}. This treatment has been further extended to the case in which a self-gravitating system of degenerate neutrons, protons and electrons is considered within the framework of relativistic quantum statistics and Einstein-Maxwell equations \\cite{letter3} and to the case in which also strong interactions are present \\cite{strong}."
    },
    "0911/0911.3658_arXiv.txt": {
        "abstract": "\\widetext In the context of Modified Newtonian Dynamics (MOND), we study how perturbation of a spherically symmetric system would affect the gravitational field. In particular, we study systems of perturbed and unperturbed spherical shells.  For a single perturbed shell, we show that the field inside the shell is much smaller than what would be expected from a naive scaling formula. The strength of the perturbation field within the shell is screened by the spherically symmetric component of the mass, and is reduced as the spherically symmetric component is increased.  For a two-shell system, we again show that the perturbed field is screened by the shells, no matter which shell's mass distribution is perturbed. The field within the inner shell is most suppressed when the inner  and outer shells coincide.  However, for a very light inner shell, the perturbation to the field can be enhanced.  The enhancement is typically larger for smaller inner shells, and the perturbed field can be amplified  by almost a factor of 2.  The relevance to the  effect of external fields on  galaxy dynamics is discussed. ",
        "introduction": "\\indent Most of the astrophysical computation concerns a local system embedded in much larger background system.  For example, Solar system is embedded in Milky Way, a galaxy in a cluster, a cluster to the supercluster etc.  In GR or Newtonian gravity, the external field effect on the local system due to much larger background mass density are often neglected based on two simplifying assumptions: that the background mass distribution is spherically symmetric and/or the background distribution is located very far from the local system.  Under the first assumption, one would invoke Birkhoff's theorem that the field inside a spherical cavity should vanish.  (Inclusion of the local mass distribution may spoil the spherical symmetry of the whole system, and in GR, the background mass distribution may exert force on the local system.  However the external field induced in this way is negligibly small.  See \\cite{Dai}).  Under the second assumption, the inverse-square law will suppress the contribution of the portion of the background mass distribution that is  far from the local system.  Therefore, even when the background mass distribution is only approximately spherically symmetric, the external field effect in GR or Newtonian gravity would not significantly alter the local dynamics.  It might be naively expected that the external field effect is more significant in MOND theory.   A small perturbation of the background mass distribution from spherical symmetry might be expected to alter the internal field configuration significantly.  This is because the field of a bounded mass distribution scales as $\\sim \\frac{1}{r}$ in MOND, and the field due to the aspherical part may therefore persist to greater distances.  Of course, the modified Possion equation in MOND is non-linear, and it does not embody the superposition principle.  One cannot treat the contribution from the spherically symmetric part of the mass distribution and the aspherical part separately. The total mass distribution has to be considered in order to infer the size of any external field effect. The external field effect in MOND has been receiving growing interest \\cite{Bruneton,WW}.  It is argued that MOND theory is consistent with the observed escape velocity in the solar neighborhood if the Milky Way is embedded in a constant external field of $\\sim 0.01a_{0}$.  (Here $a_0$ is the universal critical acceleration that characterizes MOND, $a_{0}\\simeq 1.2\\times10^{-8}{\\mathrm cm}/{\\mathrm s}^2$. On the other hand, Wu et al \\cite{Wu2} has showed that an external field of approximately $0.03a_{0}$ is needed to bind the Large Magellanic Cloud  to the Milky Way.  Since the external field effect on the local escape speed presents a strict test of MOND theory, it is of interest to understand when, if at all, the aspherical perturbation of very distant backgrounds could be neglected.  In an earlier work \\cite{Dai}, we studied  the fields inside spherical shells as probed by non-negligible test masses. We showed that when those test masses are placed off-center, thereby breaking the spherical symmetry, the acceleration of the test mass can become comparable to the (external) surface gravity of the shell. In what follows, we consider  simplified systems of spherical shells with perturbations in order to better understand how the perturbations of spherically symmetric  mass distributions translate into perturbed gravitational fields. ",
        "conclusions": "We have shown, by examining a simple toy system of a shell with a perturbation that the naive MOND equation (\\ref{naive}) overestimates the field inside the shell.  For a small perturbation on the shell signified by the perturbation paramter $\\epsilon$, the field within the shell shell is of order $\\epsilon$ and not $\\epsilon^{\\frac{1}{2}}$ as predicted by the naive MOND equation.  To properly estimate the field within the shell, the jump condition on the surface of the shell has to be taken into account. We have also shown that the field due to the perturbation within the shell can be screened by the spherical distribution of the mass.  The screening is strongest  if the spherical distribution is positioned near the shell.  This contrasts to Newtonian gravity where addition of the spherical distribution does not affect the field configuration within the distribution. Finally, the anti-screening effect is noticed when the small and light shell is added to the system.  The enhancement inside the small shell becomes significant when the generic field of the shell $\\sim\\sqrt{GMa_{0}}/R$ becomes comparable to the perturbation field near the shell.  The enhancement is generally larger for the shell with smaller radius.  Let now suppose that, in the simplest picture, a cluster is modeled by a spherical shell and the external field is produced by an aspherical perturbation on the shell.  In order to estimate the magnitude of the external field, one needs to know not only the size of perturbation but also the total mass of the shell.  Applying naive MOND equation on this system to estimate the external field would be quite inadequate.  If the galaxy can be modeled by a small shell within a large, more dominant shell (a cluster), then the external field within the galaxy is likely to be enhanced.  The enhancement depends on the ratio $M_{gal}/M_{clu}$ and the size of the galaxy, but as shown in FIG. \\ref{two}, it can be of the same size as the external field itself."
    },
    "0911/0911.1188_arXiv.txt": {
        "abstract": "The neutrino-antineutrino annihilation into electron-positron pairs near the surface of compact general relativistic  stars could play an important role in supernova explosions, neutron star collapse, or for close neutron star binaries near their last stable orbit. General relativistic effects increase the energy deposition rates due to the annihilation process. We investigate the deposition of energy and momentum due to the annihilations of neutrinos and antineutrinos in the equatorial plane of the rapidly rotating neutron and quark stars, respectively. We analyze the influence of general relativistic effects, and we obtain the general relativistic corrections to the energy and momentum  deposition rates for arbitrary stationary and axisymmetric space-times. We obtain the energy and momentum deposition rates for several classes of rapidly rotating neutron stars, described by different equations of state of the neutron matter, and for quark stars, described by the MIT bag model equation of state and in the CFL (Color-Flavor-Locked) phase, respectively. Compared to the Newtonian calculations, rotation and general relativistic effects increase the total annihilation rate measured by an observer at infinity. The differences in the equations of state for neutron and quark matter also have important effects on the spatial distribution of the energy deposition rate by neutrino-antineutrino annihilation. ",
        "introduction": "Since the pioneering works of Cooperstein et al. (1986), Cooperstein et al. (1987),  and Goodman et al. (1987), the energy deposition rate from the $\\nu +\\bar{\\nu}\\rightarrow e^{+}+e^{-}$ neutrino annihilation reaction has been intensively studied.  This reaction is of considerable importance for Type II supernova dynamics, neutron star collapse, or for close neutron star binaries near their last stable orbit. Neutrino-antineutrino annihilation into electrons and positrons can deposit more than $10^{51}$ ergs above the neutrino-sphere of a type II supernova \\citep{Go87}. This energy deposition, together with neutrino-baryon capture, significantly increases the neutrino heating in the envelope via the so-called delayed shock mechanism \\citep{BeWi85,Be90}. For large $r$ the energy deposition rate is proportional to $r^{-8}$, where $r$ is the distance from the center of the neutrino-sphere. The initial estimations of the neutrino annihilation reaction efficiencies were based on a Newtonian approach, by assuming that $2GM/c^2R<<1$, where $M$ is the gravitational mass of the star, and $R$ is the distance scale. However, for a full understanding of the effects of the neutrino annihilation in strong gravitational fields, general relativistic effects must be taken into account \\citep{SaWi99}. For a static neutron star, by adopting for the description of the gravitational field the Schwarzschild metric,  the efficiency of the $\\nu +\\bar{\\nu}\\rightarrow e^{+}+e^{-}$ process is enhanced over the Newtonian values up to a factor of more than $4$ in the regime applicable to Type II supernovae, and by up to a factor of $30$ for collapsing neutron stars \\citep{SaWi99}. The neutrino pair annihilation rate into electron pairs between two neutron stars in a binary system was calculated by Salmonson \\& Wilson (2001). A closed formula for the energy deposition rate at any point between the stars was obtained, where each neutrino of a pair derives from each star, and this result was compared with that in which all neutrinos derive from a single neutron star. An approximate generalization of this formula was also given to include the relativistic effects of gravity. The interstar neutrino annihilation is a significant contributor to the energy deposition between heated neutron star binaries.  The neutrino-antineutrino annihilation into electrons and positrons is an important candidate to explain the energy source of the gamma ray bursts (GRBs) \\citep{Pa90, MeRe92,RuJa98, RuJa99, AsIw02}. The semi-analytical study of  the gravitational effects on neutrino pair annihilation near the neutrinosphere and around the thin accretion disk were considered in Asano \\& Fukuyama (2000), by assuming that the accretion disk is isothermal, and that the gravitational field is dominated by the Schwarzschild black hole. General relativistic effects were studied only near the rotation axis. The energy deposition rate is enhanced by the effect of orbital bending toward the center. However, the effects of the redshift and gravitational trapping of the deposited energy reduce the effective energy of the gamma-ray burst's source. Although each effect is substantial, the effects partly cancel one another. As a result, the gravitational effects do not substantially change the energy deposition rate for either the spherically symmetric case or the disk case \\citep{AsFu00}. Using idealized models of the accretion disk, Asano \\& Fukuyama (2001) investigated the relativistic effects on the energy deposition rate via neutrino pair annihilation near the rotation axis of a Kerr black hole, by assuming that the neutrinos are emitted from the accretion disk. The bending of neutrino trajectories and the redshift due to the disk rotation and gravitation were also taken into consideration. The Kerr parameter, $a$, affects not only behavior of the neutrinos, but also the inner radius of the accretion disk. When the deposition energy is mainly contributed by the neutrinos coming from the central part, the redshift effect becomes dominant as a becomes large, and the energy deposition rate is reduced compared with that neglecting the relativistic effects. On the other hand, for a small $a$, the bending effect becomes dominant and makes the energy increase by factor of 2, compared with that which neglects the relativistic effects \\citep{AsFu01}.  The effect of the inclusion of the slow rotation of the star in the general relativistic treatment of the neutrino-antineutrino annihilation into electron positron pairs was considered in Prasanna \\& Goswami (2002). It was shown that the inclusion of the rotation results in a reduction in the heating rate, as compared to the no rotation case. The energy-momentum deposition rate (MDR) from the $\\nu -\\bar{\\nu}$ collisions above a rotating black hole/thin accretion disk system was calculated by Miller et al. (2003), by imaging the accretion disk at a specified observer using the full geodesic equations, and calculating the cumulative MDR from the scattering of all pairs of neutrinos and antineutrinos arriving at the observer. The dominant contribution to the MDR comes from near the surface of the disk with a tilt of approximately $\\pi/4$ in the direction of the disk's rotation. The MDR at large radii is directed outward in a conic section centered around the symmetry axis and is larger by a factor of 10-20 than the on-axis values. There is also a linear dependence of the MDR on the black hole angular momentum. The deposition of energy and momentum, due to the annihilation of neutrinos $\\nu $ and antineutrinos $\\bar{\\nu }$ in the vicinity of steady, axisymmetric accretion tori around stellar-mass black holes was investigated in Birkl et al. (2007). The influence of general relativistic effects were analyzed in combination with different neutrinosphere properties and spatial distribution of the energy deposition rate. Assuming axial symmetry, the annihilation rate 4-vector was numerically computed. The local neutrino distribution was constructed by ray-tracing neutrino trajectories in a Kerr space-time, using null geodesics. Different shapes of the neutrinospheres, spheres, thin disks, and thick accretion tori were studied, whose structure ranges from idealized tori to equilibrium non-selfgravitating matter distributions. Compared to Newtonian calculations, general relativistic effects increase the total annihilation rate measured by an observer at infinity by a factor of two when the neutrinosphere is a thin disk, but the increase is only 25\\% for toroidal and spherical neutrinospheres. Thin disk models yield the highest energy deposition rates for neutrino-antineutrino annihilation, and spherical neutrinospheres the lowest ones, independently of whether general relativistic effects are included. General relativity and rotation cause important differences in the spatial distribution of the energy deposition rate by neutrino $\\nu $ and antineutrino $\\bar{\\nu }$-annihilation \\citep{Bi07}. The study of the structure of neutron star disks based on the two-region (i.e., inner and outer) disk scenario was performed by Zhang \\& Dai (2009), who calculated the neutrino annihilation luminosity from the disk in various cases. The effects of the viscosity parameter $\\alpha $, energy parameter $\\epsilon $ (measuring the neutrino cooling efficiency of the inner disk), and outflow strength on the structure of the entire disk, as well as the effect of emission from the neutron star surface boundary emission on the total neutrino annihilation rate were investigated.  An outflow from the disk decreases the density and pressure, but increases the thickness of the disk. Moreover, compared with the black hole disk, the neutrino annihilation luminosity above the neutron star disk is higher, and the neutrino emission from the boundary layer could increase the neutrino annihilation luminosity by about one order of magnitude higher than the disk without boundary emission. The neutron star disk with the advection-dominated inner disk could produce the highest neutrino luminosity, while the disk with an outflow has the lowest\\citep{ZhDa09}. A detailed general relativistic calculation of the neutrino path for a general metric describing a rotating star was studied in Mallick \\& Majumder (2009). The neutrino path was calculated along the equatorial and polar plane. The expression for the minimum photosphere radius  was obtained and matched with the Schwarzschild limit. The minimum photosphere radius was calculated for stars with two different equations of state, each rotating with two different velocities. The results show that the minimum photosphere radius for the hadronic star is much greater than for the quark star, and that the minimum photosphere radius increases as the rotational velocity of the star decreases. The minimum photosphere radius along the polar plane is larger than that along the equatorial plane. The estimate of the energy deposition rate for neutrino pair annihilation for the neutrinos coming from the equatorial plane of a rotating neutron star was calculated along the rotation axis in Bhattacharyya et al. (2009), by using the Cook-Shapiro-Teukolsky metric. The neutrino trajectories, and hence the neutrino emitted from the disk, are affected by the redshift due to the disk rotation and gravitation. The energy deposition rate is very sensitive to the value of the temperature, and its variation along the disk. The rotation of the star has a negative effect on the energy deposition rate, it decreases with increase in rotational velocity. The standard model for Type II supernovae explosions, confirmed by the detection of neutrinos emitted during the supernova explosion, predicts the formation of a compact object, usually assumed to be a neutron star \\citep{Cha09}. However, the newly formed neutron star at the center of SN 1987A may undergo a phase transition after the neutrino trapping timescale ($\\sim10$ s). Consequently the compact remnant of SN 1987A may be a strange quark star, which has a softer equation of state than that of neutron star matter \\citep{Cha09}. Such a phase transition can induce stellar collapse and result in large amplitude stellar oscillations. A three-dimensional Newtonian hydrodynamic code was used to study the time evolution of the temperature and density at the neutrinosphere. Extremely intense pulsating neutrino fluxes, with submillisecond period and with neutrino energy (greater than 30 MeV), can be emitted because the oscillations of the temperature and density are out of phase almost $180^{\\circ}$. The dynamical evolution of a phase-transition-induced collapse of a neutron star to a hybrid star, which consists of a mixture of hadronic matter and strange quark matter, was studied in Cheng et al. (2009). It was found that both the temperature and the density at the neutrinosphere are oscillating with acoustic frequency. Consequently, extremely intense, pulsating neutrino/antineutrino fluxes will be emitted periodically. Since the energy and density of neutrinos at the peaks of the pulsating fluxes are much higher than the non-oscillating case, the electron/positron pair creation rate can be enhanced dramatically. Some mass layers on the stellar surface can be ejected, by absorbing energy of neutrinos and pairs. These mass ejecta can be further accelerated to relativistic speeds by absorbing electron/positron pairs, created by the neutrino and antineutrino annihilation outside the stellar surface.  It is the purpose of the present paper to consider a comparative systematic study of the neutrino-antineutrino annihilation process around rapidly rotating neutron and strange stars, respectively, and to obtain the basic physical parameters characterizing this process (the electron-positron energy deposition rate per unit volume and unit time, and the total emitted power, respectively), by taking into account the full general relativistic corrections. In order to obtain the electron-positron energy deposition rate for various types of neutron and quark stars we generalize the relativistic description of the neutrino-antineutrino annihilation process to the case of arbitrary stationary and axisymmetric geometries.   To compute the electron-positron energy deposition rate, the metric outside the rotating general relativistic stars must be determined. In the present study we study the equilibrium configurations of the rotating neutron and quark stars by using the RNS code, as introduced in \\cite{SteFr95}, and discussed in detail in \\cite{Sterev}. This code was used for the study of different models of rotating neutron stars in \\cite{No98} and for the study of the rapidly rotating strange stars (\\cite{Ste99}). The software provides the metric potentials for various types of compact rotating general relativistic objects, which can be used to obtain the electron-positron energy deposition rate in the equatorial plane of rapidly rotating neutron and quark stars.  The present paper is organized as follows. In Section 2 we present the basic formalism for the calculation of the electron-positron energy deposition rate from neutrino-antineutrino annihilation. The general relativistic corrections to this process are obtained in Section 3. The equations of state of dense neutron and quark matter used in the present study are presented in Section 4. In Section 5 we obtain the electron-positron energy deposition rates in the equatorial plane of the considered classes of neutron and quark stars. We discuss and conclude our results in Section 6. ",
        "conclusions": "In the present paper we have considered the general relativistic effects on the energy deposition rate from the neutrino-antineutrino annihilation process near rapidly rotating neutron and quark stars. The energy deposition rate  has been obtained numerically for several equations of state of the neutron matter, and for two types of quark stars, respectively.  All the general relativistic correction factors, related to this process, can be obtained from the metric of the central compact object. Due to the differences in the space-time structure, the quark stars  present some important differences with respect to the energy deposition rate, as compared to the neutron stars. As a general result we have found that rotation always enhances the energy deposition rate.  A possible astrophysical application of our results could be in the explanation of the physical processes that lead to the formation of the Gamma-Ray Bursts (GRBs). The so-called fireball model can basically explain the observational facts well, and thus it is strongly favored, and widely accepted today (for recent reviews on GRBs see M\\'esz\\'aros (2006) and Zhang (2007), respectively). In this model, the central engine gives birth to some energetic ejecta intermittently, like a geyser, producing a series of ultra-relativistic shells. The shells collide with each other and produce strong internal shocks. The highly variable $\\gamma$-ray emission in the main burst phase of GRBs should be produced by these internal shocks. After the main burst phase, the shells merge into one main shell and continue to expand outward. These merged shells, which moves relativistically, interacts with the surrounding medium to form an external shock. This shock will give a good fraction of its energy to the swept-up electrons and accelerate them to relativistic velocity. Similarly, a fraction of shock energy will go to the magnetic field. These shocked relativistic electrons move in the magnetic field and emit synchrotron radiation to produce a broadband electromagnetic emission  called \"afterglows\". According to the Swift observations(see Zhang (2007) for a review), the decay of the X-ray afterglow can be classified roughly into four stages. Stage I (up to $\\sim100$ s) indicates the early fast decay phase, usually with flux density decay as $t^{-3}$ to $t^{-5}$; Stage II (up to $\\sim 1000$ s) indicates the subsequent shallow decay phase, with $t^{-0.2}$ to $t^{-0.8}$ ; Stage III (up to $\\sim 10^4$ s) is the late normal decay phase, $t^{-1}$ to $t^{-1.3}$; Stage IV (beyond $10^4$ s) is the late fast decay phase, $t^{-2}$ to $t^{-2.5}$. The first stage is explained as the tail decay of the prompt emission, and the third and the fourth stages can be explained very well in terms of the standard fireball model. Although there are many different explanations of the second stage, the  \"shallow decay phase\", one of the most popular explanations is the late time continuous energy injection. Obviously, the energy deposition in the equatorial plane cannot be responsible to the prompt gamma-ray emission, which requires a fast and a very narrow beam jet. In fact,  the high energy photons detected by Fermi \\citep{Me09} imply that the Lorentz factor of fast jet is over 500.  On the other hand, most deposition energy should come from the equatorial plane. We suggest that the energy deposition along the axis will form a narrow cone jet, which encounters much less material, and hence it can maintain a very large Lorentz factor. On the other hand, the energy deposition in the equatorial plane is higher, but it also encounters much more material, and hence it moves slowly. When the fast jet slows down due to the interstellar medium, this slow but more energetic wind from the equatorial plane can inject additional energy into the fast jet, and result in the shallow decay phase. The detailed calculations will be carried out in our future study.  The neutrino annihilation and the electron-positron pair production also plays an essential role in the astrophysical processes associated with the phase transitions that could take place inside neutron stars.  For example, the sudden phase transition from normal nuclear matter to a mixed phase of quark and nuclear matter induces temperature and density oscillations at the neutrinosphere. Consequently, extremely intense, pulsating neutrino/antineutrino and leptonic pair fluxes will be emitted \\cite{Ch09}. During this stage several mass ejecta can be ejected from the stellar surface by the neutrinos and antineutrinos. These ejecta can be further accelerated to relativistic speeds by the electron/positron pairs, created by the neutrino and antineutrino annihilation outside the stellar surface. In order to produce the Gamma-Ray Bursts, a high neutrino emission rate is necessary. On the other hand, it is important to note that electron-positron pairs can deposit energy much more efficiently than the neutrinos, and the dominant energy deposition process is the neutrino-antineutrino annihilation process \\cite{Ch09}. In fact most pairs are created outside the star. A large fraction of the neutrino energy, will be absorbed by the matter very near the stellar surface. When this amount of energy exceeds the gravitational binding energy, some mass near the stellar surface will be ejected, and this mass will be further accelerated by absorbing pairs created from the neutrino and antineutrino annihilation processes outside the star. This process may be a possible mechanism for short Gamma-Ray Bursts \\cite{Ch09}.  The neutrino emission rate is also strongly dependent on the temperature. In the standard models of GRBs it is assumed that the central object is surrounded by a degenerate accretion disk, which allows super-Eddington accretion rates, of the order of one solar mass per second \\cite{Zh07}. If the central compact object is a neutron or a quark star, such super-Eddington accretion rates can maintain the compact object at very high temperatures, and thus allowing very high neutrino luminosities, as well as a high rate of electron-positron pair production. The neutrino temperature can be estimated by assuming that the accretion power $\\eta \\dot{M}c^2$, where $\\eta $ is the efficiency of the energy conversion and $\\dot{M}$ is the accretion rate, is equal to the radiation power $4\\pi R^2\\sigma T_{\\nu }^4$, which gives the temperature as $T_{\\nu }=\\left(\\eta \\dot{M}c^2/4\\pi R^2\\sigma \\right)^{1/4}$. By taking $\\eta =0.1$, $R=10^6$ cm, and an accretion rate of $\\dot{M}=1M_{\\odot}/10\\;{\\rm s}$, we obtain $T_{\\nu }=7.14\\times 10^{10}$K, a temperature which is of the order of MeV. Therefore, if the accretion disk is fed at a high rate, like, for example, by the fallback material after a supernova explosion, a high neutrino-antineutrino emission rate can be maintained, and this could explain some of the basic properties of the GRB phenomenon."
    },
    "0911/0911.2248_arXiv.txt": {
        "abstract": "Recent work has identified a population of low-redshift E/S0 galaxies that lie on the blue sequence in color vs. stellar mass parameter space, where spiral galaxies typically reside. While high-mass blue-sequence E/S0s often resemble young merger or interaction remnants likely to fade to the red sequence, we focus on blue-sequence E/S0s with lower stellar masses (${M_*} <$ a few $\\times 10^{10}\\,M_{\\odot}$), which are characterized by fairly regular morphologies and low-density field environments where fresh gas infall is possible. This population may provide an evolutionary link between early-type galaxies and spirals through disk regrowth. Focusing on atomic gas reservoirs, we present new GBT HI data for 27 E/S0s on both sequences as well as a complete tabulation of archival HI data for other galaxies in the Nearby Field Galaxy Survey. Normalized to stellar mass, the atomic gas masses for 12 of the 14 blue-sequence E/S0s range from 0.1 to $>$1.0, demonstrating that morphological transformation is possible if the detected gas can be converted into stars. These gas-to-stellar mass ratios are comparable to those of spiral and irregular galaxies and have a similar dependence on stellar mass. Assuming that the HI is accessible for star formation, we find that many of our blue-sequence E/S0s can increase in stellar mass by 10--60\\% in 3 Gyr in both of two limiting scenarios, exponentially declining star formation (i.e., closed box) and constant star formation (i.e., allowing gas infall). In a constant star formation scenario, about half of the blue-sequence E/S0s require fresh gas infall on a timescale of $\\lesssim$3 Gyr to avoid exhausting their atomic gas reservoirs and evolving to the red sequence. We present evidence that star formation in these galaxies is bursty and likely involves externally triggered gas inflows. Our analysis suggests that most blue-sequence E/S0s are indeed capable of substantial stellar disk growth on relatively short timescales. ",
        "introduction": "Current models of galaxy formation and evolution favor hierarchical growth of galaxies from smaller systems (e.g., \\citealt{white91,somerville99,bower06}). Within the paradigm of hierarchical galaxy formation, galaxies evolve along the Hubble sequence, transforming back and forth between E/S0 and spiral/irregular morphology through a series of quiescent and violent periods.  Recognition that galaxies can transform from late to early type dates back at least to \\citet{toomre72}, who proposed that elliptical galaxies can form from mergers of similar mass late-type galaxies. Recent simulations find that while similar mass (1:1--3:1) mergers of disk galaxies result in classical ellipticals, unequal mass (4.5:1--10:1) or gas-rich mergers of galaxies result in S0s and hybrid galaxies with properties of both early and late types \\citep{bekki98,naab06,bournaud05}. Substantial observational evidence confirms that early-type galaxies can form through major mergers (e.g., \\citealt{schweizer92,vandokkum05}). In agreement with the simulations, \\citet{emsellem07} argue that larger, slow-rotating elliptical galaxies form from dry major mergers while fainter, fast-rotating ellipticals and S0s seem to form from minor mergers or very gas-rich major mergers.  Can galaxies evolve in the other direction, from early to late type? Simulations suggest that galaxy evolution in this direction involves the regrowth of stellar disks \\citep{steinmetz02,governato07}. Morphological transformation from early to late type requires that several conditions be met, which recent studies have found satisfied in some early-type galaxies.  First, the transformation from early- to late-type requires a substantial reservoir of cold gas --- the raw material for disk regrowth. Historically, early-type galaxies were thought to be gas-poor \\citep{faber76,knapp78}. Subsequent surveys, however, found that the ratio of gas mass to blue luminosity (${M_{\\rm HI}/L_{B}}$) in early-type galaxies ranges from upper limits of ${M_{\\rm HI}/L_{B}} \\sim 0.009\\, {M_{\\odot}/L_{\\odot}}$ to measured ${M_{\\rm HI}/L_{B}} \\sim 2.7\\, {M_{\\odot}/L_{\\odot}}$ for large atomic gas reservoirs similar to those of spiral galaxies \\citep{hawarden81,knapp85,wardle86,sadler00,oosterloo02}. The 70\\% HI detection rate by \\citet{morganti06} suggests that atomic gas reservoirs are actually relatively common in field early-type galaxies. Recent surveys also find a significant amount of molecular gas ($10^{7}$--$10^{9} M_{\\odot}$) in 28--78\\% of early-type galaxies, depending on the survey \\citep{lees91,knapp96,welch03,sage07,combes07}. The large range in cold gas content hints that there may be distinct sub-populations of early-type galaxies --- the conventional red and dead early types, and a population that is still accreting gas and forming stars.  In addition to the {\\it presence} of cold gas, the distribution of the gas is important in assessing the potential for disk regrowth. The existence of giant HI disks and rings \\citep{morganti97,serra07}, as predicted by simulations of mergers between gas-rich galaxies \\citep{barnes02}, makes it possible for stellar disks to form from already-present gas disks. The HI structures around some early-types galaxies are known to have regular velocity fields, with a continuity between ionized and neutral gas \\citep{vangorkom86,schweizer89,schiminovich95,morganti06}. \\citet{hibbard96} find a rotationally supported HI disk in the elliptical galaxy NGC~520, which, if star formation is triggered in the disk, may generate a late-type galaxy. Kiloparsec-scale disks of regularly rotating molecular gas have also been found in the center of some early-type galaxies \\citep{young02,young05,young08,crocker08}. All these observations suggest that the gas disks around some, possibly most early-type galaxies have reached an equilibrium arrangement suitable for disk growth.  While the presence of cold gas disks in an equilibrium configuration is necessary for stellar disk growth, it does not imply that star formation is actually occurring.  A key question is whether E/S0 galaxies with such disks are currently evolving morphologically, and whether such galaxies constitute a significant fraction of the early-type galaxy population. To answer these questions, one must be able to distinguish between old, gas-poor early-type galaxies and those undergoing morphological transformation via disk regrowth.  \\citet*[hereafter KGB]{kannappan09a} recently identified a population of E/S0s that reside alongside spiral galaxies on the blue sequence in color vs. stellar mass parameter space. They argue that among these ``blue-sequence E/S0s,'' a subset with low-to-intermediate masses may represent the missing link --- i.e., galaxies that are actively (re)growing stellar disks and plausibly transitioning from early to late type. \\citetalias{kannappan09a} find that these blue-sequence E/S0s fall between spirals and red-sequence E/S0s in scaling relations (stellar mass $M_*$ vs.\\ radius and vs.\\ velocity dispersion $\\sigma$), implying that blue-sequence E/S0s might form a transitional population between the two groups. Compared to conventional red-sequence E/S0s, blue-sequence E/S0s consistently have bluer outer disks, and often bluer centers as well --- suggesting on-going star formation in both disks and disky bulges. Further supporting this disk-building picture, \\citetalias{kannappan09a} find evidence for kinematically distinct disks (i.e., counterrotating or polar) in a notable sub-population of blue-sequence E/S0s.  Unlike red-sequence E/S0s, whose mass function peaks at high masses, blue-sequence E/S0s are rare for $M_* >$ 1--2$\\,\\times 10^{11}\\,M_{\\odot}$, and are common only for ${M_* < 3 \\times 10^{10}\\,M_{\\odot}}$ \\citepalias{kannappan09a}. Their abundance increase sharply to 20--30\\% below $M_* \\lesssim 5 \\times 10^9\\,M_{\\odot}$, coincident with the mass threshold below which galaxies become notably more gas rich (\\citealt{kannappan08conf}, based on \\citealt{kannappan04}). At intermediate masses between $10^{10}$ and $10^{11}\\,M_{\\odot}$, \\citetalias{kannappan09a} find that blue-sequence E/S0s include examples of both major mergers that are likely to fade onto the red sequence after exhausting their gas, and also settled galaxies that may be evolving toward later-type morphologies. The former dominate at higher masses, and the latter at lower masses.  One outstanding question that remains is the extent of disk growth possible in low-to-intermediate mass blue-sequence E/S0s. In this paper, we approach this question from the point of view of atomic gas reservoirs. \\citetalias{kannappan09a} addressed this question with limited archival data for the Nearby Field Galaxy Survey (NFGS, \\citealt{jansen00a}), finding preliminary evidence for gas reservoirs comparable to spiral galaxies. In this paper we present more complete, higher-quality data for all NFGS E/S0s with $M_* < 4\\times10^{10}\\,M_{\\odot}$ and a sampling of more massive E/S0s. We compare the atomic gas masses (normalized to ${M_*}$) of blue-sequence E/S0s with those of other galaxies, and we examine the possible fractional growth in ${M_*}$ given current rates of star formation. The question of how efficiently the atomic gas might flow inward and condense into molecular gas is beyond the scope of this paper; however, we discuss preliminary evidence from our own and others' work that suggests efficient conversion is plausible (elaborated in a forthcoming paper: \\citealt[hereafter K09b]{kannappan09b}).  Section \\ref{section.sample} describes our statistically representative sample of red- and blue-sequence E/S0s from the NFGS, and presents new Green Bank Telescope\\footnote{The National Radio Astronomy Observatory is a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc.} (GBT) data for the galaxies with $M_* < 4\\times10^{10}\\,M_{\\odot}$, which are expected to show the most disk growth \\citepalias{kannappan09a}. We also present a tabulation of HI data for the full NFGS. Section \\ref{section.gasreservoirs} compares atomic gas masses of blue-sequence E/S0s with those of red-sequence E/S0s and galaxies of later-type morphology within the NFGS. We also compare to the sample of \\citet{sage06}. Section \\ref{section.starformation} examines the fractional stellar mass growth possible for blue-sequence E/S0s given the current star formation rate, in two limiting scenarios --- constant (i.e., allowing gas infall) and exponentially declining (i.e., closed box) star formation. In Section \\ref{section.sfmechanism}, we discuss gas exhaustion and gas inflow timescales, and we examine evidence for bursty star formation in blue-sequence E/S0s, which likely implies efficient conversion of HI to ${\\rm H_2}$. We conclude with a discussion of the evolutionary fates of blue-sequence E/S0s.  Appendix \\ref{section.appendix} details features of our new HI data for interesting individual galaxies. In this paper, we assume $H_0$ = 70 ${\\rm km\\, s^{-1}\\, Mpc^{-1}}$. ",
        "conclusions": "\\label{section.discussion} \\citet*{kannappan09a} recently identified a population of E/S0s that reside on the blue sequence in color vs. stellar mass space, where spiral galaxies typically reside. Blue-sequence E/S0s increase in numbers below the bimodality mass ($M_b \\sim 3\\times10^{10}\\,M_{\\odot}$) and especially below the gas richness threshold ($\\sim$5$\\times10^{9}\\,M_{\\odot}$). These sub-$M_b$ blue-sequence E/S0s are characterized by fairly regular morphology, and many appear to be rebuilding disks \\citepalias{kannappan09a}. Blue-sequence E/S0s also fall between red-sequence E/S0s and spirals in the $M_*$-radius and $M_*$-$\\sigma$ relations, suggesting that they may be a transitional population \\citepalias{kannappan09a}.  In this paper, we examined the atomic gas content of blue-sequence E/S0s below $\\sim$$M_b$ to determine whether they have large enough gas reservoirs to transform to later-type morphology. In a representative sample drawn from the Nearby Field Galaxy Survey, we find that blue-sequence E/S0s have substantial atomic gas masses in the range of $10^{7}$--$10^{10}\\, {M_{\\odot}}$, comparable to the gas masses of spiral and irregular galaxies in the same stellar mass range. Blue- and red-sequence E/S0s have distinctly different atomic gas-to-stellar mass ratios, with most blue-sequence E/S0s in the range of 0.1--1.0 and most red-sequence E/S0s $<$0.1. This suggests significantly greater potential for morphological transformation in blue-sequence E/S0s than red-sequence E/S0s.  Combining these atomic gas masses with current rates of star formation, we find that many of the sub-$M_b$ blue-sequence E/S0s can form new stars in the range of 10--60$\\%$ of their current stellar masses within 3 Gyr, in both constant (i.e., allowing gas infall) and exponentially declining (i.e., closed box) star formation scenarios, provided that the atomic gas reservoir is available for star formation. In the constant star formation scenario, we find that about half of the sub-$M_b$ blue-sequence E/S0 systems will exhaust their gas reservoirs in $\\lesssim$3 Gyr if no fresh gas infall is permitted.  Because of the lack of spatial resolution in our HI data, we cannot say for certain that the gas is readily available for star formation. We find evidence, however, which indicate that fresh gas may be brought inwards and made available for star formation on timescales shorter than the gas exhaustion timescale. We estimate the dynamical timescale for infall of extra-planar gas to be on average $<$1 Gyr, shorter than the gas exhaustion time for most of our galaxies. The frequency of inflow for gas trapped in diffuse disks may be dominated by the rate of minor mergers/interactions, which simulations find to occur every 1.4 Gyr or less for progenitor mass ratios of 1:100 or smaller \\citep{fakhouri08}. As evidence of such events, we find that sub-$M_b$ blue-sequence E/S0s are more often blue-centered than the general galaxy population, where blue-centeredness is measured relative to the typical color gradient of galaxies at a given luminosity. \\citet{kannappan04a} find blue-centeredness to correlate with morphological peculiarities and companions, which supports the externally triggered gas inflow scenario. For blue-sequence E/S0s, we find a relationship between blue-centeredness and variations in specific star formation rates relative to typical reference values, suggesting that such inflows may be episodic and trigger bursts of star formation. In summary, this work clearly confirms that blue-sequence E/S0s have both the gas reservoirs and the potential for sustained star formation necessary for significant disk growth, consistent with evolution toward later-type morphology if the spatial distribution of the gas is extended. While our sample of 27 blue- and red-sequence E/S0s is sufficient for the analysis presented in this paper, it is important to extend our work to a larger sample of sub-$M_b$ E/S0s for more robust statistics. The multi-wavelength RESOLVE Survey underway \\citep{kannappan08conf} would be an ideal data set for such a study.  Active follow-up includes obtaining IRAM CO(1--0) and CO(2--1) spectra to quantify the molecular gas content in these galaxies, which may be a large fraction of the gas mass content and further extend the potential for morphological transformation in blue-sequence E/S0s. We are also obtaining VLA HI and CARMA CO(1--0) maps to resolve the distribution of the atomic and molecular gas in blue-sequence E/S0s. HI maps will allow us to look for extended gas disks, small companions, and/or signs of interactions. CO(1--0) maps may reveal inner disky ``pseudobulges'' growing in tandem with extended disks, as would be expected during the formation of late type galaxies."
    },
    "0911/0911.5004_arXiv.txt": {
        "abstract": "The fact that one can associate thermodynamic properties with horizons brings together principles of quantum theory, gravitation and thermodynamics and  possibly offers a window to the nature of quantum geometry. This review discusses certain aspects of this topic concentrating on new insights gained from some recent work. After a brief introduction of the overall perspective, Sections 2 and 3 provide the pedagogical background on the geometrical features of bifurcation horizons,  path integral derivation of horizon temperature, black hole evaporation, structure of \\LL\\ models, the concept of Noether charge and its relation to horizon entropy. Section \\ref{sec:thermosecond} discusses several conceptual issues introduced by the existence of temperature and entropy of the horizons. In Section 5 we take up the connection between horizon thermodynamics and gravitational dynamics and describe several peculiar features which have no simple interpretation in the conventional approach. The next two sections  describe the recent progress achieved in an alternative perspective of gravity. In Section \\ref{sec:emergentspacetime} we provide a thermodynamic interpretation of the field equations of gravity in any diffeomorphism invariant theory and  in Section \\ref{sec:gravitystory}  we obtain the field equations of gravity from an entropy maximization principle. The last section provides a summary. ",
        "introduction": "\\label{sec:intro}  Soon after Einstein developed  the gravitational field equations, Schwarzschild found the simplest exact solution to these equations, describing a spherically symmetric spacetime. When expressed in a natural coordinate system which makes the symmetries of the solution obvious, it leads to the line interval: \\begin{equation} ds^2= -f(r)dt^2+\\frac{dr^2}{f(r)}+r^2(d\\theta^2+\\sin^2\\theta d\\phi^2) \\end{equation} where $f(r)\\equiv 1-(r_g/r)$ with $r_g \\equiv 2GM/c^2=2M$, in units with $G=c=1$. It was immediately noticed that this metric exhibits a curious pathology. One of the metric coefficients, $g_{tt}$ vanished on a surface $\\mathcal{H}$, of finite area $4\\pi r_g^2$, given by  $r=r_g$, while another metric coefficient $g_{rr}$ diverged on the same surface. After some initial confusion, it was realized that the singular behaviour of the metric is due to bad choice of coordinates and that the spacetime geometry is well behaved at $r=r_g$. However, the surface $\\mathcal{H}$ acts as a horizon blocking the propagation of information from the region $r<r_g$ to the region $r>r_g$. This leads to several new features in the theory, many of which, even after decades of investigation, defies a complete understanding. The most important amongst them is the relationship between physics involving horizons and thermodynamics.  This connection, between black hole dynamics involving the horizon and the laws of thermodynamics, became apparent as  a result of the research  in early 1970s. Hawking proved \\cite{areatheorem} that in any classical process involving the black holes, the sum of the areas of a black hole horizons cannot decrease which, with hindsight, is reminiscent of the behaviour of entropy in classical thermodynamics. This connection was exploited by Bekenstein in his response to a  conundrum raised by John Wheeler. Wheeler pointed out  that  an external observer can drop material with non-zero entropy into the inaccessible region beyond the horizon thereby reducing the entropy accessible to outside observers.\\footnote{John Wheeler posed this as a question to Bekenstein:  What happens if you mix cold and hot tea and pour it down a horizon, erasing all traces of ``crime\" in increasing the entropy of the world? This is based on what Wheeler told me in 1985, from \\textit{his} recollection of events; it is also mentioned in his book, see page 221 of Ref. \\cite{wheeler}. I have heard somewhat different versions from other sources.} Faced with this difficulty, Bekenstein came up with the idea that the black hole horizon should be attributed an entropy which is proportional to its area \\cite{Bekenstein:1972tm,Bekenstein:1973ur,Bekenstein:1974ax}. It was also realized around this time that one can formulate four laws of black hole dynamics in a manner analogous to the laws of thermodynamics \\cite{Bardeen:1973gs}. In particular, if a physical process (say, dropping of small amount of matter into the Schwarzschild black hole) changes the mass of a black hole by $\\delta M$ and the area of the event horizon by $\\delta A$, then it can be proved that \\begin{equation} \\delta M = \\frac{\\kappa}{8\\pi} \\delta A = \\frac{\\kappa}{2\\pi}\\delta\\left(\\frac{A}{4}\\right) \\label{firstlaw} \\end{equation} where $\\kappa = M/(2M)^2 = 1/4M$ is called the surface gravity of the horizon. This suggests an analogy with the thermodynamic law $\\delta E = T \\delta S$ with $S\\propto A$ and $T\\propto \\kappa$.  However, classical considerations alone cannot determine the proportionality constants. (We will see later that quatum mechanical considerations suggest $S=(A/4)$ and $T=\\kappa/2\\pi$, which is indicated in the second equality in \\eq{firstlaw}.)  In spite of this, Bekenstein's idea did \\textit{not} find favour with the community  immediately; the fact that laws of black hole dynamics have an uncanny similarity with the laws of thermodynamics was initially considered to be only a curiosity. (For a taste of history, see e.g \\cite{kst}.) The key objection at that time  was the following. If black holes possess entropy as well as energy (which they do), then they must have a non-zero temperature and must radiate --- which seemed to contradict the view that nothing can escape the black hole horizon. Investigations  by Hawking, however, led  to the discovery that a non-zero temperature should be attributed to the black hole horizon \\cite{Hawking:1975sw}. He found that black holes, formed by the collapse of matter, will radiate particles with a thermal spectrum at late times, as detected by a stationary observer at large distances. This result, obtained from the study of quantum field theory in the black hole spacetime, showed that one can consistently attribute to the black hole horizon an entropy \\textit{and} temperature.   An immediate question that arises is whether this entropy  is the same as the ``usual entropy\". If so, one should be able to show that, for any processes involving matter and black holes, we must have $d(S_{BH}+S_{matter})/dt\\geq 0$ which goes under the name generalized second law (GSL).  One simple example in which the area (and thus the entropy) of the black hole \\emph{decreases} is in the emission of Hawking radiation itself; but the GSL holds since the thermal radiation produced in the process has entropy. It is generally believed that GSL always holds though a completely general proof is difficult to obtain. Several thought experiments, when analyzed properly, uphold this law (see, for example, Ref. \\cite{Unruh:1982ic}) and a proof is possible under different sets of assumptions  \\cite{tenproofs}. All these suggest that the area of the black hole corresponds to an entropy which is  same  as the ``usual entropy''.  These ideas  can be extended to black hole solutions in more general  theories   than just Einstein's gravity. The temperature $T$  can be determined by    techniques like, for example,  analytic continuation to imaginary time (see Sec. \\ref{thermalden}) which depends only on the metric and not on the field equations which led to the metric. But the concept of entropy needs to  generalized in these models and will no longer  be one quarter of the area of horizon. This could be done by using the  first law of black hole dynamics itself, say, in the form $TdS = dM$. Since the temperature is known, this equation can be integrated to determine $S$. This was done by Wald and it turns out that the entropy can be related to a conserved charge called Noether charge which arises from the diffeomorphism invariance of the theory \\cite{Wald:1993nt}. Thus,  the notions of entropy and temperature  can be attributed to  black hole solutions in a wide class of theories.   This raises the  question: What are the degrees of freedom responsible for the black hole entropy? There have been several attempts in  the literature to answer this question  both with and  without inputs from quantum gravity models. A statistical mechanics derivation of entropy was originally attempted in \\cite{Gerlach:1976hbh}; the entropy  has been interpreted as the logarithm of: (a)  the number of ways in which black hole might have been formed \\cite{Bekenstein:1973ur,Hawking:1976de}; (b) the number of internal black hole states consistent with a single black hole  exterior  \\cite{Bekenstein:1975tw} and (c)  the number of horizon quantum states \\cite{wheeler,'tHooft:1990fr,Susskind:1993if}. There are also  other ideas which are more formal and geometrical \\cite{Jacobson:1994vj,Visser:1993nu,Banados:1994qp,Susskind:1993ws}, or  based on thermo-field theory \\cite{Frolov:1998vs,Frolov:1997up,Frolov:1997aj} just to name two possibilities. In addition, considerable amount of work has been done in calculating black hole entropy based on different candidate models for quantum gravity. (We will briefly summarize some aspects of these  in Sec. \\ref{qgbh}; more extensive discussion as well as references to original literature can be found in the reviews \\cite{Rovelli:1998yv,Das:2000su}).  Clearly, there is no agreement in literature as regards the degrees of freedom which contribute to black hole entropy and the  attempts mentioned above   sample the divergent views.  In fact, once the answer is known, it seems fairly easy to come up with very imaginative derivations of the result.  A crucial new dimension is added to this problem when we study horizons  which are \\textit{not} associated with black holes. Soon after Hawking's discovery of a temperature associated with the black hole horizon, it was realized that this result was not confined to black holes alone. The study of quantum field theory in \\textit{any} spacetime with a horizon showed that all horizons possess  temperatures \\cite{Fulling:1973md,Davies:1975th,PIdesitter}. In particular, an observer who is accelerating through the vacuum state in flat spacetime perceives a horizon and will attribute to it \\cite{Unruh:1976db} a temperature $T= \\kappa / 2\\pi$ proportional to her acceleration $\\kappa$. (For a review, see Refs.\\cite{Birrel:bkqft,mukhanovwini,Dewitt:1975ys,Takagi:1986kn,Sriramkumar:1999nw,Brout:1995rd,Wald:1999vt}.) The situation regarding entropy --- especially whether one should attribute entropy to all horizons --- remains  unclear. In fact, many of the attempts    to interpret \\textit{black hole} entropy mentioned earlier cannot be generalized to interpret the entropy of other horizons. So a unified understanding of horizon thermodynamics encompassing both entropy  and temperature still remains elusive --- which will be one of the key issues we will  discuss in detail in this review.  This is also closely related to the question of whether gravitational dynamics --- in particular, the field equations --- have any relation to the horizon thermodynamics. Given a spacetime metric with a horizon (which may or may not be a solution to Einstein's equations) one can study quantum field theory in that spacetime and discover that the horizon behaves like a black body with a given temperature. \\textit{At no stage in such an analysis do we need to invoke the gravitational field equations.} So it is reasonable to doubt whether the dynamics of gravity has anything to do with horizon thermodynamics and one may --- at first --- think that there is no connection between the two.  Several recent investigations have shown, however, that there is indeed a deeper connection between gravitational dynamics and horizon thermodynamics (for a recent review, see Ref.~\\cite{tpdialogue}). For example, studies have shown that: \\begin{itemize} \\item Gravitational field equations in \\textit{a wide variety of theories}, when evaluated on a horizon, reduce to  a thermodynamic identity $TdS=dE+PdV$. This result, first pointed out in Ref.\\cite{tdsingr}, has now been demonstrated  in several cases like the  stationary axisymmetric horizons and evolving spherically symmetric horizons in Einstein gravity, static spherically symmetric horizons and dynamical apparent horizons in Lovelock gravity, three dimensional BTZ black hole horizons, FRW cosmological models in various gravity theories and even  in the case Horava-Lifshitz gravity (see Sec.~\\ref{sec:connection} for detailed references). If horizon thermodynamics has no deep connection with gravitational dynamics, it is not possible to understand why the field equations should encode information about horizon thermodynamics. \\item  Gravitational action functionals in a wide class of theories have a a surface term and a bulk term. In the conventional  approach, we \\textit{ignore} the surface term completely (or cancel it with a counter-term) and obtain the field equation from the bulk term in the action. Therefore, any solution to the field equation obtained by this procedure is logically independent of the nature of the surface term. But  when the \\textit{surface term} (which was ignored) is evaluated at the horizon that arises in any given solution, it  gives the entropy of the horizon! Again, this result extends far beyond Einstein's theory to situations in which the entropy is \\textit{not}  proportional to horizon area. This is possible only because there is a specific holographic relationship \\cite{TPhol002,TPgravquantum,ayan} between the surface term and the bulk term which, however, is  an unexplained feature in the conventional approach to gravitational dynamics. Since the surface term has the thermodynamic interpretation as the entropy of horizons, and is related holographically to the bulk term, we are again led to suspect an indirect connection between spacetime dynamics and horizon thermodynamics. \\end{itemize}  Based on these features --- \\textit{which have no explanation in the conventional approach} --- one can argue  that there  is a  conceptual reason  to revise  our perspective towards spacetime (Sec.~\\ref{sec:emergentspacetime} and \\ref{sec:gravitystory}) and relate horizon thermodynamics with gravitational dynamics. This approach should work for a wide class of theories far more general than just Einstein gravity. This will be the new insight which we will focus on in this review.  To set the stage for this future discussion,  let us briefly describe this approach and summarize the conclusions. We begin by examining more closely the implications of the existence of temperature for horizons.  In the study of normal macroscopic systems --- like, for example, a solid or a gas --- one can \\textit{deduce} the existence of microstructure just from the fact that the \\textit{object can be heated}. The supply of energy in the form of heat needs to be stored in some form in the material which is not possible unless the material has microscopic degrees of freedom. This was  the insight of Boltzmann which led him to suggest that heat is essentially a form of motion of the microscopic constituents of matter. That is, the existence of temperature  is sufficient for us to infer the existence of microstructure without any direct experimental evidence.  The thermodynamics of the horizon shows that we can actually heat up a spacetime, just as one can heat up a solid or a gas. An unorthodox way of doing this would be to take some amount of matter and arrange it to collapse and form a black hole. The Hawking radiation emitted by the black hole can be used to heat up, say, a pan of water just as though the pan was kept inside a microwave oven. In fact the same result can be achieved by just accelerating through the inertial vacuum carrying the pan of water which will eventually be heated to a temperature proportional to the acceleration. These processes show that the temperatures of the horizons are as ``real'' as any other temperature. Since they arise in a class of hot  spacetimes, it follows \\textit{\\`{a} la} Boltzmann that the spacetimes should possess microstructure.   In the case of a solid or gas, we know the nature of this microstructure from atomic and molecular physics. Hence, in principle, we can work out the thermodynamics of these systems from the underlying statistical mechanics. This is not possible in the case of spacetime because we have no clue about its microstructure. However, one of the remarkable features of thermodynamics --- in contrast to statistical mechanics --- is that the thermodynamic description is fairly insensitive to the details of the microstructure and can be developed as a fairly broad frame work. For example, a thermodynamic identity like $TdS = dE+PdV$ has a universal validity and the information about a \\textit{given} system is only encoded in the form of the entropy functional $S(E,V)$. In the case of normal materials, this entropy arises because of our coarse graining over microscopic degrees of freedom which are not tracked in the dynamical evolution. In the case of spacetime, the existence of horizons for a particular class of observers makes it \\textit{mandatory} that these observers  integrate out degrees of freedom hidden by the horizon.  To make this notion clearer, let us start from the principle of equivalence which allows us to construct local inertial frames (LIF), around any event in an arbitrary curved spacetime. Given the LIF, we can next construct a local Rindler frame (LRF) by  boosting along one of the directions  with an acceleration $\\kappa$. The observers at rest  in the LRF will perceive  a patch of null surface in LIF as  a horizon $\\mathcal{H}$ with temperature $\\kappa/2\\pi$. These local Rindler observers and the freely falling inertial observers will attribute different thermodynamical properties to matter in the spacetime. For example, they will attribute different temperatures and entropies to the vacuum state as well as excited states of matter fields. When some matter with energy $\\delta E$ moves close to the horizon --- say, within a few Planck lengths because,  formally, it takes infinite Rindler time for matter to actually cross $\\mathcal{H}$ --- the local Rindler observer will consider it to have  transfered an entropy  $\\delta S = (2\\pi/\\kappa) \\delta E$ to the horizon degrees of freedom. We will  show (in Sec.~\\ref{sec:emergentspacetime}) that, when the metric satisfies the field equations of any diffeomorphism invariant theory, this transfer of entropy can be given \\cite{tp09papers} a  geometrical interpretation as the change in the  entropy of the horizon.   This result allows us to associate an entropy functional with the null surfaces which the local Rindler observers perceive as  horizons. We can now demand that the sum of the horizon entropy and the entropy of matter that flows across the horizons (both as perceived by the local Rindler observers), should be an extremum for all observers in the spacetime. This leads \\cite{TPgravitystory} to a constraint on the geometry of spacetime which can be stated, in $D=4$, as \\begin{equation} (G_{ab} - 8\\pi T_{ab}) n^a n^b =0 \\label{Eenn} \\end{equation} for all null vectors $n^a$ in the spacetime. The general solution to this equation is given by $G_{ab} = 8\\pi T_{ab} + \\rho_0 g_{ab}$ where $\\rho_0$ has to be a constant because of the conditions $\\nabla_a G^{ab} =0 = \\nabla_aT^{ab}$. Hence the thermodynamic principle leads uniquely to Einstein's equation with a cosmological constant in 4-dimensions. Notice, however, that \\eq{Eenn} has a new symmetry and is invariant \\cite{TPgravimmune,TPgravijtmp} under the transformation $T_{ab} \\to T_{ab} + \\lambda g_{ab}$ which the standard Einstein's theory does not posses. (This has important implications for the cosmological constant problem  \\cite{TPadvscilett} which we will discuss in Sec. \\ref{sec:cc}.) In $D>4$, the same entropy maximization leads to a more general class of theories called \\LL\\ models (see Sec. \\ref{sec:llgravity}).  We can now remedy another conceptual shortcoming of the conventional approach. An unsatisfactory feature of all theories of gravity is that the field equations do not have any direct physical interpretation. The lack of an elegant principle which can lead to the dynamics of gravity (``how matter tells spacetime to curve'') is quite striking when we compare this situation with the kinematics of gravity (``how spacetime makes the matter move''). The latter can be determined through the principle of equivalence by demanding that all freely falling observers, at all events in spacetime, must find that the equations of motion for matter  reduce to their special relativistic form.  In the alternative perspective, \\eq{Eenn} arises from our demand that the thermodynamic extremum principle should hold for \\textit{all} local Rindler observers. This is identical  to the manner in which  freely falling observers are used to determine how gravitational field influences matter. Demanding the validity of special relativistic laws for the matter variables, as determined by all the freely falling observers, allows us to determine the influence of gravity on matter. Similarly, demanding the maximization of entropy of horizons (plus matter), as measured by all local Rindler observers, leads to the dynamical equations of gravity.    In this review, we will examine several aspects of these features and will try to provide, in the latter part of the review, a synthesis of ideas which offers an interesting new perspective on the nature of gravity that makes the connection between horizon thermodynamics and gravitational dynamics obvious and natural. In fact, we will show that the thermodynamic underpinning goes far beyond  Einstein's theory and encompasses a wide class of gravitational theories.  The review is organized as follows. In Section \\ref{sec:gravhorizon} we will review several features of horizons which arise in different contexts in gravitational theories.  In particular, we will describe some generic features of the spacetimes with horizons which will be important in the later discussions. In Section \\ref{sec:thermofirst} we shall provide a simple derivation of the temperature of a (generic) horizon using path integral methods. Section \\ref{sec:pireview} introduces the basic concepts related to path integrals and Section \\ref{thermalden} applies them to a spacetime with horizon to obtain the temperature. Some alternate ways of obtaining the temperature of horizons are discussed in Section \\ref{sec:complext}. The origin of Hawking radiation from  matter that collapses to form a black hole is discussed in Section \\ref{sec:hawrad}. The next two subsections describe the generalization of these ideas to theories other than Einstein's general relativity. The relationship between horizon entropy and the Noether charge is introduced in Section \\ref{sec:noetherent} and the structure of \\LL\\ models is summarized in Section \\ref{sec:llgravity}. Section \\ref{sec:thermosecond} discusses several conceptual issues raised by the existence of temperature and entropy of the horizons, concentrating on the nature of degrees of freedom which contribute to the entropy and the observer dependence of the concept of entropy. This discussion is continued in the next section where we take up the connection between horizon thermodynamics and gravitational dynamics and describe several peculiar features which have no simple interpretation in the conventional approach. In Section \\ref{sec:emergentspacetime} we provide a thermodynamic interpretation of the field equations of gravity in any diffeomorphism invariant theory. This forms the basis for the alternative perspective of gravity described in Section \\ref{sec:gravitystory} in which we obtain the field equations of gravity from an entropy maximization principle. The last section provides a summary.  We will use the signature $(- +++)$ and units with $G=\\hbar=c=1$. The Greek superscripts and subscripts will run over the spatial coordinates while the Latin letters will cover time coordinate as well as spatial coordinates. ",
        "conclusions": "\\label{sec:conclusions}   This review concentrated, true to the title, several new insights which  have been gained regarding the thermodynamical aspects of gravity. It presented a case arguing that: (a) The conventional approach in which one considers thermodynamical aspects of gravity as an interesting but subsidiary results of, say, doing quantum field theory in a spacetime with horizon, is fundamentally flawed. (b) There are several features of the theory which should be considered as strong hints favouring a fundamental revision of our approach towards gravity and spacetime. These hints appear in the form of peculiar relationships, especially in the structure of the action functionals describing gravity, and in the generality of the thermal phenomena they represent. (c) The last part of the review presented an alternate perspective which holds the promise for providing a more natural backdrop for understanding the relationship between gravitational dynamics and horizon thermodynamics.  It is useful to distinguish clearly (i) the mathematical results which can be rigorously proved from (ii) interpretational ideas which might evolve when our understanding of these issues deepen.  (i) From a purely algebraic point of view, without bringing in any physical interpretation or motivation, we can prove the following mathematical results: \\begin{itemize} \\item Consider a functional of null vector fields $n^a(x)$ in an arbitrary spacetime given by \\eq{ent-func-2} [or, more generally, by \\eq{generalS}]. Demanding that this functional is an extremum for all null vectors $n^a$ leads to the field equations for the background geometry given by $(2\\mathcal{G}_{ab}-T_{ab})n^an^b=0$ where $\\mathcal{G}_{ab}$ is given by \\eq{genEab1} [or, more generally, by \\eq{genEab}]. Thus field equations in a wide class of theories of gravity can be obtained from an extremum principle without varying the metric as a dynamical variable. \\item These field equations are invariant under the transformation $T_{ab} \\to T_{ab} + \\rho_0 g_{ab}$, which relates to the freedom of introducing a  \\cc\\ as an integration constant in the theory. Further, this symmetry forbids the inclusion of a cosmological constant  term in the variational principle by hand as a low energy parameter. \\textit{That is, we have found a symmetry which makes the bulk \\cc\\ decouple from the gravity.} When linearized around flat spacetime, the graviton inherits this symmetry and does not couple to the \\cc . \\item On-shell, the functional in \\eq{ent-func-2} [or, more generally, by \\eq{generalS}] contributes only on the boundary of the region. When the boundary is a horizon, this terms gives precisely the Wald entropy of the theory. \\end{itemize} It is remarkable that one can derive not only Einstein's theory uniquely in $D=4$ but even \\LL\\ theory in $D>4$ from an extremum principle involving the null normals \\textit{without varying $g_{ab}$ in an action functional!}.   (ii) To provide a physical picture behind these mathematical results, it is necessary to invoke certain effective degrees of freedom which can participate in the (observer dependent) thermodynamic interactions near any local patch of a null surface that acts as a horizon for certain class of observers. At present we have no deep understanding of how this comes about but, at a qualitative level, the physical picture is made of the following  ingredients: \\begin{itemize} \\item Assume that the spacetime is endowed with certain microscopic degrees of freedom capable of exhibiting thermal phenomena. This is just the Boltzmann paradigm: \\textit{If one can heat it, it must have microstructure!}; and one can heat up a spacetime. \\item Whenever a class of observers perceive a horizon, they are ``heating up the spacetime'' and the  degrees of freedom close to a horizon  participate in a very \\textit{observer dependent} thermodynamics. Matter which flows close to the horizon (say, within a few Planck lengths of the horizon) transfers energy to these microscopic, near-horizon, degrees of freedom \\textit{as far as the observer who sees the horizon is concerned}. Just as  entropy of a normal system at temperature $T$ will change by $\\delta E/T$ when we transfer to it an energy $\\delta E$, here also an entropy change will occur. (A freely falling observer in the same neighbourhood, of course, will deny all these!) \\item We proved that when the field equations of gravity hold, one can interpret this entropy change in  a purely geometrical manner involving the Noether current. From this point of view, the  normals $n^a$ to local patches of null surfaces are related to the (unknown) degrees of freedom that can participate in the thermal phenomena involving the horizon. These degrees of freedom seem to obey standard rules of thermodynamics, including equipartition. \\item Just as demanding the validity of special relativistic laws with respect to all freely falling observers leads to the kinematics of gravity, demanding the local entropy balance in terms of the thermodynamic variables, as perceived by local Rindler observers, leads to the field equations of gravity in the form $(2\\mathcal{G}_{ab}-T_{ab})n^an^b=0$.  \\end{itemize}   As stressed in earlier sections, this involves a new layer of observer dependent thermodynamics. In particular, since observers in different states of motion will have different regions of spacetime accessible to them --- for example, an observer falling into a black hole will not perceive a horizon in the same manner as an observer who is orbiting around it --- we are forced to accept that the notion of entropy is an observer dependent concept. Fundamentally,  this is no different from the fact different freely falling observers will measure physical quantities differently comapared to non-geodesic observers; but in this case standard rules of special relativity allow us to translate the results between the observers. We do not yet have  a similar set of rules for quantum field theory in noninertial frames. It seems necessary to  integrate the entire thermodynamic machinery (involving what we usually consider to be the `real' temperature) with this notion of LRFs having their own (observer dependent) temperature.  This requires an intriguing relationship between quantum fluctuations and thermal fluctuations. As an illustration, consider the relation $\\delta E=T\\delta S$ obtained in Sec. \\ref {sec:obsdependence} for the one-particle excited state which has a curious consequence when we take the nonrelativistic limit \\cite{equipartition}. The mode function $e^{imc^2 t/\\hbar} \\langle 0| \\phi(x)|1_{\\bf k}\\rangle$ corresponding to a one-particle state in either inertial frame or Rindler frame goes over to a wave function $\\psi(x)$ in the  nonrelativistic ($c\\to \\infty$), quantum mechanical, limit such that $\\psi(x)$ satisfies a Schrodinger equation with an accelerating potential $V=mgx$ when viewed from the second frame. In describing the motion of a wave packet corresponding to such a particle, the quantum mechanical averages will satisfy the relation $\\langle \\delta E \\rangle = mg \\langle \\delta x\\rangle = F \\langle \\delta x\\rangle$. On the other hand, given the thermal description in the local Rindler frame, we would expect a relation like $\\langle \\delta E \\rangle = T\\Delta S$ to hold suggesting that the entropy gradient $\\Delta S$ (due to the gradient  $\\Delta n$ in the microscopic degrees of freedom) present over a region $\\langle \\delta x\\rangle$  to give rise to a force $F = T \\Delta S/\\langle \\delta x\\rangle$. If one assumes that (i) $\\Delta S/k_B$ has to be quantized (based on the results of ref.\\cite{entropyideas}) in units of $2\\pi$ and (ii) $\\langle \\delta x\\rangle \\approx \\hbar/mc$ for a particle of mass $m$, then we reproduce $F=mg$ on using the Rindler temperature $k_BT = \\hbar g/2\\pi c$. Alternatively, if one assumes that the force $F = T \\Delta S/\\langle \\delta x\\rangle$ should be equal to $mg$, then the universality of the Rindler temperature for bodies with different $m$ arises if we use $\\langle \\delta x\\rangle = \\hbar / mc$. In this case --- which involves the quantum mechanical limit of a one-particle state in a non-inertial frame --- we need to handle simultaneously both quantum and thermal fluctuations. The expression $F = T \\Delta S/\\langle \\delta x\\rangle$ demands an intriguing interplay between thermal fluctuations (in the numerator, $T\\Delta S$, arising from the non-zero temperature and entropy in local Rindler frame) and the quantum fluctuations (in the denominator, $\\langle \\delta x\\rangle$, related to the intrinsic position uncertainty $\\hbar/mc$) for the theory to be consistent, including the choice of numerical factors.          At a conceptual level, this may be welcome when we note that every key progress in physics involved realizing that something we thought as absolute is  not absolute. With special relativity it was the flow of time and with general relativity it was the concept of global inertial frames and when we brought in quantum fields in curved spacetime it was the notion of particles and temperature.   Many of these technical issues possibly  can be tackled in more or less straightforward manner, though the mathematics can be fairly involved. But  they may not be crucial to the alternative perspective or its further progress. The latter will depend on more serious conceptual issues, some of which are the following:  (i) How come the microstructure of spacetime  exhibits itself indirectly through the horizon temperature even at scales much larger than Planck length and obeys an equipartition law (see e.g. \\eq{idn} and ref. \\cite{TPgravcqg04,equipartition})? This is possibly because the event horizon works as some kind of magnifying glass allowing us to probe trans-Planckian physics \\cite{Padmanabhan:1998jp,Padmanabhan:1998vr} but this notion needs to be made more precise.  (ii) How does one obtain the expression for spacetime entropy density from some microscopic model? In particular, such an analysis  --- even with a toy model --- should throw more light on why normals to local patches of null surfaces play such a crucial role as effective degrees of freedom in the long wavelength limit. Of course, such a model should also determine the expression for $P^{abcd}$ and get the metric tensor and spacetime as derived concepts - a fairly tall order!. (This is somewhat like obtaining theory of elasticity starting from a microscopic model for a solid, which, incidentally, is not a simple task either.)"
    },
    "0911/0911.2412_arXiv.txt": {
        "abstract": "We report on $\\gamma$-ray observations of the Crab Pulsar and Nebula using 8 months of survey data with the \\emph{Fermi} Large Area Telescope (LAT). The high quality light curve obtained using the ephemeris provided by the Nan\\c{c}ay and Jodrell Bank radio telescopes shows two main peaks stable in phase with energy. The first $\\gamma$-ray peak leads the radio main pulse by (281 $\\pm$ 12 $\\pm$ 21) $\\mu$s, giving new constraints on the production site of non-thermal emission in pulsar magnetospheres. The first uncertainty is due to $\\gamma$-ray statistics, and the second arises from the rotation parameters. The improved sensitivity and the unprecedented statistics afforded by the LAT enable precise measurement of the Crab Pulsar spectral parameters: cut-off energy at $E_{c}$~=~($5.8~\\pm~0.5~\\pm~1.2$)~GeV, spectral index of $\\Gamma$~=~($1.97~\\pm~0.02~\\pm~0.06$) and integral photon flux above 100 MeV of $(2.09~\\pm~0.03~\\pm~0.18$)~$\\times$~10$^{-6}$ cm$^{-2}$~s$^{-1}$. The first errors represent the statistical error on the fit parameters, while the second ones are the systematic uncertainties. Pulsed $\\gamma$-ray photons are observed up to $\\sim$~20 GeV which precludes emission near the stellar surface, below altitudes of around 4 to 5 stellar radii in phase intervals encompassing the two main peaks. A detailed phase-resolved spectral analysis is also performed: the hardest emission from the Crab Pulsar comes from the bridge region between the two $\\gamma$-ray peaks while the softest comes from the falling edge of the second peak. The spectrum of the nebula in the energy range 100 MeV -- 300 GeV is well described by the sum of two power-laws of indices $\\Gamma_{sync}$~=~($3.99~\\pm~0.12~\\pm~0.08$) and $\\Gamma_{IC}$~=~($1.64~\\pm~0.05~\\pm~0.07$), corresponding to the falling edge of the synchrotron and the rising edge of the inverse Compton components, respectively. This latter, which links up naturally with the spectral data points of Cherenkov experiments, is well reproduced via inverse Compton scattering from standard Magnetohydrodynamics (MHD) nebula models, and does not require any additional radiation mechanism. ",
        "introduction": "The Crab Nebula belongs to the class of filled-center supernova remnants (SNR) \\citep{Green 2006}, i.e. without any detected shell component, and is well studied in almost all wavelength bands of the electromagnetic spectrum from the radio ($10^{-5}$~eV) to very high energy $\\gamma$-rays (nearly $10^{14}$~eV). It is held to be the archetypical pulsar wind nebula, even though its physical and spectral properties are unique. It is associated with the supernova explosion reported by Chinese astronomers in 1054 AD. Several models (\\cite{Kennel and Coroniti 1984}, \\cite{de Jager and Harding 1992}, \\cite{de Jager et al. 1996} and references therein) describe the photon production processes taking place in this nebula. Synchrotron radiation from high energy electrons in the nebular magnetic field is responsible for the observed spectrum from radio to MeV, while inverse Compton (IC) scattering of the primary accelerated electrons off the synchrotron photons, far infrared and Cosmic Microwave Background (CMB) produces high energy $\\gamma$-rays. While these two mechanisms seem to provide a reasonable description of the overall non-thermal radiation of the Crab Nebula, one cannot exclude possible deviations from this simplified picture, and \\cite{Atoyan and Aharonian 1996} proposed that significant production of high energy $\\gamma$-rays by bremsstrahlung radiation of relativistic electrons could take place in the Crab filaments. This paper, reporting the results of a precise spectral analysis of the Crab Nebula between 100~MeV and 300~GeV, adds new elements to this discussion.  At the center of the nebula lies the Crab Pulsar, PSR~J0534+2200, one of the most energetic known pulsars (spin down power of $\\dot{E}$ = 4.6 $\\times$ 10$^{38}$ erg~s$^{-1}$), located at a distance of (2.0 $\\pm$ 0.2) kpc. Estimation of its characteristic age using its rotation period ($P~=~33$~ms) and derivative ($\\dot{P}~=~4.2 \\times 10^{-13}$~s/s) yields an age of 1240 years, close to the observational value.  The Energetic Gamma-Ray Experiment Telescope (EGRET), on orbit from 1991 to 2001 on board of the Compton Gamma Ray Observatory (CGRO), reported the high energy detection of the Crab Nebula and Pulsar \\citep{Nolan et al. 1993, de Jager et al. 1996, Kuiper et al. 2001}. A more detailed study of the $\\gamma$-ray emission was then provided by \\cite{Fierro et al. 1998}, presenting a complete phase-resolved spectral analysis of the EGRET data, and more recently by \\cite{Pellizzoni et al. 2009}, describing the first AGILE timing results on $\\gamma$-ray pulsars, including the Crab. Pulsations of the Crab Pulsar were reported above 25 GeV by the Major Atmospheric Gamma-ray Imaging Cherenkov Telescope (MAGIC) collaboration \\citep{Aliu et al. 2008}, with a light curve consistent with the one measured by EGRET.  Observations of the Crab Pulsar in high energy $\\gamma$-rays can provide strong constraints on the location of the $\\gamma$-ray emitting regions: above the polar caps \\citep{Daugherty and Harding 1996}, in the intermediate models like the slot gap \\citep{Muslimov and Harding 2004}, or far from the neutron star in the outer gaps \\citep{Romani 1996}. In particular, the spectral analysis and the phase-resolved behaviour examined in this paper may be used to discriminate between these models.  Successfully launched on June 11, 2008, the Large Area Telescope (LAT), aboard the \\emph{Fermi} Gamma-ray Space Telescope, formerly GLAST, offers the unique opportunity to study the high energy behaviour of the Crab Pulsar and Nebula in great detail. In this paper, we report the results of the analysis of the Crab region using 8 months of survey observations with the \\emph{Fermi}-LAT. In Sections~\\ref{radio} and \\ref{lat}, we describe the radio and $\\gamma$-ray observations used, while Section~\\ref{results} presents the results obtained from a detailed timing and spectral analysis of the LAT data. Finally, in Sections~\\ref{discussion} and \\ref{summary}, we discuss and summarize the main implications of these results for models of both the pulsar and the nebula. ",
        "conclusions": "\\label{discussion}  \\subsection{Synchrotron and inverse Compton emission from the Crab Nebula}  The Crab Nebula is detected across the whole electromagnetic spectrum from radio to very high energy $\\gamma$-rays. The total spectral energy distribution of this source is shown in Figure~\\ref{SED_nebula}, from soft to very-high energy $\\gamma$-rays. The spectral points obtained with the LAT data analysis are also represented (red points).  With a spectral index of $\\Gamma_{IC}$~=~($1.64~\\pm~0.05~\\pm~0.07$), the LAT results on the rising edge of the inverse Compton component are consistent with EGRET ($1.85^{+ 0.65}_{-0.5}$, \\cite{de Jager et al. 1996}). As can be observed in Figure~\\ref{SED_nebula}, the highest part of the LAT spectrum links up satisfactorily to the lower energy bound of the Cherenkov data points. Using the LAT spectral parameters scaled to the full pulse phase, we obtain a flux at 77 GeV of $(1.18~\\pm~0.22~\\pm~0.37)~\\times 10^{-14}$ cm$^{-2}$~s$^{-1}$~MeV$^{-1}$ which agrees with the MAGIC differential flux at this energy of $(1.14~\\pm~0.27~\\pm~0.34)~\\times~10^{-14}$ cm$^{-2}$~s$^{-1}$~MeV$^{-1}$. The Cherenkov and \\emph{Fermi}-LAT data now cover the entire inverse Compton peak, as can be seen in Figure~\\ref{SED_nebula}, and a break is expected at $\\sim 100$~GeV. Although no significant cut-off is observed in the LAT data with the current statistics, the determination of its position with an increased \\emph{Fermi}-LAT data sample would help the calibration of Cherenkov telescopes, as discussed in \\cite{Bastieri et al. 2005}.  The inverse Compton scattering of relativistic electrons on the synchrotron, far infrared, and cosmic microwave background radiation fields is considered to be the most probable mechanism for production of $\\gamma$-rays above 1 GeV. However, using a sophisticated approach carried out in the framework of the MHD flow of \\cite{Kennel and Coroniti 1984}, \\cite{Atoyan and Aharonian 1996} have commented on the apparent deficit of GeV photons in their calculations. Taking into account both EGRET and Cherenkov results and assuming a mean magnetic field which reproduces the very high energy spectrum, they proposed that the high $\\gamma$-ray flux observed by EGRET in comparison to their model is due to the enhancement of the bremsstrahlung emission from electrons captured in dense filaments. Figure~\\ref{SED_nebula} presents the broad-band energy spectrum of the Crab Nebula together with the inverse Compton model predictions from Figure~14 of \\cite{Atoyan and Aharonian 1996} for three different values of the mean magnetic field for the nebula. In view of the results obtained with the LAT, modeling the data does not require any additional emission component. The \\emph{Fermi}-LAT, in combination with the Cherenkov observations above 100~GeV, are in good agreement with the $\\gamma$-ray flux predicted from simple IC scattering when the magnetic field lies between 100~$\\mu$G and 200~$\\mu$G, i.e. below the canonical equipartition field of the Crab Nebula of 300~$\\mu$G. This result is consistent with the estimate of the magnetic field strength B~$\\sim$~140~$\\mu$G obtained by \\cite{Horns and Aharonian 2004}.  Concerning the low energy part of the nebular spectrum, the LAT spectral points, combined with COMPTEL's (taking into account statistical errors only for the latter), can be fitted with a power-law with an exponential cut-off, following \\cite{de Jager et al. 1996}. The cut-off energy is estimated at $E_{c, sync}$~=~(97~$\\pm$~12)~MeV. The higher value of this energy compared to that of \\cite{de Jager et al. 1996} is due to the larger flux obtained with \\emph{Fermi} than by EGRET for the synchrotron component. The fit is represented with a blue dashed curve in Figure~\\ref{SED_nebula}.  \\subsection{High energy emission from the Crab Pulsar}  The high-quality statistics obtained with the \\emph{Fermi}-LAT both on the light curve and the spectrum of the Crab Pulsar, allow a more detailed comparison with theoretical models than previously possible. Currently, there are two classes of models that differ in the location of the emission region. The first comprises polar cap (PC) models which place the emission near the magnetic poles of the neutron star \\citep{Daugherty and Harding 1996}. The second class consists of the outer gap (OG) models \\citep{Romani 1996}, in which the emission extends between the null charge surface and the light cylinder, and the two-pole caustic (TPC) models \\citep{Dyks and Rudak 2003} which might be realized in slot gap (SG) acceleration models \\citep{Muslimov and Harding 2004}, in which the emission takes place between the neutron star surface and the light cylinder along the last open field lines.   \\begin{table}% \\begin{center} \\caption{\\label{lead/lag_radio}The radio delay with respect to other frequencies. } \\begin{tabular}{ccc} \\hline\\hline Spectral band\t      & Radio delay\t  & Reference\\\\ & ($\\mu$s)  & \\\\ \\hline Optical \t&255 $\\pm$ 21\t   & (1) \\\\ X-rays         &344 $\\pm$ 40\t  & (2) \\\\ Hard X-rays\t &280 $\\pm$ 40      & (3) \\\\ $\\gamma$-rays (EGRET)\t&241 $\\pm$ 29\t   & (3) \\\\ $\\gamma$-rays (LAT)   &281~$\\pm$~12~$\\pm$~21\t& (4) \\\\ \\hline \\multicolumn{3}{l}{References: (1): \\cite{Oosterbroek et al. 2008}; (2): \\cite{Rots et al. 2004};}\\\\ \\multicolumn{3}{l}{(3): \\cite{Kuiper et al. 2003}; (4): this paper.} \\end{tabular} \\end{center} \\end{table}  Observations of the time delay between emission at different wavelengths have been reported previously: Table~\\ref{lead/lag_radio} summarizes the delay of the radio main pulse with respect to the first peak seen from optical to high energy $\\gamma$-rays. The LAT has a timing accuracy better than 1 $\\mu$s \\citep{Abdo et al. 2009c} and thus enables an accurate estimation of the absolute positions of the $\\gamma$-ray peaks: it was shown in Section~\\ref{phasos} that the first $\\gamma$-ray peak leads the radio main pulse by 0.0085 $\\pm$ 0.0005 $\\pm$ 0.0006 in phase, or (281~$\\pm$~12~$\\pm$~21)~$\\mu$s in time. Taking into account the presence of the LFC at phase 0.896~$\\pm$~0.002, the first radio peak leads the $\\gamma$-ray peak by phase 0.095~$\\pm$~0.002 in phase. Observations of the evolution of the peak positions with the energy allow detailed studies of the emission regions in the magnetosphere. The delay of the radio peaks compared to other wavelengths (optical, X- and $\\gamma$-rays) gives another constraint in the modeling of the emission processes taking place in pulsar magnetospheres. In particular, \\cite{Kuiper et al. 2003} reported that, in the framework of the three-dimensional outer gap model developed by \\cite{Cheng et al. 2000}, one can reproduce the delay of the radio main pulse by shifting the production site of the radio emission inwards toward the neutron star relative to that of high energy photons.  In the PC models, $\\gamma$-rays created near the neutron star surface interact with the intense magnetic fields resulting in a sharp turnover in the few to 10~GeV energy range, while OG and SG models predict a simple exponential cut-off. Furthermore, the maximum energy of observed pulsed photons must lie below any $\\gamma$-B pair production turnover threshold, providing a lower bound to the altitude of emission. We can use the observed phase-averaged cutoff energy ($\\sim$~6 GeV) to estimate a minimum emission height as $r \\ge (\\epsilon_{max} B_{12}/1.76 \\, {\\rm GeV})^{2/7}P^{-1/7} R_\\ast$, where $\\epsilon_{max}$ is the unabsorbed photon energy, $P$ is the spin period and the surface field is $10^{12}B_{12}$\\,G \\citep{Baring 2004}. Using the parameters of the Crab Pulsar ($P~=~33$~ms, $B_{12}=3.78$), one obtains $r > 3.4 R_{\\ast}$ which precludes emission near the stellar surface. Since we see pulsed photons up to almost $\\epsilon_{max} \\approx 20$ GeV in the two main pulse peaks, emission at these phases must arise at $r > 4.8 R_{\\ast}$, with a strict lower bound of $r > 3.7 R_{\\ast}$ applying to the choice of 8 GeV, the lower energy in the highest data point window for the phase-resolved spectra in Figure~\\ref{resultats_phase_resolved}. A similar lower bound to the emission altitude was recently reported by the MAGIC collaboration using the hyper-exponential cutoff energy observed on the Crab Pulsar spectrum \\citep{Aliu et al. 2008}. We should note here that the cut-off energy derived by the MAGIC collaboration for a simple exponential cut-off ($17.7 \\pm 2.8 \\pm 5.0$)~GeV is higher than the one obtained with the \\emph{Fermi}-LAT data, $E_{c}$~=~($5.8~\\pm~0.5~\\pm~1.5$)~GeV. However, the cut-off energy obtained with the LAT using the softer EGRET spectrum ($\\gamma$~=~2.022) as done by MAGIC is within the uncertainties of the MAGIC value.  To estimate the pulsed high energy $\\gamma$-ray efficiency $\\eta$ of a pulsar, one needs to know the total luminosity radiated $L_{\\gamma}$. It can be estimated using $L_{\\gamma} = 4\\pi f_{\\Omega}F_{obs}D^2$, where $F_{obs}$~=~(1.31~$\\pm$~0.02~$\\pm$~0.02)$~\\times$~10$^{-9}$ erg~cm$^{-2}$~s$^{-1}$ is the observed phase-averaged energy flux over 100 MeV, $D$~=~(2.0~$\\pm$~0.2)~kpc is the distance to the pulsar and $f_{\\Omega}$ is a correction factor that takes into account the beaming geometry, depending upon the magnetic inclination angle $\\alpha$ and the Earth viewing angle $\\zeta$ from the rotation axis. For the Crab Pulsar, $\\zeta$ is estimated to be (63 $\\pm$ 2)$^{\\circ}$ from X-ray observations of the Crab Nebula torus \\citep{Ng and Romani 2008}. The estimated value of $\\alpha$ depends on the emission model used to interpret the data \\citep{Watters et al. 2009}. The slot gap or two pole caustic model best reproduces the observed pulse profiles for $\\alpha \\sim$ 55 -- 60$^{\\circ}$, whereas for the outer gap model $\\alpha \\sim 70^{\\circ}$ gives the best result. Optical polarization measurements yield $\\alpha$ estimates consistent with these~\\citep{Slowikowska et al. 2009}. The corresponding correction factors $f_{\\Omega}$ are then equal to 1.1 for the TPC and 1.0 for the OG models. This yields a luminosity of (6.25~$\\pm$~0.15~$\\pm$~0.15)~$\\times$~10$^{35}$ erg~s$^{-1}$ above 100 MeV. This value is consistent with the heuristic luminosity law mentioned in \\cite{Arons 1996} and \\cite{Watters et al. 2009}, according to which $\\eta \\propto \\dot{E}^{-1/2}$ and verified by several $\\gamma$-ray pulsars such as Vela, PSR J2021+3651 (assuming a distance of the order of 2 -- 4 kpc), Geminga, CTA1, etc. For a neutron star moment of inertia of 10$^{45}$ g~cm$^{2}$, the pulsed high energy $\\gamma$-ray efficiency $\\eta$ can be derived from the luminosity and the spin down power $\\dot{E}$: $\\eta~=~L_{\\gamma}~/~\\dot{E}~=~(1.36~\\pm~0.03~\\pm~0.03) \\times~10^{-3}$ above 100 MeV.  Knowing the value of the Earth viewing angle $\\zeta$ of (63 $\\pm$ 2)$^{\\circ}$, the main peak separation, which is of the order of 40\\% of the phase, would be expected to be smaller in the Polar Cap model \\citep{Watters et al. 2009}. This mismatch persists even if moderate altitudes are considered: the open field line cone opening angle enlargens to around $12.3^{\\circ}$ for the Crab at altitudes of around 6 stellar radii, the minimum bound inferred above from the observed absence of magnetic pair attenuation below 20 GeV. This opening angle is still somewhat too small according to the \\cite{Watters et al. 2009} analysis to generate the observed main peak phase separation."
    },
    "0911/0911.5718_arXiv.txt": {
        "abstract": "The results of speckle interferometric measurements of binary and multiple stars conducted in 2008 and 2009 at the Blanco and SOAR 4-m telescopes in Chile are presented. A total of 1898 measurements of 1189 resolved pairs or sub-systems and 394 observations of 285 un-resolved targets are listed. We resolved for the first time 48 new pairs, 21 of which are new sub-systems in close visual multiple stars. Typical internal measurement precision is 0.3\\,mas in both coordinates, typical companion detection capability is $\\Delta m \\sim 4.2$ at 0\\farcs15 separation. These data were obtained with a new electron-multiplication CCD camera; data processing is described in detail, including estimation of magnitude difference, observational errors, detection limits, and analysis of artifacts. We comment on some newly discovered pairs and objects of special interest. ",
        "introduction": "Speckle  interferometry at  4-m telescopes  has provided  the  bulk of binary star  measurements over the  last two decades,  giving material for calculation of orbits  and other studies. Unfortunately, access to 4-m  telescopes  has been  intermittent,  especially  in the  southern hemisphere  (speckle data from  the WIYN  telescope were  published by \\citet{Horch99,Horch02}). Here we present the results of two observing programs carried  out in  2008-2009 to help  rectify this  problem. We concentrate  on the  technique  and measurement  results, leaving  the exploitation of these data  sets for further publications.  The general neglect of  close southern hemisphere  binary stars has allowed  us to determine new orbits for a  dozen pairs and correct preliminary orbits for around one  hundred other pairs. These orbits  are currently being prepared.   Continued measurements of visual binary stars are needed for many reasons. One of the most obvious, yet most difficult tasks is to establish the orbital elements of known binaries. Long orbital periods where only a short arc is covered, on the one hand, and the lack of continuous coverage of fast systems, on the other hand, prevent calculation of orbits or cause erroneous orbits to be published. Although stellar masses derived solely from visual orbits are typically of inferior accuracy compared to other techniques, reliable orbital elements are needed nevertheless in order to be able to predict stellar positions and to study individual systems of astrophysical importance (including those with planetary companions). Multiplicity affects stellar evolution in many different ways, as illustrated e.g. by the dramatic story of the quadruple system Regulus \\citep{Rappaport2009}.   Modern hydrodynamical simulations open the way to understanding the formation of binary and multiple systems \\citep[e.g.][]{Bate08}, so good observational data on multiplicity statistics become critical for further progress. The current census of stellar multiplicity is incomplete even in the solar neighborhood. In updating the seminal work of \\citet{DM91}, \\citet{Raghavan09} determined that the fraction of triple and higher-order multiples among G-type dwarfs within 25\\,pc of the Sun was underestimated by as much as two times. These additional triples and quadruples were found around systems which were previously considered binary; i.e., the number of singles remained essentially fixed. For a well studied group, the best location to find new companions is around double or higher order multiple systems and not single stars. Therefore, we also searched for new close sub-systems in wide visual binaries, focusing on binaries in the solar vicinity.  Relative orientation of the orbits in triple stars is an indicator of their formation mechanisms and dynamical evolution \\citep{ST02,FT07}. One of the goals of the present program was to increase the number of multiple stars with known sense of relative orbital motion in sub-systems by getting additional observations.  In any statistical study, it is important not only to detect new companions, but also to establish the detection limits, so that an absence of companions can be translated into constraints on their parameters. The linearity of the detector allows us to here establish reliable detection limits for each observation, through careful data analysis and modeling. Owing to the new detector and data processing, our observations reach larger magnitude difference than previous speckle measurements. We also provide relative photometry of the components.  Speckle interferometry at the Southern Astrophysical Research (SOAR) telescope was started in 2007 with tests of a new high-resolution camera \\citep{TC08}. In the near future, this camera will work jointly with the adaptive-optics system to reach diffraction-limited resolution on faint targets in the visible. Meanwhile, it was used as a stand-alone instrument.  Observing  time for  the USNO  intensified CCD  speckle camera  at the Blanco  4-m telescope  at CTIO  was allocated  in July  2008  to cover several programs,  ranging from orbit calculation  and improvement, to observation of  nearby faint red  and white dwarfs and  subdwarfs, and also   to  the   completion  of   a   speckle  survey   of  nearby   G dwarfs. However, owing to  new export regulations, the equipment could not be sent  to CTIO in time  for the run, despite all  efforts of the NOAO  administration. The new  camera had  to be  used instead.  For a number of reasons,  it was not possible to  reach the faintest targets with  this new  camera, so  the program  had to  be modified  ``on the fly\". The  nearby, faint targets had  to be abandoned  and to preserve uniformity  of the  sample,  most of  the  nearby G  dwarfs were  also dropped.  Further,  failure   of  the  Blanco  atmospheric  dispersion corrector (ADC) restricted observing  to a smaller than desired region around the zenith. Nevertheless, a large number of useful measurements were obtained.   We present the observational technique in Sect.~2, starting with the instrument description and then detailing various data processing steps (Fig.~\\ref{fig:flow}). A new method of establishing the detection limits is described and some effects which can lead to false detections are studied. The main results are presented in Sect.~3, including comments on some objects of interest and new discoveries.  \\begin{figure*} \\epsscale{2.0} \\plotone{fig1.eps} \\caption{Schematic representation of data acquisition and processing. \\label{fig:flow}} \\end{figure*} ",
        "conclusions": ""
    },
    "0911/0911.4002_arXiv.txt": {
        "abstract": "The variation of total solar irradiance (TSI) has been measured since 1978 and that of the spectral irradiance for an even shorter amount of time. Semi-empirical models are now available that reproduce over 80\\% of the measured irradiance variations. An extension of these models into the more distant past is needed in order to serve as input to climate simulations. Here we review our most recent efforts to model solar total and spectral irradiance on time scales from days to centuries and even longer. Solar spectral irradiance has been reconstructed since 1947. Reconstruction of solar total irradiance goes back to 1610 and suggests a value of about 1--1.5~Wm$^{-2}$ for the increase in the cycle-averaged TSI since the end of the Maunder minimum, which is significantly lower than previously assumed but agrees with other modern models. First steps have also been made towards reconstructions of solar total and spectral irradiance on time scales of millennia. ",
        "introduction": "\\label{intro}  Solar irradiance is the total energy flux (or energy received per unit area and time) at the top of the Earth's atmosphere. It is strongly wavelength dependent: less than 8\\% of the solar energy is emitted at wavelengths below 400~nm, more than 60\\% come from the wavelengths between 400 and 1000~nm and roughly another 30\\% from the longer wavelengths \\citep[e.g.,][]{solanki-unruh-98,krivova-et-al-2006a}.  The small contribution of the UV part of the spectrum is, however, compensated by the spectral dependence of the transmission of the Earth's atmosphere. Thus solar radiation below 300~nm is almost completely absorbed in the atmosphere \\citep[see, e.g.,][and references therein]{haigh-2007} and is important for the chemistry of the stratosphere and overlying layers. In particular, radiation in the Ly-$\\alpha$ line (121.6~nm) and in the oxygen continuum and bands between 180 and 240~nm controls production and destruction of ozone \\citep[e.g.,][]{frederick-77,brasseur-simon-81,haigh-94,haigh-2007,% fleming-et-al-95,egorova-et-al-2004,langematz-et-al-2005b}. Solar UV radiation at 200--350~nm is the main heat source in the stratosphere and mesosphere \\citep{haigh-99,haigh-2007,rozanov-et-al-2006}.  The role of the solar UV radiation for the Earth's climate is boosted further by the fact that variations of solar irradiance are also strongly wavelength dependent. Whereas the total (integrated over all wavelengths) solar irradiance varies by only about 0.1\\% over the solar cycle, the UV irradiance varies by 1 to 3 orders of magnitude more \\citep[e.g.,][]{floyd-et-al-2003a}.  Solar near-IR radiation absorbed by water vapour and carbon dioxide is an important source of heating in the lower atmosphere \\citep{haigh-2007}. Solar variability in the IR is comparable to or lower than the TSI variations and in the range between about 1500 and 2500~nm it is reversed with respect to the solar cycle \\citep{harder-et-al-2009,krivova-et-al-2009a}.  Solid assessment of the solar forcing on the Earth's climate is still plagued, among other factors, by a shortage of reliable and sufficiently long irradiance records. The time series of direct space-borne measurements of solar irradiance is only 3 decades long (in case of spectral irradiance, even shorter). Consequently, models are required to extend the irradiance time series to earlier times. In the last years, considerable advances have been made in modelling the solar total irradiance. Also, spectral models started gathering force. Reconstructions of the past variations are quite accurate for the last decades, when sufficiently high resolution measurements of the solar photospheric magnetic flux (i.e. magnetograms) exist. On longer time scales, the secular trend remains rather uncertain and until very recently the reconstructions were limited to the telescope era, i.e. to the last 4 centuries.  \\citet{domingo-et-al-2009} have recently presented an overview of different types of irradiance models on time scales of days up to the solar cycle. Here we review the most recent progress in modelling solar total and spectral irradiance on time scales from days to centuries and even longer using the SATIRE set of models. Before discussing the model we briefly describe the available observational data (Sect.~\\ref{obs}) and main sources of irradiance variations (Sect.~\\ref{sources}). In Sect.~\\ref{short} we concentrate on the period when direct irradiance measurements are available, in Sects.~\\ref{cent} and \\ref{mill} we discuss the reconstructions for the telescope era and the first steps towards longer-term reconstructions, respectively. Section~\\ref{summary} summarises our current knowledge on irradiance variations on time scales of interest for climate research. ",
        "conclusions": "\\label{summary}  As weak as they are (about 0.1\\% over the solar cycle), solar irradiance variations affect the Earth's climate at a level of 0.2$^\\circ$K over the solar cycle \\citep{camp-tung-2007}. The role of the total solar irradiance might be significantly reinforced by the wavelength dependence of irradiance variations. The amplitude of the variations is 1 to 3 orders of magnitude higher in the UV part of the spectrum than in the visible or IR. At the same time, solar UV radiation is an important driver of chemical and physical processes in the Earth's upper atmosphere.  Reliable assessment of solar forcing on the Earth's climate is still plagued, among other factors, by a shortage of reliable and sufficiently long irradiance records. The time series of direct space-borne measurements of solar irradiance is only 3 decades long (in case of spectral irradiance, even shorter). Therefore we have recently focussed our main attention on tests of the reconstructed irradiance against the newest spectral data, improvement of the models in the UV and appraisal of the magnitude of the secular trend in irradiance.  We have reconstructed solar spectral irradiance at 115--160000~nm  back to 1947. Reconstruction of solar total irradiance goes back to 1610 and suggests a value of about 1--1.5~Wm$^{-2}$  for an increase in the cycle-averaged TSI between the end of the Maunder minimum and the last 50~years, which is significantly lower than previously accepted but agrees with other recent estimates. First steps have also been made towards reconstructions of solar total and spectral irradiance on time scales of millennia.  \\bigskip {\\bf Acknowledgments.} This work was supported by the Deutsche For\\-schungs\\-ge\\-mein\\-schaft, DFG project number SO~711/2 and by the WCU grant No.~R31-10016 funded by the Korean Ministry of Education, Science and Technology. We thank the International Space Science Institute (Bern) for hosting meetings of the international team on `Interpretation and modelling of SSI measurements'."
    },
    "0911/0911.3552_arXiv.txt": {
        "abstract": "We present projected constraints on modified gravity models from the observational technique known as 21\\,cm intensity mapping, where cosmic structure is detected without resolving individual galaxies.  The resulting map is sensitive to both BAO and weak lensing, two of the most powerful cosmological probes.  It is found that a $200\\,\\mathrm{m}\\times200\\,\\mathrm{m}$ cylindrical telescope, sensitive out to $z=2.5$, would be able to distinguish DGP from most dark energy models, and constrain the Hu \\& Sawicki \\fr{} model to $|f_{R0}|<9 \\times 10^{-6}$ at 95\\% confidence.  The latter constraint makes extensive use of the lensing spectrum in the nonlinear regime.  These results show that 21\\,cm intensity mapping is not only sensitive to modifications of the standard model's expansion history, but also to structure growth. This makes intensity mapping a powerful and economical technique, achievable on much shorter time scales than optical experiments that would probe the same era. ",
        "introduction": "\\label{s:intro}  One of the greatest open questions in cosmology is the cause of the observed late time acceleration of the universe.   Within the context of normal gravity described by Einstein's General Relativity, this phenomena can only be explained by an exotic form of matter with negative pressure.  Another possible explanation is that on cosmological scales, General Relativity fails and must be replaced by some theory of modified gravity.  Several approaches have been proposed to modify gravity at late times to explain the apparent acceleration of the universe.  The challenge in these modifications is to preserve successful predictions of the CMB at $z\\approx1000$, and also the precision tests at the present epoch in the solar system.  A generic class of theories operates with the Chameleon effect, where at sufficiently high densities General Relativity (GR) is restored, thus applying both in the solar system and the early universe.  To further understand the nature of gravity would require probing gravity on cosmological scales.  Large scales means large volume, requiring large fractions of the sky.  Gravity can be probed by gravitational lensing, which measures geodesics and thus the gravitational curvature of space, and is a sensitive probe of the growth of structure in the Universe \\cite{KnoxSongTyson,JainZhang,Tsujikawa2008,Schmidt:2008hc}.  In working out predictions for cosmology, the theoretical challenge posed by these theories are the nonlinear mechanisms in each model, necessary in order to restore Einstein Gravity locally to satisfy Solar System constraints.  We present quantitative results from nonlinear calculations for a specific \\fr{} model, and forecasted constraints for future 21\\,cm experiments.  An upcoming class of experiments propose the observation of the 21\\,cm spectral line at low resolution over a large fraction of the sky and large range of redshifts \\cite{Peterson:2009ka}.  Large scale structure is detected in three dimensions without the detection of individual galaxies.  This process is referred to as \\emph{21\\,cm intensity mapping}.  These experiments are sensitive to structures at a redshift range that is observationally difficult to observe for ground-based optical experiments due to a lack of spectral lines.  Yet these experiments are extremely economical \\cite{Seo:2009fq} since they only require limited resolution and no moving parts.  Intensity mapping is sensitive to both the Baryon Acoustic Oscillations (BAO) and to weak lensing, two of the most powerful observational methods to determine cosmological parameters. It has been shown that BAO detections from 21\\,cm intensity mapping are powerful probes of dark energy, comparing favourably with Dark Energy Task Force Stage IV projects within the figure of merit framework \\cite{Chang:2007xk,Albrecht:2006um}.  In this paper we present projected constraints on modified gravity models from 21\\,cm intensity mapping.  In Section \\ref{s:mgm} we describe the modified gravity models considered.  In Section \\ref{s:obs} we discuss the observational signatures accessible to 21\\,cm intensity mapping, and calculate the effects of modified gravity on these signatures.  In Section \\ref{s:results} we present statistical analysis and results and we conclude in Section \\ref{s:discuss}.  We assume a fiducial \\LCDM{} cosmology with WMAP5 cosmological parameters: $\\Omega_m = 0.258$, $\\Omega_b = 0.0441$, $\\Omega_\\Lambda = 0.742$, $h = 0.719$, $n_s=0.963$ and $\\log_{10}A_s = -8.65$.  We will follow the convention that $\\omega_x \\equiv h^2\\Omega_x$. ",
        "conclusions": "\\label{s:discuss}  We have shown that the first generation of 21\\,cm intensity mapping instruments will be capable of constraining the HS \\fr{} model (with $n=1$) down to a field value of $|f_{R0}| \\lesssim 2 \\times 10^{-5}$ at 95\\% confidence (Figure~\\ref{f:results:FRlowres}). This is an order of magnitude tighter than constraints currently available from galaxy cluster abundance \\cite{fRcluster}.  Furthermore, model parameters in this regime are not ruled out by Solar System tests.  In comparing Figures \\ref{f:results:FRfisher} and \\ref{f:results:FRlowres} it is clear that a more advanced experiment, with resolution improved by a factor of two, would further half the allowed value of $|f_{R0}|$.  It should be noted however, that halving of the allowed parameter space does not correspond to a factor of four increase in information.  Deviations in the lensing spectrum scale sub-linearly in the \\fr{} parameters, enhancing the narrowing of constraints as information is added (see Section \\ref{s:results}).  While we have concentrated on a particular $f(R)$ model, many viable functional forms for $f(R)$ have been proposed in the literature \\cite{Nojiri:2007as,Starobinsky:2007hu,Appleby:2007vb}.  The predictions for the growth of structure in these different models agree qualitatively: the gravitational force is enhanced by 4/3 within $\\lambda_C$, enhancing the growth on small scales.  However, there are quantitative differences in the model predictions due to the different evolution of $\\lambda_C$ over cosmic time.  Our results for the HS model with different values of $n$ should thus cover a range of different functional forms for \\fr{}.  Table~\\ref{t:results:fr} shows that our constraints do not depend very sensitively on the value of $n$. This is because the weak lensing measurements cover a wide range of scales as well as redshifts.  Furthermore, it is straightforward to map the enhancement in the linear $P(k,z)$ at given $k$ and $z$ from the HS model considered here to any other given model, to obtain approximate constraints for that model.  Future cluster constraints will almost certainly improve on the current limits of $|f_{R0}| \\lesssim\\mbox{few}\\: 10^{-4}$ \\cite{fRcluster}.  However, for smaller field values, the main effect of $f(R)$ gravity shifts to lower mass halos, since the highest mass halos are chameleon-screened (see Fig.~2 in \\cite{HPMhalopaper}).  Hence, future cluster constraints will depend on the ability to accurately measure the halo abundance at masses around few $10^{14} M_{\\odot}$ and less. Furthermore, the constraints from cluster abundances depend sensitively on the knowledge of the cluster mass scale, and are already systematics-dominated \\cite{fRcluster}.  Weak lensing constraints have a completely independent set of observational systematics, and are in principle less sensitive to baryonic or astrophysical effects.  Thus, the forecasted constraints on modified gravity presented here are quite complementary to constraints from cluster abundances.  The processes that produce the BAO feature in the matter power spectrum are understood from first principles. In addition the BAO length scale can be extracted even in the presence of large uncertainties in biases and mass calibrations.  Likewise, weak lensing on large scales is well understood, with baryonic physics being much less important than on smaller scales \\cite{Zhan:2004wq}. In addition the dominant systematics present in optical weak lensing surveys are instrumental in nature and not intrinsic to the quantities being measured. While 21\\,cm intensity mapping is as yet untested, instrumental systematics will be very different from those that affect the optical.  In the case of this study, and more generally for cosmological models which substantially modify structure formation, the motivation for higher resolution comes not from improved BAO measurements but from better weak lensing reconstruction. Higher resolution not only makes weak lensing information available at higher multi-poles, but improves the accuracy at which lensing can be reconstructed on all scales.  The inclusion of lensing information in the nonlinear regime was crucial, and largely responsible for the competitiveness of these forecasts.  As seen in Figure \\ref{f:obs:lensing}, much of the constraints come from multi-poles in the nonlinear regime.  It should be noted that for the higher resolution experiment considered, the minimum scale is limited not by the resolution at high redshift, but by the saturation of information in nonlinear source structures at low redshift \\cite{Lu:2009je}.  Our constraints from lensing are conservative since only one wide source redshift bin was considered, limited to $\\l<600$ as described above.  To maximize information, the  source redshift range could be split into multiple bins, properly considering the correlation in the lensing signal between them; a process known as lensing tomography. The low redshift bin would be limited as above, and the high redshift bin would be limited by the resolution to $\\l \\approx 850$ at $z =  2.5$. However in intermediate bins, the lensing signal could be reliably reconstructed above $\\l\\approx1000$.  Unlike most smooth dark energy models, such as quintessence, constraints on the models considered here are chiefly sensitive to structure formation, as is clear from Figure \\ref{f:results:FRfisher}. No experiment in the foreseeable future will be able to improve upon the constraints we project on these models using exclusively expansion history probes (except by breaking degeneracies). These forecasts show that 21\\,cm intensity mapping is not only sensitive to a cosmology's expansion history through the BAO, but also to structure growth through weak lensing.  The weak lensing measurements cannot compete with far off space based surveys like Euclid or JDEM, which will have galaxy densities of order $100\\,\\mathrm{arcmin}^{-2}$\\cite{Albrecht:2006um} and resolution to far greater $\\l$. However, cylindrical 21\\,cm experiments are realizable on a much shorter time scale and at a fraction of the cost. In addition, the measurements considered here are approaching the limit at which \\fr\\ models can be tested.  For $|f_{R0}|$ much less than $10^{-6}$ the chameleon mechanism becomes important before there are observable modifications to structure growth, reducing the motivation to further study these models.  It has also been shown that, for these experiments, a BAO measurement consistent with \\LCDM{} would definitively rule out DGP without a cosmological constant as a cosmological model.  Even in the case that a BAO measurement consistent with DGP is made, the model is still distinguishable from an exotic smooth dark energy model through structure growth.  The former result is not surprising given that DGP is now in conflict with current data \\cite{Fang:2008kc}.  However it is illustrative that a single experiment can precisely probe both structure formation and expansion history.  Even a dark energy model that conspires to mimic DGP is, to a large extent, distinguishable.  We have studied the effects of modified gravity theories on observational quantities for future 21\\,cm surveys.  Because these surveys measure the distribution of galaxies on large angular scales over large parts of the sky, they are well suited to measure the expected deviations relative to standard general relativity.  We have computed the predictions of modified gravity in the linear and nonlinear regimes, and compared to the sensitivity of future surveys.  We find that a large part of parameter space can be tested."
    },
    "0911/0911.1557_arXiv.txt": {
        "abstract": "Many young extra-galactic clusters have a measured velocity dispersion that is too high for the mass derived from their age and total luminosity, which has led to the suggestion that they are not in virial equilibrium. Most of these clusters are confined to a narrow age range centred around 10\\,Myr because of observational constraints. At this age the cluster light is dominated by luminous evolved stars, such as red supergiants, with initial masses of $\\sim$13$-22\\,\\msun$ for which (primordial) binarity is high.  In this study we investigate to what extent the observed excess velocity dispersion is the result of the orbital motions of binaries.  We demonstrate that estimates for the dynamical mass of young star clusters, derived from the observed velocity dispersion, exceed the photometric mass by up-to a factor of 10  and are consistent with a constant offset in the square of the velocity dispersion.  This can be reproduced by models of virialised star clusters hosting a massive star population of which $\\sim$25\\% is in binaries, with typical mass ratios of $\\sim$0.6 and periods of $\\sim$1000 days. We conclude that binaries play a pivotal role in deriving the dynamical masses of young ($\\sim$10\\,Myr) moderately massive and compact ($\\lesssim10^5\\,\\msun; \\gtrsim 1\\,$pc) star clusters. ",
        "introduction": "Young massive clusters have received considerable attention in the last decade because they trace star formation (e.g. \\citealt{1995AJ....109..960W, 1997AJ....114.2381M, 1999AJ....118..752Z}). Advances in observations enabled us to resolve such star clusters up to $\\sim 20$\\,Mpc, allowing determination of their fundamental parameters, such as mass and radius (e.g. \\citealt{2004A&A...416..537L}).  The mass of a resolved star cluster can be determined in two ways: one of them by converting the observed luminosity, age and distance directly to mass via the age dependent mass-to-light ratio ($M/L$) taken from a single stellar population (SSP) model.  We refer to the resulting mass as the photometric mass $\\mphot$. This method requires an estimate of the cluster age, which again requires estimates for the metallicity and the stellar initial mass function (IMF).  An independent mass estimate is based on the virial theorem and this mass is generally referred to as the dynamical mass \\citep{1987degc.book.....S}:  \\begin{equation} \\mdyn=\\frac{\\eta\\sigdyn^2\\reff}{G}. \\label{eq:mdyn} \\end{equation} Here $G$ is the gravitational constant, $\\sigdyn$ is the line of sight velocity dispersion in the cluster, $\\reff$ is the effective (half-light) radius\\footnote{Here we assume that the half-light radius is the same as the half-mass radius, which is not the case when the cluster is mass segregated \\citep[][]{2006MNRAS.369.1392F,2008MNRAS.391..190G}.} and $\\eta\\simeq9.75$ is a constant that depends slightly on the density profile.  Equation~(\\ref{eq:mdyn}) is valid for a cluster in virial equilibrium consisting of single stars. Since in this study we consider possible difference between $\\sigdyn$ and the observed velocity dispersion, $\\sigobs$, we will refer to the empirically derived dynamical mass, i.e. based on $\\sigobs$, as $\\mdynobs$.  A comparison between $\\mphot$ and $\\mdynobs$ serves as a check for the range of assumptions on which both mass estimates are based.  An inconsistency between $\\mphot$ and $\\mdynobs$ can be attributed to variations in the IMF, on which $\\mphot$ is in part based, or to a lack of virial equilibrium, on which $\\mdynobs$ is based.  For many young ($\\sim$10\\,Myr) star clusters $\\mdynobs>\\mphot$, with $\\mdynobs$ up to $\\sim$10 times larger than $\\mphot$ \\citep[e.g.][hereafter B06]{2006A&A...448..881B}, suggesting that these objects are super-virial. For older clusters ($ \\gtrsim100\\,$Myr) there is good agreement between $\\mdynobs$ and $\\mphot$ \\citep[e.g.][B06]{2004AJ....128.2295L}.  The alleged super-virial state of some young clusters has been attributed to the impulsive expulsion of residual gas from the parent molecular cloud in which the star cluster formed \\citep[e.g.][]{2006MNRAS.373..752G}.  Such early outgassing, driven by stellar winds of massive stars or supernovae, causes the stellar velocities to be high compared to the binding energy of the stars. This argument has been used to motivate infant mortality of young star clusters \\citep{2003ARA&A..41...57L}.  However, the gas expulsion theory has difficulties in explaining the super-virial velocity for the 10\\,Myr old clusters presented in B06. The arguments are as follows: The time needed to completely dissolve, or to find a new virial equilibrium after impulsive gas expulsion is about 20 crossing times, $\\tcr$, where $\\tcr\\propto\\rhoh^{-1/2}$ and $\\rhoh$ is the density within the half-mass radius \\citep[see for example Fig.~8 in][]{2007MNRAS.380.1589B}.  Hence, to be able to `catch' an unbound or expanding cluster at 10 Myr, $\\tcr$ should be $\\gtrsim1\\,$Myr. This corresponds to a half-mass density of stars and gas of $\\rhoh\\lesssim300\\,\\msunpc$. Clusters with shorter $\\tcr$ (higher density) have expanded into the field, or found a new equilibrium, a few Myrs after gas expulsion and are not observable as super-virial clusters at 10\\,Myr\\footnote{The models of \\citet{2006MNRAS.373..752G} start with a density of $\\sim60\\,\\msunpc$ ($\\tcr\\approx2.5\\,$Myr) in the embedded phase, and this is why they find that the effects of gas expulsion are observable for 25 Myr.}.  The density in the embedded phase of the clusters under discussion is unknown, but can be roughly estimated using their current densities. The present day densities are $\\rhoh\\approx10^{3\\pm1}\\,\\msunpc$ (Table~\\ref{tab:data}). The densities in the embedded phase were at least a factor $1/\\epsilon^4$ higher, where $\\epsilon$ is the star formation efficiency.  This because the mass of the embedded cluster has reduced by a factor $\\epsilon$ and the cluster has expanded at least by a factor of $1/\\epsilon$ as a response to it, contributing a factor $1/\\epsilon^3$ to the reduction of $\\rhoh$. The $1/\\epsilon$ expansion holds for adiabatic mass loss, for impulsive mass loss and $\\epsilon\\lesssim0.9$ the cluster expands much more \\citep{1980ApJ...235..986H}. So for the clusters at 10\\,Myr the estimated densities in the embedded phase are much too high to still have features of gas expulsion detectable in their velocity dispersion at 10\\,Myr. These arguments suggest that deviations from virial equilibrium are not a plausible explanation and an alternative explanation for the high $\\sigobs$ values is needed.  The existence of binary stars is generally ignored in the estimates for $\\mdynobs$, even though their internal velocities can lead to an over estimation of $\\mdyn$ (Kouwenhoven \\& de Grijs 2008, hereafter K08)\\nocite{2008A&A...480..103K}.  K08 studied this phenomenon in virialised star clusters with a 100\\%\\ binary fraction and a range of $\\sigdyn$. They subsequently derive $\\mdynobs$ by `measuring' $\\sigobs$ and applying equation~(\\ref{eq:mdyn}).  They found that the presence of binaries can lead to an overestimation of $\\mdyn$ by a factor of $\\sim$2 for clusters with $\\sigdyn\\simeq1\\,\\kms$.  For clusters with $\\sigdyn\\simeq10\\,\\kms$ they found only a 5\\% increase in $\\mdynobs$ due to binaries. They therefore concluded that binaries are not important for massive/dense clusters.  \\citet[][hereafter M08]{2008A&A...489.1091M} found $\\mdynobs/\\mphot \\simeq 10$ for some of the star clusters in the Antennae galaxies (NGC~4038/4039) and NGC~1487 and since these clusters have velocity dispersions of $10-20\\,\\kms$ they subsequently concluded that binaries are not important and that these star clusters are super-virial and dissolving quickly.  Here we revisit the effect of binaries on $\\mdynobs/\\mphot$ and we focus on $\\sim$10\\,Myr old star clusters.  This is motivated by our desire to incorporate the effect of the steep increase of the stellar luminosity with increasing stellar mass, which is in particular important for young clusters, an aspect not considered by K08.  Of the two approaches presented by K08, one focused on solar-type stars, the other uses a Kroupa IMF for the primary stars. They give equal weight to each binary in their computed velocity dispersion.  In addition K08 do not consider stars more massive than $20\\,\\msun$. This approach may be appropriate for studies of intermediate age and old open star clusters, but it is less suitable for young star clusters.  At an age of $\\sim$10\\,Myr the cluster light is dominated by the most massive ($\\gtrsim 15\\,\\msun$) stars for which binarity is high and ignoring them can lead to misinterpretations of observations of various astrophysical processes \\citep[e.g.][]{1998A&ARv...9...63V}. Massive binaries have a larger effect on $\\sigobs$ than low-mass binaries due to their higher orbital velocities, but also due to the more common short-periods and comparable masses \\citep[e.g.][]{1991A&A...248..485D, 2002ApJ...574..762P, 2008MNRAS.386..447S, SGE2009,2009AJ....137.3358M}.  Incorporating the massive stars in our calculation has two important effects, both of which amplify the effect of binarity on $\\sigobs$ with respect to the results of K08: massive stars dominate the cluster light and their higher masses and (intrinsically) different binary properties give rise to a larger $\\sigobs$.  In this paper we quantify the effect of the presence of (massive) binaries on $\\mdynobs/\\mphot$ and we use this ratio is a proxy of the excess dispersion. In Section~\\ref{sec:model} we discuss the properties of the binary population that is expected in young ($\\sim$10\\,Myr) clusters and we present a simple model for the additional velocity dispersion due to such binaries.  In Section~\\ref{sec:data} we summarize existing observational results to confront our model with. Our conclusions are discussed in Section~\\ref{sec:conc}. All the specific acronyms used in this study and their definitions are given in Table~\\ref{tab:acronyms}. ",
        "conclusions": "\\label{sec:conc} Several studies have found from spectroscopic analyses that for many young ($\\sim$10\\,Myr) star clusters the measured velocity dispersion is too high for the mass derived from their total luminosities and their ages. This has led several authors to conclude that these clusters are super-virial and thus dissolving. However, the conversion from velocity dispersion to mass (equation~\\ref{eq:mdyn}) does not consider the additional velocities of binaries. K08 considered this effect, but concluded that binaries are only important for clusters with low intrinsic velocity dispersion ($\\sim$1\\,$\\kms$), i.e. lower than the aforementioned clusters. K08 ignored the mass dependent mass-to-light ratio of stars and the intrinsically different binary properties of massive stars. In this study we show that taking these aspects into account makes the contribution of binarity to the dynamical mass estimates, $\\mdynobs$, of clusters in this age range non-negligable.  We present a simple analytical model that gives the 1-dimensional velocity dispersion of a virialised star cluster hosting a binary population. The model is complementary to the classical virial relation for clusters consisting of single stars (equation~\\ref{eq:mdyn}). The result is presented as a single equation that needs as input the (typical) binary fraction, mass ratio, primary mass and orbital period of the binary population and the mass and radius of the star cluster. This relation can be used to easily estimate the effect of binaries based on different parameters for the binary population and/or cluster.  The model presented here serves as a starting point for more realistic approaches using binary population synthesis models \\citep[e.g][]{2009arXiv0908.1386E}. Tentative confirmation of our results comes from the velocity dispersion of the binary population discussed in E08: $\\sim$12\\,$\\kms$ at an age of 10\\,Myr (Eldridge 2009, priv. comm.), which is close to what we find for the reference values discussed in Section~\\ref{sec:model} (see Fig.~\\ref{fig:sig}).  For 24 clusters we derive the ratio of $\\mdynobs$ over the photometric mass, $\\mphot$, and show that it decreases with increasing cluster velocity dispersion. This is also what follows from the model and most of the empirically determined $\\mdynobs/\\mphot$ ratios can be explained by binaries using a conservative binary fraction of 25\\%, a mass ratio of $0.6$ and an orbital period of a 1000 days. When allowing a spread in the binary parameters, almost all clusters are within 2-standard deviation of the model results.  The fact that $\\mdynobs$ and $\\mphot$ generally agree for older ($\\gtrsim100\\,$Myr) clusters is consistent with this binary scenario. In older clusters, we indeed expect a lower velocity contribution of binaries. The primary star will be of a later spectral type, thus $m_1$ is lower.  At 100\\,Myr the most luminous stars are roughly $5\\,\\msun$. Equation~(\\ref{eq:sigbinsq}) shows that when $m_1$ is a factor of 3 lower, $\\sigbin^2$ is a factor of $\\sim$2 lower, keeping all other parameters fixed. Also, typical periods are longer.  \\citet[][]{1991A&A...248..485D} find that the median period of solar type stars is 180\\,yr.  From equation~(\\ref{eq:sigbinsq}) we can see that the effect of such binaries on $\\sigbin^2$ is about a factor of $\\sim$15 less than the (early type) binaries considered here.  As mentioned in Section~\\ref{sec:data}, the estimated $\\mphot$ following from a comparison with SSP models, or from an IMF extrapolation from the number of RSGs as is done in resolved clusters, is also affected by binarity \\citep{2009ApJ...696.2014D} The fraction of stars that is removed from our sample due to this effect is $(1-g)f$, corresponding to 45\\%, giving rise to $\\mdynobs/\\mphot\\approx2$ for the values of Table~\\ref{tab:param}.} There is no reason, however, to expect that this would preferentially affect clusters with low ratios $\\mphot/\\reff$ and it can thus not cause the downward trend seen in Fig.~\\ref{fig:data}.  The values of the binary parameters used in the study (Table~\\ref{tab:param}) are only indirectly based on observations since we have to correct the period distribution found for massive main-sequence stars to account for the reduced RSG phase of stars in tight binaries (section~\\ref{ssec:obs_bin}).  Our assumption can be verified once the binary fraction $\\frsg$ and the associated period distribution among a statistically significant sample of resolved SGs has been determined.  This could be done spectroscopically using a long time base ($\\sim$few~$100-1000\\,\\dr$).  The recently discovered RSG clusters towards the Galactic centre \\citep{2006ApJ...643.1166F, 2007ApJ...671..781D, 2009A&A...498..109C} provide an excellent opportunity to do this. All three have approximately the same age as the extra-galactic clusters used here and their masses are relatively low (few times $10^4\\,\\msun$) and have radii of a few pc, which according to our model places them in the regime where binaries dominate the measured velocity dispersion.  The ratio $\\mdynobs/\\mphot$ was determined for two of them and is $\\sim2$ \\citep{2008ApJ...676.1016D}, lower than the extra-galactic clusters with comparable $\\mphot/\\reff$ (Fig.~\\ref{fig:data}), but still consistent with the lower 2$\\sigma$ line of our prediction. This result is very sensitive to low number statistics since the number of RSG in these clusters is $\\sim$20, so for $\\frsg=0.25$ we expect only a handful of binaries.  \\citet{2009arXiv0909.3815R} present a spectroscopic multi-epoch survey of luminous evolved stars in Westerlund~1. This cluster is slightly younger than the clusters considered here, thus its supergiants population is formed by more massive stars. They find a binary fraction in excess of $40\\%$ among the 20 most luminous supergiants. Interestingly, they also find radial velocity changes of $\\sim$$15-25\\,\\kms$ in cool hypergiants due to photospheric pulsations.  Macro turbulence dispersions of $5-10\\,\\kms$ are also found for luminosity class II and III giants by \\citet{1986ApJ...310..277G} and \\citet{2008AJ....135..892C}. This is  an additional complication in dynamical mass determinations of young star clusters containing massive giants.  Our results are an important ingredient in the discussion on the importance of the early mass independent disruption, or `infant mortality', of star clusters. The high velocity dispersions found for the clusters discussed here have been put forward as empirical evidence that many young ($\\lesssim30\\,$Myr) clusters are quickly dissolving \\citep[e.g.][M08]{2006MNRAS.373..752G}. We have provided arguments that the alleged super-virial state can largely be explained by orbital motions of binary stars.  Early dissolution due to gas expulsion can still exist, but it probably occurs on much shorter time-scales ($<<10\\,$Myr) than generally assumed.  This idea is supported by the fact that the clusters considered here have densities of $\\sim$$10^3\\,\\msunpc$, corresponding to an internal crossing times of the stars of roughly $0.5\\,$Myr. So these clusters have evolved for at least 20 crossing times. The crossing time in the embedded phase is much shorter than the crossing time at 10\\,Myr due to the nonzero star formation efficiency and the consequent expansion \\citep{2008MNRAS.389..223B}. The gas expulsion models show that clusters need about $20$ initial crossing times to find a new virial equilibrium, or completely dissolve into the field \\citep[e.g.][]{1997MNRAS.286..669G, 2001MNRAS.323..988G, 2007MNRAS.380.1589B}.  So at 10\\,Myr the super-virial state is undetectable and the clusters discussed here are therefore survivors of the gas expulsion, or `infant mortality', phase."
    },
    "0911/0911.2641_arXiv.txt": {
        "abstract": "The relatively recent insight that energy input from supermassive black holes (BHs) can have a substantial effect on the star formation rates (SFRs) of galaxies motivates us to examine the effects of BH feedback on the scale of galaxy groups.  At present, groups contain most of the galaxies and a significant fraction of the overall baryon content of the universe and, along with massive clusters, they represent the only systems for which it is possible to measure both the stellar and gaseous baryonic components directly.  To explore the effects of BH feedback on groups, we analyse two high resolution cosmological hydrodynamic simulations from the OverWhelmingly Large Simulations (OWLS) project.  While both include galactic winds driven by supernovae, only one of the models includes feedback from accreting BHs.  We compare the properties of the simulated galaxy groups to a wide range of observational data, including the entropy and temperature profiles of the intragroup medium, hot gas mass fractions, the luminosity$-$temperature and mass$-$temperature scaling relations, the K-band luminosity of the group and its central brightest galaxy (CBG), star formation rates and ages of the CBG, and gas- and stellar-phase metallicities.  Both runs yield entropy distributions similar to the data, while the run without AGN feedback yields highly peaked temperature profiles, in discord with the observations.  Energy input from supermassive BHs significantly reduces the gas mass fractions of galaxy groups with masses less than a few times $10^{14}$ M$_\\odot$, yielding a gas mass fraction and X-ray luminosity scaling with system temperature that is in excellent agreement with the data, although the detailed scatter in the $L-T$ relation is not quite correct.  The run without AGN feedback suffers from the well known overcooling problem --- the resulting stellar mass fractions are several times larger than observed and present-day cooling flows operate uninhibitedly.  By contrast, the run that includes BH feedback yields stellar mass fractions, SFRs, and stellar age distributions in excellent agreement with current estimates, thus resolving the long standing `cooling crisis' of simulations on the scale of groups.  Both runs yield very similar gas-phase metal abundance profiles that match X-ray measurements, but they predict very different stellar metallicities.  Based on the above, galaxy groups provide a compelling case that feedback from supermassive BHs is a crucial ingredient in the formation of massive galaxies. ",
        "introduction": "Galaxy formation is one of the key unsolved problems in modern astrophysics.  Traditionally, studies of galaxy formation have focused on the observable stellar properties of galaxies.  However, a large fraction of the baryonic mass of massive galaxies is believed to be in a hot diffuse form and many of the processes that regulate the formation of stars do so by influencing the properties of this hot gas.  Thus, a complete view of galaxy formation necessarily incorporates both the stars and the hot gas and an understanding of the processes by which these phases interact.  At present, galaxy groups and clusters represent the only systems in the universe for which it is possible to measure the properties of both the stellar and gaseous baryonic components out to a significant fraction of the halo viral radius.  This is possible because the gaseous component is compressed and heated to very high temperatures ($T \\sim 10^{6-8}$ K) where this material (the intragroup or intracluster media) becomes visible at X-ray wavelengths.  While clusters have higher X-ray surface brightnesses than groups, and the derived physical properties of the gas are therefore presumably more robust, clusters represent very rare and extreme environments.  Since the efficiencies of many of the environmental processes that are thought to alter galaxy formation/evolution (e.g., ram pressure stripping, strangulation, tidal stripping, galaxy harassment) increase rapidly with the mass of the main halo, one expects that galaxies in clusters have been subjected to far more violent interactions than those which reside in the lower density and more typical environment of groups.  It is therefore possible that a picture of galaxy formation derived in large part from observations of clusters is misleading.  Galaxy groups, on the other hand, contain a large fraction of present day galaxies and a significant fraction of the overall universal baryon budget (e.g., Mulchaey 2000).  They are therefore more representative hosts than clusters.  In addition, the effects of feedback from supernovae (SNe) and supermassive black holes (BHs) on the hot gas, which is thought to be crucial in shaping the properties of galaxies, are expected to be much more obvious (and therefore perhaps easier to quantify) on the scale of groups, where the energy input associated with these sources is comparable to the binding energies of these systems.  And while what could be said about the detailed properties of the hot gas was severely limited in the past (with data from the previous generation of X-ray satellites), new high-quality data from the {\\it Chandra} and {\\it XMM-Newton} X-ray telescopes has yielded a much more detailed picture of the thermodynamic state of the gas in groups.  For these reasons, galaxy groups probably represent the best objects for studying gas and stars simultaneously.  Much progress has also been made in the theoretical modeling of the formation and evolution of groups and clusters.  For example, the past two decades have seen the development of sophisticated cosmological hydrodynamic simulation codes capable of simulating the properties of the hot gas in groups and clusters {\\it ab initio} with millions to hundreds of millions of particles/cells (e.g., Evrard 1990; Thomas \\& Couchman 1992; Navarro et al.\\ 1995; Bryan \\& Norman 1998; Kravtsov et al.\\ 2000; Borgani et al.\\ 2001; Springel et al.\\ 2001; Kay et al.\\ 2002).  Broadly speaking, these simulations have had varying degrees of success in reproducing the thermodynamic properties of the hot gas but, in general, those simulations that included the effects of radiative cooling appear to have formed far too many stars relative to what is observed in local groups and clusters (see Balogh et al.\\ 2001 for further discussion of this `cooling crisis').  This may be because, until very recently, such simulations did not include feedback from supermassive BHs.  In recent years, a growing number of authors have argued, based on the results semi-analytic models of galaxy formation as well as hydrodynamic simulations (usually idealised, as opposed to cosmological), that feedback from supermassive BHs plays a crucial role in regulating the star formation rates of massive galaxies (e.g., Springel et al.\\ 2005, hereafter SDH05; Bower et al.\\ 2006, 2008; Croton et al.\\ 2006; Hopkins et al.\\ 2006; Somerville et al.\\ 2008) and suppressing the onset of catastrophic cooling of the hot gas in groups and clusters (e.g,. Churazov et al.\\ 2001; Brighenti \\& Mathews 2003; Dalla Vecchia et al.\\ 2004; Ruszkowski et al.\\ 2004; Omma et al.\\ 2004).  The insight that AGN feedback is important in the formation and evolution of massive systems motivated SDH05 (see also Thacker et al.\\ 2006) to develop a novel scheme for following the growth of supermassive BHs and feedback from AGN self-consistently in cosmological smoothed particle hydrodynamics (SPH) simulations.  Using the model of SDH05 as a backbone, Sijacki \\& Springel (2006), Sijacki et al.\\ (2007), Puchwein et al.\\ (2008), Bhattacharya et al.\\ (2009), and Fabjan et al.\\ (2009) examined the effects of AGN feedback\\footnote{Note that Sijacki et al.\\ (2007), Puchwein et al.\\ (2008), and Fabjan et al.\\ (2009) modified the original implementation of SDH05 to include both `quasar' and `radio' feedback modes.  The former is characterised by isotropic thermal heating of the gas when accretion rates are high, while the latter, which is imposed when accretion rates are low, is characterised by thermal heating of bi-polar spherical regions offset from the galaxy hosting the BH, which are meant to mimic `bubbles' observed in many nearby clusters.} on groups and clusters.  These authors indeed find that including self-consistent AGN feedback yields reduced stellar mass fractions and star formation rates while reproducing a number of key hot gas observable properties as well.  For example, Puchwein et al.\\ (2008) found that including energy input from supermassive BHs enables their simulations to reproduce the luminosity-temperature and gas mass fraction-temperature relations of groups and clusters.  Fabjan et al.\\ (2009) found similarly good matches to these relations, while demonstrating also that the simulations yield metallicity and temperature profiles in reasonable agreement with observations of galaxy groups (although the predicted stellar fractions still appear to be somewhat too high and the entropy-temperature scaling is shallower than observed).  In the present study, we examine the effects of AGN feedback on galaxy groups using a novel set of simulations.  In particular, we extract two runs from the OverWhelmingly Large Simulations ({\\small OWLS}, for short; Schaye et al.\\ 2009) and attempt to ascertain whether or not feedback from supermassive BHs is a necessary ingredient for explaining the properties of observed galaxy groups.  The simulations we analyse have very different implementations for radiative cooling, star formation, chemodynamics, and feedback from supernova and AGN from those of SDH05 and the studies mentioned above which are based upon that model.  This means a comparison of the findings can be used to check the robustness of the conclusions to differences in the parameterisations of various subgrid physics as well as inclusion of new physics.  In addition, because the \\owls simulations are performed in large periodic boxes (as opposed to being re-simulations using `zoomed initial conditions') they yield much larger samples of groups for analysis than used previously for simulations with AGN feedback, albeit at lower mass resolution (e.g., each \\owls $100$ ${\\rm h}^{-1}$ Mpc box run yields $\\approx 200$ groups with masses greater than $10^{13}$ M$_\\odot$, whereas the studies based upon re-simulations typically only analyse a handful of systems).  This is potentially important, as observed groups and clusters show a large degree of system-to-system scatter at fixed mass in the properties of the hot gas (e.g., Fabian et al.\\ 1994; Markevitch 1998; McCarthy et al.\\ 2004).  We compare the \\owls simulations to a much wider range of X-ray and optical observations of galaxy groups than previously considered.  This allows for a more thorough assessment of the successes/failures of the models.  Finally, the large suite of \\owls runs also allows us to explore the impact of additional sub-grid physics, not just for impact of AGN feedback.  For example, the effects of changing the supernova feedback prescription, switching off metal-line, and changing the stellar initial mass function.  These variations will be explored in a forthcoming paper.  The present paper is organised as follows.  In Section 2 we provide a brief description of the \\owls runs examined in this study and how we select and analyse the galaxy groups that form in these simulations.  In Section 3 we make detailed like-with-like comparisons with a wide range of observational data of both the hot gas and stellar baryonic components.  In Section 4 we summarise and discuss our findings.  We adopt a flat $\\Lambda$CDM cosmology with $h=0.73$, $\\Omega_b=0.0418$, $\\Omega_m=0.238$, and $\\sigma_8=0.74$, which are the maximum-likelihood values from fitting to the 3-year {\\it WMAP} CMB temperature anisotropy data (Spergel et al.\\ 2007).  When comparing observed and simulated metal abundances, we consistently adopt the solar abundances of Asplund et al.\\ (2005).  Finally, we note this is the first in a series of papers investigating galaxy groups in the \\owls suite of simulations.  In Paper II we examine in detail the thermal history of the intragroup medium and present a simple physical model for the excess entropy seen at intermediate/large radii in groups.  In Paper III we explore the effect of varying the parameters of the key subgrid physics models (e.g., stellar IMF, chemodynamics, supernova and BH feedback) on the properties of galaxy groups.  In Paper IV we investigate the build up of metals in the intragroup medium and the origin of the large scale gradients. ",
        "conclusions": "We have analysed two high resolution cosmological hydrodynamic simulations (the {\\it REF} and {\\it AGN} runs) from the \\owls project.  While both include galactic winds driven by supernovae, only one of the models includes feedback from accreting BHs (the {\\it AGN} run).  These runs are representative of the \\owls runs with supernovae feedback only and runs with AGN feedback, respectively, in terms of the galaxy groups properties they produce (Paper III).  We compared the runs to a wide range of observations in order to ascertain whether or not including BH feedback in cosmological simulations yields a more realistic population of galaxy groups.  This comparison has yielded the following findings:  \\begin{itemize} \\item{The two models yield similar median entropy distributions for the gas and reproduce the `excess entropy' of observed groups with respect to the self-similar scaling. However, there is a much larger degree of scatter in the central entropies of groups in the run that includes BH feedback. This, in turn, gives rise to a larger scatter in the luminosity-temperature relation of this model.} \\item{The two models produce similar temperature profiles beyond $0.1 r_{500}$, but within this radius the temperature of groups in the {\\it REF} model continue to rise, in discord with the observations.  This is due to the excessive accumulation of cold baryons in the centres of groups in the absence of AGN feedback, which significantly deepens the central potential well leading to compressional heating of the gas.} \\item{Supernovae-driven winds alone do not inject sufficient energy to reproduce the observed low gas mass fractions of galaxy groups.  The additional energy input from BHs is, however, sufficient to reduce the gas mass fractions of galaxy groups to the observed level and also reproduces the steep increase in the gas fraction with system temperature.} \\item{Both models match the observed scaling between mass and cool-corrected temperature relatively well.  This is because beyond the central regions the gravitational potential wells are similar in both models, owing to the dominance of dark matter at large radii.} \\item{The reduced gas density in the groups that have been subjected to BH feedback results in lower X-ray luminosities at fixed halo mass/mean temperature.  The resulting luminosity - cool core-corrected temperature relation is in good agreement with observations.  When the temperatures are not cool core-corrected, the scatter in the L-T relation is increased significantly for the {\\it AGN} run, with several outliers lying well off the mean relation.  This is due to recent large episodes of BH heating.  Such outliers appear to be very rare in nature.  We speculate that if the heating temperature were increased to a higher level than $10^8$ K the scatter would be reduced, as the heated particles would cease to emit most of their radiation in soft X-rays.  Also, the length of time between outbursts would likely increase, as more mass will have to be accreted by the black hole in order to have sufficient energy to heat the gas to a higher temperature.} \\item{The {\\it AGN} run yields stellar K-band luminosities of both overall group and the CBG that are in good agreement with observations.  The inclusion of feedback from BHs also yields the correct emission-weighted mean age of CBGs (as well as the scatter in age) and the fraction of CBGs that are forming stars at the present day ($\\sim 15\\%$).  In contrast, the simulation with supernova feedback alone produces CBGs that are far too massive and young (with very little scatter in mean age) and all of the CBGs in this run have significant present-day star formation rates.} \\item{Both models yield total hot gas-phase iron masses that agree with observed masses and both yield large-scale gradients that are similar to the observed profiles. However, while the simulations reproduce the observed slope and normalisation of the iron abundance profile, both appear to be over-abundant in silicon at small radii.  The discrepancy is, however, not greater than a factor of two and therefore lies within the current uncertainty in nucleosynthesis yields, so at present this is not a significant problem for the simulations.} \\item{Neither simulation yields the correct metallicity of stars, with the simulation with BH feedback producing too metal poor stars while the simulation without BH feedback produces too metal rich stars.  This may signal that the present implementation of AGN feedback is too efficient in expelling metals from star forming regions and/or that the adopted yields are incorrect in detail.  If it is the former, we note that since the mass in metals in the hot gas phase is much larger than that locked up in stars, only a relatively small fraction of metals would be needed to be removed from the intragroup medium to resolve this problem.} \\end{itemize}  Based on the above, the inclusion of feedback from supermassive black holes significantly improves the ability of our cosmological hydrodynamical simulations to yield a realistic population of galaxy groups.  This agrees qualitatively with the findings of Sijacki et al.\\ (2007), Puchwein et al.\\ (2008, 2010), and Fabjan et al.\\ (2009), who implemented BH growth and feedback in cosmological zoomed simulations.  We note, however, that our conclusions are based on a comparison of a much larger sample of simulated groups to a wider range of observations.  In addition, the \\owls implementation of BH feedback appears to have yielded better matches to several different key observables than those found in some of these previous studies.  For example, Fabjan et al.\\ (2009) and Puchwein et al.\\ (2008, 2010) report stellar mass fractions within $r_{500}$ that are a factor of $\\sim 2$ higher than observed and what we obtain from the {\\it AGN} model.  In addition, Fabjan et al.\\ (2009) find their galaxy groups have too high entropies at $r_{2500}$.  Given the many differences in the detailed implementations of cooling, star formation, chemodynamics, and feedback in the \\owls simulations and those which are based upon the model of SDH05, it is not trivial to determine what physics is most responsible for the better match to the data.  One possibility is the simulations of Sijacki et al.\\ (2007) and Puchwein et al.\\ (2008) did not include metal-line cooling which is expected to be important.  We have highlighted a number of areas where the current model of {\\it AGN} feedback could be improved, including reproducing the scatter in the (uncorrected) luminosity-temperature relation and exploring whether or not simple adjustments in the adopted SNe yields will allow the simulations to better reproduce the gas-phase silicon abundance and stellar metallicities.  We leave this for future work.  Recently, Dav{\\'e} et al.\\ (2008) argued that it is possible to reproduce the entropy and metallicity distributions of the hot gas in groups with simulations that invoke galactic winds with scalings similar to those expected for outflows driven by radiation pressure from massive stars on dust grains, but which do not include AGN feedback.  While it is not straightforward to directly compare the \\owls simulations with those of Dav{\\'e} et al.\\ (as there are differences in the implementations of  star formation, chemodynamics, cooling, and feedback), some of their results do not appear to be inconsistent with ours, as we have also demonstrated that models without AGN feedback can reproduce these quantities.  However, their model also appears to yield stellar mass fractions in reasonable agreement with the observations, whereas the \\owls {\\it REF} run does not.  Plausible reasons for this difference are that the energy injected in the momentum-driven wind implementation of Dav{\\'e} et al.\\ exceeds that available from supernovae (see Schaye et al.\\ 2009). % Thus, one must invoke alternative energy sources, namely high energy photons exerting pressure on dust grains, to motivate this model. By contrast, supermassive BHs are directly observed to be heating the hot gas in groups and clusters and the efficiency we have adopted has been constrained to yield a good match to the present-day black hole mass---halo mass relation and the cosmic black hole density.  In addition, as Dav{\\'e} et al.\\ point out, momentum-driven winds are incapable of shutting down cooling flows in the most massive haloes at the present-day, and therefore yield CBG star formation rates well in excess of what is observed.  Furthermore, this model does not appear to reproduce the large scatter seen in the X-ray luminosity-temperature relation of groups. For these reasons, we argue that feedback from supermassive BHs is likely a key ingredient in galaxy group formation and evolution.  This certainly does not preclude, however, feedback from BHs and some form of momentum-driven winds working in tandem to shape the properties of groups.  Additional insights from observations of galaxy groups would be very helpful for discriminating between contending models and to further constrain the parameters of the subgrid physics.  In particular, extending entropy measurements down to lower system masses, where the differences between the models should be largest (e.g., Fig.\\ 2), would be very useful.  Going beyond globally-averaged properties and azimuthally-averaged profiles and making detailed comparisons with the full 2D maps of groups (e.g., Finoguenov et al.\\ 2006) also represents a very promising avenue for distinguishing between AGN heating mechanisms (e.g., by looking at the structure of cavities and sound-wave ripples).  In addition, adopting well-defined group selection functions for observed samples that can be applied directly to simulated catalogs would be helpful.  At present, most samples of groups with high quality X-ray observations do not have a well-defined selection function.  They are comprised of systems that have high enough surface brightnesses to extract radial profiles, but it is presently not clear whether these are `typical' systems.  In the case of massive galaxy clusters, for example, the central entropy varies by more than an order of magnitude at fixed halo mass (Cavagnolo et al.\\ 2009).  The systems with the highest central entropies (longest cooling times) are referred to as non-cool core systems (or non-cooling flow systems, in the old jargon), and may represent more than 50\\% of the cluster population.  The analogues of non-cool core clusters on the group scale have not yet been discovered, maybe because they do not exist, or perhaps, more interestingly, because their X-ray surface brightnesses are so low that they have escaped detection.  If a population of `X-ray dark groups' could be convincingly demonstrated to exist (e.g., in optically-selected group samples with follow-up X-ray observations), this would present a very interesting challenge to current cosmological simulations."
    },
    "0911/0911.4251_arXiv.txt": {
        "abstract": "{Combined spectroscopic abundance analyses of stable and radioactive elements can be applied for deriving stellar ages. The achievable precision depends on factors related to spectroscopy, nucleosynthesis, and chemical evolution.} {We quantify the uncertainties arising from the spectroscopic analysis, and compare these to the other error sources.} {We derive formulae for the age uncertainties arising from the spectroscopic abundance analysis, and apply them to spectroscopic and nucleosynthetic data compiled from the literature for the Sun and metal-poor stars.} {We obtained ready-to-use analytic formulae of the age uncertainty for the cases of stable+unstable and unstable+unstable chronometer pairs, and discuss the optimal relation between to-be-measured age and mean lifetime of a radioactive species. Application to the literature data indicates that, for a single star, the achievable spectroscopic accuracy is limited to about $\\pm$20\\,\\% for the foreseeable future. At present, theoretical uncertainties in nucleosynthesis and chemical evolution models form the precision bottleneck. For stellar clusters, isochrone fitting provides a higher accuracy than radioactive dating, but radioactive dating becomes competitive when applied to many cluster members simultaneously, reducing the statistical errors by a factor $\\sqrt{N}$.} {Spectroscopy-based radioactive stellar dating would benefit from improvements in the theoretical understanding of nucleosynthesis and chemical evolution. Its application to clusters can provide strong constraints for nucleosynthetic models.} ",
        "introduction": "The determination of absolute ages for stars is of paramount importance in astrophysics since ages allow the timescales on which stellar populations and their chemical composition evolve to be established.  The presence of long-lived unstable elements in stellar atmospheres offers a way to date stars by investigating their radioactive decay spectroscopically. In its simplest form, one determines the photospheric abundance of a pair of elements, of which one is unstable. By means of theoretical modeling of the nucleosynthetic production channels for the two elements, the original abundance of the unstable nucleus can be estimated from the abundance of the stable one. This allows the time-span over which the unstable element decayed to be determined.  The idea emerged early after the discovery of the rapid neutron capture process (short $r$-process) as the source of production of radionuclides such as thorium and uranium \\citep[see][and references therein for an account of its early applications]{cowan91}. Since then it has been applied in a number of cases \\citep[among the most recent][]{frebel07} but, typically, as a ``side product'' of analyses of the more general n-capture enrichment patterns.  To our knowledge, the only effort in recent years to exploit radioactive dating to systematically derive the age of a significant sample of stars has been the one by \\citet{delpeloso05,delpeloso2,delpeloso3}.  One interest in such dating techniques lies in its applicability to field stars for which no distance information, hence no accurate information on surface gravity or absolute luminosity, is available, so that they cannot be placed with confidence on theoretical isochrones in the Hertzsprung-Russell diagram (HRD). For cluster stellar populations, on the other hand, HRD fitting usually provides rather precise absolute ages \\citep[e.g.][]{vandenberg96}. However, for stellar clusters (especially globular clusters), radioactive dating could produce an age estimate of comparable or even higher precision if applied to a large enough sample of stars, and would as such be very valuable as a test of the theoretical stellar isochrones.  Although potentially effective, spectroscopy-based radioactive dating is difficult to apply. The most easily measured long-lived radioactive element, thorium, is only accessible (with the exception of strongly $r$-process enriched stars) through a single, weak line in the optical spectrum of stars, in which high accuracy is needed in the measurement to provide a reliable age. Besides lying on the red wing of a Fe-Ni blend, this unique \\ion{Th}{ii} line is itself also blended by weaker lines of vanadium and cobalt.  The relative strength of the various components for the Sun are discussed in \\citet{caffau08hfth} where the line-list of \\citet{delpeloso05} was applied. Although rather easily measured, the most frequently used reference element, europium, is believed to constitute a less-than-optimal reference. In fact, its nucleosynthetic ties with Th are rather weak, and the production ratio is sensitive to the originating supernova characteristics \\citep[][]{kratz07,farouqi08}.  Besides the astrophysical problem that the $r$-process site has not been unambiguously identified, it is unclear how universal the resulting abundance pattern of the $r$-process nucleosynthesis actually is.  In view of the complex nuclear reactions shaping the $r$-process it is not obvious that the outcome is the same for a potentially wide range of the physical conditions governing the formation site(s) \\citep{goriely99,arnould07,sneden08,roederer09}.  To minimize the influence of poorly constrained physical conditions at the formation site it is desirable to use a reference element with an atomic mass as similar as possible to thorium. This should reduce the uncertainties in the formation properties to the uncertainties related to the properties of the involved nuclei as such, and environmental factors should cancel out. As pointed out before by various authors \\citep[e.g.,][]{goriely99} the Th/U pair is in this sense optimal, however, again difficult to apply since uranium is difficult to detect spectroscopically.  Recently, \\citet{kratz07} suggested the element hafnium as suitable reference element for thorium-based radioactive dating. Hafnium is closer in mass to thorium than europium, but more importantly, Kratz and collaborators showed by detailed modeling that the production ratio Th/Hf is not strongly varying over a wide range of neutron fluxes during the $r$-process. Since hafnium is a spectroscopically accessible element this makes it an interesting candidate as reference element.  Observationally, the forthcoming introduction of Extremely Large Telescopes (ELTs) is likely to make radioactive dating of stars a much more exploitable technique, since the high signal-to-noise (S/N) ratio required is one of the main difficulties associated with such studies. In this perspective, a precise assessment of the uncertainties associated with such technique is currently lacking, and attempted here.  We derive the formulae for the precision to which spectroscopy can provide age estimates. They are subsequently applied to abundance analyses of stars compiled from the literature.  This paper is a ``spin-off'' of our work \\citep{caffau08hfth} on the thorium and hafnium abundance in the Sun, and the Th/Hf chronometer pair often serves as an example in the following sections. ",
        "conclusions": "The accuracy of radioactive cosmo-chronometry is not going to be extremely high on a single star.  A realistic estimate of what future instrumentation will be able to achieve does likely not allow one to hope for a better than 2\\pun{Gyr} precision, taking into account {\\em observational} uncertainties only, leaving aside the systematics associated to production rates and chemical enrichment models. Although high-mass n-capture elements are currently believed to be produced in a fairly well understood manner by the so-called ``main'' $r$-process \\citep[][]{farouqi08,kratz07}, the quantitative determination of the production ratios $\\beta$ between the stable and the unstable elements still bear significant uncertainties.   Since in photon noise dominated spectra the measurement error on the equivalent width decreases linearly with the observational S/N ratio, the attainable dating precision increases roughly linearly with the S/N ratio. This constitutes a major practical challenge to the application of the method, since many of the involved lines are in the blue part of the optical spectrum (e.g. for Th, U, and Hf), and the natural target of the analysis are cool stars.  If possible, the decaying element should be chosen such that its mean lifetime $\\tau$ is similar to the age one wants to measure.  This applies also to the ``effective'' $\\tau$ of a chronometric pair constituted by two decaying species. From this perspective, the pairs U/Th, Th/stable, and U/stable are well suited to measure ages of old objects around 10\\pun{Gyr} of age. No other naturally occurring isotope has a mean lifetime providing a similarly close match.  Potentially important systematic spectroscopy-related effects that we feel deserve further study to fully exploit the potentials of radioactive dating are effects related to departures from thermodynamic equilibrium and gas-dynamics on the line formation of the chronometric species. For instance, the application of a 3D radiation-hydrodynamical model atmosphere by \\citet{caffau08hfth} led to a downward revision of the photospheric Solar Th abundance by 0.1\\pun{dex}. This corresponds to a decrease of Th ages by 4.7\\pun{Gyr}.   Globular clusters provide fairly large samples of metal-poor, bright, coeval and chemically homogeneous stars. Simultaneous measurement of abundances in these stars allows us to reduce the observational, statistical uncertainties. This could permit to derive $r$-process production ratios empirically by calibration against the known ``photometric'' age. To our knowledge, this has never been attempted so far."
    },
    "0911/0911.1642_arXiv.txt": {
        "abstract": "We present a new flexible estimator for comparing theoretical templates for the predicted bispectrum of the CMB anisotropy to observations. This estimator, based on binning in harmonic space, generalizes the ``optimal'' estimator of Komatsu, Spergel, and Wandelt by allowing an adjustable weighting scheme for masking possible foreground and other contaminants and detecting particular noteworthy features in the bispectrum. The utility of this estimator is illustrated by demonstrating how acoustic oscillations in the bispectrum and other details of the bispectral shape could be detected in the future PLANCK data provided that $f_{NL}$ is sufficiently large. The character and statistical weight of the acoustic oscillations and the decay tail are described in detail. ",
        "introduction": "At present the primordial blackbody component of the cosmic microwave background (CMB) appears very nearly Gaussian, as one would expect from inflationary models for the very early universe. To a first approximation, because of the weak coupling of the physics involved, inflation predicts a nearly scale-invariant spectrum of primordial density perturbations whose imprint on the CMB is completely characterized by the power spectrum $C_\\ell ^{AB}$, where $\\ell $ is the multipole number and $(A,B=T,E)$, according to the Gaussian distribution: \\ba P(\\{ a_{\\ell , m}\\} ) & = & \\prod _{\\ell , m} (2\\pi )^{-1/2} \\det {}^{-1/2} \\left[ C_\\ell \\right] \\cr && \\times \\exp \\left[ -\\frac{1}{2}a_{\\ell , m}^TC_\\ell ^{-1/2}a_{\\ell , m}\\right] \\ea where \\be a_{\\ell , m}= \\begin{pmatrix} a_{\\ell , m}^T\\cr a_{\\ell , m}^E\\cr \\end{pmatrix}, \\qquad C_\\ell = \\begin{pmatrix} C_\\ell ^{TT}  & C_\\ell ^{TE}\\\\ C_\\ell ^{TE}  & C_\\ell ^{EE}\\\\ \\end{pmatrix} \\ee with $T$ denoting temperature and $E$ the $E$-mode polarization.  It has however been noted that even in single scalar field inflationary models small nonlinear corrections do arise, at lowest order in the bispectrum \\citep{mald, acquaviva}. In models with several scalar fields even larger and potentially detectable degrees of non-Gaussianity are possible \\citep{ng_from_inflation, RSvT1, RSvT2, narsas, byrnes1, byrnes2, huang}.  Non-Gaussianity manifests itself in odd $n$-point correlation functions or in the connected even $n$-point correlation functions, from which the trivial part expressible as combinations of two-point correlation functions has been subtracted away. The extent of departures from Gaussianity can be characterized by ratios of higher-order correlation functions and the appropriate combination of two-point correlation functions \\citep{bernardeau}. The evolution, for the most part linear, of the primordial fluctuations of the inflaton field, involving both gravity and hydrodynamics, leads to CMB anisotropies whose statistical properties mirror those of the primordial fluctuations. Consequently, by studying higher-order correlation functions of the CMB anisotropies, we can detect and characterize any primordial non-Gaussianity.  The lowest order such statistic is the bispectrum, or three-point correlation function in Fourier space. The bispectrum has been shown to be an optimal statistic for measuring non-Gaussianity in the sense that the signal-to-noise squared of the non-Gaussianity estimator based on the three-point correlation function dominates over all higher-order estimators \\citep{babich}. Consequently it would also be significantly easier to constrain. Assuming that the characteristic amplitude of the bispectrum is much larger than that of the trispectrum, or the four-point correlation function in Fourier space, the non-Gaussianity of the density fluctuations from inflation for the single-field case may be approximated by the following ``local'' ansatz for the gravitational potential \\citep{fnl_expansion,acoustic} \\be \\Phi ({\\bf x}) = \\Phi _{G}({\\bf x}) +f_{NL} \\left [ {\\Phi _{G}}^2({\\bf x}) - \\langle {\\Phi _{G}}^2({\\bf x}) \\rangle \\right ], \\label{ansatz} \\ee where the nonlinearity parameter $f_{NL}$ may be considered small. While the three-point correlator of the linear or Gaussian part $\\Phi_G$ is zero (because in the quantum picture it involves the expectation value of a product of three creation and annihilation operators), this is no longer true when the quadratic correction is taken into account. The bispectrum will be proportional to $f_{NL},$ which determines the size of non-Gaussianity. While this form is only approximate for the single-field case, for the observationally more promising multi-field models for which potentially observable values of $f_{NL}$ can be generated, this ansatz is very accurate because the bulk of the non-Gaussianity is imprinted after horizon crossing.  The primordial bispectrum is characterized by an overall amplitude $f_{NL}$ and the shape. Translational and rotational invariance reduce the bispectrum to a function of the lengths only of the three wave vectors. A classification of bispectra shapes has been proposed by \\cite{shapeNGcamb}. Local and warm models peak on squeezed triangles, the equilateral models on equilateral triangles, and the folded/flat models on flattened triangles. Models in which perturbations are generated outside the horizon produce a bispectrum that peaks on squeezed triangle configurations, examples of which are multi-field inflation \\citep{curvaton,curvaton_b,RSvT2}. [For a review, see \\cite{bartolo_et_al_review}.] The bispectrum of standard single-field inflation can be viewed as a superposition of the local shape and the equilateral shape, both terms, however, being slow-roll suppressed \\citep{mald, shapeNGcamb, current_limits}. In particular, for the squeezed triangle configurations in single-field inflation the non-Gaussian signal would be proportional to the tilt of the power spectrum and thus imply a strong deviation from scale invariance \\citep{theorem,theorem_b}. A bispectrum peaking on equilateral triangles is predicted in models with non-standard kinetic terms in the inflaton Lagrangian, which is the case of multi-field inflation models such as DBI inflation \\citep{dbi}, ghost inflation \\citep{ghost} or trapped inflation \\citep{trapped}. Flattened-triangle configurations arise from initial conditions which deviate from the Bunch-Davies vacuum \\citep{folded}. The limits for $f_{NL}$ of the local shape from the WMAP5 analysis are $-9 < f_{NL}^{local} < 111$ at 95\\% confidence (i.e.\\ $2\\sigma$) \\citep{wmap5}.  After horizon crossing, non-linearities in both the gravitational and hydrodynamical evolution of the baryon-photon fluid prior to recombination \\citep{recombination_b,recombination_c,recombination_d,recombination_e}, as well as higher orders in the gravitational potential during recombination \\citep{recombination}, can generate non-Gaussianity. Other sources of non-primordial non-Gaussianity include secondary anisotropies such as weak lensing via the cross-correlation with the unlensed CMB arising from the integrated Sachs-Wolf effect \\citep{smithzaldar, lensingSW, MangilliVerde} or the Sunyaev-Zel'dovich effect \\citep{lensingZS}, as well as foregrounds such as dust, galactic synchrotron radiation and unresolved point sources. Finally there are also instrumental effects, see e.g.\\ \\cite{donzelli}. These effects contribute spurious non-Gaussian signals, thus biasing the measurement of the primordial signal. It is therefore important to develop tools to isolate the primordial signal from the contaminants.  The separation of the CMB and foreground components exploits their differing spectral and spatial distribution \\citep{wmap5_gold}. It has been shown that foreground sources distributed in an anisotropic manner can mimic a non-Gaussian signal which could be taken for primordial \\citep{carvalho}. It has been suggested that the claimed detection of $f_{NL}$ \\citep{yadav} might be due to residual foregrounds or instrumental contamination \\citep{pietrobon, cabella, smith}. Consequently it is important to develop estimators based on the bispectrum also able to mask out non-primordial parasite non-Gaussianities.  In this paper we shall focus primarily on bispectral non-Gaussianity of the ``local'' type just defined, although most of the discussion generalizes straightforwardly to other templates. See also \\cite{JamesPaul, shapeNGcamb}, whose authors are working on another bispectral estimator program, specifically aimed at dealing with theoretical primordial bispectra of general momentum dependence. In fact, their expansion of a general primordial bispectrum in terms of separable eigenfunctions is one way to easily extend our work beyond the ``local'' type of primordial non-Gaussianity. Unlike much of the literature, where the emphasis is on developing an ``optimal'' estimator, where all the data are reduced to a single number \\citep{heavens_a,koma,yadavetal,creminelli}, the emphasis in our paper is placed on developing a more flexible approach where the data can be divided in many ways and under which the effects of parasite non-Gaussianities can be masked in a well-defined way. See also \\cite{mexican_hat,needlets,needlets_b, hikage,munshia,munshib,santos,smith} for other estimators, many not based on the bispectrum but on wavelets, needlets, or Minkowski functionals.  The paper is organized as follows. In section 2 we establish some notation and discuss some general properties of estimators. A simplified and intuitive discussion of the nature and distribution of statistical weight of the \"local\" bispectral signal is given and it is shown how to modify templates in order to best mask parasite contaminants. Section 3 presents a quantitative discussion of the shape of the \"local\" signal and in particular the acoustic oscillations are examined quantitatively. In section 4 the signal-to-noise for acoustic oscillations in the vertex region (where $\\ell _1\\ll \\ell _2, \\ell _2$) is examined. Section 5 gives the computational details of a binned bispectral estimator that can be used for a general template and presents tests validating this estimator. Finally in section 6 we present some concluding remarks. ",
        "conclusions": "The search for primordial non-Gaussianity is presently in its early stages. Non-Gaussianity of the local type is motivated by multiple-field inflation and if detected would have a substantial impact on our understanding of the physics of the early universe. At present we have only an upper limit with $f_{NL}^{local}$ in the interval $[-9,111]$ at 95\\% confidence according to the official WMAP5 analysis \\citep{wmap5}, which constitutes a hint of a detection. A detection has been claimed in another analysis of the WMAP data by \\cite{yadav}, and several other alternative analyses, such as using needlets \\citep{needlets,needlets_b} or spherical Mexican hat wavelets \\citep{mexican_hat}, have been published. See also \\cite{hikage,wmapCrem,wmapCremBis,current_limits,smith}. These analyses differ in detail but are broadly consistent. Note that a confirmed detection of a primordial $f_{NL}$ of order 10 would rule out the simplest single-field slow-roll inflation models.  The PLANCK space mission promises to bring about approximately an order of magnitude improvement in sensitivity to $f_{NL},$ so if the present hint of a signal is confirmed, we will have an abundance of signal-to-noise at our disposal, allowing the data to be divided in many ways in order to confirm the predicted shape for the bispectrum and to exclude other non-primordial sources of parasite bispectral power.  In this paper, we demonstrated how a binned bispectrum estimator can be useful to evaluate the perfect template with a modest number of bins, so that applying the estimator is computationally extremely efficient. Compared to other implementations of the optimal estimator, which are likewise fast, the present estimator has the advantage that it can be used to apply other templates, either for $f_{NL}$ or to determine other bispectral characteristics. It will be important to develop models of the expected pattern of bispectral non-Gaussianity arising from foregrounds, instrumental artifacts, and foreground subtraction residuals. Even if very approximate, such templates will be invaluable for excluding non-primordial explanations for a possible future detection. We showed how using a binned estimator one should modify the template for primordial non-Gaussianity of the local type to optimally mask the contaminants.  We also investigated how the acoustic oscillations manifest themselves in the CMB bispectrum from local primordial non-Gaussianity and in particular evaluated the signal-to-noise for detecting various features. Although a rich pattern of non-trivial acoustic oscillations occurs in the central region of triangle shape space, we found that these oscillations contribute negligibly to the total available signal-to-noise. However, the oscillations having the same shape as the power spectrum near the vertices of triangle shape space offer a much larger signal. We also investigated broader features due to the damping, and these may be seen with PLANCK if $f_{NL}$ is large enough. An observation of such features in the bispectrum matching the shape of the two-point power spectrum would constitute a convincing confirmation of a first $f_{NL}$ detection. The absence of a signal in other regions where none is expected also provides a powerful check against foreground and instrumental contamination."
    },
    "0911/0911.4121_arXiv.txt": {
        "abstract": "We present resolved stellar photometry of NGC~2976 obtained with the Advanced Camera for Surveys (ACS) as part of the ACS Nearby Galaxy Survey Treasury (ANGST) program.  The data cover the radial extent of the major axis of the disk out to 6~kpc, or $\\sim$6 scale lengths. The outer disk was imaged to a depth of $M_{F606W}\\sim 1$, and an inner field was imaged to the crowding limit at a depth of $M_{F606W}\\sim -1$.  Through detailed analysis and modeling of these CMDs we have reconstructed the star formation history of the stellar populations currently residing in these portions of the galaxy, finding similar ancient populations at all radii but significantly different young populations at increasing radii. In particular, outside of the well-measured break in the disk surface brightness profile, the age of the youngest population increases with distance from the galaxy center, suggesting that star formation is shutting down from the outside-in.  We use our measured star formation history, along with H~I surface density measurements, to reconstruct the surface density profile of the disk during previous epochs. Comparisons between the recovered star formation rates and reconstructed gas densities at previous epochs are consistent with star formation following the Schmidt law during the past 0.5 Gyrs, but with a drop in star formation efficiency at low gas densities, as seen in local galaxies at the present day.  The current rate and gas density suggest that rapid star formation in NGC~2976 is currently in the process of ceasing from the outside-in due to gas depletion.  This process of outer disk gas depletion and inner disk star formation was likely triggered by an interaction with the core of the M81 group $\\gap$1~Gyr ago that stripped the gas from the galaxy halo and/or triggered gas inflow from the outer disk toward the galaxy center. ",
        "introduction": "Bursts of star formation play a significant and important role in the evolution of galaxies, particularly at low masses, such as dwarf galaxies.  Star formation histories (SFHs) of dwarfs in the Local Group show strong evidence for past or current bursts \\citep{mateo1998,dohm-palmer2002,dolphin2005,young2007,cole2007}. The SFHs of the non-tidal dwarfs in the M81 group are equally diverse \\citep{weisz2008}.  Recent observations of the entire local volume within 11 Mpc also indicate that bursts are an essential mode of star formation in low mass galaxies \\citep{lee2009}.  Furthermore, numerical simulations suggest that episodic bursts of star formation should be common in low-mass galaxies, as gas is expelled by supernovae and then falls back into the galaxy \\citep[e.g.][]{stinson2007}.  While there is little question that bursts are important in the formation and evolution of low-mass galaxies, most of the stellar mass in dwarfs may not form in bursts.  For example, several studies suggest that dwarf irregular galaxies have relatively constant star formation histories \\citep{gallagher1984,greggio1993}.  Even though bursts clearly occur, most recent measurements suggest they are responsible for the production of only about a quarter of the stellar mass in dwarfs \\citep{lee2009}. Recent studies of nearby low mass starbursts have also put constraints on their durations, suggesting they typically last 200--400 Myr \\citep{mcquinn2009}, in agreement with results from numerical simulations.  Much of the recent progress made in our understanding of the role of star formation bursts in the evolution of low-mass galaxies is due to the availability of deep resolved stellar photometry from the {\\it Hubble Space Telescope (HST)}.  Using such deep photometry, it possible to study galaxies using techniques previously available only within the Local Group.  Rather than inferring details of evolution from integrated properties such as scale length, color, or surface brightness, the star formation history (SFH) of the stars in the galaxy can be determined by fitting the distribution of stars in color-magnitude diagrams (CMDs) with model distributions determined from stellar evolution isochrones, \\citep[see e.g.][for Local Group examples]{holtzman1999,dolphin2000b,williams2002,dohm-palmer2002,harris2004,dolphin2005}.  Detailed measurements of SFHs can be compared to the results of numerical simulations, which now probe the properties of the stellar populations they produce, to shed light on the physical processes responsible for the morphological structures we observe.  Such simulations are beginning to show the importance of galaxy interactions in the evolution of disks \\citep{governato2007}, of internal dynamical interactions in distributing stars of different ages throughout the disk \\citep{roskar2008}, and of galaxy mass, feedback and gas availability in the process of star formation \\citep{stinson2007}.  However, the simulations that show the most detailed features do not yet include environmental effects and cover limited ranges of galaxy mass.  NGC~2976, which has $M_B = -17$, $W_{20} = 165$ km s$^{-1}$, $(m-M)_0 = 27.76\\pm0.23$, $A_V=0.23$, and a morphology of SAc pec, \\citep{simon2003,karachentsev2002,schlegel1998}, is an excellent specimen for studying the effects of bursts of star formation and environment on the evolution of low-mass galaxies.  This peculiar low-mass disk galaxy lies on the outskirts of the M81 group, outside of the \\citet{yun1994} H~I maps.  The inner disk of the galaxy contains many young stars, has a sharp truncation edge \\citep{simon2003}, appears to have a barred potential \\citep{spekkens2007}, and may be inside of a larger spheroidal halo \\citep{bronkalla1992}.  The galaxy's gas content and ongoing star formation \\citep{bigiel2008,leroy2008} are affecting its morphology. In order to investigate this impact, as well as potential effects of the M81 group environment, we have performed a systematic study of the stellar populations as a function of galactocentric distance in this peculiar galaxy.  Our results suggest NGC~2976 is undergoing a transition from a burst of star formation to a more quiescent state, and thus we have a unique opportunity to study the final stages of a common mode of star formation.  Herein we report the SFH, measured via CMD fitting, of several regions of NGC~2976.  We look for trends with radius, analyzing regions of constant crowding and differential extinction.  We show that the ancient population is consistent with being similar over all radii. However, the age of the young population appears to increase sharply with radius beyond the disk break.  By reconstructing the past gas surface density profile, we show that such an abrupt change in the age distribution is likely to be due to the depletion of gas in the NGC~2976 outer disk, possibly caused by an interaction with the core of the M81 group.  The paper is organized as follows.  Section 2 details our data acquisition and reduction techniques. Section 3 compares and contrasts the recovered SFHs of the different regions. Section 4 discusses and interprets potential explanations for the differences in a galaxy evolution context. Finally, section 5 summarizes our main conclusions. We assume $(m-M)_0 = 27.76\\pm0.23$ \\citep{karachentsev2002} for conversions of angular measurements to physical distances and an inclination angle $i$=64.5$^{\\circ}$ \\citep{deblok2008} for surface density measurements.  We adopt a WMAP \\citep{dunkley2009} cosmology for all conversions between time and redshift. ",
        "conclusions": "We have measured resolved stellar photometry of 2 HST/ACS fields covering $\\sim$6 scale lengths of the NGC~2976 disk.  The resulting color-magnitude diagrams show a radial trend where the fraction of main sequence stars decreases with increasing radius, even outside of the most active central region.  Modeling the CMDs to recover the age distribution of the stars reveals that, while the outer regions have undergone very little star formation for the past $\\sim$500 Myr, the inner regions have continuously formed stars to the present.  We see this lack of young stars begin at a radius of $\\sim$3~kpc, well outside of the disk break.  Inside of 3~kpc, the recent star formation rate is similar to the rate in past epochs. This trend of increasing age of young stars with galactocentric distance is similar to an age trend recently measured in the disk of the Large Magellanic Cloud, also using CMD analysis \\citep{gallart2008}.  The trend suggests a truncation of star formation from the outside-in.  The radially-dependent differences in the young populations are in contrast to the apparent uniformity of the old populations across the galaxy, which appear to be consistent with one another in proportion and metallicity.  It is possible that the old stars are dominated by a spheroidal component of the galaxy \\citep{bronkalla1992} and only the young stars are truly sampling the disk.  Alternatively, the old stars may simply be a well-mixed old disk population.  Given the correct barred potential of the galaxy \\citep{spekkens2007}, radial orbits should have led to significant radial redistribution of stars.  Detailed study of the gas in NGC~2976 shows clear evidence that it has recently been rapidly depleted. The current gas densities and gas consumption lifetimes in NGC~2976 suggest that the dominant mechanism behind the observed age gradient is the quenching of the gas disk. The lack of clustering of the few young stars at large radii does admit the possibility that radial scattering may have played a role in placing some of the young stars into the outer disk.  However, we consider the shutting down of star formation in the disk from the outside-in to be the dominant process in the current morphological transition occurring in NGC~2976.  This process may be related to ram pressure stripping of the halo gas and/or internal gas redistribution induced by tidal forces due to the most recent interaction with the core of the M81 group.   \\bigskip   We thank Gregory Stinson for helping to compare the dwarf simulations to our measurements.  Support for this work was provided by NASA through grant GO-10915 from the Space Telescope Science Institute, which is operated by the Association of Universities for Research in Astronomy, Incorporated, under NASA contract NAS5-26555."
    },
    "0911/0911.1829_arXiv.txt": {
        "abstract": "We study the formation of voids in a modified gravity model in which gravity is generically stronger or weaker on large scales.  We show that void abundances provide complementary information to halo abundances:  if normalized to the CMB, models with weaker large-scale gravity have smaller large scale power, fewer massive halos and fewer large voids, although the scalings are not completely degenerate with $\\sigma_8$.  Our results suggest that, in addition to their abundances, halo and void density profiles may also provide interesting constraints on such models:  stronger large scale gravity produces more concentrated halos, and thinner void walls.  This potentially affects the scaling relations commonly assumed to translate cluster observables to halo masses, potentially making these too, useful probes of gravity. ",
        "introduction": "A wealth of observations, from WMAP, supernovae Ia, galaxy clustering on large scales, and cross-correlations between galaxies and the CMB, suggest that we live in a spatially flat Friedmann-Robertson-Walker universe currently dominated by either a cosmological constant or repulsive dark energy. The best fit for the dimensionless energy density parameters are $\\Omega_m=0.28$ and $\\Omega_\\Lambda=1-\\Omega_m$, within the concordance $\\Lambda$CDM Model \\cite{wmap5}.  However, because this requires the vast majority of the energy density in the universe to be in two unknown substances, dark matter and dark energy, there is considerable interest in alternative interpretations of what the data imply.  The key equation of General Relativity, Einstein's equation, relates the curvature and the expansion rate of the Universe to its matter and energy content. The current paradigm is to modify the matter content of the universe, by including dark matter and dark energy, to account for observations. Instead, however, we might modify how the universe curves in response to matter, which would mean modifying our theory of gravity.  There are many ways that one could modify gravity \\cite{B-D,milgrom,bekenstein,CarrollST,fr1,fr2,fr3,Cham,dgp,dgpSims,dgpPT,dgpSim2,lss,mannheim,diaferio,huReview}, but the one that we will focus on in this paper is purely phenomenological.  The assumption is that it is merely a modification to the potential, changing the form from that of standard gravity to: \\begin{equation} \\label{potentialeq} \\phi(r)=G m \\frac{1+\\alpha (1-e^{-r/r_s}) - \\alpha (1-e^{-r/r_c})}{r} \\end{equation} This potential (with $r_c\\to\\infty$ has been studied before \\cite{ngp,shirata1,sealfon,fritz,shirata2} and in some sense, is an interesting limiting case of some $f(R)$ models \\cite{fR,fRsims}.  In what follows, we will use $r_c\\gg r_s$ but finite, following \\cite{mss09}.  The precise choice of $r_c$ does not matter for any of our conclusions.  In \\cite{mss09} we studied the formation of nonlinear objects in this model -- the first study of virialized dark matter halo abundances in any modified gravity model.  In Section~\\ref{voids}, we study the evolution of the structures that are in some ways the opposites of halos, namely voids.  Along the way, we show that the density profiles of objects in these models exhibit interesting departures from standard gravity models.  Section~\\ref{stats} shows how we use insights from the evolution model to estimate void abundances.  Section~\\ref{sims} shows that our model captures the essence of how void abundances depend on $\\alpha$ and $r_s$ in numerical simulations of structure formation with this modified potential.  A final section summarizes our results.  An Appendix shows the corresponding trends for halo profiles; these potentially allow X-ray observations to provide powerful probes of modifications to gravity. ",
        "conclusions": "We studied a particular modification to large-scale gravity, which has two free parameters:  a scale $r_s$, and an amplitude $\\alpha$ (equation~\\ref{potentialeq}).  Previous work has shown that halo abundances can be modeled well enough \\cite{mss09} to use them to constrain such models.  We present an example of this in Appendix~\\ref{alphars}.  We argued that since voids are so large, and almost fill space, they too can be useful probes. In particular, the study of the formation of large voids provides a useful counterpart to the the study of the formation of massive halos because both probe the nature of gravity on large scales.  This is timely, given that large, statistically representative void samples are only just becoming available \\cite[e.g.][]{voidp,voids}.  Whether stronger large-scale gravity produces more or fewer large voids depends on how one normalizes the models.  When normalized to have the same power spectrum at early times (the most common choice), stronger gravity produces more large voids (Figure~\\ref{fig:comparison}).  This case also produces more massive halos, so the net result is qualitatively like having standard gravity with more large scale power (Figure~\\ref{fig:pic}).  Although there are quantitative differences which allow one to distinguish between a larger normalization and modified gravity (Figures~\\ref{fig:ratio}--\\ref{fig:ratiomodICmodgravtostan}), we note that this fact alone -- a mismatch between the CMB- and later-time determinations of the amplitude of the power spectrum on 10~Mpc scales when standard gravity is assumed -- are generic signatures of modifications to large-scale gravity.  Note that this particular signature has the same sign for halos and for voids -- stronger large scale gravity means more massive voids and more large voids.  In contrast, if gravity is standard but the initial conditions were non-Gaussian, then halo and void abundances are modified in qualitatively different ways, at least for the local non-Gaussian initial conditions of current interest \\cite{lsd09}.  However, if normalized to have the same power at $z=0$, the dependence of void (and halo) abundances on whether large-scale gravity is stronger or weaker depends on how one chooses to do this.  If this is done by increasing (decreasing) the initial power spectrum in models with $\\alpha<0$ ($\\alpha>0$), then one still expects the stronger gravity models to produce larger voids (Figure~\\ref{fig:ratiomodICmodgravtostan}), although the predicted abundances are quantitatively different.  However, if one modifies the standard gravity initial conditions to match those of the $\\alpha\\ne 0$ models (increase/decrease initial large scale power to match $\\alpha>0$/$\\alpha<0$), then one predicts more (fewer) large voids when $\\alpha<0$ ($\\alpha>0$) relative to the $\\alpha=0$ case.  This dependence on how the models were normalized is qualitatively similar to the trends seen for massive halos \\cite{mss09}.  We presented a method for estimating these effects which is in good qualitative greement with the simulations, but there are quantitative differences.  However, the agreement is not as good as it was for describing halos.  Larger simulations are required to determine if this is due to the relatively small size of the simulation boxes available to us, or to some more fundamental problem with our analysis (see \\cite{smt01,ls09} for discussion of the drawbacks of the excursion set approach).  We also showed that the density profiles of voids and halos may also provide interesting probes of modified gravity.  When the initial profile is a tophat, then the evolved halo profile in the $\\alpha<0$ case generally has a cusp at the virial radius (mass flows towards the boundary), whereas when $\\alpha>0$ the halos are more centrally concentrated (Figure~\\ref{fig:densityhalos}).  The structure of the void walls also depends slightly on $\\alpha$ (Figure~\\ref{fig:density}).  We provided tentative evidence that the profiles of halos in simulations do depend on $\\alpha$ (Figure~\\ref{simprofs}), and hope that these initial measurement motivate further study of this interesting problem.  It will be interesting indeed if these trends persist in simulations with better mass and force resolution, because it is conceivable that effects of this magnitude will soon be measured by weak lensing surveys.  The more dramatic effects associated with initially tophat profiles potentially provide an interesting X-ray signature of modified gravity -- it would be interesting to simulate such models using SPH codes.  In conclusion, we expect our results will prove useful in studies which use the large scale distribution of galaxies, and the structure of galaxy clusters, to constrain large scale modifications of gravity."
    },
    "0911/0911.0314_arXiv.txt": {
        "abstract": "{Solar flare hard X-ray spectra from RHESSI are normally interpreted in terms of purely collisional electron beam propagation, ignoring spatial evolution and collective effects. In this paper we present self-consistent numerical simulations of the spatial and temporal evolution of an electron beam subject to collisional transport and beam-driven Langmuir wave turbulence. These wave-particle interactions represent the background plasma's response to the electron beam propagating from the corona to chromosphere and occur on a far faster timescale than coulomb collisions. From these simulations we derive the mean electron flux spectrum, comparable to such spectra recovered from high resolution hard X-rays observations of solar flares with RHESSI. We find that a negative spectral index (i.e. a spectrum that increases with energy), or local minima when including the expected thermal spectral component at low energies, occurs in the standard thick-target model, when coulomb collisions are only considered. The inclusion of wave-particle interactions does not produce a local minimum, maintaining a positive spectral index. These simulations are a step towards a more complete treatment of electron transport in solar flares and suggest that a flat spectrum (spectral index of 0 to 1) down to thermal energies maybe a better approximation instead of a sharp cut-off in the injected electron spectrum.} ",
        "introduction": "Hard X-ray emission has long been used as the prime diagnostic tool to study particle acceleration and energy release in solar flares. From these X-ray observations the mean electron flux spectrum (e.g. averaged over the X-ray emitting volume, see \\citet{Brown_etal2003} for details) can be determined either through forward fitting \\citep{Holman_etal2003} or more advanced inversion techniques \\citep{Piana_etal2003,Kontar_etal2004,Brown_etal2006}. At higher energies, typically above 10-20~keV, the observed hard X-ray spectrum is considered to be due to an accelerated population of electrons being stopped by the dense chromosphere via Coulomb collisions \\citep{1971SoPh...18..489B}. The spectrum below 10-20~keV normally originates from thermal coronal sources with temperatures of 10s MK \\cite[e.g.][]{KruckerLin2008}.  The Reuven Ramaty High Energy Solar Spectrometer (RHESSI) provides high resolution HXR spectra of solar flares \\citep{2002SoPh..210....3L}, greatly improving on previous measurements \\citep{JohnsLin1992}. This high energy resolution spectra has allowed, for the first time, to scrutinize the X-ray and electron spectra in search for a non-powerlaw features, revealing vital clues about electron acceleration and transport. In some RHESSI flares, the recovered mean electron flux spectrum demonstrates a local minima or \\emph{dip} between the non-thermal and the thermal components instead of a smooth transition \\citep{Piana_etal2003,Holman_etal2003,Kasparova_etal2005,Sui_etal2007,Kontar_etal2008}. The presence of the dip, as a real physical feature of the electron spectra, has been questioned as in many cases it can be attributed to photospheric albedo, e.g. Compton back-scattered X-rays \\citep{Kasparova_etal2005,Kontar_etal2008}. However a few events have been found in which after isotropic albedo correction \\citep{Kontar_etal2008}, the X-ray spectrum is still relatively flat, so they could be fitted with a thick-target model single power-law spectrum with a low energy cutoff \\citep{Sui_etal2007}. In these flares a dip was not directly observed in the mean electron spectrum, but instead inferred from forward fitting a model with low energy cutoff to the X-ray spectrum. This model has a thermal component at low energies and at higher energies purely collisional thick-target model of a single power-law of accelerated electrons above a cutoff. In this thick-target scenario \\citep{1971SoPh...18..489B,Holman2003}, the dip in the mean electron spectrum originates from a positive slope at low energies developing below the cutoff as the accelerated electrons propagate from a coronal acceleration site downwards to the chromosphere, having Coulomb collisions with the background plasma. If the dip is real it provides important insights into flare energetics since the energy in accelerated electrons is strongly dependent on the low energy cutoff.  For any reasonable X-ray producing flare, non-collisional beam-plasma interaction is much faster than that via Coulomb collisions \\citep{ZheleznyakovZaitsev1970,Karlick2009}. Such processes are inferred to occur in downward propagating electron beams from radio observations of reverse slope drift burst in flares \\citep[e.g.][]{1997A&A...320..612K,1997ApJ...480..825A}. Although generation and escape of electromagnetic radiation from Langmuir waves in a flaring plasma is not well understood. The role of wave-particle interactions in solar flares assuming stationary, time-independent injection of electrons have been considered analytically and numerically \\citep{1984ApJ...279..882E,hamilton1987,McClements1987}. \\citet{1984ApJ...279..882E} have argued that the conditionally created distribution should be constantly being flattened by quasi-linear relaxation, while \\citet{hamilton1987} and \\citet{McClements1987} suggest that although the wave-particle interactions have an important effect, the change of the electron spectra under stationary conditions should be minor. However, in a more realistic flare conditions the injection (acceleration) of electrons is likely to be highly intermittent \\citep{TsiklauriHaruki2008} with a number of short duration pulses is often observed \\citep{Kiplinger_etal1984,1994SoPh..153..403F,Aschwanden_etal1998}, so the time-dependent solution of particle transport equations accounting for wave-particle interactions should be considered. Additionally, the previous studies did not consider the spatial and temporal evolution of the beam from the coronal source down into chromosphere - a crucial aspect when considering the propagation of an electron beam in comparison to X-ray observations.  \\begin{figure}\\centering \\plotone{fig1}\\\\ \\caption{\\label{fig:density} The background plasma density $n(x)$ as a function of the height above the photosphere. Also shown is the corresponding plasma frequency for each of these densities.} \\end{figure}  The quasi-linear relaxation process of Langmuir waves has been considered in higher velocity dimensions (other than the parallel component considered here) \\citep[e.g.][]{1980SvJPP...6Q.422C} but it has only been recently that the 2D system has been fully numerically solved \\citep{2008PPCF...50h5011Z}. Even then the evolution was considered in a spatially independent manner. In these studies it was found that the parallel component (1D) is the fastest processes and likely to dominate the electron transport.  In this letter we take a step towards a more complete treatment of electron transport in solar flares by including the spatial evolution of beam-driven Langmuir wave turbulence. We numerically study the system self-consistently, simulating the propagation of an electron beam from the coronal acceleration site down to the chromosphere, considering the truncated power-law spectrum frequently used for data interpretation \\citep[e.g.][]{Holman2003,Sui_etal2007}, to investigate the evolution of the mean electron flux spectrum below this cutoff. We demonstrate that the positive slope of the mean electron flux is not present, when the response of the background plasma via Langmuir waves to the propagating electron beam is taken into account. We also show that the injected electron spectrum flattens to a decreasing distribution due to collective interaction with plasma even for weak flares \\citep[e.g.][]{hannah2008}. Furthermore, we suggest that a flat spectrum (with spectral index 0 to 1) down to thermal energies maybe a better approximation as opposed to the sharp cut-off in the injected electron spectrum. ",
        "conclusions": "Considering both the temporal and spatial evolution, the mean electron flux spectrum is very sensitive to the generation of waves. The influence of wave-particle interactions is seen to flatten the spectral index of the electron spectrum $\\delta (E)=d\\log(\\overline{n}V\\overline{F})/d\\log(E)$ below the break but also up to where the beam-plasma interaction time is faster than the electron cloud time size. The artificially introduced low-energy cutoff in the injected electron spectrum disappears not only in the local electrons distribution function  but also in the spatially integrated electron spectrum once wave-particle interactions is taken into account. The character of beam propagation is close to the simulation results of beam transport along open field lines, where collisions are normally ignored \\citep{TakakuraShibahashi1976,MagelssenSmith1977,Melnik_etal1999,Li_etal2008}. However, there are noticeable differences. In the simplistic treatment of spatially uniform beam, when all but quasilinear terms are ignored \\citep[e.g.][]{VedenovVelikhov1963}, the generation of waves leads to an exact plateau distribution with $\\delta (E)\\approx0$. Our simulations show that the spectral index, $\\delta (E)$ is more than zero, which is the result of collisions. Similarly, \\citet{KontarReid2009} show that the spatially integrated spectrum of particles will not deviate from a initial power-law, but only when processes leading to absorption of waves or removal of waves out of resonance are included.  The simulations in this letter are for typical microflare parameters \\citep{hannah2008} and given that the wave emission scales with the electron number density we would expect wave-particle interactions to have more significant, effect in large flares. Large flares could constitute multiple intermittent bursts of accelerated electrons directed along possibly different magnetic field lines. The fast time variations in hard X-ray lightcurves \\citep{Kiplinger_etal1984} indirectly support this idea. Numerical simulations of reconnection suggest ``bursty'' electron acceleration \\citep{TsiklauriHaruki2008} and spatially fragmented electron acceleration \\citep{BianBrowning2008} which could result in electron propagation along the different lines. The footpoint motion often seen in solar flares \\citep{Krucker_etal2003} also suggest that the electrons are consecutively injected onto field lines. In this scenario an ensemble of our simulations, multiple micro-beam injections, would lead to beam densities comparable to a large flare.  The convergence of the magnetic field at chromospheric heights \\citep[e.g.][]{2008A&A...489L..57K} has been ignored in this work, however in our simulations the overall evolution of the energetic particles in the top part of a relatively dense loop, $10^{10}$~cm$^{-3}$, is dominated by wave-particle interactions where the field is not converging. It is only in the denser chromosphere, where the field is likely to converge, that the collisions become dominant. The very fast flattening of the powerlaw distribution's low energy cutoff by the wave-particle interactions suggests that it is unlikely that such a cutoff could develop and is therefore an unwise initial distribution for any model of coronally accelerated electrons. The non-thermal distribution flattening at low energies as it transitions into the thermal distribution seems to be a realistic model.  Given how strongly the total energy in the accelerated elections depends on a cutoff, or its behaviour at low energies approaching the thermal distribution \\citep{2003ApJ...595L.119E,2005A&A...438.1107G}, this transition needs further study.  The work presented here is a step towards a more complete treatment of electron transport in solar flares, moving from the standard thick target model (1D velocity with collisions only) to beam-driven Langmuir turbulence (1D velocity, 1D space, collisions and wave-particle interactions). In future work we need to extend this to consider 2 or higher dimensions of velocity space, the changing magnetic field, electron scattering as well as the transition between the accelerated beam and thermal distribution."
    },
    "0911/0911.5194_arXiv.txt": {
        "abstract": "This file {\\bf spicawSample.pdf} contains the instructions for authors preparing papers for the proceedings of the workshop \\emph{`The Space Infrared Telescope for Cosmology \\& Astrophysics: Revealing the Origins of Planets and Galaxies'}. It should also be used as the starting point for writing the papers. In order to compile this \\LaTeX{} file you need the document class {\\bf spicawStyle.cls}, as well as the \\LaTeX{}\\,2$\\epsilon$ package. ! and \\verb! ",
        "introduction": "These instructions are intended to assist you in preparing your paper for the proceedings of the \\emph{'The Space Infrared Telescope for Cosmology \\& Astrophysics: Revealing the Origins of Planets and Galaxies'} Workshop in the format required by the editors. The PDF version of the file is intended to illustrate the layout of the paper, and to give general instructions on the conventions used in typesetting these proceedings. Please see the file {\\bf spicawSample.tex} for specific examples of how to implement the various \\LaTeX{} commands, and \\emph{use that file as the starting point for writing your own paper} for the proceedings.  The {\\bf spicawStyle.cls} \\LaTeX{} document class is intended for preparing papers for the proceedings of \\emph{'The Space Infrared Telescope for Cosmology \\& Astrophysics: Revealing the Origins of Planets and Galaxies'} Workshop. It makes it possible to use all standard \\LaTeX{} facilities, such as tables and encapsulated postscript figures. Make sure that you have this file, you need it, and please \\emph{don't modify it}, Once your paper is completed, and you are happy that \\LaTeX{} produces your desired output, you should ftp the complete source (i.e.\\ the {\\tt .tex}) file plus all the figures (i.e. all the {\\tt .eps} files) to us as described in the web-page.   \\label{authorf_sec:cmd}  The standard \\LaTeX{} sectioning commands \\verb!\\section! and \\\\ \\verb!\\subsection!  should be used, and will be numbered automatically. In order to make listings use the itemize or enumerate commands, e.g.  \\begin{itemize} \\item This is the first item in an itemized list \\item And this is the second one.  This list is itemized. \\end{itemize}  \\begin{enumerate} \\item This is the first item in an enumerated list \\item And this is the second one.  This list is enumerated. \\end{enumerate}   \\subsection{Subsection example} \\label{authorf_sec:ex}  This is an example of a subsection. Please do not try to modify the font and style of the section titles. ",
        "conclusions": ""
    },
    "0911/0911.2061_arXiv.txt": {
        "abstract": "The Galactic center (GC) lobe is a degree-tall shell of gas that spans the central degree of our Galaxy.  It has been cited as evidence for a mass outflow from our GC region, which has inspired diverse models for its origin.  However, most work has focused on the morphology of the GC lobe, which has made it difficult to draw strong conclusions about its nature.  Here, I present a coherent, multiwavelength analysis of new and archival observations of the GC lobe.  New radio continuum observations show that the entire structure has a similar spectral index, indicating that it has a common origin.  The radio continuum emission shows that the GC lobe has a magnetized layer with a diameter of 110 pc and an equipartition field strength ranging from 40 to 100 $\\mu$G.  I show that optical and radio recombination line emission are associated with the GC lobe and are consistent with being located in the GC region.  The recombination line emission traces an ionized shell nested within the radio continuum with diameter of 80 pc and height 165 pc.  Mid-infrared maps at 8 and 15 $\\mu$m show that the GC lobe has a third layer of warm dust and PAH-emission that surrounds the radio continuum shell with a diameter of 130 pc.  Assuming adibatic expansion of the gas in the GC lobe, its formation required an energy input of about $5\\times10^{52}$ ergs.  I compare the physical conditions of the GC lobe to several models and find best agreement with the canonical starburst outflow model.  The formation of the GC lobe is consistent with the currently observed pressure and star formation rate in the central tens of parsecs of our Galaxy.  Outflows of this scale are more typical of dwarf galaxies and would not be easily detected in nearby spiral galaxies.  Thus, the existence of such an outflow in our own Galaxy may indicate that it is relatively common phenomenon in the nuclei of spiral galaxies. ",
        "introduction": "The GC lobe is a shell of emission rising north of the Galactic plane to a height of roughly 1\\sdeg \\citep[$\\sim$140 pc at 8 kpc][]{r93}.  As shown in Figure \\ref{gcl620}, it spans the central degree of the GC region, a region noteworthy for a massive black hole, star clusters, and exotic objects \\citep{g05,f99,y88}.  It was discovered in a 10 GHz radio continuum survey of the GC region \\citep{s84}.  The discovery observations measured a thermal spectral index at 10 GHz, implying that it had an ionized mass of $4\\times10^5$ \\msol\\ \\citep{s85}.  The shape and location of the shell motivated several models for its formation, including starburst \\citep{c92}, AGN \\citep{m01}, and magnetodynamic effects \\citep{u85}.  Each model for the formation of the GC lobe has important implications for our understanding of galactic nuclei.  If it is formed by a starburst, it tells us about the star formation history of the GC region.  Is such activity episodic \\citep{st04} or persistent \\citep{f04}?  If the GC lobe was formed by a jet, it would tell us that Sgr A* can sometimes have powerful outflows, as opposed to its relatively quiescent current state \\citep{mar05}.  Knowledge of such activity would have implications for AGN feedback in galaxy evolution \\citep{h05}.  The presence of magnetodynamic effects in the GC region would alter our interpretation of the energetics of outflows in other galaxies \\citep{su87}.  However, later observations complicated the interpretation that the GC lobe was an outflow in the GC region.  Radio continuum observations measured a nonthermal spectral index in the east \\citep{r87} and thermal emission in the west \\citep{b03b}.  This was consistent with differences in the polarization fraction of the east and west sides of the lobe, which found the east side much more strongly polarized \\citep{t86,h92}.  Furthermore, there was an unusual shocked molecular cloud (named AFGL 5376) associated only with the western half of the GC lobe, implying the two halves had different origins \\citep{u90}.  These asymmetries between east and west sides of the GC lobe cast doubts on the idea that it was a coherent structure.  The idea that the GC lobe was a coherent object was reinforced with the detection of a remarkable mid-IR shell \\citep{b03}.  The \\emph{Midcourse Space Experiment} \\citep[\\emph{MSX};][]{p01} survey of the Galactic center at 8 $\\mu$m found filamentary structures coincident with the entire GC lobe.  \\citet{b03} modeled the mid-IR emission as shock-heated dust entrained in an outflow.  This model implies that the outflow has a mass of $5\\times10^5$ \\msol\\ and (assuming an expansion velocity) a kinetic energy greater than $10^{54}$ ergs.  This mass and energy are comparable to those observed in small- to moderate-sized starburst outflows \\citep{v05}.  Despite the strong morphological evidence in radio continuum and mid-IR observations, there are several lingering questions about the nature of the GC lobe.  What created the shell?  Could it be in the foreground to the GC region?  Can it be related to other GC objects?  To answer these questions in detail, I analyzed new and archival observations of the GC lobe, including new radio continuum, optical recombination line, and mid-IR observations.  Section \\ref{obs} describes the results of each observation.  Section \\ref{points} summarizes a few simple conclusions about the nature of the GC lobe.  The morphology and physical conditions in the lobe show that it is a layered, three-dimensional shell of gas in the GC region.  Section \\ref{discussion} discusses models for the formation of the GC lobe and how it relates to other GC objects.  I find that the GC lobe is best described as an outflow powered by current star formation.  Section \\ref{conclusions} presents the conclusions of this study. ",
        "conclusions": "\\label{points} Here, I synthesize my new analysis with previous work, to draw three broad conclusions about the nature of the lobe.  Each of these conclusions has been disputed before, so it is neccessary to state them clearly before moving on with the discussion.  Using these conclusions, I derive the physical conditions of the GC lobe.  \\subsection{The GC Lobe is a Single Object} \\label{single} The discovery of radio and optical recombination lines, radio continuum, and mid-IR emission in the GC lobe, all with a similar morphology, strengthens the argument that it is a single, coherent object.  The 20 cm radio continuum and optical recombination line emission show a clear bridge between the east and west halves of the GC lobe near $b\\sim1$\\sdeg, which rejects the possibility that the east and west sides are unrelated phenomena.  As described in \\S\\ \\ref{all_gbtcont}, the 6/20 cm spectral index east and west sides have similar, nonthermal values.  More importantly, the index for both edges of the GC lobe show a similar change with Galactic latitude, with significant steepening toward the highest latitudes.  This similarity suggests that they have a similar physical origin.  Finally, \\citet{gcl_recomb} note that the GC lobe has radio recombination line emission with an unusually narrow line width.  The narrow line width implies an unusually low electron temperature for the ionized gas that is not commonly seen in typical \\hii\\ regions \\citep{a96}.  The intensity ratios of the lines detected in the east and west of the GC lobe are similar, implying that all the gas has a similar density and unusually low temperature.  \\subsection{The GC Lobe is Composed of Layered Shells} \\label{shellsec} The emission from the three components of the GC lobe have a similar, limb-brightened morphology (see Fig. \\ref{all_schem}).  All three components have a similar center longitude, while each component has a different edge-to-edge diameter (see Table \\ref{all_shell}).  This gives the appearance nested shells of radio line emission, radio continuum emission, and mid-IR emission.  \\citet{b03} model the structure as a shell-like ``telescope dome'' with a radius $r$, height $h$, and thickness $\\delta$.  Table \\ref{all_shell} shows the best fit shell parameters for five of the observations presented here.\\footnote{Parameters from the optical H$\\alpha$ are not included, since the emission is extincted and more confused.}  The shell width, $\\delta$, is fit by simulating an idealized shell model and convolving it with the telescope beam, $B$;  the thickness is only weakly constrained when it is smaller than the telescope beam.  The height of the GC lobe is measured in the 20 cm radio continuum and optical H$\\alpha$ observations to be about 1\\ddeg15, or 165 pc at the GC distance.  Under the telescope dome model, the volumes of the radio recombination line, radio continuum, and mid-IR shells are roughly $1.0\\times10^6$, $1.9\\times10^6$, and $2.8\\times10^6$ pc$^{3}$ (assuming equal heights).  To test the shell model, the edge-to-center contrast is measured for each component and compared to that of an idealized model.  I used the simulation to compare the idealized, beam-convolved shell contrast $C_{\\rm{ideal}}$ to the observed contrast.  Table \\ref{all_shell} shows that three of the observations agree with the idealized shell model, suggesting that the structure seen at those wavelengths may have an intrinsically shell-like structure.  The contrast in the GBT 20 cm continuum is lower than the model predicts, perhaps indicating that the GC lobe is not hollow, but has some 20 cm emission inside.  The contrast of the GBT recombination line structure is larger than predicted, perhaps indicating that the ionized gas is clumpy or irregular.  \\subsection{The GC Lobe is in the GC Region} \\label{ingc} Perhaps the strongest morphological connection between the GC lobe and the GC region is at the Radio Arc.  The present radio continuum survey has confirmed previous work that found contiguous emission from the east of the lobe through the Radio Arc \\citep[shown schematically in Fig. \\ref{all_schem};][]{y88}.  The Radio Arc (and other nonthermal radio filaments) are known to be within the central few hundred parsecs of the Galaxy from \\hi\\ absorption measurements \\citep{l89,r03}.  This implies that the GC lobe is also in the central few hundred parsecs.  Radio recombination line observations have also found an unusually low electron temperature that implies a high metal abundance for the ionized gas in the GC lobe \\citep{gcl_recomb}.  The high metallicity for the GC lobe is consistent with the metallicity expected in the GC region, based on the Galactic abundance gradient \\citep{a96}.  The detection of optical H$\\alpha$ emission near the north of the GC lobe is consistent with a GC distance to the ionized gas.  Also, the thermal gas pressure is similar to pressures observed in the GC region \\citep{gcl_recomb}.  The fact that the ionized component of the GC lobe is inside the other components suggests it is ionized from the inside.  This is consistent with a GC location, since the ionizing flux of massive stars in the central tens of parsecs are sufficient to ionize the gas \\citep{gcl_recomb}.  Taken together, these observations show that the GC lobe is located in the central few hundred parsecs of our Galaxy.  Thus, the appearance of the GC lobe spanning the central 100 pc in projection is likely to be true physically.  \\subsection{The Physical Conditions of the GC Lobe} \\label{physcond} Given that the GC lobe is composed of three nested shells and is located in the GC region, I estimate several of its physical parameters.  As a crude minimum energy required to form the GC lobe, I estimate its gravitational potential.  The brightness-weighted mean height of the radio recombination and mid-IR emission of the GC lobe is about 50 pc north of the plane.  The gravitational potential difference between a galactocentric radius of 1--20 pc and 50 pc is $E_{gr} = 5\\times10^{51} * (T_e/3960\\ \\rm{K})^{0.61}$ ergs, for the ionized gas mass \\citep[equivalent to a velocity of about 40 \\kms;][]{b91}.  The minimum energy estimate is likely to be an underestimate because the total mass is expected to be at least twice that in the ionized gas \\citep[$M_{mol}>3\\times10^5$ \\msol;][]{b03}.  However, the minimum energy may be overestimated by assuming that all the mass originates between 1 and 20 pc.  The true minimum energy is likely to be within an order of magnitude of that calculated above.  The energy required to form the GC lobe can be estimated from the pressure observed throughout it.  The thermal pressure in the ionized gas is measured by the radio recombination line observations to be $P/k=3.8\\times10^6 (T_e/3960\\ \\rm{K})$ K cm$^{-3}$ \\citep{gcl_recomb}.  For comparison, X-ray--emitting gas in the plane has pressures of $1-5\\times10^6$ K cm$^{-3}$ \\citep{k96, m04}, molecular gas velocity dispersion implies virial pressure of $6\\times10^6$ K cm$^{-3}$, and the equipartition magnetic field implied by the radio synchrotron emission has a pressure of $1\\times10^6$ K cm$^{-3}$.  For an adiabatically expanding shell, the formation energy is $E=\\frac{\\gamma}{\\gamma-1}PV$, where $V$ is the shell volume and $\\gamma$ equals $5/3$ if the gas is nonrelativistic and $4/3$ if relativistic \\citep{d04}.  The present observations have detected relativistic (synchrotron) and nonrelativistic (recombination line) emission from inside the shell, but have not measured the pressure from that region; it is not clear that one dominates the internal pressure.  Assuming the volume enclosed by the mid-IR shell of the GC lobe, the formation of the GC lobe requires $(4-6)\\times10^{52} (T_e/3960\\ \\rm{K})$ ergs, depending if it is filled with thermal or nonthermal plasma.  This formation energy is consistent with the minimum (gravitational) energy in the GC lobe.  This energy estimate does not account for energy losses \\citep{v94} or the kinetic energy associated with the expansion of the GC lobe.  The kinetic energy estimate given in \\citet{b03} was about an order of magnitude larger than this formation energy, for an assumed expansion velocity of 100 \\kms. \\footnote{Note that this previous work calculated a total mass for the GC lobe an order of magnitude larger than the ionized mass found in the present work.  This relies on the analysis of molecular line data that covers Galactic latitudes less than $0\\ddeg3$, in which emission from the plane is highly confused with emission that may be from the GC lobe \\citep{s95}.}  The upper limit on the line-of-sight expansion of the GC lobe \\citep[$\\lesssim10$ \\kms; ][]{gcl_recomb} gives a dynamical time, $t_{\\rm{dyn}}\\gtrsim10$ Myr, assuming constant expansion velocity.  If the expansion of the gas has decelerated, the dynamical time will overestimate the actual formation time.  Assuming an initial expansion velocity, $v_{init}$, decelerating to 10 \\kms, the formation time is $t_{form} \\approx 2 (\\Delta r/50$\\ pc$)(40$\\ km s$^{-1}/v_{init})$ Myr.  The contiguous shape of the GC lobe and its modest height of $\\approx165$ pc \\citep[only roughly twice the \\hi\\ scale height in the nuclear disk;][]{r82} show that it has not blown out of the disk.  Without a measurement on the motion of the GC lobe away from the plane, it is difficult to constrain its fate.  The line-of-sight expansion is certainly far less than the escape velocity \\citep[$v_{esc}\\approx900$ \\kms;][]{b91}.  If the energy source that created the GC lobe continues to operate, then it may expand further or even accelerate, depending on the history of the energy input \\citep{v05}.  If the GC lobe is formed by the expansion of hot gas, the ``escape temperature'' is $T_{\\rm{esc}}=1.1\\times10^5 (v_{\\rm{esc}}/100$ \\kms$) \\rm{K} \\approx 9\\times10^6\\ \\rm{K}$  \\citep{w95}, which is approximately 1 keV.  However, no hot, X-ray gas has yet been found directly associated the GC lobe.  Finally, it is worth noting a possible constraint on the angular momentum of the ionized gas in the GC lobe.  Radio recombination line velocities show a trend for the ionized gas in the GC lobe to rotate like the disk gas \\citep{gcl_recomb}, suggesting that the gas came from the disk \\citep[perhaps entrained, as seen in extragalactic outflows;][]{s98,wa02}.  The gas velocity changes across the central 100 pc by 5 \\kms, as compared to molecular gas rotation of 200 \\kms\\ \\citep{b87,s04}.  In extragalactic outflows, conservation of angular momentum reduces the velocity gradient from galactic rotation during expansion \\citep{s98,s01}.  If this happens in the GC lobe, then the current angular momentum implies that the gas originated at a radius, $r_o \\approx 2.5\\ \\rm{km\\ s}^{-1} * 40\\ \\rm{pc}/100\\ \\rm{km\\ s}^{-1} = 1\\ \\rm{pc}$.  This value would be an underestimate of the initial radius if the gas velocities were less than $\\pm100$\\ \\kms \\citep[e.g., on ``x2'' orbits;][]{bi91} or angular momentum was lost during expansion.   \\label{conclusions} I have conducted a multiwavelength observing program to study the GC lobe and determine if it is a signature of an outflow from the GC region.  The new observations show that the GC lobe has a layered structure with concentric shells of emission of mid-IR, radio continuum, and optical and radio recombination lines.  The morphological, spectral, and physical properties of the GC lobe show that it is a coherent object in the GC region.  I calculate the energy and time required to create the GC lobe and compare its characteristics to models of nuclear outflows.  The canonical starburst model best explains the properties of the GC lobe, although the outflow may not be supersonic as in typial extragalactic starburst outflows.  The energy output by current star formation and supernovae can power the GC lobe.  If the GC lobe is indeed a nuclear outflow powered by stellar winds and supernovae, it is the closest example of such a phenomenon.  This opens the possibility of studying an outflow at extremely high physical resolution, potentially expanding our understanding of how they work.  New observations of X-rays or forbidden-line transitions should clarify whether the outflow is supersonic or may be described by a less energetic outflow, such as expected from super star clusters.  The relatively small size of the GC lobe also allows us to explore a type of outflow that is difficult to see in other galaxies.  The fact that the GC lobe would be difficult to detect in even the nearest spiral galaxies suggests that this type of outflow could be common.  The mechanism for periodic infall of mass to the central parsecs \\citep{st04} suggests that the GC region may undergo cycles of activity \\citep{a00,y04}.  It may be that the GC lobe is just the most recent example of a long series of Milky Way nuclear starburst outflows."
    },
    "0911/0911.0778_arXiv.txt": {
        "abstract": "{{\\bf ISM: kinematics and dynamics; globular clusters: general; open clusters and associations: general; galaxies: evolution; galaxies: star clusters; galaxies: stellar content}} Star clusters and their stellar populations play a significant role in the context of galaxy evolution, across space (from local to high redshift) and time (from currently forming to fossil remnants). We are now within reach of answering a number of fundamental questions that will have a significant impact on our understanding of key open issues in contemporary astrophysics, ranging from the formation, assembly and evolution of galaxies to the details of the star-formation process. Our improved understanding of the physics driving star cluster formation and evolution has led to the emergence of crucial new open questions that will most likely be tackled in a systematic way in the next decade. ",
        "introduction": "It is now widely accepted that stars do not form in isolation, at least for stellar masses above $\\sim 0.5$ M$_\\odot$. In fact, 70--90\\% of stars may form in a clustered mode (cf. Lada \\& Lada 2003). Star formation results from the fragmentation of molecular clouds, which in turn preferentially leads to star cluster formation. Over time, clusters dissolve or are destroyed by interactions with molecular clouds or tidal stripping by the gravitational field of their host galaxy. Their member stars become part of the general field stellar population. Star clusters are thus among the basic building blocks of galaxies. Star cluster populations, from young associations and open clusters to old globular clusters, are therefore powerful tracers of the formation, assembly and evolutionary history of their host galaxies.  Using our improved understanding of star cluster physics, we are now within reach of answering a number of fundamental questions in contemporary astrophysics, ranging from the formation and evolution of galaxies to the details of the process of star formation itself. These two issues are the backbone of research in modern astrophysics. They lead to new questions related to the make-up of the cluster and field stellar populations in a variety of galaxies and galaxy types, their relative formation timescales (and what this implies for overall and resolved galactic star-formation histories) and the relationships between local and high-$z$ stellar and cluster populations.  The refereed contributions in this issue of {\\ptra} cover a wide range of topics in contemporary astrophysics related to star cluster formation, evolution, destruction and environmental impact, particularly in the context of star formation on galactic scales. Lada (2010) and Clarke (2010) discuss the formation of star clusters from, respectively, an observer's and a simulator's point of view. Kalirai \\& Richer (2010) and van Loon (2010) then take detailed looks at stellar and chemical evolution of star clusters and their constituent stars, while Goodwin (2010) and Vesperini (2010) focus specifically on aspects related to the binary stellar populations affecting star cluster evolution and the makeup of the general field, and cluster dynamics, respectively. Larsen (2010) and Harris (2010) take wide-angle views of entire cluster populations, roughly split between young (Larsen 2010) and old (Harris 2010) samples, while Bruzual (2010) discusses their integrated evolution and issues relevant to the assumptions usually adopted for modelling these systems as `simple stellar populations' (SSPs).  While the review articles in this volume focus in detail on the wide-ranging fields touched by state-of-the-art star cluster research, here I address---non-exhaustively and roughly as a function of cluster age---some of the higher-level challenges limiting sustained progress, which will most likely be tackled by the research community in the next decade. ",
        "conclusions": "Although the importance of understanding star cluster physics (e.g., in mapping out the Milky Way or as SSP probes) had been recognized for decades, major progress has only become possible in recent years, both for Galactic and extragalactic cluster populations. We have seen a major recent investment in time and effort, largely thanks to significant new resources in theory, simulations and observations. These include breakthroughs in computational power (such as the recent embrace of graphics processing units by the computational astrophysics community), the maturing of {\\sl HST}-driven science (and the new generation of instruments as well as upgrades and refurbishment of the older facility instruments), deep and more precise (photometric, spectroscopic and astrometric) data for large numbers of Galactic clusters spaning a significant age range (including adaptive-optics-assisted imaging opportunities with the largest ground-based telescopes, and wide-field developments associated with intermediate-size apertures), and an explosion of astrometric data.  In addition, rapid progress has been facilitated by the coming online of observing facilities enabling access to hitherto unavailable spectral regions, both on the ground and in space. The latter include, among others, the {\\sl Spitzer Space Telescope} and the Japanese {\\sl AKARI} satellite at IR wavelengths, the {\\sl GALaxy Evolution eXplorer (GALEX)} in the ultraviolet, the Japanese {\\sl Suzaku} and the European {\\sl X-ray Multi-Mirror} mission {\\sl (XMM--Newton)}, and the {\\sl Chandra X-ray Observatory} at shorter wavelengths. These facilities will soon be complemented by the {\\sl Herschel Space Observatory} once commissioning has been completed, and in the intermediate and long term by the {\\sl James Webb Space Telescope} and {\\sl Gaia}, respectively. Ground-based precursors to the {\\sl ALMA} and upgrades of and extensions to existing (sub)millimetre and radio facilities worldwide are starting to enable access to the full submillimetre to metre-wave spectral range, with new facilities like the {\\sl Square Kilometer Array} due to come online in the next decade. Recent advances in instrumentation are driving a renaissance in the study of Galactic clusters (e.g., addressing the relationships among open, globular and young massive clusters, and narrowing down the systematic uncertainties hampering determinations of absolute cluster ages in the Milky Way), while extragalactic cluster studies are significantly aided by the development of new instrumentation supporting ever wider fields of view.  Putting these results and new developments into the broader context of galaxy evolution is the next logical step, which requires the combined efforts of theorists, observers and modellers working on a large variety of spatial scales, and spanning a very wide range of expertise. I am confident that the next decade will see a continued high level of research activity related to the physics driving star cluster formation and evolution, with perhaps an enhanced focus on their impact on the physics of their host galaxies. Given the exciting new observational facilities due to come online in the next few years, combined with major progress on computational and theoretical fronts, a bright future no doubt awaits this field!"
    },
    "0911/0911.5641_arXiv.txt": {
        "abstract": "Recently, high resolution observations with the help of the near-infrared adaptive optics integral field spectrograph SINFONI at the VLT proved the existence of massive and young nuclear star clusters in the centres of a sample of Seyfert galaxies. With the help of high resolution hydrodynamical simulations with the {\\sc PLUTO}-code, we follow the evolution of such clusters, especially focusing on mass and energy feedback from young stars. This leads to a filamentary inflow of gas on large scales (tens of parsec), whereas a turbulent and very dense disc builds up on the parsec scale. Here, we concentrate on the long-term evolution of the nuclear disc in NGC\\,1068 with the help of an effective viscous disc model, using the mass input from the large scale simulations and accounting for star formation in the disc. This two-stage modelling enables us to connect the tens of parsec scale region (observable with SINFONI) with the parsec scale environment (MIDI observations). At the current age of the nuclear star cluster, our simulations predict disc sizes of the order of 0.8 to 0.9\\,pc, gas masses of $10^6\\,M_{\\odot}$ and mass transfer rates through the inner boundary of $0.025\\,M_{\\odot}$/yr in good agreement with values derived from observations. ",
        "introduction": "Nuclear activity is an important phase in the evolution of galaxies. Whenever enough gas is fed onto the central supermassive black holes, galaxies become {\\it active} and their nuclear region lights up significantly, compared to non-active phases of their evolution. The infalling material builds up a fast rotating accretion disc, which heats up viscously. The emerging UV/optical radiation illuminates a gas and dust reservoir -- the so-called {\\it molecular torus}. Given the high opacity of the dust in this wavelength range, the torus morphology is able to geometrically unify two observed classes of active galaxies: type\\,1 objects, where the torus is viewed face-on (and all of the characteristics of the central region are visible) and type\\,2 objects (torus viewed edge-on), where most of the light from the accretion disc is absorbed and reemitted in the infrared regime. This is the essence of the {\\it Unified Scheme of Active Galactic Nuclei}. Despite its general success when probed with observational data, the nature of the obscuring gas and dust reservoir -- the torus, as well as the triggering mechanism of such phases are still a matter of active research. To assess the final stages of mass transport through the central tens of parsec region, nearby objects are needed, due to the limitations on resolution posed by currently available instruments. Therefore, Seyfert galaxies are an ideal testbed to study these processes. MIDI (MID-infrared Interferometer) observations recently resolved the innermost parts of the torus directly for the first time (e.~g.~\\cite{Jaffe_04,Tristram_07,Tristram_09,Burtscher_09}) and found evidence for a clumpy dust distribution. In these proceedings, we numerically model a scenario, which is able to describe the build-up and evolution of a nuclear gas disc or torus. Our simulations are put into context in Davies et al. (these proceedings). ",
        "conclusions": "In these proceedings, we show that evolving stars from a massive and young nuclear star cluster, as found in nearby Seyfert galaxies provide enough gas to assemble a parsec-sized nuclear gas disc. To this end, we combine a three-dimensional treatment of the mass-loss from an evolving nuclear star cluster with a simplified model for the innermost parsec scale region, where a nuclear disc builds up. As far as possible, we derive input parameters of our simulations from observations of the nearby and well-studied Seyfert\\,2 galaxy NGC\\,1068. This two-stage analysis enables us to (i) do a long term evolution study, (ii) link the tens of parsec scale region of galactic nuclei (observed with the SINFONI instrument) to the sub-parsec scales (probed by MIDI and in water maser emission) and (iii) test our model directly with a large number of observational results. At the current age of its nuclear starburst of 250\\,Myr, our simulations yield disc sizes of the order of 0.8 to 0.9\\,pc, gas masses of $10^6\\,M_{\\odot}$ and mass transfer rates of $0.025\\,M_{\\odot}/\\mathrm{yr}$ through the inner rim of the disc in good comparison with observed disc and torus properties. On basis of these comparisons, we conclude that the proposed scenario seems to be a reasonable model and shows that nuclear star formation activity and subsequent AGN activity are intimately related.    \\vspace{1cm}  {\\bf Acknowledgment:}\\\\ \\newline Part of the numerical simulations have been carried out on the SGI Altix 4700 HLRB\\,II of the Leibniz Computing Centre in Munich (Germany)."
    },
    "0911/0911.1124_arXiv.txt": {
        "abstract": "Recently, it was suggested that the gamma rays observed by the Fermi Gamma Ray Space Telescope from the direction of the galactic center could surprisingly well be described by a dark matter annihilation scenario, both in terms of their angular distribution and energy spectrum. Here, I point out that such a scenario would be in considerable tension with existing antiproton data, in particular the new PAMELA measurements. Radio data are even more constraining, thus disfavoring the dark matter hypothesis and making an astrophysical explanation of the observations much more plausible. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0911/0911.3904_arXiv.txt": {
        "abstract": " ",
        "introduction": "Under concordance cosmologies, galaxy mergers both major and minor are expected to play a central role in the evolution of virtually every galaxy. Starbursts triggered by tidal and/or hydrodynamic interactions between galaxies are a major driver of morphological transformation and, though rare, are prominent contributors to near- and mid-infrared emission in the Universe.  As just one example, starbursts driven by major mergers are thought to be an important driver in the creation of giant elliptical galaxies \\citep{ren06}.  Typical large galaxies generally behave as approximately self-regulated systems and exhibit a relatively smoothly varying star formation rate that depends on gas density \\citep[e.g.,][]{ken98a}, and is likely to slowly decline with time subsequent to their initial formation \\citep{lar78}. Starbursting states in large galaxies are rare events, likely triggered by mergers or strong tidal interactions \\citep[e.g.,][]{mih96,sam03}, and are often easily distinguished from a simple scaling-up of the global star-formation rate by manifesting as highly localized nuclear or circumnuclear starbursts \\citep[e.g.,][]{ken98b}.  Conversely, small galaxies are not expected to show constant or smoothly declining star formation rates, because they are far more susceptible to disruption by internal feedbacks and external perturbations \\citep[e.g.,][]{sti07}. This presents an observational challenge, because only within 10-15~Mpc can individual stars be resolved in order to study the stellar populations of starbursts in detail, and only within $\\approx$1~Mpc can stars be resolved to ages approaching a Hubble time to characterize the underlying stellar populations and search for fossil bursts.  However, within these distances luminous galaxies are rare, far outnumbered by dwarf galaxies.  Thus we have plentiful opportunity to study the burstiness of small galaxies in great detail, but the significance of observed variations in SFR is complicated by the expectation of large random fluctuations intrinsic to small galaxies.  One of the challenges of the studies of dwarf galaxy star-formation histories is to distinguish between systemic starbursts that are qualitatively different from steady state or quiescent star formation, and stochastic bursts which are merely the manifestation of normal variation.  Offsetting this difficulty is the advantage that in resolved stellar populations there is a large array of tools available to measure the timing, intensity, and metallicity of star-formation episodes of any age up to a Hubble time. This means that burstiness studies can be made of nearby galaxies regardless of their current morphology or gas content: even galaxies with no future have a history.  The study of burst histories in early-type galaxies has the potential to illuminate the processes of hierarchical assembly of large galaxies and morphological transformation of galaxies from disk-dominated to spheroidal \\citep[e.g.,][]{may07}.  Star-forming, gas-rich galaxies are of course easier targets for study, because of the higher light-to-mass ratios of young stars and the possibility to identify optically faint systems via HI surveys. Late-type galaxies in the nearby Universe exhibit a wide range of specific star formation rates ranging from nearly inactive to extreme starburst conditions \\citep{hun86}.  The only galaxies with neutral gas detections that do not seem to be forming stars are the lowest-luminosity examples \\citep{hun85}, the transition-type dwarfs such as the recently-discovered Leo~T dwarf \\citep{irw07}.  These dwarfs, which include among their number the Phoenix, Pisces, and Pegasus systems, are very faint (M$_B > -14$), isolated systems.  They may be showing the signs of a breathing mode of star formation \\citep{sti07}, in effect experiencing an ``anti-starburst'', or they could merely be forming stars at such a low rate that no star massive enough to ionize hydrogen has been produced within the past $\\approx$10$^7$ years \\citep{lee09}.  In this review paper I will not discuss the characteristics and causes of ongoing strong starbursts; an excellent summary of the subject can be found in \\citep[e.g.,][]{gal05}. The purpose of this paper is to review what is known about the role of bursts over the lifetimes of nearby galaxies, to discuss the lines of evidence that could be used to infer the presence of fossil bursts in resolved stellar populations, and to draw attention to some of the recent work on starburst statistics and durations in the Local Volume. For a thorough discussion of star formation in all modes and galaxy types, including the starburst phenomenon in context, see the review by \\citet{ken98b} and references therein.  Since the majority of the evidence for fossil starbursts must be gleaned from color-magnitude diagrams (CMDs), the most detailed results are by necessity confined to the Local Group, galaxies within about $\\sim$1~Mpc of the barycentre of the M31-Milky Way pair.  Within this group of $\\approx$50 galaxies, we find examples of late-type galaxies experiencing both booms and busts in their current specific star-formation rate; of galaxies that burst during their formation and never again; of galaxies that burst multiple times at inteverals of several Gyr; and galaxies for which there is no evidence of strong variations in SFR at all.  There even appears to be a galaxy which saved much its gas for 5 billion years after its first star formation, whereupon it experienced a major SFR event at a lookback time corresponding to a redshift $z$ $\\approx$ 1; there are hints that this type of star-formation history (SFH) may be commonplace among the most isolated small galaxies.  Comprehensive reviews of the SFH of Local Group galaxies are to be found in \\citet{mat98} and \\cite{tol09}, among others. ",
        "conclusions": "The burstiness of galaxies is a fundamental observable clue to understanding the physics of star formation and the processes which drive galaxy evolution. Starbursts, galaxies that are forming stars at such a high rate that the background light of all previously formed stars pales to insignifcance \\citep{san63}, are the extreme star formation environment in the Universe. Starbursts were likely the dominant mode of star formation at high redshift \\citep[][and references therein]{dre09}, but are not ``steady state'' phenomena in any sense \\citep{rie75}, and their prevalence is low among luminous galaxies in the nearby Universe, often tied to mergers and interactions \\citep[e.g.,][]{ken98b}.  Late-type galaxies in the nearby Universe exhibit a wide range of specific star formation rates ranging from nearly inactive to extreme starburst conditions \\citep{hun86,lee09}.  This is not surprising, as on theoretical grounds it is expected that galaxies with dynamical masses M$_{tot} < \\sim10^{9.5}$ M$_{\\odot}$ become unstable to their own stellar feedback from winds and supernovae \\citep[e.g.,][and references therein]{mac99,sti07}.  For small galaxies, this makes the dichotomy between ``starbursting'' and ``quiescent'' states somewhat artificial, because the natural state of dwarf galaxies is to show a continuum of burstiness \\citep{wei08} related to factors both internal (dynamical mass, angular momentum profile, gas content) and external (tidal interactions, ionizing background).  If a starburst is defined by a stellar birthrate that exceeds 2--3 times the long-term average, then dwarf galaxies are much more susceptible to experiencing bursty star formation histories than are giant galaxies.  While large galaxies most likely require disruptive events and extreme conditions in order to experience a starburst, dwarf galaxies should be able to meet the observational definition of a starburst without qualitatively changing the mode of star formation, i.e., by producing large numbers of massive clusters, concentrating the star formation to the circumnuclear region of the galaxy, or experiencing major reorganizations of gas content or morphology.  Within 10~Mpc, star formation `booms'' are rare, with only 6\\% of dwarf galaxies showing a current SFR more than 2.5 times their long term average \\citep{lee09}.  However, ``busts'' are equally rare, indicating that as long as neutral gas is present, some star formation occurs.  In general, the early conclusion of \\citet{hun85} suggesting that star formation in dwarf irregulars is ``down but not out'' has been borne out and put on a firm statistical footing by subsequent work.   Most recent work, typified by \\citet{lee09}, suggests that between 20--30\\% of star formation occurs during burst episodes. These episodes seem to last, very roughly, $\\sim$10$^{8.5}$ yr \\citep[e.g.,][]{mcq09}. Note that these figures do not distinguish between systemic bursts, which are not sustainable over the long term and may be signifcant in their production of star clusters and metals or for their promotion of morphological transformation, and stochastic bursts of the kind which are predicted by numerical models \\citep[e.g.,][]{sti07} and observed in nearby dwarfs \\citep[e.g.,][]{wei08}. It is interesting to note that while models predict a high degree of burstiness, increasing with decreasing galaxy mass, CMD analysis indicates relatively smooth SFHs for most Local Group galaxies.  Either some unmodelled factor is acting to suppress the burstiness seen in the models, or the CMD analyses are less sensitive than predicted to factors of 2 variation in SFR over time periods of a few hundred Myr.  The isolated dwarf galaxy Leo~A appears to have waited several Gyr before forming the vast majority of its stars in an event that itself spanned several Gyr. An upper limit of $\\approx$10\\% on the fraction of stars older than 8~Gyr was found by \\citet{col07}.  This makes it unique among galaxies that have been studied to the depth of the oldest main-sequence turnoff with the Hubble Space Telescope, but hints of similar behavior may be visible in other isolated galaxies \\citep[e.g.,][]{mcc06}.  If the delayed burst is a natural feature of dwarf galaxies, then it has not been captured by the models that predict steady ``breathing'' pulses of star formation.   One possibility is that cosmic reionization heated but did not evaporate the neutral gas from the potential well of the galaxy and long cooling times produced the delay \\citep[e.g.,][]{bul00}. On the other hand, if the late-blooming burst was triggered by a merger or accretion event, there is no hope at this late time, several Gyr later, of identifying the trigger.  The galaxy IC~10 appears to have much in common with similar size dwarf irregulars and is not currently forming stars far above its long-term average rate.  While it has repeatedly been referred to as the nearest starburst galaxy and an anomaly within the Local Group, the star formation history derived from HST imaging places it squarely within a continuum of similar-mass late-type galaxies, bearing a strong family resemblance to archetypal irregulars NGC~6822 and the Small Magellanic Cloud. However the burst age is only a few times 10$^7$ yr, indicating that the burst may be at a very early stage if IC~10 has similar properties to other dwarf starbursts, which appear to last for $\\sim$10 times longer \\citep[e.g.,][]{mcq09}.  IC~10 is similar in mass and metallicity to the SMC, but has failed to produce any massive star clusters during its lifetime, indirect evidence that the SMC has had the more tumultuous history and experienced a more intense mode of star formation at past epochs."
    },
    "0911/0911.1130_arXiv.txt": {
        "abstract": "Using the NICMOS coronagraph, we have obtained high-contrast 2.0~\\micron\\ imaging polarimetry and 1.1~\\micron\\ imaging of the circumstellar disk around AB Aurigae on angular scales of 0.3--3\\arcsec\\ (40--550~AU). Unlike previous observations, these data resolve the disk in both total and polarized intensity, allowing accurate measurement of the spatial variation of polarization fraction across the disk. Using these observations we investigate the apparent ``gap'' in the disk reported by Oppenheimer et al. 2008.  In polarized intensity, the NICMOS data closely reproduces the morphology seen by Oppenheimer et al., yet in total intensity we find no evidence for a gap in either our 1.1 or 2.0~\\micron\\ images. We find instead that region has lower polarization fraction, without a significant decrease in total scattered light, consistent with expectations for back-scattered light on the far side of an inclined disk. Radiative transfer models demonstrate this explanation fits the observations. Geometrical scattering effects are entirely sufficient to explain the observed morphology without any need to invoke a gap or protoplanet at that location. ",
        "introduction": "AB Aurigae is one of the most intensively studied of all Herbig Ae/Be stars, on account of its proximity, brightness, and youth (distance $d=144$ pc; visual magnitude $V=7.04$ mag; age $<3$ Myr; spectral type A0e). Steadily improving observational capabilities have yielded increasingly detailed views of its complex and dusty environment, providing many insights into the nature of circumstellar disks \\citep[e.g.][]{Grady:1999p2806,Fukagawa:2004p87,Pietu:2005p1365}.   In particular, \\citet[][hereafter Opp08]{Oppenheimer:2008p2388} recently presented high-angular resolution, high contrast imaging polarimetry of AB Aur at 1.6 ${\\mu}m$, obtained with the Lyot Project coronagraph on the AEOS 3.6 m telescope.  These observations resolved the disk in polarized scattered light as close as 40 AU (0.3$''$) to the star. In polarized light the disk is not axisymmetric, but instead shows an apparent gap or depleted region at a radius of $\\sim100$ AU. Such a gap may be created by dynamical perturbations from forming planets \\citep[e.g.][]{2003ApJ...588.1110K,Wyatt:2005A&A...440..937W,JangCondell:2009p2820}. Intriguingly, Opp08 report a faint point source within the gap, which they conjecture could be the perturbing object---though they are cautious with this identification due to its low statistical significance (2.8~$\\sigma$). The formation mechanism(s) of massive planets at large separations remain highly uncertain \\citep{2009arXiv0909.2662D,Nero:2009p2842}; any prospect for observing a planet forming \\textit{in situ} around AB Aur should be pursued to clarify this puzzle.  One challenge in interpreting the data from Opp08 is that the disk is visible only in polarized intensity, $P=\\sqrt{Q^2+U^2}$ (where $Q$ and $U$ are the usual Stokes parameters; see \\citealt{Tinbergen1996} for a review of polarization fundamentals and notation). In their total intensity image (Stokes $I$), the residual speckle halo of the stellar PSF completely drowns out the disk's fainter light. The origin of features seen only in polarized light is ambiguous: any observed spatial variation may be due either to variation in the \\textit{total} amount of scattered light, or to a change in the \\textit{polarization fraction} of that light.  The polarization induced by dust scattering depends strongly on the scattering angle (see Figure \\ref{scatteringprops}), allowing disk geometry or viewing angle to cause variations in the observed polarization which might be mistaken for intrinsic substructure within the disk.  In particular, for AB Aur the observed celestial position angle of the depleted region, $333\\pm2\\degr$ (Opp08) is precisely aligned with the disk's apparent rotation axis as inferred from CO emission line kinematics (330--333$\\degr$; \\citealp{Corder:2005p2693,Pietu:2005p1365}). Is this alignment coincidental?   The most direct way to answer this question is to obtain images with enough contrast to directly detect the disk in total intensity, and then calculate the polarization fraction, $p = P/I$. Such observations are best obtained with the Hubble Space Telescope, whose ACS and NICMOS coronagraphs both have provided sufficiently high contrast to precisely and accurately measure the polarization of disk-scattered light \\citep[][Hines and Schneider, in prep]{Hines:2000p332,2007lyot.confQ..22S,Graham:2007p495}.  In this paper we present new NICMOS coronagraphic imaging and polarimetry of AB Aur which clarifies the nature of the dark ``gap'' observed by Opp08. These data were obtained as part of a coronagraphic polarimetry survey of young stars across a range of masses and ages which will be reported more fully in future works \\citep[see][ for a brief overview]{Perrin2009Subaru}.  \\begin{figure} \\includegraphics[width=3.5in]{f1.eps} \\caption{The induced polarization for some possible grain compositions, demonstrating that polarization depends on both scattering angle and dust grain properties, particularly composition and porosity. Positive values indicate typical centrosymmetric linear polarization, while negative formal polarization denotes linear polarization oriented radially. Polarization is maximized for porous grains, and for scattering angles slightly above 90$\\degr$.  These models assume a power-law distribution of grain sizes from 0.03-200 $\\micron$ with slope -3.5, typical for YSO disks. See section 5.2 of \\citet{Pinte:2008p2580} for a discussion of dust model and computation details.  The overplotted points show the fractional polarization observed around AB Aur; the inferred scattering angles are not symmetric due to the flaring of the disk surface, estimated $\\sim10\\degr$. Of the models shown, the 60\\% porosity silicates provide the best fit. \\label{scatteringprops} } \\end{figure} ",
        "conclusions": "\\label{discussion}   \\subsection{Structure in AB Aurigae's Disk}  The pattern of polarization resolved around AB Aur indicates unambiguously that the spatial variation of polarized light there is due primarily to the geometry of scattering from the inclined disk's surface. These data do not support the hypothesis of significant clearing in the disk near PA=333\\degr.  Independent of but simultaneously with this study, a numerical investigation of protoplanet shadows in disks reached a similar conclusion, finding the observed morphology around AB Aur inconsistent with the presence of any protoplanet above 0.3 Jupiter masses (Jang-Condell and Kuchner, submitted).  While the previously-claimed gap does not appear to be present, there \\textit{is} real structure within AB Aur's disk which is visible in scattered light, including polarization. Many of the asymmetric spots seen in the 2.0~\\micron\\ polarization fraction image coincide with the spiral arms as seen in the 1.1~\\micron\\ image. For instance, the northeasternmost region of high polarization, near (1.7,~0) in the coordinate system of Figure 2, is precisely aligned with the brightest spiral arm. Overplotting or blinking these images shows several such alignments.  Perturbations in the disk's surface, from either localized changes in scale height or warps in midplane location, could change the scattering geometry at the optical depth $\\tau=1$ surface to produce the observed minor variations in polarization. For optically thick disks, surface features seen in scattered light do not necessarily correlate with conditions at the midplane \\citep{JangCondell:2007p447}, but in the case of AB Aur there is direct evidence from CO emission that the bulk of the gas deviates from pure Keplerian rotation \\citep{Lin:2006p2790}. These perturbations and/or warps might contribute to the difficulty in firmly establishing the system's inclination. But this just raises the inevitable next question: what causes those fluctuations?  Though we find no direct evidence for any disk gap due to a protoplanet, there still remains a case for the presence of a companion somewhere around AB Aur: Both the strong spiral structure seen in dust-scattered light and the non-Keplerian dynamics revealed by gas emission lines argue for the existence of a planetary-mass body perturbing the disk. Other perturbation mechanisms, such as gravitational instabilities in the disk or a stellar-mass companion, seem ruled out observationally \\citep{Pietu:2005p1365,Lin:2006p2790}.  The detection of this companion thus awaits future improvements in high-contrast imaging to reveal.    \\subsection{Interpreting Features in Imaging Polarimetry}  Differential polarimetry, as used by Opp08 and others, is a proven technique for observing circumstellar dust at high contrast from the ground with adaptive optics (AO). Such observations will become increasingly common with upcoming extreme AO systems, such as GPI \\citep{2006SPIE.6272E..18M} and SPHERE \\citep{2006Msngr.125...29B}. For many disks, these instruments will provide high contrast images only in polarized light.  We have shown here that care is required when interpreting such data. Variation of polarization with scattering angle cannot be neglected, even for low inclination targets like AB Aur. Yet AO differential polarimetry can still yield precise measurements of disk structure and dust properties, provided the degeneracy between polarization fraction and intensity can be broken. This may be possible through multiwavelength observations, which are differently sensitive to scattering geometry \\citep{Watson:2007p2642}, or in combination with independent constraints on geometry such as from CO emission.    For brighter disks, extreme AO should allow accurate measurement of polarization fraction from the ground, but for now, HST NICMOS offers the highest polarimetric precision for measurements of disk-scattered light.                 \\label{concl_sect}"
    },
    "0911/0911.3302.txt": {
        "abstract": "We present a statistical analysis of a sample of 20 strong lensing clusters drawn from the Local Cluster Substructure Survey (LoCuSS), based on high resolution \\emph{Hubble Space Telescope (HST)} imaging of the cluster cores and follow-up spectroscopic observations using the Keck-I telescope.  We use detailed parameterized models of the mass distribution in the cluster cores, to measure the total cluster mass and fraction of that mass associated with substructures within $R\\le250\\kpc$.  These measurements are compared with the distribution of baryons in the cores, as traced by the old stellar populations and the X-ray emitting intracluster medium.  Our main results include: (i) the distribution of Einstein radii is log-normal, with a peak and $1\\sigma$ width of $\\langle\\log_{10}\\theta_E(z=2)\\rangle=1.16\\pm0.28$; (ii) we detect an X-ray/lensing mass discrepancy of $\\langle M_{SL}/M_X\\rangle=1.3$ at $3\\sigma$ significance -- clusters with larger substructure fractions displaying greater mass discrepancies, and thus greater departures from hydrostatic equilibrium; (iii) cluster substructure fraction is also correlated with the slope of the gas density profile on small scales, implying a connection between cluster-cluster mergers and gas cooling.  Overall our results are consistent with the view that cluster-cluster mergers play a prominent role in shaping the properties of cluster cores, in particular causing departures from hydrostatic equilibrium, and possibly disturbing cool cores.  Our results do not support recent claims that large Einstein radius clusters present a challenge to the CDM paradigm. ",
        "introduction": "The evolution of galaxy clusters with cosmic time is an important cosmological probe, as it traces the gravitational growth of dark matter on large scales. In particular, the cluster mass function can be directly tested against cosmological models, as it is related to fundamental parameters such as the matter density $\\Omega_m$, the normalisation of the power spectrum $\\sigma_8$ \\citep[e.g.][]{Schuecker,Smith03}, and the dark energy equation of state parameter $w$.  Commonly used probes of the mass of clusters at $z\\sim0.1-0.5$ include the $K$-band luminosity \\citep{lin}, X-ray luminosity and temperature \\citep[e.g.][]{Vikhlinin}, Sunayev-Zeldovich effect \\citep[e.g.][]{Nagai,Bonamente}, cluster kinematics \\citep[e.g.][]{Blindert}, and gravitational lensing \\citep[e.g.][]{Bardeau,Smith05,Okabe}.  Systematic errors in these various mass probes can be calibrated by an inter-comparison of results from the various methods.  The pursuit of the most robust calibrations based on low redshift cluster samples is essential to achieve reliable results from future studies at higher redshift, and to control systematic uncertanties in cosmological experiments \\citep{Albrecht06}.  Gravitational lensing plays a central role in this effort, as it does not rely on assumptions about the symmetry or equilibrium properties of the cluster mass distribution, although parametrized lensed model implicitly assume, for instance, elliptical symmetry.  Indeed, the highest quality measurements of the dark matter mass and substructure in low redshift clusters have been obtained through the combination of strong and weak gravitational lensing.  The identification of background galaxies forming strong lensing arcs in cluster cores is a direct measurement of the enclosed mass within the Einstein radius $\\theta_E$, providing an accurate normalisation of cluster mass models both in the core, and extending to larger radii \\citep[e.g.][]{Kneib03}.  The main caveats on lensing studies are that the lensing signal is sensitive to the mass distribution projected along the line of sight through the cluster, that the commonly used parametric models must by design assume a parametric form of the mass distribution, and that gravitational lensing may prove to be impractical in samples of more distant ($z>1$) cluster samples that are starting to be discovered \\citep[e.g.][]{Rosati1, Rosati2}.  Early joint lensing/X-ray cluster studies identified large discrepancies between total cluster masses obtained from the two methods \\citep[e.g.][]{Jordi95}.  These discrepancies were subsequently eliminated, albeit within quite large uncertainties, in cool core clusters provided that a two phase gas model was used (\\citealt{Allen98}; see also \\citealt{Smail97}).  At a similar time, \\emph{HST} data started to became available for some clusters, allowing more precise strong lensing models to be constructed \\citep[e.g.][]{Kneib96,Tyson98}.  More recent strong lensing studies of cluster mass distributions have generally continued to concentrate on detailed studies of spectacular individual strong lensing clusters based on deep multi-filter \\emph{HST} observations \\citep[e.g.][]{Broadhurst05,Limousin07,Richard09}.  The interpretation of results from such studies, e.g.\\ size of Einstein radii and shape of density profile, in the context of the general cluster population is inevitably problematic.  In contrast, \\citet[][hereafter Sm05]{Smith05} studied 10 X-ray selected clusters ($L_{X,0.2-2.4\\keV}>4.1\\times10^{44}\\ergs$) at $z\\simeq0.2$ as a first step towards building large statistical samples of strong lensing clusters.  Sm05 combined strong and weak lensing constraints from moderate depth \\emph{HST} observations with X-ray spectro-imaging from \\emph{Chandra}, and near-infrared photometry of cluster galaxies from UKIRT.  The main results included the first mass-observable scaling relations to employ lensing-based mass measurements for a well-defined sample; the only previous example employed a compilation of 6 clusters from the literature \\citep{Hjorth98}.  The main limiting factor on Sm05's results were the small sample size, with just 5 of the 10 clusters containing spectroscopically confirmed strongly lensed galaxies.  The Local Cluster Substructure Survey (LoCuSS, PI: G.\\ P.\\ Smith) is an all-sky systematic survey of 165 X-ray luminous clusters at $0.15<z<0.30$ selected from the \\emph{ROSAT} All-sky Survey catalogs \\citep{Ebeling98,Ebeling00,Boehringer04}.  In addition to seeking to improve the statistical precision by enlarging the sample size of studies such as Sm05 and \\citet{Bardeau}, LoCuSS aims to incorporate new constraints on cluster baryons, most notably from observations of the Sunyaev-Zeldovich Effect \\citep{Carlstrom02} to deliver a definitive local calibration of mass-observable scaling relations that will be useful for cluster cosmology, interpretation of high-redshift cluster samples, and probing the physics of gas heating and cooling in merging clusters at low redshift.  More generally, LoCuSS aims to constrain the scatter in the baryonic properties of the local cluster population and to correlate the scatter with the recent hierarchical assembly history of the clusters, the latter being constrained by the lensing-based mass maps \\citep{Taylor}.  A wide variety of studies are therefore underway, including wide-field weak-lensing analysis with Suprime-Cam on Subaru \\citep{Okabe}, the first lensing-based mass-$Y_{SZ}$ relation \\citep{Marrone09}, and multiwavelength studies supported by space-based (\\emph{HST}, \\textit{Herschel}, \\textit{Chandra}, GALEX, \\textit{Spitzer}) as well as ground-based (Keck, VLT, Gemini, MMT, CTIO, KPNO, Palomar) facilities \\citep[e.g.][]{Zhang,Haines09b,Haines09a,Sanderson09a,Smith09b}.  This paper presents a four-fold increase in the number of clusters with spectroscopically confirmed strong lensing clusters on Sm05. Such constraints remain critical within the overall context of the studies outlined above because strong-lensing offers precise constraints on the mass distribution in cluster cores ($R\\sim100\\kpc$) which for example provide an invaluable constraint on small scales when trying to measure the shape of cluster density profiles in conjunction with wide-field weak-lensing data from Subaru.  We compare the details of the central mass distributions (integrated mass and substructure fraction) obtained from the strong-lensing constraints with near-infrared and X-ray probes of the cluster baryons.  We describe the data in \\S\\ref{data}, and present the strong-lensing analysis and models in \\S\\ref{sl}.  The main results on the mass and structure of the cluster cores are presented in \\S\\ref{results} and summarized in \\S\\ref{conc}.  Throughout the paper, we use magnitudes quoted in the AB system, and a standard $\\Lambda$-CDM model with $\\Omega_m=0.3$, $\\Omega_\\Lambda=0.7$, and $H_0=70\\kms\\Mpc^{-1}$, whenever necessary.  In this cosmology, 1\\arcsec\\ is equivalent to $3.3\\kpc$ at $z=0.2$.  We adopt the definition $e=1-b/a$ of the ellipticity, where $a$ and $b$ are the semi-major and semi-minor axis of the ellipse, respectively. ",
        "conclusions": ""
    },
    "0911/0911.3135_arXiv.txt": {
        "abstract": "In 2007, a joint IERS/IVS Working Group has been established to consider practical issues of creating the next ICRF generation, ICRF-2. The goal of the WG is to seek after ways to improve the existing ICRF. In this study we investigate a possibility of ICRF improvement by means of using combined ICRF catalogue instead of a catalogue computed in a single analysis centre, even though using most advanced models and software. In this work, we present a new version of Pulkovo combined catalogue of radio source positions computed using the method proposed in \\cite{SokMal07}. Radio source catalogues that were submitted in 2007 in the framework of the WG activity were used as input for mutual comparison and combination. Four combined catalogues have been calculated: Two first of them provide a stochastic improvement of the ICRF, and last two of them allow us to account also for systematic errors in the current ICRF version. ",
        "introduction": "The celestial reference frame (CRF), as realized by a set of coordinates for selected celestial objects, is widely used for numerous astronomy, navigation, time and other measurements. The CRF accuracy and stability are all-important for successful solutions of all these tasks. After publishing of the first VLBI radio source catalogue (RSCs), attempts were made to improve the accuracy of radio-band CRF by means of constructing combined catalogues, as it was customary for optical astronomy, where fundamental catalogues served as an international standard for astrometry and other measurements on the sky. Different methods were used to obtain a combined RSC, e.g. \\cite{Walter89a,Walter89b,Yatskiv90,Kurianova93} Also, up to 1995, IERS (International Earth Rotation Service, now International Earth Rotation and Reference Systems Service) used Derived combined RSC for maintenance of the IERS Celestial Reference Frame. In 2007, a joint IERS/IVS Working Group has been established to consider practical issues of creating the next ICRF generation, ICRF-2. The goal of the WG is to seek after ways to improve the existing ICRF. Large experience accrued by optical astrometry over centuries shows that combining catalogues of the star positions leads to better random and systematic accuracy than individual catalogues. In this work, we present a new version of Pulkovo combined catalogue of radio source positions computed using the method proposed in \\cite{SokMal07}. Radio source catalogues that were submitted in 2007 in the framework of the WG activity were used as input for mutual comparison and combination. Four combined catalogues have been calculated: \\begin{enumerate} \\item The first two catalogues (the one has been calculated using all input catalogues and the second one using only catalogues obtained with CALC/SOLVE software) provide a stochastic improvement of the ICRF \\item and the last two of them one allow us to account also for systematic errors in the current ICRF version. \\end{enumerate} All computations have been done for the set of 194 ICRF defining sources included in all input catalogues. ",
        "conclusions": ""
    },
    "0911/0911.4465_arXiv.txt": {
        "abstract": "We carry out systematic and high-resolution studies of dynamo action in a shell model for magnetohydrodynamic (MHD) turbulence over wide ranges of the magnetic Prandtl number $Pr_{\\rm M}$ and the magnetic Reynolds number $Re_{\\rm M}$. Our study suggests that it is natural to think of dynamo onset as a nonequilibrium, first-order phase transition between two different turbulent, but statistically steady, states. The ratio of the magnetic and kinetic energies is a convenient order parameter for this transition. By using this order parameter, we obtain the stability diagram (or nonequilibrium phase diagram) for dynamo formation in our MHD shell model in the $(Pr^{-1}_{\\rm M}, Re_{\\rm M})$ plane. The dynamo boundary, which separates dynamo and no-dynamo regions, appears to have a fractal character. We obtain hysteretic behavior of the order parameter across this boundary and suggestions of nucleation-type phenomena. ",
        "introduction": "}  The elucidation of dynamo action is a problem of central importance in nonlinear dynamics because it has implications for a variety of physical systems.  Dynamo instabilities, which amplify weak magnetic fields in a turbulent conducting fluid, are believed to be the principal mechanism for the generation of magnetic fields in celestial bodies and in the interstellar medium~\\cite{arnab,rudiger,goedbloed,biskamp,mkvrev,njpspecial,schekochihin}, and in liquid-metal systems~\\cite{roberts,fauve,riga,karl,lathrop,pinton} studied in laboratories.  In these situations the kinematic viscosity $\\nu$ and the magnetic diffusivity $\\eta$ can differ by several orders of magnitude, so the magnetic Prandtl number $Pr_{\\rm M}\\equiv\\nu/\\eta$ can either be very small or very large; e.g., $Pr_{\\rm M}\\simeq 10^{-2}$ at the base of the Sun's convection zone, $Pr_{\\rm M}\\simeq 10^{-5}$ in the liquid-sodium system, and $Pr_{\\rm M}\\simeq 10^{14}$ in the interstellar medium. This Prandtl number is related to the Reynolds number $Re=UL/\\nu$ and the magnetic Reynolds number $Re_{\\rm M}=UL/\\eta$ that characterize the conducting fluid; here $L$ and $U$ are typical length and velocity scales in the flow; clearly $Pr_{\\rm M}=Re_{\\rm M}/Re$.  Two dissipative scales play an important role here; they are the Kolmogorov scale $\\ell_d$ [$\\sim \\nu^{3/4}$ at the level of Kolmogorov 1941 (K41) phenomenology~\\cite{K41}] and the magnetic-resistive scale $\\ell_d^M$ [$\\sim \\eta^{3/4}$ in K41].  For large Prandtl numbers, i.e., $Pr_{\\rm M}\\gg 1$, $\\ell_d^M \\ll \\ell_d$ so the magnetic field grows predominantly in the dissipation range of the fluid till it is strong enough to affect the dynamics of the fluid through the Lorentz force.  This behavior is a characteristic of a small-scale turbulent dynamo, in which dynamo action is driven by a smooth, dissipative-scale velocity field.  In the initial stage of growth, called the kinematic stage of the dynamo, the magnetic field is not large enough to act back on the velocity field.  Dynamo action can be obtained for values of $Re_{\\rm M}$ that are large enough to overcome Joule dissipation; and the dynamo-threshold value $Re_{\\rm Mb}$ decreases as $Pr_{\\rm M}$ increases~\\cite{boldyrev,ponty}.  $Pr_{\\rm M} \\simeq 10^{-5}$ in liquid-metal flows~\\cite{riga,pinton,kenjeres} so they lie in the small-Prandtl-number region, $Pr_{\\rm M} \\ll 1$, for which the growth of the magnetic energy occurs initially in the {\\it inertial scales} of fluid turbulence, because $\\ell_d \\ll \\ell_d^M$; here the velocity field is not smooth and the local strain rate is not uniform: At the K41 level the turnover velocity of an eddy of size $\\ell$ is $v(\\ell) \\sim \\ell^{1/3}$, so the rate of shearing $(v(\\ell)/\\ell)\\sim \\ell^{-2/3}$.  Direct numerical simulations (DNS) are playing an increasingly important role in developing an understanding of such dynamo action.  Most DNS studies of MHD turbulence~\\cite{frisch,chou,ponty,minnini} have been restricted, because of computational constraints, either to low resolutions or to the case $Pr_{\\rm M}=1$; small-scale dynamos, with $Pr_{\\rm M} \\gg 1$ have also been studied via DNS~\\cite{schekochihin}.  However, given the large range spanned by $Pr_{\\rm M}$ in the physical settings mentioned above, some recent DNS studies of the MHD equations have started to explore the $Pr_{\\rm M}$ dependence of dynamo action; the range of $Pr_{\\rm M}$ covered by such pure DNS studies~\\cite{ponty,brandenburgdns} is quite modest ($10^{-2}\\lesssim Pr_{\\rm M} \\lesssim 10$). To explore the dynamo boundary in the $(Pr^{-1}_{\\rm M}, Re_{\\rm M})$ plane over a large range of $Pr_{\\rm M}$, one recent study~\\cite{ponty} has used a combination of numerical methods, some of which require small-scale models like large-eddy simulations (LES) or lagrangian-averaged MHD (LAMHD), and others, like a pseudospectral DNS, in which the only approximations are the finite number of collocation points and the finite step used in time marching; yet another DNS study~\\cite{iskakov} has introduced hyperviscosity of order $8$ to study the low $Pr_{\\rm M}$ regime; by using this combination of methods these studies has been able to cover the range $10^{-2}\\lesssim Pr_{\\rm M} \\lesssim 10^{3}$ and to obtain the boundary between dynamo and no-dynamo regions but with fairly large error bars.  We have carried out extensive, high-resolution, numerical studies that have been designed to explore in detail the boundary between the dynamo and no-dynamo regimes in the $(Pr^{-1}_{\\rm M}, Re_{\\rm M})$ plane in a shell model for three-dimensional MHD~\\cite{basu,frick,brandenburgshell,giuliani}.  This shell model allows us to explore a much larger range of  $Pr_{\\rm M}$ than is possible if we use the MHD equations. Although our study uses a simple shell model, it has the virtue that it can explore the boundary between dynamo and no-dynamo regions in great detail without resorting to the modelling of small spatial scales.  Shell-model studies of dynamo action have also been attempted in Refs.~\\cite{frick,plunian,mkvshell,benzi} but these have concentrated on aspects of the dynamo problem that are different from those we consider here.  Our study suggests that it is natural to think of the boundary between dynamo and no-dynamo regimes in the $(Pr^{-1}_{\\rm M}, Re_{\\rm M})$ plane as a first-order phase boundary that is the locus of first-order, nonequilibrium phase transitions from one nonequilibrium statistical steady state (NESS) to another. The first NESS is a turbulent, but statistically steady, conducting fluid in which the magnetic energy is negligibly small compared to the kinetic energy; the second NESS is also a statistically steady turbulent state but one in which the magnetic energy is comparable to the kinetic energy.  Indeed, the ratio of the magnetic and fluid energies $E_b/E_u$ turns out to be a convenient order parameter for this nonequilibrium phase transition since it vanishes in the no-dynamo phase and assumes a finite, nonzero value in the dynamo state. The other, intriguing result of our study is that the boundary between these phases is very intricate and might well have a fractal character; this provides an appealing explanation for the large error bars in earlier attempts to determine this boundary~\\cite{ponty,plunian}.  The analogy with first-order transitions that we have outlined above is not superficial. As in any first-order transition we find that our order parameter shows hysteretic behavior~\\cite{mrao} as we scan through the dynamo boundary by changing the forcing term at a nonzero rate. We also find some evidence of nucleation-type phenomena: the closer we are to the dynamo boundary, the longer it takes for a significant magnetic field to nucleate and thus lead to dynamo action. We compare our results with earlier studies such as Ref.~\\cite{pontyetal07}, which have suggested that dynamo action occurs because of a subcritical bifurcation.  The remaining part of this paper is organised as follows: In Sec. \\ref{models} we describe the shell model for MHD~\\cite{basu,frick,brandenburgshell} and the numerical method we employ. Sec. \\ref{results} is devoted to our results and Sec. \\ref{conclusion} contains a concluding discussion. ",
        "conclusions": "}  We have presented a detailed study of dynamo action in a shell model of turbulence~\\cite{basu,frick,brandenburgshell}. Our study has been designed to explore the nature of the boundary between dynamo and no-dynamo regimes in the $(Pr^{-1}_{\\rm M}, Re_{\\rm M})$ plane over a much wider range of $Pr_{\\rm M}$ than has been attempted in earlier numerical studies. The dynamo boundary emerges as a first-order nonequilibrium phase boundary between one turbulent, nonequilibrium statistical steady state (NESS) and another~\\cite{footnote}. This point of view is implicit in earlier work, e.g., in studies of the Kazantsev dynamo~\\cite{kazantsev} or in studies that view dynamo generation as a subcritical bifurcation~\\cite{busse77,pontyetal07,kuangetal08}.  One of these studies~\\cite{pontyetal07} has remarked that when dynamo action `` ... is obtained in a fully turbulent system, where fluctuations are of the same order of magnitude as the mean flow ... the traditional concept of amplitude equation may be ill-defined and one may have to generalize the notion of ``subcritical transition\" for turbulent flows ...\". We believe that the natural generalization is the nonequilibrium, first-order transition we suggest above. We have explored the explicit consequences of such a view in far greater detail than has been attempted hitherto. In particular, the ratio $E_b/E_u$ is a convenient order parameter for this nonequilibrium phase transition; it shows hysteresis across the dynamo boundary like order parameters at any first-order transition; and nucleation-type phenomena also seem to be associated with dynamo formation.  Last, and perhaps most interesting, we find that the dynamo boundary seems to have a fractal character; this provides a natural explanation for the large error bars in earlier attempts to determine this boundary~\\cite{schekochihin,ponty,plunian}. Furthermore, this fractal-type boundary might well be the root cause of magnetic-field reversals discussed, e.g., in Refs.~\\cite{benzi,pintonreversal}.  It is important to check, of course, that our shell-model results carry over to the MHD equations.  This requires large-scale DNS that might well be beyond present-day computing capabilities if we want to explore issues like the possible fractal nature of the dynamo boundary. However, analogs of the hysteretic behavior we mention above have been obtained in DNS studies of the MHD equations~\\cite{busse77,pontyetal07,kuangetal08}; hysteresis has also been seen in a numerical simulation that includes turbulent convection~\\cite{sim+bus09}. In some of these studies hysteretic behavior is obtained by changing the viscosity of the magnetic Prandtl number. We have obtained hysteresis by changing the forcing; this change of forcing might be easier to effect in experiments than a change of the viscosity or magnetic diffusivity.  To the best of our knowledge, earlier studies have not noted the increase in the dynamo-formation time $\\tau_c$ as the dynamo boundary is approached from the dynamo side. We have suggested that this is akin to the increase in the time required to form a critical nucleus as we approach a first-order boundary~\\cite{oxtoby}. It would be interesting to see if such an increase of $\\tau_c$ can be obtained in DNS studies of dynamo formation with the MHD equations. It is worth noting here that some DNS studies~\\cite{ponty} have suggested that simulation times comparable to the diffusion time scale $\\tau_\\eta$ are required to confirm dynamo formation; by contrast our shell-model study yields dynamo action on a much shorter time $\\tau_c$, which increases as we approach the dynamo boundary. Perhaps the large simulation times required for dynamo action in full MHD simulation might have arisen because these simulations have been carried out in the vicinity of the dynamo boundary.  To settle completely whether the dynamo boundary is of fractal-type, very long simulations might be required to make sure that the apparent fractal nature is not an artifact of long-lived metastable states. To make sure that our calculations do not suffer from such an artifact, we have carried out very long runs for representative points in the region of the dynamo boundary in Fig.~\\ref{fig:stability}; we have found that these long runs do not change our results. Furthermore, it is useful to check whether, instead of one dynamo boundary, there is a sequence of transitions, with more and more complicated temporal behaviors for the order parameter, as has been seen in the turbulence-induced melting of a nonequilibrium vortex crystal~\\cite{crystalmelt}. We have not found any conclusive evidence for this but, in the vicinity of the dynamo boundary, the order parameter can oscillate for fairly long times (see, e.g., the green full curve in Fig.~\\ref{fig:tseries}). To decide conclusively whether these oscillations characterize a new nonequilibrium oscillating state, different from the simple dynamo and no-dynamo NESSs we have mentioned, requires extensive numerical studies that lie beyond the scope of this paper.  In equilibrium statistical mechanics different ensembles are equivalent; in particular, we may determine a first-order phase boundary by using either the canonical or the grand-canonical ensemble. However, such an equivalence of ensembles does not apply to transitions between different nonequilibrium statistical steady states (NESSs); examples may be found in driven diffusive systems~\\cite{acharyya} or in the turbulence-induced melting of a nonequilibrium vortex crystal~\\cite{crystalmelt}.  Given that the dynamo boundary separates two turbulent NESSs, we might expect that this boundary might depend on precisely how the system is forced.  Evidence for this exists already: For example, the dynamo boundary depends on whether a stochastic external force is used~\\cite{schekochihin} or whether a Taylor-Green force is used~\\cite{ponty}; furthermore, this boundary is different if the fluid is helical~\\cite{brandenburgdns}, as in most astrophysical dynamos.  We hope our study of dynamo formation in a shell model for MHD will stimulate both DNS and experimental studies designed to explore the first-order nature of the dynamo transition."
    },
    "0911/0911.3590_arXiv.txt": {
        "abstract": "This work investigates the short wavelength stability of the magnetopause between a rapidly-rotating, supersonic, dense accretion disc and a slowly-rotating low-density magnetosphere of a magnetized star. The magnetopause is  a strong shear layer with rapid changes in the azimuthal velocity, the density, and the magnetic field over a short radial distance and thus the Kelvin-Helmholtz (KH) instability may be important. The plasma  dynamics is treated using non-relativistic, compressible (isentropic)  magnetohydrodynamics. It is necessary to include the displacement current  in order that plasma wave velocities remain less than the speed of light. We focus mainly on the case of a star with an aligned dipole magnetic field so that the magnetic field is axial in the disc midplane  and perpendicular to the disc flow velocity. However, we also give results for cases where the magnetic field is at an arbitrary angle to the flow velocity. For the aligned dipole case the magnetopause is most  unstable for KH waves propagating  in the azimuthal direction perpendicular to the magnetic field which tends to stabilize waves propagating parallel to it. The wave phase velocity is that of the disc matter. A quasi-linear theory of the saturation of the instability leads to a wavenumber ($k$) power spectrum $\\propto k^{-1}$ of the density and temperature fluctuations of the magnetopause, and it gives the mass accretion and angular momentum inflow rates across the magnetopause. For self-consistent conditions this mass accretion rate will be equal to the disc accretion rate at large distances from the magnetopause. ",
        "introduction": "This work investigates the short wavelength stability of the interface or magnetopause between a rapidly rotating accretion disc and the slowly-rotating, low-density magnetosphere of a magnetized star. The nature of the magnetopause is sketched in Figure 1. The rotating  disc matter is ``held off'' by the star's  magnetosphere where the magnetic field is strong and the density is small.  The disc matter rotates at approximately the Keplerian velocity which is typically much larger than the velocity of the magnetospheric plasma which corotates  with the angular velocity of the star. Thus the interface involves a strong shear layer as  sketched in the bottom part of Figure 1. Understanding the instabilities of the magnetopause is important for understanding both the transport of matter and angular momentum towards the star {\\it and} the temporal variability of the sources (van der Klis 2006).    The magnetohydrodynamic (MHD) stability of configurations such as in Figure 1 was investigated earlier by Li \\& Narayan (2003) assuming incompressible flow and perturbations independent of $z$, but with no restrictions on the azimuthal wavelength. They found both long-wavelength (i.e., $\\lesssim r$) Rayleigh-Taylor (RT) and short wavelength ($\\ll r$) Kelvin-Helmholtz (KH) instabilities for different conditions of the shear layer. Earlier, Burnard, Lea, \\& Arons (1983) studied the short wavelength KH instability of the magnetopause of a spherically accreting rotating magnetized star. Recently Tsang \\& Lai (2009) have studied the long wavelength RT stability of the sharp interface including the compressibility of the media. The MHD stability of the magnetopause for cases where the shear layer has appreciable radial width was studied by Lovelace \\& Romanova ((2007) and Lovelace, Turner, \\& Romanova (2009)   for compressible, non-barotropic perturbations independent of $z$ but no restrictions on the  azimuthal wavelength. They found a resonant long-wavelength  Rossby Wave Instability (RWI; Lovelace et al. 1999) which may contribute to the observed twin kilo-Hertz quasi-periodic oscillations of low-mass X-ray binaries (van der Klis 2006).     Here we consider a thin magnetopause and short wavelengths where the   KH instability is expected to be important. There is a vast literature on the magnetized KH instability in different astrophysical and space applications,  to the magnetopause of rotating planets in the solar wind  (e.g., Miura \\& Pritchett 1982; Roy-Choudhury \\& Lovelace 1986; Faganello, Califano, \\& Pegoraro 2008), to the stability of astrophysical jets (e.g., Hardee 2007; Osmanov et al. 2008),  to the stability of interfaces of molecular/atomic clouds in the interstellar medium (e.g., Hunter, Whitaker, \\& Lovelace 1998), and to the interface of an unmagnetized dense plasma blob falling through the strong magnetic field of a neutron star (Arons \\& Lea 1980).   Section 2 of the paper gives the basic equations where the fluid motion is assumed non-relativistic but the displacement current is retained in Maxwell's equations in order to keep the wave speeds less than the light speed. This section also describes the assumed equilibrium, the waves in the disc plasma, and the MHD waves in the magnetospheric plasma. The disc flow speed is much larger than the disc sound speed. Section 3   obtains a dispersion relation for the Kelvin-Helmholtz modes and develops an approximates solution for cases where the magnetopsheric density ($\\rho_2$)  is much less than the disc density ($\\rho_1$). Section 4 discusses the non-linear saturation of the KH modes and develops    a quasi-linear theory for the mass accretion rate across the magnetopause. Section 5 discusses briefly the case where the magnetic field is at an arbitrary angle relative to flow velocity in the disc. Section 5 gives the conclusions of this work.   \\begin{figure} \\centering \\includegraphics[scale=0.55]{KHfig1} \\caption{Sketch of the inner region of a disc around a rotating magnetized star suggested by MHD simulations of Romanova et al. (2008). Inside of $r_m$ the plasma flows along the magnetic field lines to the surface of the star in what is termed a funnel flow (denoted ff in the figure). The bottom part of the figure shows the midplane profiles of the the density $\\rho$, magnetic field $B_z$, and azimuthal velocity $v_\\phi$. } \\end{figure} ",
        "conclusions": "This work investigated the short wavelength ($\\lambda \\ll 2\\pi r_m$) stability of the magnetopause at radius $r_m$  between a rapidly-rotating, supersonic, dense ($\\rho_2$) accretion disc and a slowly-rotating low-density magnetosphere ($\\rho_1 \\gg \\rho_1$) of a magnetized star. The magnetopause is  a strong shear layer with  rapid changes in the azimuthal velocity, the density, and the magnetic field over a short radial distance ($\\Delta r_m \\ll r_m$), and thus the Kelvin-Helmholtz (KH) instability may be important. This work has focused on the case of a star with an aligned dipole magnetic field so that the magnetic field is axial in the disc midplane  and perpendicular to the disc flow velocity. For the aligned dipole case the magnetopause is most  unstable for KH waves propagating  in the azimuthal direction perpendicular to the magnetic field. Propagation not perpendicular to the magnetic field changes the magnetic field which gives a stabilizing effect. The growth rate of the instability is    $\\omega_i = (\\rho_2/\\rho_1)^{1/2} k_x v_x$, where $v_x$ is the velocity of the disc plasma relative to that of the magnetospheric plasma which corotates with the star. The wave phase velocity is that of the disc matter.    We discussed the non-linear saturation of the instability which we argued occurs for a mode  of wavelength $\\lambda=2\\pi/k$ when the displacement of the interface between the disc and magnetospheric plasmas is of the order of $\\lambda/2\\pi$. From this we developed a quasi-linear model which led to a wavenumber  power spectrum $\\propto k^{-1}$ of the density and temperature fluctuations of the magnetopause. The quasi-linear model gave  the mass accretion and angular momentum inflow rates across the magnetopause. The mass accretion rate (per unit area of the magnetopause) was found to be $2(\\rho_2/\\rho_1)^{1/2} \\rho_1 v_x \\ln(k_{\\rm max}/k_{\\rm min})$, where $k_{\\rm max, min}$ are the maximum and minimum wavenumbers discussed in \\S 4. For self-consistent conditions this mass accretion rate will be equal to the disc accretion rate at large distances from the magnetopause.  We also considered the case where the magnetic field is not perpendicular to the flow velocity but tilted by an angle $\\theta$ relative to the $z-$axis, where $\\theta=0$ corresponds to the aligned rotator case treated earlier.   We found that the maximum growth rate and the associated wave phase velocity decrease monotonically with $\\theta$ and that both approach zero as $\\theta \\rightarrow 90^\\circ$. That the growth rate goes to zero is a result of the stabilizing effect of the magnetic field. Thus  an orthogonal rotator where the star's magnetic moment $\\rvecmu$ is perpendicular to the star's rotation axis $\\rvecOmega_*$ is stable to the KH mode."
    },
    "0911/0911.1076_arXiv.txt": {
        "abstract": "We report the discovery with XMM-Newton of correlated spectral and timing behavior in the ultraluminous X-ray source (ULX) NGC 5408 X-1. An $\\approx 100$ ksec pointing with XMM/Newton obtained in January, 2008 reveals a strong 10 mHz QPO in the $> 1$ keV flux, as well as flat-topped, band limited noise breaking to a power law. The energy spectrum is again dominated by two components, a $0.16$ keV thermal disk and a power-law with an index of $\\approx 2.5$.  These new measurements, combined with results from our previous January 2006 pointing in which we first detected QPOs, show for the first time in a ULX a pattern of spectral and temporal correlations strongly analogous to that seen in Galactic black hole sources, but at much higher X-ray luminosity and longer characteristic time-scales.  We find that the QPO frequency is proportional to the inferred disk flux, while the QPO and broad-band noise amplitude (root mean squared, rms) are inversely proportional to the disk flux. Assuming that QPO frequency scales inversely with black hole mass at a given power-law spectral index we derive mass estimates using the observed QPO frequency - spectral index relations from five stellar-mass black hole systems with dynamical mass constraints.  The results from all sources are consistent with a mass range for NGC 5408 X-1 from 1000 - 9000 $M_{\\odot}$.  We argue that these are conservative limits, and a more likely range is from 2000 - 5000 $M_{\\odot}$.  Moreover, the recent relation from Gierlinski et al. that relates black hole mass to the strength of variability at high frequencies (above the break in the power spectrum) is also indicative of such a high mass for NGC 5408 X-1.  Importantly, none of the above estimates appears consistent with a black hole mass less than $\\approx 1000$ $M_{\\odot}$ for NGC 5408 X-1. We argue that these new findings strongly support the conclusion that NGC 5408 X-1 harbors an intermediate mass black hole. ",
        "introduction": "The nature of the bright X-ray sources found in nearby galaxies, the ultraluminous X-ray sources (ULXs), remains a major astrophysical puzzle.  The fundamental conundrum is that some of these objects have X-ray luminosities uncomfortably high to be stellar-mass black holes (BH) without violating standard Eddington limit arguments. Three different solutions have been proposed for the luminosity problem. 1) The objects are intermediate-mass BHs (Colbert \\& Mushotzky 1999). 2) They are stellar-mass BHs with, in some cases, substantial beaming of their X-ray radiation (King et al. 2001), or, 3) they are stellar-mass BHs emitting above their Eddington limit (Begelman 2006).  It is possible that some ULXs appear ultraluminous because of a combination of all three factors (moderately higher mass, mild beaming and mild super-Eddington emission). It may also be that these objects make up an inhomogeneous population, comprised of both a sub-sample of intermediate-mass BHs and moderately beamed stellar BHs (for recent reviews see Fabbiano \\& White 2006; Miller \\& Colbert 2004).  Due to their extragalactic nature the study of ULX counterparts at other wavebands has been difficult, and has so far precluded the use of the familiar methods of dynamical astronomy to weigh them. However, recent work has resulted in mass measurements for some Local Group stellar BHs, including IC 10 X-1 (Prestwich et al. 2007; Silverman \\& Filippenko 2008), and M33 X-7 (Orosz et al. 2007; Liu et al. 2008), but these are not ULXs.  A substantial body of work now supports the idea that timing properties can be used to constrain the masses of BHs. For example, McHardy et al. (2006) have shown that when luminosity variations are taken into account, both stellar-mass and supermassive BHs populate a ``variability plane'' linking their broad band variability time-scales, luminosities and masses (see also K{\\\"o}rding et al. 2007). Casella et al (2008) applied this relation in order to estimate the masses of two ULXs, M82 X-1 and NGC 5408 X-1, concluding that the masses of NGC 5408 X-1, and M82 X-1 lie in the range (in solar units) $115 < M_{bh} < 1300$, and $95 < M_{bh} < 1260$, respectively. However, in this case the scaling was not direct in that Casella et al. (2008) had to estimate the relevant break time-scale from the observed QPO frequencies in M82 X-1 and NGC 5408 X-1. They did this by applying a scaling relation between the two quantities derived only from observations of stellar-mass BHs.  Recently, Gierlinski, Nikolajuk \\& Czerny (2008) have shown that the strength of high frequency variability--parameterized as the root mean squared (rms) amplitude integrated above the break in the power spectrum--scales approximately linearly with BH mass.  They show that a broad correlation exists when comparing stellar and supermassive BHs, but the relation is not tight enough to predict the mass of a stellar system by comparison, for example, with another stellar-mass BH of known mass.  It is now firmly established that Galactic BHs accreting at high rates--often classified as the Intermediate State (IS) or Steep Power-Law state (SPL)--show strong correlations between their spectral and temporal parameters.  In these states a significant fraction of the X-ray luminosity is in a power-law component whose spectral index correlates well with the frequency of a QPO (so-called Type C QPOs) that is also commonly detected in such states (see Sobczak et al. 2000; Vignarca et al. 2003; Kalemci et al. 2005; Remillard et al. 2002; Casella et al. 2004). Other spectral parameters also correlate with QPO frequency, such as the disk flux.  Recent work by Shaposhnikov \\& Titarchuk (2007; 2009; hereafter ST07 and ST09) has demonstrated rather convincingly that the QPO frequency - spectral index scaling for Type C QPOs can be used as an empirical BH mass estimator. Previous efforts have been made to use this scaling argument to constrain the masses of ULXs (Fiorito \\& Titarchuk 2004; Dewangan, Titarchuk \\& Griffiths 2006; Strohmayer et al. 2007, hereafter Paper 1), however, in these previous studies, there was no direct evidence to indicate that the QPO properties seen in the ULX sources actually correlates with the spectral parameters as in stellar-mass systems.  In this paper we present the results of new timing and spectral measurements that show for the first time that NGC 5408 X-1 behaves very much like a Galactic stellar-mass BH system with the exception that its characteristic X-ray time-scales are $\\sim 100$ times longer, and its luminosity is greater by a roughly similar factor.  We argue that these new findings provide strong evidence that this system contains a few thousand solar mass BH. ",
        "conclusions": "Our new XMM-Newton observations have revealed for the first time that NGC 5408 X-1 exhibits correlated X-ray timing and spectral properties quite analogous to those exhibited by Galactic stellar-mass BHs in the ``very high'' or ``steep power-law'' state.  Its longer observed time-scales (QPO and power spectral break frequencies) at higher luminosity compared to Galactic stellar-mass BHs can be understood if the mass of NGC 5408 X-1 is a few thousand solar masses.  We have arrived at this conclusion by several independent arguments. Given our present understanding of the timing properties of BHs at all mass scales, it seems to us hard to escape the conclusion that NGC 5408 X-1 is an intermediate mass BH.  We have again found evidence for a pair of sharp, closely spaced QPOs in NGC 5408 X-1 with a frequency ratio consistent with 4:3.  We thus conclude that this is an intrinsic feature of the system and not some artifact or coincidence associated with having only a single observation of the source.  As we noted in Paper 1, it is possible that this results from detection of both a Type C QPO (the stronger feature), and a Type B QPO (Casella et al. 2004).  If the source transitions from one to another, we do not have sufficient signal to noise ratio data to ``watch'' such transitions occur. Rather, we simply see average detections of each QPO.  We have argued that the observed variations in QPO frequency, accompanied by spectral changes consistent with behavior seen in Galactic systems, allows the mass of NGC 5408 X-1 to be estimated by scaling the observed QPO frequencies to match those observed in stellar-mass systems of known mass. Essentially the same method has now been used on a good number of Galactic systems and seems empirically robust (ST09).  These arguments give mass estimates for NGC 5408 X-1 in a broad range from $1 - 9 \\times 10^{3}$ $M_{\\odot}$, however, we find that none of the QPO scaling estimates are consistent with masses below $\\approx 1000$ $M_{\\odot}$. Independent mass constraints based on the amplitude of high frequency variability also appear consistent with this range. Additional observations would help to map out the timing - spectral correlations for NGC 5408 X-1 more clearly, and thus allow more rigorous estimates.  As noted in Paper 1, the measured disk temperature in NGC 5408 X-1 is also consistent with a few thousand $M_{\\odot}$ BH if the theoretical accretion disk scaling of $kT_{disk} \\propto M^{-1/4}$ holds. Importantly, none of these new estimates appears consistent with a mass as low as even a few hundred solar masses.  We think these new findings provide strong evidence that NGC 5408 X-1 harbors an intermediate mass BH."
    },
    "0911/0911.2243_arXiv.txt": {
        "abstract": "It has been suggested that rotating black holes could serve as particle colliders with arbitrarily high center-of-mass energy. Astrophysical limitations on the maximal spin, back-reaction effects and sensitivity to the initial conditions impose severe limits on the likelihood of such collisions. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0911/0911.5713_arXiv.txt": {
        "abstract": "Previously unobservable mirror asymmetry of the solar magnetic field -- a key ingredient of the dynamo mechanism which is believed to drive the 11-year activity cycle -- has now been measured. This was achieved through systematic monitoring of solar active regions carried out for more than 20 years at observatories in Mees, Huairou, and Mitaka. In this paper we report on detailed analysis of vector magnetic field data, obtained at Huairou Solar Observing Station in China. Electric current helicity (the product of current and magnetic field component in the same direction) was estimated from the data and a latitude-time plot of solar helicity during the last two solar cycles has been produced. We find that like sunspots helicity patterns propagate equatorwards but unlike sunspot polarity helicity in each solar hemisphere does not change sign from cycle to cycle - confirming the theory. There are, however, two significant time-latitudinal domains in each cycle when the sign does briefly invert. Our findings shed new light on stellar and planetary dynamos and has yet to be included in the theory. ",
        "introduction": "``Whirling storms in the earth's atmosphere, whether cyclones or tornadoes follow a well-known law which is said to have no exceptions: the direction of whirl in the Northern hemisphere is left-handed or anti-clockwise, while in the Southern hemisphere it is right-handed or clockwise'' (Hale et al. 1919). Eruptive phenomena in the solar atmosphere are much more powerful than typhoons in the Earth's atmosphere, and strong helical magnetic fields in regions where eruptive phenomena occur store huge amounts of magnetic energy. American astronomer G.E. Hale was the first to discover the magnetic fields of sunspots and to analyse their helical magnetic configurations. He recognized a similarity with the above-cited polarity rule now known as Hale's polarity rule for sunspots: the sign of magnetic field is anti-symmetric over the solar equator and changes with every 11-year cycle. Recent observations of helical structures in the magnetic fields of solar active regions at the photospheric level encourage us to generalise Hale' law for the time-latitudinal distribution of the magnetic field helicity, which plays a key role in the mechanism of magnetic field generation operating in the solar interior.  In 1955, Parker suggested a dynamo mechanism in the solar interior based on the combined action of differential rotation and cyclonic convective vortices (Paker 1955) as a viable way to generate magnetic fields capable of driving the activity cycle. We are able to quantify the differential rotation from the motion of large-scale magnetic fields at the solar surface and from helioseismology in the solar convection zone. However, because of the opacity of the solar atmosphere, knowledge of the action of convective vortices can only be obtained from available observations of helical magnetic fields. According to mean field dynamo theory, the electromotive force ${\\bf E}$ averaged over convective eddies has a component parallel to the magnetic field, ${\\bf E}=\\alpha {\\bf B}+...$, where the pseudoscalar $\\alpha$ is related to kinetic and electric current helicities $\\alpha \\sim\\, <{\\bf V}\\cdot{\\bf {{\\bf {\\nabla}} \\times \\,}} {\\bf V}> + <{\\bf B}\\cdot{\\bf {{\\bf {\\nabla}} \\times \\,}} {\\bf B} >$. The determination of the kinetic helicity in the solar atmosphere is difficult, while the twist of magnetic fields, can be estimated from photospheric vector magnetograms of solar active regions (Seehafer 1990; Abramenko et al. 1996; Zhang 2006). The formation of twisted magnetic fields inside the Sun is a fundamental topic for solar dynamo models (Kleeorin et al. 2003; Choudhuri et al. 2004; Sokoloff et al. 2006; Zhang et al. 2006); however, for alternative interpretations see (Tanaka 1991; Zhang 1995). The achievement presented in this Report is analysis of vector magnetograms of solar active regions, obtained as a result of systematic monitoring for about 20 years at several solar observatories, such as at Mees in Hawaii, Huairou in Beijing, Mitaka in Tokyo, and also Marshall Space Flight Center in Huntsville. The information on mean electric current helicity and twist comes from vector magnetograms in the form $H_{{\\rm C}z}=({\\bf B}\\cdot{\\bf {{\\bf {\\nabla}} \\times \\,}} {\\bf B})_{z}$ or $\\alpha_{{\\rm ff}}=({\\bf B}\\cdot{\\bf {{\\bf {\\nabla}} \\times \\,}} {\\bf B})_{z}/B_{z}^{2}$, where the former is a part of the current helicity density, and the latter keeps the same sign as the former one, being a factor of the force-free magnetic field. These two quantities reflect different helical characteristics of magnetic fields (see, e.g., {Sokoloff et al. 2008; Kuzanyan et al. 2000 Zhang et al. 2002 for details) and, therefore, we consider both of them.  Obtaining values of these helical parameters averaged over an active region we may expect that they reflect the handedness of the solar magnetic field. We found that active regions in the northern (southern) solar hemisphere possess statistically mainly left (right) handedness (Seehafer 1990; Pevtsov et al. 1994, 1995; Bao \\& Zhang 1998) which is persistent over the noisy nature of the signal. Although we admit that the comparison of observational results obtained at different instruments and their interpretation performed by various research groups is not a straightforward issue (Pevtsov et al. 1994; Bao \\& Zhang, 1998; Hagino \\& Sakurai, 2004), on average, about 57-66\\% of solar active regions follow this, so called, the hemispheric helicity rule (Pevtsov et al. 1995; Bao \\& Zhang 1998; Hagino \\& Sakurai 2004; Bao, Ai \\& Zhang 2000; Pevtsov, Canfield \\& Latushko 2001; Abramenko et al. 1997). ",
        "conclusions": "Let us summarize the results. The mirror asymmetry of magnetic fields at the solar surface is related with the magnetic field generation inside the Sun. It evolves with solar cycle, in particular, it has domains of the ``wrong'' helicity sign at the beginning and end of the wings (see Fig.~2). A quantitative description of the phenomenon remains for further theoretical studies.  Our results bring up a new type of characteristic of hydromagnetic dynamos and pose a challenge for the theory of stellar and planetary magnetic fields."
    },
    "0911/0911.2419_arXiv.txt": {
        "abstract": "{We extend our theoretical computations for low-mass stars to intermediate-mass and massive stars, for which few databases exist in the literature. Evolutionary tracks and isochrones are computed for initial masses $2.50 - 20 M_{\\odot}$ for a grid of 37 chemical compositions with  metal content $Z$ between 0.0001 and 0.070 and helium content $Y$ between 0.23 and 0.40 to enable users to obtain isochrones for ages as young as about $10^7$ years  and to simulate stellar populations with different helium-to-metal enrichment laws. The Padova stellar evolution code is identical to that used in the first paper of this series. Synthetic TP-AGB models allow stellar tracks and isochrones to be extended until the end of the thermal pulses along the AGB. We provide software tools for the bidimensional interpolation (in $Y$ and $Z$) of the isochrones.  We present tracks for scaled-solar abundances and the corresponding isochrones from very old ages down to about $10^7$ years. This lower limit depends on chemical composition. The extension of the blue loops and the instability strip of Cepheid stars are compared and the Cepheid mass-discrepancy is discussed. The location of red supergiants in the H-R diagram is in good agreement with the evolutionary tracks for masses from $10$ to $20 M_{\\odot}$. Tracks and isochrones are available in tabular form for the adopted grid of chemical compositions in the extended plane $Z-Y$ in three photometric systems. An interactive web interface allows users to obtain isochrones of any chemical composition inside the provided $Z-Y$ range and also to simulate stellar populations with different $Y(Z)$ helium-to-metal enrichment laws.  } ",
        "introduction": "In general, the helium content of stellar populations is poorly known. For super-solar metallicities, the helium content is unconstrained, and has usually been derived from simple reasoning, based on naive chemical evolution models of galaxies, rather than observational grounds.  The most popular assumption in evolutionary models of stars is that $Y$ increases linearly with $Z$. This assumption is inaccurate for old globular clusters, some of which contain multiple stellar populations with evidence of significant helium enrichment (Piotto et al. 2007; Piotto 2009). Even the solar initial content of helium has been debated, since the values derived from helioseismology depend on the assumed efficiency of helium diffusion (Guzik et al. 2005, 2006; Basu \\& Antia 2008), whereas on the other hand helioseismology has been challenged by the debate about the oxygen abundance (Montalban et al. 2004; Scott et al. 2009). For these reasons it is likely that the helium content of stellar evolutionary models will in the future be revised many times. An alternative approach, followed in this paper, is that of computing stellar models for a significant range of the $Z$ - $Y$ plane, so that interpolations in the grid allow us to consider many possible $Y(Z)$ relations.  A large amount of theoretical investigation has been devoted to the evolution of low-mass stars populating old galactic globular clusters, since rich stellar clusters provide us with the opportunity to test the results of stellar evolution theories. The increasing amount of data for young populous clusters in both of the Magellanic Clouds encourages us to pay special attention to the evolutionary behaviour of  more massive stars presently evolving in those clusters. However, not so much theoretical computation has been addressed to the implementation of existing stellar model databases for more massive stars, as done for low masses.  The Padova and Geneva groups originally compiled grids of stellar models with overshooting up to $120 M_{\\odot}$ for several chemical compositions (Bressan et al. 1993; Fagotto et al. 1994a,b, and relative isochrones in Bertelli et al. 1994; Schaller et al. 1992; Schaerer et al. 1993a,b; Charbonnel et al. 1993; Meynet et al. 1994). Girardi et al. (2000) computed a new set of low- and intermediate-mass stellar models with updated opacities and equation of state. Bono et al. (2000) presented intermediate-mass standard models with different helium and metal content ($ 3 \\le M \\le 15 M_{\\odot}$), and Pietrinferni et al. (2004, 2006) database makes available stellar models and isochrones for scaled-solar  and $\\alpha$-enhanced metal distributions in the mass range between $0.5$ and $10 M_{\\odot}$ with and without overshooting.  Following the first release of the new stellar evolution models of the Padova database for low-mass stars in a large region of the $Z$-$Y$ plane by Bertelli et al. (2008) (hereinafter Paper I), in this paper we present models for stars from $2.5$ to $20 M_{\\odot}$ with the same range of chemical compositions as in Paper I. Among the astrophysical problems of interest in this range of mass, we discuss whether observations of Cepheids in the Magellanic Clouds (MCs) and the Galaxy can be reproduced by these new models. Massive red supergiants in the Galaxy and in MCs with newly derived physical parameters by Levesque et al. (2005, 2006) are compared with our models from $10$ to $20 M_{\\odot}$, evaluating the adequacy of stellar models in this mass range.  A neglected process in our massive stellar models is stellar rotation. Significant advances have been made in this field  (e.g., Heger \\& Langer 2000; Meynet \\& Maeder 2000), which demonstrated that the evolutionary paths of rapidly rotating massive stars can differ from those of non-rotating stars. A book by Maeder (2009) on the physics, formation, and evolution of rotating stars presents a very detailed description of the effects of rotation on stellar evolution. Rapid rotation can trigger strong mixing inside massive stars, extend the core hydrogen-burning lifetime, significantly increase the luminosity, and change the chemical composition at the stellar surface with time. Regardless of the success of rotating stellar models in explaining a variety of stellar phenomena, rotational mixing is still a matter of debate (de Mink et al., 2009).  Our choice of neglecting stellar rotation in this set of models has two main motivations, namely:  1) not all stars begin their life with the same rotational velocity;  2) stars have rotational axes that are not all oriented in the same direction.  These two stochastic effects introduce some dispersion in the distribution of the stars in the HR disgram, in contradiction with the deterministic character of stellar isochrones.  The ``non-rotating isochrones'' represent the starting point for any further study.  The present paper is organized as follows. In Sect. 2, we summarize the main input physics and in Sect. 3,  TP-AGB models are described. In  Sect. 4, stellar tracks are presented with a description of tables indicating changes in surface chemical composition after the first and the second dredge-up. Isochrones and relative tables are described in Sect. 5, in addition to the interpolation scheme allowing to obtain isochrones and to simulate stellar populations in a large region of the $Z-Y$ plane. Cepheids and the problem of the mass discrepancy are discussed in Sect. 6. Section 7 deals with massive red supergiants  in the Galaxy and the Magellanic Clouds, and Sect. 8 presents the concluding remarks.   \\section {Input physics and coverage of the $Z-Y$ plane}  The stellar evolution code adopted in this work is the same as in Bertelli et al. (2008, hereinafter Paper I ). The  masses described in this paper are between $2.5$ and $20 M_{\\odot}$,  to ensure that we have a complete stellar model database from $0.15$ to $20 M_{\\odot}$. In the following, we  summarize the input physics described in Paper I.   \\subsection{Initial masses and chemical compositions} \\label{sec_chemic}  \\begin{table}[!ht] \\caption{Combinations of $Z$ and $Y$ of the computed tracks} \\label{tab_comp} \\smallskip \\begin{center} {\\small \\begin{tabular}{cccccc} \\hline \\noalign{\\smallskip} Z & Y1 & Y2 & Y3 & Y4 & Y5 \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} 0.0001 & 0.23 & 0.26 & 0.30 &      & 0.40      \\\\ 0.0004 & 0.23 & 0.26 & 0.30 &      & 0.40      \\\\ 0.001  & 0.23 & 0.26 & 0.30 &      & 0.40      \\\\ 0.002  & 0.23 & 0.26 & 0.30 &      & 0.40      \\\\ 0.004  & 0.23 & 0.26 & 0.30 &      & 0.40      \\\\ 0.008  & 0.23 & 0.26 & 0.30 & 0.34 & 0.40      \\\\ 0.017  & 0.23 & 0.26 & 0.30 & 0.34 & 0.40      \\\\ 0.040  &      & 0.26 & 0.30 & 0.34 & 0.40      \\\\ 0.070  &      &      & 0.30 & 0.34 & 0.40      \\\\ \\noalign{\\smallskip} \\hline \\end{tabular} } \\end{center} \\end{table} The first release of the new evolutionary models was presented in  Paper I for masses from $0.15 $  to $2.5 M_{\\odot}$ (Bertelli et al. 2008). The stellar models were computed for initial masses $2.5, 3, 3.5, 4, 4.5, 5, 6, 7, 8, 10, 12, 15$, and $20 M_{\\odot}$, and, as in Paper I, from the ZAMS to the end of helium burning. The initial chemical composition is in the range $0.0001 \\le Z \\le 0.070$ for the metal content, and for the helium content in the range $0.23 \\le Y \\le 0.40$ as shown in Table 1. In addition, tracks were computed for masses between $4.5$ and $20 M_{\\odot}$ with $Z=0.030$  and $Y=0.26, 0.30, 0.34,$ and $0.40$ to derive more reliable interpolated isochrones for more massive stars.    \\subsection{The choice of scaled-solar abundances }  For each value of $Z$, the fractions of different metals follow a scaled solar distribution, as compiled by Grevesse \\& Noels (1993), and adopted in the OPAL opacity tables. The ratio of abundances of different isotopes is assumed to be identical to Anders \\& Grevesse (1989).  In this regard, we note that both the solar composition, and what is referred to as the alpha-enhanced distribution of metals have changed, by factors of up to 0.3 dex, in the past few years. Moreover, there is also an increasing perception that neither scaled solar, nor alpha-enhanced compositions are universal: they apply only to particular regions of spiral galaxies that are similar to our own. For instance, there are documented cases of slightly alpha-depleted composition in some areas of the LMC (e.g., Pomp\\'eia et al. 2008). For nearby dwarf galaxies, the modern observational picture is that their populations are in general slightly alpha-depleted for [Fe/H]$>-1.0$, becoming alpha-enhanced only at very low metallicities (see Fig. 11 in Tolstoy et al. 2009, and references therein). On the high-metallicity end, although it is generally believed that the most massive ellipticals are metal-rich and alpha-enhanced, there is no guarantee that they follow the same distribution of alpha-elements as  nearby halo stars.  The above-mentioned observations of dwarf galaxies demonstrate that low-metallicity models with scaled solar compositions are useful and, in many cases, appropriate, to the study of nearby galaxies. Since the observational view of  a ``typical alpha-enhanced composition'' continues to change, we prefer to be conservative and begin the modelling from the most well studied case, the scaled solar one, which is reasonably well-known, apart from the solar oxygen abundance. We will explore a few possible cases of alpha-enhancement in future papers.This approach is consistent with our choice of scaled solar abundances and the rules introduced by Salaris et al. (1993) to approximate low-Z alpha-enhanced isochrones using scaled solar ones (which can be applied to the Milky Way halo, metal-poor globular clusters, and the metal-poorest stars in dwarf galaxies).  As far as the use of  ``old solar abundances'' is concerned (Grevesse \\& Noels, 1993), we note the following: \\begin{itemize} \\item Some more recent studies (e.g., Caffau et al. 2008, and references therein) confirm  the high O abundances typical of older compilations of solar abundances, rather than the low values claimed by Asplund et al.(2004). Solar models computed with low heavy-element abundances also strongly disagree with the constraints from helioseismology (see Basu \\& Antia 2008, and references therein).  \\item We computed a few sequences of stellar tracks using several solar compositions. The difference between the models computed at the same mass, metallicity, and mixing length parameter was found to be negligible  (as discussed in Sect. 4.3.1) . \\end{itemize}  \\subsection{Opacities} \\label{sec_opac}  The radiative opacities for scaled solar mixtures are assumed to be those of the OPAL group (Iglesias \\& Rogers 1996) for temperatures higher than $\\log T = 4$, and the molecular opacities are those of Alexander \\& Ferguson (1994) for $\\log T < 4.0$ as in Salasnich et al. (2000). For very high temperatures ($\\log T \\ge 8.7$), we use the opacities of Weiss et al. (1990).  The conductive opacities of electron-degenerate matter are taken from Itoh et al. (1983). For the interpolation within the opacity tables grids in both papers, we used the two-dimensional bi-rational cubic damped-spline algorithm (see Schlattl \\& Weiss 1998; Salasnich et al. 2000, and Weiss \\& Schlattl 2000). In general, evolutionary models of low-mass stars differ very little by changing the opacity interpolation scheme. The model predictions for massive stars ($M \\ge 5 M_{\\odot}$) are far more sensitive to the interpolation algorithm as discussed in Salasnich et al. (2000)  \\subsection{Equation of state} \\label{sec_eos}  The equation of state (EOS) for temperatures higher than $10^7$~K is that of a fully-ionized gas, including electron degeneracy in the way described by Kippenhahn et al.\\ (1967). The effect of Coulomb interactions between the gas particles at high densities is taken into account as described in Girardi et al.\\ (1996). For temperatures lower than $10^7$~K, the detailed ``MHD'' EOS of Mihalas et al.\\ (1990, and references therein) is adopted. We note that the MHD EOS is critical only for stellar models of mass lower than 0.7~ $M_{\\odot}$ during their main sequence evolution.   \\subsection{Reaction rates and neutrino losses} \\label{sec_rates}  The reaction rates are taken from the compilation of Caughlan \\& Fowler (1988, hereinafter CF88), apart from $^{17}{\\rm O}({\\rm p},\\alpha)^{14}{\\rm N}$ and $^{17}{\\rm O}({\\rm p},\\gamma)^{18}{\\rm F}$, for which we use the determinations by Landr\\'e et al.\\ (1990). The uncertain $^{12}$C($\\alpha,\\gamma$)$^{16}$O rate was set to be 1.7 times the values given by Caughlan \\& Fowler (1988), as indicated by the study of Weaver \\& Woosley (1993) of the nucleosynthesis by massive stars.  The electron screening factors for all reactions are those from Graboske et al.\\ (1973). The abundances of the various elements are evaluated with the aid of a semi-implicit extrapolation scheme, as described in Marigo et al. (2001). The energy losses by neutrinos are from Haft et al. (1994).  In Sect. 4.3.2 we discuss the effects on stellar models of measurements of the $^{14}{\\rm N}({\\rm p},\\gamma)^{15}{\\rm O}$ reaction rate by the LUNA experiment (Formicola et al., 2004).  \\subsection{Convection} \\label{sec_conv}  The most important physical process for mixing in stars is convection, but at present convection remains a main  source of uncertainty in stellar models computations. In our stellar models, convection is treated with the Schwarzschild criterion for stability. Usually  energy transport in the outer convection zone of stars is described according to the mixing-length theory (MLT) of B\\\"ohm-Vitense (1958). We adopt the same value as in Paper I for the MLT parameter ($\\alpha =1.68$), calibrated by means of the solar model.  \\subsubsection{ Overshoot} The extension of the convective regions in stellar models takes into account overshooting from the borders of both core and envelope convective zones (Bressan et al. 1981; Alongi et al. 1991; Girardi et al. 2000). In the following, we adopt the formulation by Bressan et al. (1981) in which the boundary of the convective core is defined to be the layer where the velocity (rather than the acceleration) of convective elements vanishes. This non-local treatment of  convection requires the use of a free parameter related to  the mean free path $l$ of convective elements defined to be $l=\\Lambda_c H_p$, ($H_p$ being the pressure scale-height).   The choice of this parameter determines the extent of the overshooting {\\em across} the border of the classical core (determined by the Schwarzschild criterion). Other authors define the extent of the overshooting zone to be at the distance $d=\\Lambda_c H_p$ {\\em above} the border of the convective core (Schwarzschild criterion). The value 0.5 used in the Padova models is almost equivalent to the 0.2 value adopted by other groups (Meynet et al, 1994; Pietrinferni et al. 2004; Yi et al. 2001). The non-equivalence of the parameter used to describe the extension of convective overshooting by different groups has been a recurrent source of misunderstanding in the literature.  We adopt the same prescription as in Paper I (Bertelli et al. 2008) and in Girardi et al. (2000) for the parameter $\\Lambda_c$  for H- and He-burning stages. Overshooting at the lower boundary of convective envelopes is also considered. For $M>2.5 M_{\\odot}$, a value of $\\Lambda_e=0.7$ was assumed, as in Bertelli et al.\\ (1994) and Girardi et al. (2000).  \\subsubsection{Hydrogen semiconvection} During the core H-burning phase of massive stars, radiation pressure and electron scattering opacity produce a large convective core surrounded by an H-rich region, which is potentially unstable to convection if the original gradient in chemical abundance remains, but stable if suitable mixing is allowed to occur. This region of the star is assumed in theoretical models to undergo sufficient mixing until the condition of neutrality is restored, to transport negligible energy flux. The Schwarzschild condition produces smoother chemical profiles and, in some cases,  a fully intermediate convective layer. Similar instability also occurs during the early shell H-burning phase. Taking into account mass loss from massive stars, the extension of semiconvection and/or intermediate full convection at the top of the H-burning shell is much smaller than in constant-mass stellar models. The effects of H-semiconvection on the evolution of massive stars were summarized by Chiosi \\& Maeder (1986).  \\subsubsection{ Helium semiconvection} As He-burning proceeds in the convective core of low-mass stars during the early stages of the horizontal branch (HB), the size of the convective core increases. Once the central value of helium falls below $Y_c =0.7$, the temperature gradient reaches a local minimum, so that continued overshoot is no longer able to restore the neutrality condition at the border of the core. The core divides into an inner convective core and an outer convective shell. As further helium is captured by the convective shell, the outer shell tends to become stable, leaving behind a region of varying composition in conditions of neutrality. This zone is called semiconvective. Starting from the centre and going outwards, the matter in the radiatively stable region above the formal convective core is mixed layer by layer until neutrality is achieved. This picture holds during most of the central He-burning phase. The extension of the semiconvective region varies with stellar mass, being important in  low- and intermediate-mass stars up to about $5 M_\\odot$, and negligible in more massive stars. We followed the scheme of Castellani et al. (1985), as described in Bressan et al. (1993).  \\subsection {Mass loss } \\label{sec_rgbagb}  Mass loss has a crucial impact on the evolution of massive stars, affecting evolutionary tracks, lifetimes, and surface abundances. The role of radiation pressure in driving mass loss from massive stars has been firmly accepted by theoretical astrophysics since the pioneering studies by Lucy \\& Solomon (1970), and Castor et al. (1975). A relevant review on the evolution of massive stars with mass loss was given by Chiosi \\& Maeder (1986), while Kudritzki \\& Puls (2000) exhaustively describe winds from hot stars. The metallicity dependence of mass loss rates is usually included using the formula:  $\\dot M(Z) = \\dot M(Z_{\\odot}) (Z/Z_{\\odot})^{\\alpha}$,  \\noindent where the exponent $\\alpha$ varies between  values of $0.5 - 0.6$ (Kudritzki \\& Puls 2000; Kudritzki 2002), and $0.7 - 0.86$ (Vink et al. 2001; Vink \\& de Koter 2005) for O-type and WR stars. A comparison between mass loss prescriptions and observed mass loss rates, and the observational dependence of the rate on the metal content of massive stars atmospheres was presented by Mokiem et al. (2007). They found a metallicity dependence consistent with previous values and stated that those $\\dot M(Z)$ relations hold for stars more luminous than $\\sim 10^{5.2} L_{\\odot}$.  In this paper, we adopt the mass loss rates from stellar winds driven by radiation pressure according to de Jager et al.(1988) for all evolutionary stages starting from the main sequence and  from about $\\sim 6 - 7 $ to $20 M_{\\odot}$ (the upper mass limit to this evolutionary track release). The rates incorporate the dependence on metallicity given by Kudritzki et al. (1989)  $\\dot M_Z = (Z/Z_{\\odot})^{0.5} \\dot M_{Z_{\\odot}}$.  In a future paper dealing with the evolution of massive stars, for the mass-loss rates   we will adopt relationships based on more recent observations and/or determinations. ",
        "conclusions": "\\label{sec_remarks}   Observations of some globular clusters in the Galaxy and the LMC provide evidence of multiple populations and/or  significant variations in the helium content, in contrast to traditional scenarios in which globular clusters are simple stellar populations of uniform age and chemical composition. The current database of Padova tracks allows the computation of isochrone sets covering a wide range of ages and enables users to analyse stellar populations (star clusters or galaxies) with different helium enrichment laws, given the extended $Z-Y$ region of the evolutionary models.  The scaled solar database of models and isochrones, presented in Paper I, has been extended in mass to include stellar evolutionary computations from $2.5$ to $20 M_{\\odot}$ for the same large region of the $Z-Y$ plane. An important update of this database is the extension of stellar models and isochrones until the end of the TP-AGB phase by means of synthetic models, using the same relations as in Marigo \\& Girardi (2007). Revised bolometric corrections by Girardi et al. (2008) are employed to transform Padova isochrones into UBVRIJHK and ACS photometric systems.  In the range of more massive stars presented in this Paper II, our models can be tested in at least two interesting astrophysical problems: 1) Cepheids and the mass-discrepancy problem, 2) massive red supergiants and the problems of their effective temperature (Massey 2003).  Considering the observations of Cepheids in the Galaxy and the Magellanic Clouds (Tamman et al. 2003; Sandage et al. 2004; Sandage et al. 2009), we show that the observed instability strip is narrower than the theoretical one, and that the slope of the lower boundary of observed stars coincides with the slope of the faint envelope connecting each of the loops of the theoretical tracks for the chemical composition of the MCs. Because of the small number of Cepheids in open clusters and associations of the Milky Way, we cannot draw the lower boundary of the observed strip as we did for the MCs, but in this case, the loops of the theoretical models can also indicate the location of MW Cepheids in the HR-diagram.  The long-standing open question about the discrepancy between the pulsation and the evolutionary mass of Cepheid stars is considered, by determining the evolutionary mass of Cepheids in MW open clusters (Tamman et al. 2003) from our models (isochrones) and the pulsation mass from the mass dependent PLC relation for fundamental pulsators. The evolutionary mass is found to be higher than the pulsation mass for stars of mass lower than $\\sim 6 M_{\\odot}$, while for stars more massive than $10M_{\\odot}$ the pulsation mass is higher than the evolutionary one. If we consider extra convective mixing as a plausible solution of the Cepheid mass discrepancy, the parameter that determines the efficiency of the extra mixing should decrease with increasing stellar mass.  Red supergiants in both the Galaxy and the Magellanic Clouds were too cool and luminous compared with evolutionary tracks (Massey 2003; Massey \\& Olsen 2003) until when Levesque et al. (2005, 2006) derived a new effective temperature scale and new bolometric corrections, producing far closer agreement with stellar evolutionary tracks. Our tracks from $10$ to $20 M_{\\odot}$ for the appropriate chemical composition can reproduce the location in the H-R diagram of red supergiants in the MW and the MCs as given by Levesque et al. (2005, 2006).   A web site has been dedicated to make available the entire database to the scientific community.  In the section {\\bf YZVAR} of the {\\bf static databases} at {\\bf http://stev.oapd.inaf.it} users can find:  \\begin{itemize} \\item {data files with the information relative to each evolutionary track}, \\item {isochrone files from Paper I with the chemical composition of the computed grids in the Z-Y plane, and age interval in the range $\\sim 10.15 \\ge$  log age/yr $ \\ge \\sim 9.00$. } \\end{itemize} In the section {\\bf YZVAR form} of the {\\bf interactive services} a web interface allows users to obtain interpolated isochrones for the entire range of ages, from old (log age/yr $\\sim 10.15$) to young ages (log age/yr $ \\sim 7.00$) and any chemical composition in the provided range of Y and Z. The initial and final values of isochrone ages can vary a little more  or less, depending on the chemical composition. To compute the isochrones we adopt the interpolation scheme described in Sect. 5.1.  We point out that older isochrones (Paper I) are available in both the static database and the interactive service, while younger isochrones only interactively. All isochrones (old and young) are available through the interactive service for the chosen photometric system (see  Sect. 5.3).  In the future, a web interactive interface will provide data of stellar populations in terms of SFR, IMF, chemical composition, and mass loss during the RGB phase."
    },
    "0911/0911.1717.txt": {
        "abstract": "{ We present a spectroscopic campaign to follow-up red colour-selected candidate massive galaxies in two high redshift proto-clusters surrounding radio galaxies.  We observed a total of 57 galaxies in the field of \\uss0943\\ ($z=2.93$) and 33 in the field of \\pks1138\\ ($z=2.16$) with a mix of optical and near-infrared multi-object spectroscopy.  We confirm two red galaxies in the field of \\pks1138\\ at the redshift of the radio galaxy. Based on an analysis of their spectral energy distributions, and their derived star formation rates from the H$\\alpha$ and 24\\,$\\mu$m flux, one object belongs to the class of dust-obscured star-forming red galaxies, while the other is evolved with little ongoing star formation.  This result represents the first red and mainly passively evolving galaxy to be confirmed as companion galaxies in a $z>2$ proto-cluster. Both red galaxies in \\pks1138\\ are massive, of the order of $4-6 \\times 10^{11}$ M$_{\\odot}$. They lie along a Colour-Magnitude relation which implies that they formed the bulk of their stellar population around $z=4$.  In the \\uss0943 field we find no red galaxies at the redshift of the radio galaxy but we do confirm the effectiveness of our $JHK_s$ selection of galaxies at $2.3<z<3.1$, finding that 10 out of 18 (56\\%) of $JHK_s$-selected galaxies whose redshifts could be measured fall within this redshift range.  We also serendipitously identify an interesting foreground structure of 6 galaxies at $z=2.6$ in the field of \\uss0943.  This may be a proto-cluster itself, but complicates any interpretation of the red sequence build-up in \\uss0943\\ until more redshifts can be measured.  }  %{}{}{}{}{} % 5 {} token are mandatory ",
        "introduction": "At every epoch, galaxies with the reddest colours are known to trace the most massive objects \\citep{tka+05}.  The study of massive galaxies at high redshift therefore places important constraints on the mechanisms and physics of galaxy formation \\citep[e.g.,][]{hok+09}.  Studies of clusters at both low and high redshift also yield valuable insights into the assembly and evolution of large-scale structure, which itself is a sensitive cosmological probe \\citep[e.g.,][]{ecf+96}.  The most distant spectroscopically confirmed cluster lies at $z=1.45$ and has 17 confirmed members, a large fraction of which are consistent with massive, passively evolving ellipticals \\citep{srs+06, hcs+07}. High redshift proto-clusters have been inferred  through narrow-band imaging searches for Ly-$\\alpha $ emitters in the vicinity of radio galaxies \\citep[e.g.,][]{kpo+04, vrm+05}.  Such searches have been highly successful, with follow-up spectroscopy confirming $15-35$ Ly-$\\alpha$ emitters per proto-cluster.  However, the Ly-$\\alpha$ emitters are small, faint, blue star-forming galaxies, and likely constitute a small fraction of both the number of cluster galaxies and the total mass budget in the cluster.  Massive elliptical galaxies are seen to dominate cluster cores up to a redshift of $z\\approx1$ \\citep[e.g.,][]{vsh+01}, with the red sequence firmly in place by that epoch \\citep{ejd+06, bfp+03} and in fact showing very little evolution even out to $z\\sim1.4$ \\citep[e.g.,][]{lts+08}.  What remains uncertain is whether this holds true at yet higher redshift.  The fundamental questions are (1) at what epoch did massive, rich clusters first begin to assemble, and (2) from what point do they undergo mostly passive evolution? There is certainly evidence that the constituents of high redshift clusters are very different from that of rich clusters in the local Universe, with the distant clusters containing a higher fraction of both blue spirals \\citep[e.g.,][]{ejd+06} and active galaxies \\citep{gse+09}.  \\citet{vs03} found early-type galaxies with signatures of recent star formation in a $z=1.27$ cluster, perhaps symptomatic of a recently formed (or still forming) cluster.  Locating massive galaxies in higher redshift clusters will therefore yield valuable insight into the earliest stages of cluster formation and the build-up of complex structures which are observed in the local Universe.  In two previous papers, we found over-densities of red colour-selected galaxies in several proto-clusters associated with high redshift radio galaxies \\citep{kkt+06, kkk+07}. Follow up studies of distant red galaxies (DRGs) have shown that this colour selection criterion contains two populations: (1) dusty star-forming galaxies and (2) passively evolved galaxies. The exact fraction between these two populations is still a matter of debate \\citep{for04,gra06,gra07}. With our new $JHK$ colour-selection, we hope to be more efficient in identifying the evolved galaxy population. We have now embarked upon a spectroscopic follow-up investigation with the aim of confirming cluster members and determining their masses and recent star formation histories. Such a goal is ambitious as redshifts of these red galaxies are generally very difficult to obtain, especially amongst the population with no on-going star formation (i.e., lacking emission lines).  Here we present the first results from optical and near-infrared spectroscopy of targets in two of the best studied proto-clusters, \\pks1138\\ ($z=2.16$) and \\uss0943\\ ($z=2.93$).  In Section 2 we explain the target selection and in Section 3 outline the observations and data reduction steps. Section 4 presents the redshifts identified and in Section 5 we briefly analyse the properties of the two galaxies discovered at the redshift of \\pks1138, including their relative spatial location to the RG, ages, masses and star formation rates from spectral energy distribution fits and star formation rates from the H$\\alpha$ emission line flux.  In Section 6 we summarise the results and draw some wider conclusions.  Throughout this paper we assume a cosmology of H$_0$=71~km~s$^{-1}$, $\\Omega_{\\rm M}$ = 0.27, $\\Omega_{\\Lambda}$=0.73 \\citep{svp+03} and magnitudes are on the Vega system. ",
        "conclusions": "We observed a total of 57 near-infrared selected galaxies in \\uss0943\\ and 33 in \\pks1138\\ with a mix of optical and near-infrared multi-object spectroscopy to attempt the confirmation of massive galaxies in the vicinity of the central radio galaxy in these candidate forming proto-clusters.  We have determined that the $JHK_s$ selection criteria presented in \\citet{kkt+06} select mostly $2.3<z<3.1$ galaxies, thus confirming the validity of this colour selection technique (10 out of 18 confirmed redshifts of $JHK_s$-selected galaxies are in this redshift range).  Of the 57 sources in the field of the radio galaxy \\uss0943 we identify 27 spectroscopic redshifts and for the remaining 30 sources we obtained five upper limits on the redshifts, excluding them to be part of the proto-cluster.  The remaining 25 objects could not be excluded as either foreground or background. An unknown fraction of these may be at the redshift of the radio galaxy. % and are hence still plausibly consistent with being at the redshift % of the radio galaxy. However, we also pinpoint a foreground (but still quite distant) large scale structure at redshift $z\\approx 2.65$ in this field, so it is likely that some of the remaining galaxies are also associated with this structure.  Of the 33 galaxies observed in \\pks1138\\ we exclude two background objects and confirm two objects at the redshift of the radio galaxy (all on the basis of emission lines).  We cannot exclude the remaining 29 objects as being potential proto-cluster members --- we require deep near-infrared spectroscopy to locate the 4000\\AA\\ break. On the basis of SED fits, and the SFR derived from both the H$\\alpha$ line flux and the 24\\,$\\mu$m flux, we deduce that one of the two red galaxies confirmed to be at the redshift of the proto-cluster belongs to the class of dusty star-forming (but still massive) red galaxies, while the other is evolved and massive, and formed rapidly in intense bursts of star formation.  This represents the first red galaxies to be confirmed as a member of a proto-cluster above $z>2$. The only other galaxies of the same class are evolved galaxies found in an overdensity at $z=1.6$ in the Galaxy Mass Assembly ultra-deep Spectroscopic Survey \\citep{kcz+09}, and red massive galaxies in the Multi-wavelength Survey by Yale-Chile (MUSYC) sample at $z\\sim2.3$ \\citep[e.g.~][]{kvf+06}. Since their phase of high star formation has ended (the current SFR is of order 1~M$_{\\odot}$ yr$^{-1}$), we only just detect them through their emission line fluxes. However, this shows that even in these galaxies where we do not expect strong emission lines, it is still sometimes possible to derive redshifts from the H$\\alpha$ line.  It is highly likely that there are more galaxies on the red sequence in this field which have completely ceased star formation and are thus very difficult to identify spectroscopically. The next stage of our spectroscopic campaign is to obtain deep near-infrared spectroscopy in the $J$- and $H$-bands to search for objects with 4000\\AA\\ continuum breaks.  Given the inferred SFRs and stellar masses of the two confirmed $z\\sim2.15$ galaxies, they have to have formed quite early, which fits into the down-sizing scenario \\citep{cshc96}. Although this conclusion is drawn from limited numbers, the low fraction of sources with H$\\alpha$ emission lines in our sample suggests that most of the sources probably don't have much on-going star formation (SFR$<0.5 \\rm ~M_{\\odot} ~yr^{-1}$) and so their formation epoch might be quite high. This would be contrary to the idea of the giant elliptical assembly epoch being $z\\sim2-3$ \\citep[e.g.,][]{vfk+08, kvv+08}, but is consistent with results from SED age fitting of stellar populations, which point to $z_{form}>3$ (e.g., Eisenhardt et al. 2008\\nocite{ebg+08}).  The discrepancy may be resolved if subclumps form early and merge without inducing much star formation (i.e. dry merging). Alternatively, the fact that these galaxies are located in over--densities may predispose us to structures that formed early.   Finally, the proximity of these two massive galaxies to the RG implies that they will have an important impact on its future evolution. Given the crossing times, compared with the time remaining until $z=0$, it is plausible that one or both of the galaxies may eventually merge with the RG."
    },
    "0911/0911.3009_arXiv.txt": {
        "abstract": "We describe the design, performance, sensitivity and results of our recent experiments using the Australia Telescope Compact Array (ATCA) for lunar Cherenkov observations with a very wide ($600$~MHz) bandwidth and nanosecond timing, including a limit on an isotropic neutrino flux. We also make a first estimate of the effects of small-scale surface roughness on the effective experimental aperture, finding that contrary to expectations, such roughness will act to increase the detectability of near-surface events over the neutrino energy-range at which our experiment is most sensitive (though distortions to the time-domain pulse profile may make identification more difficult). The aim of our ``Lunar UHE Neutrino Astrophysics using the Square Kilometer Array'' (LUNASKA) project is to develop the lunar Cherenkov technique of using terrestrial radio telescope arrays for ultra-high energy (UHE) cosmic ray (CR) and neutrino detection, and in particular to prepare for using the Square Kilometer Array (SKA) and its path-finders such as the Australian SKA Pathfinder (ASKAP) and the Low Frequency Array (LOFAR) for lunar Cherenkov experiments. ",
        "introduction": "The origin of the ultra-high energy cosmic rays (UHE CR) --- protons and atomic nuclei with observed energies above $10^{18}$~eV and up to at least $2\\times 10^{20}$~eV --- is obscured due to their deflection and scattering in cosmic magnetic fields.  This makes the flux of all but the highest energy CR appear almost isotropic with respect to the Galaxy regardless of their source, so that measurements of arrival directions cannot reliably be used for source identification. At the highest energies, the deflection is less, and this allows the possibility of `seeing' nearby UHE CR sources.  Arrival directions of UHE CR detected by the Pierre Auger experiment above $5.6\\times 10^{19}$~eV are found to be statistically correlated with positions of nearby active galactic nuclei (AGN), which are in turn representative of the large-scale distribution of matter in the local universe \\cite{AugerScience07}. However, the flux is extremely low, and so the nature of the sources of UHE CR within this distribution remains at present unresolved.  An alternative means of exploring the origin of UHE CR is to search for UHE neutrinos.  As first noted by Greisen \\cite{Greisen} and by Zatsepin \\& Kuzmin \\cite{Zatsepin_Kuzmin}, cosmic rays of sufficient energy will interact (e.g.\\ via pion photo-production) with photons of the 2.725~K cosmic microwave background (CMB), with the resulting energy-loss producing a cut-off in the spectrum (the `GZK cut-off') at around $\\sim10^{20}$~eV from a distant source. These same interactions produce neutrinos from the decay of unstable secondaries. Several experiments \\cite{Takeda03,Bird95,Connolly06,Abu-Zayyad05, AugerSpectrum08, HiResSpectrum0809} have reported UHE CR events with energies above $10^{20}$~eV, and therefore a flux of these ``cosmogenic neutrinos'' is almost guaranteed.  Significant information on the CR spectrum at the sources is expected to be preserved in the spectrum of astrophysical neutrinos \\cite{Protheroe04} which varies significantly between different scenarios of UHE CR production. These include acceleration in the giant radio lobes of AGN, the decay of super-massive dark matter particles or topological defects, and $Z$-burst scenarios, the last of which have already been ruled out by limits placed on an isotropic flux of UHE neutrinos \\cite{Gorham04,Barwick06}. Of course, neutrinos are not deflected by magnetic fields, and so should point back to where they were produced, with even a single detection allowing the possibility of identifying the source of UHE CR. Here we emphasise that in all models of UHE CR origin we expect a flux of UHE neutrinos. See refs.~\\cite{ProtheroeClay2004, Falcke2004_SKAscienceCase} for recent reviews of UHE CR production scenarios and radio techniques for high-energy cosmic ray and neutrino astrophysics.  \\subsection{The lunar Cherenkov technique} \\label{thelunarcherenkovtechnique}  A high-energy particle interacting in a dense medium will produce a cascade of secondary particles which develops an excess negative charge by entrainment of electrons from the surrounding material and positron annihilation in flight.  The charge excess is roughly proportional to the number of particles in electromagnetic cascades, which in turn is roughly proportional to the energy deposited by the cascade.  Askaryan~\\cite{Askaryan} first noted this effect and predicted coherent Cherenkov emission in dense dielectric media at radio and microwave frequencies where the wavelength is comparable to the dimensions of the shower. At wavelengths comparable to the width of the shower, the coherent emission is in a narrow cone about the Cherenkov angle $\\theta_C=\\cos^{-1} (1/n)$ ($n$ is the refractive index), while for wavelengths comparable to the shower length, the coherent emission is nearly isotropic. The Askaryan effect has now been experimentally confirmed in a variety of media \\cite{Saltzberg_GorhamSAND01, Saltzberg_GorhamSALT05, Saltzberg_GorhamICE07}, with measurements of the radiated spectrum agreeing with theoretical predictions (e.g.\\ ref.\\ \\cite{Alvarez-Muniz02}). If the interaction medium is transparent to radio waves, the radiation can readily escape from the medium and be detected remotely. Since the power in coherent Cherenkov emission is proportional to the square of the charge excess, i.e.\\ to the square of the energy deposited, extremely high energy showers should be detectable at very large distances.  The lunar Cherenkov technique, first proposed by Dagkesamanskii and Zheleznykh~\\cite{Dagkesamanskii}, aims to utilize the outer layers of the Moon (nominally the regolith, a sandy layer of ejecta covering the Moon to a depth of $\\sim$10~m) as a suitable medium to observe the Askaryan effect. Since the regolith is transparent at radio frequencies, coherent Cherenkov emission from sufficiently high-energy particle interactions (specifically, from UHE cosmic ray and neutrino interactions) in the regolith should be detectable by Earth-based radio-telescopes (Askaryan's original idea was to place detectors on the lunar surface itself). First attempted by Hankins, Ekers \\& O'Sullivan~\\cite{Hankins96,James07} using the Parkes radio telescope, the Goldstone Lunar UHE neutrino Experiment (GLUE) at the Goldstone Deep Space Communications Complex \\cite{Gorham04}, the experiment at Kalyazin \\cite{Beresnyak05}, and the NuMoon \\cite{NuMoon20hARENA2008} experiment at the Westerbork Synthesis Radio Telescope (WSRT) have subsequently placed limits on an isotropic flux of UHE neutrinos. Our own project, LUNASKA, aims to develop the lunar Cherenkov technique for future use of the SKA.  Observations continue at both WSRT \\cite{Scholten06} and with our own project using the ATCA and using the Parkes radio telescope (to be reported elsewhere).  The technique has been the subject of several theoretical and Monte Carlo studies \\cite{ZasHalzenStanev92,Alvarez-Muniz02,GorhamRADHEP01,Beresnyak04,GayleyMutelJaeger2009} together with our own recent work \\cite{James07,JamesProtheroeLCT109,JamesProtheroe_DirectionalAperture2009}.  Future radio instruments will provide large aperture array (AA) tile clusters and arrays of small dishes with very broad bandwidths, with both factors allowing very strong discrimination against terrestrial radio frequency interference (RFI). The culmination of the next generation of radio instruments will be the Square Kilometre Array, to be completed around 2020, with smaller path-finders such as ASKAP (Australian SKA Pathfinder \\cite{ASKAP07}) to be built in the intervening period. In the meantime, we have been performing a series of experiments with the Australia Telescope Compact Array \\cite{ATCA}, an array of six 22~m dishes which were undergoing an upgrade to an eventual 2~GHz bandwidth at the time of the observations described here. Lunar Cherenkov experiments with these instruments, together with those proposed for LOFAR \\cite{Scholten06}, represent the foreseeable future of the technique.  We emphasize that the lunar Cherenkov technique is very different to conventional radio astronomy and requires non-standard hardware and signal processing as it is necessary to detect nanosecond-duration lunar Cherenkov radio pulses coming from a region too large (the apparent diameter of the Moon is $0.5^{\\circ}$) to image in conventional ways at nanosecond time resolution. Such pulses suffer dispersion in the Earth's ionosphere, and our experiment is the first to correct for this in real-time.  In the present paper we describe our recent experiments using the ATCA using the lunar Cherenkov technique.  We start by giving an overview of the experiment, the part of the moon targeted and observing times which were chosen in order to observe the Galactic Center and Centaurus A (UHE neutrino flux limits for these sources are reported in ref.~\\cite{ATCA2008data_CenA_GC_limit}), the antennas used, specialized hardware, triggering and signal processing.  We then discuss the effects of dead-time, our finite sampling rate and our approximate de-dispersion on the detection efficiency and effective observation time.  Finally, we present a new limit to an isotropic UHE neutrino flux, and discuss why it is important despite better limits existing from the Antarctic Impulsive Transient Antenna (ANITA) and the Radio Ice Cherenkov Experiment (RICE).  In appendices, we outline the effects of the two major theoretical uncertainties in our aperture and limit calculation being the UHE neutrino cross-section, and the small-scale lunar surface roughness. In the latter case, a new approximate treatment is described, which we use as well as the standard calculation methods to determine the experimental apertures and limits given in the main body of the paper. ",
        "conclusions": "\\label{ATCA_conclusion}  In 2008 we carried out observations of the Moon at the Australia Telescope Compact Array using the lunar Cherenkov technique to search for the signatures of UHE neutrinos. Although no lunar pulses of any source were positively identified, the instantaneous apertures for all the observations were the most sensitive yet achieved using the technique. The corresponding limits on an isotropic flux are not as strong as those from RICE and ANITA, but our methods to improve sensitivity to certain patches of the sky -- which, along with our limit on the UHE neutrino flux from Centaurus A and the Galactic Center \\cite{ATCA2008data_CenA_GC_limit}, will be published separately -- were a success: we found the exposure from our observations to Centaurus A to be higher than all previous experiments at neutrino energies of $10^{22}$--$10^{23}$~eV and above.  The lack of detection of nanosecond lunar pulses reported here is consistent with limits set by the ANITA experiment. Importantly, our methods to discriminate between true lunar pulses and RFI were very successful: we can be extremely confidant that no true lunar Cherenkov pulses were detected simultaneously by all three antennas, despite having of order ten million trigger events. This demonstrates primarily the power of nanosecond timing over a significant baseline, which was made possible by our use of a very wide (here, $600$~MHz) bandwidth, and was the main criterion used to discriminate against false events. Importantly, it was found in Sec.\\ \\ref{ATCAfarfield_byeye} that for events of reasonable time-structure (the category into which all candidates must fall), the automatic procedures produced times accurate to better than $1$~ns. Therefore, relying on such a procedure to search for candidate pulses in the future is justified, making a comparable analysis for a long observation run of a month or more feasible.  Our first estimates of the effects of small-scale surface roughness on the detectability of lunar Cherenkov pulses at high frequencies --- and the potential two orders of magnitude improvement in the effective aperture at the highest energies, and order-of-magnitude worsening of the threshold --- only serves to emphasize the importance of further work. Since these effects are indicative of `messy' signals caused by roughness-induced interference along the length of the cascade, it also suggests that timing criteria on signal origin (i.e.\\ that the signal comes from the Moon) be used instead of criteria on the pulse shape (i.e.\\ that it is very close to impulsive). Our use of such criteria to eliminate all our candidate events proves that this is possible. It is important to note that the energy range at which small-scale surface-roughness effects are expected to increase the strength of our limit is the range at which our ATCA experiment is most competitive.  For our observations, the use of an analog de-dispersion filter proved highly successful. The dispersion measure assumed in the filters' construction turned out to be very close to the actual values during observation periods, so that incorrect de-dispersion lost us less sensitivity than triggering inefficiency due to our non-infinite sampling rate. Though such filters must inevitably be superseded by a digital method, their continued use in the meantime will be valuable. Conversely, the finding that the loss from a finite sampling rate was greater for our simple trigger algorithm than that from incorrect de-dispersion is extremely important, and that in fact our `over'-sampling was an important factor in increasing (or rather, reducing the loss of) sensitivity. In future observations therefore the use of a smarter trigger algorithm that uses the already fully available information to reconstruct intermediate trigger values should be used, and perhaps should take as high a priority as digital de-dispersion.  Compared with an alternative single-dish experiment at Parkes, our experiment at the ATCA provided more effective area to high neutrino energies at the expense of less sensitivity to lower energy events. Since the parameter space at which the ATCA experiment is superior is best explored by low-frequency experiments such as NuMoon, we have transferred our efforts to utilising the Parkes dish, the results of which will be reported in a future contribution. We point out, however, that multi-telescope systems (such as the SKA and its pathfinders) will be more sensitive than single-dish experiments, and that our ATCA experiments would have been significantly more sensitive had we not been limited to using only three antennas due to telescope upgrade delays.  For future experiments at the ATCA or at other radio telescope arrays, further improvements such as real-time coincidence logic between three or more antennas, or even the ultimate goal of a coherent addition of the signals, would also improve the sensitivity. Without a further analysis of the typical RFI structure, it is not possible to determine the utility of real-time anti-RFI logic, though given the prevalence of RFI-triggered pulses, this too should be considered.  The lessons learned above should in all cases be applicable to any use of the lunar Cherenkov technique with an array of radio antennas. The advantage of using a giant radio array such as the SKA to search for lunar pulses has only been highlighted by these observations, especially since it will be placed in a low-RFI environment.  We have demonstrated techniques being developed by us ultimately for use with the Square Kilometer Array, and have been able with only 6 nights of observations using the ATCA to produce the lowest limit to an isotropic UHE neutrino flux below $3 \\times 10^{22}$~eV of any lunar Cherenkov experiment.  While at present the isotropic limits from lunar Cherenkov experiments are not competitive with RICE \\cite{Kravchenko06}, ANITA \\cite{GorhamANITA08} and (above $3\\times 10^{22}$~eV) NuMoon, use of the SKA in several years' time for lunar Cherenkov observations will provide a very powerful technique for UHE neutrino astronomy \\cite{JamesProtheroeLCT109}. With an estimated sensitivity to neutrinos $100$ times less energetic than those detectable with our experiment at the ATCA, the SKA will be able to probe the as-yet untested predictions for a flux of neutrinos from the GZK process. Furthermore, our current experiment has been able to access regions of the sky not accessible to ANITA and NuMoon, and with better exposure than RICE above $3\\times 10^{22}$~eV. Our flux limits to UHE neutrinos from Centaurus A will be discussed in a future paper.   \\appendix"
    },
    "0911/0911.4629_arXiv.txt": {
        "abstract": "The observational basis for asteroseismology is being dramatically strengthened, through more than two years of data from the CoRoT satellite, the flood of data coming from the Kepler mission and, in the slightly longer term, from dedicated ground-based facilities. Our ability to utilize these data depends on further development of techniques for basic data analysis, as well as on an improved understanding of the relation between the observed frequencies and the underlying properties of the stars. Also, stellar modelling must be further developed, to match the increasing diagnostic potential of the data. Here we discuss some aspects of data interpretation and modelling, focussing on the important case of stars with solar-like oscillations. ",
        "introduction": "\\label{sec:intro} In the last two decades data have become available on solar oscillations, which have allowed detailed inferences to be made on the properties of the solar interior \\citep{Christ2002}. This has led to the determination of the solar sound speed with considerable accuracy, detailed tests of the equation of state of the solar interior and a determination of the solar internal rotation, the origin of which is still poorly understood. The inferences about the solar internal structure obviously provide a strong test of modelling of stellar evolution. Although earlier solar models were in reasonable, but far from perfect, agreement with the helioseismic determinations \\citep{Gough1996}, recent redeterminations of the solar surface abundance have led to a very substantial increase in the discrepancies between the Sun and the model \\citep[for reviews, see][see also Section~\\ref{sec:soundspeed}] {Guzik2006, Guzik2008, Basu2008}.  The Sun is only one star at a specific stage of its evolution, and its structure is relatively simple, compared with other stars. Thus it is of obvious importance to extend seismic investigations to other stars. This is becoming possible, thanks to the development of ground-based instrumentation and the launch of three space missions that are providing extensive data on stellar oscillations. Thus it is important to consider the tools available for the analysis of such asteroseismic data and summarize the results that have already been obtained.    Here we concentrate on solar-like oscillations; these are characterized by being intrinsically damped and excited stochastically by near-surface convection \\citep[e.g.,][]{Houdek1999, Houdek2006}. Thus they are found in relatively cool stars, with significant outer convection zones, both on the main sequence and on the red-giant branch, as recently confirmed spectacularly by the CoRoT mission \\citep{DeRidd2009}. The modes relevant to observations of distant stars have low spherical-harmonic degree and hence are generally high-order acoustic modes; however, in evolved stars they may take on a mixed character, with a component corresponding to an internal gravity mode in the deep interior of the star, very substantially enhancing their diagnostic value \\citep[e.g.,][see also Section~\\ref{sec:mixmodes}]{Christ2004}.   At the simplest level, the oscillation frequencies are determined by the global properties of the star, more specifically the mean density $\\bar \\rho \\propto M/R^3$, where $M$ is the mass and $R$ the surface radius of the star. When combined with additional, nonseismic, data and stellar modelling, the stellar radius can be determined with substantial precision and, one may hope, accuracy \\citep{Stello2009}. Such a determination of the radius is particularly important for observations of extra-solar planets using the transit technique, where determination of the radius of the planet requires knowledge of the stellar radius \\citep{Kjelds2009}. From the point of view of stellar modelling, however, it is of substantially greater interest to obtain information about the stellar interior. From low-degree acoustic modes a measure can be obtained that is sensitive to the properties of the stellar core and hence, with additional assumptions, provides a measure of the amount of hydrogen consumed during stellar evolution and thus of the stellar age. With sufficiently detailed and accurate observations additional information can be obtained about the stellar interior, including inferences of the sound-speed and density variation in the central parts of the star \\citep[e.g.,][]{Gough1993a, Basu2002, Roxbur2004}.   An obvious requirement for asteroseismic inferences is the availability of good observational data. This is increasingly being met (see Section~\\ref{sec:obs}). Ground-based observations of Doppler velocity have been possible for a few stars, at level of sensitivity of a few ${\\rm cm \\, s^{-1}}$. Space-based asteroseismic observations were started serendipitously by the WIRE satellite \\citep{Buzasi2005, Bruntt2007}, followed by the launch in 2003 of the highly successful Canadian micro-satellite MOST \\citep{Walker2003, Matthe2007}. The French-led CoRoT satellite mission \\citep{Baglin2009}, launched at the end of 2006, has provided photometric oscillation data on a large number of stars, and high-quality data on an even more extensive stellar sample is promised by the NASA Kepler mission launched in March 2009 \\citep{Boruck2009}. Further observational projects promise a substantial increase in the quality of data available for asteroseismic analysis. However, the interpretation also depends on an adequate understanding of the relation between the observed oscillation properties, in particular the frequencies, and the underlying properties of the star, in order to design optimal diagnostics based on the observations. Furthermore, the physical interpretation of the results relies on stellar modelling, relating the physics of stellar interiors to stellar structure. Thus a crucial aspect of the analysis is the reliable computation of stellar models. This should obviously ensure numerical accuracy; an important step towards this goal is the detailed comparison of stellar models and oscillation frequencies computed by independent codes that has been undertaken in the ESTA collaboration, initially as part of the CoRoT project \\citep{Lebret2008a, Lebret2008b, Moya2008}. However, an equally important aspect is to ensure that the unavoidable approximations in the modelling are implemented in a consistent and accurate manner; only in this case can the asteroseismic analysis provide a reliable assessment of these approximations and, ideally, point to ways to improve them. Finally, the results should be presented in a statistically well-defined manner, such that the significance of the inferred properties, and of any deviations from the assumed models, can be assessed. It is probably fair to say that we are still some way from meeting these requirements.  In the present paper we discuss some aspects of the diagnostics of solar-like oscillations, based on asymptotic relations and computed properties; this in particular emphasizes the important probing of stellar cores. In addition, we consider some examples of uncertainties in stellar internal physics which have been, or may be, amenable to seismic inferences. More extensive overviews of helio- and asteroseismic diagnostics were provided by, for example, \\citet{Gough1993b} and \\citet{Christ2004}, and a comprehensive discussion of asteroseismology can be found in \\citet{Aerts2009}. ",
        "conclusions": ""
    },
    "0911/0911.2544.txt": {
        "abstract": "We investigate the stellar populations of Lyman $\\alpha$ emitters (LAEs) at $z=3.1$ and $3.7$ in $0.65$ deg$^2$ of the Subaru/\\textit{XMM-Newton} Deep Field, based on rest-frame UV-to-optical photometry obtained from the Subaru/\\textit{XMM-Newton} Deep Survey, the UKIDSS/Ultra Deep Survey, and the Spitzer legacy survey of the UKIDSS/UDS. Among a total of $302$ LAEs ($224$ for $z=3.1$ and $78$ for $z=3.7$), only $11$ are detected in the $K$ band, i.e., brighter than $K(3\\sigma)=24.1$ mag. Eight of the $11$ $K$-detected LAEs are spectroscopically confirmed. In our stellar population analysis, we treat $K$-detected objects individually, while $K$-undetected objects are stacked at each redshift. We find that the $K$-undetected objects, which should closely represent the LAE population as a whole, have low stellar masses of $\\sim 10^8$ -- $10^{8.5} M_\\odot$, modest SFRs of $1$ -- $100$ $M_\\odot$ yr$^{-1}$, and modest dust extinction of $E(B-V)_\\star < 0.2$. The $K$-detected objects are massive, $M_{\\rm star} \\sim 10^9$ -- $10^{10.5} M_\\odot$, and have significant dust extinction with a median of $E(B-V)_\\star \\simeq 0.3$. Four $K$-detected objects with the reddest spectral energy distributions, two of which are spectroscopically confirmed, are heavily obscured with $E(B-V)_\\star \\sim 0.65$, and their continua resemble those of some local ULIRGs. Interestingly, they have large Lyman $\\alpha$ equivalent widths $\\simeq 70$ -- $250${\\AA}. If these four are excluded, our sample has a weak anti-correlation between Ly$\\alpha$ equivalent width and $M_{\\rm star}$. We compare the stellar masses and the specific star formation rates (sSFR) of LAEs with those of Lyman-break galaxies (LBGs), distant red galaxies, submillimetre galaxies, and $I$- or $K$-selected galaxies with photometric redshifts of $z_{\\rm phot} \\sim 3$. We find that the LAE population is the least massive among all the galaxy populations in question, but with relatively high sSFRs, while NIR-detected LAEs have $M_{\\rm star}$ and sSFR similar to LBGs. Our reddest four LAEs have very high sSFRs in spite of large $M_{\\rm star}$, thus occupying a unique region in the $M_{\\rm star}$ versus sSFR space. ",
        "introduction": "\\label{sec:intro}    Lyman $\\alpha$ emitters (LAEs) are a galaxy population which are common in the high redshift Universe. In the last decade, many observations have succeeded in detecting LAEs from $z \\sim 2$ up to $z \\sim 7$, primarily based on narrow-band imaging to isolate Lyman $\\alpha$ emission \\citep[e.g.][]{hu1998,rhoads2000,iye2006}. Over a thousand LAEs are now photometrically or spectroscopically identified \\citep[e.g.][]{hu2002,ouchi2003,malhotra2004, taniguchi2005,shimasaku2006,kashikawa2006, dawson2007,murayama2007,gronwall2007,ouchi2008}.  Most LAEs have relatively faint, blue UV continua and large Ly$\\alpha$ equivalent widths (EWs). These properties collectively suggest that they are young star-forming galaxies with low metallicities \\citep[e.g.][]{malhotra2002}. In this sense, they may serve as building blocks of larger galaxies in hierarchical galaxy formation.   Stellar population analysis is helpful to further constrain the nature of LAEs. Recent multiwavelength observations covering near-infrared wavelengths have enabled analysis of the stellar populations of LAEs, and there is growing evidence that not all LAEs are primordial as described above. Table \\ref{tab:summary_of_earliers} summarises studies of the stellar populations of LAEs, including the work presented here. \\cite{gawiser2007} have performed a stacking analysis of $52$ LAEs at $z=3.1$, and concluded that LAEs have low stellar masses ($\\simeq 10^9 M_\\odot$), young-age components ($\\simeq 20$ Myr), and small dust extinctions ($A_V = 0$) \\citep[see also][]{gawiser2006, nilsson2007}. \\cite{pirzkal2007} have studied nine LAEs at $z \\sim 5$ found by HST/ACS slitless spectroscopy in the Hubble Ultra Deep Field, to show that these faint LAEs are all very young ($\\simeq 1$ -- $20$ Myr) with low masses ($\\simeq 10^6$ -- $10^8 M_\\odot$) and small dust extinctions ($A_V = 0$ -- $0.6$). On the other hand, \\cite{lai2007} have concentrated on three spectroscopically confirmed LAEs at $z = 5.7$ with IRAC detection, and derived relatively large stellar masses, $\\simeq 10^9$ -- $10^{10} M_\\odot$, mild dust extinctions ($E(B-V) \\simeq 0.15$ -- $0.2$), and a wide range of age, $\\simeq 5$ -- $100$ Myr. \\cite{lai2008} have studied both IRAC-detected and undetected LAEs at $z=3.1$ by the stacking method, to find that LAEs posses wide ranges of age ($\\simeq 160$ Myr -- $1.6$ Gyr) and mass ($\\simeq 10^8$ -- $10^{10} M_\\odot$) without dust extinction. \\cite{finkelstein2008b} have analysed $14$ LAEs individually, which are either IRAC-detected or undetected. Their results also show wide ranges of stellar population age ($\\simeq 2.5$ -- $500$ Myr), stellar mass ($\\simeq 10^8$ -- $10^{10} M_\\odot$), and dust extinction ($A_{1200} \\simeq 0.3$ -- $5.0$).     However, most of the samples constructed to date are not necessarily large and deep enough to constrain the average stellar populations of LAEs and to study rare, massive LAEs which may be a bridge to more massive and/or evolved objects like LBGs. In addition, correlations between stellar population parameters have not been addressed well.     Recently, \\cite{ouchi2008} have constructed the largest available sample of $z=3.1$ and $3.7$ LAEs in an about $1$ deg$^2$ of the Subaru/\\textit{XMM-Newton} Deep Field (SXDF) from deep optical broadband and narrowband data. These large survey data enable us not only to find many massive LAEs, but also to place better constraints on less massive (i.e., average) LAEs using stacking analysis. Our sample used in this study consists of $224$ ($78$) LAEs at $z = 3.1$ ($3.7$) covered by deep $JHK$ images taken with the UKIRT/WFCAM from UKIDSS Ultra Deep Survey \\citep{warren2007} and $3.6$ -- $8.0 \\mu$m images taken with the Spitzer/IRAC from the Spitzer legacy survey of the UDS field (SpUDS; PI: J. Dunlop)\\footnotemark{}\\footnotetext{\\texttt{http://ssc.spitzer.caltech.edu/legacy/abs/dunlop.html}}, among which $5$ ($6$) are detected in the $K$ band (i.e., brighter than the $3 \\sigma$ detection limit of the $K$-band image). The $K$ band corresponds to the rest-frame $V$ and $B$ bands at $z=3.1$ and $3.7$, respectively. For $K$-undetected LAEs we make median stacked images for individual bandpasses. Since the vast majority ($96${\\%}) are undetected in $K$, the spectral energy distributions (SEDs) constructed from stacking should closely represent the whole LAE population at each redshift. Since the time interval between the two redshifts is not so large (the age of the universe is $\\simeq 2.0$ Gyr at $z=3.1$ and $\\simeq 1.7$ Gyr for $z=3.7$), we treat the LAEs at $z=3.1$ and $3.7$ collectively as objects at $z \\sim 3$, without discussing possible evolution between the two redshifts. In this paper, we show the results of our stellar population analysis of these LAEs. We also compare the stellar populations of LAEs with those of other high redshift galaxy populations. The contribution of LAEs to the cosmic star formation density and the stellar mass density is also evaluated.  The outline of this paper is as follows. In Section 2, we present our data and LAE sample. The SED fitting method is described in Section 3. In Section 4, we present and discuss our SED fitting results. A summary is given in Section 5. Throughout this paper, we use magnitudes in the AB system and assume a flat universe with ($\\Omega_m$, $\\Omega_\\Lambda$, $h$) $=$ ($0.3$, $0.7$, $0.7$).   \\begin{table*} \\centering \\caption{Summary of Stellar Population Analysis of LAEs} \\label{tab:summary_of_earliers} {\\scriptsize \\begin{tabular}{cccccccc} \\hline Reference & Field & Area & redshift & $N^\\dagger$ & $m_{\\rm threshold}^\\ddagger$ & Bands$^\\ast$ & Remark \\\\  & & [arcmin$^2$] & & &  [mag] & & \\\\ \\hline This Study &  SXDF &  $2340$ &  $3.1$ &  $5(200)$ &  $K(3\\sigma) = 24.1$ &  $R, i, z, J, H, K, $[$3.6$], [$4.5$], [$5.8$], [$8.0$] &  1 \\\\ \\cite{gawiser2006} &  ECDF-S &  $992$ &  $3.1$ &  $(18)$ &  --- &  $U, B, V, R, I, z, J, K$ &  2 \\\\ \\cite{gawiser2007} &  ECDF-S &  $992$ &  $3.1$ &  $(52)$ &  $[3.6](2\\sigma) = 25.2$ &  $U, B, V, R, I, z, J, K, $[$3.6$], [$4.5$], [$5.8$], [$8.0$] &  3 \\\\ \\cite{lai2008} &  ECDF-S &  $992$ &  $3.1$ &  $18 (76)$ &  $[3.6](2\\sigma) = 25.2$ &  $B, V, R, I, z, $[$3.6$], [$4.5$] &  4 \\\\ \\cite{nilsson2007} &  GOODS-S &  $44.4$ &  $3.15$ &  $(23)$ &  $K_s(3\\sigma) = 23.4$ &  $U, B, V, i, z', J, H, K_s, $[$3.6$], [$4.5$], [$5.8$], [$8.0$] &  5 \\\\ This Study &  SXDF &  $2340$ &  $3.7$ &  $6(61)$ &  $K(3\\sigma) = 24.1$ &  $R, i, z, J, H, K, $[$3.6$], [$4.5$], [$5.8$], [$8.0$] &  1 \\\\ \\cite{pentericci2008} &  GOODS-S &  $160$ &  $3.5$ -- $6$ &  $((38))$ &  --- &  --- &  6 \\\\ \\cite{finkelstein2008b} &  GOODS-S &  $160$ &  $4.5$ &  $11 (3)$ &  --- &  $NB, V, i', z', J, H, K, $[$3.6$], [$4.5$] &  7 \\\\ \\cite{pirzkal2007} &  HUDF &  $11$ &  $4.00$ -- $5.76$ &  $5 (4)$ &  --- &  $z, J, H, K_s, $[$3.6$], [$4.5$] &  8 \\\\ \\cite{lai2007} &  GOODS-N &  $160$ &  $5.7$ &  $3$ &  $[3.6](3\\sigma) = 26.2$ &  $i', z', $[$3.6$], [$4.5$] &  9 \\\\ \\cite{chary2005} &  Abell 370 &  --- & $6.56$ &  $1$ &  --- &  --- &  10 \\\\ \\hline \\end{tabular} } %end {\\scriptsize  \\medskip \\begin{minipage}{170mm} \\begin{flushleft} $^\\dagger$ Number of objects detected in at least one rest-frame optical band. The number in the parentheses is the number of objects not detected in any of the rest-frame optical bands.   $^\\ddagger$ Rest-frame optical detection threshold for LAEs, if given.   $^\\ast$ Bands used for SED fitting, if given.   $^1$ $K$-detected objects are treated individually, while $K$-undetected objects are stacked.  $^2$ All objects are stacked.  $^3$ All objects are stacked.  $^4$ IRAC-detected and IRAC-undetected objects are stacked separately.  $^5$ All objects are stacked.  $^6$ LBGs with EW(Ly$\\alpha$) $>20$ {\\AA}, and all are treated individually irrespective of the NIR magnitude.  $^7$ All objects are treated individually irrespective of the NIR magnitude.  $^8$ All objects are treated individually irrespective of the NIR magnitude.  $^9$ All objects are treated individually.  $^{10}$ HCM 6A, a lensed LAE, detected at $3.6$ and $4.5 \\mu$m. SED fitting is performed considering the effect of H$\\alpha +$ [N\\,{\\sevensize II}] emission in the $4.5\\mu$m band.  \\end{flushleft} \\end{minipage} \\end{table*}   %--------------------------------------------------------------------- %--------------------------------------------------------------------- %--------------------------------------------------------------------- ",
        "conclusions": "\\label{sec:sum}  In this paper, we have investigated the stellar populations of LAEs at $z = 3.1$ and $3.7$ found in 0.65 deg$^2$ of the SXDF, based on deep, rest-frame UV-to-optical photometry from three surveys made with Suprime-Cam/Subaru, WFCAM/UKIRT, and IRAC/Spitzer. Among a total of $302$ LAEs ($224$ for $z=3.1$ and $78$ for $z=3.7$), only $11$ (or $4${\\%}) are detected in the $K$ band, i.e., brighter than $K(3\\sigma) = 24.1$ mag (AB). Among the $11$ $K$-detected LAEs, eight are spectroscopically confirmed. In our stellar population analysis, we treat $K$-detected objects individually, while we stack $K$-undetected objects at each redshift to derive an average SED. Since the vast majority ($96${\\%}) are undetected in $K$, the SEDs constructed from stacking should closely represent the whole LAE population at each redshift. We treat the LAEs at $z=3.1$ and $3.7$ collectively as objects at $z \\sim 3$, without discussing possible evolution between the two redshifts. Our LAE sample, based on deep optical and near-infrared data over a wide sky area, enables us not only to place strong constraints on the average properties of LAEs by stacking of many objects, but also to study rare, massive LAEs visible in $K$. We fit stellar population synthesis models to the multiband photometry of $K$-detected and $K$-undetected (stacked) objects, to derive their stellar masses, star formation rates, ages, and dust extinctions. We assume a constant star formation rate and fix the metallicity to $Z=0.2Z_\\odot$.  Our main results are as follows:  \\begin{enumerate} \\item The $K$-detected objects have stellar masses of $M_{\\rm star} \\sim 10^9$ -- $10^{10.5} M_\\odot$, while the $K$-undetected objects are 1 -- 2 orders of magnitude less massive, with $M_{\\rm star} \\sim 10^8$ -- $10^{8.5} M_\\odot$.  \\vspace{5pt} \\item The $K$-detected objects are distributed in a similar region in the $M_{\\rm star}$ versus $M_{\\rm V}$ plane to spectroscopically confirmed Lyman-break galaxies (LBGs) at $z \\sim 3$, while the $K$-undetected objects have lower $M_{\\rm star}/L_{\\rm V}$ ratios and bluer $M_{1500} - M_{\\rm V}$ colours than the $K$-detected ones (Figure \\ref{fig:R_RK_z3_all} and the left panel of Figure \\ref{fig:restOptUV_Mstar}).    \\vspace{5pt} \\item The star formation rates (SFR) of $K$-detected objects span a wide range of $10$ -- $10^4 M_\\odot$ yr$^{-1}$, while those of $K$-undetected objects fall between $1$ -- $100 M_\\odot$ yr$^{-1}$.   \\vspace{5pt} \\item There could be a bimodality in the age distribution among the $K$-detected objects (Figure \\ref{fig:histograms}). The ages of the $K$-undetected LAEs are around the median value of the ages of the $K$-detected LAEs.   \\vspace{5pt} \\item Four of the $K$-detected objects have extremely large $E(B-V)_\\star$, and thus their SFRs derived from rest-frame 1500 \\AA\\ continua and from Lyman $\\alpha$ luminosities are both underestimated by more than two orders of magnitude (Figure \\ref{fig:restUVLLyA_SFR}). On the other hand, the dust extinction of $K$-undetected objects is modest.   \\vspace{5pt} \\item The Lyman $\\alpha$ equivalent width (EW) weakly anti-correlates with $M_{\\rm star}$, except for the four dusty objects having large EWs (Figure \\ref{fig:EW_Pops}). No significant correlation is seen between EW(Ly$\\alpha$) and $E(B-V)_\\star$ (Figure \\ref{fig:EW_Pops}).   \\vspace{5pt} \\item The Lyman $\\alpha$ escape fraction decreases with $E(B-V)_{\\rm gas}$. At a fixed $E(B-V)_{\\rm gas}$, the escape fraction of our LAEs is higher than those of LBGs and local star-forming galaxies (Figure \\ref{fig:EBV_fesc}).   \\vspace{5pt} \\item The stellar masses and the specific star formation rates (sSFR) of LAEs are compared with those of LBGs, DRGs, SMGs, and galaxies with photometric redshifts of $z_{\\rm phot} \\sim 3$. It is found that the LAE population as a whole is the least massive among these galaxy populations, but with relatively high sSFRs (Figure \\ref{fig:Mstar_sSFR}).   \\vspace{5pt} \\item Four of our $K$-detected LAEs with the reddest SEDs, two of which are spectroscopically confirmed, are heavily obscured with $E(B-V)_\\star \\sim 0.65$, and their continua resemble that of Arp 220 and the local ULIRG (Figure \\ref{fig:multiwave_SEDs}). They have very high sSFRs in spite of their large stellar masses (Figure \\ref{fig:Mstar_sSFR}), and one of them are detected in the MIPS $24\\mu$m image, which supports the SED-fitting result of SFR. These facts suggest that the red LAEs are massive, dusty starburst galaxies. Our study demonstrates that wide-field Lyman $\\alpha$ surveys can detect such dusty starburst galaxies.  \\end{enumerate}"
    },
    "0911/0911.4073_arXiv.txt": {
        "abstract": "A statistical model for the equation of state and the composition of supernova matter is presented with focus on the liquid-gas phase transition of nuclear matter. It consists of an ensemble of nuclei and interacting nucleons in nuclear statistical equilibrium. A relativistic mean field model is applied for the nucleons. The masses of the nuclei are taken from nuclear structure calculations which are based on the same nuclear Lagrangian. For known nuclei experimental data is used directly. Excluded volume effects are implemented in a thermodynamic consistent way so that the transition to uniform nuclear matter at large densities can be described. Thus the model can be applied at all densities relevant for supernova simulations, i.e.~$\\rho=10^5 - 10^{15}$ g/cm$^3$, and it is possible to calculate a complete supernova equation of state table. The model allows to investigate the role of shell effects, which lead to narrow-peaked distributions around the neutron magic numbers for low temperatures. At larger temperatures the distributions become broad. The significance of the statistical treatment and the nuclear distributions for the composition is shown. We find that the contribution of light clusters is very important and is only poorly represented by $\\alpha$-particles alone. The results for the equation of state are systematically compared to two commonly used models for supernova matter which are based on the single nucleus approximation. Apart from the composition, in general only small differences of the different equations of state are found. The differences are most pronounced around the (low-density) liquid-gas phase transition line where the distribution of light and intermediate clusters has an important effect. Possible extensions and improvements of the model are discussed. ",
        "introduction": "\\label{intro}  The equation of state (EOS) of hot and dense matter plays a fundamental role in the understanding of core-collapse supernovae (SNe) and plenty of other astrophysical scenarios, like e.g.~mergers of compact stars or cooling proto-neutron stars. In hydrodynamic simulations of such systems the EOS gives the crucial information about the properties of matter under these extreme conditions. So far, even the most comprehensive numerical studies of core-collapse SNe have difficulties to achieve successful explosions within the progenitor mass range $10$ M$_\\odot  \\leq M_{prog} \\leq 15$ M$_\\odot$. Explosions in spherical symmetry where accurate three-flavor Boltzmann neutrino transport can be applied, have only been obtained for an 8.8 M$_\\odot$ O-Ne-Mg-core \\cite{kitaura06, fischer09}. In general, multidimensional simulations are expected to achieve explosions, but they are computationally very expensive \\cite{bruenn06,janka08,marek09}. In more detail, several explosion mechanisms are proposed from different groups: the neutrino-driven \\cite{bethe85} the magneto-rotational \\cite{leblanc70,bisno71} or the acoustic mechanisms \\cite{burrows06}. In addition to multidimensional effects and the aforementioned mechanisms, an improved EOS, uncertainties in the neutrino opacities and missing nuclear effects could help to revive the shock wave and finally trigger the explosion. As an extreme example, in Ref.~\\cite{sagert09} a successful explosion in spherical symmetry of a 15 M$_\\odot$ progenitor star was obtained, only due to the use of an EOS which incorporates quark degrees of freedom and an appearance of a mixed phase of quarks and nucleons in the early postbounce phase.  At present, there exist only two (hadronic) EOSs commonly used in the context of core collapse SNe. The EOS developed by Shen et al.~\\cite{shen98, shen98_2} and the EOS from Lattimer and Swesty (LS) \\cite{lattimer91}. The two EOSs are based on different models for the nuclear interactions. The former uses a relativistic mean field (RMF) approach, which we will also apply in our model. Nuclei are calculated in the Thomas-Fermi approximation. The latter is based upon a nonrelativistic parameterization of the nuclear interactions. Nuclei are described as a compressible liquid-drop including surface effects. Besides these two, there are plenty of EOSs which focus on particular aspects of nuclear matter but which are restricted to a certain range in temperature, asymmetry or density. In their range of validity they give a much more detailed description of the effects occuring there. The main difficulty in the construction of a complete EOS which is suitable for supernova simulations is the large domain in density $10^4$ g/cm$^{3} < \\rho < 10^{15}$ g/cm$^3$, temperature $0 < T < 100$ MeV and (total) proton fraction $0< Y_p< 0.6$ which in principle has to be covered. In this broad parameter range the characteristics of nuclear matter change tremendously: from nonrelativistic to ultra-relativistic, from ideal gas behavior to highly degenerate Fermi-Dirac gas, from pure neutron matter to symmetric matter. All possible compositions appear somewhere in the extended phase diagram: uniform nuclear matter, nuclei in coexistence with free nucleons, free nucleons with the formation of light clusters, or an ideal gas mixture of different nuclei, just to mention a few possibilities. Thus one needs a rather simple but reliable model which is able to describe all the different compositions and the phenomenon connected with them. Furthermore, from a numerical point of view the calculation of the EOS table itself is also not trivial.  For uniform nuclear matter plenty of different models for the EOS exist and most of them could in principle be applied in the supernova context. So far, the nuclear EOS at large densities is not fixed and is still one of the main fields of current research in high energy physics. From the previous paragraph it becomes clear, that in the application for supernova physics the main difficulties arise below saturation density, where the liquid-gas phase transition of nuclear matter takes place. Matter becomes non-uniform, as light and/or heavy nuclear clusters (nuclei) form within the free nucleon gas. The uniform nuclear matter EOS is only one of the essential input information for the construction of an EOS which is suitable for the description of all possible conditions which typically occur in a core-collapse supernova. For the non-uniform nuclear matter phase further model assumptions are necessary. In this article we will choose one particular EOS for uniform nuclear matter and then will mainly focus on the modeling and the resulting properties of matter below saturation density.  We want to emphasize that this low-density part of the EOS plays a special role in core-collapse supernovae: First of all, most of the computational time is spent in this low-density regime. Furthermore, the neutrinospheres are located at densities $\\sim 10^{11}$ g/cm$^3$. The neutrino spectra, which belong to the most important observables of a supernova, are formed here and thus carry the information about the properties of low density nuclear matter. As mentioned earlier, most supernova simulations fail to achieve a successful explosion. The shock front which is initially traveling outwards after the bounce looses energy due to nuclear dissociation and neutrino emission and finally stalls at densities $\\sim 10^{9}$ g/cm$^3$, see e.g.~\\cite{burrows85,thompson03}. In the neutrino reheating paradigm the high energetic neutrinos of the deeper layers heat the matter between the neutrinospheres and the stalled shock so that the shock is reenergized until it finally drives of the envelope of the star. Thus the EOS of subsaturation matter is of particular importance for the possible reviving of the shock wave.  The two EOSs mentioned before have been successfully applied in astrophysical simulations since many years. However, both of them are based on the single nucleus approximation (SNA) assuming that the whole distribution of different nuclei which forms at finite temperature can be represented by only one single nucleus. The single nucleus is found by a minimization procedure of the thermodynamic potential. The SNA has to be seen as a necessary assumption for any microscopic calculation based on one single unit cell with periodic boundary conditions, as e.g.~in Ref.~\\cite{bonche81,newton09}, too. In such microscopic calculations the nuclei are formed out of the nucleons which are placed in the unit cell just by the nuclear interactions, which is a very convenient aspect of these models. Already in Ref.~\\cite{burrows84} it was shown, that under most conditions the EOS is almost not affected by the SNA. But the SNA gives the composition only in an averaged way, as the spread in the distribution of nuclei can be large. In microscopic calculations quite often several similar minima of the thermodynamic potential are found, also indicating the occurrence of mixtures of different nuclei, in contrast to the SNA, see e.g.~\\cite{bonche81,newton09}. Furthermore, in Ref.~\\cite{souza08} it was presented for some particular EOS models, that the SNA leads to systematically larger nuclei, which was already shown in Ref.~\\cite{burrows84}, too. The distribution of nuclei can influence the supernova-dynamics, as e.g.~electron capture rates are modified. The electron capture rates are highly sensitive to the nuclear composition and the nuclear structure, see e.g.~\\cite{langanke00,langanke03,hix03,martinez06}. More closely connected to the present work, in Ref.~\\cite{caballero06} the authors investigated models which are based on a distribution of various nuclei, so called nuclear statistical equilibrium (NSE) models. It was shown with classical molecular dynamics simulations that these models give systematically larger neutrino cross sections, leading to shorter neutrino mean free paths. In their study also the importance of the remaining uncertainty regarding the composition was pointed out. We want to note that the previously mentioned systematic change in the composition within the SNA was not analyzed in Ref.~\\cite{caballero06}. The NSE distribution of nuclei was only compared to a SNA system with a nucleus with the mass and charge of the average of the distribution. The results of Ref.~\\cite{sawyer05} point in the same direction: it was shown that the proper treatment of a multicomponent plasma leads to greatly reduced ion-ion correlations and thus to increased neutrino opacities.  There are further limitations of the previously mentioned EOSs, \\cite{shen98, shen98_2} and \\cite{lattimer91}. Both of them do not include any nuclear shell effects. A good description of nuclei is only achieved on an averaged basis. It might be seen as part of the SNA not to attribute any certain shell structure to the single nucleus which only represents the average of the whole distribution of nuclei. But in cases where only very few different or even only one kind of nuclei appears (e.g.~at low densities and temperatures), shell effects are important and should definitely be taken into account in a self-consistent way in the composition and in the EOS. Shell effects can substantially alter the composition and are crucial for the evaluation of the weak reaction rates with neutrinos and electrons. In our model we do not want to use the liquid-drop formulation or the Thomas-Fermi approximation, but rather implement as much information gained from experiments about the nuclear masses as possible. Thus our approach is very contrary to the two other EOSs, in which the nuclear interactions are the only input information required in the physical model for non-uniform nuclear matter.  In some supernova EOSs, and also in the LS and in the Shen EOS, the distribution of light clusters is represented only by $\\alpha$-particles, in the same way as the distribution of heavy nuclei is represented by only one heavy single nucleus. In the most simple case the $\\alpha$-particles are described as a nonrelativistic, classical ideal gas, without any interactions with the surrounding nucleons. In contrast to this very simplified treatment, there are studies which focus exclusively on the role of light nuclei in supernovae and the medium effects connected with them. A model-independent description of low-density matter is given by the virial equation of state for a gas of light clusters with mass number $A\\leq 4$, \\cite{horowitz06a,horowitz06b,oconnor07}. In Ref.~\\cite{arcones08} the composition of this model was compared to the composition of a simple NSE model. Up to densities $\\rho\\sim10^{13}$ only small differences were found, mainly an increased $\\alpha$-particle fraction due to attractive nucleon-alpha interactions in the virial EOS. The more elaborated models of Refs.~\\cite{sumiyoshi08, roepke09} are based on a quantum statistical approach. Effects of the correlated medium such as Pauli blocking, Bose enhancement and self-energy are taken into account, leading e.g.~to the merging of bound states with the continuum of scattering states with increasing density (Mott effect). In both works \\cite{sumiyoshi08, roepke09} light clusters up to $A=4$ are considered and in Ref.~\\cite{sumiyoshi08} some heavy clusters in addition. In Ref.~\\cite{typel09} the quantum statistical approach is compared to a generalized RMF model which includes the light clusters as additional degrees of freedoms. In the generalized RMF model the Pauli-blocking shift is implemented in an empirical form which is derived from the quantum statistical model. In all these studies it is found that light clusters in addition to $\\alpha$-particles contribute significantly to the composition, with a particular role of the deuterons. This is also one of the results of Ref.~\\cite{heckel09}, where all stable nuclei with $A\\leq13$ are included, but medium effects are only considered on a simplified level. The inclusion of the additional degrees of freedom of the light clusters can affect the supernova-dynamics. E.g.~in Ref.~\\cite{arcones08} the influence of light nuclei on neutrino-driven supernova outflows was studied and a significant change in the energy of the emitted antineutrinos was found. As another example, in Ref.~\\cite{oconnor07} it was shown, that mass-three nuclei contribute significantly to the neutrino energy loss for $T\\geq4$ MeV. A different topic connected to light clusters is the possible formation of Bose-Einstein condensates, see e.g.~\\cite{funaki08arxiv,funaki08,heckel09}.  A lot of knowledge about the properties of hot compressed nuclear matter was gained from statistical multifragmentation models (SMM) which are used to analyze low-energy heavy-ion collisions \\cite{gross90, bondorf95}. In Refs.~\\cite{botvina04,botvina05,mishustin08, botvina08} it was pointed out that the state reached in these experiments ($T \\sim 3-6$ MeV, $\\rho \\sim 0.1 \\rho_0$, with the saturation density $\\rho_0$) is very similar to the conditions in a core-collapse supernova in the region between the proto-neutron star and the shock front. It is very attractive that the same well-established models which are used to describe matter in terrestrial experiments can be applied for matter in some of the most energetic explosions in the universe.  In an astrophysical context these models are usually called NSE models. In this chemical picture the bound states of nucleons are treated as new species of quasi-particles. NSE models are in principle extended Saha-equations, as presented in Ref.~\\cite{clifford65}. Within this approach the whole distribution of light and heavy nuclei can be included by construction. Furthermore, it is very easy to incorporate experimentally measured masses. This classical approach gives an excellent description as long as matter is sufficiently dilute that the nuclear interactions are negligible and if the temperature is so low, that the structure of the nuclei is not changed significantly, see e.g.~\\cite{ignatiuk78}. If the whole distribution of nuclei is taken into account, it becomes rather difficult to implement a proper description of the medium effects on the nuclei, which become important at large densities. Thus, especially the transition to uniform nuclear matter which happens around $1/2 \\rho_0$ leads to difficulties for NSE models. Obviously, in the high density regime close to saturation density, more microscopic SNA-models give a more reliable description. In this high density regime also very exotic nuclear structures, commonly called the \"nuclear pasta\" phases, appear \\cite{sonoda08}. In a recent 3D Skyrme-Hartree-Fock calculation \\cite{newton09} it was shown, that this frustrated state of nuclear matter is characterized by a large number of local energy minima, for different mass numbers and for different nuclear configurations. Thus one can expect that many different pasta shapes will coexist at a given temperature and density, which would require a statistical treatment beyond the present capabilities.  In addition to the subtleties at large densities, the assumption of NSE itself, in the sense that chemical equilibrium is established, is only valid for temperatures larger than $\\sim$ 0.5 MeV. For lower temperatures, the nuclear reactions are too slow compared to the typical dynamical timescales of ms within a supernova. Anyhow we will present results below $T=0.5$ MeV for completeness and for the sake of comparison with larger temperatures. For these low temperatures one has to keep in mind that they do not represent the actual conditions in a SN, but rather the groundstate of matter after a sufficiently long time.  Before presenting our own model, we want to discuss shortly the characteristics of some of the existing NSE models. In Ref.~\\cite{ishizuka03} Ishizuka et al.~included 9000 nuclei from the theoretical mass table of Myers and Swiatecki \\cite{myers90}. Excited states of the nuclei are treated with an internal partition function in the same manner as we will do. However, the model of Ishizuka et al.~does not consider any nuclear interactions or phenomenological excluded volume corrections. In this sense it is a rather pure NSE model. The SMM model of Botvina and Mishustin \\cite{botvina04,botvina05,mishustin08,botvina08} does not use any tabulated binding energies, but is based on a liquid-drop parameterization of the nuclear masses. The parameters of the SMM are extracted from nuclear phenomenology and provide a good description of multifragmentation data. Shell effects are not included, but in fact they are expected to be weak at the large temperatures for which this model is designed for. On the other hand the phenomenological SMM allows to include arbitrary heavy nuclei and is easy to modify to explore certain aspects of the nuclear interactions. E.g.~nuclear excited states do not have to be taken into account explicitly, but temperature effects are included in a temperature dependent part of the nuclear binding energies, with a separate volume and surface contribution. The SMM also contains an excluded volume correction, which is equivalent to our formulation at low densities. The NSE model of Blinnikov et al.~\\cite{blinnikov09} uses the tabulated theoretical nuclear binding energies of Ref.~\\cite{koura05}. In their model interactions of the free nucleons as well as the modification of the nuclear surface energy due to the presence of the free nucleons is taken into account. Nuclear excited states are described by the partition functions of Engelbrecht \\cite{engelbrecht91}. Besides the use of different nuclear interactions and partition functions, the major difference of this model compared to ours, which will be presented below, is, that excluded volume effects have not been implemented. In all the NSE models discussed above, the Coulomb energies are included in the Wigner-Seitz approximation, which we will also use in our model.  Despite the principle problems NSE models have to cope with at large densities, we want to construct a NSE model which allows to bridge the critical region from some percent of saturation density up to uniform nuclear matter. By using a NSE model we keep the rather simple but accurate description of low-density and low-temperature matter. For uniform nuclear matter a well-known RMF EOS will be applied. Our new NSE model shall be able to give a reasonable description of the transition to uniform nuclear matter, which in the microscopic SNA models is achieved automatically. For that we include the nuclear interactions also in the unbound nucleon contribution below saturation density. In most cases these interactions are not important, but they become crucial where the free nucleons constitute a significant fraction and the interactions are necessary for the liquid-gas phase transition. Furthermore, we develop a thermodynamic consistent description of excluded volume effects, so that we achieve the right asymptotic behavior for very dense and very dilute nuclear matter. We are aware that we apply the NSE description at densities at which we can not control all effects of the nuclear interactions any more, and our very phenomenological description becomes questionable. Anyhow, we want to explore the limits of a NSE model and compare the arising differences to other existing EOSs. So far an NSE model has never been applied close to saturation density and the existing models are not able to describe the transition to uniform nuclear matter at all. We can discuss the whole phase diagram of SN matter within one consistent model and can address all the aspects which were mentioned above, namely the distribution of heavy nuclei, the importance of the light cluster distribution, or the role of shell effects. As will be presented below, we will apply the same model which was used for the calculation of the nuclear masses for the interactions of the free nucleons, so that all nuclear interactions (apart from excluded volume effects) are based on the same Lagrangian.  The article is structured as follows. In Sec.~\\ref{sec_model} we will present our statistical model in detail. First we describe the different components, namely the RMF model for the unbound nucleons and uniform nuclear matter and the implementation of the nuclei, their excited states and the Coulomb energies. Then we derive a thermodynamic consistent description of the whole thermodynamic system, including excluded volume corrections. At the end of Sec.~\\ref{sec_model} we address the transition to uniform nuclear matter. In Sec.~\\ref{sec_res} we show and discuss our results, first for the composition and then for the EOS. The results for the EOS of our NSE model are systematically compared to the Shen and LS EOS. In Sec.~\\ref{sec_sum} we summarize and draw conclusions. In the last Sec.~\\ref{sec_out} we will give an outlook and discuss possible extensions of the model. Throughout the paper we use natural units where $\\hbar=c=k_B=1$. ",
        "conclusions": "\\label{sec_sum} In this article we presented a new statistical model for a complete supernova EOS, i.e.~a model which is in principle suitable to describe matter under all typical supernova conditions. In our EOS the interactions of the unbound nucleons are described with a relativistic mean-field model and the parameter set TMA. All nuclei are treated as separate particle species, using the experimental mass table of Ref.~\\cite{AudiWapstra} and the theoretical nuclear structure calculations of Ref.~\\cite{geng05}, which are based on the same Lagrangian density and parameter set TMA. Excited states of the nuclei are taken into account by a rather simple semi-empirical degeneracy function. Like in most other NSE models, the Coulomb energies are included by the Wigner-Seitz approximation, but are neglected for the protons for simplicity here.  The different components are set together to a thermodynamic consistent model, based on simple phenomenological and geometrical considerations: We attribute the volume of a hard uniform sphere at saturation density to the nuclei. Thus the nucleons can only reside outside of the nuclei. The nuclei feel also the presence of the unbound nucleons, so that the volume available for them is reduced by all the baryons (inside nuclei and in form of unbound nucleons) in the system. By these assumptions it is assured that nuclei can not exist above saturation density and that the unmodified RMF description is achieved if nuclei are not present. These excluded volume effects are implemented in a thermodynamic consistent way, affecting all thermodynamic quantities. Furthermore, within the model the RMF interactions of the nucleons are coupled to the nuclei via chemical equilibrium.  The thermodynamic consistent description with excluded volume corrections allows to apply the model at all densities and thus also the transition to uniform nuclear matter can be described. Still, the description of nuclei as quasi-particles requires the use of a mixed phase construction, and here we choose the most simple way of a Maxwell construction, based on locally fixed $Y_p$ and local charge neutrality. The phase transition to uniform nuclear matter occurs in most cases at the expected densities between $1/3 n_B^0$ and $2/3 n_B^0$.  Regarding the composition the following results are found. At high temperatures and low densities the clusters are completely dissolved into an ideal gas of neutrons and protons. With increasing density light clusters form which finally are replaced by heavy nuclei until uniform matter is reached. The proton fraction $Y_p$ sets the maximum fraction of nuclei. We note that in our calculations the chemical potentials of the nuclei were always low enough to describe them by a Maxwell-Boltzmann distribution. Before possible condensates of light clusters could form the light nuclei are replaced by heavy nuclei. Thus we think that for addressing the question of Bose-Einstein condensation in supernova matter it is necessary to include heavy nuclei, too.  At $T=0.1$ MeV up to the neutron drip only heavy nuclei are present, above which a dilute free neutron gas forms. We compared our composition at $T=0.1$ MeV to the one of Ref.~\\cite{ruester06}, and found a very good agreement. Thus we can conclude that a very accurate description in the low temperature and low density limit is achieved. This makes our model particular useful for the description of the evolution of the outer crust of a hot proto-neutron star to a cold deleptonized neutron star.  The nuclear distributions look qualitatively similar to the ones of Ref.~\\cite{ishizuka03}, where a different NSE model is used. At low $T$, due to neutron shell effects, the mass distributions are not the typical Gaussians as in Ref.~\\cite{botvina08}, but have rather a delta-function-like shape. Nuclei with the neutron magic number $N=$28, 50, 82, 126 or 184 appear with very high abundance. The narrow distributions lead to a stepwise change of the average mass number $<A>$ and $<Z>$ with density, in contrast to the continuous change in the models of LS \\cite{lattimer91} and Shen et al.~\\cite{shen98, shen98_2}. Consequently, also the mass fractions of the different components show a rather discontinuous behavior. Though the sharply peaked distributions support the single nucleus approximation (SNA), we find several peaks coming from different nuclei with similar abundances. Thus the average nucleus does not give a good description of the nuclear distribution. If the representative nucleus does coincide with the average of the distribution, it still would be necessary to include shell effects in the SNA. Contrary, at large temperatures the distributions can become very broad which cannot be taken into account in the SNA. In both cases the effect of the modified composition on capture rates and scattering amplitudes could be important. At $T=5$ MeV and $T=10$ MeV the shell effects become so weak, that the typical behavior of statistical models after the onset of the liquid-gas phase transition\\footnote{To avoid misunderstandings: here we are refering to the binodal surface of bulk nuclear matter shown in Fig.~\\ref{fig_bulkpd} and do not mean our Maxwell construction to uniform nuclear matter.} is recovered: the distributions change from steep exponentials to broad power-law and finally U-shape distributions with increasing density.  For such broad distributions the mass fractions of the different components change continuously. In general, the mass numbers increase with density. E.g.~at $Y_p=0.3$ and densities close to saturation all nuclei of the nuclear chart are populated up to the most heavy and super-heavy nuclei which exist in the mass table used.  In the EOSs of Refs.~\\cite{ lattimer91,shen98,shen98_2} only $\\alpha$-particles are used to represent light clusters. At low temperatures $T\\sim 1$ MeV this works well as the $\\alpha$-particles give the main contribution. Then the $\\alpha$-particle fraction is similar in the three different calculations. At high temperatures the whole ensemble of light nuclei becomes important, with a particular large contribution from deuterons, as also found in Refs.~\\cite{horowitz06a,horowitz06b,oconnor07,arcones08,sumiyoshi08,roepke09,typel09}. The deuterons surmount the alphas in many cases. At densities slightly below the abundant appearance of heavy nuclei (or the transition to uniform nuclear matter) the whole ensemble of light clusters gives the main contribution to the composition. Then also differences in the $\\alpha$-particle fractions of the NSE model compared to the SNA calculations occur and the $\\alpha$-particle fraction differs significantly (one order of magnitude) from the total light cluster fraction of the NSE model. Here we find that for very low $Y_p$ the light clusters become very neutron rich and heavier compared to large $Y_p$. However, this might also occur as a compensation of the limited mass table which does not include heavy nuclei with very low $Z/A$.  We compared the equation of state to the ones of LS and Shen et al. In general, up to densities $10^{-4}$ fm$^{-3}$ the three models agree very well. The differences of the NSE model to the two other EOSs are comparable to the differences between the LS and the Shen EOS. The NSE EOS is closer to the results of Shen et al., whose model is based on a similar RMF parameterization (TM1).  At low temperatures strong shell effects lead to narrow distributions in the NSE model so that the SNA is well applicable. In Refs.~\\cite{blinnikov09} a comparison with the Shen EOS also showed the insensitivity of thermodynamic variables to the underlying physical description. The same result was obtained in Ref.~\\cite{botvina08} in which the Statistical Multifragmentation Model (SMM) was compared to the LS EOS. However, at larger densities in our calculation some significant deviations occur.  In our NSE model the shell effects are weakened at large $T$, and significant differences around the onset of the liquid-gas phase transition appear. This happens at $n_B>10^{-3}$ fm$^{-3}$ for $T=5$ MeV and $n_B>10^{-2}$ fm$^{-3}$ for $T=10$ MeV. The large contribution of light clusters causes an increased energy density (differences up to $\\sim$10 MeV) and increased entropy per baryon (differences up to $\\sim$1.5), resulting in a lower free energy density in a narrow density range. The pressure and chemical potentials also become slightly reduced. Compared to the SNA models, the heavy clusters appear at larger densities. With the onset of heavy clusters a sudden decrease in the entropy and energy densities is observed. In this density range, one to two magnitudes below saturation density, the baryon contribution to the total EOS is large, so that the mentioned differences are also significant for the total EOS. It is important to note that the two different SNA models do not show any remarkable differences here. We think that the formalism of Ref.~\\cite{burrows84}, in which it was derived that the SNA does not influence the EOS, can not be applied in this density regime, as a Taylor-expansion around the representative nucleus is used, but the distributions are actually exponential or power-law. In a study of our NSE model without light nuclei with $A<20$ except for the alphas, we found that the differences became much weaker. Thus we conclude that it is the inclusion of light clusters (which cannot be represented by alpha-particles  in this density regime) which lead to the characteristic shape of the nuclear distributions and the appearance of heavy nuclei at larger densities. We expect that these effects cause the observed differences in the EOS compared to the two SNA models. In Ref.~\\cite{ishizuka03} a noninteracting NSE model was compared to the Shen EOS and a bulk RMF EOS. They also found that the differences are most pronounced around the onset of the liquid-gas phase transition region. They came to the similar conclusion that at the borderline of the liquid-gas phase transition region (dubbed the boiling point by the authors) fluctuations are large and that it is necessary to take into account the fragment mass and isotope distribution. Other reasons for the observed differences could be the phenomenological character of our NSE model, the description of excited states of the nuclei, the different nuclear interactions, or missing medium effects.  These results also do not take medium modifications beyond the excluded volume effects into account, which become relevant for the light clusters above $n_B=10^{-4}$ fm$^{-3}$ \\cite{roepke09,typel09}. We stress, that in the LS and Shen et al.~EOSs no medium effects on the $\\alpha$-particles are included either. However, we still achieve the dissolution of the light clusters close to saturation density by the excluded volume description. We did a comparison of the composition with detailed medium effects of Ref.~\\cite{typel09} to our NSE model in which we included only the same light cluster up to $A\\leq4$. For $T=10$ MeV and $Y_p=0.5$ we find that light clusters exist up to larger densities and appear with larger maximum fractions but the overall qualtitative behavior is similar. For $T=4$ MeV and $Y_p=0.5$ the differences are more pronounced, as our NSE model with excluded volume and interactions of the nucleons behaves similar to an ideal-gas NSE, almost until saturation density is reached. However, in the full calculation with all possible nuclei, for temperatures $T<10$ MeV and large densities heavy nuclei appear before the nucleon gas becomes very dense. A more detailed comparison of the NSE model with the works of \\cite{roepke09,typel09} would be an interesting study.  At larger densities the free energy of the NSE model becomes larger than in the EOSs of LS and Shen et al. We attribute this result to the limited description of the heavy nuclei due to the mass table used. Close to saturation density, the excluded volume effects become so strong that they can be seen directly in the EOS, e.g.~by an increased pressure. At slightly larger densities the uniform nuclear matter phase sets in, with the mixed phase given by a Maxwell construction. The chosen description of the phase transition to uniform nuclear matter leads to a continuous behavior of the free energy, the energy density and the entropy, the pressure and the total baryon chemical potential which are first derivatives of the free energy. The second derivatives behave discontinuously. We note that with the inclusion of trapped neutrinos in weak equilibrium the constant pressure plateau of the Maxwell construction is smeared out."
    },
    "0911/0911.3653_arXiv.txt": {
        "abstract": "{ Unveiling the structure of the disks around intermediate-mass pre-main-sequence stars (Herbig~Ae/Be stars) is essential for our understanding of the star and planet formation process. In particular, models predict that in the innermost AU around the star, the dust disk forms a ``puffed-up'' inner rim, which should result in a strongly asymmetric brightness distribution for disks seen under intermediate inclination. } { Our aim is to constrain the sub-AU geometry of the inner disk around the Herbig~Ae star R\\,CrA and search for the predicted asymmetries. } { Using the VLTI/AMBER long-baseline interferometer, we obtained 24 near-infrared ($H$- and $K$-band) spectro-interferometric observations on R\\,CrA. Observing with three telescopes in a linear array configuration, each data set samples three equally spaced points in the visibility function, providing direct information about the radial intensity profile. In addition, the observations cover a wide position angle range ($\\sim 97${\\deg}), also probing the position angle dependence of the source brightness distribution. } { In the derived visibility function, we detect the signatures of an extended (Gaussian FWHM $\\sim 25$~mas) and a compact component (Gaussian FWHM $\\sim 5.8$~mas), with the compact component contributing about two-thirds of the total flux (both in $H$- and $K$-band). The brightness distribution is highly asymmetric, as indicated by the strong closure phases (up to $\\sim 40${\\deg}) and the detected position angle dependence of the visibilities and closure phases. To interpret these asymmetries, we employ various geometric as well as physical models, including a binary model, a skewed ring model, and a puffed-up inner rim model with a vertical or curved rim shape. For the binary and vertical rim model, no acceptable fits could be obtained. On the other hand, the skewed ring model and the curved puffed-up inner rim model allow us to simultaneously reproduce the measured visibilities and closure phases. From these models we derive the location of the dust sublimation radius ($\\sim 0.4$~AU), the disk inclination angle ($\\sim 35${\\deg}), and a north-southern disk orientation (PA$\\sim$180-190{\\deg}). Our curved puffed-up rim model can reasonably well reproduce the interferometric observables and the SED simultaneously and suggests a luminosity of $\\sim 29~$L$_{\\sun}$ and the presence of relatively large ($\\gtrsim 1.2~\\mu$m) Silicate dust grains. Our study also reveals significant deviations between the measured interferometric observables and the employed puffed-up inner rim models, providing important constraints for future refinements of these theoretical models. Perpendicular to the disk, two bow shock-like structures appear in the associated reflection nebula NGC\\,6729, suggesting that the detected sub-AU size disk is the driving engine of a large-scale outflow. } { Detecting, for the first time, strong non-localized asymmetries in the inner regions of a Herbig~Ae disk, our study supports the existence of a puffed-up inner rim in YSO disks. } ",
        "introduction": "\\label{sec:intro}  For our understanding of the structure and physical processes in the disks around young stellar objects (YSOs), the inner-most disk regions are of special importance. Furthermore, it is believed that planet formation is a direct result of the grain aggregation and growth which should take place in the inner few AUs of the dusty disks around these stars. As the spatial scales of the inner circumstellar environment were not accessible to imaging observations until recently, most conclusions drawn on the geometry of the inner disk were based on the modeling of the spectral energy distribution (SED). Typically, the SED of intermediate-mass YSOs (Herbig~Ae/Be stars) shows a characteristic infrared excess emission, which is often interpreted as the presence of a circumstellar accretion disk, although the 3-D geometry of the innermost (AU-scale) region of these disks is still poorly known.  According to the current paradigm, most of the infrared excess emission in Herbig~Ae/Be stars originates from a passive dust disk \\citep{ada87}, whose thermal structure can be approximated with a cold disk interior and a hot surface layer \\citep{chi97}. To fulfill vertical hydrostatic equilibrium, the disk is expected to flare towards larger radii, allowing the outer disk regions to intercept more stellar light than expected for a geometrically flat disk \\citep{ken87}. At a certain distance from the star, most often referred to as the dust sublimation radius $R_{\\rm subl}$, the dust temperature will exceed the evaporation temperature of dust ($T_{\\rm subl}$), causing the truncation of the dust disk and the formation of a dust-free inner hole. This scenario also gained support from the first survey-type near-infrared interferometric observations of Herbig~Ae/Be stars. \\citet{mil01} could measure the characteristic size of many YSO disks, finding that the measured sizes scale roughly with the square-root of the luminosity $L_{\\star}$ of the stellar source \\citep{mon02}. Since this is the expected scaling-relation for the location of the dust-sublimation radius ($R_{\\rm subl} \\propto L^{1/2}_{\\star}$), this finding supports the idea that the near-infrared emission is tracing mainly hot material located in a structure at the location of the dust sublimation radius. More recently, significant deviations from the size-luminosity were detected concerning, in particular, the T~Tauri and Herbig~Be star regime \\citep{mon05}. These deviations might be explained either with contributions from scattered light \\citep[most important for T~Tauri stars; e.g.][]{pin08} or from a gaseous disk located inside of the dust sublimation radius \\citep[more important for Herbig~Be stars; see e.g.][]{ake05b,eis07a,kra08a,ise08,tan08}.  In 2001, \\citeauthor{nat01} and \\citeauthor{dul01} pointed out that the heating provided by the stellar radiation should considerably increase the disk scale height close to the dust sublimation radius, resulting in the formation of a ``puffed-up'' inner rim. This rim might cast a shadow on the more extended disk regions, possibly affecting the thermal disk structure out to hundreds of AUs. The shadowing effects of the puffed-up rim might already have been observed indirectly in the SED. \\citet{mee01} argued that the general shape of the SED of many Herbig~Ae/Be stars can be divided into two groups, where ``group~I'' sources show a pronounced 60~$\\mu$m bump, while ``group~II'' sources have a flatter SED, lacking the 60~$\\mu$m excess emission. Later, \\citet{dul04a} pointed out that this empirical classification could be explained by considering both the flaring properties of the outer disk and the shadowing effects of a puffed-up inner rim. Accordingly, group~I disks might exhibit a strongly flared shape, allowing the outer disk to step out of the shadow region, while group~II disks are fully self-shadowed.  Although these studies provided some important first insight, there are major uncertainties concerning both the magnitude of the puffing-up effect as well as the detailed rim shape. For instance, it was proposed that density-dependent dust sublimation effects \\citep[][referred to as IN05 in the following]{ise05} as well as grain growth and dust sedimentation \\citep{tan07} could affect the shape of the inner rim, resulting in a curved vertical rim geometry. In a recent study, \\citet{kam09} found that backwarming effects and the presence of highly refractory grain species might result in a dust rim which is located significantly closer to the star than anticipated in earlier studies. Using radiative transfer, they show that the rim is not an infinitely sharp wall, but has an optically thin region which might extend a significant fraction of the rim radius, resulting in a more diffuse rim morphology. Spatially resolved observations are required to directly measure the shape of the inner rim. Furthermore, various authors \\citep[e.g.][]{mir99} have pointed out that the infrared emission of many Herbig~Ae/Be stars might contain significant scattered-light contributions from optically thin circumstellar envelopes or halos. As investigated by \\citet{vin03} and others, high-angular resolution imaging observations, such as presented in this study, provide the only method to separate the disk and envelope contributions and to solve the ambiguities inherent to pure SED model fits.  One of the strongest predictions which is common to all puffed-up rim scenarios is the appearance of asymmetries in the source brightness distribution for disks seen under an intermediate inclination angle. This ``skew'' in the brightness distribution is a direct consequence of the vertical extension of the rim above the disk midplane, providing perhaps the most promising way to observationally distinguish between scenarios with and without a puffed-up rim. In order to detect these signatures, infrared interferometry provides not only the required milli-arcsecond (mas) angular resolution, but also offers a very sensitive measure for deviations from point-symmetry, namely the closure phase (CP) relation \\citep{jen58}. First closure phase measurements on YSOs were presented by \\citet{mon06}. Using the IOTA interferometer and baseline lengths up to 38\\,m, these authors measured statistically significant non-zero closure phases on six out of 14 stars (excluding binary stars). For five of these six stars, the detected closure phase signals were rather small ($\\Phi \\lesssim 5${\\deg}), while for the Be star \\object{HD\\,45677} phases of up to $\\sim 27${\\deg} could be detected.  In order to investigate the geometry of the inner circumstellar environment around a Herbig~Ae star in great detail and to obtain further evidence for or against the existence of a puffed-up inner rim, we have studied the Herbig~Ae star \\object{R\\,CrA} using ESO's Very Large Telescope Interferometer (VLTI) and the AMBER beam combination instrument. In the following, we will first summarize some earlier studies on R\\,CrA (Sect.~\\ref{sec:rcra}), followed by a description of our interferometric observations (Sect.~\\ref{sec:observations}) and the employed models (Sect.~\\ref{sec:modeling}). Finally, we will discuss our modeling results (Sect.~\\ref{sec:discussion}) and conclude with a brief summary (Sect.~\\ref{sec:conclusions}). ",
        "conclusions": "\\label{sec:conclusions}  We summarize our findings as follows: \\begin{itemize} \\item Using 24 VLTI long-baseline interferometric measurements, we studied the near-infrared emission from the Herbig~Ae star R\\,CrA and could clearly resolve the inner circumstellar environment in the $H$- and $K$-band. Besides this spectral coverage, our observations also provide a good sampling of the visibility function ($0.8 \\lesssim V \\lesssim 0.1$) towards a wide range of position angles (PA=16-113{\\deg}). \\item Even when considering only a single position angle, the measured visibility profile cannot be represented with the commonly applied one-component models, but indicates a more complex object geometry, likely including a disk and an envelope component. \\item The measured closure phase signals (up $\\Phi \\sim 40${\\deg}) clearly indicate a strongly asymmetric brightness distribution. \\item We find that the detected asymmetries are likely not related to the presence of a companion star, but directly trace the vertical structure of the disk around the dust sublimation radius ($R_{\\rm subl} \\sim 0.4$~AU). When seen under intermediate inclination, the increased disk scale height causes obscuration and shadowing effects, which result in the observed asymmetries. To constrain the precise rim geometry, we tested three geometric and physical models, including the generic skewed ring model and rim models with a vertical and a curved rim shape. Clearly, models with a smooth brightness distribution (i.e.\\ the skewed ring model and the curved rim model) provide a much better representation of our data than the sharp rim structure predicted by the vertical rim model. \\item Confronting our data with the detailed mathematical description of the rim structure presented by IN05, we find that we can reasonably well reproduce the measured SED, as well as the visibilities and closure phases with a disk inclination angle of 35{\\deg} and an incident luminosity of 29~L$_{\\sun}$. Given the absence of indications for strong active accretion, we consider that this luminosity is mainly of photospheric origin. For the dust grain size, we find that the presence of relatively large Silicate dust grains ($\\gtrsim 1.2~\\mu$m) is required to obtain agreement between the model and our data. The derived disk position angle of $\\sim 180-190${\\deg} agrees well with the orientation of the polarization disk ($\\theta=189 \\pm 5${\\deg}). Perpendicular to the disk axis, two bow shocks appear in the associated reflection nebula NGC\\,6729, suggesting that the detected disk is driving an outflow, which has shaped the bow shock-like structures. \\end{itemize}  Presenting one of the most comprehensive high-angular resolution studies on the inner structure of YSO disks, our findings provide strong evidence for the existence of a curved, puffed-up inner dust rim in Herbig~Ae stars. However, additional work is needed to further constrain the detailed rim geometry and to distinguish the rim morphology from local brightness inhomogenities, which might, for instance, be caused by hot accretion spots, spiral disk density patterns, clumpiness, or other transient phenomena. Given the complexity one likely faces on the probed sub-AU scales, this task will urgently require the model-independent aperture synthesis imaging capabilities which are just becoming available for various infrared interferometric facilities."
    },
    "0911/0911.0699_arXiv.txt": {
        "abstract": "We analyse the evolution of the properties of the low-redshift Intergalactic Medium (IGM) using high-resolution hydrodynamic simulations that include a detailed chemical evolution model. We focus on the effects that two different forms of energy feedback, strong galactic winds driven by supernova explosion and Active Galactic Nuclei (AGN) powered by gas accretion onto super-massive black holes (BHs), have on the thermo- and chemo-dynamical properties of of the low redshift IGM. We find that feedback associated to winds {(W)} and BHs leave distinct signatures in both the chemical and thermal history of the baryons, especially at redshift $z<3$. { BH feedback produces an amount of gas with temperature in the range $10^5-10^7$K, the Warm Hot Intergalactic Medium (WHIM), larger than that produced by the wind feedback. At $z=0$ the fraction of baryons in the WHIM is about 50 per cent in the runs with BH feedback and about 40 per cent in the runs with wind feedback. The amount of warm baryons ($10^4<T<10^5$K) is instead at about the same level, $\\sim 30$ per cent, in the runs with BH and wind feedback. Also, BH feedback provides a stronger and more pristine enrichment of the WHIM.}  We find that the metal--mass weighted age of WHIM enrichment at $z=0$ is on average a factor $\\sim 1.5$ smaller in the BH run than for the corresponding runs with galactic winds. We present results for the enrichment in terms of mass and metallicity distributions for the WHIM phase, both as a function of density and temperature.  Finally, we compute the evolution of the relative abundances between different heavy elements, namely Oxygen, Carbon and Iron. While both C/O and O/Fe evolve differently at high redshifts for different feedback models, their values are similar at $z=0$. We also find that changing the stellar initial mass function has a smaller effect on the evolution of the above relative abundances than changing the feedback model. The sensitivity of WHIM properties on the implemented feedback scheme could be important both for discriminating between different feedback physics and for detecting the WHIM with future far-UV and X-ray telescopes. ",
        "introduction": "Solving the problem of the missing baryons at low redshift is one of the important goals of observational cosmology \\citep[e.g.,][]{persicsalucci92,fukugita98,fukugita04}. Since the pioneering work by \\cite{cenostriker99}, cosmological hydrodynamic simulations provided a fundamental contribution in this field \\citep[see also][]{dave01}. Although the baryon budget is closed by observations at high redshift, $z\\magcir 2$, simulations indicate that about half of the baryons at low-$z$ should lie in a tenuous quite elusive phase, the so-called Warm-Hot Integalactic Medium (WHIM). The WHIM is predicted to be made of low density ($n_{\\rm H}\\sim 10^{-6}-10^{-4}$cm$^{-3}$) and relatively high temperature ($T\\sim 10^5-10^7$K) plasma, whose main constituents are ionised Hydrogen and Helium, with traces of heavier elements in high ionization states. Simulations have also shown that gas at these densities traces the filamentary structures which define the skeleton of the large-scale cosmic web, and is heated by shocks from supersonic gravitational accretion onto the forming potential wells.  Due to its low density, the collisionally ionised WHIM should be characterised in emission by a very low surface brightness in the UV and soft X--ray bands, so that its detection lies beyond the capability of available detectors and should await for instrumentation of the next generation \\citep[e.g.,][]{yoshikawa04}. Claims for the detection of the WHIM through soft X--ray emission have already been presented \\citep[e.g.,][]{zappacosta05,werner08}. However, these detections are limited to special places, such as external regions of massive galaxy clusters or large filaments connecting such clusters, where the gas density reaches values higher than expected for the bulk of the WHIM. A more promising approach to reveal the presence of the WHIM lies in the detection of absorption features in the spectra of background sources, at far ultraviolet (FUV) and soft X--ray energies, associated to atomic transitions from highly ionized elements, using GRB (Gamma Ray Burst) spectra \\citep{fiore00,branchini09}, or in the detection of Ly-$\\alpha$ absorption from the tiny fraction of neutral hydrogen revealed through FUV spectroscopy \\citep[see][ for a review]{richter08}. Although detection { of gas with} $T<10^5$K has been obtained at FUV frequencies along a fairly large number of sightlines \\citep[e.g.,][ and references therein]{tripp08}, the situation is less clear in the soft X--ray band. At these energies, claims of WHIM detection in absorption have been presented by different authors \\citep[e.g.][]{nicastro05,buote09}, although they are either controversial \\citep[e.g.][]{kaastra06,rasmussen07} or have moderate statistical significance.  Owing to the above observational difficulties in detecting and characterizing the physical properties of the WHIM, numerical simulations have played over the last years the twofold role to forecast its detectability, both in emission \\citep[e.g.,][]{roncarelli06,ursino06} and in absorption \\citep[e.g.,][]{Cen_etal01,kravtsov02,viel03whim,chen03,viel05whim}, and to predict its observational properties, also as a function of the physical processes included \\citep[see][ for a review]{bertone08}. For instance, feedback effects from star formation and accretion onto super-massive black holes (SMBHs) are expected to determine at the same time the process of galaxy formation and the physical properties of the diffuse cosmic baryons. However, as of today, the mechanisms powering feedback are still poorly understood and thus very difficult to model in a fully self-consistent way in cosmological hydrodynamical simulations. On the other hand, numerical simulations from different groups \\citep[e.g.,][]{cenostriker06,dave07,koba07,oppe08,tvtb08,wiersma_etal09} confirm that the thermo- and chemo-dynamical properties of the WHIM and, more generally, of the Intergalactic Medium (IGM) at different redshifts are sensitive to the nature and timing of the feedback.  Indeed, metals are observed in the low density IGM out to high redshift and it is likely that galactic winds, as observed in galaxies in the low redshift universe, are responsible for the IGM metal enrichment. Furthermore, SMBHs are also expected to have an impact on the gas distribution in galaxies and to blow ejecta out to large distances, especially during galaxy mergers.  In this paper, we will present an analysis of an extended set of cosmological hydrodynamical simulations with the purpose of investigating the effect that both galactic winds powered by supenova (SN) explosions and energy feedback from accretion onto SMBHs have in determining the thermal and chemical properties of the IGM. Although the main focus of the analysis is on the low-redshift IGM, we will also discuss how different feedback mechanisms leave their imprint on the evolutionary properties of diffuse baryons since $z\\sim 4$. Our simulations are based on the chemo-dynamical version of the GADGET-2 code \\citep{springel} presented by \\cite{T07}, which follows the production of different chemical species by accounting for detailed yields from Type-Ia and Type-II SN (SN-Ia and SN-II hereafter), as well as from intermediate and low mass stars in the thermally-pulsating asymptotic giant branch (TP-AGB) phase, while also accounting for the mass dependent life--times with which different stellar populations release their nucleosynthetic products. The model of galactic winds is that introduced by \\cite{springel2003}, that we consider both in its original version and in a version based on assuming that winds are never hydrodynamically decoupled from the surrounding medium \\citep[see also][]{DallaVecchia_Schaye08}. As for the BH feedback, we adopt the models originally introduced by \\cite{dimatteo05,springel05bh}. Besides investigating the effect of different feedback mechanisms, we will also consider the impact of changing the stellar initial mass function (IMF) on the resulting enrichment pattern of the IGM. Although the results of this paper have important implications for the detectability of the WHIM in view of next-generation X--ray missions, we defer an observationally oriented analysis of our simulations to a future work, aimed at investigating in detail how different instrumental capabilities will be able to characterize the WHIM properties at the level required to discriminate between different feedback models.  The plan of the paper is as follows. In Section \\ref{sim}, we describe the hydrodynamical simulations used and the feedback mechanisms adopted. Section \\ref{results} describes the main results of our analysis in terms of global gas properties (evolution of different phases over redshift), epoch of enrichment of gas particles, metallicity distributions at $z=0$ as a function of overdensity and temperatures, and evolution of different chemical elements. We provide a final discussion of our results and draw our main conclusions in Section \\ref{conclusions}. In the following, we will assume the values of the solar metallicity as reported by \\cite{asplund05}, { with $Z/X=0.0165$ ($X$: hydrogen mass; $Z$: mass contributed by all elements heavier than Helium)}. ",
        "conclusions": "\\label{conclusions} We have presented results from an extended set of chemo-dynamical simulations which have been carried out with the GADGET-2 code \\citep{springel}, including the implementation of chemical evolution described by \\cite{T07}. The analysis presented here was focussed primarily on the low-redshift properties of the inter-galactic medium (IGM) and its evolutionary properties, by considering both gas in a warm phase (gas particles with temperature { in the range $10^4$--$10^5$ K}) and in the so-called WHIM phase (gas particles with temperature in the range $10^5-10^7$ K). The main purpose of the analysis was to quantify the effect that different feedback mechanisms leave in the pattern of metal enrichment. Besides a simulation not including any efficient feedback (NW run), we performed simulations including energy feedback resulting from accretion onto super-massive black holes (BH run; \\citealt{springel05bh}; \\citealt{dimatteo05}), and from galactic winds powered by supernova (SN) explosions. We considered two implementations of the latter, based either on locally decoupling winds from hydrodynamics, so as to allow them to leave star forming regions (W runs), or on keeping gas in the winds always hydrodynamically coupled with the surrounding medium (CW runs).  The main results of our analysis can be summarized as follows. \\begin{itemize} \\item[(a)] The temperature--density and metallicity--density phase diagrams at $z=0$ are affected by the different feedback prescriptions. Including galactic winds has the effect of enriching with metals the low--density IGM. BH feedback is even more effective than winds in transporting metal-enriched gas from the star-forming regions to the WHIM phase. BH feedback also causes the presence of a non-negligible amount of relatively low--density, metal-enriched hot gas ($10^6-10^7$ K) which is not present in the simulations with galactic winds. \\item[(b)] The fraction of baryonic mass associated to warm gas, with { $T=10^4$--$10^5$K}, which should be associated to UV/Lyman-$\\alpha$ absorption systems, varies from $\\magcir 90$ per cent at high redshift, $z>3.5$, to about 30 per cent in the local universe, with a weak dependence on the adopted feedback model. This is in agreement with recent observational estimates from UV spectroscopy \\citep[e.g.][]{danforth08}. The hot phase, with $T>10^7$K, sums up to 3 per cent at $z=0$, again almost independent of the feedback model. The fraction of baryons in stars is instead the most sensitive to feedback: it ranges from about 10 per cent in the run with no feedback (NW), to about 2 per cent in run with BH feedback. \\item[(c)] The WHIM phase comprises about 35 per cent of the baryon budget at $z=0$ for the NW run, a value that increases to about 50 per cent for the run with BH feedback. This confirms that the share of cosmic baryons in the different phases does depend on the assumed feedback model. The redshift evolution of the mass fraction in the WHIM also differs in the different runs, with a stronger evolution in the run with BH feedback. \\item[(d)] The average age of enrichment of the warm and of the WHIM phases differs in the different runs, by an amount which depends on the gas overdensity $\\db$. Diffuse gas (\\db $<50$) in the warm phase at $z=0$ is typically enriched 0.5-1 Gyr earlier than in the WHIM phase. As for the gas within ``collapsed'' regions, with \\db $>50$, the warm phase is enriched 1 Gyr after the WHIM phase in all runs not including BH feedback. In the BH run, enrichment of the densest WHIM as measured at $z=0$ takes place more than $1.5$ Gyr after that of the warm medium. { BHs enrich the WHIM more promptly than W in the redshift range $2\\mincir z \\mincir 4$. On the other hand, winds are more effective at higher redshift, when BH accretion is still inefficient, and below $z\\sim 2$, when star formation in the BH run has been quenched. At all epochs, the hydrodynamically coupled winds (CW run) provide a slightly more recent enrichment of both warm and WHIM phases when compared to hydro-decoupled winds (W run). } BH feedback provides a faster enrichment at $z>2$, while below this redshift winds provide a more prompt WHIM enrichment. In particular, the model with hydrodynamically coupled winds (CW) provides at $z<3$ the most prompt enrichment of both warm and WHIM phases. \\item[(e)] In order to address the multi-phase nature of the WHIM, we compute the distribution of gas and metals in the WHIM phase as a function of overdensity. { As a result of the strong heating provided by BH feedback at $z\\simeq 2$--4, the} characteristic density of the WHIM in the BH run is a factor of a few lower than in the other simulations and, correspondingly, a smaller amount of gas is present in the dense WHIM. The same trend is also visible for the amount of metals present in the WHIM. Typically, most of the metals lie at over--densities between a few and 10 in the run with BH feedback, while the typical density of the metal-enriched gas is about two orders of magnitude higher in the run with galactic winds. { The reason for this is that most of the metals are ejected by BH heating from galaxies between $z=4$ and $z=2$. The high entropy level reached by enriched gas heated by BH feedback makes hard for it to be re-accreted within collapsed haloes at lower redshift. In turn, the drop of star formation in the BH run causes a comparable suppression of metal production at $z<2$}. \\item[(f)] Underdense WHIM regions (voids) have a very low median metallicity, however the metallicity of individual gas particles can reach values { larger than $0.1Z_\\odot$ within the 90 percentile. These enriched particles have been trasported to underdense regions both by galactic ejecta and by the action of gas-dynamical processes. Due to the intrinsic lack of diffusion in SPH, such particles spuriously preserve their high-metal content, rather than diffusing metals to surrounding metal poor particles.} \\item[(g)] The values of the [C/O] and [O/Fe] relative abundances at $z=0$ are similar for both the WHIM and warm phases, also irrespective of the feedback model. Their evolutionary pattern is instead sensitive to the adopted feedback model. Using a top--heavier IMF decreases the value of the [C/O] ratio by about 0.15 dex at all redshifts. \\end{itemize}  { The simulations presented in this paper have been analysed to study separately the effects that SN-triggered winds and BH energy feedback have on the evolution of the IGM. Clearly, in a realistic situation one expects SN and AGN feedback to be both at work and to cooperate in determining the cosmic cycle of baryons. However, we point out that our analysis was not aimed at establishing a best-fit feedback model, which is able to reproduce a variety of observational results. We aimed insted at determining the imprints that different feedback models leave of the properties of the IGM and their possible observational signatures. Furthermore, it is worth reminding that the parameters that we adopted for winds and BH feedback have been fixed by requiring each of these two feedback sources to reproduce a specific observational constraint: the cosmic star formation rate for winds \\citep{springel2003} and the $M_{BH}$--$\\sigma$ relation for BH feedback \\cite{dimatteo05}. Once the two mechanisms are allowed to be both present in a simulation, they are expected to have non--trivial interplays, which necessarely require a re-calibration of their characteristic parameters. }  The results obtained from our analysis have interesting implications for the detectability of the WHIM and for the possibility of characterising its thermal and chemical properties \\citep[e.g.,][]{stocke07}. Indeed, the possibility of detecting the WHIM through absorption lines in the spectra of background sources does not only depend on the amount of mass in this phase, but also on how this mass is enriched and distributed in density and temperature. With our analysis we have demonstrated that such distributions are rather sensitive to the adopted scheme of energy feedback. Analyses aimed at discussing the WHIM detectability from simulations, both in emission \\citep[e.g.,][]{yoshikawa04} and in absorption \\citep[e.g.,][]{Cen_etal01,viel05whim} have been so far based on approximate descriptions of the pattern of chemical enrichment. We will present in a future paper an observationally oriented analysis of our simulations, aimed at quantifying how the WHIM properties can be recovered under realistic observational conditions in the presence of different feedback schemes. There is no doubt that the possibility of performing high-resolution spectroscopy both in the X-ray (e.g., with micro-calorimeters onboard of large collecting area X-ray telescopes) and in the UV band (i.e., the now operating Cosmic Origin Spectrograph onboard of Hubble Space Telescope) will provide a leap forward in the study of the diffuse warm-hot baryons in the local universe. Cosmological hydrodynamical simulations, like those presented in this paper, offer the natural interpretive framework for these future observations, which will not only complete the census of baryons at low redshift but also characterize their physical properties."
    },
    "0911/0911.2655_arXiv.txt": {
        "abstract": "\\noindent We re-visit the production of cosmic rays by cusps on cosmic strings. If a scalar field (``Higgs'') has a linear interaction with the string world-sheet, such as would occur if there is a bosonic condensate on the string, cusps on string loops emit narrow beams of very high energy Higgses which then decay to give a flux of ultra high energy cosmic rays. The ultra-high energy flux and the gamma to proton ratio agree with observations if the string scale is $\\sim 10^{13}$ GeV. The diffuse gamma ray and proton fluxes are well below current bounds. Strings that are {\\it lighter} and have linear interactions with scalars produce an excess of direct and diffuse cosmic rays and are ruled out by observations, while heavier strings ($\\sim 10^{15}$ GeV) are constrained by their gravitational signatures. This leaves a narrow window of parameter space for the existence of cosmic strings with bosonic condensates. ",
        "introduction": "\\label{fieldtheory}  The interaction we consider is \\begin{equation} S_{\\rm int} = - \\kappa M \\int d^2\\sigma \\sqrt{-\\gamma} ~ h \\end{equation} where $\\kappa$ is a coupling constant assumed to be $\\sim 1$, $M$ is the string energy scale assumed to be at the Grand Unified scale, $\\gamma_{ab}$ the string world-sheet metric, and $h$ a scalar field.  A linear interaction can be considered quite generally. It can also arise if there is a bosonic condensate on the string. To see this explicitly, consider the model \\begin{equation} S = S_0 [\\Phi, H, \\ldots ] + \\kappa \\int d^4 x (\\Phi^\\dag \\Phi - M^2) H^\\dag H \\label{ftaction} \\end{equation} where $S_0$ is a field theory action that yields cosmic string solutions when $\\Phi$ gets a vacuum expectation value (VEV). $H$ is a scalar field that we can take to be the electroweak Higgs for concreteness.  At energy scales above the Grand Unified scale, the VEVs of $\\Phi$ and $H$ are both zero. At lower energy scales, but still above the electroweak scale, $\\Phi$ gets a VEV so that $| \\langle \\Phi \\rangle |^2 = M^2$ but $|\\langle H \\rangle |^2 =0$. At this stage, we also have string solutions whose tension is $\\mu \\sim M^2$ and width is $\\sim M^{-1}$. Inside the core of the string, where $\\Phi^\\dag \\Phi$ can become small, it may be favorable for $H$ not to vanish since the coefficient of the $H^\\dag H$ term in Eq.~(\\ref{ftaction}) becomes negative. As in the case of bosonic superconducting strings \\cite{Witten:1984eb}, there can be an $H$ condensate in the core of the string. Since $M$ is the only mass scale in the problem at this stage, the magnitude of $H$ is of order $M$ within the core of the string. At a lower energy scale, $H$ too gets a VEV. For example, if $H$ is the electroweak Higgs, this scale is $m \\sim 100$ GeV. We will assume $m \\ll M$ and hence the VEV of $H$ within the string is unaffected by the lower scale (electroweak) symmetry breaking.  The interaction term in Eq.~(\\ref{ftaction}) can now be written as \\begin{eqnarray} S_{\\rm int} &=& \\kappa \\int d^2 \\sigma \\int d^2 x_\\perp \\sqrt{-\\gamma} (\\Phi^\\dag \\Phi - M^2) H^\\dag H \\nonumber \\\\ &=& \\kappa \\int d^2\\sigma \\sqrt{-\\gamma} \\int d^2 x_\\perp (\\Phi^\\dag \\Phi - M^2) \\times \\nonumber \\\\ && \\hskip 1 in (\\langle H \\rangle_{\\rm in} +h )^\\dag (\\langle H \\rangle_{\\rm in} +h ) \\nonumber \\\\ &\\approx & - \\kappa M \\int d^2\\sigma \\sqrt{-\\gamma} ~ h + \\ldots \\end{eqnarray} In the first line we have split the integral into world-sheet and transverse integrals and included the Jacobian factor ($\\sqrt{-\\gamma}$) where $\\gamma_{ab}$ ($a,b=0,1$) denotes the induced metric on the string world-sheet. The subscript ``in'' on the angular brackets denotes that the relevant value is the VEV within the core of the string and we take $\\langle H \\rangle_{\\rm in} \\sim M$. Note that $h$ in the last line denotes the radial component of $H$ evaluated on the world-sheet.  In the next section, we will estimate the radiation of Higgses from cusps on cosmic string loops. ",
        "conclusions": "\\label{conclusions}  It is generally difficult to construct astrophysical scenarios that can accelerate protons to high energies sufficient to arrive as ultra-high energy cosmic rays. Top down models involving topological defects can naturally give the requisite energies. Depending on the precise variety of defect, however, they are likely to suffer from an excess of diffuse gamma ray production.  We have re-visited cosmic ray production from cosmic strings, taking into account the possibility of a linear interaction between a scalar field and the string world-sheet. Such an interaction arises, for example, if a scalar field acquires a VEV (``condenses'') within the string. Then beams of Higgs particles are emitted from cusps on cosmic string loops. Our analysis shows that such events can be responsible for the production of UHECR within reasonable parameters and also be consistent with measurements of the diffuse backgrounds and the photon to proton fraction in UHECR.  Our results are also interesting because they show that strings with bosonic condensates are allowed in a narrow window of energy scales. If they are much lighter than $\\sim 10^{13}$ GeV, they will produce an {\\it excess} of UHECR and diffuse fluxes, while if they are much heavier, they will cause conflicts with other cosmological constraints that depend on their gravitational effects. However, if the strings are around the Grand Unified scale, they can produce the observed flux of ultra high energy cosmic rays and also be heavy enough to be detected by their gravitational signatures in the near future.  We have focused on deriving {\\it analytical} estimates of cosmic ray fluxes under various simplifying assumptions. Further work is needed to obtain more detailed estimates that can be compared to observations."
    },
    "0911/0911.1925.txt": {
        "abstract": "The number of low-mass brown dwarfs and even free floating planetary mass objects in young nearby star-forming regions and associations is continuously increasing, offering the possibility to study the low-mass end of the IMF in greater detail. In this paper, we present six new candidates for (very) low-mass objects in the Taurus star-forming region one of which was recently discovered in parallel by \\citet{luhman2009caha6}. The underlying data we use is part of a new database from a deep near-infrared survey at the Calar Alto observatory. The survey is more than four magnitudes deeper than the 2MASS survey and covers currently $\\sim$1.5 deg$^2$. Complementary optical photometry from SDSS were available for roughly 1.0 deg$^2$. After selection of the candidates using different color indices, additional photometry from Spitzer/IRAC was included in the analysis. In greater detail we focus on two very faint objects for which we obtained $J$-band spectra. Based on comparison with reference spectra we derive a spectral type of L2$\\pm$0.5 for one object, making it the object with the latest spectral type in Taurus known today. From models we find the effective temperature to be 2080$\\pm$140 K and the mass 5-15 Jupiter masses. For the second source the $J$-band spectrum does not provide a definite proof of the young, low-mass nature of the object as the expected steep water vapor absorption at $1.33\\,\\mu$m is not present in the data. We discuss the probability that this object might be a background giant or carbon star. If it were a young Taurus member, however, a comparison to theoretical models suggests that it lies close to or even below the deuterium burning limit ($<$13 M$_{Jup}$) as well. A first proper motion analysis for both objects shows that they are good candidates for being Taurus members. %The proper motion of the third object does not seem to be in agreement with a Taurus membership but additional data are required to reduce the currently large error bars. If the object was a young Taurus member then its mass would probably be comparable to other very low-mass objects known today. Additional spectroscopic data will eventually confirm or disprove the young, low-mass nature of the candidates presented in this work. ",
        "introduction": "Taurus is certainly one of the best studied low-mass star-forming (SF) regions in the northern hemisphere. It lies at a distance of $\\sim$\\,140 pc \\citep{kenyon1994} and is very young \\citep[median age $\\sim$1 Myr,][]{briceno2003,luhman2003b}. The initial mass function (IMF) in Taurus has been studied intensively in the last years using large surveys mostly in the optical. It first seemed as if at the low-mass end the number of brown dwarfs (BDs) was too small compared to the expectations for a typical IMF \\citep{briceno2002}. While other studies found indications that this might not be the case \\citep{luhman2006spitzer,guieu2006} a recent analysis finds that a surplus of late K and early M type stars seems still to be present \\citep{luhman2009xray}. Apart from one exception the least massive objects in Taurus known today have spectral types of M9-M9.5 with corresponding effective temperatures of $\\sim$\\,2400-2500 K \\citep[][and references therein]{luhman2009caha6,luhman2006distribution,guieu2006}. Table~\\ref{bd_masses1} shows that theoretical models predict these objects to have masses above the deuterium burning limit of $\\sim$0.013 M$_{\\sun}$$\\approx$13 M$_{Jup}$ \\citep{burrows1997} which makes them very low-mass brown dwarfs. Only very recently \\citet{luhman2009xray} reported the detection of the first L0 dwarf in Taurus where the spectral type was derived from an I-band spectrum. However, this object appears to be underluminous for its assumed young age and we will discuss this object and its properties in more detail in section~\\ref{discussion}. In theory, young objects (1-3 Myr) with L and T spectral types and corresponding effective temperatures below $\\sim$2300 K could potentially be of planetary mass (Table~\\ref{bd_masses2}).  For other young star-forming regions the detection of L- and T-type objects has already been reported and their numbers are rather small but continuously increasing. In the 3-5 Myr old cluster $\\sigma$-Orionis one object, S Ori 70, has an observed spectral type of T5.5$\\pm$1 \\citep{zapateroosorio2002}, one object is of spectral type L4 and three objects have spectroscopically confirmed spectral types of L0-L1.5  \\citep{zapateroosorio2000}. If these objects are members of the young cluster then they can be considered as very good candidates for isolated planetary mass objects (i.e., with masses M$<$13 M$_{Jup}$). However, \\citet{mcgovern2004} found for one potential candidate in $\\sigma$-Orionis (S Ori 47) spectroscopic evidence that the object is rather an object from the field than from the young cluster. Recently, \\citet{bihain2009} report the detection of three additional objects in $\\sigma$-Orionis (S Ori 72-74) with theoretical masses of a few Jupiter masses (with L-T spectral types), but spectroscopic confirmation is still pending. In addition to the objects in $\\sigma$-Orionis \\citet{lodieu2008} and \\citet{luhman2008chameleon} report the finding of L-type objects in the Upper Sco Association and the Chameleon I region, respectively (median ages $\\sim$3-5 Myr). \\citet{weights2009} report the spectroscopic confirmation of 38 very low mass objects below the hydrogen burning limit in the Orion Nebular cluster (median age $\\sim$1 Myr). 10 of their objects could have masses below the deuterium burning limit and be thus planetary mass objects. \\citet{burgess2009} recently published the detection of a T-dwarf candidate in the 3 Myr-old star-forming region IC 348 based on methane-band imagery. The found a preliminary spectral type of T6 for the object and concluded that the detection of this source is consistent with the extrapolation of current lognormal IMF estimates down to the planetary mass domain given the survey area and completeness of their sample. Furthermore, there are three optically classified L dwarf candidates in the field showing some spectral signatures of youth \\citep{kirkpatrick2001,kirkpatrick2006,luhmanwilson2006}. Four additional objects in young  star-forming regions (1-3 Myr) are at the borderline between M and L-type objects but are subject to controversial discussions in the literature: \\citet{jayawardhana2006a,jayawardhana2006b} derived for four objects in Chameleon, Lupus and Ophiuchus spectral types of L0 and planetary masses while \\citet{luhman2007astroph} and \\citet{allers2007} derived at least for three of the objects significantly earlier spectral types (M7.25-M8.75) and higher masses. In the and the TW Hya association (5-10 Myr) only late M-type and thus more massive objects were found \\citep{liu2003,gizis2002}.  %Due to its young age, the disk fraction in Taurus, i.e., the fraction of young stars surrounded by protoplanetary disks, is higher than in other low-mass star-forming regions as for instance IC 348 and Chameleon I. %\\citet{luhman2006spitzer} found that for stars (spectral type M6 and earlier) the disk fraction amounts to  61$\\pm$8\\% while the disk fraction %of brown dwarfs (spectral type later than M6) is  40$\\pm$13\\%.  In this paper we present new candidates for low-mass BD in Taurus. At least two objects are good candidates for being among the least massive objects known today in this SF region possibly lying below or close to the edge of isolated planetary mass objects. For one object we derive a spectral type of L2$\\pm$0.5 and a most likely mass between 5 and 15 Jupiter masses. ",
        "conclusions": "Based on a deep NIR survey in Taurus we identified 6 new candidates for low-mass objects in this young star-forming region. One object, CAHA Tau 6 (now FU Tau B), was independently confirmed as young, low-mass BD of spectral type M9.25 by \\citet{luhman2009caha6}. For a second object CAHA Tau 1, we have spectroscopic evidence that it is indeed a young, very low-mass objects and we assigned a spectral of L2$\\pm$0.5, making it the object with the latest spectral type in Taurus. For the remaining 4 objects we thus far lack clear spectroscopic confirmation for their young, low-mass nature and we can only consider them as candidates at this point in time. However, the identification of CAHA Tau 1 and CAHA Tau 6 with spectral types L2 and M9.25, respectively, as well as the re-identification of 10 additional known low-mass Taurus members during the selection process, gives us confidence for the low-mass nature of the other objects. If they turn out to be indeed young Taurus sources then a comparison to theoretical models suggests that at least a second object (CAHA Tau 2) could be among the least massive known objects in this young star-forming region known today. In this case, its mass would be close to or below the deuterium burning limit of $\\sim$13 M$_{Jup}$ and its effective temperature would lie below $\\sim$2300 K making it also an L-type object. The findings in detail: \\begin{itemize} \\item{CAHA Tau 1:} The photometric properties, its proper motion and also the shape and spectroscopic features in the $J$-band spectrum make it very likely that CAHA Tau 1 is the first Taurus member with a spectral type as late as L2$\\pm$0.5 and an effective temperature of $\\sim$2100 K. For an age between 1 and 5 Myr, as derived from its position in a HR-diagram, this temperature corresponds to 5-15 M$_{Jup}$ and it might well be that CAHA Tau 1 is a free floating planetary mass object. Its broad band SED is consistent with model predictions for a 3 Myr-old object of 10 Jupiter masses with slight excess emission at 4.5\\,$\\mu$m possibly arising from a circum(sub-)stellar disk. \\item{CAHA Tau 2:} The optical and NIR photometric properties support the young, low-mass hypothesis for this object. Further evidence comes from its proper motion. A first $J$-band spectrum did not unambiguously prove the young, low-mass nature as the long wavelength end resembles that of evolved stars or carbon stars. However, a comparison of the photometric properties makes this rather unlikely. Assuming that CAHA Tau 2 is young, low-mass objects then theoretical models suggest that, depending on the age, it lies close to or below the deuterium burning limit making it another candidate for being among the least massive and coolest known Taurus members today. \\item{CAHA Tau 3:} For CAHA Tau 3 we lack spectroscopic information but its photometric properties are in good agreement with those of young, very low-mass objects. However, currently the proper motion analysis does not support a Taurus membership. Additional data is needed to reduce the large error bars and get a more robust assessment of the proper motion. Assuming a young, low-mass nature then theoretical models suggest that CAHA Tau 3 lies somewhere at the border between M- and L-type objects with $\\sim$15 M$\\_{Jup}$. In this case, the currently applied value for the extinction is probably overestimated and needs to be revised which would again reduce the estimated mass. \\item{CAHA Tau 4 and 5:} The optical and NIR data for CAHA Tau 5 are in good agreement with young BD while CAHA Tau 4, the brightest object in our sample, lies a little bit off in the presented color-color and color-magnitude plots. Based on its photometric properties CAHA Tau 4 could also be a (very) distant, evolved object. A fit to theoretical models for young low-mass objects shows that if CAHA Tau 4 was a young BD then it would probably have significant NIR excess emission. Fitting models to the photometric values of CAHA Tau 5 reveals that the currently applied value for the extinction (and/or possibly also the applied distance modulus) need to be refined once spectroscopic data is available. Our PM analysis does not put conclusive constraints on the Taurus membership for both sources. \\item{CAHA Tau 6:} This object was identified as FU Tau B, a young, M9.25 object in a BD binary system by \\citet{luhman2009caha6} during the referee process of this paper. We publish here the first $J$, $H$ and $K_s$ band data for this source complementing the broadband SED from \\citet{luhman2009caha6}. \\end{itemize}   So far we have not yet made an attempt to estimate the effect of our findings on the Taurus IMF. Once we know better which of the above mentioned candidates are indeed young, low-mass sources and how many additional candidates we find in the remaining and not yet here included fields of our deep NIR survey, it is certainly worth analyzing whether our results significantly influence the low-mass end of the IMF and help to fill-up the apparent deficiency of very low-mass objects compared to late K and early M type Taurus members.   %In the coming months we plan to continue this study as follows: (1) multi-filter, NIR spectroscopy of the candidates presented here will ultimately reveal the spectral type and surface gravity of the %objects and yield additional information on the individual extinction \\citep[see, e.g.,][]{guieu2006, allers2007, weights2009}. We will confirm or disprove the low-mass nature of the objects and, in %combination with the photometry and the theoretical models, refine the mass estimate for the candidates. (2) Additional Taurus fields observed during the last winter will be included in the overall %Calar Alto Taurus catalogue. (3) Matching our final NIR catalogue with additional optical catalogues covering the whole survey area will serve as basis for further candidate selections.  %CAHA Tau 4 might be an outlier in our sample as its photometric behavior can also be explained by a background carbon star. As CAHA Tau 5 and 6 appear brighter than the first three candidates they were not treated with the same attention as these objects. However, based on their photometry they are good candidates for being newly detected brown dwarfs. Additional spectroscopic observations for CAHA Tau 4, 5 and 6 will also help to confirm the true nature of these objects.  %As some of the seven candidates appear significantly fainter than the least massive young objects in Taurus known today, they are potential candidates for even less massive objects. Thus, we elaborate in this section a bit further the properties of these most interesting candidates  %In Figure~\\ref{SED} we plot the optical and IR magnitudes for CAHA Tau 1, CAHA Tau 2, and CAHA Tau 3 as listed in Table~\\ref{candidates}. However, we de-reddened the observed magnitudes  (see caption of Figure~\\ref{SED}) to compare them to theoretical models for the photospheric emission of young BD and planetary mass objects \\citep[the models are based on, e.g.,][]{burrows1997,burrows2003,burrows2006}. In the left plot we show models for objects with 10 and 20 Jupiter masses and an age of 2 Myr, while in the right plot we depict theoretical models for objects with 10 and 40 Jupiter masses and 3 Myr. We stress again that these are rather conservative age estimates for this young SF region. It shows that for the 3 Myr plot most photometric points of all three objects fall within the area defined by the theoretical models. Only the optical magnitudes of CAHA Tau 2 appear to be brighter than predicted by the model. Of course, in general the object mass should be independent of wavelength, but given the uncertainties in the models, the extinction and possible circumstellar material, it is difficult to better constrain the data. Even more interestingly, most of the photometric points of the least luminous object CAHA Tau 3 lie even within the area defined by the two very low-mass 2 Myr models. Assuming that this object might even be younger than 2 Myr and that it might be surrounded by a circum(sub)stellar disk this object could lie on the edge between young very low-mass BD and free floating planetary mass objects.  %As seen in \\citet{allers2006} and \\citet{allers2007} one has to be extremely careful to derive masses, spectral types and effective temperatures of young low-mass objects solely from photometric measurements. Hence, we complemented our data for the three least luminous objects with NIR J-band spectra. Unfortunately the data for the most interesting object (CAHA Tau 3) were taken under the least favorable observing conditions and were not of the required quality for detailed data analysis. In Figure~\\ref{spectra} we show the J-band spectra of CAHA Tau 1 and CAHA Tau 2 (both smoothed by a factor of 20) and compare them to model spectra for very young, very low-mass BD and planetary mass objects \\citep{burrows2003,burrows2006}. One can clearly see how the flux increases with wavelengths and also the deep water absorption band at the long-wavelength end. It is striking that the observed objects show an even steeper rise in the spectrum and that the water absorption sets in at longer wavelengths (around 1.4$\\,\\mu$m) than predicted by the models. Again, this might point to an extremely young age, a very low mass and/or additional circum(sub)stellar material. It is also striking that apart from the deep water band no additional spectroscopic features seem to be present in the observed spectra. The only possible detection are the Na I lines slightly shortward of 1.14$\\,\\mu$m. We point out that between 1.2$\\,\\mu$m and the 1.4$\\,\\mu$m water absorption band the theoretical spectrum for a 5 Jupiter mass object with 3 Myr appears also rather featureless. Still, the non-detection of spectral features in our objects is surprising as even despite the moderate observing conditions the signal-to-noise of the data should have been sufficient to observe any spectroscopic signature such as TiO or VO if they were background giants or KI if they were old, field L-dwarfs.  %We are aware of the fact that current theoretical models bear imperfections. Especially at these very young ages and in this low-mass regime the uncertainties are significant due to necessary assumption regarding, for instance, the accretion process. Furthermore, although we assumed a value for the visual extinction towards the objects which is higher than the typical value for Taurus \\citep[$A_V\\le3.0$,][]{kenyon1994}, this figure has still to be observationally confirmed and varies most likely from object to object. In case the extinction is higher the objects are intrinsically brighter and thus more massive. %Keeping this in mind, we still do believe that the objects presented here are extremely good candidates to be among the least massive, free-floating, young objects known today. However, the ultimate proof that one or more objects are of planetary nature with a potentially mass below the Deuterium burning limit of $\\le 13$M$_{Jup}$ has still to be made. In case of a positive confirmation, this finding might have a tremendous impact on the theory of star and planet formation as then the formation of these objects needs to be explained. Do not only stars and brown dwarfs but possibly also planetary mass objects form directly from collapsing molecular cloud material and not only inside circumstellar disks around more massive objects? Or were the objects ejected from multiple systems before they could gain more mass? Further high-resolution and proper motion studies might help to find answers to these questions.   %% If you wish to include an acknowledgments section in your paper, %% separate it off from the body of the text using the"
    },
    "0911/0911.1327_arXiv.txt": {
        "abstract": "We present a large X-ray selected serendipitous cluster survey based on a novel joint analysis of archival Chandra and XMM-Newton data. The survey provides enough depth to reach clusters of flux of $\\approx 10^{-14}~\\mbox{ergs}~\\mbox{cm}^{-2}~\\mbox{s}^{-1}$ near $z$ $\\approx$ 1 and simultaneously a large enough sample to find evidence for the strong evolution of clusters expected from structure formation theory.  We detected a total of 723 clusters of which 462 are newly discovered clusters with greater than 6$\\sigma$ significance. In addition, we also detect and measure 261 previously-known clusters and groups that can be used to calibrate the survey. The survey exploits a technique which combines the exquisite Chandra imaging quality with the high throughput of the XMM-Newton telescopes using overlapping survey regions.   A large fraction of the contamination from AGN point sources is mitigated by using this technique.  This results in a higher sensitivity for finding clusters of galaxies with relatively few photons and a large part of our survey has a flux sensitivity between $10^{-14}$ and $10^{-15}~\\mbox{ergs}~\\mbox{cm}^{-2}~\\mbox{s}^{-1}$.  The survey covers 41.2 square degrees of overlapping Chandra and XMM-Newton fields and 122.2 square degrees of non-overlapping Chandra data.  We measure the log N-log S distribution and fit it with a redshift-dependent model characterized by a luminosity distribution proportional to $e^{-\\frac{z}{z_0}}$.  We find that $z_0$ to be in the range 0.7 to 1.3, indicative of rapid cluster evolution, as expected for cosmic structure formation using parameters appropriate to the concordance cosmological model.  Confirmation of our cluster detection efficiency through optical follow-up studies currently in progress will help to strengthen this conclusion and eventually allow to use these data to derive tight contraints on cosmological parameters. ",
        "introduction": "There is a long history of compiling catalogs of nearby galaxy clusters by using optical telescopes \\citep{abell,zwicky}.  The development of focusing X-ray telescopes and the early discovery of X-rays from clusters of galaxies \\citep{byram} provided another sensitive method for finding clusters of galaxies. A number of X-ray catalogues now exist for the brightest clusters. Most of the nearby clusters are known from flux-limited surveys that cover a large fraction of the sky from ROSAT and Einstein down to $10^{-12}~\\mbox{ergs}~\\mbox{cm}^{-2}~\\mbox{s}^{-1}$ (REFLEX, \\citealt{boehringer}; BCS, \\citealt{ebeling}; EMSS, \\citealt{gioia90}). These surveys have discovered 100-450 clusters.   Deeper surveys that extend to $10^{-14}~\\mbox{ergs}~\\mbox{cm}^{-2}~\\mbox{s}^{-1}$ and therefore probe a larger range of redshifts have found 60-200 clusters serendipitously (160 square deg., \\citealt{vikhlinin}; 400 square deg. \\citealt{burenin}; ROSAT NEP, \\citealt{gioia01}, \\citealt{henry2}; WARPS, \\citealt{scharf}, \\citealt{horner}). XMM-Newton surveys have found 19 clusters serendipitously \\citep{georgantopoulos} and 12 clusters in the Large Scale Structure Survey \\citep{willis} down to deeper flux levels.  The slope of the log N-log S distribution for clusters is close to -1 down to a flux limit of $\\approx 10^{-13}~\\mbox{ergs}~\\mbox{cm}^{-2}~\\mbox{s}^{-1}$ which means that a wide survey with short exposures will approximately detect the same number of clusters as a narrow survey with longer exposures for the same total exposure time. Thus the total number of clusters is simply proportional to the product of a telescope's etendue (effective area times field of view) times the total survey exposure time divided by the minimum number of photons that one can reliably use to identify a cluster unambiguously.  A large number of the previous surveys used ROSAT which has an area of 200 cm$^2$ and a field of view of 3 square degrees. The fields of view of Chandra (0.1 square degrees) and XMM-Newton (0.2 square degrees) are not particularly high, but the effective area of Chandra (400 cm$^2$ at 1 keV) and especially XMM-Newton (2000 cm$^2$ at 1 keV) are high.  XMM-Newton, in particular, has an etendue comparable to ROSAT.  One difficulty with using XMM-Newton to perform cluster surveys compared to ROSAT is that the large area and smaller field of view will collect on average more distant groups of galaxies, with smaller angular size, because the survey will be deeper rather than wider for the same exposure time.  The fact that the clusters have smaller angular size means that it is harder to determine the difference between a cluster or group and a point source.  On the other hand, Chandra has a exquisitely sharp point-spread function (PSF) that distiguishes AGN point sources from clusters at all redshifts across its entire field of view.  Therefore, by implementing an analysis method which combines the Chandra PSF with the XMM-Newton throughput, we demonstrate that we can produce a large sensitive X-ray cluster survey.  Recently, a number of independent astronomical methods have converged on a standard cosmological model, which indicates the presence of a poorly understood dark energy component.  An extensive search for new galaxy clusters will provide new constraints on cosmological models.  In general, the abundance and distribution of clusters of galaxies in the Universe provides precision and complementary cosmological constraints to other methods \\citep{eke}. In particular, the measurements are most sensitive to $\\sigma_8 \\Omega_m$, where $\\sigma_8$ is the normalization of the matter power spectrum on 8 Mpc scales and $\\Omega_m$ is the matter density.  A large enough sample at a range of redshifts can constrain $w$, the equation of state parameter for dark energy (e.g., \\citealt{haiman}, \\citealt{wang}). The distribution of clusters as a function of mass and the spatial power spectrum constrains cosmological parameters as well.    The content of this initial paper principally concerns the exposition of the new approach for detecting a large sample of X-ray clusters based on a combined analysis of the all useful data in the full Chandra and XMM-Newton archives. We have discovered a new unbiased sample of 462 clusters that reaches the most sensitive flux limit achieved to date (~$10^{-14}~\\mbox{ergs}~\\mbox{cm}^{-2}~\\mbox{s}^{-1}$). We summarize the cosmological implications of this new information with an initial limited  approach that assumes the concordance model of the universe is correct and compares the measured Log N-Log S curve of the new X-ray cluster sample with what would be expected in a concordance Universe. The limited scope of this approach precludes a full analysis of the consistency of these new data and the concordance model. We do confirm approximate consistency with the concordance model in two independent ways. First we show that the rapid deviation of the Log N-Log S curve for fluxes less than (~$10^{-13}~\\mbox{ergs}~\\mbox{cm}^{-2}~\\mbox{s}^{-1}$) for a universe without evolution of the density of clusters is consistent with an expected exponential decrease in cluster density for clusters near z $\\approx$ 1 and beyond. This result approximately matches the expected density for clusters of mass range expected at a distance of z $\\approx$ 1 (see figure 15). Secondly, assuming a concordance model, we compute the expected Log N-Log S curve and find approximate agreement with the newly measured Log N-Log S curve (see figure 16). The detailed explanation of these two comparisons assuming a concordance model appears later in this paper. A full analysis beyond approximate consistency with the concordance model is beyond the scope of this initial paper. In particular, we must improve estimates of several possible systematic errors including the volume of the surveyed region in three dimensions, the exact detection threshold near the flux limit, the re-calibration of the temperature, size, luminosity and mass relations of X-ray selected clusters. We present the log N-log S of the unbiased sample of 462 new clusters without redshift measurements because redshifts (and thus luminosities) are not needed for this initial simple analysis. Future work will address the full cosmological implications of these measurements beyond just the Log N-log S curve for clusters. ",
        "conclusions": ""
    },
    "0911/0911.5442_arXiv.txt": {
        "abstract": "{Many questions are still open regarding the structure and the dynamics of the solar core. By constraining more this region in the solar evolution models, we can reduce the incertitudes on some physical processes and on momentum transport mechanisms. A first big step was made with the detection of the signature of the dipole-gravity modes in the Sun, giving a hint of a faster rotation rate inside the core. A deeper analysis of the GOLF/SoHO data unveils the presence of a pattern of peaks that could be interpreted as dipole gravity modes. In that case, those modes can be characterized, thus bringing better constraints on the rotation of the core as well as some structural parameters such as the density at these very deep layers of the Sun interior. } ",
        "introduction": "Over the last 40 years, the study of the solar low-degree acoustic modes has improved our knowledge of the Sun interior \\citep{TouApp1997,ThiBou2000,GarReg2001,2004A&A...413.1135S,2009MNRAS.396L.100B}. In particular, the accurate characterization of these modes has allowed us to constrain very well the radiative zone below the tachocline down to the core, for both, its structure properties \\citep[e.g.][]{2009ApJ...699.1403B} or its dynamics \\citep{ChaEls2001,GarCor2004,2008SoPh..251..119G,2004A&A...425..229M,2005A&A...440..653M}. However, the detection of gravity (g) modes in the Sun is needed to progress in our knowledge of the inner radiative zone below 0.25 $R_\\odot$ where about 50$\\%$ of the mass is concentrated. For example, the detection of a few g modes in the inversions of the rotational profile would greatly improve the inferences at 0.1~$R_\\odot$ \\citep{2008A&A...484..517M, 2009arXiv0902.4142M}.  There is a general consensus that the observational limit for low-degree, low-order p modes is around 1 mHz \\citep{BerVar2000,2008AN....329..461B,2009ApJ...696..653S}. This is due to the raising convective background \\citep[e.g.][]{2008A&A...490.1143L} and the decreasing p-mode amplitudes as the frequency decreases. In order to go towards lower frequencies, we need to improve the signal-to-noise ratio (SNR) by: reducing the convective noise \\citep[e.g.][]{GarJef1999}; combining several instruments \\citep[e.g.][]{2007MNRAS.379....2B,2008arXiv0810.1696S} or using new instruments like PICARD \\citep{2006cosp...36..170T} or GOLF-NG \\citep{2006AdSpR..38.1812T}.  The expected g-mode surface amplitudes  are also very small. However, \\citet{2009A&A...494..191B} have shown that these modes could have surface amplitudes of several mm/s  reaching the detectable regime of the GOLF instrument \\citep{2009arXiv0910.0848A}. Indeed, using this instrument as well as VIRGO -- also on board SoHO -- we have been able to detect signals in the g-mode region that are very unlikely to be due to solar or instrumental noise \\citep{2009ApJS..184..288J}. Even more, some patterns raise above the general convective noise level with a high confidence level \\citep{STCGar2004,2008AN....329..476G}. However, the most important step forward was obtained by applying the same methodology used in asteroseismology to look for the large separation in solar-like stars with a very small SNR \\citep[e.g.][]{2009arXiv0907.0608G}. Thus, by computing the power spectrum of the power spectrum (PSPS), \\citet{2007Sci...316.1591G} were able to unveil the asymptotic periodicity of dipole g modes even without seeing a clear pattern of modes in the power spectrum (as it was, for example, the case of HD175726 where \\citet{2009arXiv0908.2244M} detected the large separation without observing individual modes). The detailed study of this asymptotic periodicity reveals a higher rotation rate in the core of about 3-5 times in average faster than the rest of the radiative region and a better agreement in comparison with solar models computed with old surface abundances instead of the new ones \\citep{2008SoPh..251..135G}. In this work, we study in details the power spectrum density (PSD) of the GOLF instrument to identify the pattern of peaks that produce the detected periodicity in the PSPS. If we succeed, these peaks could be the first g modes that can be correctly identified. Even more, we could analyze their properties, as for instance, their rotational splittings. ",
        "conclusions": ""
    },
    "0911/0911.5674_arXiv.txt": {
        "abstract": "Fifty observations at frequencies between 1.4 GHz and 43 GHz of the 6.6-day O6.5-7+O5.5-6 binary \\object{Cyg OB2 \\#5} using the Very Large Array over 20 years are re-examined. The aim is to determine the location and character of the previously detected variable radio emission. The radio emission from the system consists of a primary component that is associated with the binary, and a non-thermal source (NE), $0.8\\arcsec$ to the NE of the binary that has been ascribed to a wind-collision region (WCR) between the stellar winds of the binary and that of a B-type star (Star D) to the NE .  Previous studies have not accounted for the potential contribution of NE to the total radio emission, most especially in observations where the primary and NE sources are not resolved as separate sources. NE shows no evidence of variation in 23 epochs where it is resolved separately from the primary radio component, demonstrating that the variable emission arises in the primary component. Since NE is non-variable, the radio flux from the primary can now be well determined for the first time, most especially in observations that do not resolve both the primary and NE components. The variable radio emission from the primary component has a period of $6.7\\pm0.3$~years which is described by a simple model of a non-thermal source orbiting within the stellar wind envelope of the binary. Such a model implies the presence of a third, unresolved stellar companion (Star C) orbiting the 6.6-day binary with a period of 6.7 years and independent of Star D to the NE. The variable non-thermal emission arises from either a WCR between Star C and the binary system, or possibly from Star C directly.  The model gives a mass-loss rate of $3.4\\times10^{-5}$~M$_\\odot$\\,yr$^{-1}$ for Cyg OB2 \\#5, unusually high for an Of supergiant and comparable to that of WR stars, and consistent with an unusually strong He\\,{\\sc i} 1.083-$\\mu$m emission line, also redolent of WR stars. An examination of radial velocity observations available from the literature suggests reflex motion of the binary due to Star C, for which a mass of $23^{+22}_{-14}$~M$_{\\odot}$ is deduced.  The natures of NE and Star D are also examined. If NE is a WCR, as suggested by other authors, then the required mass-loss rate is an order of magnitude higher than expected for an early B-type dwarf, and only just consistent with a supergiant. This raises the question of NE as a WCR, but its non-thermal luminosity is consistent with a WCR and a comparison of reddening between \\object{Cyg OB2 \\#5} and Star D do not rule out an association, implying \\object{Cyg OB2 \\#5} is a quadruple system. Pursuing alternative models for NE, such as an unassociated background source, would require very challenging observations. ",
        "introduction": "\\object{Cyg OB2 \\#5} (\\object{V729 Cyg}, \\object{BD $+40\\,4220$}) is an eclipsing binary system consisting of two O-type supergiants orbiting in a 6.6-day period \\citep{Hall:1974, Leung:1978, Rauw:1999}. This system is one of several luminous O-star systems in the Cyg OB2 association that shows evidence of variable radio emission \\citep{Persi:1983, Persi:1990, Bieging:1989} in observations gathered over $\\sim20$~yrs. The radio emission appears to have two states: a low-flux state of $\\sim2$~mJy at 4.8 GHz where the spectral index is consistent with thermal emission from a stellar wind, and a high-flux state of $\\sim8$~mJy at 4.8 GHz, where the spectral index is flatter than in the low state.  The variations appear to have a $\\sim7$-year period \\citep{Miralles:1994} and have been attributed to variable non-thermal emission from an expanding plasmon arising in the binary \\citep{Bieging:1989, Persi:1990, Miralles:1994}.  Observations by \\citet{Abbott:1981} with the VLA revealed two radio components: a primary component associated with the binary and a secondary radio source (hereafter NE) $\\sim0.8\\arcsec$ to the NE of the primary radio source. \\citet{Miralles:1994} confirmed the existence of NE and a $3-6$~cm spectral index of $-0.5\\pm0.3$ hinted at non-thermal emission. \\citet{Contreras:1997} argue that this emission is the result of a wind-collision region (WCR) between the stellar wind of the binary and that of a B0-2 V star (hereafter Star D), $0.9\\arcsec$ from the binary to the NE \\citep{Contreras:1997}.  All previous analyses of the radio emission from \\object{Cyg OB2 \\#5} are based on observations from the Very Large Array (VLA), obtained in all the different configurations of the array, and hence data covering different spatial frequency (i.e. baseline length) ranges.  Dependent on the spatial frequency coverage, the observations may or may not resolve the emission from both the primary and NE components. None of the previous analyses in the literature have attempted to take this into account.  Furthermore, no attempt has been made to determine either the continuum spectrum over a broad wavelength range or the location of the variable emission. As a result, not only is the nature of the radio emission from NE unknown, the evolution of the radio emission of the primary remains uncertain, which may account for the poor fits of models of the non-thermal emission to the observables \\cite[e.g.][]{Miralles:1994}.  In this paper all VLA archive radio observations of \\object{Cyg OB2 \\#5} are re-examined to produce a consistently calibrated data set. The analysis accounts for the changes in the resolution of the array and presents a consistent treatment of the emission from both the primary and NE sources in each observation.  For the first time, the nature of the emission of both the primary and NE can be determined throughout the $\\sim20$ years of the observations. ",
        "conclusions": "\\subsection{The Primary and Variable Emission} \\label{sec:primary_emission} The primary emission component is associated with the O-star binary system and is the source of all the observed variations in the radio emission. In the low state the primary radio emission is found to have a spectral index of $0.60\\pm0.04$ consistent with that expected for thermal emission arising in a steady-state radially symmetric stellar wind. The thermal emission must be reasonably constant in nature since the 350-GHz observations were not both obtained during a radio minimum yet are consistent with the stellar wind spectrum deduced from the radio minimum observations, with a best-fit spectral index $0.63\\pm0.04$ across this broad frequency range (Fig.~\\ref{fig:prim_thermal}).  For a stellar wind, the mass-loss rate can be calculated from the radio flux \\citep[e.g.][]{Wright:1975}. Assuming the stellar wind has a temperature of 10kK and a wind composition with ionic mean charge of 1, mean molecular mass of 1.5 and 1 electron/ion, and assuming all the thermal emission arises from the binary stellar wind (see below for discussion of another potential source of thermal emission), the 4.8-GHz flux at radio minimum of 2.5~mJy leads to a deduced mass-loss rate of \\begin{equation} \\begin{split} \\dot{M}=& 3.4 \\times 10^{-5} \\times \\nonumber\\\\ &\\left( \\dfrac{{\\rm v}_{\\infty}}{1500\\,{\\rm km\\,s}^{-1}} \\right) \\left(\\dfrac{d}{{1.7 \\rm kpc}}\\right)^{3/2}\\dfrac{1}{\\sqrt{F}}~~ M_\\odot \\, {\\rm yr}^{-1}, \\nonumber \\end{split} \\end{equation} where v$_{\\infty}$ is the terminal wind velocity, $F$ is the volume filling factor \\citep[e.g.][]{Abbott:1981}, and $d$ is the distance to the source. The terminal velocity of 1500 km\\,s$^{-1}$ is adopted from \\citet{Conti:1999} based on the absorption component of the P-Cygni profile of the He\\,{\\sc i} 1.083~$\\mu$m line. It should be noted that the emission component of that line is very strong, and broader than those of other supergiants observed by them or in the atlas of \\citet{Groh:2007}, more closely resembling line profiles in WN-type spectra.  The mass-loss rate derived here is consistent with other values determined for \\object{Cyg OB2 \\#5}: $3.3\\times10^{-5}$~M$_\\odot$~yr$^{-1}$ from UV lines \\citep{Howarth:1989}, $2.5\\times10^{-5}$~M$_\\odot$~yr$^{-1}$ from infrared observations \\citep{Persi:1990}, $3.7\\pm1.3\\times10^{-5}$~M$_\\odot$~yr$^{-1}$ based on the 43-GHz flux \\citep{Contreras:1996}, adjusted for the terminal velocity of \\citet{Conti:1999} and for slight differences in adopted distance. More recently the reliability of these methods has been called into question as they tend to overestimate the actual mass-loss rate if the wind is clumpy \\citep[e.g.][]{Mokiem:2007,Fullerton:2006}. The most recent estimate of mass-loss rate in \\object{Cyg OB2 \\#5} is from \\cite{Linder:2009}, who suggest a rate of $2.1\\pm0.6\\times10^{-5}$~M$_\\odot$~yr$^{-1}$ for the binary system based on the observed rate of period change. This is independent of distance and wind clumping but there are other processes which may affect the period. Nevertheless, all the mass-loss rate estimates are consistent with each other, though notably higher than values derived from model atmosphere fits for stars of similar spectral type \\citep{Mokiem:2007}, yet comparable to those of WR stars \\citep{Crowther:2007}. This is consistent with the unusual strength of the He\\,{\\sc i} 1.083-$\\mu$m line in \\object{Cyg OB2 \\#5}, indicating that the system, or one of its components, has a fast, heavy stellar wind comparable to those observed in WR stars.  As the flux increases from the low state toward the high state the continuum spectrum flattens out (see Fig.~\\ref{fig:prim_spec}). This is consistent with the findings of \\citet{Persi:1990}, who attributed the flattening to non-thermal emission from an expanding plasmon associated with the binary system. The model they produced resulted in broad agreement with the general variations in the data, though the fit to their data was poor.  Most recently, 3D-hydrodynamical models of O-star binary systems with periods of a few days have demonstrated that radio flux variations and the flatter spectral index can result from variable {\\em thermal} emission arising in a WCR between the binary stars \\citep{Pittard:2009}. \\object{Cyg OB2 \\#5} is a contact system \\citep{Leung:1978} and the nature of a wind-collision region in such a system is unclear. However, such a region will emit thermal emission that is likely variable. Whether the flux and variation amplitudes can be attained in such a region is discussed further in Sec.~\\ref{sec:thirdstar},  An alternative model for the primary radio emission component is proposed here, where the lower spectral index during high emission is the result of the addition of a non-thermal component to the thermal emission from the binary system giving a ``composite'' spectrum. Such a model has been successfully applied to describe the relatively flat continuum spectra of some Wolf-Rayet stars \\citep[e.g.][]{Chapman:1999} where the non-thermal emission arises in a WCR between the wind of the WR star and that of a massive companion star.  For a system consisting of a non-thermal source embedded in a stellar wind plasma, the total observed flux is given as a function of frequency $\\nu$ and at epoch $t$ by  \\begin{displaymath} S_{obs}(\\nu,t) = S_{th}(\\nu) + S_{nt}(\\nu,t)~~~~~{\\rm mJy}. \\label{eqn:obs_flux} \\end{displaymath} It is assumed the constant thermal emission component, $S_{th}(\\nu)$, has spectral index of $+0.6$ and a flux at 4.8~GHz of 2.5~mJy, deduced from the the primary source during the low emission state. Hence \\begin{displaymath} S_{th}(\\nu) = 2.5 \\left(\\dfrac{\\nu}{4.8}\\right)^{0.6}. \\label{eqn:th_flux} \\end{displaymath} The non-thermal emission component of the total flux, $S_{nt}(\\nu,t)$, is modelled as \\begin{displaymath} S_{nt}(\\nu,t) = S_{4.8}(t) \\left(\\dfrac{\\nu}{4.8}\\right)^{\\alpha} e^{-\\tau(\\nu,t)}, \\label{eqn:nonth_flux} \\end{displaymath} where $S_{4.8}(t)$ is the intrinsic 4.8-GHz flux of the non-thermal source at epoch $t$, $\\alpha$ is the spectral index of the non-thermal emission assumed to be constant, and $\\tau(\\nu,t)$ is the line-of-sight free-free opacity through the stellar wind to the non-thermal source at frequency $\\nu$ and epoch $t$, approximated by \\begin{displaymath} \\tau(\\nu,t) \\approx \\tau_{4.8}(t)\\left(\\dfrac{\\nu}{4.8}\\right)^{-2.1} \\end{displaymath} where $\\tau_{4.8}(t)$ is the 4.8-GHz line-of-sight free-free opacity at epoch $t$.  The line-of-sight opacity is dependent on the geometry of the line-of-sight to the non-thermal emission. Here, the case of a non-thermal source in orbit about the binary is considered. \\citet{Williams:1990} derived the varying free-free opacity along a line-of-sight to a non-thermal source orbiting in the circumbinary wind of the massive WR+O binary WR\\,140. Following Eqns.~12 and 14 in \\cite{Williams:1990}, the opacity is dependent on the orbit inclination ($i$), argument of periastron ($\\omega$), as well as the epoch-dependent true anomaly ($f$) and the separation of the orbiting source from the companion star ($r$) in units of semi-major axis distance ($a$) such that \\begin{equation} \\label{eqn:williams} \\begin{split} \\tau_{4.8}(t)=&\\dfrac{ \\xi \\sec i}{ 2\\Delta r^{3}\\cos^{3}(\\omega +f)} \\times  \\\\ & \\biggl(\\sin (\\omega +f)\\cos (\\omega +f)\\tan i  + \\\\ & (1+\\tan^{2}i)\\arctan \\biggl( \\dfrac{-\\sqrt{\\Delta }}{\\tan (\\omega +f)\\tan i}\\biggr)\\biggr), \\end{split} \\end{equation} where \\begin{displaymath} \\Delta =1+\\tan ^{2}(\\omega +f)+\\tan ^{2}(i) \\end{displaymath} and $\\xi$ is a constant proportional to the square of the ion density in the stellar wind at a radius equal to the semi-major axis $a$, and related to $C_{\\rm ff}$ in \\cite{Williams:1990} by \\begin{displaymath} C_{\\rm ff}=\\xi \\left(\\dfrac{\\nu}{4.8}\\right)^{-2.1}. \\end{displaymath}  The intrinsic non-thermal flux $S_{4.8}(t)$ is expected to depend on the local conditions e.g. electron density, which will vary as the source moves through the dense circumbinary wind. This may be approximated by assuming a simple power-law relation with separation, namely \\begin{displaymath} S_{4.8}(t) = S'_{4.8} r^{-s}, \\label{eqn:s48propto} \\end{displaymath} where $S'_{4.8}$ is the non-thermal flux when the separation is equal to $a$, and $s$ is the power-law index.  These definitions, along with the analytic solution to equation \\ref{eqn:williams} \\citep[cf.][]{Williams:1990}, allow $S_{obs}(\\nu,t)$ to be determined as a function of the orbital phase of the non-thermal source orbiting the binary system.  A standard Levenberg-Marquart $\\chi^{2}$-minimization technique was applied to both the 4.8-GHz and 8.4-GHz fluxes of the primary component to determine values for the seven free parameters in the model for each of the cases $s=0,0.5,1,$ and $2$, assuming an orbital period of 6.7 years. \\cite{Dougherty:2003} suggest $s=0.5$ for the non-thermal luminosity of a WCR. The resulting model parameters are given in Table~\\ref{tab:bestmodel_parms} and the light curves arising from these models are plotted in Fig.~\\ref{fig:spectrum_model}. Fig.~\\ref{fig:phase_model} shows the models for the $s=0$ and $s=0.5$ cases folded into the 6.7-year period.  \\begin{deluxetable*}{cccccccc} \\tablewidth{0pt} \\tabletypesize{\\scriptsize} \\tablecaption{Best model-fit parameters for the orbiting non-thermal source model.\\label{tab:bestmodel_parms}} \\tablehead{\\colhead{$s$}&\\colhead{$S'_{4.8}$}&\\colhead{$\\alpha$}&\\colhead{$\\omega$}&\\colhead{$i$}&\\colhead{$e$}&\\colhead{$T_{0}$}&\\colhead{$\\xi$}\\\\ \\colhead{}&\\colhead{(mJy $a^{s}$)}&\\colhead{}&\\colhead{$(^{\\circ})$}&\\colhead{$(^{\\circ})$}&\\colhead{}&\\colhead{(MJD)}&\\colhead{}} \\startdata \\tableline \\tableline 0  & $5.3\\pm0.5$ & $-0.18\\pm0.25$ & $319\\pm3$ & $90\\pm40$ & $0.69\\pm0.04$ & $53836\\pm35$ & $0.48\\pm0.13$\\\\ 0.5 & $6.4\\pm0.6$ & $-0.34\\pm0.26$ & $315\\pm3$ & $90\\pm25$ & $0.44\\pm0.04$ & $53516\\pm34$ & $0.62\\pm0.12$\\\\ 1 & $7.1\\pm0.6$ & $-0.42\\pm0.28$ & $352\\pm5$ & $88\\pm46$ & $0.23\\pm0.04$ & $53636\\pm30$ & $1.36\\pm0.22$\\\\ 2 & $7.6\\pm0.7$ & $-0.47\\pm0.30$ & $23\\pm5$ & $85\\pm48$ & $0.11\\pm0.05$ & $53737\\pm31$ & $1.64\\pm0.25$\\\\ \\enddata \\end{deluxetable*}  \\begin{figure}[t] \\plotone{figure9a.eps} \\plotone{figure9b.eps} \\caption{\\small The best-fit orbiting non-thermal source model is shown against the observed fluxes of the primary at 8.4 GHz (top) and 4.8 GHz (bottom) for the $s=0, 0.5, 1$, and 2 models.} \\label{fig:spectrum_model} \\end{figure}  \\begin{figure}[t] \\plotone{figure10a.eps} \\plotone{figure10b.eps} \\caption{\\small The best-fit orbiting non-thermal source model is shown for the $s=0$ (top) and $s=0.5$ (bottom) cases at 4.8 GHz (solid line) and 8.4 GHz (dashed line) against the fluxes of the primary, phased with a period of 6.7-year. Parameters for each model are given in Table~\\ref{tab:bestmodel_parms}.} \\label{fig:phase_model} \\end{figure}  Each of these simple models show a good fit to the data and are effective in matching observations at both frequencies across all four observed emission cycles through the 20 years of observation.  The reduced-$\\chi^{2}$ values range from 4.4 to 4.6 indicating that each model has its flaws, with the $s=0$ case being formally the best-fit. The $s=0$ case corresponds to intrinsic non-thermal emission that is constant throughout the orbit i.e. the variation is caused entirely by the varying line-of-sight opacity. In this case the best-fit model is a good fit to the observations except it is too low to match the peak emissions at 8.4 GHz. Introduction of non-zero values for $s$ gives models able to match these sharp emission peaks but the resulting model flux is too low to match the 4.8 GHz observations when the primary flux is in decline at MJD~46309 and MJD~46443 from the high to low emission state, just prior to radio minimum (around deduced orbital phase $\\sim0.1$).  \\citet{Williams:1990} could not fit the radio flux variations in WR\\,140 with this model and a single value of $C_{\\rm ff}$ (corresponding to our $\\xi$). This was attributed to the very different densities of the WR and O-star winds traversed by the line-of-sight at different orbital phases. The quality of the fits to the \\object{Cyg OB2 \\#5} radio-flux variations with a single value of $\\xi$ suggests that the stellar winds in this system have comparable densities, consistent with stars having comparable mass-loss rates rather than the $\\sim30$-fold ratio between the WR and O-star winds in WR\\,140.  In all four models the stellar wind is completely opaque during the low emission state. Unfortunately the radio observation at both frequencies are not evenly distributed across the 20 years of observations and there are large gaps in the phase coverage of the observed light curves. Consequently the model fitting is unevenly focused by the group of observations obtained during the low emission states. Based on the $s=0$ model, the last low emission state occurred around 2007 February, and the next high emission state will occur around 2010 November.  The inclination across all models is consistent with an eclipsing orbit, though this parameter is poorly constrained.  Likewise, even though the non-thermal spectral index tends to become more negative with increasing $s$, it is also poorly constrained. The non-thermal flux $S'_{4.8}$ increases with increasing $s$ as would be expected: a greater flux fall-off with separation would require stronger emission to match the observations. Varying $s$ ranges the time of periastron passage ($T_{0}$) by up to a year between the $s=0$ and $s=0.5$ cases, as evident in Fig.~\\ref{fig:phase_model}. The remaining $T_{0}$  % values fall within this range. By far the greatest impact of varying the $s$ parameter is on the deduced orbit eccentricity. For $s=0$, the orbit eccentricity is high with $e\\sim0.7$, and as $s$ increases the eccentricity decreases sharply, with $e\\sim0.1$ for the $s=2$ case. As noted above, many of the observations were made during the low-state when the stellar-wind is opaque and the non-thermal emission does not contribute to the observed flux being fit by the model. This contributes to the uncertainties in the fitted parameters.  \\subsection{Evidence for a Third Star?} \\label{sec:thirdstar} A non-thermal source orbiting the binary system requires a star (hereafter Star C) to be in a 6.7-year orbit around the binary. This star could contribute the non-thermal radio emission via a WCR arising from the collision of its own stellar wind with the wind from the O+O star binary \\citep[e.g.][]{Eichler:1993}. Such WCRs have been observed directly in some WR+O star and O+O star binary systems \\citep[e.g.][and references therein]{Dougherty:2006}. Alternatively, the non-thermal emission may arise from the putative third star directly, e.g. a compact object.  Given the high luminosities of the two supergiants in the binary and emission from circumstellar material, it will be very hard to detect the proposed third star directly, let alone measure its orbit. Instead, the radial velocities (RVs) of the central binary are examined to search for reflex motion due to it being in an orbit with Star C, as suggested by the radio observations.  The RVs measured by \\citet{Rauw:1999} come from four observing runs, each between 1 and 4 weeks duration and separated by about a year. As the variations should coincide with a period near 6.7 years, each of these runs is treated as a single observation. A fifth observation comes from the first five RVs measured by \\citet{Bohannan:1976} in the space of a month 23 years earlier.  For each RV observed from the primary\\footnote{The secondary has not been used as \\citet{Rauw:1999} find a significantly lower $\\gamma$-velocity for it in their orbit solution and attribute this to formation of the absorption lines in the wind}, the residual (O--C) was calculated from the orbit by \\citet{Rauw:1999} (based on all the RVs) and formed the average (O--C) for each run. These are given in Table \\ref{tab:OmC}, together with the radio orbital phases calculated using P = 6.7 years and $T_{0}$ for the $s=0$ model from Table \\ref{tab:bestmodel_parms}.   \\begin{deluxetable}{ccccc} \\tablewidth{0pt} \\tabletypesize{\\scriptsize} \\tablecaption{Mean RV deviations (O--C) from the O+O orbit as a function of radio phase for five observing runs.\\label{tab:OmC}} \\tablehead{ \\colhead{MJD}&\\colhead{$\\phi$}&\\colhead{n (RVs)}&\\colhead{mean (O--C)}&\\colhead{$\\sigma$} } \\startdata 41150    &  0.81  &   5   &  14.9  &    9.0  \\\\ 49567    &  0.26  &   11  &  -7.6  &    6.7  \\\\ 49914    &  0.40  &   7   &  -2.2  &    6.6  \\\\ 50316    &  0.56  &   3   &  -3.7  &    7.3  \\\\ 50640    &  0.69  &   4   &  8.6   &    8.7  \\\\ \\enddata \\end{deluxetable}  A systematic increase of RV between phases 0.26 and 0.89 is seen, implying that the O+O binary moves away from us more rapidly.  This implies that Star C moves towards us more rapidly in this phase interval so that the circumbinary extinction to the non-thermal radio source diminishes, consistent with it brightening during this orbital phase.  The run of mean (O--C) with phase is compared with the reflex motion of the O+O binary in orbit with Star C following the orbital elements of the embedded non-thermal radio source from the $s=0$ case (see Fig.~\\ref{fig:RVOmC}). Fitting \\begin{displaymath} {\\rm v}_{r} = \\gamma+K_{{\\rm O+O}}\\bigl( e \\cos\\omega + \\cos(f + \\omega) \\bigr) \\end{displaymath} for $K_{{\\rm O+O}}$ and systemic velocity $\\gamma$, gives $K_{{\\rm O+O}} = 32\\pm17$~km\\,s$^{-1}$ and $\\gamma=-5.9\\pm4.7$~km\\,s$^{-1}$, both very uncertain given the uncertainties in the (O--C)s and the shape of the RV curve in the phase range of the observations. The non-zero $\\gamma$ is a consequence of not weighting the mean (O--C)s by the rather unequal numbers of observations from which they were deduced. The mass function $f(m)$ can be derived from $P$ (in days) and $K$ (in km\\,s$^{-1}$) from \\begin{equation} \\begin{split} f(m) = & \\dfrac{m^3_{\\rm C}\\sin^3(i)}{(m_{{\\rm O+O}}+m_{\\rm C})^2} \\\\ = & 1.036 \\times 10^{-7} \\left(1 - e^{2}\\right)^{3/2} K^{3} P, \\nonumber \\label{eqn:massfn} \\end{split} \\end{equation} allowing an estimate of the mass, $m_{\\rm C}$, of Star C. From the data here, $f(m)=3.2^{+8.2}_{-2.8}~{\\rm M}_{\\odot}$. Assuming $\\sin(i)=1$ and adopting $m_{{\\rm O+O}} = 41.5\\pm3.4$~M$_{\\odot}$ from \\citet{Linder:2009}, this gives $m_{\\rm C} = 23^{+22}_{-14}$~M$_{\\odot}$ for Star C. The large uncertainty in the mass stems from the high relative uncertainty in $K$ and the $K^3$-dependence of $f(m)$.  \\begin{figure}[t] \\plotone{figure11.eps} \\caption{\\small Comparison of observed (O--C) residuals, plotted with 2-$\\sigma$ error bars against radio phase, and the RV curve for the reflex motion corresponding to the elements of the $s=0$ model (i.e. with $\\omega$ shifted by 180$\\degr$ to reflect the location of the binary system in the orbit, rather than the 3rd star), with $K=32\\pm17$~km\\,s$^{-1}$ and $\\gamma=-5.9\\pm4.7$~km\\,s$^{-1}$ giving the best-fit of the observables.} \\label{fig:RVOmC} \\end{figure}  With a paucity of observations, most especially at phases of the putative orbit where the radial velocity changes most dramatically, the uncertainties in this analysis are high. However, until further observations can be obtained to test this analysis, it remains a tantalizing piece of evidence that \\object{Cyg OB2 \\#5} is a tertiary system rather than a binary.  Additional support for Star C comes from X-ray data that reveal a hard component. \\cite{Linder:2009} suggest this is likely to arise in a WCR and argue that in such a compact binary like \\object{Cyg OB2 \\#5} the stellar winds of each component are far from achieving terminal velocity and the resulting X-ray emission would be rather soft. Hence they suggest a WCR would be between the binary and another star.  Such a WCR between the binary and Star C provides a ready source for the non-thermal radio emission in the model described in Sec.~\\ref{sec:primary_emission}. Such a mechanism has been widely established for WR stars and many O star systems that exhibit non-thermal emission \\citep{deBecker:2007}.  A mass of $\\sim23$~M$_\\odot$ for Star C is consistent with a late O/early B-type star, which would have a sufficiently strong stellar wind to produce a WCR with the binary system wind in an orbit of size $\\sim14$~AU.  The recent models of thermal emission in WCRs by \\cite{Pittard:2009} raise the possibility that variable plasma density in a WCR between the binary and star C might account for the variable radio emission. Given the separation of the 6.7-year orbit, a WCR between the binary and Star C is undoubtedly adiabatic (certainly away from orbit periastron) and thus any thermal free-free emission in the WCR will be optically thin, with a spectral index of $-0.1$. In combination with the stellar wind continuum from the binary system and Star C, this could result in a continuum spectrum that is flatter than a stellar wind spectrum if the thermal flux from the WCR is sufficiently high. The fluxes in the simulations of \\cite{Pittard:2009} are two to three orders of magnitude less than observed in \\object{Cyg OB2 \\#5}, attributable to the use of mass-loss rates $\\sim$two orders of magnitude less than deduced here for Cyg OB2 \\#5. The optically-thin thermal flux scales as the total number of ions in the WCR, which is $\\propto \\dot M^{2}/D$, where $D$ is the distance from the binary to the WCR. Since the deduced mass-loss rate and separation in \\object{Cyg OB2 \\#5} are respectively $\\sim100$ and $\\sim40$ times those used in the simulations \\citep[cf. model cwb2 in][]{Pittard:2009}, the optically-thin flux from the WCR could be a similar order of magnitude as the observed fluxes. However, in this adiabatic scenario the flux variations scale as $D^{-1}$, and a highly eccentric orbit, as deduced in the $s=0$ model presented above, spends the bulk of the orbit near apastron where $D$ changes little. Hence, the flux would change little, contrary to the observations. Thus variable thermal emission from the WCR alone can not account for the radio light curves in \\object{Cyg OB2 \\#5}, though free-free opacity through the circumbinary wind undoubtedly plays a role. This possibility needs to be explored further, though the types of models described by \\cite{Pittard:2009} are beyond the scope of this paper.  Lastly, an alternative source of the X-ray and non-thermal radio emission could be a compact object such as a neutron star, though the estimated mass from the reflex motion analysis implies this possibility is remote. Though it is not clear how a compact object produces the radio emission, stellar wind accretion onto a $\\sim2$~M$_\\odot$ compact star could generate the observed X-ray luminosity if gravitational potential can be converted to X-ray power efficiently. The gravitational capture radius of a neutron star of mass $M_n$ for a stellar wind of velocity v$_\\infty= 1500$~km\\,s$^{-1}$ is given by \\begin{displaymath} R_g = \\dfrac{2GM_n}{{\\rm v}_\\infty^2}\\sim 2.3\\times 10^{10}~{\\rm cm}. \\end{displaymath} Assuming the gravitational energy of the captured stellar wind is converted to X-ray emission with efficiency $\\epsilon$, then luminosity would be \\begin{displaymath} L_x=\\dfrac{GM_n \\pi R_g^2 F_m(r)\\epsilon}{R_n} \\sim 3.1 \\times 10^{41} F_m(r)\\epsilon~~{\\rm erg~s}^{-1}, \\nonumber \\end{displaymath} where $F_m(r)$ is the mass flux of stellar wind at distance $r$. For an orbit major axis of 14~AU and an eccentricity of 0.7 ($s=0$ model), periastron separation is 4.2~AU at which distance $F_m=4.3\\times10^{-8}$~g\\,cm$^{-2}$\\,s$^{-1}$ for a stellar wind mass-loss rate of $3.4 \\times 10^{-5}$~M$_\\odot$~yr$^{-1}$. This gives $L_x \\sim 3.4$~L$_\\odot$ for $\\epsilon=1$, consistent with the $L_x=1.5$~L$_\\odot$ derived by \\cite{Linder:2009} and adjusted to a distance of 1.7~kpc.  \\subsection{The Secondary Source, NE} The spectral index of $-0.50\\pm0.11$ derived here indicates the radio emission from NE is non-thermal. This was previously suggested by \\citet{Contreras:1996} from a limit to the spectral index of $-2.4\\pm0.6$ deduced from one epoch of observations at 4.8 and 8.4 GHz in 1994 April. \\citet{Contreras:1997} were the first to note that NE lies directly between the \\object{Cyg OB2 \\#5} binary and a B-type star (Star D), $\\sim0.9\\arcsec$ to the NE. This led them to propose that NE is the result of a WCR between the stellar wind of the \\object{Cyg OB2 \\#5} binary system with that from Star D. They further argued that the separation of NE from both the primary and Star D was consistent with the expected relative wind momenta of the binary and the B-type star. This assertion is re-examined here based on seven epochs of 8.4-GHz observations of the primary and NE.  The position of the WCR relative to the positions of the sources of the colliding winds is given by \\begin{displaymath} r_{\\rm O+O} = \\left ( 1 - \\dfrac{ \\eta^{1/2}}{1 + \\eta^{1/2}}\\right ) D \\end{displaymath} \\citep[e.g.][]{Eichler:1993} where $D$ is the separation between the primary O-star binary (or triple) system and the B star, $r_{\\rm O+O}$ is the distance from the WCR to the primary system and $\\eta$ is the the wind-momentum ratio of the two stellar winds given by \\begin{displaymath} \\eta = \\dfrac{\\dot{M}_{\\rm D} {\\rm v}_{\\rm D}} {\\dot{M}_{\\rm O+O} {\\rm v}_{\\rm O+O}}. \\end{displaymath} Here $\\dot{M}_{\\rm O+O}$, $\\dot{M}_{\\rm D}$, ${\\rm v}_{\\rm O+O}$ and ${\\rm v}_{\\rm D}$ are the mass-loss rates and terminal-wind velocities of the O-star binary and Star D respectively. The separation of the O-star binary and the B-type companion is $D=0.93\\pm0.02\\arcsec$ as determined from a weighted average of separations deduced from optical and IR observations (Sec.~\\ref{sec:irobs}).  Combined with $r_{\\rm O+O}=0.77\\pm0.08\\arcsec$ measured here (Sec.~\\ref{sec:proper_motion}), $\\eta = 0.04^{+0.08}_{-0.03}$, where a Monte-Carlo method was used to determine the uncertainty in $\\eta$ as it is an ill-behaved function. Adopting $\\dot{M}_{\\rm O+O} = 3.4 \\times 10^{-5}$~M$_{\\sun}$ yr$^{-1}$ with v$_{\\rm O+O} = 1500$~km\\,s$^{-1}$ \\citep{Conti:1999} leads to \\begin{displaymath} \\dot{M}_{\\rm D} = 0.5{\\rm-}6.1 \\times 10^{-6} \\left( \\dfrac{1000\\,{\\rm km}\\,{\\rm s}^{-1}}{{\\rm v}_{D}} \\right) {\\rm M}_{\\sun} \\, {\\rm yr}^{-1}. % \\end{displaymath} Keeping in mind the mass-loss rate of the binary is high compared with values deduced from model atmospheres, and the wind terminal speed for Star D is unknown, it is noted that only at the lower extremum is the mass-loss rate consistent with those expected for late-O/early-B supergiants, with early-B dwarfs having mass-loss rates around an order of magnitude lower \\citep{Mokiem:2007}, though it is noted that the IR photometry of Star D suggest a star that is more luminous than anticipated for a B0 dwarf. The need to consider the extremum mass-loss rate in order to account for the relative location of the binary, NE and Star D raises the question of whether NE is truly a WCR. However, the Wolf-Rayet system WR\\,147 provides a ready example of an early B-type dwarf star providing a sufficiently dense wind to give a readily observed WCR with the dense wind of a WN8 companion \\citep{Williams:1997}.  The proximity (in space) of Star D and the \\object{Cyg OB2 \\#5} binary can be tested by comparing their respective reddenings. Combining the $K$ magnitude determined in Sec.~\\ref{sec:irobs} with the visual magnitude $V=13.1\\pm0.4$, derived by \\citet{Contreras:1997}, gives $(V-K) = 5.5\\pm0.4$ for Star D, assuming it does not vary.  This implies a reddening of $A_V \\simeq 6.9\\pm0.5$ for Star D, greater than that ($A_V = 5.7\\pm0.3$) implied by $(B-V) = 1.6\\pm0.1$ measured by \\cite{Contreras:1997}. In spite of this apparent discrepancy between visual and visual/IR-determined reddenings, the fact that the reddening estimates bracket that of \\object{Cyg OB2 \\#5} ($A_V = 6.4$), there is no reason to rule out an association between the binary and Star D, and hence the possibility that \\object{Cyg OB2 \\#5} is a quadruple system.  It is possible to estimate the luminosity of a WCR based upon the kinetic energy of the two colliding winds. The surface area of the WCR can be approximated as a spherical cap of diameter $\\pi r_{\\rm D}$ \\citep{Eichler:1993} where $r_{\\rm D} = D - r_{\\rm O+O}$ is the distance from Star D to the WCR. Using this area combined with the kinetic luminosity of the stellar wind of the binary, $L_{\\rm O+O} = 2.4\\times10^{37}$~erg\\,s$^{-1}$, leads to luminosity $L_{WCR} = 6.4 \\times 10^{35}$~erg\\,s$^{-1}$ entering the WCR. The radio synchrotron luminosity $L_{syn}$ arising from the WCR, is estimated from $L_{syn}\\sim10^{-8} L_{WCR}$ \\citep{Chen:1994, Pittard:2006}, giving an estimated synchrotron luminosity from the WCR of $\\sim6\\times10^{27}$ erg\\,s$^{-1}$.  The radio luminosity of NE is estimated by integrating the observed continuum spectrum. Assuming a power-law spectrum between 0.1 - 100 GHz gives a synchrotron luminosity of $2.5\\times10^{26}$~erg\\,s$^{-1}$ for a distance of 1.7~kpc. Considering this is an order-of-magnitude argument, the synchrotron luminosity of the NE is closely consistent with that anticipated from a WCR.  An alternative model for NE is that of a background radio source in chance alignment with \\object{Cyg OB2 \\#5} and Star D. It is difficult to refute this possibility unequivocally, though the probability of such an alignment occurring randomly within 1 arcsecond of Cyg OB2 \\#5 is very low. Extragalactic source counts at 1.4~GHz indicate $\\sim70$~sources~deg$^{-2}$ of around 1~mJy \\citep[e.g.][]{Jackson:2005}, implying $\\sim10^{-6}$ of these sources in 1~arcsecond$^2$. If the source is at a much greater distance than \\object{Cyg OB2 \\#5} it will have a different proper motion to the binary. A distant galaxy will remain fixed relative to the reference frame of background quasars while the binary and Star D will exhibit proper motion relative to the frame. At radio wavelengths, it may be possible to measure this proper motion through astrometry. Certainly, the proper motion determined from the VLA observations discussed here is less than 100 mas, and to improve this precision would require VLBI observations. The low radio brightness of NE presents a challenge for higher precision VLBI astrometry.  Alternatively, IR/optical imaging could reveal if NE is associated directly with an object. \\object{Cyg OB2 \\#5} is bright in the optical (V=9.21) and attempts at imaging the region between the binary system and Star D are thwarted by the high contrast of the binary and the large PSF of the imaging telescopes. The IR image (Sec.~\\ref{sec:irobs}) was searched for evidence of a background source at the location of NE, but with no success.  A smaller PSF could be attained by interferometry, but the contrast could only be defeated through either adaptive nulling or coronographic imaging. This observing challenge remains to be attempted."
    },
    "0911/0911.2438_arXiv.txt": {
        "abstract": " ",
        "introduction": "The Cassini Division is a roughly 4500-km wide region in Saturn's main rings situated between the A and B rings. Far from being a completely empty gap between these two rings, this zone is actually a complex region containing an array of gaps and ringlets. The physical processes responsible for creating and maintaining most of the observed features in this region remain obscure.  In particular, there is still no definitive explanation for most of the numerous gaps present throughout the inner part of Cassini Division. As shown in Figure~\\ref{cdov}, there are eight gaps in the inner part of the Cassini Division, all of which are now named after various researchers who worked on Saturn's rings \\citep{Colbook}. The innermost gap is called the Huygens Gap and it marks the inner boundary of the Cassini Division. The inner edge of the Huygens Gap, which is also the outer edge of the massive B ring, has long been known to be associated with a 2:1 mean-motion resonance with Saturn's moon Mimas \\citep{Porco84}. This gap contains two  optically thick ringlets: the inner one, called the ``Huygens Ringlet\", is known to be eccentric, while the outer, or ``Strange\" ringlet seems to have a significant inclination \\citep{Turtle91, Spitale06, Spitale08}.  While the Mimas 2:1 resonance likely plays an important role in creating the Huygens Gap, the origins of the other seven gaps, as well as the dense ringlets in the Huygens, Herschel and Laplace gaps, are much less clear. Some have argued that each gap in the Cassini Division contains a tiny moon \\citep{Lissauer81}, just as the Encke and Keeler Gaps in the A ring are maintained by the small moons Pan and Daphnis. While some wavelike features in the Cassini Division have been interpreted as evidence for the existence of such moons~\\citep{MT86}, direct detections of the moons themselves have not yet been reported. It is also not clear why such moons would be concentrated in this particular part of the ring system, well inside Saturn's Roche limit for icy bodies.  Using the extensive occultation data obtained by the Visual and Infrared Mapping Spectrometer (VIMS) onboard the Cassini spacecraft, we have conducted an investigation of the gaps in the Cassini Division. The high resolution and accurate geometrical information possible with these data has enabled us to determine the shapes of most of the edges of these gaps and ringlets. Based on this information, we will propose alternative explanations for these gaps that do not require the existence of a small moon within each and every gap.  Section 2 of this paper described the VIMS observations used in this analysis, and how they were processed to obtain position estimates for the various edges. Section 3 presents the results of our analysis and the derived constraints on the shapes of various edges in the Cassini Division, including the B-ring edge. The B-ring edge turns out to be quite complex, so Section 4 briefly compares some of this edge's observed features to the predictions from simple dynamical models. Section 5 presents a dynamical model that could explain the location of the inner edges of the various gaps in the Cassini Divsion as the result of a series of resonances involving perturbations from the massive B ring edge. This model is only a first step towards understanding the architecture of the Cassini Division, so Section 6 discusses potential future work that could confirm and extend this model. The final section summarizes our conclusions.  \\begin{figure}[tb] \\resizebox{6in}{!}{\\includegraphics{cdgapsum_042709.eps}} \\caption{Overview of the basic architecture of the inner part of the Cassini Division. This optical depth profile is derived from the Rev 8 ingress occultation of $o$ Ceti by the rings observed by the VIMS instrument on board Cassini. The various gaps discussed in this study are labeled.} \\label{cdov} \\end{figure} ",
        "conclusions": "\\begin{itemize} \\item The outer edges of five of the eight named gaps in the Cassini Division (that is, all the  gaps that do not contain dense ringlets) are circular to within 1 km.  \\item The inner edges of six of the eight gaps in the Cassini Division (that is, all of the gap inner edges that do not lie near first-order Lindblad resonances with known satellites) are eccentric, with $m=1$ radial variations that drift around the planet at rates close to the expected apsidal precession rates.  \\item The radial excursions in the  inner edge of the Barnard Gap, which lies close to the 5:4 ILR with Prometheus, do appear to have an $m=5$ component tied to that moon.  \\item The radial excursions in the outer edge of the B-ring, which lies close to the 2:1 ILR with Mimas, have an $m=2$ component. However, the amplitude of the pattern and its orientation relative to Mimas change with time. While the available data do not yet determine a unique solution for the motion of the $m=2$ pattern, one possibility is that it librates relative to Mimas with a frequency $\\sim0.06^\\circ$/day and an amplitude of $\\sim 50^\\circ$ in longitude (or a freqeuncy of $\\sim0.12.^\\circ$/day and an amplitude of $\\sim 30^\\circ$).  \\item In addition to the $m=2$ pattern, the B-ring edge also appears to have an $m=1$ component that drifts around the planet at a rate close to the expected apsidal precession  rate at the B-ring edge of $5.06^\\circ$/day.  \\item The pattern speeds of the eccentric features in the Cassini Division, including the Huygens ringlet and the $m=1$ component of the B-ring edge, appear to form a regular series given by $\\Omega_p =\\dot{\\varpi}_B - jL$, where $j=0,1,2,3,...,7$, $\\dot{\\varpi}_B\\simeq 5.06^\\circ$/day is the apsidal precession rate at the B-ring edge, and $L\\simeq 0.06^\\circ$/day is a likely value for the libration frequency of the $m=2$ component in the B-ring edge.  \\item By combining gravitational perturbations from both components of the B-ring edge with the perturbations from the 2:1 Mimas ILR, one can find terms in the equations of particle motion of the form $\\ddot{\\varpi} \\propto \\sin(\\varpi-\\varpi_B+jLt)$. Such terms could act to align the pericenters of particle orbits at the locations of the eccentric inner edges of each of the gaps, and therefore could explain the locations of the gaps in the Cassini Division.   \\end{itemize}"
    },
    "0911/0911.2112_arXiv.txt": {
        "abstract": "Mid-infrared (IR)  imaging at resolutions of 300 mas of the central kpc region of  13 nearby, well-known active galaxies is presented. The bulk of the mid-IR emission is concentrated on an unresolved central source within a size of less   than 5 to 130 pc, depending on the object distance.  Further resolved emission is detected in 70\\% of the sample in the form of circumnuclear star-forming rings or diffuse nuclear extended emission. In the three cases with circumnuclear star formation, the stellar contribution is at least as important as that of the AGN. In those with extended nuclear emission -- a third of the sample -- this emission represents a few per cent of the total measured; however, this contribution may be underestimated because of the chopped nature of these  observations. This extended emission is generally collimated in a preferential direction often coinciding with that of the extended ionized gas or the jet. In M87 and Cen A, where the emission  extends along their respective jets, the emission is presumably synchrotron. In Circinus, NGC 1386 and NGC 3783, it can be reconciled with thermal emission from dust heated at about 100 K by the active nucleus.  In all cases, the nuclear fluxes measured  at 11.8 $\\mu$m and 18.7 $\\mu$m represent a minor contribution of the  flux levels measured by large aperture \\textit{IRAS} data at the nearest energy bands of 12 and 25 $\\mu$m. This contribution  ranges from   30\\% to less than 10\\%.  In only three cases do the active galactic nucleus (AGN) fluxes agree with \\textit{IRAS} to within a factor of 2. In the AGNs with strong circumnuclear star formation, this component can well account for most of the \\textit{IRAS} flux measured in these objects. But in all other cases, either a low surface brightness component extending over galactic scales or strong extra-nuclear IR sources -- e.g. {\\hii} regions in spiral arms -- have to be the main source of the \\textit{IRAS} emission. In either case, the contribution of these components dwarfs that of the AGN at mid-IR wavelengths. ",
        "introduction": "With the advent of interferometry and adaptive optics techniques in the infrared (IR) in large ground-based telescopes, the central regions of active galactic nuclei (AGN) are being studied in ever-increasing detail. The best details have so far been seen on the brightest and nearest AGN observed with the Very Large Telescope ( VLT). Here is a list. The nucleus of NGC 1068 is resolved into a parsec-scale disc in near-IR (Weigelt et al.\\ 2004) and mid-IR (Jaffe et al.\\ 2004) interferometry observations.  That of Circinus is also resolved into a parsec-scale disc-like structure perpendicular to the ionization cone in adaptive optics images in the near IR (Prieto et al.\\ 2004), the results being further confirmed by interferometry in the mid-IR (Tristram et al.\\ 2008). In both cases, the resolved nuclear structure fits within the characteristics of a parsec-scale torus.  The nucleus of Cen A, however, is so far unresolved down to scales of less than a parsec in adaptive optics images in the near IR (Haring-Neumayer et al. 2006) and interferometry in the mid-IR (Meisenheimer et al.\\ 2007).  Other bright AGN observed with adaptive optics in the near IR show an unresolved nucleus, so far, down to scales of tens of parsecs (Prieto et al.\\ 2007), and from those that could be targeted with interferometry in the mid-IR, the available data indicate a resolved nuclear structure, $\\sim 2$ pc in size, and NGC 4151, and possible resolved emission in NGC 1365, NGC 3783 and NGC 7469 (Tristram et al. 2009).  The torus structure may be larger than a few parsecs at mid-IR wavelengths, where dust heated by the AGN to blackbody equivalent temperatures of 100--300 K may exist at larger radii. One of the problems is that this outermost region may be fully resolved with interferometry and thus will escape detection.  This work presents difraction-limited VLT observations at 11.8 $\\mu$m and 18.7 $\\mu$m of a sample of the nearest and brigthest AGN accessible from Paranal Observatory. The resolutions achieved are a factor 10 lower than those provided by VLT-MIDI interferometry, but are more sensitive to possible extended structures around the nucleus (e.g. Perlman et al. 2001; Siebenmorgen et al. 2008); at the very least, they allow for setting an upper limit to the outer radius of the torus. At the higher resolution achieved in these observations (FWHM = 0.3 arcsecs) this upper-limit radius in the sample galaxies is 35 pc on average (ranging from 5 to 130 pc).  This paper is part of an major project focusing on the study of the central few parsecs of the nearest AGN at optical, IR and radio wavelengths, using the highest spatial resolution data available today. The sample galaxies have been extensively studied across all ranges of the electromagnetic spectrum. For all these objects VLT subarcsec resolution observations, by means of adaptive optics in the 1 to 5 $\\mu$m range \\citep{Prieto2009}, and at 11.9 and 18.7 $\\mu$m (this work) were collected.  These data are complemented with \\textit{HST} optical information available for all sources, and VLA and/or ATCA data at subarcsec resolution \\citep{Orienti2009}.  At the distance of the sample galaxies, the spatial scales at which the nuclear regions are studied range from a few pc in the optical, near-IR and radio to several tens of parsecs in the mid-IR.  A comprehensive study of the spectral energy distribution of these AGN using all the available high spatial resolution data is presented in \\citet{Prieto2009}.  Throughout this paper, $H_0$ = 71 {\\kms} Mpc$^{-1}$ is used. ",
        "conclusions": "The mid-IR observations presented in this work, reaching spatial resolutions down to FWHM $\\sim$ 0.3 arcsecs, allow us to constrain the AGN emission at mid-IR wavelengths within regions of 35 pc size in diameter on average (range form 5 to 130 pc).  Within the central kpc region, most of the emission is concentrated in the nuclear region, and in most cases the bulk of it is linked to an unresolved component. Considering the physical scales sampled in these AGN, the outer radius of the torus at mid-IR wavelengths should on average be less than 18 pc, less than 4 pc in Circinus, this being the only galaxy in the sample with a detected parsec scale disc-like structure at its centre at near- and mid-IR wavelengths. Resolved or extended emission is detected in most of the objects within a few arcsec of the centre. The measured contribution of this extended component is a lower limit as part of the emission, particularly if diffuse, may easily be subtracted out in gound-based chopped observations.  Further detection of emission from the galaxy is severely limited by the nod-throw of these observations, typically in the range of 8--10 arcsec.   On the basis of thirteen  AGN studied, the following results are found: \\begin{enumerate} \\item In three AGN (NGC 7582, NGC 7469 and NGC 1097), strong circumnuclear star-forming regions within a few arcsec from the centre are detected.  These are located at radii of $\\sim$130 pc in NGC 7582, 600 pc in NGC 1097 and 900 pc in NGC 7469. In all cases, their associated emission is comparable with, or even larger than, that of the AGN, particularly at 18.7 $\\mu$m.  \\item In six further AGN, extended emission on scales from 1 to 3 arcsec from the centre is detected. This emission is preferentially distributed along a particular direction, usually coinciding with the ionization cone or the jet direction. In all cases, this contribution represents a few per cent of the AGN flux.  \\item Only in four AGN the emission is concentrated into a central unresolved source, and in one galaxy, the Sombrero, no detection is reached.  \\item In comparing the present photometry with previous works, most centered at $\\sim$ 11 $\\mu$m and using broad filters that include PAH features, differences up to 15 to 20\\% in some cases which may be attributed to the contribution of PAH features are found; in M87 and NGC 5507 the agreement is within 5\\% indicating the absence of PHA features in these nuclei and hence of active star formation. Only in NGC 1097 and Cen A, whose nuclear spectrum does not present PAH features, genuine variability by a factor of 2 at 11.8 $\\mu$m in a time scale of 1 to 4 years respectively is indicated. \\end{enumerate}  In the six AGN with extended emission about the centre, that appears tightly collimated.  In Cen~A and M87, it is aligned with the jet (in M87, it coincides with the brightest jet knot, HST-1). In Circinus, NGC 1386, and NGC 3783 the extended emission is co-spatial with the extended ionized gas. Only in NGC 5506, the extended component spreads along the mid-plane of the galaxy being perpendicular to the optical ionization cone main axis.  The mid-IR filters used in this work, narrow and centred on windows where emission lines or PAH features are unimportant, are thus targeted to measure pure continuum emission. Thus, the extended emission in M87 and Cen A is presumably of synchrotron origin; in all other cases, its nature is more ambiguous. Possibilities include free--free emission due to cooling of the ionized gas in the ionization cone \\citep{Contini2004}, or/and emission from dust heated by the AGN. Considering the intrinsic AGN luminosities of these objects, inferred from either high spatial resolution IR spectral energy distribution \\citep{PrietoXian, Prieto2009}: $\\sim 8\\,\\times\\,10^{42}$ erg/sec in Circinus; $\\sim 4\\,\\times\\,10^{44}$ erg/sec in NGC 3783, or X-ray data; $\\sim 1.3\\,\\times\\,10^{42}$ erg/sec in NGC 1386 \\citep{Levenson2006}, and the distances at which mid-IR emission is detected -- in the range of 30 pc in Circinus, to 400 pc in NGC 3783 (see Table 3) -- the nature of this emission could be reconciled with dust directly heated by the AGN provided its equilibrium temperature is in the 100 K range (Barvainis's 1987 formalism is assumed).  The contribution of free-free emission in the mid-IR could also be important if strong shocks exciting the gas are occurring (see, for example, fig. 2c of Contini \\& Viegas-Aldrovandi 1990). Evidence for high gas velocities in the above objects is inferred from the kinematic analysis of their high ionization coronal lines. Specifically, the FWHM of [FeVII] 6087 A is $\\sim$400 km s$^{-1}$ in Circinus and $\\sim$1400 km s$^{-1}$ in NGC 1386 and NGC 3783 \\citep{Rodriguez-Ardila2006Outflow}. Moreover, the size of the extended mid-IR emission closely coincides with the observed sizes of the Fe or Si coronal region (Rodriguez-Ardila et al.) and with their spatial location, which indicates that a fraction of the mid-IR emission is due to free--free emission from predominantly shock-heated coronal gas. Estimating this contribution requires detailed modelling and is currently being explored.    Mid- and far-IR emission of galaxies is usually derived from large-aperture data obtained by IR satellites. For galaxies with an AGN, it is widely assumed that most of this emission comes from the nucleus.  The flux level of the nuclear point source in the AGN studied in this work proves the assumption to be inadequate in most cases. The estimated 11.8 and 18.7 $\\mu$m nuclear fluxes are larger, by factors of $>3$, than the fluxes measured by \\textit{IRAS} at 12 and 25 $\\mu$m in 70\\% of the sample, the largest discrepancy being more than an order of magnitude in five galaxies: NGC 1097, NGC 1566, M87, Cen A and NGC 7582 (in the latter case only at 18.7 $\\mu$m). In the remaining 30\\% of the sample, the \\textit{IRAS} flux levels are still larger but within a factor of 2.  The ``extra'' IR excess measured by \\textit{IRAS} has to come from a source other than the AGN, either strong IR sources, presumably located outside the central 20 arcsec $\\times$ 20 arcsec region -- common f.o.v.\\ in ground-based mid-IR observations -- or from a more extended low surface brightness component across the galaxy. For example, in the three AGN with circumnuclear star-forming regions, the total integrated emission (AGN plus star-forming ring) accounts for the IRAS flux levels within a factor of 2. This is the case of NGC 7582 and NGC 7469; in NGC 1097, the difference is still a factor 10, but in this case the VISIR f.o.v.\\ does not map the complete extension of the ring. In all other AGN with detected nuclear extended emission, this contribution is minor, albeit a lower limit because of the lower sensitivity imposed by chopped observations. The overall conclusion based on this small sample is that in AGN with strong circumnuclear star formation, this component can well account for most of the IRAS flux. For all other cases, either a low surface brightness component extending over galactic scales or strong extranuclear IR sources -- e.g. {\\hii} regions in spiral arms -- are the major contribution to the IR light of these galaxies. In either case, the contribution of these components surpasses by large factors to orders of magnitude that of the AGN at mid-IR wavelengths."
    },
    "0911/0911.3734_arXiv.txt": {
        "abstract": "{} {We investigate the quasar -- radio galaxy unification scenario and detect dust tori within radio galaxies of various types.} {Using VISIR on the VLT, we acquired sub-arcsecond ($\\sim0.40\\arcsec$) resolution N-band images, at a wavelength of 11.85~$\\mu$m, of the nuclei of a sample of 27 radio galaxies of four types in the redshift range $z=0.006-0.156$. The sample consists of 8 edge-darkened, low-power Fanaroff-Riley class~I (FR-I) radio galaxies, 6 edge-brightened, class II (FR-II) radio galaxies displaying low-excitation optical emission, 7 FR-IIs displaying high-excitation optical emission, and 6 FR-II broad emission line radio galaxies. Out of the sample of 27 objects, 10 nuclei are detected and several have constraining non-detections at sensitivities of 7~mJy, the limiting flux a point source has when detected with a signal-to-noise ratio of 10 in one hour of source integration.} {On the basis of the core spectral energy distributions of this sample we find clear indications that many FR-I and several low-excitation FR-II radio galaxies do not contain warm dust tori. At least $57\\pm19$ percent of the high-excitation FR-IIs and almost all of the broad line radio galaxies exhibit excess infrared emission, which must be attributed to warm dust reradiating accretion activity. The FR-I and low-excitation FR-II galaxies are all of low efficiency, which is calculated as the ratio of bolometric to Eddington luminosity $L_{\\mathrm{bol}}/L_{\\mathrm{Edd}}<10^{-3}$. This suggests that thick tori are absent at low accretion rates and/or low efficiencies. The high-excitation FR-II galaxies are a mixed population with three types of nuclei: 1) low efficiency with dust torus, 2) low efficiency with weakly emitting dust torus, and 3) high efficiency with weak dust torus. We argue that the unification viewing angle range 0-45~degrees of quasars should be increased to $\\sim$60~degrees, at least at lower luminosities.} {} ",
        "introduction": "Powerful radio galaxies are understood to host quasar-like nuclei that are obscured by circumnuclear dust tori perpendicular to the radio jets \\citep{1989ApJ...336..606B, 1995PASP..107..803U}. If this unification scenario is correct, infrared observations of the nuclei of radio galaxies may act as calorimeters that indicate the levels of accretion activity. The dust tori, believed to be clumpy and to extend out to radii of tens to hundreds of parsecs \\citep{1988ApJ...329..702K, 2002ApJ...570L...9N,1997ApJ...486..147G, 2006ApJ...653..127S}, reradiate the optical and ultraviolet emission originating from the accretion disks surrounding the supermassive black holes ($>~10^6~M_{\\odot}$). However, only radio galaxies displaying high-excitation narrow-line optical emission, referred to as high-excitation radio galaxies (HEGs), are understood to host obscured broad line and bright continuum emission, and to be in the process of active black hole accretion \\citep{1994ASPC...54..175B, 1994ASPC...54..201L}. Certain radio galaxies exhibit only low-excitation narrow-line emission \\citep{1979MNRAS.188..111H}, referred to as low-excitation radio galaxies (LEGs). Accretion in these objects may be, temporarily or not, at a low level, the line emission being to some part due to hot star ionization. Infrared observations can help to confirm whether these radio galaxies indeed currently host inactive non-thermal nuclei.  A number of mid- and far-infrared surveys of radio galaxies have detected hidden quasars among powerful radio galaxies. These suggest that the scale-height of the torus increases, that is to say the inner radius increases but the outer radius does not change, for more luminous radio galaxies \\citep{2004A&A...424..531H}. Starbursts are also infrequently found to be co-heating the dust \\citep{2007ApJ...661L..13T}. On average, quasars are more powerful than radio galaxies in the mid-infrared because of the beamed emission of the former \\citep{2007ApJ...660..117C}. Infrared surveys have furthermore identified a separate population among radio galaxies of Fanaroff-Riley type II (FR-II), which are mid-infrared weak and probably do not contain hidden quasars \\citep{2001A&A...372..719M, 2006ApJ...647..161O}. These objects are most common at lower radio luminosities and among galaxies that exhibit low-excitation optical emission features. They are predominantly found at redshifts $z<0.3$, and have weaker jets. They may correspond to the excess number of radio galaxies relative to quasars at $z<0.5$ and explain their distributions of size \\citep{1993MNRAS.262L..27S}. It is speculated that they form part of the parent population of variable active nuclei with a beamed component close to the line of sight \\citep{1994ASPC...54..227L}.  This division between radio galaxy populations, according to the radiative properties of the active nucleus, is not linked one-to-one with the well-known radio-power dependent dichotomy in radio morphology \\citep{1974MNRAS.167P..31F}. Low radio luminosity sources are edge-darkened and have radio jets that quickly decelerate and flare (Fanaroff-Riley type I sources, referred to as FR-Is). High luminosity sources are edge-brightened, with jets that remain relativistic over tens to hundreds of kpc, ending in bright hotspots (FR-IIs). The dividing line occurs around a luminosity $\\log{\\nu L_{178\\mathrm{MHz}}} \\sim 40.9~\\mathrm{erg}~\\mathrm{s}^{-1}$. Almost all FR-Is exhibit only low-excitation optical emission lines. FR-IIs have three kinds of excitation features: low-excitation or high-excitation narrow lines, and broad lines. LEGs have dominating low-excitation and/or weak/absent high-excitation emission. They are identified on the basis of one or more of the following criteria: emission-line ratios $[\\ion{O}{III}]/\\mathrm{H}\\beta<5$ and $[\\ion{O}{III}]/[\\ion{O}{II}]<1$, equivalent widths $EW([\\ion{O}{III}])<10~\\AA$ or strengths $\\log{\\nu L_{[\\ion{O}{III}]}}~<~40.7~\\mathrm{erg}~\\mathrm{s}^{-1}$ including for stellar absorption lines or the $4000~\\AA$-break, while for HEGs these emission-line ratios, equivalent widths, and strengths are higher \\citep{1992ApJ...389..208B, 1997MNRAS.286..241J, 1997MNRAS.284..541M}. Broad-line radio galaxies (BLRGs) are identified by having at least one broad emission-line ($v_{\\mathrm{FWHM}}>2000~\\mathrm{km}~\\mathrm{s}^{-1}$) \\citep{1997MNRAS.284..541M}. Since FR-IIs produce about ten to fifty times as much emission-line luminosity as FR-Is for the same radio core power, the engines of FR-Is and FR-IIs differ \\citep{1995ApJ...448..521Z}. It could be that FR-I sources are produced when the central engine is fed by matter at a lower accretion rate, leading to the creation of a source in which the ratio of radiating to jet bulk kinetic energy is low, while FR-II sources are produced when the central engine is fed at a higher accretion rate, causing the central engine to deposit a higher fraction of its energy in radiating energy \\citep{1995ApJ...451...88B}. Alternatively, \\cite{2006ApJ...648L.101E} claim that the torus is produced by an outflow of material from the centre. For a low-luminosity, low-level activity nucleus, this outflow will transform into a radio jet, similar to what is believed to happen in X-ray binaries. This could explain why for the same emission-line activity the FR-Is produce more radio core power than the FR-IIs.  \\begin{figure} \\centering \\includegraphics[width=8.8cm]{12435efg1.ps} \\caption{The core spectral energy distribution of Centaurus~A, fitted by i) a straight power-law $\\nu L_{\\nu} \\propto \\nu^{0.64}$, which cuts off exponentially above $\\nu_{\\mathrm{c}} = 8 \\times 10^{13}$~Hz and becomes optically thick below $\\nu_{\\mathrm{t}} = 5 \\times 10^{10}$~Hz and ii) a parabolic function. The luminosities calculated from observed flux values are indicated with open circles. Filled circles indicate infrared to optical extinction-corrected values assuming a foreground extinction of $A_{\\mathrm{V}}=$14~mag of a circumnuclear dust disk. See text for details.} \\label{censed} \\end{figure}  The thermal and/or non-thermal nature of the cores of low and high power radio galaxies, and their place in unification scenarios can be addressed using multiwavelength studies \\citep{2000A&A...362...75P, 2003AJ....126.2677Q}. An infrared excess above a standard synchrotron spectrum fitted to high frequency radio core data is indicative of hot circumnuclear dust. On the other hand the absence of this excess or short-period flux variablity, is indicative of a purely non-thermal nucleus. On the basis of Hubble Space Telescope (HST) optical and ultraviolet data of the cores of FR-Is, \\cite{1999A&A...349...77C, 2002ApJ...571..247C} and \\cite{2002AJ....123.1334V} show that these nuclei follow a linear correlation in the radio-optical and radio-ultraviolet luminosity plane. This supports the hypothesis that the radio, optical, and ultraviolet emission have a common synchrotron origin. Within this picture, FR-I nuclei are not covered by thick dust tori and have spectral energy distributions (SEDs) that are dominated by non-thermal emission. This is confirmed for M87 \\citep{2004ApJ...602..116W, 2007ApJ...663..808P} and for similar nuclei with low-ionization nuclear emission regions (LINERs) \\citep{2005ApJ...625..699M}. At a one-parsec-level, the FR-I radio galaxy Centaurus~A is shown to be dominated by a synchrotron source, as inferred from mid-infrared interferometric observations with the VLT \\citep{2007A&A...471..453M}. The synchrotron nucleus is argued to be surrounded by a small circumnuclear dust disk, which causes a visual foreground extinction of $A_V = 14$~mag. The overall, dereddened core spectrum of Cen~A is characterized by a $\\nu L_{\\nu} \\propto \\nu^{0.64}$ power-law that cuts off exponentially towards high frequencies at $\\nu_{\\mathrm{c}} = 8 \\times 10^{13}$~Hz and becomes optically thick at $\\nu_{\\mathrm{t}} < 50$~GHz. In Fig.~\\ref{censed}, the observed and extinction-corrected core luminosities, calculated from flux values listed in Table~\\ref{coresed}, are plotted together with the synchrotron spectrum fit, expressed by \\begin{equation} \\centering \\nu L_{\\nu} \\propto \\nu \\left(\\frac{\\nu}{\\nu_{\\mathrm{t}}}\\right)^{\\alpha_{1}} \\left[ 1 - \\exp{-(\\frac{\\nu_{\\mathrm{t}}}{\\nu})^{\\alpha_{1}-\\alpha_{2}}} \\right] \\exp{-\\left(\\frac{\\nu}{\\nu_{\\mathrm{c}}}\\right)} \\,, \\label{ot} \\end{equation} where $\\alpha_1=2.5$ and $\\alpha_2=0.64$ are the optically thick and thin spectral indices \\citep{2000A&A...362...75P}. This spectrum can also be fitted by a parabolic function, which is found to be a good approximation \\citep{1986ApJ...308...78L,2002A&A...381..389A,2009ApJ...701..891L}. Both type of fits and the nuclear mid-infrared luminosity of $\\log{\\nu L_{\\nu}}=41.8~\\mathrm{erg}~\\mathrm{s}^{-1}$ found by \\cite{2004ApJ...602..116W} do not imply a large amount of thermal emission. Although spectral energy distributions seem to indicate that thick dust tori are absent in FR-I radio galaxy nuclei, dust disks, kiloparsecs in size, surrounding FR-I nuclei have been seen in HST images \\citep{1999ApJS..122...81M}, and some contain high molecular gas masses as deduced from high resolution CO maps \\citep{2005ApJ...620..673O,2005ApJ...629..757D}.  The multiwavelength properties of FR-II galaxy cores are more diverse than those of FR-Is. Broad-line radio galaxies (BLRGs) of FR-class II are very bright in the optical: their hosts are often of the N-type. Their nuclei show a large optical excess with respect to the radio-optical correlation found for FR-I galaxies. This excess could well represent unobscured thermal emission from an accretion disk. Broad lines are seen in objects spanning many orders of optical magnitude, from LINERs to quasars. The efficiency of the mass accretion onto the disks varies enormously \\citep{2002A&A...394..791C}. The far-infrared emission of BLRGs, which is much brighter than that of LEGs and HEGs, indicates that they have a special character \\citep{1995A&A...303....8H}. Their place in unification schemes is still unclear. While some may be considered to be weak quasars, others are viewed at grazing incidence to the torus and the jets are directed farther away from us than in the case of quasars \\citep{2000A&A...364..501D}. The warm and cool dust emission indicate that BLRGs have possibly undergone a different merger evolution or have a relatively old active galactic nucleus \\citep{2001A&A...379L..21V}.  Some HEGs show evidence of harboring a hidden quasar nucleus from optical excess in the radio-optical with respect to the radio-optical correlation found for FR-I galaxies. Others do not show this excess, but sometimes spectropolarimetric studies indicate a hidden quasar. Another group of HEGs show, similar to low-excitation radio galaxies (LEG) of type II, true unobscured FR-I-like nuclei \\citep{2002A&A...394..791C}. It seems that LEGs and certain HEGs together constitute a distinct population, characterized by a low accretion rate and/or efficiency, weak or absent broad line emission, and lack of a significant nuclear absorbing structure. These LEGs and HEGs have probably undergone slower cosmic evolution \\citep{2000MNRAS.316..449W} than the HEGs with hidden quasars and BLRGs. Alternatively, it has been argued that the torus covering fraction may increase with decreasing radio luminosity \\citep{1991MNRAS.252..586L}.  To shed light on these issues, we obtained N-band (12~$\\mu$m) VISIR observations of a sample of 8~FR-I, 6 low- and 7 high-excitation FR-II galaxies, and 6 broad-line radio galaxies. We address the thermal or non-thermal nature of their nuclei and their place in the unification scheme. The sample is described in Sect.~2 and the observations in Sect.~3. Core spectral energy distributions and other results are presented in Sect.~4. Discussion and conclusions are presented in Sect.~5 and 6. Throughout this paper, we assume a $\\Lambda$ cold dark matter ($\\Lambda$CDM) cosmology with a Hubble constant $H_0~=~74~\\mathrm{km~s}^{-1}~\\mathrm{Mpc}^{-1}$ and a critical matter and cosmological constant density parameter of $\\Omega_{m}=0.24$ and $\\Omega_{\\Lambda}=0.76$ \\citep{2007ApJS..170..377S}. All luminosities are calculated in the emitted frame, using \\begin{equation} \\centering \\nu L_{\\nu} = \\nu \\frac{4\\pi d^2 F_{\\nu}}{1+z}~\\mbox{erg~s}^{-1} \\,, \\end{equation} where $\\nu$ is the observed frequency, $d$ is the distance calculated for this cosmology and redshift, $z$, and $F_{\\nu}$ is the observed flux.  \\begin{table*} \\caption{Basic parameters of the observed sample.} \\label{table1} \\centering \\begin{tabular}{lllccrccrrrcl} \\hline \\hline Source & $z$  & Class & Ref. & $\\log{\\nu L^{\\mathrm{total}}_{\\mathrm{408MHz}}}$ & $\\log{R}$ & Freq. & Ref. & $\\frac{[\\ion{O}{III}]}{[\\ion{O}{II}]}$ & $\\frac{[\\ion{O}{III}]}{[\\mathrm{H}\\beta]}$ & $\\log{\\nu L_{[\\ion{O}{III}]}}$ & Ref. & VISIR \\\\ &&&&&&&&&&&& Resolution \\\\ &&&&(erg s$^{-1}$)&&(GHz)&&&&(erg s$^{-1}$)&&(parsec) \\\\ (1)&(2)&(3)&(4)&(5)&(6)&(7)&(8)&(9)&(10)&(11)&(12)&(13) \\\\ \\hline 4C12.03      & 0.15600 & LEG    & 5     & 41.98 & $ $ $-$2.66 &   1.5 &  5      &  $\\dots$ &  $\\dots$ & $ $ 40.9 & 7      & 1400 \\\\ 3C015        & 0.07300 & LEG    & 8     & 41.64 & $ $ $-$2.02 &   2.3 &  6      &  $\\dots$ &  3.7     & $ $ 40.4 & 6      &  610 \\\\ 3C029        & 0.04503 & FR-I   & 8     & 41.26 & $ $ $-$1.79 &   2.3 &  6      &  $\\dots$ &  1.7     & $ $ 39.9 & 6      &  370 \\\\ 3C040        & 0.01800 & FR-I   & 8     & 40.63 & $ $ $-$1.61 &   2.3 &  6      &  $\\dots$ &  $\\dots$ & $<$ 38.9 & 6      &  140 \\\\ 3C078        & 0.02865 & FR-I   & 8     & 40.94 & $ $ $-$0.92 &   2.3 &  6      &  $\\dots$ &  $\\dots$ & $<$ 39.1 & 6      &  230 \\\\ Fornax~A     & 0.00587 & FR-I   & 8     & 40.72 & $ $ $-$3.40 &   2.3 &  6      &  $\\dots$ &  $\\dots$ & $<$ 38.6 & 6      &   40 \\\\ 3C093        & 0.35700 & BLRG   & 6     & 43.02 & $ $ $-$2.38 &   5.0 &  2,7    &  3.9     &  $\\dots$ & $\\dots$  & 3      & 3500 \\\\ 3C098        & 0.03045 & HEG    & 5     & 41.31 & $ $ $-$2.70 &   8.3 &  4      &  2.9     & 15       & $ $ 40.9 & 1      &  240 \\\\ 3C105        & 0.08900 & HEG    & 8     & 41.80 & $ $ $-$2.27 &   2.3 &  6      &  $\\dots$ & 16       & $ $ 40.8 & 6      &  750 \\\\ 3C120        & 0.03301 & BLRG   & 8     & 40.73 & $ $    1.01 &   2.3 &  6      &  5.7     &  $\\dots$ & $ $ 41.8 & 6      &  270 \\\\ 3C135        & 0.12738 & HEG    & 5     & 42.05 & $ $ $-$2.71 &   8.3 &  3      &  $\\dots$ &  $\\dots$ & 42.0  & 9         & 1100 \\\\ Pictor~A     & 0.03506 & BLRG   & 8     & 42.22 & $ $ $-$1.57 &   2.3 &  6      &  1.8     &  4.2     & $ $ 41.2 & 6      &  280 \\\\ 3C198        & 0.08147 & HEG    & 3     & 41.56 & $ $ $-$2.20 &   5.0 &  1,7    &  1.2     &  2.1     & $ $ 41.1 & 8      &  680 \\\\ 3C227        & 0.08583 & BLRG   & 8     & 42.13 & $ $ $-$2.55 &   2.3 &  6      & 13       &  9.0     & $ $ 41.8 & 6      &  710 \\\\ 3C270        & 0.00746 & FR-I   & 8     & 40.21 & $ $ $-$1.71 &   2.3 &  6      &  $\\dots$ &  4.7     & $ $ 38.9 & 2,10   &   55 \\\\ Centaurus~A  & 0.00095 & FR-I   & 8     & 40.30 & $ $ $-$1.35 &   2.3 &  6      &  2.5     &  6.3     & $ $ 38.3 & 5,6    &    7 \\\\ PKS1417$-$19 & 0.12000 & BLRG   & 4     & 41.79 & $ $ $-$2.20 &   5.0 &  7,8    & 22       &  2.2     &  $\\dots$ & 4      & 1000 \\\\ 3C317        & 0.03446 & FR-I   & 8     & 41.38 & $ $ $-$0.79 &   2.3 &  6      &  0.37    &  3.1     & $ $ 40.0 & 6      &  280 \\\\ 3C327        & 0.10480 & HEG    & 8     & 42.17 & $ $ $-$2.19 &   2.3 &  6      & 12       & 15       & $ $ 42.1 & 6      &  890 \\\\ 3C353        & 0.03042 & LEG    & 8     & 42.01 & $ $ $-$2.64 &   2.3 &  6      &  $\\dots$ &  0.92    & $ $ 39.3 & 6      &  240 \\\\ PKS1839$-$48 & 0.11120 & LEG    & 8     & 41.99 & $ $ $-$1.16 &   2.3 &  6      &  $\\dots$ &  $\\dots$ & $<$ 39.2 & 6      &  960 \\\\ 3C403        & 0.05900 & HEG    & 8     & 41.59 & $ $ $-$2.18 &   2.3 &  6      &  8.4     & 20       & $ $ 41.6 & 6      &  490 \\\\ 3C403.1      & 0.05540 & LEG    & 7     & 41.04 &  $\\dots$    &$\\dots$&  $\\dots$& $\\dots$  &  $\\dots$ & $\\dots$  & $\\dots$&  450 \\\\ 3C424        & 0.12699 & LEG    & 5     & 42.03 & $ $ $-$1.70 &   8.3 &  3      &  $\\dots$ &  $\\dots$ & $ $ 40.7 & 9      & 1100 \\\\ PKS2158$-$380& 0.03330 & HEG    & 2     & 40.56 & $<$ $-$1.92 &   5.0 &  7,9    & $\\dots$  &  $\\dots$ &  $\\dots$ & $\\dots$&  270 \\\\ 3C445        & 0.05620 & BLRG   & 8     & 41.66 & $ $ $-$1.57 &   2.3 &  8      & 14       & 13       & $ $ 42.0 & 6      &  460 \\\\ PKS2354$-$35 & 0.04905 & FR-I   & 1     & 41.24 & $ $ $-$1.16 &   5.0 &  7,10   & $\\dots$  &  $\\dots$ &  $\\dots$ & $\\dots$&  400 \\\\ \\hline \\end{tabular} \\begin{list}{}{} \\item[]\\textit{Column (2)} Redshifts are from the NED catalogue, except for Centaurus~A, which is calculated from the distance from \\cite{2004A&A...413..903R}. \\textit{Column (3)} Classifications: FR-I stands for Fanaroff-Riley class I radio galaxy, LEG and HEG for low- and high-excitation Fanaroff-Riley class II radio galaxy and BLRG for broad-line radio galaxy. \\textit{Column (4)} References to classifications: 1=\\cite{2006AJ....131..100B}, 2=\\cite{1992ApJ...389..208B}, 3=\\cite{1992MNRAS.256..186B}, 4=\\cite{1978ApJ...220..783G}, 5=\\cite{1998MNRAS.296..445H}, 6=\\cite{1996A&A...313..423H}, 7=\\cite{1997MNRAS.286..241J}, and 8=\\cite{1997MNRAS.284..541M}. \\textit{Column (5)} Integrated radio source luminosities, at rest 408~MHz from \\cite{1996yCat.8015....0W}. \\textit{Columns (6-8)} Radio core fraction, i.e. the ratio of the core to the extended radio flux at 1.5, 2.3, 5.0, or 8.0 GHz frequencies and references: 1=\\cite{2003yCat.8071....0B}, 2=\\cite{1994A&AS..105...91B}, 3=\\cite{1998MNRAS.296..445H}, 4=\\cite{1999MNRAS.304..135H}, 5=\\cite{1991AJ....102..537L}, 6=\\cite{1997MNRAS.284..541M}, 7=\\cite{1996yCat.8015....0W}, 8=\\cite{1995ApJ...448..521Z}, 9=\\cite{1982MNRAS.201..991F}, and 10=\\cite{1991ApJ...376..424S}.  \\textit{Columns (9-12)} Optical emission-line ratios of $[\\ion{O}{III}]$ to $[\\ion{O}{II}]$, $[\\ion{O}{III}]$ to H$\\beta$, and optical $[\\ion{O}{III}]$ emission-line luminosities calculated from flux values found in the literature. References to emission-line ratios and fluxes: 1=\\cite{1977ApJ...211..675C}, 2=\\cite{1996ApJ...470..444F}, 3=\\cite{2003ApJ...599..886E}, 4=\\cite{1978ApJ...220..783G}, 5=\\cite{2007ApJ...670L..81H}, 6=\\cite{1993MNRAS.263..999T}, 7=\\cite{2007A&A...470..531W}, 8=\\cite{1996MNRAS.281..509S}, 9=\\cite{2009A&A...495.1033B}, and 10=\\cite{1997ApJS..112..315H}. \\textit{Column (13)} Scale of the VISIR mid-infrared diffraction-limit of $0.40\\arcsec$ in parsecs. \\end{list} \\end{table*}  \\begin{table*} \\caption{Observation parameters} \\label{tablevisir} \\centering \\begin{tabular}{llllrllrrrlll} \\hline \\hline PKS &  Alias & Date & Dist. &  $F_{11.85~\\mu\\mathrm{m}}^{\\mathrm{core}}$  &  Lit. & Ref. &  $\\log{\\nu L_{11.85~\\mu\\mathrm{m}}^{\\mathrm{core}}}$  &  $\\log{\\nu L_{5\\mathrm{GHz}}^{\\mathrm{core}}}$ &  $\\log{\\nu L_{12~\\mu\\mathrm{m}}^{\\mathrm{total}}}$ & $R50$ & \\multicolumn{2}{c}{FWHM} \\\\ &&&&&&&&&&& Obj. & St. \\\\ & & & Mpc & (mJy) & (mJy) & & ($\\mathrm{erg}~\\mathrm{s}^{-1}$) & ($\\mathrm{erg}~\\mathrm{s}^{-1}$) & ($\\mathrm{erg}~\\mathrm{s}^{-1}$) & (pc) & ($\\arcsec$) & ($\\arcsec$) \\\\ (1)&(2)&(3)&(4)&(5)&(6)&(7)&(8)&(9)&(10)&(11)&(12)&(13)\\\\ \\hline 0007+12        & \\object{4C12.03}      & 2006-06-18   &   710 & $<$    3            &    $\\dots$  & $\\dots$     & $<$ 43.58      & 40.28   &    $\\dots$ &     $\\dots$ &  $\\dots$ &  $\\dots$ \\\\ 0034$-$01      & \\object{3C015}        & 2006-06-18   &   310 & $<$    3            &  6 & 4  & $<$ 42.93      & 41.21   &    $\\dots$ &     $\\dots$ &  $\\dots$ &  $\\dots$ \\\\ 0055$-$01      & \\object{3C029}        & 2006-12-27   &   190 & $<$    4            &  2 & 4  & $<$ 42.56      & 40.28   &    $\\dots$ &     $\\dots$ &  $\\dots$ &  $\\dots$ \\\\ 0123$-$01      & \\object{3C040}        & 2006-12-27   &    70 & $<$    4            &    $\\dots$  & $\\dots$     & $<$ 41.76      & 39.51   &$<$ 43.27   &     $\\dots$ &  $\\dots$ &  $\\dots$ \\\\ 0305+03        & \\object{3C078}        & 2006-12-27   &   120 & $<$    4            &    $\\dots$  & $\\dots$     & $<$ 42.18      & 40.78   &$<$ 43.57   &     $\\dots$ &  $\\dots$ &  $\\dots$ \\\\ 0320$-$37      & \\object{Fornax~A}     & 2006-12-27   &    20 & $<$    9            &    $\\dots$  & $\\dots$     & $<$ 41.03      & 37.81   &    42.44   &     $\\dots$ &  $\\dots$ &  $\\dots$ \\\\ 0340+04        & \\object{3C093 }       & 2006-12-27   &  1800 & $<$    5            &    $\\dots$  & $\\dots$     & $<$ 44.56      & 40.75   &    $\\dots$ &     $\\dots$ &  $\\dots$ &  $\\dots$ \\\\ 0356+10        & \\object{3C098}        & 2005-12-18   &   130 &       20$\\pm$1.4    &     24      & 2           & 42.93$\\pm$0.04 & 38.91   &$<$ 42.95   &   180       &  0.47    &  0.37 \\\\ 0404+03        & \\object{3C105}        & 2005-12-18   &   390 & $<$    2            &    $\\dots$  & $\\dots$     & $<$ 42.93      & 40.06   &    $\\dots$ &     $\\dots$ &  $\\dots$ &  $\\dots$ \\\\ 0430+05        & \\object{3C120}        & 2006-12-27   &   140 &      240$\\pm$15     &    109      & 1           & 44.10$\\pm$0.03 & 41.58   &    44.20   &   170       &  0.36    &  0.33 \\\\ 0511+00        & \\object{3C135}        & 2006-12-27   &   570 & $<$    4            &    $\\dots$  & $\\dots$     & $<$ 43.51      & 39.89   &    $\\dots$ &     $\\dots$ &  $\\dots$ &  $\\dots$ \\\\ 0518$-$45      & \\object{Pictor~A}     & 2006-12-27   &   150 &       68$\\pm$4.3    &    $\\dots$  & $\\dots$     & 43.59$\\pm$0.04 & 41.09   &    43.82   &   200       &  0.34    &  0.33 \\\\ 0819+06        & \\object{3C198 }       & 2006-12-27   &   350 & $<$    4            &    $\\dots$  & $\\dots$     & $<$ 43.13      & 39.13   &    $\\dots$ &     $\\dots$ &  $\\dots$ &  $\\dots$ \\\\ 0945+07        & \\object{3C227}        & 2006-12-27   &   370 &       29$\\pm$3.1    &    $\\dots$  & $\\dots$     & 44.01$\\pm$0.05 & 40.38   & $<$44.96   &   410       &  0.36    &  0.33 \\\\ 1216+06        & \\object{3C270}        & 2006-06-18   &    30 & $<$    2            &       27    & 4           & $<$ 40.72      & 39.17   & $<$42.44   &     $\\dots$ &  $\\dots$ &  $\\dots$ \\\\ 1322$-$42      & \\object{Centaurus~A}  & 2006-12-27   &   3.8 &     1200$\\pm$47     &    1800     & 3           & 41.69$\\pm$0.03 & 38.79   &    42.69   &     5       &  0.37    &  0.33 \\\\ 1417$-$19      & \\object{PKS1417$-$19} & 2006-12-27   &   530 &       26$\\pm$4.3    &    $\\dots$  & $\\dots$     & 44.27$\\pm$0.08 & 39.90   &    $\\dots$ &   770       &  0.44    &  0.33 \\\\ 1514+07        & \\object{3C317}        & 2006-06-18   &   142 & $<$    3            &  3 & 4  & $<$ 42.24      & 40.66   & $<$43.67   &     $\\dots$ &  $\\dots$ &  $\\dots$ \\\\ 1559+02        & \\object{3C327}        & 2005-07-02   &   460 &       78$\\pm$7.5    &    $\\dots$  & $\\dots$     & 44.62$\\pm$0.05 & 40.50   & $<$44.68   &  1300       &  1.25    &  0.42 \\\\ 1717$-$00      & \\object{3C353}        & 2005-07-02   &   125 & $<$    4            &    $\\dots$  & $\\dots$     & $<$ 42.28      & 40.02   &   $\\dots$  &     $\\dots$ &  $\\dots$ &  $\\dots$ \\\\ 1839$-$48      & \\object{PKS1839$-$48} & 2006-06-18   &   490 & $<$    4            &    $\\dots$  & $\\dots$     & $<$ 43.40      & 41.33   &   $\\dots$  &     $\\dots$ &  $\\dots$ &  $\\dots$ \\\\ 1949+02        & \\object{3C403}        & 2006-06-18   &   250 &       89$\\pm$6.8    &    $\\dots$  & $\\dots$     & 44.17$\\pm$0.04 & 39.55   &     44.50  &   330       &  0.34    &  0.37 \\\\ 1949$-$01      & \\object{3C403.1 }     & 2006-06-18   &   230 & $<$    4            &    $\\dots$  & $\\dots$     & $<$ 42.78      & $\\dots$ &    $\\dots$ &     $\\dots$ &  $\\dots$ &  $\\dots$ \\\\ 2045+06        & \\object{3C424  }      & 2006-06-18   &   560 & $<$    4            &  2 & 4  & $<$ 43.46      & 40.48   &    $\\dots$ &     $\\dots$ &  $\\dots$ &  $\\dots$ \\\\ 2158$-$380     & \\object{PKS2158$-$380}& 2006-06-18   &   140 &       17$\\pm$2.8    &    $\\dots$  & $\\dots$     & 42.94$\\pm$0.07 &$<$38.88 & $<$ 43.68  &   200       &  0.37    &  0.33 \\\\ 2221$-$02      & \\object{3C445   }     & 2005-07-02   &   240 &      190$\\pm$16     &    210      &  2          & 44.46$\\pm$0.04 & 40.44   &     44.44  &   390       &  0.42    &  0.42 \\\\ 2354$-$35      & \\object{PKS2354$-$35} & 2006-06-18   &   210 & $<$    3            &    $\\dots$  & $\\dots$     & $<$ 42.55      & 39.08   &    $\\dots$ &     $\\dots$ &  $\\dots$ &  $\\dots$ \\\\ \\hline \\end{tabular} \\begin{list}{}{} \\item[] \\textit{Column (1)} Parkes Radio catalogue reference. \\textit{Column (3)} Date of observation. \\textit{Column (4)} Distances under the chosen cosmology, determined from the redshifts. \\textit{Column (5-7)} VISIR mid-infrared core flux and 3-$\\sigma$ upper limits, at the SiC-band (11.9~$\\mu$m), compared to literature values: 1=\\cite{2004ApJ...605..156G} (TIMMI, 1.5~$\\arcsec$, 10.8~$\\mu$m), 2=\\cite{2004A&A...421..129S} (ISOCAM, 5~$\\arcsec$, 14.3~$\\mu$m), 3=\\cite{2004ApJ...602..116W} (Keck, 0.3~$\\arcsec$, 11.7~$\\mu$m) and 4=\\cite{2009ApJ...701..891L} (Spitzer IRS, 15~$\\mu$m). \\textit{Column (8)} VISIR mid-infrared core luminosity at rest wavelength. \\textit{Column (9)} Radio (5~GHz) core luminosity at rest wavelength from fluxes tabulated in Table~\\ref{coresed}. \\textit{Column (10)} IRAS mid-infrared galaxy luminosity at rest wavelength derived for most sources from flux from \\cite{1988AJ.....95...26G}, for 3C227 from \\cite{1995A&A...303....8H}, and for PKS2158$-$380 from \\cite{1990IRASF.C......0M}. \\textit{Columns (11-13)} The radius at $50\\%$ of the flux, the full width half maximum (FWHM) of the flux of the source and of the standard star. \\end{list} \\end{table*} ",
        "conclusions": "The main conclusions of our sub-arcsecond mid-infrared VISIR survey of the nuclei of 27 radio galaxies of various types are: \\begin{list}{}{} \\item[1.] On the basis of the core spectral energy distributions of our sample, we have found clear indications that many FR-I and several low-excitation FR-II and at maximum 43$\\pm$19~percent of the high-excitation FR-II radio galaxies have non-thermal nuclei, but lack dust tori. \\item[2.] A circumnuclear dust disk is discovered, at a distance of 2-7 parsec from the center of Centaurus~A at a position angle of 100 degrees, which contributes 6 percent of the flux. \\item[3.] The extents of the dust tori around the naked nuclei of broad-line radio galaxies have been measured. In 4 out of 5 cases, these are directed perpendicular to the jet axis. \\item[4.] The warm dust radiation in objects showing a thermal excess is almost entirely concentrated in the inner kpc. \\item[5.] The number of mid-infrared weak sources among the FR-II radio galaxies, 78$\\pm$14 percent, at redshifts $z\\sim0.06$ is not too different from the 55$\\pm$13 percent found by \\cite{2006ApJ...647..161O} at $z<1.0$. Conversely, this indicates a hidden quasar fraction among $z\\sim0.06$ FR-IIs of 22$\\pm$14 percent, compared to 45$\\pm$12 percent at redshifts $z<$1.0. \\item[6.] All FR-Is and all but one low-excitation FR-II galaxies contain nuclei with low efficiencies, $L_{\\mathrm{bol}}$~$/L_{\\mathrm{Edd}}<10^{-3}$. This suggests that the torus disappears at low efficiencies. High-excitation FR-II galaxies are a mixed population, having three types of nuclei: 1) low efficiency with dust torus, 2) low efficiency and weak dust torus, and 3) high efficiency and weak dust torus. \\item[7.] The present and other mid-infrared studies seem to indicate that roughly 50$\\pm$15~percent of powerful radio galaxies at redshifts $z=0.1-1.0$ develop a nuclear broad line region with associated circumnuclear dust. This implies that the unification viewing angle range 0-45~degrees of quasars should be increased to $\\sim$60~degrees, at least at lower luminosities. \\end{list}"
    },
    "0911/0911.3391_arXiv.txt": {
        "abstract": "{The cosmic distance scale largely depends on distance determinations to Local Group galaxies. In this sense, the Andromeda Galaxy (\\object{M\\,31}) is a key rung to better constrain the cosmic distance ladder. A project was started in 1999 to firmly establish a direct and accurate distance to M\\,31 using eclipsing binaries (EBs). After the determination of the first direct distance to M\\,31 from EBs, the second direct distance to an EB system is presented: \\object{M31V\\,J00443610+4129194}. Light and radial velocity curves were obtained and fitted to derive the masses and radii of the components. The acquired spectra were combined and disentangled to determine the temperature of the components. The analysis of the studied EB resulted in a distance determination to M\\,31 of $(m-M)_0=24.30\\pm0.11$~mag. This result, when combined with the previous distance determination to M\\,31, results in a distance modulus of $(m-M)_0=24.36\\pm0.08$~mag ($744\\pm33$ kpc), fully compatible with other distance determinations to M\\,31. With an error of only 4\\%, the obtained value firmly establishes the distance to this important galaxy and represents the fulfillment of the main goal of our project.} ",
        "introduction": "\\label{introdistm31}  Eclipsing binaries (EBs) have always been an important tool for testing and determining the physical properties of stars \\citep{A2008021,A2008029,Torres09}. They are composed of two stars that, when orbiting each other, produce periodic eclipses. The great potential of EBs is that their orbital motion, inferred from the radial velocity curves, and the shape of eclipses, obtained from the light curves, can be entirely explained by the gravitation laws and the geometry of the system \\citep[see][for details]{Hilditch01}.  The direct determination of the radii of the components of EB systems made that several authors \\citep[e.g.,][]{A2008024,A2008020} suggested the possibility of using EBs for deriving distances. The only additional requirement to determine the absolute luminosity of an EB system and, hence the distance, is the surface brightness or, equivalently, the effective temperature of the components.  The potential of EBs to derive distances encouraged several projects to obtain direct distance determinations, either within the Milky Way \\citep[e.g.,][]{A2009043}, the Magellanic Clouds \\citep[e.g.,][and references therein]{A2003004}, or M\\,31/M\\,33 \\citep[DIRECT project,][]{A2006018}. In 1999, a new project was started to obtain a direct distance determination to the Andromeda Galaxy (M\\,31) from EBs \\citep{A2003018,P2004001}, providing the first direct distance determination in \\citet[][hereafter \\citetalias{P2005002}]{P2005002}.  The main interest for an accurate distance determination to M\\,31 lies on the potential of this galaxy to be a first-class distance calibrator \\citep[e.g.,][]{A2006008}. The reasons for this are: {\\em (1)} Contrary to the Magellanic Clouds, the distance to M\\,31 is large enough so that its geometry does not introduce any systematics in the final distance determination; {\\em (2)} typically with a moderate reddening value \\citep[$E(B-V)=0.16\\pm0.01$,][]{A2009002}, it is close enough to enable the individual identification of stars suitable for distance determination (such as EBs or Cepheids); {\\em (3)} an Sb I--II giant spiral galaxy (like M\\,31) provides an appropriate local counterpart for the galaxies commonly used for distance determination \\citep[e.g.,][]{A2005021}; and {\\em (4)} M\\,31 can also provide an absolute calibration of the Tully-Fisher relationship, enabling the calibration of the furthest distance determination methods. Therefore, the characteristics of this spiral galaxy make it an important step of the cosmic distance scale.  As mentioned above, our project already provided the first direct distance determination to M\\,31 \\citepalias{P2005002}. However, the excellent dataset obtained (Sect.~\\ref{obsdistm31}) allowed the determination of additional distances in order to further constrain the distance to M\\,31. Therefore, in the present work, we are presenting the second direct distance determination to M\\,31 from an EB system (Sects.~\\ref{ssrvsb2b}--\\ref{ssevosb2b}). This result (the last one until further spectroscopic observations can be secured) enables a critical discussion on the distance to M\\,31 (Sect.~\\ref{sdistm31}). ",
        "conclusions": ""
    },
    "0911/0911.4052_arXiv.txt": {
        "abstract": "This is the second of two papers presenting a detailed long-slit spectroscopic study of the stellar populations in a sample of 36 ULIRGs. In the previous paper we presented the sample, the data and the spectral synthesis modelling. In this paper we carry out a more detailed analysis of the modelling results, with the aim of investigating the general properties of the stellar populations (i.e. age, reddening and percentage contribution) and the evolution of the host galaxies, comparing the results with other studies of ULIRGs and star forming galaxies in the high-z Universe. The characteristic age of the young stellar populations (YSPs) is $\\leq$ 100 Myr in the nuclei of the overwhelming majority of galaxies, consistent with the characteristic timescale of the major burst of star formation associated with the final stages of major galaxy mergers. However, the modelling results clearly reveal that the star formation histories of ULIRGs are complex, with at least two epochs of star formation activity. Overall, these results are consistent with models that predict an epoch of enhanced star formation coinciding with the first pass of the merging nuclei, along with a further, more intense, episode of star formation occurring as the nuclei finally merge together. It is also found that, although YSPs make a major contribution to the optical emission in most of the extended and nuclear apertures examined, they tend to be younger and more reddened in the nuclear regions of the galaxies. This is in good agreement with the merger simulations, which predict that the bulk of the star formation activity in the final stages of mergers will occur in the nuclear regions of the merging galaxies. In addition, our results show that ULIRGs have total stellar masses that are similar to, or smaller than, the break of the galaxy mass function (i.e. ULIRGs are sub-$m_{\\ast}$ or $\\sim m_{\\ast}$ systems), and that the young stellar populations detected at optical wavelengths dominate the stellar mass contents of the galaxies.  Finally, we find no significant differences between the ages of the YSP in ULIRGs with and without optically detected Seyfert nuclei, nor between those with warm and cool mid- to far-IR colours.  While this results do not entirely rule out the idea that cool ULIRGs with HII/LINER spectra evolve into warm ULIRGs with Seyfert-like spectra, it is clear that the AGN activity in local Seyfert-like ULIRGs has not been triggered a substantial period ($\\geq$100~Myr) {\\it after} the major merger-induced starbursts in the nuclear regions. ",
        "introduction": "First discovered in large numbers by the {\\it IRAS} satellite, ultraluminous infrared galaxies (ULIRGs, L$_{\\rm IR} >$ 10$^{12}$ L$_{\\odot}$) represent some of the most rapidly evolving objects in the local Universe. These objects are important in several respects, such as testing the merger models, investigating the nature of the AGN-starburst connection in merging systems, and understanding the physical processes taking place in the recently discovered high-redshift star forming galaxies, which have similar properties to the nearby ULIRGs.  A remarkable property of ULIRGs is that more than 90\\% of them are associated with galaxy mergers and interactions \\cite[see][for a review]{Sanders96}. The simulations have shown that, during galaxy collisions involving gas-rich disk galaxies, the gas loses angular momentum due to dynamical friction, tidal torques, and inflows towards the centres of the galaxies \\cite[e.g.][]{Mihos96}. Such high nuclear concentrations of gas are capable of triggering both starburst and AGN activity. In general, these simulations predict two epochs of enhanced star formation in major galaxy mergers: the first occurring around the first encounter, and the second episode when the nuclei are close to coalescence \\citep{Mihos96,Barnes96,Springel05,Cox06,PDiMatteo07} . However, both the time lag and the relative intensity of the peaks of starburst activity during the merger event depend on several factors, such as the presence of bulges, feedback effects, gas content and orbital geometry. Given that the models make specific predictions about the histories of the star formation triggered in the course of major gas-rich mergers, studies of the stellar populations in ULIRGs provide useful information about the mergers and, potentially, allow us to test the models.  ULIRGs are classified on the basis of their infrared colours as warm ({\\it f\\/}$_{25}$/{\\it f}$_{60} >$ 0.2)\\footnote{ The quantities {\\it f\\/}$_{25}$ and {\\it f}$_{60}$ represent the {\\it IRAS} flux densities in Janskys at 25$\\mu$m and 60$\\mu$m.} and cool ({\\it f\\/}$_{25}$/{\\it f}$_{60} \\leq$ 0.2) ULIRGs. Warm ULIRGs represent $\\sim$20 -- 25\\% of the total population of ULIRGs discovered by {\\it IRAS}. Most of such objects have an AGN optical spectral classfication, tend to be more compact than cool ULIRGs, and are frequently found in an advanced merger state \\citep{Surace98b}. These properties suggest that ULIRGs play an important role in the formation and evolution of QSOs (Sanders et al., 1988b) and radio galaxies \\citep{Tadhunter05}. Indeed, \\citetalias{Sanders88b} proposed a scenario in which cool ULIRGs evolve into warm ULIRGs on their way to becoming typical optically-selected QSOs. Further evidence in favour of such an evolutionary scenario was reported by \\cite{Surace98b}, \\cite{Surace99}, \\cite{Surace00a} and \\cite{Surace00b} on the basis of their imaging studies of cool and warm ULIRGs, and \\cite{Canalizo00a,Canalizo00b,Canalizo01}, based on a spectroscopic analysis of a sample of transition QSOs --- objects that, based on their infrared colors, may represent a transitionary stage between ULIRGs and QSOs. More recently, \\cite{Tadhunter07} carried out a mid- to far-IR study of powerful radio galaxies, searching for signs of hidden star formation activity. They found that the fraction of powerful radio galaxies with energetically significant star formation activity (20\\% -- 30\\%) is similar to that deduced from optical studies. The derived ages and masses of the young stellar populations in the \\cite{Wills04}, \\cite{Tadhunter05} and \\cite{Tadhunter07} samples of radio galaxies, and the fact that some of these radio galaxies are actually classified as ULIRGs, is consistent with the idea that some radio galaxies are the evolved product of ULIRGs \\citep{Tadhunter05}.  In terms of the end products of the merger events associated with ULIRGs, the kinematic studies of \\cite{Genzel01} and \\cite{Tacconi02} have shown that ULIRGs match the locations of intermediate mass, disky ellipticals and S0 galaxies in the fundamental plane. ULIRGs have effective radii and velocity dispersions smaller than those of their comparison sample of nearby radio galaxies. Furthermore, the dynamical masses of ULIRGs are 10$^{10}$ -- 10$^{11}$ M$_{\\odot}$ \\citep{Tacconi02,Colina05,Dasyra06a,Dasyra06b}, which is similar to intermediate-mass elliptical galaxies. Based on the dynamical properties of ULIRGs, \\cite{Tacconi02} concluded that not all ULIRGs can evolve into typical QSOs or powerful radio galaxies. However, the recent work of \\cite{Dasyra07}, based on near-IR spectroscopic observations of a sample of 12 (mainly Palomar Green) QSOs, has shown that it is possible for ULIRGs to have similar dynamical properties to QSOs. Therefore, it is clear that, despite the recognition of an evolutionary link between ULIRGs, QSOs/radio galaxies and elliptical galaxies, considerable uncertainties remain about the true nature of the link.  Studies of ULIRGs are also important in the context of the various populations of high-z galaxies that have been recently detected using near-IR \\cite[distant red galaxies, DRGs, ][]{Franx03}, mid-IR \\cite[24${\\mu}$m {\\it Spitzer}-selected galaxies, ][and others]{Caputi06a,Yan07,Sajina07}, and sub-millimeter wavelengths \\cite[sub-mm galaxies, SMGs, ][and others]{Smail02,Blain02,Chapman03,Pope05}. These objects have properties similar to the nearby ULIRGs (stellar mass and luminosity, for example), and make up a substantial fraction of the star formation (SF) at z $\\sim$ 1-2. As local analogues of such objects, ULIRGs provide an opportunity to study the physical processes associated with their prodigious SF activity in depth.  It is also important to emphasize that most of the studies of the stellar populations in high redshift galaxies are based on the modelling of a few photometric points and, therefore, it is hard to constrain the properties of these populations. Sensitive studies of star forming galaxies at lower redshifts are useful, not only to better understand the links between local starburst objects and their high-z counterparts, but also to gain valuable information on the uncertainties associated with the derived properties of the stellar populations in high-z star forming galaxies.  At this stage it it is clear that there are many unresolved issues concerning the nature of ULIRGs. In particular:  \\begin{itemize} \\item Do the merger simulations adequately describe the observed properties of ULIRGs? \\item Do the properties of the stellar populations in ULIRGs correlate with other properties of ULIRGs, such as infrared luminosity, interaction class or spectral classification? \\item What is the nature of the links between ULIRGs and other types of object in the local universe, such as QSOs, radio galaxies and elliptical galaxies? \\end{itemize}  Given their importance for studying the evolution of galaxies via major gar-rich mergers, it is perhaps surprising that there have been relatively few studies of stellar populations in ULIRGs. Many previous studies have concentrated on studying the SEDs and mid-IR emission line spectra of ULIRGs, with an emphasis on determining the dominant heating source for the MFIR-emitting dust \\citep[e.g.][]{farrah03}. However, such studies are not capable of determining the detailed properties of the stellar populations and how they vary with spatial location in these complex systems.  Therefore, order to address the key issues surrounding the evolution of ULIRGs we have undertaken a programme of deep optical long-slit spectroscopic observations of a substantial sample of nearby ULIRGs, aimed at investigating their star formation histories. In Rodr\\'iguez Zaur\\'in et al. (2009a), hereafter Paper I, we presented the sample, data reduction and spectral syntehis modelling techniques. In this paper we discuss the results presented in Paper I in the context of evolutionary models for merging systems and the results obtained in studies of stellar populations in high-z star forming galaxies.  Throughout this paper, we assume a cosmology with H$_{0}$ = 71 km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{0} = 0.27$, $\\Omega_{\\Lambda} = 0.73$. ",
        "conclusions": "In this paper we have presented a detailed analysis of the modelling results presented for the CS and the ES samples in Paper I. In addition, we have also discussed these results in a more general context, comparing them with other studies of ULIRGs. The conclusions can be summarized as follows.  \\begin{itemize} \\item{\\bf Age, reddening and percentage contribution}: we find that the YSPs are more significant in the nuclear than in the extended regions. The statistical analysis presented here suggests that the YSPs located in the nuclear regions tend to be younger and more reddened, although further analysis is required to confirm this result. All of these results are consistent with the merger simulations.  \\item{\\bf Correlations}: we find no evidence for correlations between the properties of the stellar populations and other properties of ULIRGs. This can be explained in terms of a selection effect, i.e. since the objects in the CS and the ES samples are ULIRGs, with high IR-luminosities and star formation rates, it is likely that we are observing them at, or after, the first encounter, or at the end of the merger event, when the nuclei are very close to coalescing. Therefore, they all have similar properties. A second possibility to explain the apparent lack of correlations is the scatter in the properties of the YSPs due to uncertainties in the modelling technique. A third possible explanation is that other variables, such as gas content of the parent galaxies and the geometry of the collision, may have a more important impact on the stellar properties of ULIRGs than the optical spectral classification, interaction class or luminosity.  \\item{\\bf Merger simulations}: in general our modelling results are consistent with the merger simulations, with the IYSPs in most systems plausibly associated with the first enhancement of activity following the first encounter of the nuclei of the merging galaxies, and the VYSPs likely related to the final enhancement as the nuclei merge together.  \\item{\\bf 100 Myr cut-off}: when using Comb I, we find a clear cut-off in luminosity weighted YSP ages at 100 Myr for the nuclear apertures, consistent with the timescale of the merger-induced starburst activity predicted by the merger simulations in the final stages of the merger event, as the nuclei coalesce. This result supports the idea that, despite the varying morphologies, we are observing most of the objects close to the peak of the star formation activity in the final stages of the merger event.  \\item{\\bf The masses of the parent galaxies}: both progenitor galaxies in the trigger events must be in the upper 25\\% of the mass range of late-type spirals for most of the ULIRGs in the ES.  \\item{\\bf The morphologies of the parent galaxies}: although in most cases the results are consistent with the idea that the progenitors are late-type spirals (albeit unusually massive late-type spirals), in 3 cases (IRAS 13451+1232, IRAS 21208-0519 and IRAS 23327+2913) there is clear evidence that at least one of the progenitors is an early-type galaxy. This demonstrates that ULIRGs can potentially be triggered by mergers between galaxies with a range of galaxy types.  \\item{\\bf Stellar mass content}: our modelling results suggest that most ULIRGs are sub-$m_{\\ast}$, or $m_{\\ast}$ systems ($m_{\\ast}$ = 1.4 $\\times 10^{11}$ M$_{\\odot}$). Such masses support the idea that these systems will eventually evolve into intermediate mass elliptical galaxies. The masses obtained from the modelling results are generally consistent with the dynamical mass estimates (when available) within a factor of $\\sim$ 2.  This result shows that the stellar populations detected in the optical dominate the stellar masses of the galaxies.  \\item{\\bf Bolometric luminosities:} our results suggest that the stellar popultions detected at optical wavelengths make a significant contribution to the mid- to far-IR luminosities of many of sources. However, due to uncertainties related to the combination of stellar populations assumed, and to the assumption that all the stellar emission is absorbed by, and heats, the dust, it is possible that the values presented in Table 10 in paper I and Figure \\ref{fig:percentage_LIR} in this paper are over-estimates. We conclude that while the visible stellar populations can make a significant contribution to the heating of the dust, it is unlikely that they dominate the heating of the dust in most cases.  \\item{\\bf The evolution of ULIRGs}: the stellar masses of the ULIRGs in the ES sample are consistent with the idea that ULIRGs evolve into intermediate mass elliptical galaxies. On the other hand, the data do not provide any clear evidence to support/refute the evolutionary sequence: cool ULIRG $\\rightarrow$ warm ULIRG $\\rightarrow$ optically selected QSO, since the stellar populations of the cold/warm, Sy/non-Sy, QSO/non-QSO ULIRGs are similar (perhaps the timescale of this transition is shorter than the timescale of the major merger-induced starburst). However, it remains possible that some ULIRGs evolve into radio galaxies (although not all radio galaxies can have this origin).  \\item{\\bf High-z counterparts}: many of the properties of the ULIRGs in our sample, such as morphology, stellar mass content, star formation histories etc., are consistent with the idea that ULIRGs are the local counterparts of the  sub-mm galaxies (SMGs) and distant-red galaxies (DRGs).  \\end{itemize}"
    },
    "0911/0911.0447_arXiv.txt": {
        "abstract": "In this spectroscopic study of infant massive star clusters, we find that continuum emission from ionized gas rivals the stellar luminosity at optical wavelengths.  In addition, we find that nebular line emission is significant in many commonly used broad-band {\\it Hubble Space Telescope (HST)} filters including the F814W $I$-band, the F555W $V$-band and the F435W $B$-band. Two young massive clusters (YMCs) in the nearby starburst galaxy \\ngc\\ were targeted for follow-up spectroscopic observations after  \\citet{Reines08a} discovered an F814W $I$-band excess in their photometric study of radio-detected clusters in the galaxy.  The spectra were obtained with the Dual Imaging Spectrograph (DIS) on the 3.5 m Apache Point Observatory (APO) telescope\\footnote{The 3.5 m Apache Point Observatory telescope is owned and operated by the Astrophysical Research Consortium.} and have a spectral range of $\\sim$3800-9800~\\AA. We supplement these data with \\hst\\ and Sloan Digital Sky Survey (SDSS) photometry of the clusters.  By comparing our data to the Starburst99 and GALEV evolutionary synthesis models, we find that nebular continuum emission competes with the stellar light in our observations and that the relative contribution from the nebular continuum is largest in the $U$- and $I$-bands, where the Balmer (3646 \\AA) and Paschen jumps (8207 \\AA) are located.  The spectra also exhibit strong line emission including the \\SIII\\ $\\lambda \\lambda 9069,9532$ lines in the \\hst\\ F814W $I$-band.  We find that the {\\it combination} of nebular continuum and line emission can account for the F814W $I$-band excess previously found by \\citet{Reines08a}.  In an effort to provide a benchmark for estimating the impact of ionized gas emission on photometric observations of young massive stellar populations, we compute the relative contributions of the stellar continuum, nebular continuum, and emission lines to the total observed flux of a 3 Myr-old cluster through various \\hst\\ filter/instrument combinations, including filters in the Wide Field Camera 3 (WFC3).  We urge caution when comparing observations of YMCs to evolutionary synthesis models since nebular continuum and line emission can have a large impact on magnitudes and colors of young ($\\lesssim 5$ Myr) clusters, significantly affecting inferred properties such as ages, masses and extinctions. ",
        "introduction": "Massive star clusters are an important mode of star formation, having an impact on a wide range of galaxy properties.  However, the earliest stages of these clusters are notoriously challenging to study since the youngest clusters are still enshrouded in remnants of their gaseous and dusty birth cocoons.  Properly accounting for the effects of gas and dust on observations of young massive clusters (YMCs) is nontrivial, rendering it difficult to understand their formation and earliest phases of evolution. Some of these difficulties, such as interstellar extinction, are commonly dealt with in studies of young star clusters.  While it is often not possible to completely disentangle the effects of extinction, the general effects are well-known and accounted for \\citep[e.g.][]{Calzetti94,Whitmore99,Johnson99,Reines08b}.  On the other hand, the treatment of ionized gas emission can be quite complicated and the effects of nebular emission on broad-band observations of YMCs are not always obvious.  Emission lines detected via spectroscopy or narrow-band imaging, however, have proved useful for both identifying young ($\\lesssim 10$ Myr) clusters (e.g. \\ha) and serving as valuable diagnostics on their physical conditions \\citep[e.g.][]{Eissner69,Castaneda92,Kewley02}. If these lines are strong enough, as is typically the case for massive star forming regions, the composite emission can even affect the integrated broad-band photometry \\citep[e.g.][]{Johnson99,Anders03}.  However, nebular emission also takes on a more subtle form via continuum emission (e.g. free-free, free-bound).  In fact, discontinuities or ``jumps'' in the nebular continuum (e.g. the Balmer and Paschen jumps) can be strong enough to serve as diagnostics in their own right \\citep[see][and references therein]{Guseva06}. This nebular continuum emission can also severely affect the broad-band photometry of very young clusters \\citep[e.g.][]{Leitherer95,Molla09}.  \\begin{figure*} \\begin{center} \\includegraphics[scale=0.8]{f1.jpg} \\caption{An SDSS 3-color image of \\ngc\\ (RGB=$ugi$) using a logarithmic scaling.  The two YMCs for which we have spectra are labeled \\citep[Sources 26 and 30 according to the notation of][]{Reines08a}.  The field of view is approximately $3\\arcmin \\times 3\\arcmin$. \\label{image_n4449}} \\end{center} \\end{figure*}  In this paper we investigate the impact of nebular continuum and line emission on observations of very young massive star clusters.  We carry out a detailed multiwavelength study of two YMCs that are still partially embedded in their birth material.  The clusters were targeted for spectroscopic follow-up after \\citet{Reines08a} found an \\hst\\ F814W $I$-band excess in their photometric study of radio-detected clusters in the nearby dwarf starburst galaxy \\ngc.  \\citet{Reines08a} considered several possible explanations for this $I$-band excess, including red supergiants, emission lines, thermal emission from hot dust, continuum emission from a sub-population of deeply embedded stars, and Extended Red Emission \\citep[ERE;][]{Witt04}.  With the observations presented here, we have determined that the {\\it combination} of nebular continuum and line emission can account for the $I$-band excess found by \\citet{Reines08a}.  Consequently, the chief goal of this study is to gain a better empirical understanding of the effects of nebular emission (both continuum and line) associated with extremely young massive star clusters.  This paper is organized as follows:  The data are described in \\S\\ref{sec_data}.  In \\S\\ref{sec_lines} we estimate the physical properties of the clusters from their \\ha\\ emission.  We model the spectral energy distributions (SEDs) in \\S\\ref{sec_model} and investigate the impact of nebular emission (line and continuum) on broad-band photometry in \\S\\ref{sec_broad-band}. A summary of our main conclusions is given in \\S\\ref{sec_summary}. ",
        "conclusions": "\\label{sec_summary}  This paper examines the importance of nebular continuum and line emission in observations of young massive stellar populations.  We have obtained spectroscopy ($\\lambda \\sim 3800-9800$ \\AA) of two young ($\\lesssim 3$ Myr) star clusters in the nearby starburst galaxy \\ngc\\ and supplement these data with archival \\hst\\ and SDSS photometry of the clusters.  We compare our data to the Starburst99 and GALEV evolutionary synthesis models to estimate the physical properties of the clusters and determine the impact of ionized gas emission (line and continuum) on broad-band photometry.  Our main results are summarized below:  \\begin{enumerate}  \\item{The {\\it combination} of nebular continuum and line emission can account for the F814W $I$-band excess found by \\citet{Reines08a} in their study of radio-detected YMCs in \\ngc.  This result is consistent with their prediction that the origin of the excess should not affect clusters older than $\\sim$~5 Myr.}  \\item{At very young ages ($\\lesssim 3$ Myr), nebular continuum emission from ionized gas can {\\it rival} the stellar luminosity of YMCs at optical wavelengths. The relative contribution from the nebular continuum emission is largest in the $U$- and $I$-bands, where the Balmer (3646 \\AA) and Paschen jumps (8207 \\AA) are located.}  \\item{Nebular line emission is significant in many commonly used broad-band \\hst\\ filters including the F814W $I$-band, the F555W $V$-band and the F435W $B$-band.  Emission lines (in addition to nebular continuum) are required to explain the distribution of radio-detected YMCs in \\ngc\\ in the [F435W$-$F550M] versus [F550M$-$F814W] color-color diagram first presented in \\citet{Reines08a}.}  \\item{SED fitting to broad-band photometry of YMCs is most successful when the photometric bands do not contain significant line emission.  In this case, evolutionary synthesis models only containing {\\it continuum} emission from stars and ionized gas are adequate. If the photometric bands include significant line emission, however, models including emission lines perform much better than those that do not.}  \\item{The impact of ionized gas emission on broad-band fluxes, magnitudes and colors can be large and should not be readily dismissed.  We urge caution when comparing observations of young ($\\lesssim 5$ Myr) clusters to synthesis models since nebular continuum and line emission can significantly affect inferred properties such as ages, masses and extinctions.}  \\end{enumerate}"
    },
    "0911/0911.0392_arXiv.txt": {
        "abstract": "The incidence and properties of Active Galactic Nuclei (AGN) in the field, groups, and clusters can provide new information about how these objects are triggered and fueled, similar to how these environments have been employed to study galaxy evolution. We have obtained new \\xmm\\ observations of seven X-ray selected groups and poor clusters with $0.02 < z < 0.06$ for comparison with previous samples that mostly included rich clusters and optically-selected groups. Our final sample has ten groups and six clusters in this low-redshift range (split at a velocity dispersion of $\\sigma = 500$ \\kms). We find that the X-ray selected AGN fraction increases from $f_A(L_X>10^{41}; M_R<M_R^*+1) = 0.047^{+0.023}_{-0.016}$ in clusters to $0.091^{+0.049}_{-0.034}$ for the groups (85\\% significance), or a factor of two, for AGN above an 0.3-8keV X-ray luminosity of $10^{41}$\\ergs\\ hosted by galaxies more luminous than $M_R^*+1$. The trend is similar, although less significant, for a lower-luminosity host threshold of $M_R = -20$ mag. For many of the groups in the sample we have also identified AGN via standard emission-line diagnostics and find that these AGN are nearly disjoint from the X-ray selected AGN. Because there are substantial differences in the morphological mix of galaxies between groups and clusters, we have also measured the AGN fraction for early-type galaxies alone to determine if the differences are directly due to environment, or indirectly due to the change in the morphological mix. We find that the AGN fraction in early-type galaxies is also lower in clusters $f_{A,n>2.5}(L_X>10^{41}; M_R<M_R^*+1) = 0.048^{+0.028}_{-0.019}$ compared to $0.119^{+0.064}_{-0.044}$ for the groups (92\\% significance), a result consistent with the hypothesis that the change in AGN fraction is directly connected to environment. ",
        "introduction": "While it is clear that the Active Galactic Nuclei are powered by accretion onto supermassive black holes (SMBHs), and that this accretion requires both a source of fuel and a trigger to remove angular momentum from the fuel, the origin of the fuel and the trigger mechanism(s) remain poorly understood. For luminous QSOs, major mergers between gas-rich galaxies are considered the dominant mechanism and are the best, and perhaps only, candidate. Many studies have shown that a large fraction of these galaxies are morphologically disturbed, with close neighbors, tidal tails, multiple nuclei, or linked by luminous matter to other galaxies \\citep{gehrens84,hutchings84,malkan84,smith86}. These results have motivated the hypothesis that lower-luminosity AGN in the local Universe are, like QSOs, triggered and fueled by galaxy mergers or at least interactions, even though there is no evidence to support the claim that mergers are the trigger for low-luminosity AGN \\citep{fuentes88,schmitt01}. If gas-rich mergers or interactions are the primary trigger for these lower-luminosity AGN, there should be higher AGN fractions in environments where galaxies have an abundant supply of gas and frequent interactions. While the cluster environment has very high number densities, the galaxies in the centers of rich clusters have less cold gas than those in less dense environments \\citep[e.g.][]{giovanellihaynes85}. In addition, the high pairwise velocity dispersions of cluster members may be too large to allow the formation of bound pairs inside the virial radius \\citep[e.g.][]{ghigna98}. In contrast, galaxies in the field have abundant supplies of cold gas, but the relatively low galaxy number density may counterbalance the abundance of fuel. Between these two extremes, the intermediate group environment could provide the ideal circumstances for the triggering and fueling of low-luminosity AGN in the nearby Universe, at least if mergers and interactions make a substantial contribution \\citep[although see][]{martini04}. Galaxies in groups have sufficiently modest velocity dispersions, numerous neighbors, and available cold gas to trigger and fuel AGN.  Many observations have shown that AGN are rarer in clusters when selected via emission-line diagnostics \\citep{gisler78,dressler85}. Recent observations with the the Sloan Digital Sky Survey have shown this holds across a wide range in galaxy number density, specifically for the most luminous AGN \\citep{miller03,kauffmann04,popesso06}. These results have readily been explained with the argument outlined above, namely that insufficient cold gas reservoirs and high velocity dispersions discourage the triggering and fueling of AGN in the cluster environment. However, these same observations have indicated that the differences between environments fade for less luminous AGN, and additionally samples of lower-luminosity AGN may vary substantially depending on the selection technique employed. The differences in member morphology and physical environment between clusters and the field, as well as the relatively unbiased nature of X-ray observations, prompted \\citet{martini06} to conduct a survey for X-ray AGN in nearby clusters. They found a factor of five higher AGN fraction than previous studies, where the AGN were identified to have X-ray luminosities $L_X \\geq 10^{41}$ \\ergs\\ in host galaxies more luminous than $M_R = -20$ mag. This dramatic increase was attributed to the greater sensitivity of X-ray observations to lower-luminosity AGN relative to visible-wavelength emission-line diagnostics \\citep{bpt81} because the contrast between X-ray emission from AGN and other physical processes (low-mass X-ray binaries, hot gas, and star formation) is higher than in the case of the emission-line diagnostics.  To date only a small number of studies have attempted to extend work on X-ray selected AGN to the lower-velocity dispersion group environment. The first of these was a study of optically-selected groups at $z\\sim0.06$ by \\citet{shen07}. These authors only identified one AGN via X-ray selection out of 140 galaxies in eight groups, yet found five based on emission lines. More recently, \\citet{sivakoff08} compared two, more massive groups and four clusters at similar redshifts and found a significantly higher X-ray selected AGN fraction in the groups compared to the clusters. One of the main goals of the present study is to dramatically improve on the small group sample in the \\citet{sivakoff08} study, as well as to provide a larger sample of more massive groups than those in the \\citet{shen07} study to better span the range of galaxy density from the field to clusters. While the groups in the \\citet{shen07} sample are representative of those found in redshift surveys, they are otherwise a fairly heterogeneous sample. In contrast, the X-ray selection of this sample strongly suggests that these are all virialized systems. This local sample also provides a valuable benchmark for observations of the AGN fraction in groups and clusters at higher redshifts \\citep[e.g.][]{jeltema07,silverman09,martini09}. For example, at $z=0.5-1$ \\citet{georgakakis08b} find a similar fraction of X-ray selected AGN in optically-selected groups as \\citet{shen07} in the local universe.  Our other motivation is to investigate the role of galaxy morphology on the observed AGN fraction, as well as on the selection of AGN. Observations of many galaxies have shown that there is a strong correlation between the mass of the SMBH at the center of a given galaxy and the velocity dispersion or luminosity of the host galaxy's spheroid component \\citep{ferrarese00,gebhardt00,marconi03}. These relationships imply the coevolution of SMBHs and spheroids, and many authors have suggested that AGN may actively impact the evolution of the host galaxy by quenching star formation \\citep{silk98,dimatteo05,hopkins05,springel05}. Since AGN are a consequence of matter accreting onto the central SMBH, and the relation between SMBH mass and bulge properties implies a connection in their evolution, there may be a relation between the incidence of AGN and galaxy morphology. Observations by \\citet{ho97} with the nearby Palomar Seyfert Survey, and more recent work with SDSS \\citep{kauffmann03}, indicate that this is the case. These studies find that Seyferts are preferentially found in early-type spirals with a significant bulge. One potential interpretation of this result is that because early-type galaxies of a given mass will have larger SMBHs compared to late-type galaxies of the same mass, if their SMBHs both accrete at the same fixed fraction of the Eddington rate, then the early-type galaxy is more likely to be detected in a luminosity-limited sample.  This is important for comparing galaxy populations between groups, clusters, and the field because galaxy morphology is observed to be a strong function of local galaxy density \\citep{dressler80}. Furthermore, many of the physical processes that are invoked to explain the relation between galaxy morphology and density may also have relevance for the available fuel supply for AGN. These include mergers and ram-pressure stripping via interactions with the hot ICM \\citep{gunn72,quilis00}, evaporation of the cold ISM by the host ICM \\citep{cowie77}, and starvation of new gas that would otherwise replenish the ISM \\citep{larson80,balogh00}. There are thus correlations between morphology and the incidence of AGN, and between morphology and galaxy environment. It is therefore necessary to take morphology into account in order to determine if any variation in AGN fraction with environment is directly due to the environment itself, or indirectly due to the change in the mix of galaxy types with environment.  To disentangle the effects of morphology and environment we present new observations of rich groups and poor clusters selected from the NORAS sample of \\citet{boehringer00}. These observations and the X-ray AGN classification procedure are described in \\S\\ref{sec:observations}. In addition, we use spectroscopic data from SDSS to classify AGN based on emission-line diagnostics and compare these AGN to those selected via X-ray emission. We then combine these new data with previous work \\citep{martini06,sivakoff08} to study the incidence of AGN as a function of environment and morphology. The morphological classification is described in detail in \\S\\ref{sec:morphology}, and the AGN fractions are presented in \\S\\ref{sec:AGNFractions}. The results, including a statistical analysis, are described in \\S\\ref{sec:discussion}. The final section contains a summary of our main results. Throughout this paper we assume $H_0 = 70$ \\kms\\ Mpc$^{-1}$. ",
        "conclusions": "\\label{sec:summary}  The distribution of AGN as a function of environment is a potentially valuable probe of the fueling and triggering of AGN, as well as the connection between galaxy and black hole evolution. We have conducted a new survey of eleven rich groups and poor clusters selected from the NORAS catalog at $0.02 < z < 0.06$ to measure the AGN fraction as a function of environment. Our group sample, defined to be environments with $\\sigma < 500$~\\kms, contains eight new groups plus two previously published in \\citet{sivakoff08} and thus represents a factor of five increase in sample size. The cluster sample contains three new clusters and three previously published, or a factor of two increase in sample size.  We identify X-ray AGN in these groups and clusters with a combination of spectral fits and flux ratio arguments to demonstrate that the X-ray emission from each X-ray AGN is inconsistent with other sources of X-ray emission. The main result of this analysis is that the X-ray AGN fraction is approximately a factor of two higher in groups than in clusters. This trend is apparent for both AGN in host galaxies more luminous than $M_R = -20$ mag and for more luminous hosts with $M_R \\leq M_R^*+1$, although the difference has greater statistical significance for the higher luminosity threshold. The X-ray AGN fractions for the higher threshold are $f_A(L_X \\geq 10^{41}; M_R \\leq M_R^* + 1) = 0.047^{+0.023}_{-0.016}$ for clusters and $0.091^{+0.049}_{-0.034}$ for groups. There is a 85\\% probability that the group AGN fraction is larger than the cluster AGN fraction. This result may be more significant for the more luminous galaxy subsample due to the larger fraction of the most luminous galaxies that are X-ray AGN.  Because the incidence of AGN in galaxies depends on host galaxy morphology, and the distribution of galaxy morphologies depends on environment, we have conducted the first quantitative morphological analysis of the AGN fraction in dense environments with these six clusters and ten groups. The morphological data for every confirmed group and cluster galaxy was obtained with the GALFIT software package by \\citet{peng02} and used to separate early-type and late-type galaxies. We then calculated the X-ray AGN fraction for just the early-type galaxy populations in the groups and clusters separately and found that the early-type AGN fraction is also a factor of two higher in groups relative to clusters. For the higher galaxy luminosity threshold we find $f_{A,n>2.5}(L_X \\geq 10^{41}; M_R \\leq M_R^*+1) = 0.048^{+0.028}_{-0.019}$ for clusters and $0.119^{+0.064}_{-0.044}$ for groups (92\\% confidence). The similar trends for early-type galaxies and all galaxies indicate that the AGN fraction is not different simply because the morphological mix of galaxies changes as a function of environment, but rather that all galaxy types have a higher probability of hosting an AGN in the group environment. In addition, the group value is also higher than the best estimate of the early-type AGN fraction in the field. This may be because group galaxies, even early-type group galaxies, are more likely to have substantial cold gas reservoirs for AGN fueling than cluster galaxies, while galaxy interactions are more likely to occur in groups than the field, or some combination of these effects.  Finally, we have also estimated the AGN fraction based on emission-line diagnostics for the subset of the galaxies with SDSS spectroscopy. There are 14 BPT AGN in this subset of groups and clusters, as well as six of our nine X-ray AGN. Strikingly, these populations are nearly completely disjoint: only one AGN meets our criterion as both an X-ray AGN and as a BPT AGN. This is a clear demonstration of how different selection techniques may identify different populations of AGN. Our morphological analysis shows that the host galaxies of these two AGN types are marginally different in that the X-ray AGN are more likely to be hosted by early-type galaxies. While the host galaxies for both AGN populations are more luminous than $M_R \\leq -20$ mag, the earlier-type hosts for the X-ray AGN imply relatively larger supermassive black holes compared to the BPT AGN. These observations are thus consistent with lower-efficiency, but relatively more X-ray bright, accretion in the X-ray AGN."
    },
    "0911/0911.5324_arXiv.txt": {
        "abstract": "The damped and sub-damped Lyman-alpha absorption line systems in quasar spectra are believed to be produced by intervening galaxies. However, the connection of quasar absorbers to galaxies is not well-understood, since attempts to image the absorbing galaxies have often failed. While most DLAs appear to be metal-poor, a population of metal-rich absorbers, mostly sub-DLAs, has been discovered in recent studies. Here we report high-resolution K-band imaging with the Keck Laser Guide Star Adaptive Optics system of the field of quasar SDSSJ1323-0021 in search of the galaxy producing the $z=0.72$ sub-DLA absorber. With a metallicity of 2-4 times the solar level, this absorber is of the most metal-rich systems found to date.  Our data show a large bright galaxy with an angular separation of only 1.25$\\arcsec \\,$ from the quasar, well-resolved from the quasar at the high resolution of our data. The galaxy has a magnitude of $K = 17.6-17.9$, which corresponds to a luminosity of $\\approx 3-6 L^{*}$.  Morphologically, the galaxy is fit with a model with an effective radius, enclosing half the total light, of $R_{e} = 4$ kpc and a bulge-to-total ratio of 0.4-1.0, indicating a substantial bulge stellar population. Based on the mass-metallicity relation of nearby galaxies, the absorber galaxy appears to have a stellar mass $\\ga 10^{11} \\, M_{\\odot}$. Given the small impact parameter (9.0  kpc at the absorber redshift), this massive galaxy appears to be responsible for the metal-rich sub-DLA.  The absorber galaxy is consistent with the metallicity-luminosity relation observed for nearby galaxies, but is near the upper end of metallicity. Our study marks the first application of laser guide star adaptive optics for study of structure of galaxies producing distant quasar absorbers. Finally, this study offers the first example of a massive galaxy with a substantial bulge producing a metal-rich absorber. ",
        "introduction": "Quasar absorption line systems provide a unique means to probe  the evolution of interstellar gas to high redshifts ($z \\le 6$).  The  damped Lyman-alpha absorbers (DLAs), with H~I column densities log $N_{\\rm H I} > 20.3$ and the sub-DLAs ($19.0 \\le $ log $N_{\\rm H I} < 20.3$),  constitute a large fraction of the neutral gas in galaxies and are thus believed to be  an important  link in understanding the star formation history of the Universe.   DLAs/subDLAs remain the only class of galaxies at cosmologically significant redshifts with detailed measurements of element abundances and can thus trace chemical evolution of galaxies over $\\ge 90 \\%$ of the cosmic history.  In order to fully use the DLA abundance data to study cosmic chemical evolution in the context of present-day galaxies, it is necessary to combine the abundance information with information on morphologies, luminosities, colors, and star formation rates from direct imaging.    However, it has proven hard to obtain this information for most DLA absorbers and as a result, the type of galaxies giving rise to the DLAs is far from clear.  DLAs are thus variously thought to arise in luminous spirals (e.g., Wolfe et al. 1986),   or gas-rich dwarf galaxies (e.g., York et al. 1986; Matteucci, Molaro, \\& Vladilo 1997),  or merging proto-galactic fragments in cold dark matter cosmologies (e.g., Haehnelt et al. 1998), or low surface brightness galaxies (e.g., Jimenez et al. 1999), or collapsing halos with merging clouds (e.g., McDonald \\& Miralda-Escude 1999).  The lack of substantial chemical evolution found in studies of element abundances in DLAs (e.g., Kulkarni et al. 2005, P\\'eroux et al. 2006b) shows that the currently known population of DLAs seems to be  dominated by metal-poor objects. On the other hand, there are  some recent indications that a few low-$z$ absorbers may be highly enriched--with  near-solar or supersolar metallicities   (see, e.g., Konig et al. 2006; P\\'eroux et al. 2006a; Gharanfoli et al. 2007; Meiring et al. 2009b and references therein).  Such metal-rich absorbers could potentially contribute substantially to the cosmic  metal budget (e.g., P\\'eroux et al. 2006b; York et al. 2006; Prochaska et al. 2006;  Kulkarni et al. 2007; Khare et al. 2007). It is therefore of special interest to image the  metal-rich absorbers.  The primary difficulty in imaging the galaxies producing DLA/sub-DLA absorbers is the  small angular separation of the galaxy from the (much brighter) background quasars.  High angular resolution imaging is thus essential to study the stellar nature of the  absorber galaxies. Adaptive optics (AO) offers a greatly promising technique to achieve the necessary angular resolution. In an earlier work (Chun et al. 2006) we reported the first AO imaging study of a few low-redshift DLAs/sub-DLAs, which uncovered several faint compact objects near the quasars.  Here, we report results from our Keck laser guide star AO observations of the field of SDSSJ1323-0021, a quasar at $z_{em} = 1.4$ with an absorber at $z=0.716$.  This absorber  was chosen for our study because it is one of the highest metallicity absorbers known.  The H I column density was derived by Khare et al. (2004) to be  log $N_{\\rm H I} = 20.21^{+0.21}_{-0.18}$ from a Voigt profile fit to the Lyman-$\\alpha$ absorption line covered in archival HST UV spectra (from HST program GO 9382, PI Rao).  Khare et al. (2004) reported a metallicity of [Zn/H]=+0.4 based on medium-resolution MMT blue channel spectra. To check this extremely high metallicity value, P\\'eroux et al. (2006a) obtained 4.7 km s$^{-1}$ resolution spectra using the Very Large Telescope (VLT)  Ultra-violet Visible Echelle Spectrograph (UVES). Based on multi-component Voigt profile fitting of the metal lines in these high-resolution spectra, P\\'eroux et al. (2006a) verified the high metallicity.  Using the $N_{\\rm H I}$ value derived by Khare et al. (2004), P\\'eroux et al. (2006a) derived [Zn/H] = $+0.61 \\pm 0.20$, [Fe/H] = $-0.51 \\pm 0.20$, [Cr/H] $< -0.52$, [Mn/H] = $-0.37 \\pm 0.20$, and [Ti/H] = $-0.61 \\pm 0.22$. We note that Prochaska et al. (2006) reported a value of log $N_{\\rm H I} = 20.3$ for the same absorber from the same HST data, while Rao et al. (2006) reported log $N_{\\rm H I} = 20.54^{+0.16}_{-0.15}$. The uncertainty in the $N_{\\rm H I}$ value stems from the low S/N in the HST UV spectrum. The metallicity estimate of course would be lower if a higher $N_{\\rm H I}$ value was adopted. In any case, it appears that this absorber is at least 2 times solar in metallicity, and thus one of the most metal-rich absorbers known. Furthermore, the large depletions of Fe, Cr, and Ti relative to Zn indicate that the absorber is dusty.  Throughout this paper, we adopt the metallicity [Zn/H] = +0.61  for the nearly undepleted element Zn derived by P\\'eroux et al. (2006a) using log $N_{\\rm H I} = 20.21^{+0.21}_{-0.18}$, which they point out gives the smallest residual with respect to the HST Lyman-alpha data.  The AO imaging data presented here were obtained with the goal of detecting and better understanding the galaxy underlying the absorber toward SDSSJ1323-0021. The Keck LGSAO system offers higher spatial resolution as well as sensitivity in the near-IR compared to previous imaging observations of DLAs, and thus offers an unprecedented look at a highly interesting absorber galaxy.  Section 2 describes the observations and data reduction.  Section 3 details the procedure adopted for the subtraction of the quasar point spread function. Section 4 describes the results, while section 5  presents a discussion. ",
        "conclusions": "We have detected a large luminous candidate galaxy that is very likely to be the sub-DLA absorber toward quasar SDSSJ1323-0021. To put the properties of galaxy G in context, in Fig. 3 we show the near-IR luminosity-metallicity relation for galaxies in the Kitt-Peak International Spectrographic Survey (KISS) from Salzer et al. (2005). The two panels show two different fits to the near IR luminosity-metallicity ($L$-$Z$) relation from Salzer et al. (2005), based on different calibrations of the relation between emission-line metallicity and the strong-line diagnostic $R_{23}$ . The left panel shows the near-IR $L$-$Z$ relation derived using the $R_{23}$ calibration of Edmunds \\& Pagel (1984) , while the right panel uses the $R_{23}$ calibration of Kennicutt, Bresolin, \\& Garnett (2003). The larger solid filled circle shows the Zn absorption metallicity of the sub-DLA  toward SDSSJ1323-0021 from P\\'eroux et al. (2006a), translated into a corresponding O/H value. We have taken the solar abundance to be log O/H + 12 = 8.66. The absorber toward SDSSJ1323-0021 appears to be at the upper end of the metallicity-luminosity relation. If we adopt the alternative K-band $L$-$Z$ relation derived by Salzer et al. (2005) with the $R_{23}$ calibration of Kennicutt, Bresolin, \\& Garnett (2003), the absorber toward SDSSJ1323-0021 appears to be more extreme compared to the KISS galaxies. In either case, the absorber appears to be a highly luminous galaxy.  This conclusion is not altered even if the metallicity of the absorber is somewhat lower.  Fig. 4 compares the stellar mass vs. metallicity ($M_{*}$-$Z$) relationship for nearby galaxies from Tremonti et al. (2004) with the metallicity of the absorber toward SDSS1323-0021. Even at the upper envelope of the $M_{*}$-$Z$ relation, the metallicity of the absorber (adopted from P\\'eroux et al. 2006a) implies a stellar mass of at least $10^{11} \\, M_{\\odot}$. (We note, however, that the stellar mass would be lower if the metallicity is lower. For example, if the metallicity is lower by 0.3 dex, the most likely stellar mass would be about $10^{10} \\, M_{\\odot}$, as inferred from the most probable  $M_{*}$-$Z$ relation.)  We conclude that the supersolar metallicity absorber toward SDSS1323-0021 arises in a massive $L_{*}$ galaxy at an impact parameter of 9 kpc. The morphological detail seen visually in Fig. 2 suggests that the galaxy has a substantial bulge population. The large bulge-to-total ratio $0.4 \\le B/T \\le 1$ indicates a classical bulge, rather than a pseudo-bulge (e.g., Gadotti 2009).  One may wonder whether a line of sight passing through a galaxy at 9 kpc from the galaxy center can have gas of high metallicity. To which we note that the metallicity at the sun's location (8.5 kpc away from the center of the Milky Way) is solar. Furthermore, the radial metallicity gradients observed in galaxies are weak, and it is not clear whether they even exist at high redshifts, so that the metallicity at a galacto-centric distance of 9 kpc need not be much lower than that near the center of the galaxy. Finally, if the galaxy G is a  galaxy with a substantial bulge (as seems to be the case from the morphological analysis), one might expect it to be substantially metal-rich, and its interstellar material to be fairly uniformly enriched.  It is interesting to compare how the clues regarding the type of the absorber galaxy learnt from our imaging study compare with those derived from element abundance studies. Of particular recent interest is the relative abundance of Mn to Fe. In the Milky Way, [Mn/Fe] is found to increase with increasing [Fe/H]. This trend could indicate that Mn is produced primarily in type Ia supernovae (e.g., Samland 1998). Alternately, the trend could arise due to metallicity-dependent yields in type II supernovae (e.g., Woosley, \\& Weaver 1995; McWilliam et al 2003; Feltzing et al. 2007). Whatever the nucleosynthetic origin of Mn, a trend of increasing [Mn/Fe] with increasing metallicity is also seen in DLAs and sub-DLAs (e.g., Meiring et al. 2009a and references therein). It is possible for dust depletion effects to mask nucleosynthetic effects (e.g. Kulkarni, Fall, \\& Truran 1997; Ledoux, Bergeron, \\& Petitjean 2002). However, as discussed in Ledoux et al. (2002) and Meiring et al. (2009a), the [Mn/Fe] trend is likely to reflect nucleosynthetic effects more, because the dust depletion levels of Mn and Fe are roughly similar in warm halo clouds and not too different in warm disk clouds (e.g., Savage \\& Sembach 1996; Jenkins 2009).  Ledoux et al. (2002) noted that the [Mn/Fe] vs. [Zn/Fe] trend in DLAs could not be fitted with any dust depletion sequence, even if non-solar values are considered for the intrinsic Mn/Fe ratio, without considering variations of the intrinsic [Mn/Fe] ratio.  As noted by Ledoux et al. (2002), the [Mn/Fe] ratio is subsolar for DLAs,  and resembles that in Galactic metal-poor halo stars. For sub-DLAs, the trend of [Mn/Fe] with increasing [Zn/H]  resembles that in Galactic bulge stars (although larger abundance samples are clearly needed to verify this trend). This may suggest that sub-DLAs have an older stellar population that got enriched earlier.   For the sub-DLA toward SDSSJ1323-0021, [Mn/Fe] = $+0.14 \\pm 0.04$ (P\\'eroux et al. 2006a). This [Mn/Fe] value is one of the highest  seen among the absorbers, and lies near the upper right end of the trend in Fig. 11 of Meiring et al. (2009a). Such high [Mn/Fe] values have been seen in metal-rich Galactic bulge stars (e.g., McWlliam et al. 2003; Alves-Brito et al. 2006). Thus, the indication from our AO imaging that the metal-rich absorber toward SDSSJ1323-0021 arises in a galaxy with a substantial bulge appears to be consistent with the relative abundances in the absorber.  Our study offers the first example of the LGSAO technique applied to the study of the structure of the stellar content of quasar absorber galaxies. The discovery of a massive, luminous, galaxy with a substantial bulge in one of the most enriched absorbers known is a striking success in finding galaxy counterparts of quasar absorbers.  In this case the galaxy is easily detected one arcsecond from the quasar even at though the galaxy is 2.5 mag fainter and extended.  Our success in detecting a galaxy counterpart to the supersolar absorber in SDSS1323-0021 is supported by other recent discoveries (though at lower resolution) of candidate galaxies producing other metal-rich absorbers (e.g., Straka et al. 2009; P\\'eroux et al. 2009). Our results are also consistent with the suggestion that most sub-DLAs (at least the metal-rich ones)  may arise in more massive galaxies than DLAs (e.g., Khare et al. 2007; Kulkarni et al. 2009).  Since the discovery of the metal-rich absorber toward SDSSJ1323-0021, several more absorbers (primarily sub-DLAs) with solar or supersolar metallicities have been found in other recent studies (e.g., Prochaska et al. 2006; Kulkarni et al. 2007, 2009; P\\'eroux et al. 2008; Meiring et al. 2009b; and references therein). In future, systematic high spatial resolution imaging of more quasar fields with absorbers of known metallicities will be critical for understanding how the stellar populations of metal-rich and metal-poor galaxies differ. A comparison of the images of galaxies underlying DLA and sub-DLA absorbers will also be crucial in understanding whether these two populations are distinct, and how they are related. Spectroscopic confirmation of the galaxy reported in this study should be possible with an integral field spectrograph using AO. Unfortunately, the sensitivity of the near-infrared integral field spectrographs currently available is not adequate at the wavelengths of key interesting lines at the redshift of 0.716. It will therefore be interesting to obtain confirming observations with an AO-supported spectrograph in the future. A combination of high-resolution imaging, emission-line spectroscopy, and absorption-line spectroscopy for a large number of DLAs/sub-DLAs is essential to fully understand where quasar absorbers fit in the big picture of galaxy evolution.  VPK and SG gratefully acknowledge partial support from the National Science Foundation grant AST-0607739 (PI Kulkarni). VPK is also grateful for partial support from the National Science Foundation grant AST-0908890 (PI Kulkarni). Finally, we thank an anonymous referee for constructive comments that helped to improve the paper."
    },
    "0911/0911.2732_arXiv.txt": {
        "abstract": "Previous \\hst\\ and \\fuse\\ observations have revealed highly ionized high-velocity clouds (HVCs) or more generally low \\hi\\ column  HVCs along extragalactic sightlines over 70--90\\% of the sky. The distances of these HVCs have remained largely unknown hampering to distinguish  a ``Galactic\" origin (e.g., outflow, inflow) from a ``Local Group\" origin (e.g., warm-hot intergalactic medium). We present the first detection of highly ionized HVCs in the Cosmic Origins Spectrograph (COS) spectrum of the early-type star HS\\,1914+7134 ($l = 103\\degr$, $b=+24 \\degr$) located in the outer region of the Galaxy at $d \\simeq 14.9$ kpc. Two HVCs are detected in absorption at $v_{\\rm LSR} = -118$ and $-180$ \\km\\ in several species, including \\civ, \\siiv, \\siiii, \\alii, \\cii, \\siii, \\oi, but \\hi\\ 21-cm emission is only seen at $-118$ \\km. Within 17\\degr\\ of  HS\\,1914+7134, we found HVC absorption of low and high ions at similar velocities toward 5 extragalactic sightlines, suggesting that these HVCs are related. The component at $-118$ \\km\\ is likely associated with the Outer Arm of the Milky Way. The highly ionized HVC at $-180$ \\km\\  is an HVC plunging at high speed onto the thick disk of the Milky Way. This is the second detection of highly ionized HVCs toward Galactic stars, supporting a ``Galactic\" origin for at least some of these clouds. ",
        "introduction": "The Milky Way is surrounded by large gaseous complexes observed at high velocities ($|\\mbox{\\vlsr}| \\ga 90$ \\km), known as HVCs. HVCs were traditionally regarded as neutral entities \\citep[e.g.,][]{wakker97}, but this view has given way to one that recognizes the importance of an ionized component to many of these clouds thanks to high quality UV observations from {\\hst} and \\fuse\\ \\citep{sembach95,sembach99,sembach03,lehner01a,lehner02,collins04,collins05,collins09,ganguly05,fox04,fox05,fox06,richter09,shull09}. Many of these  HVCs are almost completely ionized and show  absorption from the ``high ions'' \\ovi, \\civ, and \\siiv.  These are therefore deemed highly ionized HVCs, although they often show absorption  not only in the high ions, but also in the lower ions (e.g., \\cii, \\ciii) and even atoms (\\hi, \\oi).  A better designation would be multiphase HVCs or low \\hi\\ column HVCs \\citep[see also][]{fox06,collins09}. The covering factor of these low \\hi\\ column density HVCs is extremely high and much larger than that of the \\hi\\ 21-cm HVCs. \\fuse\\ surveys of AGNs and QSOs show that 60--70\\%\\ (perhaps even 85\\%) of the high-latitude sky is covered by \\ovi\\ HVCs, the majority of which ($\\sim$75\\%) have no \\hi\\ emission counterpart \\citep{sembach03,fox06}. \\citet{shull09} and \\citet{collins09} show that high-velocity \\siiii\\ absorption has a detection rate of 80--90\\%\\ at high latitude. These numbers are to be contrasted to the \\hi\\ 21-cm emission ($N($\\hi$)\\ga 10^{17.9}$ cm$^{-2}$) covering factor of $\\sim$37\\% \\citep{murphy95}.  Both galactic and larger-scale phenomena/structures have been invoked for the origin(s) of these HVCs, including galactic feedback processes, disk-halo mass exchange (e.g., through the accretion of matter condensing from an extended corona), or the shock-heated warm-hot intergalactic medium (WHIM) of the Local Group \\citep[e.g.,][]{blitz99,sembach03,nicastro03,collins05,fox06}. Pros and cons of both theories have been discussed, but they are circumstantial and tentative, relying on assumptions of poorly-constrained quantities, such as metallicities and ionization conditions. As for the larger \\hi\\ column HVCs, the distance of the HVCs is the sole {\\it direct} test for these competing models.  Previous searches of  these types of HVCs in the spectra of Galactic halo stars have mostly failed \\citep{zsargo03,sembach03}, leading some support to the hypothesis that the highly ionized HVCs could be mainly associated with the WHIM of the Local Group rather than with some galactic-scale phenomena. However, most distance limits for highly ionized HVCs are not very constraining, as the stars in the \\citet{zsargo03} survey are mostly significantly closer than those against which the higher column density \\hi\\ HVCs have been detected, $5 \\la d \\la 15$ kpc \\citep[e.g.,][]{wakker01,wakker08}. Further for the  stars at large $d$ in Zsarg\\`o et al., the signal-to-noise (S/N) levels were often not high enough and/or the stellar continua were too complicated to be able to reliably detect \\ovi\\ at high velocity.  In \\citet{zech08}, we revisited one of the Zsarg\\`o et al. stars, the PAGB star ZNG\\,1 in the globular cluster NGC\\,5904 ($z = 5.3$ kpc, see Table~\\ref{t-sum}), with additional \\fuse\\ and new {\\em HST}/STIS (E140M) observations, finding two highly ionized HVCs with no detected \\hi\\ emission. This firmly demonstrated that some of these HVCs originate near the Galactic disk. For sightlines passing through  known large \\hi\\ complexes (e.g., Complex C, Magellanic Stream), highly ionized gas is found with similar kinematics as  the neutral HVCs \\citep[e.g.,][]{collins07,fox04}.  Kinematics and metallicity arguments  also suggest some of low \\hi\\ column HVCs observed toward the Large Magellanic Cloud (LMC) are due to galactic feedback occurring within the LMC \\citep{lehner07,lehner09}.  Progress in determining the distances of the highly ionized HVCs seen in the Galactic sky is now possible owing to the recent installation of the  Cosmic Origins Spectrograph (COS) (and repair of the Space Telescope Imaging Spectrograph, STIS) onboard of \\hst. The high throughput of COS at medium resolution ($R\\sim 20,000$) allows us to efficiently observe targets much fainter in the UV than previously, and hence observe more distant stars. In this Letter, we report on the first observations from our COS program to search and characterize the high velocity \\nv, \\civ, and \\siiv\\ interstellar absorption in stars at large distances  ($4<d<21$ kpc, $3<|z|<13$ kpc). Presently, four stars out of 24 have been observed. For two of them (PG1708+142, PG1704+222, PAGB stars at $d = 10$ and 7 kpc, respectively), no HVCs were detected. For another one the signal-to-noise was lower than expected and will need a closer inspection. However, in the COS spectrum of the fourth star (HS\\,1914+7139, see Table~\\ref{t-sum}), two HVCs at $v_{\\rm LSR} = -180$ and $-118$ \\km\\ are detected in \\civ, \\siiv, \\cii, and other species.  While the HVC at $-118$ \\km\\ is detected in both absorption and \\hi\\ 21-cm emission and may be associated with the Outer Arm of the Galaxy, the other HVC is only detected in absorption.  Thus, we have found another highly ionized HVC falling toward the sun and well within the Milky Way. ",
        "conclusions": "HS\\,1914+7139 is situated at a distance of 14.9 kpc from the Sun and 6 kpc above the Galactic plane.  Three main features of the Milky Way cross its path: Orion Spur at $v_{\\rm LSR} \\sim  -40$ \\km, the Perseus arm at $v_{\\rm LSR} \\sim  -75$ \\km, and the Outer Arm at $v_{\\rm LSR} \\sim  -90,-120$ \\km. Based on the weak interstellar lines (e.g., \\sii\\ $\\lambda$1250, \\Niii\\ $\\lambda$1370), the average LSR velocity of the low velocity component is $-37 \\pm 3$ \\km, and hence the low velocity component mostly probes some relative nearby gas.  Using the rotation curve from \\citet{clemens85}, a velocity of about $-120$ \\km\\ occur in gas at Galactocentric distance $R_{\\rm G} \\sim 22 $ kpc in the direction of HS\\,1914+7139 if the gas corotates with the Galaxy. The Galactocentric distance of the star is  $R^\\star_{\\rm G} = 23.4$ kpc (assuming $R^\\sun_{G} = 8.5$ kpc). Hence combining the information from the kinematics and  distance provides support for an origin of the component at $-118$ \\km\\ from the Outer Arm, and this is therefore the first firm upper limit on the distance of the Outer Arm gas. We, however, note that both velocities and distance are also consistent with those of the nearby complex C ($d = 10 \\pm 2.5$ kpc, Thom et al. 2008, and see also Wakker et al. 2007).  Based on the rotation curve of the Milky Way, the gas moving at $-180$ \\km\\ corresponds to $R_{\\rm G} \\sim 120 $ kpc in a corotating disk/halo in the direction of HS\\,1914+7139. The limit from the distance of the star evidently rejects this hypothesis, implying that this component is a genuine HVC in the sense that its LSR velocity is inconsistent with a simple model of differential Galactic rotation. This is therefore an highly ionized HVC that is much closer than its corotating distance implies. Its high negative velocity implies it is plunging onto the Milky Way thick disk.  We already noted above that the sightline to HS\\,1914+7139 lies near complex C, but the absolute LSR velocity of the present HVC is substantially larger than those associated with complex C. Searching for other observed sightlines near HS\\,1914+7139, we found 5 QSOs/AGNs that are summarized in Table~\\ref{t-extra}. All these lines of sight show both absorption at LSR velocities lower and greater than $-150$ \\km. For the absorption centered near $-120$ \\km, sightlines with typically $b<28\\degr$ (for $l\\sim 90$--$110\\degr$) are related to the Outer Arm while HVCs at $b>30\\degr$ are associated with complex C. However, as \\citet{tripp03} also discussed, with the current observational evidence it is not clear if complex C and the Outer Arm gas are different entities or have a relationship, except for the \\hi\\ maps that indicate that the Outer Arm seems to connect more smoothly to the Galactic disk than Complex C \\citep[e.g.,][]{wakker01,tripp03}.  In what follows (except otherwise stated), we focus solely on the HVCs at $v_{\\rm LSR} \\la -150$ \\km\\ where little or no \\hi\\ 21-cm emission has been detected ($N($\\hi$) \\la 10^{18.7}$ cm$^{-2}$). Preliminary apparent column densities toward HS\\,1914+7139 for \\oi\\ and \\siii\\ are: $\\log N($\\siii$) \\simeq 13.6$, $\\log N($\\oi$) \\simeq 13.9$, giving $[$\\oi/\\siii$] = -0.9$ (where $[X/Y]= \\log N(X)/N(Y) - \\log (X/Y)_\\odot $ and the solar abundances are from Lodders et al. 2009), i.e. the gas is dominantly ionized \\citep[e.g.,][]{lehner01a}. Using the results from \\citet{tripp03}, $[$\\oi/\\siii$] \\simeq -0.53$ toward 3C\\,351, again suggesting a large fraction of ionized gas.  Only \\ovi\\ and \\civ\\ are detected toward H\\,1821+643 \\citep{savage95,oegerle00,sembach03,tripp03}, implying that the gas is highly ionized. For the other sightlines listed in Table~\\ref{t-extra}, the presence of \\siiii\\ or/and other higher ions also implies a large amount of ionized gas. So not only do the HVCs listed in Table~\\ref{t-extra} and toward HS\\,1914+7139 share similar velocities, but they all have low \\hi\\ column densities and are largely  ionized.  This suggests they trace a common structure.  HS\\,1914+7139 places these HVCs at $d<14.9$ kpc and $z<6$ kpc.  While the origin for these HVCs is still uncertain (see below), they cannot be part of the WHIM at very large distances from the Milky Way.  \\citet{tripp03} derived a metallicity for the HVC toward 3C\\,351 $[$O/H$]\\ga -1.2$ dex solar based on a comparison of \\oi\\ and \\hi. They denominated this HVC the high-velocity ridge (HVR) and argued the HVR could be the leading edge of complex C that is interacting with the lower halo of the Milky Way. This would explain the ionized nature of the HVC, which is expected according to recent hydrodynamical simulations for a HVC moving in a hot halo and approaching the Galactic plane \\citep[e.g.,][]{heitsch09}. The distance of HS\\,1914+7139 is consistent with this hypothesis, providing some support to this idea. However, several studies suggest a slight gradient in ${\\rm O/H}$ from $-0.01$ to about $-0.07$ dex\\,kpc$^{-1}$ \\citep[e.g.,][]{daflon04,esteban05}, which could accommodate low metallicity at large distances.  There is, however, a large uncertainty in the gradient and how it varies with the galactic longitude.  Based on chemical evolution models and abundance measurements, \\citet{cescutti07} found that the lowest $[{\\rm O/H}]$ has a mean and a deviation of $-0.2\\pm 0.2$ dex at $R{\\rm g} > 9.5$ kpc, which does not favor a galactic fountain phenomenon in the outer region of the Milky Way if the overall metallicity of the HVCs is low ($\\ll -0.2$ dex solar).  In summary, these early \\hst/COS results have successfully constrained the distance of two multiphase HVCs toward the direction $l = 103\\degr$ and $b=+24 \\degr$. One is related to the Outer Arm of the Milky Way and the other one traces infalling gas onto the Milky Way. It is too early to conclude if all the highly ionized HVCs are linked to galactic phenomena rather than the Local Group,  but the two detections of highly ionized HVCs in very different regions of the Galaxy (central and outer regions, see Table~\\ref{t-sum}) show that at least some of these low $N($\\hi$)$ HVCs are near the Milky Way and associated with infalling gas (likely from a Galactic foutain or  accretion of matter condensing from the Galactic corona). We note that so far only  negative high velocities  have been detected in spectra of stars, and it remains to be seen if  highly ionized HVCs with positive velocities are found at $d<10$--20 kpc from the Milky Way. Once all the observations of our program are completed, we should have a sizeable sample that will help us to better understand the origin(s) of the highly ionized HVCs that cover the Milky Way sky."
    },
    "0911/0911.2218_arXiv.txt": {
        "abstract": "We report {\\it Spitzer}/IRAC photometry of the transiting giant exoplanet HAT-P-1b during its secondary eclipse.  This planet lies near the postulated boundary between the pM and pL-class of hot Jupiters, and is important as a test of models for temperature inversions in hot Jupiter atmospheres.  We derive eclipse depths for HAT-P-1b, in units of the stellar flux, that are: $0.080\\%\\pm0.008\\%\\,[3.6\\mu$m], $0.135\\%\\pm0.022\\%\\,[4.5\\mu$m], $0.203\\%\\pm0.031\\%\\,[5.8\\mu$m], and $0.238\\%\\pm0.040\\%\\,[8.0\\mu$m]. These values are best fit using an atmosphere with a modest temperature inversion, intermediate between the archetype inverted atmosphere (HD\\,209458b) and a model without an inversion. The observations also suggest that this planet is radiating a large fraction of the available stellar irradiance on its dayside, with little available for redistribution by circulation.  This planet has sometimes been speculated to be inflated by tidal dissipation, based on its large radius in discovery observations, and on a non-zero orbital eccentricity allowed by the radial velocity data. The timing of the secondary eclipse is very sensitive to orbital eccentricity, and we find that the central phase of the eclipse is $0.4999\\,\\pm0.0005$. The difference between the expected and observed phase indicates that the orbit is close to circular, with a $3\\sigma$ limit of $|e\\,\\cos\\,{\\omega}| < 0.002$. ",
        "introduction": "Infrared light from transiting extrasolar planets can be measured using high precision photometry when the planets pass behind their stars.  Most of these secondary eclipse measurements to date have used the {\\it Spitzer Space Telescope} \\citep{charb05, deming05, deming06, deming07, demory, harrington07, charb08, machalek, machalek09, knutson08, knutson09}. {\\it Spitzer} data in the four bands of the IRAC instrument \\citep{fazio} define the planet's spectral energy distribution over the 3.6-8.0\\,$\\mu$m wavelength region, where water vapor plays a principal role in shaping the spectrum.  {\\it Spitzer} studies have shown that there are at least two classes of close-in extrasolar giant planets, differentiated by the temperature gradient with height in their atmosphere \\citep{hubeny, fortney06, fortney, burrows07, burrows08}.  One class (pM) exhibits a temperature inversion at high altitude in their atmosphere \\citep{knutson08, knutson09, machalek}.  The inversion affects the emergent spectral energy distribution by causing the water bands to appear in emission. The cause of the inversion is believed to be absorption of strong stellar irradiation in the visible, possibly by TiO/VO \\citep{hubeny, fortney06, fortney}, but other absorbers remain possible \\citep{burrows07, zahnle}.  \\citet{machalek} found that XO-1b exhibits an inverted atmospheric structure, in spite of being nominally below the stellar irradiance level projected to define the transition between the pM and the (non-inverted) pL classes. The pM/pL transition corresponds to the condensation of TiO \\citep{fortney}. Recently, \\citet{machalek09} found that XO-2b exhibits a weak temperature inversion, consistent with being near the pM/pL transition. Like XO-1b and XO-2b, HAT-P-1b \\citep{bakos} lies at the lower edge of that predicted boundary, so it provides an important additional test of the propensity toward inversion at that level of irradiation. HAT-P-1b may also exhibit a non-zero eccentricity to its orbit \\citep{johnson}, that may cause internal energy generation via the dissipation of tidal stress.  In this paper, we report {\\it Spitzer} photometry of HAT-P-1b's secondary eclipse in the four IRAC bands. We use these data to investigate the atmospheric temperature structure and energy balance of this planet, and to place more stringent limits on the eccentricity of its orbit via the timing of the secondary eclipse \\citep{charb05}. ",
        "conclusions": "\\subsection{Atmospheric Temperature Structure}  Figure~6 shows our contrast values plotted versus wavelength, and also the results from \\citet{knutson08} for HD\\,209458b, the archetype of an inverted atmosphere.  We include the contrast predicted by two 1-D models of the HAT-P-1b planetary atmosphere \\citep{fortney05, fortney}, both invoking re-emission of stellar irradiance on the dayside only. The solid line is the nominal model, without a temperature inversion. This nominal model is self-consistent and non-gray, with solar metallicity, and was calculated as in \\citet{fortney}. It has a bond albedo of 0.067 and a dayside effective temperature of 1512K. This model has no inversion because we find that it is too cool to allow gas-phase TiO.  The dashed line is a weakly inverted model, also of solar metallicity. It is {\\it not} self-consistent; it has a constant, ad-hoc temperature inversion of $dT/d\\log P = -30$K.  However, this model is constrained to have the same effective temperature (1512K) as the non-inverted model.  Since the temperature gradient of the inverted model is shallow, the emission features are weak, and it falls close to the contrast expected from a 1500K blackbody (dot-dashed line).  In all instances, we used the planetary and stellar radii from \\citet{winn}, and a Kurucz 6000/4.5/0.0 model to represent the stellar spectrum \\citep{torres}.  The rotation of HAT-P-1b should be tidally locked to its orbit, even though its orbital period is longer than for many planets in the hot Jupiter class.  Using Eq.(1) of \\citet{guillot}, we calculate a spin-down time of $8 \\times 10^6$ years, starting with Jupiter's rotation rate and adopting a very conservative Q-value ($10^6$). Since this spin-down time is much less than the age of the system \\citep{torres}, we expect tidal locking of the planet's rotation.  Based on tidally-locked rotation, the maximum dayside temperature of HAT-P-1b is 1550K, assuming zero albedo, a uniform temperature over the dayside hemisphere, and no transport to the nightside. This is too cool to produce an inversion using TiO/VO absorption, so we first investigate the non-inverted model.  Emission at the three longest IRAC wavelengths can arise from levels higher in the atmosphere than does the 3.6\\,$\\mu$m radiation \\citep{burrows07}, and these three channels exhibit a contrast that is moderately higher than the non-inverted model.  Comparing to the HD\\,209458b results \\citep{knutson08} shows good agreement at 3.6 and 8.0\\,$\\mu$m, but HAT-P-1b exhibits a contrast at 4.5 and 5.8\\,$\\mu$m that is intermediate between HD\\,209458b and the non-inverted model.  This seems qualitatively consistent with a moderate inversion, perhaps produced by an absorber other than TiO/VO.  We integrated the flux from each planetary atmospheric model and the stellar model atmosphere over the IRAC bandpass functions.  Dividing these integrated fluxes at each wavelength produces expectation values for the observations. (These are not illustrated on Figure~6, but they fall very close to the contrast values at the band-center wavelengths.) We used these expectation values to calculate the $\\chi^2$ value for the observations compared to each model.  For the inverted model (dashed line), $\\chi^2$ is $3.9$, whereas it is $11.0$ for the non-inverted model.  A value as high as $11.0$ will occur only $2.6$\\% of the time if the non-inverted model is an accurate description of HAT-P-1b's atmosphere.  This level of confidence is not sufficient to rigorously eliminate the non-inverted model, but it does indicate that the inverted model is preferable.  Within the errors, our results for HAT-P-1b can also be reproduced using a blackbody spectrum for the planet (dot-dashed line on Figure~6, $\\chi^2 =3.5$).  The luminosity of HAT-P-1A is 1.48 times solar \\citep{torres}, and it receives a stellar flux $\\sim 2/3$ of the HD\\,209458b case.  If HAT-P-1b absorbs with zero albedo and re-emits uniformly but only on the dayside hemisphere, then we expect a maximum temperature of 1550K. If the planet's emission is uniform over both hemispheres, we expect an observed temperature of 1300K. Our observations at secondary eclipse are best described by a blackbody having a temperature of $1500\\pm100$K, where we have factored in the random error of our observations as well as uncertainty in the stellar and planetary parameters. However, most of the flux is probably emitted at shorter wavelengths \\citep{barman, seager}, not directly observable using {\\it Spitzer}.  At face value our results suggest that redistribution of the stellar irradiation by dynamics may be inefficient for this planet.  \\subsection{Orbital Eccentricity}  HAT-P-1b orbits a star in a wide visual binary, and this circumstance can have significant consequences for the orbital dynamics of the planet. The inclination and eccentricity ($e$) of the planet's orbit can in principle undergo oscillations and long term evolution due to the Kozai mechanism \\citep{fabrycky}.  Misalignments between the planet's orbital inclination and the rotation axis of the star can result, and can be observed using the Rossiter-McLaughlin effect. \\citet{johnson} find that the angle between the sky projections of stellar spin axis and the orbit normal for HAT-P-1b is $3.7\\pm2.1$ degrees.  They also put an upper limit on the eccentricity of $0.067$ with 99\\% confidence.  These values suggest that any Kozai oscillations of HAT-P-1b have largely damped out. Our results for the secondary eclipse timing will further strengthen that conclusion, as discussed below.  The timing of the secondary eclipse is exquisitely sensitive to the eccentricity of the orbit, with one ambiguity being the value of $\\omega$, the argument of periastron. When our line of sight aligns with the minor axis of the planet's orbit ($\\omega =0$ or $\\pi$), then the secondary eclipse will not be centered on phase $0.5$ unless the orbit is circular.  When our line of sight aligns with the major axis of the orbit ($\\omega = {\\pi}/2$ or $3{\\pi}/2$), departures from circularity will affect the duration of the eclipse, but not the central phase.  The phase of secondary eclipse therefore constrains the value of $e\\,{\\cos}\\,\\omega$ \\citep{charb05}.  We observed two eclipses, each at two wavelengths simultaneously.  The difference in best fit phases for the same eclipse observed at different wavelengths is approximately consistent with our errors (Table~1). The larger difference occurs at 3.6 and 5.8\\,$\\mu$m, where the central phase difference is $0.5016-0.4992=0.0027$.  Because the noise at each wavelength is independent, the error on the difference in the two phases is the quadrature sum ($=0.0014$) of the phase errors at the two wavelengths. Hence the phase difference is less than a $2\\sigma$ variation.  Weighting the central phase of the eclipse at each wavelength (Table~1) by the inverse of its variance, we find a mean value of $0.4999\\pm0.0005$.  Considering the 55 seconds of light travel time across the orbit, we expect to find the eclipse at phase $0.500014$. Accounting for the light travel time and the error in the observed eclipse phase, we derive a $3\\sigma$ upper limit of $|e\\,\\cos\\,{\\omega}| < 0.002$.  If this planet's orbital eccentricity was affected by Kozai oscillations in the past, they have damped to the point where the orbit is closely circular."
    },
    "0911/0911.0665_arXiv.txt": {
        "abstract": "{We present high sensitivity H91$\\alpha$ and 3.5 cm radio continuum observations toward the planetary nebula NGC~3242. The electron temperature determined assuming local thermodynamic equilibrium is consistent within $\\sim$10\\% with that derived from optical lines and the Balmer discontinuity. The line emission and the continuum emission have very similar spatial distribution, suggesting that at this wavelength there is no other continuum process present in a significant manner. In particular, we conclude that emission from spinning dust is not important at this wavelength. In this radio recombination line the nebula presents a radial velocity structure consistent with that obtained from observations of optical lines. } ",
        "introduction": "NGC 3242 (PN G261.0+32.0) is a multiple-shell, attached-halo planetary nebula (PN), located at a distance of about 1 kpc (Stanghellini \\& Pasquali 1995). Recently, an expansion distance measured with radio continuum observations (Hajian, Phillips, \\& Terzian 1995) placed this PN at 0.42$\\pm$0.16 kpc. As it is commonly the case in planetary nebula, the individual distances show a large spread, ranging from 0.28 to 2.0 kpc for NGC~3242 (Acker et al. 1992). We will adopt the distance of 1 kpc proposed by Stanghellini \\& Pasquali (1995). The inner and outer shells of the object have diameters of 15{''} and 29{''}, respectively. NGC 3242 is a high-excitation planetary nebula (class 7; Pottasch 1984) powered by a $\\geq$1000 $L_\\odot$ (Stanghellini \\& Pasquali 1995) SD0 star (Heap 1986) with a surface temperature of $T_{eff}$ = 60,000 K (Acker et al. 1992). The post-AGB mass of the central star is estimated to be $M_{cs} = 0.56 \\pm 0.01~ M_\\odot$, corresponding to a main-sequence mass of $M_{ms} = 1.2 \\pm 0.2 ~ M_\\odot$ (Galli et al. 1997).  NGC 3242 was the first PN with a firm detection of the $^3He^+$ hyperfine line at 8.665 GHz (Balser, Rood, \\& Bania 1999). As part of these $^3He^+$ studies, the Very Large Array archive contains also high quality, unpublished data of the H91$\\alpha$ radio recombination line (project AR271), that has a nearby rest frequency of 8.584 GHz. In this paper we analyze this radio recombination line and the adjacent continuum. The motivation of this analysis is in first place to study the radio recombination line emission of this planetary nebula at interferometer angular resolutions since to our knowledge only single dish observations have been published. The second reason to study the continuum and radio recombination line emission of this PN is that recently (Casassus 2005; 2007) it was proposed as one of the planetary nebulae that may show excess anomalous microwave emission (e. g. Casassus et al. 2004; Finkbeiner et al. 2002). This new continuum radiation mechanism has been attributed to electric dipole emission from rapidly rotating dust grains (\"spinning dust\"), as predicted by Draine \\& Lazarian (1998). In a nebula that is optically thin at centimeter wavelengths one expects the free-free continuum emission and the radio recombination line emission to be very similar since both originate mostly from electron-proton interactions. The significant presence of an independent radiation mechanism such as spinning dust emission is expected to become evident in a careful comparison between the total continuum and the radio recombination line emissions. ",
        "conclusions": "We presented VLA observations of the H91$\\alpha$ radio recombination line and the adjacent continuum, as well as of H$\\alpha$ and [N~II] $\\lambda$ 6584~\\AA, toward the planetary nebula NGC 3242. Our main conclusions can be summarized as follows.  1. The LTE electron temperature derived from the line to continuum ratio agrees well with the values derived in the optical.  2. The spatial distribution of continuum and line emission are very similar, arguing against the presence of significant continuum contamination from spinning dust emission.  3. The radial velocity structure found from the H91$\\alpha$ line is consistent with that derived from the optical lines."
    },
    "0911/0911.2456_arXiv.txt": {
        "abstract": "We examine the ability of gravitational lens time delays to reveal complex structure in lens potentials.  In a previous paper, we predicted how the time delay between the bright pair of images in a ``fold'' lens scales with the image separation, for smooth lens potentials. Here we show that the proportionality constant increases with the quadrupole moment of the lens potential, and depends only weakly on the position of the source along the caustic. We use Monte Carlo simulations to determine the range of time delays that can be produced by realistic smooth lens models consisting of isothermal ellipsoid galaxies with tidal shear. We can then identify outliers as ``time delay anomalies.'' We find evidence for anomalies in close image pairs in the cusp lenses RX J1131$-$1231 and B1422+231. The anomalies in RX J1131$-$1231 provide strong evidence for substructure in the lens potential, while at this point the apparent anomalies in B1422+231 mainly indicate that the time delay measurements need to be improved.  We also find evidence for time delay anomalies in larger-separation image pairs in the fold lenses, B1608+656 and WFI 2033$-$4723, and the cusp lens RX J0911+0551. We suggest that these anomalies are caused by some combination of substructure and a complex lens environment. Finally, to assist future monitoring campaigns we use our smooth models with shear to predict the time delays for all known four-image lenses. ",
        "introduction": "Gravitational lensing has become a valuable probe of dark matter substructure in distant galaxies (see \\S B.8 of \\citealt{Kochanek_review}, and references therein). A growing body of evidence suggests that anomalous flux ratios, which are observed in many four-image quasar lenses \\citep{MetcalfZhao,Keeton-cusp,Keeton-fold}, can be explained if a few percent of the projected mass within each lens galaxy's Einstein angle (which sets the spatial scale for lensing) is contained in cold dark matter (CDM) ``clumps'' \\citep{Mao-flux,Metcalf-substruc,Chiba_substruc,Dalal-substruc,Kochanek-substruc}. It is then natural to ask whether it is possible to use lensing to measure properties beyond the mean substructure mass fraction.  Of particular interest is the clump mass function, because the discrepancy between the observed number of Local Group satellite galaxies and the theoretically predicted abundance of dark matter clumps varies strongly with mass \\citep[e.g.,][]{Klypin-missing,Moore-substructure,Strigari-VL}. Unfortunately, it is difficult to constrain the mass function with anomalous flux ratios, because there is a degeneracy such that a small clump near a lensed image (in projected distance) can produce the same flux perturbation as a large clump farther away \\citep[see][]{Dalal-substruc}.  One possible solution to this problem is to measure flux ratios at multiple wavelengths. Quasar emission regions are believed to have different sizes at different wavelengths. Roughly speaking, only clumps with Einstein radii (Einstein angles translated into units of length) larger than the size of the source can produce flux ratio anomalies \\citep[e.g.,][]{Dobler-finite}; clumps with small Einstein radii act collectively as a smooth mass component. Since the Einstein radius depends on mass, the source size effectively translates into a mass threshold. The net effect is that multi-wavelength observations provide a way to study clumps with different mass ranges.\\footnote{This is similar to the idea of ``chromatic'' microlensing, which is discussed by \\citet{Kochanek-SCMA}.}  For example, continuum \\citep{Wozniak-2237-continuum} and broad-line \\citep{Richards-1004, Keeton-0924} optical emission regions in quasars can be perturbed by objects of stellar mass and larger. Radio emission regions, by contrast, are large enough that only objects with masses $\\gtrsim 10^6 \\Msun$ are important \\citep{Dobler-finite}.  If a flux-ratio anomaly is observed at optical wavelengths but not at radio wavelengths, it is likely that microlensing by stars in the lens galaxy is the culprit.  To avoid contamination by microlensing, \\citet{Dalal-substruc} focused on radio anomalies in their study. Infrared emission regions are intermediate in size, so observations in this band provide a way to probe objects with mass scales between stars and clumps. Recent mid-IR observations with the Spitzer \\citep{Poindexter-1104-erratum, Poindexter-1104-IR} and Subaru \\citep{Chiba-IR, 0414-2237-IR} telescopes have achieved the spatial resolution necessary for strong lensing studies. In particular, \\citet{Chiba-IR} found evidence for subhalos of $\\sim 10^4 \\Msun$ in the lens system B1422+231, and possibly in PG 1115+080 as well.  In addition, \\citet{0414-2237-IR} recently found evidence of a clump with mass $> 10^5 \\Msun$ in the lens MG 0414+0534.  A related way to study the substructure mass function is to compare continuum and broad-line emission in optical lenses \\citep{Moustakas-substructure}, which originate from the central accretion disk and an extended distribution of fast-moving clouds, respectively.  These two regions differ in length scale by a factor of $\\sim$100, whence their utility for substructure lensing.  This technique has been employed to study the four-image lenses HE~0435$-$1223 \\citep{Wisotzki-0435spec}, SDSS J0924+0219 \\citep{Keeton-0924,Eigenbrod-0924}, SDSS J1004+4112 \\citep{Richards-1004}, RX~J1131$-$1231 \\citep{Sluse-1131spec}, and Q2237+0305 \\citep{Metcalf-2237,Eigenbrod-2237}, along with several two-image lenses \\citep{Burud-2149,Inada-0246,Inada-0806,Sluse-0806}  While the method of probing substructure with differential source size effects is promising, this approach does have some limitations. For one thing, the mass threshold defined by the source size and Einstein radius does not represent a sharp cutoff \\citep[see][]{Dobler-finite}, so it is difficult to obtain tight constraints on the clump mass. In addition, it may not be possible to measure flux ratios at enough wavelengths to sample the mass function well. Finally, lensing constraints on the mass function will only be as accurate as the mapping of wavelength to source size.  While the model of \\citet{Shakura-Sunyaev} has been quite successful in explaining the observed properties of accretion disks, and gives a relation between wavelength and source size, it does not apply to the full range of wavelengths relevant to strong lensing.  A second possible approach to constrain the substructure mass function is to use flux ratios in combination with other lens observables. Although they are smaller in amplitude than flux ratio perturbations, substructure effects on both image positions \\citep{Chen-imgshift} and time delays \\citep{Keeton-tdel} should be detectable. Combining different observables will be valuable because they depend on the lens potential in different ways: time delays, image positions, and flux ratios depend on the potential and its first and second derivatives, respectively. The different observables, in other words, contain different information about the substructure population. For example, \\citet{Keeton-tdel} showed that time delays are sensitive to the slope and dynamic range of the substructure mass function, which suggests that precise time delay measurements may provide a way to constrain these properties. Before we can use this method, we need to determine whether observed time delays differ from the predictions for smooth mass models in a way that may indicate the presence of substructure. Developing a method to identify such ``time delay anomalies'' is the focus of this paper.  In particular, we wish to understand the range of time delays that can be produced by reasonable smooth models. We can then classify any outlier as ``anomalous'' and use it as evidence of complexity in the lens potential. We must be careful when interpreting anomalies, however, because dark matter substructure may not be the only relevant source of complexity in the lens galaxy. Stars also constitute complex structure that is important when interpreting flux ratios (due to microlensing), but they have essentially no effect on time delays \\citep{Keeton-tdel}. This means that time delays should not depend on wavelength, so we do not need to distinguish between radio and optical time delays.\\footnote{Strictly speaking, optical and radio emission come from different regions of a quasar, so the optical and radio image positions of a lens, and hence the corresponding time delays, need not be identical.  This effect is naturally taken into account by our criterion for matching observed time delays with simulated lenses (see \\S\\ref{sec:distributions}).}  Finally, extinction by dust in the lens galaxy can perturb flux ratios, but it will not have any effect on time delays, which are measured through flux variability alone and do not depend on color information.  The other main source of complexity we need to consider is the environment of the lens galaxy. Many lens galaxies lie in groups or clusters of galaxies \\citep[e.g.,][]{Momcheva-env,Auger-env}. To lowest order, such an environment contributes a tidal shear to the lens potential \\citep[e.g.,][]{Keeton-shear-ellip}, and we therefore include shear in our analysis, but there may be higher-order terms that are non-negligible. Extreme examples of this situation include the lens B1608+656, which has two galaxies inside the Einstein angle \\citep{Koopmans-Fassnacht-1608, Koopmans-1608-model, Suyu-1608-model}, and the lens SDSS J1004+4112, which is produced by a cluster of galaxies \\citep{Inada-1004,Oguri-1004}. In general, time delay anomalies in close pairs of images potentially provide the strongest evidence of dark matter substructure, because environmental effects are fairly large-scale and should not produce dramatic differences between images separated by a distance much less than the Einstein angle. With time delay anomalies in image pairs that have larger separations, by contrast, we will need to take more care to consider environmental effects as well as substructure.  Our approach follows that of \\citet{Keeton-cusp, Keeton-fold}, who employed analytic flux-ratio relations that are generic for all lenses with fold and cusp configurations produced by smooth mass models. To be more specific, a fold lens contains a bright pair of images whose fluxes $F_A$ and $F_B$ should satisfy the relation \\begin{equation}\\label{eqn:Rfold} \\Rfold \\equiv \\frac{F_A - F_B}{F_A + F_B} \\approx \\Afold \\, d_1, \\end{equation} where $d_1$ is the image separation and $\\Afold$ depends on properties of the lens potential. A cusp lens contains a triplet of bright images whose fluxes should satisfy the relation \\begin{equation}\\label{eqn:Rcusp} \\Rcusp \\equiv \\frac{F_A - F_B + F_C}{F_A + F_B + F_C} \\approx \\Acusp \\, d_1^{2}, \\end{equation} where $d_1$ is the distance between the closest pair of images and $\\Acusp$ depends on properties of the lens potential \\citep[see, e.g.,][]{Congdon-tdel-pert}.  (Note that $\\Rfold$ and $\\Rcusp$ vanish in the limit $d_1\\rightarrow 0$.)  If these relations are violated, we may conclude that the mass distribution of the lens galaxy cannot be smooth and must contain additional structure, most likely in the form of CDM substructure or a complex lens environment. Determining in practice that a given lens is anomalous requires some care.  Equations (\\ref{eqn:Rfold}) and (\\ref{eqn:Rcusp}) show that $\\Rfold$ and $\\Rcusp$ increase with the distance between the images, so a non-zero value for one of these quantities does not automatically imply a flux ratio anomaly.  Making such an identification would require knowledge of $\\Afold$ and $\\Acusp$, which depend on the exact form of the lens potential.  Since these quantities are not directly observable, \\citet{Keeton-cusp, Keeton-fold} performed Monte Carlo simulations to determine distributions for $\\Rfold$ and $\\Rcusp$ at fixed $d_1$ using parameter values appropriate for a realistic population of lens galaxies.  With this information, it is possible to determine the probability that an observed lens is anomalous and hence contains small-scale structure.  \\citet{Oguri_tdel} employed a method similar to that of \\citet{Keeton-cusp, Keeton-fold} to study time delays, although his emphasis was different from ours.  While both studies are based on time delay distributions, Oguri's emphasis was on determining the Hubble constant, whereas we focus on small-scale structure in lens galaxies.  In fact, we explicitly remove the Hubble constant from our analysis by working with scaled time delays and time delay ratios.  Another difference is that we concentrate our attention mainly on cusp and fold lenses, which are most relevant to the questions we seek to answer.  This paper is organized as follows.  In \\S\\ref{sec:hcaus} we consider how the differential time delay depends on the parameters of the lens potential, and on the shape of the associated (astroid) caustic.  The results we present there motivate our Monte Carlo simulations, which we describe in \\S\\ref{sec:distributions}.  We apply our formalism to the 25 known four-image lenses in \\S\\ref{sec:results}.  Finally, we present our conclusions in \\S\\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  We have introduced a new method to use gravitational lens time delays to detect complex structure in the lens potential.  The complexity may be associated with CDM substructure, in which case time delays offer the chance to learn more about the substructure population than is possible with lens flux ratios; or it may be associated with the lens environment, such as a group or cluster of galaxies surrounding the lens.  The basic approach is to determine the range of time delays that can be produced by reasonable smooth lens models, so that we can identify outliers as being anomalous. To get a sense of how this could work, we first studied the dependence of the time delay between the close pair of images in a fold lens on the position of the source and the form of the lens potential. For a source near a fold point, we have found that the time delay remains approximately constant as the source moves along the caustic. For a lens modeled by an elliptical galaxy with $m=4$ multipole perturbations and tidal shear, the time delay increases with ellipticity and shear, but is not very sensitive to $m=4$ modes.  Using Monte Carlo simulations, we then constructed distributions of the time delays in four-image lenses. This approach can handle fold, cusp, and cross lenses, which comprise the three canonical four-image lens morphologies. By constructing a catalog of mock lenses based on observed populations of elliptical galaxies, we computed the range of time delays and time delay ratios that would be expected for a smooth lens potential (i.e., one with ellipticity, shear, and $m=4$ multipoles). By comparing observed time delays with the predicted ranges, we have found time delay anomalies in the systems RX J0911+0551, RX J1131$-$1231, B1422+231, B1608+656, and WFI 2033$-$4723.  It is unlikely that these anomalies can be explained by errors in our assumed values of the Hubble constant or the slope of the density profile: the Hubble constant would have to be unreasonably high or the density profile surprisingly shallow in order to explain some of the apparent anomalies (we specifically discussed RX J0911+0551 and WFI 2033$-$4723), and neither possibility could explain RX J1131$-$1231.  It is possible to further reduce sensitivity to the Hubble constant and density profile by working with time delay ratios, although in general we have found that time delay ratios have less power to reveal anomalies.  Part of the problem is that there are fewer time delay ratios known than time delays themselves, and a large uncertainty in a given time delay will lead to a correspondingly large uncertainty in the ratio, making definitive conclusions difficult.  In general, anomalies between close pairs of images in fold and cusp lenses should provide the cleanest evidence of substructure. The cusp lens RX J1131$-$1231 contains such anomalies, which is consistent with conclusions based on more detailed modeling by \\citet{Morgan_1131} and \\citet{Keeton-tdel}. The cusp lenses B1422+231 and RX J0911+0551 show evidence of time delay anomalies in the cusp triplet, but the large uncertainties in the measured time delays prevent firm conclusions at present. The only fold pair that is clearly anomalous is in B1608+656, but the peculiar nature of this system (with two lens galaxies inside the Einstein angle) makes it difficult to draw a definitive conclusion about substructure from our analysis.  Among larger-separation image pairs we have found time delay anomalies in the fold lens WFI 2033$-$4723 and the cusp lens RX J0911+0551.  Both lenses show evidence of a complex environment, but in order to distinguish between that and substructure as the origin of the time delay anomalies it will be necessary to precisely measure the time delays between the close images (the fold pair in WFI 2033$-$4723, and the cusp triplet in RX J0911+0551).  In the hope that the sample of observed precision time delays will continue to grow, we have predicted the time delays for all mixed-parity image pairs in all 25 known four-image lenses.  As new time delays are measured, it will be a simple matter to compare them with our predicted confidence intervals to determine whether they are anomalous.  If a lens galaxy contains complex structure, time delays should help reveal it.  Flux ratios provide a powerful way to find small-scale structure, but they are not unique in this; time delays hold great promise for contributing to our understanding of the role played by dark matter in the universe, especially when combined with other lensing observables."
    },
    "0911/0911.3666_arXiv.txt": {
        "abstract": "% Leading models of galaxy formation require large-scale energetic outflows to regulate the growth of distant galaxies and their central black holes. However, current observational support for this hypothesis at high redshift is mostly limited to rare $z>2$ radio galaxies. Here we present Gemini-North NIFS Intregral Field Unit (IFU) observations of the [O~{\\sc iii}]$\\lambda$5007 emission from a $z\\approx$~2 ultraluminous infrared galaxy (ULIRG; $L_{\\rm IR}>10^{12}$~$L_{\\odot}$) with an optically identified Active Galactic Nucleus (AGN). The spatial extent ($\\approx$~4--8~kpc) of the high velocity and broad [O~{\\sc iii}] emission are consistent with that found in $z>2$ radio galaxies, indicating the presence of a large-scale energetic outflow in a galaxy population potentially orders of magnitude more common than distant radio galaxies. The low radio luminosity of this system indicates that radio-bright jets are unlikely to be responsible for driving the outflow. However, the estimated energy input required to produce the large-scale outflow signatures (of order $\\approx10^{59}$~ergs over $\\approx$~30~Myrs) could be delivered by a wind radiatively driven by the AGN and/or supernovae winds from intense star formation. The energy injection required to drive the outflow is comparable to the estimated binding energy of the galaxy spheroid, suggesting that it can have a significant impact on the evolution of the galaxy. We argue that the outflow observed in this system is likely to be comparatively typical of the high-redshift ULIRG population and discuss the implications of these observations for galaxy formation models. ",
        "introduction": "Modern-day cosmology is a rich cocktail of observations and theory. Hard observational constraints guide the theoretical models, which then provide more detailed insights into the potential physical mechanisms that drive the growth of galaxies and their massive central black holes. One area where theoretical models have been particularly successful is in highlighting that star formation must have been truncated in the most massive galaxies at high redshifts (\\hbox{$z\\approx$~2--3;} e.g.,\\ Benson et~al. 2003; Di Matteo et~al. 2005; Bower et~al. 2006). The leading candidates to cause this truncation are galaxy-wide outflows and winds, associated with either star-formation activity (e.g.,\\ supernovae winds) or active galactic nuclei (AGNs; e.g.,\\ winds and jets initiated from the black-hole accretion disk). If energetic enough, these winds could heat the cool galactic gas and/or eject it from the gravitational potential of the host galaxy, effectively shutting down any further significant star formation.  Integral field units (IFUs) are the ideal tool to identify galaxy-wide energetic outflows as they provide spectro-imaging over several arcsecond fields of view (e.g.,\\ Swinbank et~al. 2005, 2006). The typical expected signature of an outflow is a broad emission-line gas component kinematically distinct from the narrow-emission line gas in the host galaxy (velocity offsets of several hundreds of kilometers/second) that is extended over kpc-scales. IFU observations of the [O~{\\sc iii}] (rest-frame 5007~ang) emission in a handful of massive radio-loud AGNs at $z\\approx$~2--3 have indeed revealed the signatures of galaxy-scale energetic outflows (e.g.,\\ Nesvadba et~al. 2006, 2007, 2008). The outflows from these systems appear to be driven by radio jets initiated by AGN activity. While these data provide evidence that energetic outflows can be identified in distant galaxies, radio-loud AGNs are rare beasts and it is far from clear how typical these outflows are in the typical distant massive galaxy population. To throw light on whether outflows are ubiquitous in the distant Universe, we have used the Gemini NIFS IFU to search for galaxy-wide energetic outflows in more typical radio-quiet systems. Here we present our initial results from this program on the $z\\approx$~2 radio-quiet AGN SMM~J1237+6203, which are published in Alexander et~al. (2009).  SMM~J1237+6203 is an optically bright $z=$~2.07 quasar that is also bright at submillimetre (850~$\\mu$m), radio (1.4~GHz), and X-ray (0.5--8 keV) wavelengths (Alexander et~al. 2003, 2005; Barger et~al. 2003; Chapman et~al. 2005). The estimated infrared luminosity of SMM~J1237+6203 is of order $6\\times10^{12}$~$L_{\\odot}$, indicating that it is a distant ultra-luminous infrared galaxy (ULIRG). The AGN in SMM~J1237+6203 is luminous (X-ray luminosity $\\approx10^{44}$~erg~s$^{-1}$) and undoubtably contributes to a considerable fraction of the infrared luminosity (Alexander et~al. 2005). However, given the bright submillimetre and radio emission, it seems likely that SMM~J1237+6203 also hosts an ultra-luminous infrared starburst in addition to the luminous AGN activity. Previously published near-infrared spectroscopy shows that SMM~J1237+6203 has bright, and possibly broad and extended, [O~{\\sc iii}] emission (Takata et~al. 2006). These are the expected signatures of a galaxy wide outflow, although high-quality IFU observations are required to provide detailed constraints.  \\setcounter{figure}{0} \\begin{figure} \\plotfiddle{dalexander_fig1.ps}{2.75in}{90.}{45.}{45.}{170}{-30} \\caption{Collapsed one-dimensional NIFS spectrum showing the emission line profiles fitted with both a broad and narrow emission-line component. Inset plot shows the narrow [O~{\\sc iii}] velocity field. Taken from Alexander et~al. (2009).} \\end{figure} ",
        "conclusions": ""
    },
    "0911/0911.1808_arXiv.txt": {
        "abstract": "{We present here new results of two-dimensional hydrodynamical simulations of the eruptive events of the 1840s (the great) and the 1890s (the minor) eruptions suffered by the massive star $\\eta$ Car. The two bipolar nebulae commonly known as the Homunculus and the little Homunculus were formed from the interaction of these eruptive events with the underlying stellar wind.  As in previous work (Gonzalez et al. 2004a, 2004b), we assume here an  interacting, nonspherical multiple-phase wind scenario to explain the shape and the kinematics of both Homunculi, but adopt a more realistic parametrization of the phases of the wind. During the 1890s eruptive event, the outflow speed {\\it decreased} for a short period of time. This fact suggests that the little Homunculus is formed when the eruption ends, from the impact of the post-outburst $\\eta$ Car wind (that follows the 1890s event) with the eruptive flow (rather than by the collision of the eruptive flow with the pre-outburst wind, as claimed in previous models; Gonzalez et al. 2004a, 2004b). Our simulations reproduce quite well the shape and the observed expansion speed of the large Homunculus. The little Homunculus (which is embedded within the large Homunculus) becomes Rayleigh-Taylor unstable and develop filamentary structures that resembles the spatial features observed in the polar caps. In addition, we find that the interior cavity between the two Homunculi is partially filled by material that is expelled during the decades following the great eruption. This result may be connected with the observed double-shell structure in the polar lobes of the $\\eta$ Car nebula. Finally, as in previous work, we find the formation of tenuous, equatorial, high-speed features that seem to be related to the observed equatorial skirt of $\\eta$ Car.} ",
        "introduction": "Located at a distance of 2.3 kpc, one of the most massive stars in our Galaxy ever discovered, $\\eta$ Car is a well-known example of the evolved and unstable luminous blue variable (LBV) stars, characterized by sporadic, violent mass-loss eruptive events (e.g., Humphreys $\\&$ Davidson 1994). The 19th century spectrograms of $\\eta$ Car (see also Walborn $\\&$ Liller 1977; Humphreys, Davidson $\\&$ Smith 1999 and references therein) provide evidence that the star underwent a giant eruption during $\\sim$20 yr around the 1840s, where a few solar masses of gas ($\\ge$10 M$_{\\odot}$) were expelled into the interstellar medium, and the total luminous output of $\\sim$10$^{49.5}$-$\\sim$10$^{50}$ erg was comparable to a supernova explosion (Smith et al. 2003a). It is still not understood how and why this outburst occurred, but it may be connected with $\\eta$ Car being a binary star system (e.g., Damineli 1996; Damineli, Conti, $\\&$ Lopes 1997). From this event, a symmetric, bipolar nebula known as the ``Homunculus\" (that extends from -8 to +8 arcsec along its major axis) was produced (Humphreys $\\&$ Davidson 1994; Currie et al. 1996; Davidson $\\&$ Humphreys 1997; Smith $\\&$ Gehrz 1998). During the 1890s, the historical light curve of $\\eta$ Car also shows a much fainter eruptive event of shorter duration ($\\sim$10 yr), but still an outburst in the sense that the mass loss rate was enhanced compared with the normal underlying wind. This eruption eruption resulted in the formation of a smaller nebula embedded within the large Homunculus (with an angular size of about $\\pm$2 arcsec) called the ``Little Homunculus\" (LH) (Ishibashi et al. 2003; Smith 2005). With a total mass of $\\sim$0.1 M$_{\\odot}$, the polar caps of the LH currently move at slower speeds ($\\sim$ 250 km s$^{-1}$) than the large Homunculus in the polar direction ($\\sim$650 km s$^{-1}$), but shares the same prolate geometry. The total kinetic energy released in the 1890s event is $\\sim$ 10$^{46.5}$ erg (a factor of $\\sim 10^3$ smaller than the kinetic energy expelled during the great eruption). Apart from both Homunculi, observations of $\\eta$ Car with the Hubble Space Telescope (e.g., Morse et al. 1998) also show the presence of an equatorial skirt (orthogonal to the axis of the Homunculus) that moves with velocities of 100-350 km s$^{-1}$ (Smith $\\&$ Gehrz 1998; Davidson et al. 2001), although high speed features, with typical velocities of $\\sim$750 km s$^{-1}$ (Smith et al. 2003b) or even larger (Weis 2005) have also been detected. These equatorial components may have been expelled from both the 1840s and the 1890s outbursts (e.g., Davidson $\\&$ Humphreys 1997). The large Homunculus is the central portion of a larger nebula (the outer ejecta) around $\\eta$ Car that extends up to a diameter of 60 arcsec (e.g., Weis 2001, 2005; Walborn 1976; Walborn et al. 1978). The outer ejecta are not symmetric and contain numerous filaments and knots that may have been produced by previous stellar wind mass loss. The sizes and morphology of such structures in the outer ejecta are manifold. Radial velocities for the knots reach up to 2000 km\\,s$^{-1}$ (Weis 2001, Weis 2005), but the average expansion velocity of the outer ejecta lies at lower values (around 750 km\\,s$^{-1}$), similar to the speeds found in the Homunculus. The outer ejecta also follow a bipolar pattern, like the large Homunculus, with approximately the same axis of symmetry.  At optical wavelengths, the large Homunculus is mainly a reflection nebula. Its spectrum shows the presence of dust scattered emission (e.g., Hillier $\\&$ Allen 1992) that allows one to see indirectly the shape of the $\\eta$ Car wind (see Smith et al. 2003b) and also low-excitation intrinsic emission. In addition, near-infrared spectra obtained by Smith (2006) confirmed the existence of a double-shell structure at the edges of the polar lobes of the Homunculus (which was previously inferred from thermal dust emission; see Smith et al. 2003a). A thin outer shell of intrinsic H$_{2}$ emission (that traces the main scattering layer seen in visual images), and a thicker inner skin of [Fe II] (which partially fills the interior of the lobes). On the other hand, the outer ejecta are an emission nebula. Smith $\\&$ Morse (2004) present optical spectra of $\\eta$ Car showing strong oxigen lines in some emission features of the outer ejecta. At high energies, Chandra has detected X-ray emission from the large-scale nebula around the star (and also from the central object, probably due to a wind collision region around the massive binary system; e.g., Corcoran et al. 2001; Pittard $\\&$ Corcoran 2002). Soft X-ray (0.1-0.8 keV) observations by Seward et al. (2001) show an extended shell clearly associated with the debris field of the outer ejecta which are compatible with a collision process between a fast wind and cometary knots of slower material. The $\\eta$ Car nebula is also a bright radio source consistent with thermal (free-free) emission. Observations by Retallack (1983) at 1.415 GHz (taken with the Fleurs synthesis telescope) show emission from an overall region as large as 40 arcsec (i.e., larger than the optical Homunculus). Gonz\\'alez et al. (2006) estimated the contribution by shocks to the total radio-continuum emission detected from the $\\eta$ Car nebula. Using observational estimates of the wind parameters of the eruptive event of the 1890s and of the stellar wind after the end of the eruption, these authors investigated the evolution of the polar caps of the LH formed as a result of the collision between these outflows. They found that the LH emits continuum radiation which is detectable at radio wavelengths and indeed, has an appreciable contribution to the total flux of the $\\eta$ Car nebula.  Different models have been proposed to explain the shaping and kinematics of the $\\eta$ Car bipolar nebulae (see, for instance, Soker 2001; Soker 2004; Matt $\\&$ Balick 2004; Gonzalez et al. 2004a, b; A. Frank et al. 1995, 1998). One possible explanation is that it is produced by the interaction of the  winds expelled by the central star at different injection velocities (e.g., Icke 1988; Frank, Balick $\\&$ Davidson 1995; Dwarkadas $\\&$ Balick 1998; Frank, Ryu $\\&$ Davidson 1998; Langer, Garc\\'ia-Segura $\\&$ Mac Low 1999; Gonz\\'alez et al. 2004a,b). Adopting a colliding wind scenario, Gonz\\'alez et al. (2004a,b) performed two-dimensional numerical simulations of the Homunculi nebulae of $\\eta$ Car. In their models, the large Homunculus is formed by the interaction of the eruptive outflow of the 1840s with the pre-outburst $\\eta$ Car wind (both with different degrees of nonspherical symmetry). A second eruption (assumed to be spherical) collides with the pre-outburst wind giving rise to the LH. These authors showed that such a scenario could explain the shape and kinematics of the Homunculi, and also the existence of the high-velocity features observed in the equatorial plane. However, their prescribed conditions to create the LH nebula do not agree with overall results suggest by the observations. During the 1890s event, the $\\eta$ Car wind slowed down and increased its mass loss rate, returning to its normal quiescent state after the eruption ended. Another observational result (as mentioned earlier) is that the Homunculus contains much more mass than that that had been previously recognized (Smith et al. 2003a). In this work, we perform new numerical simulations that incorporate these observational results and investigate  the overall evolution of the bipolar outflows of $\\eta$ Car.  The paper is organized as follows. In $\\S$ 2, we describe the model. The numerical simulations and the discussion of the results are presented in $\\S$ 3, and in $\\S$ 4 we draw our conclusions. ",
        "conclusions": "In this article, we carried out high-resolution two-dimensional gasdynamic simulations of the dramatic outbursts suffered by the star $\\eta$ Car during the 1800s, the larger of which occurred in the 1840s and resulted in the formation of the large Homunculus nebula and the smaller one in the 1890s created the inner little Homunculus. During these events, the parameters of the $\\eta$ Car wind may have drastically changed in short periods of time. In contrast with the great eruption (where both the mass loss rate and the ejection velocity were suddenly increased), during the minor event of the 1890s (much fainter and of shorter duration than the 1840s event) the mass loss rate was enhanced, but the wind velocity diminished compared with the normal $\\eta$ Car wind.  Considering a simplified interacting stellar wind scenario, we could explain the shape and the observed expansion speed (as a function of latitude) of the large Homunculus. In addition, our numerical models show that the little Homunculus is formed at the $end$ of the 1890s eruption, when the post-outburst $\\eta$ Car wind collides with the eruptive outflow. Important differences with regard to our previous models of the $\\eta$ Car nebula (Gonz\\'alez et al. 2004a,b) have been obtained. In Gonz\\'alez et al. (2004a,b), we assumed that the impact between the 1890 outburst and the pre-outburst wind was the cause of the development of the little Homunculus nebula. As a matter of fact, although the momentum flux is dominated by the 1890s eruption, the post-outburst wind has an important effect on the kinematics and morphology of the little Homunculus as it causes  the complete deposit of the material expelled during the eruption into the polar caps. At this time, the  powerful stellar wind accelerates the inner Homunculus material so that it asymptotically reaches the velocity of the post-outburst wind. Due to  Rayleigh-Taylor instabilities  generated by the low-density fast wind pushing the high-density slow wind, the polar caps of the inner Homunculus quickly develop filamentary structuring (Fig. 5) that shows some resemblance with the observed spatial structures in the polar lobes of the little Homunculus by Smith (2005) that suggest that it is not perfectly homologous to the larger Homunculus nebula.  Also, in contrast with the previous models of $\\eta$ Car (Gonz\\'alez et al. 2004a,b), it is noteworthy that in the present numerical modeling the interior cavity between the Homunculi is partially filled by material that is expelled during the decades following the end of the 1840s great eruption, rather than being almost empty. This agrees with the observed double-shell structure observed in the polar lobes by Smith (2006), consisting of a thin outer H$_2$ skin (which contains most of the material of the nebula) and a thicker [Fe II] layer with a more irregular spatial distribution.  As in the previous models (Gonz\\'alez et al. 2004a,b), the present results show the formation of an equatorial outflow with both low and high velocity features. These are probably related to the equatorial skirt of $\\eta$ Car. We note however, although it has been predicted that the observed skirt might be associated with the two outbursts (e.g. Davidson et al. 2001), our simulations indicate that only the great eruption contributes to its formation.  A final remark is in order. One could ask whether one can learn about the progenitor of Eta Car from the assumed mass loss history or what can cause such sharp changes in the mass loss rate and speed, without affecting the general shape of the Homunculi. As stressed in Gonz\\'alez et al. (2004a,b). these questions are vey probably connected with Eta Car being a binary star system and the nature of the interaction between the main and companion stars. In this work, however, we have focussed in the formation and dynamical evolution of the shock structures associated to the large and little Homunculi given the parameters (terminal velocity, mass-loss rate and/or density) of the different wind phases, without addressing the inner mechanism that first triggered it or its variability at the much smaller scales of the central source of Eta Car. The orbit of the secondary star, for instance, has apoastro and periastro distances of $\\sim$14 AU and $\\sim$3 AU, respectively, while in our simulations, the outflows are injected at much greater distances (10$^{16}$ cm). A detail study univocally connecting both the inner source scales and the outer scales was, therefore, out of the scope of the present study. Nonetheless, there have been some first efforts in this direction (e.g., Soker 2001 ; Falceta-Gon\\c{c}alvez $\\&$ Abraham 2009) that should be further explored in the future."
    },
    "0911/0911.0429_arXiv.txt": {
        "abstract": "We investigate the effects of pseudoscalar-photon mixing on electromagnetic radiation in the presence of correlated extragalactic magnetic fields. We model the Universe as a collection of magnetic domains and study the propagation of radiation through them. This leads to correlations between Stokes parameters over large scales and consistently explains the observed large-scale alignment of quasar polarizations at different redshifts within the framework of the big bang model. ",
        "introduction": "\\label{one}  Light pseudoscalar particles, such as the axion, arise naturally as pseudo-Goldstone bosons in spontaneously broken global symmetries \\cite{Peccei:1977hh,Peccei:1977ur,Weinberg:1977ma,Wilczek:1977pj,McKay:1977gd,McKay:1978wn,Kim:1979if,Dine:1981rt,Kim:1986ax}. An axion has an effective coupling to two photons, therefore in an external magnetic field it can oscillate into a photon and vice versa \\cite{Clarke:1982,Sikivie:1983ip,Sikivie:1985yu,Sikivie:1988mz,Maiani:1986md,Raffelt:1987im,Carlson:1994yqa,Bradley:2003kg,Das:2004qka,Das:2004ee,Ganguly:2005se,Ganguly:2008kh}. This mixing between axions and photons can lead to observable changes in the intensity and polarization of photons, having interesting astrophysical and cosmological implications \\cite{Harari:1992ea,Berezhiani,Mohanty:1993nh,Das:2000ph,Kar:2000ct,Kar:2001eb,Csaki:2001jk,Csaki:2001yk,Grossman:2002by,Jain:2002vx,Song:2005af,Mirizzi:2005ng,Raffelt:2006cw,Gnedin:2006fq,Mirizzi:2007hr,Finelli:2008jv}. Since the mixing depends on the frequency of radiation, it also affects the electromagnetic spectrum \\cite{Ostman:2004eh,Lai:2006af,Hooper:2007bq,Hochmuth:2007hk,Chelouche:2008ta}. Various experimental searches of pseudoscalar particles have led to significant limits on their masses and on the coupling parameters of the models \\cite{Dicus:1978fp,Vysotsskii:1978,Dearborn:1985gp,Raffelt:1987yu,Raffelt:1987yt,Turner:1987by,Mohanty:1993nh,Janka:1995ir,Keil:1996ju,Brockway:1996yr,Grifols:1996id,Raffelt:1999tx,Rosenberg:2000wb,Zioutas:2004hi,Yao:2006px,Raffelt:2006cw,Jaeckel:2006xm,Andriamonje:2007ew,Robilliard:2007bq,Zavattini:2007ee,Rubbia:2007hf}.  In the current paper we study pseudoscalar-photon mixing in the presence of correlated background magnetic fields in order to understand the results reported by Hutsemekers et al. \\cite{Hutsemekers:1998,Hutsemekers:2000fv,Hutsemekers:2005iz} pertaining to the coherent alignment of quasar polarizations over Gpc scales. It has been shown earlier that propagation of radiation through the extragalactic medium, in the presence of a background magnetic field, can affect all Stokes parameters due to pseudoscalar-photon mixing. These changes have been investigated for radio \\cite{Jain:1998kf,Jain:2002vx,Ralston:2003pf}, optical \\cite{Hutsemekers:1998,Hutsemekers:2000fv,Jain:2002vx,Jain:2003sg,Hutsemekers:2005iz,Payez:2008pm,Piotrovich:2008iy}, and cosmic microwave background (CMB) \\cite{Raffelt:1999tx,Mirizzi:2005ng,Raffelt:2006cw,Lee:2006za,Agarwal:2008ac} photons. We explicitly show that in our model it is possible to obtain nonzero correlations between the optical polarization of quasars separated by large distances in the sky. These nonzero correlations lead to an effect similar to the observed large-scale alignment of the polarizations.  We consider the extragalactic medium to be a large collection of magnetic domains, the background magnetic field in each domain being constant. This model was recently used in \\cite{Agarwal:2008ac} where the extragalactic medium was considered as a large collection of uncorrelated magnetic domains in order to study the effects of pseudoscalar-photon mixing on CMB polarization. We extend the model of \\cite{Agarwal:2008ac} to include correlations between magnetic fields in different domains, as motivated in \\cite{Subramanian:2003sh,Seshadri:2005aa,Seshadri:2009sy}, and study the propagation of optical radiation from quasars through these domains.  The paper is organized as follows. In Sec. \\ref{two} we outline the main concepts and equations of pseudoscalar-photon mixing and describe the domain propagation model. In Sec. \\ref{three} we obtain magnetic field correlations between different domains. In Sec. \\ref{four} we bring together results derived in the previous two sections and study alignments among the angles of polarization using a full numerical propagation model. We conclude with a brief discussion of our results and offer perspectives in Sec. \\ref{five}. We also discuss an approximate analytical treatment to explain correlations in quasar polarizations in the Appendix. ",
        "conclusions": "\\label{five}  A very striking alignment in the optical polarization of quasars has been observed over cosmologically large distances ($\\sim$ 1 Gpc) by Hutsemekers et al. \\cite{Hutsemekers:1998,Hutsemekers:2000fv,Hutsemekers:2005iz} at both low ($z \\sim 0.5$) and high redshifts ($z \\sim 1.5$). It is unlikely that quasar polarizations are intrinsically aligned with one another over such large distances since this appears to be in conflict with the basic assumptions of isotropy and homogeneity of the Universe. It has been proposed in the past that these observations may be explained in terms of a propagation effect related to axion-photon mixing. In the current paper we have explored this possibility in detail, and studied the propagation of electromagnetic radiation through a large number of magnetic domains.  We have shown that correlations between components of the background magnetic field in various domains induce correlations between the linear polarization of different quasars. Our analysis explicitly shows that quasar polarizations are expected to be aligned in such a scenario due to the presence of significant correlations between them, as observed by Hutsemekers et al. \\cite{Hutsemekers:1998,Hutsemekers:2000fv,Hutsemekers:2005iz}. A calculation of the dispersion of polarization with respect to nearest neighbors in a model of propagation along a given line of sight shows significant alignments compared to a random sample. A direct comparison with observations and further study of the effect clearly requires numerical simulations of the complete 3D propagation model that we have proposed. We leave such an analysis to future work.  We further note that a calculation of correlations in the circular polarization, Stokes $V$ parameter, is also expected to yield a significant result. A detailed calculation of this parameter is postponed to future research. This may prove to be an important signature of pseudoscalar-photon mixing in the observed alignment of quasar polarizations, as noted in \\cite{Hutsemekers:2008iv}."
    },
    "0911/0911.2276_arXiv.txt": {
        "abstract": "Lorentz violation (LV) is predicted by some quantum gravity theories, where photon dispersion relation is modified, and the speed of light becomes energy-dependent. Consequently, it results in a tiny time delay between high energy photons and low energy ones. Very high energy (VHE) photon emissions from cosmological distance can amplify these tiny LV effects into observable quantities. Here we analyze four VHE $\\gamma$-ray bursts (GRBs) from Fermi observations, and briefly review the constraints from three TeV flares of active galactic nuclei (AGNs) as well. One step further, we present a first robust analysis of VHE GRBs taking the intrinsic time lag caused by sources into account, and give an estimate to quantum gravity energy $\\sim 2 \\times 10^{17}$~GeV for linear energy dependence, and $\\sim 5 \\times 10^9$~GeV for quadratic dependence. However, the statistics is not sufficient due to the lack of data, and further observational results are desired to constrain LV effects better. ",
        "introduction": "The unification of standard model and general relativity is the most intriguing and desirous goal of modern physics, and it stimulates the development of many theoretical ideas such as string theory and loop gravity in the past decades. It is interesting that some of them predict Lorentz symmetry violation (LV) at low energies or being realized in a highly nonlinear form. This has arisen in the space-time foam~\\cite{am97,am98,emn99,emn08}, loop gravity~\\cite{gp99,amu00}, torsion in general gravity~\\cite{Yan}, vacuum condensate of antisymmetric tensor fields in string theory~\\cite{ks89,ks91,ck97,ck98,arnps09}, and the so-called double special relativity~\\cite{ame02c,ame02,ame02b}.  In most of LV theories, the photon dispersion relation is modified and the correction is believed to be suppressed by the large Planck scale, $E_{\\rm P} \\equiv \\sqrt{\\hbar c^5 /G} \\simeq 1.22 \\times 10^{19}$~GeV~\\cite{jlm06,km09}. These theories are mainly classified into two categories. One is effective field theory (EFT), which provides an excellent framework where the tiny LV effects are introduced through LV operators: renormalizable ones with dimension three and/or four, see, e.g., standard model extension (SME)~\\cite{ck97,ck98}, and the further extended non-renormalizable ones with dimension five and/or six~\\cite{arnps09,mp03}. However, not all quantum gravity theories can be embedded into the EFT framework, such as the quantum space-time foam model~\\cite{am97,am98,emn99,emn08} and double special relativity~\\cite{ame02c,ame02,ame02b}, and they also introduce modifications to the canonical dispersion relation. Consequently, they can result in an energy dependence of the speed of light in the vacuum, owing to the propagation through an effective gravitational medium containing space-time quantum fluctuations. Fortunately, this scenario of lower energy ``relic probe\" of Planck scale events can be tested carefully through laboratory experiments~\\cite{ge,am09} or astronomical observations~\\cite{el00,jlm03,m05,gs08,Xiao:2008yu,Bi:2008yx,as09,emn09}.  Generally, the most model-independent photon dispersion relation reads in the context of Taylor series as, \\begin{equation}\\label{ve} v(E) = c_0 \\left( 1 - \\xi \\frac{E}{E_{\\rm P}} - \\zeta \\frac{E^2}{E_{\\rm P}^2} \\right) \\, , \\end{equation} where $v(E)$ is the speed of photons with energy $E$, $c_0$ is the speed of low energy photons, and $\\xi$, $\\zeta$ are model-dependent parameters, characterizing the energy where LV occurs. Because of the suppression of Planck energy, the terms of higher orders are negligible, and the quadratic term takes effect only when the linear term vanishes. It could be possible that a more concrete form of dispersion relation may contain more specific terms of energy dependence, but we may take Eq.~(\\ref{ve}) as a general estimate in a model-independent manner. ",
        "conclusions": "Nowadays, more very high energy (VHE) data become available and they provide a rich ground to look into cosmologically accumulated effects, rooted from quantum gravity scale. In this paper, we use cosmological objects such as $\\gamma$-ray bursts (GRBs) and active galactic nuclei (AGNs) to estimate Lorentz violation (LV) effects from the time delay between photons with different energies. The quantum gravity scales from GRBs and AGNs are surprisingly compatible in some sense.  We give a crude attempt here to survey a set of VHE celestial events for LV hints and disentangle astrophysical origins from LV effects as well. Though the data are very limited in our study, even not sufficient to establish a good statistics, we get some potential clues of LV from these samples. Further, we notice that more observations are emerging and more experiment data are accumulating, e.g., Fermi LAT has detected GRB~090323, GRB~090328 and GRB~091003 with VHE emissions and their redshifts are attained from other observations as well, thus once these results are published, there will definitely be an opportunity to give a more stringent analysis. We suggest combining VHE GRBs at different redshifts to separate source effects, if data points approximately lay on two lines according to two different GRB types, it will be a supportive evidence to our approach.  Although the results in our analysis are very preliminary, they give hints for possible LV energy scales, $\\sim 2 \\times 10^{17}$~GeV for linear energy dependence and $\\sim 5 \\times 10^9$~GeV for quadratic dependence. However, the fits in Fig.~\\ref{lgrb} and Fig.~\\ref{qgrb} induce a negative intercept, $\\Delta t_{\\rm in} \\sim -20$~s for the linear dependence and $-6$ to $-10$~s for the quadratic scenario, which means that at sources, high energy photons emit earlier than low energy ones. This conflicts with common expectations based on assumptions that low energy photons are produced by electrons while high energy ones are generated by protons, and because of lighter masses, electrons are accelerated earlier and hence low energy photons come out first~\\cite{emn09}. However, because that the last word about the emission mechanism of GRBs is not addressed yet, and the negative intercepts are not very significant in our fits, they may not trouble much.  Worthy to mention that, there are several stringent restrictions coming from various tests on LV effects in the content of given theories~\\cite{arnps09,jlm06,m05,as09,emn09}. Synchrotron radiation measurements on electrons from the Crab Nebula~\\cite{jlm03,m07} place very strong constraints on the electron sector, making the linear dependence for electrons almost impossible. Most effective field theories (EFTs) predict birefringence for the photon sector~\\cite{gp99,mp03}, and astrophysical observation leads to very tight restrictions on the linear suppression~\\cite{gk01}. And the GZK cutoff originated from the ultra-high-energy cosmic rays (UHECRs) interacting with cosmic microwave background (CMB) photons also produces severe constraints for linear as well as quadratic energy dependence~\\cite{gs08,Xiao:2008yu,Bi:2008yx}. However, some scenarios can avoid the above restrictions~\\cite{emn09}, e.g., photons have different LV parameters as electrons and hadrons, thus our analysis above serves as a tentative estimate for some still surviving theories.  Finally, we stress that it would be premature to draw a rigorous conclusion on LV at the moment, and more theoretical considerations and practical data analysis are needed for a more stringent constraint and clarification on these issues."
    },
    "0911/0911.1105_arXiv.txt": {
        "abstract": "We present a new insight on NGC 6034 and UGC 842, two groups of galaxies previously reported in the literature as being fossil groups. The study is based on optical photometry and spectroscopy  obtained with the CTIO Blanco telescope and Sloan Digital Sky Survey archival data.  We find that NGC 6034 is embedded in a large structure, dominated by three  rich clusters and other small groups. Its first and next four ranked galaxies have magnitude differences in the $r$ band and projected distances which violate the optical criteria to classify it as a fossil group.  We confirm that the UGC 842 group is a fossil group, but with about half the velocity dispersion that is reported in previous works. The velocity distribution of its galaxies reveals the existence of two structures in its line of sight, one with $\\sigma_v$ $\\sim$ 223 km\\,s$^{-1}$ and another with $\\sigma_v$ $\\sim$ 235 km\\,s$^{-1}$, with a difference in velocity of  $\\sim$ 820 km\\,s$^{-1}$. The main structure is dominated by passive galaxies, while these represent $\\sim$ 60\\% of the second structure. The X-ray temperature for the intragroup medium of a group with such a velocity dispersion is expected to be $kT$ $\\sim$ 0.5--1 keV, against the observed value of $kT$ $\\sim$ 1.9  keV reported in the literature. This result makes UGC 842 a special case among fossil groups  because (1) it represents more likely the interaction between two small groups, which warms the intragroup medium and/or (2) it could constitute evidence that member galaxies lost energy in  the process of spiraling toward the group center, and decreased the velocity dispersion of the  system. As far as we know, UGC 842 is the first low-mass fossil group studied in detail. ",
        "introduction": "Fossil groups are galaxy systems optically dominated by an elliptical galaxy immersed in an extended and luminous X-ray halo (L$_{X, \\rm bol}$ $>$ 10$^{42}$ $h^{-2}_{50}$ erg\\,s$^{-1}$), in which the magnitude gap between the two brightest galaxies within half of the virial radius is greater than 2 in the $r$ band \\citep{Ponman94}.  The scarcity of L$^*$ galaxies and the evidence for the existence of a ``massive\" structure traced by the hot intragroup gas support the hypothesis of their central galaxies being formed by merging of luminous galaxies, most likely by dynamical friction \\citep{Cypriano06,MdO06,MdO09}.  The estimated time scales for dynamical friction and results of numerical simulations are consistent with an old age for fossil groups. For example, $N$-body/hydrodynamical simulations carried out by \\citet{DOnghia05} suggest a correlation between the magnitude gap between the two brightest galaxies and the formation time of the group. More recently several studies of simulated fossil groups were made using the Millennium simulations \\citep{Dariush07,Sales07,Diaz08}  which corroborated the idea that fossil groups/clusters assembled most of their masses much earlier than non-fossil systems. These studies also predicted a fraction of 3\\%--13\\% of groups with masses greater than 10$^{13}$ M$_{\\odot}$ being fossil groups, depending on the exact range of masses considered.  The true nature of fossil groups has promoted a lively debate in the past decade.  The absence of a conclusive explanation for the nature of these objects is mainly due to (1) the lack of a proper sample for statistical studies, (2) the lack of X-ray data with sufficient resolution and signal-to-noise (S/N) for a proper study of the group properties, and (3) the lack of optical spectroscopy of the group members, for membership confirmation and a study of the kinematic properties of the known groups. A growing number of fossil group candidates has been claimed in the recent literature. For example, \\citet{Santos07} identified 34 potential candidate fossil groups in a systematic search carried out using the {\\it Sloan Digital Sky Survey} (SDSS). More recently, \\citet{LaBarbera09} identified 25 other fossil groups. However, the nature of these newly identified groups still need to be confirmed. We investigate here two groups pointed in the literature as being fossil groups: UGC 842 and NGC 6034 \\citep{Voevodkin08,Yoshioka04}.  UGC 842 is a bright elliptical galaxy with a heliocentric radial velocity of 13,556$\\pm$32 km\\,s$^{-1}$ \\citep{Huchra99}. It is a weak radio source detected by the Very Large Array (VLA) telescope with flux level of S(6 cm) $\\sim$ 1.1 mJy \\citep{Gioia83}. The UGC 842 group is immersed in an extended halo with a radius of $\\sim$ 4' in the sky, which corresponds to $\\sim$ 300 kpc at its redshift. \\citet{Gastaldello07} reported an X-ray observation of UGC 842 system carried out by the {\\it Chandra} satellite. They derived a virial radius and mass of 1272$\\pm$220 $h^{-1}_{70}$\\,kpc and 12.8$\\times$10$^{13}$ $h^{-1}_{70}$\\,$M_{\\odot}$, and $r_{500}$ and $M_{500}$ of 634$\\pm$85 $h^{-1}_{70}$\\,kpc and 7.54$\\pm$3.41$\\times$10$^{13}$ $h^{-1}_{70}$\\,$M_{\\odot}$, respectively. About 1 year later, \\citet{Voevodkin08} reported a combined optical and X-ray analysis of this group from SDSS and {\\it XMM-Newton} data. According to these authors, galaxies in the UGC 842 group have $\\sigma_v$ $\\sim$ 439 km\\,s$^{-1}$ (from 16 galaxies inside $r$ = 509 $h^{-1}_{71}$\\,kpc), its intragroup gas has a temperature $kT$ of 1.90$\\pm$0.30\\,keV and metallicity of 0.34$\\pm$0.12\\,Z$_{\\odot}$, and displays a bolometric X-ray luminosity of $\\sim$ 1.63$\\times$10$^{43}$\\,$h^{-2}_{71}$\\,erg\\,s$^{-1}$. From the {\\it XMM-Newton} data, \\citet{Voevodkin08} also inferred $r_{500}$ $\\sim$ 509 $h^{-1}_{71}$\\,kpc, M$_{\\rm gas,500}$ = (3.5$\\pm$1.1)$\\times$10$^{12}$ $h^{-1}_{71}$\\,M$_{\\odot}$, and M$_{\\rm total,500}$ = (4.03$\\pm$0.69)$\\times$10$^{13}$ $h^{-1}_{71}$\\,M$_{\\odot}$. The UGC 842 group is classified as a fossil group by these authors.  \\begin{deluxetable}{cccccccrrrl} \\tabletypesize{\\scriptsize} \\tablecaption{Radial velocities and $r$-mag of the central galaxies of UGC 842 and NGC 6034 \\label{tbl:sample}} \\tablenum{1} \\tablecolumns{1} \\tablewidth{0pc} \\tablehead{ \\colhead{Group} & \\colhead{Central Galaxy} & \\colhead{Radial Velocity} & \\colhead{1st-ranked $r$-mag} & \\colhead{1st-ranked $R$-mag}\\\\ & & (CTIO-Hydra) & (SDSS) & (SDSS) } \\startdata UGC 842   & J011853.62--010007.2  & 13428$\\pm$40 km\\,s$^{-1}$ & 13.47 & -23.19\\\\ NGC 6034  & J160332.08+171155.2   & 10162$\\pm$48 km\\,s$^{-1}$ & 13.54 & -22.49\\\\ \\enddata \\end{deluxetable}  NGC 6034 is a bright E/S0 radio galaxy member of Abell 2151 and immersed in the Hercules supercluster \\citep{Corwin71,Val78}, with a heliocentric radial velocity of $\\sim$ 10,112 km\\,s$^{-1}$ \\citep{Tarenghi79}. VLA observations at $\\lambda$21-cm (H\\,{\\small I}) reported by \\citet{Dickey97} show that NGC 6034 has a head-tail morphology in the continuum, with two jets, and evidence of gas falling toward the nucleus with a velocity of about 70 km\\,s$^{-1}$ as derived from the detection of shifted narrow absorption line at 10,226$\\pm$15 km\\,s$^{-1}$. An X-ray observation of the NGC 6034 group with the {\\it Einstein} Observatory shows that $L_{\\rm X,0.5-4.5 keV}$ $<$ 1.3$\\times$10$^{42}$ $h_{50}^{-2}$ erg\\,s$^{-1}$ \\citep{Canizares87}. ASCA observation (1999 August 24) reveals that this group has a relatively cold intragroup medium with $kT$ = 0.67$\\pm$0.09 keV and abundance $Z$ = 0.08$\\pm$0.05 $Z_{\\odot}$, affected by a photoelectric absorption equivalent to $N_{H}$ $\\sim$ 3.4$\\times$10$^{20}$ cm$^{-2}$, and displaying a bolometric X-ray luminosity of about 2.8$\\times$10$^{43}$ $h_{50}^{-2}$ erg\\,s$^{-1}$ \\citep{Fukazawa04}. Curiously, also from ASCA observation, \\citet{Yoshioka04} derived $kT$ = 1.29(+0.48/-0.36) keV, $Z$ = 0.11(+0.39/-0.10) $Z_{\\odot}$, $N_{H}$ $\\sim$ 3.4$\\times$10$^{20}$ cm$^{-2}$, and $L_{\\rm X, 0.1-2.4 keV}$ $\\sim$ 7.5$\\times$10$^{41}$ $h_{75}^{-2}$ erg\\,s$^{-1}$. The main properties of the UGC 842 and NGC 6034 groups are described in Table \\ref{tbl:sample}.  We report on new optical photometry and spectroscopy of the groups UGC 842 and NGC 6034 (hereafter UGC 842 and NGC 6034, for simplicity) and their ``neighborhoods\", carried out at the CTIO-Blanco 4.0 m telescope, which are combined with data from the SDSS Data Release 6 \\citep[DR6;][]{Adelman08}. These data are used here to perform an analysis of the structure and kinematics of these two systems and also allowed an investigation about the nature of the systems as fossil groups (confirmed in the case of UGC 842 and not confirmed for NGC 6034). When needed, we adopt a lambda cold dark matter ($\\Lambda$CDM) cosmology with $H_{0}$ = 70 km\\,s$^{-1}$\\,Mpc$^{-1}$, $\\Omega_{M}$ = 0.3, and $\\Omega_{\\Lambda}$ = 0.7. ",
        "conclusions": "\\label{sct:disc}  We summarize below the main findings of this paper:  About NGC 6034:  \\begin{itemize}  \\item NGC 6034 is a group of $\\sim$ 4 L$^{*}$ galaxies with a mass of $\\sim$ 7$\\times$10$^{13}$ $h^{-1}_{70}$\\,M$_{\\odot}$ and $R_{virial}$ of $\\sim$ 954 $h^{-1}_{70}$\\,kpc (or $\\sim$ 4.7$\\times$10$^{13}$ $h^{-1}_{70}$\\,M$_{\\odot}$ and $R_{virial}$ of $\\sim$ 840 $h^{-1}_{70}$\\,kpc if only non-emission-line galaxies are considered). It is not a fossil group given that the magnitude difference in the $r$ band between the first and second (J160402.75+171656.6) ranked galaxies is only 0.79 mag and the projected distance between these two galaxies is $\\sim$ 360 $h_{70}^{-1}$\\,kpc. Other four (or five; see Section \\ref{sct:cmd}) galaxies also violate the optical criteria to classify it as a fossil group.  \\item NGC 6034 is clearly part of a much larger structure that includes at least three clusters (among them the Hercules cluster) and several other groups.  \\item The velocity distribution of NGC 6034 is fairly well represented by a Gaussian except for a high velocity tail composed of spiral galaxies. These are most probably objects falling onto the system. The velocity dispersion of only the non-emission-line galaxies (31 members with measured redshifts) is 361$\\pm$40 km\\,s$^{-1}$ and including the 13 emission-line objects it increases slightly to 410$\\pm$39 km\\,s$^{-1}$.  \\end{itemize}  About UGC 842:  \\begin{itemize}  \\item The system referred to as UGC 842 in the literature is in fact a superposition of two groups, S1 and S2, with a velocity difference of about 820 km\\,s$^{-1}$.  \\item UGC 842/S1 is dominated by passive galaxies while there is a high fraction (40\\%) of emission galaxies in S2.  \\item UGC 842/S1 is dominated by a bright elliptical galaxy and it is a low mass fossil {\\it group}. The large content in passive galaxies suggests equilibrium, and the estimated virial mass and radius are 1.1$\\times$10$^{13}$ $h^{-1}_{70}$\\,M$_{\\odot}$ and 509 $h^{-1}_{70}$\\,kpc, respectively.  \\item There is a large discrepancy between the expected temperature and X-ray luminosity expected for a group with such a low sigma and its observed values of $kT$ of 1.90$\\pm$0.30\\,keV and L$_{X,\\rm bol}$ $\\sim$ 1.63$\\times$10$^{43}$\\,$h^{-2}_{71}$ erg\\,s$^{-1}$ by \\citet{Voevodkin08}, suggesting that we have a case of interaction of the two subclumps S1 and S2 or we are whitenessing the decrease of sigma due to the central merger event.  \\end{itemize}  Other previous papers have also studied UGC 842. However, it has not been realized in these works that what is referred to as UGC 842 is in fact two systems. For example, a velocity dispersion of 439 km\\,s$^{-1}$ was derived by \\citet{Voevodkin08} from the redshifts of 16 galaxies within a radius of 509 $h^{-1}_{71}$\\,kpc around UGC 842. This value is in good agreement with our value of 471$\\pm$37 km\\,s$^{-1}$ determined from the redshifts of 35 potential member galaxies, if we ignore the double-peaked velocity distribution. However, the KMM-test \\citep{Ashman94} rejects a Gaussian velocity distribution at the 99.5\\% level, and the distribution reveals clearly two structures -- S1, the UGC 842 group, and S2 -- with a velocity dispersion of about 230 km\\,s$^{-1}$ each (see Figure \\ref{fig:veldist_ugc} and Table \\ref{tab:vel}). The fact that the bolometric X-ray luminosity of UGC 842 fits well in the L$_X$-T$_X$ relation, the measured sigma for the group is too low in the L$_X$--$\\sigma_v$ and T$_X$--$\\sigma_v$ group relations \\citep[e.g.,][]{Xue00}: UGC 842 seems to be too luminous and too hot, with a $kT$ of $\\sim$\\,1.9 keV \\citep{Voevodkin08}, for its sigma. The expected X-ray temperature for S1 and S2 from their velocity dispersion is around 0.5--1 keV, and the superposition of both plasmas in the line of sight does not imply in the detection of a temperature as high as $kT$\\,$\\sim$\\,1.9\\,keV. Although not too pronounced, similar behavior has also been observed in other fossil groups \\citep{Khosroshahi07}. In fact, this high value for the temperature of UGC 842 is in line with the suggestion of \\citet{Khosroshahi07} that both temperature and X ray luminosity has been boosted for fossil groups compared to the values for ``normal'' groups with similar velocity dispersions, given their early times of formation. Another possibility we envisage is that the galaxies have lower relative velocities in fossil groups given that they have lost energy by dynamical friction in the process of spiraling towards the group center for interacting and finally merging. The radial velocity difference between the peaks of the velocity distribution of S1 and S2 ($\\sim$ 820 km\\,s$^{-1}$; see Section \\ref{sct:veldist}) suggests that the two groups are simply overlapping in the line of sight. Although the smooth appearance of the X-ray image \\citep{Voevodkin08} agrees with this hypothesis, we think the SNR of the X-ray observation cannot rule out the presence of X-ray substructures that could indicate a recent interaction of the two sub-groups. Better S/N and spatial resolution X-ray data may reveal if there is an interaction between the two groups (S1 or S2) or not. Thus, the measured temperature of $kT$ $\\sim$ 1.9 keV for UGC 842 may either represent the intragroup medium of the most massive group, S1, or it could be the result of an interaction between S1 and S2. Both possibilities make UGC 842 especially interesting in the study of formation and evolution of fossil groups."
    },
    "0911/0911.4720_arXiv.txt": {
        "abstract": "Red supergiants (RSGs) are an evolved He-burning phase in the lifetimes of moderately high mass ($10-25M_{\\odot}$) stars. The physical properties of these stars mark them as an important and extreme stage of massive stellar evolution, but determining these properties has been a struggle for many years. The cool extended atmospheres of RSGs place them in an extreme position on the Hertzsprung-Russell diagram and present a significant challenge to the conventional assumptions of stellar atmosphere models. The dusty circumstellar environments of these stars can potentially complicate the determination of their physical properties, and unusual RSGs in the Milky Way and neighboring galaxies present a suite of enigmatic properties and behaviors that strain, and sometimes even defy, the predictions of stellar evolutionary theory. However, in recent years our understanding of RSGs, including the models and methods applied to our observations and interpretations of these stars, has changed and grown dramatically. This review looks back at some of the latest work that has progressed our understanding of RSGs, and considers the many new questions posed by our ever-evolving picture of these cool massive stars. ",
        "introduction": "Red supergiants (RSGs) \\footnotetext{Predoctoral Fellow, Smithsonian Astrophysical Observatory, 60 Garden St., Cambridge, MA 02139} are a He-burning evolutionary phase in the lifetimes of moderately massive ($10M_{\\odot} \\lesssim M \\lesssim 25M_{\\odot}$) stars. According to the Conti (1976) scenario for the evolution of massive stars, and its subsequent illustration in Massey (2003), RSGs are the end result of a nearly horizontal evolution across the Hertzsprung-Russell (H-R) diagram as their blue H-burning predecessors leave the main sequence and cross the ``yellow void\", passing through the very short-lived yellow supergiant stage. In some cases this is the terminal stage for massive stars, which spend a significant fraction of their time as RSGs before ending their lives as hydrogen-rich Type II supernovae. In other more massive cases, stars will spend a portion of their He-burning lifetimes as RSGs but then evolve back across the H-R diagram, passing once again through the brief yellow supergiant phase and exploding as either blue supergiants or Wolf-Rayet (W-R) stars depending on their initial masses and mass loss rates.  The assumed physical properties of RSGs mark them as a unique and extreme phase of massive stellar evolution. They have the largest physical size of any stars, and their very cool effective temperatures ($T_{\\rm eff}$) and extended atmospheres lead to a spectrum that is dominated by molecular absorption lines. The latter two of these characteristics both pose a challenge to the development of accurate stellar atmosphere models; the extended RSG atmospheres, with their large scale heights, invalidate the typical assumption of a plane parallel atmospheric geometry, and their cool temperatures demand a careful and complete treatment of molecular opacities. The temperatures of RSGs also result in significant negative and temperature-sensitive bolometric corrections on the order of a few magnitudes. As a result, determining these stars' luminosities is dependent on careful $T_{\\rm eff}$ measurements, making an accurate picture of their physical properties vital to any attempts at placing these stars in their proper place on the H-R diagram.  Such measurements are a challenge in and of themselves; the physical properties of RSGs have remained poorly understood, and at odds with the predictions of stellar evolutionary theory, until recently. This is due in part to the relatively sparse number of nearby RSGs in the Milky Way, difficulties in differentiating RSGs from red foreground stars when studying extragalactic samples, and significant complexities introduced as a result of these stars' mass loss rates and dusty circumstellar environments. Their extreme physical properties also make them very difficult to model in detail. Despite their importance in massive stellar evolution, RSGs were generally ignored for many years by the massive star community, with a few important exceptions (e.g., Humphreys 1978, 1979a, 1979b, Humphreys \\& McElroy 1984, Elias et al.\\ 1985). Further investigations of RSGs in recent years have made important strides towards answering many outstanding questions about these stars. These same studies have also introduced a number of new questions and revealed the true complexity of RSGs and their critical place in the grand scheme of massive stellar evolution.  In this review I look back at the latest advances in our understanding of RSG physical properties, beginning with the methods that must be used in order to photometrically and spectroscopically identify RSGs as massive stars in the Milky Way or other nearby galaxies (Section 2). Recent work has determined RSG effective temperature scales, bolometric luminosities, and their resulting placement on the Hertzsprung-Russell diagram, revealing interesting metallicity effects on the physical properties of these stars and posing new questions about their masses, luminosities, and lifetimes (Section 3). We have gained a new and more detailed understanding of the mass loss rates and mechanisms of these stars, as well as new insights into the dust produced by RSGs (Section 4). Finally, studies of unusual RSGs in the Milky Way and the Magellanic Clouds have highlighted unusual behaviors and properties of massive stars that challenge our current picture of stellar evolution (Section 5). Finally, I consider the questions that are currently being posed as a result of our improved understanding of RSG physical properties and their importance in stellar evolution (Section 6). ",
        "conclusions": ""
    },
    "0911/0911.3156_arXiv.txt": {
        "abstract": "Model color magnitude diagrams of low-metallicity globular clusters usually show a deficit of hot evolved stars with respect to observations. We investigate quantitatively the impact of such modelling inaccuracies on the significance of star formation history reconstructions obtained from optical integrated spectra. To do so, we analyse the sample of spectra of galactic globular clusters of Schiavon et al. with STECKMAP (Ocvirk et al.) and the stellar population models Vazdekis et al. and Bruzual \\& Charlot, and focus on the reconstructed stellar age distributions. Firstly, we show that background/foreground contamination correlates with E(B-V), which allows us to define a clean subsample of uncontaminated GCs, on the basis of a E(B-V) filtering.  We then identify a ``confusion zone'' where fake young bursts of star formation pop up in the star formation history although the observed population is genuinely old. These artifacts appear for 70-100\\% of cases depending on the population model used, and contribute up to 12\\% of the light in the optical. Their correlation with the horizontal branch ratio indicates that the confusion is driven by HB morphology: red horizontal branch clusters are well fitted by old stellar population models while those with a blue HB require an additional hot component. The confusion zone extends over {\\rm [Fe/H]}$=[-2,-1.2]$, although we lack the data to probe extreme high and low metallicity regimes. As a consequence, any young starburst superimposed on an old stellar population in this metallicity range could be regarded as a modeling artifact, if it weighs less than 12\\% of the optical light, and if no emission lines typical of an HII region are present. This work also provides a practical method for constraining horizontal branch morphology from high signal to noise integrated light spectroscopy in the optical. This will allow post-AGB evolution studies in a range of environments and at distances where resolving stellar populations is impossible with current and planned telescopes. ",
        "introduction": "Spectroscopy is a major tool for the study of galaxy formation and evolution. The analysis of the spectrum of galaxies provides constraints on their star formation activity, past and present. However the accuracy of these constraints  depends on the quality of the data, but also critically on the quality of the models used to interpret these spectra, as shown for example in \\cite{koleva08}. Since spectral models of stellar populations rely on stellar evolution models, it is vital the latter models be able to account for most of the observed stellar evolution stages, or at least the most luminous ones (e.g. turnoff, red giant branch and horizontal branch stars). While thermally pulsing asymptotic giant branch (TP-AGB hereafter) stars are fundamental for NIR spectral modelling, hot evolved stars are crucial in the UV for intermediate and old populations. The latter have been proposed as the origin of the UV-upturn of elliptical galaxies \\citep{han07,Yi2008}. In this paper, we will focus on these hot  horizontal branch (hereafter HB) stars. They are fairly luminous for their mass: while weighing only about 0.5 $\\Msun$, they have a typical absolute magnitude M$_{V}$=-0.7 \\citep{allens}. Their He-burning core tends to be hotter with decreasing metallicity, making the HB of low-Z galactic globular clusters bluer \\citep{lee94}, although the morphology of the HB also depends on a second parameter, so far elusive.  Accurate modelling of the HB morphology of GCs is a difficult problem. For instance, it is strongly sensitive to mass-loss on the RGB \\citep{dotter08}. Moreover, accurate fitting of blue HB in GCs usually requires a fine tuning of He abundance dispersion, as shown in NGC2808 by \\cite{lee05}. This clashes with the general treatment of spectral synthesis for galaxies, where a single stellar population (hereafter SSP) is chemically homogeneous. In effect, the parameter space allowed by popular full spectrum SSP models (as opposed to those predicting indices) such as \\cite{BC03} or \\cite{PEG-HR} is strictly (age,metallicity), and the abundances are solar-scaled, with no intrinsical dispersion. With only these parameters, it is expected that the isochrones adopted in these SSP models are not appropriate for the whole range of HB morphology. Indeed, Fig. 7 of \\cite{BC03} shows that the 'Padova 1994' and 'Padova 2000' tracks do not fit well the color-magnitude diagrams (hereafter CMD) of GCs with blue HB. In particular the HB of the models does not extend enough towards the blue. These very tracks are nonetheless used for the computation of spectral SSP models. What can be the impact of this lack of hot evolved stars when interpreting the spectra of stellar populations ? Their effect can be readily seen in the integrated colors of GCs \\citep{smith07}. Moreover, recent works indicate that space observations in the GALEX NUV and FUV are required to differentiate between blue HB stars and young hot stars \\citep{kaviraj07,schawinski07kav}. However, although GALEX has been successful at securing data for a fair fraction of the sky, a large corpus of past, present and ongoing extragalactic studies still relie on optical spectra only (e.g. \\cite{moped01,panter1,heavensnature,vespa07,vespa09,koleva09,michielsen08}).  It is now well known that HB morphology makes age-dating stellar populations from optical spectra only troublesome, using indices \\citep{beasley2000,puzia05}, but also using full spectrum fitting techniques \\citep{koleva08,sharina09}: the age inferred spectroscopically can be several Gyr younger than that measured from the turnoff. The effect on star formation history (hereafter SFH) reconstructions could be even more dramatic: do the missing hot stars show up as an additional fake burst as suggested in \\cite{coelho09}, or do they simply shift the whole age distribution to younger ages ? At the moment, no studies have been published specifically on the effect of HB morphology on the significance of SFH reconstructions from full spectrum fitting in the optical.  In this paper, we propose to investigate quantitatively this very issue. To do so, we first need a sample of objects with well-known, simple star formation histories, and a variety of well constrained HB morphologies. From this point of view, GCs are advantageous since they are mostly believed to arise from a single star formation episode, and display a whole range of HB morphologies. We will focus on the library of integrated spectra of galactic GCs of \\cite{schiavon05} with STECKMAP \\citep{STECMAP,STECKMAP}, a Bayesian method for reconstructing the star formation history of galaxies using full spectrum fitting, and a collection of popular single stellar population (SSP) models based on MILES \\citep{MILES} and also \\cite{BC03}. For each GC, the CMD is available from the literature, so that we can relate HB morphology to  possible departures of the SFH reconstructions from a unique star formation burst. We recall that the aim of the paper is not to propose a new spectral model for the sample of GCs: it is well established that interpreting HB morphology in globular clusters requires more parameters than the simple SSP population models used here \\citep{bedin04,ferraro04,norris04}. Instead, our goal is to understand how the unknown HB morphology in galaxies may complicate the interpretation of their spectra. That's why our setup reproduces the prevailing regime in which most extragalactic stellar populations studies in optical integrated light spectroscopy are carried out and therefore uses SSPs or combinations of SSPs. The plan of the paper is as follows: Sec. \\ref{s:methodology} presents the data, the SSP models and the method, and Sec. \\ref{s:results} presents the recovered stellar age distributions and their dependence on HB morphology. In Sec. \\ref{s:discussion}, we discuss the implications of our results for spectroscopic extragalactic studies. ",
        "conclusions": "The analysis of integrated spectra of galactic GCs \\cite{schiavon05} with STECKMAP yields two families of stellar age distributions: \\begin{itemize} \\item{Type I: genuinely old populations, with no contribution from components younger than $10^{9.5}$ yr.} \\item{Type II: old populations with a superimposed hot component younger than $10^{8.5}$ yr contributing a fraction $\\fy$ to the light of the GC.} \\end{itemize} GCs with red HB and a type II SAD are subject to background/foreground contamination. Since the latter correlates with color excess E(B-V), we define a clean sample (with respect to contamination) by removing GCs with E(B-V)$>0.4$. In this new sample, red HB GCs all have type I SADs, while blue HB GCs have mostly type II SADs, with up to 12\\% of the light in the hot component. This is the signature of inaccuracies in the modelling of the horizontal branch: in the model tracks used in Vazdekis et al. and \\cite{BC03}, the blue HB is often too short compared to observations. As a consequence, the observed integrated spectrum contains an additional hot component with respect to the corresponding model SSP, which gives rise to the $\\fy$-HBR correlation. This is the main result of the paper, and it is found with both SSP models Vazdekis et al. and BC03.  We have thus identified a ``confusion zone'', where spurious starbursts are very likely to appear in star formation history reconstructions, eventhough the observed population is genuinely old. These fake starbursts are young ($<10^{8.5}$ yr) and occur in 70-100\\% of cases depending on the SSP model, and can weight up to 12$\\%$ of the light in the optical. This confusion is driven by HB morphology and is thus mainly a low-metallicity effect (${\\rm [Fe/H]}<-1.2$). In this regime, any minor young burst of star formation can be an artifact, { especially if no emission lines associated with star formation are seen.} We expect this prediction to be independent of the interpretation method and models, as long as the evolutionary tracks used for the spectral synthesis are not tuned specifically to account for BHB morphology. The exact weight allowed for the fake starburst depends on the wavelength range considered. It should increase for domains bluer than in this work, and decrease towards the red.  The highest metallicity of our clean sample is {\\rm [Fe/H]}=-0.7, which prevents us from properly probing the high-Z regime. We plan to do so by studying old metal-rich extragalactic GCs or an integrated spectrum of NGC6791.  We wish to encourage the stellar evolution and extragalactic communities to consolidate the recent progresses made on HB morphology modeling by including the relevant tracks into future spectral synthesis models such as Vazdekis et al.. However we note that accurate fitting of blue HB morphology in GCs by fine tuning of He abundance dispersion such as in \\cite{lee05} introduces additional  parameters, and possibly new degeneracies in spectrum interpretation. Moreover, proper use of these tracks requires a stellar library, empirical or theoretical, for a range of non-solar abundance ratios, reliable everywhere on the CMD. Such models are still in their infancy, despite improvements such as \\cite{coelho07}.  This study is also an interesting sanity check for STECKMAP, which is found to behave very well. We propose that the exercise of analysing the E(B-V)-filtered sample and checking the results dependence on HBR as we have done be a test for all the star formation history reconstruction methods, such as MOPED \\citep{moped01}, VESPA \\citep{vespa07}, STARLIGHT \\citep{CF04-1} or \\cite{wild07}, and urge other authors to perform the same experiment.  Finally, the fact that STECKMAP  associates BHB morphology with a type II SAD with a high success rate, shows that HB morphology can be determined in integrated light through full spectrum fitting, and in the optical domain, as was realised by \\cite{schiavon04}. This technique could thus become a valuable tool for extragalactic cluster studies: by removing the need for UV (i.e. space) observations and resolved populations (i.e. relatively nearby), it gives access to samples and environments otherwise unreachable."
    },
    "0911/0911.1365_arXiv.txt": {
        "abstract": "We present the first spectroscopic analysis of the faint M31 satellite galaxies, \\andxi\\ and \\andxiii, as well as a reanalysis of existing spectroscopic data for two further faint companions, \\andix\\ (correcting for an error in earlier geometric modelling) and \\andxii. By combining data obtained using the DEep Imaging Multi-Object Spectrograph (DEIMOS) mounted on the Keck~II telescope with deep photometry from the Suprime-Cam instrument on Subaru, we have identified the most probable members for each of the satellites based on their radial velocities (precise to several $\\kms$ down to $i\\sim22$). Using both the photometric and spectroscopic data, we have also calculated global properties for the dwarfs, such as systemic velocities, metallicities and half-light radii. We find each dwarf to be very metal poor ([Fe/H]$\\sim-2$ both photometrically and spectroscopically, from their stacked spectrum), and as such, they continue to follow the luminosity-metallicity relationship established with brighter dwarfs.  We are unable to resolve dispersion for \\andxi\\ due to small sample size and low S:N, but we set a one sigma upper limit of $\\sigma_v<$4.5 kms$^{-1}$.  For \\andix, \\andxii\\ and \\andxiii\\ we resolve velocity dispersions of $\\sigma_v$=4.5$^{+3.6}_{-3.4}$, 2.6$^{+5.1}_{-2.6}$ and 9.7$^{+8.9}_{-4.5}$ kms$^{-1}$, though we note that the dispersion for \\andxiii\\ is based on just three stars. We derive masses within the half light radii for these galaxies of 6.2$^{+5.3}_{-5.1}\\times10^6\\msun$, 2.4$^{+6.5}_{-2.4}\\times10^6\\msun$ and 1.1$^{+1.4}_{-0.7}\\times10^7\\msun$ respectively. We discuss each satellite in the context of the Mateo relations for dwarf spheroidal galaxies, and in reference to the Universal halo profiles established for Milky Way dwarfs \\citep{walker09b}. Both \\andix\\ and XII fall below the universal halo profiles of \\citet{walker09b}, indicating that they are less massive than would be expected for objects of their half-light radius. When combined with the findings of \\citet{mcconnachie06b}, which reveal that the M31 satellites are twice as extended (in terms of both half-light and tidal radii) as their Milky Way counterparts, these results suggest that the satellite population of the Andromeda system could inhabit halos that with regard to their central densities -- are significantly different from those of the Milky Way. ",
        "introduction": "\\footnotetext[1]{The data presented herein were obtained at the W.M. Keck Observatory, which is operated as a scientific partnership among the California Institute of Technology, the University of California and the National Aeronautics and Space Administration. The Observatory was made possible by the generous financial support of the W.M. Keck Foundation.} \\footnotetext[2]{Based in part on data collected at Subaru Telescope, which is operated by the National Astronomical Observatory of Japan.}  In recent years, the low luminosity population of dwarf spheroidal galaxies (dSphs) has drastically increased in number, with $\\sim$30 faint satellites discovered both within the halo of our Milky Way (MW) and in that of our ``sister'' galaxy, M31 (e.g. \\citealt{zucker04,willman05b,martin06,zucker06a,zucker06b,zucker07,irwin08,mcconnachie08,belokurov08,martin09}). Whilst in principle these minuscule stellar associations should be some of the simplest objects of the galactic family - which spans roughly seven orders of magnitude in mass - to understand, the analysis of faint dSphs (M$_V>-8$) within the MW has dissolved the notion that they form in a simple, homogeneous fashion. Detailed studies (both photometric and spectroscopic) of such objects reveal several unusual properties for the population. For example, they highlight a significant departure from the well-established mass-luminosity relationship for dSphs \\citep{mateo98}. A number of these newly discovered objects are also quite controversial in their nature, possessing properties that are common to both dSphs and globular clusters and are therefore difficult to classify as either, such as the unusual objects, Willman 1 and Segue 1 \\citep{willman05a,belokurov07,niederste09}. In addition, comparisons of the MW and M31 dSph systems have revealed significant differences between the populations, where the classical M31 dwarfs (M$_V<-8$) are observed to be at least twice as extended as their Milky Way counterparts, in terms of both their half-light and tidal radii \\citep{mcconnachie06b}. In addition, the spatial distribution of the M31 satellites around their host is more extended than the Galactic satellite distribution, with approximately half of the Galactic satellites found within 100 kpc of the Galaxy, compared to roughly 200 kpc for M31 \\citep{mcconnachie06a, koch06}. Such differences indicate that environment plays a significant role in the evolution of dSphs, the full extent of which is not known.  Another significant point of interest concerning dwarf galaxies is the well known discrepancy between the observed number of these satellites, and the number predicted by $\\Lambda$CDM models in numerical simulations (e.g. \\citealt{moore99,klypin99}), which still amounts to one to two orders of magnitude. Survey completeness is likely to contribute to at least part of this deficit (see e.g. \\citealt{willman04,simon07,tollerud08,koposov08,mcconnachie09}); as in the Galactic case, observational efforts are hampered by obscuration from the disk and the bulge. It is thought that future all-sky surveys may find a wealth of ultra-faint satellites that would reduce the gap between observation and theory. Other explanations assume that many of these halos would remain dark to the present day, owing to suppression of star formation as the result of a photo-ionizing background \\citep{somerville02}. These halos would therefore represent a significant proportion of the satellite halos found in simulations, reconciling these with observations.  In an attempt to better understand this unusual population of galaxies, and to unite theory with observation, spectroscopic surveys of faint dwarf galaxies (particularly within the MW via the Sloan Digital Sky Survey (SDSS)) have been carried out. Using the velocity information from the spectra, one can measure the systemic velocities and intrinsic velocity dispersions for these objects, which can be used to estimate the mass of these systems (e.g. \\citealt{illingworth76,richstone86,walker09b,wolf09}). However, methods such as these make various assumptions about the anisotropy and virial equilibrium of the systems that are not necessarily correct. Nonetheless, such methods have been shown to be good indicators of the instantaneous mass of the system (e.g. \\citealt{oh95,piatek95,munoz08}), and can therefore provide an insight into the mass distribution of the satellites. Historically, the results of such studies indicated the existence of a mass limit of $\\sim10^7\\msun$ under which no dwarf galaxies are found, supporting the theory that these satellites would indeed only inhabit massive dark matter halos. The existence of such a mass limit would confirm the empirical relation defined by \\citet{mateo98} for satellites that are brighter than Draco ($M_V=-8.8$) and Ursa Minor ($M_V=-8.9$) in the Local Group: $M/L=2.5+ (L/10^7\\lsun)$, where $M/L$ is the mass-to-light ratio of the galaxies and $L$ is their luminosity. Work by \\citet{strigari08}, investigating the mass contained within the central 300~pc of dSphs, reinforced this relation, demonstrating that, despite the studied dSphs spanning five orders of magnitude in luminosity, they appear to show a lower dynamical mass limit of $\\sim10^7\\msun$. However, such a finding is in stark contrast with the masses derived for some of the recently discovered ultra-faint dSphs, e.g. Segue 1 and Leo V, whose masses of $\\sim10^5\\msun$ \\citep{geha09,walker09a} fall well below this bound. This discrepancy has been addressed in work by \\citet{walker09b}, where they evaluate the mass of dSphs within the half-light radius (r$_{1/2}$). Their analysis shows that these ultra faint dSphs may simply be more deeply embedded within dark matter halos that are similar to those inhabited by their larger and brighter counterparts (also discussed in \\citealt{penarrubia08a}), which could suggest the existence of a ``universal'' dark matter halo. Similar results are shown in \\citet{wolf09}, where they conclude that all the Milky Way dSphs are consistent with inhabiting dark matter halos with a mass M$_{halo}\\sim3\\times10^9\\msun$. The recently discovered faint dSphs described above populate the faint end of the satellite luminosity function, and their spectroscopic analysis would allow us to see if these objects are indeed highly dark matter dominated, and whether they conform to such a mass limit.  Until recently, only a handful of the M31 dSphs have been analysed spectroscopically (And II, \\andix, \\andxii, And XIV, And XV, And XVI and And X \\citep{cote99,chapman06,chapman07,majewski07,letarte09,kalirai09}, with only one of these (And XII by Chapman et al. 2007) probing the fainter end of the luminosity function (M$_V>$-8). A recent paper by \\citet{kalirai10} has extended the number of bright dSphs surveyed, presenting new data for And I, II and III, but the kinematic properties of M31's lowest luminosity dSphs are still unknown. Spectroscopic observations of these objects would allow us to see how they compare to both their brighter M31 brethren, and their MW cousins of a comparable luminosity.   As a step towards this end, we have initiated a kinematic survey of the trio of faint dwarfs discovered by \\citet{martin06}, \\andxi, \\andxii\\ and \\andxiii\\ using the DEep Imaging Multi-Object Spectrograph (DEIMOS) on Keck~II. We combine these kinematics with deep photometry from Suprime-Cam on the Subaru telescope. The analysis of our observations is performed here for these objects, including a re-treatment of \\andxii, whose orbital properties were modelled in \\citet{chapman07}, but whose detailed properties were not assessed. We also include a re-examination of the faint satellite, \\andix, using the same spectroscopic data as \\citet{chapman05}, as there was an error in the geometric modelling of the slit mask used to calculate the velocities, which led to a misclassification of member stars; this has since been corrected for, and we also present additional Subaru photometry for the satellite. The outline of the paper is as follows: \\S~2 presents the observing strategy, the observations and how they were reduced ; \\S~3 is dedicated to the analysis of \\andxi, \\andxii\\, \\andxiii\\ and \\andix; we discuss our findings in \\S~4 and provide conclusions in \\S~5.  \\begin{figure*} \\begin{center} \\includegraphics[angle=0,width=0.33\\hsize]{and11_verr.eps} \\includegraphics[angle=0,width=0.33\\hsize]{and12_verr.eps} \\includegraphics[angle=0,width=0.33\\hsize]{and13_verr.eps} \\includegraphics[angle=0,width=0.33\\hsize]{and9_verr.eps} \\caption{CFHT-MegaCam color-magnitude diagrams (left panel of each Figure) , where the native MegaCam g and i bands have been converted into V and I magnitudes using transformations detailed in \\citet{mcconnachie04b,ibata07}, and radial velocity uncertainties (right panel of each Figure) of the observed stars in the four objects surveyed in the paper.  The top left panel corresponds to \\andxi, the top right to \\andxii\\ and the bottom left and right correspond to \\andxiii\\ and \\andix\\ respectively. The small dots represent CFHT objects in the region, filled circles represent target stars selected in the satellites CMD Red Giant Branch features. Circled big points represent stars that were initially selected as field stars but have the radial velocity and the CMD position of dwarf stars and thus were used to derive the dwarf parameters.} \\label{cmds} \\end{center} \\end{figure*}  \\begin{table*} \\begin{center} \\caption{Details of the Keck/DEIMOS observations of \\andxi, XII, XIII and IX} \\label{masks} \\begin{tabular}{llllccccc} \\hline Date observed & Mask $\\alpha_{J2000}$  & Mask $\\delta_{J2000}$  & Object & Grating & Field P.A. & Exp. time (s) & No. targets & No. reduced  \\\\ \\hline 23/09/2006     & 00:46:21 & +33:48:21 &\\andxi    & 600 line/mm  & 0$\\deg$ & 3$\\times$1200 & 35  & 32 \\\\ 21/09/2006     & 00:47:27 & +34:22:30 & \\andxii  & 600 line/mm  & 0$\\deg$ & 3$\\times$1200 & 50 & 48 \\\\ 23/09/2006     & 00:47:27 & +34:22:30 & \\andxii  & 600 line/mm  & 0$\\deg$ & 3$\\times$1200 & 50 & 48 \\\\ 23/09/2006     & 00:51:51 & +33:00:16 & \\andxiii & 600 line/mm  & 0$\\deg$ & 3$\\times$1200 & 49 & 48 \\\\ 13/09/2004     & 00:52:51 & +43:11:48 &\\andix    & 1200 line/mm & 0$\\deg$ & 3$\\times$1200 & 648 & 277 \\\\ \\hline \\end{tabular} \\end{center} \\end{table*} ",
        "conclusions": "We have studied the four faint M31 satellites, \\andix, \\andxi, \\andxii\\ and \\andxiii, using spectroscopic data from KeckII/DEIMOS and Subaru/Suprime-Cam. Using our spectroscopic results, we have identified probable members for each galaxy and constrained their systemic velocities. Owing to our low S:N and small sample sizes, we are unable to resolve a velocity dispersion for \\andxi. Instead we use the 1$\\sigma$ confidence level from the maximum likelihood technique as an upper limit for the satellite. In the case of \\andix, \\andxii\\ and \\andxiii\\ we resolve $\\sigma_v$=4.5$^{+3.6}_{-3.4}$kms$^{-1}$, $\\sigma_v$=3.8$^{+3.9}_{-3.8}$kms$^{-1}$ and $\\sigma_v$=9.7$^{+8.9}_{-4.5}\\kms$ for each dwarf respectively, though we note that \\andxii\\ is consistent with 0~kms$^{-1}$ within its 1$\\sigma$ errors and the measurement of \\andxiii's dispersion is based on only three stars. Measuring the average spectroscopic metallicity of each satellite from their composite spectra, we find that all four satellites are metal-poor (\\FeH$\\sim$-2.0). Such metallicities are comparable to what is observed in MW dSphs of a similar luminosity, suggesting that the established relation between M$_V$ and \\FeH\\ observed in both the brighter classical MW dwarfs \\citep{mateo98} and the fainter MW dSphs \\citep{kirby08} continues down to luminosities of M$_V>-6.4$ in the M31 dSphs population.  Finally, we place the velocity dispersions and masses for these dSphs in the context of the universal dark matter halo profiles established for the MW dwarf galaxies in \\citet{walker09b} and find that both \\andix\\ and \\andxii\\ have lower dispersions than would be expected for dSphs of their size, similar to a number of dSphs studied by \\citet{kalirai10}. If these results are verified, it could imply that the dSph population of M31 is dynamically colder than that of the MW, which is unexpected, given their larger spatial extent. Such results also imply that the mass contained within the half-light radii of these dwarfs is also lower than what is typically observed in their MW counterparts of a similar size, and suggests that these dSphs may reside in lower density dark matter halos.  The results we have presented here indicate that, in terms of their dynamical properties, these satellites at the faint end of the dwarf spheroidal spectrum do not behave as expected from previous studies. Further spectroscopic observations of faint objects similar to these (e.g., those found by \\citealt{mcconnachie08}) will undoubtedly help us to better understand this low luminosity regime, and shed light onto the behaviour of these faintest of galaxies."
    },
    "0911/0911.1686_arXiv.txt": {
        "abstract": "We present the results of a search for variable stars in the globular cluster \\objectname{NGC~5286}, which has recently been suggested to be associated with the Canis Major dwarf spheroidal galaxy. 57 variable stars were detected, only 19 of which had previously been known. Among our detections one finds 52 RR Lyrae (22 \\RRc\\ and 30 \\RRab), 4 LPV's, and 1 type II Cepheid of the BL Herculis type. Periods are derived for all of the RR Lyrae as well as the Cepheid, and {\\it{BV}} light curves are provided for all the variables.  The mean period of the \\RRab\\ variables is $\\left\\langle P_{ab} \\right\\rangle = 0.656$~days, and the number fraction of \\RRc\\ stars is $N_c/N_{RR} = 0.42$, both consistent with an Oosterhoff II (OoII) type~-- thus making NGC~5286 one of the most metal-rich (${\\rm [Fe/H]} = -1.67$; \\citeauthor{har1996} \\citeyear{har1996}) OoII globulars known to date. The minimum period of the \\RRab's, namely $P_{\\rm ab,min} = 0.513$~d, while still consistent with an OoII classification, falls towards the short end of the observed  $P_{\\rm ab,min}$ distribution for OoII globular clusters. As was recently found in the case of the prototypical OoII globular cluster M15 (NGC~7078), the distribution of stars in the Bailey diagram does not strictly conform to the previously reported locus for OoII stars.  We provide Fourier decomposition parameters for all of the RR Lyrae stars detected in our survey, and discuss the physical parameters derived therefrom. The values derived for the \\RRc's are not consistent with those typically found for OoII clusters, which may be due to the cluster's relatively high metallicity~-- the latter being confirmed by our Fourier analysis of the ab-type RR Lyrae light curves. Using the recent recalibration of the RR Lyrae luminosity scale by \\citeauthor{cc08}, we derive for the cluster a revised distance modulus of $(m-M)_V = 16.04$~mag. ",
        "introduction": "\\label{sec:intro} NGC~5286 (C1343-511) is a fairly bright ($M_V = -8.26$) and dense globular cluster (GC), with a central luminosity density $\\rho_0 \\approx 14,\\!800 \\, L_{\\odot}/{\\rm pc}^3$~-- which is more than a factor of six higher than in the case of $\\omega$~Centauri (NGC~5139), according to the entries in the \\citet{har1996} catalog. In \\citet[][hereafter Paper~I]{mzea08} we presented a color-magnitude diagram (CMD) study of the cluster that reveals an unusual horizontal branch (HB) morphology in that it does not contain a prominent red HB component, contrary to what is normally found in GCs with comparable metallicity (${\\rm [Fe/H]} = -1.67$; \\citeauthor{har1996} \\citeyear{har1996}), such as M3 (NGC~5272) or M5 (NGC~5904). As a matter of fact, NGC~5286 contains blue HB stars reaching down all the way to at least the main sequence turnoff level in $V$. Yet, unlike most blue HB GCs, NGC 5286 is known to contain a sizeable population of RR Lyrae variable stars, with at least 15 such variables being known in the field of the cluster \\citep{cle2001}. In this sense, NGC~5286 resembles the case of M62 \\citep[NGC~6266;][]{rcea05}, thus possibly being yet another member of a new group of GCs with HB types intermediate between M13 (NGC~6205)-like (a very blue HB with relatively few RR Lyrae variables) and that of the Oosterhoff I (Oo I) cluster M3 (a redder HB, with a well-populated instability strip). NGC~5286 thus constitutes an example of the ``missing link'' between M3- and M13-like GCs \\citep*{vcea87}.  Previous surveys for variable stars in NGC~5286 \\citep[e.g.,][]{ll78,agea97} have turned up relatively large numbers of RR Lyrae stars. However, such studies were carried out either by photographic methods, used comparatively few observations, or utilized reduction methods that have subsequently been superseded by improved techniques, including robust multiple-frame photometry \\citep[e.g., ALLFRAME;][]{pbs94} and image subtraction \\citep[e.g., ISIS;][]{Al2000}. This, together with the large central surface brightness of the cluster, strongly suggests that a large population of variable stars remains unknown in NGC~5286, especially towards its crowded inner regions. In addition, for the known or suspected variables, it should be possible to obtain light curves of much superior quality to those available, thus leading to better defined periods, amplitudes, and Fourier decomposition parameters.  Indeed, to our knowledge, no modern variability study has ever been carried out for this cluster. A study of its variable star population appears especially interesting in view of its suggested association with the Canis Major dwarf spheroidal galaxy \\citep*{crane2003,forbes2004}, and the constraints that the ancient RR Lyrae variable stars are able to pose on the early formation history of galaxies \\citep*[e.g.,][]{mc04,mc07,cat2005,tkea04,cmea09}. Therefore, the time seems ripe for a reassessment of the variable star content of NGC~5286~-- and this is precisely the main subject of the present paper.   In \\S\\ref{sec:obs}, we describe the variable stars search techniques and the conversion from ISIS relative fluxes to standard magnitudes. In \\S\\ref{sec:var}, we show the results of our variability search, giving the positions, periods, amplitudes, magnitudes, and colors for the detected variables. We show the positions of the variables in the cluster CMD in \\S\\ref{sec:cmd}. In \\S\\ref{sec:four}, we provide the results of a Fourier decomposition of the RR Lyrae light curves, obtaining several useful physical parameters. We analyze the cluster's Oosterhoff type in \\S\\ref{sec:oos}, whereas \\S\\ref{sec:cep} is dedicated to the type II Cepheid that we found in NGC~5286. \\S\\ref{sec:summ} summarizes the main results of our investigation. All of the derived light curves are provided in an Appendix.   \\begin{center} \\begin{deluxetable*}{lccccccccccccl} \\tablewidth{0pc} \\tabletypesize{\\scriptsize} \\tablecaption{Photometric Parameters for NGC~5286 Variables} \\tablehead{\\colhead{ID} & \\colhead{RA (J2000)} & \\colhead{DEC (J2000)} & \\colhead{$P$} & \\colhead{$A_B$} & \\colhead{$A_V$} & \\colhead{$(B)_{\\rm mag}$} & \\colhead{$(V)_{\\rm mag}$} & \\colhead{$\\langle B \\rangle$} & \\colhead{$\\langle V \\rangle$} &  \\colhead{$(\\bv)_{\\rm mag}$} &  \\colhead{$\\langle B \\rangle - \\langle V \\rangle$} &  \\colhead{$(\\bv)_{\\rm st}$} & \\colhead{Comments} \\\\ \\colhead{} & \\colhead{(h:m:s)} & \\colhead{(deg:m:s)} & \\colhead{(days)} & \\colhead{(mag)} & \\colhead{(mag)} & \\colhead{(mag)} & \\colhead{(mag)} & \\colhead{(mag)} & \\colhead{(mag)} & \\colhead{(mag)} & \\colhead{(mag)} & \\colhead{(mag)} & \\colhead{} } \\startdata V01...........  \t& $13:46:21.5$& $-51:20:03.8$ & $0.635$ & $1.26$ & $0.99$ & $17.435$ & $16.743$  & $17.539$ & $16.818$ & $0.692$ & $0.721$ & $0.680$ &\\RRab\\ \\\\ V02...........  \t& $13:46:35.1$& $-51:23:14.8$ & $0.611$ & $0.82$ & $0.71$ & $17.517$ & $17.002$ & $17.578$ & $17.047$ & $0.515$ & $0.530$ & $0.507$ & \\RRab\\ \\\\ V03...........  \t& $13:46:54.1$& $-51:23:08.9$ & $0.685$ & $0.92$ & $0.73$ & $17.216$ & $16.697$ & $17.299$ & $16.744$ & $0.520$ & $0.555$ & $0.510$ & \\RRab\\ \\\\ V04...........  \t& $13:46:19.2$& $-51:23:47.2$ & $0.352$ & $0.60$ & $0.46$ & $17.067$ & $16.639$ & $17.114$ & $16.668$ & $0.428$ & $0.446$ & $0.423$ &\\RRc\\ \\\\ V05...........  \t& $13:46:33.4$& $-51:22:03.0$ & $0.5873$ & $...$  & $1.31$ & $...$ & $17.383$ & $...   $ & $17.518$ & $...  $ & $...  $ & $...  $ &\\RRab\\  \\\\ V06...........  \t& $13:46:32.9$& $-51:23:04.3$ & $0.646$ & $1.38$ & $1.08$ & $16.850$ & $16.363$ & $16.990$ & $16.485$ & $0.487$ & $0.504$ & $0.476$ &\\RRab\\ \\\\ V07...........  \t& $13:46:29.4$& $-51:23:38.4$ & $0.512$ & $...$  & $...$  & $17.08:$ & $16.58:$ & $....  $ & $....  $ & $0.50:$ & $.... $ & $.... $ &\\RRab\\ \\\\ V08......  \t& $13:46:28.5$& $-51:23:10.2$ & $2.33$ & $1.24$ & $1.15$ & $15.569$ & $15.099$ & $15.730$ & $15.259$ & $0.471$ & $0.471$ & $...  $ &BL Her \\\\ V09...........  \t& $13:46:39.0$& $-51:22:00.0$ & $0.3003$ & $0.63$ & $0.48$ & $17.266$ & $16.873$ & $17.314$ & $16.904$ & $0.393$ & $0.411$ & $0.387$ &\\RRc\\ \\\\ V10.........  \t& $13:46:24.0$& $-51:22:46.4$ & $0.569$ & $1.29$ & $1.18$ & $17.529$ & $17.141$ & $17.635$ & $17.241$ & $0.388$ & $0.394$ & $0.376$ &\\RRab\\ \\\\ V11.........  \t& $13:46:26.2$& $-51:23:35.7$ & $0.652$ & $0.94$ & $0.69$ & $17.120$ & $16.530$ & $17.194$ & $16.519$ & $0.590$ & $0.676$ & $0.580$ &\\RRab\\ \\\\ V12.........  \t& $13:46:09.1$& $-51:22:38.9$ & $0.356$ & $0.65$ & $0.48$ & $17.305$ & $16.770$ & $17.352$ & $16.803$ & $0.535$ & $0.550$ & $0.528$ &\\RRc\\ \\\\ V13.........  \t& $13:46:33.4$& $-51:23:33.8$ & $0.294$ & $0.65$ & $0.49$ & $16.871$ & $16.518$ & $16.919$ & $16.548$ & $0.353$ & $0.371$ & $0.346$ &\\RRc\\ \\\\ V14.........  \t& $13:46:23.2$& $-51:23:38.9$ & $0.415$ & $0.60$ & $0.47$ & $16.994$ & $16.535$ & $17.037$ & $16.562$ & $0.458$ & $0.474$ & $0.454$ &\\RRc\\ \\\\ V15.........  \t& $13:46:24.1$& $-51:22:56.8$ & $0.585$ & $1.78$ & $1.19$ & $17.539$ & $16.825$ & $17.768$ & $16.919$ & $0.713$ & $0.849$ & $0.689$ &\\RRab\\ \\\\ V17.........  \t& $13:46:34.6$& $-51:23:29.0$ & $0.733$ & $0.83$ & $0.23$ & $17.172$ & $16.581$ & $17.182$ & $16.584$ & $0.591$ & $0.598$ & $0.582$ &\\RRab\\ \\\\ V18.........  \t& $13:46:33.9$& $-51:23:16.0$ & $0.781$ & $0.20$ & $...$  & $17.265$ & $...$  & $17.280$ & $...$  & $...$  & $...$  & $...$  &\\RRab\\ \\\\ V20.........\t& $13:46:25.2$& $-51:21:38.6$ & $0.319$ & $0.39$ & $0.35$ & $17.077$ & $16.614$ & $17.095$ & $16.628$ & $0.463$ & $0.467$ & $0.467$ &\\RRc\\ \\\\ V21.........\t& $13:46:25.8$& $-51:24:02.7$ & $0.646$ & $0.89$ & $0.73$ & $17.176$ & $16.673$ & $17.246$ & $16.718$ & $0.503$ & $0.528$ & $0.493$ &\\RRab\\ \\\\ NV1........\t& $13:46:27.4$& $-51:25:46.1$ & $0.366$ & $0.54$ & $0.43$ & $17.156$ & $16.695$ & $17.199$ & $16.722$ & $0.461$ & $0.476$ & $0.459$ &\\RRc\\ \\\\ NV2........  \t& $13:46:17.7$& $-51:23:56.8$ & $0.354$ & $0.60$ & $0.50$ & $17.284$ & $16.758$ & $17.336$ & $16.789$ & $0.526$ & $0.547$ & $0.521$ &\\RRc\\ \\\\ NV3........  \t& $13:46:30.5$& $-51:23:32.5$ & $0.755$ & $0.55$ & $0.30$ & $17.292$ & $16.501$ & $17.321$ & $16.503$ & $0.791$ & $0.819$ & $0.788$ &\\RRab\\ \\\\ NV4........  \t& $13:46:15.9$& $-51:23:31.7$ & $0.786$ & $0.28$ & $0.21$ & $17.099$ & $16.450$ & $17.108$ & $16.455$ & $0.649$ & $0.653$ & $0.660$ &\\RRab\\ \\\\ NV5........  \t& $13:46:17.6$& $-51:20:33.0$ & $0.357$ & $0.60$ & $0.45$ & $17.166$ & $16.659$ & $17.207$ & $16.685$ & $0.507$ & $0.522$ & $0.502$ &\\RRc\\ \\\\ NV6........  \t& $13:46:26.4$& $-51:23:27.6$ & $0.566$ & $1.45$ & $1.12$ & $16.888$ & $16.511$ & $17.028$ & $16.604$ & $0.377$ & $0.425$ & $0.366$ &\\RRab\\ \\\\ NV7........  \t& $13:46:25.8$& $-51:21:34.6$ & $0.339$ & $0.48$ & $0.41$ & $16.986$ & $16.560$ & $17.017$ & $16.585$ & $0.426$ & $0.431$ & $0.427$ &\\RRc\\ \\\\ NV8........  \t& $13:46:26.2$& $-51:23:12.4$ & $0.80$ & $0.30$ & $0.30$ & $17.103$ & $16.463$ & $17.113$ & $16.473$ & $0.640$ & $0.640$ & $0.649$ &\\RRab\\ \\\\ NV9........  \t& $13:46:24.6$& $-51:23:11.6$ & $0.745$ & $... $ & $... $ & $17.24:$ & $16.58:$ & $...   $ & $... $ & $0.66:$ & $... $ & $...$ &\\RRab\\ \\\\ NV10......  \t& $13:46:25.1$& $-51:23:02.2$ & $0.339$ & $0.68$ & $0.54$ & $16.922$ & $16.436$ & $16.984$ & $16.454$ & $0.485$ & $0.529$ & $0.477$ &\\RRc\\ \\\\ NV11......  \t& $13:46:26.0$& $-51:23:02.7$ & $0.536$ & $0.85$ & $0.71$ & $16.813$ & $16.347$ & $16.889$ & $16.402$ & $0.466$ & $0.487$ & $0.457$ &\\RRab\\ \\\\ NV12......  \t& $13:46:26.7$& $-51:23:02.5$ & $0.905$ & $0.52$ & $0.41$ & $16.876$ & $16.251$ & $16.906$ & $16.269$ & $0.625$ & $0.637$ & $0.622$ &\\RRab\\ \\\\ NV13......  \t& $13:46:27.3$& $-51:22:58.2$ & $0.583$ & $0.95$ & $1.35$ & $16.514$ & $16.083$ & $16.585$ & $16.215$ & $0.431$ & $0.370$ & $0.421$ &\\RRab\\ \\\\ NV14......  \t& $13:46:29.2$& $-51:22:07.8$ & $0.284$ & $0.64$ & $0.44$ & $17.319$ & $16.803$ & $17.366$ & $16.826$ & $0.516$ & $0.540$ & $0.509$ &\\RRc\\ \\\\ NV15......  \t& $13:46:25.5$& $-51:22:53.1$ & $0.742$ & $0.34$ & $0.52$ & $17.128$ & $16.477$ & $17.138$ & $16.505$ & $0.651$ & $0.633$ & $0.657$ &\\RRab\\ \\\\ NV16......  \t& $13:46:27.0$& $-51:22:50.7$ & $0.366$ & $0.54$ & $0.32$ & $17.125$ & $16.618$ & $17.158$ & $16.629$ & $0.507$ & $0.529$ & $0.505$ &\\RRc\\ \\\\ NV17......  \t& $13:46:26.6$& $-51:22:49.3$ & $0.322$ & $0.49$ & $0.50$ & $16.923$ & $16.490$ & $16.946$ & $16.431$ & $0.433$ & $0.415$ & $0.433$ &\\RRc\\ \\\\ NV18......  \t& $13:46:29.1$& $-51:22:46.9$ & $0.362$ & $0.57$ & $0.44$ & $17.257$ & $16.410$ & $17.295$ & $16.431$ & $0.846$ & $0.864$ & $0.843$ &\\RRc\\ \\\\ NV19......  \t& $13:46:26.8$& $-51:22:44.2$ & $0.658$ & $0.94$ & $1.70$ & $16.902$ & $16.624$ & $16.924$ & $16.706$ & $0.278$ & $0.218$ & $0.268$ &\\RRab\\ \\\\ NV20......  \t& $13:46:22.5$& $-51:22:43.2$ & $0.3103$ & $1.90$ & $0.65$ & $17.555$ & $16.753$ & $18.023$ & $16.791$ & $0.802$ & $1.232$ & $0.736$ &\\RRc\\ \\\\ NV21......  \t& $13:46:29.4$& $-51:22:41.0$ & $0.570$ & $0.95$ & $1.06$ & $16.830$ & $16.583$ & $16.902$ & $16.669$ & $0.247$ & $0.233$ & $0.237$ &\\RRab\\ \\\\ NV22......  \t& $13:46:27.5$& $-51:22:39.9$ & $0.68$ & $1.15$ & $1.20$ & $16.829$ & $16.273$ & $17.022$ & $16.374$ & $0.556$ & $0.648$ & $0.544$ &\\RRab\\ \\\\ NV23......  \t& $13:46:26.7$& $-51:22:16.1$ & $0.598$ & $1.14$ & $1.30$ & $17.090$ & $16.609$ & $17.182$ & $16.729$ & $0.481$ & $0.453$ & $0.469$ &\\RRab\\ \\\\ NV24......  \t& $13:46:29.1$& $-51:22:39.1$ & $0.60$ & $0.75$ & $0.70$ & $16.927$ & $16.346$ & $16.965$ & $16.382$ & $0.581$ & $0.583$ & $0.574$ &\\RRab\\ \\\\ NV25......  \t& $13:46:24.2$& $-51:22:17.2$ & $0.550$ & $1.44$ & $1.20$ & $17.328$ & $16.724$ & $17.048$ & $16.664$ & $0.605$ & $0.384$ & $0.593$ &\\RRab\\ \\\\ NV26......  \t& $13:46:26.0$& $-51:22:36.5$ & $0.364$ & $0.46$ & $0.36$ & $17.011$ & $16.523$ & $17.040$ & $16.536$ & $0.489$ & $0.505$ & $0.490$ &\\RRc\\ \\\\ NV27......  \t& $13:46:30.1$& $-51:22:22.0$ & $0.706$ & $1.17$ & $0.80$ & $16.952$ & $16.599$ & $17.489$ & $16.811$ & $0.490$ & $0.530$ & $0.340$ &\\RRab\\ \\\\ NV28......  \t& $13:46:29.8$& $-51:22:35.4$ & $0.540$ & $0.97$ & $0.70$ & $16.880$ & $16.390$ & $16.954$ & $16.424$ & $0.490$ & $0.530$ & $0.479$ &\\RRab\\ \\\\ NV29......  \t& $13:46:28.8$& $-51:22:33.8$ & $0.301$ & $0.97$ & $0.63$ & $17.458$ & $16.963$ & $17.565$ & $17.010$ & $0.495$ & $0.555$ & $0.473$ &\\RRc\\ \\\\ NV30......  \t& $13:46:29.1$& $-51:22:24.1$ & $0.72$ & $0.65$ & $0.55$ & $17.251$ & $16.641$ & $17.292$ & $16.672$ & $0.610$ & $0.620$ & $0.605$ &\\RRab\\ \\\\ NV31......  \t& $13:46:25.5$& $-51:22:31.6$ & $0.289$ & $0.41$ & $0.40$ & $17.163$ & $16.519$ & $17.189$ & $16.535$ & $0.645$ & $0.654$ & $0.648$ &\\RRc\\ \\\\ NV32......  \t& $13:46:25.2$& $-51:22:30.0$ & $0.283$ & $0.34$ & $0.24$ & $16.804$ & $16.339$ & $16.817$ & $16.344$ & $0.465$ & $0.473$ & $0.471$ &\\RRc\\ \\\\ NV33......  \t& $13:46:27.9$& $-51:22:26.6$ & $0.294$ & $0.34$ & $0.50$ & $16.801$ & $16.457$ & $16.815$ & $16.483$ & $0.345$ & $0.333$ & $0.351$ &\\RRc\\ \\\\ NV34......  \t& $13:46:26.7$& $-51:22:24.9$ & $0.367$ & $0.42$ & $0.48$ & $16.569$ & $16.073$ & $16.591$ & $16.096$ & $0.496$ & $0.495$ & $0.499$ &\\RRc\\ \\\\ NV35......  \t& $13:46:24.9$& $-51:23:03.8$ & $...$ & $...$ & $...$ & $14.70:$ & $13.38:$ & $...$ & $...$ & $1.32:$ & $...$ & $...$ & LPV \\\\ NV36......  \t& $13:46:29.1$& $-51:22:58.4$ & $...$ & $...$ & $...$ & $15.30:$ & $13.40:$ & $...$ & $...$ & $1.90:$ & $...$ & $...$ & LPV \\\\ NV37......  \t& $13:46:28.8$& $-51:22:12.9$ & $...$ & $...$ & $...$ & $15.70:$ & $14.60:$ & $...$ & $...$ & $1.10:$ & $...$ & $...$ & LPV \\\\ NV38......  \t& $13:46:28.8$& $-51:20:37.0$ & $...$ & $...$ & $...$ & $15.30:$ & $13.74:$ & $...$ & $...$ & $1.56:$ & $...$ & $...$ & LPV \\enddata \\label{tab1} \\end{deluxetable*} \\end{center} ",
        "conclusions": "\\label{sec:summ} In this paper, we present the results of time-series photometry for NGC~5286, a GC which has been tentatively associated with the Canis Major dwarf spheroidal galaxy. 38 new variables were discovered in the cluster, and 19 previously known ones were recovered in our study (including one BL Her star that was previously catalogued as an RR Lyrae). The population of variable stars consists of 52 RR Lyrae (22 \\RRc\\ and 30 \\RRab), 4 LPV's, and 1 type II Cepheid.  From Fourier decomposition of the \\RRab\\ light curves, we obtained a value for the metalicity of the cluster of ${\\rm [Fe/H]} = -1.68  \\pm 0.21$~dex in the \\citet{zw1984} scale. We also derive a distance modulus of $(m-M)_V = 16.04$~mag for NGC~5286, based on the recent RR Lyrae $M_V - {\\rm [Fe/H]}$ calibration of \\citet{cc08}.  Using a variety of indicators, we discuss in detail the Oosterhoff type of the cluster, concluding in favor of an OoII classification. The cluster's fairly high metallicity places it among the most metal-rich OoII clusters known, which may help account for what appears to be a fairly unusual behavior for a cluster of this type, including relatively short values of $P_{\\rm ab,min}$ and $\\langle P_c\\rangle$, and unusual physical parameters, as derived for its c-type RR Lyrae stars on the basis of Fourier decomposition of their light curves.  In regard to the cluster's suggested association to the Canis Major dwarf spheroidal galaxy, we note that the metallicity and distance modulus derived in this work are very similar to the values previously accepted for the cluster \\citep{har1996}, and thus the conclusions reached by previous authors \\citep{crane2003,forbes2004} regarding its possible association with this dwarf galaxy are not significantly affected by our new metallicity and distance estimates. In addition, the position of the cluster in the HB morphology-metallicity plane is fairly similar to that found in several nearby extragalactic systems. As far as NGC~5286's RR Lyrae pulsation properties are concerned, the present study shows them to be somewhat unusual compared with bona-fide Galactic globular clusters, but still do not classify the cluster as an Oosterhoff-intermediate system, as frequently found among the Galaxy's dwarf satellites \\citep[e.g.,][and references therein]{cat2005}. It is interesting to note, in any case, that the Canis Major field, unlike what is found among other dwarf galaxies, appears to be basically devoid of RR Lyrae stars \\citep{tkea04,cmea09}, and also to be chiefly comprised of fairly high-metallicity (${\\rm [M/H]} \\geq -0.7$), young ($t \\lesssim 10$~Gyr) stars \\citep[e.g.,][]{mbea04}. The present paper, along with Paper~I, show instead that NGC~5286 is an RR Lyrae-rich, metal-poor globular cluster that is at least as old as the oldest globular clusters in the Galactic halo. It is not immediately clear that such an object as NGC~5286 would be easily formed within a galaxy with the properties observed for the Canis Major main body~-- and this should be taken into account when investigating the physical origin and formation mechanism for the Canis Major overdensity and its associated tidal ring."
    },
    "0911/0911.4295_arXiv.txt": {
        "abstract": "Correct interpretation of acoustic travel times measured by time-distance helioseismology is essential to get an accurate understanding of the solar properties that are inferred from them. It has long been observed that sunspots suppress $p$-mode amplitude, but its implications on travel times has not been fully investigated so far. It has been found in test measurements using a  'masking' procedure, in which the solar Doppler signal in a localized quiet region of the Sun is artificially suppressed by a spatial function, and using numerical simulations that the amplitude modulations in combination with the phase-speed filtering may cause systematic shifts of acoustic travel times. To understand the properties of this procedure, we derive an analytical expression for the cross-covariance of a signal that has been modulated locally by a spatial function that has azimuthal symmetry,  and then filtered by a phase speed filter typically used in time-distance helioseismology. Comparing this expression to the  Gabor wavelet fitting formula without this effect, we find that there is a shift in the travel times, that is introduced by the amplitude modulation. The analytical model presented in this paper can be useful also for interpretation of travel time measurements for non-uniform distribution of oscillation amplitude due to observational effects. ",
        "introduction": "\\label{sec:intro}  Time-distance helioseismology \\citep{duv93} is a  local helioseismological technique that has been used to study  meridional flows, flows and  sound speed perturbations beneath sunspots \\citep[e.g.,][]{kos97, gil97, zha01}. It measures the time  for a wavepacket to travel between any two points on the solar surface, by computing a temporal cross-covariance between the  Doppler time series at the two points. The travel time is then inverted to infer various  properties that are useful to map local structures of the sun. These results are interesting, as  they complement the results obtained from global helioseismology based on inversion of normal mode frequencies. However, many aspects of time-distance helioseismology are still not fully understood. In particular, it has been observed that sunspots suppress the $p$-mode amplitude appreciably, but its consequences on travel times and the properties derived by inverting them has largely remained unexplored.  The commonly used procedure of measuring the phase and group travel times of acoustic waves is based a Gabor-wavelet fitting formula derived by \\citet{kos97} for cross-covariance of randomly excited oscillation modes of the quiet Sun (represented by a spherically symmetric model). Initially, this procedure included only a broad-band frequency filtering of the solar oscillation signal, centered at the peak of acoustic power. Later, it was modified by adding a phase-speed filtering procedure in order to isolate the first-bounce signals (direct waves without additional reflections from the surface) and to improve the signa-to-noise ratio \\citep{Duvall1997}. The phase-speed filtering is important for the analysis of acoustic-wave packets traveling short travel distances, e.g. less than $1$ heliographic degree ($\\sim 12$ Mm). For larger distances the travel times can be measured without the phase-speed filtering, and it is found that the use of the phase-speed filtering does not significantly affects the measurement results. For shorter distances the influence of the phase-speed filtering may be significant, and has to be taken account.  In this paper we study the effects of the phase-speed filtering and local spatial suppression of Doppler velocity signals  from a quiet patch on the Sun, on the travel times. In this `masking' procedure suggested by \\citet{raj06} to simulate this effect in sunspots the oscillation signal of a quiet Sun region is multiplied by an inverted modulation function of spatial coordinates with azimuthal symmetry.  This function is called a mask, and is not a function of time. One should note that while this procedure models a spatial modulation of acoustic power, it does not represent the modulation observed in sunspots, where the amplitude of acoustic changes due do several physical factors, such as reduced excitation, absorption, changing wave speed, and spectral line formation observing conditions. Recently, \\citet{Chou2009} attempted to quantify these contributions using observational data. The variations of the oscillation power in sunspots have not been explained by theory or simulations \\citep[e.g.][]{Parchevsky2007}. Nevertheless, it is interesting to investigate the effects of the variations on the helioseismic travel times by using the simple `masking' model bearing in mind that it may not represent the the real situation in sunspots. One advantage of the `masking' model is that it allows relatively simple analytical investigation.  To understand the results of amplitude suppression or enhancement we theoretically model the effect of masking by computing the cross-covariance of a  masked signal that has been filtered by an appropriately designed phase speed filter. Analytical expressions for the cross-covariance are derived in terms of the mask, the filter parameters, the properties of the signal and the dispersive nature of the solar medium, when the the mask is azimuthally symmetric.  This analysis will be useful to study the effect of masking on travel time measurements and the properties that are inferred from inverting the travel times. This will be especially valuable in artificially mimicking  how the sunspots influence the travel times of  $p$ modes, and also other instrumental effects that corrupt the observed signal. ",
        "conclusions": "In this paper we give an explanation as to why the localized spatial suppression or  enhancement (masking) of the acoustic signal followed by  phase speed filtering can appreciably shift the measured travel times, and cause systematic errors in time-distance inversions. Using only a  frequency filter or reversing the operation of suppression (or enhancement) and phase speed filtering does not shift the travel times, but merely scales the cross covariance. Hence, the operations of masking and phase speed filtering are non-commutative. To explain this we develop a  model and derive a new analytical formula for the cross-covariance in terms of the masking parameters, when the mask function has azimuthal symmetry. The reason for this is due to the fact that masking changes the spatial mode eigenfunctions, which when phase speed filtered, leads to a mixing of spatial modes, and is responsible for the travel time shifts. This may be useful to mimic amplitude changes in sunspots due to pure observational reasons, such as caused by altering the height of formation of the spectral line used for helioseismology measurements, and also to correct for the shifted travel times due to such effects. However, contrary to \\citet{raj06} suggestion, the procedure of masking oscillations of the quite Sun cannot model the amplitude reduction in sunspot due to physical mechanisms (e.g. changes in emissivity or wave absorption), and thus their recommendations of correcting the amplitude reduction by reversed masking may cause artificial shifts in observed travel times, and thus must be taken with caution."
    },
    "0911/0911.3360_arXiv.txt": {
        "abstract": "A significant asymmetry in the distribution of faint blue stars in the inner Galaxy, Quadrant 1 ({\\it l} $ = 20\\arcdeg - 45\\arcdeg$) compared to Quadrant 4 was first reported by \\cite{lar96}. Parker et al (2003, 2004) greatly expanded the survey to determine its spatial extent and shape and the kinematics of the affected stars. This excess in the star counts  was subsequently confirmed by \\cite{jur08} using SDSS data. Possible explanations for the asymmetry include a merger remnant, a triaxial Thick Disk, and a possible interaction with the bar in the Disk. In this paper we describe our program of wide field photometry to map the asymmetry to fainter magnitudes and therefore larger distances. To search for the signature of triaxiality, we extended our survey to higher Galactic  longitudes. We find no evidence for an excess of faint blue stars at {\\it l} $\\ge 55\\arcdeg$ including the faintest magnitude interval. The asymmetry and star count excess in Quadrant 1 is thus not due to a triaxial Thick Disk. ",
        "introduction": "Studies of both stars and gas are revealing significant structure and asymmetries in their motions and spatial distributions including the bar of stars and gas in the Galactic bulge (\\citealt{1991ApJ...379..631B}, \\citealt{1994ApJ...429L..73S}), the evidence from infrared surveys for a larger stellar bar in the inner disk (\\citealt{1992ApJ...384...81W}, \\citealt{1997MNRAS.292L..15L}, \\citealt{2005ApJ...630L.149B}), the discovery of the Sagittarius dwarf \\citep{1994Natur.370..194I,1995MNRAS.275..591I}, and a significant asymmetry of unknown origin in the distribution of faint blue stars in Quadrant 1 (Q1) of the inner Galaxy \\citep{lar96}. Each of these observations provides a significant clue to the history of the Milky Way. When combined with the growing evidence for Galactic mergers in addition to the Sagittarius dwarf, i.e. the Monoceros stream (\\citealt{2002ApJ...569..245N}, \\citealt{2003MNRAS.340L..21I}), the Canis Major merger remnant (\\citealt{2004MNRAS.355L..33M}), the Virgo  stream (\\citealt{2001ApJ...554L..33V}, \\citealt{2005ApJ...633..205M}), we now realize that the structure and evolution of our Galaxy have been significantly altered by mergers with other systems.  Indeed the population of the Galactic Halo and possibly the Thick Disk as well, may be dominated by mergers with smaller systems.  The  asymmetry in  Q1,  $l = 20\\arcdeg - 45\\arcdeg$, first recognized by \\cite{lar96} was based on a comparison with complementary fields in the fourth  Quadrant (Q4) using  star counts from the Minnesota Automated Plate Scanner Catalog of the POSS I \\citep{cab03}\\footnote{The MAPS database is online at: http://aps.umn.edu}. To map the extent and shape of the asymmetric distribution and further identify the contributing stellar population, \\cite{par03}  extended the search to 40 contiguous fields from the digitized POSS I on each side of the Sun-Center line plus the same number of fields below the plane in Q1. They examined the star count ratios for paired fields in three color ranges: blue, intermediate and red.  Over 6 million stars were used in the star count analysis covering almost 2000 square degrees on the sky. They found a 25\\% excess in the number of probable Thick Disk stars in Q1, $l \\approx 20$ to $60\\arcdeg$ and $20\\arcdeg$ to $40\\arcdeg$ above and below the plane compared to the complementary fields ($l \\approx 340$ to $300\\arcdeg$) in Q4. While the region of the asymmetry is somewhat irregular in shape, it is also fairly uniform and covers several hundred square degrees. It is therefore a major substructure in the Galaxy due to more than small scale clumpiness. Assuming that these are primarily main sequence thick disk stars, with a typical magnitude completeness limit of $V \\approx 18$ mag, they are 1 - 2 kpc from the Sun and about 0.5 to 1.5 kpc above (and below) the plane.  \\cite{par04} also found an associated kinematic signature. Using velocities from spectra for more than 700 stars, they not only found an asymmetric distribution in the v$_{LSR}$ velocities, but the Thick Disk stars in Q1 have a much slower effective rotation rate $\\omega$ compared to the corresponding Q4 stars. A solution for the radial and tangential components of the v$_{LSR}$ velocities reveals a significant lag of 80 to 90 km s$^{-1}$ in the direction of Galactic rotation for the Thick Disk stars in Q1.  Three possible explanations for the asymmetry are the fossil remnant of a merger, a triaxial Thick Disk or inner Halo,  and interaction of the Thick Disk/inner Halo stars with the bar in the Disk. Given the  lack of any spatial overlap with the path of the Sagittarius dwarf through the Halo \\citep{2003MNRAS.340L..21I}, and the predicted path of the Canis Major dwarf \\citep{2004MNRAS.355L..33M}, its   association  with either feature is unlikely. Furthermore its spatial extent and apparent symmetry relative to the plane also do not automatically support a merger interpretation. Our  line of sight to the asymmetry is interestingly in the same direction as the stellar bar in the Disk \\citep{1992ApJ...384...81W,2005ApJ...630L.149B}, and the near end of the Galactic bar \\citep{1997MNRAS.292L..15L, 2000MNRAS.317L..45H}, but  the Disk bar is approximately 5 kpc from the Sun in this direction. Thus the stars showing the excess are between the Sun and the bar, not directly above it. However the maximum extent of the  star count excess along our line of sight is not known.  Recently \\citet{lar08} showed the stellar density distribution from the MAPS scans for the \\citet{par03} fields in Q1 and Q4 above the plane, demonstrating  that the excess was in Q1 and was not due to a ring above the plane \\citep{jur08}. Larsen et al named the asymmetry feature the Hercules Thick Disk Cloud in recognition of the  direction  on the sky where the star count excess is maximum. Thus,  interpretation of the Hercules Thick Disk Cloud is not clear-cut. While it might well be a fossil merger remnant, the star count excess is  also consistent with  a triaxial Thick Disk or inner Halo as well as a gravitational interaction with the stellar bar especially given the corresponding asymmetry in the kinematics. The distinction between a triaxial Thick Disk and interaction with a Disk bar may be difficult to discern. Indeed, it is unclear which may have formed first, and one or both may be the result of mergers.   The release of the SDSS Data Release 5 (DR5) photometry in the direction of the observed asymmetry in Q1 led to the discovery of a feature at much fainter magnitudes, the distant Hercules-Aquila cloud \\citep{bel07} and also confirmed the nearer star count excess in the inner Galaxy \\citep{jur08}.  The SDSS survey however is not well designed for a comprehensive survey  of the Thick Disk inside the Solar orbit. It extends below $b = 30\\arcdeg$ in only a few directions in Q1 and has only limited coverage in Q4.  It does not have the leverage in Galactic longitude needed to discriminate among the  possible causes of the Hercules Thick Disk Cloud.  To further explore the possible origin of the observed spatial and kinematic asymmetry, we have obtained multicolor CCD images to extend the star counts to fainter magnitudes to map the spatial extent of the asymmetry along the line of sight as a function of Galactic longitude and latitude. For example, if the Thick Disk is triaxial we would expect to observe the star count excess out to greater longitudes, but the star count excess described by \\citet{par03} appears to terminate near {\\it l} $\\sim$ 55$\\arcdeg$. However, if it is triaxial in Q1 with its major axis above the bar, the stars  would be further away at the higher longitudes. By extending the star counts to fainter magnitudes, corresponding to greater distances, we can search for the asymmetry at higher longitudes, from {\\it l} of 50$\\arcdeg$ to 75$\\arcdeg$.   In this first paper we describe our observing program and present results on a search for triaxiality of the inner Halo and Thick Disk.  Our second paper will discuss the star counts across the full range of longitudes in Q1 and Q4 containing the asymmetry feature and a third paper in the series will analyze the kinematics of the associated stars. In the next section we describe the CCD observing program and data reduction. Determination of the star count ratios and their errors for the fields at the higher longitudes are discussed in \\S 3. Our  results  do not support a triaxial interpretation of the asymmetry feature and are summarized in \\S 4. ",
        "conclusions": "Our star count ratios summarized in Table~\\ref{tbl-4} show a small excess of faint blue stars at {\\it l} of 45$\\arcdeg$ and 50$\\arcdeg$ but not at the greater longitudes including the faintest magnitude intervals. These results support Parker et al's (2003) conclusion that the asymmetry feature in Q1 extends to about 55$\\arcdeg$. We thus find {\\it no evidence that the Hercules Thick Disk Cloud and the star count excess is due to a triaxial Thick Disk.}  The ratios at {\\it l} of 45$\\arcdeg$ and 50$\\arcdeg$ return to 1.0 at the faintest magnitude bin ($ 18 < V < 19$), suggesting that we may be seeing through the asymmetry feature; a possible clue to its extent along our line of sight in this direction. At $ V \\sim 19$ however, these faint blue stars will be about 2 kpc above the plane, and may not be in the Thick Disk.   We also note that the above vs. below the plane ratios in Q1 and Q4 and the Q1/Q4 ratio below the plane are consistent with the predictions of the Galactic model. There is no star count excess. This is not in agreement with \\citet{par03} who concluded that the excess exists both above and below the plane. This may be evidence that the Hercules Thick Disk Cloud is actually a debris stream. However, considering the small number of field pairs and the fact that we are apparently at the edge of the Cloud at these longitudes, we consider this a tentative result pending the analysis of the full set of fields.  Paper II will include an analysis of the full dataset, mapping the stellar excess in Galactic longitude and latitude and along the line of sight to determine the full extent of the Cloud. Our third paper will include the spectroscopy and analysis of the kinematics of the stars showing the excess."
    },
    "0911/0911.0680_arXiv.txt": {
        "abstract": "We show that energy deposited into an expanding supernova remnant by a highly magnetic ($B\\sim 5\\times 10^{14}~{\\rm G} $)  neutron star spinning at an initial period of $P_i\\approx 2-20 \\ {\\rm ms} $ can substantially brighten the light curve.   For magnetars with parameters in  this range, the rotational energy is released on a timescale of days to weeks, which is comparable to the effective diffusion time through the supernova remnant.   The late time energy injection can then be radiated without suffering overwhelming adiabatic expansion losses.  The magnetar input also produces a central bubble which sweeps  ejecta into an internal dense shell, resulting in a prolonged period of nearly constant photospheric velocity  in the observed spectra.  We derive analytic expressions for the light curve rise time and peak luminosity as a function of $B$,  $P_i$ and the properties of the supernova ejecta that allow for direct inferences about the underlying magnetar in bright supernovae.   We  perform numerical radiation hydrodynamical calculations of a few specific instances and compare the resulting light curves to observed events.   Magnetar activity is likely to impact more than a few percent of all core collapse supernovae, and may naturally explain some of the brightest events ever seen (e.g., SN~2005ap and SN~2008es) at $L \\ga10^{44}$~ergs~s$^{-1}$. ",
        "introduction": "\\label{sec:intro }  Studies of soft gamma-ray repeaters and anomalous X-ray pulsars reveal that $\\sim10\\%$ of newly born neutron stars  \\citep{Kouv_98} have dipole magnetic fields as high as  $B\\sim 10^{14}-10^{15}~{\\rm G}$ for more than 1000 years after their birth \\citep[see][]{Woods_06}. These ``magnetars'' rotate at periods of $P=5-12\\ {\\rm s}$ at an age of  $1000-10,000 $ years.  Such highly magnetized neutron stars (NSs) were theoretically predicted \\citep{Duncan_92, Thompson_93}, and most of their activity (both sporadic and persistent) must be powered by the decay of these large magnetic fields.  What remains unknown is just how highly magnetized and rapidly rotating these magnetars may be at ``birth''. Many \\citep[see][]{Bodenheimer_74, Wheeler_00, Thompson_04} have investigated the possible impact on the central engine when the magnetar is so rapidly rotating (1-3 ms) and magnetized that its subsequent spin-down can power the explosion. Cases this extreme may also be sources for ultra-high energy cosmic rays \\citep{Arons_03} or deposit enough energy in the collapsing stellar envelope to favorably shape the deep interior \\citep{Uzdensky_07, Bucci_09} for the production of a collimated relativistic flow needed for gamma-ray bursts. Such events depend on the combination of rapid rotation and high $B$ to achieve a measurable effect during the few seconds critical to the core collapse mechanism.  Building on the work of \\cite{Gaffet_77a, Gaffet_77b}, we have found that  weaker magnetic fields and less extreme spins (that do not alter the explosion mechanism)  can dramatically impact  supernovae light curves, competing with the decay of radioactive $^{56}$Ni and thermal energy in the expanding envelope. \\cite{Maeda_07} previously  raised this possibility for the Type Ib SN~2005bf, and \\cite{Woosley_09} has independently shown their relevance as well.   We show in \\S \\ref{sec:simple} that  when the timescale of the magnetar spindown, \\tp, is comparable to the effective radiative diffusion time, \\td, the resulting peak luminosity is $\\Lpeak \\sim \\Ep \\tp /\\td^2$, where \\Ep\\ is the magnetar rotational energy. Magnetars with $10^{13}~{\\rm G} < B<10^{16} \\ {\\rm G}$ and $P_i=1-30 \\ {\\rm ms}$  can produce $L_{\\rm peak}>10^{42} \\ {\\rm erg \\ s^{-1}}$. We discuss the dynamics of the energy injection in \\S \\ref{sec:cavity} and show that the magnetar blows a central bubble in the SN ejecta, forming a dense inner shell of swept-up material which affects the spectroscopic evolution.  In \\S \\ref{sec:lightcurves}, we derive analytic expressions for the luminosity, $\\Lpeak$,  and duration, $t_{\\rm peak}$,  of magnetar powered supernovae. We confirm these formulae with numerical radiation-hydrodynamical  calculations, and show how they can be inverted to infer $B$ and $P_i$ from a given light curve. We close in \\S \\ref{sec:conclusions} by discussing observed core-collapse SNe that may be powered this way, especially the ultra-bright SN~2005ap \\citep{Quimby_07} and SN~2008es \\citep{Gezari_09, Miller_09}. ",
        "conclusions": "\\label{sec:conclusions}   \\begin{figure} \\includegraphics[width=3.5in]{f4.eps} \\caption{The dependence of $\\Lpeak$ and $\\tmax$ on the initial magnetar spin and B field. The solid lines are for fixed $B_{14}=100,30,10,3,1, 0.3$ and 0.1 and varying spin period, whereas the dashed lines are for a fixed $P_i=1,3,10$ and 30 ms  and varying $B$.  This calculation assumed $\\Esn=10^{51}~{ \\rm ergs}$ and $\\Mej=5~M_\\odot$.  \\label{Fig:waycool} } \\end{figure}    We have shown that rotational energy deposition from magnetar spin-down  with initial spin periods $<30 \\ {\\rm ms}$  can substantially modify the thermal evolution of an expanding SNe remnant. For magnetars in this range, the peak luminosity reaches $10^{42} -10^{45}~{\\rm erg \\ s^{-1}}$  ($M_{\\rm Bol}= -16.3$ to $-23.8$), substantially impacting the typical core-collapse SNe lightcurve, whether it is a Type II or a Ibc event.    The highest luminosities occur when $\\tp \\sim \\td$, in which case the total energy radiated in the light curve is $E_{\\rm rad} \\sim \\Lpeak \\td \\sim \\Ep/3$.  The maximal spin of  a NS is  around 1 ms, so $E_{\\rm rad}$ cannot exceed $\\sim 10^{52}$~ergs; supernova radiating larger energies can not be explained by this mechanism. Though we know that $\\sim 10\\%$ of core collapse events make magnetars, we do not know the distribution of initial spin periods, so the prevalence of light curve dominance is difficult to predict.  For stars with remaining hydrogen, magnetar injection may explain the brighter ($M_B \\sim -19$) subclass of Type II-L SNe  noted by \\cite{Richardson_02}, i.e. 1961F, 1979C, 1980K, 1985L. The light curves of these events are difficult to explain in standard explosion models unless extreme progenitor radii ($R > 2000~R_\\odot$)  are assumed \\citep{Blinnikov_93}.  Figure~\\ref{Fig:compare} shows that a magnetar with relatively modest rotation, $P_i = 10 $~ms, in a $\\Mej = 5~\\Msun$ supernova can reach similar luminosities.  Events brighter than $M_{\\rm Bol}=-21$ ($L>8\\times 10^{43} ~{\\rm erg \\ s^{-1}}$), such as the ultrabright Type II-L SN~2005ap \\citep{Quimby_07} and SN~2008es \\citep{Gezari_09, Miller_09} require an initial magnetar spin of $<5 \\ {\\rm ms}$. Motivated by Figure~\\ref{Fig:waycool}, we found an excellent fit to the SN~2008es light curve with  $\\Bf = 2 $, $P_i = 2$~ms and $\\Mej=5 \\Msun$. Such rapidly rotating magnetars must be rare, as \\cite{Vink_06} found that the galactic supernovae remnants of known magnetars were explained with typical explosion energies of $10^{51} \\ {\\rm ergs}$. This rarity is consistent with the specific volume rate of these events; current estimates put them at no more than $\\sim 1\\%$ \\citep{Miller_09, Quimby_09} of  the local core collapse rate.   \\begin{figure} \\includegraphics[width=3.5in]{f5.eps} \\caption{Bolometric light curve calculations of magnetar energized supernovae compared to observed events. A constant opacity $\\kappa = 0.2$~g~cm$^{-2}$ is assumed. Black circles show V-band observations of the luminous Type~IIL~SN2008es \\citep{Gezari_09} with an assumed rise time of 25~days.  Red squares show R-band observations of the Type~Ic SN~2007bi \\citep{Gal-Yam_09} with an assumed rise time of 50 days.   } \\label{Fig:compare} \\end{figure}  Debate remains \\citep[see][]{Klose_04, Gaensler_05,Davies_09}  as to whether magnetars are preferentially formed from the most massive stars that collapse to NSs. If so, then we might see a prevalence of magnetar dominated light curves amongst the Ib/c SNe, which may partially explain the wide light curve diversity noted in this class. Some extreme SN~Ic, such as SN2007bi \\citep{Gal-Yam_09, Young_09} and SN~1999as \\citep{Knop_99}, which remained very bright for a long time, have been  claimed to be the pair instability explosion of a $\\sim 100~\\Msun$ star producing nearly  $5~\\Msun$ of \\Nifs.  Figure~\\ref{Fig:compare} shows that the light curve could alternatively be explained for a supernova with  $\\Mej = 20$~\\Msun\\ forming a magnetar with $\\Bf = 2, P_i = 2.5$~ms.   The spectra of SN~1999as also revealed a slowly evolving photospheric velocity and  narrow, blueshifted absorption features, suggestive of a dense shell like that predicted here \\citep{Kasen_thesis}.  On the other hand, the magnetar model may have trouble reproducing the strong iron emission lines seen in the nebular phase spectrum of SN~2007bi.  Our initial investigations have revealed that if an appreciable fraction of highly magnetic NSs are born rapidly rotating, then we should find evidence for them in the plethora of supernovae surveys, such as the Palomar Transient Factory  \\citep{Law_09}. Many open questions remain on the theoretical side, especially how the outgoing pulsar wind thermalizes in the remnant, whether there are substantial Rayleigh-Taylor instabilities, and how these could manifest themselves in the observed spectra both at late times and during the photospheric phase. Our work has outlined the regimes of relevance, and will guide future large scale computations through parameter space in an informed manner."
    },
    "0911/0911.2366_arXiv.txt": {
        "abstract": "The area of stable motion for fictitious Trojan asteroids around Uranus' equi\\-la\\-te\\-ral equilibrium points is investigated with respect to the inclination of the asteroid's orbit to determine the size of the regions and their shape. For this task we used the results of extensive numerical integrations of orbits for a grid of initial conditions around the points $L_4$ and $L_5$, and analyzed the stability of the individual orbits. Our basic dynamical model was the Outer Solar System (Jupiter, Saturn, Uranus and Neptune). We integrated the equations of motion of fictitious Trojans in the vicinity of the stable equilibrium points for selected orbits up to the age of the Solar system of $5 \\times 10^{9}$ years. One experiment has been undertaken for cuts through the Lagrange points for fixed values of the inclinations, while the semimajor axes were varied. The extension of the stable region with respect to the initial semimajor axis lies between $19.05 \\le a \\le 19.3$ AU but depends on the initial inclination. In another run the inclination of the asteroids' orbit was varied in the range $0^\\circ < i < 60^\\circ$ and the semimajor axes were fixed. It turned out that only four 'windows' of stable orbits survive: these are the orbits for the initial inclinations $0^\\circ < i < 7^\\circ$, $9^\\circ < i < 13^\\circ$, $31^\\circ < i < 36^\\circ$ and  $38^\\circ < i < 50^\\circ$. We postulate the existence of at least some Trojans around the Uranus Lagrange points for the stability window at small and also high inclinations. ",
        "introduction": "The first discovery of a Jupiter Trojan in 1906 (Achilles by Max Wolf in Heidelberg) proved that the equilateral equilibrium points in a simplified dynamical model Sun -- planet -- massless body are not only of hypothetical interest. Ever since many of such Trojan asteroids of Jupiter have been found and now we have knowledge of several thousands of objects in the 1:1 mean motion resonance with Jupiter. Several investigations show the symmetry of these two stable equilibrium points not only in the restricted three body problem (\\citet{San2002, Erd1988}) but also in the realistic dynamical model of the Outer Solar System (OSS) consisting of Jupiter, Saturn, Uranus and Neptune, like in e.g. \\citet{Dvo2005, Fre2006, Sch2004, Nes2000, Rob2005, Tsi2005a}.  But we have also evidence for the existence of asteroids in orbit around the equilateral equilibrium points of Neptune and also of Mars. Since the discovery of the first Neptune Trojan (\\citet{Chi2003}) -- the minor planet 2001 QR322 around its equilibrium point $L_5$ -- a lot of work has been dedicated to understand the dynamics of Neptune Trojans. On one hand the stability of hypothetical bodies was studied in the model of the Outer Solar System, on the other hand long term investigations tried to understand how Trojans can survive, or even get captured in these kind of orbits when one allows the large planets to migrate. An extensive stability study of the Trojans for the four gas giants has been undertaken by \\citet{Nes2002} in different dynamical models. The results of the realistic n-body simulations for clouds of hypothetical Trojans around the Lagrange point $L_4$ showed how fast the depletion of asteroids is acting for Saturn, Uranus and Neptune. For several hundreds of bodies, different eccentricities and inclinations as well as semimajor axes the integrations were carried out over 4 Gyrs. Saturn Trojans turned out to be unstable already after several $10^5$ years, Uranus Trojans were also found to be unstable but only after several million years, whereas Neptune's Trojans mostly survived for the whole integration time of 4 Gyrs. In a similar study for Uranus and Neptune Trojans \\citet{Mar2003} determined the diffusion speed of Trojans and derived even expressions for the secular frequencies using their numerical results. Both studies were undertaken for initial inclinations of the Trojans up to $i=30^{\\circ}$ and showed quite similar results for the dependence on the inclination.  Taking into account the early evolution of the Solar System the migration processes need to be included. This has been done in a paper by \\citet{Kor2004} for the Neptune Trojans where it has been shown that only $5\\%$ of an initial Trojan population could survive this early stage when Jupiter migrated inward (5.4 to 5.2 AU), and Saturn (8.7 to 9.5 AU), Uranus (16.2 to 19.2 AU) and Neptune (23 to 29 AU) migrated outwards. Three different theories for the origin of Neptune Trojans are developed and tested by \\citet{Chi2005}, they use the results to constrain the circumstances of the formation of planets and the accretion of Trojans. In a more recent study \\citet{Nes2009} in the NICE model (e.g. \\citet{Tsi2005b}) -- where Uranus and Neptune could interchange their location in the Solar System and Jupiter and Saturn go through the 2:1 Mean Motion Resonance -- it was shown that the Neptune Trojans were captured by a process quite similar to the chaotic capture of Jupiter Trojans (\\citet{Mor2005}). A similar process can be assumed for the Uranus Trojans, but up to now no Uranus Trojans have been detected which can be explained by the results of the pure dynamical studies mentioned before (\\citet{Nes2002, Mar2003}). These investigations stopped at an inclination of $i=30^{\\circ}$; the discovery of a highly-inclined Neptune Trojan (\\citet{She2006}) raised the question of the general stability of Trojans at high inclinations. In a recent study concerning the Neptune Trojans (\\citet{Zho2009}) it was shown that especially for larger inclinations Trojans may be in stable orbits. We have therefore undertaken an extensive numerical study of the -- hypothetical -- Uranus Trojan cloud for the whole range of inclinations $0^{\\circ} \\le i \\le 60^{\\circ}$. ",
        "conclusions": "In our study we explored the regions of possible Trojan asteroids of Uranus for a large range of their inclinations up to $i= 60^\\circ$, which has never been done before. The method of determining these regions was the analysis of the results of straightforward numerical integrations of the full equations of motion in the model of the OSS. Besides the escape from the region we also checked the amplitudes of libration of these fictitious asteroids around the Lagrangian equilibrium points. To assure the survival of the 'stable' Trojans we did also integrations of three specially chosen Trojans on low inclined orbits for the age of the Solar system of $4.5 \\times 10^9$ years, where the accurate analysis of the frequencies involved was undertaken. Our results showed that close to the first instability window $6^\\circ \\leq i_{Trojan} \\leq 8^\\circ$ the strong variation of the frequency of a special secular resonance  indicates chaotic motion. Two other unstable windows with respect to the inclination are $14^\\circ \\leq i_{Trojan} \\leq 31^\\circ$ and around $i_{Trojan} = 37^\\circ$. From $i_{Trojan} \\geq 50^\\circ$ (for $L_4$) and  $i_{Trojan} \\geq 45^\\circ$ ($L_5$) no stable Trojan orbits exist for Uranus. These stable islands in inclination are surrounded by more or less small stable regions in semimajor axis.  According to the results of \\citet{Mor2005} Jupiter Trojans could be trapped during the phase of early migration of the planets in our Solar system. If that was also the case for Neptune\\footnote{up to now 6 Neptune Trojans were observed} and Uranus Trojans (and why not) we should be able to detect them. So the question is: why don't we observe them?"
    },
    "0911/0911.2150_arXiv.txt": {
        "abstract": "Markov Chain Monte Carlo (MCMC) methods have revolutionised Bayesian data analysis over the years by making the direct computation of posterior probability densities feasible on modern workstations. However, the calculation of the prior predictive, the Bayesian evidence, has proved to be notoriously difficult with standard techniques. In this work a method is presented that lets one calculate the Bayesian evidence using nothing but the results from standard MCMC algorithms, like Metropolis-Hastings. This new method is compared to other methods like MultiNest, and greatly outperforms the latter in several cases. One of the toy problems considered in this work is the analysis of mock pulsar timing data, as encountered in pulsar timing array projects. This method is expected to be useful as well in other problems in astrophysics, cosmology and particle physics. ",
        "introduction": "Bayesian inference has proved over the years to be a very powerful tool in many branches of science as it gives a very clear prescription of how to analyse datasets without loss of information, even for very complicated models. On the practical side, in performing Bayesian analysis two difficult problems often emerge:\\newline 1. Producing the posterior probability density functions (PDFs) for several interesting parameters requires marginalisation, i.e.~the integration of the full joint posterior PDF over most model parameters. In a majority of cases this must be done numerically, and it is common to employ a Markov Chain Monte-Carlo (MCMC) algorithm for performing the integration.\\newline 2. Model selection requires the calculation of the Bayes factor: the ratio of the prior predictive values of two competing models. This prior predictive, or Bayesian evidence as we will call it, is the integral of the full joint posterior PDF over all model parameters. The calculation of this value can be notoriously difficult using standard techniques.  The use of Markov Chain Monte Carlo (MCMC) algorithms, such as the Metropolis-Hastings algorithm, has become extremely popular in the calculation of the marginal posterior PDFs. MCMC algorithms allow one to sample from posterior distribution of complicated statistical models, greatly reducing the effort involved in evaluating high dimensional numerical integrals. The problem with standard MCMC algorithms lies in the fact that the normalisation constant of the marginal posterior PDFs is lost in the calculation, making it difficult to calculate the second of the above mentioned quantities.  Over the years, several methods capable of calculating the Bayesian evidence have been developed, such as using the harmonic mean \\citep{Newton1994} and parallel tempering \\citep{Earl2005}. The problems with these methods are the inaccuracy of the method for the harmonic mean, and the high computational cost for parallel tempering: this method can easily be 100 times more time-consuming than regular MCMC algorithms \\citep{Newman1999}. More recent efforts include the Nested sampling algorithm \\citep{Skilling2004}, a greatly improved version of which has been implemented in MultiNest \\citep[hereafter FHB09]{Feroz2009}.  These recent methods rely on a transformation of the integral to a 1-dimensional problem, and then use a clever trick for avoiding the calculation of the inverse of the cumulative posterior PDF. The MultiNest algorithm of FHB09 has been successfully tested on several (toy) problems, and seems to be much more efficient than traditional methods.  However, since the MultiNest algorithm by design samples from the prior - not from the posterior - we believe that the MultiNest algorithm suffers more from the curse of dimensionality than the traditional MCMC algorithms \\citep{Evans2006}.  When analysing complicated models with many parameters ($\\sim 100$ or more), the scaling with dimensionality is crucial.  In this paper we present a method to calculate the Bayesian evidence from the MCMC chains of regular MCMC algorithms. This method can be applied to chains that have already been run, provided that the posterior values, calculated with normalised prior and normalised likelihood, have been saved together with the values of the parameters at each point in the chain. For several cases, the accuracy of the new method is significantly greater than that of other methods with the same amount of sampled points, provided that the MCMC has been properly run. The error on the Bayesian evidence can be reliably calculated using a bootstrap procedure.  The outline of the paper is as follows. In section \\ref{sec:bayes} we briefly review the basic aspects of Bayesian inference for parameter estimation and model selection. Then we provide the reader with some necessary details on MCMC in section \\ref{sec:MCMC}, where we also outline the new algorithm to calculate the Bayesian evidence. In section \\ref{sec:comparison} we assess the strengths and weaknesses of the new method relative to those that use nested sampling or parallel tempering. In section \\ref{sec:applications} we test all competing algorithms on some toy problems like the analysis of pulsar timing data as encountered in pulsar timing array projects. ",
        "conclusions": "We develop and test a new algorithm that uses traditional MCMC methods to accurately evaluate numerical integrals that typically arise when evaluating the Bayesian evidence. The new method can be applied to MCMC chains that have already been run in the past so that no new samples have to be drawn from the integrand. We test the new algorithm on several toy-problems, and we compare the results to other algorithms: MultiNest and Parallel tempering. We conclude that there is no single algorithm that outperforms all others in all cases. High-dimensional, peaked problems are better tackled using a traditional MCMC method, whereas very complicated multimodal problems are probably best handled with MultiNest. When applicable, the new algorithm significantly outperforms other algorithms, provided that the MCMC has been properly executed. This new algorithm is therefore expected to be useful in astrophysics, cosmology and particle physics.  We have demonstrated that the new algorithm suffers under the use of correlated MCMC chains, produced by using the entire chain of a particular MCMC method like Metropolis-Hastings. If the new algorithm is used in combination with an uncorrelated chain, the accuracy of the numerical integral can reach the theoretical limit for stochastic methods: $\\sigma = I/\\sqrt{N}$, with $\\sigma$ the uncertainty, $I$ the value of the integral, and $N$ the amount of MCMC samples. Using correlated MCMC samples can significantly increase the integral uncertainty, and longer MCMC chains are needed for the integral to converge. Additional tests to assess convergence are required.  \\subsection{Comparison to other work} When this paper was already submitted to Monthly Notices, a preprint by \\citet{Weinberg2009} appeared on the arxiv which also attempts to construct an algorithm to calculate the Bayesian evidence using MCMC samples of the posterior. Weinberg first discusses the flaws of the harmonic mean estimator, and then proposes 2 algorithms to overcome these flaws. Weinberg then tests these interesting modifications to the harmonic mean estimator on simulated datasets. The algorithm proposed by Weinberg bears some resemblance to the algorithm in this paper in that one should only use a well-sampled subset of the parameter space when estimating the Bayesian evidence."
    },
    "0911/0911.2827_arXiv.txt": {
        "abstract": "We introduce an analytic model of the diffuse intergalactic medium in galaxy clusters based on a polytropic equation of state for the gas in hydrostatic equilibrium with the cluster gravitational potential. This model is directly applicable to the analysis of X-ray and Sunyaev-Zeldovich Effect observations from the cluster core to the virial radius, with 5 global parameters and 3 parameters describing the cluster core. We validate the model using \\chandra\\ X-ray observations of two polytropic clusters, MS~1137.5+6625 and CL~J1226.9+3332, and two cool core clusters, Abell~1835 and Abell~2204. We show that the model accurately describes the spatially resolved spectroscopic and imaging data, including the cluster core region where significant cooling of the plasma is observed. ",
        "introduction": "Galaxy cluster masses play an important role in addressing fundamental physical and cosmological problems, such as the measurement of the cluster gas mass fraction \\citep[][]{allen2008,ettori2009}, the evolution of the growth of structure \\citep[][]{mantz2008,mantz2009,vikhlinin2009},  and the gravitational sedimentation of ions \\citep{chuzhoy2003,peng2009,shtykovskiy2010}. A vital tool for the measurement of cluster masses is the diffuse hot intergalactic medium which can be detected primarily through its bright X-ray emission \\citep[][]{sarazin1988}, or through the Sunyaev-Zeldovich Effect \\citep[SZE,][]{sunyaev1972, carlstrom2002}. A variety of models are used to describe the distribution of the gas, from the simple isothermal $\\beta$ model \\citep[][]{cavaliere1976,birkinshaw1991} to more complex models that describe either the X-ray properties \\citep[gas density and temperature,][]{vikhlinin2006,cavaliere2009} or the SZE properties \\citep[gas pressure,][]{nagai2007,mroczkowski2009,arnaud2009}.  We investigate a model of galaxy clusters based on an analytic distribution for the cluster mass density inspired by the \\citet[][]{navarro1996} distribution, which we generalize following \\citet{suto1998}, \\citet{ascasibar2003} and \\citet{ascasibar2008} to include a variable asymptotic slope at large radii. This mass density is combined with a polytropic equation of state for the gas, to provide  self-consistent  density, temperature and pressure profiles for a plasma in hydrostatic equilibrium. The use of a polytropic equation of state for the cluster gas has also been recently proposed by \\citet{ascasibar2008} and \\citet{bode2009}, and tested observationally by \\citet{sanderson2009}.  In this paper we derive analytic radial profiles for the physical quantities (temperature, density and pressure) relevant to X-ray and SZE observations, and present an application of these models to high resolution \\chandra\\ observations of the galaxy clusters MS~1137.5+6625, CL~J1226.9+3332, Abell~1835 and Abell~2204. Applications of this new model include measurement of gas mass fraction from joint X-ray and Sunyaev-Zeldovich Effect observations \\citep{hasler2010} and the effect of \\he\\ sedimentation on X-ray mass estimates \\citep{bulbul2010}. This paper is organized as follows: in \\S \\ref{sec_models} we describe our model, in \\S \\ref{sec_dataAnalysis} we present the application of the model to \\chandra\\ X-ray observations of MS~1137.5+6625, CL~J1226.9+3332, Abell~2204 and Abell~1835, and in \\S  \\ref{sec:comparison} we perform a comparison between our mass measurements and the results of \\citet{mroczkowski2009}. In \\S \\ref{sec_conclusion} we present our conclusions. In the analysis of the \\chandra\\ data  we assume the cosmological parameters $h=0.73$, $\\Omega_M=0.27$ and $\\Omega_{\\Lambda}=0.73$. ",
        "conclusions": "\\label{sec_conclusion}  We introduce a new model to describe the physical properties of the hot intra-cluster medium and present an application of the model to \\chandra\\ X-ray observations of MS~1137.5+6625, CL~J1226.9+3332, Abell~1835 and Abell~2204. The model is based on a polytropic equation of state for the gas in hydrostatic equilibrium with the cluster gravitational potential. Using a function for the cluster total mass density that has the asymptotic slope as a free parameter, we obtain analytic expressions for the gas density, temperature and pressure. We also include a core taper function that accounts for the cooling of the gas in the cluster center.  This model has a number of features that make it suitable for the analysis of X-ray and SZE observations of galaxy clusters. The model is analytic, and has a limited number of parameters which describe the global properties of the cluster. For clusters which do not have a cool core, 5 parameters are sufficient to describe the distribution of density, temperature, pressure and total matter density. The gas density and temperature are linked by the polytropic equation of state, and the total matter density is related to the plasma properties by the hydrostatic equation. Therefore there is just one scale radius ($r_s$) that appears in the radial distribution of all thermodynamic quantities. The other parameters that describe the global physical properties of the cluster are the central density ($n_{e0}$) and temperature ($T_{0}$) of the gas, the polytropic index $n$, and the asymptotic slope of the total mass density ($\\beta+1$). For cool core clusters, three additional parameters allow an accurate description of the cooling of the gas in the core, and the accompanying increase in the density (\\S~\\ref{sec_models_coolcore}).  In addition to the analysis of spatially-resolved spectroscopic and imaging X-ray data (see \\S \\ref{sec_dataAnalysis}), the model is applicable to SZE observations, which require a model for the plasma pressure. A number of models suitable for SZE observations are available in the literature, for example \\citet{nagai2007} and \\citet{mroczkowski2009}. Our model has the advantage of the simultaneous applicability to both X-ray and SZE observations, and it is therefore suitable for a number of cosmological applications including the measurement of the Hubble constant \\citep[][]{bonamente2006}, the measurement of scaling relations between X-ray and SZE observables \\citep{bonamente2008}, the measurement of cluster masses independent of cosmology from joint X-ray and SZE data (Hasler et al. 2010) and the measurement of the effect of \\he sedimentation on X-ray measured masses \\citep{bulbul2010}."
    },
    "0911/0911.1403_arXiv.txt": {
        "abstract": "It is possible to provide a thermodynamic interpretation for  the field equations in any diffeomorphism invariant theory of gravity. This insight, in turn,  leads us to the possibility of deriving the gravitational field equations from another variational principle without using the metric as a dynamical variable. I review this approach and discuss its implications. ",
        "introduction": "An unsatisfactory feature of all theories of gravity is that the field equations do not have any direct physical interpretation. The lack of an elegant principle which can lead to the dynamics of gravity (``how matter tells spacetime to curve'') is quite striking when we compare this situation with the kinematics of gravity (``how spacetime makes the matter move''). The latter can be determined through the principle of equivalence by demanding that all freely falling observers, at all events in spacetime, must find that the equations of motion for matter  reduce to their special relativistic form. Our first aim will be to remedy this and provide a physical interpretation to field equations describing gravity in any diffeomorphism invariant theory. This, in turn, will lead us to the possibility of \\textit{deriving} the gravitational field equations from a thermodynamic variational principle without using the metric as a dynamical variable.  The alternative interpretation is based on the thermodynamics of horizons. Several recent investigations have shown that there is indeed a deep connection between gravitational dynamics and horizon thermodynamics (for a review, see \\cite{reviews}). For example, studies have shown that: \\begin{itemize} \\item Gravitational field equations in \\textit{a wide variety of theories}, when evaluated on a horizon, reduce to  a thermodynamic identity $TdS=dE+PdV$. This result, first pointed out in ref.\\cite{tpsstds}, has now been demonstrated \\cite{tds} in several cases like the  stationary axisymmetric horizons and evolving spherically symmetric horizons in Einstein gravity, static spherically symmetric horizons and dynamical apparent horizons in Lovelock gravity, and three dimensional BTZ black hole horizons, FRW cosmological models in various gravity theories and even \\cite{hwg} in the case Horava-Lifshitz Gravity. It is not possible to understand, in the conventional approach, why the field equations should encode information about horizon thermodynamics. \\item  Gravitational action functionals in a wide class of theories have a a surface term and a bulk term. In the conventional  approach, we \\textit{ignore} the surface term completely (or cancel it with a counter-term) and obtain the field equation from the bulk term in the action. Any solution to the field equation obtained by this procedure is logically independent of the nature of the surface term. But  when the \\textit{surface term} (which was ignored) is evaluated at the horizon that arises in any given solution, it  gives the entropy of the horizon! (Again, this result extends far beyond Einstein's theory to situations in which the entropy is not just proportional to horizon area.) This is possible only because there is a specific holographic relationship \\cite{holo,ayan} between the surface term and the bulk term which, however, is  an unexplained feature in the conventional approach to gravitational dynamics. Since the surface term has the thermodynamic interpretation as the entropy of horizons, and is related holographically to the bulk term, we are again led to an indirect connection between spacetime dynamics and horizon thermodynamics. \\end{itemize}  Based on these features --- \\textit{which have no explanation in the conventional approach} --- one can argue  that there  is a  conceptual reason why we need to relate horizon thermodynamics with gravitational dynamics (in a wide class of theories far more general than just Einstein gravity) and revise  our perspective towards spacetime. To set the stage for this future discussion,  we begin by recalling the implications of the existence of temperature for horizons.  In the study of normal macroscopic systems --- like, for example, a solid or a gas --- one can \\textit{deduce} the existence of microstructure just from the fact that the object can be heated.  This was  the insight of Boltzmann which led him to suggest that heat is essentially a form of motion of the microscopic constituents of matter. That is, the existence of temperature  is sufficient for us to infer the existence of microstructure even without any direct experimental evidence.  The non-zero temperature of a horizon  shows that we can actually heat up a spacetime, just as one can heat up a solid or a gas. An unorthodox way of doing this would be to take some amount of matter and arrange it to collapse and form a black hole. The Hawking radiation \\cite{Hawking:1975sw} emitted by the black hole can be used to heat up, say, a pan of water just as though the pan was kept inside a microwave oven. In fact the same result can be achieved by just accelerating through the inertial vacuum carrying the pan of water which will eventually be heated to a temperature proportional to the acceleration  \\cite{Davies:1975th}. These processes show that the temperatures of all horizons are as real as any other temperature  \\cite{Padmanabhan:2002ha}. Since they give rise to  a class of hot  spacetimes, it follows \\textit{\\`{a} la} Boltzmann that the spacetimes should possess microstructure.   In the case of a solid or gas, we know the nature of this microstructure from atomic and molecular physics. Hence, in principle, we can work out the thermodynamics of these systems from the underlying statistical mechanics. This is not possible in the case of spacetime because we have no clue about its microstructure. However, one of the remarkable features of thermodynamics --- in contrast to statistical mechanics --- is that the thermodynamic description is fairly insensitive to the details of the microstructure and can be developed as a fairly broad frame work. For example, a thermodynamic identity like $TdS = dE+PdV$ has a universal validity and the information about a \\textit{given} system is only encoded in the form of the entropy functional $S(E,V)$. In the case of normal materials, this entropy arises because of our coarse graining over microscopic degrees of freedom which are not tracked in the dynamical evolution. In the case of spacetime, the existence of horizons for a particular class of observers makes it mandatory that these observers  integrate out degrees of freedom hidden by the horizon.  To make this notion clearer, let us start from the principle of equivalence which allows one to construct local inertial frames (LIF) with coordinates $X^i$, around any event in an arbitrary curved spacetime. Given the LIF, we can now construct a local Rindler frame (LRF) by  boosting along one of the directions  with an acceleration $\\kappa$, thereby locally transforming the metric to a form given by: \\begin{equation} ds^2 = -dT^2 + dX^2 + d\\mathbf{x}_\\perp^2=-\\kappa^2 x^2 dt^2 + dx^2 + dL\\mathbf{x}_\\perp^2 = - 2 \\kappa l \\ dt^2 + \\frac{dl^2}{2\\kappa l}  + d\\mathbf{x}_\\perp^2 \\label{surfrind} \\end{equation} where $T=x \\sinh (\\kappa t);  X= x \\cosh (\\kappa t)$ and $l=(1/2)\\kappa x^2$. The observers at rest (with $x=$ constant) in the LRF will perceive the $X=T$ null surface as a horizon $\\mathcal{H}$ (fig 1). These local Rindler observers and the freely falling inertial observers will attribute different thermodynamical properties to matter in the spacetime. For example, they will attribute different temperatures and entropies to the vacuum state as well as excited states of matter fields. When some matter with energy $\\delta E$ moves close to the horizon --- say, within a few Planck lengths because it formally takes infinite Rindler time for matter to actually cross $\\mathcal{H}$ --- the local Rindler observer will consider it to have  transfered an entropy  $\\delta S = (2\\pi/\\kappa) \\delta E$ to the horizon degrees of freedom. We will  show (in Sec. 3) that, when the metric satisfies the field equations of any diffeomorphism invariant theory, this transfer of entropy can be given  \\cite{tp09papers} a  geometrical interpretation as the change in the  entropy of the horizon.   This result allows us to associate an entropy functional with the null surfaces which the local Rindler observers perceive as  horizons. We can now demand that the sum of the horizon entropy and the entropy of matter that flows across the horizons (both as perceived by the local Rindler observers), should be an extremum for all observers in the spacetime. This leads to a constraint on the geometry of spacetime which can be stated, in $D=4$, as \\begin{equation} (G_{ab} - 8\\pi T_{ab}) n^a n^b =0 \\label{Eenn} \\end{equation} for all null vectors $n^a$ in the spacetime. The general solution to \\eq{Eenn} is given by $G_{ab} = 8\\pi T_{ab} + \\rho_0 g_{ab}$ where $\\rho_0$ has to be a constant because of the conditions $\\nabla_a G^{ab} =0 = \\nabla_aT^{ab}$. Hence the thermodynamic principle leads uniquely to Einstein's equation with a cosmological constant in 4-dimensions. Notice, however, that \\eq{Eenn} has a new symmetry which the standard Einstein's theory does not posses; viz., it is invariant under the transformation $T_{ab} \\to T_{ab} + \\lambda g_{ab}$. This has important implications for the cosmological constant problem which we will discuss in Sec. 4.3. In $D>4$, the same entropy maximization leads to a more general class of theories called \\LL\\ models (see Sec. 4.2).  The constraint on the background geometry in \\eq{Eenn} arises from our demand that the thermodynamic extremum principle should hold for \\textit{all} local Rindler observers. This is identical  to the manner in which  freely falling observers are used to determine how gravitational field influences matter. Demanding the validity of special relativistic laws for the matter variables, as determined by all the freely falling observers, allows us to determine the influence of gravity on matter. In a similar manner, demanding the maximization of entropy of horizons (plus matter), as measured by all local Rindler observers, leads to the dynamical equations of gravity.  In this approach, both the entropy of horizons as well as the entropy of matter flowing across the horizon will be observer dependent thereby introducing a new level of observer dependence in thermodynamics. In particular, observers in different states of motion will have different regions of spacetime accessible to them; for example, an observer falling into a black hole will not perceive its horizon in the same manner as an observer who is orbiting around it. Therefore we are forced to accept that the notion of entropy is an observer dependent concept. (At a conceptual level this is no different from the fact different freely falling observers will measure physical quantities differently; but in this case, standard rules of special relativity allow us to translate the results between the observers. We do not yet have  a similar set of rules for quantum field theory in noninertial frames.) All these features  suggest a deep relationship between quantum theory, thermodynamics and gravity which forms the main theme of this article.  The rest of the article is organized as follows: In the next section, I briefly review some of the background results needed for our discussion. Section 3 describes an interpretation of gravitational field equations in a general, diffeomorphism invariant, theory of gravity. Using these results, it is possible to introduce an entropy maximization principle from which one can obtain the same field equations without using the metric as a dynamical variable. This is done in Sec. 4 and the last section summarizes the results. ",
        "conclusions": "It is useful to distinguish clearly (i) the mathematical results which can be rigorously proved from (ii) interpretational ideas which might evolve when our understanding of these issues deepen.  (i) From a purely algebraic point of view, without bringing in any physical interpretation or motivation, we can prove the following mathematical results: \\begin{itemize} \\item Consider a functional of null vector fields $n^a(x)$ in an arbitrary spacetime given by \\eq{ent-func-2} [or, more generally, by \\eq{generalS}]. Demanding that this functional is an extremum for all null vectors $n^a$ leads to the field equations for the background geometry given by $(2E_{ab}-T_{ab})n^an^b=0$ where $E_{ab}$ is given by \\eq{genEab1} [or, more generally, by \\eq{genEab}]. Thus field equations in a wide class of theories of gravity can be obtained from an extremum principle without varying the metric as a dynamical variable. \\item These field equations are invariant under the transformation $T_{ab} \\to T_{ab} + \\rho_0 g_{ab}$, which relates to the freedom of introducing a  \\cc\\ as an integration constant in the theory. Further, this symmetry forbids the inclusion of a cosmological constant  term in the variational principle by hand as a low energy parameter. \\textit{That is, we have found a symmetry which makes the bulk \\cc\\ decouple from the gravity.} When linearized around flat spacetime, the graviton inherits this symmetry and does not couple to the \\cc . \\item On-shell, the functional in \\eq{ent-func-2} [or, more generally, by \\eq{generalS}] contributes only on the boundary of the region. When the boundary is a horizon, this terms gives precisely the Wald entropy of the theory. \\end{itemize} It is remarkable that one can derive not only Einstein's theory uniquely in $D=4$ but even \\LL\\ theory in $D>4$ from an extremum principle involving the null normals \\textit{without varying $g_{ab}$ in an action functional!}.  (ii) As regards interpretation of these results,  we see from \\eq{details1} that in the case of Einstein's theory, we have a Lagrangian $n^a(\\nabla_{[a}\\nabla_{b]})n^b$ for a vector field $n^a$ (except for  a surface term) that becomes vacuous in flat spacetime in which covariant derivatives become partial derivatives. There is  no dynamics in $n^a$ (in the usual sense) but they do play a crucial role. This raises the question as to the physical meaning of these null vectors and the interpretation of our approach. As described above, this interpretation is essentially thermodynamical and the physical picture is made of the following qualitative ingredients: \\begin{itemize} \\item Assume that the spacetime is endowed with certain microscopic degrees of freedom capable of exhibiting thermal phenomena. This is just the Boltzmann paradigm: \\textit{If one can heat it, it must have microstructure!}; and one can heat up a spacetime. \\item Whenever a class of observers perceive a horizon, they are ``heating up the spacetime'' and the  degrees of freedom close to a horizon  participate in a very \\textit{observer dependent} thermodynamics. Matter which flows close to the horizon (say, within a few Planck lengths of the horizon) transfers energy to these microscopic, near-horizon, degrees of freedom \\textit{as far as the observer who sees the horizon is concerned}. Just as  entropy of a normal system at temperature $T$ will change by $\\delta E/T$ when we transfer to it an energy $\\delta E$, here also an entropy change will occur. (A freely falling observer in the same neighbourhood, of course, will deny all these!) \\item We proved that when the field equations of gravity hold, one can interpret this entropy change in  a purely geometrical manner involving the Noether current. From this point of view, the  normals $n^a$ to local patches of null surfaces are related to the (unknown) degrees of freedom that can participate in the thermal phenomena involving the horizon. \\item Just as demanding the validity of special relativistic laws with respect to all freely falling observers leads to the kinematics of gravity, demanding the local entropy balance in terms of the thermodynamic variables as perceived by local Rindler observers leads to the field equations of gravity in the form $(2E_{ab}-T_{ab})n^an^b=0$.  \\end{itemize}   As stressed in earlier sections, this involves a new layer of observer dependent thermodynamics. At a conceptual level, this may be welcome when we note that every key progress in physics involved realizing that something we thought as absolute is  not absolute. With special relativity it was the flow of time and with general relativity it was the concept of global inertial frames and when we brought in quantum fields in curved spacetime it was the notion of particles and temperature. We now know that the temperature attributed to even vacuum state depends on the observer. It seems necessary to  integrate the entire thermodynamic machinery (involving what we usually consider to be the `real' temperature) with this notion of LRFs having their own (observer dependent) temperature. There is scope for further work in this direction.    \\begin{theacknowledgments} I thank Jean-Michel Alimi for organizing a stimulating conference and providing excellent hospitality. \\end{theacknowledgments}"
    },
    "0911/0911.1310.txt": {
        "abstract": "A new wavefront sensing approach, derived from the successful curvature wavefront sensing concept but using a non-linear phase retrieval wavefront reconstruction scheme, is described. The non-linear curvature wavefront sensor (nlCWFS) approaches the theoretical sensitivity limit imposed by fundamental physics by taking full advantage of wavefront spatial coherence in the pupil plane. Interference speckles formed by natural starlight encode wavefront aberrations with the sensitivity set by the telescope's diffraction limit $\\lambda$/D rather than the seeing limit of more conventional linear WFSs. Closed-loop adaptive optics simulations show that with a nlCWFS, a 100 nm RMS wavefront error can be reached on a 8-m telescope on a $m_V  = 13$ natural guide star. The nlCWFS technique is best suited for high precision adaptive optics on bright natural guide stars. It is therefore an attractive technique to consider for direct imaging of exoplanets and disks around nearby stars, where achieved performance is set by wavefront control accuracy, and exquisite control of low order aberrations is essential for high contrast coronagraphic imaging. Performance gains derived from simulations are shown, and approaches for high speed reconstruction algorithms are briefly discussed. ",
        "introduction": "Adaptive Optics (AO) systems are now an essential part of ground-based telescopes, and routinely deliver diffraction limited images at near-IR and mid-IR wavelengths on 8 to 10 m telescopes. For several key science applications, the required wavefront quality is however still well beyond what current systems deliver. For example, obtaining high quality diffraction-limited images at visible wavelengths requires residual wavefront errors to be below 100 nm RMS. Direct imaging of exoplanets or circumstellar debris disks with a high contrast coronagraphic camera is even more demanding, and requires exquisite control of low order aberrations and mid-spatial frequencies.  In the last decade, several high performance coronagraphs concepts have been developed in the laboratory, and full AO+coronagraph laboratory systems have demonstrated that both deformable mirror and coronagraph technologies are now sufficiently mature to allow direct imaging of planets down to Earth mass provided that wavefront aberrations are small: for example, the vacuum High Contrast Imaging Imaging Testbed (HCIT) a NASA JPL has achieved better than $10^{-9}$ raw PSF contrast at 4 $\\lambda$/D separation \\citep{trau07}. This high contrast value is orders of magnitude better than, at a few $\\lambda/D$, the $\\approx10^{-3}$ best raw PSF contrast currently achieved in the near-IR on 8 m to 10 m telescopes and the expected $\\approx 10^{-4}$ raw PSF contrast for \u00d2Extreme-AO\u00d3 systems currently under assembly such as SPHERE \\citep{beuz08} on the Very Large Telescope (VLT) and Gemini Planet Imager \\citep{maci08}. The large gap between laboratory demonstrations and on-sky performance is entirely due to wavefront control accuracy. A key difference between laboratory systems and on-sky conditions is the timescale of wavefront variations, which requires fast wavefront sampling and therefore fundamentally limits the number of photons available for sensing on ground-based systems.  The wavefront information which can be extracted from a given number of photons is known from first principles, and accurately defines the \"ideal\" performance which can be reached by AO systems on ground-based telescopes \\citep{guyo05}. Direct comparison between this ideal limit and the sensitivity offered by current wavefront sensing schemes reveals a huge performance gap. Commonly used wavefront sensors such as Shack-Hartmann and Curvature are very robust and flexible, but are poorly suited to high quality wavefront measurement: for low-order aberrations on 8 to 10 m telescopes, they require $\\approx$ 100 to 1000 times more photons than more optimized wavefront sensing techniques. While they offer nearly ideal sensitivity at the spatial frequency defined by the subaperture spacing, they suffer from poor sensitivity at low spatial frequencies, which are most critical for direct imaging of exoplanets and disks.  This paper describes a high sensitivity wavefront sensing approach which is optically derived from the curvature wavefront sensing technique, and which shares the same non-linear reconstruction algorithm as used for phase diversity. In phase diversity, images near the focal plane are used for wavefront reconstruction with a non-linear algorithm, while curvature wavefront sensing uses images near the pupil plane and a linear wavefront reconstruction algorithm. Phase diversity is theoretically more sensitive than linear curvature wavefront sensing \\citep{fien98}, but, for ground-based adaptive optics applications, suffers in practice from smearing of focal plane speckles due to chromatic effects and time averaging of the fast-changing PSF. I propose in this paper a solution to improve curvature WFS sensitivity by using non-linear signals in the near-pupil images, and offer a solution to optically mitigate the chromatic smearing of the speckles which would otherwise reduce sensitivity. Unlike phase diversity, the proposed scheme is robust against time averaging of wavefront aberrations, which gradually brings the proposed technique closer to a less sensitive but stable linear curvature wavefront sensor.  The non-linear curvature wavefront sensor (nlCWFS) principle is described in \\S\\ref{sec:nlCWFS}. The wavefront reconstruction algorithm for the nlCWFS is described in \\S\\ref{sec:algo}. The WFS performance is evaluated through numerical simulations in \\S\\ref{sec:perf} and discussed in \\S\\ref{sec:disc}. ",
        "conclusions": "Simulations and analysis presented in this paper show that the non linear Curvature WFS technique is potentially able to work close to the fundamental theoretical sensitivity limit imposed by photon noise even if the PSF is not diffraction limited at the WFS wavelength. It is also a highly flexible WFS choice, as it can efficiently sense wavefront on sparse pupils, and performs an explicit wavefront reconstruction, therefore allowing open-loop operation and compensation of non-common path errors.  In the preliminary study presented in this paper, encouraging results were obtained on intermediate brightness guide stars with closed loop simulations using a relatively simple control algorithm. Further performance gains are to be expected if both modal control (optimal weighting of wavefront modes according to their measurement SNR and variance in the atmospheric turbulence) and predictive control are also included. A key challenge to fully take advantage of this technique is to achieve sufficiently high AO loop frequency, and algorithms with faster convergence than the iterative loop used in this paper should be developed.  The nlCWFS appears especially well suited for Extreme-AO systems aimed at delivering nearly flat wavefronts to a high contrast imaging camera (coronagraph) for direct imaging of exoplanets and disks. In these systems a nlCWFS can deliver for low order aberrations on a 8m telescope the same wavefront measurement accuracy as a SHWFS with $\\approx 1000$ times fewer photons. The nlCWFS can take full advantage of the telescope's diffraction limit for wavefront sensing, with a sensitivity scaling as $D^4$ (combination of a $D^2$ gain due smaller diffraction limit and a $D^2$ gain due to the telescope's collecting power) while more conventional WFSs scale only as $D^2$ (telescope's collecting power). This difference is especially large for the next generation of Extremely Large Telescopes (ELTs), where the use of high sensitivity WFS techniques such as the nlCWFS may allow direct imaging of exoplanets in reflected light."
    },
    "0911/0911.1129_arXiv.txt": {
        "abstract": "Member galaxies within galaxy clusters nowadays can be routinely identified in cosmological, hydrodynamical simulations using methods based on identifying self bound, locally over dense substructures. However, distinguishing the central galaxy from the stellar diffuse component within clusters is notoriously difficult, and in the center it is not even clear if two distinct stellar populations exist. Here, after subtracting all member galaxies, we use the velocity distribution of the remaining stars and detect two dynamically, well-distinct stellar components within simulated galaxy clusters. These differences in the dynamics can be used to apply an un-binding procedure which leads to a spatial separation of the two components into a {\\it cD} and a diffuse stellar component ({\\it DSC}). Applying our new algorithm to a cosmological, hydrodynamical simulation we find that -- in line with previous studies -- these two components have clearly distinguished spatial and velocity distributions as well as different star formation histories. We show that the {\\it DSC} fraction -- which can broadly be associated with the observed intra cluster light -- does not depend on the virial mass of the galaxy cluster and is much more sensitive to the formation history of the cluster. We conclude that the separation of the {\\it cD} and the {\\it DSC} in simulations, based on our dynamical criteria, is more physically motivated than current methods which depend on implicit assumptions on a length scale associated with the {\\it cD} galaxy and therefore represent a step forward in understanding the different stellar components within galaxy clusters. Our results also show the importance of analyzing the dynamics of the DSC to characterize its properties and understand its origin. ",
        "introduction": "\\label{sec:intro}  The existence of a diffuse stellar component in groups and clusters of galaxies is now well established. A diffuse intracluster light ({\\it ICL}) has been observed both in local universe \\citep{Feldmeier04,Mihos05,Arnaboldi04,Gerhard05,Doherty09} and at intermediate redshift \\citep{Gonzalez00,Feldmeier04b,Zibetti05,Krick07}. Such {\\it ICL} component is centrally concentrated, and typically amounts between $\\approx 10$ \\citep{Zibetti05} and $\\approx 35$ \\% \\citep{Gonzalez07} of the total stellar mass in clusters. There are some indications pointing towards a dependence of the {\\it ICL} fraction from the mass of the clusters, going from $\\approx 2$ \\% at the scale of loose groups \\citep{Castro-Rodr03} to $\\approx 5-10$ \\% of Virgo cluster \\citep{Magda03,Feldmeier04b,Mihos05,2009arXiv0908.3848C} up to $10-20$ \\% or higher in the most massive clusters \\citep{Gonzalez00,Feldmeier02,GalYam03,Feldmeier04b,Krick06}, but other studies found instead no significant dependence on the cluster richness \\citep{Zibetti05}.  {\\it ICL} is usually detected using surface photometry in various bands, sometimes stacking together observations of many clusters \\citep[see e.g.][]{Zibetti05}. In this case, light coming from galaxies is masked away, often using a fixed limit in surface brightness \\citep[e.g., as in][]{Feldmeier04}. Single {\\it ICL} stars can also be detected \\citep{Durrell02}, expecially by observing Intra-Cluster Planetary Nebulae (ICPN) \\citep{Magda96,Feldmeier04b}, which can be observed up to distances as large as 100 Mpc \\citep{Gerhard05}. On the one hand, the observational effort needed to detect single stars makes it difficult to use them to assess general properties of the {\\it ICL} distribution, like the {\\it ICL} fraction in clusters. On the other hand, the observation of single stars make it possible to gather information on the kinematic of such a component.  In particular, extended halos of bright ellipticals often overlap spatially with stars which are free-floating in the cluster gravitational potential. Efforts have been made to disentangle these two components \\citep[see e.g.][]{Doherty09}, but they are often referred together as the {\\it ICL}. In the following, we will use {\\it ICL} to indicate the observed diffuse stellar population, whose dynamical properties are difficult to obtain, and {\\it DSC} for the corresponding stellar population produced in numerical simulation, whose dynamics is known.  From a theoretical point of view, a diffuse stellar component ({\\it DSC}) has been studied in a cosmological Dark-Matter-only numerical simulation \\citep{Napolitano03}.  Later, for the first time \\cite{2004ApJ...607L..83M} (M04 hereafter) studied the properties of the {\\it DSC} in a cosmological simulation which included hydrodynamics and various astrophysical processes such as radiative cooling of gas and (self-consistent) star formation. \\cite{Willman04} and \\cite{SommerLarsen05} analyzed a {\\it DSC} in their resimulation of single galaxy clusters.  M04 analyzed a cosmological simulation of a box of 192 $h^{-1}$ Mpc, containing $\\approx 100$ galaxy cluster. Their main results were that: (i) a {\\it DSC} is clearly seen when a Sersic fit is performed on the bound, unbound and total two-dimensional radial density profiles of the stellar population in clusters; (ii) the {\\it DSC} component is more centrally concentrated than the stars bound in galaxies; (iii) on average, stars in the {\\it DSC} are older than those in galaxies; (iv) the {\\it DSC} fraction increase with the cluster mass.  \\cite{Willman04} used one cluster re-simulation, having a mass of $1.2 \\times 10^{15} $ M$_\\odot$, and confirmed results (i) of M04. They found a {\\it DSC} fraction compatible with M04 clusters in the same mass range, if slightly lower. \\cite{SommerLarsen05} used one Virgo-like and one Coma-like cluster re-simulation and confirmed results (i), (ii) and (iii) of M04. They found similar {\\it DSC} fraction, with their Coma-like cluster having an higher fraction of {\\it DSC} stars. Their average stellar age in {\\it DSC} was older than that found by M04.  Several mechanisms have been investigated to explain the formation and evolution of a {\\it DSC}: stripping and disruption of galaxies as they pass through the central regions of relaxed clusters \\citep{Byrd90,Gnedin03}; stripping of stars from galaxies during the initial formation of clusters \\citep{Merritt84}; creation of stellar haloes in galaxy groups, that later fall into massive clusters and then become unbound \\citep{Mihos04,Rudick06}; and stripping of stars during high-speed galaxy encounters in the cluster environment \\citep{Moore96}. \\citet{2007MNRAS.377....2M} (M07 hereafter) showed that the {\\it DSC} is produced during mergers in the formation history of the BCGs and of other massive galaxies, and that it grows steadily since redshift $z=1$, with no preferred epoch of formation. \\citet{Monaco06}, using a semi-analytical model of galaxy formation in clusters, showed that, in the hierarchical model of structure formation, the BCG merging activity alone between $z=1$ and $z=0$ is so intense that it is not possible to simultaneously fit the bright end of the galaxy luminosity function at both redshifts without assuming that a significant fraction of stars goes into the {\\it DSC} during such mergers. We point out that observational indications have been found that such process is in fact in action in the Coma cluster \\citep{2007A&A...468..815G}. Numerically, \\cite{Stanghellini06} studied isolated, star-only dry merger of elliptical galaxies and showed how in that case, up to $21$ the initial total stellar mass can become unbind in the process.  One major problem with the analysis of {\\it DSC} in numerical simulations is its very identification. One possibility, used e.g. by \\cite{Rudick06}, is to build simulated surface brightness maps and detect an {\\it ICL} component with a procedure similar to that used with observational data.  This has the advantage of being somewhat easier to compare with observation, at the cost of losing the intrinsic advantage of simulations, that is, the knowledge of the dynamics of each star particle. The other possibility is to attempt a dynamical distinction between stars bound to galaxies and stars which are free-floating in the group or cluster potential. This is what M04, \\cite{Willman04} and \\cite{SommerLarsen05} did; the disadvantage is that such a distinction may or may not coincide with the observational definition of {\\it ICL}. However, for doing this, it is necessary to identify structures in the simulation (groups or clusters) and substructures within them (the galaxies). Up to date, no unambiguous way to perform such identification has been formulated.  The identification of {\\it structures} is a relatively easy task. The most commonly used algorithms are based upon two methods: the Friends-of-Friends (FoF) approach \\citep{Davis85,Frenk88} and the Spherical Overdensity one (SO) \\citep{LaceyCole94}. In the former, all particles closer than a given distance $l$ (usually expressed in units of the mean inter-particle distance, typically $l\\approx0.2$ is chosen) are grouped together; then, all particles closer than $l$ to such particles are included in the group, and so on. In the latter, one identifies density peaks in the particle distribution, and puts spheres around them, choosing their radius so that they enclose a given overdensity.  The difficulty arises when {\\it sub}structures inside DM halos must be found. The density contrast of substructures against their parent halo density is lower than the density contrast of structures against the background, and this makes their identification more problematic. Even worse, the gravitational potential of a structure has a clear zero-point, given by the background density of the universe. So, it is well defined which particles are bound to the structure and what are unbound. Inside a structure, instead, such a definition is not clear: the background density of the structure varies inside them with the distance from the center, the dynamical state of the cluster, and the typical extend of the substructure itself. Therefore, telling that a particle is gravitationally bound to a given substructures involves a certain degree of arbitrariness, and an operational definition must be used.  A number of different {\\it sub}halo finders now exist in the literature. Some are based on pure geometrical measures, others construct a candidate structure and apply an unbinding criteria, e.g. comparing the kinetic energy of a particle to its potential energy. A non-exhaustive list includes: hierarchical {\\small FOF} \\citep{Klypin99},  Bound Density Maxima \\citep{Klypin99}, {\\small DENMAX} \\citep{Bert91,Gelb94}, {\\small SKID} \\citep{Weinberg97,2001PhDT........21S}, local density percolation \\citep{1995MNRAS.273..295V}, {\\small HOP} \\citep{Eisenstein98}, {\\small SUBFIND} \\citep{2001MNRAS.328..726S}, MHF \\citep{Gill04}, {\\small ADAPTAHOP} \\citep{Aubert04}, {\\small VOBOZ} \\citep{Ney05}, {\\small PSB} \\citep{Kim06}, {\\small 6D FOF} \\cite{Diemand06}, {\\small HSF} \\citep{Macie09}. {\\small AHF} \\citep{2009ApJS..182..608K}. Some of these algorithms are based on the FoF scheme, and vary the $l$ parameter inside structures or apply the scheme to the full 6D phase space (hierarchical {\\small FOF}, local density percolation, {\\small 6D FOF}); some use the SO scheme, and apply an unbinding criterion ({\\small BDM}, {\\small MHF}, {\\small AHF}); some use FoF to link particles to density maxima, and apply an unbinding ({\\small DENMAX}, {\\small SKID}); some introduce geometrical constraints, together with various unbinding procedure and phase space analysis ({\\small HOP}, {\\small SUBFIND}, {\\small ADAPTAHOP}, {\\small VOBOS}, {\\small PSB}, {\\small HFS}). Comparisons between the performance of the different algorithms in detecting substructures usually show a reasonable agreement between the different methods as can be found for example in \\citet{2009ApJS..182..608K} or \\citet{Macie09}. Already \\citet{2001MNRAS.328..726S} compared {\\small SUBFIND} and {\\small FOF}, showing that the former performs better in those cases where nearby two halos are connected by a filaments. {\\small FOF} usually joins together the two objects, while {\\small SUBFIND} can disentangle them. Such a behaviour is however common to most of the subhalo finder listed above.  As far as the {\\it DSC} detection is concerned, a further question arises. The {\\it DSC} is composed by stars {\\it bound to the gravitational potential of the cluster and not to any galaxy}, and it is centrally concentrated. Many of its general properties, therefore, are primarily determined by the stars at the center of the cluster. It is therefore important to correctly distinguish {\\it DSC} star particles from {\\it cD} star particles. But such a distinction is often somewhat arbitrary, as in {\\small SKID}, where the unbinding procedure is based upon the value of a parameter that must be given externally and whose effect must be tested {\\it a-posteriori} (see M07). Other substructure finders, like {\\small SUBFIND}, identifies all substructures but the main one and link all remaining star, gas and DM particles to the main subhalo, thus mixing {\\it cD} and {\\it DSC} stars. It is not even clear if a distinction between {\\it cD} and its extended stellar halo is physically motivated or if it is all arbitrary.  In this work, we show that such two components can be distinguished in the velocity space. We use {\\small SUBFIND} to identify substructures and disentangle galaxies from the {\\it cD} + {\\it DSC} stars. We show that the velocity distribution of the latters cannot be fit with a single Maxwellian distribution, while it is well fitted by a double Maxwellian (the sum of two Maxwellian distribution with different velocity dispersions).  We thus modify {\\small SUBFIND} to perform an unbinding procedure on the {\\it cD} + {\\it DSC} stars. We use the gravitational potential given by the matter contained in a sphere, centered on the halo center, and we find a radius which divides the star particles in two populations whose velocity distribution is separately fit by a single Maxwellian. We vary the radius until the two resulting velocity dispersions correspond to the two velocity dispersion given by the double Maxwellian fit of the velocity distribution of the whole {\\it cD} + {\\it DSC} star population. We show that this procedure is able to disentangle the central galaxy from its extended stellar halo. We thus characterize the {\\it DSC} through its phase-space distribution.  We then use our modified {\\small SUBFIND} algorithm to repeat the M04 analysis of the {\\it DSC} applied to a constrained cosmological simulation of the local universe. We also use two re-simulations of galaxy clusters, to test the behavior of our modified {\\small SUBFIND} when we increase resolution.  \\citet{2010arXiv1001.3018P} used different ways to define the {\\it DSC} in their simulations, including the one we present here. They found a broad qualitative agreement among the various methods. Here, we will focus on the comparison of our modified {\\small SUBFIND} algorithm with {\\small SKID} on our constrained cosmological simulation of the local universe, and show that they are in broad agreement, but dynamically separating the {\\it cD} and the {\\it DSC}, a dependance of the {\\it DSC} fraction on cluster mass is not found.  The plan of the paper is as follows. Section \\ref{sec:simgen} describes our simulations. Section \\ref {sec:dynamics} shows the existence of two dynamically distinct stellar components, even at the center of galaxy clusters. Section \\ref{sec:iclunbinding} presents our scheme for disentangling such components. In Section \\ref{sec:cDDSC} we apply our scheme to our local universe simulation, also comparing the results we obtain on the {\\it DSC} properties with those obtained using {\\small SKID}. In Section \\ref{sec:conc} we give our conclusions. In the Appendix we show a detailed comparison of the performance of {\\small SKID} and {\\small SUBFIND} for two galaxy clusters. ",
        "conclusions": "\\label{sec:conc}  We showed that the stellar populations of the main sub-halos of massive galaxy clusters in cosmological, hydrodynamical simulations are composed of two, dynamically clearly distinct components. Their velocity distributions can be fitted by a double Maxwellian distribution. Including this information into an un-binding procedure, we where able to spatially separate the two components in a central galaxy ({\\it cD}) and a diffuse stellar component ({\\it DSC}). The latter can be associated with the observationally very interesting intra cluster light ({\\it ICL}) component. However, we want to stress that there are significant differences in how the {\\it ICL} is measured in observations and how the {\\it DSC} is inferred in simulations. Therefore a detailed comparison of the amount of this stellar component of galaxy clusters in simulations and observations can only be done by mimicking observational strategies, e.g. inferring the component from synthetic surface brightness maps.  On the other hand our un-binding algorithm is ideally suited to study the physical processes leading to liberate the {\\it DSC} as it allows to distinguish the {\\it DSC} and the {\\it cD} with a minimum set of assumptions, namely \\begin{itemize} \\item The stars of the {\\it cD} component are those gravitationally bound to the inner most part of the halo, i.e. to the galaxy itself. \\item The stars of the {\\it cD} and the {\\it DSC} reflect two dynamical distinct populations. \\end{itemize} These two assumptions naturally reflect our expectations that the {\\it cD} galaxy is a relaxed object sitting in the center of the cluster potential, while {\\it DSC} is mainly formed by violent mergers of satellite galaxies with the {\\it cD}, is in equilibrium with the overall gravitational potential of the cluster, and still holds memory of the dynamics of the satellite galaxies.  We show that our un-binding scheme is able to disentangle two stellar populations, {\\it cD} and {\\it DSC}, whose velocity distributions separately matches the two components of the double Maxwellian velocity distribution which characterizes the stars belonging to the main sub-halos of clusters.  Applying our algorithm to a cosmological, hydrodynamical simulation of the local universe including cooling and star formation, we found that \\begin{itemize} \\item The velocity dispersion associated to the {\\it DSC} component is comparable to the velocity dispersion of the member galaxies, whereas the velocity dispersion associated to the {\\it cD} is significant smaller (by a factor of $\\approx 3$). his implies that the {\\it DSC} dynamics is determined by the general gravitational potential of the cluster, while the {\\it BCG} one feels the local galactic potential.  \\item In line with previous findings, the spatial distribution of the star particles belonging to the {\\it cD} are much more concentrated than the one of the {\\it DSC}, which starts to dominate the stellar density at radii larger than 0.03-0.08 $R_{vir}$. \\item The formation time of the star particles associated with the {\\it cD} is significant smaller than the one of the {\\it DSC}, indicating that -- in line with previous findings -- the {\\it DSC} is older (on average $\\approx 1.5$Gyears). \\item The fraction of the stars within the {\\it DSC} is only weekly dependent on the mass of the galaxy cluster, however, we find a larger fraction in clusters with a early formation time. This is in line with previous findings that the {\\it DSC} originates from the late stages of merger events of galaxies with the {\\it cD}. Therefore clusters which have had enough time to liberate stars since the last major infall show a larger {\\it DSC}. Interestingly, the mass ratio between the second brightest cluster galaxy and the {\\it cD} galaxy can be used as a proxy for the time till the last major merger. Using such a ratio, we obtain an even stronger trend: clusters with a small ratio show a large {\\it DSC} fraction, whereas clusters with a large ratio did not have yet enough time to liberate the stars contribution to the {\\it DSC}. We expect that such a trend could also be detected in {\\it ICL} observations. \\end{itemize}   We conclude that the separation of the {\\it cD} and the {\\it DSC} in simulations, based on our dynamical criteria, is more opportune than other current methods, as it depends on less numerous and physically more motivated assumptions, and leads to a stable algorithm helping to understand the different stellar components within galaxy clusters."
    },
    "0911/0911.3540_arXiv.txt": {
        "abstract": "{Turbulent deflagrations of Chandrasekhar mass White Dwarfs are commonly used to model Type Ia Supernova explosions. In this context, rapid rotation of the progenitor star is plausible but has so far been neglected.} {The aim of this work is to explore the influence of rapid rotation on the deflagration scenario.} {We use three dimensional hydrodynamical simulations to model turbulent deflagrations ignited within a variety of rapidly rotating CO WDs obeying rotation laws suggested by accretion studies.} {We find that rotation has a significant impact on the explosion. The flame develops a strong anisotropy with a preferred direction towards the stellar poles, leaving great amounts of unburnt matter along the equatorial plane.} {The large amount of unburnt matter is contrary to observed spectral features of SNe~Ia. Thus, rapid rotation of the progenitor star and the deflagration scenario are incompatible in order to explain SNe~Ia.}   \\keywords {Stars: supernovae: general -- Hydrodynamics -- Methods: numerical} ",
        "introduction": "\\label{intro}  Driven by the outstanding potential of Type Ia Supernovae (SNe Ia) to reconstruct the expansion history of the late universe, a number of observational campaigns have collected a wealth of empirical data about the lightcurves, spectra, and host populations of these powerful events \\citep{essence,hstgold}. This progress is in contrast to the sobering fact that despite many decades of theoretical and computational efforts, we are still lacking a solid understanding of the explosion mechanism, or the combination of different mechanisms, that can explain the class of SNe Ia in its entirety. While the majority of SN Ia explosions is believed to involve the rapid thermonuclear combustion of Chandrasekhar mass CO White Dwarfs (WDs) \\citep{HillebrandtNiemeyer00,2008arXiv0804.2147R}, it is still debated exactly how and where the star ignites \\citep{2004ApJ...607..921W,2006A&A...446..627S,2006A&A...448....1R}, and whether the combustion front propagates subsonically as a deflagration or supersonically as a detonation \\citep{GCDmodel,snob,RoepkeNiemeyer07}. However, it is clear that if the front is subsonic, its effective propagation rate is governed only by the turbulence produced by large-scale instabilities \\citep{1997ApJ...475..740N}.  From a purely theoretical point of view, the turbulent deflagration model is favoured compared with its competitors that invoke a spontaneous transition to a detonation \\citep{1999ApJ...523L..57N}, but large 3D simulations indicate it may be incapable to explain the full class of SNe Ia \\citep{snob}. Nevertheless, because of its conceptual simplicity and lack of free parameters, we will use it as a testbed for studying the influence of rapid rotation on the outcome of the explosion. Since the duration of thermonuclear burning is determined by the dynamical time scale of the WD, the star must rotate nearly critically in order for rotation to have any significant impact. This appears plausible at first sight, since most progenitor scenarios for SNe Ia involve the accretion of almost a solar mass of material from a binary companion before the thermonuclear runaway, allowing the WD to pick up a substantial amount of angular momentum. Indeed, some evolutionary calculations predict critical rotation of the WD at the time of the explosion \\citep{2005A&A...435..967Y}. On the other hand, the transport and dissipation of angular momentum inside the star are poorly understood and direct observations are absent. Hence, we treat the amplitude and shape of the rotation law as essentially undetermined and attempt to investigate their influence in a parameterised manner. As we will show, our general conclusion is robust with respect to the details of the rotation law: nearly critical rotation of the progenitor WD together with a pure deflagration model for the combustion front are inconsistent with the observations of SNe Ia.  The first multidimensional simulations of rapidly rotating exploding WDs were carried out by \\citet{1992A&A...254..177S}. However, in contrast to our work they concentrated on pure detonations instead of deflagrations. We will revisit the prompt detonation scenario in a follow-up paper.  This paper is structured as follows. Section 2 summarises the theoretical background on rotation of SNe~Ia progenitor stars. The motivation of the employed rotation laws is given. Different methods for the numerical initiation of the burning process are presented in section 3. Section 4 describes the computational grid, the flame modelling, treatment of the gravitational potential, the employed WD models, numerical stability, and the parameter study. The results of the numerical study are listed in section 5. The impact of rotation on the explosion is explained by considering buoyancy effects. Also, the influence of model parameters and shear motion caused by rotation is presented. An interpretation with respect to spectral features finalises the section. Section 6 concludes the paper. ",
        "conclusions": "\\label{concl}  In this paper, we investigated the influence of nearly critical rotation of the progenitor star on SN~Ia explosions. We performed three dimensional numerical simulations of thermonuclear deflagrations for different rotation laws of the white dwarf progenitors and various initial conditions. The influence of the ignition process turns out to be larger in rotating stars compared to their non-rotating counterparts, which might contribute to the observed diversity among SNe~Ia.  The main result of this study is that the amount of iron group elements is not significantly increased compared to the non-rotating case, although the rotating WD models are notably heavier than their non-rotating counterparts. Due to the centrifugal expansion, rotators contain more material at low densities in general, thus the amount of material capable of burning to iron group elements is only weakly increased if the star is burnt by a deflagration. Furthermore, due to rotationally induced anisotropic buoyancy effects and, at the same time, inhibition of effective mixing parallel to the equatorial plane, the flame preferentially propagates toward the stellar poles. This leads to comparably weak and anisotropic explosions that leave behind unburnt material at the centre and in the equatorial plane. Moreover, no significant effects of shear motion acting on the flame are observed. Deflagrations of rapid rotators do not cause high velocity features but result in overall low expansion velocities. As a consequence, the pure deflagration scenario is ruled out in the case of rapid rotation of the progenitor star. The incineration of the critical rigid rotator shows similar features as the non-rotating scenario. Therefore, deflagrations could be possible for critical rigid rotation.  In conclusion, rotation of the progenitor star is unlikely to be the parameter that causes the observed variation in peak luminosities among SNe~Ia. Rapid rotation of the progenitor star  derived by an accretion study can lead to a great variety in explosion strengths in the deflagration scenario, but only within a range that is ruled out for observational reasons. Otherwise, even the possible critical rigid rotation can not account for a significant spread in the explosion outcome."
    },
    "0911/0911.4360_arXiv.txt": {
        "abstract": "{In Spitzer observations of Tauri stars and their disks, PAH features are detected in less than 10\\% of the objects, although the stellar photosphere is sufficiently hot to excite PAHs.  To explain the deficiency, we discuss PAH destruction by photons assuming that the star has beside its photospheric emission also a FUV, an EUV and an X--ray component with fractional luminosity of 1\\%, 0.1\\% and 0.025\\%, respectively. {As PAH destruction process we consider unimolecular dissociation and present a simplified scheme to estimate the location from the star where the molecules become photo-stable. We find that soft photons with energies below $\\sim 20$\\,eV dissociate PAHs only up to short distances from the star ($r < 1$\\,AU); whereas dissociation by hard photons (EUV and X--ray) is so efficient that it would destroy all PAHs (from regions in the disk where they could be excited). } As a possible path for PAH survival we suggest turbulent motions in the disk. They can replenish PAHs or remove them from the reach of hard photons.  For standard disk models, where the surface density changes like $r^{-1}$ and the mid plane temperature like $r^{-0.5}$, the critical vertical velocity for PAH survival is proportional to $r^{-3/4}$ and equals $\\sim$5\\,m/s at 10\\,AU which is in the range of expected velocities in the surface layer.  The uncertainty in the parameters is large enough to explain both detection and non-detection of PAHs.  Our approximate treatment also takes into account the presence of gas which, at the top of the disk, is ionized and at lower levels neutral. } ",
        "introduction": "Infrared emission bands of PAHs can be used as a probe of the UV environment.  They are commonly seen in the ISM, but also in young stellar objects such as Herbig Ae/Be stars (Waelkens et al. 1996, Siebenmorgen et al. 2000, Meeus et al. 2001, Peeters et al. 2002, van Boekel et al. 2004).  The observed emission can be explained in models of an irradiated disk (Habart et al. 2004, Visser et al., 2007, Dullemond et al. 2007a).  ISO also looked at a few of the much fainter T Tauri stars but without a clear PAH detection (Siebenmorgen et al. 2000).  In the Evans et al.~(2003) legacy program which employs the more sensitive Spitzer Space Telescope (SST), 3 out of 38 T Tauri stars show PAH features (Geers et al.~2006).  This corresponds to a detection rate of only 8\\% in contrast to almost 60\\% in Herbig Ae/Be stars (Acke \\& van den Ancker 2004).  Similarly low rates for T Tauri stars are found by Furlan et al.~(2006) who present 111 SST spectra in the Taurus-Auriga star forming region and speculate that the absence of PAH resonances is due to the much weaker UV field compared to Herbig Ae/Be stars. Geers et al.~(2009), on the other hand, argue that the PAHs are simply under-abundant relative to the ISM.  They also find that variations of the disk geometry, such as flaring or gaps, have only a small effect on the strength of the PAH bands.  Clearing out gas and dust by planet formation inside the disk could effectively remove PAHs. Indeed, inner gaps in disks are observed at radii between 40--60\\,AU and at wavelengths between 20--1000$\\mu$m where the emission is dominated by large grains (Brown et al. 2008, Geers et al. 2007b).  However, in cases where PAH emission is resolved, it is extended up to 15--60\\,AU, without sub-structure and inside the inner gap region (Geers et al. 2007a).  The spatial extent of the PAH emission is also similar for T Tauri and Herbig Ae/Be stars.  We therefore suggest that PAH removal by radiative destruction is dominant.  Present radiative transfer models of the PAH emission from dusty disks consider only the stellar radiation field and no additional EUV or X--ray component (Habart et al.~2004, Geers et al.~2006, Visser et al.~2007, Dullemond et al.~2007).  Their hard photons could, according to laboratory experiments (Ruhl et al. 1989, Leach et al. 1989a,b, {Jochims et al. 1994}) and theory (Omont 1986, {Tielens 2005}, Rapacioli et al. 2006, { Micelotta et al. 2009}), destroy PAHs.  We discuss below their impact on the PAH abundance in the disks of T Tauri stars. ",
        "conclusions": ""
    },
    "0911/0911.0891_arXiv.txt": {
        "abstract": "Spiral galaxies have most of their stellar mass in a large rotating disk, and only a modest fraction in a central spheroidal bulge. This challenges present models of galaxy formation: galaxies form at the centre of dark matter halos through a combination of hierarchical merging and gas accretion along cold streams.  Cosmological simulations thus predict galaxies to rapidly grow their bulge through mergers and instabilities, and to end-up with most of their mass in the bulge and an angular momentum much below the observed level, except in dwarf galaxies. We propose that the continuous return of gas by stellar populations over cosmic times could help to solve this issue. A population of stars formed at a given instant typically returns half of its initial mass in the form of gas over 10 billion years, and the process is not dominated by supernovae explosions but by the long-term mass-loss from low- and intermediate-mass stars. Using simulations of galaxy formation, we show that this gas recycling can strongly affect the structural evolution of massive galaxies, potentially solving the bulge fraction issue, as the bulge-to-disk ratio of a massive galaxy can be divided by a factor of 3. The continuous recycling of baryons through star formation and stellar mass loss helps the growth of disks and their survival to interactions and mergers. Instead of forming only early-type, spheroid-dominated galaxies (S0 and ellipticals), the standard cosmological model can successfully account for massive late-type, disk-dominated spiral galaxies (Sb-Sc). ",
        "introduction": "The formation of disk-dominated galaxies remains a challenge for modern cosmology \\citep{white09, burkert09}. Bright spiral galaxies with stellar masses around $10^{11}$~M$_{\\sun}$ at $z=0$ typically have bulge-to-disk mass ratios (B/D) of 0.2 to 0.4 \\citep{weinzirl}. For instance, the Milky Way itself has a B/D ratio of 0.35-0.40 \\citep{robin03}. Luminosity ratios in near-infrared and optical bands are even lower \\citep{Graham2008}.  However, cosmological simulations predict much higher bulge fractions, and consequently a baryonic angular momentum much lower than observed. The hierarchical assembly of dark matter halos drives successive galaxy mergers, followed by a rapid growth of central bulges. Indeed, a single major merger, even starting with a high gas fraction, will generally end-up with more baryons in the bulge than in the disk \\citep{SH05,robertson06,hopkins09b}. Successive minor mergers also drive bulge growth (Bournaud, Jog \\& Combes 2007).  It has been realized recently that the assembly of baryons onto galaxies is far from being only driven by mergers. Rapid accretion of cold gas is another major mode of galaxy assembly \\citep{dekel09}. This cold accretion mode could apparently feed disk-dominated galaxies, but it makes the disks so massive and turbulent that they violently fragment into giant clumps (Bournaud \\& Elmegreen 2009; Dekel, Sari \\& Ceverino 2009; Agertz, Teyssier \\& Moore 2009; Burkert et al. 2009). Clump coalescence will rapidly fuel the bulge (\\citealp{noguchi}; Bournaud, Elmegreen \\& Elmegreen 2007), and signatures of this additional bulge formation process have been observed \\citep{genzel08, elmegreen09}. Disk internal secular evolution can also lead to the slow buildup of pseudo-bulges \\citep{Kormendy2004}.   A general result is that the rapid formation of massive bulges appears unavoidable because of the combination of mergers and disk instabilities. The most disky galaxies in cosmological models have B/D around 1 in the best cases and generally higher, in both the merger-driven and stream-fed dominant modes (\\citealp[e.g.,][]{governato09,white09,gibson09,scannapieco09}; Ceverino, Dekel \\& Bournaud 2009). The bulge fraction problem is related to the well-known spin crisis: galaxies could gain angular momentum from gravitational torquing by satellites (Porciani, Dekel \\& Hoffman 2002), but the early build-up of bulges quenches this process. Energy feedback from supernovae explosions can regulate star formation and keep the gas fraction higher, so that disk components survive mergers more easily. Stream-fed disks with efficient supernovae feedback can then have realistic rotation velocities, in agreement with the observed Tully-Fisher relation \\citep{maller-dekel, scannapieco08,CK09,governato09,gibson09,Piontek2009a}, but the mass fraction in these fastly rotating disks remains too low compared to their bulges, meaning that the real angular momentum is low. High bulge fractions thus remain as a separate issue that supernovae feedback cannot solve, even when many feedback models are tested in numerical simulations \\citep{scannapieco09,Piontek2009b}. Apart for dwarf galaxies \\citep{governato09b}, the standard models apparently fail to form high-spin, disk-dominated galaxies.   A stellar evolution process that is more general than supernovae feedback is the continuous return of gas from stars of any mass, in particular through stellar winds and planetary nebulae \\citep{faber76}. The gas return fraction of a coeval stellar population over the Hubble time can be as high as 50\\% for a Chabrier IMF, or around 40\\% for a Scalo IMF, with small dependencies on the metallicity (Jungwiert, Combes \\& Palou\\v{s} 2001; Lia, Portinari \\& Carraro 2002; Pozzetti et al. 2007). Supernovae account only for about 10\\% of this return fraction, which is dominated by low- and intermediate-mass stars. Half of the stellar mass forming at redshift $ z \\sim 2$ will thus be returned in the form of gas by redshift 0, which can cool down in the disk and form new stars that will in turn return gas, in a continuous recycling process. This is another type of feedback, which is more a {\\em mass} feedback than the energy feedback from supernovae explosions usually considered.  This continuous mass return has already been included in a few cosmological simulations \\citep[e.g.,][]{SC02,Wiersma2009}, but these simulations were limited to high redshift, or lacked resolution to clearly resolve disks and bulges, or the specific effects of gas return were not explored.   Here we show that this known process has unsuspected effects on regulating the growth of bulges and the survival of massive disks in spiral galaxies. Using a cosmological simulation of the formation of a massive galaxy, which in the standard model ends-up with a B/D ratio larger than unity, we show that continuous gas recycling reduces the final B/D ratio to 40\\%, and the stellar light fraction in the bulge by a factor of 3. The mass and disk rotation speed of this galaxy are consistent with observed scaling relations, and its remarkably flat rotation curve is typical for late-type, high-spin spirals. The mass return of stars regulates the bulge growth and angular momentum dissipation in such proportion that a late-type spiral (Sb-Sc) is formed instead of an early-type galaxy. This shows how strongly stellar evolution can affect the formation of galaxies, and help to solve the issue of massive late-type spiral galaxy formation. ",
        "conclusions": "Our simulation of cosmological galaxy formation shows that the continuous gas return from high-, intermediate- and low-mass stars over cosmic times can strongly regulate the growth of bulge and the dissipation of angular momentum in disk galaxies. There are at least two ways in which stellar mass loss can support the survival of large, massive disks. First, it can keep the gas fraction higher in young galaxies before a merger occurs, which helps disks survival against destruction into bulges \\citep{hopkins09a}. Second, the bulge, stellar halo and thickened disk components release fresh gas after interactions and mergers: this decreases their mass and increases that of the thin gas disk. The disk re-formed by gas returned by a bulge or halo should have a low angular momentum, but can be torqued by satellites into a large, high-spin disk that gradually forms new stars (an example of such an interaction is shown at $z=1$ on Figure~\\ref{snapshots}). The radial migration of disk stars can also arise from resonant interactions with spiral arms: \\cite{Roskar2008} have shown, in a Milky-Way like galaxy, that up to 50\\% of stars in the solar neighborhood could have migrated from the inner disk. In our model, the large disk scale length is preserved by the combination of stellar mass-loss and radial migration following external interactions and/or internal evolution.  Overall, the effect of stellar mass loss on disk survival and regulation of bulge growth is major. It reduces the bulge-to-disk ratio by a factor up to $\\sim$ 3 in both the stellar mass and light. This factor is what was typically missing in cosmological models to account for the formation of massive, disk-dominated, late-type spiral galaxies.   This does not imply that all galaxies would end-up with a spiral-like morphology at redshift zero, and we briefly presented cases of galaxies with other merger and star formation histories where the decrease of B/D is weaker. Gas return promotes disk survival against mergers and disk instabilities, but will preserve a disk-dominated galaxy only if the mergers are not too numerous, do not happen too late, and the internal instabilities not too violent. For instance, a galaxy undergoing a major merger at low redshift would still end-up as a $z=0$ elliptical or lenticular, because of a stellar population globally too old to return a significant amount of gas. Such $z=0$ ellipticals are indeed found in our cosmological simulations with the same technique \\citep[e.g. in][]{martig09}. Gas consumption, stripping and strangulation in groups, will also support the survival of early-type elliptical and S0s. Thus, gas return from stellar populations can explain the origin of late-type galaxies, without challenging the formation of early-type galaxies in systems that undergo later mergers or more violent internal instabilities.   Other sorts of feedback processes, like the energy released by supernovae or the radiation pressure from young massive stars (Murray, Quataert \\& Thompson 2009) could further regulate the bulge growth especially in low-mass galaxies, potentially forming almost bulgeless galaxies. However, we propose that the formation of massive late-type galaxies in the $\\Lambda$-CDM Universe is explained not by an unknown combination of minor factors, but mostly by one major effect, namely by the gas return from stellar populations all across their initial mass function. This process can largely help to solve one of the main challenges in the origin of modern galaxies, namely the origin of late-type, disk-dominated spiral galaxies like the Milky Way."
    },
    "0911/0911.5528_arXiv.txt": {
        "abstract": "Hyperaccreting neutron star or magnetar disks cooled via neutrino emission can be a candidate of gamma-ray burst (GRB) central engines. The strong field $\\geq10^{15}-10^{16}$ G of a magnetar can play a significant role in affecting the disk properties and even lead to the funnel accretion process. In this paper we investigate the effects of strong fields on the disks around magnetars, and discuss implications of such accreting magnetar systems for GRBs and GRB-like events. We discuss quantum effects of the strong fields on the disk thermodynamics and microphysics due to modifications of the electron distribution and energy in the strong field environment, and use the magnetohydrodynamical conservation equations to describe the behavior of the disk flow coupled with a large scale field, which is generated by the star-disk interaction. If the disk field is open, the disk properties mainly depend on the ratio between $|B_{\\phi}/B_{z}|$ and $\\Omega/\\Omega_{K}$ with $B_{\\phi}$ and $B_{z}$ being the azimuthal and vertical components of the disk field, $\\Omega$ and $\\Omega_{K}$ being the accretion flow angular velocity and Keplerian velocity respectively. On the other hand, the disk properties also depend on the magnetar spin period if the disk field is closed. In general, stronger fields give higher disk densities, pressures, temperatures and neutrino luminosity. Moreover, strong fields will change the electron fraction and degeneracy state significantly. A magnetized disk is always viscously stable outside the Alfv\\'{e}n radius, but will be thermally unstable near the Alfv\\'{e}n radius where the magnetic field plays a more important role in transferring the angular momentum and heating the disk than the viscous stress. The funnel accretion process will be only important for an extremely strong field, which creates a magnetosphere inside the Alfv\\'{e}n radius and truncates the plane disk. Because of higher temperature and more concentrated neutrino emission of a ring-like belt region on the magnetar surface covered by funnel accretion, the neutrino annihilation rate from the accreting magnetar can be much higher than that from an accreting neutron star without fields. Furthermore, the neutrino annihilation mechanism which releases the gravitational energy of the surrounding disk and the magnetically-driven pulsar wind which extracts the stellar rotational energy from the magnetar surface can work together to generate and feed an ultra-relativistic jet along the stellar magnetic poles. ",
        "introduction": "Hyperaccreting black hole systems formed in massive star collapses or compact object binary mergers have been considered as a candidate for gamma-ray burst (GRB) central engines for two decades (e.g., Eichler et al. 1989; Narayan et al. 1992; Woosley 1993). The accreting systems can drive ultra or mildly relativistic jets with huge energy via neutrino ($\\nu\\bar{\\nu}$) annihilation process (Popham et al. 1999) or magnetohydrodynamical (MHD) mechanism, such as the Blandford-Znajek mechanism (Blandford \\& Znajek 1977) or the magnetorotational instability (MRI, Balbus \\& Hawley 1991), and finally produce GRB explosions. The neutrino annihilation efficiency above hyperaccreting disks around black holes with the elaborate considerations of disk geometry, rotation, and general relativity effects have been studied by many authors (e.g., Ruffert et al. 1997, 1998; Popham et al. 1999; Asano \\& Fukuyama 2000, 2001; Miller et al. 2003; Birkl et al. 2007). Although the annihilation process can provide sufficient total energy up to $10^{50}$ ergs s$^{-1}$ above neutrino-dominated flows with accretion rate $\\sim 1M_{\\odot}$ s$^{-1}$ (Gu et al. 2006; Liu et al. 2007), it is still under debate whether such a process can successfully generate a relativistic jet from the polar region of the central black hole. For example, MacFadyen et al. (2001) discussed that the central black hole formed in the ``collapsar scenario\" is more frequent to be formed by fallback process after a mild explosion (Type II collapsar) rather than formed promptly (Type I collapsar). The Type II collapsar can establish an accretion disk with accretion rate $\\sim0.001-0.01 M_{\\odot}$ s$^{-1}$, which is not sufficient to produce a jet by neutrino annihilation. Numerical simulations showed that the MHD mechanism may be more efficient to drive an energetic magnetically-dominated wind and generate a GRB explosion (Tchekhovskoy et al. 2008, 2009; Nagataki 2009), or the annihilation process and MHD winds can work together to provide the energy for GRB explosions (Harikae et al. 2009).  On the other hand, both observational and theoretical evidences show that newborn neutron stars or magnetars rather than black holes can form in the GRB central engines (e.g., Dai \\& Lu 1998a, 1998b; Zhang \\& M\\'{e}sz\\'{a}ros 2001; Dai 2004; Dai et al. 2006; Mazzali et al. 2006; Soderberg et al. 2006; Shibata \\& Taniguchi 2006; Lee \\& Ramirez-Ruiz 2007). Highly magnetized neutron stars formed by accretion-induced-collapse, convection, $\\alpha-\\Omega$ dynamo mechanism or differential rotation can extract their rotational kinetic energy up to $\\sim 10^{51}-10^{52}$ ergs by spinning down via magnetic activities (e.g., Usov 1992; Duncan \\& Thompson 1992; Klu\\'{z}niak \\& Ruderman 1998). A thermally-driven neutrino wind is dominated during the Kelvin-Helmholtz cooling epoch after the neutron star formation, lasting from a few seconds to tens of seconds based on the strengths of surface fields. After that, magnetically-dominated or Poynting-dominated wind becomes significant along the polar region (Wheeler et al. 2000; Thompson 2003; Thompson 2004; Metzger et al. 2007). The jets from magnetars may be accelerated to higher Lorentz factor $\\sim100-10^{3}$ at large radius several tens of seconds after core bounce, and bring its magnetic energy to kinetic energy, which is finally dissipated in regular GRB internal shocks (Thompson et al. 2004; Metzger et al. 2007; Bucciantini et al. 2007, 2009; Lyubarsky 2009a, 2009b).  However, all the magnetized neutron star or magnetar models ignore the accretion process which occurs onto a protoneutron star for the first several seconds (Woosley \\& Bloom 2006). In the scenario of massive star collapse, the outgoing shock generated by a successful Type Ibc supernova explosion with a velocity of $\\sim10^{9}$ cm s$^{-1}$ can evacuate a cavity around the new born compact star in the center of a supernova remnant. However, whether the outgoing shock in the core collapse of a massive star can dominate the accretion process and turn the accretion around is still unknown. If the rotational core collapse can lead to the formation of a neutron star or a magnetar, there is possible that the prompt accretion or fallback process can make a hyperaccreting disk around the young formed star. On the other hand, recent simulations also showed the possibility of a debris disk around a massive neutron star as the outcome of a compact binary merger (e.g., Shibata \\& Taniguchi 2006, Lee \\& Ramirez-Ruiz 2007). Similar to the black hole system, the transiently existing torus around a neutron star or magnetar with sufficient angular momentum can form a neutrino-cooled disk and release its gravitational binding energy via neutrino emission.  The hyperaccreting neutron star or magnetar system is also proposed as another possible central engine of GRBs. Zhang \\& Dai (2008a, 2009, hereafter papers I and II) found that, due to the inner surface boundary of the compact star, the disk has a denser, hotter inner region with higher pressure compared to its black hole counterpart. Also, the entire disk can cool more efficiently via neutrino emission and produce a higher neutrino annihilation luminosity. A heavily thermally-driven outflow from the surface of the new-born neutron star at early times ($\\sim 100$ms) prevent the outflow from being accelerated to a high Lorentz factor (Dessart et al. 2009). However, if the disk accretion rate and the neutrino emission luminosity from the surface boundary layer are sufficiently high, an energetic ultrarelativistic jet via neutrino annihilation can be actually produced after the jet breaks out of the stellar envelope around the disk-neutron star system.  Until now, none of the above works consider the effects of magnetic fields on the hyperaccreting transient disks around the central compact stars. The MHD mechanisms in the accreting black hole systems mainly focus on energy extraction near the central black holes, and the proto-magnetar magnetically-driven wind model discusses wind generation and propagation along the polar regions. In the hyperaccreting disk the spherical Alfv\\'{e}n radius of the central star can be estimated as \\begin{equation} r_{A}\\simeq0.207\\dot{M}_{-2}^{-2/7}M_{1.4}^{-1/7} \\mu_{30}^{4/7} \\textrm{km}, \\end{equation} where $\\dot{M}=0.01\\dot{M}_{-2}M_{\\odot}$ s$^{-1}$ is the accretion rate, $M=1.4M_{1.4}M_{\\odot}$ is the central star mass, $\\mu=\\mu_{30}10^{30}$ G cm$^{3}$ is the central magnetic flux. Figure 1 shows the Alfv\\'{e}n radius $r_{A}$ as a function of accretion rate with various magnetic fields. If the surface field $B_{0}$ is less than \\begin{equation} B_{0}\\leq B_{\\rm crit}=0.89\\times10^{15}\\dot{M}_{-2}^{1/2}M_{1.4}^{1/4} r_{*,6}^{-5/4} \\textrm{G} \\end{equation} with $r=r_{*,6}10^{6}$ cm being the star radius, the accretion flow will continue to be confined in the disk plane without co-rotating with the compact object or getting funneled onto the magnetar poles. Papers I and II assumed the strength of stellar surface field is less than $10^{15}$ G, and did not consider the effects of a magnetic field in the disk. Recently, Xie et al. (2007, 2009) and Lei et al. (2009) studied the structure of magnetized hyperaccreting neutrino-cooled disks around black holes based on similar models, while the magnetic fields are generated in the disks up to $10^{15}-10^{16}$ G (see also Janiuk \\& Yuan 2009). Xie et al. (2007, 2009) showed that the magnetic braking and viscosity can drive magnetically-dominated accretion flows rather than neutrino-dominated, and turn the disk temperature to be lower than that without field. Lei et al. (2009) investigated the properties of the NDAF with the magnetic torque acted between the central black hole and the disk. The neutrino annihilation luminosity can be increased by one order of magnitude for accretion rate $\\sim0.5M_{\\odot}$ s$^{-1}$, and the disk becomes thermally and viscously unstable in the inner region. However, all of their works neglected microphysics processes and the changes of equations of state in strong magnetic fields, which may change the disk pressure and various neutrino cooling rates in the neutrino-cooled disks significantly. Furthermore, if we look into the neutron star disks, we should keep in mind that the structure of magnetic fields in disks around magnetized neutron stars or magnetars are very different from their black hole counterparts. The origin of a magnetic field in this case in the central compact star constructs the initial topology of the strong field. An interaction between the star and the surrounding disk makes the stellar magnetic field partially thread the accretion disk.  Our motivation in this paper is to investigate the effects of a strong magnetic field on the hyperaccreting disk around a magnetar formed in the collapse of a massive star or the merger of a compact object binary, and their observational phenomena related to GRBs and associated events. The accretion process onto the magnetized neutron star has been widely studied in X-ray binaries, T Tauri stars and cataclysmic variables since Ghosh \\& Lamb (1978, 1979a, 1979b), both theoretically (e.g., K\\\"{o}nigl 1991; Ostriker and Shu 1995; Lovelace et al. 1995; Wang 1995; Lai 1998), and numerically (e.g., Hayashi et al. 1996; Miller \\& Stone 1997; Romanova et al. 2005, 2008; Long et al. 2007, 2008). The neutrino-cooled hyperaccreting disks, on the other hand, are much hotter, denser with much higher pressure and accretion rates compared to the normal disks. Therefore the structure and radiation process in these magnetized disks must be very different from the normal ones. For example, the size of the magnetosphere near the central star can be ignored except for a magnetic field $\\geq10^{16}$ G. The pressure relation $B^{2}/8\\pi\\gg P_{\\rm matter}+\\rho v_{r}^{2}$ will not be satisfied in most cases. More attractively, we need to study the neutrino emission and annihilation process which is almost the most important properties of the hyperaccreting disks but never occur in the normal magnetized ones.  This paper is organized as follows. In Section 2 we consider the quantum effects of a strong magnetic field in the microphysical scale on disk density, pressure and various neutrino cooling rates. In Section 3, we discuss the disk conservation equations coupled with the global fields both from the central magnetar and generated via the magnetar-disk interaction. Also, we discuss the field topology of the central magnetar. Combining the equations in Section 2 and 3, we discuss the properties of the hyperaccreting disks in the strong magnetic fields numerically in Section 4. The structure of the magnetized disks depends on the detailed field strength and configuration. In Section 5.1, we further discuss the disk outflows, which may provide the kinetic energy of the supernova associated with a GRB (MacFadyen \\& Woosley 1999; Kohri et al. 2005). Then in Section 5.2, we describe the funnel process onto the magnetar surface when the field strength is extremely high. The funnel process will significantly affect the neutrino annihilation efficiency. Section 6 is a broad discussion. We consider the importance role played by Joule dissipation, nucleon reactions and $r$-process occurring in the magnetized disk, as well as various outcomes of the massive star collapse using a unified point of view. Moreover, we particularly focus on the application of magnetized hyperaccreting disks in the GRB and GRB-like events (e.g, X-ray flashes), and use a unified point of view to discuss the core collapse models related to GRBs. Conclusions are presented in Section 7. ",
        "conclusions": "The hyperaccreting neutron star or magnetar disk systems cooled via neutrino emission can form by the mergers of compact star binaries or the collapses of rotational massive stars. Strong fields of the magnetar can play a significant role in affecting the disk properties and even changing the accretion process. Our motivation in this paper is to investigate the influence of such strong magnetic fields on the disks, and discuss implications of the magnetar disk systems for the GRB and GRB-like events. Our conclusions are as follows.  (i) We consider the magnetar field has a dipolar vertical component $B_{z}$. The differential rotation between the disk and the magnetar will generate a toroidal field component $B_{\\phi}$, as well as a relatively weak radial component $B_{r}$. The generated field can have an open or closed configuration, depending on the disk's viscous turbulence, magnetic diffusivity, disk angular velocity as well as twist limitation. Similar to pervious works, we use the parameter $\\beta$ to measure the strength of the toroidal field (equations (\\ref{fieldstru01}) and (\\ref{fieldstru02})). The generated large-scale disk field coupled with the accretion flows will transfer the angular momentum in radial direction and heat the disk via Joule dissipation together with viscous stress outside the Alfv\\'{e}n radius $r_{A}$. On the other hand, since the distribution and energy of electrons change significantly in the strong field environment, the disk pressure and a variety of neutrino cooling processes will be different compared to the case without fields. Therefore, we discuss the quantum effects of the strong fields on the disk thermodynamic and microphysical processes in Section 2, and list the MHD conservation equations to describe the behavior of the large-scale magnetic field coupling in the disk in Section 3.  (ii) The quantum effects and field coupling in MHD equations play two competitive roles in changing the disk properties, the former to decrease the pressure, density and neutrino luminosity with increasing field strength (Figure 4), while the latter to increase them (Figure 5). However, in most cases the large-scale field coupling is more significant than the microphysical quantum effect (Figure 7). Moreover, strong fields will change the electron fraction distribution $Y_{e}$ in the disk significantly. Larger peak of electron chemical potential $\\eta_{e}$ with more degeneracy electron states is obtained by stronger fields, while the change of $\\eta_{e}$ becomes more obvious for stronger fields.  (iii) Similar to the neutron star disk without fields, the values of density, pressure, temperature, neutrino luminosity and electron chemical potential become higher for higher accretion rate, but the electron fraction $Y_{e}$ decreases with increasing the accretion rate.  (iv) The magnetized disk maintaining a plane geometry would be more favorable for an open field configuration rather than a closed one. However, we still consider the two cases for completeness (Figures 6 and 9). For the disk with an open field and $\\dot{M}=0.1M_{\\odot}$ s$^{-1}$, higher ratio of $\\beta/s$ leads to higher density and pressure in the entire disk plane, higher temperature and neutrino luminosity in the inner part of the disk, as well as lower $Y_{e}$ in the outer part. Here $s$ is the ratio of disk angular velocity and the Keplerian velocity. Also, electrons becomes more degeneracy at $\\sim20-40$ km for higher $\\beta/s$. On the other hand, the disk properties with a closed field not only depends on the values of $\\beta$ and $s$, but also the spin period of the central magnetar. Shorter period of the central star decreases the disk density, pressure, $\\eta_{e}$, but increases $Y_{e}$ at the outer part of the disk. The effects of spin period are only obvious for sufficient field strength (e.g., $\\sim10^{16}$ G for the accretion rate $0.1M_{\\odot}$ s$^{-1}$) from the central star.  (v) The accretion flow in the disk plane outside the Alfv\\'{e}n radius is viscously stable. However, whether the disk is thermally stable depends on many factors such as the disk region, magnetic field strength, disk angular velocity and accretion rate. Generally speaking, the disk will be definitely thermally unstable if its non-field disk counterpart with the same accretion rate is unstable. The disk region can also be thermally unstable even its non-field counterpart is stable, if the region is near the Alfv\\'{e}n radius where magnetic field plays a more important role in transferring the angular momentum and heating the disk than the viscous stress.  (vi) The thermally-driven outflows can also exist in the magnetar disk beyond the Alfv\\'{e}n radius for the case of $\\dot{M}<1.0M_{\\odot}$ s$^{-1}$. We assume the accretion rate as a power law in radius for simplicity (Equation (\\ref{wind01})). The outflow will take away the disk angular momentum (Equation (\\ref{wind02})) and may provide energy for a supernova associated with the GRB explosion. The disk density, pressure and neutrino luminosity decrease with increasing the outflow strength. Also, the ratio of magnetic and matter pressure $P_{\\rm B}/P_{\\rm matter}$ and the thickness of the disk become larger for stronger outflows. However, since the total energy taken by the thermal outflow depends on the size of the disk plane outside the Alfv\\'{e}n radius, the total energy injection rate from the outflow decreases significantly for stronger stellar fields (Table \\ref{tab01}). Besides the thermally-driven outflow, strong fields inside the Alfv\\'{e}n radius lead the accretion flow to co-rotating with the stellar field, which is possible to launch MHD winds along the field lines, and generate the X-type wind near the disk-magnetosphere boundary. These magnetically-dominated winds are nonrelativistic and may provide energy for a supernova explosion.  (vii) In the hyperaccreting disks, the funnel accretion can only be important for extremely strong fields, which depend on the accretion rate. The accretion process along the disk plane will be truncated in the stellar magnetosphere, so most of the accretion flow can be approximately considered as being lifted along the closed field lines onto the magnetar surface. In most cases, the hyperaccreting infalling funnel flow is unlikely to develop a shock with Mach number greater than unify (equation (\\ref{f08})). The flow is accelerated by the gravitational force and transfers the gravitational binding energy to the kinematic energy and next to the heating energy near the magnetar surface. The funnel flow will cover a ring-like belt of ``hot spot\" around the magnetar surface and emit thermal neutrinos. Because the temperature is higher and neutrino emission region is more concentrated than the hyperaccreting neutron star without fields, the funnel accretion can accumulate more powerful neutrino annihilation luminosity (Figure 13).  (viii) The neutrino annihilation process both from the magnetar surface and from the disk plane will be higher than that without fields. Moreover, the neutrino annihilation mechanism and the magnetic activity from the stellar surface (i.e., the pulsar wind mechanism) can work together to generate and feed an ultra-relativistic jet along the stellar magnetic poles. If the stellar spin period is sufficiently short (e.g., $\\sim4$ ms for the field $\\sim10^{16}$ G and $\\dot{M}=0.1M_{\\odot}$ s$^{-1}$), the jet from the magnetar will be magnetically-dominated and mainly feeded by extraction the stellar rotational energy. If the magnetar spin period is long (i.e., longer than the critical value $P_{rc}$ in Equation (\\ref{grb1})), the jet is thermally-driven and feeded by the annihilation process. In the intermediate case, on the other hand, the relativistic jet can be launched by the pulsar-wind-like process and neutrino annihilation together."
    },
    "0911/0911.2552_arXiv.txt": {
        "abstract": " ",
        "introduction": "Globular clusters (GCs) survive for a Hubble time and undergo dissipationless dynamical evolution.  They have been shown to form during major episodes of star formation    \\citep[][among others]{holtzman92,schweizer93,whitmore93,whitmore95,zepf95,schweizer96,miller97,carlson98,schweizer98,whitmore99,zepf99,chien07,goudfrooij07,trancho07}. This enables the use of GCs as probes of the formation history in their host galaxy which can tell us important information on the galaxy's assembly history and the physical conditions in which the GCs were formed.  We can use their ages and metallicities to reconstruct the star formation history of their galaxy \\citep{west04} as well as use the GCs as kinematic and dynamic tracers \\citep{bridges06}.  There are  many advantages in using  GCs as probes  of star formation and galaxy assembly.  GCs are coeval and form with a single age and a single metallicity to first approximation.  We also find  large GC  systems in early-type galaxies, providing a large basis for study. GC systems have been shown to exhibit two colour modes, both blue and red \\citep{peng06}, indicating there could have been at least two episodes of star formation within these galaxies.  Since the large majority of the colour change of the  GC light that we see happens within the first few Gyr \\citep{worthey94}, any subsequent colour difference is due to a difference in metallicity, so we refer to blue as metal-poor and red as metal-rich GCs. The advancements of multi-object spectrographs have allowed us to obtain numerous GC spectra in a homogeneous manner, providing large sample sizes.  \\subsection{The Globular Cluster System of NGC 5128} NGC 5128 is the nearest available giant elliptical for study at $3.8 \\pm0.1$ Mpc (G.L.H Harris, 2009, private communication).  Its close proximity provides the opportunity to perform a detailed study of its GC system at a level of detail that is not possible in other giant ellipticals at present. We will use the GCs as representations of the bulk stellar population to probe its formation history.  In  NGC 5128, there are  an estimated  1500 GCs within 25 arcmin \\citep{harris06}.  We currently  know of  70 GCs  within NGC 5128 that have been identified as resolved GCs from {\\it Hubble Space Telescope} images \\citep{harris06}.  There are 564 GCs also confirmed by radial velocity measurement \\citep{vandenbergh81,hesser84,hesser86,harris92,peng04b,woodley05,rejkuba07,beasley08}, including 190 new GCs our group has recently identified \\citep{woodley09a,woodley09b}. The sample now includes 605 GCs which is one of  the  largest  samples  of GCs  in  the literature. Of the total GC system with known photometry, we have identified 268 metal-poor and  271 metal-rich GCs, with a clear bimodal distribution in colour. Figure~\\ref{fig:Theta_R} shows the projected spatial distribution of the GC system which extends out to 45 arcmin in galactocentric radius, distributed mainly along the isophotal major axis of the  galaxy and heavily concentrated within 15 arcmin.  Our sample is thus spatially biased, caused primarily by the chosen field locations of previous studies.  \\begin{figure}[h] \\begin{center} \\includegraphics[scale=0.35, angle=0]{f1.eps} \\caption{The galactocentric radial distribution of GCs in NGC 5128 as function of azimuthal position, measured in degrees E of N. The {\\it squares} and the {\\it triangles} are the metal-poor and metal-rich GCs with a measured velocity, respectively, and the {\\it solid circles} are GCs that either have measured velocities but no colour information, or GCs with no measured radial velocities.}\\label{fig:Theta_R} \\end{center} \\end{figure} ",
        "conclusions": "We have studied the GC system in NGC 5128 and presented their ages and metallicities, their kinematics, their structural parameters, and we have also used them as tracers to estimate the mass of NGC 5128.  We find, using a sample of 72 GCs, that the majority of metal-rich and metal-poor GCs are old with 68\\% having ages $> 8$ Gyr.  We do find a small fraction with young ages $< 5$ Gyr, all of which are metal-rich. These results suggest there could have been a number of formation epochs within NGC 5128, which resulted in the formation of a small portion of its stellar population.  Using over 560 GCs in NGC 5128, we performed a kinematic analysis, and find the metal-rich GCs have a mild rotation of $ 41\\pm15$ km s$^{-1}$ and rotate around $191\\pm18$ deg E of N, similar to the isophotal major axis of the galaxy.  The metal-rich GCs have very mild rotation and do not appear to rotate around either isophotal major or minor axis.  The velocity dispersion profile of the GCs is generally flat in the inner regions of the galaxy.  Compared to the stellar halo population, we see a number of similarities to the metal-rich GCs, including the rotation axis and surface density profile of the PNe population, as well as approximate ages and metal enrichment of the stellar halo \\citep{rejkuba05}.  We also use the radial velocity measurements of the GCs to estimate a mass of NGC 5128 to be $5.5\\pm1.9 \\times 10^{11}$ M$_{\\odot}$ and a mass-to-light ratio of 15.35 M$_{\\odot}/$L$_{B\\odot}$ out to 20 arcmin.  We examine the structural parameters of a representative sample of normal GCs in NGC 5128 using superb IMACS images.  We examine their half-light radii, and find within 1--2 effective radii of the galaxy, the blue GCs are $\\sim30\\%$ larger than the red GCs.  Beyond this distance, however, the size difference seems to disappear.  We suggest this may be the result of their environment of formation, combined with the projection effect.  The old ages and kinematics of the GCs within NGC 5128 suggests that the majority of GCs, and presumably stars, formed early on in a rapid, or perhaps multiple, collapse(s).  This is supported by the old ages of the GCs and high [$\\alpha$/Fe] for both metal-rich and metal-poor GCs.  There is evidence for a large spread in metal-rich ages, indicating either minor merging or accretion of small neighbouring satellites in more recent times.  However the fraction of young objects is heavily biased by our selection of targets.  The youngest GCs that we find in our sample are a few Gyrs old."
    },
    "0911/0911.2764_arXiv.txt": {
        "abstract": "We study non-axisymmetric oscillations of rapidly and differentially rotating relativistic stars in the Cowling approximation. Our equilibrium models are sequences of relativistic polytropes, where the differential rotation is described by the relativistic $j$-constant law. We show that a small degree of differential rotation raises the critical rotation value for which the quadrupolar f-mode becomes prone to the CFS instability, while the critical value of $T/|W|$ at the mass-shedding limit is raised even more. For stiffer equations of state these effects are even more pronounced. When increasing differential rotation further to a high degree, the neutral point of the CFS instability first reaches a local maximum and is lowered afterwards. For stars with a rather high compactness we find that for a large degree of differential rotation the absolute value of the critical $T/|W|$ is below the corresponding value for rigid rotation. We conclude that the onset of the CFS instability is eased for a small degree of differential rotation and for a large degree at least in stars with a higher compactness. Moreover, we were able to extract the eigenfrequencies and the eigenfunctions of r-modes for differentially rotating stars and our simulations show a good qualitative agreement with previous Newtonian results. ",
        "introduction": "\\label{sec:introduction}  Rapidly spinning neutron stars can be destabilized due to the CFS \\cite{1970ApJ...161..561C,1978ApJ...221..937F,1978ApJ...222..281F} instability. This instability, for quadrupole non-axisymmetric perturbations with $m=2$, sets in for uniformly rotating polytropes when $\\beta=T/|W|\\approx 0.14$ where $T$ is the rotational kinetic energy and $W$ is the gravitational potential energy of the star. For higher values of $m$ the instability sets in for smaller rotational rates but the growth time is considerable longer and viscosity stabilizes the star. The CFS instability is generic for r-modes, i.\\,e. these modes are CFS unstable for any rotational rate while for the f-modes the previous value of $\\beta$ is approximately correct. This value of $\\beta$ corresponds to rotational frequencies of the stars which are 85-90\\% of the Kepler frequency. These high spin frequencies have not yet been observed in nature but it is expected that they may be met in newly born neutron stars.  Newly born neutron stars are expected to rotate not only fast but also differentially, and this is a factor that has to be taken into account in the study of instabilities.  Actually, differential rotation can appear in many phases of the stellar evolution of a neutron star, such as in protoneutron stars~\\cite{2002A&A...393..523D,2004ApJ...600..834O}, in the massive remnant of binary neutron star mergers~\\cite{1994ApJ...432..242R,Shibata:1999wm,Rasio:1999zr}, or as a results of stellar oscillations (r-modes) that may drive the star into differential rotation via nonliner effects~\\cite{ 2000ApJ...531L.139R, 2001MNRAS.324..917L, 2001PhRvL..86.1148S, 2004PhRvD..69h4001S}. The phase where the star rotates differentially  can last between seconds to months, depending on the dissipative mechanism that drive the star to uniform rotation, such as viscosity, magnetic braking~\\cite{2000ApJ...544..397S,2003ApJ...599.1272C} and turbulent motion~\\cite{1977ApJ...217..244H,liu:044009}. However, this very short period is met during the most violent phases of neutron star's life, as in the case of a core collapse or binary merging. This is exactly the time when we expect to get the strongest emission in gravitational waves. Since the ground based detectors are reaching sensitivities which allow the detection of gravitational wave signals from oscillating or unstable neutron stars the critical point for the onset of instabilities, the growth times of the instabilities and the exact frequencies of the emitted waves are urgently needed.  In this work we study within the general relativistic framework the oscillations and instabilities of fast and differentially rotating neutron stars in the so called Cowling approximation i.e. assuming a fixed spacetime. This is the first study of its kind since earlier works were using certain approximations e.g  Newtonian theory \\cite{2002ApJ...578..413K} or slow rotation \\cite{2007PhRvD..75f4019S,2008PhRvD..77b4029P}. Results, from fully nonlinear calculations exist  only for the axisymmetric oscillation ~\\cite{Stergioulas:2003ep,Dimmelmeier:2005zk}.  Finally, the f-mode and the secular stability limits have been investigated in~\\cite{2002ApJ...568L..41Y}.  Recently, it has been found that in stars that rotate with a high degree of differential rotation an $m=2$ dynamical instability can appear even for considerably low rotation rates ($T/|W|\\sim O(10^{-2})$) as suggested in~\\cite{2002MNRAS.334L..27S,2005PhRvD..71b4014S}.   In addition, an $m=1$ dynamical instability has been identified for high degrees of differential rotation and soft equations of state \\cite{2001ApJ...550L.193C,2003ApJ...595..352S}.  More recently, the discovery of $m=1$ and $m=2$ dynamical instability even for stiff equations of state \\cite{2006ApJ...651.1068O} has been reported . Studies based on linear analysis \\cite{2005ApJ...618L..37W,2006MNRAS.368.1429S} suggest that low $T/|W|$ instabilities might be triggered when the corotation points of the unstable modes fall within the differentially rotating structure of the star.   The structure of the paper is as follows. In Section~\\ref{sec:formulation} we present the formulation that we have used to construct the equilibrium configurations. We then derive the perturbation equations in the general case of barotropic and non-axisymmetric perturbations, which are numerically solved in Section~\\ref{sec:results} for the non-axisymmetric and barotropic case. Finally, in Section \\ref{sec:summary}, we summarize the crucial results of this study.  Throughout the paper we use geometrical units $c=G=1$. ",
        "conclusions": "\\label{sec:summary} In this work we performed time evolutions of oscillations of differentially and rapidly rotating, relativistic stars using the Cowling approximation, and we were able to investigate both axisymmetric and non-axisymmetric perturbations. We extracted eigenfrequencies and eigenfunctions by applying Fourier transforms on the obtained time-series. The agreement with literature values is excellent for both non-axisymmetric perturbations of uniformly rotating stars and axisymmetric perturbations of differentially rotating stars.  Our simulations show that a small degree of differential rotation causes the ratio of the critical $T/|W|$, where the f-mode becomes CFS unstable, to the value of $T/|W|$ at the mass-sheeding to decrease. This effect is visible for each considered equation of state and is even stronger for the stiffer ones. It implies that neutron stars can be destabilized by the CFS instability even if their spinning frequency is rather low compared to the Kepler frequency. Furthermore, when moving on to a high degree of differential rotation, the critical curve of any considered sequence exhibits a local maximum which is in agreement with Newtonian findings. Afterwards, the critical value decreases and for the sequences based on the stiffer equations of state it decreases even below the corresponding value for rigid rotation; i.\\,e. the CFS instability may act even earlier than in rigidly rotating stars as long as the star rotates with a sufficiently high degree of differential rotation.  Summarized, we showed that both a small and a high degree of differential rotation plays an important role in the evolution of a neutron star. As explained in the introduction, neutron stars rotate differentially most probably directly after their birth, but may do so also at later times. Whenever differential rotation is present, the CFS instability is a promising candidate for the destabilization of a neutron star which may finally lead to a detectable signal of gravitational waves.  Finally, we performed simulations on r-modes. Our results show a good qualitative agreement to Newtonian results. As in earlier studies which do not rely on the slow rotation approximation, we do not see any trace of a continuous spectrum.  A very important issue related to differential rotation is the freedom to construct neutron star models with much higher values of $T/|W|$. This means that nearly all available equations of state can admit models which can be CFS-unstable which is not the case for uniformly rotating stars. This enhances the probability of having the CFS-instability working in the early phases of the proto-neutron star creation which will be a significant boost in our efforts to detect gravitational waves from isolated neutron stars and via such observations to infer and constrain fundamental neutron star parameters."
    },
    "0911/0911.3281_arXiv.txt": {
        "abstract": "We attempt to constrain the black hole spin in GX 339$-$4 from spectral fitting of disc dominated data using \\textsl{RXTE} spectra from the three most recent outbursts.  We use the best current models for the disc emission, including full radiative transfer through the photosphere rather than assuming that the intrinsic emission from each radius has a (colour temperature corrected) blackbody spectrum. The results strongly depend on the poorly known binary system parameters, but we find a strict upper limit of $a_{*}<0.9$ for any distance greater than 6~kpc, assuming that the orbital inclination is the same as that of the inner disc. By contrast, the higher spin of 0.935 +/- 0.01 (statistical) +/-0.01 (systematic) claimed from fitting the iron line profile in this object requires that the inner disc is misaligned by over 20 degrees from the orbital inclination.  While some of these datasets are distorted by instrumental pileup, the same spin/inclination constraints are derived from data which are not piled up, so there is a real conflict between the two techniques to measure spin. We argue that the disc spectral fits are more likely to be robust hence that there are still issues to be understood in the iron line profile. ",
        "introduction": "An astrophysical black hole (BH) can be entirely described by two parameters in general relativity, its mass \\textsl{M} and a dimensionless spin, $a_{*}$, which is 0 for a non-rotating Schwarzschild BH and $0.998$ for a maximal Kerr BH.  Unlike mass, spin only leaves an imprint on the spacetime very close to the event horizon, so it is much more difficult to measure.  Nonetheless, it is important to constrain because it is a fundamental parameter determining the structure of the spacetime around the BH. It sets the size--scale of the last stable orbit around the BH, from $6-1.23R_g$ for $a_*=0$ and $0.998$, respectively, where $R_g=GM/c^2$. This determines the efficiency of conversion of mass to radiation for accreting objects, and may also determine the structure and power of relativistic jets. For BHs formed from stellar collapse then the resulting spin gives insight into the (poorly understood) supernovae event (Gammie, Shapiro \\& McKinney 2004) and its gravitational wave signature.  Currently there are only two methods which can be used to determine spin from accreting BHs (see Reynolds \\& Fabian 2008; McClintock \\& Remillard 2006). The first uses the luminosity and temperature of the optically thick, geometrically thin accretion disc to measure the emitting area of the inner disc, and hence its radius. This assumes that the dissipation follows that of the relativistic stress-free inner boundary condition (Novikov \\& Thorne 1973), an assumption which is now strongly supported by recent fully relativistic MHD simulations of thin discs with self-consistent magnetic turbulence as the origin of stress (Shafee et al. 2008).  To use this method requires that we can observe the spectrum at energies close to the peak temperature, which limits this method to stellar mass BH binaries (hereafter BHBs), as AGN discs typically peak in the unobservable far UV.  Additionally, the source distance, and disc inclination (assumed to be the same as that of the binary) must be constrained, as these are necessary to transform the observed disc flux into luminosity. Similarly, the BH mass is required to convert the resulting size scale into gravitational radii. Thus this technique can only be used on a small subset of systems for which this information is available. Further restrictions are that the systems should be {\\em dominated} by the disc emission. The method can still be applied to data with an increasing fraction of emission in the power law tail, but with increasingly large uncertainties in reconstructing the temperature and luminosity of the disc emission (Kubota et al. 2001; Kubota \\& Done 2004; Done \\& Kubota 2006; Steiner et al. 2009). A final restriction is that the source should not be too bright as the disc structure may change at luminosities approaching and exceeding the Eddington limit $L_{\\rm{Edd}}$. The disc can puff up to become geometrically thick, advection and winds may become important, and the spectrum may also be increasingly distorted by low temperature Comptonization which can be difficult to distinguish from disc emission (e.g. GRS 1915+105: where Middleton et al. 2006 derive $a_*\\sim 0.7$ compared to $a_*=0.98$ from McClintock et al. 2006 due to differences in Comptonisation assumptions).  Within these limitations, the models including the full physics (stress-free inner boundary condition, relativistic smearing, and modelling the non-blackbody intrinsic emission from each radial annulus) are remarkably robust to changes in the disc vertical structure from different {\\em ad hoc} stress prescriptions (Done \\& Davis 2008), and all the potential uncertainties act in the same direction which is for these models to overestimate the black hole spin (Gierli\\'nski \\& Done 2004; Done \\& Davis 2008).  The second method to measure spin uses the shape of the iron line produced by fluorescence in the X-ray illuminated accretion disc.  The width of the line is set by the line emissivity, together with the strength of the gravitational field as this determines the velocity of the disc (hence the Doppler shift, beaming and time dilation) as well as gravitational redshift. All these parameters (inclination, inner disc radius in terms of gravitational radii and emissivity) can be constrained directly from spectral fitting, so this technique can be used much more widely than disc spectral fitting. Its only restrictions are that the disc is flat and in Keplarian rotation (which again becomes increasingly uncertain at luminosities approaching/exceeding Eddington) and that there are sufficient hard X-rays illuminating the disc to produce the iron line. Thus it can be used for both (sub-Eddington) AGNs and BHBs and requires no additional information about the distance and/or inclination of the system (Fabian et al. 1989; Fabian et al. 2000).  However, unlike the disc in the disc dominated spectra, the line is only a small feature on the total spectrum, which means it can be difficult to measure. The line sits on top of a reflected continuum, and the shape of both line and reflected continuum depend on ionisation of the material (Ross, Fabian \\& Young 2001), and the radial and vertical profile of this ionisation (Nayakshin, Kazanzas \\& Kallman 2001; Done \\& Nayakshin 2007). Other key issues are the underlying continuum shape (which can be distorted by complex absorption: L. Miller et al. 2007; 2008) and disentangling the intrinsic shape of the blue wing of the line from any absorption lines from ionised iron K$\\alpha$ (Done \\& Gierli\\'nski 2006; Done \\& Kubota 2006; Young et al. 2005).  Since we have two methods to measure spin it is obviously important to compare them. We would have increased confidence in both methods if they gave the same answer for the same object. However, the very different restrictions on the two techniques mean there is a very small set of objects where both can be used.  Disc dominated spectra generally have too few photons at the high energies required to produce a strong iron fluorescence line. Similarly, disc spectral fitting cannot be used on spectra with a strong tail (carrying more than 25 per cent of the bolometric luminosity) as the uncertainties in reconstructing the intrinsic disc emission become too large (Kubota \\& Done 2004).  Thus the two methods cannot be reliably compared using the same dataset (but see Miller et al. 2009 for an attempt at this), but they can be used on the same object for a BHB with well constrained system parameters which shows spectral transitions.  To date only three objects have good spin estimates from both methods, where 'good' is defined as derived from fits with the best currently available models i.e. {\\sc bhspec} for disc spectral fitting (Davis et al. 2005) and {\\sc CDID} for ionised reflection (Ross \\& Fabian 2005). The results are not encouraging. The spin estimates match well only for XTE~J1550-564, are somewhat discrepant for 4U~1543-475, and are quite significantly different for GRO~J1655-40 (see Section~\\ref{s:compare}).  It is clearly important to expand the sample of objects for which this comparison can be made. GX 339$-$4 is one of the best studied BHBs in terms of iron line profile from three separate data sets, spanning a range of spectral states (Miller et al. 2004; 2006; 2008; Reis et al. 2008). Two of these datasets have issues with pileup, which distorts the line (Done \\& Diaz-Trigo 2010; Yamada et al. 2010), but results from the third dataset alone indicate a very high spin, with $r_{in}\\sim 1.8$ or $2$ depending on the detailed emissivity profile, requiring $a_*=0.96$ or $0.935$ (Reis et al. 2008) and an inclination of $\\sim 20^\\circ$.  Here we apply the disc spectral fitting method and find that, despite the poorly known system parameters, such high spin is unlikely to match the observed temperature and luminosity of the disc dominated spectra from this source. ",
        "conclusions": "We derive a hard upper limit for the spin of GX 339$-$4 of $a_*<0.9$ assuming that the inner disc inclination is the same as that of the binary orbit ($70^\\circ<i<45^\\circ$). This is inconsistent with the spin of $0.942^{+0.005}_{-0.004}$ and inclination of $i=18^{\\circ}\\pm 1^{\\circ}$ derived from the (non-piled up) XMM-Newton burst mode very high state data (Reis et al. 2008). This high spin/low inclination derived from the iron line is already uncomfortably extreme compared to the lower spins predicted from supernovae collapse models, and the small misalignment angles between black hole spin and binary orbit predicted from binary formation models. While these are both potentially poorly understood, they are independent constraints and the iron line profile in GX 339$-$4 conflicts with both of them.  The iron line profile itself is in subtle conflict with the X-ray continuum as the reflection smearing parameters require that the illumination pattern is highly centrally concentrated (Miller et al. 2008; Reis et al. 2008). Yet the very high state spectral shape requires that much of the inner disc is covered by optically thick Comptonizing material, making it very difficult to see strong reflection from this material (Done \\& Kubota 2006).  Thus we argue that the spin/inclination/emissivity derived from the iron line profile are uncomfortably extreme. It seems far more likely to us that the more moderate spin implied by the disc spectral fitting results where the inner disc is more or less aligned with the binary orbit are giving a more robust answer. The inescapable corollary to this is that there are systematic effects affecting spin as derived from the iron line profile that are not yet understood. This is a crucial issue in applying the iron line models with confidence to derive spin in AGNs.  \\label{lastpage}"
    },
    "0911/0911.4973_arXiv.txt": {
        "abstract": "The clustering of matter on cosmological scales is an essential probe for studying the physical origin and composition of our Universe. To date, most of the direct studies have focused on shear-shear weak lensing correlations, but it is also possible to extract the dark matter clustering by combining galaxy-clustering and galaxy-galaxy-lensing measurements. In order to extract the required information, one must relate the observable galaxy distribution to the underlying dark matter distribution. In this study we develop in detail a method that can constrain the dark matter correlation function from galaxy clustering and galaxy-galaxy-lensing measurements, by focusing on the correlation coefficient between the galaxy and matter overdensity fields. Our goal is to develop an estimator that maximally correlates the two. To generate a mock galaxy catalogue for testing purposes, we use the Halo Occupation Distribution approach applied to a large ensemble of $N$-body simulations to model pre-existing SDSS Luminous Red Galaxy sample observations. Using this mock catalogue, we show that a direct comparison between the excess surface mass density measured by lensing and its corresponding galaxy clustering quantity is not optimal. We develop a new statistic that suppresses the small-scale contributions to these observations and show that this new statistic leads to a cross-correlation coefficient that is within a few percent of unity down to $5 \\hMpc$. Furthermore, the residual incoherence between the galaxy and matter fields can be explained using a theoretical model for scale-dependent galaxy bias, giving us a final estimator that is unbiased to within 1\\%, so that we can reconstruct the dark matter clustering power spectrum at this accuracy up to $ k \\sim 1 \\ihMpc$. We also perform a comprehensive study of other physical effects that can affect the analysis, such as redshift space distortions and differences in radial windows between galaxy clustering and weak lensing observations. We apply the method to a range of cosmological models and explicitly show the viability of our new statistic to distinguish between cosmological models. ",
        "introduction": "The current paradigm for the history of our Universe, also known as the $\\Lambda$CDM cosmology, comes along with dark ingredients that have not yet been directly detected: Cold Dark Matter (hereafter CDM) and Dark Energy \\cite{Komatsu2008}. CDM particles constitute about 20\\% of the total energy budget of the Universe, and whilst there is no confirmed direct detection of them in a laboratory experiment, the indirect astrophysical evidence supporting their existence is substantial. However, even more puzzling is the existence and true physical nature of Dark Energy, which contributes roughly 75\\% of the total energy budget of the Universe and is responsible for driving the late-time accelerated expansion of spacetime.  The dark matter power spectrum and its real space equivalent, the correlation function, contain a wealth of cosmological information, e.\\,g.\\ on neutrino mass, dark energy equation of state and the initial conditions of the Universe. Thus it is a key goal of cosmology to infer these quantities from observables. However, to achieve this requires a solid understanding of the galaxy bias -- the relation between the observable galaxies and the underlying dark matter density field. This understanding is especially important for the interpretation of ongoing and upcoming surveys, such as SDSS\\footnote{\\texttt{http://www.sdss.org}}, DES\\footnote{\\texttt{http://www.darkenergysurvey.org}}, PanSTARRS\\footnote{\\texttt{http://pan-starrs.ifa.hawaii.edu/public/}} and EUCLID \\cite{Refregier2008}.  The reconstruction of the CDM distribution is usually based on the assumption, that galaxies trace the matter density field, i.\\,e.\\ that on large scales the galaxy density field equals the matter density field times a parameter known as the bias. The resulting galaxy correlation function can then be expressed as \\begin{equation} \\xi_\\text{gg}(r)=b^2 \\xi_\\text{mm}(r), \\end{equation} and similarly for the power spectrum in $k$-space. The subtlety in the standard approach is that the bias has to be determined empirically, leading to uncertainties in the amplitude of the matter correlation, which finally complicates studies of the rate of change of matter fluctuations with time (the growth factor). Furthermore there is evidence for a non-trivial scale dependence of galaxy bias \\cite{Cole2005,Smith2007,Sanchez2008}. Hence it is of great importance to devise methods that allow a direct reconstruction of the dark matter correlation function from observables. One of the most promising observational probes of dark matter on cosmological scales is the gravitational lensing.  We will focus our attention on a specific weak lensing technique, halo-galaxy lensing, which involves measurement of the shape distortions around foreground dark matter haloes in which galaxies form. Often the foreground object (lens) will be an individual galaxy, in which case this technique is called galaxy-galaxy lensing, but it can also be applied to groups and clusters. Since the first attempts to detect galaxy-galaxy lensing by \\cite{Tyson1984}, the quality of the data has been improved vastly by deeper and wider surveys. Halo-galaxy lensing has now been measured with relatively high signal-to-noise and as a function of a wide variety of properties of the lens galaxies, groups and clusters \\cite{Guzik2002,Hoekstra2003,Seljak2005,Mandelbaum2006} . It has become clear in these studies that galaxy-galaxy lensing contains much information about the mass distribution around galaxies, and has the potential to measure dark matter halo radii, shapes, concentrations and masses \\cite{Mandelbaum2006b,Mandelbaum2008,Evans2009,Okabe2009} as well as the distribution of matter within the Universe \\cite{McKay2001,Sheldon2004,Schneider2005,Sheldon2009a,Sheldon2009b}.  The interpretation of the signal in terms of the link between galaxies and dark matter is, however, complicated by the fact that (except for galaxy clusters) galaxy-galaxy lensing is only detectable by stacking the signal from many lenses. Theoretical modelling of the galaxy-galaxy lensing has been done both with numerical simulations \\cite{Hayashi2007,Hilbert2009} and with the halo model \\cite{Guzik2001,Mandelbaum2005}. The combination of lensing and clustering seems to hold the potential to put constraints on cosmological parameters \\cite{Yoo2006,Cacciato2008}.  In this paper our main objective is to develop a method that recovers a statistic closely related to the matter correlation function from a joint analysis of lensing and clustering observations. This method is presented together with a theoretical motivation and tests on simulated galaxy samples.  The starting point for these simulated galaxy samples are cosmological $N$-body simulations, which are a standard tool to investigate the non-linear evolution of the CDM density field. Despite their statistical power for describing the large scale structure of the Universe, pure dark matter simulations have the disadvantage that one must supplement them with a prescription for galaxy formation in order to reproduce the surveyed galaxy distributions. We work in the standard paradigm of hierarchical galaxy formation: {\\em galaxies only form in dark matter haloes} \\cite{WhiteRees1978}. Hence, the problem is reduced to that of relating galaxies to dark matter haloes, and we do this using the Halo Model approach and in particular the Halo Occupation Distribution (for a review see \\cite{Cooray2002}). Here we are focused on obtaining mock galaxy catalogues for the Luminous Red Galaxies (LRGs), a subset of galaxies observed with the Sloan Digital Sky Survey (SDSS). Our modelling builds on earlier approaches by \\cite{Zheng2008,Reid2008}.  The paper breaks down as follows: in \\S \\ref{sec:obs} we review the basics of weak gravitational lensing, an important probe of the dark matter on cosmological scales. Then in in \\S \\ref{sec:crosscorr} we introduce our main analysis tool, the cross-correlation coefficient.  Theoretical modelling of the latter is carried out in \\S \\ref{sec:theo}. In \\S \\ref{sec:num} we describe the simulations and the mock galaxy catalogues that we use to test our new method. The results of the numerical studies on the cross-correlation coefficient and the reconstructed matter statistic are discussed in \\S \\ref{sec:results}. The effect of redshift space distortions and radial window functions on the observational implementation of our method are explored in \\S \\ref{sec:rs}. \\S \\ref{sec:varcosm} is devoted to the cosmology-dependence of our results. Finally, in \\S \\ref{sec:discuss} we will summarise and discuss our findings. ",
        "conclusions": "\\label{sec:discuss}  In our study, we examined how well can one reconstruct the dark matter clustering from observations of galaxy clustering combined with galaxy-galaxy lensing. This reconstruction procedure could for instance be applied to the SDSS galaxy survey, in particular the Luminous Red Galaxies. In a first step, we generated realistic LRG galaxy catalogues for both a luminosity-threshold and a luminosity bin sub-sample of the LRGs. We then used these galaxy catalogues to extract information about the cross-correlation coefficient between galaxies and matter.  We introduced a new statistic $\\Upsilon(R)$, which we termed the Annular Differential Surface Density (ADSD), that removes the influence of small, non-linear scales on the excess surface mass density. This subtraction is necessary since the scales smaller than the virial radius of the haloes are dominated by the halo profile rather than the pre-shell-crossing evolution of the large scale cosmological fluid. Both numerical studies and theoretical calculations indicate that the cross-correlation coefficient of the ADSD is close to unity and that the residual scale dependence is well described by an analytic correction.  Having focused our investigations on the excess surface mass density $\\Delta\\Sigma_\\text{gm}$, our results can be directly applied to measurements of galaxy-galaxy lensing and the projected galaxy correlation function.  We also studied systematic effects that might bias the comparison of lensing and galaxy clustering measurements. In terms of the projected correlation function, both numerical studies and a linear theory treatment, following \\cite{Kaiser1987}, show that the common integration over $\\pm 50\\hMpc$ along the line of sight is still biased by redshift space distortions. We have shown the necessity of correcting for the large scale peculiar motions in any such clustering measurement. We also investigated the effect of different window functions used for lensing and clustering. These introduce additional effects, that must be accounted for in the final analysis. We have found that the ADSD statistic $\\Upsilon(R)$ is much less sensitive to both of the effects, and to long wavelength modes in general, than the usual projected correlation function $w(R)$, because of the (partially) compensated nature of its transverse window.  As our key result, we devised a method to recover the dark matter clustering from galaxy-galaxy lensing and galaxy clustering measurements using the cross-correlation coefficient for the ADSD. Assuming $r_\\text{cc}=1$ for simplicity leads to at most $8\\%$ bias on scales below $R\\approx 5 \\hMpc$ in the recovered statistic $\\Upsilon_\\text{mm}(R)$. We can however remove this bias based on our theoretical modelling of the scale dependence of the cross-correlation coefficient. The main advantage of our method is that the galaxy dependence is scaled out of the equations, since the theoretical model predictions for the cross-correlation coefficient between haloes and dark matter are relatively independent of the halo mass over a wide range of mass. Thus, we believe that the method devised here is more robust than the methods which are based on HOD fitting (e.\\,g.\\ \\cite{Yoo2009}), which fit for the cosmological parameters and the HOD parameters jointly, marginalising over the uncertainties in the HOD parameters.  A study on four other cosmological models verified the robustness of the new estimator. Varying one parameter of the fiducial $\\Lambda$CDM model at a time, we found that the cross-correlation coefficient shows a scale dependence consistent with the fiducial model. We were able to reconstruct the ADSD of the matter correlation function, and the inferred statistic $\\Upsilon_\\text{mm}$ can be used to distinguish cosmological models both, from the shape and the amplitude of the recovered statistic. This study also showed that, if one is capable of accurately distinguishing central from satellite galaxies and/or remove clusters, then one can eliminate the influence of satellite galaxies and so render the cross-correlation coefficient closer to theoretical predictions for haloes in numerical simulations.  These advantages make it worthwhile to define a clean central galaxy sample and to remove the clusters. While the non-linear matter correlation can be recovered with high fidelity, the linear correlation is only recovered at large scales. This fact strengthens the need for a well developed and tested perturbation theory of large-scale clustering that extends into the weakly non-linear regime and which can thus provide us with an estimator of $\\xi_\\text{NL}$ without having to carry out simulations for each cosmological model.  The ADSD statistic subtracts out a lensing signal at $R_0$ (Eq.~\\ref{eq:Upsilon}). This subtraction procedure decreases the signal-to-noise on the inferred statistic $\\Upsilon$ as compared to $\\Delta\\Sigma$. This price seems worth paying, since it brings the cross-correlation coefficient much closer to unity with residual deviations from unity that are well understood theoretically. An application of this method to observational data will have to address the problem of estimating $\\Delta\\Sigma(R_0)$. The cubic spline fit used for our numerical studies will not be appropriate given the large statistical fluctuations in observed lensing signal. Several alternatives to estimate $\\Delta\\Sigma(R_0)$ are explored in \\cite{Mandelbaum2009}, with the most successful being a fit with a running power-law (three parameters) to the radial bins around $R_0$.   Our numerical results are based on the SDSS spectroscopic LRG sample, i.\\,e.\\ on the galaxies living in the most massive haloes. Based on the success and generality of the theoretical model we expect that a similar behaviour for the Main spectroscopic sample galaxies in the SDSS, which live predominantly in lower mass haloes. The lower halo masses may also enable a lower cutoff radius $R_0$, especially if haloes with higher mass are effectively removed from the sample. \\tr{If the halo sample spans a wide range of masses, it should be split into mass bins and $R_0$ should be chosen appropriately for each of the mass bins.} We shall reserve this topic for future investigation."
    },
    "0911/0911.4718_arXiv.txt": {
        "abstract": "The spin and quadrupole moment of the supermassive black hole at the Galactic center can in principle be measured via astrometric monitoring of stars orbiting at milliparsec (mpc) distances, allowing tests of general relativistic ``no-hair''theorems \\cite{Will-08}. One complicating factor is the presence of perturbations from other stars, which may induce orbital precession of the same order of magnitude as that due to general relativistic effects. The expected number of stars in this region is small enough that full $N$-body simulations can be carried out. We present the results of a comprehensive set of such simulations, which include a post-Newtonian treatment of spin-orbit effects. A number of possible models for the distribution of stars and stellar remnants are considered. We find that stellar perturbations are likely to obscure the signal due to frame-dragging for stars beyond $\\sim 0.5$ mpc from the black hole, while measurement of the quadrupole moment is likely to require observation of stars inside $\\sim 0.2$ mpc. A high fraction of stellar remnants, e.g. $10\\msun$ black  holes, in this region would make tests of GR problematic at all radii. We discuss the possibility of separating the effects of stellar perturbations from those due to GR. ",
        "introduction": "Introduction}  The supermassive black hole (SBH) at the center of the Milky Way galaxy is surrounded by a compact cluster of stars that has been the target of observational surveys for more than a decade \\cite{Krabbe-95,Blum-96,Figer-00,Gezari-02,Paumard-06,Zhu-08}. Near-infrared monitoring of stellar positions using adaptive optics techniques has allowed orbital reconstruction for roughly 30 stars at distances ranging from $10^0-10^2$ milliparsecs (mpc) from the SBH \\cite{Ghez-05,Ghez-08,Gillessen-09a}. One of these stars (S2) has an orbital period of only $\\sim 15$ yr \\cite{Schoedel-02,Ghez-03} and its orbit has been followed for more than one full revolution; astrometric data for S2 yield a well-constrained mass for the SBH, $\\mh = (3.95\\pm 0.06)\\times 10^6\\msun$ (assuming a galactocentric distance of $8.0$ kpc) and a location on the plane of the sky that is consistent with that of the radio source Sgr A$^*$ \\cite{Reid-93,Eisenhauer-05,Ghez-08,Groen-08,Gillessen-09a,Gillessen-09b}.  The velocity of S2 near periapse is a few percent of the speed of light, large enough that relativistic effects like advance of the periastron become potentially measurable, even on time scales as short as a few years \\cite{Jaro-98,Fragile-00,Weinberg-05,Zucker-06}. No such effects have so far been unambiguously observed \\cite{Gillessen-09b}; one complicating factor is the likely presence of a distributed mass (stars, stellar remnants, dark matter etc.) within S2's orbit which could produce Newtonian precession of the same order of magnitude as that due to general relativity \\cite{Rubilar-01,Zakharov-07}.  If the SBH is rotating, new phenomena occur for stars orbiting at very small separations, $r\\lap 1$ mpc. Dragging of inertial frames and torques from the SBH's quadrupole moment $Q$ cause stellar orbital planes to precess, at rates that depend respectively on the first and second powers of the hole's spin angular momentum $\\mathbf{J}$. These spin-related effects are small compared with in-plane precession, but (in the absence of other non-spherically-symmetric components of the gravitational potential) they contain unambiguous information about $\\mathbf{J}$ and $Q$. In principle, observed changes in the orbital orientations of just two stars would be sufficient to independently constrain the four quantities $(\\mathbf{J}, Q)$, allowing tests of general relativistic ``no-hair'' theorems \\cite{Will-08}. The amplitude of these spin-related precessions is very small, of order microarcseconds ($\\mu$as) per year as seen from the Earth. Plans are being developed to achieve infrared astrometry at this level \\cite{GRAVITY-09,ASTRA-08}.  Such measurements will require the presence of at least a few bright stars on mpc-scale orbits around the SBH. While no such stars have yet been observed, extrapolation of the observed stellar densities at distances of $\\sim 1$ pc from the SBH suggests that of order $10^0-10^2$ stars should be present in this region. Due to their finite numbers, these stars will generate a non-spherically-symmetric component to the gravitational potential with an amplitude that scales as $\\sqrt{N}\\ms$, where $\\ms$ is the mass of a typical star and $N$ is their number. Simple arguments (\\S\\ref{sec:timescales}) suggest that such stellar perturbations might produce changes in the orbital orientations of test stars that are comparable in magnitude to the spin-related effects. This would complicate the testing of no-hair theorems by adding what is effectively a source of noise to the measured precessions.  Because the expected number of stars in this region is so small, direct $N$-body integration of the equations of motion for all $N$ stars is feasible. The major technical requirements are a high degree of accuracy in the $N$-body integrator and the inclusion of terms describing the relativistic accelerations due to the SBH, including spinless, spin-orbit, and quadrupole-orbit contributions.  Here we present the results of a comprehensive set of such simulations. Our primary goal is to evaluate the degree to which star-star perturbations might obscure the signal due to the SBH's spin; hence we focus on changes in orbital orientations rather than on the evolution of the phase-space variables ($\\mathbf{r},\\mathbf{v}$) \\cite{Kannan-09,Preto-09}. We ignore all other systematic effects that might limit the ability to carry out the high-precision astrometry for stars in crowded fiels at the Galactic center \\cite{Weinberg-05,Fritz-09}.  In \\S\\ref{sec:timescales} we summarize the relevant time scales for orbital evolution near the Milky Way SBH. \\S\\ref{sec:nbody} presents the post-Newtonian $N$-body equations of motion including the lowest-order spin-orbit terms and describes the $N$-body integrator. Observational and theoretical constraints on the distribution of stars and stellar remnants near the Galactic center SBH are summarized in \\S\\ref{sec:models}, which also describes the parametrized models used to construct the $N$-body initial conditions. \\S\\ref{sec:results} summarizes the results from the integrations, including estimates of the number of stars that can be effectively used to measure $\\mathbf{J}$ and $Q$. \\S\\ref{sec:discuss} discusses how the presence of stellar perturbations in the astrometric data can potentially be detected and removed from the GR signal. \\S\\ref{sec:conclude} sums up. ",
        "conclusions": "Conclusions}  \\noindent 1. The spin and quadrupole moment of the Galactic center supermassive black hole (SBH) can in principle be measured by observing the precession of the orbital planes of stars in the inner milliparsec (mpc) \\cite{Will-08}. However, gravitational interactions between stars in this region are likely to induce orbital precession of the same approximate amplitude as the precession due to frame dragging. The stellar perturbations manifest themselves as coherent torques over short time scales, mimicking general relativistic (GR) precession.  \\noindent 2. The number of stars and stellar remnants (e.g. stellar-mass black holes, BHs) in this region is uncertain, but small enough ($\\sim 10^0 - 10^3$) that full $N$-body simulations are feasible. A regularized post-Newtonian $N$-body algorithm is presented that includes the lowest-order spin-orbit and quadrupole-orbit terms.  \\noindent 3. Assuming near-maximal spin for the Milky Way SBH, detection of frame-dragging precession may be feasible after a few years' monitoring with an instrument like GRAVITY \\cite{GRAVITY-09} for orbits in the radial range between $\\sim 0.2$ mpc and $\\sim 1$ mpc. At smaller radii the number of stars is too small, while at larger radii the star-star and star-remnant perturbations dominate GR effects. In models where the number of stellar BHs is comparable to the number of observable stars, GR effects are almost always swamped by perturbations from the remnants.  \\noindent 4. Quadrupole-induced precession stands out clearly from stellar perturbations only in a narrow class of models for the nuclear star cluster, having moderate to high central densities and a small BH fraction, and only at radii $r\\lap 0.2$ mpc.  \\noindent 5. Because the orbit-averaged torques from stars are approximately constant in magnitude over year-long time scales, it is possible in principle to disentangle the effects of stellar perturbations from those due to GR, allowing tests of gravity even in the presence of stellar perturbations."
    },
    "0911/0911.3138_arXiv.txt": {
        "abstract": "We discuss the properties of an analytical solution for waves in radiating fluids, with a view towards its implementation as a quantitative test of radiation hydrodynamics codes. A homogeneous radiating fluid in local thermodynamic equilibrium is periodically driven at the boundary of a one-dimensional domain, and the solution describes the propagation of the waves thus excited. Two modes are excited for a given driving frequency, generally referred to as a radiative acoustic wave and a radiative diffusion wave. While the analytical solution is well known, several features are highlighted here that require care during its numerical implementation. We compare the solution in a wide range of parameter space to a numerical integration with a Lagrangian radiation hydrodynamics code. Our most significant observation is that flux-limited diffusion does not preserve causality for waves on a homogeneous background. ",
        "introduction": "\\label{}  Analytical solutions for radiation hydrodynamics are difficult to obtain due to the complexity of the equations but are a powerful tool for testing complex multi-physics codes. One of the most useful simplifying assumptions for any set of nonlinear differential equations is to set up an equilibrium state and analyze small departures from that equilibrium. Since this approach retains most of the terms in the equations, it not only provides valuable physical insight but also serves as a comprehensive and sensitive test of numerical algorithms. Many perturbation studies of radiation hydrodynamics have been performed; we follow closely Mihalas \\& Mihalas \\cite{mm83, mm84} and Bogdan et al. \\cite{bog96}, and refer the reader there for additional references. Despite the straightforward application of perturbation theory to the equations of radiation hydrodynamics, however, there appear to be few numerical tests of this type of solution in the literature (reference \\cite{ts01} is one example). Our goal here is to conduct a systematic comparison of such a solution with a numerical algorithm in a wide range of parameter space. The code that we use for comparison is the Lagrangian radiation hydrodynamics code Kull \\cite{kull00}. We begin in \\S\\ref{AE} with an overview of our assumptions and the form the equations of radiation hydrodynamics take under these assumptions. The failure of flux-limited diffusion to capture free-streaming radiation waves is highlighted in \\S\\ref{FLD}. We discuss the analytical solution in \\S\\ref{AS} and compare our results with previous work in \\S\\ref{PW}. Numerical results are given in \\S\\ref{NR} and we summarize in \\S\\ref{SD}. ",
        "conclusions": "\\label{SD}  We have conducted a systematic comparison of a perturbation analysis of the equations of radiation hydrodynamics with the Lagrangian code Kull in a wide range of parameter space. We have demonstrated that these solutions are a useful benchmark for testing any radiation hydrodynamics code. The most important issues to keep in mind when conducting such a test are 1) flux-limited diffusion does not capture these solutions in the free-streaming limit, 2) coupling to both modes can occur, making comparison with a single mode somewhat difficult and 3) care must be taken in setting the perturbation amplitude, since either the radiation or hydrodynamic perturbations can dominate the others by orders of magnitude.  The primary purpose of a numerical test like the one we have studied is to investigate the convergence properties of a numerical discretization. Since our focus has been on the test itself rather than on the convergence properties of Kull, and due to the complications of transient effects and mode coupling, we have not included any convergence results here. We have investigated the convergence properties of Kull for most of the solutions shown in Figures~\\ref{fa}-\\ref{fF} by calculating the L2 norm of the error between the numerical and analytical results (excluding the final wavelength to minimize transient effects). Kull converges at second order in most regions of parameter space, as expected, particularly when the hydrodynamics and radiation are weakly coupled. In regions of parameter space where the coupling between modes is strong, the convergence is weaker. A formal convergence test in Kull using the solutions described here would require the implementation of outflow boundary conditions.  Some additional comments on the breakdown of flux-limited diffusion are in order. Figures~\\ref{fgs1} and \\ref{fgs2} are somewhat of an academic exercise, since the radiation energy in this region of parameter space exceeds the rest mass energy of the material, and our non-relativistic treatment breaks down \\cite{bog96}. In addition, the superluminal propagation we have observed only occurs for waves whose amplitude is small compared to the background state, and it is not clear that this would have a significant impact on the energy budget of a realistic calculation. Finally, we have set up our computational domain to allow the excitation of any length and time scales we choose, and even with this freedom it was computationally expensive to access the regions of parameter space in which the flux limiters break down. A realistic calculation will only be able to resolve a small fraction of the length and time scales that we have explored, and only with a very large dynamic range will a calculation be able to see the excitation of superluminal waves. In any case, when computation has advanced to the point where one can easily resolve the length and time scales of free-streaming radiation, one should no longer be using flux-limited diffusion.  This work performed under the auspices of the U.S. Department of Energy by Lawrence Livermore National Laboratory under Contract DE-AC52-07NA27344.  \\newpage  \\begin{figure} \\psfrag{r}{$\\beta^{-1}$} \\centering \\includegraphics[width=6.0in]{acoustic_lowres.eps} \\caption{Parameter space for the radiative acoustic mode under the Eddington approximation. The phase speeds are the material sound speed (regions a and b), the isothermal sound speed (regions c-e), the radiative sound speed (region f) and $c/\\sqrt{3}$ (region g).} \\label{fra} \\end{figure}  \\begin{figure} \\centering \\includegraphics[width=6.0in]{diffusive_lowres.eps} \\caption{Parameter space for the radiative diffusion mode under the Eddington approximation. The phase speeds are $v_p \\ll a$ in the white region, $a < v_p < c$ in the light shaded region and $v_p = c\\sqrt{3}$ in the dark shaded region.} \\label{frd} \\end{figure}  \\begin{figure} \\psfrag{y}[][][1.5]{$\\frac{\\delta \\rho}{\\rho_0}$} \\psfrag{x}[][][1.5]{$\\frac{x}{\\lambda}$} \\centering \\includegraphics[width=6.0in]{region_a_acoustic.eps} \\caption{Density perturbation in region a ($r = 10^{-3}$, $\\tau_a = 10^4$) after ten wave periods. The points are the Kull results, and the solid line is the analytical solution.} \\label{fa} \\end{figure}  \\begin{figure} \\psfrag{y}[][][1.5]{$\\frac{\\delta \\rho}{\\rho_0}$} \\psfrag{x}[][][1.5]{$\\frac{x}{\\lambda}$} \\centering \\includegraphics[width=6.0in]{region_b_acoustic.eps} \\caption{Density perturbation in region b ($r = 10^{-5}$, $\\tau_a = 10^{-2}$) after ten wave periods. The points are the Kull results, and the solid line is the analytical solution.} \\label{fb} \\end{figure}  \\begin{figure} \\psfrag{y}[][][1.5]{$\\frac{\\delta \\rho}{\\rho_0}$} \\psfrag{x}[][][1.5]{$\\frac{x}{\\lambda}$} \\centering \\includegraphics[width=6.0in]{region_c_acoustic.eps} \\caption{Density perturbation in region c ($r = 10^{-3}$, $\\tau_a = 10$) after ten wave periods. The points are the Kull results, and the solid line is the analytical solution.} \\label{fc} \\end{figure}  \\begin{figure} \\psfrag{y}[][][1.5]{$\\frac{\\delta \\rho}{\\rho_0}$} \\psfrag{x}[][][1.5]{$\\frac{x}{\\lambda}$} \\centering \\includegraphics[width=6.0in]{region_d_acoustic.eps} \\caption{Density perturbation in region d ($r = 10^{-1}$, $\\tau_a = 10^{-2}$) after ten wave periods. The points are the Kull results, and the solid line is the analytical solution.} \\label{fd} \\end{figure}  \\begin{figure} \\psfrag{y}[][][1.5]{$\\frac{\\delta \\rho}{\\rho_0}$} \\psfrag{x}[][][1.5]{$\\frac{x}{\\lambda}$} \\centering \\includegraphics[width=6.0in]{region_e_acoustic.eps} \\caption{Density perturbation in region e ($r = 10^3$, $\\tau_a = 10^{-2}$) after ten wave periods. The points are the Kull results, and the solid line is the analytical solution.} \\label{fe} \\end{figure}  \\begin{figure} \\psfrag{y}[][][1.5]{$\\frac{\\delta \\rho}{\\rho_0}$} \\psfrag{x}[][][1.5]{$\\frac{x}{\\lambda}$} \\centering \\includegraphics[width=6.0in]{region_f_acoustic.eps} \\caption{Density perturbation in region f ($r = 10$, $\\tau_a = 10^4$) after ten wave periods. The points are the Kull results, and the solid line is the analytical solution.} \\label{ff} \\end{figure}  \\begin{figure} \\psfrag{y}[][][1.5]{$\\frac{\\delta T_r}{T_0}$} \\psfrag{x}[][][1.5]{$\\frac{x}{\\lambda}$} \\centering \\includegraphics[width=6.0in]{region_A_diff.eps} \\caption{Radiation temperature perturbation in region A ($r = 10^{-3}$, $\\tau_a = 10^4$) after ten wave periods. The points are the Kull results, and the solid line is the analytical solution. Only one-tenth of the computational domain is shown.} \\label{fA} \\end{figure}  \\begin{figure} \\psfrag{y}[][][1.5]{$\\frac{\\delta T_r}{T_0}$} \\psfrag{x}[][][1.5]{$\\frac{x}{\\lambda}$} \\centering \\includegraphics[width=6.0in]{region_B_diff.eps} \\caption{Radiation temperature perturbation in region B ($r = 10^{-3}$, $\\tau_a = 1$) after ten wave periods. The points are the Kull results, and the solid line is the analytical solution. Only one-tenth of the computational domain is shown.} \\label{fB} \\end{figure}  \\begin{figure} \\psfrag{y}[][][1.5]{$\\frac{\\delta T_r}{T_0}$} \\psfrag{x}[][][1.5]{$\\frac{x}{\\lambda}$} \\centering \\includegraphics[width=6.0in]{region_C_highres.eps} \\caption{Radiation temperature perturbation in region C ($r = 10^{-3}$, $\\tau_a = 10^4$) after two wave periods. The dotted line is the Kull result, and the solid line is the analytical solution.} \\label{fC} \\end{figure}  \\begin{figure} \\psfrag{y}[][][1.5]{$\\frac{\\delta T_r}{T_0}$} \\psfrag{x}[][][1.5]{$\\frac{x}{\\lambda}$} \\centering \\includegraphics[width=6.0in]{region_D_diff.eps} \\caption{Radiation temperature perturbation in region D ($r = 10^2$, $\\tau_a = 10^3$) after ten wave periods. The points are the Kull results, and the solid line is the analytical solution. Only one-tenth of the computational domain is shown.} \\label{fD} \\end{figure}  \\begin{figure} \\psfrag{y}[][][1.5]{$\\frac{\\delta T_r}{T_0}$} \\psfrag{x}[][][1.5]{$\\frac{x}{\\lambda}$} \\centering \\includegraphics[width=6.0in]{region_E_diff.eps} \\caption{Radiation temperature perturbation in region E ($r = 3 \\times 10^{-7}$, $\\tau_a = 3$) after one wave period. The points are the Kull results, and the solid line is the analytical solution.} \\label{fE} \\end{figure}  \\begin{figure} \\psfrag{y}[][][1.5]{$\\frac{\\delta T_r}{T_0}$} \\psfrag{x}[][][1.5]{$\\frac{x}{\\lambda}$} \\centering \\includegraphics[width=6.0in]{region_E_coupling.eps} \\caption{Radiation temperature perturbation in region E after ten wave periods. The solid lines are Kull results with resolution increasing from top to bottom, the dotted line is the analytical solution for the diffusion mode, and the dashed line is the analytical solution for the acoustic mode.} \\label{fEcoupling} \\end{figure}  \\begin{figure} \\psfrag{y}[][][1.5]{$\\frac{\\delta T_r}{T_0}$} \\psfrag{x}[][][1.5]{$\\frac{x}{\\lambda}$} \\centering \\includegraphics[width=6.0in]{region_F_diff.eps} \\caption{Radiation temperature perturbation in region F ($r = 3 \\times 10^{-6}$, $\\tau_a = 3 \\times 10^{-1}$) after one half of a wave period. The points are the Kull results, and the solid line is the analytical solution.} \\label{fF} \\end{figure}  \\begin{figure} \\psfrag{y}[][][1.5]{$\\frac{\\delta T_r}{T_0}$} \\psfrag{x}[][][1.5]{$\\frac{x}{\\lambda}$} \\centering \\includegraphics[width=6.0in]{region_g_superluminal1.eps} \\caption{Radiation temperature perturbation in region g ($r = 10^3$, $\\tau_a = 1, a = 0.1 c$) after one wave period. The solid lines are the Kull results (both with and without a flux limiter; differences are $O[10^{-11}]$), and the dotted line is the analytical solution under the diffusion approximation. For reference purposes, the dashed line shows a wave at the same driving frequency with a phase velocity equal to the speed of light.} \\label{fgs1} \\end{figure}  \\begin{figure} \\psfrag{y}[][][1.5]{$\\frac{\\delta T_r}{T_0}$} \\psfrag{x}[][][1.5]{$\\frac{x}{\\lambda}$} \\centering \\includegraphics[width=6.0in]{region_g_superluminal2.eps} \\caption{Same as Figure~\\ref{fgs1} with a higher perturbation amplitude. The differences between the numerical results with and without a flux limiter are $O(10^{-8})$.} \\label{fgs2} \\end{figure}  \\begin{figure} \\psfrag{y}[][][1.5]{$\\frac{\\delta T_r}{T_0}$} \\psfrag{x}[][][1.5]{$\\frac{x}{\\lambda}$} \\centering \\includegraphics[width=6.0in]{region_H_superluminal1.eps} \\caption{Radiation temperature perturbation in region H ($r = 10^2$, $\\tau_a = 10^{-6}$) after one wave period. The solid lines are the Kull results (both with and without a flux limiter; differences are $O[10^{-5}]$), and the dotted line is the analytical solution under the diffusion approximation. For reference purposes, the dashed line shows a wave at the same driving frequency with a phase velocity equal to the speed of light.} \\label{fHs1} \\end{figure}  \\clearpage  \\begin{figure} \\psfrag{y}[][][1.5]{$\\frac{\\delta T_r}{T_0}$} \\psfrag{x}[][][1.5]{$\\frac{x}{\\lambda}$} \\centering \\includegraphics[width=6.0in]{region_H_superluminal2.eps} \\caption{Same as Figure~\\ref{fHs1} with a higher perturbation amplitude. The differences between the numerical results with and without a flux limiter are $O(10^{-2})$.} \\label{fHs2} \\end{figure}  \\appendix"
    },
    "0911/0911.4004_arXiv.txt": {
        "abstract": " ",
        "introduction": "The observation of variable, non-thermal high emission from Active Galactic Nuclei (AGN) reveals that efficient particle acceleration can take place on different length scales. It is widely believed, for example, that diffusive shock acceleration of electrons can produce the power-law particle distributions that are needed to account for the observed nuclear synchrotron and inverse Compton emission features in AGN jets. While efficient electron acceleration is in most cases strongly limited by radiative losses, this is much less the case for protons and heavier nuclei, suggesting that these particles could reach much higher energy via the same acceleration process. Motivated by the indication of a possible correlation between the Pierre Auger (PAO)-measured ultra-high energy cosmic ray (UHECR) events and the nearby AGN distribution \\cite{abr07,rou09,hag09}, this contribution analyzes the conditions under which efficient cosmic ray acceleration to UHECR energies may become possible. Particular attention is given to the radio galaxy Cen~A, which, based on its proximity, could represent a promising UHECR source candidate, e.g., \\cite{rom96,gur08,kac09}.\\\\ ",
        "conclusions": "The above analysis suggests that efficient acceleration of protons to UHECR energies in Cen~A  is challenging and may require the operation of an additional acceleration mechanism like shear to further boost achievable particle energies beyond $E_c = 5 \\times 10^{19}$ eV. Efficient shear acceleration in Cen~A would require high energy seed particles which, however, could be provided by, e.g., shock acceleration. A fraction of these seed protons may then be picked up and accelerated to the maximum energy given by the confinement limit. If such a two-step process would indeed take place, spectral changes in the cosmic ray energy spectrum may not just simply be due to propagation effects. The situation is much more relaxed for heavier elements like iron nuclei, which could possibly be directly accelerated (either by shocks or within the black hole magnetosphere) to energies of $E_c$ and beyond. If Cen~A would indeed be an efficient UHECR accelerator one may thus expect the composition to become heavier above energies $\\sim 10^{19}$ eV.\\\\"
    },
    "0911/0911.3909_arXiv.txt": {
        "abstract": "Faraday rotation measurements have provided an invaluable technique with which to measure the properties of astrophysical magnetized plasmas.  Unfortunately, typical observations provide information only about the density-weighted average of the magnetic field component parallel to the line of sight.  As a result, the magnetic field geometry along the line of sight, and in many cases even the location of the rotating material, is poorly constrained. Frequently, interpretations of Faraday rotation observations are dependent upon underlying models of the magnetic field being probed (e.g., uniform, turbulent, equipartition). However, we show that at sufficiently low frequencies, specifically below roughly $13(\\RM/1\\,\\rad\\,\\m^{-2})^{1/4}(B/1\\,\\G)^{1/2}\\,\\MHz$, the character of Faraday rotation changes, entering what we term the ``super-adiabatic regime'' in which the rotation measure is proportional to the integrated {\\em absolute value} of the line-of-sight component of the field.  As a consequence, comparing rotation measures at high frequencies with those in this new regime provides direct information about the geometry of the magnetic field along the line of sight.  Furthermore, the frequency defining the transition to this new regime, $\\nuSA$, depends directly upon the {\\em local} electron density and magnetic field strength where the magnetic field is perpendicular to the line of sight, allowing the unambiguous distinction between Faraday rotation within and in front of the emission region.  Typical values of $\\nuSA$ range from $10\\,\\kHz$ (below the ionospheric cutoff, but above the heliospheric cutoff) to $10\\,\\GHz$, depending upon the details of the Faraday rotating environment.  In particular, for resolved AGN, including the black holes at the center of the Milky Way (Sgr A*) and M81, $\\nuSA$ ranges from roughly $10\\,\\MHz$ to $10\\,\\GHz$, and thus can be probed via existing and up-coming ground-based radio observatories. ",
        "introduction": "\\label{I} Magnetized plasmas are common in astrophysics, playing central roles in objects as diverse as galaxy clusters, the sites of star formation, the solar corona and accreting black holes and their ultra-relativistic outflows.  Despite their importance, however, there are few observational tools for assessing the physical conditions within them.  This is especially true for strongly ionized or diffuse regions, in which molecular line emission is absent.  In these the best evidence for significant magnetic fields, beyond theoretical arguments, is due to Faraday rotation observations of intrinsically polarized sources.  First discovered in optically active crystals \\citep{Fara:1846}, Faraday rotation has subsequently become one of the primary methods for measuring the strength and geometry of astrophysical magnetic fields.  This has been primarily via the determination of the rotation measure, defined in terms of the polarization angle, $\\Phi$, by \\begin{equation} \\RM = \\frac{\\d \\Phi}{\\d \\lambda^2} \\simeq 8.12\\times10^5\\int n_e \\bmath{B}\\cdot\\d\\bmath{\\ell}\\,\\,\\rad\\,\\m^{-2}\\,, \\label{eq:FR} \\end{equation} where $n_e$, $B$ and $\\ell$ are measured in $\\cm^{-3}$, $\\G$ and $\\pc$, respectively.  Thus, $\\RM$'s provides a line-of-sight averaged measurement of the line-of-sight magnetic field, weighted by the local plasma density.  Measured values for the $\\RM$ range from nearly zero to nearly $10^6\\,\\rad\\,\\m^{-2}$.  Rotation measure observations have played a critical role in determining the magnetic field strengths in the centers of galaxy clusters \\citep[e.g.,][]{Ge-Owen:93,Ge-Owen:94}, Galactic magnetic field strength and distribution \\citep[e.g.,][]{Men-Ferr-Han:08,Nout_etal:08}, magnetic field strength and structure within ultra-relativistic black hole jets \\citep{Zava-Tayl:04,Stir-Spec-Cawt-Para:04,Mill_etal:05,Broc_etal:07,Khar_etal:09}, magnetic field strengths near, and accretion rate of, the supermassive black hole at the center of the Milky Way and other nearby low-luminosity active galactic nuclei (AGN) \\citep[e.g.,][]{Agol:00,Quat-Gruz:00,Brun-Bowe-Falc:06,Macq_etal:06,Marr_etal:07}, and even the solar corona magnetic field \\citep[e.g.,][]{Manc-Span:00}.  However, there are a number of significant ambiguities in Faraday rotation studies.  The first is the degeneracy between density and magnetic field strength (i.e., high densities and weak fields vs. low densities and strong fields).  The second is the degeneracy between weak large-scale fields and strong tangled fields.  In light of numerical simulations of magnetohydrodynamic turbulence, and the observation of supersonic (though sub-Alfv\\'enic) turbulence in Galactic molecular clouds, these make interpreting measured $\\RM$'s highly dependent upon models for the rotating medium.  Since in many cases it is unclear if the Faraday rotation is occuring {\\em in situ} or in a foreground region distant from the polarized source, there are large uncertainties in the our understanding of the physical conditions in the source.  Here we show that at low frequencies Faraday rotation qualitatively changes, with the $\\RM$ becoming proportional to $\\int n_e|\\bmath{B}\\cdot\\d\\bmath{\\ell}|$.  Comparisons between $\\RM$'s measured within this ``super-adiabatic'' regime and at high frequencies can probe the line-of-sight geometry of the magnetic field\\footnote{Analogous phenomena include the ``MSW'' effect in neutrino astrophysics \\citep{Bahc:89}}.  Furthermore, the frequency at which a given source transitions between the standard and super-adiabatic regimes depends upon the {\\em local} plasma properties at magnetic field reversals (when the magnetic field is orthogonal to the line of sight), and can thus distinguish between different candidates for the site of the rotating medium.  In Section \\ref{sec:FRaAPMP} we review the plasma processes responsible for Faraday rotation and describe the super-adiabatic regime.  Section \\ref{sec:O} deals with the observational consequences unique to super-adiabatic Faraday rotation generally, while Sections \\ref{sec:IS}  \\& \\ref{sec:ISS} concern specific classes of sources in which ordinary Faraday rotation has been observed.  Our conclusions and implications are summarized in Section \\ref{sec:C}. The mathematical details of an {\\em ab initio} determination of the radiative transfer regimes is presented in the Appendix. ",
        "conclusions": "\\label{sec:C} Despite commonly being described in terms of the independent propagation of plasma modes, non-adiabatic effects at magnetic field reversals force the mode crossings that are critical to understanding the standard expressions for Faraday rotation.  For nearly all astrophysical systems, at frequencies above a few $\\GHz$, these effects dominate the propagation at magnetic field reversals. However, at sufficiently low frequencies it is possible to enter a ``super-adiabatic'' regime, in which the two plasma modes propagate independently throughout, resulting in rotation measures that are proportional to $\\int n_e |B_\\parallel| \\d\\ell$ instead of $\\int n_e B_\\parallel \\d\\ell$.  For large-scale ordered magnetic fields, there is little difference between the standard and super-adiabatic regimes.  However, for magnetic field geometries exhibiting many reversals along the line of sight, the rotation measure in the super-adiabatic regime can be substantially larger than that determined at high frequencies.  Thus, comparisons between the standard and super-adiabatic rotation provides a unique means to measure the magnetic field geometry along the line of sight.  The frequency at which Faraday rotation transitions between the standard and super-adiabatic regimes, $\\nuSA$, is itself dependent upon the local plasma properties at the magnetic field reversals, though only weakly dependent upon the magnetic field geometry.  Thus even the detection of super-adiabatic Faraday rotation, marked by a rapid increase in the rotation measure at low frequencies, provides localized information about the plasma density and magnetic fields.  This has the potential to remove the degeneracy between propagation path length and plasma density or magnetic field strength, allowing an unambiguous determination of the location of the Faraday screen in many systems.   For known systems, $\\nuSA$ ranges from the heliospheric cutoff, approximately $10\\,\\kHz$, to nearly $10\\,\\GHz$, with regions containing stronger magnetic fields and higher rotation measures transitioning at higher frequencies.  For resolved active galactic nuclei, such as Sgr A*, M81* and M31, this lies above the ionospheric cutoff at $15\\,\\GHz$, and hence is observable from ground based radio telescopes.  Some unresolved AGN, \\HzRG~and Pulsars may be observable from the ground as well.  Detection of super-adiabatic Faraday rotation in any of these sources would provide striking confirmation of the presence of in situ Faraday rotation. For X-ray binaries, infrared galaxies, molecular clouds, the Galactic center region, intracluster medium, interstellar medium and the solar corona, the super-adiabatic regime should be observable from space-based radio observatories.  The search for super-adiabatic Faraday rotation is especially timely given the proliferation of ultra-low frequency telescopes, such as the {\\em Murchison Widefield Array} in Australia ($80\\,\\GHz$--$300\\,\\GHz$) and the {\\em Low Frequency Array} in Europe ($10\\,\\MHz$--$240\\,\\MHz$), built ostensibly for the purpose of mapping $21\\,\\cm$ emission from the high-$z$ Universe, and the {\\em Frequency Agile Solar Radiotelescope} in North America ($50\\,\\MHz$--$21\\,\\GHz$), currently under development \\citep{Bast:03}. All of these devices will collect the full compliment of Stokes parameters, and thus will be capable of searching for super-adiabatic Faraday rotation.  In the far term, space-based radio observatories, such as {\\em Astronomical Low Frequency Array}, have been discussed for more than two decades now \\citep{Jone_etal:00}. They provide the only means to access the $10\\,\\kHz$--$15\\,\\MHz$ radio window, and beyond the confines of the Earth, can in principle do so with resolutions comparable to much shorter wavelengths at the {\\em Very Long Baseline Array}.  The ability to directly probe the geometry of astrophysical magnetic fields along the line of sight should add to the list of science motivators for such a capability."
    },
    "0911/0911.2002_arXiv.txt": {
        "abstract": "We report the discovery of a supernova (SN) with the highest apparent energy output to date and conclude that it represents an extreme example of the Type IIn subclass.  The SN, which was discovered behind the Large Magellanic Cloud at $z = 0.289$ by the \\mbox{SuperMACHO} microlensing survey, peaked at $M_{R} = -21.5$~mag and only declined by 2.9~mag over 4.7~years after the peak. Over this period, \\LMC902 had an integrated bolometric luminosity of $4\\times10^{51}$~ergs, more than any other SN to date.  The radiated energy is close to the limit allowed by conventional core-collapse explosions.  Optical spectra reveal that \\LMC902 has persistent single-peaked intermediate-width hydrogen lines, a signature of interaction between the SN and a dense circumstellar medium.  The light curves show further evidence for circumstellar interaction, including a long plateau with a shape very similar to the classic SN~IIn~1988Z -- however, \\LMC902 is ten times more luminous at all epochs.  The fast velocity measured for the intermediate-width H$\\alpha$ component (\\about 6000~\\kms) points towards an extremely energetic explosion ($> 10^{52}$~ergs), which imparts a faster blast-wave speed to the post-shock material and a higher luminosity from the interaction than is observed in typical SNe~IIn. Mid-infrared observations of \\LMC902 suggest an infrared light echo is produced by normal interstellar dust at a distance \\about 0.5~pc from the SN. ",
        "introduction": "\\label{sec:intro}  Hydrogen-rich (Type II) core-collapse supernovae (SNe) from massive stars typically have peak luminosities corresponding to absolute magnitudes $-14 > M_{R} > -18$~mag that are powered by thermal energy that is deposited into the expanding SN envelope during shock breakout \\citep{Li10}.  Depending on the mass of the hydrogen (H) envelope, a plateau can arise in the optical light curve from a cooling wave of H recombination that recedes through the ejecta layers.  After this photospheric phase, the light curve displays an exponential decline when it is predominantly powered by the radioactive decay of isotopes created in the explosion ($^{56}$Ni $\\rightarrow ^{56}$Co $\\rightarrow ^{56}$Fe).  In the presence of a dense circumstellar medium (CSM), the expanding SN shock will collide with the CSM and convert the bulk kinetic energy of the explosion into light \\citep{Chevalier94}, and produce a relatively narrow (2 -- 4 $\\times 10^{3}$~\\kms) H$\\alpha$ line that is the hallmark of SNe~IIn \\citep{Schlegel90}.  The additional luminosity from the circumstellar interaction can be extremely large and even dominate the total luminosity.  Furthermore, while the luminosity from radioactive decay will decline quickly after a few months, the luminosity from circumstellar interaction can persist at a near-constant level for several years.  Recently, several extremely luminous SNe have emerged with peak luminosities with $M_{R} < -21$~mag.  SN~2006gy had a very long rise time (70~days) and peaked at $M_{R} = -21.8$~mag \\citep{Smith07_06gy, Ofek07}.  SN~2008fz had a similar peak brightness ($M_{\\rm V} = -22.3$~mag) and light-curve shape \\citep{Drake10}. SN~2005ap was very luminous at peak ($M_{\\rm unf} = -22.7$~mag), but had a fast rise and decay.  SN~2008es was similar to SN~2005ap, peaking at $M_{V} = -22.3$~mag and having a fast rise and decline \\citep{Miller09_08es, Gezari09_08es}.  A self-obscured luminous supernova, SN~2007va, was detected by \\cite{Kozlowski10} only in the infrared with an absolute mid-IR peak magnitude of $M_{[4.5]} \\approx -24.2$.  More recently, \\citet{Quimby09} and \\citet{GalYam09} announced the discovery of three extremely luminous SNe, respectively, which they suggest are powered by a pulsational pair instability \\citep{Woosley07}. \\cite{Pastorello10} finds evidence that these extremely luminous SNe may be connected to SN~Ic.  A pulsational pair instability is expected to be important in only the most massive stars -- those exceeding 95~$M_{\\sun}$. The production of electron-positron pairs results in a contraction and then explosive nuclear burning which ejects some significant number of solar masses worth of material from the envelope.  Subsequent repetitions of this sequence of events result in ejected shells catching up with previous ejected material, now at much larger radius, producing radiated energy due to the shell collisions.  This process is estimated to be capable of producing $10^{50}$~ergs of light and, as importantly, can repeat on short timescales, providing longer-timescale luminosity.  The energetics of these events push the envelope of our understanding of stellar evolution.  The peak luminosity, if powered by radioactive decay, would require \\about10~$M_{\\sun}$ of $^{56}$Ni.  This very large amount points to an extremely massive progenitor and the possibility that these events were the result of a pair-instability SN \\citep{Barkat67, Bond84}.  Alternatively, a significant amount of the luminosity may be produced by circumstellar interaction, but the mass of the CSM necessary for the luminosity still indicates that the progenitors must have been massive stars with significant mass-loss histories.  Furthermore, there is a limit on the energy for which conventional core-collapse explosions are viable.  Under a conventional core-collapse explosion, the maximum energy emitted by a SN is equivalent to the rest mass of a neutron star, a few $10^{54}$~ergs, with 99\\% being emitted as neutrinos \\citep[e.g.,][]{Woosley05}.  The remaining energy, a few $10^{52}$~ergs, is either coupled to the baryonic material as kinetic energy, or emitted as electromagnetic radiation.  If a SN has demonstrably greater than a few $10^{52}$~ergs, the conventional core-collapse scenario must be re-examined.  This argument has been used for the extremely energetic broad-lined SNe~Ic associated with gamma-ray bursts \\citep[e.g.,][]{Woosley93, Iwamoto98}; although these constraints are placed on this class of objects using the kinetic energy, not the radiated energy.  In the local universe, there have been several well-observed SNe~IIn: SNe~1988Z \\citep{Turatto93}, 1994W \\citep{Sollerman98}, 1995N \\citep{Fransson02}, 1998S \\citep[e.g.,][]{Leonard00}, 1999el \\citep{DiCarlo02}, 2005ip \\citep{Smith09_05ip}, and 2007rt \\citep{Trundle09}.  In addition to their similar spectral evolution, SNe~IIn all have long plateaus in their light curves after an initial decline.  This plateau is from shock energy being continuously converted into visual light through circumstellar interaction.  Once the shock extends to a radius where the density of the CSM drops, the source of the luminosity is diminished and the SN fades significantly.  Observations suggest that dust can form in the ejecta of SNe \\citep{Moseley89, Kozasa89, Sugerman06, Smith08_06jc, Rho08, Fox09, Kotak09_04et}.  This process appears to be enhanced by increasing the density of the CSM; i.e., SNe~IIn appear to be better at forming dust than SNe~IIP.  An infrared light echo can appear at late times when light from the SN explosion reaches a dust shell, and the UV/optical light is reprocessed into the infrared \\citep[e.g.,][]{Gerardy02a}.  Here we present \\LMC902, discovered behind the LMC by the \\mbox{SuperMACHO} microlensing survey.  \\LMC902 has the light curve shape and spectral signatures of a SN~IIn, but with a peak luminosity comparable to the most luminous SNe ever discovered.  In Section~\\ref{sec:observations} we describe the densely-sampled \\mbox{SuperMACHO} and OGLE-III difference imaging photometry (covering a baseline of 12 years, and the SN light curve over 4.7 years), and spectroscopy taken at the peak of the SN, and 1, 2, 5, and 6 years later, and in Section~\\ref{sec:comp} we compare the observations to well studied SNe~IIn.  In Section~\\ref{sec:physical} we calculate the total bolometric luminosity and energy output of the SN, and show that it exceeds the radiative output of any other SN observed to date.  We compare possible models for the source of the energy, and show evidence for an IR echo from variable IR flux measured in archival {\\it Spitzer} IRAC observations.  In Section~\\ref{sec:conclusions}, we summarize our results. ",
        "conclusions": ""
    },
    "0911/0911.4248_arXiv.txt": {
        "abstract": "We present solar photospheric abundances for 12 elements from optical and near-infrared spectroscopy. The abundance analysis was conducted employing 3D hydrodynamical (\\COBOLD) as well as standard 1D hydrostatic model atmospheres. We compare our results to others with emphasis on discrepancies and still lingering problems, in particular exemplified by the pivotal abundance of oxygen. We argue that the thermal structure of the lower solar photosphere is very well represented by our 3D model. We obtain an excellent match of the observed center-to-limb variation of the line-blanketed continuum intensity, also at wavelengths shortward of the Balmer jump.. ",
        "introduction": "In recent years several solar abundances experienced a significant downward revision, among them major contributors to the overall solar metalicity (\\cite[Asplund et al. 2005]{AGS05}). In part, the downward revision was attributed to the application of 3D model atmospheres. Due to the importance of the solar composition as a fundamental ``yardstick'' in astronomy, the CIFIST\\footnote{Cosmological Impact of the FIrst STars, an EU funded Marie Curie Excellence project} Team and its collaborators started an independent investigation of the solar abundances applying its self-developed analysis tools, in particular its own 3D model atmosphere code dubbed \\COBOLD\\ (\\cite[Freytag et al. 2002]{freytag02}, \\cite[Wedemeyer et al 2004]{wedemeyer04}).  Table~\\ref{t1} summarizes the result for 12 elements in comparison to other works. Considering the latest compilation of Asplund and collaborators one can note a convergence towards a unique abundance set. However, there are still sizable differences present, in particular concerning the abundant element oxygen.  Formally, the overlapping error bars could be taken to basically signal correspondence. However, one must keep in mind that certain systematics (observed spectra, oscillator strength, analysis methodology) are shared among all groups, and from that perspective differences are still on a rather high level. In this contribution we want to comment on a few of the lingering problems when it comes to the spectroscopic determination of solar abundances.  \\begin{table} \\begin{center} \\caption{Abundances derived by the \\COBOLD\\ group in comparison to other compilations: AG89 \\cite{AG89}; GS98 \\cite{GS98}; AGS05 \\cite{AGS05}; AGSS09 \\cite{AGSS09}. El denotes the element, N the number of spectral lines used in our analysis. The last two rows give the total mass fraction of metals, and metals relative to hydrogen. Values set in italics refer to meteoritic abundances.}\\label{t1} {\\small \\begin{tabular}{|l|r|l|l|l|l|l|} \\hline El & N  & \\COBOLD & AG89 & GS98 & AGS05 & AGSS09 \\\\ \\hline Li & 1  & $1.03\\pm 0.03$ & $1.16\\pm 0.10$ & $1.10\\pm 0.10$        & $1.05\\pm 0.10$        & $1.05\\pm 0.10$ \\\\ C  & 43 & $8.50\\pm 0.06$ & $8.56\\pm 0.04$ & $8.52\\pm 0.06$        & $8.39\\pm 0.05$        & $8.43\\pm 0.05$ \\\\ N  & 12 & $7.86\\pm 0.12$ & $8.05\\pm 0.04$ & $7.92\\pm 0.06$        & $7.78\\pm 0.06$        & $7.83\\pm 0.05$ \\\\ O  & 10 & $8.76\\pm 0.07$ & $8.93\\pm 0.035$& $8.83\\pm 0.06$        & $8.66\\pm 0.05$        & $8.69\\pm 0.05$ \\\\ P  & 5  & $5.46\\pm 0.04$ & $5.45\\pm 0.04$ & $5.45\\pm 0.04$        & $5.36\\pm 0.04$        & $5.41\\pm 0.03$ \\\\ S  & 9  & $7.16\\pm 0.05$ & $7.21\\pm 0.06$ & $7.33\\pm 0.11$        & $7.14\\pm 0.05$        & $7.12\\pm 0.03$ \\\\ Eu & 5  & $0.52\\pm 0.03$ & $0.51\\pm 0.08$ & $0.51\\pm 0.08$        & $0.52\\pm 0.06$        & $0.52\\pm 0.04$ \\\\ Hf & 4  & $0.87\\pm 0.04$ & $0.88\\pm 0.08$ & $0.88\\pm 0.08$        & $0.88\\pm 0.08$        & $0.85\\pm 0.04$ \\\\ Th & 1  & $0.08\\pm 0.03$ & $0.12\\pm 0.06$ & ${\\it 0.09\\pm 0.02}$  & ${\\it 0.06\\pm 0.05}$  & $0.02\\pm 0.10$ \\\\ K  & 6  & $5.11\\pm 0.09$ & $5.12\\pm 0.13$ & $5.12\\pm 0.13$        & $5.08\\pm 0.07$        & $5.03\\pm 0.09$ \\\\ Fe & 15 & $7.52\\pm 0.06$ & $7.67\\pm 0.03$ & $7.50\\pm 0.05$        & $7.45\\pm 0.05$        & $7.50\\pm 0.04$ \\\\ Os & 3  & $1.36\\pm 0.19$ & $1.45\\pm 0.10$ & $1.45\\pm 0.10$        & $1.45\\pm 0.10$        & $1.25\\pm 0.07$ \\\\ &    &                         &                &                       &                       & \\\\ Z  &    & 0.0153         & 0.0189         & 0.0171                & 0.0122                & 0.0134 \\\\ Z/X&    & 0.0209         & 0.0267         & 0.0234                & 0.0165                & 0.0183 \\\\ \\hline \\end{tabular} } % \\end{center} \\end{table} ",
        "conclusions": ""
    },
    "0911/0911.4554_arXiv.txt": {
        "abstract": "New CCD photometry during 4 successive years from 2005 is presented for the eclipsing binary GW Cep, together with reasonable explanations for the light and period variations. All historical light curves, obtained over a 30-year interval, display striking light changes, and are best modeled by the simultaneous existence of a cool spot and a hot spot on the more massive cool component star. The facts that the system is magnetically active and that the hot spot has consistently existed on the inner hemisphere of the star indicate that the two spots are formed by (1) magnetic dynamo-related activity on the cool star and (2) mass transfer from the primary to the secondary component. Based on 38 light-curve timings from the Wilson-Devinney code and all other minimum epochs, a period study of GW Cep reveals that the orbital period has experienced a sinusoidal variation with a period and semi-amplitude of 32.6 yrs and 0.009 d, respectively. In principle, these may be produced either by a light-travel-time effect due to a third body or by an active magnetic cycle of at least one component star. Because we failed to find any connection between luminosity variability and the period change, that change most likely arises from the existence of an unseen third companion star with a minimum mass of 0.22 $M_\\odot$ gravitationally bound to the eclipsing pair. ",
        "introduction": "After GW Cep (CSV 5941, BV 7, 2MASS J01455862+8004553) was discovered to be a variable star by Strohmeier (Geyer et al. 1955), its light curves were made photoelectrically by Meinunger \\& Wenzel (1965), Hoffmann (1982), Landolt (1992), and Pribulla et al. (2001). They recognized it as a W-subtype (defined empirically by Binnendijk (1970)) eclipsing binary with complete eclipses. Kaluzny (1984) analyzed Hoffmann's $BV$ light curves by using Rucinski's (1976) code and concluded that the binary is an over-contact system with the characteristic parameters of $q$=2.703 and $i$=83$^\\circ.9$ and a fill-out factor $f$=11.3 \\%. Pribulla et al. analyzed their own $BV$ light curves using the Wilson \\& Devinney (1971, hereafter W-D) code. Except for a value of $f$=23.5 \\%, their results are very similar to those of the earlier study. The orbital period of GW Cep has been examined by Pribulla et al. and Qian (2003), who reported that the period is decreasing. More recently, Chochol et al. (2006) suggested that the variation of the orbital period could be produced by two possible forms: a single light-travel time (LTT) ephemeris with and without a quadratic term.  Although GW Cep has been studied photometrically these several times, intrinsic light variations due to starspots have not yet been considered. Additionally, the period variation still has not been described as conclusively as can be desired. In this paper, we present and discuss the long-term photometric behavior of the binary system from detailed studies of all available data. ",
        "conclusions": "New multiband CCD light curves plus historical ones all display complete eclipses leading to system descriptions of high weight. The O'Connell effect and superposed long-term light variability are best explained by both a cool spot and a hot one on the more massive secondary star. The long-term variability has been ascribed to changes in the size of the cool spot. This spot may be produced by magnetic dynamo-related activity because the system is rotating rapidly and has a common convective envelope. The hot spot, on the other hand, is likely caused by impact from mass transfer between the components.  Over 45 years, the orbital period change of GW Cep has shown a periodic oscillation with a cycle length of 32.6 yr and a semi-amplitude of 0.009 d. In principle, this oscillation may be produced either by an LTT effect due to an unseen third star with a scaled mass of $M_3 \\sin i_3$=0.22 $M_\\odot$ or by an active magnetic cycle in the more massive secondary star. We indicate that the latter interpretation is the less likely of the two possibilities.  With a relatively favorable light ratio and complete eclipses GW Cep clearly merits a radial velocity determination. Those eventual results will add weight to understanding the entire class of contact binaries but is there anything about the binary which appears singular at the present time?  One item which attracts attention is the unfamiliar finding that the Keplerian period has been constant over 45 years.  This may be examined in some detail from The Atlas of $O$--$C$ Diagrams of Eclipsing Binary Stars (Kreiner et al. 2001). From that source we extracted the period histories of all cool contact binaries which have periods shorter than that of GW Cep. These systems number 32 (1 A-type, 25 W-types and 6 not presently classified) and the shortest period is 0.221 d. In order to complete a sample for which GW Cep will have the median period, we next found the 32 contact binaries which show successively longer periods. These 32 systems are apportioned among 11 A-types, 17 W-types and 4 presently unclassified and the longest period is 0.365 d. Of this total sample of 65 pairs, 6 (AT Aqr, GW Cep, V906 Cyg, V743 Sgr, RZ UMi and BP Vel) appear at present to show constant periods.  The remainder show secularly variable periods of either algebraic sign upon which bounded oscillations are superposed in some cases.  Clearly, constant periods are an uncommon phenomenon in the sample and one may conjecture that they signify brief episodes of constant period behavior during themal oscillation evolutionary events. Of the 6 binaries, all but GW Cep have poor continuity in their period histories leaving that object as the best system at present to test such a possibility. Not only GW Cep, but the 5 other binaries merit continuing period monitoring."
    },
    "0911/0911.3232_arXiv.txt": {
        "abstract": "We investigate the scaling relations between the X-ray and the thermal Sunyaev--Zel'dovich Effect (SZE) properties of clusters of galaxies, using data taken during 2007 by the Y.T. Lee Array for Microwave Background Anisotropy (AMiBA) at 94~GHz for the six clusters A1689, A1995, A2142, A2163, A2261, and A2390. The scaling relations relate the integrated Compton-$y$ parameter $Y_{2500}$ to the X-ray derived gas temperature $T_\\textrm{e}$, total mass $M_{2500}$, and bolometric luminosity $L_X$ within $r_{2500}$. Our results for the power-law index and normalization are both consistent with the self-similar model and other studies in the literature except for the $Y_{2500}$--$L_X$ relation, for which a physical explanation is given though further investigation may be still needed. Our results not only provide confidence for the AMiBA project but also support our understanding of galaxy clusters. ",
        "introduction": "The Sunyaev--Zel'dovich Effect (SZE) is a powerful tool that can potentially answer long-standing questions about the large-scale distribution of matter. The SZE is a spectral distortion of the cosmic microwave background (CMB), induced when a fraction of CMB photons are scattered by hot electrons in the cores of massive galaxy clusters \\citep{SZ1970}. The redshift independence of the SZE  enables the direct detection of distant clusters without the $(1+z)^{4}$ surface brightness dimming that limits other techniques, including X-ray observations. Clusters studied via the SZE are therefore effective cosmological probes. Studying their properties in detail will lead to heightened understanding of the mass power spectrum, and should provide improved constraints on cosmological parameters.  In the simplest scenario, where gravity is assumed to be the only influence on the formation of galaxy clusters, a simple `self-similar' model can be used to relate the physical properties of clusters \\citep{Kaiser1986}. Assuming spherical collapse of the dark matter (DM) halo, and hydrostatic equilibrium of gas in the DM gravitational potential, one can derive power-law scaling relations between various X-ray and SZE quantities, e.g.~luminosity and temperature, gas mass and temperature, total mass and luminosity, entropy and temperature, and Compton-$y$ parameter and temperature. Existence of these relations in observations can be seen in, for example, \\citet{Mushotzky1997,Benson2004,Bonamente2007,Morandi2007}. These relations are also found in numerical simulations, e.g., between X-ray quantities \\citep{Nagai2007}, SZE flux and total mass or gas mass \\citep{Motl2005,Nagai2006}, and integrated Compton-$y$ and temperature or luminosity \\citep{Silva2004}. Deviations from the scaling relations should reveal the importance of non-gravitational processes for the formation of clusters \\citep[e.~g.~][]{Allen1998a,McCarthy2002,McCarthy2003a}, or constrain the mass distributions of clusters \\citep{Reiprich2002}. Furthermore, the scaling relations can be used as an utility to extract important quantities and evolution behaviors for remote clusters using SZE observables alone.  The Y.T.~Lee Array for Microwave Background Anisotropy (AMiBA) experiment \\citep{Ho2008} observed and detected the SZEs of six massive Abell clusters in the range $0.09<z<0.32$ during the year 2007 \\citep{Wu2008}. AMiBA is a coplanar interferometer that during 2007 operated at 94~GHz with seven 0.6-m antennas in a hexagonal close-packed configuration, giving a synthesized resolution of about $6\\arcmin$. The array has a sensitivity of 63~mJy/hr for on-source integration, and an overall efficiency of $0.36$ \\citep{Lin2008}. Details for the transformation of the raw data into calibrated visibilities are presented in \\citet{Wu2008}, and the checks on data integrity are described in \\citet{Nishioka2008}. At our observing frequency the SZE signal is an intensity decrement in the CMB. We fit the central (peak) decrement in the AMiBA visibilities using isothermal $\\beta$-models \\citep{CF1976}, taking account of contamination from the primary CMB and foreground emissions \\citep{Liu2008}. Other companion papers include \\citet{Chen2008} and \\citet{Koch2008a}, where the technical aspects of the instruments are described, \\citet{Umetsu2008}, where the AMiBA SZE data is combined with weak lensing data from Subaru to analyze the distributions of mass and hot baryons, \\citet{Koch2008}, where the Hubble constant is estimated from AMiBA SZE and X-ray data, and \\citet{Molnar2008}, which discusses the feasibility of further constraining the intra-cluster gas model using AMiBA upgraded to 13 antennas \\citep[AMiBA13]{Ho2008}. The consistency of our results with other observations and theoretical expectations will validate not only the performance and capability of instruments, but also the analysis methodology. Since AMiBA is one of few leading SZE instruments operating at 3-mm wavelength, we anticipate that it will fill an important role by providing 3-mm SZE data for spectral studies.  In this article we address the scaling relations between the integrated SZE Compton-$y$ parameter obtained by AMiBA and X-ray gas temperature, X-ray luminosity, and total cluster mass derived from the literature. In Section~\\ref{sec:xrpars} we discuss the cluster gas models and cluster parameters derived from the X-ray data. In Section~\\ref{sec:szepars} we calculate the integrated Compton-$y$ parameter for each of the six clusters. In Section~\\ref{sec:scalings} we investigate the scaling relations including the consideration for errors. We further discuss our results in Section~\\ref{sec:dandc} and draw conclusions in Section~\\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  In understanding cluster physics and cosmic evolution, the study of scaling relations for galaxy clusters is becoming more important since SZE observations such as those from the Sunyaev--Zel'dovich Array (SZA) and the South Pole Telescope (SPT) will detect clusters which are beyond X-ray detection limits, for which SZE scaling relations provide the first indications of cluster properties. In addition, deviations between the theoretical and observational results of the scaling relations can serve to examine the non-gravitational processes in the formation of clusters, which are not well understood at present.  As one of the few SZE instruments working at 3-mm, the AMiBA experiment observed six Abell clusters during 2007. The derived integrated Compton-$y$ parameters, $Y_{2500}$, are compared to other observations at different frequencies, as summarized in Table~\\ref{tab:SZ-obs}. Our results are consistent with those from BIMA/OVRO, but appear to show lower Comptonizations than those from SuZIE~II. We have also investigated the three scaling relations between $Y_{2500}$ and the X-ray spectroscopic temperatures, total masses within $r_{2500}$, and bolometric X-ray luminosities. Our results for the scaling relations are summarized in Table~\\ref{tab:srline}.  Our power-law indices for the three scaling relations are broadly consistent with the self-similar model and observational results in the literature, except for that the $Y_{2500}-L_{\\mathrm{X}}$ relation based on \\textit{Chandra}-derived parameters has a slope lower than the expectation of the self-similar model, and is sensitive to the chosen set of X-ray parameters and the treatment of cooling cores. These discrepancies might indicate either exotic properties for clusters or hidden flaws in our SZE data, although the scatter is still large, about a factor of two in the integrated Compton-$y$ relative to the fit line. The agreement between the normalizations found by different workers for our three scaling relations seems to support the idea that there is no strong scatter in the gas fraction (see Sec.~\\ref{sec:sub:scaletheory}).  In conclusion, the agreement between our results and those from the literature provides not only confidence for this project but also supports to our understanding of galaxy clusters. For AMiBA, significant improvements are expected following the expansion to a 13-element configuration with 1.2-m antennas \\citep[AMiBA13]{Ho2008}, which will provide better resolution and higher sensitivity. The capability of resolving SZE clusters will make it possible to measure the cluster profiles independent of the X-ray data \\citep{Molnar2008} and to estimate the properties of the clusters which currently do not have good X-ray data."
    },
    "0911/0911.4909_arXiv.txt": {
        "abstract": "An overview of selected topical problems on modelling oscillation properties in solar-like stars is presented. High-quality oscillation data from both space-borne intensity observations and ground-based spectroscopic measurements provide first tests of the still-ill-understood, superficial layers in distant stars. Emphasis will be given to modelling the pulsation dynamics of the stellar surface layers, the stochastic excitation processes and the associated dynamics of the turbulent fluxes of heat and momentum. ",
        "introduction": "\\label{s:intro} With the high-quality photometric data from the Kepler satellite \\citep[e.g.][]{Christ07} and spectroscopic data from ground-based observation campaigns, and later also from the Danish SONG network \\citep{Grundahl07}, we shall be able to address more carefully many of the current problems in modelling pulsation properties in solar-like stars. In solar-like stars the p-mode lifetimes and amplitudes are crucially affected by the processes that take place in the outer convectively unstable stellar layers. A proper modelling of the dynamics of the convective heat and momentum transfer is therefore essential. Recent 3D numerical simulations of the largest scales of the convection \\citep[e.g.][]{SteinNordlund01, Samadi03, Georgobiani04, Stein04, Jacoutot08, Miesch08} have been proven to be extremely useful for calibrating and testing semi-analytical models for convection and stochastic excitation. However, the high-Reynolds-number (and low-Prandtl-number) turbulent convection in stars still prohibits today's simulations from resolving all the required scales of stellar turbulence. Consequently such simulations have to use sub-grid-scale models, which may lead to different results \\citep[e.g.][]{Jacoutot08}. Because this situation will not change in the near future, we still need analytical models for describing the convection and pulsation dynamics in stars.  Sections 2 and 3 will discuss selected problems of our current understanding of nonadiabatic pulsation dynamics in the Sun and in the hotter F5 star Procyon. Although we can reasonably well reproduce the observed pulsation properties in the Sun, the recent Procyon observations \\citep[e.g.][and references therein]{Arentoft08} have revealed serious problems in our models for estimating the oscillation amplitudes in stars hotter than the Sun. In Section~4, therefore, I shall address the problem of selecting a proper temporal turbulence spectrum for modelling the stochastic energy supply rate for acoustic modes. ",
        "conclusions": ""
    },
    "0911/0911.5679_arXiv.txt": {
        "abstract": "The main goal of the Andyrchy-BUST experiment is to study the primary cosmic rays  spectrum and composition around the knee. The experimental data on the knee, as observed in the electromagnetic and high energy muon components, are presented. The electromagnetic component in our experiment is measured using the \"Andyrchy\" EAS array. High energy muon component (with 230 GeV threshold energy of muons) is measured using the Baksan Underground Scintillation Telescope (BUST). The location of the \"Andyrchy\" right above the BUST gives us a possibility for simultaneous measurements of both EAS components. ",
        "introduction": "In the range of primary energies of $10^{14} - 10^{15}$ eV per nucleus, direct methods for studying the energy spectrum and nuclear composition of primary cosmic rays become inefficient because of a decrease in the flux of primary particles with an increase in their energy. Therefore, at these and, of course, higher energies, indirect methods based on simultaneous measurement of the characteristics of different components of extensive air showers (EASs), which are initiated by the primary particle in the atmosphere, are used. But the interpretation of these measurements requires their comparison with EAS simulations in the atmosphere. In turn, the calculation results depend on the hadronic interaction models. The main problem is the extrapolation of these models into kinematical and energy regions still unexplored by present-day collider experiments. So, the measurements of different EAS components are now used for both studying the primary composition and testing interaction models \\cite{bib:KASCADE99} -  \\cite{bib:KASCADE09}.   In this paper three types of experimental data are analyzed: muon number spectrum, EAS size spectrum and correlation between muon number and EAS size simultaneously measured. Integral muon number spectrum has been measured using the Baksan Underground Scintillation Telescope (BUST) \\cite{bib:Petkov2008}. The  EAS size spectrum has been measured using the \"Andyrchy\" EAS array \\cite{bib:Chudakov97}. The dependence of the mean number of high energy muons on EAS size has been measured by simultaneous operation of both devices  \\cite{bib:Chudakov97a}, \\cite{bib:Petkov2003}.\\\\ ",
        "conclusions": "It is widely known that none of the present interaction models can completely describe a full set of experimental data for cosmic rays. Joint analysis of the various characteristics of different EAS components, especially measured in one and the same experiment, can be used for both studying the primary composition and testing interaction models.  In this paper three types of experimental data, taken in our experiment, were analyzed: high energy ($E_{\\mu} \\ge 230$ GeV) muon number spectrum, EAS size spectrum and dependence of the mean number of high energy muons on EAS size. CORSIKA code v.6900, with QGSJetII-03 and Fluka as the high and low energy hadronic interaction models, has been used for EAS simulations. As it was mentioned above, the calculations of these observables need only the following characteristics of EAS:\\\\ 1) $W_A(E_0,N_{r.p.},N_{\\mu})$ - the probability for primary nucleus with atomic number A and energy $E_0$ to produce EAS with size $N_{r.p.}$ and total number of high energy muons $N_{\\mu}$ at the observation level;\\\\ 2) $W_A(E_0,N_{r.p.})$ - the probability to produce EAS with size $N_{r.p.}$ at the observation level;\\\\ 3) $\\overline N_{\\mu}(E_0, A)$ -  high energy MPF;\\\\ 4) $G(A, \\overline N_{\\mu}, N_{\\mu})$ - high energy muons fluctuation function;\\\\ 5) $f(r, E_0, A)$ - high energy muons LDF.\\\\ It should be noted that characteristics 2 - 4 can be derived from 1.  \\noindent The analysis has shown that:\\\\ 1) our experimental data can be brought into good enough agreement using changed MPF (Fig. \\ref{Fig_8});\\\\ 2) in both cases the primary composition gets heavier across the knee.   \\begin{figure}[!t] \\centering \\includegraphics[width=1.0 \\columnwidth ]{icrcPetkov_fig10.EPS} \\caption{$\\overline N_{\\mu}(N_{rp})$ dependence. Points - experiment. Solid line is calculated dependencies for the first primary compositions and changed MPF. The dashed lines are calculated dependencies for pure protons and iron nuclei with changed MPF.} \\label{Fig_10} \\end{figure}   \\vspace{3mm}  {\\bf Acknowledgements.}  The author is grateful to the Organising Committee of the 31$^{st}$ ICRC for the possibility to present this talk at the highlight session. This work was supported by the \"Neutrino Physics and Astrophysics\" Program for Basic Research of the Presidium of the Russian Academy of Sciences and by \"State Program for Support of Leading Scientific Schools\" (grant no. NSh-321.2008.2). This work was also supported in part by the Russian Foundation for Basic Research, Grant No. 08-07-90400."
    },
    "0911/0911.0789_arXiv.txt": {
        "abstract": "{Star clusters, stellar evolution, white dwarfs} Open and globular star clusters have served as benchmarks for the study of stellar evolution due to their supposed nature as simple stellar populations of the same age and metallicity. After a brief review of some of the pioneering work that established the importance of imaging stars in these systems, we focus on several recent studies that have challenged our fundamental picture of star clusters.  These new studies indicate that star clusters can very well harbour multiple stellar populations, possibly formed through self-enrichment processes from the first-generation stars that evolved through post-main-sequence evolutionary phases.  Correctly interpreting stellar evolution in such systems is tied to our understanding of both chemical-enrichment mechanisms, including stellar mass loss along the giant branches, and the dynamical state of the cluster.  We illustrate recent imaging, spectroscopic and theoretical studies that have begun to shed new light on the evolutionary processes that occur within star clusters. ",
        "introduction": "Understanding stellar evolution has been one of the most important pursuits of observational astronomy over the past century.  In 1905, Ejnar Hertzsprung demonstrated that stars emitting light with a similar spectrum, and with the same parallax, could have very different luminosities.  Hertzsprung referred to these nearby stars, which therefore had the same temperature and distance, as `giants' (high luminosity) and `dwarfs'.  Just a few years after this study, Hertzsprung and Henry Norris Russell obtained and analysed new observations that characterized the basic properties of hundreds of nearby field stars (Russell 1913; 1914{\\it a}) and for the nearest co-moving groups (e.g., the Pleiades and Hyades: see figure~\\ref{fig:HRDiagrams}; Hertzsprung 1911).  This pioneering work led to one of the most important diagnostic tools in astrophysics, the luminosity versus spectral type plane, rightfully named the Hertzsprung--Russell (H--R) diagram.  The first H--R diagrams illustrated clear evidence of groupings and sequences of stars, leading to speculation of the probable order of stellar evolution, as well as some interesting correlations (Russell 1914{\\it a,b}).  For example, based on these observations, Russell remarks that ``there appears, from the rather scanty evidence at present available, to be some correlation between mass and luminosity''.  Remarkably, these initial observations were rather complete, illustrating a rich main sequence extending over 10 magnitudes from A-type stars to M dwarfs (including several eclipsing variables), an abundant population of red giants, and the first white dwarf observed (40~Eri~B, seen by W.\\ Herschel in 1783).\\footnote{Given the lack of other `faint white stars', Russell initially questioned the quality of the spectrum of this star, which is a member of a triple system (Russell 1914{\\it a}).}  \\begin{figure}[h] \\begin{center} \\includegraphics[width=7.4cm, bb= 20 145 575 575]{HertzsprungCMD.eps} \\includegraphics[width=5.1cm]{RussellCMD.eps} \\end{center} \\caption{The first published Hertzsprung--Russell (H--R) diagrams of the Pleiades and Hyades clusters ({\\it left} and {\\it middle}; Hertzsprung 1911), as well as for nearby field stars ({\\it right}; Russell 1914{\\it a}).  Plotted on the vertical axis are the magnitudes (e.g., more luminous stars are at the top) and on the horizontal axis are the peak wavelengths or spectral types of the stars (e.g., hotter, bluer stars towards the left).\\label{fig:HRDiagrams}} \\end{figure}  Star clusters were recognized as excellent laboratories to test theories of stellar evolution soon after the first H--R diagrams of the Pleiades and Hyades were published.  Over the years, it was quickly realized that each individual cluster is a simple stellar population, being coeval, cospatial and isometallic (some of this knowledge came later) and, therefore, represents a controlled testbed with well-established properties that are common to all member stars (see also Bruzual 2010).  Yet, the constituent stars in any given cluster span a large range in stellar mass, the most important parameter dictating stellar evolution as we now know.  The present observations of an H--R diagram for any star cluster therefore represents a snapshot of how stellar evolution has shaped the population at a given age and metallicity; the other properties only have a secondary influence. As we link together the observations of many of these clusters, over a wide range in age and metallicity, a complete picture of stellar evolution begins to emerge.  Fortunately, the timing of Hertzsprung and Russell's first studies of the H--R diagrams of nearby star clusters coincided with the construction of the first large reflecting telescopes.  The Mount Wilson Observatory's 1.5~m (1908) and 2.5~m (1917) facilities and the Dominion Astrophysical Observatory's 1.8~m (1918) offered the exciting possibility of undertaking the first systematic (resolved) studies of many of the clusters in Messier's famous catalog (Messier 1774).\\footnote{Available online at {\\tt http://messier.obspm.fr/xtra/history/m-cat71.html}.}  In a series of papers, Harlow Shapley catalogued photometry for hundreds of individual stars in the nearest star clusters (e.g., Shapley 1917; and related articles in the 1910s and 1920s) and provided the first clues as to the properties of these systems.  For example, Shapley measured the distances and distributions in space, sizes, luminosities and even proper motions of several nearby star clusters.  These photometric studies led directly to the first ideas on stellar physics (e.g., the static main sequence and the mass--luminosity relation; Eddington 1926), well before the source of a star's energy was known.  It was only later, in the 1930s and 1940s, that hydrogen fusion and the proton--proton chain were detailed, leading to much of our current picture of stellar evolution (e.g., Chandrasekhar 1939; Bethe \\& Marshak 1939; Hayashi 1949; Henyey \\ea 1959). ",
        "conclusions": "\\label{future}  The discussion in this paper focuses on the knowledge we have recently gained on the evolution of stars, primarily based on deep imaging and faint spectroscopic studies of star clusters in our Galaxy.  Among the topics we have discussed are multiple stellar sequences in clusters, extremely low-mass hydrogen-burning stars, post-main-sequence evolution and stellar mass loss, white dwarfs and the interplay of stellar and dynamical evolution within clusters.  As an extension of these primary studies, it is important to keep in mind a key implication of these findings as they relate to general astrophysics.  As early as the first resolved imaging studies of Milky Way clusters, it was understood that these systems are very important tracers of Milky Way structure and formation processes.  For example, Harlow Shapley used the distribution of globular clusters on the sky to infer the location of the centre of the Galaxy (Shapley 1918).  Today, the comparison of H--R diagrams of Milky Way open and globular clusters to stellar evolution models produces the calibration that commonly feeds population synthesis models, which, in turn, are used to interpret the properties of distant galaxies (e.g., Bruzual \\& Charlot 2003).  At high redshifts, the light that we measure from these unresolved sources comes from the distribution of stars along post-main-sequence evolutionary phases, such as the red-giant branch, horizontal branch and asymptotic giant branch.  Constraining stellar evolutionary processes, such as mass loss along the red-giant branch, therefore directly impacts measurements of star-formation histories, metallicities and mass-to-light ratios of these galaxies.  As we continue to refine stellar evolution theory with resolved studies of nearby systems, we must re-define our knowledge of galactic astronomy.  Even the most ambitious projects discussed above, such as the ACS Survey of Galactic Globular Clusters, represent a pencil-beam study with respect to surveys that are on the horizon for 2010--2019.  For example, {\\sl Pan-STARRS}, {\\sl SkyMapper} and {\\sl LSST} will provide multi-epoch, multifilter, deep homogeneous resolved photometry of most star clusters in our Galaxy.  It is especially important that these surveys, unlike the Sloan Digital Sky Survey ({\\sc sdss}), target regions of the Galactic disc to systematically sample open star clusters.  In this regard, the first wide-field imaging surveys of the southern hemisphere will image unexplored clusters that surely will fill empty regions of parameter space (e.g., age and metallicity). Kalirai \\ea (2009{\\it b}) present a summary of the type of science that these surveys may enable as related to star clusters.  A survey such as {\\sl LSST} will provide an unimaginable wealth of observational data to test stellar evolution models.  With a detection limit of 24--25th magnitude in multiple optical bandpasses in a {\\it single} visit, and a coadded $5\\sigma$ depth in the $r$ band of 27.8 mag, {\\sl LSST} will yield accurate turnoff photometry for all star clusters in its survey area out to beyond the edge of the Galaxy.  For a 12~Gyr globular cluster, this photometry will extend to over three magnitudes below the main-sequence turnoff at this distance.  For low-mass stars, previous surveys such as {\\sc sdss} and {\\sc 2mass} have yielded accurate photometry of faint M dwarfs out to distances of $\\sim 2$~kpc.  {\\sl LSST} will enable the first detection of such stars to beyond 10~kpc.  At this distance, the colour--magnitude relation of hundreds of star clusters will be established and permit a detailed and systematic investigation of variations in the relation with age and metallicity.  The present-day mass functions of the youngest clusters will be dynamically unevolved and, therefore, provide for new tests of the variation in the initial mass function as a function of environment, down to very-low-mass stars.  Pushing to the hydrogen-burning limit will also greatly benefit from new infrared observations, for example as proposed in missions such as {\\sl SASIR} (Bloom \\ea 2009) and {\\sl NIRSS} (Stern \\ea 2009).  These projects aim to image the entire sky in the $JHK$ (and possibly $L$) bands down to 24th magnitude.  At an age of 1~Gyr, a $M = 0.08$ M$_\\odot$ star has $M_V = 19$ mag (Baraffe \\ea 1998) and will therefore be seen by a survey such as {\\sl LSST} to 500~pc.  However, this star has a $V-H$ colour of 8, and is therefore much brighter in the near-infrared.  With {\\sl SASIR} or {\\sl NIRSS}, the complete mass function of all star clusters, down to the hydrogen-burning limit, can be characterized out to 2.5~kpc.  With the Wide Field Camera 3 (WFC3) imager on {\\sl HST}, and the {\\sl James Webb Space Telescope}, such stars can be easily probed in star clusters at distances of several tens of kpc.  These observations will also extend beyond the stellar and substellar threshold and characterize L and T dwarfs. This will make feasible a new age indicator for star clusters: the magnitude difference between the onset of brown dwarfs and the hydrogen-burning limit.  Similar to the wealth of new data that is expected for low-mass hydrogen-burning stars from these surveys, future observations of white dwarfs will allow unprecedented studies.  For the field population, a total of 20\\,000 white dwarfs are currently known (e.g., {\\sc sdss}; Eisenstein \\ea 2006).  {\\sl LSST} alone will increase the total sample size of white dwarfs in the Milky Way to over 50~million (Kalirai \\ea 2009{\\it c}).  The bright tip of the white-dwarf cooling sequence is located at $M_V \\sim 11$ mag and will be seen in clusters out to 20~kpc.  For a 1~Gyr cluster, the faintest white dwarfs have cooled to $M_V = 13$ mag and will be detected in clusters out to 8~kpc.  These white-dwarf cooling sequences not only provide direct age measurements (e.g., Hansen \\ea 2007) for the clusters (and therefore fix the primary parameter in the theoretical isochrone fitting, allowing secondary effects to be measured), but can also be followed up with current ({\\sl Keck, Gemini, VLT} and {\\sl Subaru}) and future (e.g., {\\sl TMT, GMT} and/or {\\sl E--ELT}) multi-object spectroscopic instruments to yield the connection between initial and final masses over a wide range of environments.  Finally, the synoptic nature of several future missions will permit proper-motion separation of cluster sequences from field-star contamination, permitting cleaner studies of stellar evolution in large samples of these systems."
    },
    "0911/0911.4480_arXiv.txt": {
        "abstract": "The {\\sc Galform} semi-analytic model of galaxy formation is used to explore the mechanisms primarily responsible for the three types of galaxies seen in the local universe:  bulge, bulge+disk and disk, identified with the visual morphological types E, S0/a-Sbc, and Sc-Scd, respectively. With a suitable choice of parameters the {\\sc Galform} model can accurately reproduce the observed local $K_{s}$-band luminosity function (LF) for galaxies split by visual morphological type. The successful set of model parameters is used to populate the Millennium Simulation with 9.4 million galaxies and their dark matter halos. The resulting catalogue is then used to explore the evolution of galaxies through cosmic history. The model predictions concur with recent observational results including the galaxy merger rate, the star formation rate and the seemingly anti-hierarchical evolution of ellipticals. However, the model also predicts significant evolution of the elliptical galaxy LF that is not observed. The discrepancy raises the possibility that samples of $z~{\\sim}$ 1 galaxies which have been selected using colour and morphological criteria may be contaminated with galaxies that are not actually ellipticals. ",
        "introduction": "The current paradigm, constrained in part by the $K_{s}$-band (hereafter $K$-band)  luminosity function (LF) (e.g., \\citealt{benson_what_2003}), has galaxy disks forming through a combination of cold gas accretion \\citep{weinberg_galaxy_2004} and feedback \\citep{oppenheimer_cosmological_2006-1}, with bulges resulting from mergers \\citep[e.g.,][]{barnes_dynamics_1992, hopkins_cosmological_2008, masjedi_growth_2008}. Thus, the diversity of galaxy morphologies seen today represents the culmination of multiple evolutionary paths.  These end points in galaxy evolution are captured in a taxonomy devised by  \\citet{hubble_realm_1936} and refined by \\citet{de_vaucouleurs_classification_1959} and \\cite{de_vaucouleurs_third_1991} that is based on the relative prominence of the stellar bulge and the degree of resolution of the spiral arms.  Despite attempts to emulate this classification scheme using quantitative proxies such as colours \\citep[e.g.,][]{strateva_color_2001}, concentration-asymmetry indices \\citep[e.g.,][]{watanabe_digital_1985, abraham_galaxy_1996, bell_optical_2003}, the Gini-M20 classifiers \\citep{lotz_new_2004} and parametric methods involving bulge/disk decompositions \\citep{schade_evolution_1996, ratnatunga_disk_1999, simard_deep_2002} there are no quantitative measures that uniquely segregate galaxies into the morphological groups represented by the de Vaucouleurs numerical {\\tt T}  system \\citep[e.g.,][]{graham_inclination-_2008}. Thus, proxies are merely a useful stopgap until the difficult and time consuming task of assigning visible morphologies to the ever-growing number of cataloged galaxies can be accomplished\\footnote{See the Galaxy Zoo Project, \\href{http://www.galaxyzoo.org/}{\\tt http://www.galaxyzoo.org/}}.  Fortunately, the time-consuming task of assigning visible morphologies has now been completed for the vast majority of nearby galaxies.  The principal aim of this paper is therefore to use the $K$-band LF for nearby galaxies \\citep{devereux_nearby_2009} to constrain models of their formation and evolution. The $K$-band LF  is of particular interest because of its relevance to understanding galaxy evolution in the context of Lambda Cold Dark Matter (${\\Lambda}$CDM) cosmology; at $z=0$, the $K$-band LF traces the stellar mass accumulated in galaxies at a wavelength where interstellar extinction is minimal \\citep{devereux_infrared_1987, bell_stellar_2000, bell_optical_2003}. Additionally, the functional forms for the $K$-band LFs distinguish between {\\it bulge} dominated and  {\\it disk}  dominated systems which suggests that at least two quite distinct galaxy formation mechanisms are at work to produce the diversity of morphological types seen in the local universe. Previous attempts to model the LFs in the context of hierarchical clustering scenarios \\citep[e.g.,][]{cole_hierarchical_2000,benson_what_2003} have focussed on the {\\it Universal} or {\\it Total} $K$-band LF, which is the sum over all galaxy types, because they are based on samples for which there are no pre-existing morphological assignments. In this paper, the $K$-band LFs are modeled for galaxies segregated by visible morphology in order to better understand: 1) the origin and evolution of bulgeless disk galaxies, represented in the observable universe by Sc-Scd galaxies; 2) the role of secular processes in building the bulges of spiral (S0/a---Sbc) and lenticular galaxies and; 3) the galaxy merger history that led to the formation of elliptical galaxies. Our analysis differs from \\cite{parry_galaxy_2009} in that the GALFORM model predictions are constrained by the observed $K$-band LFs for galaxies of different morphological type. Those model predictions are then compared with recent observational results including the type averaged galaxy merger rate and the type specific star formation rate, both quantified as a function of redshift.   Observationally, there is a weak dependence of $K$-band B/T on visual morphology \\pcite{graham_inclination-_2008} which is used in this paper to segregate the model galaxies into three broad groups representing the three types of galaxies seen in the local universe, namely bulge, bulge+disk and disk.\\footnote{In principle, other outputs from the model could be used to provide additional information on morphological type. However \\protect\\cite{graham_inclination-_2008} show that bulge/disk flux ratio is the most useful discriminant as it varies more, as a function of morphological type, than other indicators, such as bulge/disk size ratio, for example.} These three groups are identified with the visual morphological types E, S0/a-Sbc, and Sc-Scd, respectively. Following \\cite{graham_inclination-_2008}, these three broad groups are identified in the models by adopting the range of B/T ratio listed in Table~\\ref{tb:BTMap}. The table also indicates the distribution of B/T within each morphological class, illustrating that most disk galaxies (including S0's) have a rather low B/T ratio. For example, 80\\% of S0 and later type galaxies brighter than $M_{\\rm K}-5\\log_{10}{\\it h} = -21$ have B/T $\\le$ 0.54. If the magnitude limit is lowered to $M_{\\rm K}-5\\log_{10}{\\it h} = -19$, we find that 80\\% of such galaxies have B/T  $\\le$ 0.35. This is in agreement with the broad conclusion of \\cite{graham_inclination-_2008} that most disk galaxies (including S0's) have low B/T ratios (e.g. B/T $<$1/3).  \\begin{table} \\label{tb:BTMap} \\begin{center} \\begin{tabular}{lcccc} \\hline {\\bf Morphological} & & \\multicolumn{3}{c}{{\\bf B/T distribution}} \\\\ {\\bf class} & {\\bf B/T Range} & 20\\% & 50\\% & 80\\% \\\\ \\hline E      & $0.92 <   B/T_{\\rm K} \\le 1.00$ & 0.99 & 1.00 & 1.00 \\\\ S0-Sbc & $0.11 <   B/T_{\\rm K} \\le 0.92$ & 0.22 & 0.39 & 0.65 \\\\ Sc-Scd & $0.00 \\le B/T_{\\rm K} \\le 0.11$ & 0.00 & 0.01 & 0.05 \\\\ \\hline \\end{tabular} \\end{center} \\caption{The Table presents the mapping between morphological type and K-band bulge-to-total ratio required for our best-fit model to reproduce the relative abundances of each class in the interval $-23.5 < M_{\\rm K}-5\\log_{10}h \\le -23$. Additionally, the final three columns show the percentage of galaxies in each class with $M_{\\rm K}-5\\log h < -21$ and  B/T ratios less than the value indicated. For example, 80\\% of bright S0-Sbc galaxies have a B/T less than 0.65.} \\end{table} ",
        "conclusions": "The {\\sc Galform} semi-analytic model of galaxy formation \\citep{bower_breakinghierarchy_2006} has been employed to reproduce and explain trends in galaxy morphology seen across a wide range of luminosities. Small adjustments to the model parameters permit excellent agreement with the observed morphologically selected $K$-band LF of \\cite{devereux_nearby_2009}, although the unmodified \\cite{bower_breakinghierarchy_2006} model also performs very well. While predicting morphology from semi-analytic models (and numerical simulations) remains challenging, a conservative approach has been adopted by categorising model galaxies into just three broad classes: E (``pure spheroid''), S0-Sbc (``intermediate'') and Sc-Scd (``pure disk'') based on their $K$-band bulge-to-total luminosity ratio. Within this conservative categorisation, the results presented are converged with respect to the resolution of the N-body-derived merger trees (see Appendix~\\ref{app:ResStudy}).  Our conclusions follow.  \\begin{itemize} \\item The Sc-Scd disk galaxies, with ${\\sim}$ 10\\% or less of their K-band light coming from the bulge, can form in abundance even in a hierarchical cold dark matter cosmology in which merging is, in general, commonplace. The median merger history of the halos that these galaxies inhabit is not significantly different from that of elliptical galaxies which do experience significant merging. However, the variation around that median history is much less for the Sc-Scd galaxies than for ellipticals. The Sc-Scd class typically experience a very simple accretion history, deviating little from the median and exhibiting relatively few large merging events. The formation history of these systems is therefore determined by the rate of growth of their dark matter halo and the corresponding rate of gas supply, cooling and infall. Galaxies in this class have typically been forming stars in their disks over large fractions of cosmic history and continue to do so at the present day. \\cite{steinmetz_2002} and, more recently, \\cite{dutton_origin_2009} have suggested that producing bulge-less disk galaxies in sufficient numbers could be challenging in a CDM universe. In fact, because there is a sufficiently large number of dark matter halos with relatively quiet merging histories, bulge-less disk galaxies can prevail, in the numbers observed. Additionally, strong feedback prevents the formation of relatively massive satellites around these galaxies which are predominantly low mass. \\item Luminous elliptical galaxies form through multiple mergers of smaller progenitors. For the most massive ellipticals these mergers tend to be ``dry'' and so contribute no new star formation. Interestingly, the {\\sc Galform} model does not produce the dwarf elliptical sequence which should emerge at $M_{\\rm K}-5\\log_{10}h \\sim -21$ mag. Thus, the model prediction concerning the shape of the LF for  low luminosity ellipticals (with $M_{\\rm K}-5\\log_{10}h> -21$ mag) and the model prediction that they form through disk instability events is uncertain. \\item The intermediate morphological types, S0-Sbc, are characterized by an early period of merging and disk instability events which form the bulge, followed by an extended period of steady, uninterrupted disk growth. \\item The model results concur with several observational results including the type averaged galaxy merger rate, measurements of the specific star formation rate, the secular evolution of spiral bulges, the seemingly anti-heirarchical evolution of elliptical galaxies and the factor of ${\\sim}$ 2 growth in the luminosity density of elliptical galaxies since z ${\\sim}$ 1. With regard to the last point however, the model predicts significant evolution at the high luminosity end of the $K$-band elliptical galaxy LF that has yet to be confirmed observationally. \\end{itemize}"
    },
    "0911/0911.4355_arXiv.txt": {
        "abstract": "Our adaptive optics observations of nearby AGN at spatial resolutions as small as 0.085\\arcsec\\ show strong evidence for recent, but no longer active, nuclear star formation. We begin by describing observations that highlight two contrasting methods by which gas can flow into the central tens of parsecs. Gas accumulation in this region will inevitably lead to a starburst, and we discuss the evidence for such events. We then turn to the impact of stellar evolution on the further inflow of gas by combining a phenomenological approach with analytical modelling and hydrodynamic simulations. These complementary perspectives paint a picture in which all the processes are ultimately regulated by the mass accretion rate into the central hundred parsecs, and the ensuing starburst that occurs there. The resulting supernovae delay accretion by generating a starburst wind, which leaves behind a clumpy interstellar medium. This provides an ideal environment for slower stellar outflows to accrete inwards and form a dense turbulent disk on scales of a few parsecs. Such a scenario may resolve the discrepancy between the larger scale structure seen with adaptive optics and the small scale structure seen with VLTI. ",
        "introduction": "Many studies of active galactic nuclei (AGN), particularly those concerned with understanding the co-evolution of black holes and their host galaxies through cosmic time, rely on observations of quasars. The reason is simply that these are the most luminous of such objects and can be studied at all redshifts. While they form an important component of the AGN population, the luminosity function (\\cite[Hao et al. 2005]{hao05}) indicates that the number density of lower luminosity AGN is orders of magnitude greater. This is particularly true in the local universe because of the effects of downsizing, which means that at low redshift black hole growth occurs primarily in AGN with black hole masses $<10^8$\\,M$_\\odot$ (\\cite[Best et al. 2005]{bes05}) and Seyfert-like luminosities (\\cite[Hasinger et al. 2005]{has05}). As a result, with only a few exceptions such as Mkn\\,231 (\\cite[Davies et al. 2004]{dav04b}), detailed studies of nearby AGN have to focus on the lower luminosity AGN such as Seyferts.  This has important implications. Although most quasar activity may be fuelled via mergers of gas rich galaxies, Seyfert nuclei typically reside in (preferentially early type) spiral galaxies (\\cite[Ho 2008]{ho08}). As a result, Seyfert activity is likely to be fuelled via secular processes associated with disk evolution such as those outlined by \\cite{shl90} and \\cite{wad04}. One obvious part of the overall scheme is the role played by bars, which are known to drive gas inwards from the large-scale disk. Yet despite the many attempts to find a correlation between the presence of a bar and an AGN, the statistical outcome is, at best, for only a marginal link between these phenomena. This, perhaps, should not be surprising. Bars drive large amounts of gas inwards (of order 0.1\\,M$_\\odot$\\,yr$^{-1}$, \\cite[Regan \\& Teuben 2004]{reg04}), working on spatial scales of kpcs and timescales of Gyrs. In contrast, a Seyfert nucleus requires typically $\\lesssim0.01$\\,M$_\\odot$\\,yr$^{-1}$ of gas (\\cite[Jogee et al. 2006]{jog06}), and the relevant temporal and spatial scales are Myr and pc. The two regimes are orders of magnitude different in every respect. Thus, while it is well understood how bars can create a gas reservoir and a starburst in the central kiloparsec (and there is plenty of observational evidence to support the link between bars and both of these phenomena: \\cite[Sakomoto et al. 1999]{sak99}, \\cite[Laurikainen et al. 2004]{lau04}, and \\cite[Jogee et al. 2005]{jog05} are a few exmaples among many), this is only a single possible step in the process that leads to AGN fuelling.  It is now understood that the properties of the central kpc -- the circumnuclear region -- can play an important role in driving gas further inwards. But although observations have yielded many tantalising results, reaching definitive conclusions is difficult. For example, \\cite{mar03} looked at the detailed circumnuclear dust morphologies of matched samples of active and inactive galaxies. One of their results was that AGN are only found in galaxies that exhibit dusty spiral patterns. But they also found that spiral patterns of any sort were associated equally often with active and inactive galaxies. This, again, is likely a result of the differing spatial and temporal scales of the phenomena. The aspect of time variability is one of the issues addressed by NUGA (\\cite[Garc\\'ia-Burillo et al. 2007]{gar07}), a study of the distribution and kinematics of the molecular gas in nearby low luminosity AGN. One of their key results was that gravitational torques can not only drive gas in to form a circumnuclear ring, but also tend to drive gas out from the nuclear region towards the ring. This led \\cite{gar05} to propose a scenario in which further inward flow was possible episodically if two conditions were fulfilled: the bar should be weakened by dynamical feedback, and the viscosity should increase due to a sufficiently steep density gradient in the circumnuclear ring.  These examples show that there is still much we do not understand about the conditions under which gas can be transported inward from kpc scales to the nuclear region, and it is one of the topics addressed here: we describe two highly contrasting case studies where gas is being driven inwards to the central tens of parsecs. We also show evidence that, once there, the gas is likely to trigger a starburst. And we discuss the impact of the starburst on the fate of the gas, and a possible link to the sub-parsec scales. ",
        "conclusions": "We have presented a brief review of the current status of our work on the impact that star formation has on gas inflow to Seyfert nuclei. We are not yet able to follow gas all the way down to the AGN itself, and instead we have focussed on spatial scales from a few hundred parsecs down to the central parsec. Our main conclusions are:  \\begin{itemize} \\item We have seen 2 contrasting methods of gas inflow that feed a nuclear starburst. While showing how gas might reach the central few to tens of parsecs, this also raises questions about the sustainability of inflows, how far in they reach, and whether it is possible to define what a typical inflow mechanism or rate is.  \\item There are short intense starbursts in the central tens of parsecs around Seyfert nuclei. In general, starbursts will be an inevitable consequence of gas inflow.  \\item Stellar outflows play a key role in the evolution of the ISM on these scales. Specifically, fast outflows tend to blow out the diffuse ISM, while slow outflows can lead to efficient accretion of gas onto smaller scales. This may be a link between what is seen on 10--50\\,pc scales with adaptive optics, and what is observed on 0.1--1\\,pc scales with infrared interferometry.  \\item The molecular obscuring torus most likely consists of multiple structures on different scales that fulfil different roles. Since the state of the torus is strongly affected by the episodic star formation that occurs within it, we need to think of it as a dynamic entity.  \\end{itemize}"
    },
    "0911/0911.2234.txt": {
        "abstract": "{%----------------context (optional: leave void) } {%---------------------aims In this paper we study density cusps that may contain central black holes. The actual co-eval self-similar growth would not distinguish between the central object and the surroundings. } {%---------------------methods To study the environment of a growing black hole we seek descriptions of steady `cusps' that may contain a black hole and that retain at least a memory of self-similarity. We refer to the environment in brief as the `bulge' and on smaller scales, the `halo'.% This depends on whether the black hole is a true singularity or rather only %a core mass concentration. } {%---------------------results We find simple descriptions of the simulations of %simulated collisionless matter %in the process of examining the presence and growth of central masses by comparing predicted densities, velocity dispersions and distribution functions with the simulations. %However one remarkably simple distribution function that has a filled loss %cone describes a bulge that transits from a near black hole domain to an %outer `zero flux' regime. The transition passes through a region having a %$1/r$ density profile. The structure is likely to be developed at an early %stage in the growth of a galaxy. A central black hole is shown to grow %exponentially in this background with an e-folding time of a few million %years. In some cases central point masses may be included by iteration. We emphasize that the co-eval self-similar growth allows an explanation of the black hole bulge mass correlation between approximately similar collisionless systems.} {%-------------------conclusions (optional: leave void) (to be structured?) We have derived our results from first principles assuming adiabatic self-similarity and either self-similar virialisation or normal steady virialisation. We conclude that distribution functions that retain a memory of self-similar evolution provide an understanding of collisionless systems. The implied energy relaxation of the collisionless matter is due to the time dependence. Phase mixing relaxation may be enhanced by clump-clump interactions.} ",
        "introduction": "In a previous work (Henriksen et al., referred to here as paper \\cite{HLeDMcM09a}, and references hereafter) we discussed the % \\begin{comment} importance of the \\end{comment} {} relation between the formation of black holes (hereafter BH) and of galaxies. We presented distribution functions (DF) for the co-eval formation of spherically symmetric cusps and bulges, as formed by radially infalling, collisionless matter (e.g. dark matter or stars). We also discussed the influence of a centrally dominant mass, including a point mass BH.  Observations (\\cite{KR1995,Ma98,FM2000,Geb2000}) have established a strong correlation between the BH mass and the surrounding stellar bulge mass (or velocity dispersion), which we take to be an indication of co-eval growth. Such growth occurs in part during the dissipative baryon accretion by BH `seeds' in the AGN (Active Galactic Nuclei) phase, but there is as yet no generally accepted scenario for the origin of the seeds. Moreover, recent suggestions of very early very supermassive BHs (e.g. \\cite{Kurk2007}), together with changes in the normalization of the BH mass-bulge mass proportionality (relatively larger black holes at high red shift) (e.g. \\cite{Mai2007}), suggest an alternate early growth mechanism. Recently \\cite{PFP2008} have studied the possible size of the dark matter component in BH masses. They deduce that between 1\\% and 10\\% of the black hole mass could be due to dark matter.  We explored this possibility in the case of spherically symmetric radial infall (paper \\cite{HLeDMcM09a}) and in this paper we will extend it to anisotropic infall. We use the same technique of inferring reasonable distribution functions for collisionless matter from limits of the time dependent Collisionless Boltzmann (CBE) and Poisson set, and are aided by some high resolution simulations of the evolution without a dominant central mass. We add a dominant central mass either analytically or by iteration. Just as in the radial case, the loss cones are not empty for substantial growth. The relaxation of collisionless matter is due to the temporal evolution (including the radial orbit instability) and, in addition, to possible `clump-clump' (two clump) interactions (\\cite{H2009,MH2003}).  We use, in this paper as in the previous paper (\\cite{HLeDMcM09a}), the Carter-Henriksen (\\cite{CH91}) procedure to obtain a quasi-self-similar system of coordinates (Henriksen 2006a,2006b, hereafter \\cite{H2006,H2006A}). This allows writing the CBE-Poisson set with explicit reference to possible transient self-similar dynamical relaxation. In this way we can remain `close' to self-similarity just as the numerical simulations appear to do.  Studies of BH-density cusps originated with the problem of BH feeding (\\cite{P1972,BW76}), and with the notion of adiabatic growth (\\cite{Y1980,Q1995,MH2002}). Observations of the central Milky Way have detected a mainly isotropic density cusp with logarithmic power in the range $-1.1\\pm0.3$ (\\cite{G2009}). This is flatter than the adiabatic limit (\\cite{MH2002}) and in any case the adiabatic growth scenario does not produce the BH bulge correlations (ibid). All this has spurred the investigation of co-eval dynamical growth instead of adiabatic growth. Central cusps flatter than $-1.5$ can be created by tight binary BH systems formed in mergers that `scour' the stellar environment (e.g. \\cite{MS2006,NM99}). This process can produce log slopes as flat as  $-1$ or even $-0.5$, and it is supported by a strong correlation between nuclear BH mass and central luminosity deficit (\\cite{KB2009}).  However the correlation in itself only implicates the influence of the BH. It does not necessarily require the merger history, which in any case is unlikely to be the same for different galaxies. Consequently we explore in this paper whether cusps as flat as those resulting from scouring might also be produced during the dynamical formation of the black hole.  We begin the next section with a summary of the general formulation in spherical symmetry% \\begin{comment} and of useful results discussed in paper \\cite{HLeDMcM09a} \\end{comment} {}. Subsequently we study a system comprised of anisotropic non-radial orbits in spherical symmetry. % \\begin{comment} The results may apply to the outer (beyond the NFW scale radius) regions of dark matter halos. \\end{comment} {}Finally we give our conclusions. ",
        "conclusions": "In this paper, we have developed distribution functions that describe both dark matter bulges and a central black hole or at least a central mass concentration. We succeed mainly in describing the dark matter bulges.  In the discussion of cusps and bulges based on purely radial orbits (paper \\cite{HLeDMcM09a}), we were able to distinguish the Distribution function of Fridmann and Polyachenko % \\begin{comment} (\\ref{eq:FPDF}) \\end{comment} {} from that of Henriksen and Widrow% \\begin{comment} (\\ref{eq:steadyF}) \\end{comment} {}. The FPDF was found to describe accurately the purely radial simulations of isolated collisionless halos carried out in (\\cite{MacM2006}). Moreover a point mass could be added without changing the form of the DF, which allows a self-similar growth of a central bulge or black hole.  The final result concerning steady, self-similar radial orbits concerned the special case $a=1$. The DF is a Gaussian that has been found previously in coarse graining (\\cite{HLeD2002}). We included it there as a second example of a radial DF that produces an $r^{-2}$ density profile (\\cite{Mutka09}).  The inclusion of angular momentum here leads to more realistic situations. We re-derived the steady self-similar DFs from first principles in equation (\\ref{eq:anisopsteady}). We showed that these can be used to describe the simulated collisionless halos calculated in (\\cite{MacM2006}), if we use the value of $a\\approx0.72$ identified in (\\cite{H2007}) and take two limits. In one (\\ref{eq:DFE}) the DF is isotropic and describes approximately the most relaxed central region of the bulge. The other limit (\\ref{eq:DFJ}) describes the outer region and the transition across the NFW scale radius. The qualitative correspondence supports the relevancy of this family. The potential of a central mass can not be included exactly in these distribution functions, but iteration is possible.% \\begin{comment} along the lines discussed for radial bulges, but since the growth of a black hole in an isotropic central bulge is negligible, we have not pursued this avenue. \\end{comment} {} An example was given for the DF (\\ref{eq:DFE}).  In particular, velocity dispersions have been calculated and compared qualitatively with the results of a high resolution numerical simulation of an isolated halo. These correspond to reasonable densities when adiabatic self-similarity is employed. Unfortunately these do not apply to the vicinity of the black hole itself. However the eventual dominance of the black hole through the $r^{-1}$ potential becomes obvious, as seen in the iterated potential (\\ref{eq:phiiter}).  Simple forms for the DF have been suggested for the various regions of the halo. And given the dynamic co-eval growth of black-hole/halo and bulge, we explain the black-hole bulge mass correlation simply (\\ref{eq:bhmass}). This relies on the approximate similarity between different halos.  The forms that we find follow from the assumed self-similar path to equilibrium. Remarkably they are in accord with the general forms for the DF as presented in Stiavelli and Bertin (\\cite{StiavelliBertin}) and as modified in \\cite{MTJ}. That is, these distribution functions contain an inner inverted distribution in energy (referred to as negative temperature in \\cite{MTJ}) and a cut-off in the distribution in angular momentum farther from the centre. The scale of the inner region that is argued to coincide with NFW scale radius in this paper is the parameter $r_{a}$ in the earlier papers, so inner and outer have similar meanings.  The variation is a power-law generally in both energy and angular momentum, although in the very relevant case of an inverse square density profile, an exponential variation is possible. These conclusions follow from the assumed adiabatically self-similar evolution towards equilibrium. As such they tend to answer questions posed in \\cite{MTJ} such as whether these distributions arise naturally and if the variations are exponential or power-law. We now know however that part of the answer lies in the anisotropy produced by the radial orbit instability (\\cite{MWH2006}).  We have also found a self-similar generalization of the radial FPDF in equation (\\ref{eq:genFPDF}). Unfortunately this DF does not have the property of yielding a density that is independent of the potential and so a point mass potential is incompatible with self-similarity. It might describe a system at large radii with a large central mass.  As always $a=1$ must be treated separately and we give a derivation from first principles. The result is new and it generalizes the FPDF in the sense that $\\rho\\propto r^{-2}$ always. Although we can not place a central mass inside this bulge exactly, it allows an anisotropic DF in the $r^{-2}$ bulge region.  Generally we find that we can not describe elegantly anisotropic bulges containing black holes with self-similar DFs. % \\begin{comment} We point out in a non self-similar context that a sharp cut-off of any kind (whether due to black hole binary scouring or not) can yield a density profile as flat as $-1/2$. \\end{comment} {}  In the next paper in this series (paper \\cite{HLeDMcM09c}), we shall widen our scope to the study and production of non-self-similar cusps and DF."
    },
    "0911/0911.3269_arXiv.txt": {
        "abstract": "Coupled dark matter-dark energy systems can suffer from non-adiabatic instabilities at early times and  large scales. In these proceedings, we consider two parameterizations of the dark sector interaction. In the first one the energy-momentum transfer 4-vector is parallel to the dark matter 4-velocity and in the second one  to the dark energy 4-velocity. In these cases, coupled models which suffer from non-adiabatic instabilities can be identified as a function of a generic coupling $Q$ and of the dark energy equation of state $w$. In our analysis,  we do not refer to any particular cosmic field. We  confront then a viable class of models in which the interaction  is directly proportional to the dark energy density  and to the Hubble rate parameter to recent cosmological data. In that framework, we show that correlations between  the dark coupling and several cosmological parameters allow for a larger neutrino mass than in uncoupled models. ",
        "introduction": "\\label{subsec:intro}  Interactions between dark matter and dark energy  are still allowed by observational data today. At the level of the background evolution equations, one can generally introduce a coupling between these two sectors   as follows: \\begin{eqnarray} \\label{eq:EOMm} \\dot\\rho_{dm}+ 3\\mathcal{H}\\rho_{dm} &=& Q\\,,\\\\ \\label{eq:EOMe} \\dot\\rho_{de}+ 3 \\mathcal{H}\\rho_{de}(1+ w)&=&- Q\\,. \\end{eqnarray} $\\rho_{dm} (\\rho_{de})$ denotes the dark matter (dark energy) energy density,  the dot indicates derivative with respect to conformal time $d\\tau = dt/a$, $\\mathcal{H}= {\\dot a}/a$ and  $w=P_{de}/\\rho_{de}$ is the dark-energy equation of state ($P$ denotes the pressure). We work with the Friedman-Robertson-Walker (FRW) metric, assuming a flat universe and  pressureless dark matter $w_{dm} = P_{dm}/\\rho_{dm}=0$.   Q encodes the dark coupling and drives the energy exchange between dark matter and dark energy. For {\\it e.g.} $Q<0$ the energy flows from dark matter to dark energy. It also changes the  dark matter and dark energy redshift dependence acting as an extra contribution to their effective equation of state. For {\\it e.g.} $Q<0$, dark matter redshifts faster, as a consequence, there is  more dark matter in the past compared to uncoupled scenarios  assuming that the dark matter density today is the same in the two models (see also the discussion in Ref.~\\cite{CalderaCabral:2009ja}). This general feature of coupled models is sketched in Fig.~\\ref{fig:Q} (for a particular form of $Q$, see also Fig.~\\ref{fig:fishesplots}). \\begin{figure}[t] \\includegraphics[width=6cm]{plots/rho-09QN.ps} \\caption{\\it This figure illustrates roughly the effect of a negative dark coupling term $Q$ on dark matter and dark energy evolution. $a$ is the scale factor, $\\rho_i$ with $i=m,r,de $ corresponds  to matter (dark matter plus baryon), radiation  and dark energy densities respectively. Solid curve account for an uncoupled model with constant dark energy equation of state $w>-1$ and dotted curve for coupled models with $Q<0$ with $w>-1$.      } \\label{fig:Q} \\end{figure}     Also notice that a universe in accelerated expansion today requires $w<-1/3$ even in the presence of a dark coupling. Indeed, the deceleration parameter satisfies \\begin{equation} \\label{eq:decel} q=-  \\frac{\\dot {\\mathcal H} }{{\\mathcal H}^2}=\\frac12 (1+3\\,w\\,\\Omega_{de})~, \\end{equation} either with or without dark coupling. The curvature contribution has been neglected in this equation.  In order to deduce the evolution of  density and velocity perturbations in  coupled models, we need an expression of the energy transfer in terms of the stress-energy tensor: \\begin{eqnarray} \\nabla_\\mu T^\\mu_{(dm)\\nu} =Q_\\nu \\quad\\mbox{and}\\quad \\nabla_\\mu T^\\mu_{(de)\\nu} =-Q_\\nu. \\label{eq:conservDMDE} \\end{eqnarray} The  4-vector $Q_\\nu$ governs the  energy-momentum transfer between the dark components and  $T^\\mu_{(dm)\\nu}$ and $T^\\mu_{(de)\\nu}$ are the energy-momentum tensors for the dark matter and dark energy, respectively. Equation~(\\ref{eq:conservDMDE}) guaranties the conservation of the total energy-momentum tensor. In the following we consider two scenarios which reproduce at the background level Eqs.~(\\ref{eq:EOMm}) and~(\\ref{eq:EOMe}): \\begin{enumerate} \\item $Q_\\nu$ is parallel to the dark matter four velocity $u_{\\nu}^{(dm)}$: \\begin{equation} Q_\\nu = Q u_{\\nu}^{(dm)}/a. \\label{eq:um} \\end{equation} This choice of parameterization, which was first proposed in Ref.~\\cite{Valiviita:2008iv},  avoids momentum transfer in the rest frame of dark matter ({\\it i.e} $v^i_{(dm)}=0$), \\item $Q_\\nu$ is parallel to the dark energy four velocity $u_{\\nu}^{(de)}$: \\begin{equation} Q_\\nu = Q u_{\\nu}^{(de)}/a. \\label{eq:ue} \\end{equation} In this case the dark matter velocity (Euler) equation is modified by the coupling and the weak equivalence principle is violated (see {\\it e.g.} Ref~\\cite{Koyama:2009gd} for a recent discussion). Notice that one example for this interaction  is $Q_\\nu= \\beta \\rho_{dm} \\nabla_\\nu \\phi$ \\cite{Wetterich:1994bg,Amendola:1999qq}, where $\\phi$ would be  a coupled  dark energy scalar field and $u_{\\nu}^{(de)}\\propto \\nabla_\\nu \\phi$. \\end{enumerate}   It was first pointed out in Ref.~\\cite{Valiviita:2008iv} that the dark coupling terms which appears in the non-adiabatic dark energy pressure perturbations are a source for  early time instabilities  at large scales for coupled models satisfying Eqs.~(\\ref{eq:conservDMDE}) and~(\\ref{eq:um}). Their analysis was however restricted to $Q>0$ proportional to the dark matter density, in which case the coupled model is particularly unstable for a constant dark energy equation of state. Non-adiabatic instabilities were subsequently  analyzed  by several authors, and it was proved that the stability depends on the type of dark coupling $Q$, on the dark energy equation of state $w$ and on the $Q_\\nu$ 4-velocity dependence,   see Ref.~\\cite{Jackson:2009mz,He:2008si,Gavela:2009cy,Majerotto:2009np} (for similar instabilities pointed out in coupled quintessence models, see Ref.~\\cite{Corasaniti:2008kx,Chongchitnan:2008ry}).  Among these references, we proposed in Ref.~\\cite{Gavela:2009cy}, a criteria associated to the dubbed {\\bf \\it doom factor}  to identify the stability region of coupled models with a constant dark energy equation of state $w$ satisfying Eqs.~(\\ref{eq:conservDMDE}) and~(\\ref{eq:um}). The doom factor is a function of the model parameters such as $Q$ and $w$, but it is defined independently of the explicit form of the coupling $Q$. We review this result in the next section and extend it to  models satisfying Eqs.~(\\ref{eq:conservDMDE}) and~(\\ref{eq:ue}).  Based on this analysis, we study the compatibility  of  a successful class of models, in which $Q$ is proportional to the dark energy density, with a fixed dataset including WMAP 5 year~\\cite{Dunkley:2008ie,Komatsu:2008hk} , HST~\\cite{Freedman:2000cf}, SN~\\cite{Kowalski:2008ez}, $H(z)$~\\cite{Simon:2004tf} and LSS~\\cite{Tegmark:2006az} data. This work was done using the publicly available \\texttt{CAMB} code~\\cite{Lewis:1999bs} and \\texttt{cosmomc} package~\\cite{Lewis:2002ah}. The latter were  modified in order to include the interaction between the dark matter and dark energy components and the modified dark matter velocity equation  in the case of  coupled models satisfying Eq.~(\\ref{eq:ue}) (see Eq.~(\\ref{eq:thetaees})). Present data will be shown to allow for a sizeable interaction strength and to imply weaker cosmological limits  on neutrino masses with respect to non-interacting scenarios. ",
        "conclusions": "\\label{sec:concl} \\begin{figure}[t] \\vspace{-0.1cm} \\begin{tabular}{cc} \\includegraphics[width=8cm]{plots/oldplotnuw_v2.ps} & \\includegraphics[width=8cm]{plots/newplotnuw_v2.ps} \\\\ \\end{tabular} \\caption{\\it Scenario with $Q\\propto \\rho_{de}$. Left (right) panel: 1$\\sigma$ and 2$\\sigma$ marginalized contours in the $\\xi$--$w$  plane for $Q_\\nu\\propto u_\\nu^{(dm)}$  ($Q_\\nu\\propto u_\\nu^{(dm)}$). The largest, green contours show the current constraints from WMAP (5 year data), HST, SN and $H(z)$ data. The smallest, red contours show the current constraints from WMAP (5 year data), HST, SN, $H(z)$ and LSS data.} \\label{fig:fig2} \\end{figure}   In these proceedings, we discuss the origin of non-adiabatic instabilities in coupled models satisfying \\\\ $\\nabla_\\mu T^\\mu_{dm\\,\\nu}= Q u_\\nu^{dm\\,(de)}/a $ and  $\\nabla_\\mu T^\\mu_{de\\,\\nu}= Q u_\\nu^{dm\\,(de)}/a $. We show that the sign of the doom factor \\begin{equation} {\\bf d} \\equiv \\frac{Q}{3\\mathcal H\\rho_{de}(1+w)} \\end{equation} reveals the presence of instabilities. In particular, when ${\\bf d }$ is positive and sizeable, ${\\bf d }>1$, the dark-coupling dependent terms may dominate the evolution of the  dark energy perturbation, and drive an  unstable growth regime. Notice that {\\bf d} has been defined for a constant dark energy equation of state $w$ (see {\\it e.g.} Ref.~\\cite{Majerotto:2009np, Valiviita:2009nu} for $w=w(a)$).  We have  studied the constraints from current cosmological data on the dimensionless coupling $\\xi$ in a class of viable models in which $Q=\\xi {\\mathcal H} \\rho_{de}$, independently of the dark interaction term 4-velocity dependence. The analyses presented here were carried out in the  $\\xi<0$ and positive $(1+w)$ region of the parameter space, which offers the best agreement with data on large scale structure formation. From the results of our fits we find that both $w$ and $\\xi$ are not very constrained from data: it can be noticed from Fig.~\\ref{fig:fig2} that substantial values for both parameters, near -0.5, are easily allowed. Furthermore, $\\xi$ turns out to be positively correlated with  $\\Omega_{dm} h^2$ and a larger neutrino fraction $f_{\\nu}$ than in uncoupled models is allowed for negative values of the coupling $\\xi$.    \\begin{theacknowledgments} The work reported associated to the  $Q_\\nu\\propto u_\\nu^{(dm)}$ models has been done in collaboration with B.~Gavela, D.~Hernandez, O.~Mena and S.~Rigolin, see Ref.~\\cite{Gavela:2009cy}. L.~L.~H thanks the organizers of the Invisible Universe conference (Paris) for giving her the opportunity to give a talk on ``dark couplings'' and the I.F.I.C. group (Valencia), where part of the $Q_\\nu\\propto u_\\nu^{(de)}$ work was carried out, for its hospitality .  We also acknowledge R. Jim\\'enez B.~Reid and L. Verde for useful comments and discussions. L.~L.~H was partially supported by CICYT through the project FPA2006-05423, by CAM through the project HEPHACOS, P-ESP-00346,  by the PAU (Physics of the accelerating universe) Consolider Ingenio 2010, by the F.N.R.S. and  the I.I.S.N.. O.~M. work is supported by a Ram\\'on y Cajal contract from MEC, Spain  \\end{theacknowledgments}   \\appendix"
    },
    "0911/0911.3743_arXiv.txt": {
        "abstract": "We present the long-term monitoring of the High Mass X-ray Binary GX~301-2 performed  with the SuperAGILE instrument on-board the {\\it AGILE} mission. The source was monitored in the 20--60~keV energy band during the first year of the mission from 2007 July 17 to 2008 August 31, covering about one whole orbital period and three more pre-periastron passages for a total net observation time of about 3.7~Ms. The SuperAGILE dataset represents one of the most continuous and complete monitoring at hard X-ray energies of the 41.5~day long binary period available to date.  The source behavior was characterized at all orbital phases in terms of hard X-ray flux, spectral hardness, spin period history, pulsed fraction and pulse shape profile.  We also complemented the SuperAGILE observations with the soft X-ray data of the RossiXTE/ASM. Our analysis shows a clear orbital modulation of the spectral hardness, with peaks in correspondence with the pre-periastron flare and near phase 0.25. The hardness peaks we found could be related with the wind-plus-stream accretion model proposed in order to explain the orbital light curve modulation of GX~301-2.  Timing analysis of the pulsar spin period shows that the secular trend of the $\\sim$680~s pulse period is consistent with the previous observations, although there is evidence of a slight decrease in the spin-down rate. The analysis of the hard X-ray pulsed emission also showed a variable  pulse shape profile as a function of the orbital phase, with substructures detected near the passage at the periastron, and a clear modulation of the pulsed fraction, which appears in turn strongly anti-correlated with the source intensity. ",
        "introduction": "GX~301-2 is a High Mass X-ray Binary (HMXB) system containing an X-ray pulsar, with a $\\sim$685~s rotation period, orbiting the B-emission line hypergiant star (B1 Ia+) Wray 977. The X-ray source was discovered in 1969 April \\citep{Lewin1971} with a balloon observation, subsequently \\cite{White1976} discovered the 700~s X-ray pulsation. The pulsar orbit has an eccentricity of  0.462 \\citep{Koh1997} and a period of $\\sim$41.5~days \\citep{White1978, Sato1986, Koh1997, Doroshenko2008}. The mass function of the system is 31.1~M$_\\odot$ with a companion mass in the range 39~$<$~M~$<$~53~M$_\\odot$, estimated by the measurement of optical radial velocity amplitude \\citep{Kaper2006}. The radius of Wray 977 obtained by  \\cite{Kaper2006} by fitting atmosphere model is 62~$R_\\odot$, while the effective temperature is about 1.8~$\\times$~10$^4$~${\\rm K}$.  GX~301-2 shows regular X-ray flares about two days before the periastron passage, while a secondary intensity peak was sometimes observed near the apastron passage \\citep{Pravdo1995, Koh1997, Pravdo2001}. Assuming a distance of 3~kpc, as derived by \\cite{Kaper2006}, the source luminosity during pre-periastron (PP) flares reaches values of $\\sim 10^{37}$~erg$ \\: $s$^{-1}$, about 25 times larger than the luminosity showed by the source in the other orbital phases.  Since its discovery, GX~301-2 exhibited unpredictable torque variations superimposed on a secular spin trend, which passed from a spin-up to a spin-down state between MJD~48500 and MJD~49500. Two rapid spin-up episodes with $\\dot{\\nu} \\sim 2\\mbox{--}5 \\times 10^{-12}\\, {\\rm Hz}\\, {\\rm s}^{-1}$ were also found by \\cite{Koh1997} analyzing BATSE data, thus confirming the erratic behavior of the pulsar spin period. \\cite{Koh1997} suggested that the long-term spin-up trend observed between 1984 and 1994 might not be regular and continuous, but due to several rapid spin-up episodes, which in turn could be related to the formation of a temporary accretion disk. \\cite{Pravdo2001}, using BATSE data, found evidence that the hard  X-ray (20-50~keV) intensity of the binary system during quiescent orbital phases is correlated with the pulse period, suggesting a possible relation between the spin period behavior and the mass accretion rate.  Several wind accretion models were proposed to explain both the regular PP and apastron flaring activities of the source. A dense disk around Wray~977, which intercepts the neutron star (NS) orbit near the periastron and the apastron, was proposed by \\cite{Pravdo1995} and \\cite{Pravdo2001}. However, this model seems unable to provide a satisfactory modeling of the observed orbital X-ray light curve of the source \\citep{Leahy2002} and optical observations by \\cite{Kaper2006} found no confirmation of the presence of a disc. The existence of a high density stream of matter in addition to the stellar wind was proposed by \\cite{Haberl1991, Leahy1991, Leahy2002,  Leahy2008}. The stream outflows from Wray~977 due to tidal interaction intercepts the NS orbit near the periastron and the apastron, causing a large increase in the mass accretion rate with a sudden enhancement of the X-ray luminosity from the neutron star. Signatures of the presence of such a stream were found by \\cite{Saraswat1996} analyzing the low-energy excess in the {\\it ASCA} spectrum of GX~301-2, and also by \\cite{Kaper2006} from the orbital modulation of the spectral lines detected in the stellar wind of the optical companion. The wind-plus-stream model predicts also a large variation in the column density as a function of the orbital phase. The $N_{\\mathrm{H}}$ is expected to change from a value of $\\sim 10^{23}\\,$cm$^{-2}$ at orbital phase 0.8 to values greater than 10$^{24}\\,$cm$^{-2}$ during the PP flares and between orbital phases 0.2--0.3.  We observed the HMXB GX~301-2 with the SuperAGILE (SA) experiment onboard the  Astro-rivelatore Gamma ad Immagini LEggero ({\\it AGILE}) satellite during the first year of the mission. {\\it AGILE}, launched on 2007 April 23, is a scientific mission of the Italian Space Agency dedicated to high-energy astrophysics \\citep{Tavani2009}. In the following sections we report on the analysis of the SuperAGILE data, in the energy range 20--60 keV. In Section~\\ref{sec:obs_and_data} we describe the SuperAGILE experiment and the observations of GX~301-2 we performed, in Section~\\ref{sec:timing} we present the timing analysis of the pulsar emission and in Section~\\ref{sec:orb_lc} we analyze the orbital light curve of GX~301-2 using the hard X-ray SuperAGILE and the soft X-ray {\\it RXTE}/ASM data and the orbital variability of the source pulsed fraction. Discussion of the data analysis and the interpretation of the results are presented in Section~\\ref{sec:discussion}, while in Section~\\ref{sec:conclusion} we draw our conclusions. ",
        "conclusions": "\\label{sec:conclusion}  We studied the long term behavior of the HMXB GX~301-2 using the hard X-ray data collected by SuperAGILE during the first year of the {\\it AGILE} mission. The observations, carried out in the 20--60~keV energy band, have a large net exposure, about 3.7~Ms, and cover all the orbital phases of the binary system. The SuperAGILE data offer one of the most complete hard X-ray monitoring of the 41.5-day long binary period available to date. The data span covers large fractions of six orbital cycles. The secular trend of the $\\sim$680 s pulse period is consistent with the previous observations, although the SuperAGILE data show a decrease in the spin-down rate with respect to the observations reported by \\cite{Doroshenko2008}. The spin-down trend is approximately constant, but there is indication of a significant timing noise, of the order of $\\sim$0.2\\%.  Complementing the SuperAGILE data with those of {\\it RossiXTE}/ASM, the source behavior was characterized at all the phases of the binary orbit, in terms of its soft and hard X-ray flux, spectral hardness, spin period history, pulsed fraction and pulse shape. Compared with the soft X-ray data, the hard X-ray lightcurve exhibits a smaller modulation of the emission during the orbital phases 0.1--0.8 and a larger intensity of the pre-periastron flare. This results in a spectral hardness which presents clear orbital modulation, with peaks in correspondence with the pre-periastron flare and near phase 0.25. The timing analysis of the hard X-ray emission showed a variable pulse shape profile as a function of the orbital phase, with substructures detected near the passage at the periastron, and a clear modulation of the pulsed fraction, which appears in turn strongly anti-correlated with the source intensity.  Overall, the scenario depicted by the SA data appears generally consistent with the model for this binary system proposed by \\cite{Leahy2008}. The predictions of the model are based on the variability in the accretion rate caused by the orbital motion of the NS in the environment of the matter outflow from the companion star. The model is suited to the low energy observations by {\\it RXTE}/ASM, for which it reproduces the peaked flare at the pre-periastron and the broader peak at the apastron. They both are related to the NS crossing of the accretion stream, the difference between the two being mostly in the local density and geometrical size of the stream at the crossing point. The higher accretion rate causes a soft X-ray brightening, and scattering of the radiation by the same matter would cause loss of coherence and then a decrease in the pulsed fraction. In the hard X-rays we do not detect the broad peak at the apastron phase (although other hard X-ray observations with BATSE and {\\it Swift}/BAT suggest time variability for it). Also, our observation of the folded hardness ratio (20--50~keV to 2--12~keV) along the orbit is not fully consistent with the model expectations. Indeed the model predicts an increase in the column density near the apastron phase, but it does not predict a similar increase near the pre-periastron, that we instead clearly detect.  More in general, our results fit in the overall scenario of the binary X-ray pulsars \\citep[e.g.][]{Parmar1989}, where the timing properties are related to the source luminosity. As a matter of fact, our analysis shows that the pulsed fraction is clearly anti-correlated with the source intensity, allowing us to attribute the accretion rate the role of driving parameter in the system. However, while the timing properties appear rather well correlated with the source intensity, the same is not true for the hardness ratio, for which we cannot identify a trend with the flux. This implies that the emitting process and/or geometry change along the orbit in a way that is not immediately associated to the variation of the accretion rate only."
    },
    "0911/0911.4814_arXiv.txt": {
        "abstract": "We have predicted theoretically and detected in laboratory experiments a new type of particle clustering (namely, tangling clustering of inertial particles) in a stably stratified turbulence with imposed mean vertical temperature gradient. In the stratified turbulence a spatial distribution of the mean particle number density is nonuniform due to the phenomenon of turbulent thermal diffusion, i.e., the inertial particles are accumulated in the vicinity of the minimum of the mean temperature of the surrounding fluid, and a non-zero gradient of the mean particle number density, $\\bec\\nabla N$,  is formed. It causes  generation of fluctuations of the particle number density by tangling of the large-scale gradient, $\\bec\\nabla N$, by velocity fluctuations. In addition, the mean temperature gradient, $\\bec\\nabla T$, produces the temperature fluctuations by tangling of the large-scale gradient, $\\bec\\nabla T$, by velocity fluctuations. The anisotropic temperature fluctuations contribute to the two-point correlation function of the divergence of the particle velocity field, i.e., they increase the rate of formation of the particle clusters in small scales. We have demonstrated that in the laboratory stratified turbulence this tangling clustering is much more effective than a pure inertial clustering (preferential concentration) that has been observed in isothermal turbulence. In particular, in our experiments in oscillating grid isothermal turbulence in air without imposed mean temperature gradient, the inertial clustering is very weak for solid particles with the diameter $\\approx 10 \\, \\mu$m and Reynolds numbers based on turbulent length scale and rms velocity, ${\\rm Re} =250$. In the experiments the correlation function for the inertial clustering in isothermal turbulence is much smaller than that for the tangling clustering in non-isothermal turbulence.  The size of the tangling clusters is of the order of several Kolmogorov length scales. The clustering described in our study is found for inertial particles with small Stokes numbers and with the material density that is much larger than the fluid density. Our theoretical predictions are in a good agreement with the obtained experimental results. ",
        "introduction": "Laboratory experiments \\cite{EF94,FK94,WA00}, numerical simulations \\cite{WM93,MC96,SC97,S01} and atmospheric observations \\cite{KM93,VY00,KS01,S03,BH03,KPE07,LS07} revealed small-scale long-living inhomogeneities (clusters) in spatial distribution of particles in different turbulent flows. The origin of these inhomogeneities is not always clear but their influence on particle transport and mixing can be hardly overestimated. In particular, turbulence causes formation of small-scale droplet inhomogeneities, increases the relative droplet velocity and affects the hydrodynamic droplet interactions. All these effects can enhance the rate of droplet collisions (see reviews \\cite{VY00,S03,KPE07}). It is known that atmospheric clouds are regions with strong turbulence, and the turbulence effects are of a great importance for understanding of rain formation in atmospheric clouds, e.g., these effects can cause the droplet spectrum broadening and acceleration of raindrop formation \\cite{S03,VY00,KPE07,RC00}.  Different kinds of particle clustering, i.e., the preferential concentration of inertial particles have been studied in a number of numerical simulations \\cite{HC01,B03,BL04,CK04,FP04,CC05,BE05,CG06,BC07,BB07,YG07,AC08,ACW08,GP09,BB10}, laboratory experiments (see review \\cite{WA09} and \\cite{AC02,WH05,AG06,DF06,SA08,XB08,SSA08}) and analytical investigations \\cite{EKRC96,EKRC00,BF01,EKR02,ZA03,DM05,MW05,WM05,WM06,GDH05,EKR07,OV07,FM07,FH08,OL10}.  Inertial clustering in isothermal and nearly isotropic turbulence with the Reynolds numbers ${\\rm Re} = u_0 \\, \\ell_0 / \\nu \\sim 10^3$ produced by fans at the corners of the box has been observed in laboratory experiments \\cite{SA08} with hollow glass spheres with a mean diameter of $6 \\, \\mu$m in air. Here $u_0$ is the characteristic turbulent velocity in the maximum scale of turbulent motions $\\ell_0$ and $\\nu$ is the kinematic viscosity. The three-dimensional radial distribution function (RDF) as a statistical measure of clustering, has been determined from the particle position field using three-dimensional digital holographic particle imaging. These experiments reveled inertial particle clustering at the dissipation scales. The detailed comparisons between experiments and direct numerical simulations of inertial particle clustering in \\cite{SA08} have demonstrated a good agreement.  The quantitative measurements of inertial clustering have been also performed in \\cite{WH05,SSA08}. In particular, the two-dimensional RDF of particles suspended in a turbulence flow with the Reynolds numbers ${\\rm Re} \\sim 10^3$ has been measured in \\cite{WH05} by shining a laser sheet at the particles (the glass particles with the sizes $20 \\, \\mu$m and $50 \\, \\mu$m, the lycopodium particles with the size $25 \\, \\mu$m), whereby particle locations have been determined by a CCD camera. The experiments have been conducted in the spherical turbulence chamber with eight synthetic jet actuators which create homogeneous and isotropic turbulence with no mean flow at the center of the chamber. The experiments conducted in \\cite{WH05} have shown that the inertial clustering is most effective for particles with Stokes numbers near unity and the size of the clusters is of the order of ten Kolmogorov length scales.  A one-dimensional RDF has been measured in \\cite{SSA08} by sampling droplet arrivals at a fixed volume in a wind tunnel, whereby the arrival statistics of water droplets with the mean size $22 \\, \\mu$m in turbulence with the Reynolds numbers ${\\rm Re} \\sim 10^4$ has been used to determine the one-dimensional RDF. In the experiments described in \\cite{SSA08} a phase Doppler interferometer downstream of the active grid and a spray system have been used, and strong inertial clustering has been observed. Remarkably, the similar experimental findings have been reported in \\cite{WH05,SA08} using completely different experimental set-ups.  The mechanism of inertial clustering of particles in isothermal turbulence is as follows \\cite{M87}. The particles inside the turbulent eddies are carried out to the boundary regions between the eddies by the inertial forces. This mechanism of the inertial clustering  acts in all scales of turbulence, and is more pronounced in small scales.  The goal of this study is to investigate experimentally and theoretically a new type of particle clustering, namely tangling clustering of inertial particles in stably stratified turbulence with imposed mean vertical temperature gradient. In experimental study of particle clustering we use Particle Image Velocimetry to determine the turbulent velocity field, an image processing technique to determine the spatial distribution of particles, and a specially designed temperature probe with twelve sensitive thermocouples for measurements of the temperature field. The clustering described in our study is found for inertial particles with small Stokes numbers and with material density that is much larger than the fluid density.  In the stratified turbulence a spatial distribution of the mean particle number density is nonuniform due to phenomenon of turbulent thermal diffusion \\cite{EKR96,EKR00}. In particular, the inertial particles are accumulated in the vicinity of the minimum of the mean temperature of the surrounding fluid, which causes formation a non-zero gradient of the mean particle number density, $\\bec\\nabla N$. The phenomenon of turbulent thermal diffusion has been predicted theoretically in \\cite{EKR96}, detected in the laboratory experiments in stably and unstably stratified turbulent flows in \\cite{EEKR04,BEE04,EEKR06a,EEKR06b} and observed in atmospheric turbulence in \\cite{SEKR09}.  Fluctuations of the particle number density can be generated by tangling of the gradient, $\\bec\\nabla N$, of the mean particle number density by velocity fluctuations \\cite{EKR95}. On the other hand, the imposed mean temperature gradient, $\\bec\\nabla T$, results in generation of the anisotropic temperature fluctuations by tangling of this large-scale gradient, $\\bec\\nabla T$, by velocity fluctuations. These temperature fluctuations may contribute to the two-point correlation function of the divergence of the particle velocity field. The latter enhances the rate of formation of particle clusters in small scales by the tangling mechanism.  The tangling mechanism is universal and independent of the way of generation of turbulence for large Reynolds numbers. For instance, tangling of the gradient of the large-scale velocity shear produces anisotropic velocity fluctuations \\cite{L67,WC72} which are responsible for different phenomena: formation of large-scale coherent structures in a turbulent convection \\cite{EKRZ02}, excitation of the large-scale inertial waves in a rotating inhomogeneous turbulence \\cite{EGKR05}, generation of large-scale vorticity \\cite{EKR03} and large-scale magnetic field \\cite{RK03} in a sheared turbulence.  The paper is organized as follows. Section II describes the experimental set-up for a laboratory study of the tangling clustering of inertial particles in stably stratified turbulence. The data processing and experimental results are presented in Section III. The theoretical analysis and comparison with experimental results are performed in Section IV. Finally, conclusions are drawn in Section V. ",
        "conclusions": "In the present experimental and theoretical study we have found a new type of particle clustering, tangling clustering of inertial particles, that occurs in a stably stratified turbulence with imposed mean vertical temperature gradient. Fluctuations of the particle number density are generated by tangling of the large-scale gradient, $\\bec\\nabla N$, by velocity fluctuations. The gradient of the mean particle number density is formed in the stratified turbulence with imposed mean vertical temperature gradient due to the phenomenon of turbulent thermal diffusion. The tangling clustering of inertial particles is also enhanced by the additional tangling of the mean temperature gradient by velocity fluctuations, that results in generation of the temperature fluctuations. Therefore, the heat flux due to the imposed mean temperature gradient, plays a twofold role, i.e., it causes formation of (i) the gradient of the mean particle number density and (ii) a non-zero two-point correlation function of the divergence of the particle velocity field, $B({\\bf R}) = \\langle \\tau b({\\bf x}) \\, b({\\bf y}) \\rangle$, where $b =$ div $\\, {\\bf v}$. The latter is the main source of the particle clustering in small scales.  There are two contributions to the correlation function $B({\\bf R})$. The first contribution is caused by particle inertia, so that the particles inside the turbulent eddies are carried out to the boundary regions between the eddies by the inertial forces. This is the main mechanism of the inertial clustering in isothermal turbulence. The second contribution to the correlation function $B({\\bf R})$ is due to the temperature fluctuations produced by tangling of the mean temperature gradient by velocity fluctuations. The latter results in the tangling clustering of inertial particles. It must be emphasized that the heat flux and particle inertia are two important ingredients which cause the tangling clustering.  We have shown that in the laboratory stratified turbulence the tangling clustering is much more stronger than both, the inertial clustering and the gravitational-tangling clustering occurring in isothermal turbulence. In particular, in our experiments the correlation function for the particle clustering in isothermal turbulence is much smaller than that for the tangling clustering in non-isothermal turbulence. In the stably stratified turbulence with imposed mean temperature gradient, we have found particle clusters with the size that is of the order of several Kolmogorov length scales.  The clustering described in our study is found for inertial particles with small Stokes numbers and with material density that is much larger than the fluid density. The Reynolds numbers, based on turbulent length scale and rms velocity, ${\\rm Re} =250$, in our experiments is smaller than that in experiments in isothermal turbulence described in \\cite{SA08,WH05,SSA08}. Probably these are the reasons for weak inertial clustering which has been observed in our experiments. It must be emphasized that inertial and tangling clusterings in our study are understood in the statistical sense, i.e., we employ an ensemble averaging over many images with the instantaneous particle distributions rather than analyze a single instantaneous image.  In the present study we have developed a theory of the tangling clustering of inertial particles that is based on the analysis of the two-point second-order correlation function of the particle number density. The theoretical predictions are in a good agreement with the obtained experimental results.  \\medskip"
    },
    "0911/0911.3928_arXiv.txt": {
        "abstract": "We focus on a comparison of the space densities of FRI and FRII extended radio sources at different epochs, and find that FRI and FRII sources show similar space density enhancements in various redshift ranges, possibly implying a common evolution. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0911/0911.5603_arXiv.txt": {
        "abstract": "{} {Deviations from axial symmetry are necessary to maintain self-sustained MRI-turbulence. We define the parameters region where nonaxisymmetric MRI is excited and study dependence of the unstable modes structure and growth rates on the relevant parameters.} {We solve numerically the linear eigenvalue problem for global axisymmetric and nonaxisymmetric modes of standard-MRI in Keplerian disks with finite diffusion.} {For  small magnetic Prandtl number the microscopic viscosity  completely drops out from the analysis so that the stability maps and the growth rates expressed in terms of the magnetic Reynolds number Rm and the Lundquist number S do not depend on the magnetic Prandtl number Pm. The minimum magnetic field for onset of nonaxisymmetric MRI grows with Rm. For given S all nonaxisymmetric modes disappear for sufficiently high Rm. This behavior is a consequence of the radial fine-structure of the nonaxisymmetric modes resulting from the winding effect of differential rotation. It is this fine-structure which presents severe resolution problems for the numerical simulation of MRI at large Rm.} {For weak supercritical magnetic fields only axisymmetric modes are unstable. Nonaxisymmetric modes need stronger fields and not too fast rotation. If Pm is small its real value does not play any role in MRI.} ",
        "introduction": "The leading concept for the origin of turbulence in accretion disks is the turbulence driving by the magnetorotational instability (MRI). The instability can be excited by even a very weak magnetic field provided that rotation with outward decreasing angular velocity is present (see Balbus \\& Hawley \\cite{BH98}).  MRI is expected to possess the  remarkable property of being self-sustained, i.e. to support the destabilizing magnetic field via its own dynamo (Brandenburg et al. \\cite{BNST95}; Hawley et al. \\cite{HGB96}).  The MRI-driven dynamo needs a finite but small initial field to start.  Differential rotation alone cannot drive a dynamo (Elsasser \\cite{E46}). Deviations from axial symmetry are necessary for dynamo and the deviations have to be produced by the MRI itself. The excitation of nonaxisymmetric modes of MRI is thus necessary for the self-sustained turbulence.  The present paper focuses on the nonaxisymmetric MRI. A model of differentially rotating disk with finite diffusivities and axial background field is applied to perform linear analysis of global stability. As in the axisymmetric case, the nonaxisymmetric MRI can be found in a range between some minimum $B_\\mathrm{min}$ and maximum $B_\\mathrm{max}$ values of the background field. The instability range depends on rotation rate (parameterized by the magnetic Reynolds number Rm). In contrast to the axisymmetric case, $B_\\mathrm{min}$ for nonaxisymmetric modes does not approach a (small) constant value with increasing rotation but grows in linear proportion to Rm. The larger Rm, the stronger background field is required to maintain the  nonaxisymmetric instability.  This behavior is increasingly difficult to follow numerically as Rm grows because of the winding effect of differential rotation. The shearing of nonaxisymmetric fields by differential rotation produces fine radial structure when the field is too weak to resist the winding. This shearing effect can simultaneously be the reason for the increase of $B_\\mathrm{min}$ with Rm and for the high resolution required to resolve the nonaxisymmetric MRI numerically (Fromang \\& Papaloizou \\cite{FP07}). ",
        "conclusions": ""
    },
    "0911/0911.0483_arXiv.txt": {
        "abstract": "We calculate the properties of neutron matter and highlight the physics of chiral three-nucleon forces. For neutrons, only the long-range $2 \\pi$-exchange interactions of the leading chiral three-nucleon forces contribute, and we derive density-dependent two-body interactions by summing the third particle over occupied states in the Fermi sea. Our results for the energy suggest that neutron matter is perturbative at nuclear densities. We study in detail the theoretical uncertainties of the neutron matter energy, provide constraints for the symmetry energy and its density dependence, and explore the impact of chiral three-nucleon forces on the $S$-wave superfluid pairing gap. ",
        "introduction": "The physics of neutron matter ranges over exciting extremes: from universal properties at low densities~\\cite{dEFT,Gezerlis} that can be probed in experiments with ultracold atoms~\\cite{coldatoms}; to using neutron matter properties at nuclear densities to guide the development of a universal density functional~\\cite{DFT1,DFT2} and to constrain the physics of neutron-rich nuclei; to higher densities involved in the structure of neutron stars~\\cite{LP}. In the theory of nuclear matter, recent advances~\\cite{nucmatt,chiralnm} are based on systematic chiral effective field theory (EFT) interactions~\\cite{chiral,RMP} combined with a renormalization group (RG) evolution to low momenta~\\cite{Vlowk,smooth}. This evolution improves the convergence of many-body calculations~\\cite{nucmatt,% NCSM,Sonia} and the nuclear matter energy shows saturation with controlled uncertainties~\\cite{chiralnm}. In this paper, we extend theses developments to neutron matter with a focus on three-nucleon (3N) forces.  Our studies are based on evolved nucleon-nucleon (NN) interactions at next-to-next-to-next-to-leading order (N$^3$LO)~\\cite{N3LO,N3LOEGM} and on the leading N$^2$LO 3N forces~\\cite{chiral3N1,chiral3N2}. In Sect.~\\ref{densdep}, we show that only the long-range $2\\pi$-exchange 3N interactions contribute in pure neutron matter. We then construct density-dependent two-body interactions $\\vbar$ by summing the third particle over occupied states in the Fermi sea. Effective interactions of this sort have been studied in the past using 3N potential models and approximate treatments (see for example, Refs.~\\cite{3Nold1,3Nold2}). We derive a general operator and momentum structure of $\\vbar$ and analyze the partial-wave contributions and the density dependence of $\\vbar$. This provides insights to the role of chiral 3N forces in neutron matter.  In Sect.~\\ref{results}, we apply $\\vbar$ to calculate the properties of neutron matter as a function of Fermi momentum $\\kf$ (or the density $\\rho = \\kf^3/(3\\pi^2)$) based on a loop expansion around the Hartree-Fock energy. Our second-order results for the energy suggest that neutron matter is perturbative at nuclear densities, where N$^2$LO 3N forces provide a repulsive contribution. We study in detail the theoretical uncertainties of the neutron matter energy and find that the uncertainty in the $c_3$ coefficient of 3N forces dominates. Other recent neutron matter calculations lie within the resulting energy band. In addition, the energy band provides constraints for the symmetry energy and its density dependence. Finally, we study the impact of chiral 3N forces on the $^1$S$_0$ superfluid pairing gap at the BCS level. We conclude and give an outlook in Sect.~\\ref{concl}. ",
        "conclusions": "\\label{concl}  In summary, we have shown that only the $c_1$ and $c_3$ terms of the long-range $2\\pi$-exchange part of the leading chiral 3N forces contribute in neutron matter. Based on these parts and including all exchange terms, we derived density-dependent two-body interactions $\\vbar$ for general momentum and spin configurations by summing the third particle over occupied states in the Fermi sea. The resulting $\\vbar$ was found to be dominated by the repulsive central part. The comparison to the exact Hartree-Fock energy demonstrated that the $P$ dependence is very weak and we have therefore taken $P=0$ in $\\vbar$. In addition, the partial-wave matrix elements of our full $\\vbar$ approximately agree with the density-dependent two-body forces derived recently in the framework of in-medium chiral perturbation theory with certain approximations in Ref.~\\cite{Holt}.  The density-dependent two-body interactions $\\vbar$ correspond to the normal-ordered two-body part of 3N forces and we stressed that there are different normal ordering or symmetry factors when $\\vbar$ is used at the Hartree-Fock, one- or two-body level. Moreover, it is important that the summation of the third particle over occupied states is performed in the basis consistent with the one used in the many-body calculation.  We have presented the first results for neutron matter based on chiral EFT interactions and including N$^2$LO 3N forces. The RG evolution of N$^3$LO NN potentials to low momenta rendered the many-body calculation more controlled and our results for the energy suggest that neutron matter is perturbative at nuclear densities. This is based on small second-order contributions, self-energy corrections being negligible, and a generally weak cutoff dependence. We have found that N$^2$LO 3N forces provide a repulsive contribution to the energy due to the repulsive central part in $\\vbar$.  An important direction for future work is understanding the connection to the repulsive 3N contributions to the oxygen isotopes discovered in Ref.~\\cite{oxygen}.  We have studied in detail the theoretical uncertainties of the neutron matter energy and found that the uncertainty in the $c_3$ coefficient of 3N forces dominates compared to other many-body uncertainties. This resulted in an energy band in Fig.~\\ref{fig:EOS_compare}. Other recent neutron matter calculations were found to lie within this band. Moreover, the energy band provides microscopic constraints for the symmetry energy and its density dependence, and thus for the neutron skin in $^{208}$Pb~\\cite{skin}. Finally, we have obtained a significant reduction of the $^1$S$_0$ superfluid pairing gap due to 3N forces for densities where the gap is decreasing. Our results show that chiral EFT and RG interactions can provide useful constraints for developing a universal density functional based on microscopic interactions~\\cite{DFT1,DFT2,Wolfram}."
    },
    "0911/0911.5573_arXiv.txt": {
        "abstract": "We investigate strongly interacting dense matter and neutron stars using a flavor-SU(3) approach based on a non-linear realization of chiral symmetry as well as a hadronic flavor-SU(2) parity-doublet model. We study chiral symmetry restoration and the equation of state of stellar matter and determine neutron star properties using different sets of degrees of freedom. Finally, we include quarks in the model approach. We show the resulting phase diagram as well as hybrid star solutions for this model. ",
        "introduction": "The study of strongly interacting matter at extreme conditions of temperature and density is at the forefront of modern nuclear physics. This regime covers the physics of ultrarelativistic heavy-ion collisions as well as important aspects of nuclear astrophysics. Whereas the determination of the phase structure of excited strongly interacting matter at high temperature and baryon densities, as pursued in heavy-ion collisions, involves relatively high temperatures, the complementary study of the structure of compact stars is directly related to the properties of very dense and rather cold matter.  In order to study all these regimes in a unified model approach we developed an effective chiral SU(3) model that can be studied over the whole relevant range of chemical potentials and temperatures. Reflecting the correct degrees of freedom of the system after the transition to a quark-gluon plasma state we include quarks and the Polyakov loop as order parameter for the deconfinement phase transition. ",
        "conclusions": ""
    },
    "0911/0911.0951_arXiv.txt": {
        "abstract": " ",
        "introduction": "How does a galaxy such as our Milky Way form? This remains a pressing question for modern astrophysics. The seminal work of \\citet{ELS} proposed that galaxy formation occurred in the wake of a monolithic collapse of an isolated proto-galactic cloud. \\citet{SZ} challenged this picture of isolated formation by proposing instead that the halo was assembled by the accretion of numerous small entities over an extended period.  Recent observational evidence suggests both formation scenarios have a role to play. \\citet{Carollo07} presents evidence that the inner halo (that of galactocentric radius $<15$kpc) is dominated by highly eccentric, prograde orbits and a metallicity of [Fe/H] $\\sim -1.6$ due to in-situ formation, whereas the outer halo shows a more uniform distribution of eccentricities, includes highly retrograde orbits, and lower mean metallicity ([Fe/H] $\\sim -2.2$) that, they argue, has an accretion origin. The \\citet{Kinman07} study of the local halo, similarly identifies a population that possesses retrograde rotation and streaming motions (i.e.\\ low velocity dispersion) and another population that exhibits negligible rotation and less indication of streaming. \\citet{Kinman07} show that the horizontal branch morphologies of these two halo components are consistent with those seen in the young and old globular cluster systems respectively. This reinforces the conclusion of \\citet{Zinn85} that the inner halo is consistent with formation from rapid collapse and that the outer halo was formed from ongoing accretion of dwarf galaxies.  Simulations of galaxy formation in the context of the prevailing $\\Lambda$CDM cosmology predict that the majority of stars in the outer halo formed in progenitor dwarf galaxies that were subsequently accreted (see for example, \\citet{Freeman02}). \\citet{Bullock05} show that spatially coherent substructures formed of tidal debris should survive in the outer halo for several Gyrs and that their kinematic signatures (in velocity space) should remain detectable well after spatial substructure is no longer evident \\citep{Helmi03}. On the other hand, a halo that formed from in-situ star formation would possess, at the current epoch many dynamical times later, negligible substructure. \\citet{Bell07} compare the level of substructure seen in the SDSS data with cosmological simulations (over galactocentric distances 1-40 kpc) and find that the Milky Way's halo is consistent with assembly from dwarf galaxies.  The most prominent contribution to outer halo substructure results from the on-going disruption of the Sagittarius Dwarf Spheroidal galaxy \\citep[Sgr, ][]{Ibata94}. The leading and trailing arms of Sgr are seen to encircle the sky \\citep{Majewski03, Newberg02, Newberg07, fieldofstreams, Keller08}. The extensive debris of Sgr offer an important probe of the mass and shape of the Galaxy \\citep{Ibata01, Helmi01, Helmi04, Johnston05, Law05, Fellhauer06}.  Another outer halo substructure, the Virgo Over-density (VOD) was found as an over-density in RR Lyrae variables \\citep[RRLs, ][]{Vivas01, Vivas04, VivasZinn06, Keller08} and in SDSS as a diffuse structure of main-sequence stars \\citep{Juric08, Newberg02, Newberg07} and blue horizontal branch stars \\citep[BHBs, ][]{Ivezic05}. Radial velocity measurements by \\citet{Duffau06}, \\citet{Newberg07}, \\citet{Vivas08}, and \\citet{Prior09} reveal that a subset of associated RRLs share a common spatial velocity. This moving group is termed the Virgo Stellar Stream (VSS) to distinguish it from the general stellar over-density and potentially other velocity groups that may spatially coincide.  The study of \\citet{Duffau06} explored the spatial extent of the VOD by considering the excess of the observed luminosity function over an off-source control field. This study discerned a size for the VOD of at least 106 square degrees. A similar technique is implemented by \\citet{Prior09} to derive a spatial extent for the VOD of $\\sim760$ square degrees. The study by \\citet{Juric08} utilized the photometric parallaxes of F-type main-sequence stars to describe a large diffuse over-density covering $\\sim 1000$ square degrees and spanning a range of distance from 6-20 kpc. \\citet{Vivas08} find a small concentration of stars with similar radial velocities to the VSS reaching to 12 kpc from a previous detection at 19 kpc.  \\citet{MartinezDelgado04,MartinezDelgado07} suggest that the VOD is the result of the confluence of the leading and trailing arms of Sgr. However, as pointed out by \\citet{Newberg07} and \\citet{Vivas08}, models of the Sgr debris stream predict highly negative radial velocities contrary to the $+$100-130 kms$^{-1}$ seen in the VSS. Extant models also fail to match the observed density enhancement of the VOD. \\citet{Newberg07} and \\citet{Vivas08} conclude that the VOD is most likely tidal debris from a separate and distinct accretion event. Recently, \\citet{Casetti-Dinescu09} have presented a preliminary orbit for the VSS based on one RRL that places the system near pericentre on an eccentric and highly destructive orbit.  In light of the considerable uncertainty regarding the apparent size, distance and origin of the VOD we have presented a series of new observational studies targeting this region. \\citet{Keller08} has examined the spatial distribution of RRL candidates in the VOD region and finds in addition to the over-density seen by the QUEST team at 17 kpc, a second region of over-density 15 degrees to the south-east and at a distance of 20 kpc. The study of \\citet{Prior09} examines the kinematics of a sample of RRL candidates from \\citeauthor{Keller99} et al. (2008). In contrast to previous studies, this study finds RRLs that possess radial velocities consistent with membership of Sgr trailing arm debris. They conclude that a significant contribution from Sgr can not be ruled out at the present time. \\citet{Keller09_TNO} utilizes the luminosity function excess over that expected from the Besancon galaxy model \\citep{Robin03} to map the extent of the VOD. This study finds that the region consists of three spatially distinct over-densities that extend from declinations $+$2$^{\\circ}$ to $-$15$^{\\circ}$.  In the present study, we utilize the sub-giant population (those stars between the main-sequence turn-off and the base of the red giant branch) to perform tomography of the outer halo in order to describe the VOD in more detail and to look for additional substructure. In Section \\ref{section:subgiant_tomography} we present our method for determining the significance of, and distance to, halo substructure. Section \\ref{section:results} presents our results on the extent and luminosity of a series of halo sub-structures with comparison, where possible to existing literature. Finally, in Section \\ref{section:origins} we discuss our map of substructure in relation to the tidal debris of Sgr. ",
        "conclusions": "In addition to the Sgr leading arm, four spatially distinct regions of over-density are apparent in Figure \\ref{figure:nogcs}. We identify these as the Virgo Over-density (VOD) that has been known from previous studies discussed above and new candidate stellar over-densities I term the 160$^\\circ$ and 180$^\\circ$ features, and the Virgo Equatorial Stream (VES).  The first issue we must consider is if these features could result from a background population of red clump stars in our sub-giant colour selection window, due, for instance, to the Sgr stream.  In order to produce the Virgo Equatorial Stream, 160$^{\\circ}$ and 180$^{\\circ}$ over-densities the Sgr stream red clump stars would have to lie at $\\sim110$, $\\sim65$ and $\\sim75$kpc respectively. As evident from numerous studies of Sgr (see above discussion, and shown here in Figure \\ref{figure:nogcs}) the leading arm is a factor of two closer at these positions and so can not account for the observed over-density via its population of red clump stars. In the case of the VOD, the distance to its associated RRL over-density is in concordance the distances derived here.  Similarly, to bring about the observed over-densities be means of lower main-sequence stars falling in the sub-giant colour selection window would require distances of approximately 7, 4, and 5kpc for the Virgo Equatorial Stream, 160$^{\\circ}$ and 180$^{\\circ}$ over-densities. There are no indications from the distribution of RRLs at these distances that such over-densities do exist \\citep{Miceli08}.  Secondly, the features do appear to be artefacts of sampling to a coarse spatial grid. They are not expunged by a different choice of spatial pixel centroids. We have experimented with sub-pixel spatial shifts of the grid and recover the over-densities at similar significance.  The region examined in the present study has been the focus of a number of previous studies targeting halo substructure. In the following section we discuss our results in the context of these previous studies where possible. To define the extent of each feature we implement a kernel discriminant analysis. The technique is a non-parametric smoothing of the spatial distribution by an optimally selected multivariate kernel. The result is a probability density function for the data that retains maximum information content afforded by the data (i.e.\\ not over/under smoothed). Such probability density function were constructed (using the {\\tt{ks}} package \\citep{Duong07} in {\\tt{R}}) for all four features. Contours that contain 75\\% and 50\\% of feature data (drawn from a broad rectangular volume containing the feature) are shown projected in two dimensions in Figure \\ref{figure:pos}. The corresponding areas of the features are given in Table \\ref{table:mvs}.  \\begin{figure*}[h] \\begin{center} \\includegraphics[scale=0.4, angle=0]{Keller_fig5.pdf} \\caption{{\\bf{Top}}: Dashed lines show the distance at which sectional slices through the significance data cube are made to reveal the spatial orientation and morphology of the identified substructures. {\\bf{Bottom}}: RA, Dec slices through the excess data cube at the mean distances of each sub-structure. The distance range averaged in each panel is given in the top right. The contours around each sub-structure shows the extent of the structure. The inner contour contains 50\\% of the sub-giant over-density, and the outer contour contains 75\\% of the over-density (see Section \\ref{section:discussion_of_substructure} for details). Slices through the Sgr leading arm can be seen in the projected sky significance maps for the VOD, the more distant 160$^{\\circ}$ feature and the Virgo Equatorial Stream.}\\label{figure:pos} \\end{center} \\end{figure*}  We draw the reader's attention to the fact that these are sectional slices through the data cube and so as such only show those features visible at a specific distance. For this reason the Sgr leading arm does not appear as a contiguous element but as a 6-9 degree wide section in these panels. The Monoceros ring \\citep{Newberg02} (or warp, \\citet{Conn08}) falls outside the field under consideration here at lower galactic latitudes off the right of Figure \\ref{figure:nogcs}.  In Figure \\ref{figure:cmds} we show the Hess diagrams for each sub-structure population. These diagrams are derived by taking the on-source population CMD density and subtracting an adjacent off-source CMD density, appropriately scaled to match the number of stars in the on-source field. Here, $g$ and $r$ magnitudes have been dereddened according to the reddening maps of \\citet{Schlegel98} together with a colour-specific reddening index of E($g$$-$$r$) = 1.042E($B$$-$$V$) and $A_g $/E($B$$-$$V$) = 3.793. The Hess diagrams reveal low density main-sequence turnoff features for each sub-structure. For example, Figure \\ref{figure:cmds} (top left) shows the $g_0 \\sim 19$ turnoff for the VOD. In this example, the more distant Sgr leading arm material can be seen in the sub-giants at $g_0 \\sim 21$.  \\begin{figure}[h] \\begin{center} \\includegraphics[scale=1, angle=0]{Keller_fig6.pdf} \\caption{Hess diagrams of the population of the sub-structures discussed in the present study. We have subtracted from the on-source field a population taken from an adjacent off-source field, appropriately scaled to match the total number of stars in the on-source field. Lighter greyscale indicates population excess. The arrows indicate the approximate magnitude of the main-sequence turn-off in each case.}\\label{figure:cmds} \\end{center} \\end{figure}  \\subsubsection{The Virgo over-density} \\label{section:VOD}  The VOD is seen to be a diffuse structure centred at RA=194.6$^{\\circ}$, Dec=+1.8$^{\\circ}$ covering 144 square degrees and distances from $\\sim12$-$18$ kpc. With a peak significance at a distance of 16 kpc, the projected dimensions of the VOD in RA is 5 kpc and 1.7 kpc in Dec. We note that \\citet{Keller09_TNO, Prior09, Newberg07} (discussed below) have shown that the VOD extends outside the bounds of the SDSS survey to the south-west and hence the spatial extent derived here is a lower limit. Further more, the significance measured along the line of sight for the VOD is also underestimated due to the depth of the structure. This underestimation occurs because our method for determining significance requires the subtraction of a low order polynomial fit to the general background. Our simulations show that this has the effect of diminishing the significance of, but not erasing, features that exhibit a line of sight depth of $g>0.4$ mag.  The VOD was first detected as an over-density in RR Lyrae variables (RRLs) by \\citet{Vivas01} and independently detected as an over-density of F-type main-sequence stars by \\citet{Newberg02}. Both studies describe the centre of the VOD at RA=190$^{\\circ}$, Dec=0$^{\\circ}$ and a heliocentric distance of 18 kpc. From a 2.3$^{\\circ}$-wide band centred on Dec=-1.18$^{\\circ}$, \\citet{VivasZinn06} find a peak density for the VOD at RA=186$^{\\circ}$ and a distance of 17kpc. In the study of RRL candidates by \\citet{Ivezic05} a distance of 18 kpc is derived for the VOD with a center at RA$\\sim190$$^{\\circ}$ in a 2.5$^{\\circ}$-wide band centred on Dec=0$^{\\circ}$.  \\citet{Duffau06} presents radial velocities for a sample of \\citet{Vivas01} RRLs and \\citet{Sirko04} BHBs.  They find six of the nine stars in the densest region of the VOD (centred at RA=186$^{\\circ}$ and Dec=-1$^{\\circ}$) share a common radial velocity of $+100$ kms$^{-1}$ with a velocity dispersion less than measurement uncertainties. This moving group was termed by \\citeauthor{Duffau06} the Virgo Stellar Stream to differentiate it from the general region of over-density. Recently, \\citet{Vivas08} finds members of the VOD that share this common radial velocity (i.e.\\ the VSS) that extend as close as 12 kpc. By quantifying the excess in the luminosity function in VOD fields relative to offset fields, \\citet{Duffau06} reports that the VOD covers an area in excess of 106 square degrees.  It was demonstrated by \\citet{Newberg07} that the VOD extends beyond the southern edge of the SDSS survey. \\citeauthor{Newberg07} shows significant over-density that appears a continuation of the VOD in one of the three SEGUE `outrigger' scans. \\citet{Newberg07} examines SDSS spectroscopy in a 1.5$^{\\circ}$ radius field centred on the peak of this over-density at RA=191$^{\\circ}$ and Dec=-7.8$^{\\circ}$. They find a marginally significant peak in the radial velocity distribution in common with the VSS. This suggests that the VSS extends across this area.  \\citet{Juric08} presents a study of the density excess of main-sequence F-type stars towards the VOD. They present evidence that the VOD covers $\\sim$1000 square degrees centred on RA=192$^{\\circ}$ and Dec=+2$^{\\circ}$ and has extensive line-of-sight depth (6-20 kpc). The study of \\citet{Prior09} implements a technique similar to that of \\citet{Duffau06} that utilizes the excess in the observed luminosity function. In this case, the excess is relative to Besancon galaxy model \\citep{Robin03} predictions. This study estimates an areal extent of the VOD of 760 square degrees.  \\citet{Keller09_TNO} matches the technique of \\citet{Prior09} over 1$^{\\circ} \\times 1^{\\circ}$ fields from the SDSS DR6 \\citep{Adelman-McCarthy08} and SEKBO \\citep{Keller08} datasets. The SEKBO data enables an extension to the south-west that reveals that the region of the VOD encompasses three regions of over-density. \\citet{Keller09_TNO} Feature $A$ is a linear feature approximately 3$^{\\circ}$ wide and 14.5$^{\\circ}$ long centred at (198$^{\\circ}$, -10$^{\\circ}$) with a distance of 20 kpc (hence linear size of 1$\\times$5 kpc). Feature $A$ is not seen here as it does not extend into the SDSS sample. Feature $B$ at (192$^{\\circ}$, -2$^{\\circ}$) is 3$^{\\circ}$ wide and 10$^{\\circ}$ long, with a heliocentric distance of 17 kpc (linear size of 1$\\times$3 kpc). It is coincident with, and hence identified as, the VOD. A third over-density, Feature $C$, is identified with the 180$^{\\circ}$ feature of the present study (discussed below).  The extent of the VOD seen in the present study, and that of \\citet{Keller09_TNO}, is substantially smaller than that of \\citet{Juric08} or \\citet{Prior09}. The study of \\citet{Juric08} derives the spatial extent of the VOD from the density of stars within 0.2$<g-r<$0.3 and 20$<r<$21, so chosen to select main-sequence stars at a distance appropriate for the VOD. \\citet{Newberg07} applies identical selection criteria, but for bracketing distances ($r$ magnitudes one magnitude brighter and fainter; see \\citet{Newberg07} figures 5 and 6). The results of Newberg et al.\\ show that the Sgr stream remains a significant contributor to the stellar density over 20$<r<21$. We propose that the northern portion of the stellar density seen the Juric et al.\\ result (their figure 37) is due to the Sgr leading arm rather than the VOD, and that this leads to Juric et al.\\ to greatly overestimate the spatial extent of the VOD. In addition, the \\citet{Juric08} study utilizes photometric parallax to determine the distance to the target main-sequence stars. As discussed previously, the inherently large uncertainties associated with this technique lead to an {\\it{over}}estimation of the depth of the VOD by \\citeauthor{Juric08}. In the case of the \\citet{Prior09} study, the spatial sampling proved insufficient to resolve the VOD region into the series of distinct components seen by \\citet{Keller09_TNO}.  The VOD possesses an apparent line-of-sight inclination on the sky such that the closest region at 12 kpc is to the eastern extremity (RA$\\sim 184^\\circ$) and the furthest region at 18 kpc is at the westerly extremity (RA=$210^\\circ$). This is in line with the findings of \\citet{Keller09_TNO} who, from distances to the RR Lyrae members, finds the westerly portion to be at a distance of 20 kpc from a central region (at RA=194$^{\\circ}$) of 17 kpc. Since \\citet{An09} finds no significant metallicity gradient across the VOD region, this is likely due to a true spatial gradient.  \\subsubsection{180$^{\\circ}$ over-density}  The 180$^{\\circ}$ over-density (centred at RA=179.8$^{\\circ}$, Dec=+2.0$^{\\circ}$) is of unresolved depth ($<2$ kpc) and covers 45 square degrees of sky. At a distances of 19 kpc, the projected radius for this feature is 1.5 kpc. The 180$^{\\circ}$ over-density was first seen in the \\citet{Keller09_TNO} study (their Feature $C$) but without significant RRL population a distance was not able to be ascribed for the feature.  \\subsubsection{160$^{\\circ}$ over-density}  The 160$^{\\circ}$ over-density is centred at RA=160.9$^{\\circ}$, Dec=13.3$^{\\circ}$ and has an extent of 36 square degrees. It is located at a heliocentic distance of 17 kpc, at which distance the approximately circular feature presents a radius of 1.3 kpc. The 160$^{\\circ}$ over-density is seen projected in front of the Sgr leading arm. In Figure \\ref{figure:cmds} (top right) this Sgr material is apparent as the sub-giant population seen at $g_0 \\sim 20.2$.  \\subsubsection{The Virgo Equatorial Stream}  The Virgo Equatorial Stream (VES) is centred at (RA=203.3$^{\\circ}$, Dec=+2.7$^{\\circ}$, r=28 kpc) and covers at least 162 square degrees (since it is not clear that it is contained within the southern boundary of the SDSS field). It possess a small line-of-sight depth of $<$4 kpc. The corresponding projected dimensions of the VES at a mean distance of 28 kpc are 8.9$\\times$4.4 kpc.  The VES is situated slightly westward of the VOD but 10kpc more distant ($\\Delta g = 1.6$ magnitudes).  Could the VES be a metal-rich sub-population of the VOD? The above derived distances assume an old (12Gyr), metal-poor population ([Fe/H]=-1.7). This is analogous to stellar population seen in present-day faint dwarf spheroidals in the local neighbourhood. If the observed structures formed from a much larger ($M>10^8M_{\\odot}$) gas-rich system such as Sgr we might expect a broad metallicity distribution function for the debris. This is seen in the case of the Andromeda Giant Stellar Stream believed to have been recently derived from a $M\\sim10^8M_{\\odot}$ satellite \\citep{Ibata07}. However, if we assume [Fe/H]=-1.7 for the VOD, the implied VES metallicity would be greatly super-solar. In addition, the work of \\citet{An09} does not support the presence of a substantial spatial metallicity variation in this region. Similarly, it is not possible that the VES could result from the lower MS of the VOD since the VES is only 1.6 mag.\\ more distant that the VOD and not the $\\sim3$ mag.\\ expected for the lower MS. For these reasons we conclude that the VOD and VES are separate spatial structures.    An examination of the excess of sub-giant stars over a subset of the SDSS DR6 reveals a series of extended halo over-densities in the survey heliocentric distance range of 10 - 40 kpc. The most prominent feature is that of the Sagittarius dwarf galaxy (Sgr) leading tidal debris arm. We derive a spatial configuration for the Sgr leading arm material that is in agreement with previous studies. The Virgo Over-density is seen to extend from 12-18 kpc, with a mean distance of 16 kpc, and to cover an area of 144 square degrees. This is considerably smaller extent than found in previous studies. We contend this is due to our resolution of the region into a series of sub-structures. We determine the surface brightness of the VOD to be 30.5 mag/arcsec$^{2}$.  Three new candidate halo substructures are detected. They consist of two roughly spherical over-densities, the 160$^{\\circ}$ and 180$^{\\circ}$ features. These features lie at 17 and 19 kpc respectively and present a surface brightness of 32.3 and 33.2 mag/arcsec$^{2}$ respectively. The third feature, the Virgo Equatorial Stream, lies at 28 kpc and extends over at least 162 deg$^{2}$. It possesses a stream-like morphology and a surface brightness of 31.7 mag/arcsec$^{2}$.  We compare our observations of halo substructure to the predictions of Sgr tidal debris. We highlight a critical deficiency of current modelling of Sgr, namely, the discordance between the observations and the model predictions for the distance to the leading arm material. We conclude that further modelling effort and radial velocities are required to clarify the origins of the substructures observed."
    },
    "0911/0911.3163_arXiv.txt": {
        "abstract": "The standard model of cosmology relies on the existence of two components, ``dark matter'' and ``dark energy'', which dominate the expansion of the Universe. There is no direct proof of their existence, and their nature is still unknown. Many alternative models suggest other cosmological scenarios, and in particular the {\\it dark fluid} model replace the dark matter and dark energy components by a unique dark component able to mimic the behaviour of both components. The current cosmological constraints on the unifying dark fluid model is discussed, and a dark fluid model based on a complex scalar field is presented. Finally the consequences of quantum corrections on the scalar field potential are investigated. ",
        "introduction": "According to the standard model of cosmology, the total energy density of the Universe is presently dominated by two component densities: the {\\it dark matter} component, which is a pressureless matter fluid having attracting gravitational effects, and {\\it dark energy}, whose main properties are a negative pressure and a nearly constant energy density today, which evoke the idea of vacuum energy. The nature of both components remains unknown, and in the near future we can hope that the Large Hadron Collider (LHC) will be able to provide hints on the nature of dark matter. In spite of the mysteries of the dark components, it is generally considered that dark matter can be modeled as a system of collisionless particles, whereas the most usual models of dark energy are the scalar field based quintessence models. Nevertheless, many difficulties in usual dark energy and dark matter models still cast doubts on the cosmological standard model bases, leaving room for other models to be investigated. Here we consider a unifying model in which the dark matter and dark energy components can be considered as two different aspects of a same component, the {\\it dark fluid}. We will first review how such a unifying model can be constrained by the cosmological observations. Then we will consider a dark fluid model based on a complex scalar field, and we will discuss the question of the choice of the scalar field potential through an analysis of quantum corrections. ",
        "conclusions": "\\noindent The standard model of cosmology assumes the existence of two unknown dark components, dark matter and dark energy. We have seen that it is possible to replace both components with a unique component, the dark fluid, and be consistent with the cosmological data. The properties of the dark fluid are severely constrained because of the dual dark energy/dark matter conditions, but it is nevertheless possible to model the dark fluid with a simple and usual complex scalar field. The main question of the model is the choice of the scalar field potential. An adequate choice can lead to a fluid explaining at the same time the dark matter observations at local scales, and the dark energy behaviour at cosmological scales. The possibility that a dark fluid can lead to a correct structure formation has still to be tested, and methods to simulate scalar field condensation are being developed \\cite{Bernal:2006ci,Bernal:2009zy}. However, the relations between such a model and quantum theories are not obvious, and still need to be worked out. In this context, a new interesting potential is currently under investigation: \\begin{equation} V(\\Phi) = \\alpha \\, \\mbox{cotanh}\\left( \\frac{\\beta}{\\Phi^\\dagger \\Phi} \\right)\\;\\;. \\end{equation} It has the advantage of finding roots in brane theories, and yet no big difference in behaviour is expected from the one determined by the potential in Eq. (\\ref{potential}). Another way to find an adequate potential would be to study the possibility of an even larger unification, {\\it i.e.} dark matter, dark energy and inflation, as proposed in \\cite{Liddle:2006qz,PerezLorenzana:2007qv,Liddle:2008bm,Bose:2008ew}.  To conclude, dark fluid models appear as viable alternatives to the standard cosmological model, and need to be further investigated."
    },
    "0911/0911.2988.txt": {
        "abstract": "We study the dependence of the \\Mbh{}--\\Mhost{} relation on the redshift up to $z=3$ for a sample of 96 quasars the host galaxy luminosities of which are known. Black hole masses were estimated assuming virial equilibrium in the broad line regions (Paper I), while the host galaxy masses were inferred from their luminosities. With this data we are able to pin down the redshift dependence of the \\Mbh{}--\\Mhost{} relation along 85 per cent of the Universe age. We show that, in the sampled redshift range, the \\Mbh{}--\\Lhost{} relation remains nearly unchanged. Once we take into account the aging of the stellar population, we find that the \\Mbh{}/\\Mhost{} ratio ($\\Gamma$) increases by a factor $\\sim 7$ from $z=0$ to $z=3$. We show that $\\Gamma$ evolves with $z$ regardless of the radio loudness and of the quasar luminosity. We propose that most massive black holes, living their quasar phase at high-redshift, become extremely rare objects in host galaxies of similar mass in the Local Universe. ",
        "introduction": " ",
        "conclusions": "\\subsection{Comparison with previous results} \\begin{figure} \\begin{center} \\includegraphics[width=0.49\\textwidth]{fig_gammalit.ps}\\\\ \\caption{The average values of $\\Gamma$ for the quasars in our sample (filled symbols), together with the linear best fit (dotted line). Error bars are the 2-$\\sigma$ uncertainties in the average values for each data bin. For a comparison, the trend found in the study of radio loud AGN by \\citet{mclure06} is plotted as a dashed line. We also plot the average $\\Gamma$ values of the 51 lensed and non-lensed quasars from \\citet[empty circles]{peng06a,peng06b} and of the 89 AGN from \\citet[empty triangles]{merloni09}. The horizontal solid line represents the constant $\\Gamma=0.002$ case. All together the data depict a clear increase of $\\Gamma$ with $z$. }\\label{fig_gammalit} \\end{center} \\end{figure}   As a key result of this analysis, we find that the \\Mbh{}/\\Mhost{} ratio significantly increases with redshift. Hereafter we compare these findings with those of other studies available in the literature.  \\citet{mclure06} match the average trend of \\Mbh{} observed in 38 Radio Loud quasars at $z<2$ with the typical stellar masses of massive radio galaxies in the same redshift bins. This approach relies on the assumption that quasar host galaxies are comparable, at any redshift, with massive radio galaxies. The major caveat here is that their results may be biased by the different histories of quasar and radio galaxies (for instance, note that the luminosity functions of AGN evolves differently for various luminosity subclasses). Nevertheless, \\citet{mclure06} find an increase of $\\Gamma$ comparable to the one observed in the present study (see Figure \\ref{fig_gammalit}). Our results extend these findings to RQQs and beyond the peak age of quasar activity.  \\citet{peng06a,peng06b} address the evolution of the \\Mbh{}--\\Mhost{} relation as a function of redshift in $\\sim20$ low-$z$ quasars and in $\\sim 30$ high-$z$ lensed quasars imaged with the HST. Their data show a larger scatter than ours, possibly due to uncertainties in the modelling of the lens mass distribution and the lens light subtraction. They find that there is practically no evolution in the \\Mbh{}--\\Lhost{} relation. On the other hand, when correcting for the evolution of the stellar population, they find an excess (a factor 3--6) in the \\Mbh{} values at high-$z$ with respect to the prediction from the local \\Mbh{}--\\Mhost{} relation, in qualitative agreement with our findings.  \\citet{merloni09} study the \\Mbh{}--\\Mhost{} ratio in a sample of 89 type-1 AGN with $1<z<2.2$ from the zCOSMOS survey. Black hole masses are derived through the standard virial assumption, while the host galaxy luminosities and stellar masses are inferred from multi-wavelength fitting of the spectral energy distribution of the targets (with no direct information about the morphology of the galaxies). This technique is effective with intermediate to low-luminosity AGN, while it cannot be applied to quasars as bright as ours, where the nuclear light overwhelms the galaxy contribution. They find that the average $\\Gamma$ is higher than what observed in the Local Universe, the excess scaling as $(1+z)^{0.74}$, consistently with the trend observed in our data in the same redshift bin. \\citet{jahnke09} observed 10 of the targets in Merloni's sample with the HST and independently derived host galaxies luminosities with a procedure similar to the one adopted in our data sources \\citep[see, e.g.,][]{kotilainen09}. They find no evolution in the \\Mbh{}--\\Mhost(total) ratio. However, clues of the occurrence of discs are present, thus the \\Mbh{}--\\Mhost{}(bulge) ratio is expected to evolve as $(1+z)^{1.2}$, in agreement with our findings.  \\citet{bennert09} address the \\Mbh{}--\\Lhost{} relation in a sample of 23  Seyfert galaxies with $0.3<z<0.6$. They study the morphology of the host galaxies of these objects using NIR observations from the HST. A careful modelling is adopted in order to disentangle nuclei, bulges and disc components. Black hole masses are derived in a way similar to that presented in paper I. Once corrected for the evolution of the stellar population, they find $\\Gamma \\propto (1+z)^{(1.4\\pm0.2)}$, in good agreement with our results.   Additional indication of an evolution of $\\Gamma$ comes from the relation of \\Mbh{} with the stellar velocity dispersion, $\\sigma_*$, of the host galaxy. In particular, it is remarkable that these works suggest that the higher is the redshift, the more massive is the black hole for a given $\\sigma_*$. For instance, \\citet{salviander07} use the width of the \\Oiii{} narrow emission line as a proxy of the stellar velocity dispersion and study the \\Mbh{}--$\\sigma_*$ relation in a sample of $\\sim 1600$ quasars up to $z=1.2$ taken from the SDSS. They find that \\Mbh{} at high redshift are $\\sim 0.2$ dex larger than what expected from local \\Mbh{}--$\\sigma_*$ relation. A smaller evolution ($\\lsim0.1$ dex), albeit with small significance, is also proposed by \\citet{shen08b}, based on a sample of 900 Type-1 AGN with $z\\lsim0.4$. More recently, \\citet{woo08} and Woo et al. (in preparation) address the \\Mbh{}--$\\sigma_*$ relation in Seyfert galaxies up to $z\\sim0.6$, and find an overall \\Mbh{} excess at high redshift with respect to the prediction from low-$z$ relationships. These findings support our results, notwithstanding the different characteristic luminosities, morphologies and stellar contents of the sampled targets with respect to those examined in our analysis.  It is interesting to note that the $z=6.42$ quasar SDSS J1148+5251 has a black hole mass of few $\\times10^{9}$ \\Msun{} \\citep{willott03} and a dynamical mass of the host galaxy of $\\sim5\\times10^{10}$ \\Msun{} \\citep{walter03,walter04}, yielding $\\Gamma\\sim0.1$, which is in agreement with the extrapolation of our results at that redshift ($\\Gamma\\approx0.13$) and well beyond the $\\Gamma=0.002$ value observed in the Local Universe.  These results as a whole support a picture where, for a given quasar host galaxy, its central black hole at high redshift is `over-massive' with respect to its low-$z$ counter-parts. This picture is also consistent with the constrains on the \\Mbh{}--\\Mhost{} evolution derived from the comparison between the galaxy stellar mass function and the quasar luminosity function \\citep{somerville09}.  %Only two exceptions are found in the literature: \\citet{shields03} %do not find any significant evolution in the \\Mbh{}--$\\sigma_*$ %relation in a sample of $\\sim75$ quasars from $z=0$ to $z=3$. Their %findings disagree with the subsequent work by \\citet{salviander07}, %who apply similar techniques on a sample of thousands of objects %from the SDSS database. %The only exception is provided by \\citet{alexander08}, who study the %BH--host galaxy relation in a sample of 6 submillimeter-emitting galaxies %at $z\\approx2$ + 6 ULIRG at low redshift. They find no redshift dependence %of log $\\Gamma$, which is below the low-$z$ relation in both subsample %(log $\\Gamma\\approx-3.5$). Such a distinction may be is probably due to %the peculiar stage of formation/interaction that these galaxies are %experiencing, which may place them outside the expectations for typical %luminous AGN.  \\subsection{Why does $\\Gamma$ evolve?}  The interpretation of the observed evolution in the \\Mbh{}--\\Mhost{} ratio is challenging. Exotic scenarios involving black hole ejection from their host galaxies due to gravitational wave recoil or to 3-body scatter may be applicable for few peculiar targets \\citep[e.g., see][but see also Bogdanovic et al., 2009, Heckman et al., 2009 and Dotti et al., 2009 for alternative explanations]{komossa08}, but they are not applicable to the general case. Thus, {\\it if} high-$z$ quasars are destined to move towards the local \\Mbh{}--\\Mhost{}, the unavoidable consequence of our results is that, at a given \\Mbh{}, galaxy masses increase from $z=3$ to the present age. Hereafter we sketch three possible basic pictures for that. We also present an alternative scenario, in which the fate of high-$z$ quasars may be different, the remnants of high-$z$ quasars keeping high $\\Gamma$ values down to the present age. \\begin{description} \\item[{\\it Galaxy growth by mergers} --] A first scenario involves substantial mass growth of quasar host galaxies through merger events. It is remarkable that strong gravitational interactions may trigger intense gas infall in the centre of galaxies and may even lead to the activation of BH accretion. This is observed in a number of relatively low-redshift AGN \\citep[e.g.,][]{bennert08,bennert09} showing dense close environments or disturbed morphologies, and confirmed by the presence of young stellar populations in some quasar host galaxies \\citep{jahnke04,jahnke07}. Two arguments disfavour this scenario. First, theoretical models based on the structure evolution in a $\\Lambda$CDM cosmology \\citep[e.g.,][]{volonteri03} predict that a massive galaxy experience only few (1--2) major merger events from $z=3$ to $z=0$. However, our study shows that a factor $\\sim7$ increase of the stellar mass of the host galaxies is required from $z=3$ to $z=0$, which means that the host galaxies have to suffer $\\geq3$ major mergers in this redshift range. Secondly, several pieces of evidence suggest that massive inactive galaxies, as well as quasar host galaxies, have already formed/assembled the majority of their mass in very remote Cosmic epochs \\citep[$z\\gsim3$; see, e.g.,] [and references therein]{kotilainen09}. The stellar population may experience episodic rejuvenation, but this only marginally affects the mean age of the stellar content: The stellar shells observed, e.g., by \\citet{canalizo07} and \\citet{bennert08} in low-redshift quasar host galaxies account for 5 -- 10 per cent of the total stellar population. Similarly, in a comparison with inactive galaxies of similar mass, \\citet{jahnke04} find that quasar host galaxies are on average only $0.3$ mag bluer. If the galaxies enter the quasar phase $\\sim1$ Gyr after the activation of the starburst, as suggested by the authors of that study, then the involved mass is $\\sim30$ per cent of the initial mass of the galaxy. Moreover, if the quasar host galaxies contain a significant fraction of young stellar populations, then the Mass-to-Light ratio would be smaller. Therefore, young host galaxies at high-$z$ would yield a $\\Gamma$--$z$ relation even steeper than that reported in Figure \\ref{fig_gamma}.\\\\  \\item[{\\it Stellar mass growth through gas consumption} --] Another possible interpretation is that high redshift quasar host galaxies are gas rich, and form a significant fraction of their stellar content in relatively recent Cosmic epochs. This is consistent with a picture in which the black hole mass is somehow sensitive to the energetic budget of the galaxy or its dynamical mass rather than its stellar mass \\citep[see for instance][]{hopkins07}. This scenario is disfavoured as all the quasar host galaxies in our sample are massive elliptical, and the stellar content of these galaxies is usually old. Moreover, if significant star formation occurred in quasar host galaxies in the redshift range explored in this work, the evolution of $\\Gamma$ would be much steeper, making this scenario even less realistic. \\\\  \\item[{\\it Evolution of the fundamental plane} --] A number of studies suggest that inactive, massive elliptical galaxies were more compact in the high redshift than in the Local Universe \\citep[Trujillo et al. 2006; but see also][]{cappellari09}. In particular, an evolution of the Faber--Jackson relation is predicted, implying that the higher is $z$, the higher is the galaxy velocity dispersion $\\sigma_*$. If \\Mbh{} constantly regulates the host galaxy $\\sigma_*$ \\citep[e.g.,][]{silk98}, so that the \\Mbh{}--$\\sigma_*$ relation does not evolve significantly, then even a small (a factor $\\sim 1.6$) increase of $\\sigma_*$ for a given galaxy would yield to an excess of a factor 7 in $\\Gamma$. The main limit of this picture is that studies of the evolution of the \\Mbh{}--$\\sigma_*$ relation through redshift do find an increase of the average \\Mbh{} for a given $\\sigma_*$, when moving from the Local to high-$z$ Universe \\citep{salviander07,woo08,bennert09}. \\\\  \\item[{\\it Remnants of high-$z$ quasars as rare outliers} --] We propose a scenario in which the local counter-parts of high-$z$ quasars are high-mass outliers above the \\Mbh{}--\\Mhost{} relation. The more massive is the BH, the earlier it experiences its quasar phase \\citep{marconi04,merloni04}. Our study shows that these objects have higher expected $\\Gamma$, but they are extremely rare, and contribute marginally to the presently known \\Mbh{}--\\Mhost{} relation. In particular, the $2<z<3$ quasars should appear as inactive massive galaxies with \\Mbh{}$\\sim10^{9.5}$ \\Msun{} in the nearby Universe. In order to quantify the occurrence of such objects in the Local Universe, we assume the mass function of quasars: $\\Phi({\\cal M}_{\\rm BH}\\sim10^{9.5} {\\rm M_{\\odot}},2<z<3)\\approx4\\times10^{-9}$ Gpc$^{-3}$ \\Msun$^{-1}$ \\citep{vestergaard09}, and take a volume corresponding to the most distant inactive BH for which a direct measurement of the mass is available\\footnote{The BH at the centre of the brightest cluster galaxy in Abell 1836, at 147 Mpc; see \\citet{dallabonta09}}. Under these assumptions, we expect virtually no objects ($0.2$ in the whole volume) with such high values of $\\Gamma$.\\\\ Using the same arguments for targets at intermediate redshift ($1.0<z<1.5$), we expect few ($\\sim 2$) objects. Quasars at $z\\sim1.2$ have $\\Gamma$ values $\\sim0.3$ dex larger than that at $z=0$. This offset is close to the one observed for the handful of objects populating the high-mass end of the local \\Mbh{}--\\Mhost{} relation \\citep[e.g.][]{marconi03}. This scenario is thus consistent both with the $\\Gamma$--$z$ relation of quasars and with the observed shape of the local \\Mbh{}--\\Mhost{} galaxy relation of nearby inactive galaxies. \\end{description}"
    },
    "0911/0911.4743_arXiv.txt": {
        "abstract": "Here we present HST/ACS imaging, in the B and I bands, of the edge-on Sb/Sc galaxy NGC 5170. Excluding the central disk region region, we detect a 142 objects with colours and sizes typical of globular clusters (GCs). Our main result is the discovery of a `blue tilt' (a mass-metallicity relation), at the 3$\\sigma$ level, in the metal-poor GC subpopulation of this Milky Way like galaxy. The tilt is consistent with that seen in massive elliptical galaxies and with the self enrichment model of Bailin \\& Harris. For a linear mass-metallicity relation, the tilt has the form Z $\\sim$ L$^{0.42 \\pm 0.13}$. We derive a total GC system population of 600 $\\pm$ 100, making it much richer than the Milky Way. However when this number is normalised by the host galaxy luminosity or stellar mass it is similar to that of M31. Finally, we report the presence of a potential Ultra Compact Dwarf of size $\\sim$ 6 pc and luminosity M$_I$ $\\sim$ --12.5, assuming it is physically associated with NGC 5170. ",
        "introduction": "Studies of the globular cluster (GC) system of the Milky Way have benefited from the wealth of information that is available. However such studies are also fundamentally limited by the small sample size of about 150 objects. Discerning more subtle trends in GC properties, or finding GCs in more extreme astrophysical states, requires larger samples which can only be found in the GC systems of other galaxies.  One such trend which was revealed first in extragalactic GC systems, and not seen in the Milky Way GC system, is the so-called `blue tilt'. The blue tilt (a trend for the blue GC subpopulation to have redder colours at brighter magnitudes) has been found using the Advanced Camera for Surveys (ACS) on the {\\it Hubble Space Telescope} (HST) in a variety of galaxies from the most massive ellipticals (Harris et al. 2006), to lower mass ellipticals (Strader et al. 2006) and even dwarfs (Mieske et al. 2006a). It has also been detected in an early-type spiral galaxy, the Sombrero, by Spitler et al. (2006).  It has not been detected in any HST study of late-type spirals, but the few studies to date (e.g. Goudfrooij et al. 2003) have generally used the WFPC2 camera which has a more limited field-of-view compared to the ACS camera and hence such studies have focused on the bulge region and their associated red GCs. Interestingly, at this stage, it has not been detected in the rich GC system of the massive elliptical M49 (Strader et al. 2006; Mieske et al. 2006a).   Given the evidence that extragalactic GCs are mostly old (Brodie \\& Strader 2006), the blue tilt implies a mass-metallicity relation for the metal-poor subpopulation of GCs (although this has yet to be confirmed spectroscopically). Various explanations have been proposed (see Bekki et al. 2007; Mieske et al. 2006a; Mieske 2009), but perhaps the most plausible explanation to date is one of self enrichment, whereby more massive GCs are enriched with more heavier metals during their brief formation period.  In the self enrichment models of Strader \\& Smith (2008) and Bailin \\& Harris (2009), a single generation of star formation and resulting supernovae is responsible for the self enrichment of each GC, which only becomes important above a threshold mass of $\\sim$ 10$^6$ M$_\\odot$. Both models predict heavy element abundance variations within each massive GC, which may be the situation in some extragalactic GC systems (e.g. Cenarro et al. 2007). The Galactic GC system includes only a few GCs with masses $\\ge$ 10$^6$ M$_\\odot$, so a clear signature of variation in colour (or metallicity) with magnitude (a blue tilt) is difficult to detect. It is thus important to determine whether galaxies similar to the Milky Way reveal a GC blue tilt or not.  Recently the ACS was used to study the star cluster system of the Sc galaxy NGC 3370 by Cantiello et al. (2009). NGC 3370 is a near face-on spiral at a somewhat larger distance of 28.7 Mpc. These two properties make it more difficult to study its GC system.  Cantiello et al. detected 35 GC candidates with a mean colour B--I = 1.60 $\\pm$ 0.27 but no evidence of a blue tilt.   Here we present an analysis of the globular cluster system of NGC 5170 using the ACS camera on board the HST. NGC 5170 is a potential analogue of the Milky Way Galaxy with a Hubble type of Sb (Sandage \\& Bedke 1994) to Sc (NASA Extragalactic Database; NED). Fischer et al. (1990) derive a bulge-to-disk ratio of 0.5 which is more consistent with an Sb, like M31, than an Sc.  It is a relatively isolated, edge-on ($i$ $\\sim$ 86$^o$; Bottema et al. 1987) spiral with a total V band magnitude of 9.84 (NED). It lies at a Galactic latitude of 43$^{o}$, making the surrounding field relatively free of Galactic stars and hence a good target for a GC study. A wide range of distance estimates are found in the literature. Ferrarese et al. (2000) report the GC luminosity function distance of Fischer et al. (1990) of 21.5 Mpc.  However, Fischer et al. did not reach beyond the GC turnover magnitude so this value should be considered somewhat uncertain.  The Tully-Fisher based distance estimates of Willick et al. (1997) range from 24 to 43 Mpc. The kinematic study of NGC 5170 by Bottema et al. (1987) assumed a distance of 20 Mpc. The Virgo infall corrected velocity (NGC 5170 lies behind but close to the Virgo cluster southern extension) from NED is 1665 $\\pm$ 27 km $^{-1}$, which with H$_o$ = 73 $\\pm$ 5 km s$^{-1}$ Mpc$^{-1}$ corresponds to a distance of 22.8 $\\pm$ 1.6 Mpc (m--M = 31.79 $\\pm$ 0.15). The final error is the quadrature sum of the velocity error combined with the error on the Hubble constant. For this distance, the galaxy has an absolute magnitude of M$_V$ = --21.95. We initially adopt this latter value, for which the pixel scale of the ACS (0.05$^{''}$) equals $\\sim$5 pc, but we reinvestigate this issue below and derive a new distance estimate based on measured sizes of the NGC 5170 GCs. ",
        "conclusions": "Using two pointings of the HST/ACS camera, in the B and I filters, we have explored the GC system surrounding the edge-on spiral galaxy NGC 5170. We employed the aperture-correction method of Harris (2009a) for partially resolved objects to obtain accurate total magnitudes.  The initial selection of candidate GCs is based on (B--I)$_o$ colour and signal-to-noise ratio (which closely tracks magnitude).  We used the object size, as measured by ISHAPE, to further refine the selection process and to provide an estimate of the distance to NGC 5170 (under the assumption of a universal average size for GCs). This distance is 19.5 Mpc and is within $\\sim$2$\\sigma$ of the Hubble flow distance.  The colour-magnitude diagram reveals a well-populated blue subpopulation and a sparse red subpopulation, in a ratio of roughly 3 to 1. The GC subpopulation colours are similar to those seen in the Milky Way and M31 GC systems. We also find a blue tilt that is similar, within our restricted magnitude range, to that for massive ellipticals found by Harris (2009a) and to the self enrichment model described in Harris et al. (2009a). These findings suggest that normal late-type spiral galaxies, along with galaxies of earlier types, host GC systems with a blue tilt.  The blue tilts are most likely evidence for a mass-metallicity relation within the metal-poor GC subpopulation. The reason for the lack of an observed blue tilt in the Milky Way GC system appears to be simply due to the small number of relatively massive GCs (Bailin \\& Harris 2009; Harris 2009a). We would expect that new, high precision photometry of a large sample of the M31 GC system should reveal a blue tilt.  We note that the reality of blue tilts in general has been disputed by Kundu (2008), and in the case of M87 in particular by Waters et al. (2009). Subsequent work by Harris (2009a,b), Peng et al. (2009) and Madrid et al. (2009) has addressed the issues raised by these workers and has strongly confirmed the astrophysical reality of blue tilts.   Although of similar Hubble type and total magnitude to the Milky Way (Sbc, M$_V$ $\\sim$ --21.3), we find that the GC system of NGC 5170 is much richer with a total GC population of 600 $\\pm$ 100 GCs. In this sense NGC 5170 bears more resemblance to the Sb galaxy M31 with $\\sim$450 GCs (Barmby et al. 2000), or the Sa Sombrero galaxy with 1900 GCs (Rhode \\& Zepf 2004).  A commonly-used measure to directly compare GC systems is the specific frequency S$_N$, which normalises GC number by M$_V$. Assuming M$_V$ = --21.6 for NGC 5170, we estimate S$_N$ = 1.37 $\\pm$ 0.23.  This is higher than the Local Group spirals (i.e. for the Milky Way S$_N$ = 0.6 and M31 S$_N$ = 0.9) and the four Sb-Sc spirals studied by Rhode et al. (2007) who found a range of 0.5 $<$ S$_N$ $<$ 0.9, but less than S$_N$ = 2.1 for the Sombrero galaxy (Rhode \\& Zepf 2004). An alternative measure for comparing GC systems is the stellar mass normalised GC frequency, T$_N$ (Rhode \\& Zepf 2004; Spitler et al. 2008). Using the 2MASS K band magnitude of 7.63 and a Salpeter IMF we derive a stellar mass of 7.6 $\\times$ 10$^{10}$ M$_{\\odot}$. This gives T$_N$ $\\sim$ 8. However this likely an upper limit given that some fraction of the K band flux will be obscured due to the highly edge-on ($i$ $\\sim$ 86$^o$) nature of NGC 5170. Spitler et al. showed that T$_N$ varies with the stellar mass of the host, and that the Sombrero galaxy had the highest T$_N$ value ($\\sim$ 9) of the 8 spiral galaxies listed. Thus NGC 5170 appears to have a high GC frequency, placing it between the Local Group spirals and the Sombrero galaxy.  Using the scaling relation of Spitler \\& Forbes (2009) between total GC system mass with halo mass, we estimate a halo mass for NGC 5170 to be 3.4 $\\times$  10$^{12}$ M$_{\\odot}$. This is a factor of 2-3$\\times$ that of the Milky Way halo mass.   We also report the discovery of a possible Ultra Compact Dwarf associated with NGC 5170.  If confirmed by its radial velocity, it will be another rare example of a UCD located near a spiral galaxy outside that of a galaxy cluster (see Hau et al. 2009)."
    },
    "0911/0911.1811_arXiv.txt": {
        "abstract": "We study the evolution of cosmological perturbations in $f(\\GB)$ gravity, where the Lagrangian is the sum of a Ricci scalar $R$ and an arbitrary function $f$ in terms of a Gauss-Bonnet term $\\GB$. We derive the equations for perturbations assuming matter to be described by a perfect fluid with a constant equation of state $w$. We show that density perturbations in perfect fluids exhibit negative instabilities during both the radiation and the matter domination, irrespective of the form of $f(\\GB)$. This growth of perturbations gets stronger on smaller scales, which is difficult to be compatible with the observed galaxy spectrum unless the deviation from General Relativity is very small. Thus $f(\\GB)$ cosmological models are effectively ruled out from this Ultra-Violet instability, even though they can be compatible with the late-time cosmic acceleration and local gravity constraints. ",
        "introduction": "Independent observational evidence for dark energy has stimulated the idea that General Relativity (GR) may be modified on large distances to give rise to a late-time cosmic acceleration \\cite{review}.  A simple dark energy scenario constructed in this vein is so-called $f(R)$ gravity in which $f$ is a function of the Ricci scalar $R$ \\cite{fRori}.  Although there are some restrictions to the functional form of $f(R)$ to satisfy both cosmological and local gravity constraints, it is possible to design viable models \\cite{fRviable} that can be distinguished from GR at least in the metric formalism of $f(R)$ gravity.  The $f(R)$ gravity in the metric formalism corresponds to the so-called Brans-Dicke theory with a parameter $\\omega_{\\rm BD}=0$ in the presence of a potential of gravitational origin \\cite{Chiba}.  One can generalize this to scalar-tensor theories with an arbitrary Brans-Dicke parameter $\\omega_{\\rm BD}$. In fact it is possible to construct scalar-field potentials that can be responsible for the cosmic acceleration, while at the same time satisfying local gravity constraints \\cite{TUMTY}.  These models, including $f(R)$ gravity, exhibit several interesting observational signatures such as the phantom equation of state \\cite{AT07}, the modified matter power spectrum \\cite{fRmatter}, and the modified weak lensing spectrum \\cite{WL}.  The Ricci scalar $R$ is not the only scalar quantity which is used to change gravity, as we can easily construct other scalar quantities such as $R_{\\mu \\nu}R^{\\mu \\nu}$ and $R_{\\mu \\nu \\rho \\sigma}R^{\\mu \\nu \\rho \\sigma}$ from the Ricci tensor $R_{\\mu \\nu}$ and the Riemann tensor $R_{\\mu \\nu \\rho \\sigma}$ \\cite{PQR}.  However, for the Gauss-Bonnet (GB) curvature invariant \\begin{equation} \\GB \\equiv R^2-4R_{\\mu \\nu}R^{\\mu \\nu} +R_{\\mu \\nu \\rho \\sigma}R^{\\mu \\nu \\rho \\sigma}\\,, \\end{equation} one can avoid the appearance of spurious spin-2 ghosts \\cite{ghost1,ghost2}.  If a scalar field $\\phi$ with an exponential potential $V(\\phi)=V_0 e^{-\\lambda \\phi}$ couples to the GB term \\cite{NO05}, a scaling matter era can be followed by a late-time de Sitter solution for the exponential GB coupling $F(\\phi) \\propto e^{\\mu \\phi}$ with $\\mu>\\lambda$ \\cite{KMota1,TS07}.  While this GB coupling is well motivated by low-energy effective string theory \\cite{Gas} the joint likelihood analysis using observational data of big bang nucleosynthesis, large scale structure, and baryon acoustic oscillations disfavors such a model, provided that it aims to account for dark energy \\cite{KMota2}. In addition the energy contribution coming from the GB term needs to be strongly suppressed for consistency with solar-system experiments \\cite{Amen,Thomas}. This cannot be compatible with the requirement of cosmic acceleration today (see also Ref.~\\cite{Davis}),  at least in the presence of a kinetic term for the scalar field and some forms of the potential.  The instability of tensor perturbations is also present in those models during the epoch of cosmic acceleration \\cite{Kawai,TS07,Guo}.  However, it is possible to explain the late-time cosmic acceleration for the modified gravity scenario in which the Lagrangian density is given by ${\\cal L}=R+f(\\GB)$, where $f(\\GB)$ is an arbitrary function in terms of $\\GB$ \\cite{fGO}, provided the function $f$ satisfies some conditions \\cite{DT08}. This is equivalent to a theory with a scalar field coupled to the GB term in the absence of a kinetic term \\cite{fGO,Soti07,Uddin}. A number of $f(\\GB)$ models that have a matter era followed by a de Sitter (dS) attractor have been proposed in Ref.~\\cite{DT08} (see also Refs.~\\cite{Cognola,Li,Zhou}).  These models can be also consistent with local gravity constraints for a wide range of parameter space \\cite{DT09}.  In order to test the cosmological viability of $f(\\GB)$ dark energy models, it is important to study cosmological perturbations responsible for structure formation. In Ref.~\\cite{Li} the evolution of density perturbations has been discussed for non-relativistic matter with an equation of state $w=0$, under the approximation that the background cosmological evolution mimics that of the $\\Lambda$CDM model. In this paper we derive the equation for density perturbations with a general constant equation of state $w$. Therefore our analysis includes the perturbations in radiation ($w=1/3$) as well as those in non-relativistic matter. Moreover we use concrete $f(\\GB)$ models that satisfy both local gravity constraints and cosmological constraints at the background level. This is particularly important when we discuss the evolution of perturbations at late times, because the deviation from the $\\Lambda$CDM model can be significant.  We will show that, in the Universe dominated by a single perfect fluid with an equation of state  parameter $w$, the perturbations in the fluid exhibit violent instabilities for $w>-1/2$ in the small-scale limit. This is associated with a negative speed squared $c_s^2$ of one eigenvector-mode. The perturbations of non-relativistic matter as well as radiation, at some scale, will start to grow exponentially during the matter/radiation domination, unless the deviation from GR is very small. In GR, this same mode does not exist, so that there is no smooth limit from one theory to the other.  This paper is organized as follows. In Sec.~\\ref{sec:tm} we review $f(\\GB)$ models that satisfy cosmological constraints at the background level as well as local gravity constraints. Sec.~\\ref{sec:pgwm} is devoted to the analysis of cosmological perturbations in $f(\\GB)$ gravity in the presence of a perfect fluid with the equation of state $w$. In this section we discuss the presence of an instability at small scales due to a negative speed squared for the propagating mode. In Sec.~\\ref{sec:mg} we analyze more in detail for perturbations in non-relativistic matter in order to describe the growth of large-scale structure. We shall numerically integrate the perturbation equations for concrete $f(\\GB)$ models and estimate how much deviation from GR can be allowed by the observations of galaxy clustering in the linear regime. We conclude in Sec.~\\ref{sec:c}. ",
        "conclusions": "\\label{sec:c}  In this paper we have studied cosmological perturbations in $f(\\GB)$ gravity, in the presence of a perfect fluid with a constant equation of state $w$.  In the Universe dominated by a single fluid with $w > -1/2$, we have shown the presence of an instability associated with a negative speed squared of one eigenvector-mode.  Hence the perturbations in radiation and non-relativistic matter are affected by this instability during the radiation and matter domination, respectively.  Our results are more general than those given in Ref.~\\cite{Li} in the sense that we have considered a general equation of state $w$ and that we have not assumed the $\\Lambda$CDM-like background evolution.  A useful quantity that characterizes the deviation from the $\\Lambda$CDM model is $\\mu=H \\dot{f}_{,\\GB}$.  In the limit that $\\mu \\to 0$ (i.e. the $\\Lambda$CDM model) one can avoid the appearance of the negative instability.  If $\\mu \\ne 0$, the instability of perturbations appears for $\\mu \\gtrsim (aH/k)^2$.  Even for tiny values of $\\mu$ much smaller than 1, there are small scale modes that satisfy this condition.  We have studied the evolution of non-relativistic matter perturbations numerically and confirmed that the perturbations are strongly amplified once they enter the regime $\\mu \\gtrsim (aH/k)^2$. {}From the requirement that the matter power spectrum in the linear regime is not affected by this violent instability, we have found that the maximum value of the deviation parameter is constrained to be $\\mu_{\\rm max} \\lesssim 10^{-6}$.  Nonetheless the UV limit of this theory remains unsatisfactory, as perturbation theory would break down eventually at some scale, and the background is not under control any longer at these scales. When this happens, there is not even an easy way, for the cosmological background of this theory to be checked against observations. It is mostly this UV unpredictable behavior which sets the strongest bound. This feature will always remain unless one sets $\\mu$ identically to 0 at any time, as in this case the theory reduces to GR and the unstable eigenvector-mode automatically disappears. In this sense we believe that this theory is ruled out by our analysis, which explicitly shows the existence of eigenmodes with negative squared speed in the past.  While we have focused on linear perturbations, non-linear effects become important once $\\delta_m$ grows to the order of 1. It may be of interest to see whether such non-linear effects strengthen or weaken the growth of perturbations."
    },
    "0911/0911.1202_arXiv.txt": {
        "abstract": "We use time-dependent, one-dimensional disc models to investigate the evolution of protostellar discs that form through the collapse of molecular cloud cores and in which the primary transport mechanism is self-gravity. We assume that these discs settle into a state of thermal equilibrium with $Q = 2$ and that the strength of the angular momentum transport is set by the cooling rate of the disc. The results suggest that these discs will attain a quasi-steady state that persists for a number of free-fall times and in which most of the mass within $100$ au is located inside $10-20$ au. This pile-up of mass in the inner disc could result in temperatures that are high enough for the growth of MHD turbulence which could rapidly drain the inner disc and lead to FU Orionis-like outbursts. In all our simulations, the inner regions of the discs ($r < 40$ au) were stable against fragmentation, while fragmentation was possible in the outer regions ($r > 40$ au) of discs that formed from cores that had enough initial angular momentum to deposit sufficient mass in these outer regions. The large amounts of mass in these outer regions, however, suggests that fragmentation will lead to the formation of sub-stellar and stellar mass companions, rather than planetary mass objects.  Although mass accretion rates were largely consistent with observations, the large disc masses suggest that an additional transport mechanism (such as MRI occuring in the upper layers of the disc) must operate in order to drain the remaining disc material within observed disc lifetimes. ",
        "introduction": "The formation of low-mass stars occurs through the collapse of cold, dense molecular cloud cores \\citep{terebey84}.  These cores, however, generally contain amounts of angular momentum far in excess of the rotational angular momentum of a single star \\citep{caselli02}. Most of the mass must therefore first pass through a circumstellar disc before accreting onto the central star. One of the open problems in star and planet formation is what mechanism acts to transport the angular momentum outwards, allowing this accretion to take place.  It is now generally accepted that in most astrophysical discs, angular momentum transport is driven by magnetohydrodynamic (MHD) turbulence initiated by the magneto-rotational instability (MRI) \\citep{balbus91,papaloizou03}. During the earliest stages of star formation protoplanetary discs are, however, cold and dense and probably do not have even the very small degree of ionization needed to sustain MHD turbulence \\citep{blaes94}.  At these early times, the disc is, however, likely to be massive and disc self-gravity may then provide an alternate and possibly dominant transport mechanism through the growth of the gravitational instability \\citep{toomre64,lin87,laughlin94}.  Although it has been suggested that gravitationally unstable discs may fragment to form gas giant planets or substellar companions \\citep{boss98,boss02}, it is now generally accepted that the conditions for fragmentation are difficult to achieve in the inner planet forming regions of protostellar discs \\citep{matzner05,rafikov05,boley06,whitworth06,stamatellos08, forgan09}. A self-gravitating protostellar disc is more likely to settle into a quasi-steady state in which the instability acts to transport angular momentum outwards \\citep{gammie01, rice03, lodato04, vorobyov07}. It has been shown \\citep{lodato04,lodato05} that in this quasi-steady state, the transport can be approximated as a local viscous process for all but the most massive discs.  Since the disc will also be in thermal equilibrium, the strength of the angular momentum transport is also set by the cooling rate of the disc \\citep{gammie01,rice03}. Using the above results, \\citet{rice09} have shown that if self-gravity is the dominant transport mechanism, protostellar discs will settle into quasi-steady states that are largely independent of the initial conditions. The resulting surface density profiles are quite steep with $\\sim 80$ \\% of the mass within $50$ au located inside $10 -20$ au. \\citet{rice09} also found that the quasi-steady mass accretion rates depend strongly on the disc mass and on the mass of the central star. Their simulations also suggested that no regions of the disc (which extended to $50$ au) were susceptible to fragmentation with the inner $10 - 20$ au being particularly stable, consistent with other recent calculations \\citep{clarke09,rafikov09}.  The simulations carried out by \\citet{rice09} assumed initial power-law surface density profiles that were evolved until a quasi-steady state - which we define as a state in which the mass accretion rate is approximately independent of radius - was achieved. Although this gives information about the quasi-steady nature of such systems, it doesn't tell us if such a state can be achieved under realistic conditions. Here, we extend the work of \\citet{rice09} by considering the evolution of self-gravitating, circumstellar discs formed through infall from a molecular cloud core.  We find that these disc do settle rapidly into quasi-steady states, with properties similar to that found by \\citet{rice09}. The mass accretion rates are also generally consistent with observations \\citep{muzerolle05}, although at early times the simulated accretion rates are somewhat higher than observed.  Although the disc masses are generally higher than observed, a significant fraction of the mass is located in the optically thick inner regions and would not be easily detected using current techniques. These results are consistent with the suggestion \\citep{armitage01,zhu09} that mass will pile-up in the inner regions of the disc and will drain episodically onto the star - potentially explaining FU Orionis outbursts \\citep{hartmann96} - when the temperature becomes sufficiently high for MRI to operate.  The paper is organised as follows. In Section 2, we describe the basic model that builds the disc from an infalling molecular cloud core and evolves the disc self-consistently through, primarily, self-gravity.  In Section 3, we describe the results, and in Section 4 we summarize our conclusions. ",
        "conclusions": "We consider here the evolution of protostellar discs that form through the collapse of molecular cloud cores and in which the primary transport mechanism is self-gravity. We assume that the discs settle quickly into a marginally gravitationally stable state with $Q = 2$ and are then in thermal equilibrium. In this state angular momentum transport is driven by an effective gravitational viscosity that is calculated using the assumption that energy is dissipated at a rate equal to the rate at which the disc loses energy through radiative cooling. Following the work of \\citet{armitage01} and \\citet{zhu09} we also assume that MRI \\citep{balbus91} will operate if the disc is hot enough ($T > 1400$ K) to be partially ionised. This occurs only in the inner disc and is episodic - draining material from the inner disc and then turning off until the inner disc has been replensished by material from the outer disc - and has therefore been proposed as a mechanism for explaining FU Orionis outbursts \\citep{hartmann96}.  Our simulations consider a range of cores masses ($0.25 - 5$ M$_\\odot$) and a range of rotation rates ($0.05 \\le f \\le 1.3$ with $f = \\Omega_c / \\sqrt{G \\rho_c}$). The primary results are as follows \\begin{itemize} \\item{The discs settle quickly into a quasi-steady state with a reasonably steep surface density profile and with a lot ($\\sim 50$ \\%) of the mass within $100$ au located inside $10 - 20$ au. The disc-to-star mass ratio decreases with increasing protostar mass, consistent with observations in the Orion Nebula cluster \\citep{eisner08}. Although the disc masses tend to be higher than observed, with so much mass located in the inner optically thick regions of the disc, most of this mass would be difficult to detect with current techniques.} \\item{Although the mass accretion rates are initially higher than those observed, these high accretion rates only persist for $\\sim 1$ free-fall time and quickly drop to values similar to that observed.  The simulations also suggest that the accretion rate varies less strongly with protostar mass ($\\dot{M} \\propto M_*$) than suggested by observations ($\\dot{M} \\propto M_*^2$) \\citep{muzerolle05, natta06}, but is consistent with the upper boundary of the observed accretion rates \\citep{hartmann06}.} \\item{In none of our simulations did the inner disc ($r < 40$ au) have conditions suitable for fragmentation ($\\alpha_{\\rm g} > 0.06$). In some cases, the outer disc was susceptible to fragmentation, with the primary factor determining if a disc could fragment being the rotation rate of the molecular core. The large amount of mass in these discs, however, suggests that fragmentation is more likely to result in sub-stellar or stellar companions that planetary mass objects \\citep{kratter09}. None of our simulations with $f \\le 0.1$ satisfied the fragmentation conditions while it was satisifed for all those with $f \\ge 0.25$. The inner radius of the fragmentation region increased from $\\sim 40$ au for $M_{\\rm core} = 0.25$ M$_\\odot$ to $\\sim 100$ au for $M_{\\rm core} = 5$ M$_\\odot$ as did the radial range over which fragmentation could occur, although the outer regions of some of these discs could be stabilised against fragmentation by cosmic ray ionisation \\citep{rafikov09}.} \\item{The mass accretion rate depends strongly on the disc mass and in all cases drops below $10^{-7}$ M$_\\odot$ yr$^{-1}$ when the disc mass is still quite substantial ($M_{\\rm disc} > 0.3$ M$_\\odot$). If self-gravity were to remain the primary transport mechanism, disc clearing timescales would be significantly longer than that observed. This suggests that an additional tranport mechanism, such as MRI occuring in the upper layers of the disc \\citep{gammie96}, must also operate.  We propose that this additional mechanism is likely to be negligible initially and become more effective with time. This is probably the case for layered accretion which should become more efficient as the disc mass decreases and is consistent with the requirement that smalls dust grains must be depleted before MRI can operate effectively \\citep{sano00,sano02,ilgner06}, but the pile-up of material in the inner regions of the disc (ultimately resulting in episodic FU Orionis-like outbursts) would also not occur if the additional transport mechanism were too efficient at early times.} \\end{itemize}  Ultimately it appears that transport driven by self-gravity can explain protostar formation and disc evolution at early times. Our simulations suggest that disc masses may - at early times - be higher than suggested by observations but these large disc masses also suggest that an additional transport mechanism must dominate at later times ($t > 1$ Myr) to remove the remaining disc material within observed disc lifetimes. The pile-up of material in the inner regions of the disc - which may explain FU Orionis outburst - may also play an important role in the subsequent formation of planets."
    },
    "0911/0911.1728_arXiv.txt": {
        "abstract": "We analyze the Mass Varying Neutrino (MaVaN) scenario. We consider a minimal model of massless Dirac fermions coupled to a scalar field, mainly in the framework of finite temperature quantum field theory. We demonstrate that the mass equation we find has non-trivial solutions only for special classes of potentials, and only within certain temperature intervals. We give most of our results for the Ratra-Peebles Dark Energy (DE) potential. The thermal (temporal) evolution of the model is analyzed. Following the time arrow, the stable, metastable and unstable phases are predicted. The model predicts that the present Universe is below its critical temperature and accelerates. At the critical point the Universe undergoes a first-order phase transition from the (meta)stable oscillatory regime to the unstable rolling regime of the DE field. This conclusion agrees with the original idea of quintessence as a force making the Universe roll towards its true vacuum with zero $\\Lambda$-term. The present MaVaN scenario is free from the coincidence problem, since both the DE density and the neutrino mass are determined by the scale $M$ of the potential. Choosing $M \\sim 10^{-3}~$eV to match the present DE density, we can obtain the present neutrino mass in the range $m \\sim 10^{-2}-1~$eV and consistent estimates for other parameters of the Universe. ",
        "introduction": "Neutrino mass related questions are of great interest for particle physics as well as for cosmology (for reviews see Ref. \\cite{Dolgov} and references therein). Current upper limits on the sum of neutrino masses from cosmological observations are of the order of 1 eV \\cite{Lesgourgues:2006nd,Strumia:2006db,Hannestad:2006zg}, while neutrino oscillations give a lower bound of roughly 0.01 eV \\cite{Eidelman:2004wy,Abele:2008zz}, making neutrino mass an established element of particle physics. Furthermore, understanding the origin of neutrino mass opens a window into understanding physical processes beyond the standard model of particle physics \\cite{Mohapatra07,Avignone:2007fu,Weinberg08,Dodelson}.  It is now well established that about seventy four percent of the Universe is comprised of dark energy (DE) (for reviews see Ref. \\cite{Peebles:2002gy} and citation therein). The present stage of evolution of the Universe is governed by this dominant DE contribution, and the Universe experiences an accelerating expansion \\cite{Carroll:2003qq,Copeland:2006wr}. The nature of DE is still unknown, and it is one of the major questions of modern cosmology. There are, broadly speaking, three major possibilities proposed to explain the DE \\cite{Peebles:2002gy}. Most straightforwardly, and in good agreement with the current observational data, it can be present just as the cosmological constant \\cite{Peebles:2002gy}. Secondly, the DE can be accommodated in some framework of the modified non-Einsteinian gravity theories (see, e.g., Refs.~\\cite{Nojiri:2003ft,Carroll:2004de}). And lastly, following the original proposals \\cite{Ratra:1988,Wetterich:1988} on the DE originating from a scalar field action similar to the inflaton field, there has been a lot of activity in constructing and analyzing various trial scalar field Lagrangians to model the DE \\cite{Copeland:2006wr}. Note, that it is even unclear what kind of scalar field potential governs the inflationary expansion of the Universe \\cite{Linde}, and as the result, the effective quantum field that adequately describes inflation is still under debate \\cite{Weinberg:2008hq}. A similar observation can be drawn from analyzing many potentials proposed for the DE action \\cite{Copeland:2006wr}.  On the other hand, several cosmological and astrophysical observations imply that about twenty two percent of the Universe consists of dark matter (DM) \\cite{Peebles:2002gy}, if we admit the general relativity theory of gravity. Most probably DM is formed through massive weakly interacting particles (WIMPs), and the nature of these particles is also still unknown. There are several recent observations performed by PAMELA \\cite{Mocchiutti:2009sj} and GLAST missions which indicate DM particle annihilations \\cite{Baltz:2008wd}. Recently it was proposed that both these observations could be used to test baryogenesis \\cite{Kohri:2009yn} which is one of the important problems of the standard particle physics model.  Another puzzling question in modern cosmology is the coincidence problem - the density of DE is comparable to the present energy density of DM. In turn, the latter is comparable (within the order of magnitude), to the energy density of cosmological neutrinos \\cite{Dolgov,Lesgourgues:2006nd}). Is there a mechanism explaining this coincidence? A very convincing answer to this question is given by the mechanism of DM mass generation via various types of DM-DE couplings, ranging from Yukawa to more exotic ones. \\cite{Anderson:1997un,Hoffman:2003ru,Farrar:2003uw,Kawasaki:1991gn,GarciaBellido:1992de,Comelli:2003cv} The mass of the DM particle in this approach is naturally time-dependent, and they were coined Varying Mass Particles (VAMPs). Various DE--DM interaction models have been constrained by observations of Supernovae type Ia \\cite{SuNo}, the age of the Universe \\cite{Franca:2003zg,Franca:2009xp,Huey:2004qv}, Cosmic Microwave Background (CMB) anisotropies \\cite{Wang:2006qw,Mainini:2007nq}, and Large Scale Structure (LSS) formation \\cite{LSS}.     Fardon, Nelson and Weiner elaborated on the VAMP mechanism in the context of neutrinos. \\cite{Fardon:2003eh}\\footnote{The DE-neutrino coupling and the baryogenesis constraints have been also studied also in Ref. \\cite{Gu:2003er}}. In their model the relic neutrinos, i.e., fermionic field(s), interact with a scalar field via the Yukawa coupling. If the decoupled neutrino field is initially massless, then the coupling generates a (varying) mass of neutrinos in this DE-neutrinos model. This mass varying neutrino (MaVaN) scenario is quite compelling, since it connects the origin of neutrino mass to the DE, and solves the additional coincidence problem of why the neutrino mass and DE are of comparable scales \\cite{Peccei:2004sz}. (For more on the coincidence, see, e.g. \\cite{SolaEtalCoincidence}). To consider neutrinos as particles which get their mass through the coupling is attractive for particle physics, as well as for its cosmological consequences. However there are significant issues that have to be resolved for the sake of viability of the MaVaN scenario. Most notably, it has been shown \\cite{Afshordi:2005ym} that the model of Ref. \\cite{Fardon:2003eh} suffers from a strong instability due to the negative sound speed squared of the DE-neutrino fluid (see also \\cite{Kaplinghat:2006jk}).   Any DM-DE coupling induces observable changes in large scale structure formation \\cite{Pettorino:2008ez}. The main reason for this is due to the presence of additional DM contributions (perturbations) in the equation of motion which determines the dynamics of the scalar field. The changes in the dynamics are drastic when massive neutrinos are coupled to DE \\cite{Afshordi:2005ym}. In this case the squared sound speed of the DE-neutrino fluid defined as $c_s^2 = {\\delta P}/{\\delta \\rho}$, (where $\\delta$ represents the variation, and $P$ and $\\rho$ are pressure and energy density of the DE-neutrino fluid) is negative. The negative squared sound speed results in an exponential growth of scalar perturbations. \\cite{Bjaelde:2007ki,Bean,Valiviita:2008iv,Wetterich}     After the critique in Ref. \\cite{Afshordi:2005ym}, the issue of stability of the DE-neutrinos fluid  has been addressed by many authors \\cite{Kaplinghat:2006jk,Wetterich,Brookfield,Bernardini,Takahashi:2005kw,Takahashi:2006jt,Spitzer:2006hm, Fardon:2005wc,Ichiki:2007ng}. Various physical assumptions were made in those references in order to avoid the exponential clustering of neutrinos. In particular, to achieve stability, proposals were put forward to make the DE-DM model more complicated, e.g., by extending it to a multi-component scalar field, or by promoting its supersymmetry. \\cite{Takahashi:2005kw,Fardon:2005wc} We however are not inclined to pursue this line of thought and will explore the simplest possible ``minimal'' model. As we will demonstrate, the occurrence of the instability in the coupled DE-neutrinos model is meaningful, and we will explore the physical implications of this phenomenon. Note that Wetterich and co-workers \\cite{Wetterich} have already analyzed various implications of the instability in the MaVaN model on the dynamics of neutrino clustering.     In this paper we re-address the analysis of the DE-neutrinos coupled model. What is really new in our results, to the best of our knowledge, apart from a consistent equation for the equilibrium condition, is the analysis of the thermal (i.e. temporal) evolution of the MaVaN model and prediction of its stable, metastable and unstable phases. The analysis of the dynamics in the unstable phase results in, for the first time in the framework of the MaVaN scenario,  a picture of the present-time Universe totally consistent with observations. Our findings are in line with the original proposal \\cite{Ratra:1988,Wetterich:1988} of the DE potential (quintessence) to model the Universe slowly rolling towards its true vacuum ($\\Lambda=0$). As it turns out, the present Universe, seen as a system of the coupled DE (quintessence) field and fermions (neutrinos) is below its critical temperature. It is similar to a supercooled liquid which has not crystallized yet: its high temperature (meta)stable phase became unstable, but the new low-temperature stable phase ($\\Lambda=0$) is still to be reached. The Afshordi-Zaldarriaga-Kohri instability corresponding to $c_S^2 <0$ is just telling us this.   The rest of the paper is organized as follows: In Section~\\ref{M&F} we give the outlook of the model and formalism applied and derive the basic equations for the coupled model. In Section~\\ref{MasEq} we present the qualitative analysis of the equation which yields the fermionic (neutrino) mass. Section~\\ref{PRPot} contains analysis of the coupled model with the Ratra-Peebles DE potential at equilibrium. The dynamics of the model applied to the whole Universe is studied in Section~\\ref{Dynamics}. The results are summarized in the concluding Section~\\ref{Concl}. ",
        "conclusions": "In this paper we analyzed the MaVaN scenario in a framework of a simple ``minimal'' model with only one species of the (initially) massless Dirac fermions coupled to the scalar quintessence field. By using the methods of thermal quantum field theory we derived for the first time (in the context of the MaVaN or, even more broadly, the VAMP models) a consistent equation for fermionic mass generation in the coupled model.  We demonstrated that the mass equation has non-trivial solutions only for special classes of potentials and only within certain temperature intervals. It appears that these results have not been reported in the literature on VAMPs before now.  We gave most of the results for the particular choice of a trial DE potential -- the Ratra-Peebles quintessence potential. This potential has all the necessary properties we needed for our task: it is simple, it satisfies the criteria we found for non-trivial solutions of the mass equation to exist, and it has only one dimensionfull parameter- the energy scale $M$ to tune. Also, at small values of the exponent $\\alpha$ it effectively crosses over to the case of a logarithmic potential. We have checked that other potentials, e.g., exponential, lead to a qualitatively similar picture, but they have at least one more energy scale to handle, which we consider as an unnecessary complication at this point.   We analyzed the thermal (i.e. temporal) evolution of the model, following the time arrow. Contrary to what one might expect from analogies with other contexts, like, e.g., condensed matter, the model does not generate the mass via a conventional spontaneous symmetry breaking below a certain temperature. Instead it has a non-trivial solution for the fermionic  mass evolving ``smoothly''  from zero at the ``point'' $T=\\infty$. The scalar field is infinitely heavy at the same point. More realistically, we assumed the model is applicable starting at the temperatures somewhere in the beginning of the radiation-dominated era. We found that the DE contribution in this regime is subleading, and the model behaves as an ultra-relativistic Fermi gas at those temperatures.  This regime corresponds to a stable phase of the model given by a global minimum of the thermodynamic potential $\\Omega(\\varphi)$.  The temperature/time dependent minimum $\\langle  \\varphi \\rangle$ generates the varying fermionic mass $m \\propto \\langle  \\varphi \\rangle$.   With increase in time, as the temperature decreases, the model reaches the point of metastability where its pressure ($P$) vanishes. From our estimates of the model's scales, we showed that this happens during the matter-dominated era of the Universe. At this point the system's ground state becomes doubly degenerate, and the potential $\\Omega=0$ at the non-trivial (finite) minimum  $\\langle  \\varphi \\rangle$ as well as at the trivial vacuum $ \\varphi= \\infty$.  Further on, at lower temperatures the system stays in the metastable (supercooled) state until it reaches the critical point where the local minimum of the thermodynamic potential disappears and it becomes an inflexion point. At this critical temperature the model undergoes a first-order (discontinuous) phase transition. At the critical point \\textit{the equilibrium values} of the fermionic and the scalar field masses discontinuously jump to the `doomsday'' vacuum  state values $m= \\infty$ and $m_\\phi=0$, respectively. The square of the sound velocity and equation of state parameter $w$ have \\textit{the equilibrium values} corresponding to the de Sitter Universe with a cosmological constant, i.e. $c^2_s=w=-1$. It is worth pointing out that $c^2_s>0$ in both the stable and metastable phases, and the sound velocity vanishes reaching the critical temperature from above.  Since the equilibrium approach is not applicable below the critical temperature, we find parameters of the model from  direct numerical solution of the equation of motion and the Friedmann equations. The single scale $M$ of the quintessence potential is chosen to match the present DE density, then other parameters of the Universe are determined. We obtain a consistent picture: the phase transition has occurred rather recently at $z_{\\mathrm{crit}} \\lesssim 5$  during the matter-dominated era, and the Universe is now being driven towards the stable vacuum with zero $\\Lambda$-term. The expansion of the Universe accelerates starting from $z^* \\approx 0.83$. Setting  $\\alpha =0.01$ for $M \\approx 2.4 \\cdot 10^{-3}~\\mathrm{eV}$, we end up with the neutrino mass $m \\approx 0.27~\\mathrm{eV}$.   The present results allow us to propose a completely new viewpoint not only on the MaVaN, but on the quintessence scenario for the Universe as well. The common concerns about the slow-rolling mechanism for the DE relaxation toward the $\\Lambda=0$ vacuum are related to the question of what is the mechanism to set the initial value of the scalar field $\\varphi$ where it evolves (rolls down) from. Our results demonstrate that up to recent times (i.e. above the critical temperature) the quintessence field was locked around its average (classical) value $\\langle  \\varphi \\rangle$. Its value is determined by the scale $M$ and the temperature. The average $\\langle  \\varphi \\rangle$ gives the fermionic mass at the same time. The scalar field is rigid (i.e. massive), although it softens (i.e., its mass decreases) as the system approaches the critical temperature. Above the critical temperature the scalar field can only oscillate around its equilibrium value $\\langle \\varphi \\rangle$. At the critical point the minimum of the thermodynamic potential becomes the inflexion point, the scalar field looses its rigidity (mass). Then the field can only roll down towards the new stable ground state $\\Omega=0$ at $\\varphi=\\infty$. So physically, the critical point corresponds to the transition of the Universe from the stable oscillatory to the unstable rolling regime.  A more sophisticated numerical study of the kinetics after the critical point is warranted in order to address such issues as the detailed description of the crossover between different regimes, and the clustering of neutrinos. These and some other questions are relegated to our future work."
    },
    "0911/0911.4611_arXiv.txt": {
        "abstract": "{ The $S^3$-Tools are a set of Python-based routines and interfaces whose purpose is to provide user-friendly access to the SKA Simulated Skies ($S^3$) set of simulations, an effort led by the University of Oxford in the framework of the European Union's SKADS program ({\\tt http://www.skads-eu.org}). The databases built from the $S^3$ simulations are hosted by the Oxford e-Research Center (OeRC), and can be accessed through a web portal at {\\tt http://s-cubed.physics.ox.ac.uk}. This paper focuses on the practical steps involved to make radio images from the $S^3$-SEX and $S^3$-SAX simulations using the $S^3$-Map tool and should be taken as a broad overview. For a more complete description, the interested reader should look up the user's guide. The output images can then be used as input to instrument simulators, e.g. to assess technical designs and observational strategies for the SKA and SKA pathfinders. } ",
        "introduction": "The SKA Simulated Skies project ($S^3$) aims at building mock radio skies suitable for planning science with the SKA and the many pathfinder experiments. The project comprises a number of simulations, of which essentially two are relevant for this paper:  $\\bullet$ $S^3$-SEX ({\\bf S}emi-{\\bf E}mpirical e{\\bf X}tragalactic) : a large-scale simulation of the extragalactic radio continuum sky covering a sky area of $20^{\\circ} \\times 20^{\\circ}$, out to redshift $z=20$, and down to 10 nJy. It also includes {\\sc Hi} gas masses for star-forming galaxies. For a complete description, see \\cite{wilman08}.  $\\bullet$ $S^3$-SAX ({\\bf S}emi-{\\bf A}nalytical e{\\bf X}tragalactic) : a smaller-scale simulation of the extragalactic {\\sc Hi} and CO line emissions, derived from the Millenium simulation (\\cite{springel05}), which comes in two flavours. $S^3$-SAX-Sky is a skyfield simulation giving the apparent properties of galaxies, with a field of view that depends on the maximum redshift, while $S^3$-SAX-Box is a simulation of a cubic volume giving their intrinsic properties. In all of this paper, $S^3$-SAX shall be understood as meaning $S^3$-SAX-Sky. For a complete description, see \\cite{obreschkow09}. ",
        "conclusions": ""
    },
    "0911/0911.3177_arXiv.txt": {
        "abstract": "We report on a single star, B030D,  observed as part of a large survey of objects in M31, which has the unusual radial velocity of  $-780$ \\kmsc.  Based on details of its spectrum, we find that the star is an F supergiant, with a circumstellar shell.  The evolutionary status of the star could be one of a post-mainsequence close binary, a symbiotic nova, or less likely, a post-AGB star, which additional observations could help sort out.  Membership of the star in the Andromeda Giant Stream can explain its highly negative velocity. ",
        "introduction": "In large surveys, often a few objects will be found that are unusual in one respect or another. Our spectroscopic survey of candidate star clusters in the field of M31 \\citep{PaperI} has provided such a case. From the 1400 objects observed, we expected (and found) a large number of misclassified stars, due to the heterogeneous quality of the imaging material used to identify cluster candidates.  All but a few of these stars are in the foreground, as revealed by their dwarf spectral types.  Our catalog magnitude limit of V=20 required that any M31 stars observed would have to be bright giants or supergiants.  Of the few stars we observed in that area that were clearly not from the foreground Milky Way disk, one star proved to be unusual in two ways: its spectrum is that of an F supergiant with weak Balmer emission, and its heliocentric radial velocity is $-780$ \\kmsc, which is the most negative velocity yet found for a single star.  The star is called B030D in the Bologna catalog of M31 objects \\citep{galleti2} , and ``Bol D30'' in SIMBAD.  This star is unlikely to be a member of M31's disk: it lies close to the center of M31 on the SW side, at a place where the heliocentric velocity of the disk is $-425 $ \\kms (the minimum disk velocity for M31 reaches $-550$ \\kms).  A difference of even 230 \\kms with respect to the disk circular velocity is much too high for any disk or thick disk star in M31. While M31's halo velocity distribution is not well known, a velocity of $-480$ \\kms with respect to M31's center (which has V= $-300$ \\kms) is very high, even for a halo star.  Could it be an unusual star in the halo of the Milky Way?  M31 lies at $l$ = 121.2, $b$ = $-21.6$. Accounting for the Sun's motion with respect to the LSR and the LSR's motion around the galactic center\\footnote{We have assumed that the Sun's motion with respect to the LSR is (U,V,W) = ($-9$, 12, 7) \\kms and the LSR velocity is 220 \\kms \\citep{mihalas}} reduces the velocity to $-$614 \\kms with respect to the Milky Way center.  A recent estimate of the Galaxy's escape speed \\citep[544 \\kmsc,][]{martinsmith} suggests that it is unlikely that this star is bound to the Milky Way; and its highly {\\it negative} velocity and galactic longitude imply that it cannot be a hypervelocity star ejected from the galactic center such as those identified by \\citet{wbrown}.  This paper describes the analysis of the unusual spectrum of the star B030D and the reasoning that led us to conclude that the star is most likely a member of the M31 halo's ``giant stream'' \\citep{ibata01} close to the pericenter of its orbit. ",
        "conclusions": "B030D is a curiosity. Not only does it have an extremely negative velocity, but it is an unusual and rare kind of star as well.  In our research into spectroscopic databases, we could not find any exact analogues that matched the details of its spectrum with regard to the Balmer emission. That problem was certainly mostly due to the lack of modern, digital spectra of the classes of unusual stars described above, but the short life of this kind of F shell star must also be part of the story.  In addition to IR observations, a further piece of data that may help to determine the exact classification would be the star's variability.  We examined the star on both epochs of the Palomar Sky Survey, but could not find a significant change in brightness over that period (36 years), though the different passbands used for the two surveys makes a comparison difficult. Modern digital images of M31 could of course provide more stringent limits on variability.   \\begin{figure*}[ht] \\includegraphics[scale=1.0,clip=true]{f1.eps} \\caption{Location of B030D in M31. The left image comes from a mosaic made from J band Digitized Sky Survey images, and covers 2.5\\degr . The right image comes from a V band image of the Local Group Survey of \\cite{massey}, and the field size is 30\\arcs . North is up, and east is to the left. \\label{m31}} \\end{figure*}    \\begin{figure*}[ht] \\includegraphics[scale=0.4,clip=true]{f2.eps} \\caption{UBV colors of stars with V$<19.5$ in the M31 field, photometry from \\cite{massey}. No reddening corrections have been applied. Stars observed spectroscopically to be members of M31 are highlighted. Of those members, the ones whose spectra are shown in Fig \\ref{members_blue} and succeeding figures are further indicated with larger symbols. B030D is shown as a large filled circle.  The sequence of dwarf stars (class V) and supergiants (Ia) are shown as dashed and solid lines, respectively, the data coming from \\cite{fitz} to which a reddening of E(B$-$V)=0.13 has been applied.  The black arrow indicates the effect of a reddening of  E(B$-$V)=0.25. \\label{ubv}} \\end{figure*}  \\begin{figure*}[ht] \\includegraphics[scale=0.4,clip=true]{f3.eps} \\caption{Blue spectra of M31 stars, membership being determined by velocity, Here we show spectra with Ca II ratios and/or color similar to B030D. Note the similar narrow widths of the Balmer lines, but also the weak H$\\beta$ in B030D compared to  its own other Balmer lines.  The spectra in this figure and subsequent figures have had the continua removed for ease of display. \\label{members_blue}} \\end{figure*}  \\begin{figure*}[ht] \\includegraphics[scale=0.4,clip=true]{f4.eps} \\caption{Foreground galactic stars, also determined by velocity, with  Ca II ratios and colors similar to B030D. Note the much broader Balmer lines in the foreground stars. \\label{nonmembers_blue}} \\end{figure*}  \\begin{figure*}[ht] \\includegraphics[scale=0.4,clip=true]{f5.eps} \\caption{Same as Fig. \\ref{members_blue} in the red. Note the  strong O I line in all these, indicative of very low surface gravity.  Also note the filled-in H$\\alpha$ line of B030D, and the H$\\alpha$ emission for one other star plotted right above it. Neither star shows forbidden emission lines that would indicate a local HII region, such as is found in the third spectrum from the top, which reveals [NII] emission as well as  H$\\alpha$. \\label{members_red}} \\end{figure*}  \\begin{figure*}[ht] \\includegraphics[scale=0.4,clip=true]{f6.eps} \\caption{ Same as Fig. \\ref{nonmembers_blue} in the red. B030D has much stronger O I than the foreground stars. \\label{nonmembers_red}} \\end{figure*}   \\begin{figure*}[ht] \\includegraphics[scale=0.5,clip=true]{f7.eps} \\caption{Equivalent width of the O I$\\lambda7774$ line vs the ratio of the Ca II H\\&K lines. The former is sensitive to surface gravity and the latter is a useful temperature index. \\label{OI}} \\end{figure*}  \\begin{figure*}[ht] \\includegraphics[scale=0.5,clip=true]{f8.eps} \\caption{H$\\alpha$ spectrum, showing weak emission in B030D compared to other M31 member stars which have similar color and luminosity \\label{halpha}} \\end{figure*}   \\begin{figure*}[ht] \\includegraphics[scale=0.7,clip=true]{f9.eps} \\caption{Velocities of PNe and B030D (dark square) plotted against distance along the major axis, for objects within 1 kpc of the major axis ($|$Y$|<$ 1 kpc). The orbit for the M31 Giant stream proposed by \\cite{merrett03} is also shown as a solid line, revealing that B030D can be considered part of the stream.  \\label{stream}} \\end{figure*}"
    },
    "0911/0911.2940_arXiv.txt": {
        "abstract": "We present new high-resolution infrared echelle spectra of V1647 Ori, the young star that illuminates McNeil's nebula.  From the start, V1647 Ori has been an enigmatic source that has defied classification, in some ways resembling eruptive stars of the FUor class and in other respects the EXor variables. V1647 Ori underwent an outburst in 2003 before fading back to its pre-outburst brightness in 2006. In 2008, it underwent a new outburst.  In this paper we present high-resolution $K-$band and $M-$band spectra from the W. M. Keck Observatory that were acquired during the 2008 outburst. We compare the spectra to spectra acquired during the previous outburst and quiescent phases. We find that the luminosity and full width at half maximum power of Br $\\gamma$ increased as the star has brightened and decreased when the star faded indicating that these phases are driven by variations in the accretion rate.  We also show that the temperature of the CO emission has varied with the stellar accretion rate confirming suggestions from modeling of the heating mechanisms of the inner disk (e.g. Glassgold et al. 2004). Finally we find that the lowest energy blue-shifted CO absorption lines originally reported in 2007 are no longer detected. The absence of these lines confirms the short-lived nature of the outflow launched at the start of the quiescent phase in 2006. ",
        "introduction": "In November 2003 V1647 Ori began to brighten and illuminate a reflection nebula, which became known as McNeil's Nebular Object (McNeil et al. 2004). Photometry compiled by Acosta-Polido et al. (2007) shows that V1647 Ori brightened by four magnitudes in the Cousins $I_C$ band in about three months. We refer to this event as the 2003 outburst. The star faded by $\\sim$ 1 magnitude over the next 600 days, then faded an additional three magnitudes in about 100 days. For a detailed summary of the 2003 outburst, see Aspin\\& Reipurth (2009). In August 2008, Itagaki (2008) reported the renewed brightening of V1647 Ori, suggesting that the star had undergone another outburst. We refer to this event as the 2008 outburst. This is the third recorded outburst as archival data indicates that it brightened in 1966 (Aspin et al. 2006).  During the onset of the 2003 outburst, the hydrogen and helium lines had P-Cygni profiles. The absorption component of these lines was shifted by several hundred kilometers per second - indicative of a hot wind with an expansion velocity of several hundred kilometers per second (Brice\\~{n}o et al. 2004; Reipurth \\& Aspin 2004; Vacca et al. 2004, Fedele et al. 2007).  The HI absorption features varied considerably in strength throughout 2004 and 2005 (Ojha et al. 2006; Gibb et al. 2006; Fedele et al. 2007), and these blue-shifted features have reappeared in the same lines at the outset of the 2008 outburst (Aspin et al. 2009).  Over the course of the 2003 outburst, the infrared lines of CO also underwent significant changes.  At the start of the outburst, both the overtone bandheads (Vacca et al. 2004) and the fundamental ro-vibrational CO lines appeared in emission (Rettig et al. 2005).  Following the end of the 2003 outburst, February 2006, an absorption component with a blue shift of -35 km s$^{-1}$ appeared superimposed on the emission components in the fundamental lines. The relative strengths of these post-outburst absorption features were consistent with 700 K gas (Brittain et al. 2007).  The absorption feature was absent in observations of the v=1-0 P30 line nearly a year later, suggesting that the mechanism driving the outflow had been brief (Brittain et al. 2007).  It is known that the truncation radius of an accretion disk scales with the stellar accretion rate (e.g. Bouvier et al. 2007 and references therein) and that rapid changes in the truncation radius can lead to outflows (Balsara 2004). With that in mind, Brittain et al. (2007) speculated that the post-outburst outflow was driven by the realignment of the stellar magnetic field as the accretion rate suddenly dropped.  Aspin et al. (2008; hereafter ABR08), on the other hand, have suggested that the fading was due to a combination of increased extinction and a more modest decrease in the accretion rate. If so, then it is possible that the truncation radius of the disk did not shift as radically as Brittain et al. (2007) supposed and that the relatively brief CO outflow has a different origin.  In order to better characterize the evolution of the inner disk, we present high-resolution $K-$band and $M-$band spectra of V1647 Ori. We compare the evolution of Br $\\gamma$, the v=2-0 bandhead of CO, and the fundamental CO lines spanning the years from 2004 to 2009. We also compare the entire $K-$band spectrum during the 2008 outburst to the period immediately preceding the outburst. Finally we discuss the evolution of the line profiles, gas temperature, and truncation radius of the inner disk. ",
        "conclusions": "We find that the veiling of the photospheric lines in the $K-$band increased modestly but were otherwise unchanged. Thus the underlying reason for the strong correlation between these lines and those in FUors is independent of the accretion state of the system. Further, the accretion rate of the system during the 2008 outburst is a factor of 16 greater than the smallest accretion rate observed during the quiescent phase and nearly as high as the peak accretion rate observed during the 2003 outburst. The width of Br $\\gamma$ has varied with the accretion rate.  The fundamental CO spectrum indicates that the mechanism that gave rise to the blue-shifted CO lines observed as the star returned to the quiescent phase was indeed short-lived. There is no evidence of remnant cool blue-shifted gas nor has the column of cold CO absorption increased indicating that the column density of intervening material has not changed appreciably."
    },
    "0911/0911.3937.txt": {
        "abstract": " ",
        "introduction": "The collisional evolution of small body populations, such as the main belt asteroids, is generally studied by using numerical models that compute the evolution of the size and velocity distributions of objects as a result of both collisional and dynamical processes. The asteroid disruption and fragmentation algorithms used in such models contain significant uncertainties. In particular, the scaling parameter is commonly defined as the critical specific impact energy $Q^*_D$, which results in the escape of half of the target's mass (hence the label $D$ for dispersal), called also the catastrophic impact energy threshold. Thus, $Q^*_D$ is a critical function needed in all codes that include fragmentation between rocky bodies. In this paper, by numerically simulating the collisional process between small bodies, we look for the value of this catastrophic threshold as a function of the target's diameter, internal structure, as well as impact speed, and we characterize the outcome in terms of fragment size and ejection velocity distributions. Ideally, a wide range of outcomes should be determined depending on the collisional energy (i.e. not limited to the catastrophic threshold), but this requires very massive computations that will be undertaken in further studies.  The specific impact energy is defined as $Q=0.5 m_p v^2_p/M_t$, where $m_p$, $v_p$ and $M_t$ are the mass and speed of the projectile and the target's mass, respectively. The catastrophic disruption threshold $Q^*_D$ is defined as the specific impact energy leading to a largest fragment containing $50 \\%$ of the original target's mass. Up to now, some {\\it scaling laws} have been used in collisional models to define the outcome properties of a collision. These scaling laws have been developed in order to extrapolate the results of laboratory experiments on centimeter-size targets to asteroid scales. They generally employ non-dimensional ratios involving projectile's size, impact speed and target's strength, and they have led to the characterization of a relationship between $Q^*_D$ and the target's diameter $D$ over a wide range of values (from centimeters to several hundreds kilometers). Unfortunately, the relations derived from these scaling laws assume a uniformity of process, structural continuity and other idealizations that put their applicability into question. As a result, depending on the assumptions made, the relationships obtained between $Q^*_D$ and $D$ can differ over several orders of magnitude. Nevertheless, despite these discrepancies, some systematic trends arise. In particular, impacts on small objects take place in the so-called ``strength-scaling'' regime, where the fragmentation of an object is essentially governed by its tensile strength, whereas impacts on large asteroids take place in the ``gravity-scaling'' regime, where the gravity is generally assumed to control the outcome. Benz  and Asphaug (1999) found that the transition between the two regimes may occur in the range of target's diameters between $100$ and $200$ meters. Values of $Q^*_D$ have been estimated using both laboratory and numerical hydrocode experiments (see recent reviews on these topics by Holsapple et al. 2002, and Asphaug et al. 2002).  The first {\\it self consistent} study aimed at characterizing the catastrophic disruption threshold was performed by Benz and Asphaug (1999), who used a smoothed particle hydrodynamics (SPH) code to simulate the break-up of basalt and icy bodies from centimeters-scale to hundreds kilometers in diameter. Their simulations included in a simplified way the combined effects of material strength and self-gravitation, which allowed covering both the strength and the gravity regimes. However, in the latter regime, their study was still limited by the fact that the gravitational phase was not explicitly simulated; rather, an iterative procedure based on energy balance was used to identify the largest fragment formed by reaccumulation of smaller ones (see Benz and Asphaug 1999, Section 3.3 for details). Nevertheless, the mass of large fragments can be determined quite accurately using this method.  Recently, Leinhardt and Stewart (2008) started to investigate the dependency of $Q^*_D$ on the strength of the body using the hydrocode CTH (McGlaun et al. 1990) to compute the fragmentation phase and the $N$-body code {\\it pkdgrav} to compute the subsequent gravitational evolution of the fragments. They found that the value of $Q^*_D$ can vary by up to a factor of three between strong crystalline and weak aggregates. However, they did not characterize the outcome of the disruptions at $Q^*_D$ in terms of fragment size and velocity distributions and their model was not appropriate to address the case of porous materials.   In our study, we use an improved version of the SPH hydrocode of Benz and Asphaug (1999), which now includes a model adapted to porous materials (Jutzi et al. 2008) validated at small scales by a successful test against laboratory experiments using pumice targets (Jutzi et al. 2009a). Then, in order to explicitly characterize the fragments produced by gravitational reaccumulation in the gravity regime, we combine it with the gravitational $N$-body code {\\it pkdgrav}, as was done to reproduce asteroid families (e.g. Michel et al. 2001, 2003, Jutzi et al. 2009b). We next look for the catastrophic specific energy threshold as a function of target diameter in both the strength and gravity regime, and we determine in more detail in the gravity regime the outcome properties of the disruption as a function of the model used to characterize the target.  By explicitly accounting for the two main processes involved in large-scale collisions (fragmentation and gravitational interaction), our method allows us to determine not only the value of $Q^*_D$ but also the full size and ejection velocity distributions of fragments down to the resolution limit imposed by the numerical techniques. Moreover, we can model different kinds of internal structure of the target and analyze the dependency of the value of $Q^*_D$ and the fragments' properties on the target's properties. This is an important aspect since, as described for instance by Holsapple et al. (2002) and Asphaug et al. (2002), researchers developing collisional evolution codes continue to debate over which values of $Q^*_D$ are appropriate for particular material properties, internal structures, impact speeds and object diameters. In particular, the effect of the internal structure on the outcome and on the impact energy required for disruption is a crucial information, as it constrains the collisional lifetime of the small body under consideration. It also has many implications in the framework of impact risk assessment and mitigation strategies. Indeed, it is important to make sure that a planned deflection does not lead to a disruption, and therefore the energy threshold for disruption is important information.  Here, thanks to the implementation of a model of fragmentation of porous materials, we consider two kinds of parent bodies, either non-porous or porous, which are generally believed to represent bright and dark asteroids, respectively.  In the following, we present the results of our investigations, starting in Section 2 by a description of the two kinds of target model for which we will provide the specific impact energy threshold for disruption. The numerical method is then briefly detailed in Section 3. The catastrophic disruption energy threshold of non-porous and porous targets as a function of projectile's velocity and target strengths is investigated in Section 4. The outcomes in terms of fragment size and speed distributions for both models are described in Section 5 and 6, respectively. Section 7 presents a comparison between these results allowing us to assess the sensitivity of the outcome properties and impact energy values upon the internal structure. Conclusions and perspectives are presented in Section 8. ",
        "conclusions": "In this paper, we presented the results of complete simulations of disruptions that allowed the determination of the relationships between the specific impact energy threshold for disruption, called $Q^*_D$, and both the target diameter and its internal structure represented by porous and non-porous materials made of pumice and basalt. We confirmed the results from previous studies indicating that $Q^*_D$ first decreases with target size in the strength regime and then increases with target size in the gravity regime. Moreover, we found that a porous body (as defined by our nominal model) requires more energy to be disrupted than its non-porous counterpart. In the gravity regime, the situation is reversed but the difference remains small. This might explain why, to first order, collisional evolution models could get away with only a single scaling law to reproduce the main characteristics of the asteroid populations (e.g. Bottke et al. 2005), despite the wide variety of internal properties that they can have.  However, by changing the nominal values of the tensile and in particular the shear strength of our targets, we found that in the gravity regime, the value of $Q^*_D$ is not greatly influenced by the assumed strength for porous targets. Conversely, in the case of non-porous targets, the value of $Q^*_D$ decreases significantly with decreasing strength. Therefore, in the gravity regime one cannot reasonably assume systematically that porous targets are easier or more difficult to disrupt than non-porous ones, as it depends on the assumed strengths of the latter. We plan to investigate the effect of other material properties as well as other kinds of materials in the future.  The value of $Q^*_D$ also depends on the impact velocity and we find similar tendency as previous studies, namely an increase of $Q^*_D$ with the impact speed. We then propose a scaling with speed for our porous targets, and find that it requires parameters that are consistent with what is expected for porous materials.  We next determined the outcome properties of the disruptions at $Q^*_D$, limited here to the fragment size and ejection velocity distributions in the gravity regime. We found that the size distributions keep the same qualitative aspect independent of the parent body's size and impact speed in the investigated range. For both porous and non-porous parent bodies, they can be represented by a power-law whose exponent can be used to approximate the distributions produced from all target's sizes. The average and median ejection speeds show also some systematic trends, i.e. they increase systematically with the target's diameter. Moreover, they are of the same order for both kinds of parent bodies, although slightly higher in the porous case.  These results (although limited to a particular specific impact energy) can be easily implemented in numerical algorithms aimed at studying the collisional evolutions of small body populations. They provide systematic trends in the outcome properties and scaling laws for $Q^*_D$, at least for the two kinds of target internal structure that we investigated, using nominal values of their material properties. Obviously, the real internal structures of small bodies cannot be limited to these two models and we plan to develop other models and use different material properties in order to study their resistance as well as the outcome properties of their disruption. While we limited our study to monolithic bodies, pre-shattered bodies and rubble piles, which may contain macroscopic voids and/or some microporosity, are likely to be present in the asteroid population. It is then essential to understand how macroscopic voids alone or combined with microporous properties as considered in this paper can influence the critical specific impact energy for disruption and whether they show some signature in the outcome properties. Then, using an improved version of our $N$-body code (Richardson et al. 2009), we shall also be able to characterize the spin and shape distributions of fragments, in addition to the size and velocity distributions.  It is already clear that a deep understanding of collisions between small bodies requires the investigation of a huge parameter space. As a conclusion, there is room for a great number of studies in order to characterize the impact energies and outcome properties that will allow us to provide to collisional evolution models the different recipes that are valid for all impact conditions and kinds of real small bodies. Furthermore, this information is crucial to assess the efficiency of mitigation techniques aimed at deflecting a potential impactor with the Earth, as a first step requires a characterization of which impact conditions prevent the disruption of the body and rather permit its deflection."
    },
    "0911/0911.1614_arXiv.txt": {
        "abstract": "We have studied the various properties of magnetically deformed atoms, replaced by deformed Wigner-Seitz cells, at the outer crust region of strongly magnetized neutron stars (magnetars) using a relativistic version of Thomas-Fermi model in cylindrical coordinates. ",
        "introduction": "From the observational evidence of a few strongly magnetized neutron stars, which are the sources of anomalous X-rays and soft gamma rays, also called magnetars \\cite{R1,R2,R3,R4}, the study of crustal matter, in particular the outer crust of such compact stellar objects have gotten a new dimension. These exotic objects are also called anomalous X-ray pulsars (AXP) and soft gamma repeaters (SGR). The outer crust of a typical neutron star in general, is mainly composed of dense crystalline metallic iron. The density of such metallic crystalline matter is $\\sim 10^{-4}\\rho_0$, where $\\rho_0$ is the normal nuclear density ($\\sim 2.8\\times 10^{14}$gm/cc). Therefore, it is absolutely impossible to investigate the properties of such ultra-dense matter in material science laboratories, even at zero magnetic field. The observed surface magnetic field of the magnetars is $\\sim 10^{15}$G, which is again too high to achieve in the terrestrial laboratories. Also, it is quite possible that the interior field of such exotic objects can go up to $\\sim 10^{18}$G (which can be shown theoretically by Virial theorem). If the magnetic field at the interior is really so high, then most of the physical and chemical properties of the dense neutron matter should change significantly from the conventional neutron star (radio pulsars) scenario (see the recent article by one of the co-authors of this article \\cite{R5} for necessary references). Even though the magnetic field strength at the crustal region is slightly higher than $10^{15}$G, it must change significantly most of the properties of dense matter, both in the outer crust and the inner crust regions of the magnetars \\cite{R5,R6}. It is believed that strong magnetic field can cause a structural deformation of the metallic atoms present in the outer crust of a neutron star. The spherical symmetry of the atoms will be destroyed and becomes cigar shape with the elongated axis along the direction of strong magnetic field. The atoms may even become almost an one dimensional string like object, i.e., needle shape, if the magnetic field strength is extremely strong. In a future article we shall present the problem related to structural deformation of atoms in a strong quantizing magnetic field using a completely different approach \\cite{RR6}.  In this article we shall study the properties of outer crust matter composed of magnetically deformed metallic (iron) atoms. In section 2 we have developed the basic formalism and discuss the numerical results, whereas in the last section we have given the conclusion and the future prospect of this work. In this article, for the sake of simplicity, we shall assume a cylindrical deformation of the atoms in the outer crust region and use cylindrical coordinate system with azimuthal symmetry. In reality, to investigate the cigar like deformed atoms in presence of strong magnetic field one has to use prolate spheroidal coordinate system. ",
        "conclusions": "In this article we have investigated various physical properties of non-uniform electron gas within the cylindrically deformed atoms of metallic iron at the outer crust of a strongly magnetized neutron star (magnetar). Because of extremely strong magnetic field, we have assumed a cylindrical type deformation of the atoms, replaced by WS cells with the same kind of deformation. The axes of the cylinders are along the direction of magnetic lines of forces. The curved surfaces of these cylinders are therefore parallel to the boundary surface of the neutron stars in the region far away from the magnetic poles. We have studied the variation of electron number density, kinetic energy density, electron-nucleus interaction energy per unit volume of electron gas and also kinetic pressure of non-uniform electron gas within the cylindrically deformed WS cells with the strength of magnetic field  for a given spatial coordinate ($r,z$) and also with the radial and axial distances within the cell for a given magnetic field value. We have also investigated the variations of longitudinal and transverse dimensions of deformed WS cells with the strength of magnetic field. We have noticed that the transverse dimension of a cylindrically deformed WS cell becomes extremely thin in presence of ultra-strong magnetic field. It is believed that the atoms become cigar shape in presence of strong magnetic field. In our future study, the properties of neutron star crustal matter with cigar shape atoms in the metallic crystal in presence of strong magnetic field. In future, we shall also evaluate the electron-electron direct interaction energy and the exchange part of electron-electron interaction."
    },
    "0911/0911.1087_arXiv.txt": {
        "abstract": "{We explore cosmological consequences of two quintessence models in which the current cosmic acceleration is a transient phenomenon. We argue that one of them (in which the EoS parameter switches from freezing to thawing regimes) may reconcile the slight preference of observational data for freezing potentials with the impossibility of defining observables in String/M-theory due to the existence of a cosmological event horizon in asymptotically de Sitter universes. ",
        "introduction": "Dark energy seems to be the first observational piece of evidence for new physics beyond the domain of the Standard Model of particle physics and may constitute a link between cosmological observations and a more fundamental theory of gravity. Among the many possible candidates for this dark component, the energy density associated with the quantum vacuum or the cosmological constant ($\\Lambda$) emerges as the simplest and the most natural possibility. From the observational point of view, flat models with a very small cosmological term ($\\rho_{\\Lambda} \\simeq 10^{-47}$ ${\\rm{GeV}}^4$) are in good agreement with almost all sets of cosmological observations, which makes them an excellent description of the observed Universe~(For recent reviews, see \\citep{review,paddy,rp,cop,alcanizrev,friedcce}).  From the theoretical viewpoint, however, the situation is rather different and some questions still remain unanswered. First, it is the unsettled situation in the particle physics/cosmology interface [the so-called cosmological constant problem (CCP)~\\citep{weinberg}], in which the cosmological upper bound differs from theoretical expectations ($\\rho_{\\Lambda} \\sim 10^{71}$ ${\\rm{GeV}}^4$) by more than 100 orders of magnitude. The second is that, although a very small (but non-zero) value for $\\Lambda$ could conceivably be explained by some unknown physical symmetry being broken by a small amount, one should be able to explain not only why it is so small but also why it is exactly the right value that is just beginning to dominate the energy density of the Universe now. Since both components (dark matter and dark energy) are usually assumed to be independent and, therefore, scale in different ways, this would require an unbelievable coincidence, the so-called coincidence problem (CP).  A third question also arises if we try to reconcile the $\\Lambda$CDM description of the current cosmic acceleration with the only candidate for a consistent quantum theory of gravity we have today, i.e., Superstring (or M) theory. As is well known, in the standard cosmological scenario, after radiation and matter dominance, the Universe asymptotically enters a de Sitter phase with the scale factor $a(t)$ growing exponentially, which results in an eternal cosmic acceleration.  In such a background, the cosmological event horizon \\begin{equation} \\Delta = \\int_{t_0}^{\\infty}{\\frac{dt}{a(t)}} \\quad \\mbox{converges}, \\end{equation} and this is particularly troublesome for the formulation of String/M theory because local observers inside their horizon are not able to isolate particles to be scattered, which implies that a conventional S-matrix cannot be built. Since the only known formulation of String theory is in terms of S-matrices (which require infinitely separated noninteracting in and out states), we are faced with an important and challenging task of finding alternatives to the conventional S-matrix or, equivalently, defining observables in a string theory described by a finite dimensional Hilbert space~[see \\cite{fischler,suss,haylo} for more on this subject]\\footnote{It is worth noticing that this problem is not strictly related to $\\Lambda$, but a consequence of eternal acceleration.}.  A possible way out of this dark energy/String theory conflict is to construct a dark energy scenario that predicts the possibility of a transient cosmic acceleration. In fact, this possibility can be achieved in the context of the so-called thawing~\\cite{pngb,jsa} and hybrid~\\cite{hybrid} scalar field models in which a new deceleration period will take place in the future\\footnote{\\emph{Thawing} models describe a scalar field whose the equation-of-state (EoS) parameter increases from $w \\sim -1$, as it rolls down toward the minimum of its potential, whereas \\emph{freezing} scenarios describe an initially $w > -1$ EoS decreasing to more negative values~\\cite{caldwell,bs,sch}. \\emph{Hybrid} models are characterized by EoS that switches from freezing to thawing regimes or vice-versa.}. Examples of transient cosmic acceleration can also be found in brane-world cosmologies~\\cite{sahni}, as well as in models of coupled quintessence (dark matter and dark energy), as recently discussed in Refs.~\\cite{ernandes,nelson}.  In this short contribution, I will focus on two specific scenarios of transient acceleration that are derived from two different \\emph{ans\\\"atze} on the scale factor derivative of the field energy density~\\cite{jsa,hybrid}. Both scenarios are natural generalizations of the exponential potential studied by~\\cite{RatraPeebles} and admit a wider range of solutions. ",
        "conclusions": ""
    },
    "0911/0911.4946_arXiv.txt": {
        "abstract": "The first detection of ammonia (\\amm) is reported from the Magellanic Clouds. Using the Australia Telescope Compact Array, we present a targeted search for the ($J$,$K$) = (1,1) and (2,2) inversion lines towards seven prominent star-forming regions in the Large Magellanic Cloud (LMC). Both lines are detected in the massive star-forming region \\nw, which is located in the peculiar molecular ridge south of 30\\,Doradus, a site of extreme star formation strongly influenced by an interaction with the Milky Way halo. Using the ammonia lines, we derive a kinetic temperature of $\\sim16$\\,K, which is 2-3 times below the previously derived dust temperature. The ammonia column density, averaged over $\\sim17\\arcsec$ is $\\sim6\\times10^{12}$\\,\\cden\\ ($<1.5\\times 10^{13}$\\,\\cden\\ over $9\\arcsec$ in the other six sources) and we derive an ammonia abundance of $\\sim4\\times10^{-10}$ with respect to molecular hydrogen. This fractional abundance is 1.5-5 orders of magnitude below those observed in Galactic star-forming regions. The nitrogen abundance in the LMC ($\\sim10$\\% solar) and the high UV flux, which can photo-dissociate the particularly fragile \\amm\\ molecule, must both contribute to the low fractional \\amm\\ abundance, and we likely only see the molecule in an ensemble of the densest, best shielded cores of the LMC. ",
        "introduction": "Our position within the disk of the Milky Way makes it very difficult to understand galaxy--wide effects that may influence the formation of molecular clouds, their stability against gravitational collapse, and the subsequent formation of stars. Ideally, one would like to study a more face-on, nearby galaxy and fortunately such objects exist, the Magellanic Clouds. The Large and Small Magellanic Clouds (LMC, SMC) are at distances of $\\sim50$\\,kpc and $\\sim60$\\,kpc, respectively \\citep[][this LMC distance is used throughout the paper]{mad91,gie98,kel06}. The proximity is ideal for observations to provide the high detail needed to study the formation and evolution of giant molecular clouds and even to ``zoom'' into individual prestellar cores (at the distance of the LMC 1\\arcsec\\ corresponds to 0.24\\,pc). The LMC exhibits a favorable inclination of $\\sim35\\degr$ \\citep{vdm01}. This value is ideal for both, a detailed census of all molecular clouds and their distribution as well as to determine their kinematics from radial velocities. The LMC may therefore be the very best object to study the influence that galaxy-wide processes exert on the star formation (SF) properties of an entire galaxy.   The ``fuel'' for SF is molecular gas. Good knowledge of its physical and chemical condition, such as temperature and density, is indispensable if we are to understand under which conditions, and through which processes, and phases the molecular clouds collapse to protostars. Molecular lines are a rich source for such information. The sites of actual SF, the prestellar cores, are typically very dense and cold and they may be surrounded by a warmer, more tenuous molecular envelope. However, rotational lines from linear molecules, such as the very abundant CO, have only limited power to distinguish between the two phases. Radiative transfer models, such as ``Large Velocity Gradient'' (LVG) codes \\citep[e.g.][]{sco74,kwa75,vdt07}, show that cold and dense gas on one hand, and hot and tenuous gas on the other hand, produce similar rotational line ratios. This degeneracy can, in principle, be broken by observing a large number of rotational lines, including a variety of isotopomeres or even different molecular species \\citep[e.g.,][]{wei01a,wei07}. Such an approach requires extremely well-calibrated spectra at millimeter and sub-millimeter wavelengths, which are difficult to obtain. Different beam sizes lead to additional complexity and resulting parameters also depend on the detailed properties of the applied radiative transfer model.  A much more direct and therefore more reliable path is to observe lines from molecules with non--linear structures. One of the most abundant non--linear molecule in the Universe is ammonia (\\amm). \\amm\\ is a tetrahedral symmetric--top with the exceptional ability for the nitrogen atom to tunnel through the plane defined by the three hydrogen atoms. This peculiar feature produces a series of energy doublets, and the transitions between the levels of such pairs are known as inversion transitions. As a symmetric-top, the level scheme of \\amm\\ is not, as in the case of CO, characterized by a single ($K$=0) rotational ladder. Instead, there is a multitude of $K$-ladders, whose relative populations are determined by approximately Boltzmann-distributed collisional processes, providing direct information on kinetic temperatures.  Despite previous searches for ammonia in the Magellanic Clouds \\citep[e.g.][]{ost97}, the molecule has not been detected in the past. In this paper we provide deeper data and present the first detection of ammonia in the Magellanic system. ",
        "conclusions": "\\label{sec:discuss}   \\subsection{The Kinetic Temperature of  \\nw} \\label{sec:T}  Assuming the optically thin case (see above), the column densities of the upper ammonia levels of the metastable ($J$,$J$) inversion doublets, $N_{\\rm u}$($J$,$J$) (in units of \\cden), can be derived using    \\begin{equation} N_{u}(J,J)=\\frac{7.77\\times10^{13}}{\\nu}\\frac{J(J+1)}{J^{2}}\\,\\,\\,\\int T_{\\rm mb}\\,\\, dv \\label{eq:n} \\end{equation} \\citep[e.g.][with $T_{\\rm mb}$ denoting the main beam brightness temperature of the $(J,J)$ inversion line in Kelvin, $v$ being the velocity in \\kms, and $\\nu$ representing the frequency in GHz]{hen00}. Our values for \\nw\\ are listed in Table\\,\\ref{tab:params}. For a single ortho- or para- ammonia species (the observed lines in \\nw\\ all belong to para-ammonia) the Boltzmann distribution for a rotational temperature $T_{\\rm rot,JJ'}$ (in K) is described by  \\begin{equation} \\frac{N_{u}(J',J')}{N_{u}(J,J)}=\\frac{2J'+1}{2J+1}\\,\\exp\\left(\\frac{-\\Delta E}{T_{\\rm rot,JJ'}}\\right). \\label{eq:t} \\end{equation} Here, the energy difference between the $(J,J)$ and $(J',J')$ level is given by $\\Delta E$ [41.2\\,K between \\amm\\ (1,1) and (2,2)]. We derive a rotational temperature $T_{\\rm rot,12}$ of $(17\\pm2)$\\,K for the high-velocity \\amm\\ component and a very similar $(15\\pm5)$\\,K for the low-velocity component (uncertainties are based on the $\\pm1\\sigma$ values for $N_{\\rm u}$ derived in Eq.\\,\\ref{eq:n}, not quadratically propagated but using the extremes). Rotational temperatures are always lower limits to the kinetic gas temperatures. However, for temperatures $\\lesssim20$\\,K the differences between $T_{\\rm kin}$ and $T_{\\rm rot,12}$ are less than 10\\% and thefore negligible \\citep[e.g.][]{wal83,dan88}. We can safely assume $T_{\\rm kin} = T_{\\rm rot,12}$ for both components of \\nw.  Thermal linewidths of gas at such cold kinetic temperatures are $\\sim$0.2\\,km\\,s$^{-1}$. The measured \\amm\\ linewidths are at least an order of magnitude larger than this figure and deconvolving the lines by the instrumental velocity resolution and the applied boxcar smoothing is not sufficient to reach the predicted thermal widths. The gas kinematics must thus be dominated by other processes such as turbulence or by systematic motions of spatially unresolved cloud components. This is supported by the line widths of other molecules which are similar to those of \\amm\\ \\citep[e.g.][]{hei99}.  The kinetic temperature of the gas traced by \\amm\\ is colder than the 30-40\\,K dust temperature that is measured in \\nw\\ \\citep[][]{hei99,bol00}. In their LVG analysis of high excitation molecular gas, \\citet{bol05} establish a two-component model for \\nw, which comprises a gas phase with $\\sim20$\\,K kinetic temperature and a volume density of $\\sim10^{5}$\\,\\vden, as well as a second phase with $\\sim100$\\,K and $\\sim10^{2}$\\,\\cden. Likely, the detected ammonia traces the low temperature component, albeit at larger volume densities (for volume density estimates based on \\amm, see Sect.\\,\\ref{sec:mass}). This could be in the form of cold cores surrounded by a warmer envelope. As mentioned above, dust temperatures in the region have been estimated to be $\\sim2-3$ times that of \\amm\\ and it is likely that the dust is not only present in the cold cores but also occupies the interface to the warm envelope, covering a rather wide range of densities.   The detailed dust and gas morphologies have implications for the shielding against UV radiation. The \\amm\\ molecule is quite susceptible to photo-dissociation \\citep{sut83}. \\nw\\ itself is a rather strong source of ionizing photons \\citep[e.g.][also see Sect.\\,\\ref{sec:sf}]{jon05} and a smaller contribution comes from the most luminous star-forming region in the Local Group, 30\\,Doradus, which is about $\\sim600$\\,pc away from \\nw\\ \\citep[see, e.g.][]{pin09}. The flux of energetic UV photons may penetrate into the envelope and destroy ammonia. It requires dense cores to provide enough shielding to keep ammonia intact and it is therefore likely that the molecule resides within such dense, cold, compact cores. In addition, unlike CO, nitrogen-bearing species such as \\amm\\ and N$_{2}$H$^{+}$ are less prone to depletion on dust grains at densities typical for cold dense cores such as protostellar cores \\citep[$>10^{5}$\\,\\vden; e.g.][for a density estimate of \\nw\\ see Sect.\\,\\ref{sec:mass}]{taf04,flo05} and \\amm\\ is expected to be abundant in such objects.    \\subsection{Densities, Masses, and Abundances}  \\label{sec:mass}   For a given single rotational temperature, total \\amm\\ column densities, $N({\\rm NH_{3}}, \\rm total)$, can be derived by extrapolating the level distribution over all energies. \\citet{ung86} provide this extrapolation in their Eq.\\,A15, which only requires $T_{\\rm rot, 12}$ and $N_{\\rm u}(1,1)$ to be determined. For \\nw\\ we derive a total, 17\\arcsec\\ beam-averaged, ammonia column density of $(5.8\\pm0.6)\\times10^{12}$\\,\\cden. Toward the same position, the integrated CO($1\\rightarrow0$) intensity is $\\sim45$\\,K\\,\\kms\\ over a 45\\arcsec\\ beam \\citep[Mopra data taken from][see also Fig.\\,\\ref{fig:map}]{ott08}. Assuming a CO to \\mh\\ conversion factor of $X_{\\rm CO}\\sim4\\times 10^{20}$\\,\\cden\\,(K\\,\\kms)$^{-1}$ for the LMC \\citep{pin09}, the \\mh\\ column density can be estimated to be $\\sim1.8\\times10^{22}$\\,\\cden. With this value, the fractional ammonia abundance in the \\nw\\ complex becomes $\\sim 3.3\\times10^{-10}$. This number is based on a 45\\arcsec\\ average of the \\mh\\ column density. The \\mh\\ distribution within that area, however, is most likely not smooth but exhibits some variation. To excite \\amm, a critical density needs to be exceeded. We do not observe \\amm\\ outside the smaller, 17\\arcsec\\ ATCA beam, which indicates that the \\mh\\ density is likely higher at the position of the \\amm\\ emission than further away from it. Consequently, the ammonia abundance of $\\sim 3.3\\times10^{-10}$ may be considered an {\\it upper} limit and any \\mh\\ column exceeding the average over 45\\arcsec\\ will dilute the ammonia abundance. Furthermore, larger $X_{\\rm CO}$ factors, such as reported by \\citet{isr03} or \\citet{fuk08}, would also drive the fractional ammonia abundance to smaller values.     The $X_{\\rm CO}$ factor is just one way to estimate the total \\mh\\ content of the molecular cloud. A different approach is to derive virial masses $M_{\\rm vir}$, assuming that the system is virialized and the masses of the clouds are dominated by \\mh. For a cloud with a Gaussian density profile, the virial mass can be estimated using  \\begin{equation} M_{\\rm vir}=444\\,r\\,\\Delta v_{\\rm 1/2}^{2} \\label{eq:virial} \\end{equation} \\citep[][]{pro08} with the radius $r$ in pc, the full width at half maximum (FWHM) velocity $\\Delta v_{\\rm 1/2}$ in \\kms, and the virial mass $M_{\\rm vir}$ in M$_{\\sun}$ \\citep[note that the numerical factor changes with the shape of the radial density profile; e.g. it is $\\sim 2$ times lower for $\\rho=\\rm constant$, and $\\sim3.5$ times lower for a $\\rho\\propto r^{-2}$ profile;][]{mla88}. The \\amm\\ emission position angle may show some deviation from a pure point source but, given the faintness of the source, we use half of the FWHM ATCA beam size as an upper limit to the radius of the cloud, 8\\arcsec\\ which corresponds to $\\sim2$\\,pc. Together with $\\Delta v_{\\rm 1/2}\\sim5$\\,\\kms, the virial mass amounts to an upper limit of $<2.2\\times10^{4}$\\,\\msun\\ (for comparison, the virial mass based on the Mopra CO data over the much larger Mopra beam would be $\\sim10^{5}$\\,\\msun). This limit also accounts for the boxcar smoothing of the data; a more narrow line would lower the upper limit and our estimate is therefore a more conservative one. In addition, this estimate also accounts for material other than \\mh\\ that is contributing to the virial mass, justifying the upper limit nature of the \\mh\\ mass based on the virial method.   Using a spherical geometry of the \\amm\\ emission and assuming that all the virial mass is due to \\mh, the mean volume density of the molecular gas then can be estimated to $\\lesssim 1.3\\times10^{4}$\\,\\vden. This value is fairly typical for dense gas in Galactic star-forming regions, spread over the physical size of the beam, $\\sim4$\\,pc \\citep[e.g.][]{fos09}. Finally, if we take the column density of ammonia of $(5.8\\pm0.6)\\times10^{12}$\\,\\cden\\ (see above), we derive a beam averaged ammonia mass of $\\sim 1\\times10^{-5}$\\,\\msun, or $\\sim3.3$\\,M$_{\\earth}$. With \\mh\\ masses and densities being upper limits, the fractional abundance of ammonia (molecule number density ratio) then follows as a {\\it lower limit} of $\\gtrsim4.5\\times 10^{-10}$. This value is in agreement with the upper limit derived via $X_{\\rm CO}$ above.   We cannot infer any robust variation in fractional ammonia abundance between the two velocity components. The CO(1$\\rightarrow$0) emission is about two orders of magnitude brighter (cf. Fig.\\,\\ref{fig:spec}) at the velocity of the ammonia high-velocity component. The ratio appears to be somewhat smaller for the low-velocity component, which would indicate some variation in the abundance. But, given the uncertainties in beam filling factors for the CO and \\amm\\ observations (see also Table\\,\\ref{tab:params}), better data is needed to reliably establish such a variation.  Using the above values for temperature and upper volume density limit for \\nw, we estimate the pressure of the clumps in the \\nw\\ region to be $P/k\\gtrsim2\\times10^{5}$\\,K\\,\\vden. This value is likely to increase substantially when the average value for \\mh\\ densities is resolved into individual, smaller clumps with relatively little \\mh\\ in between. For comparison, \\citet{won09} calculate the midplane pressure for the positions where molecular gas is observed in the LMC. Their highest value is close to our lower limit, which may indicate that the dense \\amm\\ clumps are not in pressure equilibrium with their surroundings, a property that is typical for self-gravitating clouds. We like to note though that the volume density of our calculation is actually based on the virial theorem, i.e. a gravitational equilibrium of the molecular gas that is independent of its environment.  We can now revisit the upper limits of the ammonia survey toward the other six positions. If we assume the same linewidth and kinetic temperature that we find for the two components of \\nw, the $3\\sigma$ upper limits to the non--detections translate to ammonia columns $N_{\\rm u} (1,1)$ and $N ({\\rm NH_{3}}, \\rm total)$ of $<5\\times10^{12}$\\,\\cden, and $<1.5\\times10^{13}$\\,\\cden\\ for a $9\\arcsec$-sized region, respectively.    \\subsection{Star Formation Properties of \\nw}  \\label{sec:sf}  If we take the upper limit to the virial mass of the ammonia emitting region of $\\sim2\\times10^{4}$\\,\\msun\\ for the mass of the molecular clump (see Sect.\\,\\ref{sec:mass} above), the typical SF efficiency of a few percent \\citep[][]{eva09} would convert the gas to stellar masses of $\\sim10^{3}$\\,\\msun. This falls within the mass range of open stellar clusters. To be independent of extinction effects \\citep[the very nearby ionizing source \\nw\\ AN is deeply enshrouded by dust as shown by][]{jon05}, we use the 1.4\\,GHz radio continuum map presented in \\citet{hug07} to estimate the current SF rate (SFR) via the conversions provided by \\citet{haa00}. For \\nw\\ AN, we measure a flux of $\\sim260$\\,mJy which corresponds to a current SFR of $\\sim8\\times10^{-4}$\\,\\sfr. The radio emission at that position, however, cannot be entirely separated from other nearby sources and we estimate an uncertainty to the flux and the SFR of up to 50\\%. If SF continues at that level, the available molecular gas will be fully converted into stars within the next $\\sim10$\\,Myr.   \\subsection{Comparison of  \\nw\\ with Galactic and Extragalactic Star Forming Regions}  \\label{sec:galactic}    The LMC has a lower mass and metallicity than the Galaxy. In addition, the \\nw\\ region is located in the molecular ridge south of the extremely active 30\\,Doradus star-forming region, at a location where the LMC's regular disk becomes disturbed by tidal and ram pressure forces that act as the galaxy moves through the Milky Way halo \\citep[e.g. ][]{ott08}. Do these peculiarities change the properties of the molecular clumps formed at this location? In order to address this question we compare our results to the properties of dense molecular cores in the Galaxy. A homogenized sample of molecular cores, drawn from Galactic star-forming regions such as Orion, Ophiuchus, Perseus, Taurus, Serpens, Cepheus, etc. is presented in \\citet{jij99}. The authors have collected ammonia (1,1) and (2,2) observations for all of their sample and can therefore provide data that may act as a reference for comparison with clouds in other galaxies such as the LMC. More recently, \\citet{ros08} \\citep[see also][]{stu92,fos09} published a comprehensive ammonia survey of 193 dense molecular cores toward the Perseus molecular cloud, a cloud that may also be representative for typical star-forming regions in the Galaxy. The studies agree that Galactic dense molecular cores exhibit typical kinetic gas temperatures of $\\sim10$\\,K, with sizes of $\\sim0.1$\\,pc, velocity widths of $\\sim0.1-0.7$\\,\\kms, ammonia column densities of $\\sim10^{14.5}$\\,\\cden, and virial masses of $\\sim5$\\,\\msun. In most aspects these values are clearly different from what we measure toward the LMC cloud \\nw. The difference can be explained by observing an ensemble of cores within the ATCA beam. Such a site could potentially be the birth of an entire stellar cluster (see also Sects.\\,\\ref{sec:mass} and \\ref{sec:sf}). Along these lines, the \\nw\\ region may consist of individual, dense molecular cores that are at temperatures very similar to those found in the Galaxy along with some inter-clump material. Together, the line profiles of the clumps and the more diffuse gas blend to form the measured ammonia lines with a width of $\\sim5$\\,\\kms\\ (note that \\citealt{ros08} also largely attribute the upper end of their measured linewidths to blending of a number of cores along the line of sight). We would like to note that the \\nw\\ SF timescales as derived in Sect.\\,\\ref{sec:sf} also match those that are typical in Galactic SF regions \\citep{eva09}.    The differences in ammonia column densities and fractional abundances between Galactic cores and the emission in \\nw, however, cannot be reconciled by beam filling arguments. The \\amm\\ column densities in \\nw\\ are about 50 times lower ($\\log(N[NH_{3}])\\sim12.8$ for \\nw, $\\sim14.5$ for Galactic SF regions). If beam filling would be responsible for the difference, the ATCA beam would be diluted by about the same factor. This implies that of the 20\\,pc$^{2}$ ATCA beam area only 0.4\\,pc$^{2}$ is covered by clouds. If these clouds are similar to Galactic cores (with typical sizes of 0.1\\,pc and masses of $\\sim5$\\,\\msun, see above), the coverage corresponds to $\\sim50$ molecular cores adding up to a total mass of $\\sim250$\\,\\msun. In Sect.\\,\\ref{sec:mass} we derive an upper limit of $\\sim2\\times 10^{4}$\\,\\msun\\ to the virial mass in \\nw, which is about two orders of magnitude larger. This indicates that the basic assumption for the calculation, the equality of Galactic and \\nw\\ ammonia column density, is wrong.  The virial mass is an upper limit due to unknown mass components other than \\mh\\ within the beam and due to the fact that the ATCA beam was taken as the virial radius of the unresolved emission. Can the parameters play together to create the $250$\\,\\msun\\ mass as predicted from Galactic clumps? We measure a velocity width of 5\\,\\kms\\ within \\nw\\ and, according to Eq.\\,\\ref{eq:virial}, this determines the size of the gaseous emission for the $250$\\,\\msun\\ to be $0.05$\\,pc if the clouds are in virial equilibrium. Such a small value is clearly unrealistic and we conclude that it is not possible to reconcile the ammonia column densities with those of Galactic SF regions based on beam filling arguments alone. It is likely that intra-clump gas covers a larger region of the beam and/or a larger number of cores is present in \\nw. But this implies that the ammonia column densities in \\nw\\ are intrinsically lower than those in Galactic cores.  Similarly, the low fractional abundances of ammonia in \\nw\\ cannot be fully explained by clumping of the gas within the beam (note that the the abundances are independent of the beam size when applying the virial mass as in Sect.\\,\\ref{sec:mass}). \\citet{ben83} derive for Galactic Dark Clouds on average fractional ammonia abundances of $\\sim 10^{-7}$ and \\citet{mau87} obtain $\\sim10^{-5}$ for a hot core. The resolved Galactic clumps as listed in the more recent \\citet{ros08}, \\citet{fos09}, and \\citet{jij99} studies have \\amm\\ abundances of order $\\sim10^{-8}$. In contrast, we estimate a fractional ammonia abundance in \\nw\\ of $\\sim4\\times10^{-10}$ (see Sect.\\,\\ref{sec:mass}). \\nw\\ thus exhibits at least $\\sim1.5$ orders of magnitude lower fractional ammonia abundance than seen in individual cores embedded in Galactic clouds. But \\nw\\ is not an exception: the upper limits towards the other LMC star-forming regions in our survey exhibit a limit of at least one order of magnitude below the Galactic ammonia abundances. Such a difference cannot be an effect of the overall lower metallicity of the LMC alone, which is $\\sim45$\\% solar \\citep[based mostly on iron, carbon, oxygen and silicon]{luc98,rol02}. A combination of the $\\sim5$ times lower abundance of nitrogen in the LMC, $\\sim 10$\\% solar \\citep{hun07}, and the high UV flux at the \\nw\\ location \\citep[see Sec.\\,\\ref{sec:T} and also][]{deh92,bol00,pin09} that can potentially photo-dissociate the fragile \\amm\\ \\citep{sut83} must therefore contribute to its low fractional abundance.   We cannot exclude that the fraction of clump to interclump \\amm\\ emission might be very different in \\nw\\ as compared to Galactic SF regions. In such a scenario, hardly any of the molecular gas parameters can be reconciled between the two galaxies. Beam dilution effects can explain most of the differences, except for the ammonia abundance and column density, and is therefore a more natural interpretation. To fully resolve this issue, sensitive observations at better spatial resolution are required.  Among extragalactic objects, only M\\,82 is known to exhibit a similarly low fractional ammonia abundance like \\nw. \\citet{wei01b} derive a value of $\\sim5\\times10^{-10}$ for this prominent starburst galaxy. The gas in M\\,82, at $T_{\\rm kin}\\approx 60$\\,K, is warmer than in \\nw\\ and the metallicity of M\\,82 may also exceed that of the LMC, in particular close to the central starburst region \\citep[see, e.g.][]{mcl93,mar97}. The presence of nearby, strongly ionizing sources in both M\\,82 and \\nw\\ suggests that photo-dissociation plays an important role in regulating the ammonia abundances in both galaxies.        In this paper, we present a search for ammonia toward seven prominent star forming regions as well as deep observations towards the star-forming \\nw\\ region close to 30\\,Doradus. Observing the ($J$,$K$)=(1,1) and (2,2) \\amm\\ inversion lines we find the following:  \\begin{itemize}   \\item Two ammonia velocity components are observed toward \\nw, both with linewidths of $\\sim5$\\,\\kms. The components are $\\sim 8$\\,\\kms\\ apart. At least for the brighter line, we can exclude large optical depths. The linewidths and derived virial masses (of order $\\sim 10^{4}$\\,\\msun) indicate that we may see the molecular gas that forms an ensemble of stars similar to open clusters in the Galaxy.   \\item We derive the kinetic temperature of the gas to be $\\sim16$\\,K for both components, $\\sim$2-3 times colder than the dust in the region.   \\item The ammonia column densities in all of the observed LMC star-forming regions are lower than typical values in our Galaxy, even after correcting for beam dilution effects.   \\item In the absence of a large, diffuse and hot component in \\nw, the fractional ammonia abundance relative to \\mh\\ is with $\\sim4\\times10^{-10}$ within a $17\\arcsec$ ($\\sim4$\\,pc) beam, about 1.5-5 orders of magnitude lower than in Galactic star-forming clouds and prestellar cores. The difference in abundances (based on this cloud) is therefore much larger than the metallicity difference of the LMC with respect to the Galaxy (LMC metallicity: $\\sim45$\\% solar). \\amm\\ photo-dissociation from the high UV flux in the LMC as well as its lower nitrogen abundance ($\\sim10$\\% solar) potentially explain the \\amm\\ deficiency.  \\item If the gas is converted into stars with the current, radio continuum-based star formation rate, SF will cease within the typical SF timescale of a few Myr.   \\item The multi-position survey did not reveal any detection down to a limit of $\\sim120$\\,mK in a 0.8\\,\\kms\\ velocity bin. If there are molecular clumps with properties similar to those in \\nw, we derive $3\\,\\sigma$ upper limits to the total column densities at these positions to be $<1.5\\times10^{13}$\\,\\cden\\ within a 9\\arcsec\\ (2\\,pc) beam.   \\end{itemize}  The Magellanic Clouds are unique as they allow us to perform analyses similar to that possible for Galactic star forming regions, but without the confusion that occur in edge-on systems. They are also ideal to trace the effects that low metallicity has on the state and chemistry of the dense molecular gas, the material that eventually forms the stars. The kinetic temperature of the gas is a very important parameter to understand the dense gas. It will be the focus of future studies to extend the use of ammonia as a thermometer for a larger number of gas clumps in order to obtain a more representative sample."
    },
    "0911/0911.2084_arXiv.txt": {
        "abstract": "{}{We determine the steady-state of an axisymmetric thin accretion disc with an internal dynamo around a magnetised star.} {Starting from the vertically integrated equations of magnetohydrodynamics we derive a single ordinary differential equation for a thin accretion disc around a massive magnetic dipole and integrate this equation numerically from the outside inwards.} {Our numerical solution shows that the torque between the star and the accretion disc is dominated by the contribution from the dynamo in the disc. The location of the inner edge of the accretion disc varies between $R_{\\rm A}$ and  $10R_{\\rm A}$ depending mainly on the strength and direction of the magnetic field generated by the dynamo in the disc}{} ",
        "introduction": "In this paper we present a new solution for an accretion disc around a magnetic star.  This star could be a neutron star, a white dwarf, or a T Tauri-star, but we assume that it is a neutron star since it is easy to measure the torque between the neutron star and the accretion disc by timing the X-ray pulses from the neutron star.  The new feature of our solution is to include the effect of an internal dynamo in the accretion disc.  By doing this we hope to be able to explain the torque reversals that have been observed in some X-ray pulsars.  \\cite{shakura} formulated the standard model of a geometrically thin, optically thick accretion disc. They were able to obtain an analytical solution of the height-integrated hydrodynamic equations, after having introduced the $\\alpha$-prescription for the turbulent stress, which transports the angular momentum outwards through the disc; however, they did not explain why the disc is turbulent in the first place, since a disc in Keplerian rotation is stable according to Rayleigh's criterion. \\cite{balbus1} instead showed that it is unstable if there is a weak magnetic field in the disc. Subsequent numerical simulations (e.g. \\cite{hawley};  \\cite{balbus}) confirmed that this instability generates turbulence and that the resulting turbulent stresses transport angular momentum outwards.  The interaction between a magnetised  star and a surrounding accretion disc  is one of the most poorly understood aspects of accretion. At the same time, it is central for our understanding of the spin evolution of objects as diverse as T Tauri stars and  X-ray pulsars.  The magnetic field of the star  penetrates the surrounding accretion disc and couples the two. According to the Ghosh \\& Lamb (\\cite{ghosh}) model, the part of the accretion disc that is located inside the corotation radius provides a spin up torque on the star, since it is rotating faster than the star, while the more slowly rotating outer part of the accretion disc brakes the star. The net torque is determined by the location of the inner edge of the disc, which moves inwards as the accretion rate increases, thereby increasing the spin up-torque on the star.  \\cite{campbell3} proposes  physical descriptions of the magnetic diffusivity in terms of turbulence or buoyancy, and  \\cite{campbell2} solve the resulting equations numerically. For both forms of diffusivity, the magnetic coupling between the disc and the star leads to an enhanced dissipation in the inner part of the accretion disc compared to the standard \\cite{shakura} model. This raises the temperature such that electron scattering dominates Kramer's opacity at larger radii than is otherwise the case, thus increasing the fraction of the disc that is subject to the (\\cite{lightman}) instability. \\cite{brandenburg} considered a form of magnetic diffusivity that allows further analytical progress to be made, but the qualitative results remain the same.  All the  models predict a positive correlation between the accretion rate and the torque on the neutron star and even predict a negative torque on the neutron star at very low accretion rates. Timing of X-ray pulsars during outbursts of Be/X-ray transients have provided at least some qualitative support for such a correlation (e.g. \\cite{parmar}). The BATSE instrument on the Compton Gamma Ray Observatory made it possible to extend this database significantly (\\cite{bildsten}).  In particular there are a few X-ray pulsars with permanent discs that are oscillating between phases of constant spin-up and constant spin-down  without a significant difference in the X-ray luminosity between these states, which appears to contradict the standard model for a disc-accreting X-ray pulsar.  \\cite{nelson} propose that these torque-reversals can be the result of transitions between co-rotating and counter-rotating accretion discs, though several other models have also been proposed. \\cite{torkelsson} argue that the torque between an accreting star and its disc can be enhanced by the presence of a magnetic field generated by the turbulence in the accretion disc. The torque reversals are then the result of a reversal of the magnetic field generated by this dynamo. However, he did not construct a self-consistent model of the accretion disc. The aim of this paper is to construct such a model of an accretion disc with an internal dynamo around a magnetic star. We work in the spirit of \\cite{shakura} and assume that the disc is geometrically thin. In Sect. 2 we start from the equations of magnetohydrodynamics (MHD) and derive a single ordinary differential equation for the radial structure of the accretion disc. We then present numerical solutions of this equation in Sect. 3 and discuss the properties of these solutions in Sect. 4. Finally we summarise our conclusions in Sect. 5. ",
        "conclusions": "We have investigated the interaction between a magnetic neutron star and its surrounding accretion disc in the case where the accretion disc is supporting an internal dynamo.  The magnetic field that is produced by the dynamo can lead to a significant enhancement of the magnetic torque between the neutron star and the accretion disc, compared to what is seen in the model by Ghosh \\& Lamb (\\cite{ghosh}). This extra magnetic torque can explain the large variations in spin frequency \\object{Cen X-3} and \\object{OAO 1657-415} (Bildsten et al.1997). Furthermore, a reversal of  the magnetic field that is generated by the dynamo, similar to the reversals that we see of the magnetic fields on the Sun, could explain the torque reversals in these objects.  From the way that we calculate the structure of the accretion disc, we find two kinds of solutions with different behaviours at the inner edge. A few of our solutions have case D boundaries at which the density and temperature go to 0 at finite radius, while most of our solutions have case V boundaries at which the accretion is driven entirely by the magnetic tension between the accreting matter and the neutron star.  In this case there is a viscous stress between the accretion disc and the boundary layer, which can transfer angular momentum outwards at a rate that is comparable to the one at which it is advected inwards by the accreting matter itself.  We have also found that the dynamo leads to that the inner edge of the accretion disc occurs at a radius that is larger than the traditional Alfv\\'en radius.  This effect is weak, though, for a realistic value of the dynamo-generated magnetic field."
    },
    "0911/0911.2798_arXiv.txt": {
        "abstract": "We investigate the relationship between brightest cluster galaxies (BCGs) and their host clusters using a sample of nearby galaxy clusters from the Representative \\xmm\\ Cluster Structure Survey (\\rexcess).  The sample was imaged with the Southern Observatory for Astrophysical Research (SOAR) in R~band to investigate the mass of the old stellar population.  Using a metric radius of 12\\unit{h^{-1}}~kpc, we found that the BCG luminosity depends weakly on overall cluster mass as $L_{BCG} \\propto M_{cl}^{0.18\\pm0.07}$, consistent with previous work.  We found that 90\\% of the BCGs are located within 0.035~\\rfive\\ of the peak of the X-ray emission, including all of the cool core (CC) clusters.  We also found an unexpected correlation between the BCG metric luminosity and the core gas density for non-cool core (non-CC) clusters, following a power law of $n_{e} \\propto L_{BCG}^{2.7\\pm0.4}$ (where $n_e$ is measured at 0.008~\\rfive).  The correlation is not easily explained by star formation (which is weak in non-CC clusters) or overall cluster mass (which is not correlated with core gas density).  The trend persists even when the BCG is not located near the peak of the X-ray emission, so proximity is not necessary.  We suggest that, for non-CC clusters, this correlation implies that the same process that sets the central entropy of the cluster gas also determines the central stellar density of the BCG, and that this underlying physical process is likely to be mergers. ",
        "introduction": "\\label{intro}  The intracluster gas in some galaxy clusters shows a central concentration in a cool core.  The higher density of this core gas allows more rapid energy loss in the form of X-ray emission.  In the absence of other energy sources, the classic ``cooling flow'' model suggests that the central gas should cool.  The cool gas would rapidly condense and form stars at rates greater than 100$\\msun$\\unit{yr^{-1}}, but optical colors and spectral lines indicate much lower rates of less than $\\sim$10\\unit{\\msun yr^{-1}} (reviewed in \\citealp{donahue04a}).  Thus, the gas must be heated by other processes, such as AGN activity or other forms of feedback that regulate the thermal properties of the gas (reviewed in \\citealp{mcnamara07a}).  Yet the identity of the feedback process is still widely debated, and current simulations tend to overpredict the fraction of clusters with cool cores (\\eg \\citealp{kay07a}).  The gas core is often located near the Brightest Cluster Galaxy (BCG), particularly in cool core clusters \\citep{jones84a,bildfell08a,rafferty08a,sanderson09b}.  The BCG is not only luminous and massive, but has a more extended optical light profile than other ellipticals, due to its unique merger history at the bottom of the gravitational potential well (\\eg\\ \\citealp{hausman78a,vale08a}).  Models predict that the BCG mass is correlated with the total cluster mass (\\eg \\citealp{delucia07b}), but the observed correlation is typically weaker \\citep{lin04a,popesso07a,brough08a,whiley08a,yang08a,mittal09a}.  To better understand the connection between BCGs and cool cores, we investigated the BCGs in 31 southern hemisphere clusters from the Representative \\xmm\\ Cluster Structure Survey (\\rexcess) \\citep{boehringer07a}.  The \\rexcess\\ sample was chosen to evenly represent the range of X-ray luminosities in the local ($0.06<z<0.18$) population, with no bias regarding X-ray morphology or central surface brightness.  Thus, it provides an ideal laboratory for testing the connections between BCGs and their host clusters.  \\cite{pratt07a}, \\cite{croston08a}, and \\cite{pratt09a} measure the X-ray luminosity, temperature, mass, core gas density, and cooling time. Here we present ground-based CCD (optical R~band) imaging of the cluster BCGs and calculate colors relative to 2MASS K~band magnitudes. We investigate the connections between the old stellar population and the properties of the X-ray gas.  Interestingly, the strongest correlation we found was not between the BCG mass and the total cluster mass, but between the central BCG stellar density and the core gas density.  We describe the cluster sample and X-ray measurements in \\S\\ref{x.obs}, and the R~band observations and BCG magnitude calculations in \\S\\ref{opt}.  In \\S\\ref{results} and \\ref{special} we examine this new correlation, as well as other connections between BCGs and their host clusters.  Our conclusions are summarized in \\S\\ref{summary}.  We use $h=0.7$, $\\Omega_m=0.3$, and $\\Omega_\\Lambda=0.7$ throughout. ",
        "conclusions": "\\label{summary}  Using the \\rexcess\\ sample of galaxy clusters, we investigated the relationships between the BCG stellar mass and the properties of the X-ray emitting gas.  We confirmed the weak correlation seen by others between the BCG luminosity and the total cluster mass, and the close proximity of the BCG to the X-ray peak.  We detected a trend among the non-CC clusters in which the core gas density increases with the BCG aperture luminosity, a proxy for the BCG central stellar mass density.  This trend is much clearer when the BCG luminosity is determined from well-masked aperture magnitudes rather than from unmasked or isophotal magnitudes.  This trend holds even in cases where the gas is disturbed or the BCG is located far from the central region.  We argue that the core gas density is an indicator of the deviation of the central entropy from the self-similar value.  Thus, the core gas entropy and the central BCG stellar density appear to be more closely related than previously thought.  We suggest that cluster mergers could be the underlying cause, since mergers can both increase gas entropy and decrease BCG central stellar density.  If so, this trend could set important constraints on models of the central gas in clusters and of BCG formation.  These results are based on a sample of only 31 clusters, and need to be confirmed using other, larger cluster samples.  For example, \\cite{cavagnolo09a} recently fit entropy profiles to 239 clusters from the \\Chandra X-ray Observatory archive.  The optical data corresponding to this or other large samples of X-ray selected clusters would allow an investigation of the trend with a statistically significant sample size. If the trend is real and our interpretation is correct, it should be present independent of how the clusters are selected.    \\vskip .2in  \\small DBH acknowledges the support of a Calvin Research Fellowship. LRL acknowledges the support of a Sid Jansma Summer Research Fellowship. MD and GMV acknowledge the support of an LTSA grant NASA NNG-05GD82G. MD and SB acknowledge the support of MSU research funds. DBH and LRL thank the AstrW10 class of Calvin College for their assistance with observations of RXCJ1044.  We present data obtained with the Southern Observatory for Astrophysical Research, which is a joint project of Conselho Nacional de Pesquisas Cientificas e Technol\\'ogicas Brazil, the University of North Carolina at Chapel Hill, Michigan State University, and the National Optical Astronomy Observatory. This research used data obtained with \\xmm, an ESA science mission with instruments and contributions directly funded by ESA Member States and the USA (NASA). This research used data from the Two Micron All Sky Survey, a joint project of the University of Massachusetts and the Infrared Processing and Analysis Center/California Institute of Technology, funded by the National Aeronautics and Space Administration and the National Science Foundation. This research used data from the NASA/IPAC Extragalactic Database (NED) which is operated by the Jet Propulsion Laboratory, California Institute of Technology, under contract with the National Aeronautics and Space Administration. This research used data from the Sloan Digital Sky Survey, funded by the Alfred P. Sloan Foundation, the participating institutions, the National Science Foundation, the U.S. Department of Energy, the National Aeronautics and Space Administration, the Japanese Monbukagakusho, the Max Planck Society, and the Higher Education Funding Council for England.  \\normalsize"
    },
    "0911/0911.0424_arXiv.txt": {
        "abstract": "We report the results of the analysis of high resolution photospheric line spectra obtained with the UVES instrument on the VLT for a sample of 15 solar-type stars selected from a recent survey of the distribution of H and K chromospheric line strengths in the solar-age open cluster M67. \u00a0We find upper limits to the projected rotation velocities that are consistent with solar-like rotation (i.e., $v\\,\\sin{i} \\la$ 2--3\\,km\\,s$^{-1}$) for objects with Ca II chromospheric activity within the range of the contemporary solar cycle. Two solar-type stars in our sample exhibit chromospheric emission well in excess of even solar maximum values. In one case, Sanders~1452, we measure a minimum rotational velocity of $v\\,\\sin{i} = 4 \\pm 0.5$\\,km\\,s$^{-1}$, or over twice the solar equatorial rotational velocity. The other star with enhanced activity, Sanders~747, is a spectroscopic binary. \u00a0We conclude that high activity in solar-type stars in M67 that exceeds solar levels is likely due to more rapid rotation rather than an excursion in solar-like activity cycles to unusually high levels. \u00a0We estimate an upper limit of 0.2\\,\\% for the range of brightness changes occurring as a result of chromospheric activity in solar-type stars and, by inference, in the Sun itself. We discuss possible implications for our understanding of angular momentum evolution in solar-type stars, and we tentatively attribute the rapid rotation in Sanders~1452 to a reduced braking efficiency. ",
        "introduction": "In a survey of the Ca II H and K core strengths of a sample of 60 solar-type stars in the solar-age, solar-metallicity open cluster M67, \\citet{Giampapa06} found that the distribution of the HK index -- a measure of the strength of the chromospheric H and K cores -- is broader than the distribution seen in the contemporary solar cycle. Significant overlap between the HK distribution of the solar cycle and that for the sun-like stars in M67 is seen with over 70\\% of the solar analogs exhibiting Ca II H+K strengths within the range of the modern solar cycle. About $\\sim$ 10\\% are characterized by high activity in excess of solar maximum values while approximately 17\\% have values of the HK index less than solar minimum.  In view of these results, a key question that arises is whether the distribution of the HK index in M67, and especially the occurrence of solar-type stars in this cluster with levels of activity exceeding solar maximum, is due to excursions in the amplitude of otherwise solar-like cycles or whether the high HK values in some solar-type members are simply the result of their more rapid rotation. The former possibility is reminiscent of the prolonged episode of quiescence known as the Maunder Minimum where sunspots vanished, coinciding with the Little Ice Age \\citep{Foukal90}, as well as the extended period of high solar activity, as inferred from isotopes in the terrestrial record, known as the Medieval Solar Maximum during the so-called Medieval Warm Epoch \\citep{Jirikowic94}. In the case of the latter possibility of more rapid rotation, we know that atmospheric activity associated with surface magnetic fields, in general, increases with increasing equatorial rotation velocity in late-type stars (e.g., Skumanich 1972; Pallavicini et al. 1981; Baliunas et al. 1995; Pizzolato et al. 2003).  \\begin{figure*} \\centering \\includegraphics[width=.9\\textwidth]{Spectrum.eps} \\caption{\\label{fig:Spectrum}Spectral window of our Ganymede spectrum.} \\end{figure*}  The presence of active solar-type stars in the solar-age cluster M67 is particularly surprising. It is therefore crucial to determine whether it is rapid rotation that is the origin of the relatively enhanced activity, or if the rotation of these objects is solar-like but their cycle variations extend to relatively higher amplitudes than seen in the Sun. If it is more rapid rotation that is the natural origin of the higher Ca II core emission then this result would invite further investigation of the angular momentum evolution of M67 in contrast to that of the more slowly rotating Sun and other solar-age G dwarfs. However, if rotation velocities are solar-like at $\\sim$2\\,km\\,s$^{-1}$ then this would suggest that excursions in the cycle variation of solar-type stars that are significantly in excess of contemporary solar maximum values can occur in Sun-like stars and, by implication, in the Sun itself. We therefore discuss in this investigation the results of our attempt to measure projected rotation velocities in selected M67 solar-type stars from the sample of \\citet{Giampapa06}. The sample selection and the acquisition of the observations are discussed in $\\S\\S$2-3 while the novel approach to the analysis is given in $\\S$4. The results are presented in $\\S5$ and a discussion followed by our conclusions are given in $\\S\\S$6 and 7, respectively. ",
        "conclusions": "We find on the basis of the analysis of high resolution spectra of a subset of solar-type stars from the H \\& K survey of these objects in M67 by \\citet{Giampapa06} that stars with levels of chromospheric activity within the range of the contemporary cycle appear to have rotational velocities that are solar-like. At least one solar-type star with an HK index well in excess of the modern solar maximum is rotating more rapidly. Specifically, we measured a value for the projected rotation velocity of Sanders~1452 of $v\\,\\sin{i} = 4 \\pm 0.5$\\,km\\,s$^{-1}$, or more than twice the solar equatorial rotation velocity. Another object in our sample, Sanders~747, is found to be a spectroscopic binary and is likely characterized by more rapid rotation than the Sun in at least one of its components. In view of these results, we conclude that the occurrence of high-activity levels in excess of solar maximum values for solar-type stars in M67 is most likely due to relatively faster rotation rather than to an excursion of a solar-like cycle to high values at solar rotation rates. Based on the correlation found between relative chromospheric emission and mean brightness variations we estimate an upper limit of about 0.2\\% in the level of broad-band variability in solar-type stars at solar ages and, by inference, in the Sun itself.  After a consideration of current models for angular momentum evolution for low mass stars, we tentatively attribute the more rapid rotation of solar-age stars such as Sanders~1452 to a reduced efficiency of braking due to a magnetized wind. Alternatively, the rapid rotation also could result simply from the particular mass-dependence of rotational braking in M67, as is seen in other clusters. A more comprehensive $v\\,\\sin{i}$ survey in M67 with specific emphasis on stars slightly warmer than the Sun will have to be performed to explore this hypothesis further."
    },
    "0911/0911.2617_arXiv.txt": {
        "abstract": "We combine results from interferometry, asteroseismology and spectroscopic analyses to determine accurate fundamental parameters (mass, radius and effective temperature) of 10 bright solar-type stars covering the H-R diagram from spectral type F5 to K1. Using ``direct'' techniques that are only weakly model-dependent we determine the mass, radius and effective temperature. We demonstrate that model-dependent or ``indirect'' methods can be reliably used even for relatively faint single stars for which direct methods are not applicable. This is important for the characterization of the targets of the CoRoT and {\\em Kepler} space missions. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0911/0911.0612_arXiv.txt": {
        "abstract": "The cosmic microwave background (CMB) represents a unique source for the study of gravitational lensing. It is extended across the entire sky, partially polarized, located at the extreme distance of $z=1100$, and is thought to have the simple, underlying statistics of a Gaussian random field. Here we review the weak lensing of the CMB, highlighting the aspects which differentiate it from the weak lensing of other sources, such as galaxies. We discuss the statistics of the lensing deflection field which remaps the CMB, and the corresponding effect on the power spectra. We then focus on methods for reconstructing the lensing deflections, describing efficient quadratic maximum-likelihood estimators and delensing. We end by reviewing recent detections and observational prospects. ",
        "introduction": "\\label{sec:introduction} \\subsection{\\textit{What is the CMB?}} The cosmic microwave background (CMB) is a blackbody radiation which permeates the Universe. It was released approximately $400,000$ years after the Big Bang, when the primordial plasma had cooled enough that neutral hydrogen atoms could form, and the Universe became effectively transparent to radiation. To a very good degree of approximation, the CMB temperature is uniform across the sky, with a value today of $2.725\\,\\mathrm{K}$~\\cite{2002ApJ...581..817F}. At the level of one part in $10^{5}$, however, the CMB temperature has small anisotropies which are due to fluctuations of density and bulk velocity in the primordial plasma. The measurement of these anisotropies is clearly a difficult task, but a number of very successful experiments have driven rapid progress in this field over the past two decades. Modern temperature experiments such as ACT\\footnote{\\url{http://www.physics.princeton.edu/act/}}, SPT\\footnote{\\url{http://pole.uchicago.edu/}}, and Planck\\footnote{\\url{http://www.rssd.esa.int/index.php?project=planck}} are now pushing to robust arcminute-scale measurements, while polarization experiments like Spider\\footnote{\\url{http://www.astro.caltech.edu/~lgg/spider_front.htm}}, QUIET\\footnote{\\url{http://quiet.uchicago.edu/}}, BICEP2 and Planck\\footnote{\\url{http://www.rssd.esa.int/index.php?project=planck}} will soon obtain low-noise maps of the CMB polarization. Gravitational lensing effects should become readily apparent in both types of observation, requiring detailed modelling to interpret the data correctly, but also offering additional information about the Universe at intermediate redshifts.  \\subsection{\\textit{What is special about the CMB in the context of lensing?}} In the context of lensing, the CMB can be thought of as a single source ``plane'' with several unique features: \\begin{enumerate} \\item{} It is extended across the entire sky. \\item{} It probes the matter distribution at relatively higher redshifts ($z \\sim 2$) than can be studied with e.g. galaxy lensing: the last-scattering surface from which the CMB was released ($z\\sim 1100$) is the most distant light source we will ever be able to see. \\item{} It is partially linearly polarized, which provides two additional lensed observables. The local quadrupole of the radiation field at the last scattering surface generates linear polarization through the mechanism of Thomson scattering, at a level of $10\\%$ of the temperature anisotropies \\cite{Hu:1997hv}. The observable CMB is therefore characterized by three numbers as a function of position on the sky -- the total radiation intensity, and two Stokes parameters describing the amplitude and direction of the linear polarization. \\item{ To a good degree of approximation the source fluctuations may be treated as those of a Gaussian random field. The signal along any particular direction is essentially random, but the correlations between directions at different angular separations is encoded in the power spectrum, which completely characterizes the signal. A typical CMB temperature realization, for example, is plotted in Fig.~\\ref{fig:lensed_unlensed_map}}. \\end{enumerate}   \\begin{figure} \\begin{center} \\psfrag{tT}{$\\tilde{T}$} \\psfrag{tU}{$\\tilde{T} - T\\ (\\!\\times 5)$} \\psfrag{thtdg}{$\\theta\\ (\\rm deg.)$} \\psfrag{psi}{$\\psi$} \\includegraphics[width=\\figwidth]{lensed_unlensed_map.eps} \\caption{Left: lensed CMB realization. Middle: difference map between lensed and unlensed CMB. Right: realization of $\\psi$ for the lensing, and an overlay of its gradient, the deflection angle.} \\label{fig:lensed_unlensed_map} \\end{center} \\end{figure} \\subsection{\\textit{Why study CMB lensing?}} To someone interested in the primary CMB, the lensing of the CMB is important as a contaminant. As we will discuss in Sec.~\\ref{sec:lensing_effects_on_cmb_observables}, lensing has a small but non-negligible effect on the CMB power spectra, and must be accounted for in modern parameter analyses to obtain unbiased constraints. Perhaps more importantly, lensing effects generate a curl-like ($B$ mode) polarization pattern on the sky which acts as a limiting source of confusion for low-noise polarization experiments targeting the signal from primordial gravitational waves~\\cite{Knox:2002pe,Seljak:2003pn}. This confusion can be reduced with an accurate cleaning of the lensing-induced signal, which we will discuss in Sec.~\\ref{sec:lens_reconstruction}.  Apart from being a nuisance for traditional observables, the lensing of the CMB can act as an additional source of information. A typical analysis of the CMB assumes Gaussianity and statistical isotropy, in which case the power spectrum is the only quantity of interest. As we shall discuss, lensing can be thought of as introducing into the CMB small amounts of non-Gaussianity (when marginalized over realizations of the lenses) or statistical anisotropy (for a fixed distribution of lenses). This effectively \\textit{introduces} information into the CMB, contained in the higher-order statistics (for the non-Gaussian viewpoint) and the off-diagonal elements of its covariance matrix (for the anisotropy viewpoint). With only one CMB sky to observe and interpret, both these viewpoints are useful. The additional information from lensing probes the state of the Universe at intermediate redshifts ($z \\sim 2$). This can be used to break parameter degeneracies and place improved constraints on quantities that affect the geometry or density perturbations at late times, such as the dark energy equation of state and portion of the energy budget in massive neutrinos. An optimal analysis of lensing effects with the data from the Planck satellite, for example, will enable us to measure the sum of neutrino masses to $\\sim 0.1\\, {\\rm eV}$, while lens reconstruction with a next-generation polarization mission such as EPIC/CMBPol can constrain the sum to $0.05\\, {\\rm eV}$ or better~\\cite{Kaplinghat:2003bh,Lesgourgues:2005yv,dePutter:2009kn}. This is an interesting limit, close to the minimum value for the sum of the masses suggested by terrestrial oscillation measurements in the normal hierarchy. ",
        "conclusions": ""
    },
    "0911/0911.0239_arXiv.txt": {
        "abstract": "Porosity evolution of dust aggregates is crucial in understanding dust evolution in protoplanetary disks. In this study, we present useful tools to study the coagulation and porosity evolution of dust aggregates. First, we present a new numerical method for simulating dust coagulation and porosity evolution as an extension of the conventional Smoluchowski equation. This method follows the evolution of the mean porosity for each aggregate mass simultaneously with the evolution of the mass distribution function. This method reproduces the results of previous Monte Carlo simulations with much less computational expense. Second, we propose a new collision model for porous dust aggregates on the basis of our $N$-body experiments on aggregate collisions. As the first step, we focus on ``hit-and-stick'' collisions, which involve neither compression nor fragmentation of aggregates. We first obtain empirical data on porosity changes between the classical limits of ballistic cluster--cluster and particle--cluster aggregation. Using the data, we construct a recipe for the porosity change due to general hit-and-stick collisions as well as formulae for the aerodynamical and collisional cross sections. Our collision model is thus more realistic than a previous model of Ormel et al. based on the classical aggregation limits only. Simple coagulation simulations using the extended Smoluchowski method show that our collision model explains the fractal dimensions of porous aggregates observed in a full $N$-body simulation and a laboratory experiment. By contrast,  similar simulations using the collision model of Ormel et al. result in much less porous aggregates, meaning that this model underestimates the porosity increase upon unequal-sized collisions. Besides, we discover that aggregates at the high-mass end of the distribution can have a considerably small aerodynamical cross section per unit mass compared with aggregates of lower masses. This occurs when aggregates drift under uniform acceleration (e.g., gravity) and their collision is induced by the difference in their terminal velocities. We point out an important implication of this discovery for dust growth in protoplanetary disks. ",
        "introduction": "It is a widely accepted idea that the first step of planet formation in protoplanetary disks involves the collisional growth of submicron/micron dust grains into macroscopic aggregates. A standard scenario is that dust grains coagulate by mutual sticking, gradually settle toward the midplane of the disk, and form a dense dust layer (e.g., \\citealt*{WC93}). It is still an open issue whether subsequent planetesimal formation is achieved by the gravitational instability of the layer \\citep{S69,GW73,S98,YS02} or by the direct collisional growth of the aggregates \\citep{WC93,SV97,BDH08}. To address this issue, further understanding on earlier evolutionary stages is needed.  Most of the previous studies have simplified dust aggregates as compact, nonporous spheres. However, both numerical and laboratory experiments have revealed that aggregates are not at all compact, but have an open, fluffy structure (for a review, see \\citealt*{Blum04,DBCW07}). This is particularly true for aggregates formed at an early growth stage where the collisional velocity is so low that collisional compression is negligible. It has been observed in $N$-body \\citep{KPH99} as well as experimental \\citep{WB98,Blum+98,Blum+00} studies that the outcome is an ensemble of fractal aggregates with the fractal dimension of $D \\la 2$. This fractal growth lasts until the impact energy becomes high enough to cause collisional compaction \\citep{SWT08}. The porosity change due to compressive as well as destructive collisions are also studied by some authors \\citep[e.g.,][]{Wada+07,Wada+08,Wada+09,SWT08,PD09} using $N$-body simulations including monomer surface interactions  \\citep{DT97}.  There are several strong reasons why the porosity evolution of aggregates is important in studying dust evolution in protoplanetary disks. These include the following. \\begin{enumerate} \\setlength{\\itemsep}{-0pt} \\item The porosity affects the aerodynamical property of aggregates. In a gas disk, the motion of a small aggregate is controlled by the friction force from the ambient gas. The friction coefficient can vary by many orders of magnitude depending on whether the aggregate is compact or fluffy. For example, a fractal ($D\\la2$) aggregate made of one thousand grains receives a friction force an order of magnitude stronger than a compact aggregate of the same mass. This difference is significant when one considers the formation of sedimentary dust layer on the disk midplane. \\item The porosity even affects the turbulence of protoplanetary disks. Magnetorotational instability (MRI; \\citealt*{BW91}), which is recognized as the most likely mechanism driving disk turbulence, operates only in a region where the ionization degree is sufficiently high \\citep{Gammie96}. The ionization degree strongly depends on the surface area of dust aggregates, since the recombination of ionized gases efficiently takes place on dust surfaces. Again, a fractal ($D\\la2$) aggregate made of one thousand grains absorbs ionized gases an order magnitude more efficiently than a compact grain of the same mass. For this reason, compact dust growth tends to make the gas environment turbulent \\citep{Sano+00}, while fractal dust growth tends to keep the environment laminar \\citep{Okuzumi09}. This difference is critical to the evolution of dust itself, since the turbulence causes the diffusion and collisional fragmentation of aggregates \\citep{BDH08,J+08}. \\item The fractal nature of aggregates matters when we try to interpret observed emission spectra of protoplanetary disks. For example, the $10\\micron$ spectral feature characteristic of submicron silicate grains quickly vanishes in compact aggregation, but only slowly in fractal $(D\\sim 2)$ aggregation \\citep{Min+06}. This means that the conventional compact dust model generally underestimates the degree of dust growth in the observed objects. \\end{enumerate} Thus, the consistent modeling of the growth and porosity evolution of dust aggregates is needed to better understand and predict dust evolution in protoplanetary disks.  Theoretically, however, the consistent modeling is not an easy task. First of all, we must consider the evolution of aggregates in the two-dimensional (2D) parameter space of mass and porosity. However, such 2D calculation is generally much more elaborate than one-dimensional (1D; i.e., mass only) one. Furthermore, such calculation requires a reliable {\\it collision model} for porous aggregates. Here, a ``collision model'' refers to a set of (1) the definition of the ``porosity,'' or ``volume,'' of an aggregate, (2) a recipe that determines the porosity change due to collisions, and (3) formulae for collisional and aerodynamical cross sections that determine the collision rate of porous aggregate pairs. To obtain a reliable collision model, one needs a sufficient amount of empirical data on aggregate collisions.  Because of the above theoretical difficulty, there are only a few studies addressing the coagulation of porous dust aggregates in the literature. A 2D simulation on porous dust growth was first done by \\citet{Ossenkopf93} in the context of dust extinction in molecular clouds. He used a collision model based on the $N$-body simulations of the ballistic cluster--cluster and particle--cluster aggregation (BCCA and BPCA; \\citealt*{Meakin91}). Ormel et al. (2007; hereafter, \\citetalias{OST07}) have presented a Monte Carlo method for studying porosity evolution of dust aggregates in protoplanetary disks. In this method, porous aggregates are represented by the same number of computational particles with mass and porosity, and the collisions among the particles are successively calculated using random numbers. The change in porosity on each individual collision is determined from a recipe based on previous $N$-body experiments on aggregate collisions \\citep{Ossenkopf93,DT97}. The Monte Carlo approach has been further developed independently by \\citet{OS08} and \\citet{ZD08} to achieve a high dynamical range in the parameter space. The compression/fragmentation model of \\citetalias{OST07} has been recently revised by \\citet{OPDT09} on the basis of numerical experiments on low-mass aggregate collisions \\citep{PD09}.  In this study, we provide new theoretical tools useful for studying the collisional evolution of porous dust aggregates. First, we present a new numerical method to solve the growth and porosity evolution of aggregates. This method is an extension of the conventional Smoluchowski method, a method commonly used for 1D simulations of dust aggregates \\citep[e.g.,][]{THI05,DD05,BDH08}. The core of our new method is the approximation that the porosity distribution is narrow for each aggregate mass. With this approximation, we derive a closed set of moment equations of the 2D Smoluchowski equation for the mass distribution function and the averaged mass-porosity relation. These equations are formally similar to the conventional 1D Smoluchowski equation, and makes it possible to adopt well-established algorithms for 1D problems. We confirm that this approach indeed reproduces the results of the Monte Carlo simulations by \\citetalias{OST07} with much less computational effort. Thus, our method will be particularly useful for adding the porosity evolution to global simulations including radial drift \\citep[e.g.,][]{BDH08} or coupled simulations involving hydrodynamic calculation \\citep[e.g.,][]{J+08}. Second, we provide a new collision model based on our $N$-body experiments for various types of aggregate collisions. As the first step of the modeling, we focus on ``hit-and-stick'' collision, i.e., low-velocity collision involving neither compression nor fragmentation. In a typical protoplanetary disk, the hit-and-stick picture well represents the growth of dust aggregates up to several centimeters \\citep{SWT08}. In contrast to previous studies, we first use empirical data on the porosity evolution between the BCCA and BPCA limits. We will see that  the numerical simulations of porosity evolution using our collision model well agree with the results of previous full $N$-body and laboratory experiments. Extension of our model to compressive and destructive collisions will be made in the future work.  This paper is organized as follows. In Section 2, we describe our extended Smoluchowski method and its numerical implementation. In Section 3, we use this method to try to reproduce the results of the Monte Carlo simulations by \\citetalias{OST07}. In Section 4, we present a new porosity model as well as the results of our $N$-body simulations from which we construct the porosity increase recipe. In Section 5, we compare this collision model with that of \\citetalias{OST07}. Section 6 is devoted to the summary. ",
        "conclusions": "In this study, we have presented new numerical tools for studying the coagulation and porosity evolution of dust aggregates. Our findings are summarized as follows.  \\begin{enumerate} \\setlength{\\itemsep}{-0pt} \\item We have presented a new numerical method for simulating the coagulation and porosity evolution of dust aggregates as an extension of the conventional Smoluchowski method (Section 2). The new method treats the averaged volume of aggregates of the same mass as dynamical, and follows its evolution as well as the evolution of the mass distribution function consistently (Equations \\eqref{eq:coag20b} and \\eqref{eq:coag21b}). This method enables to treat the 2D (i.e., mass and porosity being dynamical) coagulation problems in a very efficient way. We have confirmed that our method well reproduces the results of previous full 2D Monte Carlo simulations with much fewer computational expense (Section~3).  \\item We have presented a new collision model based on our numerical experiments on aggregates collisions (Section~4). A collision model refers to a set of definitions and formulae for the properties of porous aggregates, including a recipe for the porosity change upon a collision. As the first step, we have focused on ``hit-and-stick'' collisions, i.e., collisions involving no restructuring or fragmentation. In the numerical experiments, we have first considered successive collisions between aggregates of a fixed mass ratio, which we have referred as QBCCA. This allows to fill the gap between the classical BCCA and BPCA. Using the results of the $N$-body experiments on BCCA, BPCA, and QBCCA, we have constructed a recipe for the porosity change due to a general hit-and-stick collision (Equation~\\eqref{eq:chi}) as well as formulae for the aerodynamical and collisional cross sections. In a forthcoming paper, we will include the effects of collisional compression and fragmentation into our collision model.  \\item To clarify the validity of our collision model and the difference from previous models, we have performed a couple of simple simulations on porous dust coagulation with the extended Smoluchowski method (Section 5). We considered two cases where the coagulation is driven by Brownian motion and differential drift, respectively. For both cases the simulations using our model result in fractal aggregates of $D\\approx2$, which is consistent with the findings of previous full $N$-body and laboratory experiments \\citep{KPH99,WB98}. By contrast, the simulations using the hit-and-stick model of \\citetalias{OST07} result in much less porous aggregates. This is because this model underestimates the porosity increase upon unequal-sized collisions.  \\item We have discovered that, when the coagulation is driven by differential drift, aggregates at the high-mass end of the mass distribution obey a flat density--mass relation, in marked contrast to aggregates of lower masses, which obey a fractal (i.e., power-law) density--mass relation (Section 5.2). This is explained by the difference in the dominant growth modes: typical aggregates grow by colliding with similar ones, while the most massive aggregates grow by colliding with much smaller ones. As a consequence of this tendency, the massive aggregates can have a drift velocity ($\\propto$ mass-to-area ratio) considerably higher than smaller ones. This finding may be crucially important in protoplanetary disks where the growth of slowly moving aggregates is suppressed due to the negative charging. \\end{enumerate}"
    },
    "0911/0911.4901_arXiv.txt": {
        "abstract": "We study the phase mixing and dissipation of a packet of standing shear Alfv\\'{e}n waves localized in a region with non-uniform Alfv\\'{e}n background velocity. We investigate the validity of the exponential damping law in time, $\\exp(-At^3)$, presented by Heyvaerts \\& Priest (1983) for different ranges of Lundquist, $S$, and Reynolds, $R$, numbers. Our numerical results shows that it is valid for $(R,S)\\geq 10^7$. ",
        "introduction": "\\label{intro}  Since the discovery of the hot solar corona by Edl\\'{e}n (1943), different theories of coronal heating have been put forward and debated. Heyvaerts \\& Priest (1983), hereafter HP83, were first to suggest that the phase-mixing of Alfv\\'{e}n waves in coronal plasmas could be a primary mechanism in coronal heating. They showed that the phase-mixing occurs due to inhomogeneity of the local Alfv\\'{e}n phase speed across the background magnetic field. HP83 analytically showed that in both the strong phase-mixing limit and the weak damping approximation, the amplitude of standing Alfv\\'{e}n waves decays with time as $\\exp(-At^3)$ where $A$ is a function of the coordinate corresponding to the inhomogeneous direction ($x$, in this paper). Since then, much analytical and numerical work has been done on the subject. Nocera et al. (1984) studied the phase mixing of propagating Alfv\\'{e}n waves in an inhomogeneous medium. They pointed out if transverse gradients are smeared out as soon as they are formed, this yields to weak phase mixing where damping laws differ from solutions of HP83. Parker (1991) investigated the effect of a density and/or temperature gradient in the direction of vibration of a transverse Alfv\\'{e}n wave. The result was a strong coupling of the waves on different lines of force, producing a coordinated mode that was not subject to simple phase mixing. Hood et al.(1997a) derived a self similar solution of the Alfv\\'{e}n wave phase-mixing equations for heating of coronal holes. They showed that the damping of the waves with height follows the scaling predicted by HP83 at low heights, before switching to an algebraic decay at large heights. Hood et al. (1997b) obtained a simple, self similar solution for the heating of coronal loops by phase mixing. They showed the HP83 model still does work well in a certain class of coronal loops and the phase mixing can supply heating at large Lundquist number at timescales shorter than or comparable with the radiative cooling timescale. Nakariakov et al. (1997) considered the nonlinear excitation of fast magnetosonic waves by phase mixing Alfv\\'{e}n waves in a cold plasma with a smooth inhomogeneity of density across a uniform magnetic field. They suggested this nonlinear process as a possible mechanism of indirect plasma heating by phase mixing through the excitation of fast waves. But Botha et al. (2000) showed that the nonlinear generation of fast modes by Alfv\\'{e}n waves has little effect on classical phase mixing. De Moortel et al. (1999) elaborated the effect of density stratification on phase-mixing. They remarked that when the inhomogeneity in the horizontal direction in the plasma is sufficiently large, so the phase mixing is strong, stratification is unimportant. In the other words due to the rapid phase mixing, energy can be dissipated before the effects of stratification build up. De Moortel et al. (2000) studied the combined effect of a gravitationally stratified density and a radially diverging background magnetic field on phase mixing of Alfv\\'{e}n waves. They found that: i) The efficiency of phase mixing depends strongly on the particular geometry of the configuration. ii) Depending on the value of the scale height the wave amplitudes can damp either slower or faster than in the uniform non-diverging model.  Hood et al. (2002) showed that the amplitude of single pulse and bipolar pulse traveling in the z direction, contrary to infinite wavetrain, have slower algebraic damping of the form $t^{-3/2}$ and $t^{-3}$, respectively, rather than exponential in time. Tsiklauri et al. (2003) cleared that the decay rate of the Alfv\\'{e}nic part of a compressible 3D MHD pulse is affected linearly by the degree of localization of the pulse in the homogeneous transverse direction, but the dynamic of Alfv\\'{e}n waves can still be obtained from the previous 2.5D models, e.g., Hood et al. (2002). Smith et al. (2007) found that in presence of the both density stratification and magnetic field divergence, the enhanced phase mixing mechanism can dissipate Alfv\\'{e}n waves at heights less than half that was predicted by the previous analytical solutions. They stated that if phase-mixing takes place in strongly divergent magnetic fields, it is not necessary to invoke anomalous viscosity in corona.  In the present work, we study the phase mixing of a packet of standing shear Alfv\\'{e}n waves in presence of the both viscous and resistive dissipations. To do this, we numerically solve the linearized MHD equations and obtain the damping time of the oscillations. Our aim is to test the validity of HP83's damping law for different ranges of the Reynolds and Lundquist numbers. This paper is organized as follows. In Sect. 2, we introduce the basic equations of motion and introduce the model. In Sect. 3, the numerical results are reported, while the conclusions are given in Sect. 4. ",
        "conclusions": "Phase mixing of a packet of standing Alfv\\'{e}nic pulses in fundamental mode is studied. Using a finite difference method, the linearized MHD equations for a zero-$\\beta$ plasma are solved, numerically. The damping times of oscillations in presence of the both viscous and resistive dissipations are calculated, numerically. They are in good agreement with the TRACE observations. The exponential damping law in time of HP83, $\\exp(-At^3)$, for the different ranges of the Reynolds and the Lundquist numbers are examined. Our numerical results shows that it is valid for $(R,S)\\geq 10^7$. \\\\ \\noindent{{\\bf \\"
    },
    "0911/0911.1465_arXiv.txt": {
        "abstract": "We discuss Preheating after an inflationary stage driven by the Standard Model (SM) Higgs field non-minimally coupled to gravity. We find that Preheating is driven by a complex process in which perturbative and non-perturbative effects occur simultaneously. The Higgs field, initially an oscillating coherent condensate, produces non-perturbatively $W$ and $Z$ gauge fields. These decay very rapidly into fermions, thus preventing gauge bosons to accumulate and, consequently, blocking the usual parametric resonance. The energy transferred into the fermionic species is, nevertheless, not enough to reheat the Universe, and resonant effects are eventually developed. Soon after resonance becomes effective, also backreaction from the gauge bosons into the Higgs condensate becomes relevant. We have determined the time evolution of the energy distribution among the remnant Higgs condensate and the non-thermal distribution of the SM fermions and gauge fields, until the moment in which backreaction becomes important. Beyond backreaction our approximations break down and numerical simulations and theoretical considerations beyond this work are required, in order to study the evolution of the system until thermalization. ",
        "introduction": "If the Universe went through an inflationary expansion in an early stage of its evolution, it must also have undergone a Reheating period afterwards, during which (almost) all the matter of the Universe was created. A simple consequence of Inflation is that the number density of any particle species exponentially dies away with the inflationary expansion. Therefore, at the end of inflation, the Universe has no particles at all. Only the homogeneous energy density responsible for inflation (usually the potential energy of a scalar field, the \\textit{inflaton}) has survived. So, what?  The hot Big Bang (hBB) theory describes the early Universe as an expanding space filled up with relativistic species in thermal equilibrium. Thus, we are left on one hand, with an empty universe at the end of inflation and, on the other hand, with a universe full of relativistic particles according to the hBB. So, how can those two periods be matched? Obviously, the energy density responsible for inflation - the potential energy of the inflaton - had to get converted somehow into particles. The precise epoch in which the primordial inflationary energy was converted into (almost) all the particles of the Universe, is known as Reheating.  How can particles be created out of the energy of one field (the inflaton)? The idea is that by coupling the inflaton to other fields of matter, then the potential energy of the inflaton gets converted into quanta of those fields. A specific scenario of Reheating will consist of a specific model of particle physics, specifying a particular form of the inflaton couplings to other matter fields. Depending on the model, there will be specific mechanisms of particle production taking place. For instance, the main mechanism of production might be \\textit{perturbative decays}~\\cite{Perturbative}, \\textit{parametric resonance}~\\cite{Preheating}, \\textit{tachyonic instabilities}~\\cite{Hybrid}, etc. Whatever the mechanism of particle production, the created particles interact among themselves and eventually reach a thermal equilibrium state with a common temperature. That is the \\textit{Reheating Temperature}, $T_{\\rm RH}$, which represents the energy scale at the end of Reheating and determines the first moment in which the description of the evolution of the Universe can finally be put within the hBB theory framework.  In practice, in order to analyze Reheating, we will study the dynamics of a system described by a potential \\begin{equation} V = V_{\\rm inf}(\\chi) + V_{\\rm int}(\\chi,\\phi_a,\\psi_b,A^c_\\mu,...)\\,, \\end{equation} where $V_{\\rm inf}$ is the inflationary potential and $V_{\\rm int}$ represents the interaction terms between the inflaton $\\chi$ and the rest of matter fields of the model: scalar fields ($\\lbrace\\phi_a\\rbrace$), fermions ($\\lbrace\\psi_b\\rbrace$), vector fields ($\\lbrace A^c_\\mu\\rbrace$), etc. In the last few years, simplified models based only on scalar fields or, at most, including (non-gauge) fermionic species, have been the main target of study. For example, a simple model in which the inflaton $\\chi$ is coupled to another scalar field $\\phi$ and to some fermionic field $\\psi$, can be described by \\begin{equation}\\label{eq:potentials} V = V_{\\rm inf} + V_{\\rm int} = V_{\\rm inf}(\\chi) + g^2\\chi^2\\phi^2 + y^2\\chi\\bar\\psi\\psi\\,, \\end{equation} where $g^2$ and $y^2$ are dimensionless couplings. Unfortunately, even dealing only with two or three fields, the aspects of Reheating are extremely complex~\\cite{Preheating},~\\cite{FermionicPreheating}. Consequently, most phenomena of particle creation have been discovered in simplified scenarios like the one described by~(\\ref{eq:potentials}), where the Standard Model (SM) particles or Dark Matter (DM) candidates are simply absent. Models incorporating a gauge principle in some sector of scalar fields (ignoring fermions), have also been considered. For instance, Hybrid Inflation models where the inflaton $\\chi$ is a singlet of the SM, and the symmetry breaking field $\\Phi$ coupled to the inflaton, is the SM Higgs doublet. Such models have been studied~\\cite{BellidoTonyMarga} through lattice numerical integration of the classical equations of motion obtained from the Lagrangian \\begin{equation} -\\mathcal{L} = \\mathcal{D}_\\mu\\Phi^\\dag\\mathcal{D}^\\mu\\Phi + \\lambda(\\Phi^\\dag\\Phi - v^2)^2 + g^2\\chi^2\\Phi^\\dag\\Phi + \\frac{1}{2}{\\rm Tr}\\lbrace F_{\\mn}F^{\\mn}\\rbrace\\,, \\end{equation} with $g,\\lambda$ dimensionless couplings, $F_{\\mn}$ the non-abelian field strength and $\\mathcal{D}_\\mu = \\partial_\\mu - ieA_\\mu$ the gauge derivative, coupling $\\Phi$ with the gauge fields $A_\\mu$ with strength $e$.  A realistic scenario of Reheating should account for (almost) all the matter of the Universe. Therefore, it should really incorporate a complete gauge theory with all kind of fermions and bosons, since our actual understanding of particle physics relies on the Standard Model (SM), based on the $SU(3)\\times SU(2)\\times U(1)$ group involving scalar, spinor and gauge fields. Of course, the SM is known to be incomplete and there are several extensions of which one can think about. But whatever new ingredients might be added, this does not change the fact that SM particles had to be produced during Reheating, since the posterior evolution of the Universe cannot be understood if those particles were not already present in an early epoch. Let's then discuss a potential way to realize Reheating within the SM. ",
        "conclusions": ""
    },
    "0911/0911.1186_arXiv.txt": {
        "abstract": "{Results from new \\hi synthesis observations of the nearby blue compact dwarf galaxy, NGC 2915, carried out on the Australian Telescope Compact Array are presented.  High resolution \\hi moment maps for the galaxy reveal complex \\hi distributions and kinematics.  The presence of large non-circular velocity components within the gas at inner radii is revealed.  The central gas dynamics are consistent with the simple scenario in which winds from high-mass stars are expelling the gas outwards from the centre of the galaxy.  A model in which the central region is treated as a rotating, expanding gas torus is able to reproduce the inner \\hi morphology and kinematics of NGC 2915.  We also find intriguing evidence for a gas infall event which, if confirmed, would be the first such evidence for a low-mass system.}   \\FullConference{Panoramic Radio Astronomy: Wide-field 1-2 GHz research on galaxy evolution - PRA2009\\\\ June 02 - 05 2009\\\\ Groningen, the Netherlands}  \\begin{document} ",
        "introduction": "NGC 2915 is classified as a nearby \\citep[3.78 Mpc,][]{karachentsev_catalog} blue compact dwarf galaxy.  Photometry carried out by \\citet{meurer1} revealed two dominant stellar populations: a compact blue stellar core which is the location of current high-mass star formation, surrounded by a diffuse, older stellar population.  \\citet{meurer1} measured the average star formation rate to be $\\sim$ 0.05 M$_{\\odot}$ yr$^{-1}$.  Interesting is the observation that the stellar component of the galaxy is embedded in a huge ($\\sim$ 22 \\textit{B}-band scale lengths), low surface brightness \\hi disk with well-defined spiral structure, but with apparently no significant star formation in the outer parts.  The absence of stars in the outer disk does, however, make NGC 2915 an ideal candidate for dark matter studies.  \\citet{meurer2} were the first to study the dark matter distribution of a blue compact dwarf galaxy by mass-modeling the observed rotation curve of NGC 2915.  They showed that NGC 2915 is dark-matter-dominated at nearly all radii with a very dense and compact dark matter core.  New \\hi synthesis observations of NGC 2915 using the Australian Telescope Compact Array have been carried out.  These data are used to study the gas dynamics of this galaxy in detail.  Previous investigators have focused on explaining the spiral structure of the outer disk of NGC 2915.  In this work we focus on the innermost \\hi morphology of the galaxy.  In particular, an effort is made to link the energy output of the central high-mass stellar population to the central \\hi morphology which, for the first time, is clearly resolved by our observations.  Details of the results will be presented by Elson et al. in prep. ",
        "conclusions": ""
    },
    "0911/0911.5716_arXiv.txt": {
        "abstract": "Over the past 15 years, our knowledge of the interior of the Sun has tremendously progressed by the use of helioseismic measurements. However, to go further in our understanding of the solar core, we need to measure gravity (g) modes. Thanks to the high quality of the Doppler-velocity signal measured by GOLF/SoHO, it has been possible to unveil the signature of the asymptotic properties of the solar g modes, thus obtaining a hint of the rotation rate in the core (Garc\\'\\i a \\textit{et al.} 2007, 2008a). However, the quest for the detection of individual g modes is not yet over. In this work, we apply the latest theoretical developments to guide our research using GOLF velocity time series. In contrary to what was thought till now, we are maybe starting to identify individual low-frequency g modes... ",
        "introduction": "Gravity modes are very sensitive to the structure (e.g. \\cite{Bas09,Gar08c}) and the dynamics  (e.g. \\cite{Mat08}) of the radiative zone and, in particular, to the inner core of the Sun. There have been many attempts to look for them without so far an undisputed detection of such modes (\\cite{App09}), although some interesting peaks and patterns have been detected with GOLF and VIRGO (e.g. \\cite{Stc04,Jim09}) with a high confidence level.  A 4500-day GOLF time series (\\cite{Gab95}) starting on April 11, 1996 and calibrated into velocity (\\cite{Gar05})  has been used to compute a single, full resolution power spectrum density (PSD) in spite of the different sensitivity to the visible solar disk between the blue- and the red-wing GOLF measurements (see for further details \\cite{Gar98,Ulr00}). To increase the signal-to-noise ratio, we have also smoothed the PSD  with a 41-nHz boxcar function as it is commonly done in asteroseismology when the signal is weak (e.g. \\cite{Mic08}). ",
        "conclusions": "Several of the highest peaks between 60 and 140 $\\mu$Hz are located around the theoretical frequencies of the dipole modes obtained by the Saclay seismic model (e.g. \\cite{Mat07}). Moreover, varying the splitting of the modes from 1 to 5 times the rotation rate of the radiative region, $\\Omega_{\\mathrm{rad}}$,  we notice that there is a quasi complete sequence of peaks matching the model when the splitting of these modes is around 4.5 $\\Omega_{\\mathrm{rad}}$ (see Fig.\\ref{fig1}). This could be the first time that individual g modes are identified in the Sun. For example, the candidate mode $\\ell$=1, n=$-4$ has an amplitude of 1.8 $\\pm$ 0.4 mm/s (with a signal-to-noise ratio of $\\sim$ 4) which is close to the latest theoretical predictions (\\cite{Bel09}).   \\begin{figure}[!htb] \\begin{center} \\includegraphics[angle=90,width=0.8\\textwidth]{fig1.eps} \\caption{GOLF PSD. The dotted vertical lines are the central frequencies of the dipole g modes computed by the Saclay seismic model. The vertical dashed and dot-dashed lines are the rotational split components with a rotational splitting 4.5 times larger than in the radiative zone above 0.2~$R_{\\odot}$.} \\label{fig1} \\end{center} \\end{figure}   The splittings can be estimated for each individual g-mode candidate. However, it is very difficult to obtain an accurate inference of the rotational profile of the core (\\cite{Gar08b,Mat09}). The rotation rate in the core could be up to $\\sim$7 $\\Omega_{\\mathrm{rad}}$, since $\\sim$60$\\%$ of the rotational kernel of these g modes is located in the inner core.%"
    },
    "0911/0911.0419_arXiv.txt": {
        "abstract": "{The traditional view that Ly$\\alpha$ emission and dust should be mutually exclusive has been questioned more and more often; most notably, the observations of Ly$\\alpha$ emission from ULIRGs seem to counter this view. In this paper we seek to address the reverse question. How large a fraction of Ly$\\alpha$ selected galaxies are ULIRGs? Using two samples of 24/25 Ly$\\alpha$ emitting galaxies at $z = 0.3/2.3$, we perform this test, including results at $z = 3.1$, and find that, whereas the ULIRG fraction at $z = 3.1$ is very small, it systematically increases towards lower redshifts. There is a hint that this evolution may be quite sudden and that it happens around a redshift of $z \\sim 2.5$. After measuring the infrared luminosities of the Ly$\\alpha$ emitters, we find that they are in the normal to ULIRG range in the lower redshift sample, while the higher redshift galaxies all have luminosities in the ULIRG category. The Ly$\\alpha$ escape fractions for these infrared bright galaxies are in the range $1-100$\\% in the $z = 0.3$ galaxies, but are very low in the $z = 2.3$ galaxies, 0.4\\% on average. The unobscured star formation rates are very high, ranging from 500 to more than 5000 M$_{\\odot}$~yr$^{-1}$, and the dust attenuation derived are in the range $0.0 < A_V < 3.5$. } ",
        "introduction": "Galaxies at high redshift come in all kinds of flavours. Depending on the search criteria, the sources found may be dusty or dust-free, more or less massive, star forming or quiescent. Some of the most impressive beasts of the high redshift zoo are the sub-mm and ultra-luminous infrared galaxies: SMGs (Ivison et al. 1998, 2000, Blain et al. 2002, Chapman et al. 2005) and ULIRGs (Sanders \\& Mirabel 1996). These galaxies show star formation rates in the hundreds or thousands of solar masses per year, with most of their light re-processed into the infrared or sub-mm wavelengths due to large amounts of dust. On the other end of the scale, Ly$\\alpha$ emitters are found (LAEs; M{\\o}ller \\& Warren 1998, Gronwall et al. 2007, Nilsson et al. 2007, Finkelstein et al. 2009, Ouchi et al. 2008, Grove et al. 2009) which have hitherto been generally considered blue and dust-free galaxies with moderate star formation rates of a few to a few times ten solar masses per year.  Only a few publications have reported finding red and/or dusty Ly$\\alpha$ emitters (Stiavelli et al.~2001, Colbert et al.~2006, Lai et al.~2007, Finkelstein et al.~2009, Nilsson et al.~2009). Considering the traditional view that even small amounts of dust would quench any Ly$\\alpha$ emission (for a discussion, see Pritchet 1994), the detection of the Ly$\\alpha$ line in dusty galaxies would seem surprising, but even more so the detection of Ly$\\alpha$ in a sample of sub-mm galaxies (Chapman et al.~2003, 2005). Likely explanations for the existence of dusty galaxies with Ly$\\alpha$ emission are either special geometrical alignments or strong local variations in the dust-to-gas ratios in a clumpy medium (Neufeld 1991, Hansen \\& Oh 2006). Observational evidence that Ly$\\alpha$ emission is commonly seen among sub-mm galaxies does not, however, constitute evidence for the opposite, since sub-mm galaxies are rare and because there are few ways other than Ly$\\alpha$ to determine the redshift of very red galaxies. It follows that the overlap reported so far could well be an observational selection bias. The fundamental questions that remain are therefore whether {\\it i)} sub-mm galaxies and ULIRGs are common among Ly$\\alpha$ emitters and {\\it ii)} whether the ULIRG fraction evolves with redshift. This \\emph{Letter} aims to address those two questions. We here present the infrared properties of two samples of galaxies found through their Ly$\\alpha$ emission at $z = 0.3$~and~$2.3$. We show that they have infrared fluxes well into the ULIRG regime. This sample is unique, bridging the gap between low-mass star forming and large star-bursting galaxies.  Throughout this paper, we assume a cosmology with $H_0=72$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega _{\\rm m}=0.3$ and $\\Omega _\\Lambda=0.7$. ",
        "conclusions": "We have here presented the infrared properties of two samples of Ly$\\alpha$ emitters, at $z = 0.3$ (24 objects) and $z = 2.3$ (25 objects). The samples were originally Ly$\\alpha$-selected but here we considered only the subsamples that were also detected with \\emph{Spitzer}. The two subsamples were selected to determine the fraction of powerful dusty starburst galaxies (ULIRGs) among LAEs and the redshift evolution of this fraction. Our most important conclusions are the following.  \\emph{i)} At all redshifts below $z = 3$, a non-zero fraction of LAEs are found to be ULIRGs. This result directly contradicts the classic view that dust and Ly$\\alpha$ emission are mutually exclusive, and it holds together with the finding that Ly$\\alpha$ emission at high redshifts correlate with metallicity (M{\\o}ller et al.~2004) makes the case that we must re-think the importance of dust for the Ly$\\alpha$ escape fraction.  \\emph{ii)} There is evidence for a strong evolution in ULIRG fraction from redshift 3.1 to the present universe. The fractions derived have been shown to be very robust against different selection criteria. This explains why there was severe disagreement about the colour of Ly$\\alpha$ galaxies originally. At redshifts around three they were reported to be young, blue, low-dust starbursts (Warren \\& M{\\o}ller 1996, Fynbo et al.~2000), while at redshifts closer to two they were reported to be red/dusty (Stiavelli et al.~2001).  \\emph{iii)} There may be evidence that the evolution of the ULIRG fraction is not linear in time, but rather that there is a jump from redshift three to two, and then a more gradual evolution to the present time. Even though the evidence is merely suggestive at present, this jump does coincide with the number density evolution of quasars and suggests a connection between those classes of objects. This question bears strongly on the connection between starbursts and the formation of quasars, so it should be investigated more thoroughly via dedicated surveys of LAEs at several different redshifts in the range $z = 2$ to $z = 3.5$. We provide a simple formalism in the form of a three parameter function which is well suited to future tests of how sharp the transition is.  From the accumulation of evidence, it is clear that some Ly$\\alpha$ emitting galaxies are very red and dusty, and vice versa that at least some ULIRG galaxies exhibit Ly$\\alpha$ emission. As was shown in e.g. Finkelstein et al.~(2009), it is possible that for some LAEs the Ly$\\alpha$ emission is not affected by the dust at all, or at least not along the line-of-sight towards us. The possibility that Ly$\\alpha$ equivalent widths may be enhanced in dusty environments has also been discussed in, e.g., Neufeld~(1991), Hansen \\& Oh~(2006), and Finkelstein et al.~(2008, 2009). Interestingly, one of the ULIRG LAEs found at high redshift in the Nilsson et al.~(2009) sample also has by far the largest equivalent width of their sample, with a restframe Ly$\\alpha$ equivalent width of nearly $800$~{\\AA}. Considering that samples of LAEs now have for several years consistently included red, dusty galaxies, it is advisable to stop considering Ly$\\alpha$ emitters as one type of galaxy, as we have seen in this and previous publications (Nilsson et al. 2009) that Ly$\\alpha$ emitters exhibit a range of properties and that this range increases with decreasing redshift."
    },
    "0911/0911.3901_arXiv.txt": {
        "abstract": "{X-shooter is the first second-generation instrument to become operative at the ESO Very Large Telescope (VLT). It is a broad-band medium-resolution spectrograph designed with gamma-ray burst (GRB) afterglow spectroscopy as one of its main science drivers.} {During the first commissioning night on sky with the instrument fully assembled, X-shooter observed the afterglow of GRB\\,090313 as a demonstration of the instrument's capabilities.} {GRB\\,090313 was observed almost two days after the burst onset, when the object had already faded to \\textit{R}$\\sim$21.6. Furthermore, the 90\\% illuminated Moon was just 30 degrees away from the field. In spite of the adverse conditions, we obtained a spectrum that, for the first time in GRB research, covers simultaneously the range from 5\\,700 to 23\\,000$~\\AA$.} {The spectrum shows multiple absorption features at a redshift of 3.3736, which we identify as the redshift of the GRB. These features are composed of 3 components with different ionisation levels and velocities. Some of the features have never been observed before in a GRB at such a high redshift. Furthermore, we detect two intervening systems at redshifts of 1.8005  and 1.9597.} {These results demonstrate the potential of X-shooter in the GRB field, as it was capable of observing a GRB down to a magnitude limit that would include 72\\% of long GRB afterglows 2 hours after the burst onset. Coupled with the rapid response mode available at VLT, allowing reaction times of just a few minutes, X-shooter constitutes an important leap forward on medium resolution spectroscopic studies of GRBs, their host galaxies and intervening systems, probing the early history of the Universe.} ",
        "introduction": "During the first hours after the onset of long gamma-ray bursts \\citep[GRBs, for a short-long classification see][]{kou93}, their optical/near infrared (nIR) counterparts shine as the brightest beacons in the Universe \\citep{kan07, rac08, blo09}. Being generally produced by the collapse of a massive star \\citep[][and references therein]{woo06}, they give us the opportunity to study the environment of massive star forming regions at cosmological distances \\citep[the average redshift of long GRBs observed by \\textit{Swift} is $\\sim$2.2;][]{fyn09} through the use of medium to high resolution spectroscopy. Time variability of absorption features corresponding to fine-structure and metastable level transitions has been sometimes identified \\citep{des06,vre07,del09}, which makes it possible to constrain the distance from the GRB to the absorbing gas, at typical scales of kpc. Only few such cases have been spotted up to now, due the required spectral resolution and sensitivity. Furthermore, being bright sources that disappear over time, they allow us to study intervening systems through their absorption signatures in the afterglow spectrum, and to search for the associated galaxies when the afterglow has faded away \\citep{jak04, che09}.  The family of short GRBs, although still poorly studied, presents on average fainter counterparts and has been suggested to be produced by the coalescence of compact objects \\citep[][and references therein]{nak07}, although some of them could come from other progenitors, such as extragalactic magnetars \\citep{hur05}. Up to now, there have been few attempts to acquire a spectrum of a short GRB, most of them having been unsuccessful due to late observations or faint afterglows. There is still no medium or high resolution spectrum of a short GRB.  Within this context, X-shooter \\citep{dod06} is presented to the GRB community. It is the first of the second-generation instruments at ESO's Very Large Telescope (VLT), at Paranal Observatory (Chile). It is a single target spectrograph capable of obtaining a medium resolution spectrum ($\\Re$=4\\,000 - 14\\,000, depending mainly on the slit width and wavelength) covering the complete range from 3\\,000 to 24\\,800 $\\AA$ in a single exposure. It has been designed to maximise efficiency by splitting the light with dichroics into three arms: ultraviolet/blue (UVB), visible (VIS) and near-infrared (NIR). Each arm has an echelle spectrograph with optimised optics, coatings, dispersive elements and detectors. Since March 2009, it is installed in the Cassegrain focus of \\textit{Kueyen}, the second 8.2m Unit Telescope of the VLT, where it began routine scientific observations in October 2009.  One of the key science drivers of X-shooter is the study of optical and nIR afterglows of GRBs as cosmological probes of the circumstellar, interstellar and intergalactic medium back to the epoch of first star formation in the Universe. During its first commissioning night on sky with its full setup, X-shooter was used to observe the afterglow of GRB\\,090313, which had exploded two days before, as a test target. In this article we present the results of this observation of a GRB with X-shooter, giving a hint of the full potential that the instrument will have during regular observations.  The paper is organised as follows: Sect. 2 describes our observations in the context of the GRB evolution. In Sect. 3 we present our results, focusing on the spectral features local to the GRB and on the intervening systems. Finally, in Sect. 4 we discuss the results and present our conclusions. ",
        "conclusions": "In this paper we present the first spectrum of a GRB obtained with X-shooter, covering for the first time a range from optical to NIR simultaneously. Although observed during an early commissioning phase and under unfavourable conditions, the results allow us to have an optimistic view of X-shooter's potential.  For the GRB, we determine a redshift of 3.3736, at which we detect absorption features with three velocity components. From the values in Table~\\ref{table:lines} and from Fig.~\\ref{FigGRB} it is inferred that component II has the highest column densities, and we used it to calculate the redshift. Component I is the least ionised (high ionisation lines are only weakly detected) while in component III we detect mainly highly ionised species. Although not conclusive, this could indicate that component III is the closest to the GRB, while component I is the furthest. The relative velocities between the 3 different components are due to relative motion of the gas within the host galaxy, in principle unrelated to the GRB itself. The large values of [Si/Fe] measured in the GRB host galaxy may be indicative of enhanced abundance of Si (produced by extensive star formation) or a higher depletion level of Fe (evidence of dust). In fact, \\citet{kan09} measured an extinction corresponding to $A_V$=0.34$\\pm$0.15 for this burst, which as they note, is atypically large for such a high redshift event. They also note that GRB\\,090313 is among the most intrinsically luminous optical afterglows detected. However, if we look at the EW$_{\\mathrm{rest}}$ measured for the absorption features in this spectrum (see Table~\\ref{table:ewlines}), we find that they are consistent with the typical values found in the sample of \\citet{fyn09}. We note that, to our knowledge, some features such as \\ion{Mg}{II}, \\ion{Mg}{I} and \\ion{Ca}{II} had not been observed before in GRB spectra at such high redshift.  The surveys of strong  \\ion{Mg}{II} (EW$_{\\mathrm{rest}}$ $>$1~\\AA) intervening absorbers along GRB lines of sight gave the surprising result of an excess of these systems compared to QSO lines of sight \\citep{pro06,ver09}. We have measured the EW$_{\\mathrm{rest}}$ of the \\ion{Mg}{II}$\\lambda$2796 line of both intervening systems, finding a combined value (coadd of the three components) of  1.35$\\pm$0.10 $\\AA$ for the $z$=1.80 and 0.50$\\pm$0.03 $\\AA$ for the $z$=1.96. The number density of strong \\ion{Mg}{II} systems for this spectrum is ${\\rm d}n/{\\rm d}z = 0.88$, in agreement with the excess of a factor of $\\sim 2$ compared to QSO lines of sight found by \\cite{ver09}. The same authors also considered the statistics of weak systems (0.3~\\AA$<$EW$_{\\mathrm{rest}}$ $<$1~\\AA) and conclude that their number density is in agreement with what is found along QSO lines of sight, as confirmed by \\citet{tej09}. The S/N of the GRB\\,090313 X-shooter spectrum is too low to cover a significant redshift path for a rest frame equivalent width limit of 0.3\\AA, on the other hand this limit will be normally reached for the guaranteed time programme for observations of GRB afterglows that will be performed with X-shooter. This program will collect about 100 spectra within 3 years, hugely increasing the redshift path of the GRB surveys and therefore the significance of the statistics, and bringing useful information to explain the unexpected excess of strong absorbers. Moreover, X-shooter will cover simultaneously a larger redshift path compared to the instruments presently used for QSO and GRB afterglow spectroscopic observations, up to $\\Delta z\\sim 5$ (the total redshift path for strong \\ion{Mg}{II} systems for the spectrum presented here is $\\Delta z=2.1$). Therefore X-shooter QSO and GRB surveys will be able to investigate systematically the presence and nature of \\ion{Mg}{II} systems up to a much higher redshift than currently done.  No significant emission lines have been detected at the redshift of the GRB or at any of the intervening systems. \\citet{ber09} pointed out the presence of two galaxies near the afterglow (G1 and G2, see Fig.~\\ref{FigAcq}). The brightest one (G1, $r$=15.6) has a known redshift of $z$=0.0235 and is located at 17\\farcs8 from the afterglow, equivalent to 8.3 kpc projected at the redshift of the galaxy. We do not see any features at this redshift and, in particular, nothing is detected at the wavelength of the most prominent absorption features that can be seen in the SDSS spectrum of G1 \\citep{aba09}, which are the \\ion{Ca}{II}$\\lambda$8500, 8544, 8664 and \\ion{Na}{I}$\\lambda$5891, 5897.  \\ion{Ca}{II} features are not detected with 3-$\\sigma$ limits of EW$_{\\mathrm{rest}}<$0.6 $\\AA$ and  \\ion{Na}{I} with limits of EW$_{\\mathrm{rest}}<$1.2 $\\AA$. The closest galaxy, G2 at 2\\farcs3, with $r$=21.6 does not have an identified redshift, so its relation with any of the observed absorption features will need further investigation.    To conclude, we point out some facts that show the potential that X-shooter has in the GRB field. Given a 10-$\\sigma$ limiting magnitude (per spectral bin) of \\textit{R}$\\sim$21.0 with 1 hr exposure, it will be able to study 33\\% (72\\%) of long bursts 12 (2) hours after the onset \\citep[using optical fluxes and decay indices from][]{nys09}. As an example, X-shooter would have been able to obtain a spectrum similar to or better than the one presented here of the $z\\sim8.2$ GRB\\,090423 \\citep{tan09,sal09} during the first 24 hours and to follow the nearby GRB\\,030329 for over a month. If we consider short GRBs, we expect to be able to study 3\\% (17\\%) within 12 (2) hours of the burst onset. Together with the rapid response mode available at VLT, that allows reaction times of just a few minutes, the possibilities increase. With its intermediate resolution, we will distinguish components with differential velocities of $\\sim$30 km s$^{-1}$. As shown here, this resolution will allow to derive, through line fitting, abundances of the different element species, although with some limitations as compared to higher resolution spectrographs when fitting blended components or marginally saturated lines. Thanks to the wide wavelength coverage it will be capable of observing afterglows up to redshifts of $\\sim$18, assuming that they exist (at redshift 18, Ly-$\\alpha$ would lie at $\\sim$23\\,000~$\\AA$). It will allow systematic studies of lines over wide redshift ranges (\\ion{Mg}{II}$\\lambda$2796, 2803 will be detectable from $z$=0.1 up to $z$=7.5) as well as measurements of ratios with widely separated lines. Furthermore, through the wide spectral coverage we will have a larger amount of lines to determine metallicities with better accuracy. It will also allow the study of dust extinction profiles in GRB environments. Thanks to the increased sensitivity compared to other instruments of similar resolution, we will be able to study feature variability with higher time resolution and up to later times. All this places X-shooter in a position to lead breakthrough advances in GRB research in the next years."
    },
    "0911/0911.4710_arXiv.txt": {
        "abstract": "Estimators of outer scales of atmospheric turbulence usually fit the phase screen snapshots derived from local wave front sensors to a Zernike basis, and then compare the spectrum of expansion coefficients in this basis with a narrowing associated with decreasing outer scales.  This manuscript discusses aspects that arise if the Zernike basis is exchanged for a Karhunen-Lo\\`eve basis of statistically independent modes. Data acquisition turns out to be more demanding because sensing the tip-tilt mode \\emph{is} required. The data reduction methodology is replaced by fitting of variance ratios of an entire (long exposure) data set. Statistical testing of hypotheses on outer scale models can be applied to a set of modes---supposing other noise originating from detector readout and the optical train can be disentangled. ",
        "introduction": "\\label{sec.intro}  \\subsection{Measuring and Interpreting Pupil Phases} A contemporary adaptive optics (AO) system of an astronomic telescope measures angles of arrival of the incoming wave front sampled across segmented areas that subdivide the full entrance pupil. The task is to weave the two principal parameters---motion of the image either represented by two Cartesian directions on the detector or by an angle of arrival and azimuth \\cite{GlindemannJOSAA11}---of the sensor array into a phase function covering the entire pupil. For the purpose of the present work we leave aside aspects of phase unwrapping \\cite{MearaJOSA67,LloydHartJOSAA11,FriedJOSA67}, finite element discretization \\cite{PearsonJOSA67}, or noise besides the atmospheric turbulence itself \\cite{BechetJOSAA26}.          As the information obtained represents the phase in the pupil, optical path differences integrated along paths through the entire atmosphere and altitudes dependencies are basically not available within a single AO system. Thicknesses and altitude variation of the structure constants may vary \\cite{CoulmanAO30}. Understanding of the origin of finite outer scales of turbulence remains poor for that reason, and the scattering of outer scale lengths at wavelengths in the visible is significant \\cite{AvilaJOSAA14,ConanSPIE2000,TakatoJOSAA12,CoulmanAO27}.           Optical path differences from stellar interferometers with baseline lengths of the order of the outer scales promise higher sensitivity. If the interferometric subsystems do not share any data streams with the AO systems of the individual telescopes, the time series of the fringe tracker may be the only information available, and then mediation by some wind velocity with its own set of unknowns is needed to fit the shape of temporal spectra \\cite{CliffordJOSA61,ColavitaAO26_4106,BuscherAO34} to outer scales. Results from the Palomar Testbed Interferometer merged with tip-tilt information of the telescopes show outer scales of some tens of meters---compatible with data from single-telescope apertures \\cite{LinfieldApJ554}.                  Auxiliary data products of this kind are not a standard of reduction pipelines of interferometric observations for now, so we discuss only the case of a single circular aperture. The remainder of Section \\ref{sec.intro} summarizes the literature on the von K\\'arm\\'an model of outer scales with a single layer with isotropic turbulence. Section \\ref{sec.2d} introduces its modal composition for the circular pupil. On the technical side, the matrix elements that arise in the Zernike basis are reduced to Hypergeometric Functions (Section \\ref{sec.p2d}), and central obscurations common to on-axis mirror trains are handled in a hybrid form that separates Fourier and real space with the aid of a Neuman series (Section \\ref{sec.obsc}). Benefits and problems compared to the use of Zernike covariances are shortly discussed in Section \\ref{sec.meth}.   \\subsection{The von K\\'arm\\'an Model of Outer Scales}  The refractive index structure function ${\\cal D}_n$ as a function of the 3D separation $\\Delta {\\bf r}$, and its Fourier representation $\\Phi_n$ at spatial frequencies $f$ are related via \\begin{equation} {\\cal D}_n(\\Delta {\\bf r}) = 2\\int \\Phi_n(f)[1-\\cos(2\\pi{\\bf f}\\cdot \\Delta {\\bf r})]d^3f. \\end{equation} The von K\\'arm\\'an model of the power spectrum in particular is \\begin{equation} \\Phi_n(f) = c_n C_n^2(f^2+f_0^2)^{-(\\gamma+3)/2} \\label{eq.vK} \\end{equation} where $C_n^2$ is the structure constant which characterizes the global strength of the turbulence, and where $f_0$ is a wavenumber equal to the inverse of the outer scale. $\\gamma$ is a power spectral index which for our analysis will be set to the Kolmogorov value of $2/3$; formulas will generally not be reduced with this value to support independent tests of this scaling \\cite{WaltersJOSA69,LewisJOSA67,BuserJOSA61,DaytonOL17,BoremanJOSAA13}. The constant \\begin{equation} c_n = -\\frac{\\Gamma(\\frac{3+\\gamma}{2})}{2\\pi^{3/2+\\gamma}\\Gamma(\\frac{-\\gamma}{2})} \\approx 0.00969315\\quad (\\gamma=2/3) \\label{eq.c3d} \\end{equation} is settled by demanding that at small $\\Delta {\\bf r}$ \\begin{equation} {\\cal D}_n(\\Delta r) = C_n^2 |\\Delta r|^\\gamma. \\end{equation}   This is a model of isotropic turbulence, so the information can be bound to the radial dependencies in real and Fourier space by integration over polar and azimuth angles, \\begin{equation} {\\cal D}_n(r)= 8\\pi\\int_0^\\infty f^2[1-j_0(2\\pi fr)]\\Phi_n(f)df = 8\\pi c_n C_n^2\\int_0^\\infty f^2[1-j_0(2\\pi fr)]\\frac{1}{(f^2+f_0^2)^{(3+\\gamma)/2}}df \\label{eq.Dnr} . \\end{equation} Performing the integration leads to a representation in terms of Modified Bessel Functions $K$ \\cite{VoitsekhoJOSAA12a,ConanJOSAA25,Conan00}, \\begin{equation} \\frac{1}{C_n^2}{\\cal D}_n(r) r^{-\\gamma} = -\\frac{(\\pi rf_0)^{-\\gamma}\\Gamma(\\gamma/2)}{\\Gamma(-\\gamma/2)} \\left[ 1 - \\frac{2}{\\Gamma(\\gamma/2)} (\\pi rf_0)^{\\gamma/2}K_{\\gamma/2}(2\\pi rf_0) \\right] . \\label{eq.Dn} \\end{equation} The limit of infinite outer scale recovers the Kolmogorov power law \\begin{equation} \\lim_{f_0\\to 0} \\frac{1}{C_n^2}{\\cal D}_n(r) r^{-\\gamma} = 1, \\end{equation} which is the value reached by the curve of $r=R$ in the left graph of Fig.\\ \\ref{fig.tst}.  \\begin{figure}[hbt] \\includegraphics[scale=0.5]{tstDr.eps} \\includegraphics[scale=0.5]{tstDr2.eps} \\caption{The refractive index structure function ${\\cal D}_n(r)/C_n^2$ of (\\ref{eq.Dn}) at $\\gamma=2/3$ as a function of outer scale with the radial distance $r$ as the parameter, or with the roles of these two dependencies swapped. ${\\cal D}_n$ depends essentially on the product $rf_0$, so scaling with some anonymous reference length $R$ leads to the universal curves shown. } \\label{fig.tst} \\end{figure}  \\subsection{Phase Structure Functions}    In most applications, the observable is not the three-dimensional distribution of the refractive indices but the phase difference $\\varphi$ between light rays at an optical wavenumber $\\bar k$ (which is $2\\pi$ divided by the wavelength). We write ${\\cal D}_\\varphi$ for the structure function of the phase if observed over a projected distance $P$ in an input pupil of the instrument after two rays have traveled along parallel lines through a layer of height $h$ (which includes the air mass factor) in which the refractive index structure function modeled in the previous section defines the statistics. This type of projection onto a two-dimensional plane introduces a weighting of the refractive index power spectrum $\\Phi_n$ with Bessel Functions $J_0$ \\cite{Goodman,ChesnokovOC141,ConanJOSAA25,EllerbroekAO36}, \\begin{equation} {\\cal D}_\\varphi(P) =4\\pi \\bar k^2 h\\int_0^\\infty [1-J_0(2\\pi fP)]\\Phi_n(f)f df, \\label{eq.Dphase} \\end{equation} for the von K\\'arm\\'an spectrum (\\ref{eq.vK}) condensed into the well-known \\begin{equation} {\\cal D}_\\varphi(P) = - \\bar k^2 h C_n^2 P^{1+\\gamma} \\frac{\\sqrt{\\pi}}{\\Gamma(-\\gamma/2)(\\pi Pf_0)^{1+\\gamma}} \\left[\\Gamma\\left(\\frac{1+\\gamma}{2}\\right) - 2(\\pi Pf_0)^{(1+\\gamma)/2}K_{(1+\\gamma)/2}(2\\pi Pf_0)\\right] . \\end{equation} The Kolmogorov limit of infinite outer scale is the familiar \\begin{equation} \\lim_{f_0\\to 0} {\\cal D}_\\varphi(P) = \\bar k^2 h C_n^2 P^{1+\\gamma} g(\\gamma) \\label{eq.dKlim} \\end{equation} where \\begin{equation} g(\\gamma)\\equiv \\sqrt{\\pi} \\frac{\\Gamma(-1/2-\\gamma/2)}{\\Gamma(-\\gamma/2)} \\approx 2.9143808, \\quad (\\gamma=2/3). \\label{eq.gofgamm} \\end{equation}  Collecting variances of optical path differences from this model---obtained by striking the factor $\\bar k^2$ of $\\cal D_\\varphi$--- we notice that the strength of the fluctuations remains smaller than predicted by the factor $g(\\gamma)$ in (\\ref{eq.gofgamm}), because this has been derived in the limit $h/P\\to \\infty$ of infinite layer height \\cite{Goodman}. For finite $h/P$, \\begin{eqnarray} {\\cal D}_\\varphi &=& 2\\bar k^2 \\int_0^h (h-\\Delta z)\\left[{\\cal D}_n(\\sqrt{(\\Delta z)^2+P^2})-{\\cal D}_n(\\Delta z)\\right]d(\\Delta z) \\\\ &=& 2\\bar k^2 C_n^2 P^{\\gamma+1} h \\int_0^{h/P} \\left(1-\\frac{t}{h/P}\\right) \\left[(t^2+1)^{\\gamma/2}-t^\\gamma \\right]dt , \\label{eq.gofP} \\end{eqnarray} where the integral can be rephrased as \\cite[(3.243.1),(3.251.1)]{GR} \\begin{equation} g(\\gamma)= 2\\int_0^{h/P} [1-t/(h/P)]\\left[(1+t^2)^{\\gamma/2}-t^\\gamma\\right]dt= \\frac{h}{P} \\,_3F_2\\left(\\begin{array}{c}-\\gamma/2,1/2,1 \\\\ 3/2,2\\end{array}\\mid -\\left(\\frac{h}{P}\\right)^2\\right) -\\frac{(h/P)^{1+\\gamma}}{(\\gamma+1)(1+\\gamma/2)} , \\label{eq.gambar} \\end{equation} with \\cite{GottschalkJPA21,RainvilleBAMS51} \\begin{equation} \\,_3F_2\\left(\\begin{array}{c}-\\gamma/2,1/2,1 \\\\ 3/2,2\\end{array}\\mid z\\right) = 2\\,_2F_1\\left(\\begin{array}{c}-\\gamma/2,1/2 \\\\ 3/2\\end{array}\\mid z\\right) - \\,_2F_1\\left(\\begin{array}{c}-\\gamma/2,1 \\\\ 2\\end{array}\\mid z\\right) . \\end{equation} This factor cannot reach $2.91$ in practise: the layer height $h$ multiplied by the air mass is limited to the order of 10 km representing the entire atmosphere, and $P$ is of the order of $2$ m for a single telescope up $100$ m for an interferometer. Then $h/P$ is in the range 5000 to 100, and Figure \\ref{fig.joInt} demonstrates that this finite layer thickness places $g(\\gamma)$ in the range 2.2 to 2.6\\@. Long baseline interferometers ``miss'' some power of ${\\cal D}_\\varphi$ through this geometric sampling effect. \\begin{figure}[hbt] \\includegraphics[scale=0.5]{joInt.eps} \\caption{The undersampling factor (\\ref{eq.gambar}) as a function of the ratio $h/P$. } \\label{fig.joInt} \\end{figure} ",
        "conclusions": "The text demonstrates computation of the statistically independent modes of phase screens in the von K\\'arm\\'an class of outer scales sampled over spherical receiver pupils, optionally taking into account masking of a circular region in the center. The main value in this work resides in techniques to compute integrals in the eigenvalue problem.  These modes can be employed to compute outer scales from measured phase functions. The goodness-of-fit of the von K\\'arm\\'an modal spectrum can be addressed. The number of degrees of freedom depends on how many of the modes remain free from other genuine sources of noise besides atmospheric turbulence.  \\appendix"
    },
    "0911/0911.3240_arXiv.txt": {
        "abstract": "An increasing number of large molecules have been positively identified in space. Many of these molecules are of biological interest and thus provide insight into prebiotic organic chemistry in the protoplanetary nebula. Among these molecules, acetic acid is of particular importance due to its structural proximity to glycine, the simplest amino acid. We compute electronic and vibrational properties of acetic acid and its isomers, methyl formate and glycolaldehyde, using density functional theory. From computed photo-absorption cross-sections, we obtain the corresponding photo-absorption rates for solar radiation at 1 AU and find them in good agreement with previous estimates. We also discuss glycolaldehyde diffuse emission in Sgr~B2(N), as opposite to emissions from methyl formate and acetic acid that appear to be concentrate in the compact region Sgr B2(N-LMH). ",
        "introduction": "The richness of interstellar chemistry has been growing steadily with the variety of objects and regions observed. The excitation and abundances of molecules contain key information on the physical structure and evolution of the host regions. Through molecules, we can trace the cycle of matter for interstellar space into stars and planets, and back again into the interstellar space \\citep{HW98}.  There is an increasing body of evidence for the existence of large molecules in the interstellar medium and in the interplanetary space. Interferometric observations of high-mass star forming regions in molecular clouds have revealed hot molecular cores, short-lived remnants of clouds not incorporated into the newly born massive stars. These hot cores contain within them the evaporated material of ices deposited on dust grain surfaces during the collapse, and observations show a very interesting chemistry. In particular, there exist substantial column densities of large partly hydrogen-saturated molecules, many of them being of pre-biotic interest. Some of these species have been also observed in comets and embedded in minor bodies of the solar system \\citep{C05}. The study of biologically interesting large interstellar molecules offers the exciting opportunity of learning more about the chemical evolution preceding the onset of life on the early Earth 4.5 billion years ago. Comets may be important carriers of prebiotic chemistry, and relevant agents in the delivery of complex organics to early Earth, as well as to newly formed planets.  The Sgr B2 molecular cloud complex is the prime target in the search for complex species, in particular in a hot core, the so-called Large Molecule Heimat, Sgr B2(N-LMH), within the more extended molecular cloud. In this compact source, smaller than the Oort cloud ($\\sim0.08$~pc) with a mass of several thousands M$_\\odot$ \\citep{MS97}, an extraordinary number of complex organics have been observed to exhibit very high column densities (e.g. amino acetonitrile, \\citealt{B08}). Large partly hydrogen-saturated species challenge the completeness of the standard ion-neutral scheme in interstellar chemistry, suggesting that reactions on dust grains are involved in their formation (e.g. \\citealt{BK07}).  Of the chemical species detected so far, particular attention has been paid to the formation of different isomer groups. In this work, we focus on C$_2$H$_4$O$_2$, i.e. acetic acid (CH$_3$COOH), glycolaldehyde (HCOCH$_{2}$OH), and methyl formate (HCOOCH$_3$), because of their potential role in the origin of life \\citep[e.g.,][]{W00,C05}. Glycolaldehyde, the simplest of the monosaccharide sugars, has first been detected by \\citet{H00} towards Sgr B2(N-LMH), and its most recently determined column density in that source is $5.9 \\times 10^{13}$~cm$^{-2}$ \\citep{H06}. It has been recently observed outside the galactic center by \\citet{B09} towards the hot molecular core G31.41+0.31, with the emission coming from the hot and dense region closest to the protostars. In addition, \\citet{C04} and \\citet{D05} presented an upper limit for glycolaldehyde abundance in the comet Hale-Bopp. Acetic acid shares the C$-$C$-$O backbone with glycine, from which it differs by an amino group (NH$_2$). First detected by \\cite{M97} in Sgr B2(N-LMH), acetic acid shows a column density of $6.1 \\times 10^{15}$~cm$^{-2}$ in that region \\citep{R02}. Finally, methyl formate, discovered towards Sgr B2(N) by \\citet{B75}, has been also observed in other star forming regions of both high \\citep{Mc96,G00} and low \\citep{RH06} mass, towards a proto-planetary nebula \\citep{R05}, and in comets \\citep{BM00,D05,R06}. The column density of methyl formate in  Sgr B2(N-LMH) is  $1.1 \\times 10^{17}$~cm$^{-2}$ \\citep{L01}.  Sgr B2(N-LMH) is the only source where all three of these isomers have been observed. However, while acetic acid and methyl formate are concentrate in Sgr B2(N-LMH), glycolaldehyde appears to be more diffusely spread through Sgr B2(N) \\citep{H01}. This behaviour is generally shared by other aldehydes \\citep{S06}. In comets only methyl formate has been observed so far.  An understanding of molecular structure, spectroscopy, and photo-absorption processes may be of critical importance in interpreting current observations. In particular, the lifetime of cometary molecules versus photo-destruction is a basic parameter for all cometary studies: as a matter of fact any error in the photodissociation rate translates linearly into an error on the abundance derived in the cometary nucleus. It is also needed for chemical modeling of planetary atmospheres. In this work, we derive electronic and vibrational properties of the isomer triplet C$_2$H$_4$O$_2$ described above using the Density Functional Theory (DFT). Photo-destruction rates for acetic acid and methyl formate were derived by \\citet{C94} using old laboratory absorption data published by \\citet{S88}, while the rate for glycolaldehyde is just an estimate. In section~\\ref{theory} we present a brief outline of the method, together with a description of computational settings and results. Section~\\ref{comets} contains the application of these results to cometary photochemistry, while discussion and conclusions are in the last section. ",
        "conclusions": "In this work we studied the photo\\textendash physics of the acetic acid and its isomers, glycolaldehyde and methyl formate. Computed photo\\textendash absorption rates are consistent with literature data for acetic acid and methyl formate. In the case of glycolaldehyde, for which the photo-absorption was completely missing up to the far\\textendash UV, our calculations indicate that photodestruction is slower than in the other two members of the isomer triplet. However, the overall photo\\textendash absorption rate is much larger, and likely to produce either resonant scattering or IR emission powered by near UV solar photons.  Sgr B2 observations of glycolaldehyde show that, unlike acetic acid and methyl formate, its emission is extended in the surrounding molecular cloud. Such a behaviour, that has not been understood so far (e.g. \\citealt{chenga}), is part of a more general problem involving differentiation in isomers, such as e.g., isocyanide isomers (CH$_3$CN, CH$_3$NC). The behaviour of this isomer triplet has also recently been discussed by \\citet{lat09} as a notable exception to what was called the ``Minimum Energy Principle'' (MEP) whereby whenever several isomers are possible for a given formula, their observed abundances are in order of binding energy. The authors made the educated guess that this anomaly is first created by differences in the chemical pathways leading to the formation of glycolhaldehyde, acetic acid and methyl formate on the surfaces of dust grains, and then preserved due to large energy barriers for the conversion among them. Our results may help to shed some more light on this problem. Although both hot core and the embedding molecular cloud are dark regions, cosmic\\textendash rays provide a source of UV photons at high visual extinctions by exciting molecular hydrogen in the Lyman and Werner bands \\citep{taraf}. This locally generated photon flux typically has fluences lower than $10.000$ photons cm$^{-2}$ s$^{-1}$ \\citep{ccp92}, and may produce important chemical effects \\citep{gredel,BK07}. However, the similarity in both the photo\\textendash absorption cross\\textendash sections and ionisation potentials of the three species makes chemical differentiation due to selective photo\\textendash destruction unlikely. As a possible explanation of the extended spatial scale of glycolhaldehyde, we consider the possibility of a slow, selective isomerisation, i.e. the possibility that a species may convert itself into another member of the C$_2$H$_4$O$_2$ triplet by interacting with the radiation field. This would imply an isomerisation mechanism which operates on a timescale which is much longer than the typical lifetime of a hot core \\citep[few times $10^4$~years,][]{wilner} but still short enough to be effective on the timescale of the lifetime of a molecular cloud \\citep[$\\sim10^8$~years,][]{blitz00}. In this framework, hot cores would reflect the relative abundances among the isomers as created by their production mechanisms, whereas in the molecular cloud the isomers would ``relax'' to the most stable one, namely glycolaldehyde, in agreement with the MEP.  Isomerisation may be induced by the absorption of radiation essentially in two main ways; following the absorption of a UV\\textendash visible photon the molecule could \\begin{enumerate} \\item[\\it (i)] move to an electronic state whose energy surface presents a minimum close to the equilibrium configuration of another isomer; \\item[\\it (ii)] convert, via one or more non\\textendash radiative transitions, a substantial part of the electronic excitation energy into  vibrational energy allowing the overcoming of the isomerisation potential barriers. \\end{enumerate} The present data do not allow a discrimination between the two cases. Moreover, the analysis of the first isomerisation channel would require a detailed study of the energy hypersurfaces in the excited states accessible with photons generated by the \\citet{taraf} mechanism. Therefore, we will discuss qualitatively the second isomerisation process. We assume that every absorption heats up the molecule in a time scale characteristic of electronic transitions ($\\sim 10^{-8}$~s), that then decades via non\\textendash radiative transitions ($\\lesssim 10^{-10}$~s). The electronic excitation energy can be uniquely released by a cascade of vibrational transitions. We also assume, as simplifying hypotheses, that (\\textit{i}) all the excitation energy is converted in vibrational excitation, (\\textit{ii}) that there exist only one isomerisation barrier, (\\textit{iii}) that the excitation energy is far greater than this barrier and (\\textit{iv}) that the isomerisation rate is high above the threshold and zero below it. Then every time a molecule absorbs an UV\\textendash visible photon, it will be vibrationally heated. If the energy of the absorbed photon is above the isomerisation threshold, the molecule will establish a statistical equilibrium, in which the probability to find it in one of the isomeric configurations will be proportional to the density of vibrational states at the given energy. The molecule will then cool down, in timescales of the order of a second, with a vibrational cascade (all vibrational modes are IR\\textendash active for all three isomers). Its energy will thus eventually fall below the barrier for the isomerisation. When this happens, the proportion among the isomers is frozen because the conversion rate among the species drops to zero, the abundances are therefore those given by the ratios among the densities of vibrational states. Since we are here dealing with three conformations, one should consider at least three different isomerisation channels, each with its different barrier(s). Whatever they are, however, as long as they are easily overcome with the energy of a single UV photon, the ratio of the abundances of the isomers, if they are exposed to UV light, should be fixed at the ratios of the densities of states at the threshold(s). \\begin{figure} \\includegraphics[width=\\hsize]{fig5.eps} \\caption{Densities of the vibrational states for the three species as a function of energy.\\label{dos}} \\end{figure}  \\begin{table} \\caption{Total energies for the three species considered. The second column shows the energy difference $\\delta$ in comparison with glycolaldehyde.} \\label{tab:rate} \\begin{center} \\begin{tabular}{ccc} \\hline \\hline & Total energy (eV) & $\\delta$ (eV)\\\\ \\hline Glycolaldehyde & -6231.83 & /\\\\ Metyl formate & -6232.32 &  0.49\\\\ Acetic acid & -6233.04 & 1.21\\\\ \\hline \\hline \\end{tabular} \\end{center} \\end{table}  The densities of vibrational states for the three species, calculated in harmonic approximation through the algorithm of \\cite{stein}, is shown in Fig.~\\ref{dos}. The distances among the curves are due to the differences in the total energies shown in Table~\\ref{tab:rate} and the differences in the frequencies of the vibrational modes.  The naive model drawn above would essentially turn all isomers into glycolaldehyde as soon as they absorb an UV photon. However, UV absorption, in the weak radiation field produced by to cosmic\\textendash ray induced H$_2$ fluorescence, occurs on timescales of the order of $\\ga 50.000$ years, that are comparable to the lifetime of hot cores. As consequence, MEP should not operate effectively in the hot core phase.  Therefore, one of the simplistic assumption in the naive model must be incorrect, slowing down isomerisation. Either internal conversion has a low quantum yield for these molecules, meaning that they get vibrationally heated only a small fraction of the times they absorb an UV photon; or isomerisation is \\emph{not} fast whenever vibrational excitation is sufficient to overcome reaction barriers, due to the morphology of the molecular potential energy surface. The latter can easily be the case if the unimolecular reaction path(s) to isomerisation are narrow and complicated in terms of phase space of the ions.  Of course, conversion among isomers needs not to proceed \\emph{only} via unimolecular reactions induced by radiation: any conversion reaction including chemical reactions with abundant enough partners and without activation barriers could equally well operate on the right timescales to fulfil the MEP in molecular clouds but not in hot cores. A thorough study of the potential energy surface of these three isomers, including the reaction paths connecting them, is called for in order to further progress in the understanding of this observational puzzle."
    },
    "0911/0911.3853_arXiv.txt": {
        "abstract": "{ Spatial and temporal variations in the electron-to-proton mass ratio, $\\mu$, and in the fine-structure constant, $\\alpha$, are predicted in non-Standard models aimed to explain the nature of dark energy. Among them the so-called chameleon-like scalar field models predict strong dependence of masses and coupling constants on the local matter density. To explore such models we estimated the parameters $\\Delta \\mu/\\mu \\equiv (\\mu_{\\rm obs} - \\mu_{\\rm lab})/\\mu_{\\rm lab}$ and $\\Delta \\alpha/\\alpha \\equiv (\\alpha_{\\rm obs} - \\alpha_{\\rm lab})/\\alpha_{\\rm lab}$ in two essentially different environments,~-- terrestrial (high density) and interstellar (low density),~-- from radio astronomical observations of cold prestellar molecular cores in the disk of the Milky Way. We found that $\\Delta\\mu/\\mu = (22\\pm4_{\\rm stat}\\pm3_{\\rm sys})\\times10^{-9}$, and $|\\Delta \\alpha/\\alpha| < 1.1\\times10^{-7}$. If only a conservative upper limit is considered, then $|\\Delta\\mu/\\mu| \\leq 3\\times10^{-8}$. We also reviewed and re-analyzed the available data on the cosmological variation of $\\alpha$ obtained from Fe\\,{\\sc i} and Fe\\,{\\sc ii} systems in optical spectra of quasars. We show that statistically significant evidence for the changing $\\alpha$ at the level of $10^{-6}$ has not been provided so far. The most stringent constraint on $|\\Delta \\alpha/\\alpha| < 2\\times10^{-6}$ was found from the Fe\\,{\\sc ii} system at $z = 1.15$ towards the bright quasar HE 0515--4414. The limit of $2\\times10^{-6}$ corresponds to the utmost accuracy which can be reached with available to date optical facilities. ",
        "introduction": "Spatial and temporal variations in the electron-to-proton mass ratio, $\\mu \\equiv m_{\\rm e}/m_{\\rm p}$, and in the fine-structure constant, $\\alpha \\equiv e^2/(\\hbar c)$, are not present in the Standard Model of particle physics but they arise quite naturally in grant unification theories, multidimensional theories and in general when a coupling of light scalar fields to baryonic matter is considered \\citep{uz03,mar08,ch09}. The light scalar fields are usually attributed to a negative pressure substance permeating the entire visible Universe and known as dark energy \\citep{cds98}. This substance is thought to be responsible for a cosmic acceleration at low redshifts, $z < 1$ \\citep{pr03}. However, scalar fields cause a problem since they could violate the equivalence principle which has never been detected in local tests \\citep{tu04}. A plausible solution of this dissension was suggested with a so-called `chameleon' model which assumes that a light scalar field acquires an effective potential and an effective mass due to its coupling to matter that strongly depends on the ambient matter density \\citep{kw04, ms07, op08}. In such a way the scalar field may evade local tests of the equivalence principle since the range of the scalar-mediated force is too short ($\\lambda_{\\rm eff} \\la 1$ mm) to be revealed at the terrestrial matter densities. This is not the case, however, for the space based tests where the matter density is considerably lower, an effective mass of the scalar field is negligible, and an effective range for the scalar-mediated force is very large ($\\lambda_{\\rm eff} \\ga 1$ pc).  Calculations of atomic and molecular spectra show that different transitions have different sensitivities to changes in fundamental constants \\citep{vl93,dz99,ko08}. Thus, measuring relative radial velocities, $\\Delta V$,  between such transitions one can probe the hypothetical variation of physical constants. For instance, a spatial dependence of $\\mu$ and $\\alpha$ on the ambient matter density can be tested locally using terrestrial measurements of molecular transitions and comparing them with radio astronomical observations of molecular clouds in the disk of the Milky Way (here we have a typical difference between the environmental gas densities of $\\sim$19 orders of magnitude). Complementary to these measurements, a temporal dependence of $\\mu$ and $\\alpha$ on cosmic time can be tested through observations of high-redshift intervening absorbers seen in quasar spectra.  In this presentation we summarize our recent results on spatial and temporal variations of $\\mu$ and $\\alpha$ obtained at high redshifts $z > 0$ (QSO absorption systems) and at $z=0$ (the Milky Way disk). ",
        "conclusions": "Our current results of astrophysical tests on spatial and temporal variations of fundamental constants can be summarized as follows. \\begin{enumerate} \\item[1.] The relative radial velocities of the rotational transitions in HC$_3$N $(J=2-1)$ and the inversion transition in NH$_3$ $(J,K)=(1,1)$ measured with the 100-m Effelsberg radio telescope reveal a velocity offset of $\\Delta V = 23\\pm4_{\\rm stat}\\pm3_{\\rm sys}$ \\ms. This offset is regularly reproduced in observations of different cold molecular cores with different facilities at the 32-m Medicina and 45-m Nobeyama radio telescopes. \\item[2.] Being interpreted in terms of the electron-to-proton mass ratio variation, the found velocity offset corresponds to \\dmm\\ = $(22\\pm4_{\\rm stat}\\pm3_{\\rm sys})\\times10^{-9}$. To cope with negative results obtained in laboratory experiments, a chameleon-like scalar field is required to explain the non-zero $\\Delta\\mu$ value. \\item[3.] A new approach to probe $\\alpha$-variations using the fine-structure transitions in atomic carbon [C\\,{\\sc i}] and low-$J$ rotational transitions in $^{13}$CO is suggested, and a strong constraint on the dependence of $\\alpha$ on the ambient matter density is deduced at $z=0$: $|\\Delta\\alpha/\\alpha| < 1.1\\times10^{-7}$. \\item[4.] Another method to measure the cosmological variability of $\\alpha$ from resonance transitions of neutral iron Fe\\,{\\sc i} is discussed and applied to the Fe\\,{\\sc i} system at \\zabs\\ = 0.45. The resulting constraint is \\daa\\ = $(7\\pm7)\\times10^{-6}$. \\item[5.] The Fe\\,{\\sc ii} system at \\zabs\\ = 1.84 is re-analyzed and a corrected value of \\daa\\ is obtained: \\daa\\ = $(4.0\\pm2.8)\\times10^{-6}$. \\item[6.] For the Fe\\,{\\sc ii} system at \\zabs\\ = 1.15 the overlooked correlations between individual \\daa\\ values are now taken into account that gives \\daa\\ = $(-0.12\\pm1.79)\\times10^{-6}$. \\item[7.] It is shown that spectral radio observations in the Milky Way are at present the most sensitive tool to probe the values of fundamental constants at different ambient physical conditions. The internal (statistical) error of our Effelsberg data is $\\sigma_{\\Delta\\mu/\\mu} = 4\\times10^{-9}$ which is about 1000 times lower as compared with the error in optical observations of quasars. \\item[8.] Since extragalactic molecular clouds have gas densities similar to those in the interstellar medium, the value of \\dmm\\ in high-$z$  molecular systems is expected to be at the same level as in the Milky Way, i.e. \\dmm$\\sim$10$^{-8}$~--- providing no temporal dependence of the electron-to-proton mass ratio is present. \\end{enumerate}  To conclude, we note that in order to be completely confident that the revealed velocity shift between rotational and inversion lines is not due to kinematic effects in the molecular clouds but reflects a density-modulated variation of \\dmm, new high precision radio-astronomical observations are needed for a wider range of objects. Besides, such observations should include essentially different molecules with tunneling transitions sensitive to the changes in $\\mu$ and molecules with $\\Lambda$-doublet lines which also exhibit enhanced sensitivity to variations of $\\mu$ and $\\alpha$ (Kozlov 2009). It is also very important to measure the rest frequencies of molecular transitions with an accuracy of about 1 \\ms\\ in laboratory experiments~--- in some cases the present uncertainties in the rest frequencies are larger than the errors of radio-astronomical measurements. In addition, search for variations of the fine-structure constant $\\alpha$ in the Milky Way disk using mid- and far-infrared fine-structure transitions in atoms and ions, or search for variations of the combination of $\\alpha^2/\\mu$ using the [{C}\\,{\\sc ii}] and/or [{C}\\,{\\sc i}] fine-structure transitions and low-lying rotational lines of CO would be of great importance for cross-checking the results."
    },
    "0911/0911.3124_arXiv.txt": {
        "abstract": "In this paper, several issues are considered, related to the construction of the next ICRF generation, ICRF-2. Between them, the following points are touched: ICRF-2 structure, ICRF Core sources selection, and some expected user's requirements. ",
        "introduction": "By now, more than 6 million geodetic and astrometric VLBI observations were made, much more than were available for the construction of ICRF, and methods of data analysis were substantially advanced. This opens up new opportunities for ICRF improvement. For this purpose, two ad hoc Working Groups were organized by IAU, and jointly IERS and IVS\\footnote{http://rorf.usno.navy.mil/ICRF2/}, and preparation to ICRF-2 creation was started \\cite{Ma08}. In particular, one of the main tasks of the preparation campaign is to investigate possible strategies that are or may be used for ICRF-2, the second ICRF realization. In this paper several issues related to construction of ICRF-2 are touched. ",
        "conclusions": ""
    },
    "0911/0911.4532_arXiv.txt": {
        "abstract": "We present a simplified model to study the orbital evolution of a young hot Jupiter inside the magnetospheric cavity of a proto-planetary disk. The model takes into account the disk locking of stellar spin as well as the tidal and magnetic interactions between the star and the planet. We focus on the orbital evolution starting from the orbit in the 2:1 resonance with the inner edge of the disk, followed by the inward and then outward orbital migration driven by the tidal and magnetic torques as well as the Roche-lobe overflow of the tidally inflated planet. The goal in this paper is to study how the orbital evolution inside the magnetospheric cavity depends on the cavity size, planet mass, and orbital eccentricity. In the present work, we only target the mass range from 0.7 to 2 Jupiter masses. In the case of the large cavity corresponding to the rotational period $\\approx 7$ days, the planet of mass $>1$ Jupiter mass with moderate initial eccentricities ($\\gtrsim 0.3$) can move to the region $< 0.03$ AU from its central star in $10^7$ years, while the planet of mass $<1$ Jupiter mass cannot. We estimate the critical eccentricity beyond which the planet of a given mass will overflow its Roche radius and finally lose all of its gas onto the star due to runaway mass loss. In the case of the small cavity corresponding to the rotational period $\\approx 3$ days, all of the simulated planets lose all of their gas even in circular orbits. Our results for the orbital evolution of young hot Jupiters may have the potential to explain the absence of low-mass giant planets inside $\\sim 0.03$ AU from their dwarf stars revealed by transit surveys. ",
        "introduction": "The number of discovered extrasolar planets has been increasing quickly\\footnote{ See the Extrasolar Planets Encyclopaedia (http://exoplanet.eu) for detail.} since the first extrasolar planet, 51 Pegasi b, was detected around a solar-type star \\citep{Mayor95}. 51 Pegasi b is the prototype of a class of exoplanets known as ``hot Jupiters\" because they are generally Jupiter-mass planets located $\\lesssim 0.1$ AU from their parent stars and therefore their equilibrium temperatures are much higher than that of Jupiter. The high temperature of the protoplanetary disk at the orbital radii where they are now provides an environment extremely different from where giant planets are expected to form via either the core-accretion model \\citep{Pollack96,KI02} or the gravitational instability model \\citep{Gammie01,Boss04,Rafikov05}. One of the most commonly adopted solutions is that hot Jupiters have formed at larger orbital radii in the protoplanetary disk and then moved inward to their current locations via the disk-planet interactions \\citep[e.g.][]{Lin96}.  The magnetic fields of Classical T Tauri stars (hereafter CTTS) are typically strong enough to truncate the inner regions of protoplanetary disks to the corotation radius, where the Keplerian angular velocity of the disk is the same as the stellar-spin angular velocity, and create an inner magnetospheric cavity. Besides, the migration due to the planet-disk interaction is expected to slow dramatically once a young hot Jupiter passes the inner disk edge \\citep{Lin96}. The pile-up population of hot Jupiters at $\\sim 0.04$ AU from their solar-type parent stars revealed from radial-velocity searches \\citep{Gaudi05} may be attributed to the existence of the cavity \\citep[][and references therein]{Carr07}.   On one hand, an attempt has been made by \\citet{Setiawan08} to search for a hot Jupiter around a CTTS, although \\citet{Huelamo08} substituted an explanation of a cool stellar spot for an orbiting hot Jupiter. Future observations will provide more evidence for planets of early ages. On the other hand, the simulated orbital evolution of giant planets as they migrate into the cavity was carried out by \\citet{Rice08}. Their results suggest that high-mass hot Jupiters would move closer to their central stars and even be destroyed because planet's entry into the magnetospheric cavity results in a growth in orbital eccentricity $e$. Although the eccentricity excitation is subject to the uncertain properties such as disk mass and viscosity, this model may provide a mechanism to pump up the eccentricity of a hot Jupiter inside the cavity and in turn affect the orbital evolution of the planet via the tidal interactions between the star and planet. In addition, the magnetic interactions between the star and planet in the magnetospheric cavity would be important as well, which may be an upscale analogy to the magnetic interactions between Jupiter and the Galilean satellites \\citep[][and references therein]{Zarka07}.    The coupled tidal evolution of the radius and orbit of hot Jupiters in the absence of a disk has been investigated \\citep{Gu03,Mardling07,IB09}. To obtain noticeable influence on $R_p$, the tidal heating is assumed to be deposited deep in the planet in these models \\citep{Gu04}. The purpose of \\citet{Mardling07} and \\citet{IB09} is to explain the large radii of some transiting hot Jupiters using tidal heating. On the other hand, \\citet{Gu03} showed that for modest eccentricities, young hot Jupiters of radius $\\approx 2$ Jupiter radii at $< 0.04$ AU swell beyond two Jupiter radii and their internal degeneracy is partially lifted\\footnote{For comparison, the initial orbits and eccentricities of a young hot Jupiter in \\citet{IB09} are larger. Therefore the tidal heating in their model can proceed at a modest rate to attain the large $R_p$ on the timescale of $\\sim$ Gyrs.}. Thereafter, their thermal equilibria become unstable and they undergo runaway inflation until their radii exceed the Roche radius. Then Roche lobe overflow ensues. Under the assumption that the overflowing gas immediately loses its orbital angular momentum to the planet and plunges into the central star, these mass-losing planets migrate outwards, such that their semi-major axes and Roche radii increase while their mass, eccentricity, and tidal dissipation rate decrease until the mass loss is quenched. Of course, including a disk would change the orbital evolutions of Roche-lobe-overflowing planets, as has been calculated by \\citet{Trilling98} who do not consider the tidal expansion of $R_p$ and the magnetospheric cavity though.  As more and more hot Jupiters located $\\lesssim 0.04$ AU have been discovered by transit surveys for recent years, it is noteworthy that almost no giant planets of mass $<1$ Jupiter mass ($M_j$) have been detected within $\\sim 0.03$ AU from their solar-type parent stars \\citep[see Figure \\ref{correl}; also see][and references therein]{US07}. This puzzle may be attributed to the subsequent evolutions of hot Jupiters after they enter the magnetospheric cavity of the protoplanetary disk. Since thermal evaporation \\citep[e.g.][]{DW09} has not been undoubtedly identified to be the major cause \\citep[][and references therein]{Murray09}, we shall discuss the potential implication of our model to explain the observation in the end of the paper.  For the above reasons, in this work we concentrate on the orbital evolution of a young hot Jupiter inside the magnetospheric cavity of a protoplanetary disk by taking into account the tidal and magnetic interactions between the star and planet. To restrict to the problem without significantly involving the planet-disk interaction for simplicity, we consider the evolution after the hot Jupiter enters the 2:1 orbital resonance with the inner edge of a disk. The initial eccentricity, which is expected to be excited in the cavity before the planet migrates to the resonance location, is treated as one of the free parameters of the problem. The equations for the orbital and interior evolution of a hot Jupiter inside a magnetospheric cavity are described in \\S2. The results are presented in \\S3, and are summarized and discussed in \\S4. ",
        "conclusions": "We construct a simple model to study the orbital evolution of a young hot Jupiter in an eccentric orbit inside a magnetospheric cavity of a proto-planetary disk around a CTTS. For the sake of simplicity, we assume that the magnetospheric cavity of the protoplanetary disk is truncated by a stellar dipole field. Through the magnetic linkage between the central star and the disk, the rotational period of the star is quickly locked by the inner edge of disk. We then introduce a young hot Jupiter at $a_{2:1}$ and restrict ourselves to the orbital evolution of a hot Jupiter inside the cavity. Assuming that the planet behaves like a plasma, we apply the same formulism for the star-disk magnetic interaction to that for the star-planet interactions. We adopt the equilibrium-tide equations with $Q_*=3\\times 10^5$ and $Q_p=10^6$ to model the star-planet tidal interactions. In so doing, our model focuses on the planet migration due to tidal and magnetic interactions between the planet and the star, and does not consider any interactions between the planet and the disk. However, as the size of the magnetospheric cavity evolves, the planet is artificially pushed to $a_{2:1}$ if it lies beyond $a_{2:1}$. We vary three parameters $B_0$ ($=500$ and 1500 G), initial $M_p$ ($=0.7$, 1, 1.5, 2 $M_j$), and initial $e$ ($e_i=0$, $e_c-0.01$, and $e_c$) in our simulations to investigate the fate of the planet migration under the influence of the cavity size, planet mass, and orbital eccentricity.  Two sizes of magnetospheric cavities are considered to approximately match the bimodal distribution of spin periods of young stars in the Orion Nebula Cluster. The initial $a_{2:1}$ corresponding to the large (small) cavity opened up by the stronger (weaker) stellar fields is $\\sim 0.045$ AU ($\\sim 0.032$ AU). In the case of the large cavity ($B_0=1500$ G), the planets require nonzero $e_i$ to enhance their tidal and magnetic interactions with their central stars and make significant inward migration. The migration rate increases with $e_i$ for a given planet mass. When $e_i$ is as large as the critical value $e_c$, $R_p$ exceeds $R_l$ and the Roche overflow occurs. For planets with different masses, massive planets ($M_p>1M_j$) have higher $e_c$ (i.e. more intense tidal heating) than less massive ones ($M_p\\le M_j$) because massive planets are more difficult to thermally inflate due to their greater gravitational binding energies and cooling rates. As a result, stronger stellar tides raised by the planet with higher $M_p$ and $e_i$ allow for faster migration. Overall, massive planets can therefore migrate further in than less massive ones, as shown in the simulations for $e_i=e_c-0.01$. When $e_i=e_c$, low-mass planets ($M_p=0.7$ and 1 $M_j$) inflate by tidal heating faster than inward migration. Hence they overflow their large Roche lobes at large $a$, resulting in the low density and therefore the runaway mass loss. On the other hand, the high-mass planets ($M_p=1.5$ and 2 $M_j$) migrate inwards fast without significant tidal inflation. In contrast to the low-mass planets, once these planets fill their small Roche lobe at small $a$, their density is high enough that they can undergo stable L1 overflow. During this phase, tidal and magnetic interactions instead of tidal inflation drive them to lose mass, expand, and migrate outward. This stable mass-loss phase proceeds until these outward migrating planets are large enough to become unstable against mass loss. Finally, they lose all gas to their central stars.  In the case of the small cavity ($B_0=500$ G), the simulated planets in circular orbits all quickly migrate in due to the fierce tidal and magnetic interactions until they stably overflow their Roche radii. After that, they can move outwards due to the mass loss and move inwards again by the tidal and magnetic torques. As their mass goes down and their degeneracy is lifted, all of these planets are finally destroyed, suffering from the runaway mass loss as in the large-cavity case.  To study the significance of the magnetic interactions between the star and the planet in our model, we also simulate the cases by shutting down the interaction (i.e. $\\epsilon=0$). We found that the migration of less massive planets is more sensitive to the magnetic interaction, which is enhanced by their easily inflated radii. Setting $\\epsilon=0$ makes the migration rate of low-mass planets slower, as shown in the case for $M_p= 0.7M_j$ with $e_i=0.21$ in the large-cavity case. By contrast, the inward migration of massive planets is less sensitive to the magnetic interaction but more to the tidal interaction. As has been summarized in the preceding paragraph, all planets in the small cavity destruct under both the tidal and magnetic torques at such close-in initial positions from the stars even if $e_i=0$. When $\\epsilon=0$, only the planet of $0.7M_j$ can survive unless the low-mass planet starts with an eccentric orbit with $e_c=0.18$.    The model presented here for the orbital evolution of young hot Jupiters is different from the works by \\citet{Gu03} and \\citet{Trilling98}. In this work, we employ simple models including the star-planet interactions and disk locking, which are not considered by \\citet{Gu03}. Furthermore, we take into account the planet expansion due to the adiabatic mass loss and model the mass loss rate using the approach in \\citet{Kolb90}, whereas \\citet{Gu03} ignore the radius adjustment to the mass loss and use a free parameter to control the mass loss rate. Ignoring the mass-radius relation during the Roche-lobe overflow phase as done in \\citet{Gu03}, a Roche-lobe filled giant planet with $e_i=e_c$ continuously loses its mass via tidal inflation and migrates outwards until the tidal heating rate is weaker than the cooling rate, leading to a survived planet of lower mass. On the other hand, \\citet{Trilling98} consider young hot Jupiters moving inwards via the Type II migration without a magnetospheric cavity and tidal inflation. In their work including the mass-radius relation, once the planets of mass $< 3.36 M_j$ get close enough to their central stars and therefore overflow their Roche lobes, they migrate outwards and finally lose all their mass onto the star. In this regard, our results are compatible with theirs. One of the major differences is that the mass loss is driven by the Type II migration in the model by \\citet{Trilling98}, while in our work the mass loss is driven by the loss of orbital angular momentum via the tidal and magnetic interactions with the parent star whose spin is almost locked by the disk. Moreover, the planets of mass $\\leq 2 M_j$ with $e_i<e_c$ survive in the cavity during the CTTS phase in our model.        For the last few years, transit surveys have revealed the peculiar $M_p-a$ correlation for hot Jupiters: less massive hot Jupiters ($< 1$ Jupiter mass) are almost absent within $\\sim0.03$ AU from their parent dwarf stars of the ages $\\sim$ Gyrs. Our results show that during the CTTS phase, planets of $M_p\\ge1M_j$ with modest initial eccentricities ($e_i\\gtrsim 0.3$) and sufficient magnetic interactions in a large cavity can migrate to $a < 0.03$ AU, while planets of $M_p<1M_j$ cannot. More specifically, given $e_i<e_c$, the planet of $2M_j$ can safely arrive at $a\\lesssim 0.024$ AU and the planet of $1.5M_j$ can get to $a\\lesssim 0.025$ AU in the end of our simulations. Whether or not these planets can further migrate toward their spun-down parent stars during the main-sequence phase depends on $Q'_*$, which is not yet well understood.  \\citet{OL07} suggested that $Q'_*$ induced by hot Jupiters may become large (i.e. $\\gg 10^6$) as CTTSs evolve to main-sequence stars based on their dynamical-tide model. As they and others \\citep{Patzold02,Jiang03} pointed out, we would be extremely fortunate to observe these very close-in hot Jupiters if the current $Q'_*\\sim 10^6$, implying that $Q'_*$ of the solar-type main-sequence host stars may be quite large. However, it should be noted that these very close-in planets were all discovered by transit surveys (see Figure \\ref{correl}). Although transit surveys overall prefer to detect close-in planets \\citep{Gaudi05}, the detectability of a hot Jupiter as a function of the semi-major axis of interest here (i.e. $a\\lesssim 0.05$ AU) is extremely ill-defined as the various transit surveys are subject to different observational limitations such as the transit depth, sampling rate, red noise, and detection threshold \\citep{Pont06}. If $Q'_*$ indeed gradually becomes $\\gg 10^6$ as CTTSs evolve to main-sequence stars, therefore allowing the planets to move a little bit inward during the short transition after $t=10^7$ years while halting the further planet migration for most of the main-sequence phase, our model may have the potential to explain the absence of low-mass hot Jupiters within $a\\lesssim 0.3$ AU.  On the other hand, a large body of studies have suggested that $Q'_\\ast$ associated with the tidal dissipation induced by the hot Jupiters in main-sequence dwarfs could be as small as $\\sim 10^{5-6}$. As a result of the efficient tidal dissipation, the orbits of the hot Jupiters can decay so significantly that the planets may plunge into their host stars during the main-sequence phase \\citep{Jiang03,Jackson08}. This ``accretion\" model may account for the paucity of the extremely close-in hot Jupiters concluded from radial-velocity surveys \\citep[][also see Figure \\ref{correl}]{Gaudi05}. In addition, the relatively young ages of the extremely close-in hot Jupiters revealed by transit surveys may also lend support to the model \\citep{Jackson09}. Although the planet-metallicity correlation independent of the spectral type does not seem to favor the accretion model \\citep{Fischer05}, excessive heavy elements in the shallow convection zone of high-mass dwarfs may be able to sink down to the radiative zone via the double diffusive effect, thereby eliminating the metallicity enhancement and resolving the problem for the model \\citep{Vauclair}. Furthermore, \\citet{Pont09} found an excess spin of the host stars of hot Jupiters, which may arise from tidal spin-up, although the uncertainty of the stellar ages and the limitation of a small number of samples demand more detections in the future to confirm the conclusion. The study by \\citet{Jackson09} for the accretion model is carried out based only on one-mass case (i.e. $1M_j$). If the accretion model is the main cause responsible for the observed $M_p-a$ relation, it would be expected that the extremely close-in giant planets of mass $<1M_j$ of younger ages should also be detected by transit surveys. The orbital decay timescale during the main-sequence phase due only to the tidal dissipation is given by $t_{tide} \\approx (J_0)/(2\\dot{J_*})$ (see equation(\\ref{eq:J*dot})). In the case of a planet of 0.7$M_J$, the planet can survive in $\\sim10^9$ yrs from the distance of $\\sim0.03$AU at $t=10$ Myrs. In contrast, the orbit of a planet of $2M_J$ will decay on the timescale of $\\sim10^7$ yrs from a distance of $\\sim0.025$AU at $t=10$ Myrs. In other words, during the main-sequence phase, less massive planets excite weaker tides on their parent stars and therefore still stay outside $a\\sim0.03$ AU, while the massive planets keep falling into the stars in the way as interpreted by \\citet{Patzold02} and \\citet{Jiang03}. The recently discovered hot Jupiter WASP-18b with an orbital period of 0.94 days and a mass of $10M_J$ may provide a contraint on $Q_*'$ of a main-sequence star in the next few years \\citep{Hellier}.          In this paper, we have attempted to bring several physical components to the context of the evolution of young hot Jupiters. Nonetheless, a large number of assumptions have been made to simplify the physical processes included in our model. The structure of the truncated disk is not modelled. How a planet migrates to $a_{2:1}$ and how the eccentricity evolves during its entry to the cavity \\citep[e.g.][]{Rice08} are not addressed. The disk accretion rate is assumed to be quasi-steady to determine the cavity sizes, while observations show otherwise \\citep[e.g.][]{Baraffe09}. The dipole configuration of the stellar fields connecting to the disk ignores any types of winds even though the effect of disk locking is introduced in a simplified manner to imitate the sink of the total angular momentum in our model. The star-planet magnetic interactions are parameterized to scale with the Poynting flux at the magnetopause of the planet rather than appealing to any specific models, such as the dissipation due to the Alfv\\'en waves and unipolar inductor \\citep[e.g.][]{Zarka07} or induced by tilted stellar dipole fields (e.g. Laine et al. 2008; cf. Papaloizou 2007). The values of $Q'_*$ and $Q'_p$ chosen for the calculations are quite arbitrary; we have not explored a wide range of the values. The interior structure of an inflated planet and the mass-radius relation of a hot Jupiter are inferred from the 1-D numerical simulation \\citep{Bodenheimer01,Gu03} even for a Roche-lobe filled planet. Also, the code ignores thermal properties of the solid core of a giant planet, leaving a question as to what would happen to the core under the intense tidal heating. In a hot Jupiter-star system, the mass loss may occur via the Lagrangian 2 point as well \\citep{Gu03}, which is not considered in the present work. Besides, the mass loss from a planet in an eccentric orbit may be intermittent and nonconservative, while the mass loss rate in our calculation is simply estimated from the L1 overflow at $r=a$. To use a more appropriate expression for the Roche radius and to consider the more realistic mass loss in asynchronous eccentric planetary systems, a more careful treatment \\citep[e.g.][]{Sepinsky} is desired to refine the results. In light of these limitations and possible uncertainties, the simulations covering a wide range of the parameter space along with more realistic modelling on individual issues will be explored in a future work."
    },
    "0911/0911.4018_arXiv.txt": {
        "abstract": "{In this contribution we give a brief overview of the Panoramic Radio Astronomy (PRA) conference held on 2-5 June 2009 in Groningen, the Netherlands. The conference was motivated by the on-going development of a large number of new radio telescopes and instruments which, within a few years, will bring a major improvement over current facilities. Interferometers such as the Expanded Very Large Array (EVLA), Australian SKA Pathfinder (ASKAP), Allen Telescope Array (ATA), Karoo Array Telescope (MeerKAT) and the Aperture Tile In Focus (APERTIF) upgrade to the Westerbork Synthesis Radio Telescope (WSRT) will provide a combination of larger field of view and increased simultaneous bandwidth, while maintaining good collecting area and angular resolution. They will achieve a survey speed 10-50 times larger at 1-2 GHz than the current possibilities, allowing for the first time optical-like all-sky extra-galactic surveys at these frequencies.  Significant progress will be made in many fields of radio astronomy. In this conference we focused on research into the evolution of galaxies over the past few Gyr. In particular, wide-field observations at 1-2 GHz will provide an unprecedented panoramic view of the gas properties and star formation in galaxies, embedded in their environment, from $z\\sim$0.2-0.5 to the present. Within the framework of our current knowledge of the galaxy population at $z<0.5$, we discussed: the key science questions that the new telescopes will permit us to answer in combination with complimentary work at other wavelengths; the observing modes and analysis strategies which will allow us to most efficiently exploit the data; and the techniques for most effectively coping with the huge volume of survey products, so far unusual for the radio community. Emphasis was placed on the complementarity of the upcoming facilities and on their role in paving the way for the technological development and science goals of the Square Kilometre Array.}  \\FullConference{Panoramic Radio Astronomy: Wide-field 1-2 GHz research on galaxy evolution - PRA2009\\\\ June 02 - 05 2009\\\\ Groningen, the Netherlands}  \\begin{document} ",
        "introduction": " ",
        "conclusions": ""
    },
    "0911/0911.4829_arXiv.txt": {
        "abstract": "The quadruple quasar H1413+117 ($z_s$ = 2.56) has been monitored with the 2.0 m Liverpool Telescope in the $r$ Sloan band from 2008 February to July. This optical follow--up leads to accurate light curves of the four quasar images (A--D), which are defined by 33 epochs of observation and an average photometric error of $\\sim$ 15 mmag. We then use the observed (intrinsic) variations of $\\sim$ 50$-$100 mmag to measure the three time delays for the lens system for the first time (1$\\sigma$ confidence intervals): $\\Delta \\tau_{AB}$ = $-$17 $\\pm$ 3, $\\Delta \\tau_{AC}$ = $-$20 $\\pm$ 4, and $\\Delta \\tau_{AD}$ = 23 $\\pm$ 4 days ($\\Delta \\tau_{ij} = \\tau_j - \\tau_i$; B and C are leading, while D is trailing). Although time delays for lens systems are often used to obtain the Hubble constant ($H_0$), the unavailability of the spectroscopic lens redshift ($z_l$) in the system H1413+117 prevents a determination of $H_0$ from the measured delays. In this paper, the new time delay constraints and a concordance expansion rate ($H_0$ = 70 km s$^{-1}$ Mpc$^{-1}$) allow us to improve the lens model and to estimate the previously unknown $z_l$. Our 1$\\sigma$ estimate $z_l$ = 1.88$^{+0.09}_{-0.11}$ is an example of how to infer the redshift of very distant galaxies via gravitational lensing. ",
        "introduction": "The time delay between two images of a gravitationally lensed source depends on the distribution of mass in the lens and the current expansion rate of the Universe \\citep{refsdal64,refsdal66}. This expansion rate is quantified by the Hubble constant $H_0$. If the source is a quasar, the intrinsic quasar variability may be used to determine the time delays between its multiple images. Thus, observed delays for lensed quasars lead to valuable information on $H_0$, provided lensing mass distributions can be constrained by observational data \\citep[e.g.,][]{kochanek04,schechter04,saha06,jackson07,oguri07}. Future large samples of lens systems could be useful tools to obtain accurate estimates of the main cosmological parameters \\citep[e.g.,][]{dobke09}.  For a given lensed quasar, each time delay between two of its images is indeed scaled by a factor containing $H_0$, the lens redshift $z_l$, and additional physical parameters. If $z_l$ is known (this is the usual situation), the Hubble constant and the lensing mass distribution can be simultaneously deduced by using a parametric lens scenario and a set of observational constraints \\citep[including information on the time delay(s); e.g.,][]{schneider06}. However, a large set of constraints is required to reliably determine both cosmological and galactic properties.  Accurate observations of a quadruply imaged quasar in a well--studied galaxy field bring an excellent opportunity to study in detail $H_0$ and the mass distribution of the gravitational lens. Apart from the three independent time delays, and the positions and fluxes of the four images, data on the neighbour galaxies also constraint the lens scenario \\citep[e.g.,][]{jackson07}. Alternatively, if $z_l$ is very uncertain or unknown, one may infer a lens redshift value from the time delay measurements and the rest of constraints (using a value of $H_0$ from other experiments).  In this paper we focus on the quadruple quasar \\object{H1413+117} \\citep[Cloverleaf;][]{magain88}, lying at a redshift $z_s$ = 2.558 \\citep[e.g.,][]{barbainis97}. Although the lens redshift is currently unknown, several neighbour galaxies were detected by \\citet{kneib98}. \\citet{kneib98} found the main lensing galaxy G1 amid the four quasar images, and presented astrometric and photometric data for additional galaxies surrounding \\object{H1413+117}. They analysed a secondary lensing galaxy (G2) close to G1, as well as some other objects probably belonging to an overdensity (galaxy group/cluster) at photometric redshift $z_{ove} \\sim$ 0.9. \\citet{faure04} also found evidence for the presence of two different overdensities at $z_{ove}$ = 0.8 $\\pm$ 0.3 \\citep[corresponding to the structure discovered by][]{kneib98} and $z_{ove}$ = 1.75 $\\pm$ 0.2, which could contribute noticeably to the lensing potential.  Very recently, \\citet{macleod09} have used the positions of G1, G2, and other candidate lensing galaxies, the quasar image positions \\citep{turnshek97}, new mid--IR flux ratios, and some priors to constrain the lensing mass distribution. They have concluded that the galaxy pair G1--G2 and an external shear (likely related to the observed galaxy overdensities) are required to explain the observations. From an optical monitoring over the 1987--94 period, \\citet{ostensen97} also reported a quasi--simultaneous variability of the four quasar images. Unfortunately, the scarce sampling did not permit them to estimate the time delays in the system. Hence, the measurement of time delays for the Cloverleaf quasar should improve knowledge about the lens (mass and redshift), as well as allow completion of new studies (e.g., microlensing variability).  In Section 2 we present new optical light curves of \\object{H1413+117}. In Section 3, from these light curves, we estimate the time delays between quasar images. In Section 4 we compare the most recent lens scenario with all relevant observations. In Section 5 we discuss our results and put them into perspective. ",
        "conclusions": "In the lens system \\object{H1413+117}, the main lensing galaxy G1 is a very faint object \\citep[$R >$ 22.7 mag;][]{kneib98}, surrounded by four close and relatively bright quasar images ($R \\sim$ 18 mag). Thus, it is very difficult to separate the spectrum of the galaxy from those of the quasar images. Moreover, the quasar spectra indicate the presence of intervening objects (absorption lines) at different redshifts less than $z_s$ = 2.56 \\citep[e.g.,][and references therein]{monier98}, so one cannot decide on the redshift of G1. The secondary lensing galaxy G2 and other galaxies in the vicinity of the lensed quasar are also very faint objects \\citep[$R >$ 23 mag;][]{kneib98}, and spectroscopic redshifts of these galaxies are not available yet. Photometric data of field galaxies are consistent with the presence of two galaxy overdensities at $z_{ove}$ = 0.8 $\\pm$ 0.3 and $z_{ove}$ = 1.75 $\\pm$ 0.2 \\citep{kneib98,faure04}. For example, the galaxy G2 has a photometric redshift of about 2 \\citep{kneib98}. Our gravitational lensing estimate of the redshift of G1--G2: $z_l$ = 1.88$^{+0.09}_{-0.11}$ (1$\\sigma$ interval), is in reasonable agreement with the photometric redshift of G2 and the most distant overdensity, as well as the absorption system at $z_{abs}$ = 1.87. However, the nearest group/cluster is far away from the principal gravitational deflector (galaxy pair G1--G2).  So far it is not considered a possible uniform external convergence, i.e., $\\kappa_{ext}$ = 0. If $\\kappa_{ext} \\neq$ 0, then fits to $\\kappa_{ext}$ = 0 lens scenarios (lens models with $\\kappa_{ext}$ = 0) must be conveniently rescaled. Original estimates of $H_0$ (if that were the case) or $z_l$ (in our case) also require suitable corrections. This is because of the so--called mass sheet degeneracy \\citep[e.g.,][]{falco85,gorenstein88,saha00,nakajima09}. For \\object{H1413+117}, it is unclear what is the main perturber producing external shear, and perhaps external convergence. For example, infrared photometry of the neighbour galaxy H2 is consistent with a redshift of about 2 \\citep{kneib98}, so it could be located in the principal lens plane. Moreover, the orientation of the external shear, $\\theta_{\\gamma_{ext}} = 45\\fdg 4$, is in the direction of this galaxy \\citep[e.g., see Fig. 4 of][]{macleod09}. Thus, H2 and some other related objects (belonging to the very distant overdensity) might produce most of the external shear and a negligible convergence. Apart from this optimistic perspective, one may also consider that the main perturber is the most distant group/cluster as a whole. If this were true, the external convergence at the quasar positions would be $\\kappa_{ext}$ = $\\gamma_{ext} \\sim$ 0.1 for a singular isothermal sphere. Taking $\\kappa_{ext} \\sim$ 0.1 into account, the derived lens redshift increases by $\\sim$ 3\\% (+ 0.05). This increase represents only one half of our 1$\\sigma$ uncertainty in $z_l$ ($\\sim$ 0.10; see above).  We can also quantify the gravitational influence of the nearest overdensity. A weak lensing analysis of this structure indicated a shear direction of $45\\degr$ \\citep[see the last column in Table 5 of][]{faure04}. This shear direction coincides with our strong lensing determination in Table~\\ref{tbl3}, which suggests that the nearest overdensity might play a noticeble role in the lens phenomenon. \\citet{faure04} also determined an upper limit on the shear strength: $\\gamma <$ 0.17 \\citep[see Fig. 7 of][]{faure04}. \\citet{keeton03} and \\citet{momcheva06} discussed the effective convergence and shear when a perturber does not lie at $z_l$, e.g., $z_{per} < z_l$. Assuming a singular isothermal sphere to describe the perturber, $\\kappa_{eff} = \\gamma_{eff} \\sim (1 - \\beta) \\gamma$, where $\\beta \\sim$ 0 if $z_{per} \\sim z_l$ and $\\beta >$ 0 if $z_{per} < z_l$ \\citep[see Appendix of][]{momcheva06}. From the involved redshifts and the upper limit on the shear strength, we infer that $\\kappa_{eff} = \\gamma_{eff} <$ 0.03. Thus, the nearest structure cannot account for the external shear in the lens system, and it likely generates no more than 10\\% of the total shear strength. The constraint on the effective convergence leads to a negligible increase in $z_l$, i.e., $\\Delta z_l$ is below + 0.01. In some other lens systems, there are also perturbers at redshifts less than those of the principal lensing objects, which produce a few hundredths of external convergence and shear \\citep[e.g.,][]{fassnacht02,fassnacht06}.  Our lens scenario incorporates singular isothermal mass profiles. Despite the fact that these profiles lead to an acceptable fit ($\\chi^2$/dof = 7.5/7), other distributions of mass (including a core, some deviation from the isothermal behaviour or both ingredients) could also lead to good fits for the observational data. This is the well--known profile degeneracy \\citep[e.g.,][]{jackson07}. An exhaustive study of mass distributions consistent with the available observational constraints is out of the scope of this paper. However, we check the influence of non--isothermal profiles of G1 on $z_l$ estimations. The power--law index of G1 ($\\alpha$) is assumed to be around 1 (isothermal index), and additional fits with singular $\\alpha \\neq 1$ profiles are done. For $\\alpha$ = 1.1, the solution is characterized by $\\Delta \\chi^2 = \\chi^2 - \\chi^2(\\alpha = 1)$ = 0.1. For $\\alpha$ = 0.8--0.9, we find a very modest improvement in $\\chi^2$, $\\Delta \\chi^2$ = $-$0.1. The full range $\\alpha$ = 0.8--1.1 leads to best--fit values of $z_l$ within our 1$\\sigma$ \"isothermal\" estimate in Section 4.  Microlensing variations in \\object{H1413+117} may shed light on the nature and structure of the source quasar \\citep[e.g.,][]{lewis98,popovic06}. The three time delay measurements in Section 3 and the new data on the lens (mass and redshift) in Section 4 are useful tools for analyses of microlensing variability. Once the delays are known, it is possible to properly compare quasar light curves and to search for microlensing signals. Moreover, the (lens and source) redshifts and the improved lens model allow construction of microlensing magnification patterns and simulated microlensing light curves \\citep[e.g.,][]{wambsganss90}. The key parameters for microlensing simulations are the convergence and shear strength at the positions of the quasar images. To address the space distribution of both convergence and shear, we consider our lens model in the third column of Table~\\ref{tbl3}. This gives ($\\kappa_A$, $\\gamma_A$) = (0.51, 0.58), ($\\kappa_B$, $\\gamma_B$) = (0.52, 0.32), ($\\kappa_C$, $\\gamma_C$) = (0.48, 0.37), and ($\\kappa_D$, $\\gamma_D$) = (0.57, 0.65). Although we cannot rule out the existence of an external convergence at a level of 0.1, this scenario is uncertain (see above). Hence, we do not take into account any external convergence due to galaxy overdensities along the line of sight to the lensed quasar.  Our lens model causes different magnifications at the four positions of the quasar images. The model magnification ratios are $B/A$ = 0.82, $C/A$ = 0.78, and $D/A$ = 0.45. All these ratios agree, as expected, with the mid--IR measurements by \\citet{macleod09}. However, the model ratios are not included in the error bars of our optical ($r$--band) flux ratios in Section 3. These differences between optical and model ratios are probably due to dust extinction and microlensing magnification. If microlensing is currently playing a role \\citep[e.g., the optical continuum of the image D could be magnified by microlensing;][and references therein]{anguita08}, it should be a long--term effect that induces small (optical) flux variations on time scales of several months. This microlensing variability scenario is supported by recent studies for other systems with non--local lens galaxies \\citep[e.g.,][]{gaynullina05,fohlmeis07,shalyapin09}."
    },
    "0911/0911.1121_arXiv.txt": {
        "abstract": "We critically examine the constraints on internal angular momentum transport which can be inferred from the spin down of open cluster stars.  The rotation distribution inferred from rotation velocities and periods are consistent for larger and more recent samples, but smaller samples of rotation periods appear biased relative to $v\\sin i$ studies.  We therefore focus on whether the rotation period distributions observed in star forming regions can be evolved into the observed ones in the Pleiades, NGC\\,2516, M\\,34, M\\,35, M\\,37, and M\\,50 with plausible assumptions about star-disk coupling and angular momentum loss from magnetized solar-like winds.  Solid body models are consistent with the data for low mass fully convective stars but highly inconsistent for higher mass stars where the surface convection zone can decouple for angular momentum purposes from the radiative interior.  The Tayler-Spruit magnetic angular momentum transport mechanism, commonly employed in models of high mass stars, predicts solid-body rotation on extremely short timescales and is therefore unlikely to operate in solar-type pre-MS and MS stars at the predicted rate. Models with core-envelope decoupling can explain the spin down of 1.0 and 0.8 solar mass slow rotators with characteristic coupling timescales of $55\\pm 25$ Myr and $175\\pm 25$ Myr respectively.  The upper envelope of the rotation distribution is more strongly coupled than the lower envelope of the rotation distribution, in accord with theoretical predictions that the angular momentum transport timescale should be shorter for more rapidly rotating stars. Constraints imposed by the solar rotation curve are also discussed.  We argue that neither hydrodynamic mechanisms nor our revised and less efficient prescription for the Tayler-Spruit dynamo can reproduce both spin down and the internal solar rotation profile by themselves. It is likely that a successful model of angular momentum evolution will involve more than one mechanism.  Further observational studies, especially of clusters younger than 100 Myr, will provide important additional constraints on the internal rotation of stars and could firmly rule out or confirm the operation of major classes of theoretical mechanisms. ",
        "introduction": "\\label{sec:intro}  Rotation is an important attribute of the life of a star. When it is fast enough, rotation can trigger various (magneto-)hydrodynamic instabilities that will drive mixing and angular momentum transport in the star (e.g., \\citealt{z92,s99}). Unfortunately, except for the Sun, observations give information only about surface rotation of stars. Therefore, indirect methods have to be used to study rotation-driven transport processes in stellar interiors. When stars reach the main sequence (MS) their surface response to the torque from a magnetized wind depends on the timescale for internal angular momentum transport. If open clusters are treated as an evolutionary sequence, the assumptions of solid body (SB) and differential rotation (DR) lead to statistically distinguishable differences in the time evolution of the distributions of cluster star rotation rates (see, for example, \\citealt{kmc95,kea97,a98}). At late ages, models also have to be consistent with the strong coupling evident in the solar internal rotation profile, with nearly SB rotation in the radiative core down to $\\sim$0.2$\\,R_\\odot$ (\\citealt{tst95,cea03}).  The prior estimates of the timescale for core-envelope coupling in solar-type stars usually agreed on a short coupling time of order one Myr for the fastest rotators. However, they considerably disagreed on the coupling time for slowly rotating solar analogs for which the estimated values varied from 10\\,--\\,100 Myr (e.g., \\citealt{kmc95,a98}) to 0.5 Gyr (e.g., \\citealt{iea07}). This discrepancy has been caused, on the one hand, by not carefully treating the selection effects and statistics and, on the other hand, by making particular assumptions about such things as the disk-locking time, initial rotation period distribution or the relative cluster ages which have since been updated. For example, the early work assumed that $\\alpha$\\,Per was 50 Myr and the Pleiades was 70 Myr old (\\citealt{mb91}). In this work we apply rigorous methods of statistical analysis to the most recent large observational datasets on rotation periods of low-mass stars in open clusters and extensive computational modeling of their rotational evolution to constrain the timescale for core-envelope coupling in the slowest rotators.  In this paper we find that the timescale for core-envelope coupling depends on both rotation rate and mass, but it is significantly longer than would be predicted if SB rotation was enforced on a short timescale. Angular momentum transport by magnetic torques generated by the Tayler-Spruit dynamo (\\citealt{s99,s02}) does not pass this key test because it always enforces SB rotation in a solar-type star, no matter how slowly it rotates. This disagrees with the observational evidence that slow rotators in young clusters are most likely to possess DR (their cores rotate faster than their envelopes) rather than follow the $P$-age relations computed using SB rotation models. Furthermore, we show that even a revised prescription for the Tayler-Spruit dynamo proposed by \\cite{dp07}, which reduces the effective magnetic viscosity by nearly two orders of magnitude, is still too efficient in redistributing angular momentum in the outer parts of the radiative core; this is inconsistent with the observed evolution of slowly rotating solar analogues in the $P$-age plane. The revised prescription comes to a better agreement with observations when it is supplemented by an additional angular momentum transport mechanism that should be able to operate in the inner core on a longer timescale. Such a mechanism is also needed to assist the revised prescription in shaping the solar SB rotation.  In our work, three different theoretical models are employed to study the rotational evolution of solar-type stars: a simple double-zone model and two full stellar evolutionary models, one with a constant viscosity and the other with the effective magnetic viscosities provided by the original and revised prescriptions for the Tayler-Spruit dynamo. The models are described in detail and compared with each other in Section 2. Because the choice of basic model parameters is necessarily constrained by observations, we find it helpful to briefly introduce the samples of rotation period and $v\\sin i$ data, that will be used for a more detailed analysis later in Section~3, already at the beginning of Section~2. In Section 3, we present the results of our computations of rotation period distributions for solar-type stars, as they evolve from the pre-main sequence (pre-MS) deuterium birth line to the solar age, and compare them with observations of rotation periods in open clusters. A brief discussion and our main conclusions are given in Section 4. ",
        "conclusions": "One of the main conclusions made in this paper is that the observed rotation period distributions of low-mass stars in open clusters do not seem to support the hypothesis that all solar-type stars evolve as SB rotators. Instead, statistical analyses of these data show that only the fastest rotators among solar-type stars can be considered to possess SB rotation during their entire evolution, whereas their slowly rotating counterparts are most likely to manifest internal DR between the ages of $\\sim$\\,30 Myr and several hundred Myr. This is not a new result though because, e.g., \\cite{iea07} came to a similar conclusion. A novelty of our work is that we have used the entire period distributions of solar-type stars in a number of open clusters of different ages, not just some of their statistics, and extensive Monte-Carlo simulations to put this conclusion on a more rigorous quantitative basis. This seems to be quite a reasonable approach to the solution of the problem, given that it has a large number of poorly constrained parameters. In particular, we have found that a star with $M = 1.0\\pm 0.1\\,M_\\odot$ that starts its rotational evolution with a period $P_0 \\ga 2$\\,--\\,4 days should have rotation of its convective envelope and radiative core coupled on a timescale of order $\\tau_{\\rm c} = 55\\pm 25$ Myr, where the systematic uncertainty of this estimate takes into account the anticipated decrease of the coupling time with an increase of the disk-locking time. For a slightly less massive star with $M = 0.8\\pm 0.1\\,M_\\odot$, the coupling time increases to $\\tau_{\\rm c} = 175\\pm 25$ Myr. Given that the initial period distributions, those for the youngest open clusters, have rather densely occupied bins up to $P_0 \\approx 12$ days (Fig.~\\ref{fig:f2}), it turns out that quite large fractions of solar-type stars (up to 50\\%) should go through a phase of DR evolution.  It is important to note that solar-type stars in open clusters older than a few hundred Myr are not suitable for a type of statistical analysis employed by us. The problem with the older clusters is that rotation periods of solar-type stars in them have already converged too close to each other, all aiming eventually to approach the solar rotation. Therefore, period distributions for the older clusters do not allow to make an unambiguous conclusion about the rotational evolution of solar type stars. To illustrate this, we have evolved the NGC\\,2264 period distribution of stars with $M = 1.0\\pm 0.1\\,M_\\odot$ to a period distribution at $t = 550$ Myr corresponding to the age of M\\,37 (Fig.~\\ref{fig:f14}). We see that at this old age the period distribution mapping admits two branches of solutions, one with DR and another with SB rotation. This bimodality is caused by the very narrow range of the mapping at old ages which finally degenerates into a point at the solar age.  The second novelty of our investigation is that we have related the coupling time from the double-zone model to its corresponding value of the constant viscosity from our full stellar evolutionary model (Fig.~\\ref{fig:f7}). In particular, the minimum coupling time of 30 Myr obtained in our analysis of DR of solar-type stars roughly corresponds to $\\nu_0 = 7.5\\times 10^4$ cm$^2$\\,s$^{-1}$, whereas the longest coupling time of $\\sim$80 Myr for stars with $M = 1.0\\pm 0.1\\,M_\\odot$ implies that in some of them the internal transport of angular momentum can be as slow as that modeled with a viscosity $\\nu_0\\approx 3\\times 10^4$ cm$^2$\\,s$^{-1}$.  Finally, we have shown that original Spruit's prescription always enforces SB rotation in a solar-type star, no matter how slowly it rotates. Therefore, it cannot be accepted as a physical mechanism for the internal angular momentum redistribution in radiative cores of solar-type stars many of which are indirectly proved to possess a degree of DR.  We have also found that the revised prescription for the Tayler-Spruit dynamo (\\citealt{dp07}) results in the rotational evolution superficially resembling that obtained using the original prescription, although the former leaves a large differentially rotating radiative core in the present-day solar model (Fig.~\\ref{fig:f6}). Hence, the revised prescription appears to be in conflict both  with the rotation period data for solar-type stars in open clusters (like the original prescription) and with the helioseismic data (unlike the original prescription). It is worth seeing if other possible angular momentum transport  mechanisms can assist the revised prescription in bringing the rotating solar model closer to observations. To test this idea, we have supplemented the effective magnetic viscosity (\\ref{eq:nueminrev}) with viscosities arising from the secular shear instability (denoted by the subscript ``ss'' below), Goldreich-Schubert-Fricke instability (GSF), as well as with the molecular viscosity (mol) and a viscosity $\\nu_{\\rm mc}=|rU(r)|/5$ that approximately describes the transport of angular momentum by meridional circulation as a diffusion process: \\bea \\nu = (\\nu_{\\rm e})_{\\rm min}^{(\\ref{eq:nueminrev})} + \\nu_{\\rm ss} + \\nu_{\\rm GSF} + \\nu_{\\rm mc} + \\nu_{\\rm mol}. \\label{eq:allvisc} \\eea To calculate the quantities $\\nu_{\\rm ss}$, $\\nu_{\\rm GSF}$, and the meridional circulation velocity $U(r)$, we have used the corresponding equations from Appendix to the paper of \\cite{chpt05}. Note that their expression for $U(r)$ neglects terms depending on the $\\mu$-gradient and its derivatives. Therefore, the meridional circulation in this approximation is expected to penetrate deeper into the radiative core than in the case of its fully consistent implementation originally proposed by \\cite{mz98} that had later been applied to study the rotational mixing in solar-metallicity MS stars with $M\\geq 1.35\\,M_\\odot$ by \\cite{pea03}.  In Fig.~\\ref{fig:f15}, the dotted curves represent results of our computations with the combined viscosity (\\ref{eq:allvisc}) substituted into equation (\\ref{eq:amt}). Note that a slightly increased amount of angular momentum that can be transported from the core to the envelope on a longer hydrodynamical timescale has changed our results in the right directions: first, after the wind constant $K_{\\rm w}$ is properly recalibrated, the $P$-age relation has shifted toward longer periods, and, second,  the degree of DR in the core has slightly been reduced. To find out if these changes will further grow in the same directions when the efficiency of supplementary angular momentum transport mechanisms increases, we have multiplied $\\nu_{\\rm GSF}$, the dominating viscosity in the core, by a factor of 10. Results of this artificial viscosity enhancement are plotted in Fig.~\\ref{fig:f15} with dashed curves. We see that both the $P$-age relation and the core rotation profile have indeed continued to change in the right directions. This exercise shows that the revised prescription cannot be rejected as easily as the original one. There remains a possibility that, when being assisted by other transport processes, it may reproduce the observational data yet.  Our main conclusion is critically based on a comparison of theoretical predictions and observations of rotation periods for the slowest rotators in the intermediate-age open clusters. Therefore, if the observational data were biased toward the longest periods that would undermine confidence in our results. The real situation turns out to be opposite. Observations for some of our used open clusters (for those with smaller data samples) are actually biased toward the shortest periods, as is expected from the photometric period measurement procedure that needs longer observational times to accumulate data sufficient for extracting longer periods. The following example illustrates this bias. Filled circles in upper panels of Fig.~\\ref{fig:f16} represent the Pleiades $v\\sin i$ unbiased data from \\cite{aps03}. For comparison, open circles in the upper left panel show the Pleiades data used in our work that do appear to be biased toward lower rotation periods. On the other hand, open circles in the upper right panel are the M\\,35 data from \\cite{mms09}. Their original periods have been transformed into $v\\sin i$ values using a randomly generated angle $0\\leq i\\leq \\pi/2$ and $R = R_\\odot$. The lower right panel shows that the unbiased data for the two clusters of similar age look alike ($P_{\\rm KS} = 0.112$, the medians are 6.8 and 8.5, the first and third quartiles are 5.4 and 13 for the Pleiades, and 4.7 and 18 for M\\,35). On the contrary, the biased Pleiades data in the lower left panel have very different median and third quartile values, 15 and 39, respectively.  To sum up, our main conclusions can concisely be formulated as follows. Whereas the period distributions for the fastest rotators among solar-type stars in open clusters are very well reproduced assuming their SB rotation, the period distributions for the slowest rotators are better described by stellar models with DR in their radiative cores. This conclusion is in agreement with previous results reported by \\cite{iea07}. Our new result is that the original prescription for the Tayler-Spruit dynamo always enforces SB rotation in a solar-type star even if the star is a slow rotator. Therefore, this angular momentum transport mechanism is unlikely to operate in solar-type early MS stars. The revised prescription for the Tayler-Spruit dynamo cannot explain the observed period distributions either. Besides, it leaves a large rapidly rotating radiative core in the present-day solar model. To be consistent with observations, the revised prescription needs to be assisted by other angular momentum transport mechanisms that must be able to penetrate into the deep core and operate on a longer timescale."
    },
    "0911/0911.2314_arXiv.txt": {
        "abstract": "We present the results of a spectral principal component analysis on 9046 broad-line AGN from the Sloan Digital Sky Survey.  We examine correlations between spectral regions within various eigenspectra (e.g., between Fe II strength and H$\\beta$ width) and confirm that the same trends are apparent in spectral measurements, as validation of our technique.  Because we found that our sample had a large range in the equivalent width of [\\ion{O}{3}] $\\lambda5007$, we divided the data into three subsets based on [\\ion{O}{3}] strength.  Of these, only in the sample with the weakest equivalent width of [\\ion{O}{3}] were we able to recover the known correlation between [\\ion{O}{3}] strength and full width at half maximum of H$\\beta$ and their anticorrelation with \\ion{Fe}{2} strength.  At the low luminosities considered here ($L_{5100 \\mbox{\\AA}}$ of $10^{42}-10^{46}$ erg s$^{-1}$), interpretation of the principal components is considerably complicated particularly because of the wide range in [\\ion{O}{3}] equivalent width.  We speculate that variations in covering factor are responsible for this wide range in [\\ion{O}{3}] strength. ",
        "introduction": "\\label{sec:intro}  Active galactic nuclei (AGN) display an incredibly diverse set of observed properties.  They reside in host galaxies of many morphologies, and have a range of radio luminosities and x-ray spectra. There are obscured AGN, known as Type II AGN, that exhibit only narrow forbidden emission lines emerging from the narrow line region (NLR), with scales $\\gtrsim$ 100 pcs.  Also, there are broad-line AGN that exhibit both NLR emission and broad ($>1000$ km s$^{-1}$) emission lines from the broad line region (BLR) with scales of light days.  There is great variation in the strengths, shapes, and widths of emission lines. This diversity provides many opportunities for investigating the physics at work in these systems.  A powerful tool for studying AGN phenomenology is principal component analysis (PCA).  PCA finds the eigenvectors, or most dominant relationships, within a set of input data.  \\citet{1992ApJS...80..109B}, hereafter BG92, performed a principal component analysis on a sample of $z < 0.5$ quasars from the Palomar-Green Bright Quasar Survey \\citep[BQS; ][]{1983ApJ...269..352S}, a UV-selected sample of point-like quasars (where most of the light fell within 1\" on blue Palomar Sky Survey prints).   Spectral parameters that they used in the PCA included the the EW and FWHM of H$\\beta$, optical \\ion{Fe}{2}, [\\ion{O}{3}] $\\lambda$5007, and He II $\\lambda$4686;  ratios of the EWs of optical \\ion{Fe}{2}, [\\ion{O}{3}], and He II relative to H$\\beta$; and magnitude and peak flux of [\\ion{O}{3}] relative to H$\\beta$, absolute magnitude, optical-to-X-ray spectral slope ($\\alpha_{\\rm{ox}}$), radio-loudness, and line shift, shape, and asymmetry of H$\\beta$.  The eigenvector that distinguished most quasars from one another was dominated by the anticorrelation between the broad-line \\ion{Fe}{2} strength relative to broad H$\\beta$ versus the peak flux of [\\ion{O}{3}].  They found a second eigenvector that included a correlation among optical luminosity $L$, He II, and $\\alpha_{\\rm{ox}}$.  The BG92 anticorrelation between \\ion{Fe}{2}/H$\\beta$ and [\\ion{O}{3}]/H$\\beta$, called eigenvector 1 (EV1), revealed a tie between the BLR and NLR, even though these regions occupy very different physical scales.  At one end of this relationship lie narrow-line Seyfert 1 galaxies \\citep[NLS1; ][]{1985ApJ...297..166O}, with relatively narrow broad lines, strong \\ion{Fe}{2}/H$\\beta$, weak [\\ion{O}{3}]/H$\\beta$, quiet radio emission, and a steep slope between the x-ray and optical continuum \\citep{1994ApJ...420..110L, 1997ApJ...477...93L}.  At the other end of EV1 lie radio galaxies with broad ($> 4000$ km s$^{-1}$) BLR emission, strong [\\ion{O}{3}]/H$\\beta$, very little \\ion{Fe}{2}, and flat slopes between the optical and x-ray regions.  EV1 includes a correlation with the linewidth of H$\\beta$ \\citep{1996A&A...305...53B}, which is used to estimate virial black hole masses.  This led to the idea that EV1 could be driven by black hole mass ($M_{\\rm{BH}}$) or Eddington ratio \\citep[$L_{\\rm{bol}}/L_{\\rm{Edd}}$; BG92, ][]{1994ApJ...420..110L, 1997ApJ...477...93L}, so that high accretion ratios correspond to the NLS1 end of EV1. A correlation with $L_{\\rm{bol}}/L_{\\rm{Edd}}$ provides an explanation for the observed x-ray properties in AGN such that a higher accretion state, or higher $L_{\\rm{bol}}/L_{\\rm{Edd}}$, causes the accretion disk to become thicker and produce more soft x-rays \\citep[e.g.][]{2004AJ....127.1799G}.  Physically, this thicker disk could be related to disk outflows that additionally excite \\ion{Fe}{2} \\citep{2000NewAR..44..531C}.   It is also possible that the covering factor is higher for the BLR in strong \\ion{Fe}{2} objects, resulting in less continuum reaching the NLR gas, which would explain the \\ion{Fe}{2} - [\\ion{O}{3}] anticorrelation.  Over time, many parameters have been examined for correlations with EV1, such as the slope of the x-ray spectrum \\citep{1994ApJ...420..110L,1997ApJ...477...93L,1995MNRAS.277L...5P}, black hole mass and $L_{\\rm{bol}}/L_{\\rm{Edd}}$ or $L/M$ \\citep{1994ApJ...420..110L,1997ApJ...477...93L,1992ApJS...80..109B,2002ApJ...565...78B,2003MNRAS.345.1133M}, radio properties \\citep{2002ApJ...565...78B} and the \\ion{Fe}{2} emission \\citep{2003ApJ...586...52S}.  Also, many authors have utilized PCA to analyze correlations among measured quantities \\citep{2004AJ....127.1799G, 2008ApJ...678...22H}, or used PCA to find and investigate outliers with unusual properties \\citep{1999ASPC..162..363F}.  Still others have chosen to apply a PCA directly to the spectra \\citep{1994ApJ...430..495B, 2003ApJ...586...52S, 2004AJ....128.2603Y}.  These investigations have provided important insights, but we still lack a complete physical understanding of the connection between the inner BLR and the far-ranging NLR.  The Sloan Digital Sky Survey \\citep[SDSS; ][]{2000AJ....120.1579Y} has provided an immense data set of broad-line AGN spectra that can be useful for understanding whether the BG92 EV1 relationships persist in large samples.  The SDSS is especially appropriate:  large samples with well-defined selection criteria, larger luminosity ranges with reasonable completeness, and better spectral resolution than has been available for most broad-line AGN samples in the past.  The sheer number of objects also provides a greater diversity of broad-line AGN properties available for study.  From such a large sample, we can also identify extreme objects for further study.  In this paper we report an SPCA of broad-line AGN from the SDSS.  In \\S ~\\ref{sec:method} we discuss our sample and SPCA methodology.  In \\S ~\\ref{sec:results} we present the results of performing spectral PCA on our data set, and find that, over the broad range of properties in the sample, we cannot find simple physical interpretations for the principal components, perhaps because of nonlinear relationships between the physical parameters and the observables.  We further probe the relationships in the sample by dividing it into three subsets.  In \\S ~\\ref{sec:bg92} we investigate the differences between our work and that of BG92.  In \\S ~\\ref{sec:outliers}, we discuss properties of our most extreme subset, the strong NLR objects, and examine possible reasons for their extraordinarily strong [\\ion{O}{3}] emission.  We find that high covering factors could be a plausible explanation for the strong NLR emission.  We present a summary of our results in \\S ~\\ref{sec:summary}.  All luminosities in this paper were calculated using the cosmological parameters $H_{0}=70$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{\\rm{M}}=0.3$, and $\\Omega_{\\Lambda}=0.7$. ",
        "conclusions": "\\label{sec:summary} We have presented results of an SPCA of 9046 broad-line AGN from the SDSS DR5 with $0.1 < $z$ < 0.56$ and $10^{42} < L_{5100 \\mbox{\\AA}} < 10^{46}$ erg s$^{-1}$. We investigated the properties of the rest-frame 4000 - 6000 \\AA{}{} spectrum and found that a large range in EW$_{\\mathrm{[O III]}}$ dominates the behavior of the sample.  This range causes difficulty in interpreting the PCs.  To try to understand the behavior of [\\ion{O}{3}], we divided our sample into three subsets to restrict the broad range of EW$_{\\mathrm{[O III]}}$.  Our most extreme [\\ion{O}{3}] subset had principal components dominated by NLR emission.  Traditionally, the first principal component in eigenvector analysis (EV1) of quasar spectra has been an anticorrelation between the strengths of \\ion{Fe}{2}/H$\\beta$ and [\\ion{O}{3}]/H$\\beta$, and a correlation between the strength of \\ion{O}{3}/H$\\beta$ with H$\\beta$ linewidth.  While the entire sample and subsets show evidence for the anticorrelation between EW$_{\\mathrm{[O III]}}$ and \\ion{Fe}{2} in their components and mean spectra, we find the H$\\beta$ linewidth to be related in only the lowest EW$_{\\mathrm{[O III]}}$ subset.  Overall, the width of the Balmer lines seems to have little dependence on the strength of [\\ion{O}{3}] or \\ion{Fe}{2}.  We find that the differences between our work and BG92 and S03 are caused by the substantially lower luminosities included in our sample.  Also, we find that the usual anticorrelation between EW$_{\\mathrm{[O III]}}$ and $L$ only applies for the very high-$L$ end of our sample, and thus the importance of relationships among \\ion{Fe}{2}, [\\ion{O}{3}], and H$\\beta$ linewidth is a strong function of $L$ (or perhaps $L_{\\rm{bol}}/L_{\\rm{Edd}}$). Our results are consistent with the ``disappearing NLR'' effect suggested by \\cite{2004ApJ...614..558N} where high-$L$ AGN experience a decrease in NLR emission due to loss of NLR gas from the system.  The strong EW$_{\\mathrm{[O III]}}$ objects in our sample have dominant NLR lines that do not seem to correlate with any other properties of the broad-line AGN, including $L_{5100 \\mbox{\\AA}}$, $L_{\\rm{bol}}/L_{\\rm{Edd}}$, color, BLR properties, or radio emission. We did not find conclusive evidence for the cause of the high EW$_{\\mathrm{[O III]}}$ in these objects, although EW$_{\\mathrm{[O III]}}$ could be enhanced by suppression of the AGN continuum.  We suggest that the intrinsic range in [\\ion{O}{3}] emission is primarily caused by a covering factor that is highly variable among AGN.  We thank Sarah Salviander for access to her data.  Without her cooperation this work would not have been feasible.  Also we acknowledge the helpful comments and discussion provided by Greg Shields and John Barentine.  We thank Todd Boroson for a helpful referee report."
    },
    "0911/0911.2855_arXiv.txt": {
        "abstract": "s{Underground laboratories host two kind of experiments at the frontier of our knowledge in Particle Physics, Astrophysics and Cosmology: the direct detection of the Dark Matter of the Universe and the search for the Neutrinoless Double Beta Decay of the nuclei. Both experimental quests pose great technical challenges which are being addressed in different ways by an important number of groups. Here a updated review of the efforts being done to detect Dark Matter particles is presented, emphasizing latest achievements.} ",
        "introduction": "Since first suggested by Zwicky in the 1930s, the existence of an invisible and non-conventional matter as a dominant part of our Universe has been supported by an ever increasing body of observational data. The latest precision cosmology measurements \\cite{Spergel:2003cb,Spergel:2006hy,Dunkley:2008ie} further constrain the geometry of the Universe to be flat ($\\Omega \\sim 1\\pm0.04$), and its composition to be mostly dark energy ($\\Omega_\\Lambda = 74 \\pm 3\\%$) and non-baryonic cold dark matter ($\\Omega_{c} \\sim 21.4 \\pm 3\\%$), leaving about $\\sim 4.4\\%$ for ordinary baryonic matter. Dark energy is a theoretical concept related to Einstein's cosmological constant, the nature of which is essentially unknown. Dark matter, on the contrary, could be composed by elementary particles with relatively known properties, and which could be searched for by a variety of means. These particles must have mass, be electrically neutral and interact very weakly with the rest of matter. They must provide a way of being copiously produced in the early stages of the Universe life, so they fill the above-mentioned $\\sim$23\\% of the Universe contents. Neutrinos are the only standard particles fitting in that scheme, but the hypothesis of neutrinos being the sole component of dark matter fails to reproduce part of the cosmological observations, in particular the current structure of the universe. The dark matter problem is therefore solved only by going into models beyond the standard model of elementary particles, among which two generic categories emerge as the best motivated for the task: WIMPs and axions.  WIMP is a generic denomination for any Weakly Interacting Massive Particle. A typical example of WIMP is the lightest supersymmetric particle (LSP) of SUSY extensions of the standard model, usually the neutralino. They would have been thermally produced after the Big Bang, cooled down and then frozen out of equilibrium providing a relic density\\cite{Primack:1988zm,Jungman:1995df}. The interesting mass window for the WIMPs spans from a few GeV up to the $\\sim$ TeV scale, but can be further constrained for specific models and considerations.  Axions, on the contrary, are light pseudoscalar particles that are introduced in extensions of the Standard Model including the Peccei-Quinn symmetry as a solution to the strong CP problem\\cite{Peccei:1977hh}. This symmetry is spontaneously broken at some unknown scale $f_a$, and the axion is the associated pseudo-Goldstone boson \\cite{Weinberg:1977ma,Wilczek:1977pj}. The axion framework provides several ways for them to be produced copiously in the early stages of the Universe, which makes it a leading candidate to also solve the dark matter problem\\cite{Raffelt:1990yz,Turner:1989vc}.  The hypothesis of axions or WIMPs composing partially or totally the missing matter of the Universe is specially appealing because it comes as an additional bonus to what these particles were originally thought for, i.e. they are not designed to solve the dark matter problem, but they may solve it. In addition, the existence of WIMPs or axions could be at reach of the sensitivity of current or near future experiments, and this has triggered a very important experimental activity in the last years. The detecion of WIMPs by direct means, i. e. aiming at their direct interaction with terrestrial detectors is being pursued by a large variety of techniques. All of them, however, share the feature that they are carried out in underground sites, to reduce backgrounds induced by cosmic rays. In fact, WIMP experiments are, together with double beta decay, one of the main pillars of \\emph{underground physics}. In the following pages a review is given of the current experimental efforts to directly detect these particles. Axions, on the other hand, are not necessarily searched in underground labs (although some axion detection techniques do use underground setups), and because of this they are sometimes dropped out of dark matter reviews. The fact that the \"dark matter community\" is substantially polarized towards WIMPs (probably due to the boost of underground physics in the last years) should not be regarded as a justification to ignore or minimize the importance of the axion as a potential candidate for dark matter. Much on the contrary, for many the existence of the axion is better motivated by theory arguments other than being a good dark matter candidate. Axions are in fact being searched for in experiments \\emph{per se}, without the assumption of them composing the dark matter, something that does not happen with WIMPs (excepting supersymmetry searches in accelerators). To compensate a bit this prejudice, and in spite of this review being focused on underground physics, we will devote a section on axion searches at the end of it.   Finally, indirect methods, like those looking for the decay products of WIMPs or axions in astronomical or cosmic rays observations may also put constraints on the properties of these particles\\cite{Fornengo:2006yy,Hooper:2007vy}, although they suffer from extra degrees of uncertainties, like the phenomenology driving the accumulation of dark matter particles in astrophysical bodies and their decay into other particles, and they are left out of the scope of the present review. ",
        "conclusions": "A review of the current status of the experimental searches for WIMPs and axions has been given. The field lives a moment of great activity, triggered by the fact that very well motivated theoretical candidates could be within reach of present technologies. The next years will witness the results of many very interesting developments currently ongoing to define and operate a new generation of experiments, well into the region of interest for both WIMPs and axions."
    },
    "0911/0911.4489.txt": {
        "abstract": "{A simple theory of supersymmetric dark matter (DM) naturally linked to neutrino flavour physics is studied. The DM sector comprises a spectrum of mixed lhd-rhd sneutrino states where both the sneutrino flavour structure and mass splittings are determined by the associated neutrino masses and mixings.  Prospects for indirect detection from solar capture are good due to a large sneutrino-nucleon cross-section afforded by the inelastic splitting (solar capture limits exclude an explanation of DAMA/LIBRA).  We find parameter regions where all heavier states will have decayed, leaving only one flavour mixture of sneutrino as the candidate DM.  Such regions have a unique `smoking gun' signature---sneutrino annihilation in the Sun produces a pair of neutrino mass eigenstates free from vacuum oscillations, with the potential for detection at neutrino telescopes through the observation of a hard spectrum of $\\nu_\\mu$ and $\\nu_\\tau$ (for a normal neutrino hierarchy).  Next generation direct detection experiments can explore much of the parameter space through both elastic and inelastic scattering.  We show in detail that the observed neutrino masses and mixings can arise as a consequence of supersymmetry breaking effects in the sneutrino DM sector, consistent with all experimental constraints. } ",
        "introduction": "The substantial observational evidence for non-baryonic dark matter (DM) is one of the clearest calls for new physics beyond-the-Standard-Model, specifically the existence of an exotic stable or very-long-lived massive particle or particles.   All this evidence, however, is indirect in form, and to-date does not fix the nature of the DM particle beyond requiring it to be colour and electromagnetically neutral and relatively weakly self-interacting. Despite this ignorance, one particular candidate has tended, until recently, to dominate theoretical speculation---the neutralino of the minimal supersymmetric (SUSY) extension of the Standard Model, the MSSM.   Although the neutralino is the lightest supersymmetric particle (LSP) over substantial regions of parameter space, and thus stable if $R$-parity conservation is assumed, the increasing reach of both direct DM searches and collider constraints on the MSSM now pushes MSSM neutralino DM into uncomfortable corners, as has been widely recognized.  The problems are two-fold: First, direct search experiments probe WIMP-nucleon cross sections orders of magnitude smaller than canonical weak scale cross sections; and second, given the direct search and collider constraints, the standard freeze-out mechanism of dark-matter-genesis generically generates a substantial excess of MSSM neutralino DM, limiting the allowed neutralino parameter space to special `fine-tuned' regions of parameter space with coannihilations or resonances.  Overall, neutralino DM matter works much less naturally than it did in the early 1990's after the excitement of the LEP-I results pointing to SUSY unification.  An independent reason to possibly doubt the standard MSSM neutralino story is that it is based on an assumption of simplicity of the dark sector that is probably not warranted given our experience of normal matter.  In the observable sector there is an extraordinary richness of stable or very long-lived massive states, including more than one hundred essentially stable nuclear isotopes (on the timescale of the Hubble time), a stable charged lepton, and three essentially stable neutrino mass eigenstates of various flavour compositions.  This is despite the fact that all global discrete and continuous symmetries such as individual lepton number, as well as total lepton and baryon numbers are violated (by at least neutrino mixing plus the electroweak anomaly, and likely GUT-suppressed interactions too which violate $B-L$ also).   There is no reason why we should expect the DM sector to be any less complicated than the observed world, and in a sense the observation of massive neutrinos already directly supports this hypothesis---the WIMP dark sector is already composed of at least a four-component cocktail of the three massive neutrino species plus one more WIMP.  Motivated by these arguments we, in this paper, investigate the physics of a very slight modification of the usual SUSY LSP hypothesis that naturally and simply accommodates a much more structured and varied DM sector (and in fact one that has attractive and potentially testable connections to neutrino flavour physics).  Specifically we study what we consider to be one of the simplest alternatives to standard SUSY neutralino DM, though maintaining the SUSY framework that naturally possesses the well-known successes of weak-scale SUSY theories, such as precision gauge coupling unification and a solution to the hierarchy problem, namely, mixed lhd-rhd sneutrinos.  Sneutrinos are potentially interesting alternate SUSY DM candidates since, unlike Majorana neutralinos, they can carry flavour quantum numbers while being charge and colour neutral, In addition they are the LSP in reasonably large regions of parameter space. \\footnote{Another attractive feature of considering such sneutrino DM is that it might afford a dynamical explanation of the baryon-to-dark-matter ratio $\\Omega_b/\\Omega_{DM} \\simeq 1/5$ as the DM can now possess a lepton number asymmetry connected to the baryon-number-asymmetry, in contradistinction to Majorana neutralino DM matter which can carry no such asymmetry.  We will not pursue these ideas in this work, but see \\cite{Kaplan:1991ah,Hooper:2004dc,Farrar:2004qy,Kitano:2004sv,Cosme:2005sb, Suematsu:2005zc,Abel:2006hr, Gopalakrishna:2006kr,McDonald:2006if,Page:2007sh,Deppisch:2008bp,Kaplan:2009ag,Kribs:2009fy,Cohen:2009fz} for studies along these lines.} The traditional reason why sneutrinos have not been considered a good DM candidate is that pure lhd, ie, $SU(2)_L$-active, sneutrinos have too large an annihilation cross section, and thus too small a freeze-out relic density to be the observed DM.   However this problem is easily solved if one posits the existence of weak-scale rhd, ie, $SU(2)_L$-sterile, sneutrinos which mix with the lhd states once electroweak symmetry and SUSY is broken.    Because of the observed family replication it is most natural to posit three rhd sneutrinos, one associated with each lepton flavour, in addition to the three standard lhd sneutrinos.      Therefore we are led to a model with six weak-scale sneutrinos, that mix with each other and form, as we will argue in detail, a structured DM sector carrying lepton flavour quantum numbers identical to that carried by the neutrino mass eigenstates observed in neutrino oscillation experiments.   In addition, each of these complex scalars splits into CP-odd and CP-even components, so the full spectrum of DM states is thus twelve real degrees of freedom with a variety of masses.  Shortening the story of the subsequent sections, what we find is a successful model comprising an entire sector of mixed sneutrino DM with the novel feature that the sneutrino flavour structure is identical to that of the neutrinos.  Further to this the sneutrinos have mass splittings between scalar and pseudoscalar components that are proportional to the neutrino masses. (See Figure \\ref{MassSpec}.) \\begin{figure}[h] \\centering \\includegraphics[height=3.5in]{spectrumdiag.pdf}\\caption{The spectrum of sneutrino masses after lifting of scalar and pseudoscalar mass degeneracy by the mass term $M_B^2$.  The red, yellow and blue lines represent, respectively, the electron, muon and tau lepton number flavour components, which mirrors the neutrino flavour structure (here we have assumed the normal neutrino mass hierarchy). Note that the 12 real degrees of freedom split into a heavy sector and a light sector of 6 states each, with further fine structure splittings among these two groups, and the splittings between scalar and pseudoscalar components are proportional to the neutrino masses.  This diagram is not to scale.} \\label{MassSpec} \\end{figure} In addition to these interesting features we find regions of parameter space where the splittings are large enough that all states will have decayed down, leaving only one flavour of sneutrino as the candidate DM.  These regions of parameter space potentially have a unique `smoking gun' signature as pair annihilations of these sneutrinos through neutralino exchange produce just one neutrino mass eigenstate $\\nu_3$ (for the normal neutrino mass hierarchy), free from vacuum oscillations, that has the potential for detection at future neutrino telescopes through the observation of a hard spectrum of $\\nu_\\mu$ and $\\nu_\\tau$ but not $\\nu_e$.  These neutrinos would have energy equal to the sneutrino mass and could arise from annihilations in the Sun.  Since we require rhd sneutrinos at the weak scale our model also possesses sterile rhd neutrinos at this scale or below.   Naively one may therefore think that, as the rhd neutrinos are not superheavy, it is impossible to generate, eg, by the see-saw mechanism, an acceptable spectrum of light neutrinos as observed in neutrino oscillation experiments.  However as originally show by Arkani-Hamed \\etal \\cite{ArkaniHamed:2000bq, ArkaniHamed:2000kj} and Borzumati \\etal \\cite{Borzumati:2000mc, Borzumati:2000ya} and amplified on and extended by later authors (see, eg, \\cite{TuckerSmith:2001hy, TuckerSmith:2004jv, MarchRussell:2004uf, Hambye:2004jf, West:2004me}) this is not the case.  In fact there occurs a new and extremely attractive mechanism of neutrino mass generation linked to supersymmetry breaking, in a sense generalizing the Giudice-Masiero mechanism for generating the $\\mu$-term---the Higgsino mass term---in the MSSM.  This model of neutrino masses allows for new explanations of the origin of the neutrino mass and flavour structure \\cite{ArkaniHamed:2000bq, ArkaniHamed:2000kj,MarchRussell:2004uf,West:2004me} and elegantly accommodates such nice features as weak-scale resonant leptogenesis \\cite{Hambye:2004jf}.  Since our main focus in this work is the novel sneutrino DM phenomenology we only very quickly summarize the physics of neutrino mass generation in as far as it impacts the sneutrino sector, and we strongly encourage the reader to refer to these other papers for more details.  Before we start, one other feature that deserves discussion is the fact that this class of models was the original and motivating example of inelastic DM \\cite{TuckerSmith:2001hy, TuckerSmith:2004jv}, providing a possible explanation of the reported DAMA/LIBRA \\cite{Bernabei:2008yi} signals consistent with the exclusions reported by other direct DM detection experiments. Since these original models involving sneutrinos, the idea of inelastic DM has been implemented in many more set ups, and has been much studied (see, eg, \\cite{Cui:2009xq} and references therein). This inelastic scattering also changes the phenomenology of DM solar capture and we include the exclusion limits from Refs. \\cite{Nussinov:2009ft, Menon:2009qj} on the nucleon scattering cross-section.  These can be more constraining than the current generation of direct detection experiments over large regions of parameter space.  There are many previous studies of sneutrino DM, see eg, \\cite{Hall:1997ah,Kolb:1999wx,Asaka:2005cn,Asaka:2006fs,Lee:2007mt,Arina:2007tm,Thomas:2007bu,Arina:2008yh,Arina:2008bb,Cerdeno:2008ep,Allahverdi:2009ae,Cerdeno:2009dv,Demir:2009kc,Allahverdi:2009se,Cerdeno:2009zz,Allahverdi:2009kr,Kumar:2009sf}, though most do not consider connecting the sneutrino and neutrino flavour structure. (For treatments of sneutrino DM taking limits from the generation of neutrino masses we refer the reader to \\cite{Arina:2007tm,Thomas:2007bu,Kumar:2009sf}.)  The previous work most closely related to that considered here is Ref.\\cite{Kumar:2009sf} where the flavour structure of the sneutrinos is also considered.  However our model has the novel feature that the neutrino and sneutrino flavour structures are exactly the same, allowing a predictive `smoking gun' signature for sneutrino annihilations into neutrinos, as discussed in Section~\\ref{indirect_detection}.  Another important difference is that in our model there is an interplay between two terms contributing to the neutrino mass generation (see the off-diagonal entries of eqs.(\\ref{massmat2})) that allow us to fit the observed neutrino masses consistent with successful sneutrino DM over a significant region of parameter space. Finally, our analysis includes the important limits arising from solar capture, which excludes the possibility of an inelastic DM explanation of DAMA/LIBRA using mixed sneutrinos.  In Section~\\ref{model} we first discuss the specific model, outlining the generation of neutrino masses and the origin of the identical neutrino and sneutrino flavour structure, as well as the resulting mass spectrum of the sneutrino states. We then go on in Section~\\ref{dark_matter} to present the decay lifetimes for the different flavoured sneutrino states and consider the DM phenomenology of the sneutrinos, showing the regions of parameter space where they are good DM candidates.  In Section~\\ref{detection} the potential unique signatures of this model are considered, while our conclusions are contained in Section~\\ref{conclusions}. Two short appendices contain technical details. ",
        "conclusions": "The model detailed herein constitutes a viable model of both DM and neutrino mass generation in which the DM is comprised of mixed lhd-rhd sneutrinos in the same flavour mixtures as the neutrino mass eigenstates.   We have shown in detail that it is possible that the observed neutrino masses and mixings arise as a consequence of supersymmetry breaking effects in the sneutrino sector, consistent with all experimental constraints.   The prospects for indirect detection of the associate sneutrino DM from solar capture are good, owing to the large sneutrino-nucleon cross-section afforded by the inelastic splitting of Z-coupled states as shown in Figures \\ref{Limits}, \\ref{ml200&bf}, and \\ref{VaryDe}.  For some regions of the allowed parameter space a potential `smoking gun' signature of this model would be monochromatic, non-oscillating, neutrinos comprised of just one mass eigenstate which originate from sneutrino annihilation in the Sun, see Figure \\ref{spectrum}.  The current generation of direct detection experiments place strong limits at low sneutrino masses and at high mass if the inelastic splitting is small, and next generation experiments can explore much of the parameter space, see Figures \\ref{Limits}, \\ref{ml200&bf}, and \\ref{VaryDe}.  There is a prospect of seeing both elastic and inelastic scattering in these experiments, allowing the size of the sneutrino splitting $\\delta_{\\alpha, L}$ to be directly measured.   Despite being one of the earliest examples of inelastic DM, limits from solar capture completely exclude this model as an explanation of DAMA/LIBRA.   The mixed lhd-rhd sneutrinos can have signatures at the LHC, and can change aspects of Higgs physics. Finally, we have shown that a simple extension of the MSSM leads to a much richer structure in the DM sector linking different areas of beyond-the-Standard-Model physics, namely dark matter and neutrino physics."
    },
    "0911/0911.2252_arXiv.txt": {
        "abstract": "We measure the spatial clustering of galaxies as a function of their morphological type at $z\\simeq0.8$, for the first time in a deep redshift survey with full morphological information. This is obtained by combining high-resolution HST imaging and VLT spectroscopy for about $8,500$ galaxies to $I_{AB}=22.5$ with accurate spectroscopic redshifts from the zCOSMOS-Bright redshift survey. At this epoch, early-type galaxies already show a significantly stronger clustering than late-type galaxies on all probed scales. A comparison to the SDSS at $z\\simeq0.1$, shows that the relative clustering strength between early and late morphological classes tends to increase with cosmic time at small separations, while on large scales it shows no significant evolution since $z\\simeq0.8$. This suggests that most early-type galaxies had already formed in intermediate and dense environments at this epoch. Our results are consistent with a picture in which the relative clustering of different morphological types between $z\\simeq1$ and $z\\simeq0$, reflects the evolving role of environment in the morphological transformation of galaxies, on top of the global mass-driven evolution. ",
        "introduction": "In the local Universe, the clustering properties of galaxies depend on luminosity \\citep[e.g.][]{davis88,hamilton88,white88,park94,loveday95,benoist96, guzzo00,norberg01,zehavi02,zehavi05,li06,skibba06,wang07,swanson08}, colour \\citep[e.g.][]{willmer98,zehavi02,zehavi05,li06,swanson08,skibba09c}, spectral type \\citep[e.g.][]{loveday99,norberg02,madgwick03,wang07}, and environment \\citep{abbas06,abbas07}. Differences in the clustering of the various galaxy morphological types are also observed \\citep[e.g.][]{davis76,giovanelli86,iovino93,loveday95,hermit96,guzzo97, willmer98,zehavi02,skibba09}. Elliptical galaxies are more strongly clustered than spiral and irregular galaxies. Another manifestation of this phenomenon is the existence of a morphology-density relation \\citep{dressler80,postman84}, which implies that a higher fraction of ellipticals than either spirals or irregulars, reside in denser environments. Spiral and irregular galaxies are more likely to populate less dense regions. However, the origin and relation between these dependences are still not fully understood. They are usually discussed in terms of a \\emph{nature} or \\emph{nurture} scenario, i.e. being intrinsic properties of galaxies at formation or originated from the interaction of galaxies with their environment all along their cosmic evolution.  At higher redshifts, our knowledge of galaxy clustering properties remains fragmentary. We know little about the clustering of the various morphological types. Deep surveys have detected an evolution in the global galaxy clustering with redshift up to $z\\sim1-1.5$ \\citep{coil04,lefevre05,mccracken08}. Additional results indicate that the luminosity, colour, and spectral type dependences of clustering evolve with redshift. A less significant segregation is evident when we consider the dependence of clustering on either colour or spectral type, the difference in clustering strength between red/early-type and blue/late-type galaxies being shallower at $z\\sim1-1.5$ than at $z\\simeq0$ \\citep{meneux06,coil08}.  These observations are consistent with the accepted picture of hierarchical clustering growth and galaxy evolution, in which cold dark matter haloes form from the gravitational collapse of dark matter around peaks in the density field. Haloes evolve hierarchically, such that smaller haloes assemble to form larger and more massive haloes in high-density regions \\citep{mo96,sheth02}. In parallel, galaxies form within haloes, by means of the cooling of hot baryonic gas \\citep{white87}. In this framework, luminous and massive galaxies are expected to be more strongly clustered than fainter and less massive ones that tend to form in less clustered haloes, which are less biased with respect to the underlying mass distribution. Moreover, early-type galaxies, which are in general brighter than late-type galaxies, are more strongly clustered. This difference is thought to be related to their different formation histories. Early-type galaxies are understood to be the product of interaction processes (e.g. galaxy merging) and/or physical mechanisms that take place in dense environments, although mergers are more common in less dense group environments than at galaxy cluster cores \\citep{ellison10}. Instead, late-type galaxies may not experience such dramatic events, being more likely to reside in less dense environments. Mergers are expected to affect galaxy morphology by creating or growing a galaxy bulge component, albeit a large gas fraction may result in a galaxy temporarily re-growing a disk. They contribute to the evolving number densities \\citep[e.g.][]{deravel09} and clustering properties of different types of galaxies.  Galaxy colours or spectral types are often used as proxies for galaxy morphologies, but are affected by uncertainties in tracing the same underlying galaxy populations across cosmic time. For instance, a disk galaxy containing an old stellar population would be classified as an early-type object. In a similar way, a starburst galaxy heavily obscured by dust would be classified as a red, early-type galaxy. In contrast, an old and red stellar population may dominate the measured stellar mass of an object, but the galaxy colour can be blue because of a recent burst of star formation. Colours are primarily related to galaxy recent star formation history at variance with morphologies, which may be the result of the interaction of galaxies with their surroundings. Selecting galaxies by colour or spectral type, or by morphology may give one different views on galaxy evolution and environmental effects at work. Studying how galaxy clustering depends on morphology and cosmic epoch may thus provide important and complementary clues to mechanisms through which galaxies developed into the population we see today. Such a study has not yet been possible at $z>0.3$ because of the lack of high resolution imaging surveys of sufficient size. The Cosmic Evolution Survey \\citep[COSMOS,][]{scoville07} provides us with a unique sample in this respect.  In this work, we use the zCOSMOS-Bright spectroscopic sample of galaxies \\citep{lilly09} and measure for the first time the morphological dependence of galaxy clustering at $z\\simeq0.8$.  This paper is part of a series of three using the first epoch zCOSMOS-Bright sample, that investigate galaxy clustering as a function of galaxy physical properties: morphology (this paper), luminosity and stellar mass \\citep{meneux09}, and colour (Porciani et al., in preparation).  The paper is organised as follows. In Sect. 2, we present the galaxy sample and its basic properties. In Sect. 3, we describe the method used to measure galaxy clustering. In Sect. 4, we provide the measurements of early- and late-type galaxy clustering. In Sect. 5, we summarise and discuss our results. Throughout this paper, we assume a flat $\\Lambda CDM$ cosmology with $\\Omega_M=0.25$, $\\Omega_\\Lambda=0.75$, and $H_0=100~\\rm{h~km \\cdot s^{-1} \\cdot Mpc^{-1}}$. The magnitudes are quoted in the AB system and for simplicity, we denote the absolute magnitude in $B$-band $M_B-5\\log{\\rm (h)}$ as $M_B$. ",
        "conclusions": "We have measured the dependence of clustering on morphology at $0.6<z<1.0$ for galaxies of luminosity greater than $L^*$.  For this purpose, we have used the first epoch zCOSMOS-Bright spectroscopic sample of galaxies, our study benefiting from the availability of high resolution HST imaging for the COSMOS field. We have computed the projected correlations function in volume-limited samples of two broad morphological classes, early types (elliptical/S0) and late types (spiral/irregular) and have compared them to $z\\simeq0.1$ SDSS measurements. Our two main results can be summarised as follows: \\begin{enumerate} \\item We find that at $z\\simeq0.8$, early-type galaxies exhibit a stronger clustering strength than late-type galaxies on scales from $0.1$\\hmpc to $10$\\hmpc. \\item Comparing our results to those for the SDSS for galaxies with comparable luminosities, shows that while the relative difference in clustering between early and late morphological classes seems to increase with cosmic time on scales smaller than a few \\hmpc, the large-scale difference does not evolve significantly since $z\\simeq0.8$. This indicates that a large fraction of early-type galaxies were already formed in intermediate and dense environments at this epoch. \\end{enumerate}  The observed difference in shape of the correlation function that we observe for early- and late-type galaxies, can be interpreted within the framework of halo occupation distribution models \\citep[e.g.][]{cooray02}. In these models, the galaxy correlation function is the sum of two contributions, one dominating the small scales that characterises the clustering of galaxies inside haloes (1-halo term), and a large-scale contribution, which characterises the clustering of galaxies belonging to different haloes (2-halo term). The prominence of the 1-halo term observed for early-type galaxies, i.e. the enhancement of the observed correlation function on scales smaller than or of the order of the typical halo radius ($1-2$\\hmpc), implies that these galaxies are on average hosted by massive haloes with larger virial radii in dense environments \\citep{abbas06}. The weaker but steeper 1-halo term of late-type galaxies indicates that these galaxies instead may be hosted on average by less massive haloes inhabiting low-density environments. This trend is qualitatively consistent with $z\\simeq0$ measurements of the clustering of red and blue galaxies, that dominate respectively the early- and late-type populations \\citep[e.g.][]{zehavi05,skibba09c}.  Our work is complementary to those of \\citet{tasca09} and \\citet{kovac10} who study within the same sample, the evolution of the fraction of early and late morphological types with environment, as defined either by the continuous overdensity field or group/field environments, respectively. The measurements presented here are consistent with the picture emerging from these works. In particular, \\citet{tasca09} observed a flattening in the morphology-density relation with increasing redshift at fixed luminosity, extending to morphology similar evolutionary trends observed in the colour- and spectral-type-density relations \\citep{cucciati06,cooper06,grutzbauch10}. The finding that early- and late-type galaxies tend to have increasingly similar probabilities of populating high-density/group and low-density/field environments with increasing redshift, is qualitatively consistent with the evolution in the morphological dependence of clustering discussed here, where we find a corresponding decrease in the relative small-scale clustering of early- to late-type galaxies. The mass and environmental physical processes play an important role in shaping the morphology-density relation and its evolution with cosmic time \\citep[e.g.][]{kovac10}. It is found that mass has a dominant contribution at $z\\simeq1$ in particular for relatively bright luminosity-selected sample \\citep{tasca09}. In contrast, in the local Universe, one finds that only a part of this relation can be attributed to the variation in the stellar-mass function with environment, with the dense environment-related processes becoming more important \\citep{bamford09}. Our results corroborate this evolutionary picture. While the large-scale shape of the relative bias has remained constant since $z\\simeq1$, the small-scale shape exhibits a significant evolution. This is due to an increase in the number of small-scale early-type pairs with cosmic time, as a result of that of the relative contribution of environmental physical processes in transforming early- to late-type galaxies, especially in dense environments."
    },
    "0911/0911.5644_arXiv.txt": {
        "abstract": "Gamma-ray bursts (GRBs) have been regarded as standard candles at very high redshift for cosmology research. We have proposed a new method to calibrate GRB distance indicators with Type Ia supernova (SNe Ia) data in a completely cosmology-independent way to avoid the circularity problem that had limited the direct use of GRBs to probe cosmology [N. Liang, W. K. Xiao, Y. Liu, and S. N. Zhang, Astrophys. J. 685, 354 (2008).]. In this paper, a simple method is provided to combine GRB data into the joint observational data analysis to constrain cosmological models; in this method those SNe Ia data points used for calibrating the GRB data are not used to avoid any correlation between them. We find that the $\\Lambda$CDM model is consistent with the joint data in the 1-$\\sigma$ confidence region, using the GRB data at high redshift calibrated with the interpolating method, the Constitution set of SNe Ia, the cosmic microwave background radiation from Wilkinson Microwave Anisotropy Probe five year observation, the baryonic acoustic oscillation from the spectroscopic Sloan Digital Sky Survey Data Release 7 galaxy sample, the x-ray baryon mass fraction in clusters of galaxies, and the observational Hubble parameter versus redshift data. Comparing to the joint constraints with GRBs and without GRBs, we find that the contribution of GRBs to the joint cosmological constraints is a slight shift in the confidence regions of cosmological parameters to better enclose the $\\Lambda$CDM model. Finally, we reconstruct the acceleration history of the Universe up to $z>6$ with the distance moduli of SNe Ia and GRBs and find some features that deviate from the $\\Lambda$CDM model and seem to favor oscillatory cosmology models; however further investigations are needed to better understand the situation. ",
        "introduction": "Gamma-ray bursts (GRBs) are the most intense explosions observed so far and likely to occur in high redshift range. Recently, GRB 090423 has been  observed at a very high redshift, $z>8$ \\cite{GRB0904}. The early universe can be explored by using GRBs at the high redshift which is hardly achievable by Type Ia supernova (SNe Ia). In recent years, several GRB luminosity relations between measurable properties of the prompt gamma-ray emission with the luminosity or energy have been proposed as distance indicators, many of which are the two-dimensional (2D) luminosity relations, such as the isotropic energy ($E_{\\rm iso}$) - peak spectral energy ($E_{\\rm peak}$) relation (i.e., the so-called Amati relation) \\cite{Amati}, the luminosity ($L$) - spectral lag ($\\tau_{\\rm lag}$) relation \\cite{Nor00}, the $L$ - variability ($V$) relation \\cite{VL}, the $L$ - $E_{\\rm peak}$ relation \\cite{EpL}, the $L$ - minimum rise time ($\\tau_{\\rm RT}$) relation \\cite{Sch02}, and the collimation-corrected energy ($E_{\\gamma}$) - $E_{\\rm peak}$ relation (i.e., the so-called Ghirlanda relation) with small scattering \\cite{GGL04}. In order to reduce the large scattering in some of these 2D relations, several 3D luminosity relations are also presented \\cite{LZ05,Fir06,Yu09}. Liang and Zhang first proposed the relation among $E_{\\rm iso}$, $E_{\\rm peak}$ and $t_{\\rm b}$ (i.e., the so-called Liang-Zhang relation) \\cite{LZ05}, where $t_{\\rm b}$ is the break time of the optical afterglow light curves; and Firmani et al. found the relation among $L$, $E_{\\rm peak}$ and $T_{0.45}$ (i.e., the so-called Firmani relation) \\cite{Fir06}, where $T_{0.45}$ is the ``high-signal'' timescale of the prompt emission. Both of the 3D luminosity relations above could reduce the scattering in the luminosity or energy \\emph{versus} the $E_{\\rm peak}$ relation. More recently, Yu et al. find that, for the 3D luminosity relations between the luminosity and an energy scale $E_{\\rm peak}$ and a timescale ($\\tau_{\\rm lag}$ or $\\tau_{\\rm RT}$), the intrinsic scattering is considerably smaller than that of those corresponding 2D luminosity relations \\cite{Yu09}. Other GRB luminosity relations have been proposed in many works \\cite{Hakkila08,Willingale07,Tsutsui09J}. For reviews of GRB luminosity relations, see e.g. \\cite{zhang07,SC07,GGF06,Sch07}.  However, all the luminosity relations of GRBs presented above are always empirical, but still without solid physical interpretations \\cite{GGF06}. Moreover, the luminosity relations can not be calibrated with a sufficiently large low-redshift GRB sample. Therefore previous calibrations have been usually obtained by assuming a particular cosmological model. Once these GRB luminosity relations are calibrated, GRB data could be considered as standard candles at very high redshift for cosmology research \\cite{GGF06,Sch07,Sch03,Bloom03,Dai04,GGLF04,LZ05,Xu05,Fir05,Fir06b,Fir07,Friedman05,Wang06,Cuesta08,BP08,Wri07,Wang07,Gong07,Qi08,Taka03,Berto06} (see e.g. \\cite{GGF06} and \\cite{Sch07} for reviews). In 2003, Schaefer derived the luminosity distances of nine GRBs with known redshifts by using two GRB luminosity relations ($L$ - $\\tau_{\\rm lag}$, $L$ - $V$) to construct the first GRB Hubble diagram \\cite{Sch03}. For 16 GRBs with redshift measurements, Bloom et al. found a narrow clustering of geometrically-corrected Gamma-ray energies within the framework of the uniform conical jet model \\cite{Bloom03}. Dai et al.  proposed an approach to consider the Ghirlanda relation to constrain cosmological parameters and dark energy \\cite{Dai04}. Because of the small scatter, the 3D luminosity relations have been used for cosmological constraint, including the Liang-Zhang relation \\cite{LZ05}, and the Firmani relation \\cite{Fir06b,Fir07}. In 2007, Schaefer used five 2D relations ($L$ - $\\tau_{\\rm lag}$, $L$ - $V$, $L$ - $E_{\\rm peak}$, $L$ - $\\tau_{\\rm RT}$ and $E_{\\gamma}$ - $E_{\\rm peak}$) to construct the Hubble diagram of 69 GRBs \\cite{Sch07}. These GRB data have been widely used to constrain cosmology and dark energy models in many recent works \\cite{Cuesta08,BP08}, and joint constraints by combining these GRB data with SNe Ia and the other cosmological probes have been derived in \\cite{Wri07,Wang07,Gong07,Qi08}. Instead of a hybrid sample over the whole redshift range of GRBs, Takahashi et al. \\cite{Taka03} firstly calibrated GRB relations ($L$ - $\\tau_{\\rm lag}$, $L$ - $V$) at low redshift where distance-redshift relations have been already determined from SNe Ia. This method was adopted by Bertolami and Silva for considering the use of GRBs at $1.5<z<5$ calibrated with the bursts at $z\\le1.5$ as distance markers to study the generalized Chaplygin gas model \\cite{Berto06}. All the calibration methods above carried out have been derived usually from the $\\Lambda$CDM model with particular model parameters according to the concordance cosmology.   An important point related to the use of GRBs for cosmology is the dependence on the cosmological model in the calibration of GRB relations, which had limited the direct use of GRBs for cosmology research. In order to investigate cosmology, the relations of standard candles should be calibrated in a cosmology-independent way; otherwise, the circularity problem cannot be avoided easily \\cite{GGF06}. Recently, the possibility of calibrating GRBs in a low-dispersion in redshift near a fiducial redshift has been proposed \\cite{Lamb05,Ghir06}, which has been developed further based on Bayesian theory \\cite{LZ06}. However, the GRB sample available now is far from what is needed to calibrate the relations in this way \\cite{GGF06,LZ06}. Many of the works treat the circularity problem with statistical approaches \\cite{Sch03,Fir05,Li08,Amati08,WangY08}. A simultaneous fit of the parameters in the calibration curves and the cosmology is carried out to find the optimal GRB relation and the optimal cosmological model in the sense of a minimum scattering in both the luminosity relations and the Hubble diagram \\cite{Sch03}. Firmani et al. proposed a Bayesian method to get around the circularity problem \\cite{Fir05}. Li et al. presented a global fitting analysis for the Ghirlanda relation to deal with the problem \\cite{Li08}. Amati et al. used the $E_{\\rm iso}$ - $E_{\\rm peak}$ relation  to measure the cosmological parameter by adopting a maximum likelihood approach \\cite{Amati08}. More recently, Wang has shown that the current GRB data can be summarized by a set of model independent distance measurements, with negligible loss of information \\cite{WangY08}, which is followed by \\cite{Qi09,SR09,WangY09}. However, an input cosmological model is still required in doing the joint fitting, the circularity problem can not be circumvented completely by means of these statistical approaches.  In our previous paper \\cite{liang08a}, we presented a new method to calibrate several GRB luminosity relations in a completely cosmology-independent manner. Our method avoids the circularity problem more clearly than previous cosmology-dependent calibration methods. It is obvious that objects at the same redshift should have the same luminosity distance in any cosmology. Therefore, the luminosity distance at any redshift in the redshift range of SNe Ia can be obtained by interpolating (or by other mathematical approach) directly from the SNe Ia Hubble diagram. Using the interpolation method, we calibrated seven GRB luminosity relations ($L$ - $\\tau_{\\rm lag}$, $L$ - $V$, $L$ - $E_{\\rm peak}$, $L$ - $\\tau_{\\rm RT}$, $E_{\\gamma}$ - $E_{\\rm peak}$, $E_{\\textrm{iso}}$ - $E_{\\rm peak}$, and $E_{\\rm iso}$-$E_{\\rm peak}$-$t_{\\rm b}$) with the data compiled by Schaefer \\cite{Sch07}. Then if further assuming these calibrated GRB relations valid for all long GRB data, we can use the standard Hubble diagram method to constrain the cosmological parameters from the GRB data at high redshift obtained by utilizing the relations \\cite{liang08a}. These distance data of GRBs are so far the most cosmology-independent GRB distance indicators. Following this cosmology-independent GRB calibration method from SN Ia \\cite{liang08a}, the derived GRB data at high redshift range can be used to constrain cosmological models without circularity problem \\cite{liang08a,Capozziello08,Izzo09,wei09a,wei09b,WLiang09}. Capozziello and Izzo first used the $E_{\\rm iso}$-$E_{\\rm peak}$-$t_{\\rm b}$ relation and the $E_{\\gamma}$ - $E_{\\rm peak}$ relation calibrated with the so-called Liang method to derive the related cosmography parameters [deceleration parameters ($q$), jerk parameters ($j$), and snap parameters ($s$) \\cite{Visser04}], which are only related to the derivatives of the scale factor  without any a priori assumption \\cite{Capozziello08}. Wei and Zhang used the Amati relation calibrated with the interpolation method to reconstruct the acceleration history of the Universe and constrain cosmological models \\cite{wei09a,wei09b}.  Furthermore, we proposed another approach to calibrate GRB relations by using an iterative procedure \\cite{liang08b}, which is a nonparametric method in a model independent manner to reconstruct the luminosity distance at any redshift in the redshift range of SNe Ia \\cite{Shafieloo06,Shafieloo07,Wu08}. Similar to the interpolation method, Cardone et al. constructed an updated GRBs Hubble diagram on six 2D relations calibrated by local regression from SNe Ia \\cite{CCD09}. This method and GRB data have been used to cosmological constraints with SNe Ia in some following works \\cite{Qi09,CDC09}. Kodama et al. presented that the $L$ - $E_{\\rm peak}$ relation can be calibrated with the empirical formula fitted from the luminosity distance of SNe Ia \\cite{kod08}. This method has been used to constrain cosmological parameters by combining these GRB data with SNe Ia in a following work \\cite{Tsu09a}. However, it is noted that this calibration procedure depends seriously on the choice of the formula and various possible formulas can be fitted from the SNe Ia data that could give different calibration results of GRBs. As the cosmological constraints from GRBs are sensitive to GRBs calibration results \\cite{WangY08}, the reliability of this method should be tested carefully. Moreover, as pointed out in \\cite{WangY08}, the GRB luminosity relations which are calibrated by this way are no longer completely independent of all the SNe Ia data points; therefore these GRB data can not be used to directly combine with the whole SNe Ia dataset to constrain cosmological parameters and dark energy.  It is easy to find that the number of SNe Ia data points that have been used in the linear interpolating procedure to obtain the GRB data is relatively small compared to the whole SNe Ia sample. In order to combine GRB data obtained by our interpolation method into the joint observational data analysis to constrain cosmological models, we can exclude the SNe points that have been used in the interpolating procedure from the SNe Ia sample used to the joint constraints. In this work, with the updated distance moduli of the 42 GRBs at $z>1.4$ obtained by the interpolating method \\cite{liang08a} from the Union set of 307 SNe Ia \\cite{sn307}, we provide a simple method to ensure that the calibrated GRB data are independent of the SNe Ia data used in the joint data fitting to constrain cosmological models, by using the Constitution set of 397 SNe Ia  \\cite{sn397}, along with the cosmic microwave background radiation (CMB) observation  from the five-year Wilkinson Microwave Anisotropy Probe (WMAP5) result \\cite{wmap5}, ($R$, $l_a$, $\\Omega_bh^2$); the baryonic acoustic oscillation (BAO) \\cite{bao} observation from the spectroscopic Sloan Digital Sky Survey (SDSS) Data Release 7 (DR7) galaxy sample \\cite{SDSS7} ($d_{0.2}$, $d_{0.35}$); the 26 baryon mass fraction in clusters of galaxies (CBF) from the x-ray gas observation ($f_{\\textrm{gas}}$) \\cite{Allen04}; and the 11 Hubble parameter versus redshift data, $H(z)$ \\cite{Simon05,Gazta08}. Our goal is to determine the contribution of GRBs to the joint cosmological constraints in the confidence regions of cosmological parameters by comparing to the joint constraints with GRBs and without GRBs. Finally we also reconstruct the acceleration history of the Universe up to $z>6$ with the distance moduli of SNe Ia and GRBs. ",
        "conclusions": "Because of the lack of enough low red-shift GRBs to calibrate the luminosity relation, GRBs could not be used reliably and extensively in cosmology until recently. In our previous paper, we have proposed a new method to calibrate GRB luminosity relations in a completely cosmology-independent manner to avoid the well-known circularity problem \\cite{liang08a}.  In this work, with the recent GRB data at high redshift whose distance moduli are calibrated with the interpolating method \\cite{liang08a} from the Union set of 307 SNe Ia \\cite{sn307}, as well as the Constitution set of SNe Ia \\cite{sn397},the CMB observation  from the WMAP5 result ($R$, $l_a$, $\\Omega_bh^2$) \\cite{wmap5}, the BAO observation from the spectroscopic SDSS DR7 galaxy sample ($d_{0.2}$, $d_{0.35}$) \\cite{SDSS7} , the x-ray baryon mass fraction in clusters of galaxies \\cite{Allen04}, and the observational $H(z)$ data \\cite{Simon05,Gazta08}, we find that the $\\Lambda$CDM model is consistent with the joint data in 1-$\\sigma$ confidence region; this confirms the conclusion of many previous investigations. We also find that the current GRB data are substantially less accurate than the SNe Ia data, if each data set is used alone, consistent with previous investigations.  The new results and insights we have obtained in this work are briefly summarized as follows:  (1) In order to combine GRB data into the joint observational data analysis to constrain cosmological models, we provide a simple method to avoid any correlation between the SNe Ia data and the GRB data; in this method those SNe Ia data points used for calibrating the GRB data are not used.  (2) Comparing to the joint constraints with GRBs and without GRBs, we find that the contribution of GRBs to the joint cosmological constraints is a slight shift in the confidence regions of cosmological parameters to better enclose the $\\Lambda$CDM model.  (3) Finally we reconstruct the Hubble parameter $H(z)$, the deceleration parameter $q(z)$, and the \\textsl{EoS} of dark energy $w(z)$ of the acceleration Universe up to $z>6$ with the distance moduli of the Constitution set of SNe and GRBs and find some features that seem to favor oscillatory cosmology models; however further investigations are needed to better understand the situation.  For considering the use of GRBs for cosmology, the gravitational lensing effect may need to be considered \\cite{Oguri06}. However Schaefer found that the gravitational biases of GRBs are small \\cite{Sch07}. Recently, some possible observational selection bias \\cite{But07,Fiore07,Campana07,Sha09} and evolution effects \\cite{Li07,Tsu08,Salva08,Petro09} in GRB relations have also been discussed. However,  Ghirlanda et al. confirmed the spectral-energy relations of GRBs  observed by \\emph{Swift} \\cite{Ghi07}. Moreover, it is found  that instrumental selection effects do not dominate for the Amati relation \\cite{Ghi08} and for the $E_{\\rm peak}-L$ relation \\cite{Ghi09}, as well as no strong evolution with redshift of the Amati relation can be found  \\cite{Ghi08,WDQ08}. Nevertheless, for considering GRBs as standard candles to constrain cosmology, further examinations of possible selection bias and evolution effects should be required in a larger GRB sample.  Different dark energy models may have very different Hubble diagrams at high redshifts \\cite{Sch07}. Therefore, the best plan to investigate the property of dark energy is measuring the dark energy over a wide range of redshifts. GRBs can extend the Hubble diagram to much higher redshifts beyond SNe Ia data \\cite{Sch07}. It is worth noticing that GRBs are important potential probes for cosmic history up to $z>6$. However, the contribution of GRBs to the cosmological constraints would not be sufficiently significant at present, due to the large statistical scatters of the relations and the small dataset of GRBs. As discussed in \\cite{Sch07}, GRBs are almost immune to dust extinction, whereas in the case of SNe Ia observations, there is extinction from the interstellar medium when optical photons propagate towards us. On the other hand, SNe Ia are substantially more accurate standard candles than GRBs, which will lead to tight constraints on cosmological parameters; whereas a single GRB at high redshift will provide more information than a single maximal redshift SNe Ia \\cite{Sch07}. Through Monte Carlo simulations, future prospect of probing dark energy parameters with a larger sample of GRBs has been investigated \\cite{Taka03,Xu05,LZ05,Tsu09a}. It has been found that cosmological constraints would improve substantially with more simulated GRBs expected by future observations. Recently, Xiao and Schaefer have compiled 107 long GRBs with their spectroscopic/photometric redshifts measured \\cite{XiaoSchaefer09}, observed by BATSE, Konus, HETE, and \\textit{Swift}. Along with more and more GRBs observed from \\emph{Fermi Gamma-ray Space Telescope} (formerly known as \\emph{GLAST}) with much smaller scatters, and its combination with the increasing \\emph{Swift} data, GRBs could be used as an additional choice to set tighter constraints on cosmological parameters of dark energy models.    \\appendix"
    },
    "0911/0911.3368_arXiv.txt": {
        "abstract": "We study the clustering properties of the first galaxies formed in the Universe. We find that, due to chemical enrichment of the inter-stellar medium by isolated Population III stars formed in mini-halos at redshift $z \\gtrsim 30$, the (chronologically) first galaxies are composed of metal-poor Population II stars and are highly clustered  on small scales. In contrast, chemically pristine galaxies in halos with mass $M\\sim 10^8 \\mathrm{M_{\\sun}}$ may form  at $z<20$ in relatively underdense regions of the Universe. This occurs once self-enrichment by Population III in mini-halos is quenched by the build-up of an $H_2$ photo-dissociating radiative background in the Lyman-Werner bands. We find that these chemically pristine galaxies are spatially uncorrelated. Thus,   we expect that deep fields with the James Webb Space Telescope may detect clusters of chemically enriched galaxies but individual chemically pristine objects.  We predict that  metal-free galaxies at $10 \\lesssim z \\lesssim 15$ have surface densities of about 80 per square arcmin and per unit redshift but most of them will be too faint even for JWST. However, the predicted density makes these objects interesting targets for searches behind lensing clusters. ",
        "introduction": "The concept of first galaxies is not as easy to define as that of first stars and it's partly a matter of convention. A first, somewhat arbitrary, decision, is the mass scale where an object becomes a galaxy. A common convention in the context of the first galaxies is describing galaxies as those objects forming within dark halos massive enough that their gaseous content can cool by atomic hydrogen (Ly$\\alpha$ cooling), i.e. characterized by a virial temperature greater than $10^4$ K \\citep{ricotti08,WiseAbel08,Greifetal08}. The advantage of this definition is that it is based on a physical process and makes galaxy formation relatively independent of the presence of a Lyman-Werner (LW) background that could hamper the formation of Population III stars in mini-halos by photo-dissociating molecular hydrogen \\citep{lepp83,tegmark97,haiman97,machacek03,ts09,trenti09b}.  Once we have defined ``galaxy'' we need to define ``first''.  The usual assumption is that ``first'' refers to a chronological sequence of events, so that the first galaxies are those that formed at the earliest redshifts. It is sometimes assumed that the chronologically ``first'' galaxies will also be the least chemically evolved ones. In fact, these objects are most likely to form in highly clustered areas and they are also very likely to have been polluted by metals produced by Population III stars forming in mini-halos \\citep{WiseAbel08,Greifetal08}. Thus, these chronologically first galaxies would not have a primordial chemical composition. However, we can imagine another class of galaxies that would form after a LW background is in place and would then not be polluted by stars in mini-halos \\citep{jimenez06,trenti09b,Johnsonetal08}. These objects could conceivably form stars out of pristine primordial gas and would then be galaxies made of Population III stars. Assuming that they can form Population III stars in significant numbers, these objects would be chemically less evolved (i.e. ``younger'') despite forming later in a chronological sense.  This paper is devoted to exploring the validity of the two scenarios described above and to studying the redshift distributions and clustering properties of these two classes of objects. These quantities are relevant for planning surveys to study the first galaxies with the James Webb Space Telescope and other (future) facilities, such as thirty-meter class telescopes used for imaging behind lensing clusters.  The structure of the paper is the following. In Sec. 2 we discuss our model for galaxy formation during the Dark Ages, in Sec. 3 we present our results on star formation rates and clustering. Sec. 4 concludes with prospects for future detections and speculates on the possibility that metal-free galaxies have already been observed behind a lensing cluster by \\citet{starketal07}. ",
        "conclusions": "Fig.~\\ref{fig:tpcf} predicts that galaxies containing Population III stars will be essentially uncorrelated on scales larger than $\\sim150$ comoving kpc. At redshift $z=9.5$, 1 arcsec corresponds to roughly 50 comoving kpc in the transverse direction. This implies that the typical arcmin scales of the imaging instruments on the James Webb Space Telescope correspond to a few comoving Mpc and on those scales we should not detect any significant clustering for galaxies with primordial chemical composition. In contrast, enriched protogalaxies could be very clustered on small scales, as they live in high density regions. Thus, in an ultradeep field with JWST we expect to detect individual unevolved galaxies at $z\\simeq10$, but small clusters of chemically evolved objects.  From Fig.~\\ref{fig:first_gal} we see that the rest frame rate of galaxy formation per comoving volume at $z\\simeq10$ is $\\sim 2\\times 10^{-8}$ Mpc$^{-3}$ yr$^{-1}$ (estimating the halo formation rate from the simulations gives a similar number). Assuming that the burst of Population III stars is short lived, we expect that the galaxy will be visible as Population III galaxy only for a time of the order of the stellar lifetime (assumed to be $2\\times 10^6$ yrs). Considering that the comoving volume per unit redshift at $z\\sim 10$ is $\\sim2000$ comoving Mpc$^3$ over an area of one arcmin squared, we predict a surface density of $~\\sim 80$ galaxies per square arcmin  and per unit redshift. However, there are large uncertainties as to the luminosity such a galaxy would have. In a deep exposure, the James Webb Space Telescope should be able to detect a young cluster with a stellar mass of $10^6$ M$_\\odot$\\footnote{In contrast, typical Lyman Break Galaxies currently found at $z>7$ are older (stellar age $\\gtrsim 100$ Myr \\citealt{labbe10}) and reside in dark-matter halos with $M_{halo} \\gtrsim 2 \\times 10^{10} M_{\\sun}$}. Most of the pristine galaxies will have lower stellar masses. At a virial temperature $T = 10^4$ K at $z\\sim 10$ a galaxy has a halo mass of $\\sim3.7 \\times 10^7$ M$_\\odot$ and only $\\sim 6 \\times 10^6$ M$_\\odot$ in baryons. It is at this stage unclear whether the Population III stellar burst in a proto-galaxy would produce many stars and with what efficiency $\\epsilon$, but $\\epsilon \\gtrsim 0.15$ is unlikely in the first burst. Even if the majority of these objects will be too faint, their large number density and the expected lack of strong correlations are encouraging as they make it more likely that we could detect some of these objects through a lensing cluster. A Population III galaxy of $10^6$ M$_\\odot$ of stars formed over a stellar lifetime forms stars at about 0.3 M$_\\odot$ yr$^{-1}$ and would have a Ly$\\alpha$ luminosity of $\\sim3 \\times 10^{42}$ erg s$^{-1}$ \\citep{schaerer03} and a Ly$\\alpha$ line flux at $z=10$ of $\\sim 1.9 \\times 10^{-18}$ erg s$^{-1}$ cm$^{-2}$. These fluxes are close to what can be detected today in a lensing cluster after modest amplification and is in fact in agreement with some reports of detections \\citep{starketal07}.  The (chronologically) first galaxies at the same redshift and the extremely poor galaxies are instead expected to be highly correlated at least on scales of a few arcsec. Thus, to first order our models predict that one should observe small clusters of chemically evolved galaxies but only isolated primordial galaxies."
    },
    "0911/0911.3642_arXiv.txt": {
        "abstract": "We present \\textit{Suzaku} observations of the Galactic black hole candidate \\swt~in the low-hard state. The broadband coverage of Suzaku enables us to detect the source over the energy range 0.6 -- 250 keV.  The broadband spectrum (2 -- 250 keV) is found to be consistent with a simple power-law ($\\Gamma \\sim$ 1.63). In agreement with previous observations of this system, a significant excess of soft X-ray flux is detected consistent with the presence of a cool accretion disc. Estimates of the disc inner radius infer a value consistent with the ISCO ($\\rm R_{in} \\lesssim 6~R_g$, for certain values of, e.g. $\\rm N_H,~ i$), although we cannot conclusively rule out the presence of an accretion disc truncated at larger radii ($\\rm R_{in} \\sim 10 - 50~R_g$). A weak, relativistically-broadened iron line is also detected, in addition to disc reflection at higher energy. However, the iron-K line profile favours an inner radius larger than the ISCO (R$\\rm _{in} \\sim 10 - 20~R_g$). The implications of these observations for models of the accretion flow in the low-hard state are discussed. ",
        "introduction": "The study of accretion processes is of fundamental astrophysical importance, being ubiquitous on both the largest and smallest scales, i.e. ranging from accretion by supermassive black holes in the center of galaxies at one end to star and planet formation at the other.  Here we are interested in studying accretion phenomena in the presence of a large gravitational potential such as that observed in X-ray binaries (XRBs) and active galactic nuclei (AGN). In these sources, the fundamental timescales of interest, governing the accretion process, scale with the mass of the central object ($\\rm \\propto M_x$). Hence, studies of the Galactic XRB population provide an excellent laboratory for detailed examination of the process of accretion on humanly accessible timescales.  The very-high, high-soft and low-hard states (hereafter VHS, HSS \\& LHS) are the primary active accretion states observed in XRBs (See the review by \\citealt{I2} for a detailed description of accretion states in black hole binaries). Despite being the most common mode of accretion in black hole X-ray binaries, the nature of the accretion flow in the low-hard state remains uncertain. Emission in the LHS is characterized by a hard power-law spectrum ($\\Gamma \\sim$ 1.4 -- 1.7) and strong X-ray variability (30\\% -- 40\\% RMS). Correlated variability at radio/NIR/optical wavelengths has also been observed from the black hole systems while in the hard state \\citep{a15}. The hard power-law component may be the result of Comptonization by a hot optically thin plasma of soft seed photons from a thermal disc or magnetic structures (through cyclo-synchrotron processes - see, e.g., \\citealt{a35}). However, the geometry of this plasma is not well understood. A popular model for the accretion geometry in the low-hard state was given by \\citet{a36}. In that model the standard thin accretion disc, that dominates in the spectrally soft states is radially truncated and replaced by an advection dominated accretion flow (ADAF: see \\citealt{a16}). The fundamental assumption of a radially recessed accretion disc may find some support in the low disc reflection fractions which are sometimes measured in the hard state (e.g., \\citealt{a37}).  A number of models have been developed which do not require a recessed disc in the low-hard state. \\citet{a38} proposed that black hole states are driven by the height and bulk velocity of magnetic flares above a disc, which remains at the innermost stable circular orbit (ISCO). These flares would serve to feed a mildly relativistically outflowing corona. Low disc reflection fractions do not signal a recessed disc in this model, but result from mild beaming of the hard X-ray flux away from the disc. A similar model for the LHS, based on magnetically dominated coronae, has been proposed by \\citet{a39}. These outflowing coronae share some properties with jet based models for the hard component \\citep{a40}, in the sense that the latter also produce low reflection fractions \\citep{a25} without the need for a recessed accretion disc.  A number of recent observations call into question the assumption of a recessed disc in the LHS. In particular GX 339-4 was observed during its 2004 outburst at a luminosity of $\\rm L \\sim 0.05~L_{Edd}$ by both {\\it RXTE} \\& {\\it XMM} \\citep{a20}. The observations by {\\it XMM} are critical here as they provide coverage at energies below 3 keV whereas {\\it RXTE} is limited to energies above this. It was found that a cool disc blackbody ($\\sim$ 0.35 keV) consistent with an optically thick geometrically thin accretion disc extending to the innermost stable circular orbit (ISCO) existed, in contrast to theoretical expectations. Fits with reflection models revealed reflection fractions $\\sim$ 0.2 -- 0.3. Previous observations of Swift J1753.5-0127 yielded similar results for the cool disc component, in this case a disc temperature of $\\sim$ 0.2 keV was required \\citet{a18}; however, no significant disc reflection features were detected.  \\begin{figure} \\begin{center} \\includegraphics[height=0.33\\textheight,width=0.29\\textwidth,angle=-90]{fig1_colour.eps} \\caption{{\\it Swift} BAT hard X-ray lightcurve for \\swt~from May 2005 to November 2008. The times of the {\\it Integral} \\citep{a19}, {\\it XMM} \\citep{a18}, {\\it RXTE} \\citep{a76}, \\textit{Swift} (see \\S 4.1) and \\textit{Suzaku} (this paper) observations are indicted.} \\label{asmlc} \\end{center} \\end{figure}  \\swt~was discovered by the \\textit{Swift} burst alert telescope (BAT) at X-ray and $\\gamma$-ray energies on 2005 May 30 (\\citealt{a1}). Subsequent observations with the X-ray telescope (XRT) revealed a hard power-law spectrum \\citep{a2,a3}, while pointed \\textit{RXTE} observations detected 0.6 Hz quasi periodic oscillations (QPO; \\citealt{a7}). The system was also detected at UV \\citep{a8} and optical wavelengths \\citep{a4}. Radio observations with the \\textit{MERLIN} array also detected a variable counterpart \\citep{a5}. Observations at optical wavelengths by \\citet{a69} have also detected a significant modulation with a period of 3.2hrs, which they identify as a superhump period slightly larger than the actual orbital period. This would make \\swt~the black hole binary with the shortest known orbital period.  In addition to the observations of \\citet{a18} above, a number of other high energy studies of this source have been published, which we summarize below (see Fig. \\ref{asmlc}). An analysis of \\textit{RXTE} observations of the outburst was reported in \\citet{a6,a17}. The X-ray spectrum was found to be consistent with a power-law ($1.6 \\leq \\Gamma \\leq 1.8$), while low-frequency quasi-periodic oscillations (QPOs) were also detected (up to 0.9 Hz).  \\citet{a19} analysed simultaneous \\textit{RXTE} \\& \\textit{INTEGRAL} data, which were also obtained during the 2005 outburst. The combined spectrum (3 -- 400 keV) could be fit with a model consisting of thermal Comptonization modified by disc reflection ($\\rm kT_0 \\sim 0.5~keV, kT_e \\sim 150~keV, \\tau \\sim 1, f \\equiv \\Omega/2\\pi \\sim 0.3$). QPOs were also detected here, but at a lower frequency than during the outburst peak (0.24 Hz).  To place meaningful constraints on the accretion flow in the LHS one requires sensitivity to both the soft X-rays ($<$ 2 keV), in order to detect the cool accretion disc, and higher energies in order to detect the most prominent disc reflection features ( 5 -- 7 keV and 20 -- 30 keV). \\textit{Suzaku} with its large bandpass and low background, is ideally equipped to carry out these observations.  In this paper, we describe observations undertaken with the \\textit{Suzaku} X-ray observatory in 2007, while \\swt~was in the low-hard state.  In \\S2, we describe our observation and extraction of source spectra and lightcurves. We proceed to analyze the data in \\S3, where both phenomenological and more physically motivated models are considered. The broadband spectrum (2 -- 250 keV) is consistent with a simple power-law model, although there are also contributions from the accretion disc. In \\S4, these results are discussed in the context of models for the accretion flow in the low-hard state, and finally our conclusions are presented in \\S5. ",
        "conclusions": "We have presented \\textit{Suzaku} broadband X-ray observations of the candidate black hole X-ray binary \\swt. The broadband spectrum (2 -- 250 keV) is observed to be consistent with a simple power-law model ($\\Gamma \\sim 1.63$). Confirming previous observations, we detect the presence of an excess at soft X-ray energies. In addition, a weak relativistic iron line and curvature consistent with a reflection hump at 20 -- 30 keV are detected.  These observations point towards the persistence of the accretion disc at a much lower radius than previously appreciated in the low-hard state. Estimates of the disc inner radius with both a simple {\\tt diskbb+po} model and a detailed reflection model reveal values consistent with the ISCO ($\\rm R_{in} \\lesssim 6~R_g$) for certain values of both the column density and inclination. In contrast modelling the spectra with a Comptonization model, while revealing a truncated inner disc ($\\rm R_g \\sim 50~R_g$), implies a value for the Hydrogen column density in disagreement with all previous estimates."
    },
    "0911/0911.4853_arXiv.txt": {
        "abstract": "We have collected  continuum data of  a sample of D-type symbiotic stars. By modelling their spectral energy distribution in a colliding-wind theoretical scenario we have found  the common characteristics to all the systems: 1) at least two dust shells are clearly present, one  at  $\\sim$ 1000 K and the other at $\\sim$ 400 K; they dominate the emission in the IR; 2) the radio data are explained by  thermal self-absorbed emission from the reverse shock between the stars; while 3)  the data in  the  long wavelength tail come from the expanding shock outwards the system; 4) in some  symbiotic stars, the contribution from the WD in the UV is directly seen. Finally, 5) for some objects soft X-ray emitted by bremsstrahlung  downstream of the reverse-shock between the stars are predicted. The results thus confirm the validity of the colliding wind model and the important role of the shocks. The comparison of the fluxes calculated at the nebula with those observed at Earth   reveals the  distribution throughout the system of the different components, in particular the  nebulae and the dust shells. The correlation of  shell radii with the orbital period shows that larger radii are found at larger periods. Moreover, the temperatures of the dust shells regarding the sample are  found  at $\\sim$ 1000 K and $\\leq$ 400 K, while, in the case of late giants, they spread more uniformly throughout the same range. ",
        "introduction": "\\begin{table*} \\centering \\caption{The sample in the literature. \\label{tab:lit}} \\begin{footnotesize} \\begin{tabular}{c c c c c c c c c}\\\\ \\hline \\hline Object & Coordinates & distance & $P_{orb}$ & $P_{Mira}^\\textit{a}$ & \\.{M}$_{Mira}$ & $T_{Mira}$ & $T_{dust}$ \\\\ & $\\alpha \\; [J2000] \\; \\delta$ &[kpc] & [years]& [days]& [M$_{\\odot}$/yr] & [K] & [K] \\\\ \\hline BI Cru & 12:23:27.4 $\\;$ -62:38:16.5 & 2 & - & 280 & $3 \\, 10^{-6}$& - & 1300\\\\ SS73 38 & 12:51:26.2 $\\;$ -64:59:58.8 &4.8& - & 463 & $2.5 \\, 10^{-6}$ & - & -\\\\ V835 Cen & 14:14:09 $\\;$ -63:25:45.3 & 2.8 & - & 450 & - & 2250 & 1000\\\\ H1-36 & 17:49:48.2 $\\;$ -37:01:27.0 & 4.5 & $>$95 & 450-500 & $10^{-5} d^{3/2}$ & 2500 & 700 \\\\ HM Sge & 19:41:57.1 $\\;$ 16:44:39.6 & 2.3 & 90$\\pm$20 & 527 & $\\approx10^{-5}$ & 3000 & 1400\\\\ V1016 Cyg & 19:57:05 $\\;$ 39:49:36.0 & 2.93$\\pm$0.75 & 80$\\pm$25 & 474 & $\\geq 3 \\, 10^{-7}$ & 2450$\\pm$150 & 600 \\\\ RR Tel & 20:04:18.52 $\\;$ -55:43:33.4 & 2.5 &-& 387 & $1.6\\, 10^{-6}$ & 2300 & 1000 \\\\ V627 Cas & 22:57:42 $\\;$ 58:49:14.0 & $<$0.8 & - & 466 & - & 2650$\\pm$50 & 1000$\\pm$50 \\\\ R Aqr & 23:43:49.4 $\\;$ -15:17:40.3 & 0.2 & 44 & 395 & $10^{-6}$ & 2800 & 1000 \\\\ \\hline \\end{tabular} \\begin{scriptsize} \\flushleft References: \\textit{a}: Belczy{\\' n}ski et al. (2000); BI Cru: Schwarz \\& Corradi (1992), Contini et al. (2008c); SS73 38: Munari (1992), Gromadzki et al. (2009); V835 Cen: Whitelock (1987), Feast et al. (1983); H1-36: Allen (1983), Angeloni et al. (2007b); HM Sge: Whitelock (1987), Richards et al. (1999), Bogdanov \\& Taranova (2001), Schild et al. (2001); V1016 Cyg: Parimucha (2003), Schild \\& Schmid (1996), Lee et al. (2003), Taranova \\& Shenavrin (2000); RR Tel: Feast et al. (1983), Kotnik-Karuza et al. (2002), Gromadzki et al. (2009); V627 Cas: Bergner et al. (1988); R Aqr: Gromadzki et al. (2009); Hollis et al. (1997), Contini \\& Formiggini (2003), Dougherty et al. (1995). \\end{scriptsize} \\end{footnotesize} \\end{table*}  {\\it The word \"symbiotic\" was introduced by Merrill in 1944 and this term is currently used for the category of variable stars with composite spectrum. The  main spectral features of these objects are: 1) the presence of a red continuum typical of a cool star , 2) the rich emission line spectrum, and 3) the UV excess. [...] In addition to the peculiar spectrum, the very irregular photometric and spectroscopic variability is the major feature of symbiotic stars}. So Viotti in 1993 introduced the symbiotic phenomenon.  Since then, a large number of symbiotic stars (hereafter SSs) was observed. In 1995 we started the analysis of SSs by interpreting the line spectra of RS Ophiuci (Contini et al. 1995)  by means of shock-dominated models. At that time, it  was established that both the WD and red giant stars loose winds which collide within and outwards the system (Nussbaumer 2000 and references therein). Therefore, to model the SS systems on a large scale, we adopted the colliding-wind scenario of Girard \\& Willson (1987) which leads to different shock fronts throughout the SS. The nebulae downstream are  illuminated by the flux from the  stars, therefore, a composite model (shock + photoionization) was needed, and successfully applied to HM Sge (Formiggini, Contini, Leibowitz 1995).  In certain objects accretion disks are formed  at certain epochs. They blow up during the outburst of the hot component star. The presence of a  disk  is revealed by the jets that were observed and modelled, for instance, in BI Cru (Contini et al 2009b), He2-104 (Contini \\& Formiggini 2001), R Aqr (Contini \\& Formiggini 2003),  CH Cyg (Contini et al. 2009c), etc. The line ratios observed in the  jet regions  show lower densities ($\\sim$ 200 \\cm3)  and velocities (FWHM  of 150-200 \\kms) than those corresponding to the nebulae created by the wind collision ($\\geq$ 10$^7$ \\cm3 and $\\leq$ 1000 \\kms, respectively). Therefore, also the continuum flux intensity is lower and  does not significantly affect the SED.  We realized that the results of the line spectrum analysis  of SSs should be used to constrain the spectral energy distribution (SED) of the continuum. The observations in the different wavelength ranges provided sufficiently reliable   data. Therefore, not only the red giant star and some times even the WD  could be recognized throughout the SED, but today the plural nature of the emitting nebulae and of the dust shells  within the SS systems is well established (Contini 1997; Contini \\& Formiggini 2001, 2003; Angeloni et al. 2007a, hereafter Paper I; Angeloni et al. 2007b,c and references therein; Contini et al. 2009a,b,c and references therein).  Moreover, it is difficult to explain the observed IR SED of dusty SSs by a single temperature, and nowadays theoretical models and observations too have been confirming that several dust temperatures should be combined in order to reproduce the NIR-MIR data (Anandarao et al. 1988, Schild et al. 2001, Paper I). This will be definitively demonstrated  by   further data coming from the MIR-FIR windows that ISO, Spitzer, and AKARI have opened.  The  present analysis  focuses on a sample of D-type (i.e. dusty) SSs, producing a detailed picture of each system. The SS continua are   analysed in a  quantitative  way mainly through the light-curve variability,  through  the cool and hot star spectra and through the emission line spectra in the optical and UV range. In particular, by modelling the data in the continuum, we could clearly recognise in  each  object the relative  contribution of the dust shells and of the  nebulae. The shells are ejected as a consequence of pulsations of the Mira, while the nebulae  appear  downstream of the shock fronts created by the collision of the star winds.  As mentioned above, the calculation of the spectra  must account consistently for shocks and photoionization; therefore, we use for the calculation of the spectra the code SUMA\\footnote{http://wise-obs.tau.ac.il/$\\sim$marcel/suma/index.htm}, which simulates the physical conditions of an emitting gaseous nebula under the coupled effect of photoionization from an external source and shocks, and in which line and continuum emission from gas are calculated consistently with dust reprocessed radiation as well as with grain heating and sputtering. SUMA  has been  applied to interpret several SSs, as extensively reported in previous papers (e.g. Angeloni et al. 2007b; Contini et al. 2009b and references therein), as well as nova stars (V1974, Contini et al. 1997; T Pyx, Contini \\& Prialnik 1997) and supernova remnants (e.g. Kepler's SNR, Contini 2004).   We present the  D-type SS sample  in Table 1. We gathered for each object the observational data from the literature and the  results  obtained by modelling the continuum SED,  which were constrained by the analysis of the line spectra in each object.  Our  aim is to characterize the SED of D-type SS; therefore we chose our targets on the basis of the best  data available, in order to create a  SED on a large wavelength range for each object. This has allowed to better constrain several physical parameters essential for a meaningful and comprehensive interpretation, taking  into account some issues, e.g.  the effects of intrinsic flux variability on our modelling method. In this way, we are able to compare the  features common to the dusty SSs. Moreover  the correlations of the characteristic physical properties such as luminosity, orbital period, dust shell sizes, etc.,  is  discussed on the basis of reliable parameters.  In Sect. 2 the data are presented and discussed in the light of the photometric and spectroscopic variability. Sect. 3 deals with theoretical aspects: the colliding wind scenario, the calculation code, and the modelling method. In Sect. 4 the SEDs  are analysed. The results are shown for the single SS in Sect. 5. Discussion and conclusion remarks follow in Sect. 6. ",
        "conclusions": "We have collected continuum data in the IR for a sample of D-type SSs. For most of the objects the detailed modelling of the line spectra and the continuum SED in previous specific works leads to an accurate description of the physical characteristics and of the morphological structure throughout each system. For the other objects we have guessed the constraining parameters from a   first analysis of the line ratios and/or by similarity with other well-known objects.  The study of dust grains  and, in general, of dust shell properties in SSs is interesting not only in itself, but also in the light of SSs as colliding-wind systems, namely systems where shocks are at work. We thus have the unique opportunity of studying in a relatively small regions both dust condensation and destruction processes, as well as the mutual interaction of grains with a very dynamic ambient gas.  By modelling the SED continuum for each object, we have found  the  common characteristics to all the systems: 1) at least two dust shells are clearly present at about 1000 K and 400 K, 2) the radio data are explained by  thermal self-absorbed emission from the reverse shock between the stars, while 3)  the data in  the  long wavelength tail come from the expanding shock outwards the system. 4) In some system the contribution from the WD in the UV can be directly seen. Finally, 5) soft X-ray produced by bremsstrahlung from the reverse-shock nebula between the stars, are predicted for some objects.  The results thus confirm the validity of the colliding wind model and the important role of the shocks. The dust shells which were so important in explaining some features in the light curve at different phases of the  CH Cyg (Contini et al. 2009c) are well recognizable in all the D-type objects. We wonder whether their presence will have an impact on the IR light curve at certain phases.  The flux peak of the 1000 K dust shells are inversely correlated with the square distance to Earth. This permitted  to guess out the distance of SS73 38 and to select the appropriate distance of V835 Cen.  By consistent modelling of the dust shell at 1000 K in R Aqr, we have found that the dust-to-gas ratio by mass is $\\geq$ 0.1, even higher than that evaluated for supernovae.  The comparison of the flux calculated at the nebula with those observed at Earth   reveals the  distribution throughout the system of the different components, both nebulae and dust shells. The correlation of  shell radii with the orbital period show that  larger radii are found at larger periods. Moreover, the temperatures of the dust shells regarding the sample, are  generally  concentrated at $\\sim$ 1000 K and $\\sim$ 400 K, while in the case of late giants, they spread more uniformly throughout the same range because the WD radiation most probably  leads to evaporation of the grains in its proximity."
    },
    "0911/0911.1701_arXiv.txt": {
        "abstract": "The accretion disk around a compact object is a nonlinear general relativistic system involving magnetohydrodynamics. Naturally the question arises whether such a system is chaotic (deterministic) or stochastic (random) which might be related to the associated transport properties whose origin is still not confirmed. Earlier, the black hole system GRS~1915+105 was shown to be low dimensional chaos in certain temporal classes. However, so far such nonlinear phenomena have not been studied fairly well for neutron stars which are unique for their magnetosphere and kHz quasi-periodic oscillation (QPO). On the other hand, it was argued that the QPO is a result of nonlinear magnetohydrodynamic effects in accretion disks. If a neutron star exhibits chaotic signature, then what is the chaotic/correlation dimension? We analyze RXTE/PCA data of neutron stars Sco~X-1 and Cyg~X-2, along with the black hole Cyg~X-1 and the unknown source Cyg~X-3, and show that while Sco~X-1 and Cyg~X-2 are low dimensional chaotic systems, Cyg~X-1 and Cyg~X-3 are stochastic sources. Based on our analysis, we argue that Cyg~X-3 may be a black hole. ",
        "introduction": "X-ray binary systems vary on timescales ranging from months to milli-seconds (see, e.g., \\cite{chen, paul97, nowak, cui, gleis, axel0}). Detailed analysis of their temporal variability and fluctuation provides important insights into the geometry and physics of emitting regions and the accretion process. However, the origin of variability is still not clear. It could be due to varying external parameters, like the infalling mass accretion rate. It could also be due to possible instabilities in the inner regions of the accretion disk where the flow is expected to be nonlinear and turbulent. Uttley et al. (2005) (see also Timmer et al. 2000 and Thiel et al. 2001) argued that the non-linear behavior of a system can be understood from the log-normal distribution of the fluxes and the $rms$-flux relation. This implies that the temporal behavior of the system may be driven by underlying stochastic variations. By studying the underlying nonlinear behavior, important constraints can be obtained on these various possibilities.  An elegant way of obtaining the constraint is to perform the nonlinear time series analysis of observed data and to compute the correlation dimension $D_2$ in a non-subjective manner. This technique has already been used to diverse situations (\\citep{grass0,grass,schr,aber,skb,bmc,misra,hari}, and references therein). By obtaining $D_2$ as a function of the embedding dimension $M$, one can infer the origin of the variability. For example, $D_2\\approx M$ for all $M$ corresponds to the system having stochastic fluctuation which favors the idea that X-ray variations are driven by variations of some external parameters. On the other hand, a saturated $D_2$ to a finite (low) value, beyond a certain M, implies a deterministic chaos which argues in favor of inner disk instability. However, to implement the algorithm successfully, the system in question should provide enough data.  The technique was used earlier to understand the nonlinear nature of a black hole system Cyg~X-1 \\citep{unno} and an Active Galactic Nucleus (AGN) Ark~564 \\citep{agn}, but due to insufficient data points the analyses were hampered and no concrete conclusions were made about $D_2$. Later on, another black hole system GRS~1915+105 was analyzed \\cite{bmc,misra,hari} which was shown to display low dimensional chaos in certain temporal classes, while stochastic in other classes.  However, so far none of the neutron star systems have been analyzed in detail in order to understand the origin of nonlinearity. Decades back, Voges et al. (1987) attempted to understand the chaotic nature of Her~X-1, but the analysis was hampered by low signal to noise ratio \\cite{noris}. Since then the investigation of chaotic signature in neutron stars remains unattended. Can a neutron star system not be deterministic? Indeed several features of X-ray binary systems consisting of a neutron star, such as their magnetosphere and kHz Quasi-Periodic Oscillation (QPO) and its possible relation to the spin frequency of the neutron star, favor the idea that they exhibit nonlinear resonance (e.g. \\cite{blaes,bm}). While the QPO itself is a mysterious feature whose origin is still unclear, its possible link to the spin frequency of the neutron star\\footnote{However, some authors (Mendez \\& Belloni 2007) suggested that the kHz QPOs may not be related to the spin.} indicates the origin of QPO to be from nonlinear phenomena. Several LMXBs having a neutron star exhibit twin kHz QPOs \\cite{mendez, vander}. For some of them, e.g. 4U~1636-53 \\cite{jonker}, KS~1731-260 \\cite{smith}, 4U~1702-429 \\cite{markwardt}, 4U~1728-34 \\cite{vans}, the spin frequency of the neutron star has been predicted from observed data. However, for the source Sco~X-1, which exhibits noticeable time variability \\cite{menvan}, while we observe twin kHz QPOs, we do not know the spin frequency yet (but see \\cite{bm}). For another neutron star Cyg~X-2, we do observe kHz QPOs \\cite{wijnands} as well. Several black holes also exhibit QPOs, e.g. GRS~1915+105 \\cite{belloniqpo, mcclintock}, Cyg~X-1 \\cite{angelini}.  In the present paper, we first aim at analyzing the time series of two neutron star sources Sco~X-1 and Cyg~X-2 to understand if a neutron star is a deterministic nonlinear (chaotic) system. Then we try to manifest the knowledge of nonlinear (chaotic/random) property of compact sources to distinguish a black hole from a neutron star. Subsequently, knowing their difference based on the said property, we try to identify the nature of a unknown source (whether it is a black hole or a neutron star). While the nature of some sources, as mentioned above, has already been predicted based on alternate method, for some others, e.g. Cyg~X-3, SS433, it has not yet been confirmed.  For the present purpose, we therefore concentrate on three additional sources Cyg~X-1, Cyg~X-2 and Cyg~X-3. While Cyg~X-1 has been predicted to be a black hole and Cyg~X-2 be a neutron star, the nature of Cyg~X-3 is not confirmed yet. Some authors \\cite{ergma98,schmutz96,szostek08} argued for a black hole nature of Cyg X-3, on the basis of its jet, the time variations in the infrared emission lines, the BeppoSAX X-ray spectra and so on. However, earlier it was argued for a neutron star \\cite{chadwick85} by measuring its $1000$ GeV $\\gamma$-rays which suggests a pulsar period of $12.5908\\pm 0.0003$ ms. By analyzing the time series and computing the correlation dimension $D_2$, here we aim at pinpointing the nature of Cyg~X-3: whether a black hole or a neutron star.  In the next section, we briefly outline the procedure to be followed in understanding the nonlinear nature of a compact object from observed data and to implement it to analyze the neutron star source Sco~X-1. In \\S3, we then describe nonlinear behaviors of Cyg~X-1, Cyg~X-2 and Cyg~X-3. Subsequently, in \\S4, we compare all the results and argue for a black hole nature of Cyg~X-3. Finally, we summarize in \\S5. ",
        "conclusions": "The source Cyg~X-3, whose nature is not confirmed yet, seems to be a black hole based on the analysis of its nonlinear behavior. On the other hand, we have shown, for the first time to best of our knowledge, that neutron star systems could be chaotic in nature. The signature of deterministic chaos, which argues in favor of inner disk instability, into an accreting system has implications in understanding its transport properties particularly in Keplerian accretion disk \\cite{win}. Note that in Keplerian accretion disks transport is necessarily due to turbulence in absence of significant molecular viscosity. The signature of chaos confirms instability and then plausible turbulence. On the other hand, for a rotating neutron star having a magnetosphere, signature of chaos suggests their QPOs to be nonlinear resonance phenomena \\cite{bm}. The absence of chaos and related/plausible signature of instability in Cyg X-1 and Cyg X-3 suggests the underlying accretion disk to be sub-Keplerian \\cite{ny95,chak96} in nature which is dominated significantly by gravitational force."
    },
    "0911/0911.0531_arXiv.txt": {
        "abstract": "\\noindent Young stellar systems are known to undergo outbursts, where the star experiences an increased accretion rate, and the system's luminosity increases accordingly.  The archetype is the FU Orionis (FU Ori) outburst, where the accretion rate can increase by three orders of magnitude (and the brightness of the system by five magnitudes).  The cause appears to be instability in the circumstellar disc, but there is currently some debate as to the nature of this instability (e.g. thermal, gravitational, magneto-rotational).  This paper details high resolution Smoothed Particle Hydrodynamics (SPH) simulations that were carried out to investigate the influence of stellar encounters on disc dynamics.  Star-star encounters (where the primary has a self-gravitating, marginally stable protostellar disc) were simulated with various orbital parameters to investigate the resulting disc structure and dynamics.  Crucially, the simulations include the effects of radiative transfer to realistically model the resulting thermodynamics.  Our results show that the accretion history and luminosity of the system during the encounter displays many of the features of outburst phenomena.  In particular, the magnitudes and decay times seen are comparable to those of FU Ori.  There are two caveats to this assertion: the first is that these events are not expected to occur frequently enough to explain all FU Ori or EX Lupi; the second is that the inner discs of these simulations are subject to numerical viscosity, which will act to reduce the accretion rate (although it has less of an effect on the total mass accreted).  In short, these results cannot rule out binary interactions as a potential source of some FU Ori-esque outbursts. ",
        "introduction": "The traditional picture of star formation describes the free-fall collapse of a protostellar molecular cloud (e.g. \\citealt{Shu_et_al_87}).  Conservation of angular momentum will in general ensure that, within typical free-fall times of \\(\\sim 10^5\\) yr, the collapse produces a protostar with protostellar disc.  These formation rates are consistent with the observed statistics of protostellar objects, for example those in Taurus \\citep{Kenyon_et_al_90} and with numerical simulations \\citep{Bate_cluster_03,Stam_frag,Bate_09}.  If these formation timescales are converted to average mass accretion rates, then it appears that the standard picture will form a star at the rate of \\(10^{-5} M_{\\odot}\\, yr^{-1}\\), whereas current observations suggests an average infall rate of \\(10^{-6} M_{\\odot}\\, yr^{-1}\\) or lower.  This leads to the realisation that accretion rates in protostars are not constant, which has been confirmed observationally \\citep{Herbig_77, Armitage_et_al_01,Zhu_et_al_09}, and that short periods of increased accretion, (accompanied by periods of mass pile-up where the infalling matter is not accreted) can solve the apparent inconsistencies between observation and theory.  This also provides an explanation for the FU Orionis (FU Ori) outburst objects, Class 0 to Class II protostellar objects which undergo a characteristic rapid rise in luminosity spanning up to five magnitudes, and then decay on timescales of a few hundred years \\citep{Hartmann_Kenyon_96}.  Calculated accretion rates for FU Ori objects are typically \\(10^{-4} M_{\\odot}\\, yr^{-1}\\) or more \\citep{Herbig_77,Hartmann_Kenyon_96}, which show a strong increase over typical infall rates \\citep{Kenyon_et_al_93,Furlan_et_al_08}.  The precise origin of these FU Ori outbursts is not known, although almost all theories involve a protostellar disc, as these discs are expected to be present around the majority of early-type stars.  The theories range from thermal instability \\citep{Lin_Papa_85,Bell_and_Lin}, gravitational instability \\citep{Vorobyov_Basu_05,Vorobyov_Basu_06,Vorobyov_Basu_08,Boley_and_Durisen_09}, a combination of gravitational instability and magnetorotational instability \\citep{Armitage_et_al_01,Zhu_et_al_09}, even the presence of a planet controlling accretion onto the central star \\citep{Clarke_Syer_96,Lodato_Clarke_04}.  All these theories agree that triggering a disc instability event is crucial to providing the observed accretion rates that cause the outburst.  The aim of this paper is not to identify the correct theory of FU Ori objects, but to comment on what was once considered a potential cause of FU Ori: the encounter of a primary star plus protostellar disc with a discless secondary \\citep{Kenyon_et_al_88,Bonnell_Bastien_92}.  Observations are beginning to reveal that outburst phenomena appear to belong to different classes, e.g. FU Ori (FUors), EX Lupi (EXors) \\citep{Herbig_07}, and others.  This work proposes that stellar encounters could produce outbursts that share many observational features with FUors and EXors (as was suggested by \\citet{Pfalzner_08}, combining treecode simulations with cluster dynamics), and yet may belong to a different class.   Results are presented from a series of high resolution Smoothed Particle Hydrodynamics (SPH) simulations with radiative transfer \\citep{intro_hybrid} of the encounter of a star-disc system with a discless companion.  The luminosity of these systems during the encounter is studied, and their potential for observation as a separate subclass of eruptive variable is discussed. ",
        "conclusions": "\\label{sec:conclusions}  \\noindent This paper has investigated the possibility of a stellar encounter (where one participant has a protostellar disc) being the progenitor of an outburst phenomenon.  The outbursts described here have several key features in common, originating from two distinct accretion events (for the primary and secondary respectively), independent of the secondary's orbital parameters.  The secondary's accretion rate grows and fluctuates as it traverses the disc's spiral structure: the peak occurs when the secondary disc is being formed from infalling stripped primary disc material, and the accretion rate then decays over several hundred years.  The primary accretion rate (perhaps the easiest to observe) increases slowly with a rise timescale of tens of years as the primary disc readjusts to mass stripping, and decays on a longer timescale than the secondary.  It has been established that these stellar encounters can enhance accretion rates to levels corresponding to FU Ori and EX Lupi outbursts, and that they have  observational signatures, that although low in probability to detect, are nonetheless possible in principle to see.  But, it is important to emphasise that these encounters cannot be responsible for \\emph{all} FU Ori phenomena (or EX Lupi for that matter), for the following reasons:  \\begin{enumerate} \\item This type of stellar encounter is too infrequent to explain the catalogue of outbursts currently observed, \\item Such encounters cannot maintain the rapid periodicity or repeatability required of some outbursts without significant decay of the outburst strength, \\item This origin would predict the detection of a companion (perhaps in the infrared or submillimetre) for every outburst host.  This is not the case; out of at least twenty FUors, only seven have a confirmed companion \\citep{Pfalzner_08}. \\end{enumerate}  \\noindent Despite this, outbursts from stellar encounters can mimic FU Ori well, with the correct general behavioural trends.  They may also potentially have the same triggering mechanism - mass pile up leading to MRI activation \\citep{Armitage_et_al_01,Zhu_et_al_09}, although higher resolution simulations (which resolve the inner region and reduce the artificial viscosity) are required to confirm this. Taking the fraction of FUors with a companion as a guide, it is estimated that at most 30\\% of outbursts detected could be due to an encounter or binary (with the actual figure presumably much smaller). The outbursts identified here represent a subtly different type of object, that although not currently detected, may be detected in future large-scale surveys of star-forming regions at far-infrared and sub-mm wavelength."
    },
    "0911/0911.5408_arXiv.txt": {
        "abstract": "We use the Novikov-Thorne thin disk model to fit the thermal continuum X-ray spectra of black hole X-ray binaries, and thereby extract the dimensionless spin parameter $a_* = a/M$ of the black hole as a parameter of the fit.  We summarize the results obtained to date for six systems and describe work in progress on additional systems.  We also describe recent methodological advances, our current efforts to make our analysis software fully available to others, and our theoretical efforts to validate the Novikov-Thorne model. ",
        "introduction": "We now know of about 40 stellar-mass black holes (BHs) in X-ray binaries in the Milky Way and neighboring galaxies with masses ranging from $\\sim 5-20$\\msun~[1]. Astrophysical BHs are completely described by the two numbers that specify their mass and spin.  BH spin is commonly expressed in terms of the dimensionless quantity $a_* \\equiv cJ/GM^2$ with $|a_*| \\le 1$, where $M$ and $J$ are respectively the BH mass and angular momentum.  Currently, there are two techniques that are delivering measurements of spin, namely fitting the thermal X-ray continuum [2--6] and modeling the profile of the Fe K line.  Our group is engaged in using the continuum-fitting (CF) method, which is the focus of this paper.  (For a discussion of the Fe K method see [7]).  Knowledge of BH spin is crucial for answering many key questions.  For example: Are relativistic jets powered by spin?  What role does spin play in producing a gamma-ray burst?  What constraints can be placed on models of supernovae, BH formation, and BH binary evolution?  What distribution of BH spins should LIGO waveform modelers be considering? For supermassive BHs, is the distribution of spins of the merging partners consistent with hierarchical models for their growth?  In the following two sections, we describe the CF method of determining spin and present the results we have obtained to date, while describing our current work on four additional sources.  The next section describes recent methodological advances, including our efforts to make all of our fitting software publicly available.  The penultimate section describes our work aimed at validating the theoretical underpinning of our work -- the Novikov-Thorne disk model -- via GRMHD simulations.  We conclude with a number of questions that motivate us. ",
        "conclusions": "We conclude with a list of questions that motivate us as we work toward our goal of measuring the spins of a dozen or more BHs.  What range and distribution of spins will we find?  Will GRS 1915+105 stand alone, or will we find other examples of extreme spin?  As we continue to refine our models and our measurements of $M$, $i$ and $D$, will we consistently find values of $a_* < 1$, or will we be challenged by apparent and unphysical values of the spin parameter that exceed unity? Will all our spin estimates be positive, or will we find some BHs with $a_* < 0$, corresponding to a counter-rotating disk?  Will there be large differences in spin between the class of young, persistent systems with their massive secondaries (Cyg X-1, M33 X-7, LMC X-1 and LMC X-3) and the ancient transient systems with their low-mass secondaries?  What constraints will these spin results place on BH formation, evolutionary models of BH binaries, models of relativistic jets and gamma-ray bursts, etc.?  What will be the implications of these spin measurements for the emerging field of gravitational-wave astronomy in the Advanced LIGO era? How will this new knowledge help shape the observing programs of IXO, LISA and other future space missions?"
    },
    "0911/0911.5122_arXiv.txt": {
        "abstract": "{Properties of candidate stars, forming out of molecular clouds, depend on the ambient conditions of the parent cloud. We present a series of 2D and 3D simulations of fragmentation of molecular clouds in starburst regions as well as clouds under conditions in dwarf galaxies, leading to the formation of protostellar cores.} {We explore in particular the metallicity dependence of molecular cloud fragmentation and the possible variations in the dense core mass function, as the expression of a multi-phase ISM, due to dynamic and thermodynamic effects in starburst and metal-poor environments.} {To study the level of fragmentation during the collapse, the adaptive mesh refinement code FLASH is used. With this code, including self-gravity, thermal balance, turbulence and shocks, collapse is simulated with four different metallicity dependent cooling functions. Turbulent and rotational energies are considered as well. During the simulations, number densities of 10$^8$-10$^9 \\rm ~cm^{-3}$ are reached. The influences of dust and cosmic ray heating are investigated and a comparison to isothermal cases is made.} {The presented results indicate that fragmentation increases with metallicity, while cosmic ray and gas-grain collisional heating counteract this. We also find that modest rotation and turbulence can impact the cloud evolution as far as fragmentation is concerned. In this light, we conclude that radiative feedback, in starburst regions, will inhibit fragmentation, while low-metallicity dwarfs also should enjoy a star formation mode in which fragmentation is suppressed.} {} ",
        "introduction": "The initial mass function (IMF) in our local neighborhood is observed to be a power-law function, nicely following a Salpeter slope. Recent studies \\citep{2005ASSL..327...89F, 2005Natur.434..192F, 2007AAS...210.1601S, 2008ApJ...681..365E} confirm this universality at low and high masses. When we turn to observations of more radical environments like the Galaxy center, a 'universal' shape is again confirmed and only hints for different IMFs appear. Measurements of abundance patterns \\citep{2007A&A...467..117B, 2008IAUS..245..339C} indicate a flatter IMF. Observations of the Arches cluster initially showed a shallow IMF \\citep{1999ASPC..186..329F, 2002A&A...394..459S} and a turn-over mass, and thus a typical mass, at a few solar masses \\citep[6-7M$_{\\odot}$,][]{2005ApJ...628L.113S}. However, \\cite{2006ApJ...653L.113K} do detect low mass stars with their deep photometry and merely find the presence of a local bump at $\\sim$6.3M$_{\\odot}$ and a somewhat shallower IMF ($\\Gamma = -1.0$ to $-1.1$). \\cite{2007MNRAS.381L..40D} confirm these results and claim a top-heavy IMF. Again, there appear to be perfectly reasonable explanations for the higher mass stars detected in the Arches cluster as \\cite{2007MNRAS.378L..29P} argue with their idea of a still collapsing cluster core. Of course, the idea that the conditions and the environment play an important role in shaping the IMF is worth studying. Looking at the mass-to-light ratios of ultra-compact dwarf galaxies, \\cite{2009arXiv0901.0915D} conclude from their models that it is most likely due to a top-heavy stellar IMF. If so, then they attribute this to feedback effects induced by massive stars. Still, these studies do not benefit from resolved stellar population data. One can associate the conditions and the environment of dwarf galaxies with those in the early Universe \\citep{2008IAUS..255..226H, 2009MNRAS.tmpL.207S}. It is also believed that the primordial, high redshift and zero metallicity, IMF must have been top-heavy, $\\sim$100M$_{\\odot}$ \\citep{2000ApJ...540...39A, 2001ApJ...561L..55O, 2002ApJ...564...23B, 2002Sci...295...93A, 2003ApJ...589..677O, 2008arXiv0810.1867J}. Hence, it appears that the IMF is universal, but that it is worthwhile to explore strong environmental (feedback) influences, e.g., starburst and dwarf galaxies, to understand the IMF better \\citep{2007MNRAS.374L..29K}. Since star formation, fragmentation and the IMF all depend on initial conditions \\citep{2003ApJ...596..253S, 2008IAUS..255...33G}, the question is how strong the impact of these initial conditions is. Also, how different should the environment be, or how strong the feedback, to cause a significant deviation from a Salpeter IMF?  Molecular clouds are the birth grounds of stars. These clouds are formed (rapidly) when interstellar gas and dust gather and cool down under the influence of hydrodynamics, thermodynamics and gravity \\citep{1999ApJ...527..285B, 2001ApJ...562..852H, 2005ApJ...620..786K, 2008ApJ...689..290H, 2008MNRAS.385.1893D, 2008MNRAS.391..844D}. Molecules form in the gas phase and on dust grains, and allow low temperatures to be reached \\citep{2005ApJ...626..627O, 2005JPhCS...6..155C, 2009A&A...496..365C, 2009arXiv0903.3120D}. Dense regions, like the center of the cloud, are able to form many stars, with a modest efficiency \\citep[$\\sim$5\\%;][]{2000ApJ...539..342E, 2005MNRAS.359..809C}. The distribution of the dense regions is thought to be the precursor of a stellar initial mass function \\citep{1998A&A...336..150M, 2001ASPC..243..301M, 2008A&A...477..823G, 2008MNRAS.391..205S}, and perhaps the stellar IMF is directly linked to this core mass function as it seems for the Pipe nebula \\citep{2008ApJ...672..410L}. Following a typical molecular cloud from its infancy to a mature state, for different starting conditions, gives insight into the processes that determine stellar masses and is the focus of this work. Specifically, we consider low metallicity environments (dwarfs) and warm dusty systems (starbursts).  In the next section we present the initial conditions used for our 38 simulations and present the results in the following section. Twenty of these simulations form the basis of this research. They comprise 4 different metallicities with 5 combinations of turbulent and rotational energies each. We add to these basis simulations comparison studies that include extra heating sources and cases where the primary heating and cooling terms are removed. ",
        "conclusions": "We have performed 38 hydrodynamical simulations where we have rigorously tested the dependence of molecular cloud fragmentation on the ambient conditions as they pertain to starburst and dwarf galaxy regions. We have compared these with the well known conditions of the Milky Way. Our results on the metallicity simulations show that the amount of fragmentation and the compressibility of the gas scales with increasing metallicity. In this, the cooling by gas phase and grain surface chemical species plays a subtle but fundamental role. We find that turbulence and rotation also play an important role. Turbulence of 0.3$E_{\\rm grav}$, following the Larson laws, is a crucial ingredient in the development of fragments. Although rotation is observed for sub-pc scale molecular clouds, it is unlikely that clouds of these sizes (parsec) have rotational support. All simulations have therefore been performed for $\\beta_0$ values of much less than unity. These results agree with expectations and are comparable to the work of \\cite{Machida1} and \\cite{2009arXiv0907.3257M}, albeit at lower density. The metallicity has a greater impact on cloud fragmentation than $\\beta_0$. For the highest metallicity model, fragmentation (into a binary) is still possible for even the lowest (or zero) rotational energy.  Regarding the initial temperature of the simulations, we note that starting with a high temperature of 100 K, as for starbursts, does not seem to be too important. When we do not consider the extra heating sources, radiative cooling can quickly overcome the high initial temperature and cool the cloud down to $\\sim$10 K. This is true for most cases except when the initial rotational energy is very low. In this situation, the adiabatic heating and the heating due to shocks, which are more prominent now, can overcome the cooling. Fragments have higher temperatures $>$100 K under these circumstances, which suppresses fragmentation. The surrounding gas, though, stays cold ($\\sim$10-20 K).  When we use low metallicity, similar to dwarf galaxies or the early universe, the average fragment mass increases in our simulations. This leads to massive cloud cores and eventually can induce massive star formation \\citep{2009arXiv0901.0915D, 2008ApJ...672..410L}, relevant to the origin of the IMF. \\\\  We summarize our main conclusions as follows: \\begin{itemize} \\item Increased metallicity promotes fragmentation. There is a clear and strong dependence of fragmentation on metallicity, where the average fragment mass decreases with increasing metallicity.  \\item Heating by cosmic rays and dust, for high metallicity initial conditions, are strong enough to slow down cooling and reduce fragmentation. When the gas and dust are thermally coupled, we find that the gas temperature stabilizes close to the dust temperature, T$_g$ = T$_d$, at densities $>$10$^5 \\rm ~cm^{-3}$.  \\item Fragmentation is suppressed for isothermal clouds. In almost all of the near-isothermal simulations that we did, we see a smaller number of fragments compared to their non-isothermal counterparts, except when adiabatic heating is very strong. Although the temperatures are exactly the same in several cases, even lower compared to the non-isothermal runs for the smallest $\\beta_{0}$ ratios, and the Jeans masses are comparable, the evolution and fragmentation of the cloud differs. We attribute this to the stiffer effective equation of state, i.e. $\\gamma=1$, that evidently has a great impact \\citep{2005IAUS..227..337K}. \\end{itemize}  For the main part in this research, we have performed 2D simulations. One can argue that 2D simulations are not representative for a 3D cloud. Certain physics, like MHD, is indeed not possible to do in 2D, but we have not used any kind of physics or initial condition that requires a third dimension. We find that a higher merger rate occurs for 2D simulations due to one less degree of freedom and that our fragments are of somewhat higher mass. Still, we do not see many mergers in general. The 3D simulations that we have run, show that a disk forms in a time that is comparable to or is less than the fragmentation time scale if even a modest rotation exists, effectively making the problem a two-dimensional one. The results are slightly different, perhaps due to the lower resolution, or possibly owing to the shorter simulation time, but show comparable structures. A disk does not form when there is complete lack of initial rotation. This, however, does not change our results, since the non-rotating 3D simulation also fragments into two compact objects with comparable masses and enjoys a similar temperature evolution (figure \\ref{fig:phasediag-suppl}) as its 2D counterpart.  Changes in the density profile do not affect the evolution of the cloud as far as the impact of metallicity is concerned. Testing a power-law density profile (eq. \\ref{eq:profile}), as observed in active galactic regions \\citep{2003ApJ...594..812G, 2008arXiv0805.0185M}, against a flat profile and keeping the mean density constant, we find that the results are similar. The results (not shown) are the same for the number of fragments, but also alike in their temperature evolution. Only the fragment masses are marginally smaller when using a flat profile, due to competitive accretion \\citep{2005ASSL..327..425B}.  \\begin{equation} \\rho = \\rho_0 \\left( \\frac{1} {1 + R/R_c} \\right)^{\\alpha} $~~cm$^{-3}. \\label{eq:profile} \\end{equation}  \\noindent $\\rho_0$ is the central density of $n=10^4 \\rm ~cm^{-3}$, R$_{c}$ is the characteristic length of 1 pc. This is the radius within which the central density remains approximately constant. $R$ is the specific cloud radius and $\\alpha=1.4$ is the power of the profile. \\\\  We do not reach star densities, but the fragmentation and formation of dense cores is determined at the much earlier phases that we probe. We have not made use of sink particles, as this is usually done to make a jump in resolution and form protostars from dense cores, but our fragments should be able to form many stars. Opacity effects beyond 10$^{22.5} \\rm ~cm^{-2}$ are not considered in our simulations simply because we do not reach very high densities. The maximum densities that we reach are n$\\sim$10$^8$-10$^9$ cm$^{-3}$. Molecular opacity at high densities can actually have a quite different effect on fragmentation compared to densities below 10$^{3} \\rm ~cm^{-3}$ \\citep{2005A&A...440..559P, 2006A&A...453..615P}. A cloud can become optically thick sooner and trap the lines when there are more metals to form molecules with. This in turn will heat the system and prevent further fragmentation. Incorporating radiative transfer will be the next step in the continuation of this research."
    },
    "0911/0911.0707_arXiv.txt": {
        "abstract": "X-ray afterglow light curves have been collected for over 400 Swift gamma-ray bursts with nearly half of them having X-ray flares superimposed on the regular afterglow decay. Evidence suggests that gamma-ray prompt emission and X-ray flares share a common origin and that at least some flares can only be explained by long-lasting central engine activity.  We have developed a shell model code to address the question of how X-ray flares are produced within the framework of the internal shock model. The shell model creates randomized GRB explosions from a central engine with multiple shells and follows those shells as they collide, merge and spread, producing prompt emission and X-ray flares. We pay special attention to the time history of central engine activity, internal shocks, and observed flares, but do not calculate the shock dynamics and radiation processes in detail.  Using the empirical $E_p-E_{iso}$ (Amati) relation with an assumed Band function spectrum for each collision and an empirical flare temporal profile, we calculate the gamma-ray (Swift/BAT band) and X-ray (Swift/XRT band) lightcurves for arbitrary central engine activity and compare the model results with the observational data.  We show that the observed X-ray flare phenomenology can be explained within the internal shock model. The number, width and occurring time of flares are then used to diagnose the central engine activity, putting constraints on the energy, ejection time, width and number of ejected shells. We find that the observed X-ray flare time history generally reflects the time history of the central engine, which reactivates multiple times after the prompt emission phase with progressively reduced energy. The same shell model predicts an external shock X-ray afterglow component, which has a shallow decay phase due to the initial pile-up of shells onto the blast wave. However, the predicted X-ray afterglow is too bright as compared with the observed flux level, unless $\\epsilon_e$ is as low as $10^{-3}$. ",
        "introduction": "The study of gamma-ray bursts (GRBs) has been greatly advanced following the launch of the Swift Gamma-Ray Explorer (Gehrels et al. 2004) on November 20, 2004. Thanks to its rapid slewing capability, multi-wavelength observations of GRB afterglows, usually as soon as $<100$ seconds after the burst trigger, have been performed regularly for many bursts. The X-ray telescope (XRT, Burrows et al. 2005a) aboard this satellite has given us unprecedented access to early afterglows in the X-ray band. For most GRBs, XRT has observed smoothly decaying power-law or broken power-law afterglows (Nousek et al. 2006; O'Brien et al. 2006; Zhang et al. 2007; Liang et al. 2007, 2008; Evans et al. 2009).  In about half of GRBs, superimposed on the background X-ray afterglow, one or more X-ray flares are seen (Burrows et al. 2005b; Romano et al. 2006; Falcone et al. 2006, 2007; Chincarini et al. 2007).   Flares among the bursts share some common properties, which include the following:  \\begin{itemize}  \\item The morphology of flares is similar: they all show a smooth, rapid rise and a rapid fall (Romano et al. 2006).  \\item They are superimposed on a background power-law decay afterglow component with the slope before the flaring equal to that after the flare (Burrows et al. 2005c).  \\item The width of flares is typically narrow, with $\\delta t / t \\sim 1/10$ on average, where $t$ is the emission time of the flare.  However, flares can become wider ($\\delta t / t$ becomes larger) as the time of their emission increases, and no sharp flares are seen at late times (Chincarini et al. 2007; Kocevski et al. 2007).  \\item Similar flaring activity has been detected in both types of GRBs: those believed to be of massive star origin (Type II, typically long duration) and those believed to be of compact star merger origin (Type I, typically short, e.g. in GRB 050724, Barthelmy et al. 2005) \\footnote{For a full discussion of the two physically distinct types of GRBs, see Zhang et al. (2009).}. This suggests that the phenomenon is insensitive to the progenitor type (Perna et al. 2006).  \\item Flares are spectrally harder than the underlying afterglow (Burrows et al. 2005b; Romano et al. 2006; Falcone et al. 2006).  \\item For a small sample of bright flares whose time-dependent spectral analysis can be performed, flares are found to soften as they decay (Burrows et al. 2005b,c; Falcone et al. 2006, 2007), reminiscent of GRB prompt emission.  \\item Average X-ray flare luminosity decays with time as a power-law with slope $\\sim$ -1.5 for a sample of flares from many GRB's.  This temporal relationship also seems to hold among flares from a single multi-flared burst (Lazzati, Perna \\& Begelman 2008). \\end{itemize}  Flares also vary from burst to burst, from flare to flare, in the following ways: \\begin{itemize}  \\item The shape of flares can be best fit with a number of different profiles: gaussian, log-normal distributions, exponential or power-law rise or decay with differing slopes for both the rising and falling portion of the flare (Chincarini et al. 2007).  \\item The number of flares seen per burst varies from modal value of 1, mean value of $\\sim$ 2.5 and maximum value of 8 (Chincarini et al. 2007).  \\item The fluence seen in flares can vary from a few percent of the burst fluence to larger than the burst itself (e.g. for GRB050502B, Burrows et al. 2005b,c; Falcone et al. 2006).  \\item The emission time of flares varies from perhaps before the slew of XRT (e.g. GRB060607) to late flares at $10^4$ seconds (e.g. GRB050904, Cusumano et al. 2006) to $10^5$ seconds (for the Type I GRB050724, Barthelmy et al. 2005). Some flares are superimposed on other flares (e.g. GRB050916, Burrows et al. 2005b,c; Chincarini et al. 2007). Although most flares happen early, from 100 to 1000 seconds, the distribution of $t$ tails off to $10^6$ seconds (Chincarini et al. 2007). \\end{itemize}  A few properties of X-ray flares suggest that they are connected to internal emission processes due to a restarting of the central engine at late times (Burrows et al. 2005b; Falcone et al. 2006; Romano et al. 2006; Zhang 2007). The arguments in favor of such a ``late internal'' model (in contrast to the external shock model) are the following: First, the rapid rise and fall of flares with $\\delta t / t \\sim 1/10$ strongly disfavors external shock models that involve a large angular area over which radiation would be emitted (Zhang et al. 2006, see also Ioka et al. 2005; Fan \\& Wei 2005; Lazzati \\& Perna 2007).  Second, the connected, underlying continuum which has the same slope both before and after the flare suggests that flares are not related to the canonical afterglow and are instead superimposed on top of this regular afterglow decay (Chincarini et al. 2007).  Third, given the same observed X-ray flare amplitude, the internal model requires a much smaller energy budget than the external shock model (Zhang et al. 2006). While the internal model can produce significant X-ray flares with energy much less than that during the prompt emission, the external shock model requires an energy budget at least comparable to that of the initial blast wave in order to make a noticeable change in flux (Zhang \\& M{\\'e}sz{\\'a}ros, 2002). Internal models are much more ``economical'' as far as energy budget is concerned. Next, Liang et al. (2006) have shown that as long as the central engine clock is reset to zero at the beginning of each flaring episode, the curvature effect can naturally explain the spectral index and the temporal decay index of the decay phase of flares. Finally, although there is no correlation between the number of prompt emission burst pulses and X-ray flares in any given burst (Chincarini et al. 2007), X-ray flares do exhibit the same spectral softening as GRB prompt emission (Burrows et al. 2005b,c; Falcone et al. 2006, 2007). This suggests that the properties of X-ray flares make them likely to be caused by the same mechanism as prompt emission seen in gamma-rays.  Any X-ray flare model must be able to explain the bulk similarities and differences among flares enumerated above. The leading ``internal'' model is the internal shock model (Rees \\& M\\'esz\\'aros 1994), which invokes internal collisions of shells within an unsteady central engine wind. Previous internal shock models have focused on interpreting prompt gamma-ray emission properties (e.g. Kobayashi et al. 1997; Panaitescu et al. 1999; Spada et al. 2000; Guetta et al. 2001). Kobayashi et al. (1997) have shown that internal shocks can explain the highly variable profiles seen in GRBs.   In their model, shells of matter with varying Lorentz factors ejected from the central engine produce about as many collisions as the number of shells, and the time of ejection of the shell is highly correlated to the time of collision of the shells, indicating that there is essentially no delay between energy ejection by central engine emission and the time when that emission is seen. Recently, using a model where collisions are not perfectly inelastic, Li \\& Waxman (2008) studied the effect of residual collisions in optical emission which would be slightly delayed from gamma-ray prompt emission.  Within the context of X-ray flares, the internal shock model has been discussed by a number of authors. Zhang et al. (2006) and Fan \\& Wei (2005) discussed how a late internal shock may produce a softer flare than prompt gamma-ray emission. Wu et al. (2005) discussed both late internal and external shock models to study X-ray flares and concluded that at least some flares have to be produced by late internal shocks. Lazzati \\& Perna (2007, see also Zhang 2007) proved the suggestions of Burrows et al. (2005b) and Zhang et al. (2006) that late flares must require late injection of shells and cannot be produced by late collisions of shells ejected during the prompt phase. Yu \\& Dai (2009) studied the shock physics and radiation processes of late internal shocks that may be responsible for X-ray flares.  In this paper, we focus on another aspect of the internal shock model for X-ray flares. Extending the work of Kobayashi et al. (1997) to concentrate on X-ray flares, we have created a GRB fireball model with a central engine that can eject multiple episodes of matter shells with any Lorentz factor, thickness and mass distributions. This code is used to address the question of what kind of central engine activities are demanded in order to reproduce the observed properties of X-ray flares. The conclusion drawn from this study may be taken as the requirements of some central engine models on X-ray flares (e.g. King et al. 2005; Perna et al. 2006; Proga \\& Zhang 2006; Dai et al. 2006; Lee et al. 2009)\\footnote{Other internal dissipation models for X-ray flares (e.g. Panaitescu 2008) may not subject to these requirements.}. In \\S2, we first examine the physics of two-shell collisions and describe our phenomenological treatment of the spectrum and lightcurve of individual collisions. In \\S3, we discuss blast wave evolution, which sets an outer boundary for the internal collisions that are relevant to X-ray flares. We then dedicate our discussion (\\S4) to multiple collisions for multiple ejection episodes and use the observations to diagnose the required central engine activities. Our results are summarized in \\S5 with some discussion. ",
        "conclusions": "We have developed a numerical code to model the internal collisions of an unsteady wind with arbitrary central engine activities. This is an extension of the internal shock models previously developed to model GRB prompt emission (e.g. Kobayashi et al. 1997), with a focus on X-ray flares that are commonly observed in GRB X-ray afterglows (e.g. Chincarini et al. 2007). Our motivation is to diagnose the required central engine activities based on the observational properties of X-ray flares as summarized in \\S1. The following conclusions can be reached:  \\begin{itemize}  \\item The internal shock model with multiple ejection episodes can generally reproduce the properties of X-ray flares. Our shell model naturally explains both prompt emission and X-ray flares, suggesting that they originate from very similar mechanisms. \\item We find that the number of pulses/flares is directly related to the number of shells released from the central engine. More shells means more flares.  In general, we find that the number of collisions is slightly smaller than the number of the shells ejected (e.g. the most probable values are: $5\\%$ less for 100 simulated shells, $10\\%$ less for 50 simulated shells, 8 collisions for 10 simulated shells and 1 collision for 4 simulated shells, see Fig.\\ref{stats} for distributions). For a modal number of observed flares of 1, between 2 and 5 shells need to be released on average. \\item The correlation between $t_{ej}$ and $t_{\\oplus,col}$ ensures that the shells that are responsible for creating flares must be ejected just prior to being seen. This is because the second term in the right hand side of Eq.(\\ref{time}) is typically much shorter than the first term. This means that not only the central engine activity must be prolonged, but it must also be episodic. Steady energy injection cannot produce flare-like features.  In other words, early flares are created by shells ejected early, while late flares by shells ejected late.  Since flares are seen as late as $10^6$ seconds, this means that the central engine can be active for a long time after the prompt emission. These conclusions are consistent with Zhang et al. (2006), Liang et al. (2006) and Lazzati \\& Perna (2007). \\item The large variance of fluences seen in flares can be explained by the different energies ($E~\\sim \\gamma m c^2$) of the ejected shells. The peak luminosity of a flare also depends on the width of the pulse, which is either related to shell spreading or intrinsically different durations of shell ejection. \\item This study seems to rule out uniform (or narrow distribution) thickness shells.  The observed widths of flares cannot be reproduced solely by spreading effects. One requires that later ejection episodes eject shells with larger widths to explain the typical flare width of $\\delta t_e/t \\sim 1/10$.  Since most shells collide before spreading, typical late collisions would be too narrow if later ejected shells had the same width as the initial prompt ejection episodes. Sharp flares where the width is near 1/10 of the emission time of the flare are common, but ``fat'' flares are seen occasionally as well.  For example, the giant flare of GRB050502B has a $\\delta t_e/t \\sim 1$ (Chincarini et al. 2007) implying that shells may either have spread before colliding or may simply have had a large width when ejected.  Shells which have spread significantly before colliding will have energies spread over a large log-Gaussian shape and may be too dim to reach above the afterglow decay.   \\item For goldilocks type shells, flares are seen in the XRT band only.  The gamma-ray component of the X-ray flare (Band-function extension) is below the BAT detector threshold.  Figure \\ref{fluxlightcurves} shows light curves from a few typical simulations and these features can be seen. \\item Flares superimposed on other flares are simply shells that collide near the same time with different widths and fluences. The observed X-ray afterglow is a superposition of flares due to prolonged central engine activity and a background afterglow radiation, whose origin is not addressed in our paper.  \\item The decrease in average flare luminosity as a function of time indicates shells released at later times must create less energetic collisions (Lazzati et al. 2008). In order to produce the results seen, late released shells must be wider, less dense, slower or a combination of the above in order to create less energetic collisions seen.  \\item The same shell model can give a prediction on the X-ray afterglow emission from the blast wave. Due to the continuous piling up of late-ejection shells onto the blast wave, the lightcurve can show a shallow decay phase that is commonly observed. However, if a standard value $\\epsilon_e=0.1$ is adopted, the predicted X-ray lightcurve is much brighter than what is seen, by about three orders of magnitudes. In order to reproduce the observed data, either a much lower $\\epsilon_e$ (as low as $10^{-3}$) needs to be introduced, or the internal emission that powers the prompt emission and X-ray flares has to be much more efficient than internal shock predictions. \\end{itemize}  Recently there has been interest in optical flares (Kr\\\"{u}hler et al. 2009, Melandri et al. 2006). These flares can be naturally interpreted in our model by invoking collisions between low energy shells or wide shells. Such collisions could be seen in optical bands but may be missed in higher energy bands. A direct expectation from this model is that optical flares should on average have lower energies and broader profiles than X-ray flares. An internal shock origin of optical flares was also proposed by Wei (2007).  Finally, we want to emphasize that observationally flares have been seen in both Type I (e.g. GRB 050724) and type II GRBs. This requires that both types of progenitor have a similar central engine, which can eject an episodic wind to power late central engine activities which is a requirement of any central engine models for X-ray flares."
    },
    "0911/0911.2228_arXiv.txt": {
        "abstract": "Radio loud active galactic nuclei (AGN) are on average 1000 times brighter in the radio band compared to radio quiet AGN. We investigate whether this radio loud/quiet dichotomy can be due to differences in the spin of the central black holes that power the radio-emitting jets.  Using general relativistic magnetohydrodynamic simulations, we construct steady state axisymmetric numerical models for a wide range of black hole spins (dimensionless spin parameter $0.1\\le \\astar \\le 0.9999$) and a variety of jet geometries.  We assume that the total magnetic flux through the black hole horizon at radius $r_{\\rm H}(\\astar)$ is held constant. If the black hole is surrounded by a thin accretion disk, we find that the total black hole power output depends approximately quadratically on the angular frequency of the hole, $P\\propto\\OmegaH^2\\propto (\\astar/r_{\\rm H})^2$.  We conclude that, in this scenario, differences in the black hole spin can produce power variations of only a few tens at most.  However, if the disk is thick such that the jet subtends a narrow solid angle around the polar axis, then the power dependence becomes much steeper, $P\\propto\\OmegaH^4$ or even $\\propto \\OmegaH^6$.  Power variations of $1000$ are then possible for realistic black hole spin distributions. We derive an analytic solution that accurately reproduces the steeper scaling of jet power with $\\OmegaH$, and we provide a numerical fitting formula that reproduces all our simulation results. We discuss other physical effects that might contribute to the observed radio loud/quiet dichotomy of AGN. ",
        "introduction": "\\label{sec_intro}  The first active galactic nuclei (AGN) were discovered through radio emission associated with their relativistic jets.  However, it soon became clear that not all AGN\\footnote{We use the generic term AGN to refer to both luminous quasars and less luminous active nuclei such as Seyferts, LINERs, etc.} have powerful radio jets; in fact, only about 10\\% of quasars do.  The evidence for a dichotomy between radio loud and radio quiet AGN has become stronger over the years, culminating in the impressive demonstration by \\citet*{ssl07} that two very distinct populations of AGN are clearly visible when radio luminosities $L_R$ of AGN are plotted against optical luminosities $L_B$.  For a given value of $L_B$, these authors show that $L_R$ of radio loud AGN is $\\sim10^3{-}10^4$ times greater than that of radio quiet AGN.  Also, the two populations follow two well-separated tracks on the plot.  The origin of the radio loud/quiet dichotomy has been much discussed in the literature.  At Eddington ratios $\\lambda = L_{\\rm bol}/L_{\\rm Edd} \\sim 10 L_B/L_{\\rm Edd} > 0.01$, where $L_{\\rm bol}$ is the bolometric luminosity of the AGN and $L_{\\rm Edd}$ is its Eddington luminosity, a likely explanation for the dichotomy \\citep{ho00,ssl07} is that these systems accrete via a standard thin accretion disk.  Jet production is then expected to be intermittent, as found to be the case in black hole (BH) X-ray binaries (\\citealt{fend04a}).  However, the existence of two distinct populations for $\\lambda<0.01$ is harder to explain.  Even at these low luminosities, the radio loudness parameter $R=L_{\\rm 5\\ GHz}/L_B$ of the radio loud population is at least a factor of $10^3$ times larger than that of the radio quiet population. However, at low values of $\\lambda$, BH X-ray binaries typically are in a hard spectral state associated with an advection-dominated accretion flow (ADAF, \\citealt{nm08}), and in this state, all BH X-ray binaries are radio loud \\citep{fend04a}.  Why then do AGN with similar values of $\\lambda$ have a radio loud/quiet dichotomy?   One possible explanation is that radio loud objects are driven by a central BH with a large spin which produces a jet by the Blandford-Znajek (BZ) mechanism \\citep[hereafter, \\citetalias{bz77}]{bz77}.  This is referred to as the spin paradigm \\citep{blandford1990,wc95,blandford99}, which is in contrast to the accretion paradigm which states that the BH mass and mass accretion rate determine the jet power \\citep{blandfordrees74,blandfordrees92}.  These different paradigms plausibly operate together to some degree \\citep{bbr84,meier02}.  The spin paradigm has been invoked to explain the observed correlation between jet and accretion power in elliptical galaxies \\citep{allen2006} by combining an ADAF accretion model with the BZ effect \\citep{nemmen07,benson09}.  In terms of the dimensionless spin parameter $\\astar = J/GM^2$, where $M$ and $J$ are the mass and angular momentum of the BH, it is found that one requires $\\astar\\gtrsim 0.9$ to explain the correlation.  The dichotomy in the power and spatial distribution of emission in Fanaroff-Riley classes 1 and 2 (FR 1 and FR 2) radio galaxies may also be explained by the spin paradigm \\citep{baum95}.   The radio loud/quiet dichotomy in AGN \\citep{kellerman89,msl98,ivezic04} has been explained in terms of an in situ trigger for relativistic jets \\citep{meier97}. It could also be explained by the differences in the evolutionary stages at which we observe the objects \\citep{blundell08} coupled with the episodic activity of the AGN evidenced by the change or precession of jet orientation \\citep{sb08,sbd08,shb10}. However, another possibility is that the two AGN populations have different merger and accretion histories which lead to different BH spins. Recent observations show that, for $\\lambda<0.01$, all radio loud AGN reside in elliptical galaxies, whereas radio quiet AGN live mostly in spirals \\citep{ssl07}. \\citet*{volonteri07} explored a number of scenarios for the formation of ellipticals and spirals and showed that it is plausible for the nuclear BHs in spirals to have lower spins than those in ellipticals. This suggests that the spin paradigm may explain the radio loud/quiet dichotomy.   The main problem with this explanation is that the difference in radio loudness between the radio loud and radio quiet populations is a factor of $10^3$ \\citep{ssl07}. This is a strikingly large difference. According to the accretion or spin paradigms, relativistic jets are produced by magnetic outflows from either the inner region of the disk or the spinning BH.  The power in the disk outflow is expected to be proportional to the disk luminosity, which leaves no room for a radio loud/quiet dichotomy, so we will ignore this possibility\\footnote{The BH spin also drives power into the disk causing a more powerful disk outflow, but this still requires BH spin to introduce a dichotomy.}. The power from the BH does depend on the spin parameter, and we will focus on this\\footnote{There are arguments to suggest that the luminosity of the disk outflow should be greater than that from the central spinning BH \\citep{ga97,lop99}.  However these arguments assume low values of turbulent viscosity and weak magnetic fields near the BH, and also do not account for the effects of the general relativistic plunging region (see, e.g., \\citealt{mck07a} for a discussion). \\citet{mck05} finds that the luminosities of the jet and the disk are similar (however, see \\citealt{nagataki09}). For the purposes of this paper, we ignore the disk wind. }. However, the analytical model of \\citetalias{bz77} indicates that the jet power $P$ varies only as $\\astar^2$ for fixed magnetic flux threading the horizon. If such a weak variation is to produce a difference of $10^3$ in radio power\\footnote{We assume that the radio luminosity is proportional to the jet power.}, we need $\\astar$ in the two populations to differ by a factor $\\sim30$, which does not seem plausible given likely merger histories \\citep{hughesblandford03,gammie_bh_spin_evolution_2004}. More plausible is a factor $\\sim3$ difference in the median values of $\\astar$ in the two populations (e.g., see \\citealt{volonteri07}), but this will produce only a factor $\\sim10$ difference in jet power, not $10^3$.   The \\citetalias{bz77} scaling for power, $P\\propto \\astar^2$, was derived in the limit $\\astar\\ll1$, for a razor-thin disk, assuming the magnetic flux threading the BH is independent of $\\astar$. Recent analytical and numerical work by \\citet{tn08} show that, for a BH threaded by a split monopole magnetic field, $P$ increases as $\\astar^4$ at large values of $\\astar$ when higher-order corrections are included. However, the analytical model worked out by these authors only gives a factor of two increase in power above the \\citetalias{bz77} result even at $\\astar=1$. Their numerical simulations achieve a slightly steeper scaling, but still only a factor of four above the \\citetalias{bz77} result at $\\astar=1$. In any case, their work hints that a much steeper dependence of power on $\\astar$ occurs as $\\astar\\to 1$. Are there other effects that can introduce an even steeper dependence on $\\astar$?   General relativistic magnetohydrodynamic simulations of accretion disks by \\citet{mck05} showed that for $\\astar\\gtrsim 0.5$ the jet power varies as steeply as the fourth power of the BH angular rotation rate, i.e., $P\\propto \\OmegaH^4$, where $\\OmegaH\\propto \\astar/r_{\\rm H}$ and $r_{\\rm H}$ is the radius of the horizon. Compared to the scaling $P\\propto\\astar^4$, the scaling $P\\propto \\OmegaH^4$ introduces an additional factor of $16$ due to the division by $r_{\\rm H}$, since $r_{\\rm H}$ decreases from $2M$ to $M$ as $\\astar$ varies from $0$ to $1$. \\citet{mck05} also finds that the power output of the entire BH has a shallower dependence on $\\OmegaH$ compared to the power output of the jet, which subtends a small solid angle above the disk and corona\\footnote{\\citeauthor{mck05}'s models have an accretion disk with a disk+corona+wind of angular extent $H/R\\sim 0.6$. For $\\astar\\gtrsim 0.5$ the power per unit mass accretion rate scales as $\\propto \\OmegaH^5$ for the polar jet and $\\propto \\OmegaH^4$ for the entire horizon. However, in these models, the mass accretion rate through the BH horizon per unit fixed mass accretion rate at large radius scales as $1/\\OmegaH$ for $\\astar\\gtrsim 0.5$, because of the ejection of a massive wind (as also seen in \\citet{hk06}). Hence, expressed in terms of $\\dot{M}$ at large radius, the polar jet power scales $\\propto \\OmegaH^4$ and the power output from the entire horizon scales $\\propto \\OmegaH^3$.}. This suggests that changes in the solid angle subtended by the jet (via changes in the disk thickness) could change the steepness of the power output as a function of $\\astar$.   Since the scaling of $P$ with BH spin is important for jet studies, we have explored this issue in detail using general relativistic magnetohydrodynamic numerical simulations. We consider a variety of geometries for the shape of the jet to see if we can come up with any scenario in which jet power could change by a large factor for a modest variation in $\\astar$. We show that the most favorable scenario is a BH surrounded by a thick accretion flow with an angular thickness $H/R\\sim1$.  We show that in this case the power output into a polar jet has a steep dependence on the spin, $P\\propto \\OmegaH^4$, and that the scaling steepens to $P\\propto \\OmegaH^6$ for even thicker disks. Hence, we confirm the basic result found by \\citet{mck05} of a steep dependence of jet power on $\\astar$ at high latitudes above the disk. We suggest that this strong dependence may explain the radio loud/quiet dichotomy in AGN.  Our numerical setup is described in \\S\\ref{sec:num}.  The results for BHs with razor-thin disks are presented in \\S\\ref{sec:thindisk}, and for jets from BHs with thick disks in \\S\\ref{sec:thickdisk}.  We discuss the results in the context of the AGN radio loud/quiet dichotomy in \\S\\ref{sec:discussion}, and conclude in \\S\\ref{sec:conclusions}.  We work with Heaviside-Lorentzian units and set $c=G=1$. ",
        "conclusions": "\\label{sec:conclusions}  We set out in this paper to explain the radio loud/quiet dichotomy of AGN in the context of the BH spin paradigm. For razor-thin disks, we found that BH spin alone is insufficient to explain the observations even if the radio loud and radio quiet populations have very different merger and accretion histories. However, we found that the presence of a thick disk, such as an ADAF, can significantly enhance the spin dependence of the power output, to the extent that it can reasonably account for the observed radio loud/quiet dichotomy. Our only modification to the revised spin paradigm of \\citet{ssl07} is that both populations should contain a BH surrounded by a thick disk such that the jet subtends  a small solid angle around the polar axis.   These results were obtained by performing general relativistic numerical simulations of collimated force-free and MHD jets from spinning magnetized BHs for a wide range of spins (up to $a=0.9999$) and jet confinement geometries.  We showed that, regardless of the geometry, a BH threaded with a magnetic flux $\\Phi_{\\rm tot}$ and surrounded by a razor-thin disk produces a jet with power $P\\approx k\\Phi_{\\rm tot}^2\\OmegaH^2$, where $k$ is a known constant factor which depends only weakly on the field geometry, $r_{\\rm H} = M[1+(1-a^2)^{1/2}]$ is the radius of the BH horizon, and $\\OmegaH=a/(2r_{\\rm H})$ is the angular frequency of the BH. This result gives a somewhat steeper dependence of jet power on $\\astar$ compared to the original scaling $P\\propto \\astar^2$ obtained by \\citetalias{bz77}.  Nevertheless, we conclude that for a fixed magnetic flux $\\Phi_{\\rm tot}$, even this revised scaling is much too shallow to explain the radio loud/quiet dichotomy of AGN.  Our goal, therefore, was to identify any other effect that may cause the jet power to depend more steeply on BH spin.  We found that such an effect naturally exists.  We showed that the power output of a BH surrounded by a thick accretion disk with $H/R\\sim1$ (this is the effective thickness of the disk, corona and mass-loaded disk wind), as expected in systems with advection-dominated accretion flows (ADAFs, \\citealt{nm08}), is $P\\propto \\OmegaH^4$, and even $\\propto\\OmegaH^6$ for very thick disks (\\S\\ref{sec:thickdisk}).  In this case we can explain the radio loud/quiet dichotomy by having two different populations of galaxies with modestly different BH spins.  For the case $H/R=1$, the radio loud population needs to have large spins $a\\simeq1$ while the radio quiet AGN population needs to have $a\\simeq0.15$. Such spin values may plausibly result from differences in the merger and accretion histories of supermassive BHs in elliptical and spiral galaxies (\\S\\ref{sec:discussion}).  We worked out in the Appendices a first principles analytic model which accurately reproduces our numerical results for the jet power over a wide range of BH spin and disk thickness (Figures~\\ref{fig3}, \\ref{fig6})."
    },
    "0911/0911.0477_arXiv.txt": {
        "abstract": "We present a closed system of coupled first order differential equations governing the secular linear momentum loss of a compact binary due to emitted gravitational waves, with the leading order relativistic and spin-orbit perturbations included. In order to close the system, the secular evolution equations of the linear momentum derived from the dissipative dynamics are supplemented with the secular evolutions of the coupled angular variables, as derived from the conservative dynamics. ",
        "introduction": "The inspiral of compact objects in binary systems is driven by the emitted gravitational radiation which carries energy, angular and linear momenta away from the source. Global conservation of these quantities implies a radiative orbital evolution, including a possible recoil of the system. The energy and magnitude of orbital angular momentum of the binary determine the quasi-Keplerian orbit in the post-Newtonian (PN) regime. These quantities are known to high accuracy for binaries on eccentric orbits, with PN, spin-orbit (SO), spin-spin (SS), mass quadrupole - mass dipole (QM) and magnetic dipole - magnetic dipole (DD) coupling terms \\cite{Kepler}. Moreover, the orientation of the (quasi-precessing) plan of motion, defined by the direction of the Newtonian orbital angular momentum, is known to high accuracy with all the above-mentioned contributions.  Due to asymmetries in the configuration of the source the radiation is often emitted anisotropically. The linear momentum loss of the binary leads to the recoil of the center of mass in the opposite direction, which in extreme situations results the kick-off of the binary from its host galaxy. This effect, however, vanishes for equal mass binaries due to the fact that the leading order relativistic contribution to the linear momentum loss scales with the mass difference $\\delta m=m_{1}-m_{2}$ of the orbiting bodies.  The gravitational recoil of a binary system and the classical linear momentum loss of the final black hole were discussed in \\cite% {Peres,Bekenstein} with the inclusion of the lowest multipoles needed for the computation of momentum ejection. The first quasi-Keplerian analytic studies were given in \\cite{Fitchett}, where the linear momentum flux of gravitational waves from a binary system of two point masses in Keplerian orbit were calculated. 1\\ PN corrections to the gravitational recoil were discussed in \\cite{Wiseman}, while 2PN recoil effects for binaries on quasicircular orbits in \\cite{Blanchet}. Several numerical estimates for the kick velocity of non-spinning binaries were given, e.g. in \\cite% {Baker,HHind,GonzalezNS}, with the maximum recoil in this case found as $% \\sim 175$ km/s \\cite{Gonzalez}. Recently it has been shown \\cite{TBW}, that the ringdown phase acts as an anti-kick as compared to the inspiral and plunge, the global result being consistent with the numerical estimates.  The rotation of the components adds however additional structure to the emitted radiation and recoil. Spin contributions to the linear momentum loss are analyzed in \\cite{Kidder}, and more recently in \\cite{SchBuon}. The SO contributions scale with the magnitude of spins which could be high for galactic black holes. Numerical analyses indicate that significant gravitational recoil can be obtained in spinning binaries \\cite% {Herrmann,Koppitz,Campanelli,Schnittman1,Lousto} even with a high degree of symmetry in the configuration, i.e. for equal-mass binaries with antialigned initial spins in the orbital plane \\cite{Brugmann}. In a detailed review \\cite{Racine}, Racine, Buonanno and Kidder have computed the instantaneous linear momentum flux emitted by spinning binaries at 2PN order with the inclusion of the next-to-leading order SO, SO tail and SS terms. Moreover, the recoil velocity as a function of the orbital frequency was given for quasicircular orbits.  In the present work we give the previously unknown \\textit{secular expressions} for the linear momentum loss of an inspiralling binary system, with the inclusion of leading order relativistic and SO contributions. The analytic approach presented here is suitable for characterize analytically the dependence of the recoil on binary parameters.  Section 2 contains elements of the conservative dynamics with the inclusion of first post-Newtonian and spin-orbit contributions. At the end of the section the secular conservative evolution of the relevant angle variables is given. In Section 3 we start from the expressions of the SO contributions to the instantaneous linear momentum loss given in \\cite{Kidder}, then we derive the dissipative secular linear momentum evolution. The two sets of secular evolutions couple to a closed system of first order differential equations.  The secular evolutions are derived as follows: first we rewrite the instantaneous evolutions in terms of the radial parametrization of quasi-Keplerian orbits \\cite{Kepler}. As a result, all the radial functional dependencies are expressed in terms of a single variable $\\chi $, the generalized true anomaly. Averaging these expressions over a radial period becomes particularly straightforward in terms of the complex version of $% \\chi $, which renders the problem to the computation of residues in the origin of the complex parameter plane \\cite{Param,Gergely}. This procedure smears out short timescale effects and we obtain a simpler dynamics, suitable for monitoring the secular changes. The procedure is similar to our previous computations of the SO-induced secular changes of the energy, magnitude of orbital angular momentum and relative angles among the orbital angular momentum and spins \\cite{GPV3}. An additional difficulty in the present computation arises from the vectorial character of the linear momentum, and the subsequent system of coupled differential equations.  Our description is generic, being valid for generic (non-circular, non-spherical) orbits.  \\textit{Notation.} For any vector $\\mathbf{V}$ its magnitude is denoted as $% V $ and its direction (a unit vector) as $\\mathbf{\\hat{V}}$. ",
        "conclusions": "We have derived secular evolution equations, Eqs. (\\ref{angledotave}) and (% \\ref{psipdotave}) for the angular variables $\\alpha,\\phi_n$ and $\\psi_p$ characterizing the orientation of the orbit, with the inclusion of the leading order relativistic and spin-orbit coupling contributions. Together with the evolution of the angles $\\kappa_{i}$ given by the first two Eqs. (2.17) of \\cite{GPV3} and the projections of $\\mathbf{J=L+S}_{\\mathbf{1}}+% \\mathbf{S}_{\\mathbf{2}}$, determining $\\psi_i$ in terms of the other variables, these form a closed system of differential equations for the conservative evolution of the angular variables.  We have complemented these with the equations expressing the respective secular losses of the linear momentum components, Eqs. (\\ref{Pdotave}). Then we have a closed system of coupled first order differential equations, giving the linear momentum loss during the inspiral in analytic~form.  This system of equations is also suitable for numerical evolution, which in principle can lead to the value of the recoil velocity right before the plunge and also allows to determine the shift in the position of the binary during the inspiral. Our analytic approach is suitable for discussing the dependence of the recoil on various parameters characterizing the binary.  In order to gain higher accuracy, the results derived in this paper can be further generalized by the inclusion of higher order corrections: the spin-spin, mass quadrupole - mass monopole, and second order relativistic contributions."
    },
    "0911/0911.1270_arXiv.txt": {
        "abstract": "{} { % We searched for long-term period changes in the polar \\ekuma\\ using new optical data and archival X-ray/EUV data.} { % An optical ephemeris was derived from data taken remotely with the MONET/N telescope and compared with the X-ray ephemeris based on \\einstein, \\rosat, and \\euve\\ data. A three-parameter fit to the combined data sets yields the epoch, the period, and the phase offset between the optical minima and the X-ray absorption dips.  An added quadratic term is insignificant and sets a limit to the  period change.} { % The derived linear ephemeris is valid over 30 years and the common optical and X-ray period is $P=0.0795440225(24)$\\,days. There is no evidence of long-term $O-C$ variations or a period change over the past 17 years ($\\Delta P = -0.14\\pm 0.50$\\,ms). We suggest that the observed period is the orbital period and that the system is tightly synchronized. The limit on $\\Delta P$ and the phase constancy of the bright part of the light curve indicate that $O-C$ variations of the type seen in the polars DP\\,Leo and HU\\,Aqr or the pre-CV NN\\,Ser do not seem to occur in \\ekuma. The X-ray dips lag the optical minima by $9.5^\\circ\\pm0.7^\\circ$ in azimuth, providing some insight into the accretion geometry.} {} \\keywords {stars: binaries: novae,cataclysmic variables -- stars: individual: \\ekuma\\ -- X-rays: stars} ",
        "introduction": "Long-term studies of cataclysmic variables and related objects have shown that variations in the orbital period occur in some of them, which are not understood or badly so \\citep[e.g.][]{beuermannpakull84,pandeletal02,schwopeetal02,brinkworthetal06,schwarzetal09}. We set out to search for period changes in other eclipsing and non-eclipsing CVs using remote observations with the MONET/N telescope as part of a secondary school research project.  The polar EK~UMa is the 19 mag optical counterpart of the bright {\\it EINSTEIN\\/} X-ray source 1E1048.5+5421 discovered by \\citet{morrisetal87}. No optical photometry besides that in the discovery paper is in the literature, but an accurate ephemeris based on \\rosat\\ X-ray data was derived by \\citet{schwopeetal95}. We present new optical photometry and supplement the \\rosat\\ X-ray timings by an analysis of the \\euve\\ deep survey DS/S light curves taken in 1994, 1998, and 1999. The combined data allow us to establish a common long-term ephemeris that extends 30 years back to the \\einstein\\ era and accounts for a systematic phase offset between the X-ray absorption dips and the optical cyclotron minima. No period change was found, but we derive a tight limit on $O-C$ variations and associated period changes over the past 17 years. ",
        "conclusions": "We find that the optical and X-ray data of \\ekuma\\ define an ephemeris that is linear over at least 17 years. A quadratic term is insignificant and limits the period variation over this time span to $-1.1<\\Delta P<0.9$\\,ms (2-$\\sigma$ limit). With long-term coverage becoming available for an increasing number of CVs, measuring period variations in the ms or sub-ms range becomes feasible. In the eclipsing pre-CV NN~Ser, \\citet{brinkworthetal06} find $\\dot P=-(9.06\\pm0.06)\\,10^{-12}$\\,s\\,s$^{-1}$ over 45\\,000 cycles and a still higher $\\dot P=-(2.85\\pm0.15)\\,10^{-11}$\\,s\\,s$^{-1}$ over some 3\\,000 cycles. The eclipsing polars DP~Leo \\citep{pandeletal02,schwopeetal02} and HU~Aqr \\citep{schwarzetal09} showed decreases in the orbital periods with $\\dot P\\simeq-5\\,10^{-12}$\\,s\\,s$^{-1}$ over 120\\,000 cycles and $\\dot P=-(7.3\\pm0.5)\\,10^{-13}$\\,s\\,s$^{-1}$ over 60\\,000 cycles, with superposed possibly sinusoidal variations in the latter. A sinusoidal modulation in the $O-C$ values has previously been seen in the eclipsing dwarf nova U~Gem \\citep[e.g.][]{beuermannpakull84}. While star spot cycles may account for the quasi-periodic variations, there is no entirely convincing explanation of the secular period changes emphasizing the need for further studies meant to differentiate between typical behavior and idiosyncrasies. We have demonstrated that subtle period variations -- or their absence -- are easily detected using a small, remotely operated telescope."
    },
    "0911/0911.1336_arXiv.txt": {
        "abstract": "The rotation rate in pre-supernova cores is an important ingredient which can profoundly affect the post-collapse evolution and associated energy release in supernovae and long gamma ray bursts (LGRBs). Previous work has focused on whether the specific angular momentum is above or below the critical value required for the creation of a centrifugally supported disk around a black hole. Here, we explore the effect of the {\\em distribution} of angular momentum with radius in the star, and show that qualitative transitions between high and low angular momentum flow, corresponding to high and low luminosity accretion states, can effectively be reflected in the energy output, leading to variability and the possibility of quiescent times in LGRBs. ",
        "introduction": "\\label{sec:intro}  The collapse of massive stellar cores producing Supernovae (SN) has been clearly linked to the production of Long Gamma Ray Bursts (LGRBs) in the last ten years (see \\citet{wb06,grrf09} for reviews), providing tantalizing clues to the progenitors and the environments in which they occur \\citep{fruchter06}, as well as providing fresh glimpses of the high-redshift universe \\citep[e.g.,][]{prochaska09}. A key ingredient in this context is the stellar rotation rate, impacting the energy release and general outcome of the event following the implosion of the Chandrasekhar-mass iron core. Stellar evolution considerations have shown that it is non-trivial to have rapid rotation in the inner regions of the star, as mass loss and magnetic fields both conspire to reduce the pre-SN angular momentum. Calculations \\citep[hereafter WH06]{spruit02,hws05,yoon05,wh06} indicate that the final distributions of specific angular momentum may not easily lead to the production of a centrifugally supported accretion disk following collapse, with the envelope experiencing essentially radial infall. This has generally been believed to be insufficient to power LGRBs since the lack of shocks and the associated dissipation are unable to transform gravitational binding energy into radiation effectively. A possible evolutionary channel that avoids some of the pitfalls associated with the giant phase was proposed in WH06 and by \\citet{yoon05}, in which thorough mixing on the main sequence in very massive stars reduces mass and angular momentum losses. However, a large class of stars remains in which rotation will be substantially slowed in late evolutionary stages, thus influencing the outcome of the collapse.  The low angular momentum regime in collapsars has not been fully explored. In most cases, studies of the attending neutrino cooled accretion flows have considered specific rotation laws that guarantee by a large margin the formation of a centrifugally supported disk, either because the angular velocity is assumed to be nearly Keplerian, or because the absolute value of the angular momentum given implies a circularization radius much larger than the radius of the innermost stable circular orbit (ISCO), which in General Relativity is $r_{\\rm ISCO}=3 r_{\\rm g}$ for a Schwarzschild black hole, where $M_{\\rm BH}$ is the black hole mass and $r_{\\rm g}=2GM_{\\rm BH}/c^{2}$ \\citep{mw99,pwf99,h00,npk01,pb03,pmab03,lrrp05,fujimoto06,nagataki07}. We previously studied \\citep[hereafter Paper~I]{lrr06,lclrr09}, cases at the threshold for disk formation, where a centrifugally supported disk gives way to nearly radial inflow due to relativistic effects, finding that even in the case of slow rotation, enough energy may be available through the formation of {\\em dwarf} disks \\citep{ib01,zb09} in near free fall to power LGRBs. However, the distributions of specific angular momentum we considered were constant in the equatorial plane, and smoothly decreasing towards the rotation axis. This is unrealistic, as the specific angular momentum {\\em increases} outwards in the core and envelope, with marked transitions at the boundaries between different layers in the star (see, e.g., Figure~2 in WH06).  Further, as the newborn black hole (BH) accretes infalling matter, both its mass and angular momentum increase, raising the threshold value of the critical angular momentum required for the formation of a centrifugally supported disk, $J_{\\rm crit}=2 r_{\\rm g} c$. The competition with the increase of angular momentum in the infalling material determines the large scale properties of the flow. This was recently pointed out and addressed by \\citet{jp08} in simple form, generically by \\citet{knj08a,knj08b} in studying the effect of the distribution of angular momentum in the star on the light curve of LGRBs, and more recently by \\citet{lindner09} in two-dimensional simulations of collapsar accretion.  In this study, we explore how different distributions of angular momentum as a function of radius, can have important effects on the qualitative properties of the accretion flow, and hence on the accretion rate and global energy release, which we quantify through the neutrino luminosity, $L_{\\nu}$. We pay particular attention to the general form and rate of increase of specific angular momentum with radius in the star in this respect, and show that state transitions may in principle produce observable consequences in LGRBs relevant to variability and periods of quiescence. In section \\S~\\ref{sec:input} we describe the assumptions and approximations made in our calculations and identify the numerical setup. Section \\S~\\ref{sec:results} is devoted to presenting the flow transitions which result from different angular momentum distributions, and \\S~\\ref{sec:states} to the conditions necessary for them to occur. In \\S~\\ref{sec:critical} we combine these, given realistic angular momentum distributions, to investigate episodic energy release from accretion. Our conclusions and prospects for observability are discussed in \\S~\\ref{sec:discussion}. ",
        "conclusions": "\\label{sec:discussion}  We wish to stress that while we have presented the neutrino luminosity as a measure of energy output, it is by no means the only one possible, and should be viewed here as a proxy for central engine activity, like the mass accretion rate with which it is very closely correlated. One could equally use $\\dot{M}$ (as long as it occurs in state where energy is dissipated efficiently), or the power output in magnetic outflows as a measure of the ability to drive relativistic outflows. Our numerical scheme is geared towards appropriate handling of thermodynamics and the associated neutrino emission, so it is natural to rely on these properties when making quantitative statements.  It is one thing to have a variation in the energy release close to the BH, and another to have it manifest itself in the light curve of a LGRB. Now, the minimum time required for activity to lead to high energy emission is about 2~s, the crossing time for a relativistic outflow originating in the inner regions to pierce the stellar envelope \\citep[see, e.g.,][]{mw99,aloy00,mwh01,aloy02,ramirezruiz02,zhang03}. The crossing time is estimated here simply as $t_{\\rm cross}=R_{\\rm env}/(c/ \\sqrt 3)$. In Table~\\ref{tab:models}, we present a summary of our results for the models shown in Figure~\\ref{fig:Jdist} and Table~\\ref{tab:inmodels}, computed at the threshold where $\\mu=1/3$ and a quiescent period is marginally present.  After an initial delay ($t_{\\rm delay}$) between the onset of core collapse and the initial rise in neutrino luminosity, a first disk is active for an interval $t_{\\rm d}$, followed by quiescence and late disk creation. Only models with $t_{\\rm d} \\geq t_{\\rm cross}$, in boldface in Table~\\ref{tab:models}, are potentially observable as producing a quiescent period in the high-energy light curve under this scheme. Note that $f$ is of order unity, but was scaled both to larger and smaller values in order to find a condition at the threshold value, $\\mu=1/3$.  \\begin{table}\\centering \\caption{Model evolution time scales. Events with quiescent periods.} \\label{tab:models} \\begin{tabular}{cccccc} \\hline \\hline Model & $t_{\\rm delay}[s] \\ \\tablenotemark{a} $ & $t_{\\rm d}[s]$ & $t_{\\rm q}[s]$ & $t_{\\rm cross}[s]$ & $f$ \\\\ \\hline {\\bf 16TB} & 0.5 & 2.3 & 9.6 &  1.8 & 0.6 \\\\ 16TE & 1.3 & 0.9 & 3.0 &  1.8 & 2.0\\\\ 12ON & 3.5& 0.6 & 1.6 &  1.5 & 1.9\\\\ {\\bf 16TI}  &  2.3 & 1.1 & 4.2 &  1.0 & 2.0\\\\ 12TJ  & 0.7 & 0.7 & 4.1 &  1.9 & 3.9 \\\\ {\\bf 35OA} & 0.5 & 2.4 & 12.6 &  1.5 & 80\\\\ {\\bf 16OH} & 0.4 & 2.1 & 1.4 &  2.0 & 39\\\\ \\hline \\tablenotetext{a}{Delay between onset of collapse and initial rise in $L_{\\nu}$.} \\end{tabular} \\end{table}  The most obvious limitation of the current study is the range of time scales directly modeled. Nevertheless, we believe the current analysis to be of relevance to the generic behavior in collapsing cores. Within this range, the current calculations clearly show that the characteristics of the energy output are closely correlated to the distribution of specific angular momentum in the progenitor star. The state transitions are abrupt, with the luminosity rising or dropping by more than one order of magnitude, and reflect the naturally short (ms) time scales in the vicinity of the accretor. The scaling to larger radii and longer time scales can lead to quiescent period such as those studied by \\citet{ramirezruiz01a}, which in this case would be related to dormant periods in the central engine \\citep{ramirezruiz01b}. \\citet{ramirezruiz01a} found that a small fraction, $\\simeq$~15\\%, of long GRBs exhibit at least one quiescent interval, with about one quarter of these showing two such episodes. Furthermore, the durations of the quiescent and subsequent period of activity were found to be directly correlated. We believe it is possible to account generically for this behavior under the present picture in at least two relevant aspects.  First, the number of shell-like jumps in the distribution of angular momentum within the star is small (one or two), accounting roughly for the number of transitions observed. The second point is related to the correlation itself. A short quiescent interval arises when the low-angular momentum shell has a narrow radial extent. Its inner limit at $r_1$ is fixed by the transition at the edge of the CO core and remains approximately at the same radius, independently of the normalization of angular momentum through the parameter $f$. The outer boundary at $r_2$ on the other hand, varies with $f$ (see Figure~\\ref{fig:Jdist16TI}). Due to the latter and to the fact that $J_{\\rm crit}$ is a monotonically increasing function of radius, higher overall angular momentum will lead to: (i) a short quiescent interval, and (ii) a smaller circularization radius for the bulk of the innermost matter within the centrifugally supported disk, $r_{\\rm circ}[J(r_2)]$. Conversely, when the angular momentum normalization is low, the quiescence period will be longer, and the circularization radius for the majority of the innermost material will be greater. If the evolution of the late--time disk is driven by the viscous transport of angular momentum, the corresponding viscous time scale and the duration of the associated accretion episode will scale essentially with the initial disk radius. Thus, a longer quiescent period will be followed by a longer period of disk activity. The crucial point is that a lower global angular momentum normalization will place the bulk of the initial accreting matter capable of forming a centrifugally supported disk at a larger circularization radius. There is in addition an effective upper limit for this radius if the gas is to release its binding energy effectively, because if it is too large, the density and temperature will not be sufficient for neutrino cooling to operate.  As a final point, we may add that just as not all LGRBs exhibit this behavior, likewise it is clear from this work that not all progenitors are capable of producing the state transitions presented here. In fact, \\citet{pm10} have recently argued that the lack of flaring activity in most LGRBs on long time scales is a signature of complete mixing in the progenitor, which is fully in agreement with the line of argument presented here. Once the deposition of energy in the inner regions of the star has produced a relativistic jet capable of traversing the stellar envelope, further variability may be reproduced in the overall light curve \\citep{zhang03}. Determining whether this can power precursor activity is another matter, requiring the initial episode of accretion to create a low density polar funnel in the star, which remains to be addressed."
    },
    "0911/0911.0795_arXiv.txt": {
        "abstract": "{\\bf binaries: general; stars: general; stars: formation; galaxies: star clusters}Many, possibly most, stars form in binary and higher-order multiple systems. Therefore, the properties and frequency of binary systems provide strong clues to the star-formation process, and constraints on star-formation models. However, the majority of stars also form in star clusters in which the birth binary properties and frequency can be altered rapidly by dynamical processing. Thus, we almost never see the birth population, which makes it very difficult to know if star formation (as traced by binaries, at least) is universal, or if it depends on environment. In addition, the field population consists of a mixture of systems from different clusters which have all been processed in different ways. ",
        "introduction": "Observations suggest that a significant fraction of stars (perhaps most) {\\em in the field} are in binary or multiple systems\\footnote{Stars appear to have multiplicities from binaries to septuples (e.g., Eggleton \\& Tokovinin 2008). For brevity we shall use `binary' to mean systems of any multiplicity, only drawing a distinction when it is required.} (see, e.g., Duquennoy \\& Mayor 1991; Fischer \\& Marcy 1992; Lada 2006; Eggleton \\& Tokovinin 2008). It seems to be impossible to dynamically produce binaries in anywhere near the numbers observed, and so the vast majority of binaries must have formed as binaries (Goodman \\& Hut 1993).  As we have seen earlier in this volume (see Clarke 2010; de Grijs 2010; Lada 2010), a significant fraction of stars appear to form in star clusters (see also Lada \\& Lada 2003). Because of their high densities, clusters are regions in which binary systems are likely to be altered (either destroyed or changed) by dynamical processes. Therefore, it is highly likely that far more binaries were formed than are now observed, and even those binaries that survive may have different properties at the present time compared to when they formed.  The field is the sum of all star formation. As much of that star formation was clustered, we expect that the field binary population has undergone some (very probably significant) dynamical processing in their birth clusters. Therefore, it is important to remember that the field binary population is {\\em not} the birth binary population. The field is a mixture of systems from different environments, each of which will have been processed to some degree.  The dynamical processing of binaries in clusters will also alter the numbers and properties of various types of interesting astrophysical systems such as blue stragglers, low- and high-mass X-ray binaries, type Ia supernovae, and intermediate-mass black holes (Hut \\ea 1992). However, in this review we will concentrate on `typical' star formation that results in the bulk of the field, i.e., relatively low-mass clusters (10--$10^5$ M$_\\odot$) that disperse (are destroyed?) within a few~Myr (see Lada \\& Lada 2003). Thus, we will ignore the extremely interesting, but relatively rare cases (at least in the Galaxy) of extremely massive young clusters, or extremely long-lived clusters in which many interesting dynamical processes involving binaries occur.  In this contribution we will review our current understanding of two key astrophysical problems: the universality of star formation and the origin of the field. Binaries are an excellent tool with which to attempt to answer these questions, as binary formation is a fundamental and very common (possibly universal?) outcome of star formation and so similarities and differences between binary populations are indicators of the similarities and differences in the star-formation process in different environments.  First, does star formation care about its environment (see also Clarke 2010; Lada 2010)? Is the outcome of star formation in cores of a particular mass (at the end of the class 0/I phase) always (statistically) the same? The initial mass functions (IMFs) of different regions often appear very similar, but binary properties are probably a far more detailed indicator of the similarity or otherwise of star formation in different regions (Goodwin \\& Kouwenhoven 2009).  Second, what is the origin of the field-star population? The field is the sum of star formation in high- and low-mass clusters and isolated star formation. Do we understand the origin of the field?  To attempt to answer these questions we will concentrate on local regions and clusters for which we have detailed observations down to low (often substellar) masses. Unfortunately, such clusters are of low mass and often low density and the applicability of extending the conclusions drawn from local star-forming regions to more extreme star formation in massive young clusters and starbursts is debatable.  In \\S2 we examine the observations of binary systems in the field and in different clusters. In \\S3 we discuss how binaries can be dynamically processed in clusters and how this alters the birth binary population. In \\S4 we will investigate what we can infer about the birth binary populations and if they are universal and discuss the origin of the field binary population. We conclude in \\S5. ",
        "conclusions": "We initially asked two fundamental questions related to star formation:  \\noindent 1. Is star formation (as probed by binary properties, at least) universal, or does it depend on environment?\\\\ \\noindent 2. What is the origin of the field binary population?  As we have seen, the properties of binaries in different environments {\\em are} different. In particular, dense clusters have fewer binaries and those that they have tend to be close or intermediate binaries ($<1000$~au). However, this difference can be explained by dynamical processing of the initial binary population in a cluster.  If a cluster such as the ONC did form a significant wide binary population, it would have been destroyed by now. Low-density star-forming regions may provide a clue to the birth properties of stars, but only if one believes that star formation is universal. The difference between Taurus and the ONC is striking, but can be explained by both a different birth population or dynamical processing, or a mixture of both (dynamical processing {\\em must} have occured in the ONC). Of particular interest in this regard are observations of USco, especially the differences between USco A and B, both of which are thought to have formed at low density but have very different binary properties. This may indicate differences in the initial conditions of the cloud. Could this be due to triggered star formation (see Preibisch \\ea 2002)?  It is crucial to remember that we will {\\em never} see an intact birth population, even in a young cluster. The only way to see a potentially intact birth system is to examine the deeply embedded phases of star formation. But even then, significant accretion and fragmentation to form binaries is ongoing during the class 0 phase, so that we will see `unfinished' systems. However, by the class I phase, dynamical decay can have occurred. Does this mean that the `birth population' is a meaningless phrase?  One of the few statements that we can make without argument is that {\\em the field is not the birth population}. The field is a mixture of potentially different birth populations which have been processed to different degrees in their birth clusters and then mixed.  Dynamical processing destroys binary systems. Therefore, there {\\em must} have been more binaries formed than we see in the field (of course, if star formation is not universal, then some regions may form like the field, but not all).  Without an understanding of the universality or otherwise of star formation in different environments, it is impossible to constrain the origin of the field population. As we have seen, it is possible to explain the field binary population as the result of a universal mode of star formation that has been processed differently in different-density environments. Equally, it is possible to explain it as the sum of many different modes of star formation, each of which was then processed in different ways (if this is the case, then we have a vast parameter space of possible answers). For simplicity, we would probably prefer that (binary) star formation is universal. However, this might well not be the case."
    },
    "0911/0911.5453_arXiv.txt": {
        "abstract": "We report results from tests of \\Kr\\, as a calibration source in liquid argon and liquid neon. \\Kr\\, atoms are produced in the decay of $^{83}$Rb, and a clear \\Kr\\, scintillation peak at 41.5 keV appears in both liquids when filling our detector through a piece of zeolite coated with $^{83}$Rb. Based on this scintillation peak, we observe 6.0 photoelectrons/keV in liquid argon with a resolution of $6\\%$ ($\\sigma$/E) and 3.0 photoelectrons/keV in liquid neon with a resolution of $19\\%$ ($\\sigma$/E). The observed peak intensity subsequently decays with the \\Kr\\, half-life after stopping the fill, and we find evidence that the spatial location of \\Kr\\, atoms in the chamber can be resolved. \\Kr\\, will be a useful calibration source for liquid argon and neon dark matter and solar neutrino detectors. ",
        "introduction": "\\label{sec:intro} Liquefied nobles gases are widely used as targets in low background searches, particularly in direct dark matter searches where a weakly interacting massive particle (WIMP) may scatter elastically with a nucleus to produce a nuclear recoil in the liquid. Several WIMP-nucleon cross-section limits have been set in recent years using liquid argon and xenon, and several larger, more sensitive argon and xenon detectors are currently under construction~\\cite{Angle:2007,Alner:2005,Lebedenko:2009,Brunetti:2005}. Liquid neon is another target attracting interest, having been proposed as a potential target for {\\it pp}-solar neutrinos in addition to dark matter~\\cite{McKinsey:2000,Boulay:2006}. In particular, the MiniCLEAN detector which will be deployed at SNOLAB in the Creighton mine in Sudbury, Ontario will run in both liquid argon and liquid neon phases~\\cite{McKinsey:2007}.  Because the expected dark matter signal has energies of tens of keV and drops exponentially with increasing energy, a dark matter detector must be calibrated at low energy to precisely determine its energy threshold and ultimate WIMP sensitivity. Low energy calibrations are also important for a liquid neon {\\it pp} neutrino detector because any uncertainty in the energy scale produces a systematic error in the observed neutrino energy spectrum and flux. As liquid noble gas detectors get larger, self-shielding will render it increasingly difficult to illuminate the central volume of the liquid with external gamma rays, particularly at low energies. Therefore, a low-energy radioactive source that can be distributed throughout the detector volume is highly desirable for calibrating the new generation of liquid noble gas detectors. One example of such a calibration was the introduction of activated xenon isotopes into the XENON10 detector. A sample of xenon was irradiated at Yale University to produce $^{129}$Xe$^{\\mathrm{m}}$ and $^{131}$Xe$^{\\mathrm{m}}$, which emit 236 keV and 164 keV gamma rays with half-lives of 8.9 and 11.8 days, respectively. This sample was shipped to Gran Sasso National Laboratory in Italy and introduced into the detector to provide a uniform energy calibration~\\cite{Ni:2007}. While this calibration was successful, the energies of these isotopes are higher than those expected from a WIMP signal and their relatively long half-lives limit the frequency of deployment.  An alternative is the use of \\Kr, which is used as a diagnostic tool for studying the beta decay of tritium and has been proposed for use with the KATRIN detector~\\cite{Katrin:2001}. \\Kr\\ is produced in the decay of $^{83}$Rb, which has a half-life of 86.2 days. In turn, the \\Kr\\, decays with a half-life of $(1.83\\pm0.02)$ hours, emitting 32.1 keV and 9.4 keV conversion electrons (see Fig.~\\ref{fig:KrLevels})~\\cite{Venos:2005,Venos:2006,Wu:2001}. As a noble gas, \\Kr\\, is easily introduced into noble liquid detectors without compromising the purity of the liquid, and it is not expected to create any long-lived radioisotopes.  Recently, \\Kr\\, has been successfully used to calibrate liquid xenon detectors~\\cite{Kastens:2009,Manalaysay:2009}. Because argon and neon are liquids at temperatures below the freezing point of krypton, any \\Kr\\, atoms in the liquid could potentially freeze out on detector surfaces before decaying. This paper describes successful tests of \\Kr\\, as a calibration source in both liquid argon and liquid neon.  While some \\Kr\\, atoms may be freezing out, enough reach the center of the liquid volume to provide a good energy calibration.  \\begin{figure}[htbp] \\centerline{ \\hbox{\\psfig{figure=Fig1.eps,width=6.5cm,angle=270,% clip=}}} \\caption[Scintillation cell] {Energy levels in keV of $^{83}$Rb. $75\\%$ of the time, $^{83}$Rb decays to \\Kr\\,, which in turns emits two conversion electrons at 32.1 and 9.4 keV. The half-life of \\Kr\\, to decay via the first electron to the $7/2^{-}$ state is 1.83 hours and the half-life for the subsequent decay to the stable $9/2^{+}$ state of $^{83}$Kr is 154 ns.} \\label{fig:KrLevels} \\end{figure} ",
        "conclusions": "The results discussed here show that \\Kr\\, is readily introduced into both liquid argon and liquid neon volumes and could serve as a useful calibration to characterize the scintillation signal yield of liquid argon and neon detectors at low energies. While some \\Kr\\, may be freezing onto the walls, enough atoms reach the central volume to be clearly observed. A detector with x-y position reconstruction could potentially observe \\Kr\\, atoms frozen to the walls as an exterior ring of activity in the detector. The current design for MiniCLEAN calls for a continuous purification loop, and a $^{83}$Rb trap could be easily included in that design. Alternatively, it appears possible to run without directly flowing gas through the trap, although this mode is less efficient.  In addition, the asymmetry results show that it might be possible to spatially resolve krypton atoms as they enter the detector, providing a handle on flow, mixing and the spatial resolution of a large detector."
    },
    "0911/0911.4872_arXiv.txt": {
        "abstract": "Proper characterization of the host star to a planet is a key element to the understanding of its overall properties. The star has a direct impact through the modification of the structure and evolution of the planet atmosphere by being the overwhelmingly larger source of energy. The star plays a central role in shaping the structure, evolution, and even determining the mere existence of planetary atmospheres. The vast majority of the stellar flux is well understood thanks to the impressive progress made in the modeling of stellar atmospheres. At short wavelengths (X-rays to UV), however, the information is scarcer since the stellar emission does not originate in the photosphere but in the chromospheric and coronal regions, which are much less understood. The same can be said about particle emissions, with a strong impact on planetary atmospheres, because a detailed description of the time-evolution of stellar wind is still lacking. Here we review our current understanding of the flux and particle emissions of the Sun and low-mass stars and briefly address their impact in the context of planetary atmospheres. ",
        "introduction": "Our knowledge of planets, including our Earth, is directly driven by our knowledge of the parent star. Indeed, the host star to a planet is the overwhelmingly larger source of energy and its emissions critically affect the composition, thermal properties and the mere existence of planetary atmospheres. Our current description of the internal structure, evolution and radiative flux of stars has reached a relatively mature level and theoretical models can now predict stellar attributes to fairly good accuracy. While this refers to the bulk properties of stars, some other components related to stellar emissions still defy detailed understanding. This is the case of the emissions related to magnetic activity, a common feature of all low-mass stars. High-energy emissions of solar-type stars and their variations over different timescales are now quite well characterized but not so active stars in general. The situation for particle winds is even more frustrating, with a lack of direct detections and large uncertainties on the predicted wind of active stars from indirect measurements. However, all evidence collected so far supports the fact that low-mass stars undergo an early phase of strong magnetic activity (lasting from tens of Myrs to Gyrs) in which their high-energy and particle emissions are enhanced by orders of magnitude with respect to our current Sun. If the effects of the active Sun are even noticeable today on our Earth, it is undeniable that such effects {\\em must} have been even stronger for our planet and all Solar System planets in the past. In general, all planets orbiting low-mass stars are affected by the magnetic evolution of their parent star.  In the following sections we review the current state of knowledge of the properties of stars from the point of view of planet hosts, including their bulk emissions and those related with magnetic activity. Further, we discuss the possible impact that long-term variability of high-energy radiations and particle fluxes has on the Solar System planets and exoplanets around solar-type and lower-mass stars. ",
        "conclusions": "Our planet Earth orbits around a dependable star that has provided a suitable environment for life to thrive. But the apparent stability of the Sun is just an illusion. Thanks to our vantage point we are able to scrutinize the solar properties with exquisite detail. Centuries of observation have shown a variety of effects related to the magnetic activity that reveal our Sun as a variable star. Sunspots, faculae, coronal mass ejections, are all different facets of the active Sun. Firsthand experience on Earth has demonstrated that such phenomena have a powerful influence on our planet's upper atmosphere and, perhaps, even on our climate. But consider the current Sun with an increase of its high-energy and particle emissions by two or three orders of magnitude. This would have been the infant Sun. Our planet and the rest of its Solar System siblings had to endure an epoch of heavy irradiations early in their history. The atmosphere of Mars, once it lost its magnetic shielding, succumbed to the intense erosion. The strong emissions also left their imprint in the isotopic ratios of the atmospheres of Venus and Titan, for example. Fortunately, the mass of the Earth and its magnetic protection helped our planet to keep its volatile inventory and eventually allowed for the development of complex life. While our Sun's wild ages lasted for a few hundred million years, other less massive stars may have kept the strong emissions for much longer. Future research will tell us what could happen to a terrestrial planet orbiting a red dwarf star, which stays active for billions of years, and whether habitability can exist in such hostile environment."
    },
    "0911/0911.4569_arXiv.txt": {
        "abstract": "The  preliminary results of new photometric observations of the $\\delta$ Scuti stars 7 Aql and 8 Aql are reported. 51 hr of $uvby-\\beta$ photoelectric photometric data were obtained over the period June and July 2007 at the San Pedro M\\'artir Observatory, Mexico. Period analyses confirm the three  pulsation modes discovered in 8 Aql in the framework of the STEPHI 2003 multisite campaign. For the star 7 Aql we were able to detect only the main pulsation modes. The standard magnitudes of both stars are obtained. The frequency, amplitude and phases of the frequency modes in different filters are presented. ",
        "introduction": "7~Aql (HD~174532, SAO~142696, HIP~92501) is a $\\delta$~Scuti variable discovered in a systematic search and characterization of new variables in preparation for the COROT mission [1] and was selected as the main target of the STEPHI XII multisite campaign in 2003. In this campaign 8~Aql (HD~174589, SAO~142706, HIP~92524) was used as the only comparison star because there are no other bright stars in the field-of-view (FOV) of the 4-channel photometer used in the STEPHI network.  \\medskip Although no check comparison star was implemented in that campaign, by carefully analyzing the derived light curves, differential and non-differential,  it was demonstrated that 8 Aql is a new $\\delta$ Scuti variable. Moreover, it was shown that the amplitude spectrum of both stars are not superposed. Three and seven frequency peaks were unambiguously detected with a 99 \\% confidence level  in 8 Aql and 7 Aql respectively and a possible identification of the observed modes in terms of radial order was performed [2]. CCD photometric observations of 7 Aql and 8 Aql were reported in [3].  \\medskip In the present paper,  preliminary result of new  photoelectric photometric observations of 7 Aql and 8 Aql are reported. ",
        "conclusions": ""
    },
    "0911/0911.1982_arXiv.txt": {
        "abstract": "We explore the biological damage initiated in the environments of F, G, K, and M-type main-sequence stars due to photospheric, chromospheric and flare radiation.  The amount of chromospheric radiation is, in a statistical sense, directly coupled to the stellar age as well as the presence of significant stellar magnetic fields and dynamo activity.  With respect to photospheric radiation, we also consider detailed synthetic models, taking into account millions or hundred of millions of lines for atoms and molecules. Chromospheric UV radiation is increased in young stars in regard to all stellar spectral types.  Flare activity is most pronounced in K and M-type stars, which also has the potential of stripping the planetary atmospheres of close-in planets, including planets located in the stellar habitable zone. For our studies, we take DNA as a proxy for carbon-based macromolecules, guided by the paradigm that carbon might constitute the biochemical centerpiece of extraterrestrial life forms.  Planetary atmospheric attenuation is considered in an approximate manner. ",
        "introduction": "The centerpiece of all life on Earth is carbon-based biochemistry. It has repeatedly been surmised that biochemistry based on carbon may also play a pivotal role in extraterrestrial life forms, if existent. This is due to the pronounced advantages of carbon, especially compared to its closest competitor (i.e., silicon), which include: its relatively high abundance, its bonding properties, and its ability to form very large molecules as it can combine with hydrogen and other molecules as, e.g., nitrogen and oxygen in a very large number of ways (\\cite[Goldsmith \\& Owen 2002]{gol02}).  In the following, we explore the relative damage to carbon-based macromolecules in the environments of late-type main-sequence stars using DNA as a proxy. We focus on the effects of both photospheric and chromospheric radiation and comment on the significance of flare activity.  Previous studies indicating the importance of UV radiation concerning the suitability of extrasolar planets for the existence of life and biological evolution has been given by, e.g., \\cite{coc99}, \\cite{gui02}, \\cite{gui03}, \\cite{rib05}, \\cite{buc06}, and \\cite{gui09}.  With respect to the Sun-Earth system, a detailed investigation about the UV effects on biological systems has been pursued by \\cite{dif91}. This line of studies also includes the detailed modeling of the atmospheric attenuation of UV (especially due to an Earth-type ozone layer, e.g., \\cite[Segura et al. 2003]{seg03}), photobiological effects and the effects of UV on biomolecules, particularly DNA, and microorganisms.  \\begin{figure*} \\vspace*{+3.5 cm} \\centering \\epsfig{file=fig1.eps,width=0.5\\linewidth,height=0.45\\linewidth} \\vspace*{-3.0 cm} \\caption{DNA action spectrum following \\cite{hor95} and \\cite{coc99}. } \\end{figure*}  \\begin{table} \\caption{Stellar Data and Habitable Zones} \\centering \\begin{tabular}{l c c c c c c c c} \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} Sp. Type &  $T_{\\rm eff}$ &  $R_\\ast$ & $L_\\ast$ & HZ-i (gen.) & HZ-i (con.) & HZ (E.) & HZ-o (con.) & HZ-i (gen.) \\\\ \\noalign{\\smallskip} ...   &  (K) & ($R_\\odot$) & ($L_\\odot$) & (AU) & (AU) & (AU) & (AU) & (AU) \\\\ \\hline \\noalign{\\smallskip} F0~V  &   7200  &   1.62  &  6.33  &  1.826 &  2.251 &  2.516 &  3.222 &  3.710  \\\\ F5~V  &   6440  &   1.40  &  3.03  &  1.370 &  1.614 &  1.740 &  2.316 &  2.741  \\\\ G0~V  &   6030  &   1.12  &  1.49  &  1.002 &  1.152 &  1.220 &  1.657 &  1.991  \\\\ G5~V  &   5770  &   0.95  &  0.90  &  0.799 &  0.904 &  0.948 &  1.302 &  1.580  \\\\ K0~V  &   5250  &   0.83  &  0.47  &  0.604 &  0.665 &  0.685 &  0.961 &  1.188  \\\\ K5~V  &   4350  &   0.64  &  0.13  &  0.342 &  0.359 &  0.363 &  0.525 &  0.670  \\\\ M0~V  &   3850  &   0.48  &  0.05  &  0.207 &  0.213 &  0.213 &  0.313 &  0.407  \\\\ \\hline \\noalign{\\smallskip} \\multicolumn{9}{l}{Abbreviations: con. = conservative, gen. = general, E. = Earth-equivalent} \\\\ \\end{tabular} \\end{table}  We consider the radiative effects on DNA by applying a DNA action spectrum (\\cite[Horneck 1995]{hor95}) that shows that the damage is strongly wavelength-dependent, increasing by more than seven orders of magnitude between 400 and 200~nm (see Fig.~1).  The different regimes are commonly referred to as UV-A (400--320 nm), UV-B (320--290 nm), and UV-C (290--200 nm).  The test planets are assumed to be located in the stellar habitable zone (HZ). Following the analyses by \\cite{kas93} and \\cite{und03}, we distinguish between the conservative and generalized HZ (see Table~1).  For stars of different spectral types, we also define planetary Earth-equivalent positions given as $R_{\\rm E} = \\sqrt{L/L_\\odot}$.  For the conservative HZ, the inner limit of habitability is given by the onset of water loss that occurs in an atmosphere warm enough to have a wet stratosphere, resulting in a gradual loss of water by photodissociation and subsequent hydrogen loss to space.  Moreover, the outer limit of habitability is given by the first CO$_2$ condensation where for a surface temperature of 273~K, CO$_2$ begin to form.  Concerning the general HZ, the inner limit of habitability is given by the runaway greenhouse effect.  In the following, we will present our results for different positions of Earth-type planets while considering both stellar photospheric and different levels of chromospheric radiation.  We will also discuss the likely consequences of stellar flares.  \\begin{figure*} \\begin{tabular}{cc} \\epsfig{file=fig21.eps,width=0.48\\linewidth,height=0.38\\linewidth} & \\epsfig{file=fig22.eps,width=0.48\\linewidth,height=0.38\\linewidth} \\\\ \\epsfig{file=fig23.eps,width=0.48\\linewidth,height=0.38\\linewidth} & \\epsfig{file=fig24.eps,width=0.48\\linewidth,height=0.38\\linewidth} \\end{tabular} \\caption{Kurucz models for the F0~V, G0~V, K0~V, and M0~V stars used in this study.  For the sake of display, the spectral resolution has artificially been decreased to ${\\lambda}/{\\Delta\\lambda} \\simeq 300$, implemented by running means. Note the different scales of the $y$-axes. } \\end{figure*}  \\begin{figure*} \\centering \\epsfig{file=fig3.eps,width=0.5\\linewidth,height=0.45\\linewidth} \\caption{Biological damage to DNA for a planet (no atmosphere) at an Earth-equivalent position (solid line), and at the inner and outer limits of the conservative (dashed lines) and generalized HZ (dotted lines). } \\end{figure*}  \\begin{figure*} \\centering \\epsfig{file=fig4.eps,width=0.5\\linewidth,height=0.45\\linewidth} \\caption{Biological damage to DNA for a planet at an Earth-equivalent position without an atmosphere (solid line), and an atmosphere akin to Earth 3.5 Gyr ago (dashed line) and today (dotted line).  The dash-dotted line refers to a planet without an atmosphere at 1~AU from its host star, irrespectively of the position of the stellar HZ. } \\end{figure*} ",
        "conclusions": "We attempted to quantify the biological damage expected to occur in the environments of late-type main-sequence stars due to photospheric, chromospheric and flare radiation.  Stellar photospheric radiation has been considered by utilizing spectral models given by R.~L.~Kurucz and collaborators, which take into account millions or hundred of millions of lines for atoms and molecules; see, e.g., \\cite{cas04}.  Concerning chromospheric radiation, we considered models exemplifying basal chromospheric emission (e.g., \\cite[Schrijver 1987]{sch87}, \\cite[Rutten et al. 1991]{rut91}) and significantly increased chromospheric emission (e.g., \\cite[Vilhu \\& Walter 1987]{vil87}). Note that basal emission is usually attributed to acoustic heating in the limiting case of magnetic energy dissipation to be small or negligible.  Detailed models of one-component and two-component chromospheric heating for late-type stars have been given by, e.g., \\cite{buc98}, \\cite[Cuntz et al. (1998, 1999)]{cun98,cun99}, \\cite{faw02}, and \\cite{ram03}.  It is highly noteworthy that there is a deeper underlying connection between the level of chromospheric heating and emission, on one hand, and more fundamental stellar physics involving processes concerning the stellar interior and phenomena associated with stellar winds, on the other hand.  In case of the Sun, this connection has been explored in detail by, e.g., \\cite{gui02}, \\cite{gui03}, \\cite{rib05}, and \\cite{gue07}.  A highly crucial question, as gauged by its implied planetary, biological and societal percussions, is whether the Sun has always been a relatively inactive star or, in contrast, has experienced some periods of stronger magnetic activity.  Compelling observational evidence (\\cite[G\\\"udel et al. 1997]{gue97}) shows that zero-age main-sequence (ZAMS) solar-type stars rotate over 10 times faster than today's Sun.  As a consequence of this, young solar-type stars, including the young Sun, have vigorous magnetic dynamos and correspondingly strong high energy emissions.  From the study of solar-type stars of different ages, \\cite{sku72}, \\cite{sim85}, \\cite{cha93}, \\cite{mac94} and others showed that the Sun loses angular momentum with time via magnetized winds (magnetic breaking), thus leading to a secular increase in its rotational period (e.g., \\cite[Durney 1972]{dur72}, \\cite[Keppens et al. 1995]{kep95}).  This rotation slowdown is well fitted by a Skumanich-type power law roughly proportional to $t^{-1/2}$ (e.g., \\cite[Skumanich 1972]{sku72}, \\cite[Soderblom 1982]{sod82}, \\cite[Ayres 1997]{ayr97}).  In response to slower rotation, the solar dynamo strength diminishes with time, causing the Sun's high energy emissions and mass loss also to undergo significant decreases.  A direct consequence of this behavior concerns the generation of photospheric and chromospheric magnetic flux, which constitute the physical reason for the varying level of chromospheric heating and associated generation of UV and EUV radiation with stellar age or rotation period (e.g., \\cite[Mathioudakis et al. 1995]{mat95}). Consequently, it is thus also possible to relate the photospheric magnetic flux to the stellar rotation period (\\cite[Noyes et al. 1984]{noy84}; \\cite[Marcy \\& Basri 1989]{mar}; \\cite[Montesinos \\& Jordan 1993]{mon93}; \\cite[Saar 1996a]{saa96a}) as well as to the emergent chromospheric emission (\\cite[Saar \\& Schrijver 1987]{saa87}; \\cite[Schrijver et al. 1989]{sch89}; \\cite[Montesinos \\& Jordan 1993]{mon93}; \\cite[Saar 1996b]{saa96b}; \\cite[Jordan 1997]{jor97}).  The acquisition, interpretation and modeling of data for any of these phenomena is the main focus of the ``Sun in Time\" project (originally solely focused on G0--G5~V stars) pursued under the leadership of E.~F. Guinan, which relates, in a statistical sense, the strength of the emergent UV, EUV and X-ray flux and the stellar mass loss rate to the stellar age.  More recently, this project has been extended to K-type and M-type dwarfs with the latter element of study now referred to ``Living with a Red Dwarf\". The critical ultimate connection to be targeted in the future, see \\cite{gui09} for the dissemination of various intermediate results, is to relate the type and strength of stellar activity for different types of stars to planetary climatology, particularly for planets in stellar HZs, and to exosolar planetary biology, if existing."
    },
    "0911/0911.1666_arXiv.txt": {
        "abstract": "{The metal content of the spiral galaxy M33 is analyzed through spectroscopic observations of its emission-line populations, \\hii\\ regions and Planetary Nebulae (PNe): their abundance gradients are identical within the errors. The 2D metallicity map is presented, finding an off-center peak, located in the southern arm. A chemical evolution model of M33 described by Magrini et al. (2007a) is updated at the  light of recent results on the Schmidt law and the metallicity gradients. ",
        "introduction": "M33 (NGC~598) is a spiral galaxy whose closeness (840 kpc, Freedman et al. 1991; optical size 53' $\\times$53', Holmberg 1958) and inclination (i=53 \\deg) allow detailed studies of its stellar populations and ionized nebulae. Planetary Nebulae (PNe) and \\hii\\ regions represent two very different stages in the lifetime of a galaxy, and their comparison provides insight on the chemical evolution of the host galaxy. PNe are indeed the final ejecta of evolved low- and intermediate-mass stars with mass between 1 and 8 M$_{\\odot}$, which must have formed between 3$\\times$10$^7$ yr and 10 Gyr ago, e.g., Maraston \\cite{maraston05}, while \\hii\\ regions belong to a very young population.  One of the most debated questions about the chemical evolution of galaxies is the metallicity gradient behaviour with time.  Chemical evolution models predict different temporal behaviors of the metallicity gradient, depending on the assumptions one makes on gas inflow and outflow rates, and the star and cloud formation efficiencies.  Observations are needed to constrain these assumptions, but so far they have been insufficient, especially for the older populations. The idea behind the observations leading to our work is to study the chemical and physical properties of a large number of PNe and \\hii\\ regions in M33, using the same set of observations, the same data reduction and analysis techniques, and identical abundance determination methods, to avoid all biases due to the stellar vs. nebular analysis. The aim is to derive abundances of the $\\alpha$-elements for as many PNe and \\hii\\ regions as possible, and to study the variation of the metallicity gradient and the average abundances in M33.   In this paper, several aspects of the study of M33 will be described. In Section \\ref{sect_mmt} the observations of a sample of PNe and \\hii\\ regions with MMT will be presented. These data, together with a large literature data-set, allow us to recompute the PN and \\hii\\ region metallicity radial gradients. In Section \\ref{sect_distr} the spatial distribution of the metallicity will be shown. The off-center of the metallicity maps obtained from \\hii\\ regions and from PNe  will be discussed. In Section \\ref{sect_model} the chemical evolution model of M33 (Magrini et al. 2007a, hereafter M07) will be revised, including a star formation process, parametrized by a Schmidt law, consistent with the observations. The consequences, in particular the time evolution of the metallicity gradient, will be analyzed. Finally, our conclusions and summary will be given in Section \\ref{sect_conclu}. ",
        "conclusions": "\\label{sect_conclu} The chemical evolution of M33 is studied by means of new spectroscopic observations of PNe and \\hii\\ regions. Their 2D metallicity maps have been found both off-centered, with a peak in the southern arm, at 1-2 kpc from the center. This might be related to the absence  of a dominant gravitational source in the center. The slow evolution of the metallicity gradient from the present time to the birth of the PNe progenitors is explained with an {\\em accretion} model where the SF process is driven by the Schmidt law."
    },
    "0911/0911.3380_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\setcounter{equation}{0}  Inflation \\cite{Guth:1980zm} has become the leading paradigm of the early universe. However, the detailed dynamics of inflation is still a mystery. A major theme in cosmology is to build inflationary models and compare their predictions with experimental data. The theoretical predictions of general single field inflation models have been well understood. Nonetheless, there exists another important possibility that the inflationary dynamics can involve multiple fields. This leads to a variety of new models and interesting phenomenologies.  When multiple fields are involved, the field space can be decomposed into the inflationary direction and isocurvature directions. The quanta in these field directions are called inflaton and isocurvatons, respectively. In this paper, we investigate a class of models where there is one flat slow-roll direction, and all the other isocurvature directions have mass at least of order the Hubble parameter $H$. We call this class of models {\\em quasi-single field inflation}. If the inflaton decouples from the isocurvatons or the isocurvaton mass are all much larger than $\\CO(H)$, quasi-single field inflation makes the same prediction as the single field inflation. However, once large couplings exist and the mass are of order $H$, we will show that these massive isocurvatons can have important effects on density perturbations.  In this paper, we shall study a simple model of quasi-single field inflation \\cite{Chen:2009we}. In this model, the coupling between the inflaton and the massive isocurvaton is introduced by a turning trajectory. The tangential direction of this turning trajectory is the usual slow-roll direction, while the orthogonal direction is lifted by a mass of order $H$. See Fig.~\\ref{Fig:turn}.   \\begin{figure}[htpb] \\begin{center} \\epsfig{file=figs/turn.eps, width=8cm} \\end{center} \\caption{\\label{Fig:turn}This figure illustrates a model of quasi-single field inflation in terms of turning trajectory. The $\\theta$ direction is the inflationary direction, with a slow-roll potential. The $\\sigma$ direction denotes the isocurvature direction, which typically has mass of order $H$.} \\end{figure}   The motivations for investigating quasi-single field inflation are as follows.  $\\bullet$ {\\em UV completion and fine-tuning in inflation models.} To satisfy the conditions for inflation, fine-tunings or symmetries should generally be evoked. At least this is found to be the case for models that have reasonable UV completion in string theory and supergravity\\cite{Copeland:1994vg,Chen:2008hz}. For slow-roll inflation, this means that, in the inflationary background, the light fields will typically acquire mass of order the Hubble parameter $H$, which is too heavy to be the inflaton candidates. On the other hand, in a UV completed theory, multiple light fields arise naturally. Taking these facts into consideration, a natural picture of inflation emerges: There is one inflation direction with mass $m\\ll H$, and some other directions in the field space with $m \\sim H$. In contrast to the slow-roll potential, large higher order terms in the potential such as $V'''\\sim H$ and $V''''\\sim 1$ can arise naturally in these non-flat directions. To have more than one flat direction needs extra fine-tuning. The above picture for inflation suggests the quasi-single field inflationary models.  $\\bullet$ {\\em Inflationary phenomenology.} When the inflaton trajectory turns in the field space, isocurvature perturbation is converted to curvature perturbation. Precisely understanding the effect of such a conversion on density perturbations is in general an important but difficult question. Quasi-single field inflation provides explicit examples where such a conversion can be calculated from first principles and give non-trivial predictions, such as new shapes of large non-Gaussianities and running of the density perturbations.  $\\bullet$ {\\em Filling the gap between single and multiple field inflation.} Quasi-single field inflation fills the gap between single field inflation and multi-field inflation, with new observational consequences. The relation between single field inflation, quasi-single field inflation and multiple field inflation is illustrated in Fig. \\ref{quasicompare}.  \\begin{figure}[thpb] \\centering \\epsfig{file=figs/regimes.eps, width=8cm} \\caption{\\label{quasicompare} The relationship between single field inflation, quasi-single field inflation and multi-field inflation. Here we take inflation evolving two fields as an example. Quasi-single field inflation fills the gap between single field inflation and multiple field inflation.} \\end{figure}   $\\bullet$ {\\em Non-Gaussianity.} Non-Gaussianity is potentially one of the most promising probes of inflation. Generally speaking, observing shapes and running of non-Gaussianities can tell us about the details of inflatons and isocurvatons during inflation. In the most well-studied models, scale-invariant non-Gaussianities usually take either local shape or equilateral shape. In quasi-single field inflation, large non-Gaussianities with a family of new shapes lying between the local and equilateral shape can arise naturally, and they are potentially observable in the near future. This provides new opportunities for interactions between model building and data analyses.  $\\bullet$ {\\em Methodology.} We compute the correlation functions of the curvature perturbation using the in-in formalism \\cite{inin,Weinberg:2005vy}. The methodology used here is inspiring in several aspects. Firstly, we use the transfer vertex to account for the transformation from the isocurvature mode to curvature mode, and use Feynman diagrams to compute the correlation functions perturbatively. As a result, the leading order contribution comes from the fourth (for bispectrum) or higher (for trispectrum) order perturbation theory. Such methods can have broader applications in, for example, multi-field inflation models. Secondly, we will use two representations of the in-in formalism, i.e.~the factorized form and commutator form. We find that each has its computational advantages and disadvantages and they are complementary to each other. Based on this, we develop a ``mixed form'' of the in-in formalism, which has the advantages of both forms.   This paper is organized as follows.  In Sec.~\\ref{Sec:model}, we study in detail an explicit model of quasi-single field inflation with one isocurvature direction, proposed in \\cite{Chen:2009we}. We investigate the zero-mode solution, the perturbation theory, solve the mode functions and discuss the transfer vertex (as illustrated in Fig.~\\ref{Fig:fdiag}(a)) between the inflaton and isocurvaton.  In Sec.~\\ref{Sec:power}, we study the power spectrum of this model. The power spectrum receives correction from two transfer vertices, as illustrated in Fig.~\\ref{Fig:fdiag}(b). In the slow-turn case, we shall show that the correction is of order \\begin{equation} \\delta P_\\zeta \\sim \\left(\\dot\\theta/H\\right)^2 P_\\zeta~, \\end{equation} where $P_\\zeta$ is the power spectrum, $\\dot\\theta/H$ is the turning angular velocity in Hubble time.  \\begin{figure}[thpb] \\center \\epsfig{file=figs/Fdiagrams.eps, width=0.7\\textwidth} \\caption{\\label{Fig:fdiag} Feynman diagrams for the transfer vertex (a), corrections to the power spectrum from isocurvature modes (b), and the leading bispectrum (c).} \\end{figure}    In Sec.~\\ref{Sec:bispectra}, we calculate the bispectra of quasi-single field inflation. Since the slow-roll condition is not required in the isocurvature direction, there is a sharp contrast between $V'''$ for the massive isocurvaton potential and $V_{\\rm sr}'''$ for the slow-roll potential. The slow-roll condition requires $V_{\\rm sr}{'''}\\sim {\\cal O}(\\epsilon^2) P_\\zeta^{1/2} H$, (where $\\epsilon$ collectively denotes slow-roll parameters), which contributes an ${\\cal O}(\\epsilon^2)$ term to $f_{NL}$ \\cite{Maldacena:2002vr}. However $V'''$ is not subject to such conditions. For example, $V'''$ can naturally become of order $H$. To estimate the size of the bispectrum, we first note that $V'''$ contributes a factor of $f_{NL}^{\\rm iso}\\sim P_{\\zeta}^{-1/2}(V'''/H)$ to $f_{NL}$. We also note that what we observe is not $f_{NL}^{\\rm iso}$, but its projections onto the curvature perturbation (Fig.~\\ref{Fig:turnphys}). The efficiency of this projection is determined by the turning angular velocity $\\dot\\theta/H$. For bispectra, we need three transfer vertices (Fig.~\\ref{Fig:fdiag}(c)), so the bispectrum is of order \\begin{equation}\\label{fnlestimate} f_{NL} \\sim P_\\zeta^{-1/2} \\left(\\dot\\theta/H\\right)^3 \\left(V'''/H\\right) ~. \\end{equation} In this section, we will derive Eq. \\eqref{fnlestimate} rigorously with explicit numerical coefficients. We will also compute the full shape functions using the mixed form of the in-in formalism.   \\begin{figure}[thpb] \\begin{center} \\epsfig{file=figs/turnphys.eps, width=7cm} \\end{center} \\caption{ The physical interpretation of non-Gaussianity in quasi-single field inflation. The non-Gaussianity in the isocurvature ($\\sigma$) direction(s) is not suppressed by slow roll parameters. This large non-Gaussianity is projected onto the inflationary ($\\theta$) direction if the inflaton trajectory turns.} \\label{Fig:turnphys} \\end{figure}   In Sec.~\\ref{Sec:squeezed}, we investigate the shapes of bispectra analytically by taking the squeezed limit, $p_3\\ll p_1=p_2$, where $p_i$'s are the momenta in the three-point function. We find that the three-point function (up to a momentum conservation delta function) scales as \\begin{equation} \\langle\\zeta^3\\rangle \\propto p_3^{-3/2-\\nu}~, \\end{equation} where $\\nu=\\sqrt{9/4-m^2/H^2}$, and $m$ is the effective mass of the isocurvaton. Near $\\nu=0$, the shape changes smoothly to \\begin{equation} \\langle\\zeta^3\\rangle \\propto \\ln(p_3/p_1) p_3^{-3/2}~. \\end{equation} This scaling behavior lies between that of the local shape ($\\langle\\zeta^3\\rangle \\propto p_3^{-3}$) and the equilateral shape ($\\langle\\zeta^3\\rangle \\propto p_3^{-1}$). We call these shapes the {\\em intermediate shape}. We shall explain the underlying mechanisms that determine these different shapes.    \\begin{figure}[thpb] \\center \\epsfig{file=figs/4pt.eps, width=0.7\\textwidth} \\caption{\\label{Fig:fdiag4} Feynman diagrams for the scalar-exchange four-point correlation function (a) and contact-interaction four-point correlation function (b).} \\end{figure}   In Sec.~\\ref{Sec:shapeAns}, we approximate the shape functions by simple analytical expressions, to facilitate future data analyses.  In Sec.~\\ref{Sec:trispectra}, we discuss the trispectra. As illustrated in Fig.~\\ref{Fig:fdiag4}, we show that the trispectra are of order \\begin{equation} t_{NL} \\sim \\max \\left\\{P_\\zeta^{-1} \\left(\\dot\\theta/H\\right)^4 \\left(V'''/H\\right)^2~,~ P_\\zeta^{-1} \\left(\\dot\\theta/H\\right)^4 V''''\\right\\}~. \\end{equation} Note that $t_{NL}\\gg f_{NL}^2$ for $\\dot\\theta/H \\ll 1$. So in future experiments where angular multipoles $l$ can be measured as large as thousands, trispectra may become a better probe for quasi-single field inflation.  We conclude in Sec.~\\ref{Sec:conclusion} and discuss many future directions.   In Appendix \\ref{App:thirdorderaction}, we present all terms in the Lagrangian up to the cubic order in two different gauges, and justify the calculation in the main text. In Appendix \\ref{App:allterms}, we present all terms used in our calculations in terms of the factorized, commutator, and mixed form, respectively. In Appendix \\ref{App:wick}, we explain the Wick rotation and its several applications in this paper. In Appendix \\ref{App:inin}, we give general expressions for the series expansion in the in-in formalism. ",
        "conclusions": "\\label{Sec:conclusion} \\setcounter{equation}{0}  To conclude, in this paper, we have investigated in detail a quasi-single field inflation model. We find fields with mass of order $H$ can have important impacts on density perturbations through, for example, turning trajectories. These effects can be computed perturbatively using transfer vertex and Feynman diagrams in the in-in formalism. A one-parameter family of potentially observable large bispectra arise. These new shapes are controlled sensitively by the mass of the isocurvaton and lie between the equilateral and local shape. We also note that the sizes of the trispectra are even larger than those of the bispectra squared in this model.  There are a lot of issues remaining to be investigated in the quasi-single field inflation models. For example,  $\\bullet$ {\\em Data analysis}. Experimental constraints on non-Gaussianities depend on their detailed shapes and running. A variety of experimental methods have been applied to the local and equilateral bispectra \\cite{CMB,LSS,otherCMB,lspace,running,Komatsu:2009kd}. It will be very interesting to constrain the family of new shapes that we find here, or to fit the parameter $\\nu$, using the observational data.   $\\bullet$ {\\em Running of density perturbations}. In this paper we have only considered the constant turn case. More realistically, we expect parameters to vary along the trajectory. As we noted, such a non-constant turn results in a running of the spectral index and non-Gaussianities. It is worth to investigate this running effect in more details.  $\\bullet$ {\\em Trispectra}. In this paper, we obtained the order of magnitude estimate of the trispectra in quasi-single field inflation. However, the $\\nu$-dependent coefficients in the trispectra, and more importantly the shape of the trispectra still remain to be calculated.   $\\bullet$ {\\em Multiple isocurvature directions}. In this paper, we considered one massive isocurvaton. If there exist multiple massive isocurvature directions, instantaneously along the inflaton trajectory, one can find a two-dimensional hyper-surface where there is only one effective isocurvaton. If this hyper-surface is not changing with time, the model belongs to the two-field model such as the one we considered here. If the hyper-surface is changing with time, the situation becomes more complicated. It is worth to investigate the observational consequences of such a case.  $\\bullet$ {\\em Other quasi-single field inflaton models}. The turning trajectory model that we studied is one simple example of the quasi-single field inflation models. The coupling between the inflaton and isocurvatons can be introduced through the kinetic terms in a more general way which may or may not be described by turning trajectories.\\footnote{We thank Jim Cline for helpful discussions on this point.} Such couplings can even be introduced through other types of couplings that do not involve the kinetic terms. Which correlation functions/non-Gaussianities are enhanced is determined by the structure of these couplings.  $\\bullet$ {\\em String cosmology}. As we discussed in the Introduction, quasi-single field inflation is a natural picture for inflation in string theory and supergravity. Fields with mass of order $H$ are ubiquitous and in fact are a common hazard to inflation model building. We have seen that, as isocurvatons, such fields can have important consequences on density perturbations. It thus becomes very interesting to build explicit quasi-single field inflation models from string theory, in terms of either turning trajectories or more general couplings. Because the shapes of the non-Gaussianities depend very sensitively on the mass, and the sizes depend on the strength of the couplings, we have the opportunities to probe such fields through experiments.  $\\bullet$ {\\em Generalizations}. As the number of fields increases, it becomes a logical possibility that we can have more than one inflaton fields. It is worth to generalize the formalism developed in the current work to these more general multiple field models. It is also worth to apply the perturbative method used here to multifield inflation models.    \\medskip"
    },
    "0911/0911.4333_arXiv.txt": {
        "abstract": "{Very high-energy (VHE; E$>$100~GeV) $\\gamma$-rays have been detected from a wide range of astronomical objects, such as pulsar wind nebulae (PWNe), supernova remnants (SNRs), giant molecular clouds, $\\gamma$-ray binaries, the Galactic Center, active galactic nuclei (AGN), radio galaxies, starburst galaxies, and possibly star-forming regions as well. At lower energies, observations using the Large Area Telescope (LAT) onboard \\emph{Fermi} provide a rich set of data which can be used to study the behavior of cosmic accelerators in the MeV to TeV energy bands. In particular, the improved angular resolution of current telescopes in both bands compared to previous instruments significantly reduces source confusion and facilitates the identification of associated counterparts at lower energies. In this paper, a comprehensive search for VHE $\\gamma$-ray sources which are spatially coincident with Galactic \\emph{Fermi}/LAT bright sources is performed, and the available MeV to TeV spectra of coincident sources are compared. It is found that bright LAT GeV sources are correlated with TeV sources, in contrast to previous studies using EGRET data. Moreover, a single spectral component seems unable to describe the MeV to TeV spectra of many coincident GeV/TeV sources. It has been suggested that $\\gamma$-ray pulsars may be accompanied by VHE $\\gamma$-ray emitting nebulae, a hypothesis that can be tested with VHE observations of these pulsars.} ",
        "introduction": "Our understanding of the very high-energy (VHE; E $>$100 GeV) sky has greatly improved during the last few years, thanks to the high sensitivity of current imaging atmospheric Cherenkov telescopes (IACTs), e.g. H.E.S.S., MAGIC, and VERITAS. They typically cover the energy range of $\\sim$100~GeV up to several tens of TeV, and provide an angular resolution of $\\sim 6\\arcmin$. This allows spectral and morphological studies of the various types of VHE sources: pulsar wind nebulae (PWNe), supernova remnants (SNRs), giant molecular clouds, $\\gamma$-ray binaries, the Galactic Center, active galactic nuclei (AGN), radio galaxies, starburst galaxies, and possibly star-forming regions as well. See~\\citet{VHE_Review_08} for a review of the field in 2008, with more recent updates given by the H.E.S.S. \\citep{hess_survey09}, MAGIC \\citep{MAGIC_results_Fermi_Sym09}, and VERITAS collaborations \\citep{VERITAS_results_Fermi_Sym09,VERITAS_cygnus_survey09}. However, many of the sources are not yet identified at other wavelengths, e.g., nearly a third of the Galactic H.E.S.S. sources have no firm identification; in many cases, there are multiple plausible counterparts while, in other cases, no viable counterparts have yet been identified.  Gamma-ray observations of Galactic sources can help us to solve a number of important astrophysical questions, including (1) the physics of pulsars, PWN and SNR; and (2) the origin of cosmic rays. Our Galaxy contains a large number of cosmic accelerators, where particles are accelerated to highly-relativistic energies (up to at least $10^{14}$~eV). The origin of cosmic rays is still not well known, largely because of the lack of directional information of these particles. These very energetic particles can be traced within our Galaxy by a combination of non-thermal X-ray emission and $\\gamma$-ray emission via leptonic (such as inverse Compton scattering of electrons, Bremsstrahlung and synchrotron radiation) or hadronic (via the decay of charged and neutral pions, due to interactions of energetic hadrons) processes. Therefore, observations of $\\gamma$-rays at energies $\\ga$100~MeV can probe the sources of particle acceleration.  The Large Area Telescope (LAT), onboard the \\emph{Fermi Gamma-ray Space Telescope}, provides the best information of the non-thermal sky in the energy range from 20 MeV to 300 GeV. The point-source sensitivity of LAT is $\\sim$10$^{-8}\\,\\mathrm{ph}\\,\\mathrm{cm}^{-2}\\,\\mathrm{s}^{-1}$ above 100~MeV in one year of survey-mode observations~\\citep{lat_technical}, which is an order of magnitude better than that of its predecessor, the Energetic Gamma Ray Experiment Telescope (EGRET). Its angular resolution is $\\la 0\\fdg6$ above 1~GeV, which is particularly important for identifying $\\gamma$-ray sources with multi-wavelength counterparts and revealing their nature~\\citep{lat_technical}. As an important step towards the first source catalog, the LAT collaboration has published a bright source list (BSL) that includes 205 sources, designated with the prefix 0FGL, using data taken during the first three months of observations~\\citep{bsl_lat}. Among them, 121 sources are identified with AGN and one with the Large Magellanic Cloud. Most of the remaining 83 sources are believed to originate in our own Galaxy. It is natural to investigate which of them also have been detected at energies $\\ga$100~GeV.  The search for VHE counterparts of LAT sources is important for the following reasons: \\begin{enumerate} \\item it aids the identification of the true nature of the LAT sources through their VHE counterparts; \\item for pulsars, it helps us to identify their VHE-emitting nebulae; \\item it may provide us with broad-band $\\gamma$-ray spectra, thereby better constraining the emission mechanisms (e.g. distinguish between hadronic and leptonic scenarios). \\end{enumerate}  \\citet{funk08} compare $\\gamma$-ray sources in the third EGRET (3EG) catalog~\\citep{3rd_egret_cat} and the 22 H.E.S.S. sources known at the time within the region of $l=-$30$\\degr$ to 30$\\degr$, $b=-$3$\\degr$ to 3$\\degr$~\\citep{hess_survey}. They do not find spatial correlation between the two populations. Though some coincidence cases are found, the authors conclude that these few cases can be explained by chance coincidence. However, due to the capabilities of EGRET, this study suffers from the following limitations: (1) The sensitivity of EGRET is lower than that of LAT. The lack of photon statistics leads to poorly constrained spectral indices and the spectra terminate $\\la$10~GeV at the upper end for a typical source; (2) EGRET sources are only localized at degree scales, which is much larger than the angular resolution of IACTs. The second point is the instrumental reason which explains the weak correlation of EGRET and H.E.S.S. sources~\\citep{funk08}. These shortcomings are now largely overcome due to the enhanced performance of LAT over EGRET. In addition to the above caveats, they do not consider the extension of the VHE $\\gamma$-ray sources in their analysis. As such, the full potential of this search is not realized for very extended sources like the SNR RX~J1713.7$-$3946, as pointed out by \\citet{HowCanFermiHelp_SciNeGHe08}. After the launch of LAT, one largely benefits from the increased LAT angular resolution over previous studies. As noted in~\\citet{lat_technical}, EGRET could not distinguish the GeV emission of RX~J1713.7$-$3946 from 3EG~J1714$-$3857, while the capabilities of LAT allow one to study individual sources in this region, which contains three VHE $\\gamma$-ray sources~\\citep[See Fig. 1 in][]{CTB_37A}.  The water Cherenkov detector MILAGRO covers the energies above $\\sim$1~TeV and its angular resolution can reach $<$ 1$\\degr$. Using MILAGRO, a search for $\\gamma$-rays from the Galactic LAT BSL has been performed by~\\citet{milagro_bsl}. They found that 14 sources (of the selected 34) show evidence for multi-TeV $\\gamma$-ray emission at a significance of $\\geq3\\sigma$, although most of the source candidates cannot be established as firm detection on an individual basis~\\citep{milagro_bsl}.  In this paper, a search for VHE counterparts of all the presumed Galactic sources in~\\citet{bsl_lat} and \\citet{lat_psr_cat} is performed, with spatial coincidence as the primary criterium for association. The extensions of the VHE $\\gamma$-ray sources are taken into account and the search is not limited to the H.E.S.S. Galactic plane survey region. The broad-band MeV to TeV spectra of coincident sources are then presented. ",
        "conclusions": "In this work, we search for VHE counterparts of each Galactic GeV source in the 0FGL catalog~\\citet{bsl_lat}, based on spatial coincidence. This study benefits significantly from the increased LAT angular resolution and its better sensitivity over previous instruments.  Compared to the EGRET era, not only are there more coincident sources (improvement in quantity), but improvements in quality start to emerge. With the much better sensitivity of LAT, weaker sources are detected which were unknown in the EGRET era. New relations between the GeV and TeV spectra are revealed. A single spectral component is unable to describe some sources detected at both GeV and TeV energies. Two spectral components may be needed in these cases to accommodate the SEDs, where the VHE flux is higher than a PL extrapolation from GeV energies.  A high fraction of \\emph{Fermi} bright sources are found to be spatially coincident with a VHE $\\gamma$-ray source. This shows that there exists a common GeV/TeV source population, a conclusion that is in stark disagreement with that of~\\citet{funk08}."
    },
    "0911/0911.5335_arXiv.txt": {
        "abstract": "Observations in the sub-THz range of large solar flares have revealed a mysterious spectral component increasing with frequency and hence distinct from the microwave component commonly accepted to be produced by gyrosynchrotron (GS) emission from accelerated electrons. Evidently, having a distinct sub-THz component requires either a distinct emission mechanism (compared to the GS one), or different properties of electrons and location, or both. We find, however, that the list of possible emission mechanisms is incomplete. This Letter proposes a more complete list of emission mechanisms, capable of producing a sub-THz component, both well-known and new in this context and  calculates a representative set of their spectra % produced by a) free-free emission, b) gyrosynchrotron emission, c) synchrotron emission from relativistic positrons/electrons, d) diffusive radiation, and e) Cherenkov emission. We discuss the possible role of the mechanisms in forming the sub-THz emission and emphasize their diagnostics potential for flares. ",
        "introduction": "Solar flares, being manifestations of prompt energy releases, produce electromagnetic radiation throughout the entire spectrum of electromagnetic emission from radio waves to gamma rays \\citep[e.g.][]{DennisSchwartz1989,BrownKontar2005}. Although the radio, UV, X-ray, and gamma-ray emission is generally well observed with high resolution from huge number of events, the sub-THz radiation has only recently been observed from a few large events, at a small number of frequencies \\citep{Kaufmann_etal2001,Trottet_etal_2002,Luethi_etal_2004a, Luethi_etal_2004b,Kaufmann_etal_2004,Cristiani_etal_2008,Trottet_etal_2008, Kaufmann_etal_2009a,Kaufmann_etal_2009fast,Silva_etal_2007}. % The available observational tools are unable to measure polarization and are clearly insufficient to provide detailed spectral and positional information about the sub-THz bursts. Nevertheless, available sub-THz observations have already provided us with puzzling questions which have yet to be answered.  The observations suggest that, on top of quiet Sun % emission, at least two kinds of sub-THz emission can be produced. The first kind looks like a natural extension of the microwave spectrum at higher frequencies and so can reasonably be interpreted as synchrotron radiation from accelerated electrons, which are also responsible for microwave and hard X-ray emission. The second kind % looks like a distinct spectral component rising with frequency in the sub-THz range in contrast to the microwave spectrum, which falls with frequency. The origin of this component is unclear; there is no consensus about the possible emission mechanism producing it, moreover, it is ambiguous whether the emission is of thermal or nonthermal origin.  The main observational characteristics of this component are: relatively large radiation peak flux  of the order of 10$^4$~sfu \\citep{Kaufmann_etal_2004};  radiation spectrum  rising with frequency $F(f){\\propto} f^\\delta$;  spectral index  varying with time within $\\delta\\sim1\\dots6$;  sub-THz component can display a sub-second time variability with the modulation about $5\\%$ \\citep{Kaufmann_etal_2009fast}; the source size is believed to be less than $20''$, however, this conclusion is based on a multi-beam observation of a few antennas with $\\sim4'$ resolution each, rather than on true imaging; therefore, the size estimate must be considered with caution. In addition, observations \\citep{Luethi_etal_2004a,Luethi_etal_2004b} suggest the existence of both compact $\\sim 10''$ and extended $\\sim 60''$ components, and source sizes increasing with frequency.  The emission mechanisms proposed so far to account for the sub-THz component are thermal free-free emission, GS emission from flare accelerated electrons, and synchrotron emission from nuclear decay-generated relativistic positrons \\citep[e.g.][as a review]{Nindos_etal_2008}. None of these mechanisms is readily consistent with the full list of the sub-THz component properties and/or with available context observations at other wavelength. Below we analyze these options % and also consider two other emission mechanisms capable of producing a spectrum rising with frequency---diffusive radiation in Langmuir (DRL) waves \\citep[e.g.][]{Fl_Topt_2007_MNRAS, Fl_Topt_2007_PhRvE} and Vavilov-Cherenkov emission \\citep[e.g.,][]{Bazylev_Zhevago_1987} from  chromospheric layers, which make the list of options more complete. Calculating the spectrum from all of the above mentioned models, we establish the range of main source parameters for each model emphasizing the strengths and weaknesses of each emission mechanism.  \\begin{figure*}[t] \\begin{center} \\epsscale{0.9} \\plottwo{f1a.eps}{f1b.eps} \\caption{Radio spectra produced by thermal free-free emission from a uniform cubic source with a linear size of $20''$ for $n_e=10^{11}\\dots4\\cdot10^{12}$~cm$^{-3}$ and $T_e=0.5\\dots 5$~MK. } \\label{FIG_ff} \\end{center} \\end{figure*} ",
        "conclusions": "We have analyzed a number of emission processes, which are capable of producing radiation at the sub-millimeter wavelengths: free-free emission, gyrosynchrotron and synchrotron processes, diffusive radiation, and Vavilov-Cherenkov process. Having in mind the characteristic ranges of parameters of a solar flare region, we calculated the spectra for each process discussed in the paper. Although current observations at two frequencies only do not provide us with detailed spectral shape, we estimate the range of flux levels and spectral indices for each model and compare them with the observations. %  It is likely that the sub-THz emission originates from more than a single source and  more than one mechanism is involved. Free-free emission is a plausible candidate in many cases, at least for large sources observed by \\citet{Luethi_etal_2004a,Luethi_etal_2004b}. The free-free emission is clearly always present, so other mechanisms build additional contributions on top of the free-free component. Gyrosynchrotron/synchrotron emission is likely to play a role in moderate events and also as the extension of normal microwave bursts, which falls with frequency. The role of DRL is less clear since the level of long-wavelength Langmuir waves is  yet unknown in flares. Finally, the Vavilov-Cherenkov emission from compact sources located at the chromospheric level seems to be a plausible process to account for the rising with frequency sub-mm component of large flares. Indeed, the dielectric permittivity of partially ionized plasma is known to fluctuate around unity due to atomic and molecular transition contributions. When the flare-accelerated electrons are present, they will produce Cherenkov radiation at all frequency windows (including IR, viz, and UV bands) where the dielectric permittivity is (even marginally) above unity, giving rise to a radiation spectrum raising with frequency for (mean) dielectric permittivity increasing with frequency.   The available constraints on sub-THz source sizes, time variability, and spectral index are not yet fully reliable particularly because the ground-based sub-THz observations are extremely difficult to calibrate due to strongly variable atmospheric opacity. Further progress in understanding the physics of sub-THz emission from flares requires observations a) with a more complete spectral coverage at the sub-mm range (preferably well-calibrated space-based observations) and b) polarization measurements. Nevertheless, the sub-THz spectral window can be extremely informative, e.g., for diagnostics of the chromospheric chemical composition if the role of the Vavilov-Cherenkov emission is confirmed. This calls for a new project mounting a sub-THz receivers/interferometers, combining good sensitivity with high spatial resolution, on a space mission, complementing on-going efforts in the microwave range."
    },
    "0911/0911.5273_arXiv.txt": {
        "abstract": "{ We investigate the leptophilic properties of Dirac gauginos in an R--symmetric N=2 supersymmetric model with extended gauge and Higgs sectors. The annihilation of Dirac gauginos to leptons requires no chirality flip in the final states so that it is not suppressed as in the Majorana case. This implies that it can be sizable enough to explain the positron excess observed by the PAMELA experiment with moderate or no boost factors. When squark masses are heavy, the annihilation of Dirac gauginos to hadrons is controlled by their Higgsino fraction and is driven by the $hZ$ and $W^+W^-$ final states. Moreover, at variance with the Majorana case, Dirac gauginos with a non-vanishing Higgsino fraction can also have a vector coupling with the $Z$ gauge boson leading to a sizable spin--independent scattering cross section off nuclei. Saturating the current antiproton limit, we show that Dirac gauginos can leave a signal in direct detection experiments at the level of the sensitivity of dark matter searches at present and in the near future.} ",
        "introduction": "N=1 supersymmetry broken around the weak scale would be the prime new physics candidate to be searched at the LHC. It provides not only an appealing explanation for the origin of the electroweak symmetry breaking but also a natural dark matter (DM) particle of the Universe. While N=1 supersymmetry is enforced by the chiral structure of matter fields in the Standard Model, the gauge sector may be extended to have N=2 supersymmetry in which gauginos can be Dirac particles \\cite{hall91,Randall92,fox02,nelson02,fox04,benakli05,benakli06, benakli08,amigo08} (for a review of N=2 supersymmetry, see \\cite{N=2}). Determining the Majorana/Dirac nature of gauginos will be an interesting task for future experiments looking for supersymmetric CP and flavor violation~\\cite{hisano06,kribs07}, collider signatures \\cite{nojiri07,drees08}, and dark matter properties~\\cite{ullio06,hsieh07,kribs08,benakli09}.  There are some important distinctions between Majorana and Dirac gauginos. First, the annihilation of a Dirac gaugino pair into a fermion--antifermion pair, $\\chi\\bar{\\chi} \\to f\\bar{f}$, has a non-vanishing $s$--wave contribution even in the limit of vanishing fermion masses, and thus the leptonic final states are not suppressed. This implies that a Dirac gaugino can provide a viable DM explanation for the abundance of energetic electrons and positrons recently observed in cosmic rays~\\cite{HEAT,AMS-01,PAMELA,PPB-BETS,ATIC,HESS,FERMI}. Second, Dirac gauginos can have a vector coupling with the $Z$ gauge boson, $\\bar{\\chi}\\gamma_\\mu\\chi\\, Z^\\mu$, leading to sizable spin--independent scattering off nuclei.  As a consequence, the Higgsino component in the DM gets constrained by the direct detection data \\cite{CDMS} and also by the non-observation of antiproton excesses in PAMELA \\cite{pbar} as will be discussed in detail in the following.  In order to ensure such Dirac nature of DM, Majorana mass terms, which provide the mass splitting between the Dirac components, have to be highly suppressed. Otherwise, the heavier component of the two quasi-degenerate Majorana gauginos will decay to the lighter one, and the galactic DM will consist of a pure Majorana gaugino.  If one assumes the N=2 structure in the Higgs sector, i.e. that the two Higgses $H_u$ and $H_d$ form an N=2 hypermultiplet \\cite{benakli06,benakli08}, the Dirac structure imposed at the tree level is spoiled by a large amount due to radiative corrections with Higgs--Higgsino and fermion--sfermion in the loop.  A way to enforce an (almost) pure Dirac property for gauginos is then to assume a continuous R symmetry which forbids the usual $\\mu$ term for the Higgs bilinear coupling $H_u H_d$ and the soft breaking trilinear $A$ term~\\cite{kribs07}.  In this paper, we will work out a framework for R--symmetric Dirac gauginos, which are introduced by extending the gauge/Higgs sector of a supersymmetric model. We assume the Dirac bino as the main component of the lightest supersymmetric particle. The annihilation of two Dirac binos into a lepton--antilepton pair usually dominates over the annihilation into a quark--antiquark pair, as sleptons are typically much lighter than squarks. As we will see, the annihilation to a CP even final or intermediate state turns out to be velocity--suppressed, which is a property that is also true in the usual Majorana gaugino case.  Then, the Dirac bino annihilation to the two $W$ bosons ($W^+W^-$) and that to a CP even Higgs boson and a $Z$ boson ($hZ$), both controlled by the Higgsino component in the DM particle, turn out to be the next important sources of the cosmic antiproton flux.  Quantifying the suppression factor for the Higgsino component, one can make a prediction for the direct detection rate.  In Section 2, an R--symmetric N=2 extension of the Minimal Supersymmetric Standard Model is introduced to enforce an almost pure Dirac gaugino structure. Based on this framework, we derive interaction vertices for the Dirac bino, which is assumed to be the dark matter, and then provide a qualitative analysis for its indirect and direct detection properties in Section 3. A numerical analysis is performed in Section 4 to obtain various fits and constraints of the leptophilic Dirac bino dark matter from the current indirect and direct detection data. We give our conclusions in Section 5. ",
        "conclusions": "Determining the Majorana/Dirac nature of gauginos will be an interesting task for future experiments to look for supersymmetry. Moreover, as an interesting variance of the standard Majorana case, Dirac gauginos can be natural realizations of leptophilic DM, which has been recently proposed to explain the rising positron flux observed in cosmic rays by PAMELA without producing excesses in the antiproton signal. In this paper, we have analyzed the phenomenology of Dirac gauginos in a specific supersymmetric realization, where a continuous R symmetry is assumed to protect vanishingly small Majorana masses. We have shown that a light Dirac gaugino ($m_{\\chi}\\lsim$ 500 GeV) can explain the PAMELA data with moderate or no boost factor for light ($\\gsim m_{\\chi}$) slepton masses. Furthermore, in the case of a non-vanishing Higgsino fraction, Dirac gauginos can have a vector coupling with the $Z$ gauge boson leading to a sizable spin--independent scattering off nuclei. In some specific examples, we have shown that present constraints on the Higgsino fraction of Dirac gauginos from direct detection experiments are at the same level of those coming from antiproton data. This implies that a Dirac gaugino signal is potentially at the level of the sensitivity of direct detection experiments at present and in the near future.  \\medskip  {\\bf Acknowlegement:} E.J.C. was supported by Korea Neutrino Research Center through National Research Foundation of Korea Grant (2009-0083526). S.S. was supported by the WCU program (R32-2008-000-10155-0) of National Research Foundation of Korea."
    },
    "0911/0911.4920_arXiv.txt": {
        "abstract": "\\vspace{1cm} Geiger-mode avalanche photodiodes (G-APD) are promising new sensors for light detection in atmospheric Cherenkov telescopes. In this paper, the design and commissioning of a 36-pixel G-APD prototype camera is presented. The data acquisition is based on the Domino Ring Sampling (DRS2) chip. A sub-nanosecond time resolution has been achieved. Cosmic-ray induced air showers have been recorded using an imaging mirror setup, in a self-triggered mode. This is the first time that such measurements have been carried out with a complete G-APD camera. ",
        "introduction": "Imaging atmospheric Cherenkov telescopes (IACT) have been very successful in detecting very high energy ($\\sim$ 0.1--30\\,TeV) $\\gamma$-rays from cosmic sources \\cite{aharonian08}. The key component is a pixelated camera which has to resolve flashes of Cherenkov light from air showers (duration 1--5\\,ns for $\\gamma$-induced air showers, main wavelength range 300--650\\,nm). High-sensitivity photo-sensors are needed, even using light-collecting optics, since e.\\,g. a 1\\,TeV primary photon hitting the atmosphere only results in about one hundred Cherenkov photons per square meter. Until now, matrices of photomultiplier tubes (PMT) have always been employed for this task. This is a well-known technology which, however, comes with some intrinsic disadvantages for telescope applications. PMTs are rather heavy and bulky, but at the same time fragile. They require high voltages of several 100\\,V or even kV and are damaged when exposed to sunlight. Typical PMTs furthermore have photon detection efficiencies (PDE) of only 20--30$\\%$.  Since a few years, a new type of semiconductor light-sensor is being developed, the so-called Geiger-mode avalanche photodiode (G-APD\\footnote{Also referred to as silicon photomultiplier (SiPM) or multi-pixel photon counter (MPPC).}) \\cite{renker09}. These light-weight devices are built up from multiple APD cells operated in Geiger-mode. All the cells are connected in parallel and the overall signal is the sum of all simultaneously fired cells. They need bias voltages of 50--100\\,V, usually 1--5\\,V above the breakdown voltage. A high gain similar to PMTs is reached ($10^5$--$10^6$) and, potentially, higher PDEs of up to 50$\\%$. The market for G-APDs is continuously growing, and several manufacturers are working on their improvement. For an IACT application under realistic ambient conditions, several technical challenges have to be met. This mostly concerns the necessity to compensate for gain variations due to changes in temperature or night-sky background light (NSB). While the former is an intrinsic feature of G-APDs, the latter is a general problem and applies to other photo-sensors as well. An advantage of G-APDs is that they can be operated at NSB rates up to several GHz per sensor, thus allowing measurements under twilight and moonlight conditions.  First tests to detect Cherenkov light with small G-APD arrays have been performed in the past \\cite{biland08}. In order to develop a complete G-APD camera and find solutions for the technical challenges, the FACT project (First G-APD Camera Test) \\cite{braun09} has been launched. The prototype presented in this paper marks the first step towards a large camera with a field of view of about $5^{\\circ}$ (0.1--0.2$^{\\circ}$ per pixel). This device is foreseen for the DWARF telescope (Dedicated multi-Wavelength AGN Research Facility) \\cite{bretz08} which will perform monitoring of strong and varying $\\gamma$-ray sources. ",
        "conclusions": "A 36-pixel prototype G-APD camera for Cherenkov astronomy has successfully been constructed and commissioned. A DAQ system based on the DRS2 chip has been set up. In a self-triggered mode, images of air showers induced by cosmic rays have been recorded. For the first time, this has been achieved with a complete G-APD camera. Stable operation at room temperature and under high NSB light conditions is possible. In summary, these photo-sensors have proven to fulfill the requirements of IACT applications, with the potential to replace or complement PMTs for future projects like the planned Cherenkov Telescope Array (CTA) \\cite{cta}."
    },
    "0911/0911.3507_arXiv.txt": {
        "abstract": "{ The elements technetium and lithium are two important indicators of internal nucleosynthesis and mixing in late type stars. Studying their occurrence and abundance can give deep insight into the structure and evolution in the late phases of the stellar life cycle. } { The aims of this paper are: 1) to revisit the Tc content of a sample of oxygen-rich asymptotic giant branch (AGB) variables and 2) to increase the number of such stars for which the Li abundance has been measured to provide constraints on the theoretical models of extra-mixing processes. } { To this end, we analysed high-resolution spectra of 18 sample stars for the presence of absorption lines of Tc and Li. The abundance of the latter was determined by comparing the observed spectra to hydrostatic MARCS model spectra. Bolometric magnitudes were established from near-IR photometry and pulsation periods. } { We reclassify the star V441~Cyg as Tc-rich, and the unusual Mira star R~Hya, as well as W~Eri, as Tc-poor. The abundance of Li, or an upper limit to it, was determined for all of the sample stars. In all stars with Tc we also detected Li. Most of them have a Li content slightly below the solar photospheric value, except for V441~Cyg, which has $\\sim$1000 times the solar abundance. We also found that, similar to Tc, a lower luminosity limit seems to exist for the presence of Li. } { We conclude that the higher Li abundance found in the cooler and higher luminosity objects could stem from a production mechanism of Li operating on the thermally pulsing AGB. The stellar mass might have a crucial influence on this (extra mixing) production mechanism. It was speculated that the declining pulsation period of R~Hya is caused by a recent thermal pulse (TP). While not detecting Tc does not rule out a TP, it indicates that the TPs are not strong enough to drive dredge-up in R~Hya. V441~Cyg, on the other hand, could either be a low-mass, intrinsic S-star that produced its large amount of Li by extra-mixing processes, or an intermediate-mass star ($M\\gtrsim4\\,\\rm{M}_{\\sun}$) undergoing Li production due to hot bottom burning. } ",
        "introduction": "The asymptotic giant branch (AGB) phase is the last stage of evolution with nuclear burning for stars of low and intermediate mass ($1 - 8\\rm{M}_{\\sun}$). In the most luminous part of the AGB, the behaviour of an AGB star is characterised by the so-called thermal pulses (TP), thermal instabilities of the He shell accompanied by changes in luminosity, temperature, pulsation period, and internal structure \\citep[see e.g.][for reviews]{Busso99,Herwig}. Between the repeated events of explosive He-burning, heavy elements can be produced via the ``slow neutron capture process'' \\citep[s-process, see e.g.][]{Wallerstein97} in the region between the hydrogen- and the helium-burning shells. The processed material may then be brought to the stellar surface by the convective envelope that temporarily extends to these very deep layers. This mixing event is called the third dredge-up (3DUP) and is the cause of the eventual metamorphosis of an oxygen-rich M-star into a carbon-rich C-star.  The element technetium (Tc, $Z=43$) is a well-established diagnostic tool in the study of the thermally pulsing AGB (TP-AGB) phase. Because of the short half life time of the longest lived Tc isotope produced in the s-process ($^{99}$Tc, $\\tau_{1/2}=2.1\\times10^{5}$\\,yr), any Tc that is detected in the star's atmosphere must have been produced in-situ. This requirement has been used to study the evolution of AGB stars ever since the pioneering discovery of Tc absorption lines in the spectra of late type giants by \\citet{Mer52b}. Several classes of objects have been investigated for the presence of Tc in the past \\citep[e.g.][]{DW86,Lit87,VEJ99}. A key question in this research is if and when in the evolution on the AGB, in terms of luminosity, the 3DUP sets in. This question was addressed by \\citet{LH99,LH03} for a sample of solar-neighbourhood semi-regular pulsating AGB stars and for a larger sample of Galactic disk AGB stars, including several Mira variables. Recently, \\citet{Utt07a} has investigated a sample of Galactic bulge AGB variables, for which the distance and hence luminosity is more precisely known than for their Galactic disk counterparts. The general conclusion of these works is that the lower luminosity limit for 3DUP to set in, as predicted by theoretical evolutionary models of the AGB, is consistent with the lowest luminosity Tc-rich stars. However, being above this limit is not sufficient for a star to have Tc in its atmosphere \\citep{LH03,Utt07a}. In the present paper we revisit the Tc content of a sample of oxygen-rich AGB stars, with important consequences for a few individual stars.  The element lithium (Li, $Z=3$) is very fragile and temperature-sensitive. Once it is mixed into stellar layers with temperatures exceeding $3\\times10^6$\\,K, it is destroyed via proton captures to form helium. Its fragility makes Li an important diagnostic tool for stellar evolution \\citep{Reb91}, and also has a high significance in the determination of cosmological parameters \\citep{Spi82,Kor06}. The standard model \\citep[e.g.][]{Michaud91} predicts a strong reduction in the Li surface abundance after the first dredge up to values $\\log \\epsilon(\\rm{Li}) < 1.5$, or even less if depletion effects on the main sequence are considered. Most giants for which a Li abundance has been determined \\citep[see a summary by][]{Michaud91} agree with this result.  Despite its fragility, Li can be produced rather than destroyed under certain condition in stars via the so-called Cameron-Fowler or $^7$Be-transport mechanism \\citep[$^3$He($\\alpha$,$\\gamma$)$^7$Be($e^-$, $\\nu$)$^7$Li; ][]{CF71}. The production of Li is always connected to mixing mechanisms at or below the bottom of the convective envelope, and requires $^{3}$He to be present in the stellar atmosphere. One stage in stellar evolution where Li production occurs is the luminosity bump of the red giant branch (RGB) when the H-burning shell erases the molecular weight discontinuity left behind by the first dredge up. A small fraction of giants in this phase are known to be very Li-rich \\citep{delaReza,CharBal}. The mixing process below the convective envelope has been named extra-mixing or cool bottom processing \\citep[CBP;][]{Was95}. Its effects are also reflected in the carbon isotopic ratio measured from the spectra of giant stars \\citep{gil91}, as well as in other isotopic ratios measured in meteoritic pre-solar grains of RGB or AGB origin \\citep{Nol03}.  The physical cause of the extra-mixing below the convective envelope is hitherto unknown. \\citet{Egg06}, on the basis of a 3D simulation, find that Rayleigh-Taylor instabilities below the convective envelope can develop from the inversion of the mean molecular weight gradient induced by $^{3}$He burning. Alternatively, \\citet{ChaZah} suggest that the double-diffusive mechanism called ``thermohaline instability'' should be at play. In principle, both these phenomena might also occur on the AGB, though detailed models are just emerging \\citep{Can08}. Finally, \\citet{Busso07} explore the possibility that circulation of partially processed matter can be accounted for by the magnetic buoyancy induced by a stellar dynamo operating on the RGB and on the AGB. That CBP is probably active on the AGB has been recently demonstrated by \\citet{Utt07b} and \\citet{Leb08}.  Among the AGB stars, those with carbon-rich chemistry (C-type, C/O$>$1 due to the dredge-up of freshly synthesised $^{12}$C) have been studied more intensively for their Li content than those with oxygen-rich chemistry (M-type, C/O$<$1) in the past. This may partly have a historical reason because of the early detection of carbon stars with a strong Li resonance line \\citep{McK40,Tor64}, but also an observational reason. Although the molecular line blanketing mainly caused by the CN and C$_2$ molecules in the spectral vicinity of the Li line in C-stars is strong, that of the TiO molecule in oxygen-rich stars of the same effective temperature is even more severe. The lower limit of Li abundance that can be measured is thus higher for O-rich stars than for C-rich stars. In both cases, a Li abundance determination can only be achieved with the help of spectral synthesis techniques including large numbers of molecular lines. Thus, the counterparts to the catalogues of Li abundance measurements in C-stars \\citep{Den91,Bof93} include only relatively small tables of M-type giants \\citep{Luc82}. Even in this last work, a number of the observed stars are supergiants and non-variable giants that are rather in the RGB than in the AGB phase of evolution.  There are, however, M-type AGB stars that are known to be very rich in Li. This type of star has been observed in the Magellanic Clouds \\citep{Smith95}, and very recently also in the Milky Way galaxy \\citep{GH07}. These objects are intermediate-mass ($M \\gtrsim 4 - 8\\,\\rm{M}_{\\sun}$) TP-AGB stars with a convective envelope penetrating the H-burning shell, thus nuclear-burning occurs partly under convective conditions. This process is generally known as hot bottom burning \\citep[HBB;][]{Iben1973}. Under these conditions, carbon, on the one hand, is converted into nitrogen, keeping the C/O ratio below 1, and on the other hand $^7$Li is produced very efficiently by the Cameron-Fowler mechanism. The Li abundance on the surface of HBB stars can be higher than the initial abundance of the interstellar medium from which they formed. We do not intend to deal with this kind of stars in the present paper, but rather with the less massive AGB stars that have not (yet) dredged up enough carbon to have C/O$>$1 and are not massive enough to remain O-rich due to the operation of HBB. Nevertheless, by inspecting their bolometric magnitudes, we will check whether our Li-rich objects belong to this group of stars.  The correlation of Tc and Li was studied for a sample of Galactic S-type (C/O$\\simeq$1) stars by \\citet{Van07}. \\citet{Utt07b} present Li abundances and their correlation with Tc in a homogeneous sample of O-rich AGB stars located in the Galactic bulge. Here, we want to increase the sample of O-rich AGB stars whose Li abundance has been measured, and study its correlation with the presence of Tc in this type of stars. These measurements will be useful for understanding the interaction between CBP mechanisms and 3DUP in AGB stars.  \\smallskip The paper is structured in the following way. In Sect.~\\ref{samobs}, the sample and the observations are described and bolometric magnitudes of the sample stars derived; in Sect.~\\ref{results} we present the analysis of the spectra, results concerning the presence of Tc and the abundance of Li; in Sect.~\\ref{discussion} the results are discussed in view of a possible enrichment mechanism of Li, and we take a closer look at two of the sample stars; Sect.~\\ref{summary} summarises the work. ",
        "conclusions": "\\label{summary}  We analysed high-resolution optical spectra of a sample of 17 M-type (oxygen-rich) and one C-type (carbon-rich) AGB variables for the presence of Tc and the abundance of Li. We employed a flux-ratio method to decide on the presence of the 3$^{\\rm{rd}}$ dredge-up indicator Tc in the stellar atmosphere, and measured the abundance of the temperature-sensitive element Li with spectral synthesis techniques of the 670.8\\,nm Li resonance line using MARCS model atmospheres. Bolometric magnitudes were established from near-IR photometry and pulsation periods. The lower luminosity limit for 3DUP to occur, established in previous works, has been confirmed. We re-classified three sample stars with respect to their Tc content, namely that \\object{R Hya} and \\object{W Eri} are Tc-poor, and \\object{V441 Cyg} is Tc-rich. We thus suggest that R~Hya has not yet experienced a dredge-up of freshly produced s-elements, but, regarding the possibility that it recently underwent a TP, might have a 3DUP episode in the future. V441~Cyg, on the other hand, is an intrinsic S-star.  Lithium is detected in ten sample stars (with a $\\log \\epsilon(\\rm{Li}) > 0$). After combining these abundances with bolometric magnitude estimates, we find that Li is more abundant in more luminous and cool stars. All Tc-rich sample stars are found to contain also some Li in their atmosphere. We thus suggest that this fragile element could be replenished in the stellar atmosphere by some (extra-mixing) process that is only active in the more luminous objects. This could also be connected to a mass effect within the sample. The S-star \\object{V441 Cyg} is found to be super-Li rich, with a higher Li abundance than the solar photospheric value probably by a factor 1000. The gathered evidence does not allow us to clearly conclude whether this object is an intermediate-mass AGB star ($M \\gtrsim 4\\,\\rm{M}_{\\sun}$) producing Li by the HBB mechanism, or a low-mass star with extremely efficient extra-mixing processes. The data obtained in this study may be useful for constraining theoretical models of such extra-mixing processes."
    },
    "0911/0911.1335.txt": {
        "abstract": "The following  parameters are determined for 63 Galactic supergiants in the solar neighbourhood: effective temperature $T_{\\rm eff}$, surface gravity $\\log g$,  iron abundance $\\log\\epsilon$(Fe), microturbulent parameter $V_t$, mass $M/M_\\odot$, age $t$ and distance $d$. A significant improvement in the accuracy of the determination of $\\log g$ and, all parameters dependent on it, is obtained through application of van Leeuwen's (2007) rereduction of the {\\it Hipparcos} parallaxes. The typical error in the $\\log g$ values  is now $\\pm$0.06 dex for supergiants with distances $d < 300$ pc and $\\pm$0.12 dex for supergiants with $d$ between 300 and 700 pc;  the mean error in $T_{\\rm eff}$ for these stars is $\\pm$120 K. For supergiants with $d > 700$ pc parallaxes are uncertain or unmeasurable, so typical errors in their $\\log g$ values are 0.2--0.3 dex.  A new $T_{\\rm eff}$ scale for A5--G5 stars of luminosity classes Ib--II is presented. Spectral subtypes and luminosity classes of several stars are corrected. Combining  the $T_{\\rm eff}$ and $\\log g$ with evolutionary tracks, stellar masses and ages are determined; a majority of the sample has masses between 4$M_\\odot$ and 15$M_\\odot$ and, hence, their progenitors were early to middle B-type main sequence stars.  Using Fe\\,{\\sc ii} lines, which are insensitive to departures from LTE, the microturbulent parameter $V_t$ and the iron abundance $\\log\\epsilon$(Fe) are determined from high-resolution spectra. The parameter $V_t$ is correlated with gravity: $V_t$ increases with decreasing $\\log g$. The mean iron abundance for the 48 supergiants with distances $d$ $<$ 700 pc is $\\log\\epsilon$(Fe)=7.48$\\pm$0.09, a value close to the solar value of 7.45$\\pm$0.05, and thus the local supergiants and the Sun have the same metallicity. ",
        "introduction": "In this new series of papers, we shall obtain and comment upon chemical compositions of a large sample of A-,F-, and G-type supergiants. An initial step in an abundance analysis is the determination of the principal stellar fundamental parameters: the effective temperature $T_{\\rm eff}$, the surface gravity $\\log g$, the iron abundance $\\log\\epsilon$(Fe), and the microturbulent parameter $V_t$. In now traditional fashion, a model atmosphere is computed for a given set of fundamental parameters and applied to the analysis of the observed stellar spectrum. Iteration is essentially necessary because the derived composition may differ from that assumed in the construction of the model atmosphere. Additionally, analysis of the spectrum may provide new information on the fundamental parameters.  In this paper, the sample of supergiants is introduced and their fundamental parameters derived. For many stars, this is not the first presentation of estimates of their fundamental parameters. A new discussion appeared vital because inspection of the literature   shows generally divergent results; large uncertainties about $T_{\\rm eff}$ and $\\log g$ (particularly) translate to inaccurate estimates of the elemental abundances. As one illustrative example, we show in Table 1 published $T_{\\rm eff}$ and $\\log g$ for $\\alpha$ Per (HR 1017) where the ranges are  340 K and 1.4 dex.  Lyubimkov et al. (2009) note that the scatter for the F0 Ib star $\\alpha$ Lep (HR 1865) is 500 K in $T_{\\rm eff}$ and 1.2 dex in $\\log g$. An even more extreme spread of 2400 K in $T_{\\rm eff}$ is reported by Schiller \\& Przybilla (2008) for the A2 Ia supergiant $\\alpha$ Cyg (Deneb, HR 7924): the spread in $\\log g$ from their collated references was 0.5 dex. A salutary lesson to be drawn from these examples of bright well-studied supergiants is that fundamental parameters of supergiants in the literature deserve a careful redetermination if one is seeking accurate abundances.  \\begin{table} \\centering \\begin{minipage}{70mm} \\caption{The $T_{\\rm eff}$ and $\\log g$ values derived by various authors for the bright supergiant $\\alpha$ Per (F5 Ib)} \\begin{tabular}{lll} \\hline $T_{\\rm eff}$, K & $\\log g$ & Authors\\\\ \\hline 6250 & 1.8 & Parsons (1967) \\\\ 6250 & 0.90 & Luck \\& Lambert (1985) \\\\ 6500 & 1.5 & Klochkova \\& Panchuk (1988)\\\\ 6250 & 1.20 & Gonzalez \\& Lambert (1996)\\\\ 6270 & -- & Evans et al. (1996)\\\\ 6200 & 0.60 & Andreievsky et al. (2002)\\\\ 6541 & 2.0 & Kovtyukh et al. (2008)\\\\ 6350 & 1.90 & present work \\\\ \\hline \\end{tabular} \\end{minipage} \\end{table}  Here, we consider not only familiar spectroscopic and photometric methods for determining the fundamental parameters but also a determination of the gravity $\\log g$ that is based on the stellar parallax. The {\\it Hipparcos} parallaxes from the original catalogue (ESA 1997) have been significantly improved by van Leeuwen (2007). In particular, the parallaxes for many bright supergiants  have their errors reduced by 4--5 times, a reduction that allows the stellar parallax to be a competitive method for estimating $\\log g$.  Our initial sample of 67 stars was reduced to 63 after $T_{\\rm eff}$ and $\\log g$ estimates had been obtained and four stars shown to be misclassified  in {\\it The Bright Star Catalogue} (Hoffleit \\& Warren 1991). For these 63 supergiants, we determined the parameters -- $T_{\\rm eff}$, $\\log g$, $V_t$, and $\\log\\epsilon$(Fe) -- as well as the stellar mass $M/M_\\odot$, age $t$, and distance $d$.  Determination of these fundamental parameters is a prerequisite for the abundance analyses pursued in this series of papers to examine questions primarily of stellar evolution (say, mixing arising from rotation and deep convection) and secondarily of Galactic chemical evolution. Several abundance anomalies in atmospheric composition of supergiants have been found and widely attributed to mixing between the interior and the atmosphere. For example, the carbon abundance shows a tendency to be underabundant (e.g., Luck \\& Lambert 1985; Venn 1995a,b; Venn \\& Przybilla 2003). Nitrogen, as might be expected for C-depleted stars, is overabundant. Sodium appears to be overabundant (Boyarchuk \\& Lyubimkov 1983; Lyubimkov 1994; Andrievsky et al. 2002). The lithium abundance varies greatly from star-to-star (Luck 1977). These anomalies with respect to a star's initial composition are presumed traceable to the mixing into the atmosphere of nuclear-processed material from the stellar interior. There are two possible episodes of mixing: (i) rotationally-induced mixing in the rapidly-rotating main sequence (MS) B-type progenitor of the supergiant; and (ii) the deep convective mixing, the so-called first dredge-up, that occurs towards the end of the first crossing the Hertzsrpung gap when the star is a K- or M-type supergiant and before a return to earlier spectral types. Since rotationally-induced mixing in the MS progenitors should be confined to the most rapidly-rotating stars and effects of the first dredge-up seen only in stars that have completed a first-crossing of the Hertzsrpung gap, one expects the strength of the abundance anomalies (Li, C, N, Na etc.)  to vary from absent to marked  across a large sample of A, F, and G supergiants. If this expectation is not confirmed, i.e., an anomaly is present in all stars, the conclusion must be that episodes (i) and (ii) are incompletely understood or another process exists for affecting the surface compositions. ",
        "conclusions": "In this paper, we have set the stage for a new abundance analysis of AFG supergiants of luminosity classes Ib and II from the solar neighbourhood. The major contribution of this paper is our determination  of accurate fundamental parameters characterizing the stellar atmospheres: $T_{\\rm eff}, \\log g$, $V_t$, and $\\log\\epsilon$(Fe). Of  particular note is our use of  stellar parallaxes from the rereduction of the {\\it Hipparcos} parallaxes (van Leeuwen 2007) in determining the surface gravity $\\log g$ with what is claimed to be an unprecedented precision: the error is typically $\\pm$0.06 dex for supergiants with distances $d < 300$pc increasing to $\\pm0.12$ dex for the supergiants with $d$ between 300 pc and 700 pc. In the case of  supergiants with distances greater that 700 pc, the parallaxes are too small or unmeasureable for the stellar parallax to provide a useful constraint on the gravity and, then, the typical error in $\\log g$ is 0.2--0.3 dex, the commonly declared accuracy for published $\\log g$ values.  The primary $T_{\\rm eff}$ indicators are the Balmer lines and the indices $[c_1]$, $Q$ and $\\beta$ which provide loci in the plane $(T_{\\rm eff},\\log g$) plane with the degeneracy broken by the  $T_{\\rm eff}$-insensitive  $\\log g$ estimate from the stellar parallax. The mean error in $T_{\\rm eff}$ for stars with $d < 700 $ pc is $\\pm120$ K. Our accurate $T_{\\rm eff}$ for  63 supergiants, some of which are reclassified on the basis of our effective temperatures and gravities,  are the basis for a new $T_{\\rm eff}$ scale for A5--G5 stars of luminosity classes Ib--II. (Three stars  selected from the {\\it Bright Star Catalogue} as A or F supergiants are shown to be dwarfs or subgiants. A fourth star listed as a G supergiants is shown to be a  K supergiant. This quartet  will not be considered in subsequent papers.)  The turbulent parameter $V_t$ and the iron abundance $\\log\\epsilon$(Fe) were determined from Fe\\,{\\sc ii} lines on account of their insensitivity to non-LTE effects. The parameter $V_t$ is seen to be gravity dependent: $V_t$  increases with decreasing $\\log g$.  The iron abundance is independent of stellar distance $d$ and equal to the solar Fe abundance: the mean abundance for the 48 supergiants with distance $d < 700$ pc is $\\log\\epsilon$(Fe) = 7.48$\\pm$0.09 dex, a value coincident with the solar abundance $\\log\\epsilon$(Fe) = 7.45$\\pm$0.05 dex (Grevesse et al. 2007).  The coincidence between the iron abundance of young nearby stars and that of the 4.5~Gyr old Sun is interesting from the viewpoint of models of Galactic Chemical Evolution (GCE). It is important to note that there are other data as well, which confirm a closeness of the metallicity of these stars to the solar one. We may mention our earlier conclusion based on the Mg abundance in local B stars, the progenitors of the AFG supergiants. We reported the abundance $\\log\\epsilon$(Mg) = 7.59$\\pm$0.15 (Lyubimkov et al. 2005), a value equal within the uncertainties to the solar value of $\\log\\epsilon$(Mg) = 7.53$\\pm$0.09.  Moreover, studies of young stars, both hot and cool,  by a variety of authors have shown that their compositions are very similar to that of the Sun (see, e.g., the iron abundance obtained by Luck, Kovtyukh \\& Andrievsky (2006) for Cepheids in the solar neighbourhood, the iron and magnesium abundances found by Fuhrmann (2004) for nearby F, G and K dwarfs and subgiants of the thin Galactic disk, and the sulfur abundance derived dy Daflon et al. (2009) for B-type main-sequence stars in the Orion association).  The question arises: may these results be reconciled with models of GCE? One may cite the recent work of Spitoni et al. (2009), where an enrichment of the solar neighborhood by various metals is studied, in particular, by Fe and Mg (see their Fig. 16 and 17). One sees from these results that during the Sun's life the Fe and Mg abundances in its neighborhood are predicted to increase by about 0.15 dex and 0.05-0.07 dex, respectively. The ordinary accuracy of observed abundances in stars seems to be insufficient to detect such a small enrichment. In particular, according to our above-mentioned finding, the enrichment $[Fe/H]$ = 0.03$\\pm$0.09 and $[Mg/H]$ = 0.06$\\pm$0.15 takes place, that does not contradict Spitoni et al.'s theoretical estimations but is negligible in comparison with uncertainties in the derived Fe and Mg abundances.  The fundamental parameters and Claret's (2004) evolutionary tracks provide estimates of stellar mass and age.  The great majority, 57 of the 63 stars, have masses between 4$M_\\odot$ and 15$M_\\odot$  showing that their progenitors were early to middle B-type main sequence stars. Five of the remaining six stars have lower mass, $M =2.8-3.9M_\\odot$, and their progenitors were late B-type main sequence stars.  The final star in the minority of six is HR 825 with an estimated mass of 24$M_\\odot$ and its progenitor was an O-type main sequence star.  In subsequent papers, we shall derive detailed chemical compositions  for the supergiants.  A primary aim of this future work, as noted in the Introduction, is to quantify accurately the various signatures of internal mixing that are anticipated to occur in preceding stages of evolution  from the main sequence to the dredge-ups occurring in supergiants.  In particular, we shall compare the compositions of the supergiants with the results for early and middle B-type stars from our continuing collaboration (Lyubimkov et al. 2000, 2002, 2004, 2005)."
    },
    "0911/0911.3394_arXiv.txt": {
        "abstract": "We present a unified picture for the evolution of star clusters on the two-body relaxation timescale. We use direct N-body simulations of star clusters in a galactic tidal field starting from different multi-mass King models, up to 10\\% of primordial binaries and up to $N_{tot}=65536$ particles. An additional run also includes a central Intermediate Mass Black Hole. We find that for the broad range of initial conditions we have studied the stellar mass function of these systems presents a universal evolution which depends only on the fractional mass loss. The structure of the system, as measured by the core to half mass radius ratio, also evolves toward a universal state, which is set by the efficiency of heating on the visible population of stars induced by dynamical interactions in the core of the system. Interactions with dark remnants (white dwarfs, neutron stars and stellar mass black holes) are dominant over the heating induced by a moderate population of primordial binaries (3-5\\%), especially under the assumption that most of the neutron stars and black holes are retained in the system. All our models without primordial binaries undergo a deep gravothermal collapse in the radial mass profile. However their projected light distribution can be well fitted by medium concentration King models (with parameter $W_0 \\sim 8$), even though there tends to be an excess over the best fit for the innermost points of the surface brightness. This excess is consistent with a shallow cusp in the surface brightness ($\\mu \\sim R^{-\\nu}$ with $\\nu \\sim 0.4-0.7$), like it has been observed for many globular clusters from high-resolution HST imaging. Generally fitting a King profile to derive the structural parameters yields to larger fluctuations in the core size than defining the core as the radius where the surface brightness is one half of its central value. Classification of core-collapsed globular clusters based on their surface brightness profile may thus fail in systems that appear to have already bounced back to lower concentrations, particularly if the angular resolution of the observations is limited and the core is not well resolved. ",
        "introduction": "\\label{sec:intro} Until a few years ago, the standard picture of globular clusters stellar population and dynamical evolution emerging from theoretical and observational studies was yielding a consistent and well understood framework. Globular clusters were thought to be 'simple stellar population', composed of stars with the same age and chemical composition. Globular cluster dynamical evolution had been thoroughly investigated by a large number of numerical studies showing that clusters surviving the early expansion triggered by primordial gas expulsion and mass loss due to stellar evolution evolve, due mainly to two-body relaxation, toward core collapse and higher central concentrations while losing stars through the boundary set by the tidal field of their host galaxy (see e.g. \\citealt{hh}).  In the last few years, however, a wealth of observational data mostly from exquisite space based observations have revealed the actual complexity of globular cluster dynamics and stellar populations. It is now clear that there is a close link and interplay between dynamical evolution and the stellar content of clusters, the structure of clusters and the abundances of exotic objects (such as blue stragglers, X-ray sources, pulsars etc.) (see e.g. \\citealt{bel06,shara06,hut06}).  Recently, a number of photometric and spectroscopic observations have also shown evidence of the presence of multiple stellar populations and challenged the commonly held view according to which globular cluster are 'simple stellar populations'. It appears that essentially any globular cluster that has high quality photometric and spectroscopic data available presents evidence against a single star-formation burst \\citep{gratton2004,bedin2004,piotto05,piotto07,carretta09}. The presence of multiple stellar populations has major implications for the formation and the early evolution of globular clusters (see e.g. \\citealt{ercole08}).  Furthermore, our understanding of the late-time evolution of globular clusters has been called into question. A recent observational study \\citet{dem07},  focused on the stellar mass function of Galactic globular clusters, found that clusters with lower concentration index of the surface brightness profile tend to have flatter stellar mass functions. The observed trend is at odds with what is expected from the standard dynamical evolution scenario according to which as clusters evolve toward core collapse and higher concentrations they lose mass and the preferential loss of low-mass stars flattens the stellar mass function; dynamically older clusters should be more centrally concentrated and have flatter stellar mass functions. \\citet{bau08} have suggested that strong primordial mass segregation might be required to explain the flat stellar mass function of some of the low-concentration cluster in the observed sample of \\citet{dem07}. However a strong primordial mass segregation might lead to the rapid dissolution of the cluster due to stellar evolution mass-losses \\citep{vesperini09}.  High-resolution photometry of the center of globular clusters has also revealed that the standard \\citet{king66} models do not capture the details of the surface brightness profile, which often shows the presence of shallow cusps within a relatively large core \\citep{noyola06}. The shallow cusps contrast with the classic expectation that a strong, isothermal cusp develops as a result of core-collapse. The presence of a central IMBH is a possible explanation for the presence of those cusps \\citep{bau04,tre07b}, but this interpretation appears in contrast with the lack of strong evidence for IMBHs in globular clusters \\citep[e.g. see][and references therein]{gill08}. NGC2298, for example, is one of the clusters showing a central excess in the surface brightness profile from the \\citet{noyola06} data, but a central IMBH is ruled out at high confidence level \\citep{pas09}.  The goal of this paper is to study in detail the relation between a cluster structural properties and the dynamical phases of its evolutionary path as well as the relation between the properties of a cluster stellar mass function and its structure. We focus our attention on the long-term evolution of star clusters driven by the effects of two-body relaxation. We analyze our numerical models following the procedures used in the observational studies of the structure and stellar content of star clusters and by directly comparing our results with observational data we aim at clarifying their interpretation in the context of dynamic evolution of relaxed stellar systems.  We have carried out a survey of multimass simulations of star clusters evolving in a tidal field and starting from a broad range of different initial conditions spanning different initial density profiles, primordial binary fractions and initial stellar mass function. This paper is organized as follows. In Section~\\ref{sec:sim} we present our sample of simulations. In Section~\\ref{sec:defs} we introduce our framework for analyzing simulations consistently with the observational derivation of structural quantities. In Section~\\ref{sec:results} we presents the results of our analysis, concluding in Section~\\ref{sec:conc}. ",
        "conclusions": "\\label{sec:conc}  In this paper we discussed the evolution of the global mass function index $\\alpha$ and of the structural parameters for multimass direct N-body models of star clusters in a galactic tidal field with and without primordial binaries using a total number of particles up to $N=65536$. We adopted a \\citet{ms} initial mass function evolved with an instantaneous step of stellar evolution to a turn-off mass of $0.8$ M$_{\\sun}$. The simulation results have been interpreted with great care in order to use definitions of structural quantities consistent with observations of Galactic star clusters.  The initial density profile of the cluster or the tidal density field is not important in the long term, as the memory of the starting concentration is erased on a relaxation timescale, with a similar subsequent evolution for models starting from different initial conditions (see figs.~\\ref{fig:alpha_m} to \\ref{fig:c_m_64}). $r_c/r_{hP}$ appears to be stable also over different IMFs, but its value depends on the retention fraction of neutron stars and stellar mass black holes. In particular a deep core collapse can be identified from the core radius as determined following the observational procedures, $r_c$, only if few dark remnants are present in the system. Even a significant fraction (f=10\\%) of primordial binaries does not appear to influence much the observational determination of the core radius. This is likely related to the heating induced by mass segregation of dark remnants in the core of the system.  These findings highlight that the core radius $r_c$ as obtained from King model fits over the complete radial extent of the surface brightness profile are unable to reliably identify a core-collapse cluster. Post-core collapsed systems are in fact relatively well fitted by medium concentration King models, although high residuals are present at the center of the system (see Fig.~\\ref{fig:sb_fit}). It is suggestive to note that these residuals appear similar to shallow cusps that have been observed in several galactic globular clusters with HST photometry \\citep{noyola06}. In our simulations a shallow cusp ($\\mu \\sim R^{-\\nu}$ with $\\nu \\sim 0.4-0.7$) in the surface brightness profile appears after core collapse (see Fig.~\\ref{fig:sb_fit}). The presence of such cusps may thus provide a hint that a system is past core-collapse, but other indicators are needed. One possibility if high angular resolution data are available, is to use $r_{\\mu}$ as a measure of the core, that is the radius at which the surface brightness falls to half of its central value. $r_{\\mu}$ does in fact track the evolution of the tri-dimensional core radius defined in simulations, i.e. $\\sd{r_c}$ (see Fig.~\\ref{fig:rc_defs}). In alternative, one could explore the use of the mass-to-light ratio if spectroscopic information are available. Shallow cusps at the center of medium concentration surface brightness profiles appears to be so common in the post core-collapsed evolution of a globular cluster that this indicator is not particularly helpful in suggesting clusters likely to harbor a central IMBH, contrary to past suggestions \\citep{baum05,miocchi07}.  Our investigation on the evolution of structural parameters ``observed'' from simulations shows an improved agreement with the $c-\\alpha$ correlation identified by \\citet{dem07} for Galactic globular clusters, but still the simulated clusters are too concentrated in their final stages of their evolution. This might be related to the definition of the tidal radius $r_t$, and thus of the concentration, from the $\\gtrsim 15$ years old, ground-based data used by \\citet{trager95}. The light profile is in fact dominated by contributions from stars along the giant branch, more massive than average, and therefore more mass-segregated than the main-sequence stars we used in this analysis. In any case, from the observed $c-\\alpha$ correlation we can expect, based on our analysis, that some of the systems with the most depleted stellar mass function can indeed be core-collapsed systems lurking in the sample of objects fitted by regular King profiles. This highlights the importance of carrying out a uniform survey to construct the structural parameters of Milky Way globular clusters, combining large area coverage with the high angular-resolution data at the center, like those obtained for about 50 globulars by \\citet{sara07} with HST."
    },
    "0911/0911.0262_arXiv.txt": {
        "abstract": "A selected overview of stellar effects and reaction mechanisms with relevance to the prediction of astrophysical reaction rates far off stability is provided. ",
        "introduction": "Nuclear theory is important in the determination of reaction cross sections and rates for astrophysics in several respects. Firstly, a large number of reactions in astrophysical environments involve highly unstable nuclei which are unaccessible in laboratory measurements. This is especially true for the high temperature plasmas of stellar explosions. Secondly, despite of high plasma temperatures, the relevant particle interaction energy is low by nuclear physics standards. This is a challenge for experimentalists, especially when studying charged-particle reactions at or below the Coulomb barrier. Very small cross sections at low energy may even prevent a measurement and have to be predicted in this case. Thirdly, the nuclei occur in excited states because they are in thermal equilibrium with the stellar plasma. This modifies the reaction cross sections and has important consequences for the determination of stellar reaction rates. Laboratory measurements only account for a fraction of the possible transitions and the actual plasma corrections (stellar enhancement factors, SEF) can only be provided by theory. (Electron screening is also important but will not be discussed here.)  Large-scale predictions of reaction cross sections across the chart of nuclei are required to provide the rates needed for reaction networks applied to nucleosynthesis. In astrophysical applications usually different aspects are emphasized than in pure nuclear physics investigations. Reliable predictions of the nuclear properties and optical model ingredients, like optical potentials for particle and $\\alpha$ transmission coefficients \\cite{kissprl,rausupp,yalcin,Moalpha,silver,raubranch}, nuclear level densities (NLD) \\cite{rau97,mocelj}, resonance energies, and $\\gamma$-transition strengths \\cite{rau08,litvi}, are needed far from stability. Even at stability, there are still considerable uncertainties, especially in the optical potential for charged particles which is not well constrained at low energy. In addition to the well-known difficulty in predicting $\\alpha$-potentials \\cite{yalcin,mohr,somor,koehl,avri}, it was recently found that also the optical potential for protons may require special attention \\cite{kissprl,rausupp,kissold,raujpconf}. ",
        "conclusions": ""
    },
    "0911/0911.3327_arXiv.txt": {
        "abstract": "We present self-consistent high-resolution simulations of NGC4038/4039 (the ''Antennae galaxies``) including star formation, supernova feedback and magnetic fields performed with the \\textsl{N}-body/SPH code $\\textsc{Gadget}$, in which magnetohydrodynamics are followed with the SPH method. We vary the initial magnetic field in the progenitor disks from $10^{-9}$ to $10^{-4}$ G. At the time of the best match with the central region of the Antennae system the magnetic field has been amplified  by compression and shear flows to an equilibrium field value of $\\approx 10$ $\\mu$G, independent of the initial seed field. These simulations are a proof of the principle that galaxy mergers are efficient drivers for the cosmic evolution of magnetic fields. We present a detailed analysis of the magnetic field structure in the central overlap region. Simulated radio and polarization maps are in good morphological and quantitative agreement with the observations. In particular, the two cores with the highest synchrotron intensity and ridges of regular magnetic fields between the cores and at the root of the southern tidal arm develop naturally in our simulations. This indicates that the simulations are capable of realistically following the evolution of the magnetic fields in a highly non-linear environment. We also discuss the relevance of the amplification effect for present day magnetic fields in the context of hierarchical structure formation. ",
        "introduction": "Within the framework of the hierarchical galaxy formation picture, galaxies assemble and evolve via mergers of smaller progenitor galaxies (e.g. \\citealp{White&Rees1978}, \\citealp{White&Frenk1991}). Thus, galaxy interactions are essential for the understanding of structure formation. In the bottom-up-picture of structure formation dwarf galaxies merge to form the first galaxies at an early epoch of the universe. Later, there is still a continuous merging of fully evolved galaxies. The further growth of galaxies progresses through a combination of mergers and diffuse accretion of gas.  Interactions of galaxies change their dynamics drastically (see e.g. \\citealp{Toomre&Toomre1972}, \\citealp{Barnes1992}, \\citealp{Hernquist1994}, \\citealp{Barnes1999} and \\citealp{Burkert&Naab2003}, \\citealp{Gonzalez-Gracia2006}) as the gravitational potential is changing rapidly during the interaction.  Since the gas component is dissipative and most sensitive to changes of the gravitational potential, it is strongly affected during the interaction and driven to the galaxy centers, eventually causing bursts of star formation (\\citealp{Barnes&Hernquist1992}, \\citealp{Mihos&Hernquist1994}, \\citealp{Barnes&Hernquist1996}, \\citealp{Bekki&Shioya1998}, \\citealp{Springel2000}, \\citealp{Barnes2002}, \\citealp{Bournaud2005}, \\citealp{Cox2006}, \\citealp{Naab&Jesseit2006}, \\citealp{Robertson2006}, \\citealp{Cox2008}, \\citealp{Hopkins2008}). So far simulations of interactions and mergers of disk galaxies have only been investigated with respect to changes in stellar dynamics, gas flows, star formation (SF) or formation of central supermassive black holes (\\citealp{DiMatteo2005Natur}, \\citealp{Springel2005}, \\citealp{Springel2005BlackHoles}, \\citealp{Robertson2006}, \\citealp{DeBuhr2009}, \\citealp{Johansson2009}, \\citealp{Johansson2009_2}). However, the dramatic impact of mergers on the gas flows will directly affect the magnetic fields of the systems (and vice versa) via the induction equation of magnetohydrodynamics (MHD) and the Lorentz force. The magnetic fields will change their morphology following the motion of the gas and will be amplified by shocks and gas inflow.  Changes in the magnetic field structure, on the other hand, might influence gas flows, local collapse and the morphology as well as the star formation activity. Local, interacting galaxies are the perfect laboratories for investigating the effects associated with their magnetic fields. However, the timescales for galaxy mergers are far too long to observe these processes directly. The only observational possibility to study the time evolution of mergers is to consider different systems at different evolutionary stages. However, the available sample of interacting nearby galaxies is too small to investigate the evolution of magnetic fields in detail. Thus, numerical simulations pose a promising tool to study the magnetic field evolution in interacting systems.  The structure of an interacting system strongly depends on the properties of the progenitor galaxies. Thus, matching observed nearby interacting systems with simulations in space and time can give us an idea of the properties of their progenitors, e.g. their sizes, gas fractions and relative velocities. Furthermore, comparing simulated systems with observations helps to asses the performance of the applied numerical method. Numerical methods supported by these comparisons can then be used to study processes in the early universe, when galaxies were very different from present-day galaxies. High resolution simulations of the formation of individual galaxies in a full cosmological context (see e.g. \\citealp{Naab2007}) including magnetic fields could help us in understanding the processes leading to the magnetization of the universe. This type of study would complement earlier semi-analytical studies that investigated the possibility of magnetic field seeding by galactic winds in a cosmological context (\\citealp{Bertone2006}).  The standard theory of magnetic field amplification in galaxies is the so called Galactic Dynamo based on the mean field theory (see \\citealp{Kulsrud1999} for a review). Within this theory, turbulent motions of the ionized gas driven by stellar activity lead to the generation of a random magnetic field ($\\alpha$-effect). This random magnetic field (particularly its radial component) can then be amplified by the differential rotation of the galaxy ($\\Omega$-effect), leading to an efficient dynamo action which results in an exponential growth of the magnetic field (the $\\alpha\\Omega$-dynamo). However, dynamos may probably only work efficiently if magnetic helicity is transported away from the differentially rotating disc (\\citealp{Brandenburg2005}). \\citet{Gressel2008} performed high-resolution box-simulations which demonstrate that a dynamo may operate if supernova explosions release magnetic helicity from the disc. However, for an efficient magnetic helicity transport out from a galactic disk, galactic winds or galactic fountains may be required. This might be a problem particularly for massive galaxies due to the deeper potential well. The fact that it is difficult to get an efficient dynamo is generally addressed as the dynamo problem. Different solutions, e.g. turbulence driven by large-scale SN-bubbles (\\citealp{Ferriere1992}) or the Cosmic Ray Dynamo (\\citealp{Hanasz2009}) have been proposed. These solutions describe the exponential growth of a small-scale magnetic seed field which is amplified up to present-day values within several Gyr. However, recent observations indicate that magnetic fields in galaxies have been already very strong (comparable to present-day galactic magnetic fields) at very high redshifts, at a time when the universe was only $t \\approx$ 6 Gyr old ($z\\approx 1$) (\\citealp{Bernet2008}). Former observations of damped Ly-$\\alpha$ systems by \\citet{Wolfe1992} indicate that progenitors of galactic discs had magnetic fields of a few $\\mu$G even at $z\\approx2$ ($t\\approx 3$ Gyr). The very fast amplification required to generate the strong magnetic fields at high redshifts can probably not be achieved with any Galactic Dynamo model (see e.g. \\citealp{Arshakian2009}). Thus, alternative possibilities for the amplification of galactic magnetic fields on shorter timescales need to be explored. \\citet{Lesch&Chiba1995} have shown analytically that strong magnetic fields in high redshift objects can be explained by the combined action of an evolving protogalactic fluctuation and electrodynamic processes providing magnetic seed fields. \\citet{Wang&Abel2009} performed numerical simulations of the formation of disc galaxies within an collapsing halo imposing a uniform initial magnetic field of $10^{-9}$ G. The initial field was amplified by three orders of magnitude within approximately 500 Myr of evolution. The amplification might be due to the combined effects of magnetic field compression during the collapse and amplification of the uniform initial field by differential rotation as studied also in \\citet{me2009}. These studies indicate, that the amplification of magnetic fields might be a natural part of the galaxy formation process. However, interactions of galaxies, which were more frequent at earlier times, pose another promising possibility.  Although it would be worthwile to consider cosmological studies of structure formation including magnetic fields in the long run, numerical studies of interacting magnetized systems in the local universe may serve as a first step towards a more complete scenario. These studies help us in understanding the morphological evolution of galaxies, their star formation histories (\\citealp{Springel2005Natur}, \\citealp{Cox2008}, \\citealp{Bournaud2007}, \\citealp{DiMatteo2008}, \\citealp{Jesseit2009}, \\citealp{Naab&Ostriker2009}), and as we will show in this paper also their magnetic histories. The system NGC 4038/39, also known as the Antennae galaxies, is one of the best examples for an ongoing local merger.  It is also the by far best observed interacting galaxy system.  In this paper we present further steps towards a more complete numerical representation of the Antennae system as a prototype for interacting galaxies. For the first time we will follow the evolution of the magnetic field in a galaxy interaction simulation. We also address the general question whether smoothed particle magnetohydrodynamics (SPMHD) is capable of following the evolution of magnetic fields in interacting systems at reasonable accuracy. In a previous paper we have shown that SPMHD is well suited for following the evolution of magnetic fields in isolated disk galaxies (\\citealp{me2009}) so the study of interacting systems is a natural extension of this earlier study.  The paper is organized as follows: Section \\ref{ANT_PROP} summarizes the properties of the Antennae system as known from observations and theory. In section \\ref{SIMULATIONS} we describe our numerical methods (section \\ref{NUMERICS}), the setup of the isolated disks (section \\ref{SETUP_ISOLATED}) and the match to the observed Antennae system (section \\ref{SETUP_MATCH}). A detailed analysis of the evolution of the system is presented in section \\ref{EVOUTION_OF_ANTENNAE}, where we discuss the evolution of the magnetic field (section \\ref{EVOUTION_OF_ANTENNAE_MAG}) the numerical stability of our simulations (section \\ref{NUMERICAL_STABILITY}) and the self-regulation of the magnetic field amplification (section \\ref{PRESSURES}). In section \\ref{RADIO_MAPS} we describe our method to calculate artificial radio maps (section \\ref{RADIO_METHOD}) and present applications to the isolated disk and the Antennae simulations (section \\ref{RADIO_APPLICATIONS}). The artificial radio maps derived from the simulations can be compared against the radio observations of the system, thus providing a further tool for constraining the numerical model and method. We conclude in section \\ref{SUMMARY} and briefly discuss the relevance of our simulations in the context of hierarchical structure formation. ",
        "conclusions": "\\label{SUMMARY} We have presented the first fully self-consistent \\textsl{N}-body/SPH simulations of the interacting Antennae galaxy system including magnetic fields. We show that weak magnetic seed fields in the isolated disk galaxies are amplified by the gravitational interaction throughout the two galactic encounters. Thereby the magnetic pressure saturates at a level corresponding to equipartition between the turbulent and the magnetic pressure, independent of the initial field strength. Particularly, magnetic fields with an initial value higher than the equipartition value diminish during the evolution, demonstrating that the state of equipartition is the natural state for magnetized galactic systems. An analysis of artificial total synchrotron emission and polarization maps provides a convincing agreement with the observations. Summarizing, the method of \\textsl{N}-body/SPH simulations including magnetic fields reproduces quite conclusively the complicated dynamics of the amplification and spatial design of magnetic fields in interacting galaxies.  Moreover, a detailed discussion of the numerical divergence of the magnetic field in SPH simulations has been presented in section \\ref{NUMERICAL_STABILITY}. Our analysis strongly suggests that numerical divergence measures which are smaller than a certain threshold can be considered as measures of sub-resolution fluctuations which do not affect the overall evolution of the magnetic field. Considering simulations using the Euler-Potentials, which pose a $\\nabla\\cdot\\textbf{B}$-free prescription by definition, this threshold can be assessed to be $h\\nabla\\cdot\\textbf{B}/|\\textbf{B}| \\approx 1$ (see also \\citealp{me2009}).  What can we learn from these simulations for the global evolution of cosmic magnetic fields? Within the framework of standard CDM hierarchical clustering models the formation of large disk galaxies as well as elliptical galaxies is characterized by more or less intense merging of smaller galactic subunits, e.g. dwarf galaxies, collapsing gas clouds or globular clusters. If we assume that at least some of the accreted subunits have been magnetized by stellar activity (e.g. supernova explosions, stellar winds or T-Tauri-jets), the merging of such subunits to larger galaxies must have been accompanied by a significant amplification and restructuring of the magnetic field on galactic scales. The amplification and ordering of small-scale magnetic fields to a toroidal configuration during the evolution of isolated galaxies was recently shown by \\citet{Hanasz2009} and \\citet{Dubois&Teyssier2009} independently. \\citet{Hanasz2009} considered an axially symmetric galactic disk in which stellar seed fields were amplified by a cosmic ray driven dynamo. \\citet{Dubois&Teyssier2009} demonstrated the amplification and ordering of small-scale fields seeded by SF activity in the context of the formation of a dwarf galaxy with significant galactic winds. Complementary to these findings, our simulations prove that amplification via non-axisymmetric three dimensional gravitational interaction alone may provide an alternative channel for galactic as well as intergalactic magnetic field evolution. In other words, given that the structure formation is characterized by a galactic bottom-up architecture, we would expect that within one or two Giga-years the Universe has been globally magnetized by the combination of dynamo action in isolated galaxies and dynamical amplification by interacting galactic objects. However, dynamo action is supposed not to be very efficient in dwarf galaxies since their differential rotation is not strong enough (\\citealp{Gressel2008}). Thus, at an early epoch of the universe, when most of the galaxy population consists of dwarfs, magnetic field amplification due to interactions may be even more significant.  With their study of the formation of dwarf galaxies including magnetic fields and galactic winds, \\citet{Dubois&Teyssier2009} demonstrate an alternative scenario based on the ideas of \\citet{Bertone2006} which they call the \"Cosmic Dynamo\". According to their findings, galactic winds from young dwarf galaxies eject magnetic field energy into the intergalactic medium, leading to a mean intergalactic field $B_\\mathrm{IGM}$ of $10^{-11}$ to $10^{-10}$ G.  The preceding amplification of the magnetic field inside the dwarf galaxy by the combined action of stellar activity and differential rotation (i.e. the Galactic Dynamo) is thereby restricted by the IGM magnetic field already present at the formation time of the galaxy. For an IGM magnetic field of $B_\\mathrm{IGM}\\approx 10^{-10}$ G, the Lorentz force may prevent the formation of a new generation of dwarf galaxies and subsequent star formation. As a consequence, the IGM magnetic field never grows significantly above $10^{-10}$ G. Since dwarf galaxies are characteristic in the early phase of the evolution of the universe, this Cosmic Dynamo may have been very efficient in magnetizing the IGM. However, besides the accretion of IGM material previously enriched with magnetic fields, \\citet{Dubois&Teyssier2009} also point out the importance of accretion of satellite galaxies for the evolution and amplification of the magnetic field in galaxies at later times.  Our simulations emphasize the Cosmic Dynamo scenario proposed by \\citet{Dubois&Teyssier2009}. The efficient amplification of the magnetic field during the equal-mass-merger presented here clearly shows that interactions of galaxies should be taken into account in studies of the magnetic evolution of the universe. We would also expect an intergalactic medium which is not only enriched with heavy elements by stellar activity, but also magnetized on large scales by galaxy interactions. Our simulations may help to understand the observationally well established facts that very young galaxies already exhibit magnetic field strengths comparable with nearby fully developed spiral galaxies and the rotation measure estimates of intergalactic magnetic fields (e.g. \\citealp{Bernet2008})."
    },
    "0911/0911.1608_arXiv.txt": {
        "abstract": "New $UBV$-photometry obtained in 2000-2008 is presented for the post-AGB star IRAS 22272+ 5435=V354 Lac. The star showed semi-regular light variations with varying amplitudes. The maximal amplitude did not exceed: $\\Delta V$=0.$^{m}$5, $\\Delta B$=0.$^{m}$7 and $\\Delta U$=1.$^{m}$0. For 2000-2008, we have found a photometric period near 128 days. The analysis of long-term observations in 1990-2008 reveals variations with two close periods: 128 and 131 days, causing amplitude modulation. The $V$-$(B-V)$ diagram shows a clear correlation: the star is generally bluer when brighter. From our $UBV$ data, we derive $E(B-V)$=0.5 and conclude that the spectral type of the star varies between K1 to K7 during pulsations. The mean $UBV$-data of V354 Lac have not changed during the past 19 years: $V=8.^{m}60$, $B-V=2.^{m}06$ and $U-B=2.^{m}14$. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0911/0911.3111_arXiv.txt": {
        "abstract": "{Cosmological structure formation offers insights on all of the universe's components: baryonic matter, dark matter, and, notably, dark energy. The mass function of galaxy clusters at high redshifts is a particularly useful probe to learn about the history of structure formation and constrain cosmological parameters.} {We aim at deriving reliable masses for a high-redshift, high-luminosity sample of clusters of galaxies selected from the \\survey ~survey of X-ray selected clusters. Weak gravitational lensing allows a comparison of the mass estimates derived from X-rays, forming an independent test. Here, we will focus on a particular object, \\cl00 ~at $z\\!=\\!0.50$.} {Using deep imaging in three passbands with the \\megacam ~instrument at MMT, we show that \\megacam ~is well-suited for measuring gravitational shear, i.e.\\ the shapes of faint galaxies. A catalogue of background galaxies is constructed by analysing photometric properties of galaxies in the $g'r'i'$ bands.} {We detect the weak lensing signal of \\cl00 ~at $5.8\\sigma$ significance, using the aperture mass technique. Furthermore, we find significant tangential alignment of galaxies out to $\\sim\\!10\\arcmin$ or $>\\!2~r_{200}$ distance from the cluster centre. The weak lensing centre of \\cl00 ~agrees with several X-ray measurements and the position of the brightest cluster galaxy. Finally, we infer a weak lensing virial mass of $M_{200}\\!=\\!7.2^{+3.6+2.3}_{-2.9-2.5}\\times 10^{14}\\mathrm{M}_{\\sun}$ for \\cl00 .} {Despite complications by a tentative foreground galaxy group in the line of sight, the X-ray and weak lensing estimates for \\cl00 ~are in remarkable agreement. } ",
        "introduction": "\\label{sec:intro} The mass function $n(M,z)$ of galaxy clusters is a sensitive probe to both cosmic expansion and evolution of structure by gravitational collapse \\citep[cf.\\ e.g.,][]{2002ARA&A..40..539R,2005RvMP...77..207V,2005RvMA...18...76S}. Therefore, mass functions derived from statistically well-understood cluster samples can be and are frequently used to determine cosmological parameters like $\\Omega_{\\mathrm{m}}$, the total matter density of the universe in terms of the critical density, and $\\sigma_{8}$, the dispersion of the matter density contrast. In addition, measurements of the mass function at different redshifts have the potential to constrain the possible evolution of the dark energy component of the universe \\citep{2006ewg3.rept.....P,2006astro.ph..9591A}, expressed by the value and change with time of the equation-of-state parameter  Because the abundance and mass function of clusters are sensitive functions of these cosmological parameters, they have been studied intensively both theoretically \\citep{1974ApJ...187..425P,1999MNRAS.308..119S,2001MNRAS.321..372J,2008ApJ...688..709T} and observationally.  Several methods to measure cluster masses have been employed to determine their mass function. Assuming hydrostatic equilibrium of the intracluster medium (ICM), the mass of a cluster can be computed, once its X-ray gas density and temperature profiles are known. If the quality of the X-ray data does not allow the determination of profiles for individual clusters, for instance at high redshift, X-ray \\emph{scaling relations} of the X-ray luminosity \\citep[$L_{\\mathrm{X}}$--$M_{\\mathrm{tot}}$; e.g.,][]{2002ApJ...567..716R,2008MNRAS.387.1179M}, temperature ($T_{\\mathrm{X}}$--$M_{\\mathrm{tot}}$), gas mass ($M_{\\mathrm{gas}}$--$M_{\\mathrm{tot}}$) or $Y_{\\mathrm{X}}\\!=\\!T_{\\mathrm{X}}M_{\\mathrm{gas}}$ with the total mass are used as proxies \\citep{2009ApJ...692.1033V}. Simultaneously constraining cosmological parameters and X-ray cluster scaling relations, \\citet{2009arXiv0909.3098M} found $w_{\\mathrm{DE}}\\!=\\!-1.01\\!\\pm\\!0.20$ for the dark energy equation-of-state parameter, compiling data from a large sample of galaxy clusters.  Weak gravitational lensing provides a completely independent probe of a cluster's mass, as it is sensitive to baryonic and dark matter alike, not relying on assumptions of the thermodynamic state of the gas. The fact that sources of systematic errors in the lensing and X-ray methods are unrelated opens the possibility to compare and cross-calibrate X-ray and weak lensing masses.  Several studies have been recently undertaken, comparing X-ray and weak lensing cluster observables: % \\citet{2006ApJ...653..954D} found the weak lensing mass to scale with X-ray luminosity like $\\propto~L_{\\mathrm{X}}^{1.04\\pm0.46}$ and constrain a combination of $\\Omega_{\\mathrm{m}}$ and $\\sigma_{8}$. \\citet{2007MNRAS.379..317H} established a proportionality between the weak lensing mass $M_{\\mathrm{wl}}$ within the radius $r_{2500}$ inside which the density exceeds the critical density by a factor $2500$ and $T_{\\mathrm{X}}^{\\alpha}$, with an exponent $\\alpha\\!=\\!1.34^{+0.30}_{-0.28}$. For the same radius, \\citet{2008MNRAS.384.1567M} quoted a ratio $M_{\\mathrm{X}}/M_{\\mathrm{wl}}\\!=\\!1.03\\pm0.07$ along with a decrease towards smaller radii. Directly aiming at the ratio of weak lensing to X-ray mass, \\citet{2008A&A...482..451Z} measured at a radius of $M_{\\mathrm{wl}}/M_{\\mathrm{X}}\\!=\\!1.09\\!\\pm\\!0.08$, for a radius $r_{500}$, and later confirmed this value and the trend with radius \\citep{2010arXiv1001.0780Z}. \\citet{2009MNRAS.396..315C} investigated the role of the mass estimator on the systematics of the weak lensing mass function.   In order to % make progress, it is particularly important to determine the masses of more clusters to a high accuracy, especially at high ($z\\!>\\!0.3$) redshifts. To date, only a few studies have been undertaken in this regime which provides the strongest leverage on structure formation and is thus crucial for tackling the problem of dark energy.  In the redshift range $0.3\\!\\lesssim\\!z\\!\\lesssim\\!0.8$, for two reasons, weak gravitational lensing becomes more advantageous over methods relying on the X-ray emission from the ICM with increasing redshift. First, X-ray temperature profiles are becoming increasingly difficult to determine with cluster redshift. Second, the fraction of clusters undergoing a merger increases with $z$, in accordance with the paradigm of hierarchical structure formation \\citep[e.g.][]{2005APh....24..316C}, rendering the assumption of hydrostatic equilibrium more problematic at higher $z$.  With this paper, we report the first results of the largest weak lensing follow-up of an X-ray selected cluster sample at high redshifts. ",
        "conclusions": "With this study, we report the first results for the largest weak lensing survey of X-ray selected, high-redshift clusters, the \\emph{400d cosmological sample} defined by \\citet{2009ApJ...692.1033V} and determine a weak lensing mass for an interesting cluster of galaxies, \\cl00 , which had not been studied with deep optical observations before. We observed \\cl00 , along with other clusters of our sample, using the \\megacam ~$\\sim\\!24\\arcmin\\times24\\arcmin$ imager at the MMT, obtaining deep $g'r'i'$ exposures. Employing an adaptation of the \\citet{2005AN....326..432E} pipeline, \\texttt{THELI}, and the ``TS'' KSB shape measurement pipeline presented by T.\\ Schrabback in \\citet{2006MNRAS.368.1323H}, we, for the first time, measure weak gravitational shear with \\megacam , showing its PSF properties to be well suited for such venture.  The lensing catalogue of background galaxies is selected by a photometric method, using $g'r'i'$ colour information. Despite similar number count statistics, we find different photometric properties in our \\megacam ~field than in the CFHTLS Deep fields used to estimate the redshift distribution of the lensed galaxies. The photometric measurements establish the galaxy we name G1, for which \\citet{1997MNRAS.285..511B} determined a redshift $z\\!=\\!0.516$ as the BCG of \\cl00 , ruling out a slightly brighter source found inconsistent in its colours with the cluster redshift $z\\!=\\!0.50$. We find additional evidence for the presence of a foreground structure at $z\\!\\approx\\!0.25$ from photometry but find it does neither significantly affect the lensing nor the X-ray mass estimate of \\cl00 .  Having applied several consistency checks to the lensing catalogue and optimising the aperture mass map of the cluster, we detect \\cl00 ~at $5.84\\sigma$ significance. The weak lensing centre obtained by bootstrapping this map is in good agreement with the BCG position and the X-ray detections by \\textsc{Rosat}, \\textsc{Chandra}, and \\textsc{XMM-Newton}. Two tentative strong lensing arcs are detected in \\cl00 .  Tangential alignment of galactic ellipticities is found to extend out to $10\\arcmin$ separation and well modelled by an NFW profile out to $>\\!2r_{200}$. The low concentration parameter found by least-squares fitting to the shear profile is confirmed by the likelihood method with which we determine \\cl00 ~to be parametrised by $r_{200}^{\\mathrm{min}}\\!=\\!1.52^{+0.22}_{-0.24}\\,\\mbox{Mpc}$ and $c_{\\mathrm{NFW}}^{\\mathrm{min}}\\!=\\!2.0^{+1.8}_{-1.2}$. Modifying the likelihood analysis for the fiducial case, we estimate systematic errors due to shear calibration, the redshift distribution of the background galaxies, and the likely non-sphericity of the cluster. Further, we confirm the best model to change little if the BCG is chosen as cluster centre rather than the weak lensing centre. We arrive at a virial weak lensing mass for \\cl00 ~with statistical and systematic uncertainties of $M_{200}\\!=\\!7.2^{+3.6+2.3}_{-2.9-2.5}\\times 10^{14}\\mathrm{M}_{\\sun}$, in excellent agreement with the virial mass obtained using \\textsc{Chandra} and the hydrostatic equation, $M_{\\mathrm{hyd}}\\!=\\!6.26\\pm1.26\\times 10^{14}\\mathrm{M}_{\\sun}$.  Noting that statistical errors on the lensing mass are still high, we conclude that high-quality data and well-calibrated analysis techniques are essential to exploit the full available cosmological information from the mass function of galaxy clusters with weak lensing. Nevertheless, once lensing masses for all the $36$ clusters in the sample are available, these statistical errors are going to be averaged out and reduced by a factor of 6 by combining all clusters when testing for cosmological parameters. Thus, understanding and controlling systematic errors remain important issues. We are going to proceed our survey with the analysis of further high-redshift clusters from the \\emph{400d cosmological sample} observed with \\megacam ."
    },
    "0911/0911.1787_arXiv.txt": {
        "abstract": "This is the third of a series of papers in which we derive simultaneous constraints on cosmological parameters and X-ray scaling relations using observations of the growth of massive, X-ray flux-selected galaxy clusters. Our data set consists of 238 clusters drawn from the {\\it ROSAT} All-Sky Survey, and incorporates extensive follow-up observations using the {\\it Chandra} X-ray Observatory. Here we present improved constraints on departures from General Relativity (GR) on cosmological scales, using the growth index, $\\gamma$, to parameterize the linear growth rate of cosmic structure. Using the method of \\cite{Mantz:09a}, we simultaneously and self-consistently model the growth of X-ray luminous clusters and their observable-mass scaling relations, accounting for survey biases, parameter degeneracies and the impact of systematic uncertainties. Such analysis of the survey and follow-up data is crucial, else spurious constraints may be obtained. We combine the X-ray cluster growth data with cluster gas mass fraction, type Ia supernova, baryon acoustic oscillation and cosmic microwave background data. We find that the combination of these data leads to a tight correlation between $\\gamma$ and the normalization of the matter power spectrum, $\\sigma_{8}$. Consistency with GR requires a measured growth index of $\\gamma\\sim 0.55$. Under the assumption of self-similar evolution and constant scatter in the cluster observable-mass scaling relations, and for a spatially flat model with a cosmological constant, we measure $\\gamma(\\sigma_{8}/0.8)^{6.8}=0.55^{+0.13}_{-0.10}$, with allowed values for $\\sigma_{8}$ in the range $0.79$ to $0.89$ (68.3 per cent confidence limits). Relaxing the assumptions on the scaling relations by introducing two additional parameters to model possible evolution in the normalization and scatter of the luminosity-mass relation, we obtain consistent constraints on $\\gamma$ that are only $\\sim 20$ per cent weaker than those above. Allowing the dark energy equation of state, $w$, to take any constant value, we simultaneously constrain the growth and expansion histories, and find no evidence for departures from either GR or the cosmological constant plus cold dark matter paradigm. Our results represent the most robust consistency test of General Relativity on cosmological scales to date. ",
        "introduction": "\\label{sec:introduction}  Since \\cite{Riess:98} and \\cite{Perlmutter:99} first demonstrated, using observations of type Ia supernovae (SNIa), that the expansion history of the Universe is accelerating at late times, improved SNIa studies have continuously corroborated these results \\citep{Garnavich:98, Schmidt:98, Knop:03, Tonry:03, Barris:04, Riess:04, Riess:07, Astier:06, Wood-Vasey:07, Kowalski:08, Hicken:09}. At the same time, measurements of anisotropies in the cosmic microwave background (CMB) with the Wilkinson Microwave Anisotropy Probe \\citep[WMAP][and companion papers]{Spergel:03,Spergel:07,Komatsu:09, Dunkley:09} and other CMB experiments \\citep{Readhead:04, Jones:06, Reichardt:09, Chiang:09, Gupta:09} have tightly constrained other key cosmological parameters. A variety of other cosmological data sets have also been used to independently measure cosmic acceleration. Measurements of the gas mass fraction ($f_{\\rm gas}$) in galaxy clusters, having earlier shown that the mean matter density of the Universe is low, $\\Omega_{\\rm m}\\sim 0.25$ \\cite[see e.g.][]{White:93}, have directly confirmed the effects of cosmic acceleration, at comparable significance to that seen in SNIa data \\citep[see e.g.][]{Allen:04, Allen:08}. Measurements of baryon acoustic oscillations (BAO) in galaxy surveys \\citep[see e.g.][]{Eisenstein:05, Percival:07, Percival:09}, the correlation of the low multipoles of the CMB with large scale structure observations from various surveys \\citep[see e.g.][]{Scranton:03, Fosalba:03, Cabre:06, Giannantonio:08, Ho:08}, and weak gravitational lensing by large scale structure \\citep{Schrabback:09} have also shown evidence for cosmic acceleration. The combination of subsets and/or all of these data have led to the establishment of the $\\Lambda$CDM paradigm, in which the energy density of the Universe is currently composed of $\\sim 5$ per cent baryons, $\\sim 20$ per cent cold dark matter (CDM) and $\\sim 75$ per cent dark energy, the last of which drives cosmic acceleration and has identical characteristics to Einstein's cosmological constant, $\\Lambda$.  Together, the above experiments robustly show that the Universe is accelerating. However, the ability of such data to probe the underlying cause of this acceleration is limited. All of the above constraints on dark energy are primarily driven by its effects on the background geometry. Such data alone cannot distinguish acceleration due to a true dark energy component with negative pressure from modifications to the standard theory of gravity, i.e. Einstein's General Relativity (GR). Recently, however, \\citet[hereafter \\depaper{}]{Mantz:08, Mantz:09a}, and \\cite{Vikhlinin:09} have presented new constraints on dark energy from measurements of the growth of cosmic structure, as evidenced in X-ray flux-selected cluster samples. These experiments are sensitive not only to the evolution of the mean energy density background, but also to the evolution of the density perturbations with respect to this background. This work has provided clear, independent confirmation of the effects of dark energy in slowing the growth of X-ray luminous galaxy clusters.  The observed evolution of density perturbations provides a sensitive probe of the underlying theory of gravity and the clustering\\footnote{Note that for models other than the cosmological constant, dark energy is expected to couple with gravity, i.e. cluster \\cite[see e.g.][]{Hu:05, Mota:07}. For specific dark energy and modified gravity models, see the reviews of \\cite{Copeland:06} and \\cite{Frieman:08}.} properties of dark energy. \\citet[hereafter \\citetalias{Rapetti:09}]{Rapetti:09} exploited this sensitivity to constrain departures from the predicted cosmic growth rate for GR using the convenient parameterization $\\Omega_{\\rm m}(z)^{\\gamma}$, for which GR has $\\gamma\\sim 0.55$ \\citep{Peebles:80, Wang:98, Linder:05, Huterer:06, DiPorto:07, Sapone:07, Gannouji:08, Wei:08,Linder:07, Nesseris:07}. Using the growth of structure analysis of \\cite{Mantz:08}, \\citetalias{Rapetti:09} found no evidence for deviations from either GR or $\\Lambda$CDM. Recent results from the combination of other cosmological data sets are also consistent with GR \\citep{Daniel:10, Zhao:10, Reyes:10}.  This is the third of a series of papers in which, for the first time, we employ a fully self-consistent analysis of the growth of massive clusters that combines current X-ray cluster surveys and deep, pointed, follow-up observations to simultaneously constrain both cosmological parameters and observable-mass (luminosity-mass and temperature-mass) scaling relations. Importantly, the improved method, which is described in \\depaper{}, properly accounts for all selection biases, covariances and parameter degeneracies, and models fully the impact of systematic uncertainties. Here, we utilize this improved method (hereafter XLF, as defined in \\depaper{}) to re-investigate the simultaneous constraints that can be placed on $\\gamma$ and the background evolution (expansion history). For the background, we use two reference expansion models: flat $\\Lambda$CDM, parameterized by the mean matter density, $\\Omega_{\\rm m}$; and flat $w$CDM, parameterized by $\\Omega_{\\rm m}$ and a constant dark energy equation of state, $w$. Our constraints on $\\gamma$ arise primarily from the XLF experiment, which uses the cluster samples of \\citet[][\\ROSAT{} Brightest Cluster Sample; BCS]{Ebeling:98}, \\citet[][\\ROSAT{}-ESO Flux Limited X-ray sample; REFLEX]{Bohringer:04}, \\citet[][2009 in preparation, Bright MAssive Cluster Survey; Bright MACS]{Ebeling:01}, plus extensive X-ray follow-up data from the {\\it Chandra} X-ray Observatory and {\\it ROSAT} \\cite[for details see][hereafter \\scpaper]{Mantz:09b}, although we also utilize the (currently) small additional constraining power available from measurements of the Integrated Sachs-Wolfe (ISW) effect at low multipoles in the CMB.  We emphasize that the models used for the evolution of the background and density perturbations are purely phenomenological, encompassing both modified gravity and clustering dark energy as possible sources for cosmic acceleration. Our analysis allows for a simple, but elegant test for departures from the standard GR and $\\Lambda$CDM paradigms. Using the combination of XLF, $f_{\\rm gas}$, SNIa, BAO and CMB, we demonstrate good agreement with the standard GR+$\\Lambda$CDM model and provide tight constraints on the model parameters. We show that our results are robust against reasonable assumptions regarding the evolution of galaxy cluster scaling relations. We stress that in order to obtain robust results from current and future XLF studies, it is essential to employ a consistent analysis method, such as the one used here (see details in \\depaper{}), in order to properly account for selection biases, covariances, parameter degeneracies and systematic uncertainties. Otherwise, spuriously tight constraints may be obtained. ",
        "conclusions": "\\label{sec:conclusions}  We have used measurements of the evolution of massive galaxy clusters from wide-area X-ray cluster surveys (spanning the redshift range $z<0.5$), and simultaneous measurements of the observable-mass scaling relations, to test the consistency of the observed growth rate of clusters with that predicted for GR. We allow departures from GR as described by the parameterization $\\Omega_{\\rm m}(z)^{\\gamma}$, for which $\\gamma\\sim0.55$ corresponds to GR. We find that the current data are consistent with GR and give tight constraints on departures from it. Assuming constant scatter and purely self-similar evolution of the cluster observable-mass scaling relations, and a flat $\\Lambda$CDM background evolution model, the combination of XLF, $f_{\\rm gas}$, SNIa, BAO, and CMB data gives $\\gamma(\\sigma_{8}/0.8)^{6.8}=0.55^{+0.13}_{-0.10}$. Allowing $w$ to vary, we simultaneously demonstrate consistency with the expansion history of $\\Lambda$CDM (i.e. $w=-1$).  Our analysis employs the method described in \\depaper{}, which self-consistently models X-ray cluster survey data and follow-up observations to simultaneously constrain both cosmological and scaling relation parameters. This allows us to constrain $\\gamma$ even when allowing for departures from self-similar evolution and/or constant scatter in the luminosity-mass relation. We caution that the adoption of such a self-consistent approach to model the survey+follow-up data is important, else spuriously tight, and potentially biased, constraints may be obtained.  Our measurements of the growth of the most X-ray luminous, most massive galaxy clusters within $z<0.5$ constrain the growth of cosmic structure during the epoch of cosmic acceleration. It is primarily this redshift dependence that provides our constraints on $\\gamma$, although we have also utilized the small contribution from the ISW effect, which probes larger scales at $z<2$. The inclusion of CMB and other data tightens significantly the constraints on the background expansion model and other cosmological parameters leading to a clear correlation between $\\gamma$ and $\\sigma_8$. This highlights the potential for further improvements in $\\gamma$ with the incorporation of precise, independent measurements of $\\sigma_8$."
    },
    "0911/0911.3147_arXiv.txt": {
        "abstract": "We develop a new formalism for modeling the formation and evolution of galaxies  within  a  hierarchical  universe.  Similarly  to standard semi-analytical   models   we   trace  galaxies   inside dark-matter merger-trees  which are  extracted from  a large $N$-body simulation. The formalism includes treatment of feedback, star-formation, cooling, smooth  accretion, gas  stripping  in satellite galaxies,  and merger-induced  star bursts. However, unlike  in other models, each process  is assumed to have  an efficiency which  depends only on the  host halo  mass  and redshift.  This  allows us  to describe  the various components of the model in a simple and transparent way. By allowing the efficiencies to have any value for a given halo mass and redshift, we can easily encompass a large range of scenarios. To demonstrate this point, we examine several different galaxy formation models, which are all consistent with  the  observational  data.  Each  model  is  characterized  by a different  unique feature:  cold accretion  in low  mass haloes, zero feedback, stars  formed only in merger-induced  bursts, and shutdown of star-formation after mergers. Using these  models we are able  to examine the  degeneracy inherent in galaxy  formation models, and  look for  observational data  that will help to break this degeneracy. We  show that the full distribution of star-formation rates in  a given  stellar mass  bin is promising in constraining  the models. We compare our  approach in detail to the  semi-analytical model of De Lucia \\&  Blaizot. It is shown  that our formalism is  able to produce a very similar population of  galaxies once the same median efficiencies per halo mass and redshift  are being used. Consequently, our approach may  be useful for comparing various published  semi-analytical models within the same framework. We provide a public version of the model galaxies on our web-page, along with a tool for running models with user-defined parameters. Our model is able to provide results for a 62.5 $h^{-1}$ Mpc box within just a few seconds. ",
        "introduction": "\\label{sec:intro}  The physics of galaxy formation combines a wide range of astrophysical phenomena.  On one  side it  relies heavily  on cosmology  and  on the hierarchical formation  of dark matter structures in our Universe. On  the other hand galaxy formation  is greatly affected by the  detailed physics of gas  dynamics, star-formation (SF)  processes, stellar  evolution, and black-hole physics.  A model that  will combine all the  above physics into a well-defined, coherent  methodology should thus include a large dynamical range of scales, from the formation of  individual stars to the formation of the largest clusters of galaxies.  Historically, a  detailed understanding  of the formation  of galaxies has  grown   alongside  with  the  development  of   the concepts  of dark-matter clustering \\citep[e.g.][]{Rees77,Blumenthal84,White91}. It became clear that galaxy  formation is  a `two-stage  process', where  first dark matter clumps collapse and virialize into `haloes'. This is then followed by a  second stage  of  gas cooling into  the central regions of  these haloes, which can condense and  form stars \\citep{White78}.  In the last few decades there   have  been substantial   improvements in   our  quantitative understanding  of structure  formation of dark-matter, both  from the analytical point of view \\citep{Press74,Bond91,Sheth02} and also using numerical simulations \\citep[e.g.][]{Lacey94,Navarro97,Springel05}.  The current level of accuracy  in modeling  dark-matter  structure growth  has reached the level of few percents. Some larger uncertainties are still inherent in the  halo definition  and  its virial mass,  but the way these uncertainties affect galaxies is partially related to the virialization of gas inside haloes, and thus cannot be fully understood using dark-matter only.  Unlike  the physics  of  dark matter, the  processes  that govern gas dynamics and SF still remain highly uncertain, usually at the level of at least an order of magnitude. For example, cooling rates may be dramatically affected by the assumed shape of the density profiles or  different assumptions made in computing the cooling function \\citep[e.g.][]{Efstathiou92,Landi99,Maio07,Gnat07, Kaufmann09, Wiersma09}. Also, both star formation and supernova feedback are very complex processes whose basic characteristics are still under debate \\citep[e.g][]{McKee77,Dekel86,Elmegreen97, Maclow04,Murray05,Scannapieco06,Leroy08,Fumagalli09}.    In  view of  the  above  difficulties to  model  galaxies, three main methodologies were developed in  the literature. The first approach is a statistical one, attempting to interpret  the data with a minimum set of assumptions.  This includes  variants of  the halo  model, where the halo and galaxy abundances are linked \\citep[see e.g.][]{Jing98,Peacock00,Seljak00,Cooray02,Kravtsov04,Zehavi05,vdBosch07}. The approach has been useful in  quantifying the relationships between  galaxies and haloes, the clustering  properties of both populations, and  the properties of satellite galaxies.  However, by definition, this formalism does not allow to follow the evolution of  galaxies  with  time,  and  to  study their formation  histories \\citep[but see][]{Drory08,Conroy09}.  A totally different approach is the detailed  modeling  of individual galaxies  using  high  resolution simulations which include  gas   hydrodynamics,   and  a   simple prescriptions for SF, feedback and black hole growth \\citep[e.g.][]{Abadi03,Scannapieco09}.  This approach  is  valuable for testing specific  ingredients of  the models (i.e.  different feedback mechanisms)   and   it   is   critical  to   deepen our physical understanding of these processes. However, it is  still difficult to obtain a statistical sample of galaxies using such simulations, or to span a large range of possible models \\citep[see a recent attempt by][]{Schaye09}. In addition,  processes like  star formation and supernova (SN) feedback occur far below the resolution limits of these simulations, which makes  it necessary to describe them with analytical and approximate recipes. The finite  number  of particles representing the baryonic content of galaxies, and the presence of numerical artifacts make it difficult to find appropriate recipes which lead to numerical convergence while producing realistic model galaxies.  The  most widely  used methodology  to interpret  observations  and to provide  a coherent  picture for  the  physics of  galaxy formation  is implemented by  Semi-Analytical Models (hereafter SAMs). The approach was  introduced  by  \\citet{White91}  and then implemented  by \\citet{Cole91,Kauffmann93b}.    More recent    works    include \\citet{Somerville99b,Cole00,Benson02,Hatton03,Bower06,Cattaneo06, Croton06,DeLucia07,Monaco07,Kang05,Somerville08}. These models provide a detailed picture of galaxy formation processes and produce a population of simulated galaxies whose global statistical properties   resemble   the    observations in   many aspects. Consequently, the models can then be used to examine specific issues in the formation of  galaxies, like the impact of different processes on the shape of the luminosity function  \\citep{Benson03},   the  formation history   of  elliptical galaxies    \\citep{DeLucia06},      and galaxy      merger-rates \\citep{Guo08}. SAMs are also valuable for developing new observational techniques \\citep{Chen09}.  The SAM approach is in an intermediate position between the detailed hydrodynamical simulations and variants of the halo model discussed above. While in contrast to hydrodynamical simulations, SAMs include no 3D information on individual galaxies, they follow the evolution of galaxies within their dark-matter haloes, which cannot be done using the `halo model'. In SAMs, halo merger-trees  are  used  as  a skeleton for the  evolution  of galaxies. The first generation of models used Monte-Carlo realizations based on the extended Press-Schechter approach, while modern SAMs use trees extracted from $N$-body simulations. Recent versions of these models include merger-trees which follow the evolution of substructure inside haloes. Within these  haloes galaxies are assumed to grow,  with each  galaxy being described by a small number  of parameters. This means that only a few mass components are followed with time: hot gas, cold gas, stars, black-hole mass, gas which was ejected out  of the  halo, heavy elements and the fraction of mass contained in the bulge. SAMs include recipes that are used to compute the rate of change in each component, according to the properties of the galaxy and its host halo. For example, the SF rate is usually assumed to be proportional to the mass of cold gas in the disk, divided by the disk time-scale. These recipes then enable  a fast and detailed evolution  of galaxies within the model. The benefit of SAMs over other approaches is that using a few simple assumptions  the model can reproduce the detailed formation histories of galaxies.   Although SAMs have been  quite successful in generating model galaxies that resemble  observations, there are a few limitations to this method that we would like to  address. First, it seems that the degeneracy inherent in these models has not been sufficiently explored. Often a fixed functional form is used to describe a given process, and tuning is only done by varying parameters. In addition, it is not always tested whether the more complex models can be replaced by simple recipes. For example, the physics of AGN feedback is poorly understood, but nevertheless implemented with a relatively detailed recipe in certain SAMs \\citep[e.g.][]{Croton06,Monaco07}. \\citet{Cattaneo06} and \\citet{Bower06} have shown that simply suppressing cooling within massive haloes gives results which are very similar to the complex models mentioned above \\citep[see however the discussion in][]{DeLucia06}. This means that the recipe for AGN feedback can be easily summarized by a range of halo mass and redshift for which the suppression is active. This fact cannot directly be seen from the detailed equations being used, and the effective range  in halo mass  where this  suppression  is active is not  well examined.   A different limitation of SAMs is  that some recipes are hard-wired into a complex simulation code and inter-connected in non-trivial ways. It can thus be difficult for users to change the functional form of these recipes. For example, cooling efficiencies are computed using information on the density profile of hot gas within haloes, combined with the physics of radiative cooling processes. There is no practical way to change the effective cooling efficiencies for ranges in halo mass and redshift, and to test its effect on the population of galaxies. In addition, testing a large number of very different recipes is nearly impossible, and it can be very hard or impossible to tune a SAM in order to match observations if no freedom in the functional forms is allowed. For example, as  was shown by \\citet{Fontanot09}, various  SAMs are unable to reproduce the decrease of  the specific SF  rate as a  function of stellar mass. Up to now, it has been unclear if this limitation is fundamental to the SAM approach  or not. We will show below  that once the functional shape of the  recipes is relaxed, this  observational feature can easily be reproduced.   Another  minor  disadvantage is that SAM  recipes depend occasionally  on  a  large chain  of computational  steps, which  do not  allow  the reader  to extract  the actual  values being used. For example,  it is difficult to estimate the actual cooling and SF efficiencies used by a given SAM, which makes it nearly impossible to compare different SAMs easily.   In this work we develop a new approach for following a galaxy within a given merger-tree.  We present   a formalism that  is more simple  and transparent  than the  usual  implementation of  a set of recipes, but on the other hand  general and flexible enough so that  model  galaxies  will  reproduce the observational constraints accurately. We will show that it is possible to simplify the SAM considerably by only allowing the recipes to depend on halo mass and redshift. We find that this  simplification does not reduce the  level of complexity in the  modeled galaxies, but it enables a very simple and transparent description of the various physical processes. As a second step, in  order to keep the  model general and flexible enough, we do not restrict  the functional dependence of each process on  the halo mass  and redshift. This generalization greatly facilitates the tuning process. The formalism we  propose is more compact  than usual  SAMs  and can be represented  by a  few simple functions.  The main assumption in our formalism is that each physical process can be described well  enough by its dependence on the  host halo mass and redshift.  Few SAM recipes  are already following this assumption. For example, the SN feedback, including reheating, ejection and reincorporation of disk gas, only depends on the virial velocity and virial mass of the host halo in \\citet{DeLucia07}. There are  other recipes which  are  in general not a  simple function of  halo mass and time. For example, the cooling  efficiency is usually modeled  in SAMs by following  the density profile  of the hot gas component, and  computing the radius within which  gas has had time to cool. By extracting cooling efficiencies and star formation efficiencies directly from the SAM of \\citet[][hereafter DLB07]{DeLucia07} we test our main assumption and show that it is relatively accurate for all the recipes used in this model.   In this paper we use our  formalism to explore the level of degeneracy inherent in current galaxy formation models.  We try to examine very different models  of   galaxy  formation  that  are  all consistent  with  the observational data.   We deliberately consider also  some very extreme and physically less well motivated models in order  to span a large  range of solutions  allowed by the observations.   This will help us better to understand the degeneracies involved in SAMs.  Specific questions  can then be addressed by comparing  the  different  modeled  galaxies. For example, we derive the minimum cooling efficiencies which are consistent with the observed population of galaxies; we examine how much SN feedback is required in a model with instantaneous cold accretion at low halo masses; and we investigate a model in which star formation and cooling terminates after major mergers.   This paper is organized  as follows. We start with a review of the current SAM ingredients and their uncertainties in section \\ref{sec:motivation}. In section \\ref{sec:formalism} we present  our  general formalism  and explain  where  it differs from previous works. Our approach is then compared with a state-of-the-art SAM in \\S\\ref{sec:test}. Section \\ref{sec:models} summarizes the various ingredients of each specific model we develop here. These models are further discussed in \\S\\ref{sec:results} where their results are compared to observational data. We discuss our results and summarize them  in \\S\\ref{sec:discuss}. Throughout this paper  we use the cosmological model adopted  by the Millennium simulation  with $(\\omm,\\,\\oml,\\,h,\\,\\sigma_8) =(0.25,\\,0.75,\\,0.73,\\,0.9)$.  We  use $\\log$ to designated $\\log_{10}$, $t$ is the time in Gyr since the big-bang. ",
        "conclusions": "\\label{sec:discuss}  In this paper we develop a new formalism for the formation and evolution of galaxies within a hierarchical universe. Our approach is similar to semi-analytical models (SAMs), although we try to be more simple and schematic in order to encompass a large range of possible models.  We claim that a significant step towards more simple, general, and transparent models is achieved by parameterizing the basic processes of galaxy formation as functions of host halo mass and redshift alone. This language enables an easy and compact summary of the model, and results in a well defined set of differential equations. Once each recipe is just a function of halo mass and redshift it is also simple to check the dependencies between recipes, and to determine those that will result in a population of galaxies which matches the observational constraints.  We compare our approach to the SAM of \\citet{DeLucia07} and show that also within this specific model, the recipes are effectively a function of halo mass and redshift. The only recipe which deviates from this approximation to some degree is the radiative gas cooling. However, for the range where the deviations are important (i.e. small mass haloes), it is still not clear how accurate the cooling approach adopted by this SAM is. Using the median values per halo mass and redshift for each recipe calculated from the original SAM, our model can produce a very similar population of galaxies. This demonstrates that forcing the physical recipes to depend only on halo mass and redshift does not affect the results of the SAM substantially, at least for the properties examined here.  The general language we develop here might be useful for comparing the various SAMs currently in use and for highlighting differences between models. It can also serve as a tool for comparing hydrodynamical simulations against SAMs. For example, it might be possible to obtain average feedback efficiencies as a function of halo mass and redshift using detailed simulations, and then to insert this information into the SAMs. This simple language might help to improve the communication between the different methodologies.  The approach adopted here is very flexible as we can allow any efficiency value for a given host halo mass and redshift. We use this flexibility to explore very different scenarios of galaxy formation. We present one standard model, which is based on \\citet{DeLucia07} but without an ejected phase or a threshold mass for star formation. We then introduce four scenarios, each characterized by a unique feature: zero feedback, cold-accretion in low mass haloes, stars formed only in merger-induced bursts, and shutdown of star-formation after mergers. In each model we tune the other processes so that the resulting galaxies match observations. We do not argue that all these models are fully plausible, but  show how different processes like feedback, cooling, and merger-induced bursts can compensate for each other. This also provides some insight into the range of efficiencies allowed by observations of the global properties of the galaxy population over time. For example, with our Model I, we are able to put a lower limit on the allowed cooling efficiencies given the general properties of the galaxy population over time.  We test the five different models presented here against various observations. We show that all of them reproduce the observed stellar mass functions at all redshifts reasonably well. We explore inconsistencies related to the universal SF rate density, and show the level of agreement that can be achieved between these two observations. Surprisingly, we find that all our five models reproduce the observed tilt in the relation of specific SF rates versus stellar mass, unlike other current SAMs \\citep[see e.g.][]{Somerville08,Fontanot09}. This is likely related to the way that ejection and cooling is implemented in current SAMs. We plan to investigate this issue in future work.   It is still difficult to assess what it really means if a given SAM manages to reproduce observations, as degeneracies are not well understood. How can we know whether or not a given SAM presents a relatively unique solution to the problem of galaxy formation and evolution? We try to improve on these limitations by presenting a set of models which are very different from each other. This can help us revisit basic issues of galaxy formation models. For example, we plan to investigate how galaxy merger-rates behave in each model, what causes the distinction between red and blue galaxies, and what shapes the stellar mass function.  We use the set of models developed here to investigate which observational quantity will break the degeneracy and allow us to differentiate between the models. The most promising such quantity is the full distribution of SF rates in different stellar mass bins. Specifically, SF rates for passive galaxies can give us information on the physics of SF although they do not affect stellar mass function, or the cosmic SF density. On the other hand, at low SF rates it might be that the processes which regulate SF are very different from the main mode of SF. We find that it is surprisingly difficult for different models to reproduce the passive fractions of galaxies as a function of stellar mass, which indicates that this observations may also be useful for constraining the models. In addition, cold gas fractions might be useful for discriminating between different models which all produce a similar galaxy population in terms of stellar mass and SF. Unfortunately, up to now published datasets do not take into account contributions from the atomic and molecular gas components \\citep[for estimates of gas mass ratios see for example][]{Bothwell09,Obreschkow09}.  The formalism presented here can be viewed as an intermediate step between the halo model and SAMs. This is because all recipes depend on halo mass only, but the galaxies are followed inside the complex structure of merger-trees, and their properties will depend strongly on the halo formation history. Consequently, galaxies which live inside haloes of the same mass might have very different SF histories. Similarly to the halo model, our model does not need to be parameterized a-priori, and the functional dependence of the recipes is obtained by matching the observations. However, unlike in the halo model, the ingredients we match (efficiencies of cooling, SF, and feedback) are directly connected to the underlying physical processes, and we can follow galaxies in detail through time. This level of modelling is more complex than the halo model, demands more computational time, and involves more assumptions.  The current version of our model does not include a computation of luminosity and color. This limits potential comparisons between our model galaxies and observations. For example, other SAMs usually compare the model against luminosity functions, color-magnitude diagrams, and mass-metallicity relations. On the other hand, comparing to these observations will require additional model ingredients like dust obscuration, metallicity enrichment, and the initial stellar mass function (IMF). It seems that from the above additional observational measurements, only the mass-metallicity relation is adding a new constraint to our model. However, the number of ingredients that would need to be added to the model is quite large, so we suspect that the level of degeneracy will only increase. An additional shortcoming of our model is that it does not include any morphological information on galaxies. Since the establishment of the Hubble sequence it is well known that correlations between morphologies and SF rates do exist \\citep[see a recent work by][]{Schiminovich07}. However, it is unclear if this correlation is a consequence of the underlying physics, or if it plays an active role in shaping the SF histories of galaxies as suggested by \\citet{Martig09}. We plan to examine galaxy morphology in a future work.  We do not attempt to provide improved recipes for SF, cooling or feedback, even though finding such recipes, which are physically motivated as well as in agreement with observations both on global and local scales, is one of the ultimate goals of semi-analytical modelling. We however believe that our method can help finding such recipes, and studying the interconnection between them, by opening up the allowed parameter space.  The usual philosophy behind SAMs is that assumptions on the relevant physical processes are made, and it is then checked whether the properties of observed galaxies can be reproduced by tuning a certain set of parameters connected with these physical processes. We try to solve the inverse problem: how the observed population of galaxies can constrain the physical mechanisms which drive galaxy formation. The outcome of our models are average efficiencies per halo mass and redshift, which need a further interpretation in order to have a physical meaning. For example, our different results for the cooling efficiencies need to be confronted with physical models of cooling, in order to find a scenario which is able to predict such deviations from the standard model. It might be that deviations due to e.g. metallicity variations and gas density profiles would not be able to give the required dependencies on halo mass and redshift. This dependence is more important than our results on the absolute value of the same quantity, which we have shown is only poorly constrained by observations. Lastly, the fact that our recipes do not rely on detailed physical assumptions means that they generally should not be extrapolated to ranges in halo mass and redshift for which they have not been tuned. Such problems may also occur in standard SAMs, although in these cases some guidance for the extrapolation of each recipe is given.   We note that the cosmological model used here is slightly different in its parameters from the most recent estimate by \\citet{Komatsu09}. It might be possible to scale the merger-trees from the Millennium cosmology into the recent cosmology, by changing the time variable as was suggested by \\citet{Neistein09}. If scaling of the subhalo merger-trees will only require changing the time variable, then it might be possible to scale our galaxy formation models by simply transforming only the dependence of the \\emph{recipes} on time. This should leave the results of the models unchanged. However, while \\citet{Neistein09} found that time scaling should work well for \\fof merger-trees, we use merger-trees based on subhaloes in this work, which might scale somewhat differently.   We hope that the models developed in this work will help to interpret both observational data and different galaxy formation models. To serve this goal we provide a public online access to catalogs of our model galaxies. Our online web page\\footnote{\\texttt{http://www.mpa-garching.mpg.de/galform/sesam}} is also able to run various galaxy formation models with user-defined parameters over a box size of 62.5 Mpc $h^{-1}$. This volume represents the results obtained from the full Millennium simulation with a sufficient accuracy for many applications. We allow users to upload any kind of recipe which depends on halo mass and redshift. Results of the model are obtained within a few seconds and can be downloaded for further analysis."
    },
    "0911/0911.5347.txt": {
        "abstract": "Correlations between the intrinsic shapes of galaxy pairs, and between the intrinsic shapes of galaxies and the large-scale density field, may be induced by tidal fields.  These correlations, which have been detected at low redshifts ($z<0.35$) for bright red galaxies in the Sloan Digital Sky Survey (SDSS), and for which upper limits exist for blue galaxies at $z\\sim 0.1$, provide a window into galaxy formation and evolution, and are also an important contaminant for current and future weak lensing surveys. Measurements of these alignments at intermediate redshifts ($z\\sim 0.6$) that are more relevant for cosmic shear observations are very important for understanding the origin and redshift evolution of these alignments, and for minimising their impact on weak lensing measurements.  We present the first such intermediate-redshift measurement for blue galaxies, using galaxy shape measurements from SDSS and spectroscopic redshifts from the WiggleZ Dark Energy Survey.  Our null detection allows us to place upper limits on the contamination of weak lensing measurements by blue galaxy intrinsic alignments that, for the first time, do not require significant model-dependent extrapolation from the $z\\sim 0.1$ SDSS observations.  Also, combining the SDSS and WiggleZ constraints gives us a long redshift baseline with which to constrain intrinsic alignment models and contamination of the cosmic shear power spectrum.  Assuming that the alignments can be explained by linear alignment with the smoothed local density field, we find that a measurement of $\\sigma_8$ in a blue-galaxy dominated, CFHTLS-like survey would be contaminated by at most $^{+0.02}_{-0.03}$ (95 per cent confidence level, SDSS and WiggleZ) or $\\pm 0.03$ (WiggleZ alone) due to intrinsic alignments.  We also allow additional power-law redshift evolution of the intrinsic alignments, due to (for example) effects like interactions and mergers that are not included in the linear alignment model, and find that our constraints on cosmic shear contamination are not significantly weakened if the power-law index is less than $\\sim 2$.  The WiggleZ sample (unlike SDSS) has a long enough redshift baseline that the data can rule out the possibility of very strong additional evolution. ",
        "introduction": "\\label{S:introduction} %=====================================================================  Gravitational lensing, the deflection of light due to matter between the source and the observer, is sensitive to all matter (including dark matter).  As a result, in the past decade, weak gravitational lensing \\citep{2001PhR...340..291B,2003ARA&A..41..645R} has become a powerful tool for addressing outstanding questions related to cosmology and galaxy formation.  Its scientific applications include measurement of the amplitude of matter fluctuations using the auto-correlation of galaxy shapes \\citep[e.g., most recently, ][]{2006ApJ...647..116H,2006A&A...452...51S,2007MNRAS.381..702B,2007ApJS..172..239M,2009arXiv0911.0053S}, known as cosmic shear; and determination of the relationship between the baryonic content and the dark matter content of galaxies, using the cross-correlation between background galaxy shapes and foreground galaxy positions \\citep[e.g., ][]{2005ApJ...635...73H,2006MNRAS.371L..60H,2006MNRAS.368..715M}, known as galaxy-galaxy lensing.  Because of the utility of these applications of lensing, plus its potential to constrain models of dark energy by splitting the sample of source galaxies into redshift slices \\citep[tomography: ][]{2002PhRvD..66h3515H,2002PhRvD..65f3001H}, future surveys are being planned to measure the lensing signal with sub-per cent statistical errors.  There is a large body of work devoted to solving the technical problems in measuring the weak lensing signal, primarily related to unbiased shear estimation \\citep{2006MNRAS.368.1323H,2007MNRAS.376...13M,2009arXiv0908.0945B} and to photometric redshifts \\citep{2004ApJ...600...17B,2005PhRvD..71b3002I,2006MNRAS.366..101H,2008MNRAS.387..969A}. In this work, we focus on a source of astrophysical uncertainty, intrinsic alignments of galaxy shapes.  When measuring the lensing signal, it is assumed that in the absence of lensing, galaxy shapes are uncorrelated.  Intrinsic alignments are alignments of galaxy shapes that violate that assumption, for example due to the alignment of galaxy shapes with a local tidal field.  These alignments can therefore contaminate the gravitational lensing signal.  One type of intrinsic alignment is the correlation between the intrinsic ellipticities of two galaxies (II correlations) that reside in the same local or large-scale structure.  Cosmological $N$-body simulations robustly predict alignments between the shapes of dark matter halos that are a declining function of separation \\citep{1997ApJ...479..632S,2000MNRAS.319..614O,2002A&A...395....1F,2005ApJ...618....1H,2008MNRAS.389.1266L}. However, the true observational impact of these alignments are difficult to estimate using $N$-body simulations, because the observed alignments depend on the shape of the baryonic component of the galaxy rather than on dark matter alone.  While several analytical models for these alignments have also been developed \\citep{2000ApJ...545..561C,2000MNRAS.319..649H,2001MNRAS.320L...7C,2001ApJ...559..552C,2002MNRAS.335L..89J}, they predict wildly varying levels of alignment, so observational constraints are necessary.  More recently, \\cite{2004PhRvD..70f3526H} pointed out that the correlation of galaxy shapes with large-scale density fields can also contaminate lensing measurements.  These alignments, known as GI correlations, are caused by a lower redshift tidal field that both causes gravitational shear experienced by a higher redshift galaxy, and intrinsically aligns the shape of galaxies that are in the tidal field.  GI correlations are also predicted to have very different magnitudes depending on the model used to estimate them \\citep{2002astro.ph..5512H,2004PhRvD..70f3526H,2006MNRAS.371..750H}, and have been detected using dark matter halos at many different mass scales in $N$-body simulations \\citep{2005ApJ...627..647B,2006MNRAS.370.1422A,2006MNRAS.365..539B,2006MNRAS.371..750H,2007ApJ...671.1135K}.  The relevant signature of these GI correlation detections in $N$-body simulations is that dark matter halos align so that they point preferentially towards other halos that are part of the same large-scale structure.  When considering GI correlations of galaxies that contaminate cosmological weak lensing measurements, the effect is manifested as galaxy shapes that point preferentially towards other galaxies (both locally, within a halo, and on cosmological scales).  These alignments of galaxy shapes are anti-correlated with the gravitational shear due to large-scale structure, so GI correlations reduce the measured cosmic shear signal, unlike the II correlations which increase it.  Also unlike the II correlations, the GI correlations are not due to the inclusion of galaxy pairs at the same redshift; pairs at different redshifts are affected when the higher redshift galaxy of the pair is lensed by a structure that has caused an intrinsic alignment of the lower redshift galaxy.  Several different schemes have been proposed to remove intrinsic alignment contamination from weak lensing measurements, including the removal of galaxy pairs that are close in redshift space (to remove II: \\citealt{2002A&A...396..411K,2003A&A...398...23K,2003MNRAS.339..711H,2004ApJ...601L...1T}); projecting out both types of intrinsic alignments using their known scalings with the redshifts of the galaxy pair \\citep{2004PhRvD..70f3526H,2008A&A...488..829J,2008arXiv0811.0613Z,2009A&A...507..105J}; and modeling them jointly with the lensing signal using some parametric models, the parameters of which are then marginalised over \\citep{2005A&A...441...47K,2007NJPh....9..444B}.  A common feature of these methods is a loss of information, and therefore weakening of cosmological constraints from the weak lensing signal.  Measurements of intrinsic alignments can place strong priors on the intrinsic alignment model, which would minimise the loss of cosmological information from future surveys.  Direct intrinsic alignment measurements will also constrain the impact of intrinsic alignments on previous lensing measurements that did not explicitly account for them.  This measurement is particularly important given the aforementioned difficulty in theoretical predictions; however, the observations can then be used to refine the theory and, in turn, learn something about galaxy formation and evolution.  To observe intrinsic alignments, we require a source of data with robust galaxy shape measurements free of contamination from the PSF, and a way of isolating nearby (in all three dimensions) galaxy pairs. The GI correlations are then measured by calculating, statistically, the tendency for galaxies to point towards other galaxies that are relatively nearby (on tens of \\hmpc\\ scales).  Alternatively, it is possible to measure II correlations at low redshift without any redshift information, given that the cosmic shear signal below $z\\sim 0.2$ is vanishingly small \\citep{2002MNRAS.333..501B}.  On the large scales used for cosmological lensing analyses, the first measurement of GI correlations used SDSS data \\citep{2006MNRAS.367..611M}, with a follow-up analysis by \\cite{2007MNRAS.381.1197H} that also included redshifts of SDSS Luminous Red Galaxies (LRGs) and from the 2dF-SDSS LRG and QSO survey (2SLAQ, \\citealt{2006MNRAS.372..425C}) to constrain intrinsic alignments of red galaxies up to intermediate redshifts, $z\\sim 0.5$--$0.6$.  While GI correlations were detected in these works at $z\\sim 0.1$--$0.4$ for bright red galaxies (and II correlations for the same galaxy sample were found by \\citealt{2009ApJ...694..214O}), with a weak ($2\\sigma$) detection at intermediate redshifts (due to the small size of the 2SLAQ sample), further work at intermediate to high redshift is crucial for constraining the impact of intrinsic alignments on cosmological lensing analyses.  It is difficult to extrapolate these low-redshift analyses to higher redshift, because different dynamical scenarios might entail very different redshift evolution.  For example, blue galaxies at $z\\approx 0.1$ that have no measurable GI alignment in SDSS may have been very highly aligned with the density field at intermediate-high redshift, with mergers and interactions serving to disrupt those alignments, leading to the null detection that we see at low redshift; or, these alignments might be very small at all redshifts.  Due to its overlap with the SDSS, which provides galaxy shape measurements, the WiggleZ Dark Energy survey \\citep{2010MNRAS.401.1429D} is an ideal source of spectroscopic redshifts for constraining intrinsic alignments of UV-selected blue galaxies at intermediate redshift.  In this work, we use that sample to attempt the first measurement of galaxy intrinsic alignments for blue galaxies at intermediate redshift, which fills in a very important gap in our knowledge of intrinsic alignments.  At higher redshift, we expect that blue galaxies will dominate the galaxy samples used for weak lensing. Thus, these observations will facilitate further development in the fields of weak lensing and galaxy dynamics and evolution.  Here we note the cosmological model and units used throughout this work. Pair separations are measured in comoving $h^{-1}\\,$Mpc (where $H_0=100h\\,$km$\\,$s$^{-1}\\,$Mpc$^{-1}$), with the angular diameter distance computed in a spatially flat $\\Lambda$CDM cosmology with $\\Omega_m=0.3$.  %For the redshift of the WiggleZ %sample, varying $\\Omega_m$ within the $1\\sigma$ allowed range of %values according to the WMAP5 results \\citep{2009ApJS..180..330K} %changes the calculation of the %transverse separations by at most 3 per cent, which leads to changes %in interpretation that are far below the statistical error on the %results. %cospars_vect_base=[0.3 -1 0 0.70 0.75 0.05 1 0 0.05 0 1 1 0 0 0.5]; %For the estimates of the contamination of the cosmic shear signal, we adopt the following set of cosmological parameters: (\\rsm{Sarah: give numbers and reference}). For the bias and cosmic shear calculations, we additionally normalise the matter power spectrum using $\\sigma_8=0.75$, set the baryon density $\\Omega_b=0.05$ and scalar primordial spectral index $n_s=1$, and use the transfer function from \\cite{1996ApJ...471...13M}.  We begin in Section~\\ref{S:formalism} with a summary of the intrinsic alignment and cosmic shear formalism used in this work. Section~\\ref{S:data} contains descriptions of data used for the analysis.  The methodology used for the data analysis is described in Section~\\ref{S:methodology}.  We present the results of the analysis in Section~\\ref{S:results}, including systematics tests and a comparison with previous observations.  The interpretation of these results, including an estimate of  contamination of the cosmic shear signal, is given in Section~\\ref{S:interpretation}, and we conclude in Section~\\ref{S:conclusions}. ",
        "conclusions": "\\label{S:conclusions}  In this paper, we have placed the first direct observational constraints on the intrinsic alignments of blue galaxies at intermediate redshift ($z\\sim 0.6$), using the WiggleZ spectroscopic redshifts and galaxy shape measurements from SDSS.  We followed a comparable procedure as has been used before in SDSS at low redshifts ($z\\sim 0.1$--0.3) for blue and red galaxies in \\cite{2006MNRAS.367..611M} and \\cite{2007MNRAS.381.1197H}.  This procedure relies on finding pairs of galaxies that are physically associated in terms of their three-dimensional separation, and calculating the correlation between their shapes, and between the shape of each galaxy with the line connecting their positions on the sky.  Our result was a null measurement for the full WiggleZ sample and for two redshift subsamples.  This null measurement can in turn be used to constrain parameters of physically-motivated intrinsic alignment models, and to constrain the contamination of cosmic shear observations due to intrinsic alignments of galaxies that are comparable to this sample.  We have found that if we assume a model involving linear alignment with the smoothed local density field, then we can constrain the intrinsic alignment contamination for a CFHTLS-like survey dominated by WiggleZ-like galaxies to be small enough that $\\sigma_8$ is biased by an amount %SLB [I think this phrasing about CLs is OK] %SLB ($\\pm 0.03$ at 95 per cent CL) %($\\pm 0.02$ at 95 per cent CL) that is smaller than the statistical errors.  If we allow additional power-law redshift evolution in these alignments on top of the redshift evolution that is encoded in the linear alignment model, then we see that the constraints for $|\\eta_{\\rm other}| < 2$--$4$ do not significantly weaken, and the models with very large $\\eta_{\\rm other}$ can be ruled out because the data cover a fairly long redshift baseline (but see footnote 2 in section~\\ref{SS:lamodelfits} for a cautionary note on interpreting these $\\eta_{\\rm other}$ values). % that they can rule out %model combinations with both large power-law index and intrinsic %alignment amplitude. Combination with low-redshift SDSS results, which may be valid if the UV-selected WiggleZ galaxies have comparable formation histories to optically-selected $L_*$ blue cloud galaxies in SDSS, allows for tightening of these constraints, particularly due to the ability to rule out models with strongly increasing intrinsic alignments at low redshift.  As previously noted, theoretical models of intrinsic alignments are currently poorly constrained by the data, so direct measurements of these alignments are necessary to estimate how %SLB important serious %SLB [overuse of \"important\" in this paragraph] the alignments are for current and future cosmic shear surveys.  These observations of blue galaxies at intermediate redshift fill in an important gap in our knowledge of intrinsic alignments.  Our constraints were predominantly phrased in terms of the linear alignment model.  However, we note that our null measurement if II alignments could also be used to constrain other intrinsic alignments models, such as the quadratic alignment model that predicts no GI alignments but potentially significant II.  While blue galaxies tend to dominate cosmic shear samples, strong alignments of red galaxies at intermediate-high redshift could still be %SLB important. a significant contaminant. As a result, it will be important to obtain similar constraints of intrinsic alignments of red galaxies, which are poorly constrained for $z>0.4$.  Future work with our measurements of intrinsic alignments in WiggleZ could also focus on considering other physically-motivated intrinsic alignments models, and on combining pre-existing constraints for red galaxies at low redshift with our constraints for blue galaxies to come up with an estimate of intrinsic alignment contamination to cosmic shear surveys with a more realistic blue plus red galaxy sample.  Several methods have been proposed to remove the intrinsic alignment signal from future cosmic shear surveys \\citep[e.g.][]{2005A&A...441...47K,2007NJPh....9..444B,2008A&A...488..829J,2009A&A...507..105J,2008arXiv0811.0613Z,bernstein08,2009arXiv0911.2454J}.  In general, these methods rely on using some of the weak lensing signals (auto- and cross-correlations) to constrain parameters of the intrinsic alignment models, resulting in a loss of cosmological information.  Our measurements in this paper using the WiggleZ dataset will allow for the placement of stronger priors on the intrinsic alignment models, and therefore minimise this loss of cosmological information, preserving the cosmological constraining power of future datasets."
    },
    "0911/0911.3001_arXiv.txt": {
        "abstract": "We perform axisymmetric relativistic magnetohydrodynamic (MHD) simulations to investigate the acceleration and collimation of jets and outflows from disks around compact objects. Newtonian gravity is added to the relativistic treatment in order to establish the physical boundary condition of an underlying accretion disk in centrifugal and pressure equilibrium. The fiducial disk surface (respectively a slow disk wind) is prescribed as boundary condition for the outflow. We apply this technique for the first time in the context of relativistic jets.  The strength of this approach is that it allows us to run a parameter study in order to investigate how the accretion disk conditions govern the outflow formation. Substantial effort has been made to implement a current-free, numerical outflow boundary condition in order to avoid artificial collimation present in the standard outflow conditions. Our simulations using the PLUTO code run for 500 inner disk rotations and on a physical grid size of 100x200 inner disk radii. The simulations evolve from an initial state in hydrostatic equilibrium and an initially force-free magnetic field configuration. Two options for the initial field geometries are applied - a hourglass-shaped potential magnetic field and a split monopole field. Most of our parameter runs evolve into a steady state solution which can be further analyzed concerning the physical mechanism at work. In general, we obtain collimated beams of mildly relativistic speed with Lorentz factors up to 6 and mass-weighted half-opening angles of 3-7 degrees. The split-monopole initial setup usually results in less collimated outflows. The light surface of the outflow magnetosphere tends to align vertically - implying three relativistically distinct regimes in the flow - an inner sub-relativistic domain close to the jet axis, a (rather narrow) relativistic jet and a surrounding subrelativistic outflow launched from the outer disk surface - similar to the spine-sheath structure currently discussed for asymptotic jet propagation and stability. The outer subrelativistic disk-wind is a promising candidate for the X-ray absorption winds that are observed in many radio-quiet AGN. The hot winds under investigation acquire only low Lorentz factors due to the rather high plasma-$\\beta$ we have applied in order to provide an initial force-balance in the disk-corona. When we increase the outflow Poynting flux by injecting an additional disk toroidal field into the outflow, the jet velocities achieved are higher.  These flow gains super-magnetosonic speed and remains Poynting flux dominated. ",
        "introduction": "Astrophysical jets emanate from sources spanning a huge range in energy output or length scale - among them young stellar objects (YSO), stellar mass compact objects as X-ray binaries or $\\mu$-quasars, or the powerhouses of some active galactic nuclei (AGN) which host a super-massive black hole. In particular for radio-loud quasars, for which synchrotron emission dominates the radio spectrum, relativistic jets are a generic feature. Due to the omnipresent angular momentum conservation, mass accretion to all of these objects features a disk structure around the central mass. It is commonly believed that jets are launched as disk winds, which are further accelerated and collimated by magnetic forces (see \\citet{1982MNRAS.199..883B, 1983ApJ...274..677P, 1986A&A...162...32C, 1997SvPhU..40..659B, 2003ApJ...596.1240H, 2007prpl.conf..277P}. Relativistic jets may gain further energy by interaction with the black hole magnetosphere \\citep{1977MNRAS.179..433B, 1997MNRAS.292..887G, 2005MNRAS.359..801K}.  The magnetohydrodynamic (MHD) self-collimation of non-relativistic jets has been proven in general by time-dependent simulations \\citep{1995ApJ...439L..39U,1997ApJ...482..712O} and have been investigated in further detail considering additional physical effects as magnetic diffusivity by \\cite{2002A&A...395.1045F}, a variation in \\cite{1999MNRAS.309..233O}, non-axisymmetric instabilities in the launching region \\citep{2003ApJ...582..292O}, or a variation in the mass flow profile or the magnetic field geometries \\citep{2006ApJ...651..272F, 2006MNRAS.365.1131P}, or the influence of a central magnetic field \\citep{2009ApJ...692..346F,2008A&A...477..521M}.  In the case of relativistic jets the efficiency of MHD self-collimation is under debate. The main reason is the existence of electric fields which are negligible for non-relativistic MHD and which are commonly thought to have a net de-collimating effect on the jet. Essentially, \\cite{1991ApJ...377..462C} have demonstrated the current carrying relativistic jet can be highly collimated. However, the actual structure of these jets still remains unclear - mainly due to the need for simplifying assumptions to solve the corresponding set of MHD equations.  So far, a variety of theoretical models have been developed for the case of {\\em self-similar} jets \\citep{1992ApJ...394..459L, 1994ApJ...432..508C, 1995ApJ...446...67C, 2003ApJ...596.1104V,2006A&A...447..797M}, although it seems clear that relativity does not obey self-similarity. Fully 2.5D theoretical solutions for the internal magnetic jet structure could be obtained by neglecting matter inertia \\citep{1997A&A...319.1025F, 2001A&A...365..631F}. These force-free solutions for the field structure can in principle be coupled to the dynamical wind solution along the field lines \\citep{1996A&A...313..591F, 2001A&A...369..308F, 2004ApJ...608..378F}. \\cite{1997A&A...319.1025F} obtained solutions for the internal jet force balance in Kerr metric with an asymptotically cylindrical jet emerging from a disk-like structure around the central rotating black hole. The shape of the collimating jet boundary was obtained as result of the internal force equilibrium, in particular considering the regularity condition along the jet outer light surface.  Time-dependent simulations of relativistic MHD jet formation have been performed considering a general relativistic metric, including also the evolution of the underlying accretion disk. Early - seminal - simulations did last for a few inner disk rotations only \\citep{1998ApJ...495L..63K, 1999ApJ...522..727K}, which is sufficient time to demonstrate the launching of an outflow, but hardly sufficient in order to investigate the long-term dynamical evolution of the emerging jet.  More recent simulations were able to follow several 100 disk rotations and show the formation of a so-called funnel flow origination in the shear layer between the horizon and the inner disk radius \\citep{2005ApJ...620..878D, 2007MNRAS.375..513M, 2008MNRAS.388..551T, 2009MNRAS.394L.126M}. These simulations indicate a highly time-variable mass ejections of rather low degree of collimation. However, the funnel flow achieves Lorentz factors of up to 50. General relativistic MHD simulations are also able to determine the interrelation between jet formation and the Blandford-Znajek mechanism \\citep{2005ApJ...630L...5M, 2007MNRAS.377L..49K}.  Ultra-relativistic MHD simulations of accelerating and collimating jets have been presented by \\citet{2007MNRAS.380...51K}, spanning over a huge range of length scale and providing jets of large Lorentz factor $\\Gamma \\sim 10$. Their simulations, however, did not start from the very base of the jet - the accretion disk, but at some fiducial boundary above the equatorial plane. Since the jet has been launched already with super-escape speed, gravity has not been considered. The jet flow has been confined within a rigid wall of predefined shape which naturally affects the opening angle of the MHD jet nozzle and thus jet collimation and acceleration.  The focus of our present paper is i) to concentrate on the formation and acceleration of a relativistic MHD jet right from the launching area the accretion disk surface, ii) to investigate the (self-) collimation of relativistic MHD jets under the influence of de-collimating electric forces and an \"open\" boundary condition for the outflow iii) to consider gravity as en essential gradient to provide a realistic disk boundary condition in equilibrium, iv) to run long-term simulations lasting more than 1000 inner disk rotations until the jet reaches steady state, v) to concentrate on MHD disk jets as disks are the natural origin for the mass load for AGN jets.  The outline of this paper is as follows: In section \\ref{sec:concepts} we discuss the concepts of ideal special relativistic magnetohydrodynamics in the perspective of jet formation. Section \\ref{sec:model setup} is devoted to the initial- and boundary conditions of the numerical simulations, whose results are shown in section \\ref{sec:results}. We conclude in section \\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:results}  We now present the results of our numerical simulations considering the formation of relativistic MHD jets from accretion disks. Each simulation consumed approximately 48 hours on 16 processors. The overall goal is to test whether the paradigm of MHD self-collimation of non-relativistic jets established from numerical simulations \\cite{1995ApJ...439L..39U, 1997ApJ...482..712O, 1999ApJ...526..631K, 2002A&A...395.1045F} also holds in the relativistic case.   \\begin{deluxetable*}{ccccccl||crrrrr} \\tablecaption{Parameter summary of our disk-wind simulations. \\label{tab_all} } \\tabletypesize{\\scriptsize} \\tablewidth{\\textwidth} \\tablehead{ \\colhead{ID}                      & \\colhead{Top}               & \\colhead{$\\beta$}                     & \\colhead{$\\epsilon$}       & \\colhead{$v_{\\rm inj} $}        & \\colhead{$\\eta$}         & \\colhead{Remarks}      & \\colhead{$\\Gamma_{\\rm max}$} & \\colhead{$\\mu_{\\rm max}$} & \\colhead{$\\xi$} & \\colhead{$v_{p,\\rm max}$} & \\colhead{$r_{\\rm jet}$} & \\colhead{$\\dot{M}$} } \\startdata        WA01 & A & 0.2 & 2/3 & var. & var. & - & 1.23&         1.33 &        18.84 &       0.54 &        20.26 &        23.77 \\\\  WA02 & A & 1 & 2/3 & var. & var. & - & 1.27 &         1.37 &        11.47 &       0.58 &        21.86 &        26.62 \\\\ WA03 & A & 2 & 2/3 & var. & var. & - &  1.26&         1.34 &        10.69 &       0.57 & 22.21 &        24.15 \\\\ WB01 & B & 1 & 2/3 & var. & var. & $\\theta=77^\\circ$ & 1.33&         1.41 &        8.27 &       0.64 &        24.19 &        16.48 \\\\ WB02 & B & 1 & 2/3 & var. & var. & $\\theta=85^\\circ$ & 1.29&         1.33 &        8.19 &       0.60 & 21.27 &        15.66 \\\\ WA04 & A & 0.2 & 1/6 & var. & var. & - & 1.27&         1.38 &        13.40 &       0.58 &        21.14 &        36.08 \\\\ WA05 & A & 1 & 1/6 & var. & var. & - & 1.25&         1.33 &        9.84 &       0.57 &         22.77 &        32.67 \\\\ WA06 & A & 2 & 1/6 & var. & var. & - & 1.25&         1.32 &        9.34 &       0.57 & 22.92 &        29.27       \\enddata \\tablecomments{Columns are from left to right: simulation ID; initial magnetic field distribution (Top): potential field (A) or split monopole (B); plasma-$\\beta$; accretion disk temperature parameter $\\epsilon$; injection speed $v_{\\rm inj}$ as a fraction of $v_{\\phi}(r)$; toroidal disk-field scaling $\\eta$;  specific remarks; to the right we show values of the steady state, the maximum Lorentz factor $\\Gamma_{\\rm max}$ collimation degree $\\xi$, maximal poloidal velocity $v_{p,\\rm max}$, jet radius $r_{\\rm jet}$ and the total mass flux out of the domain $\\dot{M}$. } \\end{deluxetable*}    \\subsection{Overall evolution of the outflow}  The initial evolution of the disk corona is governed by the propagation of toroidal Alfv\\'en waves launched due to the rotation of the field line foot-points. The initial force-free magnetic field structure is adapted to a new dynamic equilibrium according to a rotating wind magnetosphere.  A wind is launched from the disk boundary and is continuously accelerated driving a shock front through the initial hydrostatic corona and sweeping this material out of the computational domain (Fig.\\,\\ref{fig:simul}). The disk wind evolves into a collimated outflow of super-magnetosonic speed. Along the symmetry axis the hydrostatic initial condition is very well preserved. Once the bow shock has passed through the domain, the jet mass flux declines to a value which is solely governed by the internal outflow dynamics and the injection boundary conditions. Similarly, the post-shock magnetic field distribution follows as well from the internal outflow dynamics and has in principle little in common with the initial setup. Certain combinations of boundary conditions for mass flux and magnetic field will result in a quasi-stationary\\footnote{We denote the dynamical state as {\\em quasi} stationary as due to Keplerian disk rotation the disk outflow a large radii has evolved for a considerably lower number of disk rotations. Therefore, a slight change in the dynamical state of the outer outflow can be expected after another few outer, thus 1000s of inner rotations.} state of the outflow evolution (see next section). From this point onwards we can start our investigations of collimation and acceleration. In this paper we concentrate on analysis when the flow has reached a quasi-steady state. We usually terminate our simulations after 500 inner disk-rotations $P$, while a quasi-steady state is established over most of the domain after about 200 rotations.  Figure \\ref{fig:simul} shows the time evolution for two exemplary simulations with an initial hourglass-shaped potential field distribution (case A) and a split-monopole field distribution (case B), each for the parameter choice $(\\beta,v_{K},\\epsilon) = (1,0.5,2/3)$.  \\begin{figure*}[htbp] \\centering \\includegraphics[width=0.7\\textwidth]{f4}      \\caption{ Formation of a relativistic MHD jet. Shown is the Lorentz factor (color gradient) in terms of $log_{10}(\\Gamma-1)$ at the time of 25, 250 and 500 (left to right) inner disk rotations. {\\em Top:} Field distribution case A, hourglass-shape potential field (run WA02). {\\em Bottom:} Split monopole initial field distribution case B (run WB01). Shown are poloidal magnetic field lines (solid white lines); the critical MHD surfaces (solid black lines); the light surface $r\\cdot \\Omega = c$ (solid black line). For time $T/P=250$ electric current flow lines are added (solid green line). Time step $T/P=500$ also show the initial field lines for comparison (dashed white lines). \\label{fig:simul}} \\end{figure*} The figure shows the Lorentz factor, the poloidal magnetic field lines, poloidal electric current flow lines and the critical MHD surfaces. In addition the light surface is drawn.  Phenomenologically, the solutions form a magnetic nozzle with, depending on the disk flux distribution, considerable difference in the width, but comparable final opening angles of the fast component.  A broader initial field distribution (case B) also results in broader and faster winds where the material originating from the inner disk is more effectively thinned out. In analogy to hydrodynamic nozzles, the flow reaches the slow-magnetosonic speed directly above the throat. Collimation happens mainly before the fast-magnetosonic surface is reached. Afterwards, the opening angle of a given field line is approximately conserved.  Of particular interest is the electric current distribution (shown for the time step $T/P = 250$). The electric current distribution is a consequence of the dynamical evolution of the outflow and therefore a direct outcome of the disk boundary magnetic flux profile and the Keplerian field line rotation.  In general, the electric current leaves the outer disk to return within the fast component of the outflow. It is expected to enter the inner disk and then flow radially outwards closing with the outgoing current. Such butterfly-shaped circuits are expected in Keplerian disks while the $j_{r}$ plays a leading role in the disk-jet feedback \\citep{1997A&A...319..340F}. A positive radial electric current in the disk-corona supports accretion by braking the disk material due to its magnetic torque $j_{r}\\times B_{z}$ similar to a Barlow-wheel\\footnote{ However, this region is not resolved within our numerical domain, as it is located below our injection boundary as part of the underlying non-ideal MHD accretion disk. See \\cite{2002ApJ...581..988C, 2007A&A...469..811Z} for non-relativistic simulations of the disk-jet interaction}.  The inclination between the poloidal current vector and the magnetic field line indicates the direction of (de-) collimating magnetic forces acting on the flow. When the inclination becomes less than $90^{\\circ}$, the Lorentz force $j_{p}\\times B_{\\phi}$ changes from collimation to de-collimation. This can be clearly seen in the snapshots at $T/P=250$ of the case A simulation where actual field lines are indicated in white and initial field lines in red. In the actual field distribution, the field lines are somewhat pushed away from the surface $\\mathbf{j_{p} \\perp B_{p}}$ (this is also where magnetic acceleration is most effective). For case B this happens beyond the light surface. As a result, both electric and magnetic forces deflect the flow towards the disk boundary which leads to a highly unstable layer just above the outer disk. We will provide an in-depth analysis including all forces acting on the flow in section \\ref{sec:trans-field}.  The locations of the characteristic surfaces are signatures for the MHD flow. Depending on the initial magnetic flux distribution (case A,B) and the mass flux profile (see also \\citet{2006ApJ...651..272F}), this location may vary a great deal. In our case B simulations, we generally observe surfaces which leave the domain in radial direction (parallel to the disk surface). For the case A simulations these surfaces tend to \"collimate\" leaving the domain in vertical direction. The latter implies a two-layered structure of the jet - a central super-fast magnetosonic jet surrounded by a sub-Alfv\\'enic outflow. This is an interesting aspect for observational modeling and for stability analysis of sheath-spine jets \\citep{2005MNRAS.356..859P, 2007ApJ...662..835M, 2007ApJ...668L..27K, 2007ApJ...664...26H, 2009MNRAS.397.1486B}. The broad wind launched from the outer regions of the disk has much lower velocities, decreasing continously with increasing launching-radius. For example, the terminal velocity of the flow originating from $r_{\\rm fp}>32$ of the case A simulations drops below $0.2c$, consistent with the X-ray absorption features observed in a mounting number of AGN \\citep{2006AN....327.1012C, 2009A&ARv..17...47T}. In principle, our dynamical models can provide basic ingredients (e.g. flow geometries and velocity gradients) for the modeling of spectral line profiles of disk winds \\citep{1995MNRAS.273..225K, 2008MNRAS.388..611S}.   Following  \\cite{2006ApJ...651..272F}, we may define an average collimation degree $\\xi$ of the outflow measured as the fraction of vertical and radial mass flux through equal-area surfaces at a certain height (here at $z=z_{m}$), \\begin{equation} \\xi = \\frac{ \\int_{0}^{r_{m}}r \\Gamma v_{z} \\rho |_{z_{m}}\\ dr }{ \\int_{z_{m}-r_{m}/2}^{z_{m}}r_{m}\\Gamma v_{r}\\rho |_{r_{m}}\\ dz }. \\end{equation} Similarly, we define a mass flux weighted jet radius, \\begin{equation} r_{jet} = \\frac{ \\int_{0}^{r_{m}}r \\Gamma v_{z} \\rho |_{z_{m}} \\ dr }{ \\int_{0}^{r_{m}} \\Gamma v_{z} \\rho |_{z_{m}}\\ dr }. \\end{equation} The corresponding values for $\\xi$ and $r_{\\rm jet}$ derived for $z_{m}=200$ at the upper end of the domain and for time $t/P=500$ are given in Tab.\\,\\ref{tab_all} along with the maximum Lorentz factor $\\Gamma_{\\rm max}$, the maximum poloidal velocity $v_{p,max}$, and the total mass flux $\\dot{M}$. Figure \\ref{fig:diag} shows the time evolution of these quantities in the top panel.  In general, we observe that the collimation degree $\\xi$ is the most sensitive tracer for secular trends among the observables mentioned. In the lower panel, we show the evolution of jet-power in the individual energy channels leaving the computational domain (radial and vertical). \\begin{figure}[htbp] \\centering \\includegraphics[width=\\columnwidth]{f5} \\caption{ \\red Time evolution of characteristic quantities.  \\textit{Top panel:} After an initial adjustment till $t/P\\simeq200$, mass flux $\\dot{M}$, jet radius $r_{jet}$, collimation degree $\\xi$, maximal Lorentz-factor $\\Gamma_{\\rm max}$  and poloidal velocity  $v_{p, \\rm max}$ cease to evolve. \\textit{Lower panel}: Power in the individual energy chanels out of  the domain. Thermal power ($T$) peaks when the shock reaches the upper boundary and is negligible otherwise. The total outgoing power (labeled accordingly) is dominated by rest-mass ($M$) and Poynting flux ($S$). Also shown are the gravitational ($G$) and (purely) kinetic ($K$) contributions (simulation run WA02). \\black \\label{fig:diag}} \\end{figure} After the re-configuration of the initial stationary-state to the dynamical solution, the partitioning of energies is completed at around $t/P=100$.  Thermal energy-flux peaks when the hot bow-shock passes through the upper boundary.  Far away from the central object, gravitational and thermal energy-flux are negligible % The integrated energy-flux is dominated by rest-mass, reflecting the fact that only the inner component reaches significant Lorentz-factors. The balance between Poynting and kinetic flux is of particular interest. Figure \\ref{fig:diag} shows merely the end result of the spatial conversion history with the remaining electromagnetic energy $\\mathcal{S}$ above the purely kinetic part $\\mathcal{K}$.  % More detailed insight into how this is established is provided in the following section using an individual field line. \\subsection{Stationary state analysis} Simulations starting from an initial field distribution A evolve into a quasi-stationary flow solution after about 200 inner disk rotations. Figure \\ref{fig:stationary_flow} shows our reference simulation WA04 at time $t/P=250$, including an enlarged subgrid of the innermost area of the domain.  Steady state solutions are helpful to understand the flow structure for a number of reasons. Firstly, by using MHD conservation laws, the conserved quantities (see Sect.\\,\\ref{sec:flconst}) allow to identify the momentum and {\\em energy channels} of the flow during acceleration and collimation. Secondly, by using the force-balance equations \\ref{eq:trans-field},\\ref{eq:parallel-field} we may identify the leading forces on the material along the outflow. Thirdly, the cross-check for conserved quantities provides another test for the quality of our setup and the numerical approach. A secondary indicator of stationarity is the alignment of poloidal velocities with the poloidal magnetic field lines, $E_{\\phi} = 0$. Figure \\ref{fig:stationary_flow} shows corresponding velocity vectors confirming this picture.  This is confirmed by checking in detail the complete set of integrals of motion of the MHD-flow $k,\\Omega^F,Q,l,\\mu$ as defined in equations $\\ref{eq:kf}-\\ref{eq:mu}$. Figure \\ref{fig:const} shows the relative deviation of these quantities from their average value along a given field line after $t/P=500$. The integrals are conserved within $1\\%$-accuracy already right above the injection boundary - clearly demonstrating the quality the choice of our numerical setup, in particular the injection boundary conditions carefully constructed from an equilibrium of Keplerian rotation and gas pressure.  Due to the differential rotation law, the number of Keplerian rotations $t/P(r)$ scales with radius as $t/P(r) = (t/P) r^{-3/2}$, implying that at the end of our simulations ($t/P=500$), we have performed roughly one rotation at $r=64$ and half a rotation at $r=100$. Nonetheless the integrals of motion for the field line $r_{\\rm fp}=64$ are conserved within $0.1\\%$.  \\begin{figure}[htbp] \\centering \\includegraphics[width=\\columnwidth]{f6} \\caption{Field line constants (conserved quantities). Shown is the deviation from the average value along the field line rooted at $r_{\\rm fp}=2$ (simulation run WA02). \\label{fig:const}} \\end{figure}   \\begin{figure}[htbp] \\centering \\includegraphics[width=8cm]{f7a} \\includegraphics[width=8cm]{f7b} \\caption{Logarithmic (rest-frame) density of the stationary flow (simulation run WA04). Shown are poloidal magnetic field lines (solid white), electric current flow lines (solid black), characteristic MHD surfaces (various dot-dashed green), surface of escape velocity (dotted green), light surface (solid green). Arrows in the top plot indicate the velocity field. The bottom figure is an enlarged picture of the central region indicating the three regimes defined by the light surface. \\label{fig:stationary_flow}} \\end{figure}  \\subsubsection{Collimating and accelerating forces}\\label{sec:trans-field} In this section we identify the forces responsible for jet acceleration and collimation applying the steady-state parallel and transversal force-equilibrium Eqs.\\,\\ref{eq:trans-field},\\ref{eq:parallel-field}.  Fig.\\,\\ref{fig:f-hot} compares these forces for a number of reference simulations (WB01, WA02, WA05) along a field line rooted at $r_{\\rm fp}=2$. As check for consistency, we also show the gradient of the Mach number $a \\equiv B_{p}^2/4\\pi\\nabla_{||}M^2$ which just coincides with the summation of the parallel forces, indicating a steady state (see yellow solid and black dashed line). \\begin{figure*}[htbp] \\centering \\includegraphics[width=0.49\\textwidth]{f9a} \\includegraphics[width=0.49\\textwidth]{f9b} \\includegraphics[width=0.49\\textwidth]{f9c} \\includegraphics[width=0.49\\textwidth]{f9d} \\includegraphics[width=0.49\\textwidth]{f9e} \\includegraphics[width=0.49\\textwidth]{f9f} \\caption{Accelerating (left column) and collimating (right column) forces along the field line rooted at $r_{\\rm fp}=2$ for the three reference simulations (from top to bottom): {\\it WB01} (split monopole), {\\it WA02} (hourglass potential field, hot case), and {\\it WA05} (hourglass potential field, cool case). Shown are the contributions from gravity, gas pressure gradient, centrifugal force, poloidal magnetic field pressure gradient, poloidal field tension and pressure gradient, toroidal magnetic field pressure gradient and tension, and forces due to the electric field. Vertical lines indicate the critical surfaces - the Alfv\\'en point along this field line, denoted by 'A', the light cylinder radius denoted by 'lc', and the fast magnetosonic point denoted by 'F'. In this logarithmic representation a change of sign in the force direction is indicated by the singularities along the graphs. \\label{fig:f-hot}} \\end{figure*}  In general, the outflow starts with sonic speed and is first launched by thermal pressure in the hot disk corona, respectively the centrifugal force in the colder version. Until the Alfv\\'en point, the Lorentz-force of the poloidal electric current ($F_{\\rm pbphi}+F_{\\rm pinch}$) is the main magnetic driver. Ultimately the poloidal tension ($F_{\\rm curv}$) keeps the acceleration up even above the fast surface.  Concerning the transverse force, we reproduce the expected sign-change of the curvature (tension) force (first collimating until the Alfv\\'en surface, de-collimating beyond) and the poloidal pressure force (de-collimating until the light cylinder, collimating beyond). For the cross-field balance, we observe the following three regimes:  Just on top of the inlet, the main de-collimating forces besides poloidal magnetic pressure are thermal pressure in the hot case (WB02, WA02) and centrifugal support in the colder case (WA05).  Gravity is here the strongest force towards the origin and the situation just reflects the radial force-equilibrium we have applied for the inlet boundary.  This is the hydrodynamic regime.  At the Alfv\\'en point, the residual of the pinch- and toroidal pressure-force ($\\mathbf{j_{p}}\\times B_{\\phi}$) is the main collimator, balanced by the centrifugal term.  Thermal pressure quickly looses importance. This is the magneto-hydrodynamic regime.  In the asymptotic region beyond the light-cylinder, de-collimation by electric forces overcomes the centrifugal force and is balanced by the poloidal magnetic pressure that changes its sign at the light-cylinder (best seen in WB01). This is the relativistic regime.  To get a global impression on the relative importance of the individual forces we show a radial cut throughout the asymptotic jet in figure \\ref{fig:force-balance}. The strongest forces arise across the inner asymptotic light surface which separates field lines of high angular velocity from those in the non-rotating corona along the axis. Here the electric de-collimation is essential. The $B_{\\phi}(r)$ profile is curled up from the inner disk radius along the outflow - resulting in a magnetic pressure gradient that works in unison with the toroidal field pinch force until at some radius the toroidal field surpasses its maximum and decreases (negative gradient).  The strong gradients in toroidal field and rotation induce a current sheet and give rise to an electric charge. The space charge $\\rho_{\\rm e}=(1/4)\\pi\\nabla\\cdot \\mathbf{E}$ is positive close to the axis and changes its sign at a critical line as defined by \\cite{1969ApJ...157..869G}.   \\begin{figure}[htbp] \\centering \\includegraphics[width=\\columnwidth]{f10} \\caption{ Trans-field force cut at $z=200$, inner ($lc_{i}$) and outer ($lc_{o}$) light-cylinder.  The differentially rotating field lines are fastest at the inner disk-radius, resulting in the inner light-cylinder - here electric de-collimation is important.  The $B_{\\phi}(r)$ profile is curled up from the inner disk radius onwards and results in a magnetic pressure gradient that works in unison with the pinch force until the toroidal field surpasses its maximum. Within $r<3$, we omit the curves for thermal and poloidal pressure. These terms fluctuate around $\\pm10^{-5}$ while balancing each other. (run WA02) \\label{fig:force-balance}} \\end{figure}   \\subsubsection{Energy conversion} In the simulations where injection is sub-magnetoslow, the energy flux is not a free parameter, but is consistently determined by the simulation of the disk-wind. It is hence of interest how the partitioning and conversion is realized. From the values of $\\mu_{\\rm max}$ given in Tab.\\,\\ref{tab_all} it is obvious that our disk-corona supports only mildly relativistic flows below $\\Gamma=1.5$ (section \\ref{sec:flconst}).  In Fig.\\,\\ref{fig:vp} (bottom left panel) we show the efficiency $\\sigma$ of Poynting flux to kinetic flux conversion along the field line with $r_{\\rm fp}=2$ in the fast component of the jet. \\begin{figure*} \\centering \\includegraphics[width=0.49\\textwidth]{f8a} \\includegraphics[width=0.49\\textwidth]{f8b} \\caption{ Dynamical quantities as a function of the distance along the field line rooted at $r_{\\rm fp}=2$ for the model WA02. Vertical lines indicate the slow-magnetosonic, Alfv\\'en, light surface and fast magnetosonic transitions (from left to right). \\textit{Left:} Isorotation parameter $\\Omega^F$, poloidal and toroidal velocity are shown in the top panel. Lorentz factor $\\Gamma$, energy conversion efficiency $\\sigma$, and normalized total energy flux $\\mu$ in the bottom panel.  The flow is below equipartition already at the base of the jet. \\textit{Right:} Complete energy flux ratios.  In the asymptotic region, thermal and gravitational fluxes are obviously negligible. \\label{fig:vp}} \\end{figure*} Here, $\\sigma$ is below equipartition already at the inlet and it further decreases as $\\Gamma$ approaches $\\mu$. In Fig.~\\ref{fig:vp} (top left) the toroidal velocity shows that the flow decouples from co-rotation with the magnetic field at the Alfv\\'en point. Beyond the Alfv\\'en point, angular momentum is then carried predominantly by the magnetic field. The poloidal velocity increases from low injection value (sonic velocity) to $\\sim0.5c$. Further acceleration can not be expected as the bulk of the energy is already in kinetic form. The right panel of Figure \\ref{fig:vp} shows the individual energy channels compared to the rest-mass flux for the same field line. At the base of the jet, the strong poloidal electric currents (a strong toroidal field) give rise to an outflow with $\\mathcal{K}<\\mathcal{T}<\\mathcal{-G}<\\mathcal{S}<\\mathcal{M}$, predominantly transporting energy via rest-mass and Poynting-flux.  The kinetic energy flux surpasses the thermal flux at the Alfv\\'en point and further overcomes the gravitational binding energy term shortly thereafter.  This is not surprising, since the escape-surface can be close to the Alfv\\'en surface at least for the inner field lines (see also Fig. \\ref{fig:stationary_flow}).  Only then, the cold limit $\\mu = \\Gamma(\\sigma+1)$ is applicable - it is certainly valid in the asymptotical outflow where thermal and gravitational energy fluxes are negligible.   \\subsection{Dependence on the launching environment} For the simulations described up to now, we have performed in addition several parameter runs in order to investigate how the resulting jet dynamics depends on the (prescribed) launching conditions - the disk corona (see Tab.~\\ref{tab_all}). We now focus on the impact of the plasma $\\beta$ and the disk temperature parameter $\\epsilon$.  In general, a low $\\beta$ (a stronger magnetic field) we find that the outflow tends to collimate more, as indicated by the higher average collimation degree $\\xi$ and a lower momentum weighted jet radius $r_{\\rm jet}$. This in principle decreases the MHD acceleration efficiency which critically depends on the divergence of flux surfaces. It is straightforward to define the mass flux-weighted (half-) opening angle of outflow, \\eqon \\theta_{\\dot{M}} = \\rm atan\\ \\xi^{-1}, \\eqoff which translates to angles of $3^{\\circ}<\\theta_{\\dot{M}}<7^\\circ$ for the outflows under consideration.  The impact of the magnetic field strength on the amount of mass flux is not clearly visible, as the two simulations with $\\epsilon=1/6, 2/3$ show a different trend. As the $\\epsilon$-parameter is simply a proxy for the disk corona density, it will affect the collimation in the following manner: A higher inflow density lowers the Alfv\\'en surface towards the disk surface which in turn broadens the current topology and therefore widens the flow.  This is also the trend that we observe in the indicators $\\xi$ and $r_{\\rm jet}$. Given that the injection speed calculated iteratively from the outflow simulation approaches the slow-magnetosonic speed, we expect the mass flux to scale as $\\dot{M}\\propto \\sqrt{p \\rho}$. In fact, this is approximately realized since we have $\\dot{M}(\\epsilon=1/6)/\\dot{M}(\\epsilon=2/3) =\\sqrt{2.5}\\simeq 1.6$.  The change of the initial split-monopole inclination $\\theta$ has little effect on the overall jet collimation angle. In particular, we observe an opposite trend as the wider initial field with $\\theta=77^\\circ$ ends up slightly more collimated than the one with $\\theta=85^\\circ$. Clearly a wider initial field leads to a larger jet radius $r_{\\rm jet}$. Here, we like to stress the point that for the final steady state solutions in our simulations the initial field structure is important only insofar as it also prescribes the poloidal magnetic field profile along the outflow launching boundary. The field structure is completely changed from the initial steady structure to a new dynamic equilibrium. Thus it makes no sense to compare the collimation of the initial field with the collimation of the outflow field distribution.   \\par Having pointed out the crucial role of the \\red vertical energy flux from the disk surface \\black $\\mu$ and the closely related quantity $\\sigma = \\mathcal{S/(K+M)}$, we now study the two most promising handles in increasing $\\mu$. That is (i) a decrease in mass flux $\\mathcal{M}$ and (ii) an increase in Poynting-flux $\\mathcal{S}$. We first focus on (i) and describe (ii) thereafter.  \\subsubsection{Towards low mass loading}\\label{sec:light-winds} \\red A way to obtain high-speed jets seems to be a lower mass load injected into a similarly-strong magnetic flux. According to the well-known Michel-scaling \\citep{1969ApJ...158..727M}, the asymptotic outflow velocity depends on the mass flux $u_{\\infty}\\propto \\dot{M}^{-1/3}$. We investigate this interrelation running another set of simulations where we change the mass flux by prescribing a low injection velocity (instead of lowering the density of the injected gas, see Tab.~\\ref{tab:extended1})  \\begin{deluxetable*}{ccccccl||crrrrr} \\tablecaption{Effect of mass-loading \\label{tab:extended1}} \\tabletypesize{\\scriptsize} \\tablewidth{\\textwidth} \\tablehead{ \\colhead{ID}                      & \\colhead{Top}               & \\colhead{$\\beta$}                     & \\colhead{$\\epsilon$}       & \\colhead{$v_{\\rm inj} $}        & \\colhead{$\\eta$}         & \\colhead{Remarks}      & \\colhead{$\\Gamma_{\\rm max}$} & \\colhead{$\\mu_{\\rm max}$} & \\colhead{$\\xi$} & \\colhead{$v_{p,\\rm max}$} & \\colhead{$r_{\\rm jet}$} & \\colhead{$\\dot{M}$} } \\startdata I001 & A & 0.2  & 2/3  &  0.01  &  var.  & - & 1.66&         2.01 &        25.61 &       0.75 & 21.11 &        10.63 \\\\  I01 & A & 0.2  & 2/3  &  0.1  &  var.  & - & 1.52&         1.76 &        25.58 &       0.71 & 20.49 &        13.78 \\\\  I02 & A & 0.2  & 2/3  &  0.2  &  var.  & - & 1.38&         1.55 &        26.02 &       0.65 & 19.34 &        17.53 \\\\  I04 & A & 0.2  & 2/3  &  0.4  &  var.  & - & 1.31&         1.45 &        25.30 &       0.60 &        17.07 &        25.84 \\\\ I08 & A & 0.2  & 2/3  &  0.8  &  var.  & - & 1.16&         1.24 &        40.05 &       0.47 &        14.74 &        51.30 \\\\ I12 & A & 0.2  & 2/3  &  1.2  &  var.  & - & 1.21&         1.21 &        33.77 &       0.50 &        14.10 &        81.88 \\enddata \\tablecomments{As in table \\ref{tab_all}, but for the extended parameter study with a given $v_{\\rm inj}$.} \\end{deluxetable*}   However, we observe that for low injection speeds $v_{z}<0.4\\ v_{\\phi}$ the resulting (numerical) mass flux does not follow the expected linear relation $\\dot{M} \\propto v_{\\rm inj}$. Instead, for lower and lower injection speed it approaches some seemingly unphysical offset value (Fig.~\\ref{fig:diag-comp}). This value is unphysical in the sense that it departs from the prescribed value at the boundary condition $d\\dot{M}_{\\rm inj}/dr = 2\\pi r \\rho_{\\rm inj} v_{\\rm inj}$. This difficulty arises because the injected flow is sub-magnetosonic and thus overdetermined by simultaneously assigning a profile in $\\rho$ and $v_{z}$.   Although this is in principle problematic, it is not necessarely fatal for the investigation of low mass flux outflows. The mass flux is governed by the magneto-slow point and is thus re-arranged along the flow. Indeed we find that above the magneto-slow surface the field line constants are very well conserved. This indicates that the flow dynamics transcends at the magneto-slow surface from a seemingly unphysical state into a MHD flow that satisfies the critical conditions at the magnetosonic surfaces.  This technique of self-adjusting the sub-slow mass inflow, however, turned out to be limited towards lower mass fluxes. The resulting mass fluxes are unfortunately, still too high and do not allow substantially higher magnetisation (e.g. $\\mu_{\\rm max}=2.01$ for simulation I001 compared to $\\mu_{\\rm max}=1.33$ in run WA01). Figure \\ref{fig:diag-comp} shows the trends concerning collimation $\\xi$, jet radius $r_{\\rm jet}$, and integral mass flux $\\dot{M}$ in the asymptotic flow. \\begin{figure}[htbp] \\centering \\includegraphics[width=\\columnwidth]{f11} \\caption{Jet collimation and dynamics against injection speed parameter $v_{\\rm inj}$, The horizontal line indicates the mass flux when $v_{\\rm inj}$ is not specified (WA01). For $v_{\\rm inj}<0.4$, the mass flux approaches an obviously unphysical offset value. At higher $v_{\\rm inj}$, the mass flux follows the expected linear dependence on the injection parameter (thin solid line). \\label{fig:diag-comp}} \\end{figure} Increasing the mass flux enhances collimation while $\\mu_{\\rm max}$ decreases accordingly. For high injection speed, we observe that the wind originating from the very inner disk evolves into a thin ballistic flow layer that has little in common with the jets we are interested in. \\black \\subsubsection{Poynting dominated flows}\\label{sec:extended} Given the limitations mentioned above, we prescribe {\\it a priori} the limiting energy flux parameter $\\mu$ by constraining both the mass flux and the Poynting flux along the injection boundary - with the hope of thus providing a sufficiently energetic disk wind.  To achieve this, we adopted a fixed-in-time toroidal magnetic field distribution, $B_{\\phi}\\propto -\\eta/r$ which necessarily changes $\\Omega^F(r)$\\footnote{ In case of a central black hole causality requires that $r\\Omega^F(r)< 0.6$ which corresponds to the ISCO velocity for the Schwarzschild case with $r_{ISCO}=1$ in our scaling.}. The toroidal field distribution following a $1/r$ profile corresponds to $j_{z}=0$ and has a profound physical motivation as it anticipates the radial currents expected in the disk-corona.  Table \\ref{tab:extended2} summarizes the simulation runs performed within this setup. These simulations have a considerably stronger toroidal magnetic field at the injection point. We apply up to $B_{\\phi}\\sim 4 B_{p}$. \\begin{deluxetable*}{ccccccl||crrrrr} \\tablecaption{Poynting dominated flows \\label{tab:extended2}} \\tabletypesize{\\scriptsize} \\tablewidth{\\textwidth} \\tablehead{ \\colhead{ID}                      & \\colhead{Top}               & \\colhead{$\\beta$}                     & \\colhead{$\\epsilon$}       & \\colhead{$v_{\\rm inj} $}        & \\colhead{$\\eta$}         & \\colhead{Remarks}      & \\colhead{$\\Gamma_{\\rm max}$} & \\colhead{$\\mu_{\\rm max}$} & \\colhead{$\\xi$} & \\colhead{$v_{p,\\rm max}$} & \\colhead{$r_{\\rm jet}$} & \\colhead{$\\dot{M}$} } \\startdata  M02 & A & 0.2  & 2/3  &  0.1  &  2  & - & 2.18&         3.29 &        27.55 &       0.86 &        23.71 &        22.48 \\\\  M04 & A & 0.2  & 2/3  &  0.1  &  4  & - & 3.51&         7.46 &        8.96 &       0.94 &        29.00 &        48.85 \\\\  M08 & A & 0.2  & 2/3  &  0.1  &  8  & - & 6.11&         25.61 &        6.8377 &       0.97 &        35.89 &        80.40 \\enddata \\tablecomments{As in table \\ref{tab_all}, but for the extended parameter study with a given $v_{\\rm inj}$ and $\\eta$.} \\end{deluxetable*}  An exemplary process enhancing large-scale toroidal fields in the disk corona could be the MRI-driven dynamo under current investigation by many authors \\citep[e.g.][]{2000ApJ...534..398M, 2003A&A...398..825V}. The jet eventually evolving from these disks is not propelled by the \\citet{1982MNRAS.199..883B} mechanism, but driven by the toroidal magnetic pressure \\citep{1996LNP...471...74C, 2003ApJ...599.1238D, 2004ApJ...605..307K}. These so called Tower jets have initially been proposed by \\cite{1996MNRAS.279..389L} and directly extract Poynting-flux from the dist rather than first converting rotational energy into the twisted magnetosphere that is present at the Alfv\\'en surface.\\footnote{See \\cite{2007Ap&SS.307...11K} for a review. Note also that these jets have successfully been reproduced by \\cite{2005MNRAS.361...97L} in laboratory experiments with purely radial current distributions at the base.}  The simulations described here are of a mixed type, since they combine large-scale open field lines with a toroidal field emerging from a radial current. For an example simulation of this type, we show the conversion of energy in Fig.~ \\ref{fig:vp-ext2} similar to Fig.~\\ref{fig:vp} for the low-energy case. By design, the injected Poynting flux surpasses the rest mass flux with $\\sigma\\simeq5$ within the fast outflow component.  \\begin{figure}[htbp] \\centering \\includegraphics[width=\\columnwidth]{f12} \\caption{As in Fig.~\\ref{fig:vp}, however calculated for simulation run M04.  The injected flow is strongly magnetized and remains Poynting-dominated when leaving the computational domain. \\label{fig:vp-ext2}} \\end{figure}  The outflow, which is initially Poynting flux dominated, does not reach equipartition at the fast magnetosonic surface $r=r_{\\rm F}$, where we merely find $\\Gamma\\sim\\mu^{1/3}$ following \\cite{1969ApJ...158..727M,1998MNRAS.299..341B}. In the asymptotic region $r\\gg r_{F}$, the length scales for additional flow acceleration and collimation would increase exponentially with $\\Gamma\\propto (\\mu \\ln r)^{1/3}$ (e.g. \\cite{1994PASJ...46..123T}), which is clearly beyond the reach of our numerical method.   Our simulations indicate that, given sufficient Poynting flux, bulk Lorentz factors derived from AGN jet observations can be obtained within several hundred Schwarzschild radii, $\\Gamma \\simeq 6$ for model M08. Since acceleration has proven to be most effective around the Alfv\\'en point which is expected to be very close to the central object ($r_{\\rm A}<r_{\\rm lc}$), we expect this conclusion to remain valid also for higher $\\mu$ and thus higher terminal $\\Gamma$ flows. However, we note that when increasing $\\sigma$, the energy conversion efficiency decreases, a situation commonly denoted as $\\sigma$-problem in pulsar winds \\citep{1974MNRAS.167....1R, 1984ApJ...283..694K}. Several authors have recently addressed this issue with partly controversial results \\citep{2009MNRAS.394.1182K, 2009ApJ...699.1789T, 2009ApJ...698.1570L}, so that after the successful acceleration towards relativistic speeds, Poynting-flux could still remain and one is tempted to ask: Is there a $\\sigma$-problem for AGN-jets? The simulations presented here are not fit to answer this question satisfactory. However, we certainly know that the mildly relativistic disk winds presented earlier do not suffer from this, as they are launched already in sub-equipartition."
    },
    "0911/0911.0002_arXiv.txt": {
        "abstract": "We present the full source catalogue from the Australia Telescope 20~GHz (AT20G) Survey. The AT20G is a blind radio survey carried out at 20~GHz with the Australia Telescope Compact Array (ATCA) from 2004 to 2008, and covers the whole sky south of declination 0\\degr. The AT20G source catalogue presented here is an order of magnitude larger than any previous catalogue of high-frequency radio sources, and includes 5890 sources above a 20~GHz flux-density limit of 40\\,mJy.  All AT20G sources have total intensity and polarisation measured at 20~GHz, and most sources south of declination $-15\\degr$ also have near-simultaneous flux-density measurements at 5 and 8~GHz. A total of 1559 sources were detected in polarised total intensity at one or more of the three frequencies.  The completeness of the AT20G source catalogue is 91 per cent above 100~mJy\\,beam$^{-1}$ and 79 per cent above 50~mJy\\,beam$^{-1}$ in regions south of declination $-15\\degr$. North of $-15\\degr$, some observations of sources between $14-20$~hr in right ascension were lost due to bad weather and could not be repeated, so the catalogue completeness is lower in this region. Each detected source was visually inspected as part of our quality control process, and so the reliability of the final catalogue is essentially 100 per cent.  We detect a small but significant population of non-thermal sources that are either undetected or have only weak detections in low-frequency catalogues. We introduce the term Ultra-Inverted Spectrum (UIS) to describe these radio sources, which have a spectral index $\\alpha(5,20)>+0.7$ and which constitute roughly 1.2 per cent of the AT20G sample.  The 20~GHz flux densities measured for the strongest AT20G sources are in excellent agreement with the WMAP 5-year source catalogue of \\citet{wright09}, and we find that the WMAP source catalogue is close to complete for sources stronger than 1.5~Jy at 23~GHz. ",
        "introduction": "Large-area high-frequency radio surveys are time-consuming, and as a result relatively few large scale surveys have been carried out. Those that have been completed are either deep but covering small areas, or shallow all-sky surveys. As a result, our knowledge of the high frequency ($>10$~GHz) radio source population is poor.  In addition to the scientific benefits of studying the radio source population at high frequencies, large scale surveys are useful for foreground subtraction for Cosmic Microwave Background (CMB) experiments. Measurement of CMB anisotropies is limited by contamination from astronomical foregrounds, both Galactic and extragalactic. To improve the efficiency of the component separation techniques, observations such as those of ESA's Planck mission are performed on a broad spectral region ranging from a few tenths to hundreds of GHz. At frequencies up to $\\sim100$~GHz, extragalactic radio sources are the major contaminants on angular scales smaller than 30 arcminutes \\citep{dezotti05}, so identification of radio sources at high frequencies is critical.  The first blind radio survey above 8~GHz was carried out by \\citet{taylor01} with the Ryle telescope at 15.2~GHz. The survey covered a 63 deg$^2$ region and detected 66 sources to a limiting flux density of 20 mJy. \\citet{waldram03} extended the survey to 520 deg$^2$, detecting 465 sources to a flux density limit of 25\\,mJy (the 9C survey). Both surveys found that the existence and flux density of sources at 15~GHz cannot be accurately predicted by extrapolation from lower frequency radio surveys such as the NRAO VLA Sky Survey (NVSS) at 1.4~GHz \\citep{condon98}, further demonstrating the need for large scale high frequency surveys for population characterisation and source subtraction in CMB studies.  The Wilkinson Microwave Anisotropy Probe (WMAP) survey, which covers the whole sky at 23, 33, 41, 61 and 94~GHz \\citep{bennett03}, is the first all-sky radio survey above 5~GHz. The WMAP Point Source Catalogue constructed from the 5~yr maps contains 390 sources \\citep{wright09}, compared to 323 and 208 sources found in the previous 3~yr and 1~yr maps respectively \\citep{bennett03,hinshaw07}. Recently, \\citet{massardi09} have detected 516 sources in the 5~yr WMAP maps by exploiting a combination of blind and non-blind detection approaches. Section~\\ref{s_wmap} presents a detailed comparison of our AT20G results with the 5~yr WMAP point-source catalogue.  Table \\ref{t_surveys} summarises some earlier large-area high-frequency radio surveys, showing the context in which the AT20G survey was designed. A pilot survey for the AT20G at 18.5~GHz was carried out in 2002 and 2003 with the Australia Telescope Compact Array (ATCA) \\citep{ricci04,sadler06} and detected 173 sources stronger than 100\\,mJy in the declination range $-60$\\degr to $-70$\\degr. \\citet{ricci04} confirmed that ATCA (with custom hardware) had the capability to rapidly survey the sky at high frequencies. The differential source counts for extragalactic sources from the pilot survey  were found to be in good agreement with the \\citet{waldram03} 15~GHz survey.  \\begin{table*} \\centering \\caption{Comparison of the AT20G with other high-frequency radio surveys. Note that while the WMAP images cover the whole sky, regions at low Galactic latitude ($|b|<5\\degr$) are excluded from the point-source search. A reanalysis of the WMAP survey by \\citet{massardi09} resulted in a catalogue of 516 sources above a limit of 695 mJy.\\label{t_surveys}} \\begin{tabular}{llrrrll}\\hline Survey & Frequency & Sky area & Flux limit & N sources & Reference/s \\\\ & (GHz)     & (deg$^2$) & (mJy)  & & \\\\ \\hline Ryle & 15                 & 60      & 20   & 66 & \\citep{taylor01} \\\\ 9C   & 15, 43             & 520     & 25   & 465 & \\citep{waldram03} \\\\ 9C   & 15                 & 115     & 10   &     & \\citep{waldram09} \\\\ 9C   & 15                 & 29      & 5.5   &     & \\citep{waldram09} \\\\ WMAP & 23, 33, 41, 61, 94 & 32,177  & 1000 & 390 & \\citep{bennett03,hinshaw07,wright09} \\\\ AT20G pilot & 18          & 1216    & 100  & 126 & \\citep{ricci04} \\\\ AT20G & 5, 8, 20          & 20,086  & 40   & 5867 & \\citep[This work;][]{massardi08} \\\\ \\hline \\end{tabular} \\end{table*}  In this paper we present the full catalogue from the Australia Telescope 20~GHz Survey (AT20G). The survey covers 20,086 square degrees (the complete Southern sky to declination $0\\degr$) to a limiting flux density of 40~mJy\\,beam$^{-1}$. We followed up candidate sources detected in the survey at 20~GHz and also have near-simultaneous follow-up observations at 5 and 8~GHz for AT20G sources south of declination $-15^\\circ$. An accompanying paper (Massardi et al., in preparation) will provide more detailed statistical analysis of the AT20G sample. A subset of the 320 brightest ($S_{20} > 0.5$~Jy) extragalactic ($|b| > 1.5$\\degr) AT20G sources were presented and discussed by \\cite{massardi08}. The Galactic plane was included in our scanning survey but no follow-up observations were carried out at $|b|<1.5\\degr$, except for a blind survey of optically thick compact HII regions \\citep{murphy09}.  In Section~\\ref{s_obs} we describe the survey and follow-up observations, and in Section~\\ref{s_data} we describe the data reduction process. In Section~\\ref{s_errors} we calculate the accuracy of our measured positions and flux densities. Section~\\ref{s_cat} presents the source catalogue and defines its format, with the completeness and reliability of the catalogue discussed in Section~\\ref{s_comp}. As an additional investigation into the completeness, Section~\\ref{s_wmap} compares our catalogue with the 5~yr WMAP results. Finally, Section~\\ref{s_prop} discusses some statistical properties of the sample and Section~\\ref{s_conc} presents our conclusions. ",
        "conclusions": "\\label{s_conc} We present a catalogue of 5890 sources from the Australia Telescope 20~GHz survey, the deepest large-area survey at high radio frequency. For 3766 of these sources we have near-simultaneous 5 and 8~GHz measurements, and 1559 sources have a detection in polarized total intensity at one or more of the three frequencies.  The 20~GHz flux densities measured for the strongest AT20G sources are in excellent agreement with the WMAP 5-year source catalogue recently published by \\citet{wright09}, and we find that the WMAP source catalogue is close to complete (and highly reliable) for sources stronger than 1.5~Jy at 23~GHz.  We identify a population of Ultra-Inverted Spectrum radio sources with a spectral index of $\\alpha(5,20)>+0.7$. These are rare sources, comprising roughly 1.2 per cent of the AT20G population.  There are several ongoing projects as part of the AT20G. Massardi et al.~(in preparation) will present a statistical analysis of the AT20G sources, and Hancock et al.~(in preparation) will present results from the original scanning survey. Mahony et al.~(in preparation) is carrying out an analysis of the optical identifications of AT20G sources. \\citet{sadler08} are following up subsamples of the AT20G sources at 95~GHz. \\citet{murphy09} is targeting a subset of Galactic sources that were excluded from the main follow-up survey. Chhetri et al.~(in preparation) has carried out a search for gravitational lens candidates and planetary nebulae using data from the 6~km baseline. In a related project, Sadler et al.~(in preparation) is conducting a deeper survey at 20~GHz to explore the high frequency radio population at much lower flux densities."
    },
    "0911/0911.2007_arXiv.txt": {
        "abstract": "We report on a broader evaluation of statistical bootstrap resampling methods as a tool for pixel-level calibration and imaging fidelity assessment in radio interferometry. Pixel-level imaging fidelity assessment is a challenging problem, important for the value it holds in robust scientific interpretation of interferometric images, enhancement of automated pipeline reduction systems needed to broaden the user community for these instruments, and understanding leading-edge direction-dependent calibration and imaging challenges for future telescopes such as the Square Kilometer Array. This new computational approach is now possible because of advances in statistical resampling for data with long-range dependence and the available performance of contemporary high-performance computing resources. We expand our earlier numerical evaluation to span a broader domain subset in simulated image fidelity and source brightness distribution morphologies. As before, we evaluate the statistical performance of the bootstrap resampling methods against direct Monte Carlo simulation. We find both model-based and subsample bootstrap methods to continue to show significant promise for the challenging problem of interferometric imaging fidelity assessment, when evaluated over the broader domain subset. We report on their measured statistical performance and guidelines for their use and application in practice. We also examine the performance of the underlying polarization self-calibration algorithm used in this study over a range of parallactic angle coverage. ",
        "introduction": "Radio-interferometric image formation requires a solution for both the source brightness distribution over the image field and the interferometric array instrumental and signal propagation effects, estimated jointly from measurements of the electric vector spatial coherence function of the incident radiation field measured on each interferometer baseline \\citep[and references therein]{tho01}. The coherence data are sparsely sampled, leading to an ill-posed inverse imaging problem that requires regularization for convergent solution \\citep{cor99}. This regularization is typically imposed as a constraint on the properties of the source brightness distributions during deconvolution, including positivity and compact support \\citep{hog74}, or via information or entropy measures \\citep{nar84,cor85}.  The fidelity of the resulting source brightness distribution cannot be readily estimated given the analytic intractability of the coupled, non-linear calibration and imaging equation, combined with the fact that the parent probability distribution of the measured spatial coherence function is parametrized by the source brightness distribution and the instrumental array calibration, both unknown a prior at the time of observation. Instead, achieved image quality is typically estimated heuristically \\citep{eke86} or by approximate global measures. Common metrics include the ratio of the image brightness root-mean-square (rms) measured in regions of low brightness, $\\sigma_{off}$, to the thermal noise limit, $\\sigma_{th}$, as calculated for the array from the known antenna sensitivities and receiver and system thermal noise levels, assuming idealized observations of an unresolved point source with perfect array calibration \\citep{wro99}. Other measures include the achieved dynamic range $d_r$ (which is not a measure of image fidelity per se), expressed as the ratio of the peak brightness, $I_{peak}$ to the off-source rms, ${I_{peak}\\over{\\sigma_{off}}}$ \\citep{per86}, or the use of the deepest negative in a total intensity image (where Stokes $I > 0$) to derive a scaling relation to reduce the typical under-estimation of $\\sigma_{off}$ \\citep{kem93}.  These gross measures of image quality are however compromised by their idealized underlying assumptions, including that of implied direction-independence of image fidelity across the field. In practice image fidelity is not constant across the field. It is direction-dependent (equivalently, pixel-dependent) due to residual, unmodeled instrumental calibration errors jointly or separately in the visibility- or image-plane, and the interaction of these residual errors with non-linear deconvolution effects, amongst other factors. As a result, the assessment of pixel-dependent image fidelity is a general problem in radio interferometry. It is not confined to the important problem of calibrating direction-dependent instrumental errors specifically, as described for example by \\citet{bha08}.  Although a challenging and largely unsolved problem, understanding pixel-level radio-interferometric imaging fidelity is essential for current and future telescope arrays. Current instruments need to be made more accessible to the larger astronomical community; in particular their user base needs to expand beyond those who have invested the substantial time and effort required to acquire expertise in effective radio-interferometric calibration and imaging data reduction processes (heuristics as summarized for example by \\citet{per89}). This is best achieved through the provision of automated pipeline reduction systems; these require associated estimates of image fidelity to allow effective, broad-based community scientific interpretation and analysis.  For leading-edge future arrays, such as the Square Kilometer Array (SKA\\footnote{http://www.skatelescope.org}), solving the problem of quantitative pixel-level image fidelity assessment is key to their effective design. Any interferometer that cannot reach its target thermal noise limit in a representative integration needed to meet its key science goals is dynamic-range limited. This is important for many interferometers but is particularly acute for the SKA given its high sensitivity. As a result, the SKA has stringent imaging dynamic range requirements \\citep{sch07}, particularly in continuum observing modes where $d_r \\sim 10^6$ will be required routinely across wide image fields in rapid survey modes and $d_r \\sim 10^7$ for individual, targeted fields. This dynamic range is not achieved routinely by contemporary radio-interferometric arrays except for the innermost pixels in a handful of images, and then only after extensive custom reduction by the most skilled radio interferometry practitioners. With increasing projected radial distance from the field center the dynamic range may typically decline by several orders of magnitude, for reasons noted above. Contemporary examples of high dynamic-range and high-sensitivity observations in radio interferometry, and their associated challenges, are provided by \\citet{gel00}, \\citet{deb05}, and \\citet{nor05}.  Exponential advances in currently available and anticipated high-performance computing (HPC) capabilities and resources \\citep{bad08} allow fundamentally new approaches to the problem of pixel-level radio-interferometric imaging fidelity assessment. Two complementary approaches have recently been reported: a frequentist statistical resampling method \\citep[Paper I]{kem05} and a Bayesian imaging technique \\citep{sut06}. In Paper I, we described the first application of statistical bootstrap resampling techniques to the problem of interferometric imaging fidelity assessment. Statistical resampling is an active area in contemporary statistics research \\citep{efr03,dav03}, and general reviews are provided in recent monographs by \\citet{dav97}, \\citet{che99}, \\citet{pol99}, \\citet{lah03}, and \\citet{zou04}.  In the context of future interferometer arrays, such as the SKA discussed above, the statistical fidelity assessment methods discussed here are particularly important in understanding the optimal design of these telescopes and the contributing factors to their dynamic-range imaging performance. Specifically, although resampling methods may prove too computationally expensive for real-time calibration and imaging at a telescope as computationally demanding as the SKA, they are very valuable tools during SKA design and development. These methods can also be used in off-line analysis of data from interferometer arrays with smaller numbers of elements and lower associated computational costs for calibration and imaging.  We note that resampling techniques as described here are not strongly sensitive to the specific calibration and imaging algorithm in use; their goal is instead to provide a measure of the statistical properties of the underlying calibration and imaging estimator. Our initial evaluation in Paper I considered fidelity assessment of polarization calibration for Very Long Baseline Interferometry (VLBI) arrays, a representative radio-interferometric calibration and imaging problem of interest in its own right \\citep{kem99}. The statistical performance of both model-based and subsample bootstrap resampling was evaluated by inter-comparison of the bootstrap results with those obtained by direct Monte Carlo simulation for a single fixed two-component polarized source model and array configuration. Our initial study found both bootstrap resampling techniques to be computationally tractable with modern HPC resources and to have good statistical performance for image variance estimation for the single source model and array configuration considered in that initial study.  In the current paper we describe results from broader evaluation of the applicability of statistical resampling to radio-interferometric imaging fidelity assessment. We extend the scope of our original evaluation in Paper I by expanding the problem domain along two axes, namely source model structure and array parallactic angle coverage. Both parameters vary the expected degree of systematic error in the polarimetric calibration and imaging estimator used in the study and therefore the pixel-level morphology and magnitude of the image fidelity distribution over which the bootstrap resampling methods can be evaluated.  In the expanded evaluation reported here, we find here that the model-based and sub-sample bootstrap resampling methods for data with long-range statistical dependence, used in Paper I, continue to show good statistical performance when applied over a wider range of observing configurations and source models. We refine guidelines and rules for their general use in radio interferometry and report initial results on their applicability to image bias estimation.  The paper is structured as follows. In Section 2 the theory of bootstrap resampling is summarized, as applied to the interferometric calibration and imaging problem under examination. Section 3 describes the numerical simulation methods used to measure the statistical performance of the bootstrap resampling techniques under evaluation. Simulation results are presented in Section 4 and discussed in Section 5. In Section 6, we summarize the conclusions from the current work. ",
        "conclusions": "We draw the following conclusions from the current evaluation of bootstrap resampling for radio-interferometric imaging fidelity assessment over a broader domain subset:  \\begin{enumerate}  \\item{Statistical resampling techniques for dependent data remain a promising approach to the important problem of pixel-level fidelity assessment in radio-interferometric calibration and imaging. They show good statistical performance as assessed against direct Monte Carlo simulation, and are computationally affordable and scalable on contemporary and future HPC resources. They have the desirable statistical properties of independence on the detailed functional form or parametrization of the parent PDF of the observed visibility data and of not requiring an analytic solution for the calibration and imaging inference problem.}  \\item{Model-based bootstrap techniques outperform subsampling bootstrap methods in general in our study, although both are of value in this application. The subsample bootstrap performs well for delete fraction $f_s \\sim 0.75$, is straight-forward to implement, and has no adjustable parameters if $f_s$ is adopted as this value. The model-based bootstrap has superior statistical performance in general. It has free parameters $(N_p,\\ \\dtb)$ that can in principle be deduced from the observed data and against which the method is relatively robust to the exact values used.}  \\item{The subsample bootstrap variance scaling relation $v_f=f(f_s)$ measured in the current study shows relatively limited variability across the broader domain subset considered here. The current data are consistent with an exponential function: $v_f=be^{-a(1-f_s)}$.}  \\item{The polarization self-calibration algorithm introduced in Paper I is found to converge at a faster than exponential rate for good parallactic coverage $\\ta$. Its performance for lower parallactic angle ranges is found to deteriorate for $\\ta < 100^\\circ$ and to break down in the limit as $\\ta \\to 0$ as expected.}  \\item{An initial examination of statistical resampling for bias estimation shows qualified promise in the high-bias regime in this discipline, but further evaluation is required.}  \\end{enumerate}"
    },
    "0911/0911.2231_arXiv.txt": {
        "abstract": "We have observationally quantified the effect of gravitational torques on stars in disk galaxies due to the stellar distribution itself and explored whether these torques are efficient at transporting angular momentum within a Hubble Time.  We derive instantaneous torque maps for a sample of 24 spiral galaxies, based on stellar mass maps that were derived using the pixel-by-pixel mass-to-light estimator by Zibetti, Rix and Charlot.  In conjunction with an estimate of the rotation velocity, the mass maps allow us to determine the torque-induced instantaneous angular momentum flow across different radii, resulting from the overall stellar distributions for each galaxy in the sample.  By stacking the sample, which effectively replaces a time average by an ensemble average, we find that the torques due to the stellar disk act to transport angular momentum outward over much of the disk (within 3 disk scale lengths).  The strength of the ensemble-averaged gravitational torques within one disk scale length have a timescale of  ~ 4 Gyr for angular momentum redistribution.  The individual torque profiles show that only a third of our sample exhibit torques strong enough to redistribute angular momentum within a Hubble Time, mostly those with strong bars.  However, advective angular momentum transport is another source of angular momentum redistribution, especially in the presence of long-lived spiral arms, but is not accessible to direct observations.  The torque-driven angular momentum redistribution is thus observed to be effective, either in one third of disk galaxies or in most disk galaxies one third of the time and should lead to either changes in the mass density profile or the orbital shapes.  We use a set of self-consistent disk, bulge and halo simulated isolated disk galaxies with realistic cold gas fractions to verify that the torques exerted by the stellar distribution, such as spiral arms or a bar, exceed those of the gas and halo, as assumed in the analysis of the observations.  This study is the first to observationally determine the strength of torque-driven angular momentum flow of stars for a sample of spiral galaxies,  providing an important empirical constraint on secular evolution. ",
        "introduction": "Three primary aspects control the observable features of present-day galaxies: (1) the particular density fluctuations from which they originated, (2) successive mergers and various interactions with the environment and (3) the (internal) secular evolution.  While it is clear that the first and second point play a decisive role in shaping and reshaping a galaxy over time, the significance of the third mechanism continues to be debated.   In particular, it has not been empirically demonstrated how effective gravitational torques (induced by non-axisymmetric features like bars and spiral arms) are at redistributing angular momentum, which can lead to re-shaping the dynamics and stellar density profile of a galaxy over time.  In this present study we develop a framework on how to combine simulations, 'stellar mass density' images of spiral galaxies and sample averaging to address this question.  Specifically, we examine 24 galaxies with photometry from SDSS in order to assess how strong gravitational torques are at any given instant on the stellar distribution.  Unless otherwise stated, we consider torques due to the stellar distribution acting solely on the stellar distribution.   This modest galaxy sample serves to a) determine the limitations of both the method and the data; b) determine if gravitational torques are significant for a set of individual galaxies and c) explore how samples of galaxies should be 'stacked' in order to establish an average measure over time of the secular evolution for galaxy disks.  Subsequent papers will include an analysis of a much larger sample of galaxies.  In \\S 2 we describe our motivation and in \\S 3 we describe how the torques are calculated.  Tests completed with N-body/SPH simulations are outlined in \\S 4.  The observations and analysis are described in \\S 5 and \\S 6 and \\S 7 summarizes the conclusions. ",
        "conclusions": "Previous studies that have addressed the significance of secular evolution due to gravitational torques have taken either a computational or theoretical approach ({\\it i.e.} Lynden-Bell \\& Kalnajs, 1972, Bertin, 1983, Sellwood \\& Binney, 2002, etc.). We present here a direct observation-based estimate for the long-term torque-driven angular momentum transport .  We have extended and improved upon the work of G95, by using a larger sample of 24 galaxies and constructing accurate mass maps using $g$ and $i$-band images from SDSS and a pixel-by-pixel mass-to-light ratio (Z09).  These mass maps allowed us to derive instantaneous torque maps based on the stellar mass distribution and compute radial matter inflow profiles.  We stacked the 24 inflow profiles scaled by disk scale length in order to use an ensemble average to bypass the time integral required to investigate the long-term strength of secular evolution.  We tested the validity of this approach using N-body/SPH simulations, where we verified that the torques due to the stellar distribution are stronger than those from the dark matter distribution.  Thus, torques derived only from the stellar distribution are a conservative (lower) bound for the total torque and we derived a correction for the torques due to the dark matter distribution.  We also verified that the uncertainties in deprojecting the galaxies, finding the centroid and determining the disk scale length were moderate.  For each galaxy in the sample we derived instantaneous torque profiles scaled by disk scale length.  We translated these into angular momentum flow profiles.  However, we did not include the effects of advective transport, since there is no way to measure these effects observationally.  For the individual galaxies we found that there was consistently angular momentum outflow in the inner regions and the peak of each torque curve was within $1.5~r_{\\rmn{exp}}$.  The ensemble properties of the stacked sample of 24 galaxies show that for typical present day spiral galaxies: \\renewcommand{\\labelenumi}{\\roman{enumi})} \\begin{enumerate} \\item angular momentum is transported outward within $3~r_{\\rmn{exp}}$; \\item within $1~r_{\\rmn{exp}}$ the torques lead to an angular momentum outflow timescale much less than a Hubble Time; \\item  the average angular momentum flow timescale in this inner region is $\\sim$ 4 Gyr; \\item beyond $1~r_{\\rmn{exp}}$ the timescale is much longer than a Hubble Time, implying that the stellar radial profile has largely remained unchanged. \\end{enumerate} In conducting this study we have found that accurate mass maps are central in properly determining the torques due to the stellar distribution.  Global mass-to-light ratios ({\\it e.g.} scaled $i$-band images) are not sufficient and pixel-by-pixel mass-to-light ratios ensure that the total mass is properly estimated and that regions with young, blue, luminous populations, like spiral arms are not over-weighted (Z09).  Our results show that, for spiral galaxies as a class, secular evolution is {\\it observed} to be important within one disk scale length and may lead to the formation of pseudo-bulges.  Beyond one disk scale length however, the timescales are such that little to no change in the radial profiles should be expected. We stress that our simple picture must be qualified since advective transport and the fact that angular momentum flow does not directly translate into matter flow has not been treated.  This has been a preliminary pilot study and we are extending this project to a much larger sample of galaxies.  We will make comparisons between different spiral types and barred and unbarred galaxies in subsequent papers."
    },
    "0911/0911.0691_arXiv.txt": {
        "abstract": "We give a brief report on spectroscopic properties of V\\,393 Scorpii. H$\\alpha$ emission and shape and radial velocity of  He\\,I\\,5875 are modulated with the long cycle. The long cycle is explained as a relaxation cycle in the circumprimary disc,  that cumulates the mass transferred  from the donor until certain instability produces disc depletion. ",
        "introduction": "After the discovery of a long periodicity in the V\\,393 Sco ASAS light curve by \\cite[Pilecki \\& Szczygiel (2007)] {2007IBVS.5768....1P}, we realized that this star is a bright Galactic member of the new class of interacting binaries Double Periodic Variables (DPVs), that are characterized by two closely related photometric periodicities (\\cite[Mennickent et al. 2003] {2003A&A...399L..47M}, \\cite[2009c]{2009arXiv0908.3900M}). They have been interpreted as intermediate mass semi-detached binaries experiencing cyclic mass loss into the interstellar medium (\\cite[Mennickent et al. 2008]{2008MNRAS.389.1605M}, \\cite[Mennickent \\& Ko{\\l}aczkowski 2009 a,b,c]{2009RMxAC..35..166M, 2009arXiv0904.1539M, 2009arXiv0908.3900M}). The mechanisms for cyclic mass loss still is a matter of research, and a study of V393 Sco could yield  important insigths on the DPV phenomenon.  \\begin{figure}[b] \\begin{center} \\includegraphics[width=2.6in]{MennickentRVall.eps} \\includegraphics[width=2.6in]{Mennickentstrongweak.eps} \\caption{(Left) Radial velocities of the donor star yield $K_{2}$ = 169.7 $\\pm$ 0.6 km/s (crosses and pluses) whereas He\\,I\\,5875 radial velocities at different long cycle phases, 0.2 $< \\Phi_{l} \\leq $ 0.8 (asterisks) and 0.8 $< \\Phi_{l} \\leq $  0.2 (squares), are peculiar.  Orbital ephemeris is from Kreiner (2004). (Right) H$\\alpha$ emission at  $\\Phi_{l}$= 0.94 ($\\Phi_{o}$= 0.00, above) and  $\\Phi_{l}$= 0.57 ($\\Phi_{o}$= 0.97, down). } \\label{fig1} \\end{center} \\end{figure} ",
        "conclusions": ""
    },
    "0911/0911.1281_arXiv.txt": {
        "abstract": "We present a quantum-mechanical study of the exothermic $^7$LiH reaction with H. Accurate reactive probabilities and rate coefficients are obtained by solving the Schr\\\"{o}dinger equation for the motion of the three nuclei on a single Born-Oppenheimer potential energy surface (PES) and using a coupled-channel hyperspherical coordinate method. Our new rates indeed confirm earlier, qualitative predictions and some previous theoretical calculations, as discussed in the main text. In the astrophysical domain we find that the depletion process largely dominates for redshift (z) between 400 and 100, a range significant  for early Universe models. This new result from first-principle calculations leads us to definitively surmise that LiH should be already destroyed when the survival processes become important. Because of this very rapid depletion reaction, the fractional abundance of LiH is found to be drastically reduced, so that it should be very difficult to manage to observe it as an imprinted species in the cosmic background radiation (CBR). The present findings appear to settle the question of LiH observability in the early Universe. We further report several state-to-state computed reaction rates in the same range of temperatures of interest for the present problem. ",
        "introduction": "The chemistry of the early Universe plays an important role in our understanding of the formation of the first chemical objects and on the birth and evolution of galaxies and interstellar clusters.  Observation of molecular absorption and emission lines have provided important information about the physical conditions in which molecules are first formed then excited and destroyed. Lepp et \\emph{al} (2002) identified 200 reactions considered to contribute to the abundance of 23 atomic and molecular species. Molecular formation in the early Universe began when the temperature was low enough that the newly formed atoms could survive for further evolution. In this era, termed the recombination era, at a redshift around z = 2500 the Universe was about 100,000 years old and the temperature was about 4000 K. After recombination, the density was still very low and three-body (3B) reactions were still very inefficient: however, it was then that the first molecular species were postulated to be formed through radiative association. In spite of the low fractional abundance which is expected to exist for LiH ($\\sim$10$^{-10}$-10$^{-19}$), this molecule has nevertheless been considered one of the most likely  species to be important in that domain, due to its large permanent dipole moment and to the light masses of the involved atoms. Its search has therefore spurred several experimental studies in the recent past (Combes and Wiklind, 1998; de Bernardis et \\emph{al}, 1993). Indeed, it was further suggested by Dubrovich (1993) that primordial molecules with large dipole moments might be detectable by their imprint left on the CBR and arising from the resonant enhancement of the Thomson scattering cross sections which occur at the transition frequencies of the molecules. LiH was thus considered to be very relevant because of the role it may have played as a possible coolant during the late stages of  the gravitational collapse of the first cosmological objects, as this very molecule may have survived even H$_2$. In their works Lepp and Shull (1984), and later Puy et \\emph{al} (1993) and Palla et \\emph{al} (1995), concluded that the major formation mechanism of lithium hydride is likely to be via the following radiative association process: \\begin{equation} \\rm Li + H \\rightarrow LiH + \\nu \\end{equation} The rate coefficients for reaction (1) were later calculated by Dalgarno et \\emph{al} (1996) and Gianturco and Gori-Giorgi (1997), who considered also  the reduction of the abundance of LiH as a consequence of the exchange reaction, the latter being the strongly exothermic channel of the LiH destruction indicated below: \\begin{equation} \\rm LiH + H \\rightarrow Li + H_2 \\end{equation} Since the observational detection of LiH is still an  ongoing possibility among the options of experimental efforts, the present work intends to clearly, and accurately assess the feasibility of such measures vis \\'a vis the outcomes of rigorous computational data. We therefore intend to study the very reaction (2) in the redshift range 1000 $<$ z $<$ 20 to try to provide a more detailed evaluation of the possible role of this molecule in the early Universe evolution. Quantum coupled-channel equations have been solved to first calculate accurate rate coefficients as described in section 2, while the astrophysical consequences of our results will be detailed in sections 3 and 4, together with our conclusions. ",
        "conclusions": "As is well known by now, the first galaxies and stars were formed from an atomic gas of chiefly H and $^4$He with trace amounts of D, $^3$He and $^7$Li. Since there was a clear absence of heavier elements, the necessary radiative cooling kinetics below 8000 K must have been controlled by a small fraction of that gas which was molecular, in order to make more efficient the collapse of the primordial clouds. In the present work we have therefore revisited the chemistry of the $^7$LiH molecule, expected to be formed via radiative recombination via eq. (1), to now understand from ab-initio calculations how likely it would be for it to survive the possible reactions with the very abundant hydrogen gas existing under those conditions. We have carried out quantum calculations of the reactive processes that are allowed at the existing conditions of those clouds and which essentially involve the exothermic channels of reaction (2), i.e. either the hydrogen replacement process LiH$^a$ + H$^b$ $\\rightarrow$ LiH$^b$ + H$^a$, that allows the LiH molecule to survive, or the H$_2$ formation processes (into oH$_2$ and pH$_2$) LiH + H $\\rightarrow$ Li + o,pH$_2$, which describe the destruction of the initial molecular partner.  The calculations employ a newly obtained PES (Wernli et \\emph{al}, 2009) and carry out a coupled-channel study of the quantum reaction, summing over all possible final states of H$_2$, and further analysing the possible changes due to having the initial LiH molecule formed in some of its excited vibrational levels. The final rates were obtained by employing the uniform J-shifting procedure (Zhang and Zhang, 1999) and were extended to higher temperatures following the scheme outlined in section 2. The rates for both events (i.e. H replacement and H$_2$ formation) were thus obtained over a very broad range of redshift values, covering those that pertain to the recombination era.  The present calculations are the first quantum results in regions of astrophysical interest which employ both the \"exact\" coupled-channel dynamics and the best, thus far, reactive potential energy surface obtained from ab-initio computations (Wernli et \\emph{al}, 2009). They allow us to make the following points: \\begin{itemize} \\item over the range of relevant z  the destruction reaction dominates over the LiH \"survival\" of the hydrogen replacement channel (see figs 3 and 4); \\item to have the initial molecule still excited in any of its internal vibrational states does not modify the above result, although showing an increase of the simple H-replacement rates as $\\nu$ increases (see fig. 1); \\item the analysis of the rates of H$_2$ formation into different rotovibrational levels (fig. 2) indicates that such species is preferentially formed into its lower ($\\nu'$, $j'$) states, although the formation of internally \"hot\" H$_2$ molecules is far from negligible; \\item formation of either oH$_2$ or pH$_2$ changes very little the computed rate values, thereby not influencing the above findings; \\item the computed rates with the present PES turn out to be larger than the earlier calculations with a less accurate form of reactive interaction (Padmanaban, 2004; Petrongolo, 2005) and similar in size to the earlier, empirical estimates by Stancil et \\emph{al}, 1996. Such accurate findings therefore indicate as being very hard to have and to observe any possible LiH survival under early Universe conditions. \\end{itemize} In conclusion, the present new data strongly indicate rapid destruction of LiH molecules by exothermic reaction with the ubiquitous hydrogen atom in the clouds, a process very likely to occur at the time of the recombination era. As a consequence, it would be very difficult for such short-lived species to leave their observational imprinting in the CBR. One should also keep in mind that the present, accurate new depletion rates are directly affecting the use of cosmological models for estimating LiH abundances. As a consequence of such new values, the ensuing optical depth of the Universe due to elastic Thomson scattering of CBR photons from the present LiH species should be lower, and thus certainly different from that estimated in earlier studies (e.g. E. Bougleux and D. Galli, 1997): hence settling in the negative the question of possible LiH detection."
    },
    "0911/0911.1554_arXiv.txt": {
        "abstract": "The dynamical interactions that occur in newly formed planetary systems may reflect the conditions occurring in the protoplanetary disk out of which they formed. With this in mind, we explore the attainment and maintenance of orbital resonances by migrating planets in the terrestrial mass range. Migration time scales varying between $\\sim 10^6 $~yr and $\\sim 10^3 $~yr are considered. In the former case, for which the migration time is comparable to the lifetime of the  protoplanetary gas disk, a 2:1 resonance may be formed. In the latter, relatively rapid migration regime commensurabilities of high degree such as 8:7 or 11:10 may be formed. However, in any one large-scale migration  several different commensurabilities may be formed sequentially, each being associated with significant orbital evolution. We also use a simple analytic theory to develop conditions for first order commensurabilities to be formed. These depend on the degree of the commensurability, the imposed migration and circularization rates, and the planet mass ratios. These conditions are found to be consistent with the results of our simulations. ",
        "introduction": "The increasing number of extrasolar multi-planet systems have stimulated studies  of their  origin, evolution and  stability.  An important feature in  planetary system evolution is the occurrence of mean-motion resonances, which may relate to conditions at the time of or just after the process of  formation. There are  some well known examples of systems such as   Gliese 876 (\\cite{mbfv00}), HD 82943 (\\cite{munp00}) and 55 Cancri (\\cite{mmfq00}) involving  planets with masses in the Jovian  range. More recently a system of three superearths orbiting HD40307  has been announced for which the period ratios are not strictly commensurable but are close to a 2:1 commensurability (\\cite{mulp08}).  It is therefore  important to establish the main  features of the evolution of low mass planets embedded in  a gaseous disc and  in particular  to determine the types of resonant  configurations that might arise when a pair of such planets evolves together. The disc-planet interaction naturally produces orbital migration through the action of tidal torques (\\cite{gt80}, \\cite{lpap86}) which in turn  may lead to an orbital resonance in a many planet system (eg. \\cite{spn01}, \\cite{lp02}). This is because it is expected  that two planets with different masses will migrate at different rates. This has the consequence that their period ratio will evolve with time. In the situation where the migration is such that the orbits converge,   they tend to enter and  become  locked in a mean-motion resonance (eg. \\cite{np02}, \\cite{kpb04}) and then subsequently migrate together.  \\noindent In the simplest case of nearly circular and coplanar orbits the resonances that are formed  are the first-order resonances which occur at locations where the ratio of the two orbital periods can be expressed as $(p+1) / p$,  with $p$ being an integer.  Papaloizou \\& Szuszkiewicz (2005) performed a recent analytic and numerical study of the formation of first order commensurabilities by a system of two planets in the earth mass range migrating in a laminar disc. Here we extend these studies using $N$-body simulations to a larger range of migration rates and commensurabilities and compare the numerical work to  analytically derived conditions for particular commensurabilities to form. We begin by deriving the analytic criteria and go on to present the results of simulations to which they are reconciled. We consider migration time scales ranging between $\\sim 10^3 $~yr and $\\sim 10^6 $~yr. For the longest time scales,  2:1 commensurabilities may be set up while for the shortest, commensurabilities as high as 11:10 have been found. Finally we summarize our results. ",
        "conclusions": "We have studied  the development  of orbital resonances by migrating planets in the terrestrial mass range for  migration time scales varying between $\\sim 10^6 $~yr and $\\sim 10^3 $~yr. In a typical simulation, a sequence of resonances  occurred, each of which could survive significant orbital evolution before being lost. For the slowest migration rates considered, a 2:1 commensurability was able to form. At the fastest rates only commensurabilities with large $p$ persisted. We also found analytic conditions for first order commensurabilities to be formed.  These were consistent with our numerical simulations."
    },
    "0911/0911.0015.txt": {
        "abstract": "Very-high energy photons emitted by distant cosmic sources are absorbed on the extragalactic background light (EBL) during their propagation. This effect can be characterized in terms of a photon transfer function at Earth. The presence of  extragalactic magnetic fields could also induce conversions between very high-energy photons and hypothetical axion-like particles (ALPs). The turbulent structure of the extragalactic magnetic fields  would produce a stochastic behaviour in these conversions, leading to a statistical distribution of the photon transfer functions for the different realizations of the random magnetic fields. To characterize this effect, we derive new equations to calculate  the mean and the variance of this  distribution. We find  that, in presence of ALP conversions, the photon transfer functions on different lines of sight  could have relevant deviations with respect to the mean value, producing both an enhancement or a   suppression in the observable photon flux with respect to the expectations with only absorption. As a consequence, the most striking signature of the mixing  with ALPs would be a reconstructed EBL density from TeV photon observations which appears to vary over different directions of the sky: consistent with standard expectations in some regions, but inconsistent in others. \\\\ \\noindent {\\em Keywords}: axions, very high-energy gamma-rays. ",
        "introduction": "Axion-like particles (ALPs) with a two-photon vertex are predicted in many extensions of the Standard Model~\\cite{Svrcek:2006yi,Arvanitaki:2009fg,Masso:2006id}. The $a\\gamma\\gamma$ coupling allows for ALP-photon conversions in electric or magnetic field. This effect  is exploited by the ADMX experiment to search for   axion dark matter~\\cite{Duffy:2006aa}, by CAST to search for solar axions~\\cite{Zioutas:2004hi,Andriamonje:2007ew,Arik:2008mq}, and by regeneration laser experiments~\\cite{Robilliard:2007bq,Chou:2007zzc,% Fouche:2008jk,Afanasev:2008jt,ringwald}.  ALPs  also play an intriguing role in astrophysics. Indeed,   photons emitted by distant sources and propagating  through  cosmic magnetic fields can oscillate into ALPs.  The consequences of this effect have been studied in different situations~\\cite{Csaki:2001yk,Mirizzi:2005ng,Mirizzi:2006zy,Mirizzi:2009nq,Csaki:2003ef,Mirizzi:2007hr,Dupays:2005xs,% Hooper:2007bq,De Angelis:2007yu,Fairbairn:2009zi,Burrage:2009mj}. In particular, in the last recent years photon-ALP conversions have been proposed as a mechanism to avoid the opacity of the extragalactic sky to high-energy radiation due to pair production on the Extragalactic Background Light (EBL). At this regard, recent observations of cosmologically distant gamma-ray sources by ground-based gamma-ray telescopes have revealed a surprising degree of transparency of the universe to very high-energy (VHE) photons ($E\\gtrsim 100$~GeV)~\\cite{Aharonian:2005gh,Mazin:2007pn}. Surprisingly,  data seem to require a lower density of the EBL than expected and/or considerably harder injection spectra than initially thought~\\cite{Stecker:2007jq,Stecker:2007zj}. Oscillations between very high-energy  photons and ALPs could represent an intriguing possibility to explain this puzzle through a sort of ``cosmic light-shining through wall'' effect. Infact, if VHE  photons are converted into ALPS and then regenerated, they should not suffer absorption effects while they propagate as ALPs. In this sense, two complementary mechanisms have been proposed: a)  VHE photon-ALP conversions in the magnetic fields around gamma-ray sources~\\cite{Hooper:2007bq,Hochmuth:2007hk} and then  further back-conversions in the magnetic field of the Milky Way~\\cite{Simet:2007sa} (see also~\\cite{Bassan:2009gy}); b) oscillations of VHE photons into ALPs in the random extragalactic magnetic fields~\\cite{De Angelis:2007dy,DeAngelis:2008sk}. In principle, both the mechanisms can be combined together, as shown in~\\cite{SanchezConde:2009wu}. Currently, the inference of EBL from   VHE photons emitted by sources at high redshift~\\cite{alberto} is still object of debate, and it is not clear how robust are the conclusions on the absorption effects obtained  from the recent gamma data~\\cite{Costamante:2009gz}. In this sense, it is also possible that the observed transparency of the universe to VHE photons could be explained without the need of introducing nonstandard mechanisms (see, e.g.,~\\cite{Aharonian:2008su,Essey:2009zg}). Nevertheless, the seminal works mentioned before have pointed out the nice connection between VHE gamma-astronomy and ALP searches. Therefore, it seems  worthwhile to further explore  the consequences of this exciting possibility.  In this context, the treatment of the oscillations of VHE photons into ALPs in the extragalactic medium in presence  of the  absorption on the EBL presents a certain degree of complexity. Extragalactic magnetic fields are supposed to have a turbulent structure which can dramatically affect the development of photon-ALP conversions. Brute force numerical simulations  get the solution of the mixing equations  along a given photon line of sight by iterating the equations in each domain in which the magnetic field is assumed constant~\\cite{Csaki:2003ef,De Angelis:2007dy}. Since one cannot know the given configuration of magnetic domains crossed by VHE photons during their propagation, in the previous literature VHE photon-ALP conversions were usually characterized in terms of the mean conversion probability, obtained averaging the resulting conversion probabilities over an ensemble of magnetic field configurations along the photon line of sight. Such a procedure  slows down  the solution of our problem, since to obtain stable results typically the average has to be performed over more than $10^{3}$ realizations of the magnetic fields~\\cite{De Angelis:2007dy}. Moreover, the use of the mean probability as representative value does not appear a priori completely justified, since the variance of the probability distribution could produce relevant deviations from the mean value for  conversions occurring  in different realizations of the random magnetic fields.   In order to overcome these previous limitations, in this paper we perform a new study of the mixing equations of VHE photons in the turbulent magnetic fields and we provide a user-friendly calculation  of the mean and of the variance for the distribution of the photon transfer functions. The plan of our work is as follows. In Section~2 we  discuss about the  absorption of VHE photons on the extragalactic background light. In Section~3 we  characterize  VHE photon-ALP mixing in presence of absorption and we  present our calculation of the mean photon transfer function averaged over the ensemble of all the possible realizations of random  magnetic field configurations. We also show how to calculate the variance of the statistical distribution of the transfer functions. In Section~4 we  present our results for the photon transfer function of  VHE gamma-rays emitted from distant sources with and without the mixing with  ALPs. Contrarily to  previous predictions,  the presence of a  broad  variance in the statistical distribution of the photon transfer functions could produce both an enhancement or a suppression of the observed VHE photon flux with respect to the case with only absorption. The resulting photon flux would depend on the particular random magnetic field configuration along the photon  line of sight. As a consequence, photon-ALP mixing can not provide an universal mechanism to obtain the transparency of the universe to VHE radiation, but instead they would produce a strong direction-dependent behaviour in the flux of VHE photons from distant sources. In Section~5 we  discuss about possible developments of our study and we conclude. There follow two Appendices, in which we present some details for the derivation of the mean photon transfer function (Appendix~A) and for the variance of the distribution (Appendix~B).    %........................................................ ",
        "conclusions": "Very high-energy gamma-ray observations would open the possibility to probe the existence of  axion-like particles (ALPs) predicted in many theories beyond the Standard Model. Recent gamma observations of cosmologically distant gamma-ray sources  have revealed a surprising degree of transparency of the universe to very high-energy photons. The oscillations between  high-energy  photons and ALPs in the random extragalactic magnetic fields have been proposed as an intriguing possibility to explain these observations~\\cite{De Angelis:2007dy,DeAngelis:2008sk}. Apart from the original proposal, the consequences of this mechanism are testable with the measurements of the new generation of Imaging Atmospheric  Cherenkov Telescopes, like MAGIC~\\cite{magic}, HESS~\\cite{hess}, VERITAS~\\cite{Weekes:2001pd} or CANGAROO-III~\\cite{Enomoto:2001xu}, covering energies in the range 0.1-20~TeV, and hopefully with the future Cherenkov Telescope Array, reaching energies of 100~TeV~\\cite{cta}.  In order to perform a systematic study of these effects, without recurring to brute-force time-consuming numerical simulations, in this paper we have presented a simple calculations of the  photon mean transfer function and of its variance in presence of absorption on the EBL and mixing with ALPs. We have found that our prescription is enough accurate for the case under study and we have shown some numerical examples of VHE photon transfer functions. It results that  VHE photon-ALP mixing would produce peculiar deformations in the energy spectra of very high-energy gamma-rays emitted at high redshift. We have also found that the photon transfer function in presence of  VHE photon-ALP conversions presents a relevant dispersion around the mean value  due to the randomness of the extragalactic magnetic fields crossed by the photons.  Due to this fact, the measured  flux for VHE photons traveling along different lines of sight can be strongly suppressed or enhanced  with respect to the case with only absorption, depending on the particular configuration of random magnetic fields encountered. This would suggest that photon-ALP conversions could not be an universal mechanism to produce the transparency of universe to VHE photons. Conversely, the most striking signature of the mixing  with ALPs would be a reconstructed EBL density from TeV photon observations which appears to vary over different directions of the sky: consistent with standard expectations in some regions, but inconsistent in others. To test this effect  we would  need to collect data from  sources along different directions in the sky in order to perform a study of the photon energy distributions, from which we could hope to  infer possible hints of ALPs. A further signature of these stochastic conversions would be the detection of peculiar direction-dependent dimming effects in the diffuse photon radiation observable in GeV range, testable with the FERMI (previously called GLAST) experiment~\\cite{Gehrels:1999ri}.  As further developments, we plan to use our calculation to  perform a systematic study of ALP signatures in very high-energy gamma-rays, analyzing in details the spectral deformations expected for observed sources at different redshifts and for different models of the extragalactic background light. Thanks to our calculation, this task now appears more doable than before. After that, it will remain  to see if  elusive  ALPs will show up from the sky."
    },
    "0911/0911.3883_arXiv.txt": {
        "abstract": "{ By creating and analyzing the two dimensional gas temperature and abundance maps of the RGH 80 compact galaxy group with the high-quality \\Chandra\\ data, we detect a high-abundance ($\\simeq 0.7$ $Z_\\odot$) arc, where the metal abundance is significantly higher than the surrounding regions by $\\simeq 0.3$ $Z_\\odot$. This structure shows tight spatial correlations with the member galaxy PGC 046529, as well as with the arm-like feature identified on the X-ray image in the previous work of Randall et al. (2009).  Since no apparent signature of AGN activity is found associated with PGC 046529 in multi-band observations, and the gas temperature, metallicity, and mass of the high-abundance arc resemble those of the ISM of typical early-type galaxies, we conclude that this high-abundance structure is the remnant of the ISM of PGC 046529, which was stripped out of the galaxy by ram pressure stripping due to the motion of PGC 046529 in RGH 80. This novel case shows that ram pressure stripping can work efficiently in the metal enrichment process in galaxy groups, as it can in galaxy clusters. ",
        "introduction": "In galaxy groups and clusters, the inter-galactic medium (IGM) is enriched with the metals synthesized in the stellar interior, which were released into the vast inter-galactic space mainly via mergers (e.g., Durret et al. 2005), galactic wind (e.g., Nath \\& Trentham 1997), and AGN activities (e.g., Simionescu et al. 2009). As of today, the details of how these mechanisms work and how effective they are at different evolution stages of the IGM are still unclear (e.g., Aguirre et al. 2001; Hayakawa et al. 2004, 2006). Although AGN feedback has been regarded as the most popular and effective process to enrich the IGM, the works of e.g., Hayakawa et al. (2004, 2006) still showed that ram pressure stripping caused by merger may have played a crucial role in removing metals out of galaxies.   Very few cases of ram pressure stripping of metal-enriched gas contained in galaxies have been reported in literature (e.g., Sun et al. 2007; Hayakawa et al. 2004, 2006), and all these cases were found in galaxy clusters, where the dynamical environment is complicated. However, in compact galaxy groups where there exist much fewer member galaxies as compared with the clusters, the galaxy number density is actually as high as that of the core region of a rich cluster, so that strong interactions between member galaxies are often observed (e.g., Carlberg et al. 1994). This indicates that compact galaxy groups are ideal sites for us to study the effects of ram pressure stripping of metal-enriched gas held in the member galaxies.   The compact galaxy group RGH 80 ($z=0.0377$) may be such an ideal place. In 2009, Randall et al. reported that there exist interactions between the dominating galaxy PGC 046515 (E, Vorontsov-Velyaminov et al. 2001; R.A.=13h20m14.8s, Dec=+33d08m37.7s, J2000, Brinkmann et al. 2000) and the member galaxy PGC 046529 (E, Vorontsov-Velyaminov et al. 2001; R.A.=13h20m17.8, Dec=+33d08m41.2, J2000, Rines et al. 2003), and ascribed the origin of an arm-like X-ray feature identified on the X-ray map to the ram pressure stripping. However, a detailed analysis of the two dimensional distribution of gas metallicity, which is of fundamental importance to study the origin of the X-ray arm was not provided by the authors.   In this work, we calculate and study the two dimensional spatial distributions of both gas temperature and abundance of the hot IGM of RGH 80 with the high-quality \\Chandra\\ data archived at the \\Chandra\\ X-ray Observatory Science Center that is operated by the Smithsonian Astrophysical Observatory. We find firm evidence for the IGM of this group being enriched via the ram pressure stripping caused by the motion of the member galaxy PGC 046529. Throughout the paper, we quote errors at 90\\% confidence level unless mentioned otherwise. We adopt the solar abundance standard of Grevesse and Sauval (1998), where the iron abundance relative to hydrogen is $3.16\\times10^{-5}$ in number, and employ the cosmological parameters $H_0=71h_{71}$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_m=0.27$, $\\Omega_\\Lambda=0.73$, so that $1^\\prime$ corresponds to $44.2h_{71}^{-1}$ kpc. ",
        "conclusions": "We detect a high-abundance arc in the RGH 80 compact group, which shows tight spatial correlations with both the arm-like X-ray feature and the member galaxy PGC 046529. A reasonable and natural way to explain these correlations is to assume that the excess emission in the arm-like X-ray feature is due to the excess iron contained in the high-abundance arc, which originated from PGC 046529. But can this be real? Firstly, we note that the gas temperature ($1.1-1.3$ keV) and abundance ($\\simeq0.7$ $Z_\\odot$) of the high-abundance arc fall into the temperature and abundance range of early-type galaxies. On the other hand, considering that the metal abundance in the arc is about $0.3$ $Z_\\odot$ higher than its surrounding regions, and using the gas density profile given in Xue et al. (2004), we estimate that the total gas mass of the arm-like feature is $M_{\\rm gas}\\simeq5.4\\pm 0.5\\times10^{10}$ $M_\\odot$, so that the excess iron mass therein is $M_{\\rm Fe}\\simeq4.7_{-0.7}^{+6.5}\\times10^6$ $M_\\odot$. Assuming the time dependent type Ia supernova rate in Dahlen et al. (2004) and the wind mass loss rate in Ciotti et al. (1991), these amounts of gas and iron are close to the lower limit of gas and iron mass of the inter-stellar medium (ISM) of early-type galaxies, and can easily be produced in PGC 046529 within only $\\simeq 0.4$ Gyr, which strongly supports the idea that the high-abundance gas was originally contained in PGC 046529. Since there is no signatures of AGN activity (e.g., X-ray cavity and radio substructures) in PGC 046529, the possibility of the asymmetric high-abundance arc being formed by AGN activity can be excluded immediately. Hence the ram pressure stripping due to the motion of PGC 046529 becomes the most hopeful mechanism.  By examining Figure 2, we also note that the gas temperature in the high-abundance arc varies from $\\simeq1.3$ keV in the north (region B) to $\\simeq1.1$ keV in the east (region A). Using the best-fit spectral parameters from region A that are presented in Table 1 and Eq. (6) of Gu et al. (2009), we estimate that the time needed for the low temperature gas stripped out of PGC 046529 to achieve a thermal equilibrium with the surrounding IGM via thermal conduction is $t_{\\rm cond}\\simeq 0.2$ Gyr. On the other hand, based on the velocity dispersion ($450$ km s$^{-1}$; Ramella et al. 2002) of this group, we roughly estimate that the time needed for PGC 046529 to travel a distance equal to the separation between region B and region A is $t_{\\rm B\\to A}\\simeq0.15$ Gyr. Since $t_{\\rm cond}\\simeq t_{\\rm B\\to A}$ and PGC 046529 is located in region A, the spatial variation of gas temperature from region B to region A strongly supports the idea that the high abundance gas was stripped out of PGC 046529, and has been conductively heated since then.   Compared with the two straight gas tails behind the infalling galaxy ESO 137-001 in Abell 3627, as found by Sun et al. (2009), the spatial distribution of the high-abundance gas in RGH 80 forms a curved shape around PGC 046529, and possesses a width much wider than the gas tails behind ESO 137-001. This could be ascribed to the low spatial resolution ($\\simeq0.8^\\prime$ on the high-abundance arc) of our abundance map. Besides, since the location of PGC 046529 from the group center ($30.5h_{71}^{-1}$ kpc) is significantly less than the Roche limit \\begin{equation} d\\simeq2.44\\left(\\frac{M}{\\frac{4}{3}\\pi\\rho_m}\\right)^{1/3}\\simeq182~h_{71}^{-1}~{\\rm kpc}, \\end{equation} where $\\rho_m\\simeq5\\times10^{-26}$ g cm$^{-3}$ is the density of the gas stripped from PGC 046529, and $M\\simeq 1.5\\times10^{12}$ $M_\\odot$ is the total gravitational mass contained inside the radius of PGC 046529 (Xue et al. 2004), it is possible that the high-abundance arc has also been significantly broadened by the tidal force. Such an off-center enrichment process may help explain the central metal abundance dip observed in many galaxy clusters and groups (e.g., Abell 3266, Henriksen \\& Tittley 2002; the NGC 1550 group, Sun et al. 2003).   Based on above analysis and discussion, we conclude that the high-abundance gas around the member galaxy PGC 046529 in RGH 80 is the remnant of the ISM of this galaxy, which was removed outwards by the ram pressure due to the motion of PGC 046529. This novel case shows that the ram pressure can serve as an efficient metal enrichment mechanism in galaxy groups, just as in galaxy clusters."
    },
    "0911/0911.4515.txt": {
        "abstract": "We present a large robust sample of 1503 reliable and unconfused 70\\ts${\\mu}$m selected sources from the multiwavelength data set of the Cosmic Evolution Survey (COSMOS). Using the {\\it Spitzer} IRAC and MIPS photometry, we estimate the total infrared luminosity, $L_{\\rm IR}$ (8--1000\\ts${\\mu}$m), by finding the best fit template from several different template libraries. The long wavelength 70 and 160\\ts${\\mu}$m data allow us to obtain a reliable estimate of  $L_{\\rm IR}$, accurate to within 0.2 and 0.05\\ts dex, respectively. The 70\\ts$\\mu$m data point enables a significant improvement over the luminosity estimates possible with only a 24\\ts$\\mu$m detection. The full sample spans a wide range in $L_{\\rm IR}$, $L_{\\rm IR}\\approx 10^{8}-10^{14}\\ts L_{\\odot}$, with a median luminosity of $10^{11.4}\\ts L_{\\odot}$. We identify a total of 687 luminous, 303 ultraluminous, and 31 hyperluminous infrared galaxies (LIRGs, ULIRGs, and HyLIRGs) over the redshift range $0.01<z<3.5$ with a median redshift of 0.5. Presented here are the full spectral energy distributions for each of the sources compiled from the extensive multiwavelength data set from the ultraviolet (UV) to the far-infrared (FIR). A catalog of the general properties of the sample (including the photometry, redshifts, and $L_{\\rm IR}$) are included with this paper. We find that the overall shape of the spectral energy distribution (SED) and trends with $L_{\\rm IR}$ (e.g., IR color temperatures and optical-IR ratios) are similar to what has been seen in studies of local objects, however, our large sample allows us to see the extreme spread in UV to near-infrared (NIR) colors spanning nearly three orders of magnitude. In addition, using SED fits we find possible evidence for a subset of cooler ultraluminous objects than observed locally. However, until direct observations at longer wavelengths are obtained, the peak of emission and the dust temperature cannot be well constrained. We use these SEDs, along with the deep radio and X-ray coverage of the field, to identify a large sample of candidate active galactic nuclei (AGN). We find that the fraction of AGN increases strongly with $L_{\\rm IR}$, as it does in the local universe, and that nearly 70\\% of ULIRGs and all HyLIRGs likely host a powerful AGN. ",
        "introduction": "Luminous and ultraluminous infrared galaxies (LIRGs, $L_{\\rm IR} = 10^{11}-10^{12}\\ts L_{\\odot}$ and ULIRGs, $L_{\\rm IR} = 10^{12}-10^{13}\\ts L_{\\odot}$) have played a major role in the study of galaxy formation and evolution since their initial discovery. They were first identified in large numbers in the local universe by the {\\it Infrared Astronomical Satellite} (IRAS) and follow-up redshift surveys revealed that though they are rare in the local universe, their number density increases strongly with redshift (e.g., the Bright Galaxy Survey (BGS): \\citealt{Soifer:1989p2523}; the Revised Bright Galaxy Survey (RBGS): \\citealt{Sanders:2003p1575}; 1\\ts Jy ULIRG Survey: \\citealt{Kim:1998p3280}).  Since the launch of the {\\it Spitzer Space Telescope},  many deep infrared surveys (particularly at 24\\ts${\\mu}$m where the MIPS detector is most sensitive) have been undertaken and have shaped our view of the significance of IR galaxies out to $z\\sim1$. We now know that the cosmic star formation rate is dominated by LIRGs beyond $z>0.7$ \\citep{LeFloch:2005p2544} and ULIRGs start to dominate by $z\\sim2$ \\citep{Caputi:2007p2597,Magnelli:2009p2619}. However, one limitation in all 24\\ts${\\mu}$m selected studies is the difficulty in obtaining an accurate and reliable measurement of the total infrared luminosity of these sources since the 24\\ts${\\mu}$m selection wavelength represents rest-frame 12\\ts${\\mu}$m by $z=1$ and 8\\ts${\\mu}$m by $z=2$. Since the peak of the infrared emission for star-forming galaxies and galaxies with AGN is typically at $\\sim50-200$\\ts${\\mu}$m, at higher redshifts the selection wavelength moves further from the peak and becomes heavily influenced by the presence of polycyclic aromatic hydrocarbon (PAH) features.  In order to study how various galaxy properties are related to luminosity, an accurate estimate of the total infrared luminosity is needed. Previous studies \\citep{Bavouzet:2008p2620, Symeonidis:2008p2625} have shown that rest-frame monochromatic luminosities in the mid-infrared (MIR) correlate very well with the total $L_{\\rm IR}$, but they show a considerable amount of scatter. Of course, this problem becomes worse at higher redshifts as observed MIR probes shorter wavelengths. Ideally, one would use observations in the far-infrared (FIR) and submillimeter in order to sample the peak of the infrared emission in these sources, but current IR and submillimeter detectors lack the sensitivity to reach the necessary depths and only the most luminous sources are detected. Studies at long wavelengths are also hampered by confusion due to the large beam size of most detectors.  Previous studies \\citep[e.g.,][]{Huynh:2007p62, Symeonidis:2008p2625, Magnelli:2009p2619} based on 70\\ts${\\mu}$m selected sources have been hampered by this sensitivity problem and have lacked the depth or area necessary to detect a large sample of objects over a wide range of luminosities and redshifts. Deep surveys over small areas allow the detection of fainter, lower-luminosity sources, but they suffer from cosmic variance and poor statistics. Shallow surveys over wide areas only sample the bright end of the luminosity function. In order to overcome these problems, we have obtained deep imaging at 24, 70, and 160\\ts${\\mu}$m over the entire $\\sim 2 \\rm\\ts deg^2$ of the COSMOS field (Le Floc'h \\etal 2009, Frayer \\etal 2009).  The Cosmic Evolution Survey (COSMOS: \\citealt{Scoville:2007p1776}) was designed to study galaxy evolution over a large range of redshifts and environments, particularly with respect to large scale structure. The initial HST imaging \\citep{Koekemoer:2007p2299} has been complemented with wide multiwavelength coverage from space and the ground.  Along with many spectroscopic redshifts, the deep multiwavelength photometry has allowed us to obtain photometric redshifts with unprecedented accuracy, both for normal galaxies \\citep{Ilbert:2009p2146} and AGN \\citep{Salvato:2009p2142}. This unique data set provides the first opportunity to study the properties of a large sample of 70\\ts${\\mu}$m selected sources in detail over a wide range in redshift, luminosity, and environment.  This paper is the first of a two part series on the properties of a large sample of $\\sim$1500 70\\ts${\\mu}$m selected galaxies in the COSMOS field. \\S2 summarizes the COSMOS data sets used for this study and our sample selection is described in \\S3 along with a summary of general properties of the sources. We present the full UV--FIR SEDs and describe the estimate of their total infrared luminosity in \\S4. We discuss our results and their implications in \\S5 and present our conclusions in \\S6. Paper II will discuss the morphological and optical color properties of these sources as a function of redshift and luminosity.  Throughout this paper we assume a  $\\Lambda$CDM cosmology with $\\rm H_0=70\\ts \\rm km\\ts s^{-1} \\ts Mpc^{-1}$, $\\Omega_{\\Lambda}=0.7$, and $\\Omega_{m}=0.3$. All magnitudes are in the AB system unless otherwise stated. ",
        "conclusions": "In this paper we have presented a sample of 1503 70\\ts${\\mu}$m sources in the COSMOS field and determined a reliable measure of their total infrared luminosity. We believe that this is the largest sample of 70\\ts${\\mu}$m sources to date. When the source is detected at 160\\ts${\\mu}$m we obtain an estimate of $L_{\\rm IR}$ accurate to within 0.05 dex on average and 0.2 dex for those without a 160\\ts${\\mu}$m detection. This result is much better than can be obtained from a 24\\ts${\\mu}$m data point alone ($\\sim 0.5$ dex). In this section, we discuss in more detail the UV-to-FIR SEDs and the variations of the mean SED shape with $L_{\\rm IR}$. In particular we explore three spectral regions (the UV, NIR, and MIR) since these have been found to be the most sensitive to the underlying power source (i.e., starbursts and AGN). We make use of the multiwavelength properties of the sample to identify a large population of AGN candidates.  \\subsection{Analysis of SEDs}  Figures~\\ref{sed_z} and \\ref{sed_L} show the full range of spectral energy distributions that we observe for the 70\\ts ${\\mu}$m selected sample of galaxies binned by redshift and $L_{\\rm IR}$ with the median SED in each bin over plotted as a blue line.  The median SEDs are similar to sources observed locally for sources at low luminosity (e.g., Dale \\etal 2007: SINGS, $L_{\\rm IR} < 10^{11}\\ts L_{\\odot}$) and at high luminosity  (e.g., U \\etal 2009, in preparation: RBGS, $L_{\\rm IR} > 10^{11}\\ts L_{\\odot}$). All of these studies have shown that the median UV-to-FIR SEDs have two prominent thermal peaks (the optical stellar peak and the FIR dust peak) and, in addition, a range of spectral shapes, particularly in the UV (0.2--0.5\\ts${\\mu}$m), NIR (2--5\\ts${\\mu}$m), and MIR (10--30\\ts${\\mu}$m). With the large number of sources in our 70\\ts${\\mu}$m sample we can finally begin to quantify the full extent of these variations and explore how different parts of the SED correlate with each other and with $L_{\\rm IR}$.  As should already be known from studies of smaller samples of local objects (e.g., RBGS: \\citealt{Sanders:2003p1575}, 1 Jy ULIRG Survey: \\citealt{Kim:1998p3280}), it is clear that a single archetypical template for a given infrared luminosity (e.g., Arp 220 or Mrk 231 for ULIRGs) would not be applicable for all objects with that luminosity.   The median SEDs clearly show several general features. First, most of them show the 1.6\\ts ${\\mu}$m peak of the stellar bump. This feature becomes less pronounced at higher luminosities. The SEDs also display a prominent 5\\ts ${\\mu}$m dip at low luminosities which becomes less clear at high luminosities.  The 7.7\\ts ${\\mu}$m PAH emission feature is also evident at $L_{\\rm IR} <10^{11}\\ts L_{\\odot}$ near 8\\ts${\\mu}$m but becomes less evident in the higher luminosity bins as rest frame 7.7\\ts${\\mu}$m is no longer probed. However, in the highest luminosity bins, 24\\ts${\\mu}$m approaches rest-frame 7.7\\ts${\\mu}$m and this enhancement is still not observed, possibly due to the median SED being dominated by AGN at those luminosities. Many of the SEDs show an increase in emission toward the UV, indicative of the big blue bump \\citep{Malkan:1982p3097, Sanders:1989p3181,Elvis:1994p5581} seen in many AGN. The GALEX data points, particularly at high redshift where they probe further into the FUV, point to a broad maximum for many of these sources.  To provide a more quantitative measure of the spread in SED shapes, we compute three spectral indices (logarithmic rest-frame $\\nu L_{\\nu}$ ratios) probing three different wavelength regimes: UV/optical from 0.2 to 0.5\\ts$\\mu$m ($\\alpha^{0.5}_{0.2}$), NIR from 2 to 5\\ts$\\mu$m ($\\alpha^{5}_{2}$), and MIR from 10 to 30\\ts$\\mu$m ($\\alpha^{30}_{10}$). These spectral indices are plotted against each other in Figures~\\ref{color1}--\\ref{color3} binned by $L_{\\rm IR}$.  All three of these plots show strong trends with $L_{\\rm IR}$.   $\\alpha^{5}_{2}$ increases with $L_{\\rm IR}$ as sources become dominated by power-law SEDs (see \\S5.2). This is also evident in Figure~\\ref{color3} where the sources with high  $\\alpha^{5}_{2}$ form a separate distinct branch at intermediate values of  $\\alpha^{30}_{10}$. For galaxies with low values of  $\\alpha^{5}_{2}$,  $\\alpha^{30}_{10}$ is low at low luminosities and steadily increases with $L_{\\rm IR}$.  $\\alpha^{0.5}_{0.2}$ spans a wide range at all $L_{\\rm IR}$ (a range of nearly 3 orders of magnitude) and has lower values on average at the most extreme luminosities.  \\begin{figure*} \\epsscale{1.1} \\hspace*{-0.5in} \\plotone{f25} \\caption{Observed far-infrared colors ($\\nu F_{\\nu}(160)/ \\nu F_{\\nu}(70)$) as a function of redshift (left) and $L_{\\rm IR}$ (right). A typical $1\\ts\\sigma$ error bar is shown in the corner. The colored points represent the stacked data points (described in \\S4.4) for those sources not detected at 160\\ts $\\mu$m divided into several flux and redshift bins. Local infrared sources from the RBGS are over plotted with their FIR colors obtained from MIPS observations from SINGS ($10^{9}<L_{\\rm IR}<10^{11}\\ts L_{\\odot}$: open diamonds) and GOALS ($10^{11}<L_{\\rm IR}<10^{11.5}\\ts L_{\\odot}$: open triangles; $10^{11.5}<L_{\\rm IR}<10^{12.5}\\ts L_{\\odot}$: open squares). The gray lines represent the redshift evolution of empirical templates of local objects from \\cite{Dale:2002p2130} in the luminosity range $10^{8.3}<L_{\\rm IR}<10^{13.5}\\ts L_{\\odot}$. while the black lines highlight the luminosities at $10^{9}, 10^{10}, 10^{11}$, and $10^{12}\\ts L_{\\odot}$ from top to bottom.} \\label{fir_obs} \\end{figure*}  \\begin{figure*} \\epsscale{1.1} \\hspace*{-0.5in} \\plotone{f26} \\caption{Same as Fig.~\\ref{fir_obs} but converted into rest-frame colors using the best fit template to the infrared. Overall, the COSMOS sources span the same range in FIR colors as the local sources, however, at high luminosities there are a number of sources with colder SEDs than would be expected based on local sources. This may be evidence for a population of sources with colder SEDs at high redshift, however, at these redshifts the 160\\ts${\\mu}$m data point is not long enough to constrain the peak of the SED nor to obtain a precise dust temperature. The large range in SED templates that fit these data points is indicated by the fact that these sources are the same ones with a large scatter in Fig.~\\ref{lir_sym}.} \\label{fir_rf} \\end{figure*}  \\subsubsection{160/70\\ts${\\mu}$m Flux Ratios}  In addition to the SED shapes described above, our large 70\\ts${\\mu}$m sample and the substantial number of detections at 160\\ts${\\mu}$m makes it particularly useful to investigate the 160/70\\ts${\\mu}$m flux ratios as a crude measure of the characteristic dust temperature for the FIR peak emission. Figure~\\ref{fir_obs} shows the observed FIR colors ($\\nu F_{\\nu}(160)/ \\nu F_{\\nu}(70)$) as a function of redshift and $L_{\\rm IR}$ for all 70\\ts${\\mu}$m sources detected at 160\\ts${\\mu}$m. For those not detected at 160\\ts${\\mu}$m, the colored points on the plot represent the colors obtained from their stacked 160\\ts${\\mu}$m fluxes. To put these values in context we also over plot the flux ratios for local IRAS 60\\ts${\\mu}$m selected sources from the RBGS over a wide range of $L_{\\rm IR}$ ($10^8-10^{12.5}\\ts L_{\\odot}$) with MIPS observations obtained as a part of GOALS or SINGS (the Spitzer Infrared Nearby Galaxies Survey: \\citealt{Kennicutt:2003p5199}; \\citealt{Dale:2007p5190}). A typical error bar for the sources detected at 160\\ts${\\mu}$m is shown in the corner. The overall scatter of the data points remains the same even with a cut at larger $S/N$, indicating that the observed trends are likely real. The local sources span a large range in $\\nu F_{\\nu}(160) / \\nu F_{\\nu}(70)$ from 0.2--5 and show a strong trend with $L_{\\rm IR}$. Also over plotted in gray are a set of values obtained from the \\cite{Dale:2002p2130} templates spanning an $L_{\\rm IR}$ range of  $10^{8.3}-10^{13.5}$ from top to bottom as a function of redshift. These templates span most of the range occupied by the COSMOS 70\\ts${\\mu}$m selected sources (the full set of libraries covers the entire range of colors observed at each redshift but we show just this one library for illustration). The sources not detected at 160\\ts${\\mu}$m span the lower portion of the range occupied by the sources detected at 160\\ts${\\mu}$m. This result is expected if the 160\\ts${\\mu}$m detection preferentially selects colder sources than the rest of the sample.  We obtained rest-frame $\\nu F_{\\nu}(160) /  \\nu F_{\\nu}(70)$ values by using the 70 and 160\\ts${\\mu}$m fluxes determined from the best-fit IR template used to estimate $L_{\\rm IR}$. These values are plotted in Figure~\\ref{fir_rf} as a function of redshift and luminosity. As a note of caution, since these values are dependent on the best-fit SED template this adds an extra source of uncertainty, particularly at the high redshift end (where shorter rest-frame wavelengths are observed). We find that the COSMOS sources span the same range of colors (and thus dust temperatures) as the local sources but that there is an apparent excess of colder sources at higher luminosities, similar to the result of \\cite{Symeonidis:2009p5203}.  However, our interpretation is that this excess is likely a result of the longer wavelength selection (rest-frame 80\\ts${\\mu}$m at $z\\sim1$ versus 60\\ts${\\mu}$m locally) and the fact that the peak of emission is not well constrained at these redshifts and thus a wide range of templates can fit the data points. Future surveys at FIR and submillimeter wavelengths with {\\it Herschel} and SCUBA2 will be able to disentangle the temperature dependent selection effects and investigate the dust temperature distributions of high redshift galaxies.    \\subsection{AGN Diagnostics}  Both X-ray and radio detections are traditional methods used to identify potential AGN in high-redshift galaxies.  The vast majority of the 70\\ts${\\mu}$m selected galaxies that are detected in the X-ray (86\\%) have an X-ray luminosity greater than $10^{42} \\ts \\rm erg\\ts s^{-1}$ and are therefore bona fide AGN. The 22 lower luminosity sources are all at low redshifts ($z<0.35$).  Most of the X-ray detected AGN have power law SEDs (though not all), particularly at high $L_{\\rm IR}$. Many sources that are not detected as AGN in the X-ray also have a power-law SED. These are possibly heavily obscured (Compton thick) AGN that are opaque even to hard X-ray emission \\citep{Polletta:2006p5580}.  \\begin{figure*} \\epsscale{0.9} \\plotone{f27} \\caption{Radio/infrared flux ratio ($q$) vs. $L_{\\rm IR}$ (same as Figure~\\ref{q}) with X-ray detected AGN (red), objects with power-law SEDs (blue), IRAC color selected AGN candidates (magenta), and obscured AGN candidates (cyan) over plotted from top to bottom. The solid line indicates the median value of $q$ (2.34) for the entire sample and the dotted lines represent the dividing lines for radio ($q<1.64$) and infrared ($q>3.0$) excess sources following Yun \\etal (2001). The panels on the right show the distribution of the entire sample with radio detections and the different AGN selections over plotted in color.} \\label{q_compare} \\end{figure*}    Since $\\alpha^{5}_{2}$ is a measure of the $\\nu\\ts F_{\\nu}$ slope between 2 and 5\\ts${\\mu}$m it is therefore a good way to select power law galaxies. Anything above zero on the y-axis of Figure~\\ref{color1} or the x-axis of Figure~\\ref{color3} has the shape of a power-law in the NIR, i.e., these objects do not have a strong 1.6\\ts${\\mu}$m bump nor a dip at 5\\ts${\\mu}$m. The  $\\alpha^{5}_{2}>0$ cutoff here corresponds to a $F_{\\nu}$ spectral slope of 0.4 (similar to what is used as a power law slope cutoff in the literature, e.g., \\citealt{AlonsoHerrero:2006p2013} and \\citealt{Donley:2007p3092}). X-ray detected AGN span almost the entire area of the spectral index plots rather than being clearly separated the way the power law galaxies are. However, they do on average have lower $\\alpha^{30}_{10}$ as AGN in the local universe do.  We select all sources with $\\alpha^{5}_{2}> 0$ to be power-law AGN candidates. A total of 166 (11.2\\%) objects satisfy this criterion. Forty-five of these are detected in X-ray ($\\sim 27\\%$; a slightly lower fraction than found by \\citealt{AlonsoHerrero:2006p2013} and \\citealt{Donley:2007p3092} but consistent with the shallower COSMOS X-ray depths).  A few differences between power-law galaxies detected in the X-ray and those not detected in the X-ray can be seen in their SEDs. Those sources detected in the X-ray tend to display the full range of properties in the UV where only one or two sources that are not X-ray detected appear bright in the UV. The prominence of the UV emission in the X-ray detected power-law galaxies could be evidence that these sources tend to be less obscured than the non X-ray detected sources and hence are transparent to UV (and X-ray) emission. The SEDs of the non-X-ray detected power-law galaxies also tend to have a steeper slope.   We use the infrared-radio correlation as another way to investigate the energy source. Using the limits defined in \\cite{Yun:2001p2885} (5 times larger radio or IR flux than expected from the infrared-radio correlation) for radio excess ($q<1.64$) and infrared excess ($q>3.0$) objects we find many objects that fall into these two categories. In total there are 27 radio excess sources, 6 of which (22\\%) are also detected as X-ray AGN (as shown in Fig.~\\ref{q_compare}). The fraction of objects that this population represents increases with $L_{\\rm IR}$ from 4 (2\\%) at $<10^{11}\\ts L_{\\odot}$, to 5 (2\\%)  LIRGs and 16 (14\\%) ULIRGs.  Radio excess sources are believed to be potential AGN hosts where the extra radio emission seen comes from either a compact radio core or from associated radio jets or lobes (Sanders \\& Mirabel 1996).  All 14 of the infrared excess objects are ULIRGs or HyLIRGs and represent 11\\% of the sample at these luminosities. 6 (43\\%) of them are X-ray detected AGN. A total of 25 of the obscured AGN candidates (64\\%) described in \\S5.2.2 are detected in the radio and have $<q>=2.54\\pm0.63$. 5 of these (2 X-ray AGN) are infrared excess objects.  It is interesting to note that the IR excess objects contain a higher fraction of X-ray detected AGN than the radio excess objects. It is also interesting to note that many of the radio and infrared excess sources also have a power-law SED (see Fig.~\\ref{q_compare}).   \\subsubsection{IRAC Color Selected AGN}  \\begin{figure*} \\epsscale{0.8} \\plotone{f28} \\caption{\\cite{Stern:2005p3000} IRAC color-color plot for selecting AGN candidateswith X-ray detected AGN (red), objects with power-law SEDs (blue), radio excess objects (green), and obscured AGN candidates (cyan) over plotted from top to bottom. The \\cite{Stern:2005p3000} selection box is represented by the dashed lines.} \\label{irac} \\end{figure*}  Color selection in the mid-infrared has also been shown to be an effective way to select potential AGN \\citep{Stern:2005p3000,Lacy:2004p3062} as these colors can separate star-forming galaxies with a stellar bump at 1.6\\ts${\\mu}$m from those with a power-law like SED. Also, this method can potentially find dust obscured AGN since the MIR wavelengths are not as affected by extinction as the UV and optical. Following the criteria of \\cite{Stern:2005p3000} we used the color-color IRAC plot (see Fig.~\\ref{irac}) to select potential AGN and compare these sources with those found using four other methods. A total of 232 sources meet the Stern \\etal criteria, representing $\\sim 15\\%$ of the entire 70\\ts${\\mu}$m sample. Of these 232, 82 (35\\%) are detected as AGN in the X-ray. This is 62\\% of the total X-ray AGN population, which means the other 50 X-ray AGN are not selected by this method. According to Stern \\etal, this method is good at selecting broad line AGN (selects $\\sim 91\\%$) and not as good at selecting narrow line AGN ($\\sim 40\\%$) with an overall contamination rate of $\\sim 20\\%$. \\cite{Donley:2007p3092} show that this region of the plot can be contaminated by starburst dominated ULIRGs (like Arp 220) and normal star-forming galaxies. If we assume that 20\\% of the 232 objects are potentially contaminants, this leaves us with 186 actual AGN, or about 1.5 times the number detected in the X-ray.  Of the potentially obscured objects described in the next section ($F(24\\ts {\\mu} {\\rm m})/F(R) > 10^3$), 36 out of the 60 match these color criteria as do 33 out of the 56 that also have $R-K>4.5$. Most of the power law galaxies (62\\%) fall in a narrow region in the \\cite{Stern:2005p3000} selection box but many ($\\sim 35\\%$) fall outside of this region. Very few of these objects fall in the lower left corner of the selection region, indicating that this is possibly the region that contributes most of the contamination in the selection (consistent with the results of \\citealt{Donley:2008p3069}).  13 of the radio excess sources are also selected by the Stern \\etal criteria.  \\begin{figure*} \\epsscale{1} \\plotone{f29} \\caption{MIR (24\\ts$\\mu$m) to optical (R-band) flux ratio as a function of $R-K$ color with X-ray detected AGN (red), objects with power-law SEDs (blue), radio excess objects (green), and IRAC color selected AGN candidates (magenta) over plotted from top to bottom. The dotted lines represent the $F(24)/F(R)>10^{3}$ and $R-K>4.5$ selection cuts from Fiore \\etal 2008 for obscured AGN candidates.} \\label{obscure} \\end{figure*}   \\subsubsection{Obscured AGN Candidates}  Fiore \\etal (\\citeyear{Fiore:2008p2961}, \\citeyear{Fiore:2009p2143}) have shown that obscured AGN tend to have high MIR to optical flux ratios and red $R-K$ colors. The 24\\ts${\\mu}$m to R-band flux ratio for each of the 70\\ts${\\mu}$m sources in the sample is shown in Figure~\\ref{obscure} as a function of $R-K$ color with AGN selected through four other methods over plotted. The dotted lines show the cuts of $(R-K)_{\\rm Vega} > 4.5$ and $\\rm F(24\\ts {\\mu} m)/F(R) > 10^3$ defined in \\cite{Fiore:2008p2961} for selecting candidate obscured AGN at high redshift.  56 of the galaxies in our sample make both of these cuts and 14 of them are detected in the X-rays.  Three are detected only in the hard band and one only in the soft band. The remaining sources have a median hardness ratio of 0.2. The redshifts of these galaxies range from 0.7 to 2.5 with the median being $z=1.5$, confirming the results of \\cite{Fiore:2008p2961}. Additionally, four galaxies match the MIR to optical flux ratio criteria but are bluer than the $(R-K)_{\\rm Vega} > 4.5$ magnitude cutoff. Of these, one is detected as an AGN in the X-ray with a detection only in the hard band. These four sources have redshifts that range from $z=1.2$ to 2.4 with a median redshift of 2.1 -- a similar distribution to the sources with red $R-K$ colors. The 24\\ts${\\mu}$m to R-band flux ratio correlates with the total infrared luminosity (as expected for obscured AGN where the luminosity from the nuclear regions is greatly reduced by the obscuring material) and those above the cutoff of $10^3$ are predominantly ULIRGs and HyLIRGs. We consider all 61 of these galaxies to be obscured AGN candidates. %and include their numbers in Table~\\ref{frac}.  It is also worth noting that this selection is similar to other IR/optical color selections in the literature, such as dust obscured galaxies \\citep{Dey:2008p2965} and infrared bright, optically faint galaxies \\citep{Yan:2004p2979,Houck:2005p2986, Weedman:2006p2991}. The dust obscured galaxy color cut ($R-[24] > 14$) is actually the same as the one used here and is expected to select high redshift ($z\\sim2$) extremely luminous galaxies with large column densities of dust.   \\subsubsection{AGN Fraction Among Luminous and Ultraluminous Infrared Galaxies}  In the local universe, the AGN fraction of galaxies is known to increase very strongly with $L_{\\rm IR}$. Using spectroscopic diagnostics of a sample of 200 luminous infrared galaxies from IRAS, \\cite{Veilleux:1995p2081} found that 62\\% of objects at the highest luminosities were AGN. A similar study of the 1\\ts Jy sample of ULIRGs \\citep{Veilleux:1999p2073} found that above $L_{\\rm IR} > 10^{12.3}$ $\\sim 50\\%$ of objects are AGN.  \\cite{Tran:2001p3271} found a similar transition luminosity of $L_{\\rm IR} > 10^{12.4}\\ts L_{\\odot}$ between objects whose luminosities are dominated by star-burst emission and those dominated by AGN. Above this transition they find that ``most ULIRGs appear AGN-like\". With the large sample of objects over the wide range of luminosities probed by our 70\\ts${\\mu}$m selected sample, we can address the question of whether this trend holds out to higher redshifts.  Table~\\ref{xray_area} shows the total number of galaxies in each $L_{\\rm IR}$ bin along with the number of X-ray detected sources over the full area, in the area covered by {\\it Chandra}, and in the area with the deepest {\\it Chandra} coverage. The total fraction of X-ray detected AGN increases with increasing $L_{\\rm IR}$, from 5\\% at $L_{\\rm IR} < 10^{10}\\ts L_{\\odot}$ to $\\sim 15\\%$ for ULIRGs and $\\sim 42\\%$ for HyLIRGs, however, these numbers are lower limits due to the uneven X-ray coverage across the field. The fraction increases when limited to the smaller areas with deeper X-ray coverage, though the small numbers make a statistical analysis difficult. This indicates that as many as 12\\% of LIRGs, 30\\% of ULIRGs, and 43\\% of HyLIRGs are AGN with detectable X-ray emission. For the rest of the discussion here, we use the fractions obtained over the entire area.  Table~\\ref{frac} shows the fraction of candidate AGN found by each method (including X-ray detected, radio selected, IRAC color selected, power-law galaxies, and potentially obscured objects). If we add up all of the galaxies that fall into any of the possible AGN candidate categories (total column in Table~\\ref{frac}) then we find that the AGN fraction ranges from 3\\% at $L_{\\rm IR} < 10^{10}\\ts L_{\\odot}$ to $\\sim 70-100\\%$ for ULIRGs and HyLIRGs (illustrated in Fig.~\\ref{fraction}). There is likely to be some contamination in this estimate, particularly from the IRAC color selection method.  If we leave out the IRAC selected objects, the AGN fraction goes from 2\\% for objects with $L_{\\rm IR} < 10^{10}\\ts L_{\\odot}$ to 10\\% for LIRGs, 51\\% for ULIRGs, and 97\\% for HyLIRGs. This result is in excellent agreement with the AGN fraction in the local universe, indicating that high redshift ULIRGs are indeed similar to local sources in this respect.  \\begin{figure} \\epsscale{1.2} \\plotone{f30} \\caption{Fraction of 70\\ts${\\mu}$m sources selected as AGN candidates by any one of the five different selection methods illustrated in Figs.~\\ref{q_compare}--\\ref{obscure} as given in Column 9 of Table~\\ref{frac}. $1\\ts\\sigma$ Poissonian error bars are shown. Many of these sources are selected by more than one method though each source is counted only once. The fraction goes from $<5\\%$ at $L_{\\rm IR} < 10^{11}\\ts L_{\\odot}$ to 100\\% at $L_{\\rm IR} > 10^{13}\\ts L_{\\odot}$.} \\label{fraction} \\end{figure}  \\begin{figure*} \\plotone{f31} \\caption{Redshift (left) and infrared luminosity (right) histograms for AGN selected by each of the five different selection methods: (from top down) X-ray selected, radio excess sources, IRAC color selected sources, sources with power law SEDs, and obscured AGN candidates. The filled histograms represent the objects in each method that are also detected in the X-ray. The median $L_{\\rm IR}$ is plotted for each method.} \\label{hists} \\end{figure*}     The redshift and $L_{\\rm IR}$ distribution for the AGN candidates found using each method is shown in Figure~\\ref{hists}. Those objects that are also selected in the X-ray are shown by the filled histogram.  It is interesting to note that while the X-ray selection finds objects at a wide range in redshifts and $L_{\\rm IR}$, each of the other methods has a narrower range and typically identifies objects at higher luminosities.  The IRAC color selection finds two peaks in the redshift distribution, one at $z\\sim0.35$ and the other at $z\\sim1.15$. For the radio, IRAC, and power-law selections the X-ray seems to detect objects over the full range of redshifts and $L_{\\rm IR}$.  The majority of the power-law objects are ULIRGs or HyLIRGs with only a few ($\\sim 10$) with LIRG luminosities."
    },
    "0911/0911.3417_arXiv.txt": {
        "abstract": "The discovery of decay products of a short-lived radioisotope (SLRI) in the Allende meteorite led to the hypothesis that a supernova shock wave transported freshly synthesized SLRI to the presolar dense cloud core, triggered its self-gravitational collapse, and injected the SLRI into the core. Previous multidimensional numerical calculations of the shock-cloud collision process showed that this hypothesis is plausible when the shock wave and dense cloud core are assumed to remain isothermal at $\\sim 10$ K, but not when compressional heating to $\\sim 1000$ K is assumed. Our two-dimensional models (Boss et al. 2008) with the FLASH2.5 adaptive mesh refinement (AMR) hydrodynamics code have shown that a 20 km/sec shock front can simultaneously trigger collapse of a 1 $M_\\odot$ core and inject shock wave material, provided that cooling by molecular species such as H$_2$O, CO, and H$_2$ is included. Here we present the results for similar calculations with shock speeds ranging from 1 km/sec to 100 km/sec. We find that shock speeds in the range from 5 km/sec to 70 km/sec are able to trigger the collapse of a 2.2 $M_\\odot$ cloud while simultaneously injecting shock wave material: lower speed shocks do not achieve injection, while higher speed shocks do not trigger sustained collapse. The calculations continue to support the shock-wave trigger hypothesis for the formation of the solar system, though the injection efficiencies in the present models are lower than desired. ",
        "introduction": "Triggering the collapse of the presolar cloud with an interstellar shock wave propagating away from a site of stellar nucleosynthesis is a favored explanation for the widespread evidence of short-lived radioisotopes (SLRI) in chondritic refractory inclusions (Lee et al. 1976; Cameron and Truran 1977; MacPherson, Davis and Zinner 1995) and, much more rarely, in chondrules (Russell et al. 1996) found in primitive meteorites. The goal of this paper is to continue the theoretical exploration of the trigggering and injection scenario for SLRIs, in the context of shock waves striking a dense molecular cloud core that could have collapsed to form the solar system.  \\subsection{Short-Lived Radioisotopes}  The dozen or so confirmed ($^{10}$Be, $^{41}$Ca, $^{26}$Al, $^{60}$Fe, $^{53}$Mn, $^{107}$Pd, $^{182}$Hf, $^{129}$I, and $^{244}$Pu) or suspected ($^{99}$Tc, $^{36}$Cl, $^{205}$Pb, and $^{92}$Nb) short-lived radioisotopes may require a fairly involved history for their complete explanation (Goswami and Vanhala 2000; Meyer and Clayton 2000; McKeegan \\& Davis 2003; Wadhwa et al. 2007). A stellar nucleosynthetic source (a supernova or an AGB star) has been the leading explanation for most of these nuclei (Cameron 1993, 2001; Wasserburg et al. 1994, 1995, 1998; Trigo-Rodr\\'iguez et al. 2009; Huss et al. 2009), though there are other possibilities as well, in particular local production (i.e., in the solar nebula) of short-lived radioisotopes produced during spallation reactions involving energetic particles emanating from protosolar flares (Shu et al. 1997). Such local irradiation models appear to have a problem with being able to match the observed abundance ratio of $^{26}$Al to $^{41}$Ca (Srinivasan et al. 1996; Sahijpal et al. 1998; Lee et al. 1998). The observed abundance of $^{26}$Al thus seems to require its production by stellar nucleosynthesis (McKeegan et al. 2000). However, this problem can be avoided by assuming that the refractory inclusions are shielded by a less refractory mantle during the irradiation, with the mantle being lost later on during heating of the inclusions, thereby yielding the approximate abundance ratio of $^{26}$Al to $^{41}$Ca observed in certain meteorites (Gounelle et al. 2001). On the other hand, producing $^{26}$Al by multiple episodes of local irradiation would negate its use as a precise chronometer for the early solar system (Bizzarro et al. 2004; Halliday 2004; Krot et al. 2005; Thrane et al. 2006), which seems to be ruled out by the agreement of $^{26}$Al ages with those derived from the Pb-Pb dating system (Connelly et al. 2008).  Evidence has also appeared for the presence of the short-lived isotope $^{10}$Be in an Allende inclusion (McKeegan et al. 2000). Because $^{10}$Be is thought to be produced only by nuclear spallation reactions, its existence has been used to argue strongly in favour of local irradiation (McKeegan et al. 2000; Gounelle et al. 2001). Sahijpal \\& Gupta (2009) have calculated that even if all of the $^{10}$Be was produced by local irradiation, then the amount of $^{26}$Al also produced by local irradiation was about 10\\% of the total amount of $^{26}$Al, so that the bulk of the $^{26}$Al was probably synthesized by a massive star. However, if irradiation is responsible for the $^{10}$Be, it is unclear if the irradiation occurred in the solar nebula, or in an earlier phase of evolution. Arnould et al. (2000) point out that spallation can occur in the winds ejected from H-depleted Wolf-Rayet (WR) stars. Desch et al. (2004) showed that the $^{10}$Be might well have originated from $^{10}$Be galactic cosmic rays that were stopped in the presolar cloud. Evidence has also been advanced for the presence of live $^7$Be in Ca,Al-rich, refractory inclusions (CAIs), which are believed to represent the earliest solids formed in the solar nebula that have survived relatively unaltered (Chaussidon et al. 2006). Because of the extremely short half-life of $^7$Be of 53 days, this evidence, if correct, would require formation of $^7$Be in the solar nebula. Desch \\& Ouellette (2006) have disputed the $^7$Be claim, which has not been confirmed by other groups to date.  The short-lived isotope $^{60}$Fe (Goswami and Vanhala 2000) cannot be produced in the appropriate amounts by spallation, and requires a stellar nucleosynthetic source (Tachibana \\& Huss 2003), as does the bulk of the $^{26}$Al observed to be polluting the interstellar medium. The half-life of $^{60}$Fe is 2.6 million years (Rugel et al. 2009), roughly four times that of $^{26}$Al, so in any case, the evidence for live short-lived isotopes in refractory inclusions seems to require that no more than about 1 million years elapsed between the nucleosynthesis of some of the short-lived isotopes in a star and the formation of refractory inclusions in the solar nebula.  Solar-type stars are believed to form from the collapse of dense molecular cloud cores, which are supported against collapse primarily by magnetic fields (e.g., Mouschovias, Tassis, \\& Kunz 2006), though turbulence also plays a role (e.g., Kudoh \\& Basu 2008). Collapse of a quiescent cloud core begins once the magnetic field support decreases sufficiently through the process of ambipolar diffusion. Recently Kunz \\& Mouschovias (2009) have shown that ambipolar diffusion leading to collapse and fragmentation is able to reproduce the observed distribution of molecular cloud core masses, i.e., the initial core mass function, suggesting the importance of magnetic fields for star formation in general. Ambipolar diffusion is estimated to require of order 10 Myr to lead to collapse (Mouschovias, Tassis, \\& Kunz 2006), a period considerably longer than the half-life of $^{26}$Al of 0.73 Myr. If the Solar System's $^{26}$Al was produced in a massive star, it may have been injected promptly into the protosolar cloud, which must have then collapsed and formed cm-sized solids, all within $\\sim 1$ Myr. This constraint assumes that the same stellar outflow that carried the $^{26}$Al may have triggered the collapse of the protosolar cloud and injected other newly-synthesized elements, including other short-lived isotopes (Cameron \\& Truran 1977; Boss 1995). Abundant observational support exists for the triggering of star formation by expanding supernova shells in Upper Scorpius (Preibisch \\& Zinnecker 1999; Preibisch et al. 2002) and the Cygnus Loop (Patnaude et al. 2002), by superbubbles in OB associations (Oey et al. 2005; Lee \\& Chen 2009), by ionization fronts associated with HII regions (Leppanen, Liljestrom, \\& Diamond 1998; Healey, Hester, \\& Claussen 2004; Hester \\& Desch 2005; Snider et al. 2009), by generic external shocks (Tachihara et al. 2002), and by protostellar outflows (Barsony et al. 1998; Sandell \\& Knee 2001; Yokogawa et al. 2003). Here we consider generic shocks, with a special emphasis on supernovae and AGB stars as shock sources.  \\subsection{Injection Scenarios}  Boss (1995) showed that shock fronts from a nearby AGB star or a relatively distant supernova could trigger the collapse of a 3D dense cloud core and inject shock front material into the collapsing cloud. Foster \\& Boss (1996, 1997) studied this process in greater detail for axisymmetric, 2D clouds, and pointed out the crucial role of the assumed isothermal shock front for achieving both goals of triggered collapse and injection. A supernova shock passes through three phases: ejecta-dominated, Sedov blast wave, and radiative (Chevalier 1974). The latter phase occurs at distances of about 10 pc, after which the shock front sweeps up a cool shell of gas and dust as it propagates. Several recent AMR studies (Nakamura et al. 2006; Melioli et al. 2006) have confirmed the results of Boss (1995) and Foster \\& Boss (1996, 1997) that shock-triggered star formation is likely to occur when the supernova shock front has evolved into a radiative shock, i.e., the shock front is able to cool so rapidly by radiation that the thin shock front gas is essentially at the same temperature as the ambient gas, which is the same situation as in the isothermal shocks considered by Boss (1995) and Foster \\& Boss (1996, 1997).  Rayleigh-Taylor (R-T) fingers were identified as the physical mechanism for achieving injection of dust grains and gas into the collapsing presolar cloud (Foster \\& Boss 1997; Vanhala \\& Boss 2000, 2002). Because the R-T fingers strike the outermost layers of the presolar cloud, inducing collapse, the R-T fingers do not reach the central regions until shortly after the central protosun and the early solar nebula have formed, possibly explaining the lack of $^{26}$Al in certain (Fraction Unknown Nuclear = FUN) refractory inclusions  (Sahijpal \\& Goswami 1998) that may have formed before the R-T fingers arrived. Boss (2007) modeled the R-T finger injection process in the context of this scenario simply by imagining spraying the $^{26}$Al onto the surface of an existing solar nebula.  A related but alternative scenario involves having a nearby ($\\sim$ 0.1 pc) supernova inject $^{26}$Al directly into the solar nebula (rather than into the presolar cloud), as studied by Ouellette, Desch \\& Hester (2007). They found that the gas from a shock front could not be injected efficiently into a protoplanetary disk because of the disk's much higher density compared to the presolar cloud. Chevalier (2000) had found the same result and attributed it to the hot shock gas not having enough time to cool down. As a result, Ouellette et al. (2007) suggested that the $^{26}$Al resided primarily in dust grains that shot through the stalled shock-front gas and thereby penetrated into the disk. Dust grains smaller than 0.1 micron would be deflected, but micron-sized and larger dust grains would be injected with a high efficiency (Ouellette et al. 2009).  Supernovae are known to produce large amounts of dust grains, but theoretical models suggest that the newly condensed grains are essentially all smaller than 0.1 micron, and are sputtered to even smaller sizes in the reverse shock front driven into the expanding supernova remnant (SNR) by the interstellar medium that the SNR encounters (Bianchi \\& Schneider 2007). The models are in accord with dust extinction estimates for an observed reddened QSO. An insignificant fraction of the total dust grain mass is contained in grains larger than 0.1 micron, presenting a problem for scenarios that rely on large dust grains for injection. However, Nittler (2007) has argued that a sub-class of presolar grains appears to have been formed in a single supernova, conceivably the same supernova that produced many SLRIs. These presolar grains have sizes of 0.1 to 10 micron, large enough to be injected into the disk, or the presolar cloud core. What fraction of the mass of the initial dust grain population these relatively large grains represent is difficult to determine, given the processing associated with detecting presolar grains in meterorites and the typical limitation to the study of grains larger than about 0.1 micron (e.g., Amari, Lewis, \\& Anders 1994).  Recently it has been suggested that $^{60}$Fe was injected by a supernova directly into the solar nebula roughly 1 Myr after Solar System formation (Bizzarro et al. 2007), as in the Ouellette et al. (2007) models. In this scenario, the $^{26}$Al was derived from the wind from a Wolf-Rayet (WR) star, a massive star that would later become a supernova and inject the $^{60}$Fe. WR winds do indeed carry large amounts of $^{26}$Al (Arnould, Goriely, \\& Meynet 2006), with winds that are comparable in speed ($\\sim$ 1500 km/sec) to supernova shocks (Marchenko et al. 2006), meaning that such winds would need to be slowed down by sweeping up interstellar gas and dust to speeds less than $\\sim 40$ km/sec if they are to trigger cloud collapse rather than simply shred clouds to pieces (Foster \\& Boss 1996, 1997). However, other groups have not been able to replicate the Ni isotope data that forms the basis for the $^{60}$Fe scenario (Dauphas et al. 2008; Regelous et al. 2008). Nevertheless, WR stars should be considered as a possible source of $^{26}$Al in addition to that obtained from a supernova (e.g., Gaidos et al. 2009).  Williams \\& Gaidos (2007) estimated that the likelihood of a protoplanetary disk being struck by a supernova shock was less than 1\\%, but considering that we do not know if the Solar System's inventory of short-lived isotopes is rare or not, such an argument cannot be considered decisive. In fact, it has been argued recently that significant $^{26}$Al is necessary for the development of technological civilizations (Gilmour \\& Middleton 2009). However, Gounelle \\& Meibom (2007) and Krot et al. (2008) argued that injection could not have occurred directly onto a relatively late-phase, low mass solar nebula, as $^{26}$Al from a massive star supernova would have been accompanied by sufficient oxygen to lead to an oxygen isotope distribution in the solar nebula that would be distinct from that inferred for the sun based on mass-independent fractionation of carbonaceous chondrites and from that recently measured by the Genesis Mission (McKeegan et al. 2008). Krot et al. (2008) therefore argued that injection must have instead occurred into the presolar cloud, so that the sun and the solar nebula shared a common reservoir of oxygen isotopes that could then undergo fractionation. Ellinger, Young, \\& Desch (2009), however, pointed out that supernova explosions are not spherically symmetric, and so some degree of anisotropy in the ejecta is to be expected, perhaps enough to permit $^{26}$Al injection into the presolar nebula without disturbing the oxygen isotope ratios.  \\subsection{Shock Thermodynamics}  While isothermal shock fronts are capable of simultaneous triggering and injection (Boss 1995; Foster \\& Boss 1996, 1997; Vanhala \\& Boss 2000, 2002), it has been less clear what happens when detailed heating and cooling processes in the shock front are considered. Vanhala and Cameron (1998) found that when they allowed nonisothermal shocks in their models, they could not find a combination of target cloud and shock wave parameters that permitted both triggered collapse and injection to occur: they could trigger cloud collapse, or they could inject particles, but not both in the same simulation. Such an outcome would be fatal to the triggering and injection hypothesis if definitive. However, Vanhala \\& Cameron's (1998) models employed a smoothed-particle hydrodynamics (SPH) code, which has since been shown to be poor at resolving dynamical instabilities such as the Rayleigh-Taylor or Kelvin-Helmholtz instabilities (Agertz et al. 2007). Furthermore, Vanhala and Cameron's (1998) thermodynamical routines led to post-shock thermal profiles that were quite different from those of Kaufman \\& Neufeld (1996), who found that for a 40 km/sec shock, the post-shock gas cooled down from a maximum of over 3000 K to less than 100 K within a distance of 0.001 pc. With a 25 km/sec shock in the Vanhala \\& Cameron (1996) SPH code, however, the post-shock temperature rose to 3000 K and showed no signs of decreasing over a distance on the order of 0.5 pc.  Kaufmann \\& Neufeld (1996) studied C-type shock fronts, which result in the most successful models of the shock emission from the Kleinmann-Low nebula in the Orion molecular cloud. Kaufman \\& Neufeld (1996) assumed preshock magnetic field strengths in their C-shock models that are consistent with magnetic field strengths measured by Zeeman splitting in molecular clouds (Crutcher 1999). Such nondissociative, magnetohydrodynamic C-type shocks appear to be the correct analogue for the shock speeds ($\\sim$ 5 to 40 km s$^{-1}$) that we expect will be necessary to simultaneously achieve triggered collapse and injection. The relatively low postshock temperatures in C-type shocks are crucial for this scenario. Postshock cooling depends sensitively on the detailed microphysics, e.g., on the emission from rotational states of H$_2$O, H$_2$, CO, and OH, and hence on quantities such as the ratio of ortho- to para-hydrogen molecules (Neufeld \\& Kaufmann 1993; Kaufmann \\& Neufeld 1996). Atomic species are also important coolants, as are dust grains. More recently Morris, Desch, \\& Ciesla (2009) have reexamined the question of cooling shocked gas by H$_2$O line emission, finding that the cooling rates in the optically thin limit are at least as high as those calculated by Neufeld \\& Kaufmann (1993).  Boss et al. (2008) found that simultaneous triggering and injection was possible for a 20 km/sec shock striking a 1 $M_\\odot$ cloud, provided that Neufeld \\& Kaufmann's (1993) molecular cooling functions were employed. Gounelle et al. (2009) interpreted this result as meaning that only a very narrow range of shock speeds was consistent with the triggering and injection hypothesis. In this paper we examine what happens for a wide range of shock speeds (and different mass clouds) compared to the single shock speed considered by Boss et al. (2008), in order to determine the robustness of the shock wave trigger hypothesis. ",
        "conclusions": "When cooling by appropriate molecular species is included, shocks with speeds in the range from 5 km/sec to 70 km/sec are able to trigger the gravitational collapse of otherwise stable, dense cloud cores, as well as to inject shock wave material into the collapsing cloud cores. This injected material consists of shock wave gas as well as dust grains small enough to remain coupled to the gas, i.e., sub-micron-sized grains, which are expected to characterize supernova shock waves (Bianchi \\& Schneider 2007) and to carry the SLRI whose decay products have been found in refractory inclusions of chondritic meteorites. Evidently a radiative-phase supernova shock wave (Chevalier 1974) is able to cool sufficiently rapidly to behave in much the same way as a shock wave that is assumed to remain isothermal with the target cloud (e.g., Boss 1995). Given that Wolf-Rayet star winds and supernova shocks both are launched with shock speeds on the order of $10^3$ km/sec, these shock waves can only trigger collapse after they have travelled some distance (typically about 10 pc) and been slowed down to 5 km/sec to 70 km/sec by the snowplowing of intervening interstellar cloud gas and dust. The distance a fast shock must travel in order to slow down to speeds consistent with simultaneous triggering and injection is inversely proportional to the mean density of the intervening material (assuming this material to be moving much less than the shock speed); typical interstellar medium densities lead to distances of a few pc, depending on the desired shock speed. An AGB star wind with a typical speed of 10 to 30 km/sec could also have triggered collapse and injection without the need for snowplowing. The low injection efficiencies of the present models, however, point to the need to consider shock fronts and target clouds with different parameters, in order to learn if the injection efficiencies can be increased to the levels thought to be necessary to explain the observed abundances of fossil SLRIs in meteorites.  We are currently running three-dimensional (3D) models on the Xenia cluster at DTM. The need for a 3D treatment of the shock triggering process is evident from previous 3D studies of the R-T instability (Stone \\& Gardiner 2007) and shock fronts (Stone \\& Norman 1992; Whalen \\& Norman 2008). The R-T ``sheets'' that form in axisymmetric 2D models become true R-T fingers in 3D, allowing better penetration into the target dense cloud cores. We will present the results of these ongoing 3D models as well as of models with varied shock fronts and target clouds in future papers."
    },
    "0911/0911.0414_arXiv.txt": {
        "abstract": "We develop a theoretical framework for describing the hierarchical structure of the phase space of cold dark matter haloes, due to gravitationally bound substructures. Because it includes the full hierarchy of the cold dark matter initial conditions and is hence complementary to the halo model, the stable clustering hypothesis is applied for the first time here to the small-scale phase space structure. As an application, we show that the particle dark matter annihilation signal could be up to two orders of magnitude larger than that of the smooth halo within the Galactic virial radius. The local boost is inversely proportional to the smooth halo density, and thus is ${\\cal O} (1)$ within the solar radius, which could translate into interesting signatures for dark matter direct detection experiments: The temporal correlation of dark matter detection can change by a factor of $2$ in the span of 10 years, while there will be significant correlations in the velocity space of dark matter particles. This can introduce ${\\cal O} (1)$ uncertainty in the direction of local dark matter wind, which was believed to be a benchmark of directional dark matter searches or the annual modulation signal. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0911/0911.0622_arXiv.txt": {
        "abstract": "The problem of  cosmological particle creation for a spatially flat, homogeneous and isotropic Universes is discussed in the context of $f(R)$ theories of gravity. Different from cosmological models based on general relativity theory, it is found that a conformal invariant metric does not forbid the creation of massless particles during the early stages (radiation era) of the Universe. ",
        "introduction": "It is widely believed that the possible emergence of space and time trough a cosmic singularity and the presence of horizons in the Friedman-Robertson-Walker (FRW) models strongly suggest that matter and radiation need to be somehow created in  order to overcome some basic conceptual difficulties of big-bang cosmology \\cite{GWS83}. In these connections, the process of  matter creation in an expanding universe has been extensively discussed in the last four decades either from a macroscopic, as well as from  microscopic viewpoints [2-18].  Macroscopically, the matter creation was extensively investigated as a byproduct of bulk viscosity processes near the Planck era as well as in the slow-rollover phase of the new inflationary scenario \\cite{Zeld70,Murphy73,Hu82,Barrow86}. Later on,  the first self-consistent macroscopic formulation of the matter creation process was put forward by Prigogine and coworkers \\cite{Prigogine} and formulated in a covariant way by Calv\\~{a}o {\\it et al.} \\cite{CLW} with basis on the relativistic nonequilibrium thermodynamics. In comparison to the standard equilibrium equations, the process of creation  at the expense of the gravitational field is described by two new ingredients: a balance equation for the particle  number density and a negative pressure term in the stress tensor. Such quantities are related to each  other in a very definite way by the second law of thermodynamics. In  particular, the creation pressure depends on the creation rate and may operate, at level of Einstein's equations, to prevent either a spacetime singularity \\cite{Prigogine,AGL} or to generate an early inflationary phase \\cite{LGA}.  More recently, such formulation has also been applied to explain the late time accelerating stage of the Universe and other complementary observations \\cite{zimd,LSS08}.   Microscopically,  after the pioneering work of Parker \\cite{park68}, the quantum process of particle creation in the course of the cosmological expansion has also been studied by several authors \\cite{park69,books,mukh,partcrea,grav,gribmama}.  These `Parker Particles' are created because, according to the covariant equations of the fields in the Heisenberg picture, the positive and negative frequency parts of the fields become mixed during the universe expansion, so that the creation and annihilation operators at one time $t_1$ are linear combinations of those ones at an earlier time $t_2$, resulting at first, in a particle production.  One of the most interesting results from Parker's work is that in a radiation dominated Universe, the positive- and negative- frequency mode functions for massless fields are not mixed in the course of the expansion due to the conformal invariance of the metric. In particular, this means that there is no creation of massless particles, either of zero or non-zero spin. Viewed in another way, the dispersion relation in Fourier space  is exactly the same one of the flat Minkowski space, and, as such,  a wave propagation of a massless particle is not accompanied by particle production. As a consequence, in the GR framework, no photon, graviton or any other kind of massless particles are produced by purely expanding effects at early times \\cite{grav}.  Usually, the calculations of particle production  deal with comparing the particle number at asymptotically early and late times, or with respect to the vacuum states defined in two different frames and do not involve any loop calculation. However, the main problem one encounters when treating quantization in expanding backgrounds concerns the interpretation of the field theory in terms of particles. The absence of Poincar\\'e group symmetry in curved space-time leads to the problem of the definition of particles and vacuum states. The problem  may be solved by using the method of the diagonalization of instantaneous Hamiltonian  by a Bogoliubov transformation, which leads to finite results for the number of created particles \\cite{gribmama}.  On the other hand, the major developments concerning the study of particle production in curved background are concentrated in the standard general relativity. Recently, non-standard gravity theories have been proposed as an alternative to explain the present accelerating stage of the universe only with cold dark matter, that is, with no appealing to the existence of dark energy. This reduction of the so-called dark sector is naturally obtained in $f(R)$ gravity theories. In such an approach, the curvature scalar  $R$ in the Einstein-Hilbert action is replaced by a general function $f(R)$, so that the Einstein field equation is recovered as a particular case \\cite{fR,allemandi}.  A basic  motivation for this kind of generalization is the fact that higher order terms in curvature invariants have to be added to the effective Lagrangian of the gravitational field when quantum corrections are taken into account \\cite{gaspe,vilk,staro}. Two different formalisms are usually adopted to study $f(R)$ theories. In the so-called metric approach, the equations of motion are obtained by varying the metric tensor, while in the Palatini formalism, one considers the metric and the affine connection to be independent of each other, and one has to vary the action with respect to both of them. These two methods lead to the same equations only if $f(R)$ is a linear function of $R$ (GR case). The metric approach leads to fourth-order equations driving the evolution of the scale factor while the Palatini method leads to second-order equations, so that it is appealing because of its simplicity. Actually, for many interesting examples, it has been argued that some problems of instability are present in the metric approach \\cite{dolgov}, although some recent works have cast doubts on the existence of such instabilities \\cite{cembranos}. In what follows, we will focus our attention on the Palatini formalism and a simple power law type of $f(R)$. As shown by Vollick \\cite{fR} using the Palatine approach,  the inclusion  of a  $R^{-1}$ term in the action leads to a theory of gravity that predicts late time accelerated expansion of the Universe. Sotiriou \\cite{sotiriou} also demonstrated  that a $R^3$ term can account for an early time inflation, contrary to a $R^2$ type discussed by many authors (see Meng and Wang \\cite{meng} and Refs. there in).  More recently, Miranda {\\it et al.} discussed  a viable $f(R)$ cosmology based  on a logarithm contribution of the scale factor which emerges naturally as a limiting case of a general power law \\cite{waga09}.  In this letter we will present only the basic ideas and results on the scalar particle creation on expanding universe in the view of a $R^{n}$ type $f(R)$ theory. In brief, our basic result can be stated as follows:  massless particles in a radiation dominated Universe can be produced by purely expanding effects in generic $f(R)$ theories. In other words, Parker's result is valid only in the context of general relativity. ",
        "conclusions": "\\label{s5}  Due to the  discovery of the accelerating Universe at low redshifts, an increasing attention has been paid to $f(R)$ gravity as a possibility to reduce the so-called cosmological dark sector.  In this kind of cosmology the Universe can be accelerating driven only by cold dark matter.  In this work, assuming that the deviations from general relativity are described by a power law, we have addressed the problem of particle creation in $f(R)$ homogeneous and isotropic cosmologies dominated by a  fluid obeying a  $\\omega$-type equation of state. The mode equation in such background was derived and their general solution explicitly obtained. As an illustration, we have calculated the spectrum of created particles for a radiation dominated universe considering the cases of massive and massless particles.  Unlike the Parker results which are based on the standard general relativity, we have also explicitly shown that massless particles can be produced during the radiation dominated phase in the framework of  $f(R)$ cosmologies.  This result is not valid only for the case of a scalar field as discussed here. Actually,  due to the new contributions to the FRW differential equation,  one may show that the positive- and negative-frequency parts  of the fields become mixed in an expanding $f(R)$ cosmology thereby leading to the creation of massless particles of nonzero spin even at early times.         \\centerline{\\bf Acknowledgments} JASL is partially supported by CNPq and FAPESP No. 04/13668-0 and CHGB is supported by CNPq No 150429/2009-6 (Brazilian Research Agency)."
    },
    "0911/0911.5188_arXiv.txt": {
        "abstract": "Dark matter decaying or annihilating into $\\mu^+\\mu^-$ or $\\tau^+\\tau^-$ has been proposed as an explanation for the $e^\\pm$ anomalies reported by PAMELA and Fermi.  Recent analyses show that IceCube, supplemented by DeepCore, will be able to significantly constrain the parameter space of decays to $\\mu^+\\mu^-$, and rule out decays to $\\tau^+\\tau^-$ and annihilations to $\\mu^+\\mu^-$ in less than five years of running. These analyses rely on measuring track-like events in IceCube+DeepCore from down-going $\\nu_\\mu$. In this paper we show that by instead measuring cascade events, which are induced by all neutrino flavors, IceCube+DeepCore can rule out decays to $\\mu^+\\mu^-$ in only three years of running, and rule out decays to $\\tau^+\\tau^-$ and annihilation to $\\mu^+\\mu^-$ in only one year of running.  These constraints are highly robust to the choice of dark matter halo profile and independent of dark matter-nucleon cross section. ",
        "introduction": "The existence of dark matter has been established by numerous observations. Although it constitutes most of the matter in the universe~\\cite{WMAP5}, the nature of dark matter remains largely unknown.  One widely held possibility is that it is a new fundamental particle produced in the early universe and present today as a thermal relic.  Among the best motivated of these are the so-called weakly interacting massive particles, or ``WIMPs'' (for reviews, see Refs.~\\cite{review:jungman,review:bertone}) which are predicted to be undergoing annihilations~\\cite{ann:sun,ann:earth,ann:halo,ann:dsph} and possibly decays~\\cite{wimp:decay} in the current epoch.  Recently, the instruments PAMELA~\\cite{Adriani:2008zr}, ATIC~\\cite{ATIC}, PPB-BETS~\\cite{Torii:2008xu}, and Fermi~\\cite{Abdo:2009zk} have observed features in the cosmic-ray $e^\\pm$ spectrum and a positron fraction that are inconsistent with known backgrounds.  While these anomalies may be due to unidentified astrophysical sources~\\cite{pulsars}, one exciting possibility is that they are due to the decay~\\cite{Chen:2008yi,Chen:2008dh,Yin:2008bs, Ishiwata:2008cv,Shirai:2009kh,Ibarra:2009bm,Ibe:2009en} or annihilation~\\cite{Barger:pmodel,Cholis:pmodel,Bergstrom:pmodel, Sommerfeld:cirelli, Various:pmodels,Sommerfeld:nima} of dark matter particles into standard model states.  Even in the case that anomalies are not caused by new physics in the dark sector, the constraints are generally applicable to dark matter models.  In order for dark matter to explain the anomalies, the products of decay or annihilation must be primarily leptonic.  In either case, the Fermi observation gives the most precise preferred region in mass and lifetime/cross-section; for decays, only $\\mu^+\\mu^-$ and $\\tau^+\\tau^-$ final states fit the data, while for annihilations a small region of $e^+e^-$ is also permitted~\\cite{Meade:fermi}. For the allowed parameter space, decays or annihilations into hadrons and weak and Higgs bosons are severely limited.  Decays are constrained by the Fermi observation of the isotropic diffuse gamma-ray flux~\\cite{Chen:fermiiso}, and annihilations are constrained by the production of antiprotons~\\cite{Adriani:2008zq}.  Explaining the anomalies with annihilations has an additional challenge. In order to match the observed rates, the cross-section required is $10^3$--$10^4$ times that expected for thermal production of the dark matter in the early universe. This necessitates nonthermal production mechanisms or low-velocity enhancements to the cross-section, such as the Sommerfeld enhancement~\\cite{Sommerfeld:cirelli,Sommerfeld:nima} or the Breit-Wigner enhancement~\\cite{Feldman:breit,Ibe:breit}. Moreover, high cross section annihilations to leptonic states are tightly constrained by synchrotron radio emissions from the galactic center, although this constraint can be evaded if the true galactic dark matter halo profile is much less steep at the galactic center than benchmark profiles~\\cite{Bertone:gcradio,Meade:fermi}.  For either decays or annihilations, neutrino observations may provide the strongest constraints --- or the most promising corroborating signatures --- for dark matter to be the source of the anomalies. With the exception of the disfavored $e^+e^-$ channel, the required leptonic final states decay into neutrinos. These travel undeflected from their sources, eliminating any uncertainties in modeling their propagation. Moreover, astrophysical sources that may explain the anomalies are not expected to produce a large flux of neutrinos.  Pulsars, for example, would generate at most ${\\cal O}(1)$ events/year at a full km-scale detector, and only at energies greater than 10 TeV~\\cite{pulsarbkg}.  There are, however, challenges to constraining dark matter models with neutrinos. As they are observed only by collecting Cherenkov light from induced particle showers or from secondary muons, angular resolution is poor compared to gamma-ray observatories.  Also, there are significant backgrounds from atmospheric muon and neutrino fluxes.  However, if these backgrounds can be controlled, the poor angular resolution need not be a barrier; indeed by integrating the flux over a large area of the sky, the resulting constraint is much less sensitive to the choice of dark matter profile~\\cite{beacom:anncosm,yuksel:anngal}. Moreover, by observing the galactic and extragalactic diffuse dark matter, rather than any that may have been captured by the Sun or the Earth, any constraints will be independent of the dark matter-nucleon cross sections, which are related to final states in a model-specific way~\\cite{dmcapture}.  Recent analyses show the power of neutrino constraints, using various strategies to reduce the effect of atmospheric backgrounds.  The Super-Kamiokande observatory resides in the northern hemisphere, facing away from the galactic center, minimizing atmospheric backgrounds.  Measurements of upward-going muons place a limit on the flux of galactic $\\nu_\\mu$, providing a robust constraint that eliminates annihilations to $\\tau^+\\tau^-$ as a source of the $e^\\pm$ anomalies~\\cite{superk,Meade:fermi}. The IceCube observatory, on the other hand, resides at the South Pole where down-going atmospheric fluxes are coincident with the neutrinos from the galactic center.  The overwhelming background of atmospheric muons can be suppressed by event selection to establish an isotropic diffuse flux limit~\\cite{kowalski:diff,gerhardt:amanda,kiryluk:icecube}, but this limit only starts at a high energy threshold ${\\cal O}(10\\;{\\rm TeV})$, and yields a relatively weak constraint as we will show.  Currently under construction is DeepCore~\\cite{deepcore}, an in-fill of the IceCube detector which will use the outer instrumented volume as a veto on downward-going muons to a level of one part in $10^6$~\\cite{deepcorestatus}. This will allow the galactic neutrino flux to be measured and compared against the atmospheric neutrino flux, providing a constraint on dark matter decays and annihilations. Recent work~\\cite{dc:ann,dc:decay} by some of the authors shows that IceCube+DeepCore will be able to significantly constrain the parameter space of decays to $\\mu^+\\mu^-$, and rule out decays to $\\tau^+\\tau^-$ and annihilation to $\\mu^+\\mu^-$ as possible sources of the anomalies in less than five years of running.  Recently, IceCube+DeepCore has demonstrated in simulations the ability to distinguish between track-like events, which are due to the charged-current interactions of $\\nu_\\mu$, and cascade events, which are induced by $\\nu_{e,\\tau}$ through charged-current interactions and by all neutrino flavors through neutral-current interactions~\\cite{middell:casc,grant:nutau}. This is very useful for constraining dark matter neutrino fluxes because $\\nu_\\mu$ is the dominant flavor of atmospheric neutrinos above 40~GeV~\\cite{numunuebkg,nutaubkg}. The neutrino-nucleon cross sections are the same for all flavors, so signal would be considerably enhanced, while background would be reduced because $\\nu_\\mu$ only creates cascade events through the neutral-current interaction, which is lower in cross section than the charged-current interaction~\\cite{nuNcs,beacom:casc,beacom:anncosm} (see Fig.~\\ref{nuNcs}).  Moreover, cascade events are easy to distinguish from the tracks caused by any atmospheric muons that are not vetoed by the outer volume. \\begin{figure}[t!] \\begin{center} \\includegraphics[width=0.49\\linewidth]{nuNcs.pdf} \\end{center} \\caption{Cross sections for (anti)neutrino-nucleon interactions, given by Ref.~\\cite{nuNcs}.  The blue lines are for charged-current interactions, and the red lines are for neutral-current interactions.  Solid lines are for neutrinos and dashed lines are for antineutrinos.} \\label{nuNcs} \\end{figure}  In this paper we show that observation of cascade events at IceCube+DeepCore can enhance the neutrino constraints on dark matter, and rule out (or corroborate) dark matter decays to $\\mu^+\\mu^-$ or $\\tau^+\\tau^-$ and annihilation to $\\mu^+\\mu^-$ as sources of the observed $e^\\pm$ cosmic-ray anomalies in a much shorter time compared to searches that rely solely on track-like events. ",
        "conclusions": "We have shown that, by using cascade events, IceCube+DeepCore can more quickly establish constraints on dark matter models that would explain the reported $e^\\pm$ anomalies, and over time establish stronger constraints than from track-like events. Specifically, track-like events will be able to significantly constrain the parameter space of decays to $\\mu^+\\mu^-$, and rule out decays to $\\tau^+\\tau^-$ and annihilations to $\\mu^+\\mu^-$ in less than five years of running.  In comparison, cascade events can rule out decays to $\\mu^+\\mu^-$ in only three years, and rule out decays to $\\tau^+\\tau^-$ and annihilation to $\\mu^+\\mu^-$ after only one year.  Moreover, these constraints are highly robust to the choice of dark matter halo profile and independent of dark matter-nucleon cross section.  In closing, we note two interesting possibilities for future work.  First, if the pointing accuracy for track-like events at IceCube+DeepCore is established to be less than $10^\\circ$, the signal-to-noise for annihilations may be significantly enhanced by observing a smaller region around the galactic center, possibly out-performing cascade searches (albeit with greater dependence on the choice of dark matter halo profile).  This would strengthen the discovery potential for dark matter because the galactic center could be identified as a localized source of excess neutrinos.   Second, because the $\\nu_\\tau$ atmospheric background is so low at energies above 40 GeV and at low zenith angles, if IceCube+DeepCore can demonstrate efficient $\\nu_\\tau$ discrimination~\\cite{grant:nutau}, signal to noise could be increased by a factor $\\sim 100$.  This would put leptophilic dark matter to a severe test."
    },
    "0911/0911.3285.txt": {
        "abstract": "% context heading (optional) {} % aims heading (mandatory) {We demonstrate the application of our 3D radiative transfer framework  in the model atmosphere code \\phx\\ for a number of spectrum synthesis calculations for very different  conditions. } % methods heading (mandatory) {The 3DRT framework discussed in the previous papers of this series was added to our general-purpose model atmosphere code \\phxO\\ and an extended 3D version \\phxT\\ was created.  The \\phxT\\ code is parallelized via the  MPI library using  a hierarchical domain decomposition and displays very good strong scaling. } % results heading (mandatory) {We present the results of several test cases for widely different atmosphere conditions and compare the 3D calculations with equivalent 1D models to assess the internal accuracy of the 3D modeling. In addition, we show the results for a number of parameterized 3D structures. } %conclusions heading (optional)%,leave it empty if necessary {With presently available computational resources it is possible to solve the full 3D radiative transfer (including scattering) problem with the same micro-physics as included in 1D modeling.} ",
        "introduction": "In a series of papers \\citet*[][hereafter: Papers I--V]{3drt_paper1, 3drt_paper2, 3drt_paper3, 3drt_paper4, 3drt_paper5}, we have described a framework for the solution of the radiative transfer equation in 3D systems (3DRT), including a detailed treatment of scattering in continua and lines with a non-local operator splitting method. These papers deal solely with the radiation transport problem and its numerical solution for test cases designed to stress-test the algorithms and codes. It is important, however, to apply the radiative transfer codes to `real' problems, e.g., model atmosphere simulations and to compare the results to 1D equivalents. We have extended our general purpose model atmosphere code \\phx\\ to use the 3DRT framework so that the new version of \\phx\\ can calculate both 1D (\\phxO) and 3D (\\phxT) models and spectra. In this paper we will describe the implementation and the results of \\phx\\ calculations comparing the results of 1D and 3D spectrum syntheses for different model parameters. ",
        "conclusions": "We have described first results we have obtained by incorporating the 3D radiative transfer framework we have discussed in Papers I-V into our general purpose model atmosphere package \\phx, thus allowing both 1D models (\\phxO) and 3D models (\\phxT) with the same micro-physics. We have verified and tested \\phxT\\ by computing a number of test spectra for 1D conditions and comparing the results to the corresponding \\phxO\\ calculations. The conditions range from M dwarfs, solar type stars to A stars and Type II supernovae with relativistic expansions speeds. In addition, we have calculated spectra for a 3D hydrodynamical simulation of solar atmosphere convection. These tests demonstrate the it is now possible to calculate realistic spectra for 3D configurations including complex micro-physics. \\phxT\\ can be used to calculate synthetic spectra for a number of complex 3D atmosphere model, including irradiated stars or planets, novae, and supernovae. We are currently working on extensions of the 3D radiative transfer framework to arbitrary velocity fields in the Euler (for low velocities, e.g., in convection simulations or planetary winds) and the Lagrangian (for Supernovae, accretion disks and matter flow in the vicinity of black holes) frames, which will extend the applications of \\phxT\\ significantly."
    },
    "0911/0911.5680_arXiv.txt": {
        "abstract": "{The \\object{Sagittarius dwarf Spheroidal Galaxy} (\\sgr) provides us with a unique possibility of studying a dwarf galaxy merging event while still in progress.  Moving along a short-period, quasi-polar orbit in the Milky Way Halo, \\sgr~ is being tidally dispersed along a huge stellar stream.  Due to its low distance (25 kpc), the main body of \\sgr~ covers a vast area in the sky (roughly $15 \\times 7$ degrees). Available photometric and spectroscopic studies have concentrated either on the central part of the galaxy or on the stellar stream, but the overwhelming majority of the galaxy body has never been probed. } {The aim of the present study is twofold. On the one hand, to produce color magnitude diagrams across the extension of \\sgr~ to study its stellar populations, searching for age and/or composition gradients (or lack thereof). On the other hand, to derive spectroscopic low-resolution radial velocities for a subsample of stars to determine membership to \\sgr~ for the purpose of high resolution spectroscopic follow-up.  } {We used VIMOS@VLT to produce V and I photometry on 7 fields across the \\sgr~ minor and major axis, plus 3 more centered on the associated globular clusters \\ters, \\tere~ and \\arpt. A last field has been centered on \\mfifo, lying in the center of \\sgr. VIMOS high resolution spectroscopic mode has then been used to derive radial velocities for a subsample of the observed stars, concentrating on objects having colors and magnitudes compatible with the \\sgr~ red giant branch.} {We present photometry for 320,000 stars across the main body of \\sgr, one of the richest, and safely the most wide-angle sampling ever produced for this fundamental object. We also provide robust memberships for more than one hundred stars, whose high resolution spectroscopic analysis will be the object of forthcoming papers. \\sgr~appears remarkably uniform among the observed fields. We confirm the presence of a main \\sgr~ population characterized roughly by the same metallicity of 47 Tuc, but we also found the presence of multiple populations on the peripheral fields of the galaxy, with a metallicity spanning from [Fe/H]=-2.3 to a nearly solar value.  } {}  \\titlerunning{VIMOS photometry of \\sgr} ",
        "introduction": "Dwarf galaxy mergers are believed to play a key role in the buildup of the Milky Way (MW) Halo \\citep[][]{searle78,carollo07,diemand07}. These are also invoked as suspect responsible of the appearance of the Thick Disk \\citep[][]{kroupa02}, of the Bulge/Bar system \\citep[][]{heller07}, and of the Disk warp \\citep[][]{ibata97}. The discovery of the \\object{Sagittarius dwarf Spheroidal galaxy} \\citep[\\sgr;][]{ibata94,ibata95} allowed for the first time to observe the MW in the act of tidally accreting a dwarf galaxy. Immediately after the discovery, theoretical predictions \\citep[][]{ibata97,helmi99} as well as observational hints \\citep[][]{ibata97,Ng1,ngs97,Ng2,majewski99,ivezic00} concurred in indicating that \\sgr~ should be releasing its stellar content in the Halo along a massive, kinematically cold stream. The existence of the stream was finally confirmed by  \\citet[][]{ibata01} and \\citet[][]{majewski03}, and further branches of it were also identified in the stellar sample of the Sloan Digital Sky Survey \\citep[][]{newberg03,martinez04,belokurov06}. As a matter of fact, the \\sgr~ stream is clearly the most prominent feature visible in the galactic Halo.  All such discoveries increased the importance of \\sgr~ not only as a ``test case'' of tidal merging, but as the source of a significant portion of the stellar population in the Halo. \\citet{majewski03} estimated that roughly 75\\% of high-latitude Halo M-giants originated from \\sgr~ debris, which also provided almost all Halo AGB C-stars \\citep[see][]{mauron05}. \\citet{zijlstra06} estimated that up to 10\\% of the Halo material could come from \\sgr~ debris. A thorough reconstruction of the \\sgr~ stream(s), moreover, would pose very strong constraints on the  shape of the MW dark matter Halo \\citep[e.g.][]{fellhauer06}. With the discovery of \\sgr, 4 globular clusters (GC) which were before considered belonging to the halo have been recognized as being associated with the \\sgr. While three of them (\\ters, \\tere, and \\arpt) lay at the outskirts of \\sgr, the position of the fourth one, \\mfifo,  coincides with the center of the dwarf spheroidal \\citep[][]{monaco05b} although it likely formed elsewhere and fell subsequently into the galaxy core \\citep[][]{bellazzini08}. It became subsequently clear that a number of other globular clusters, currently lying far from \\sgr, have kinematical properties compatible with an origin within the \\sgr~ system, and a subsequent stripping \\citep[][]{bellazzini03}, or were actually embedded into the stream material \\citep[][]{martinez02}. For at least one of them, \\object{Pal 12}, the origin within the \\sgr~ system has now been firmly established \\citep[][]{cohen04,sbordone2007}. The large \\sgr~size (roughly 15$\\times$7 square degrees in the sky), a consequence of its modest distance \\citep[25 kpc, see][]{monaco04}, is likely the reason why only the central region of the galaxy, and the associated globular clusters, have been studied so far. \\mfifo~ and the surrounding area have been examined photometrically \\citep[e.g.][]{M98,bella99,bellazzini99B,LS00,monaco04,bella06, siegel07} as well as by means of high resolution spectroscopy \\citep[]{bonifacio00,bonifacio2004,tarantella04,zaggia2004,sbordone2005,caffau05,monaco2005,mcwilliam05,sbordone2007,mottini08}. While ground based photometry suggests for the \\sgr~ main body population a mean metallicity of [Fe/H]$\\sim$-1, spectroscopy shows a surprising abundance spread in the main body, with a significant population around [Fe/H]=0, a mean value around [Fe/H]=-0.5 and a weak metal poor tail possibly extending below [Fe/H]=-2.5. The spectroscopic discovery of the metal rich population in the \\sgr~core was then confirmed by the detection of its Turn Off superimposed over the M54 CMD in the HST ACS photometry by \\citep[][]{siegel07}. Globular clusters considered belonging to \\sgr ~ spans a similar wide range in metallicity, from [Fe/H]=-0.6 for \\ters~ to [Fe/H]=-2.3 for \\tere.  Stars in the \\sgr~ stream appear to be slightly less metal rich than those in the main body, with metallicity gradients  along the stream \\citep[][]{bellazzini06b,clewley06,monaco07,chou07}. This would not be unexpected if the most metal poor stars were located in the outskirts of \\sgr~ and then subject to the tidal interaction with the MW.  \\begin{table}[tp]\\tiny \\caption{List of observed fields} \\label{radec} \\centering \\begin{tabular}{lcccccc} \\hline\\hline &&&&&&\\\\[-5pt] Field & Pos. & RA        & DEC       & l      & b       & VIMOS \\\\ Name  &      & (J2000.0) & (J2000.0) &  (deg) & (deg)   & cameras \\\\[5pt] \\hline &&&&&&\\\\[-5pt] M 54 & Cen & 18:55:03.3 & -30:28:41 &  5.954 &  -14.692 & 4 \\\\ Sgr0 & Min & 18:56:12.9 & -30:01:48 &  6.482 &  -14.743 & 4 \\\\ Sgr1 & Maj & 18:43:00.0 & -28:58:34 &  6.258 &  -11.689 & 4 \\\\ Sgr2 & Maj & 18:48:56.1 & -29:41:49 &  6.139 &  -13.165 & 3 \\\\ Sgr3 & Maj & 19:00:57.6 & -31:12:00 &  5.779 &  -16.140 & 3 \\\\ Sgr4 & Maj & 19:07:00.0 & -31:57:04 &  5.561 &  -17.615 & 3 \\\\ Sgr5 & Min & 18:57:13.6 & -29:35:24 &  6.990 &  -14.770 & 3 \\\\ Sgr6 & Min & 18:52:44.6 & -31:20:31 &  4.931 &  -14.581 & 4 \\\\[5pt] \\hline \\end{tabular} \\end{table}   The detailed chemical analysis of \\sgr~ population showed a highly peculiar abundance pattern, with anomalous abundances for a number of elements \\citep[][and references therein]{sbordone2007}. The detection of the same signature in \\object{Pal 12} was the concluding argument to assign the cluster to the \\sgr~ system.  Although it seems now established \\citep[e. g.][]{lanfra07,lanfra06A,lanfra06B} that galactic winds triggered by starbursts are required to explain the low [$\\alpha$/Fe] ratios observed in many dSph, the complete history of the  \\sgr~ is far from being clarified. For instance, the extent to which \\sgr~ has been altered by the interaction with the MW is currently difficult to assess. The original mass of \\sgr~ is uncertain, as well as its original gas content. The large mass of the stream, the unusually high metallicity among the dSph, and the large population of GCs, are all indications that we are dealing with an object which was in the past more a dwarf elliptical than a ``classic'' dwarf spheroidal galaxy.  There is thus a number of open questions regarding this complex stellar system, some of which can only be addressed by the analysis of stellar populations {\\em across} the main body of \\sgr. Among these open questions we will only  mention the following: how did star formation proceed in \\sgr? How uniform is its chemical composition? Which is the lower limit of metallicity in \\sgr and does it show gradients in composition or age ? What is the dynamics of the main body?   \\begin{figure*}[t] \\centering \\includegraphics[width=18.5cm]{mappacampi.eps} \\caption{ Map of the \\sgr~ galaxy obtained from 2MASS and UCAC catalog selected in K, J$-$K, and E(B$-$V) (E(B$-$V) $< 0.555$, $0.95 <$ (J$-$K)$_0 < 1.10$, $10.5 <$ K$_0< 12$, see \\citet{majewski03}. The boxes shows the position of our fields} \\label{mappa} \\end{figure*} ",
        "conclusions": "In this paper we present a V, I VIMOS photometric catalog for $\\sim$320000 sources across eight fields in the direction of the Sagittarius dwarf Spheroidal galaxy, covering roughly 0.43 square degrees. The main purpose of this survey was to probe \\sgr~ stellar populations over a significant fraction of its spatial extension, from its center to its periphery, and to provide targets for high resolution spectroscopical analysis.  Our survey extends over 6.2 degrees along the galaxy major axis, and 2 degrees along the minor axis, i.e. about one third of \\sgr~ extension. Two occurrences limit the extension of this sort of studies: first the presence of the MW disk which makes  the contamination too severe at the position of our field Sgr 1, second the fact that  the density of \\sgr~ stars drops quickly when moving away from the center, making a spectroscopic follow-up highly inefficient. During the FLAMES follow-up observations, in fields such as Sgr 1, Sgr4, Sgr 5, and Sgr 6 we were able to place only a few dozens of fibers on ``radial velocity members'' detected by VIMOS-MOS. VIMOS spectroscopy provided radial velocities adequate for FLAMES candidates selection, but not good enough for dynamical analysis. The uncertainties in the MOS wavelength calibration, together with the limited S/N of the spectra led to a final uncertainty of about 10-20 km s$^{-1}$, comparable to the overall \\sgr~ velocity dispersion (see Fig. \\ref{allflames}).  Within the precision allowed by the mentioned MW contamination, the population of  \\sgr~ appears remarkably uniform among the observed fields. We thus confirm the presence of a main \\sgr~ population characterized roughly by the same metallicity of 47 Tuc, but we also found the presence of multiple populations on the peripheral fields Sgr0-Sgr6, with a metallicity spanning from [Fe/H]=-2.3 to a nearly solar value.  Our research  thus confirms the extremely complex evolution of this object.  Using our FLAMES data we plan to better characterize these populations both from a chemical and a dynamical point of view, adding new pieces to this intriguing puzzle."
    },
    "0911/0911.0293_arXiv.txt": {
        "abstract": "{Stellar encounters potentially affect the evolution of the protoplanetary discs in the Orion Nebula Cluster (ONC). However, the role of encounters in other cluster environments is less known.} {We investigate the effect of the encounter-induced disc-mass loss in different cluster environments.} {Starting from an ONC-like cluster we vary the cluster size and density to determine the correlation of collision time scale and disc-mass loss. We use the {\\textsc{\\mbox{nbody6\\raise.2ex\\hbox{\\tiny{++}}}}} code to model the dynamics of these clusters and analyze the disc-mass loss due to encounters.} {We find that the encounter rate depends strongly on the cluster density but remains rather unaffected by the size of the stellar population. This dependency translates directly into the effect on the encounter-induced disc-mass loss. The essential outcome of the simulations are: i) Even in clusters four times sparser than the ONC the effect of encounters is still apparent. ii) The density of the ONC itself marks a threshold: in less dense and less massive clusters it is the massive stars that dominate the encounter-induced disc-mass loss whereas in denser and more massive clusters the low-mass stars play the major role for the disc mass removal.} {It seems that in the central regions of young dense star clusters -- the common sites of star formation -- stellar encounters do affect the evolution of the protoplanetary discs. With higher cluster density low-mass stars become more heavily involved in this process. This finding allows for the extrapolation towards extreme stellar systems: in case of the Arches cluster one would expect stellar encounters to destroy the discs of most of the low- and high-mass stars in several hundred thousand years, whereas intermediate mass stars are able to retain to some extant their discs even under these harsh environmental conditions.} ",
        "introduction": "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). Though it is known that discs disperse over time \\citep{2001ApJ...553L.153H,2002astro.ph.10520H,2006ApJ...638..897S,2008ApJ...672..558C} and that in dense clusters ($n$\\,$\\gtrsim$\\,$10^3$\\,pc$^{-3}$) the disc frequency seems to be lower in the core \\citep[e.g.][]{2007ApJ...660.1532B}, it is an open question as to how far interactions with the surrounding stars influence planet formation in clusters of different densities.  Two major mechanisms are potentially able to strongly affect the evolution of protoplanetary discs and planets in a cluster environment: photoevaporation and gravitational interactions. Photoevaporation causes the heating and evaporation of disc matter by the intense UV radiation from massive stars. First models of photoevaporation were developed by \\citet{1998ApJ...499..758J} and \\citet{1999ApJ...515..669S} (see also references in Hollenbach, Yorke \\& Johnstone 2000), and have been much improved in the past years \\citep{2001MNRAS.328..485C,2003ApJ...582..893M,2004RMxAC..22...38J,2005MNRAS.358..283A,2006MNRAS.369..229A,2008ApJ...688..398E,2009ApJ...690.1539G,2009ApJ...699L..35D}. Gravitational interactions are another important effect on the population of stars, discs, and planets in a cluster environment. An encounter between a circumstellar disc and a nearby passing star can lead to significant loss of mass and angular momentum from the disc. While such isolated encounters have been studied in a large variety \\citep{1993ApJ...408..337H,1993MNRAS.261..190C,1994ApJ...424..292O,1995ApJ...455..252H,1996MNRAS.278..303H,1997MNRAS.287..148H,2004ApJ...602..356P,2005ApJ...629..526P,2006ApJ...653..437M,2007ApJ...656..275M,2008A&A...487..671K}, only few numerical studies have investigated the effect of stellar encounters on circumstellar discs in a dense cluster environment directly \\citep{2001MNRAS.325..449S,2006ApJ...641..504A}.  Only recently has it been shown from numerical simulations 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 massive stars in the centre of such a stellar cluster act as gravitational foci for the lower mass stars \\citep{2006A&A...454..811P}. Discs are most affected when the masses of the stars involved in an encounter are unequal \\citep{2006ApJ...642.1140O,2007ApJ...656..275M}, so it is the massive stars that dominate the encounter-induced disc-mass loss in young dense clusters \\citep{2006ApJ...642.1140O}. Observational evidence for this effect has been found by \\cite{2008A&A...488..191O}.  The numerical results obtained in our previous investigations are based on a dynamical model of the Orion Nebula Cluster (ONC) -- one of the observationally most intensively studied young star cluster. It was demonstrated that in the ONC stellar encounters can have a significant impact on the evolution of the young stars and their surrounding discs \\citep{2006ApJ...652L.129P,2006ApJ...642.1140O,2006A&A...454..811P,2007A&A...462..193P,2008A&A...488..191O,2008A&A...487L..45P}. However, investigating one model star cluster is not sufficient to draw general conclusions -- in fact, one could not answer questions as: How would things change in a \\emph{denser} cluster? Would a higher density inevitably imply that stellar encounters play a more important role in the star and planet formation process? And what would be the situation in more \\emph{massive} clusters? Would the larger number of stars -- in particular \\emph{massive} stars -- play a role?  Conclusive answers to these questions demand further numerical investigations covering a larger parameter space in cluster parameters. Here we realise this by modelling scaled versions of the standard ONC model -- clusters with varying stellar numbers and sizes. The focus of this investigation is on the encounter-induced disc-mass loss. 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 Section~\\ref{sec:observations_onc} we briefly outline the observationally determined basic properties of the ONC, that serves as a reference model for the other cluster models. The computational method is described in Section~\\ref{sec:numerical_method} and the construction of the numerical models is outlined in Section~\\ref{sec:numerical_models}. Afterwards we present results from a numerical approach to this problem in Section~\\ref{sec:numerical_results} and compare with analytical estimates in Section~\\ref{sec:analytical_results}. The conclusion and discussion mark the last section of this paper. ",
        "conclusions": "\\label{sec:conclusions}  The influence of different cluster environments on the encounter-induced disc-mass loss has been investigated by scaling the size, density and stellar number of the basic dynamical model of the ONC. The findings can be summarized as follows: \\begin{enumerate} \\item The disc-mass loss increases with cluster density but remains rather unaffected by the size of the stellar population. \\item The density of the ONC itself marks a threshold: \\begin{enumerate} \\item in less dense and less massive clusters it is the massive stars that dominate the encounter-induced disc-mass loss via gravitational focusing of low-mass stars whereas \\item in denser and more massive clusters the interactions of low-mass stars with equal-mass perturbers play the major role for the removal of disc mass. \\end{enumerate}  \\item In clusters four times sparser than the ONC the effect of encounters is still apparent. \\end{enumerate}  These findings from numerical simulations are well confirmed via observations. It is widely known that the disc frequency of even very young clusters (e.g. NGC~2024) is significantly below 100\\,\\% \\citep{2001ApJ...553L.153H,2005astro.ph.11083H}. This implies a very short time-scale for the disc destruction of a part of the cluster population. In fact, stellar interactions are a potential candidate for a very rapid physical process and thus the encounter-induced disc-mass loss is potentially a vital mechanism for disc destruction at the earliest stages of cluster evolution. Observational evidence for the role of this mechanism has also been provided for the $\\sim$1\\,Myr old Orion Nebula Cluster \\citep{2008A&A...488..191O}. Moreover, observations of NGC~2024, ONC, and NGC~3603 confirm also the correlation of decreasing cluster disc fraction (CDF) with increasing cluster density and even provide evidence for a critical density $\\rho_{\\mathrm{crit}} \\approx \\rho_{\\mathrm{ONC}}$. A compilation of observational data shows that these similar aged ($\\sim$1\\,Myr) clusters with peak densities of roughly $10^3\\,\\mathrm{pc}^{-3}$, $10^4\\,\\mathrm{pc}^{-3}$, and $10^5\\,\\mathrm{pc}^{-3}$, respectively, show CDFs of $85\\,\\%$, $80\\,\\%$, and $40\\,\\%$ \\citep{2000AJ....120.1396H,2000AJ....120.3162L,2004AJ....128..765S}. The fact that NGC~2024 is sparser than the ONC but has a similar CDF is in good agreement with our simulations.  \\begin{figure} \\centering \\includegraphics[width=1.0\\linewidth]{figures/ONC_range.pdf} \\caption{Cluster density as a function of cluster size for clusters more massive than $10^3\\,\\Msun$ and embedded clusters with more than 200 observed members from \\cite{2009A&A...498L..37P}. The parameter space covered by the simulations in this work is indicated by the large pink cross centered on the ONC.} \\label{fig:cluster_sequences__simulated_parameter_space} \\end{figure}  Our results have several important implications for the general picture of star and cluster formation. Very recently, \\cite{2009A&A...498L..37P} has shown that there exist two cluster sequences evolving in time along pre-defined tracks in the density-radius plane, the ``leaky'' and the ``star-burst'' clusters. The simulations performed in the present investigation cover the parameter space of the ``leaky'' clusters in their embedded stage (see Fig.~\\ref{fig:cluster_sequences__simulated_parameter_space}). Comparison with our findings shows that at the earliest evolutionary stage leaky clusters have densities above the critical density. Hence in leaky clusters star-disc systems are initially efficiently destroyed via encounters that occur preferentially between low-mass stars. The ONC corresponds to an intermediate stage in the embedded phase of leaky clusters with the transition towards a preferred disc-mass loss of high-mass stars via gravitational focusing of low-mass stars. At the final stage of the embedded phase the encounter-induced disc-mass loss in leaky clusters ceases. Gravitational focusing by high-mass stars may still affect single objects yet at this age is most probably exceeded by other disc-destructive processes like photoevaporation or planet formation. The effects in star-burst clusters would be similar yet much more pronounced. In case of the Arches cluster one could expect stellar encounters to destroy the discs of most of the low- and high-mass stars in several hundred thousand years. Combining our results with the finding of \\cite{2003ARA&A..41...57L} that most stars are born in clusters, it becomes evident that a significant fraction of all stars must have been affected by stellar encounters at the early evolutionary stages of their hosting environment.  The application of our results to the dynamics of embedded clusters -- though obtained from simulations that do not contain gas -- is justified for three reasons: \\begin{enumerate} \\item Rather than simulating the \\emph{evolution} of leaky clusters (which would explicitly require the treatment of gas), we use our cluster models to map certain \\emph{evolutionary stages} of the sequence of leaky clusters in terms of \"dynamical snapshots\". The dynamical effects of these numerical models are then used to estimate the effect of encounters in the observed clusters at their current dynamical state. \\item The effect of gas in an embedded cluster is to lower the frequency of close encounters (due to the smoother cluster potential), yet unless the gas mass is dominating cluster dynamics - as is not the case for the only partially embedded leaky clusters shown in Fig.~\\ref{fig:cluster_sequences__simulated_parameter_space} (e.g. NGC~2024) - the effect is minor. \\item Gas expulsion causes the clusters to expand and thus their density to decrease much faster than in our simulations. Consequently, when mapping our results to the current dynamical state, we \\emph{underestimate} the initial cluster density and thus the effect of encounters in the early evolutionary phases. Hence, our results are least accurate for older clusters, yet well applicable for the young evolutionary stages were encounters have the largest effect. \\end{enumerate}   The effect of the encounter-induced disc-mass loss in the early evolution of stellar systems has important implications for the formation and evolution of planets. From our findings we conclude that the probability to find planets around intermediate mass stars is highest. While the high initial densities of leaky clusters imply that planets around low-mass stars are expected to be less frequent, hardly any are expected in star-burst clusters. Independent of the cluster environment, planet formation around high-mass stars seems to be completely hindered via interactions with other cluster members."
    },
    "0911/0911.2069_arXiv.txt": {
        "abstract": "We study the effect on the primordial cosmological perturbations of a sharp transition from inflationary to a radiation and matter dominated epoch respectively. We assume that the perturbations are generated by the vacuum fluctuations of a scalar field slowly rolling down its potential, and that the transition into the subsequent epoch takes place much faster than a Hubble time. The behaviour of the superhorizon perturbations corresponding to cosmological scales in this case is well known. However, it is not clear how perturbations on scales of and smaller than the Hubble horizon scale at the end of inflation may evolve through such a transition. We derive the evolution equation for the gravitational potential $\\Psi$, which allows us to study  the evolution of the perturbations on all scales under these circumstances. We show that for a certain range of scales inside the horizon at the end of inflation, the amplitude of the perturbations are enhanced relative to the superhorizon scales. This enhancement may lead to the overproduction of Primordial Black Holes (PBHs), and therefore constrain the dynamics of the transitions that take place at the end of inflation. ",
        "introduction": "Inflation has become the main paradigm to understand the presence of small inhomogeneities during the early universe. The observed Cosmic Microwave Background (CMB) and Large Scale Structure (LSS) spectrum of primordial perturbations, reveals that simple inflationary models are consistent with the analysis of the current data. Nevertheless, CMB and LSS observations only probe primordial perturbations on scales that leave the horizon long before the end of inflation, and possible cosmological implications of the perturbations on subhorizon scales have not been studied much.  In this paper, we investigate the behavior of curvature perturbations on scales of and smaller than the horizon size at the end of inflation. We assume standard slow-roll inflation and that the perturbations are generated from the vacuum fluctuations of the inflaton field. It is expected that the universe enters either a stage of radiation-domination or matter-domination right after inflation, depending on the efficiency of preheating or reheating, and the transition takes place within a timescale shorter than the Hubble time. We therefore model this transition as a change of the equation of state from $w\\approx -1$ to $w=1/3$ or to $w=0$ within a time scale $\\Delta t\\ll H_e^{-1}$ where $H_e$ is the Hubble parameter at the end of inflation.  On large scales where the wavelengths are much larger than $\\Delta t$, we may regard the transition as instantaneous, occuring on a given spacelike hypersurface. Then we can match the perturbations between inflation and the succeeding epoch by requiring that the intrinsic and extrinsic curvatures are continuous on the hypersurface at which we do the matching~\\cite{Israel:1966rt,Deruelle:1995kd,Martin:1997zd}. It is natural to assume that the end of inflation occurs when the inflaton field reaches a critical value, $\\phi_c(x,\\tau)={\\rm constant}$. This implies that the transition occurs on a comoving surface on which $\\phi$ is spatially homogeneous. Then one can show that we have the following junction conditions, \\bea \\[\\calr_{\\bm k}\\]^+_-&=&0 \\label{junc1}\\,, \\\\ \\[\\Psi_{\\bm k}\\]^+_-&=&0 \\label{junc2}\\,, \\eea where $\\calr$ is the curvature perturbation on comoving hypersurfaces, and $\\Psi$ is the generalized gravitational potential. The junction conditions give the initial conditions for the following radiation- or matter-dominated stage. On superhorizon scales, the potential $\\Psi$ then reaches a constant value, $\\Psi\\sim\\calr$, a few Hubble times after the end of inflation~\\cite{Deruelle:1995kd,Martin:1997zd,Lyth:2005ze}.  What we are interested in here is perturbations on scales of or smaller than the horizon scale. On subhorizons scales, but still larger than the scale $\\Delta t$, the resulting perturbations are  the same as in the case of superhorizon perturbations if the universe is matter-dominated after the transition. On the other hand, the result is not so simple if the universe is radiation-dominated after the transion. In Refs.~\\cite{Lyth:2005ze, Zaballa:2006kh}, we investigated this case, and argued there that the relative large amplitude of the potential $\\Psi$ can be achieved on subhorizon scales and may lead to the overproduction of PBHs at the end of inflation. The constraint on the amplitude of the curvature perturbation is then tighter than the current bound from the PBHs formed throughout the radiation epoch \\cite{Zaballa:2006kh} (for the constraints on the \"standard\" formation of PBHs see~\\cite{Carr:1993,Carr:1994ar,Kim:1996hr,Green:1997sz, Leach:2000ea,Bugaev:2006fe,Kohri:2007qn,Peiris:2008be,Josan:2009qn, Bugaev:2008gw,Alabidi:2009bk}).  In this paper we investigate the effect of the transition on curvature perturbations on all subhorizon scales, including those on sufficiently small scales on which the transition cannot be regarded instantaneous any more. The paper is organized as follows. In Sec.~(\\ref{sec:ginv}) we derive the governing equation for the evolution of the potential $\\Psi$ for an arbitrary equation of state and velocity of sound. We review the generation of the perturbations during slow-roll inflation in Sec.~(\\ref{sec:inflation}). We study the evolution of $\\Psi$ through a transition into radiation and matter domination in Secs.~(\\ref{sec:rad}) and (\\ref{sec:mat}) before concluding. In Appendices \\ref{appA} and \\ref{appB}, we present analytical approximations that we have employed to reproduce and understand certain features that we have found in our numerical estimations.  Thoughout the paper, we use the Planck units $8\\pi G=M_P^{-2}=1$. ",
        "conclusions": "We have studied the behaviour of the primordial perturbations through a sharp transition from inflation into the radiation and matter dominated epochs. We have found that for transitions that occur much faster than a Hubble time, there is a range of scales inside the horizon for which the amplitude of the primordial perturbations is enhanced, relative to the amplitude of perturbations that exit the horizon a few efolds before the end of inflation, if slow-roll conditions hold during the final stages of inflation. For a Gaussian distribution of the perturbations, this relatively large amplitude increase the probability of strong perturbations on these small scales, that may  lead to a significant production of PBHs.    For a transition from inflation into a radiation dominated universe, the maximum amplitude of the peculiar gravitational potential is about $\\Psi\\sim2\\sqrt{3}H^2_{\\rm e}/\\dot\\phi_{\\rm e}$. Assuming a Gaussian distribution for the primordial perturbations, one can estimate the PBH abundance produced at the end of inflation, using for example Press-Schechter approach. (See however \\cite{Hidalgo:2007vk} for the effect of non-Gaussian perturbations.) Given that the observed amplitude of the curvature perturbation on cosmological scales is roughly given by $\\calr\\sim10^{-5}$, it turns out that unless the spectrum of $\\calr$ is considerably larger at the end of inflation, the amount of PBH production would not have any cosmological significance \\cite{Zaballa:2006kh}. However, a much larger amplitude on scales $k\\sim\\calh_{\\rm e}$, is still compatible with current observational data even within the slow-roll paradigm \\cite{Kohri:2007qn} . Futhermore, PBHs may be overproduced as well with a more complicated scale dependence if slow-roll conditions do not hold throughout the whole period of inflation.  In the case of a transition into a matter dominated universe the production of PBHs could be more dramatic even with a very flat spectrum of $\\calr$ on all scales. The maximum amplitude for $\\Psi$ now depends on delta as $\\Psi\\propto 1/\\delta$. Therefore if the velocity of sound does vanish on a very small scale, such that $c_s=0$ for $k\\sim\\calh_{\\rm e}/\\delta$, there may be a significant production of small mass PBHs even if the transition time is not very rapid.  Finally we would like to stress that a detailed analysis of specific models, in which the transition from inflation occurs very rapidly, may reveal the overproduction of PBHs, placing constraints on the spectrum of the primordial perturbations on scales that are too small for conventional observations."
    },
    "0911/0911.4479_arXiv.txt": {
        "abstract": "We have undertaken a Parkes ammonia spectral line study, in the lowest two inversion transitions, of southern massive star formation regions, including young massive candidate protostars, with the aim of characterising the earliest stages of massive star formation. 138 sources from the submillimetre continuum emission studies of Hill \\etal\\, were found to have robust (1,1) detections, including two sources with two velocity components, and 102 in the (2,2) transition.  We determine the ammonia line properties of the sources: linewidth, flux density, kinetic temperature, \\nh\\, column density and opacity, and revisit our SED modelling procedure to derive the mass for 52 of the sources. By combining the continuum emission information with ammonia observations we substantially constrain the physical properties of the high-mass clumps. There is clear complementarity between ammonia and continuum observations for derivations of physical parameters.  The MM-only class, identified in the continuum studies of Hill \\etal, display smaller sizes, mass and velocity dispersion and/or turbulence than star-forming clumps, suggesting a quiescent prestellar stage and/or the formation of less massive stars. ",
        "introduction": "\\label{sec:intro}  Massive stars are dynamical and enigmatic powerhouses that shape and drive both their local stellar neighbourhood and the ecology and evolution of their host galaxy. Despite this heavy influence, the formation and evolution of a massive star is not well understood. Particularly perplexing are the earliest evolutionary stages of massive star formation (MSF). The difficulty lies in the unambiguous detection, identification and characterisation of these stages, snapshots of which are difficult to obtain as a result of the general rarity of candidates and the rapidity of their evolution \\citep{garay99}.  Whether or not massive stars are scaled-up analogs of low-mass stars is still uncertain. Massive stars exert considerable radiative pressure on the surrounding dust and gas, which in principle could halt and reverse spherical infall of the collapsing protostar. Therefore, a `simple' scaled-up version of low mass star formation is insufficient for massive stars. There are a number of different approaches to address this dilemma as outlined in the reviews by \\citet{evans02, menten05, zinnecker07, beuther07}.  As their sophistication increases numerical simulations in two and three dimensions show that the radiation pressure problems associated with spherical symmetry can be overcome \\citep[e.g.][]{krumholz09}. It is still not clear however, whether the bulk of the mass that finally ends up on the massive star comes from the monolithic collapse of a single dense core \\citep[e.g.][]{mckee03} or is accreted from surrounding lower density gas which is funneled to the centre of a cluster's gravitational potential where the most massive stars are forming \\citep[e.g][]{bonnell98}.  The natal molecular cloud, from which high-mass stars are formed, is expected to be dense, massive and cold, detectable only at (sub)millimetre wavelengths. Within the molecular cloud, core collapse will be triggered. As the collapsing protostar gains mass, the gravitational energy will serve to heat the core and ionise the surrounding material, causing an increase in the luminosity of the core. In terms of parameter evolution, the initial protostar will be massive, cool and of low luminosity. As the core evolves, it will accumulate more mass, and the temperature and luminosity of the core will increase. This evolution is seen in low mass protostars \\citep{evans02}.  In a search for cold cores that would mark the earliest stages of massive star formation \\citet{hill05} undertook a (1.2) millimetre SIMBA\\footnote{The SEST IMaging Bolometer Array (SIMBA) was a 37 channel hexagonal array on the Swedish ESO Submillimetre Telescope (SEST).} continuum emission study of sources exhibiting signs of methanol maser and/or radio continuum emission - both of which have previously proven successful tracers of the earliest stages of massive star formation \\citep[e.g.][]{minier00, williams04, pestalozzi05}. This SIMBA survey revealed a large number of millimetre continuum sources (255) devoid of the maser and \\uchii\\, sources targeted. Subsequent submillimetre observations of these sources, which were dubbed `MM-only cores', unveiled submillimetre equivalents for each - often revealing the multiplicity of individual MM-only sources, and confirmed their association with cold, deeply embedded objects \\citep{hill06}.  The aforementioned SIMBA survey detected a total of 405 millimetre continuum sources, which (based on their star formation association) could be broken into four classes of source. \\citet{hill05} proposed that each of these classes of source could represent a different phase of massive star evolution, with the MM-only class a possible example of the very earliest stages of massive star formation.  A caveat however, is whether the MM-only cores are currently undergoing or will/can support massive star formation.  In order to determine the characteristic properties of the sources in the SIMBA sample, in particular the previously unknown and unstudied MM-only cores, \\citet{hill09} performed spectral energy distribution (SED) modelling of a significant fraction of the sample.  SED diagrams are useful tools from which physical quantities such as the luminosity, mass and temperature can be derived \\citep[cf.][]{whitney03, robitaille06, hill09}. If the observational data are well constrained \\citep[cf.][]{minier05}, SED modelling can provide useful estimates of each of these parameters, for each star-forming core, which allows estimation of  the evolutionary phase of a young massive star.  The luminosity, mass and temperature of an astronomical source are fundamental physical attributes that may be used to clarify and characterise the nature of a source, possibly providing insight into their evolutionary status similarly as they do for low mass stars \\citep{andre00}. Indeed, the temperature of a source places important physical constraints on the chemical composition of the core, including which chemical species are present in the core, as well as the size and types of grains present in the central star-forming cores.  The method employed for the SED fitting was that of Bayesian inference, which enabled a statistically probable range of suitable values for the luminosity, mass and temperature, for each source modelled. As SED modelling is heavily reliant upon robust data, it was not possible to usefully constrain each of these parameters for a thorough assessment of massive star formation scenarios. In the absence of reliable far-infrared data, which would serve to constrain the peak of the SED and thus parameters resultant from the fit, additional means of constraining the source parameters, such as temperature, are necessary.  Ammonia is an excellent molecular cloud thermometer \\citep{menten05} from which the rotational temperature and ultimately the kinetic temperature of the source can be determined. Ammonia is readily detectable in quiescent dark clouds and regions of low luminosity \\citep{ho83} making it perfectly suited to study the birthplace of massive stars. The ratio of the hyperfine components of ammonia's signature five-fingered spectrum provides the optical depth information of the molecular cloud environment. As \\nh\\, is a particularly good probe of high density gas, and is more resilient to the effects of depletion than other high density tracers (e.g. CS), ammonia observations also provide information pertaining to the density of the core \\citep[cf.][and references therein]{longmore07}.  We have undertaken an ammonia spectral line survey of sources from the SIMBA sample of \\citet{hill05} in order to obtain an independent determination of their temperature and simultaneously examine the robustness of our SED fitting \\citep{hill09}. In addition to facilitating and constraining further SED fitting, these ammonia observations provide information pertaining to the molecular abundance of ammonia in the cores, as well as the physical parameters such as linewidth, rest velocity, column density and virial mass. With this information, we seek to address formation scenarios for massive stars and identify the star formation status/phase of the MM-only core. ",
        "conclusions": "We have undertaken an ammonia molecular line study, in the lowest two inversion transitions, of young massive star formation regions in the southern hemisphere as part of an effort to characterise the earliest stages of high-mass star formation. The sample targeted was derived from the millimetre continuum emission studies of \\citet{hill05, hill06} and included sources with and without signatures of high mass star formation. In total, 244 sources were observed in both ammonia transitions, using the Parkes radio telescope. Of these, 138 had detections ($>$3-$\\sigma$) in \\nhone, including two sources with two velocity components, and 102 in \\nhtwo. The MM-only sources are not more or less likely to be detected in either transition than the sources with a methanol maser and/or radio continuum association.  The spectral line properties of the sources have been determined from gaussian fits to the lines: linewidth, flux density, opacity. In addition, physical parameters, such as the rotational and kinetic temperatures as well as the column density, were derived from the spectra. From the kinetic temperature, and revisiting previous SED modelling techniques \\citep{hill09}, we have determined the mass and luminosity of the sources in the sample. We have used the Bayesian inference method of SED modelling, which provides robust estimates of the parameters as well as good estimates of the uncertainties associated with each parameter, further allowing robust statistical conclusions.  Combining continuum and line observations has proven to be quite powerful. In our case, ammonia observations and SED modelling are very complementary in terms of parameter determinations. It is clear, that the kinetic temperature, as derived from ammonia, used in combination with SED modelling has constrained the range of validity for the temperature, which in turn has constrained the range of validity for the mass and luminosity for each of the sources in our sample.   The MM-only sources, those with no overt signs of massive star formation, have smaller \\nhone\\, and (2,2) linewidths on average compared with sources associated with a methanol maser and/or radio continuum source. They are also less massive and smaller on average but reach the same upper values for both parameters. Because the MM-only sources have  similar brightness/flux, column densities and similar temperatures as star forming sources, they have the potential to form stars. The least massive MM-only sources, cannot form high mass stars and will likely proceed to form intermediate mass stars, whilst the more massive clumps, are more interestingly  strong candidates for early stage massive protostars. The different hypotheses will result in different internal structures.  Higher spatial resolution observations with ALMA will allow resolution of the internal structures of the MM-only cores, which may then allow us to distinguish between starless cores that will begat massive stars and those that are transient.   \\vspace{-0.2cm}"
    },
    "0911/0911.3545_arXiv.txt": {
        "abstract": "{Superluminous type Ia supernovae (SNe Ia) may be explained by super-Chandrasekhar-mass explosions of rapidly rotating white dwarfs (WDs). In a preceding paper, we showed that the deflagration scenario applied to rapidly rotating WDs generates explosions that cannot explain the majority of SNe Ia.} {Rotation of the progenitor star allows super-Chandrasekhar mass WDs to form that have a shallower density stratification. We use simple estimates of the production of intermediate and iron group elements in pure detonations of rapidly rotating WDs to assess their viability in explaining rare SNe Ia.} {We numerically construct WDs in hydrostatic equilibrium that rotate according to a variety of rotation laws. The explosion products are estimated by considering the density stratification and by evaluating the result of hydrodynamics simulations.} {We show that a significant amount of intermediate mass elements is produced for theoretically motivated rotation laws, even for prompt detonations of WDs.} {Rapidly rotating WDs that detonate may provide an explanation of rare superluminous SNe Ia in terms of both burning species and explosion kinematics.}   \\keywords {Stars: supernovae: general -- Hydrodynamics -- Methods: numerical} ",
        "introduction": "\\label{intro}  Type Ia supernovae (SNe Ia) were long believed to form a relatively homogeneous class of events in terms of their spectra, peak luminosities, and light curves. Surveys, however, have shown that the distribution of SNe Ia properties is substantially broader than previously anticipated, including those of highly peculiar events such as SN 2007ax \\citep{Kasliwal_etal}, SN 2005hk \\citep{Stanishev_etal}, and SN 2003fg \\citep{Howell_etal}. The latter, in particular, has been interpreted in terms of a super-Chandrasekhar-mass explosion of a rapidly rotating white dwarf (WD) \\citep{Jeffery_etal}.  Motivated by the goal of explaining the presumed homogeneity of SNe Ia within a single explosion scenario, most multi-dimensional hydrodynamical simulations so far have focused on turbulent deflagrations \\citep{2006A&A...446..627S,snob} or delayed detonations \\citep{GamKhok05,RoepkeNiemeyer07} of Chandrasekhar-mass WDs (this class includes the gravitationally confined detonation model proposed by \\citet{GCDmodel}). The delayed detonation model yields an explosion consistent with kinematics as well as spectra, light curves, and nucleosynthesis of type Ia supernovae. So far, a major disadvantage of this model has been that the conditions for a deflagration-to-detonation transition (DDT) as proposed by \\citet{1991A&A...245..114K} and \\citet{WW94}, are questionable in the context of WD matter \\citep{1999ApJ...523L..57N}. Although \\citet{2008ApJ...681..470P} proposed a theoretical explanation of delayed detonations, the numerical studies by \\citet{WoosKer09} and \\citet{SchmCir09} nevertheless impose severe constraints on DDTs. Prompt detonations, on the other hand, have been considered to be infeasible for explaining SNe Ia since the pioneering work of \\citet{Arn69}, as too little material is burned at sufficiently low densities to produce intermediate-mass elements (IMEs).  The opposite problem occurs in the case of rapidly rotating WDs, which burn as turbulent deflagrations (\\citet{PaperI}, hereafter referred to as Paper~I). Here, too much stellar material remains unburned as a consequence of the lower density at large radii. Combining the fast propagation of a detonation front with the shallower stratification of the stellar fuel, rapidly rotating initial models have already been proposed as a possible means of repairing the prompt detonation mechanism by \\citet{1992A&A...254..177S}. However, in their models the bulk of WD matter was nevertheless burnt to form iron-group elements (IGEs) and, accordingly, even rapid rotation did not change the situation. Whenever a significant amount of IMEs was produced by rotation in their study, the amount of IGEs was simultaneously too high.  In this work, we revisit the scenario of promptly detonating carbon-oxygen (CO) WDs. In contrast to \\citet{1992A&A...254..177S}, we investigate a variety of rotation laws including differentially rotating models inspired by the results of \\citet{2005A&A...435..967Y}. Some of the models in our sample significantly exceed the Chandrasekhar mass, which is indeed suggested by some observations, for instance, \\citet{Howell_etal}. For this reason, we do not attempt to interpret our results as those for normal SNe Ia but instead look for possible connections with observed peculiar supernovae.  It is important to emphasize that we only attempt to predict the expected range of masses of the produced IMEs and IGEs, to determine whether these types of explosions might account for rare SN Ia events. To this end, we first provide two simple estimates based on the equilibrium stratification of the WD and different choices of the density threshold for burning into nuclear statistical equilibrium (NSE). These numbers can be interpreted as plausible upper and lower bounds on the produced IME masses in these events. This interpretation is confirmed by the nucleosynthetic yields obtained by post-processing a full explosion simulation, which lie in between these bounds. Going beyond these estimates would require a parameter study of rotation laws and ignition conditions of sufficient resolution to capture the detonation front. Given the speculative nature of this model, we do not think that this effort is warranted at present.  Based on these simplifying assumptions, we find that a significant amount of IMEs (0.1 to 0.4~$M_{\\sun}$) is being produced accompanied by a Super-Chandrasekhar mass amount of IGEs (1.5 to 1.8~$M_{\\sun}$) and a marginal amount of unburnt stellar material. The ejecta expand at higher radial velocities because of the greater amount of nuclear energy released as a consequence of the detonation. Unburnt stellar material is leftover only in stellar regions close to the stellar surface at the equatorial plane, whereas iron group elements are predominant within the star and at the stellar poles as a consequence of the density stratification that is shaped by rapid rotation. For the same reason, a bulge of IMEs is generated in the outer regions of the equatorial plane.  The discussion will proceed as follows. In Sect. 2, we introduce our numerical method. The procedure that allows us to estimate the composition of burning species is discussed in Sect. 3. The process of a typical prompt detonation initiated within a rapidly rotating WD is described in Sect. 4. Section 5 contains an interpretation with respect to spectral features. The detailed stellar composition is investigated in Sect. 6. Section 7 concludes the paper. ",
        "conclusions": ""
    },
    "0911/0911.3897_arXiv.txt": {
        "abstract": "Previous studies of the \\nad interstellar absorption line doublet have shown that galactic winds occur in most galaxies with high infrared luminosities.  However, in infrared-bright composite systems where a starburst coexists with an active galactic nucleus (AGN), it is unclear whether the starburst, the AGN, or both are driving the outflows.  The present paper describes the results from a search for outflows in 35 infrared-faint Seyferts with 10$^{9.9}$ $<$ \\lir/\\lsun $<$ 10$^{11}$, or, equivalently, star formation rates (SFR) of $\\sim$0.4 -- 9 \\smpy, to attempt to isolate the source of the outflow.  We find that the outflow detection rates for the infrared-faint Seyfert 1s (6\\%) and Seyfert 2s (18\\%) are lower than previously reported for infrared-luminous Seyfert 1s (50\\%) and Seyfert 2s (45\\%). The outflow kinematics of infrared-faint and infrared-bright Seyfert 2 galaxies resemble those of starburst galaxies, while the outflow velocities in Seyfert 1 galaxies are significantly larger.  Taken together, these results suggest that the AGN does not play a significant role in driving the outflows in most infrared-faint and infrared-bright systems, except the high-velocity outflows seen in Seyfert 1 galaxies.  Another striking result of this study is the high rate of detection of inflows in infrared-faint galaxies (39\\% of Seyfert 1s, 35\\% of Seyfert 2s), significantly larger than in infrared-luminous Seyferts (15\\%).  This inflow may be contributing to the feeding of the AGN in these galaxies, and potentially provides more than enough material to power the observed nuclear activity over typical AGN lifetimes. ",
        "introduction": "\\label{intro}  Galactic-scale gas outflows appear to play an important role in the evolution of the universe, possibly providing an explanation for a number of cosmological questions (\\citealt{vcb05} and references therein).  These galactic outflows are frequently observed in local galaxies with high star formation rates and/or an active nucleus, often extending on a scale of order of a few kiloparsecs or larger.  The high frequency of outflows measured in high-redshift, actively star-forming objects such as Lyman break galaxies suggests that outflows are a common stage in galactic evolution, with wind velocities decreasing over time \\citep{ss03,fr06}.  Such outflows may be responsible for the deficit of baryons seen in our own Milky Way galaxy and the mass-metallicity relation in external galaxies, if winds are capable of preferentially ejecting metals into the intergalactic medium (\\citealt{l74,g02,thk04}). The intergalactic medium can be heated as well as enriched with metals by these galactic outflows, since energy inputs into the IGM on the order of 10$^{56}$ ergs have been measured from galaxies with strong outflows \\citep{chk08}.  Additionally, outflows may quench star formation by heating up cold gas and ejecting it from the host \\citep{b04,ssb05}.  Specifically, AGN-powered outflows have been proposed as the cause of the drop in AGN luminosity at low redshift and the cutoff at high luminosity of the galaxy luminosity function \\citep{sp99,c00,tsc06}.  These winds may also limit black hole and spheroid growth, and be responsible for the observed tight correlation between black hole mass and galactic spheroid mass \\citep{fm00,md02,gzs04}. Finally, winds driven by AGN and/or starbursts may help remove enough nuclear material with low angular momentum early on in the evolution of gas-rich systems to aid in the formation of large disks with high specific angular momentum, more in line with the observations \\citep{bgs01,kk04}.  Studies performed on actively star-forming galaxies, near and far, have revealed that outflow energetics increase with infrared (8-1000 $\\mu m$) luminosity, \\lir, or equivalently, star formation rate (\\citealt{m05,rvs05a}, hereafter RVS05a; \\citealt{rvs05b}, hereafter RVS05b; \\citealt{smnkl09,w09}).  Outflows have also been detected in composite galaxies with both starbursts and AGN \\citep[e.g.,][]{is03,hs06}.  Seyfert 1 composites show significantly higher outflow velocities than pure starbursts and Seyfert 2 composites (\\citealt{rvs05c}, hereafter RVS05c).  Hydrodynamical simulations of starburst-driven outflows \\citep[e.g.,][]{tsc06,cbsbh08} have difficulties reproducing the very high velocities detected in Seyfert 1 composites.  The AGN in these objects thus appear to play a significant role in driving these outflows, at least on small scales \\citep{ckg03}.  The situation in Seyfert 2 composites is more ambiguous -- the AGN or starburst or both could be powering the outflow in these systems \\citep{sch85}.  The focus of this paper is on ``pure'' Seyfert galaxies with weak infrared starbursts.  Our sample consists of 35 galaxies that are faint in the infrared (10$^{9.9}$ $<$ \\lir /\\lsun~$<$~10$^{11.2}$), equally split between Seyfert 1s and Seyfert 2s.  We compare the results from our study with those from previous studies of starburst and Seyfert ultraluminous infrared galaxies (ULIRGs; \\lir/\\lsun~$\\ge$ 10$^{12}$) and luminous infrared galaxies (LIRGs; 10$^{11}$ $<$ \\lir/\\lsun~$<$ 10$^{12}$) to isolate the role of the AGN in powering these outflows.  As in previous studies, the IR-faint objects were observed in the \\nad $\\lambda\\lambda$5890,5896 doublet absorption feature.  This Na feature has a low ionization potential (5.1 eV), so it probes neutral gas, and it can be used to study the ISM due to its high interstellar abundance.  All objects in the sample have $z$ $<$ 0.05, and thus the \\nad feature is found in the optical.  Velocity components in \\nad absorption are unambiguous indicators of outflowing (blueshifted velocities) or inflowing (redshifted velocities) gas, as there must be a continuum source behind the absorber.  The organization of this paper is as follows.  The sample is described in Section \\ref{sample}, including methods of selection and sample properties in terms of redshifts, star formation rates, and spectral types. The observations are discussed in Section \\ref{obs}.  Section \\ref{analysis} describes the line fitting, the derivations of the velocities and column densities, and estimates of the expected stellar contributions to the measured \\nad absorption.  In Section \\ref{outsec} we present the results on outflows, starting with the Seyfert 2s and followed by the Seyfert 1s.  Next, Section \\ref{insec} describes the results on inflows in the same fashion as Section \\ref{outsec}.  This is followed by a discussion of the implications of our findings in Section \\ref{disc}, including comparisons with previous studies, dynamical estimates, and possible connections between inflows and nuclear structures.  We conclude in Section \\ref{conc} with a summary of our findings.  All calculations in this study assume present-day cosmological parameter values, with $H_0$ $=$ 75 \\kms~Mpc$^{-1}$, $\\Omega_m$ $=$ 0.3, and $\\Omega_\\Lambda$ $=$ 0.7.  All wavelengths quoted are vacuum wavelengths. ",
        "conclusions": "\\label{disc}  \\subsection{Outflows} \\label{outdisc}  \\subsubsection{Comparison with Previous Studies}  Since the purpose of this study is to determine whether starbursts or AGN are the primary mechanism behind galactic outflows, it is important to compare the results determined here to those of RVS05c for the ULIRG Seyferts -- galaxies with both starburst and AGN. As seen in Table~\\ref{avgpropout}, the detection rates for the IR-faint Seyfert 2s and Seyfert 1s are only 18\\% $\\pm$ 9\\% and 6\\% $\\pm$ 6\\%, respectively, but rates were as high as 45\\% $\\pm$ 11\\% and 50\\% $\\pm$ 20\\% for the IR-luminous Seyfert 2s and 1s (RVS05c).  Our detection rates are even lower in comparison to non-Seyfert ULIRGs, since rates of 80\\% $\\pm$ 7\\% and 46\\% $\\pm$ 13\\% were measured for low-$z$ and high-$z$ ULIRGs in RVS05b, and a rate of 83\\% was reported in Martin (2005).  Poststarburst galaxies have also been found to have high outflow detection rates on the order of 70\\% \\citep{tmd07}.  The measured outflow detection rates in the IR-faint Seyfert galaxies are therefore considerably smaller than those of pure starbursting galaxies, poststarbursts, and AGN + starburst composites.  Previous observations have indicated that the outflow detection rate increases with infrared luminosity (RVS05b, \\citealt{smnkl09}). Figure~\\ref{detrateplots} shows the outflow detection rates of the IR-faint Seyferts in our sample as well as the the IR-luminous Seyferts and starburst ULIRGs and LIRGs of RVS05b and RVS05c as a function of \\lfir .  The increase in outflow detection rate with \\lfir~is apparent, despite the limited range in \\lfir.  As \\lfir~is correlated with the star formation rate, this increase in the detection rate with \\lfir~suggests that star formation in all of these systems is the main driver of neutral outflows detected in \\nags.  Geometry could also be playing a role, since AGN winds with smaller opening angles (higher collimation) would reduce the detection rate (see Section \\ref{mme}).  Histograms showing the distributions of all (negative and positive) velocities in our current data as well as those of RVS05b and RVS05c provide another point of comparison (Figure~\\ref{velhistplots}).  The left panel compares the velocities of IR-faint Seyfert galaxy components from this study to the Seyfert and starburst ULIRGs \\& LIRGs of RVS05b and RVS05c.  We can safely rule out rotation as being a primary cause for the motion of the \\nad gas based on the fact that we do not observe a symmetry about zero velocity in these histograms. Were rotation the dominant gas motion, we should observe roughly equal amounts of blueshifted and redshifted gas in each object and in the overall distributions shown in Figure 8, assuming the \\nad gas is distributed more or less symmetrically around the center of rotation and is unaffected by severe differential obscuration.  Negative velocities are found in a much higher percentage of IR-luminous objects (Seyferts of RVS05c and starbursts of RVS05b) than IR-faint objects, again indicating that outflows are both stronger and more frequent on average in objects with high SFR.  The right panel in Figure~\\ref{velhistplots} combines together the IR-faint and IR-luminous Seyferts from this study and RVS05c and compares them to the starbursts of RVS05b.  In terms of the negative velocities, the results for Seyfert 1 and Seyfert 2 galaxies are quite different, with Seyfert 2 galaxies showing similar outflow percentages to those of the ULIRGs \\& LIRGs.  This suggests a physical connection between the mechanisms that drive outflows in Seyfert 2s and starbursting galaxies and a physical difference between the mechanisms that cause outflows in Seyfert 1s and Seyfert 2s. Interestingly, studies of {\\em ionized} gas outflows have not found such a dichotomy in velocities between Seyfert 1s and Seyfert 2s (e.g., \\citealt{rck01,vcb05,rck05,dck07}, and references therein).  This suggests that the neutral gas probed by our \\nad observations is not kinematically related to the ionized material of the NLR.  This difference is even more obvious when we also consider our results on positive (inflow) velocities (Section \\ref{incomp}).  Rigorous statistical tests confirm these results.  Kolmogorov-Smirnov (K-S) and Kuiper tests were both used, as the K-S test has an inherent bias in terms of differences in the mean and the Kuiper test does not. Low values reported by both of these tests can help rule out two sets of data having the same parent distribution. Results of these tests are listed in Table~\\ref{stattests}.  The small values of $P$(null) for $\\Delta v$ and \\dvmax~when comparing IR-faint and IR-luminous Seyferts suggest that they do not come from the same parent distribution.  A more significant comparison comes from combining together the results for the IR-faint Seyfert galaxies from this study and those for the IR-luminous Seyfert galaxies from RVS05c, and comparing them with those for the starburst ULIRGs and LIRGs from RVS05b.  The same statistical tests were performed on these distributions, first using all velocities, then the outflowing components only.  The results are listed in Tables~\\ref{stattests2a} and \\ref{stattests2b}, respectively.  The results in Table \\ref{stattests2a} indicate that Seyferts and starbursts do not share the same parent velocity distributions; we return to this point in Section \\ref{indisc}.  When considering only the outflowing components, all comparisons show low probability of originating from the same parent distribution except when the Seyfert 2 galaxies are compared with the starburst ULIRGs \\& LIRGs.  In that case, both the K-S and Kuiper tests return large $P$(null) values. This confirms quantitatively that the outflows in these two classes of objects may arise from the same physical process.  Figure~\\ref{dvmaxplots} displays plots showing \\dvmax~as a function of far-IR luminosity (correlated with star formation rate) and of galactic circular velocity (correlated with galactic mass) for Seyferts from this study, IR-luminous Seyferts from RVS05c, starburst LIRGs \\& ULIRGs from RVS05b, and four starburst dwarf galaxies from Schwartz \\& Martin (2004).  The four dwarfs were added in order to see if low-mass galaxies follow the same trends as high-mass systems.  The IR-faint Seyfert 2 galaxies seem to follow the same trends as the IR-bright galaxies, with \\dvmax~increasing with both SFR and galactic mass, while the outflow velocities of all Seyfert 1 galaxies lie above these trends.  This again suggests a fundamental difference in the way the winds in IR-faint/bright Seyfert 1 galaxies are powered compared with the winds in IR-faint/bright Seyfert 2 and starburst galaxies. We return to this issue in Section \\ref{mme}.  \\subsubsection{Outflow Dynamics} \\label{mme}  Another useful comparison between IR-faint and IR-luminous Seyferts is to look at the dynamical properties in addition to the kinematics. Using calculated covering fraction, column density, and velocities, we can estimate the mass, momentum, and kinetic energy of the neutral gas phase of the ISM being probed by those winds.  We follow the method outlined in the original study (RVS05b), which made the assumption that the outflows are spherically symmetric mass-conserving free winds, with a velocity and instantaneous mass outflow rate which do not depend on radius within the wind and which are zero outside the wind.  This method also assumed that the wind is a thin shell with a uniform radius of 5 kpc. This value was based on actual spatial measurements in some of the objects of RVS05c, but such spatial measurements are not available for our IR-faint Seyfert galaxies.  In the present study, we use a 5 kpc radius for both the Seyfert 2 and Seyfert 1 galaxies.  This is different from the radius of 10 pc that was used for the Seyfert 1 galaxies in the original study (RVS05c); the radius in IR-faint Seyfert 1s was chosen to be the same as for IR-faint Seyfert 2s to facilitate comparisons within the IR-faint sample, and to that end, the IR-luminous Seyfert 1 dynamical values from RVS05c have been scaled up to a 5 kpc radius as well.  These results should be considered order-of-magnitude estimates since they are based on a number of largely unproven assumptions.  We have not listed all results here; selected results can be seen in Figure~\\ref{dmdtplots}.  For the objects in the present study, we assume a modest value of 0.3 for the large-scale opening angle (\\co), which is the typical value for local disk winds \\citep{vcb05}.  Following the method of RVS05b, the covering fraction is used to parameterize the clumpiness of the wind, or it may reflect the global solid angle subtended by the wind when viewed from the galactic center.  Using the value of $\\langle$\\cf$\\rangle$ as listed in Table~\\ref{avgpropout}, this yields a global covering factor of $\\Omega$ $\\sim$ 0.1 for both Seyfert 1s and Seyfert 2s.  These values are not well constrained, especially for the Seyfert 1s since only one outflow was detected in this type of Seyfert.  The $\\Omega$ value is inconsistent with values of $\\Omega$ $\\sim$ 0.5-1.0 calculated by Crenshaw et al.~(2003b) from UV absorption lines in local Seyfert 1s.  This low $\\Omega$ value could imply that the winds we are seeing are collimated rather than wide-angle, and thus our lack of outflow detection in some of these objects may be due to their orientation relative to us rather than a complete lack of neutral gas outflow.  The influence of host galaxy inclination on column density was explored using extinction-corrected axis ratios from the de Vaucouleurs Third Reference Catalog \\citep{rc3}, but no conclusion could be drawn due to small-number statistics.  Additionally, these $\\Omega$ values will be low in comparison to values from other studies; there is contribution here from the background galaxy, and many of the outflows we have found are small-scale. Comparisons with the results of UV absorption line studies should be considered with caution, though, since UV absorbers are typically of much higher ionization (\\ion{N}{5}, \\ion{C}{4}) and likely located much closer to the AGN than \\nad absorbing material.  The \\nad absorbers may also be affected more strongly by host galaxy contamination than the UV absorbers.  Blueshifted absorption lines have also been detected in X-ray spectra for a number of Seyfert galaxies, including Mrk 6, with column densities of significantly higher order of magnitude than we have found here for \\nad in the optical (e.g., \\citealt{mew95,fbe99,mmwe01,knb03,vu08}).  Again, we caution comparing these results with our own, as it has been noted that these X-ray lines are much higher ionization (\\ion{O}{7}, \\ion{O}{8}) and are often intrinsically related to the aforementioned UV absorption lines.  One particularly interesting dynamical quantity is the mass outflow rate, since galactic outflows may contribute to the IGM enrichment and are possibly a quencher of star formation \\citep{tmd07}.  Plots of $dM/dt$ for the IR-faint Seyferts, IR-luminous Seyferts, and starbursting LIRGs and ULIRGS from RVS05b and RVS05c, and the dwarfs of Schwartz \\& Martin (all calculations are based on a absorber radius of 5 kpc) are presented in Figure~\\ref{dmdtplots}.  We see a general trend of mass outflow rate increasing with both \\lfir~and galactic mass.  There is a considerable difference in the mean $dM/dt$ rates between IR-luminous Seyfert 1s and IR-luminous Seyfert 2s.  The momentum and energy, as well as outflow rates for those quantities, are also significantly higher for IR-luminous Seyfert 1s than IR-luminous Seyfert 2s and all IR-faint Seyferts.  The overall good agreement in detection rate and kinematics between IR-faint and IR-bright Seyfert 2s and starburst ULIRGs \\& LIRGs (Figures~\\ref{detrateplots}-\\ref{dmdtplots}) suggests that the outflows in all of these objects are powered by star formation, while the marked differences when IR-faint Seyfert 1 galaxies are considered suggest that the AGN plays an important role in driving the (generally high velocity) outflows in Seyfert 1 galaxies.  Escape fractions were not calculated here since too few objects have outflow velocity components.  We can also look at energy outflow rates in order to determine what role these outflows could play in galactic feedback.  The rates that we have calculated can be found in Table~\\ref{outindiv}, though it must be cautioned that these numbers are a function of our uncertain $\\Omega$ value, as well as our assumed absorber radius of 5 kpc.  For the Seyfert 2 outflows, the average energy outflow rate was found to be $\\sim$ 10$^{41.1}$ \\omone \\rfive ergs s$^{-1}$, or $\\sim$ 10$^{7}$ \\omone \\rfive \\lsun.  In comparison to the average bolometric luminosity for these objects ($\\sim$ 10$^{10}$ \\lsun, taken from \\citealt{wu02}), the energy outflow rates are only $\\sim$1\\% of the host galaxy luminosity.  This indicates that outflow energetics in our Seyfert 2 galaxies are not strong enough to play a large role in galactic feedback.  The findings of Schlesinger et al.~(2009) for Mrk 573 are in agreement with our conclusion. For the only Seyfert 1 galaxy that we have found to have an outflow, Mrk 6, the energy outflow rate of its lower velocity component (10$^{41.6}$ \\omone \\rfive ergs s$^{-1}$) is again $\\sim$1\\% of the host bolometric luminosity.  This indicates that this Seyfert 1 outflow is also not energetic enough to play a vital role in galactic feedback, and thus previous findings are consistent with our results \\citep{kne07,smkn09}.  If we look at the higher velocity outflow found in Mrk 6, its energy outflow rate (10$^{42.9}$ \\omone \\rfive ergs s$^{-1}$) is $\\sim$5\\% of the host bolometric luminosity, which is higher than for the Seyfert 2 galaxies but again likely not strong enough to influence the evolution of its environment.  However, one should note that these values are all highly dependent on the calculated value of $\\Omega$ and assumed value of $r$ and are thus uncertain in comparison to our measured velocities.  \\subsection{Inflows} \\label{indisc}  \\subsubsection{Comparison with Previous Studies} \\label{incomp}  An unexpected result of the present study is the high detection rate of inflows in IR-faint Seyferts (39\\% $\\pm$ 12\\% for Seyfert 1s, 35\\% $\\pm$ 11\\% for Seyfert 2s). In contrast, only $\\sim$15\\% of the IR-luminous objects in RVS05b and RVS05c showed redshifted \\nad absorption.  This difference is clearly seen in the left panel of Figure~\\ref{velhistplots}. Interestingly, a recent search for outflows in the AEGIS database has also revealed an excess of inflows among AGN-powered systems \\citep{smnkl09}.  Inflow has been observed in at least one Seyfert 1 galaxy, NGC 5548, using ionized gas detected in the UV, though detections have not been reported for such a large number of objects as we have found here \\citep{mew99}.  There has been one tentative observation of redshifted X-ray absorption in a Seyfert 1 galaxy, but the authors caution that the significance of the absorption line they have measured is highly uncertain \\citep{dcm05}.  When we examine the inflow data on the IR-faint galaxies alone, we find no significant correlation between \\dvmax~and galactic mass (left panel in Figure~\\ref{inplots}).  Neither is there any obvious trend with Seyfert type or far-infrared luminosity (right panel of Figure~\\ref{inplots}), in contrast to the trends seen for the outflowing gas.  Additionally, we see inflows in nearly the same fraction of Seyfert 1s and Seyfert 2s, indicating no particular trend with Seyfert type.  As for the detected outflows, the influence of host galaxy inclination on column density was explored using extinction-corrected axis ratios \\citep{rc3}, but again, no conclusion could be drawn due to small-number statistics.  Next, the velocity distributions of the inflowing components for the IR-faint and IR-luminous Seyfert 1 galaxies were combined together and compared with the combined distribution of IR-faint and IR-luminous Seyfert 2 galaxies. K-S and Kuiper tests were performed on $\\Delta v$, \\dvmax, and Doppler parameter, following the same procedure as in Section \\ref{outdisc}, and the results are listed in Table~\\ref{stattests2c}.  They confirm the lack of obvious differences in the inflow properties between the two types of Seyfert galaxies.  \\subsubsection{Search for Connection between Inflows and Nuclear Structures} \\label{nuc}  The \\nags\\ absorption infall velocities often extend to relatively high values so we favor a nuclear location for this gas rather than a galactic origin (e.g. galactic fountains, \\citealt{fr06}). Nuclear accretion of cool gas like \\nags\\ with T~$\\sim$~100 K can provide fuel not only for star formation but also for nuclear activity \\citep{smo08}.  Various mechanisms have been proposed to help reduce the large angular momentum of the gas in the nuclei of galaxies. These include nuclear ($\\lesssim$ 1 kpc) bars and spirals \\citep{pm02,m03a,m03b,dm09} and gravitational interactions with neighboring galaxies \\citep{ckga03}.  Nuclear bars are thought to form when a large galactic bar forces gas inwards, creating a gaseous disk, and instability causes formation of a small gas bar near the nucleus (\\citealt{s89,mp01}).  Stellar bars could also accomplish the same thing.  Both bar types are capable of removing angular momentum from gas rotating near the nucleus \\citep{mp01}.  Nuclear spirals have been proposed as another AGN fueling mechanism since shock fronts that occur at their boundaries can take away angular momentum from local gas and cause material to fall in towards the black hole \\citep{m03b}. However, these nuclear structures do not necessarily lead to AGN fueling since they are present in a equally large fraction of non-active galaxies \\citep{pm02}.  Star formation may occur in these objects and disrupt AGN-fueling inflows \\citep{dm09}.  We revisit this issue here by looking for the presence of nuclear spirals or bars in the IR-faint Seyfert galaxies with inflow signature.  We have compared our results to nuclear structure surveys done by three different groups \\citep{mgt98,m03a,dck06}.  Nuclear dust structures in objects in common with these studies are classified into five distinct morphological categories (see Column (11) of Table~\\ref{objprop}): irregular dust, dust filaments, nuclear dust ring, nuclear dust bar, and nuclear dust spiral.  Of the thirteen objects which show inflows, five show evidence for nuclear dust spirals, bars, or rings, one shows evidence for dust filaments, one for irregular dust, and six show no sign of nuclear dust structure. The significant fraction of objects that show both inflows and nuclear dust spirals, bars, rings, and/or filaments lends credence to the idea of a connection between morphology and kinematics.  Of the seven objects with irregular dust or no dust, five have nearby companion galaxies, so tidal forces due to interactions with these companions may cause AGN fueling \\citep{rvb95,hc99,ssh07}.  This leaves only two objects with inflow which do not show nuclear structure or a companion: Akn~202 and Mrk~1018.  There are also two objects in our study with nuclear dust structure that show measurable {\\em outflow} rather than inflow (NGC~3786 and NGC~7319), and thus whether we measure outflow rather than inflow may be a consequence of our line of sight to the nucleus (inflow and outflow may be occurring in different planes, or in the same plane but over a different range of azimuthal angles), rather than a lack of inflow in the object.  Measurements of the line-of-sight velocity field with a resolution of $\\sim$10s of pc will be needed to disentangle the geometry of the inflows/outflows detected in our data (e.g., the study of NGC~1097 by \\citealt{dm09, vdv09}).  \\subsubsection{Inflow Dynamics} \\label{mmein}  The same method used to calculate mass, momentum, and kinetic energy for outflows (Section \\ref{mme}) was used for inflows, but the characteristic absorber radius was reduced to 1 kpc, a rough upper limit to the scale of the circumnuclear structures (nuclear bars/spirals) believed to be responsible for feeding the AGN (Section \\ref{nuc}, \\citealt{m03a}).  Again, these dynamical quantities are rather uncertain, but are calculated to find out at least roughly how these inflows compare with the mass accretion rates necessary to power the AGN in these systems.  Table~\\ref{inindiv} lists the mass, momentum, and kinetic energy calculated for all objects with measured inflows.  The mass accretion rate for these inflows ranges from just under 1 \\rone \\smpy~to just under 5 \\rone \\smpy.  For comparison, the mass accretion rate needed to power an AGN is $\\dot{M} = L_{\\mathrm{bol}} / c^2 \\eta,$ where $L_{bol}$ is the bolometric luminosity of the AGN, $\\dot{M}$ is the mass accretion rate, $c$ is the speed of light, and $\\eta$ is an efficiency factor dictating how much of the rest mass of the material being accreted is turned into radiation.  If we take $\\eta$ to be $\\approx$ 0.1 \\citep{rif08} and allow the bolometric luminosity to be $\\sim$ 10$^{44}$ ergs s$^{-1}$, typical for Seyfert galaxies \\citep{pr88, ckg03}, then we find that a mass accretion rate of $\\dot{M}$~$\\sim$~10$^{-2}$ \\smpy~is required. Even with our rough order of magnitude estimates, the mass accretion rates of all of our observed \\nad inflows are well above the amount necessary to power the AGN.  Thus the inflows that we are measuring carry enough material to fuel the AGN in these objects, even if only $\\sim$1\\% of this material makes its way down to the AGN.  The total infalling mass of $\\sim$~10$^{7}$ \\rone \\msun, estimated from our data, is enough to sustain nuclear activity over typical AGN lifetimes ($\\sim 10^7 - 10^8$ yrs; \\citealt{mt02,cr04})."
    },
    "0911/0911.1776_arXiv.txt": {
        "abstract": "Active galactic nuclei (AGNs) have been one of the most widely discussed sources of ultrahigh-energy cosmic rays (UHECRs).  The recent results of Pierre Auger observatory (PAO) have indicated a possible composition change of UHECRs above $\\sim10^{18.5}$~eV towards heavy nuclei.  We show here that if indeed UHECRs are largely heavy nuclei, then nearby radio quiet AGNs can also be viable sources of UHECRs.  We derive constraints on the acceleration sites which enable acceleration of UHECRs to $10^{20}$~eV without suffering losses.  We show that the acceleration of UHECRs and the survival of energetic heavy nuclei are possible in the parsec scale weak jets that are typically observed in these objects, the main energy loss channel being photodisintegration. On this scale, energy dissipation by shock waves resulting from interactions inside a jet or of the jet with surrounding material are expected, which may accelerate the particles up to very high energies.  We discuss the possible contribution of radio-quiet AGNs to the observed UHECR flux, and show that the required energy production rate in UHECRs by a single object could be as low as $\\approx 3 \\times 10^{39} \\rm{\\, erg \\, s^{-1}}$, which is less than a percent of the bolometric luminosity, and thus energetically consistent. We discuss consequences of this model, the main one being the difficulty in detecting energetic secondaries ($\\gamma$-rays and neutrinos) from the same sources. ",
        "introduction": "\\label{sec:intro}  The origin of ultrahigh-energy cosmic rays (UHECRs), cosmic rays with energies above $\\gtrsim 10^{18.5}$~eV, is still a mystery. It is commonly believed that the sources of UHECR's are extragalactic, since at these energies the cosmic rays cannot be confined by the magnetic field of our galaxy (for recent reviews, see Refs. \\cite{Der07, BEH09}).  Due to energy losses by photo-meson production with the cosmic microwave background (CMB), the sources of UHECR's are further limited to a distance $\\lsim 100$~Mpc from the observer (this limit is known as the `Greisen-Zatsepin-Kuzmin (GZK) cutoff'; Refs. \\cite{Greisen66, ZK66}).  While early observations \\cite{Tak98} suggested a violation of the GZK cutoff at high energies, recent, high statistics observations by the High-Resolution Fly's eye (HiRes) experiment and the Pierre Auger observatory (PAO) finds a downturn in the spectrum consistent with the GZK predictions \\cite{Abbasi08, Abraham08b}.   Extracting information from observations is, however, not easy. For example, observations of the spectrum itself did not lead, so far, to a firm conclusion about the origin of the steepening at high energies (i.e., whether it originates from the GZK cutoff or not), and therefore other possibilities are still allowed.  One of the consequences of this difficulty is the lack of a clear theoretical picture of the extragalactic UHECR sources.  Possible candidates of UHECR sources are limited among the known astrophysical objects, because of the difficulty in satisfying the two nearly contradictory requirements: A strong magnetic field in the source is needed to confine the accelerated cosmic rays, while the magnetic and photon fields cannot be too strong in order to avoid too much synchrotron radiation and photo-meson energy losses \\cite{Hillas84}. Several possible sources of UHECRs that fulfill the constraints mentioned above are often discussed in the literature. The most widely discussed candidates are gamma ray bursts (GRB )\\cite{MU95, W95, Vietri95, Der02, W04, MINN08}, low luminosity GRBs and hypernovae \\cite{MINN06, WRMD07, MINN08}, and jets in radio-loud (RL) AGNs \\cite{BS87, RB93, NMA95, Aha02, BGG06, Berezhko08, DRFA08}.  UHECR acceleration in the vicinity of black holes has also been considered for AGNs, including radio-quiet (RQ) AGNs (see, e.g., Refs. \\cite{PS92,BG99}).  Additional suggestions include magnetars \\cite{Arons03,MMZ09} and clusters of galaxies \\cite{KRJ96,ISMA07}.  An important step towards understanding the origin of UHECRs was taken recently with the discovery of a correlation between the arrival direction of the highest energy cosmic rays and nearby galaxies in the \\cite{VCV06} (VCV) catalog of active galactic nuclei (AGNs) [12th edition] \\cite{Abraham07, Abraham08a}. These results, however, should be taken with great care. It was noted already by Refs. \\cite{Abraham07, Abraham08a} that the observed correlation alone is insufficient to conclude that the sources of UHECRs are AGNs, because AGNs themselves are clustered within galaxies, and thus may only be the tracers of the true sources.  Indeed, a correlation was found between the arrival directions of UHECRs and galaxies \\cite{KS08,Tak08,Ghis08}, which suggests a correlation with the large scale structure (LSS) in the nearby universe.  The emerging picture gets further complicated by the recent results of HiRes experiment which do not confirm the correlation found by PAO \\cite{Abbasi08b}.  Moreover, the latest PAO results \\cite{Abraham09} suggest that the correlation degree is reduced compared to that obtained by earlier measurements.  An additional clue on the nature of UHECRs sources may come from measurements of their chemical composition. However, deducing the chemical composition of UHECRs from current observations is also difficult due to our poor knowledge on hadronic interactions. PAO results indicate a heavy composition at the highest energies \\cite{Unger07, Belido09}, while HiRes results suggest a proton composition (see Ref. \\cite{BEH09}, and references therein). The results of these two experiments are therefore clearly in contradiction. Moreover, the heavy element composition suggested by PAO results may be in contradiction with the anisotropy result of the same experiment (see below).  The chemical composition of UHECRs may be crucial in determining the nature of their source. If UHECRs are protons, the theoretical restrictions on the physical conditions inside potential sources are rather extreme, so that if indeed AGNs are the sources of UHECRs, then only powerful AGNs such as RL AGNs can be viable sources.  In addition, the known values of the galactic magnetic field, as well as observational upper limits on the intergalactic magnetic fields (IGMFs) in voids \\cite{Kro94} and theoretical investigations about IGMFs in the structured region \\cite{RKCD08} allow one to expect a correlation between the arrival directions of UHECRs and AGNs \\cite{DKRC08}, which is tentatively suggested by PAO.  On the other hand, a drawback of the idea that AGNs are sources of UHECRs is the fact that a strong correlation with RL AGNs is not found, except for very few cases, such as Cen A or Cen B \\cite{MSPC09}. In addition, a correlation with the most powerful AGNs, FR II galaxies, is currently not confirmed. Moreover, the number density of FR II galaxies seems too small to explain the observed small scale anisotropy \\cite{TS09}. Furthermore, the observed bolometric luminosity of AGNs which are correlated with the arrival direction of UHECRs seems too small to satisfy the minimum conditions for UHECRs acceleration in continuous jets \\cite{ZFG09}. These results indicate that only very few nearby AGNs are sufficiently powerful to accelerate particles to the observed ultrahigh energies.  In order to overcome this problem, it was suggested by Ref. \\cite{FG09} that UHECR production may occur transiently during giant AGN flares caused by tidal disruptions.  Although this may be a valid scenario, several restrictions on the duration and rate of such flares can be derived \\cite{Sig09,MT09}.  Since the actual activity of AGNs is not known, concrete conclusions cannot be drawn at this stage (see also Ref. \\cite{WL08}).   An alternative picture emerges if UHECRs are composed of heavier elements, such as iron nuclei. The ambiguity in the interpretation of the data led \\cite{Allard05, Allard08} to study a mixed composition scenario above $\\sim 10^{18.5}$~eV. A similar model (albeit with somewhat lower maximum energy of UHECR) was studied recently by \\cite{ABG09}.  In these models, significant correlation between the arrival directions of UHECRs and galaxies is not expected, because of the strong deflection of heavy nuclei in galactic magnetic field. In addition, protons dominate the cosmic rays spectra only up to energies $\\sim {10}^{18.5}$~eV, which implies that the requirement on the physical conditions at the acceleration site of potential sources is significantly loosened.   As we show here, a heavy element composition of UHECRs opens the possibility for RQ AGNs, which compose the majority of the AGN population within $\\sim 100$~Mpc, to be the sources of UHECRs.  As opposed to RL AGNs, in which evidence exist for strong jets on $\\gtrsim$~kpc scale, in RQ AGNs there are no evidence of such strong jets \\footnote{Note though that \\cite{PC09} showed recently that radio emission from jets may be suppressed if the emitting particles cool rapidly enough close to the base of the jet, so that synchrotron emission from outer parts of the jet is self absorbed. Although this work was done in the context of jets in microquasars, a similar idea can be easily implemented in the case of RQ AGNs as well. However, a general discussion on the role of synchrotron-self absorption on the distinguishing between RL and RQ AGN's is outside the scope of this paper, and is left for future work.}. Nonetheless, in recent years there has been an accumulation of evidence for weak, parsec-scale jets in these objects (Refs. \\cite{MWPG03,GBO04,UWTGM05,Gal+06, MARK07,Ho08}, and references therein).  Internal interactions of blobs inside a jet, or of the jets with their environment, produce shock waves, which can accelerate particles.  Alternatively, cosmic rays may be accelerated by the dissipation of magnetic energy inside these jets (see, e.g., Ref. \\cite{LO07}). Moreover, as was pointed out by Ref. \\cite{AM04}, particle acceleration could take place in shocks within the corona, which could be caused by abortive jets.  In this work we thus discuss weak jets in RQ AGNs as possible sources of UHECRs, under the assumption that at very high energies the composition of UHECRs is dominated by heavy (possibly, but not necessarily, iron) nuclei. As we will show below, RQ AGNs can fulfill the requirements needed from sources of UHECRs under this assumption. We assume here that particles acceleration takes place inside the jets, so that this work is different from previous works that considered UHECR production in the vicinity of black holes that do not have jets (e.g., Ref. \\cite{BG99}).  This paper is organized as follows. In \\S\\ref{sec:2} we present the basic constraints on the physical conditions at the acceleration sites of heavy nuclei. In particular, in \\S\\ref{sec:photodisintegration} we derive the constraints from photodisintegration and photomeson production. We show that these conditions can be fulfilled easily. In \\S\\ref{sec:3} we discuss the efficiency of particle acceleration required to explain the observed flux of UHECRs on earth, under the assumption that RQ AGNs are the only sources of UHECRs.  We further derive constraints on the value of the magnetic field at the acceleration site, and show that a lower limit of few percents of equipartition value is required. We summarize and conclude in \\S\\ref{sec:summary}. ",
        "conclusions": "\\label{sec:summary}  In this paper we have considered radio quiet AGN's as possible sources of UHECRs. So far, jets in these objects were not considered as sources of the highest energy cosmic rays around $10^{20}$~eV, since they are not luminous enough to support acceleration of protons to these energies. However, the recent results of the AUGER collaboration shows indications that at the highest energies the composition of cosmic rays may be dominated by heavy nuclei \\cite{Unger07,Belido09}. The assumption that UHECRs are heavy nuclei eases the constraint on the source luminosity (see eq. \\ref{eq:2}). Thus it allows, in principle, radio-quiet AGNs to be the sources of UHECRs.  We have calculated in \\S\\ref{sec:2} the constraints on the acceleration site which enable the acceleration of heavy nuclei to high energies. The most restrictive constraint arises from the photo-disintegration process. We showed in \\S\\ref{sec:photodisintegration}, that for typical nearby AGN's with bolometric luminosity $L_{bol}^{ob} \\approx 10^{43} \\; {\\rm erg \\, s^{-1}}$, energetic nucleus can survive photo-disintegration if the acceleration takes place on a parsec scale (see eqs. \\ref{eq:survive}). Interestingly, this is the same scale at which weak jets were seen in these objects \\cite{MWPG03,GBO04,UWTGM05,MARK07}. This fact further supports our idea, since interaction of jets with the surrounding material inevitably leads to creation of shock waves, which are the most plausible acceleration site of particles to high energies.  The question of the acceleration of particles to ultrahigh energies may depend on the shock velocity. In the Bohm approximation $\\eta \\propto \\beta^{-2}$, some assumptions about the shock velocities in the jets are needed. For example, a shock velocity of $\\beta \\sim 0.1$ requires magnetic luminosity of $\\epsilon_B L \\gsim 10^{43.5} \\; {\\rm erg \\, s^{-1}}$ in order to enable acceleration of iron nuclei to the maximum observed energy (see \\S\\ref{sec:2}, equation \\ref{eq:2}), which may be too low. However, somewhat higher values, $\\beta \\sim 0.3$ may be sufficient. Since currently there is a high uncertainty in the determination of the shock velocities in RQ AGNs, it is not possible to put further constraints. We have also showed in \\S\\ref{sec:3} that the minimum inferred values of the magnetic field required to confine the accelerated particles at parsec scale jets is sub-Gauss. This value was shown to be smaller than the equipartition value, which is also consistent with models of particle acceleration in shock waves that may be produced on this scale.   We further showed in \\S\\ref{sec:3}, that the required luminosity in UHECRs could be less than a percent of the total bolometric luminosity of more abundant nearby AGNs. If the spectrum of produced cosmic rays is a power law with index close to $d \\log N/d\\log E \\approx -2$, as suggested by models of particle acceleration in non-relativistic shock waves \\cite{BE87}, then the total luminosity in energetic particles (at all energies) could still be much smaller than the total bolometric luminosity. Even for higher power law index, $p=2.3$, we find that the total energy requirement for acceleration of cosmic rays above $\\GeV$ is $\\approx (10^{10})^{0.3} = 10^3$ times higher than the energetic requirement from UHECRs alone, which is still (marginally) consistent with the total energy budget in RQ AGNs (see discussion in \\S\\ref{sec:3}).  Assuming that the main radiative source in these objects is the accretion disk, this implies a high efficiency, of the order of tens of percents, in the conversion of accretion energy to acceleration of particles. Since large uncertainties exist in the efficiencies of both the energy conversion in the jet and the acceleration of particles, we can only conclude that our model is consistent with the energetic requirements, provided a high efficiency is achieved in both these processes.   Spectral synthesis as well as chemical enrichment models predict that AGN's are metal rich. The metallicities in the broad emission line region is typically $\\sim 1$ to $\\gsim 10$ times the solar metallicity \\cite{HF93, Ham02}, and grows with the luminosity \\cite{Shemmer04}. Thus, abundant existence of heavy nuclei in AGN's disks is expected. As the plasma jet is composed of material from the disk, abundant population of heavy nuclei in the jet is expected. Thus, AGN's may be one of the natural sources of high energy heavy nuclei.    If indeed the acceleration takes place on a parsec scale, then the main energy loss channel is photodisintegration, rather than photomeson production. As discussed in \\S\\ref{sec:implications}, this fact implies that we do not expect abundant production of energetic neutrinos, that may result from the decay of energetic pions. On the contrary: the results in equations \\ref{eq:survive} indicate that the relative contribution of photomeson production may be no more than few percents of photodisintegration as an energy loss channel of UHECRs. Thus, we do not expect a copious production of neutrinos under the conditions which enable the acceleration and survival of heavy nuclei, as considered in this paper.  After they escape from their sources, UHECR nuclei can be subject to photodisintegration in the intergalactic space. UHE nuclei with energy $\\gsim 10^{19}$~eV mainly interact with the cosmic infrared background (CIB) photons \\cite{SS99,Allard05, Allard08}. UHE iron nuclei with energy $\\lsim 10^{20}$~eV have a mean free path of $\\gsim 500$~Mpc, while the mean free path of oxygen nuclei at a similar energy is only $\\sim 30$~Mpc \\cite{HST07}. Both increase rapidly as the energy of the nuclei decreases. Thus, the GZK horizon of UHE nuclei is comparable to, or even somewhat larger than the GZK horizon of energetic protons (see also Ref. \\cite{Der07}).  To summarize, we have pointed out here that the assumption that UHECRs are composed of heavy nuclei, as is suggested by recent AUGER results \\cite{Unger07, Belido09}, enables radio quiet AGN's to be their sources. This picture is consistent with what is currently known about the existence of parsec scale jets seen in some of these objects, as well as the energy constraints. Moreover, this picture is supported by the recent analysis carried by Ref. \\cite{George+08}, in which a strong correlation between the arrival directions of UHECRs and the sky coordinates of AGN's detected by the {\\it Swift}-BAT within the GZK cutoff (which are largely radio-quiet) was found.  As more data are collected by the AUGER collaboration, a clearer picture of the composition of UHECR's will become available, allowing further constraints on possible sources."
    },
    "0911/0911.0685_arXiv.txt": {
        "abstract": "We re-examine claims for redshift evolution in black hole-bulge scaling relations based on lensed quasars.  In particular, we refine the black hole mass estimates using measurements of Balmer lines from near-infrared spectroscopy obtained with Triplespec at Apache Point Observatory.  In support of previous work, we find a large scatter between Balmer and UV line widths, both \\mgii$~\\lambda \\lambda 2796, 2803$ and \\civ$~\\lambda \\lambda 1548, 1550$.  There is tentative evidence that \\ciii$~\\lambda 1909$, despite being a blend of multiple transitions, may correlate well with \\mgii, although a larger sample is needed for a real calibration.  Most importantly, we find no systematic changes in the estimated BH masses for the lensed sample based on Balmer lines, providing additional support to the interpretation that black holes were overly massive compared to their host galaxies at high redshift. ",
        "introduction": "Locally, we observe tight correlations between the properties of bulge-dominated galaxies and the masses of their central supermassive black holes \\citep[BHs; e.g.,][] {kormendyrichstone1995,tremaineetal2002, gultekinetal2009}.  The mechanisms that establish and maintain these relations are uncertain, despite innumerable suggestions in the literature \\citep[e.g.,][]{silkrees1998,murrayetal2005, miraldakollmeier2005,hopkinsetal2006,peng2007}.  In principle, evaluating the demographics of nuclear BHs as a function of redshift should observationally constrain the processes that lead to the tight scaling relations observed today.  Unfortunately, it is currently prohibitive to obtain dynamical BH masses for systems beyond tens of Mpc.  Thus, estimates of BH mass at large distance are necessarily based on very indirect methods linked to accretion processes in active BHs \\citep[e.g.,][]{vestergaard2002}.  Several studies have used active galaxies to probe evolution in BH-bulge scaling relations, probing redshifts from $0.4 < z < 0.6$ \\citep{wooetal2006,treuetal2007,wooetal2008} and 1$\\lesssim z \\lesssim$ 4 \\citep{pengetal2006a,pengetal2006b,salvianderetal2007, shieldsetal2006,ho2007,jahnkeetal2009,mcleodbechtold2009} all the way to $z=6.4$ \\citep{walteretal2004}.  Generally speaking, a wide variety of observations suggest that BH-bulge relations do evolve with redshift, in the sense that the ratio of BH to bulge mass was higher at early times \\citep[but see also][]{shieldsetal2003,alexanderetal2005}, although we have observational constraints only for the most massive systems (\\mbh$>10^8$~\\msun) at high redshift.  This counterintuitive result has stimulated vigorous discussion both about the ramifications for the coevolution of BHs and bulges \\citep[e.g.,][]{robertsonetal2006,croton2006} and about whether there are built-in biases in the measurement techniques \\citep{laueretal2007}.  Unfortunately, even apart from potential population biases, our interpretation of the observations are prone to significant uncertainty.  On the one hand, it is only possible to obtain BH mass estimates at cosmological distances using active galaxies (typically luminous quasars at high redshift).  Immediately it becomes very challenging to characterize the host galaxy properties, when the quasar outshines the underlying galaxy starlight by factors of $\\sim 10-30$ \\citep[e.g.,][]{pengetal2006a,kimetal2008a}.  Apart from direct imaging, some groups have used gas measurements (predominantly CO) to obtain dynamical masses \\citep{shieldsetal2006,hoetal2008b,walteretal2004,riechersetal2008,riechersetal2009}, which may or may not provide a reliable tracer of the galaxy mass \\citep{ho2007}.  The width of narrow emission lines, particularly \\oiii~$\\lambda 5007$, have also been substituted for the galaxy velocity dispersion \\citep[e.g.,][]{shieldsetal2003,boroson2003,salvianderetal2007,gaskell2009}. While there is a strong correlation between stellar and gaseous velocity dispersion in low-luminosity sources \\citep[e.g.,][]{heckmanetal1981,nelsonwhittle1996,greeneho2005o3,ho2009o3}, there is good reason to suspect that it does not hold at high luminosity \\citep[e.g.,][]{greeneetal2009}.  \\citet[][P06 hereafter]{pengetal2006b} mitigated the host galaxy contrast problem by focusing on lensed quasars.  The quasars are lensed differently from the underlying (resolved) host galaxies, reducing the contrast problem discussed above.  For that reason, we focus on the lensed quasar sample in this paper.  Daunting as measuring high-redshift galaxy mass and stellar velocity dispersion may be, particularly in the presence of a luminous quasar, the BH mass measurements are equally problematic.  The techniques are indirect and model-dependent.  Briefly, active galaxies contain dense gas orbiting at distances of light days to months from the central BH that gives rise to broad emission lines with widths of thousands of \\kms.  By combining a size scale with the line width of the emitting region, the broad-line region (BLR) gas can be used as a dynamical tracer of the BH mass \\citep[e.g.,][]{dibai1980}.  Direct size estimates are obtained by measuring the time lag between variability in the continuum and line emission \\citep[reverberation mapping;][]{blandfordmckee1982,petersonetal2004}. \\hskip 0.in \\psfig{file=tableobsv2.eps,width=0.45\\textwidth,keepaspectratio=true,angle=0} \\vskip 4mm \\noindent A tight, empirically determined correlation between BLR radii and the luminosity of the active galaxy \\citep[the radius-luminosity relation; ][] {kaspietal2005,bentzetal2006,bentzetal2009a} can be used to obtain approximate BLR radii for the quasar population in general \\citep[e.g.,][]{vestergaard2002}.  The current data support a slope of $R_{\\rm BLR} \\propto L^{0.5}$, as is expected if neither the density structure of the BLR nor the shape of the ionizing continuum depends on luminosity \\citep[e.g.,][]{bentzetal2006}.  Unfortunately, the radius-luminosity relation has not been calibrated extensively for luminosities greater than \\lf$\\approx 10^{46}$~erg~s$^{-1}$ \\citep[see][]{kaspietal2007}.  Because the kinematic structure and inclination of the BLR are unknown, the derived virial ``masses'' (\\mbh$\\propto \\upsilon^2 R_{\\rm BLR}/G$) have no physically motivated normalization.  Recent reverberation mapping experiments reveal signatures of inflow, rotation, and outflow in individual objects, but such two-dimensional maps remain scarce \\citep{welshhorne1991,bentzetal2009b,denneyetal2009rm}.  Our current practice is to compare the virial masses with independent estimates of BH mass, typically using the \\msigma\\ relation, to derive an average scale factor \\citep[e.g.,][]{gebhardtetal2000b,ferrareseetal2001,onkenetal2004, nelsonetal2004,greeneho2006msig,shenetal2008a}.  The fact that any correlation is seen between virial masses and the mass inferred from the \\msigma\\ relation is encouraging, but leaves much room for large systematic uncertainties \\citep[e.g.,][]{krolik2001,collinetal2006}.  At higher redshift, only rest-frame UV spectra are readily available for large samples of quasars.  The scaling relations for the UV lines (e.g.,~\\mgii~$\\lambda 2800$~\\AA; \\civ~$\\lambda 1550$~\\AA) include additional layers of uncertainty.  Virtually no reverberation mapping has been performed with \\mgii, and the scaling relations for this line are simply scaled to match \\hbeta\\ \\citep[e.g.,][]{mcluredunlop2004,onkenkollmeier2008}.  In the case of \\civ, reverberation mapping has been done \\citep[e.g.,][]{petersonetal2005}, but there are strong reasons to suspect that the \\civ\\ line width is not dominated by virial motions \\citep[e.g.,][]{gaskell1982,baldwinetal1996,richardsetal2002a,leighlymoore2004, baskinlaor2005,sulenticetal2007,shenetal2008b} although debate continues on this point \\citep[e.g.,][]{vestergaardpeterson2006,kellybechtold2007, gavignaudetal2008}.  For these reasons, virial masses based on Balmer lines \\citep[preferably \\halpha;][]{greeneho2005cal} have the most credibility, since these have been directly compared with alternate estimates of \\mbh.  We focus specifically on obtaining Balmer-based virial masses for the high-redshift lensed quasar sample from P06.  Our primary goal is to determine whether the masses presented in P06 are {\\it systematically} biased by the use of UV line transitions.  We follow P06 and assume a standard cosmology with $H_0 = 100~h = 70$~\\kms~Mpc$^{-1}$, $\\Omega_{\\rm m} = 0.30$, and $\\Omega_{\\Lambda} = 0.70$. ",
        "conclusions": "We revisit the BH mass estimates for a sample of lensed quasars with high-fidelity host-galaxy luminosities from \\citet{pengetal2006b}. While the Balmer-based masses presented here are arguably more robust than the UV-based estimates, we find no evidence for a systematic difference in BH masses based on the two methods.  If we can take the Balmer-based virial masses at face value, then we confirm the result of \\citet{pengetal2006b} that BHs appear to be overly massive relative to their hosts at high redshift.  Intriguingly, the minor mergers that are invoked to puff up compact elliptical galaxies at late times would be very effective at boosting the galaxy to BH mass ratio as well.  However, as pointed out many times, the persistently narrow range in observed line width for luminous quasars at high redshift provides substantial cause for concern in the veracity of the virial masses.  Furthermore, it is not clear that we have yet assembled consistent comparison quasar samples across cosmic time with which to compare the BH-bulge scaling relations.  Our work is a necessary, but not sufficient, step in determining the true evolution of BH-bulge relations with cosmic time."
    },
    "0911/0911.0150_arXiv.txt": {
        "abstract": "} \\nc{\\eab}{  We revisit the anisotropic universe model previously developed by Ackerman, Carroll and Wise (ACW), and generalize both the theoretical and computational framework to include polarization and various forms of systematic effects. We apply our new tools to simulated WMAP data in order to understand the potential impact of asymmetric beams, noise mis-estimation and potential Zodiacal light emission. We find that neither has any significant impact on the results. We next show that the previously reported ACW signal is also present in the 1-year WMAP temperature sky map presented by \\cite{liu2009}, where data cuts are more aggressive. Finally, we reanalyze the 5-year WMAP data taking into account a previously neglected $(-i)^{l-l'}$-term in the signal covariance matrix. We still find a strong detection of a preferred direction in the temperature map. Including multipoles up to $\\ell=400$, the anisotropy amplitude for the W-band is found to be $g = 0.29 \\pm 0.031$, nonzero at $9 \\sigma$. However, the corresponding preferred direction is also shifted very close to the ecliptic poles at $(l,b)= (96,30)$, in agreement with the analysis of \\cite{hanson:2009}, indicating that the signal is aligned along the plane of the solar system. This strongly suggests that the signal is not of cosmological origin, but most likely is a product of an unknown systematic effect. Determining the nature of the systematic effect is of vital importance, as it might affect other cosmological conclusions from the WMAP experiment. Finally, we provide a forecast for the Planck experiment including polarization. ",
        "introduction": "\\label{sec:introduction} In recent years, the study of the cosmic microwave background (CMB) has proved to be the most fruitful addition to our understanding of the early universe. Observations of the CMB anisotropies, like those obtained by the Wilkinson Microwave Anisotropy Probe (WMAP) experiment \\citep{bennett:2003, hinshaw:2007}, have provided us with incomparable insight on the composition of structure in our universe. Combined with previous experimental knowledge and a sound theoretical framework, the concordance model of $\\Lambda$CDM has been established.  The $\\Lambda$CDM model relies on the framework of inflation. Inflation was initially proposed as a solution to the horizon and flatness problem \\citep{guth:1981}. Additionally, it established a highly successful theory for the formation of primordial density perturbations, providing the required seeds for the large-scale structures (LSS).  Eventually, these later gave rise to the temperature anisotropies in the cosmic microwave background radiation that we observe today \\citep{guth:1981,linde:1982, muhkanov:1981, starobinsky:1982, linde:1983, linde:1994, smoot:1992, ruhl:2003, runyan:2003, scott:2003}.  One of the predictions from inflation is that the observed universe should be nearly isotropic on large scales. However, anomalies found in the CMB during the recent years \\citep{de Oliveira-Costa:2004, vielva:2004, eriksen:2004a} suggest that anisotropic inflationary models should be considered.  A specific example is the generalized model presented by \\citet{ackerman:2007}, which considers violation of rotational invariance in the early universe. A general framework for describing similar models was presented by \\citet{pullen:2007}.  \\cite{himmel:2009a, himmel:2009b} showed that the anisotropic inflationary background of the ACW model characterized by a fixed-norm vector field ultimately is unstable. However, the parametrization of the signal covariance matrix is independent of that unstable model, and is very useful. It represents general correlations induced by rotations in the CMB at a phenomenological level. Several papers have recently investigated the properties of the ACW model with extensions \\citep{hou:2009, karviauskas:2009, dimopoulos:2009, carroll:2008}.   Work in this field suggests that the 5-year WMAP data contains a significant ACW anisotropic signal, corresponding to a 3.8$\\sigma$ detection in the W-band \\citep{groeneboom:2008b}. A more recent paper by \\cite{hanson:2009} points out that the direction is incorrect due to a neglected factor of $(-i)^{l-l'}$ corrected in a later version of \\cite{ackerman:2007}, yielding an ACW-signal in which the preferred direction is located very close to the ecliptic poles.      In this paper, we re-analyze the 5-year WMAP data including the previously neglected $(-i)^{l-l'}$-factor, and investigate whether traces of the ACW anisotropic contribution signal are still evident. The analysis will, as previously, be performed with the CMB Gibbs sampling framework \\citep{jewell:2004, wandelt:2004, eriksen:2004b}, which by \\cite{groeneboom:2008b} was included to allow for non-diagonal, but sparse covariance matrices. This framework allows for exact Bayesian analysis of high-resolution CMB data with a non-diagonal CMB signal covariance matrix. The isotropic method has already been applied several times to the WMAP data \\citep{odwyer:2004, eriksen:2007a, eriksen:2007b, eriksen:2008b}, and has already been extended to take into account polarization \\citep{larson:2007} and internal component separation \\citep{eriksen:2008a}.  In our re-analysis of the 5-year WMAP temperature data, we confirm that the direction is shifted to the ecliptic poles, at a greatly increased significance. As the north and south ecliptic poles are aligned with our solar system, the ACW signal in the WMAP data is therefore most likely a systematic effect and not of cosmological origin. This is in complete agreement with \\cite{hanson:2009}.    It has not yet been possible to fully rule out whether any known systematic effect could have contributed  to the signal. In theory, either asymmetric beams, mis-estimated noise or even the Zodiacal light could have affected the detection of the ACW-signal. In this paper, we consider these three effects and conclude that neither have any effect on the ACW-signal.   Until now, the framework has only supported temperature-temperature correlations (TT). Here, we extend the mechanics to include E-mode correlations (EE), including cross-mode correlations (TE). We then provide a forecast for the upcoming Planck experiment, considering simulated T+E maps. The Planck data will hopefully be able to rule out all doubts about the origin of the ACW signal. ",
        "conclusions": "\\label{sec:conclusion} We have generalized a previously developed Bayesian framework to allow for exact analysis of any general anisotropic universe models that predicts a sparse signal harmonic space covariance matrix, including polarization data. This generalization involved incorporation of a sparse matrix library into the existing Gibbs sampling code called ``Commander''. We implemented support for this model in our codes, before demonstrating and validating the new tools with appropriate simulations including polarization data. First, we compared the results from the Gibbs sampler with brute-force likelihood evaluations, and then verified that the input parameters were faithfully reproduced in realistic WMAP simulations.  We then considered a special case of anisotropic universe models, namely the \\citet{ackerman:2007} model which generalizes the primordial power spectrum $P(k)$ to include a dependence on direction, $P(\\mathbf{k})$. The equations were however not complete, and the analysis performed by \\cite{groeneboom:2008b} has been re-done including the previously neglected $(-i)^{l-l'}$-term.  We then analyzed the five-year WMAP temperature sky maps, and presented the updated WMAP posteriors of the ACW model. The results from this analysis are in accordance with the results from \\cite{hanson:2009}, showing that the preferred direction is now located at the ecliptic poles. This suggests that the signal is most likely not of cosmological origin, and its origin must be either from within the solar system or systematics.  We have investigated four cases of systematic effects that share similar structures with the ACW signal. We have shown that neither asymmetric beams, the Zodiacal light, noise RMS mis-estimation nor possible pixel anti-correlations in the WMAP data could have given rise to the observed signal.  To summarize, we have shown that there exist a strong anisotropic signal corresponding to the ACW signal in all the WMAP data that is aligned with the north and south ecliptic poles. The probability that the axis should correspond so closely to the ecliptic poles is very low, indicating that the signal is due to a systematic effect. The signal makes up more than $5 \\%$ of the total power of the temperature fluctuations in the CMB. We have excluded some of the possible candidates as source of the ACW signal. Determining the nature of the systematic effect will be of vital importance, as it might affect other cosmological conclusions from the WMAP experiment, and the upcoming Planck data will clearly be invaluable for understanding the nature of this feature."
    },
    "0911/0911.0220_arXiv.txt": {
        "abstract": "In the past several years high resolution kinematic data sets from Milky Way satellite galaxies have confirmed earlier indications that these systems are dark matter dominated objects. Further understanding of what these galaxies reveal about cosmology and the small scale structure of dark matter relies in large part on a more detailed interpretation of their internal kinematics. This article discusses a likelihood formalism that extracts important quantities from the kinematic data, including  the amplitude of rotation, proper motion, and the mass distribution. In the simplest model the projected error on the rotational amplitude is shown to be $\\sim 0.5 $ km s$^{-1}$ with $\\sim 10^3$ stars from either classical or ultra-faint satellites. The galaxy Sculptor is analyzed for the presence of a rotational signal; no significant detection of rotation is found, and given this result limits are derived on the Sculptor proper motion. A criteria for model selection is discussed that determines the parameters required to describe the dark matter halo density profiles and the stellar velocity anisotropy. Applied to four data sets with a wide range of velocities, the likelihood is found to be more sensitive to variations in the slope of the dark matter density profile than variations in the velocity anisotropy. Models with variable radial velocity anisotropy are shown to be preferred relative to those in which this quantity is constant at all radii in the galaxy. ",
        "introduction": "Since their initial discovery~\\cite{Shapley}, dwarf spheroidals (dSphs) have offered a unique insight into the formation of galaxies and structure on the smallest scales. Initially characterized as unusual and ghostly stellar systems, photometric studies tended to find that these systems contained old stellar populations with no recent signature of star formation activity~\\cite{Mateo:1998wg}. Though photometrically well-studied since their discovery over seventy years ago, as late as nearly 30 years ago minimal was known on the internal kinematic properties of their stellar populations or on the kinematic properties of these objects in the Milky Way (MW) halo.  Aaronson \\cite{Aaronson1983} provided the first measurement of the line-of-sight velocities of stars in Milky Way dSphs. From the spectra of merely three carbon stars, Aaronson suggested a mass-to-light ratio for the Draco dSph nearly an order of magnitude greater than that of Galactic globular clusters. Follow-up studies of several dSphs, including Sextans, Fornax, Ursa Minor, Sculptor, increased the velocity samples by an order of magnitude, and in the process established these systems to be dark matter dominated ~\\cite{Suntzeff1993,Mateo1991,Mateo1993}. It was further suggested that all of these systems share a similar dark matter halo mass of $\\sim [1-5] \\times 10^7$ M$_\\odot$~\\cite{Mateo1993}. Even at the time of these early measurements, it was understood that the mass distributions of these systems provide strong constraints on the properties of the particle nature of dark matter, including its mass and primordial phase-space density ~\\cite{Lake1990,Gerhard1992,Faber1983}.  With the advent of high resolution, multi-object spectroscopy, the velocity samples from the brightest dSphs initially studied in Refs.~\\cite{Suntzeff1993,Mateo1991,Mateo1993} have now increased by up to three orders of magnitude~\\cite{Gilmore2007,Walker2007,Walker:2008ji}. These new data sets have revealed that the velocity dispersions of the systems are all $\\sim 10$ km s$^{-1}$, and in all cases the dispersions remain constant even out to the projected radius of the outermost velocity measurements ~\\cite{Walker2007}. Though the data sets have increased by more than ten-fold, the more modern analysis of these systems still confirms the global conclusion established from the initial observations that dSphs are strongly dark matter dominated ~\\cite{Gilmore2007,Walker2007,Strigari:2008ib,Lokas:2009cp}.  Not only has the past several years seen an increase in the kinematic data sets for the brightest dSphs, the number of known Milky Way satellites has more than doubled due to the Sloan Digital Sky Survey (SDSS). As of the writing of this article, the SDSS has discovered 14 new Galactic satellites ~\\cite{Willman:2004kk,Willman:2005cd,Belokurov:2006ph}. The new SDSS systems have lower luminosities and surface brightnesses than the 11 classical Milky Way satellites that were known prior to SDSS. The half-light radius for several of these new objects is less than $\\sim 100$ pc; this radius is smaller than the typical half-light radius of the classical satellites but still somewhat larger than the typical globular cluster half-light radius of $\\sim 1-10$ pc.  Several kinematics studies on the ultra-faint population of SDSS satellites have been undertaken in the past several years~\\cite{Munoz:2006vg,Martin:2007ic,Simon:2007dq,Geha:2008zr}. Using spectra from eight of the SDSS satellites,  Simon and Geha~\\cite{Simon:2007dq} concluded that these objects are strongly dark matter dominated. Several of the ultra-faint satellites have velocity dispersions as low as $\\sim 5 $ km s$^{-1}$, making them the most-promising systems to study the phase space limits of the dark matter. It has additionally been observed that the ultra-faint satellites are the most metal-poor systems known, and that they form a continuation of the luminosity metallicity trend set by the brightest dSphs ~\\cite{Kirby2008,Geha:2008zr}.  With the above data sets now available, it is becoming increasingly necessary to develop better theoretical tools to interpret them. An important aspect of the theoretical modeling will necessarily require an interpretation of the kinematic data sets for the population of MW satellites; a detailed understanding of these kinematic data sets will be important not only for determining the mass distributions of each individual system, but for a global comparison to theories of Cold Dark Matter (CDM) ~\\cite{Strigari:2007ma,Strigari:2008ib}. Understanding the mass distributions will also be important for interpretation of limits on particle dark matter masses and annihilation cross sections in high-energy gamma-ray experiments~\\cite{Strigari:2006rd,Essig:2009jx,Martinez:2009jh}. Further, understanding the kinematics of these systems may eventually reveal whether they have dark matter cusps or cores, which would in itself provide a stringent test of the CDM paradigm~\\cite{Hogan:2000bv}.  The primary aim of this article is to discuss a maximum likelihood formalism that is used for extracting important physical quantities from dSph kinematic data sets. Section~\\ref{sec:data} begins by reviewing the properties of the kinematic data sets and defining the likelihood. Section~\\ref{sec:tracers} then uses the likelihood to extract rotational and proper motion signals. Section~\\ref{sec:mass} discusses mass modeling and a new calculation for model selection. Section~\\ref{sec:conclusion} presents the conclusions. ",
        "conclusions": "\\label{sec:conclusion} This article has discussed the analysis of kinematic data from Milky Way dwarf spheroidals, with a primary motivation of 1) understanding physical quantities that are well-constrained by the data and 2) understanding the systematics that underly the determination of the dark matter masses of these systems, given the simplest assumption that the dSphs are purely pressure supported systems. Of the possible systematics perhaps the most significant and observationally- accessible is the determination of a velocity gradient in the data sample, which may be indicative of tidal disruption from the potential well of the Milky Way. The results in the literature indicate that, based on the kinematic data alone, velocity gradients due to tidal disruption or rotation are not conclusively present in any of the dSphs. This article has provided an example, using a simple parameterization, of how to search for rotation in the kinematic data sets using a maximum likelihood analysis. The kinematic sample of Sculptor was analyzed, and it was found that the maximum likelihood rotational amplitude is zero, with an upper limit of $\\sim 2$ km s$^{-1}$ at 90\\% c.l. The magnitude of these errors are consistent with the projected magnitude of the errors from theoretical modeling.  When modeling the mass distribution of the dark matter halos of the dSphs, degeneracies between model parameters affect the determination of the total mass profiles, even in the context of the simplest spherical models. To shed light on these degeneracies, this article has discussed a new criteria for model selection applied to the dSph kinematic data sets, taking a step towards determining how many parameters are needed to describe the mass distribution of spherical halos. For the four dSphs studied here, chosen because they have a wide range of available line-of-sight velocities, it is shown that, assuming CDM-motivated Einasto profiles for the dark matter halos, models with variable velocity anisotropy are slightly preferred relative to those with constant velocity anisotropy. Further, central slopes for the dark matter profile that are found in CDM simulations are not a unique description of the data sets; both more cuspy and less cuspy models are allowed for the central slope. This is primarily due to the degeneracy between the central dark matter slope with the central stellar profile and the velocity anisotropy distribution ~\\cite{Strigari:2007vn}.  Future photometric and kinematic data sets promise to further pin down the mass distributions of the dSph dark matter halos. Upcoming data for the ultra-faint satellites will be particularly important, and may be able to show whether any tidal effects are present in these galaxies. Further, development of non-spherical distributions for both the light and dark matter should be considered given these data sets (for initial results along these lines see Ref~\\cite{Lokas:2009ty}). Controlling systematics in these data sets will prove to be important step towards further testing the currently favored $\\Lambda$CDM theory of structure formation."
    },
    "0911/0911.2013_arXiv.txt": {
        "abstract": "Steady meridional flow makes no first-order perturbation to the frequencies of helioseismic normal modes. It does, however, Doppler shift the local wavenumber, thereby distorting the eigenfunctions. For high-degree modes, whose peaks in a power spectrum are blended into continuous ridges, the effect of the distortion is to shift the locations of those ridges. From this blended superposition of modes, one can isolate oppositely directed wave components with the same local horizontal wavenumber and measure a frequency difference which can be safely used to infer the subsurface background flow. But such a procedure fails for the components of the more-deeply-penetrating low-degree modes that are not blended into ridges. Instead, one must analyze the spatial distortions explicitly. With a simple toy model, we illustrate one method by which that might be accomplished by measuring the spatial variation of the oscillation phase. We estimate that by this procedure it might be possible to infer meridional flow deep in the solar convection zone. ",
        "introduction": "\\label{sec:introduction}  The character of the meridional flow in the solar convection zone is a fundamental component of modern mean-field and flux-transport dynamo models \\citep[e.g.,][]{Charbonneau:2005, Dikpati:2006, Rempel:2006a, Rempel:2006b}. In the upper convection zone, meridional flows influence how magnetic fields are transported from low to high latitudes, and how the fields are subducted and subsequently processed in deeper layers. In some models the flows at the base of the convection zone are believed to establish the period of the activity cycle through the equatorward advection of the active-region belts.  Many dynamo models assume a simple structure for the meridional circulation with a single cell in each hemisphere. However, global simulations of the solar convection zone often exhibit complex multicellular patterns \\citep[e.g.,][]{Brun:2002, Brun:2004} which lack the simple structure sometimes posited. Even those simulations that are dominated by a single cell in each hemisphere \\citep{Miesch:2008} usually possess countercells in thin layers at the top and the base of the convection zone. Furthermore, while these 3-D simulations generate meridional flows with amplitudes of about 15 to 20 m s$^{-1}$ close to the surface, at greater depths the flow velocity is about 10 m s$^{-1}$, quite unlike the far smaller values that have been conjectured from simplified mass-conservation arguments. The difference arises from systematic coupling between the radial flow and the meridional circulation which results from turbulent entrainment and detrainment processes.  Observations are not yet able to winnow the multitude of theoretical and numerical models successfully. The meridional circulation deep below the surface is presently poorly constrained, owing to a distressing paucity of reliable information. A variety of techniques---including direct Doppler measurement \\citep[e.g.,][]{LaBonte:1982, Hathaway:1996}, magnetic feature tracking \\citep{Komm:1993, Svanda:2007}, and local helioseismology \\citep[e.g.,][]{Giles:1997, Haber:2002, Basu:2003, Zhao:2004}---have had great success as indicators of meridional flow in the upper layers of the convection zone. The concensus is that, in the near-surface layers, the meridional flow is largely poleward and remarkably constant with depth, with amplitude roughly 20 m s$^{-1}$. However, the flow is quite variable with longitude, with active regions being sites of local inflow \\citep{Haber:2002, Gonzalez-Hernandez:2008, Hindman:2009}. While local-helioseismological analyses have determined in exquisite detail the flows in the upper 15\\% of the convection zone, the detection of meridional flow in deeper layers has been elusive. A variety of attempts to measure the deep meridional flow have been made using both time--distance procedures \\citep{Giles:2000, Duvall:2003} and spectral procedures that seek frequency shifts between northward- and southward-propagating waves \\citep{Braun:1998, Krieger:2007, Mitra-Kraev:2007}. But, stymied by the weak signal and the large systematic errors, the flow below a depth of roughly 30 Mm remains elusive \\cite[for a detailed discussion of the inherent difficulties see][]{Duvall:2009}.  Helioseismology determines flows below the solar surface from measurements of the difference in the Doppler shift between acoustic waves propagating in opposite directions. This is fundamentally true whether the method being used is a spectral technique (such as ring analysis or global-mode analysis) or a time--distance procedure. In spectral techniques, the Doppler shift has traditionally been inferred by assessing the temporal frequency shift of the acoustic-wave dispersion relation. In the presence of a steady flow, waves propagating in opposite directions suffer opposite wavenumber shifts which are often interpreted as Doppler frequency shifts at a given wavenumber in a power spectrum of the wave field. However, this Doppler-frequency procedure is not universally applicable, and, as we shall see in subsequent sections, the measurement of meridional circulations in the lower half of the convection zone is just such an instance where it should not be applied blindly. In fact, in this instance, ``traditional\" techniques will fail, and inversions of so-called frequency splittings to determine the flow field are prone to misinterpretation.  In this paper we present a 1-D analog of the effects of meridional flow on acoustic modes. We use this simple model to illustrate the pitfalls that may occur when using spectral techniques to measure meridional circulation in the lower reaches of the convection zone and in the tachocline. We attempt to dispel several misconceptions that have recently arisen concerning the nature of the acoustic spectra in the presence of meridional flow, and suggest a novel approach by which progress might be made. Explicitly, in Section 2, we discuss the changing nature of the acoustic-mode spectrum from low to high degree. In Section 3, we present the 1-D analog model and the artificial data that the model generates. These data are used in Section 4 to illustrate an analysis scheme designed to measure the flow speed from helioseismic observations. In Section 5 we consider briefly how the model and the analysis scheme generalize to the Sun, and we discuss implications of our results. ",
        "conclusions": "\\label{sec:Discussion}  The principles we have illustrated with our 1-D model can be applied to the Sun. The most significant difference is that, for each mode, $c^{-1}U$ in Equation~\\eqnref{eqn:UfromPhase} should be replaced by its appropriately weighted vertical average $\\left<c^{-1}U_{ln}\\right>$. For deeply penetrating modes, the high-order asymptotic approximation to the oscillation eigenfunctions provides an adequate first guide \\citep{Gough:1993}, yielding a weighting function $W(r,\\omega/L)$ that changes form at the lower turning point $r_{\\rm t}$ given implicitly by $c(r_{\\rm t})/r_{\\rm t}=\\omega/L$ where $L = l + 1/2$. Below the lower turning point the weighting function vanishes, $W = 0$, and above the lower turning point it is given by  \\begin{equation} W\\left(r,\\frac{\\omega}{L}\\right) \\sim \\left(1-\\frac{L^2c^2}{\\omega^2 r^2}\\right)^{-1/2}. \\end{equation}  \\noindent Approximate weighting functions for higher-order modes are displayed by \\cite{Gough:1983}.  The analogous quantity to the return travel time in the one-dimensional model, $2a/c$, is the circumundulation time $T_{ln} \\sim 4 l\\pi/\\omega_{ln}$, the time an acoustic-wave packet takes to traverse a great circle around the Sun. If the damping time is short compared with the circumundulation time, $\\eta_{ln} T_{ln} \\gg 1$ (as is true for high-degree modes), the wave suffers little self interference, the mode power blends with nearby modes forming a continuous ridge, and measuring the apparent frequency shift $\\delta\\omega$ at given $l$ works well. But once $\\eta_{ln} T_{ln} \\ll 1$ (low-degree modes), resonances dominate the wave field, and the modes are well separated in the power spectrum, each with a frequency that is unshifted by the meridional circulation.  Attempting to measure a shift by fitting a Lorentzian to the frequency profile, or cross-correlating the power from poleward- and equatorward-propagating waves, cannot therefore result in a meaningful nonzero value. Instead, a background flow shifts and broadens the distribution of power in degree $l$. As our discussion indicates, one expects a null frequency shift for waves of low $l$: once $\\eta_{ln} T_{ln} \\lesssim 1$, the apparent shift $\\delta\\omega$ must decline with decreasing $l$. Indeed that is just what appears to have been observed \\citep{Braun:1998, Mitra-Kraev:2007}, as is illustrated in Figure~\\ref{fig:MitraKraev}.  Misinterpreting this decline as a true decline in frequency has led to an erroneous inference about the depth of the reversal of the Sun's meridional circulation. The circulation is well known to be directed poleward near the surface, and if the radial average $\\left<c^{-1}U_{ln}\\right>$ over most of the convection zone (as sampled by low-$l$ modes) vanishes, then the poleward surface flow must be counterbalanced by a deep equatorward flow within the convection zone. If $\\delta\\omega$ were a measure of that radial average, the flow reversal would be located at the lower turning point of the modes near where $\\delta\\omega$ starts to decline with decreasing $l$. That is just where the continuous ridges in the power spectrum of the modes give way to distinguishable discrete peaks. The transition occurs at $\\nu/l \\lesssim 20 \\mu$Hz (see Figure~\\ref{fig:MitraKraev}) which \\cite{Mitra-Kraev:2007} acknowledge corresponds to a turning point at about 40 Mm beneath photosphere. However, we have argued that this depth has actually little, if anything, to do with the location of the reversal of the real flow.  Since the eigenfrequencies are insensitive to the meridional flow, the only unambiguous way to detect meridional flow in the lower half of the convection zone is via the structure of the oscillation eigenfunctions. One telling property of that structure is that separability in space and time is destroyed by the flow, causing latitudinal variation in temporal phase. We propose here that the structure of the flow can in principle be determined from direct measurements of that phase. We recognize that the phase variation can equivalently be interpreted as a spatial variation at fixed time, which together with an accompanying amplitude variation, distorts the eigenfunction from its flow-free, spherical-harmonic counterpart. In a decomposition into spherical harmonics, this distortion appears as a leakage of power across degree $l$, that could conceivably be measured.  However, we suspect that any attempt to measure that distortion directly is likely to be much more susceptible to noise; it would also need to be separated from the much larger contributions to the total distortion resulting from shear in the zonal flow.  Finally, it behooves us to address the likelihood that a reversal in the flow direction with depth will actually be detectable. Our simulations with 20 modes indicates that, if our error estimates are realistic, the lower limit on the Mach number for a measurable flow is between $10^{-4}$ and $10^{-3}$. That corresponds to a flow velocity between about 20 and 200 m s$^{-1}$ at a depth of about 150 Mm. The lower value is comparable with the flow speeds in the convection-zone simulations by \\cite{Miesch:2008}. So perhaps the detection of the flow reversal in the Sun is almost in sight."
    },
    "0911/0911.4129_arXiv.txt": {
        "abstract": "We describe comprehensive calculations of the formation of icy planets and debris disks at 30--150~AU around 1--3 \\msun\\ stars. Disks composed of large, strong planetesimals produce more massive planets than disks composed of small, weak planetesimals. The maximum radius of icy planets ranges from $\\sim$ 1500~km to 11,500~km.  The formation rate of 1000~km objects -- `Plutos' -- is a useful proxy for the efficiency of icy planet formation. Plutos form more efficiently in massive disks, in disks with small planetesimals, and in disks with a range of planetesimal sizes. Although Plutos form throughout massive disks, Pluto production is usually concentrated in the inner disk. Despite the large number of Plutos produced in many calculations, icy planet formation is inefficient.  At the end of the main sequence lifetime of the central star, Plutos contain less than 10\\% of the initial mass in solid material. This conclusion is independent of the initial mass in the disk or the properties of the planetesimals. Debris disk formation coincides with the formation of planetary systems containing Plutos.  As Plutos form, they stir leftover planetesimals to large velocities. A cascade of collisions then grinds the leftovers to dust, forming an observable debris disk. In disks with small ($\\lesssim$ 1--10~km) planetesimals, collisional cascades produce luminous debris disks with maximum luminosity $\\sim 10^{-2}$ times the stellar luminosity. Disks with larger planetesimals produce debris disks with maximum luminosity $\\sim 5 \\times 10^{-4}$ (10~km) to $5 \\times 10^{-5}$ (100~km) times the stellar luminosity.  Following peak luminosity, the evolution of the debris disk emission is roughly a power law, $ f \\propto t^{-n}$ with $n \\approx$ 0.6--0.8. Observations of debris disks around A-type and G-type stars strongly favor models with small planetesimals. In these models, our predictions for the time evolution and detection frequency of debris disks agree with published observations. We suggest several critical observations that can test key features of our calculations. ",
        "introduction": "Dusty disks of debris surround many main sequence stars \\citep{bac93,che05,rie05,moor06,rhe07a,wya08}. Among young stars, the frequency of debris disks ranges from $\\sim$ 50\\% for B-type and A-type stars \\citep[][2009]{su06,cur08b} to $\\sim$ 10\\% to 20\\% for solar-type stars \\citep{tri08,mey08} to $\\sim$ 5\\% for M-type stars \\citep{pla09,les09}. Binary stars and single stars are equally likely to have debris disks \\citep{stau05,su06,bry06,gor07,sie07,tri07}. Among stars with masses \\mstar\\ $\\gtrsim$ 0.8--1~\\msun, the debris disk frequency declines with stellar age \\citep{rie05,cur08,car09b}.  Debris disks are signposts of planet formation. Infrared (IR) and radio observations indicate that grains with typical sizes of 1--100 $\\mu$m produce most of the dust emission. Radiation processes remove these grains on timescales much shorter than the age of the central star \\citep{bac93,wya08}.  To maintain emission from small grains throughout the main sequence lifetime, some process must replenish the dust. The simplest replenishment mechanism invokes a 10--100 \\mearth\\ reservoir of $\\sim$ 1~km solid objects which continuously collide at high velocities and fragment into smaller objects \\citep{bac93,hab01,kb04b,wya08,heng10}. Although this mass is plausible \\citep[e.g., ][2007b]{and05}, collisional damping among an ensemble of 1~$\\mu$m to 1 m objects rapidly reduces collision velocities to low values \\citep[e.g.,][]{kb01}. Thus, debris disks require a mechanism to maintain the high velocities of the solids. Gravitational stirring by massive planets is the most successful mechanism \\citep{wya08}. Thus, in the current picture, maintenance of debris disks requires massive planets.  Two broad classes of planetary systems can explain the general observations of debris disks. In the Solar System, external perturbations power the local debris disk. Jupiter's gravity excites the orbits of the asteroids and produces the Jupiter family comets; collisions among asteroids and mass loss from comets produce the Zodiacal light \\citep[e.g.,][2008, 2009]{nes06}. Beyond 30~AU, Neptune's gravity plays a similar role for the Kuiper belt and the scattered disk \\citep[e.g.,][]{mor04,cha07}.  The recent discovery of gas giant planets associated with the debris disks in HR 8799 and Fomalhaut \\citep{kal08,mar08,che09,su09} suggests similar processes occur in other planetary systems \\citep[see also][and references therein]{wil02,moran04,del05,mor05,qui06,wya06,fab07}. In particular, \\citet{must09} show that the gravitational perturbations of newly-formed planets at 5--10~AU can produce dusty debris at 30--100~AU. Their results suggest that these perturbation may yield the most luminous debris disks.  Smaller planets can also produce debris disks. In a series of papers, we show that the collisional evolution of solid material in a gaseous protostellar disk naturally leads to the formation of planets along with copious amounts of dusty debris \\citep[e.g.,][2004b, 2005]{kb02b}. In our picture, mergers of~km-sized objects first produce 500--1000~km protoplanets. These protoplanets stir up leftover planetesimals along their orbits. Destructive collisions among the leftovers initiate a collisional cascade, which produces a dusty debris disk \\citep[for different approaches to this problem, see also][]{kri08,the08,wya08,heng10,kenn10}. In the terrestrial zone at a few~AU from the central star, rocky protoplanets reach masses of 0.5--2~\\mearth\\ and rapidly remove the debris \\citep[][2005, 2006; see also Bottke et al 2007]{kb04a}. At 5--20~AU, the gas entrains the debris, allowing icy protoplanets to grow rapidly into the cores of gas giant planets \\citep{kb09}. At 30--150~AU, icy protoplanets reach maximum sizes of only $\\sim$ 1500--2000~km \\citep[][hereafter KB08]{kb08}. The collisional cascade among the leftover planetesimals produces a luminous debris disk; the maximum brightness and evolution of the debris matches the observations of known debris disks reasonably well (KB08).  Here, we continue our exploration of the formation and evolution of icy planets and debris disks. In KB08, we described a suite of calculations for disks at 30--150~AU around 1--3~\\msun\\ stars. We considered disks with a single surface density law, a single initial size distribution for planetesimals, and a single set of fragmentation parameters for solid objects. For the calculations discussed here, we examine planet formation in disks with an expanded set of initial conditions for the initial surface density of the disk and for the initial properties of the planetesimals. These results yield new predictions for the maximum sizes of icy planets as a function of initial planetesimal size. Results for the long-term evolution of debris disks continue to account for many fundamental aspects of the data.  Our new analysis demonstrates that luminous debris disks at 30--150~AU require planetesimals with initial sizes of 1--10~km instead of 100--1000~km.  Our calculations suggest that the minimum stable grain size, the slope of the size distribution for small grains, and the slope of the IR emissivity law are also critical parameters.  Spatially resolved images of debris disks around A-type and solar-type stars can improve our understanding of the minimum stable grain size. Larger samples of debris disks with high quality submm data from the Atacama Large Millimeter/Submillimeter Array (ALMA), the {\\it Herschel Space Observatory}, and the Stratospheric Observatory for Infrared Astronomy (SOFIA) can place better constraints on the size distribution for small objects and the slope of the emissivity law.  Together, these data can test our predictions for the time evolution of debris disk emission around 1--3~\\msun\\ stars and can provide input for more complete calculations that include the formation and dynamical evolution of giant planets.  We outline the numerical model in \\S2. We describe the formation of icy planets in \\S3 and the evolution of debris disks in \\S4; these sections include `highlights' (\\S3.3 and \\S4.3) which summarize the main results and conclusions. In \\S5, we consider applications of our results to the observed time evolution (\\S5.1) and frequency (\\S5.2) of debris disks around A-type and solar-type stars.  We conclude with a brief summary in \\S6. ",
        "conclusions": "Together with the discussion in KB08, the results described here provide a robust theoretical picture for the formation of icy planets and debris disks from a disk of icy planetesimals surrounding 1--3~\\msun\\ stars.  These calculations set the context for the evolution of dusty debris in a dynamic system of planets and establish  a framework for interpreting existing observations of debris disks around intermediate mass stars. Our analysis also suggests new observational tests of this picture.  We describe a suite of numerical calculations of planets growing from ensembles of icy planetesimals at 30--150~AU in disks around 1--3~\\msun\\ stars. Using our hybrid multiannulus coagulation code, we solve for the evolution of sizes and orbits of objects with radii of $\\sim$ 1~m to $\\gtrsim$ 1000~km over the main sequence lifetime of the central star.  These results allow us to constrain the growth of planets as a function of disk mass, planetesimal size, surface density law, stellar mass, and semimajor axis.  In disks with small planetesimals, debris disk formation is coincident with the formation of a planetary system.  All calculations of icy planet formation with 1--10~km planetesimals at 30--150~AU lead to a collisional cascade which produces copious amounts of dust on timescales of 5--30~Myr. This dust is observable throughout the lifetime of the central star. Because we consider a broad range of input parameters, we derive the time evolution of (i) dust produced in the collisional cascade and (ii) the IR and submm emission from this dust as a function of the initial parameters.  In disks with larger planetesimals, debris disks form near the end of the main sequence lifetime of the central star. The collisional cascade is weak and produces little dust. This dust is barely observable with current technology. As long as the planetesimals are initially large, this conclusion is independent of the initial mass or surface density of the disk, the semimajor axis, or the stellar mass.  We divide the rest of this section into (i) theoretical considerations, (ii) observable consequences, and (iii) observational tests. The theoretical considerations build on the highlights of icy planet formation in \\S3.3. Observable consequences and tests of the calculations follow from the discussions in \\S4.3 and \\S5.  \\subsection{Theoretical Considerations}  \\begin{enumerate}  \\item Icy planet formation at 30--150~AU is self-limiting. Starting with a swarm of $\\lesssim$ 10~km planetesimals, runaway growth produces a set of 100--500~km protoplanets. As the protoplanets grow, they stir up leftover planetesimals along their orbits. When the leftovers reach high $e$, collisions produce debris instead of mergers. Protoplanets cannot accrete leftovers rapidly; a cascade of destructive collisions grinds the leftovers to dust. Poynting-Robertson drag and radiation pressure then remove the dust from the disk.  \\item In disks composed of $\\lesssim$ 10~km planetesimals, the maximum sizes of icy planets at 30--150~AU depend primarily on the initial disk mass.  The largest icy planets around 1--3~\\msun\\ stars have $r_{max} \\sim$ 1500--3500~km and $m_{max} \\sim$ 0.003--0.05~\\mearth.  These objects contain $\\lesssim$ 3\\%--4\\% of the initial disk mass.  Plausible ranges in the fragmentation parameters can reduce the maximum mass by a factor of $\\sim$ 10. In addition to these robust limits on the maximum size of icy planets, the finite main sequence lifetimes of 1--3~\\msun\\ stars limits the formation of many large planets in the outer disk. Thus, the inner disk produces many more Pluto-mass planets than the outer disk (Tables \\ref{tab:rad1000.1}--\\ref{tab:rad1000.6}).  \\item In disks composed of $\\gtrsim$ 10~km planetesimals, icy planets reach larger sizes.  Our results suggest $r_{max} \\sim$ 6000--11500~km and $m_{max} \\sim$ 0.2--1.5 \\mearth. These planets form slowly. By the time the central star evolves off the main sequence, large planets contain $\\lesssim$ 10\\% of the initial disk mass. The weak collisional cascade produces little dust. Thus, Poynting-Robertson drag and radiation pressure remove little dust from the disk.  \\item At 30--50~AU, collisional cascades remove a substantial fraction of small planetesimals in the massive disks around 1--3~\\msun\\ stars.  In disks with $x_m \\gtrsim$ 0.1, $\\lesssim$ 10\\% of the initial mass in small objects remains when the central star evolves off the main sequence. At larger disk radii ($\\gtrsim$ 100~AU), more than 50\\% of the initial disk mass is still in small planetesimals at $t = t_{ms}$. Throughout lower mass disks with $x_m \\lesssim$ 0.1, most of the mass remains in small planetesimals.  \\item Collisional cascades produce copious amounts of dust. Dust starts to form during the transition from runaway to oligarchic growth, peaks when the largest objects first reach sizes of $\\sim$ 1000~km, and slowly declines as objects reach their maximum sizes. This evolution is more sensitive to the initial sizes of the planetesimals than on the initial surface density law or the fragmentation parameters. In massive disks with small planetesimals, this evolution begins as the central star approaches the main sequence; in lower mass disks and in disks with large planetesimals, peak dust formation occurs as the central star evolves off the main sequence. In ensembles of disks with $x_m$ = 0.01--1, the peak mass in 0.001--1 mm (0.001--1 m) particles ranges from 0.002 to 10 lunar masses (0.002 to 10 \\mearth).  \\item Radiative processes can remove large amounts of mass from the disk. During the early stages of the collisional cascade, radiation pressure produces a wind of small particles. This wind contains 60\\% to 100\\% of the mass lost from the disk. During the end stages of the cascade, Poynting-Robertson drag pulls the rest of the lost mass into regions with $a <$ 30~AU.  \\item The impact of gas giant planets on debris disks is a major uncertainty in our calculations. Gas giants at 5--20~AU \\citep[e.g.,][]{kenn08} produce small-scale fluctuations in the orbits of planetesimals at 30--150~AU \\citep[e.g.,][]{must09} and large-scale dynamical events throughout the disk \\citep[e.g.,][]{gom05}. These perturbations modify the evolution of the collisional cascade and impose structure in the radial distribution of small objects \\citep[e.g.,][]{wil02,mor05,qui06,boo09}. Current observational constraints on the frequency of gas giants around G stars \\citep{cum08} and A-type stars \\citep{john07} suggest gas giants probably form only in the most massive disks. Thus, gas giants probably cannot change our predictions for debris disks in low mass disks ($x_m \\lesssim$ 0.1).  In massive disks around 2--3~\\msun\\ stars, the rapid rise in dust production at 30--50~AU roughly coincides with the timescale for gas giant planet formation. Thus, gas giants probably enhance the dust production rate and may sometimes change the outcome of planet formation at 30--150~AU.  In massive disks around 1--1.5~\\msun\\ stars, gas giants form well before the rise in the dust production.  Here, gas giants probably play a major role in the evolution of the debris disk. Quantifying this role requires simulations that include gas giant planet formation in disks with dimensions of 3--150~AU. These simulations are now possible on high performance computers.  \\item Planetesimal formation is another uncertainty. Current simulations of the growth of small, micron-sized dust grains in a turbulent disk require accurate treatments of collisions, dust dynamics, and gas dynamics \\citep[e.g.,][]{ric06,gar07,kre07,bra08,cuz08,lai08,you08,bir09}. Recent results are inconclusive on the typical planetesimal size, with expected sizes ranging from meters to more than a thousand kilometers. Expected formation times are generally shorter than the typical evolution time of $\\gtrsim$ 1~Myr. Thus, our assumption of initial ensembles of 1~m to 100~km planetesimals is probably reasonable.  \\item Our models provide some guidance for models of planetesimal formation.  Several recent numerical simulations of streaming instabilities in turbulent disks yield planetesimals with $r \\gtrsim$ 100--1000~km \\citep[e.g.,][2009]{joh07}.  \\citet{mor09} demonstrate that an initial ensemble of large planetesimals may resolve several outstanding issues in the formation of the asteroid belt in the Solar System. In the context of our calculations, however, the observed evolution of debris disks around A-type and G-type stars strongly favors disks with a significant amount of solid material in small ($\\lesssim$ 1--10~km) planetesimals at 30--150~AU. Although forming small planetesimals in a turbulent disk is challenge, the dead zones of protoplanetary disks are attractive locations for producing planetesimals with $r \\sim$ 0.1--1~km on short timescales \\citep[e.g.,][]{bra08}. Based on our results, identifying similar mechanisms to produce small planetesimals at 30--150~AU may provide important insights into the formation of debris disks and planetary systems.  \\end{enumerate}  \\subsection{Observational Consequences}  Our calculations yield robust observational consequences of the collisional cascade.  \\begin{enumerate}  \\item The dusty debris from the collisional cascade is directly observable throughout the main sequence lifetime of the central star. For disks around 1--3~\\msun\\ stars, the maximum fractional dust luminosity of $L_d / L_{\\star} \\sim 10^{-2}$ is comparable to the maximum dust luminosities of known debris disks \\citep{bac93,rie05,su06,rhe07a}. With typical radial optical depths of 1--10 at maximum, this limit is roughly an order of magnitude smaller than theoretical maximum \\citep[e.g.,][]{wya07a,heng10}. The dust temperature at the inner edge of a 30--150~AU disk scales with the temperature of the central star; thus, the predicted 24~$\\mu$m excess is very sensitive to the stellar mass. At 70~$\\mu$m, the predicted excesses scale roughly linearly with disk mass and stellar mass. The predicted 160--850~$\\mu$m excesses depend on the disk mass but are nearly independent of the stellar mass.  \\item The amount of dusty debris depends on the initial sizes of planetesimals. In our calculations, bright debris disks at 5--30~Myr require planetesimals with initial sizes $\\lesssim$ 1~km. Calculations with 10--100~km planetesimals produce no debris around young stars and little debris around older stars. Thus, our results suggest planetesimals with initial sizes $\\sim$ 1~km.  \\item The amount of debris depends weakly on the initial surface density law and the fragmentation parameters. Disks with similar total masses produce similar amounts of dust emission. Disks with stronger planetesimals produce more dust emission than disks with weaker planetesimals.  \\item Stellar evolution limits the brightness evolution of debris disks. Among coeval stars with ages $\\gtrsim$ 100~Myr (e.g., star clusters), lower mass stars have brighter debris disks. Among older stars near the main sequence turn-off (e.g., field stars), our results suggest that debris disks around massive stars are relatively brighter than debris disks around lower mass stars.  \\item The time evolution of \\ldlstar, IR fluxes, and IR colors provide useful diagnostics of the source of the debris. When the collisional cascade begins, \\ldlstar\\ and IR excesses increase with time, reach a plateau, and then decline with time. IR colors generally become redder when IR excesses rise and bluer when IR excesses fall. This evolution contrasts with the expected evolution of a pre-existing disk of small particles responding to changes in the properties of the central star, where \\ldlstar\\ always declines with time.  \\end{enumerate}  \\subsection{Observational Tests}  Observations of debris disks around A-type and G-type stars provide good tests of our predictions.  More comprehensive comparisons with the data require calculations with a broader range of grain parameters than discussed here. For example, \\citet{kenn10} develop an elegant semi-analytic model that captures many aspects of our planet formation calculations and apply this model to observations of A-type stars. By deriving the best set of fragmentation parameters and dust properties to fit the data, they infer robust constraints on the initial mass, radius, and surface density in the disk.  Here, we focus on comparisons that test the ability of our calculations to explain the observed evolution and frequency of debris disks around A-type and G-type stars.  \\begin{enumerate}  \\item Our calculations explain the observed time evolution of debris disk fluxes for solar-type stars (Figure \\ref{fig: f70-gstars}). For an initial ensemble of 1~km planetesimals, the predicted evolution of the 70 $\\mu$m excess follows the apparent rise in source fluxes at 10--100~Myr, the peak at $\\sim$ 100--300~Myr, and the decline from $\\sim$ 300~Myr to $\\sim$ 10~Gyr.  Larger samples of young G stars \\citep[e.g.,][2007b]{cur07a} would provide a better discrimination among models with different initial conditions.  Although predicted fluxes for models with a smallest stable grain size $r_2 \\lesssim$ 1 $\\mu$m, a slope of the emissivity law for small grains $q \\lesssim$ 1, and a slope for the size distribution of small grains $p \\gtrsim$ 3.5 improve our match to the data, our standard models explain 75\\% of the observed fluxes. These models also predict an undiscovered population of low luminosity debris disks. More sensitive surveys with {\\it Herschel} can test this prediction and provide better guidance on our adopted values for $r_2$ and $q$. Better constraints on $r_2$ and $q$ should also yield better guidance on $p$.  \\item Our models predict very small 24 $\\mu$m excesses from dust at 30--150~AU around G-type stars.  Matching the observed levels of 24 $\\mu$m emission requires another source of dust emission.  Dust in the terrestrial zone is the best candidate. For young stars, predicted fluxes from terrestrial debris disks typically exceed observed fluxes \\citep[e.g.,][2005]{kb04b}. Although the average flux in model debris disks drops rapidly, catastrophic collisions can produce observable levels of debris around stars with ages exceeding $\\sim$ 100 Myr \\citep[e.g.,][]{gro01,wya02,kb05,wya07a,smi08}. Resolved observations of debris disk emission around G-type stars at 24 $\\mu$m would test model predictions for dust emission from the terrestrial zone and from material at 30--150~AU.  \\item Our calculations also match the observed time evolution for debris disks around A-type stars (Figures \\ref{fig: f24-astars1}--\\ref{fig: f24-astars2}). For ensembles of planetesimals with initial sizes $\\lesssim$ 1~km, our results match the apparent rise in fluxes at 5--10~Myr, the plateau at 10--30~Myr, and the decline from $\\sim$ 30 Myr to $\\sim$ 1 Gyr.  The predicted rate of the decline, \\ldlstar\\ $\\propto \\lambda^{-n}$ with $n \\approx$ 0.6--0.8, agrees with the observed rate of decline \\citep[$n \\approx$ 0.5--1;][]{gre03,rie05,rhe07a,wya08}. Models for disks around 2.5~\\msun\\ stars match the peak fluxes of A-type stars at 10--30~Myr; models for 2~\\msun\\ stars provide a better match to fluxes during the decline. Accurate estimates of stellar mass would improve our evaluation of these matches.  \\item Predicted colors for debris disks around A-type stars also agree with the observations reasonably well (Figure \\ref{fig: c2470-astars}). Colors for disks with initial surface densities $\\Sigma \\propto a^{-1}$ match the data somewhat better than colors for disks with steeper initial surface density laws. Models around 2~\\msun\\ stars also provide better matches to the data. These models cannot match the colors of the reddest sources; colors for debris disks with ages of 30--100 Myr can test the models in more detail. These models yield clear predictions for the fluxes at longer wavelengths as a function of the stellar mass and the initial disk surface density relation. Large surveys -- such as the proposed JCMT Legacy Survey \\citep{mat07}, several Herschel key programs, and submm observations with ALMA and SOFIA -- can test these predictions.  \\item Our predictions provide reasonable matches to the observed detection probabilities of debris disks around A-type and solar-type stars. For solar-type stars, our predictions for $p_{d,70}$ and $p_{d,850}$ agree with {\\it Spitzer} and submm observations. Disks with outer radii of 75--100~AU may match the data better than disks with outer radii of 150~AU. Deeper surveys with {\\it Herschel} will provide better constraints on the outer disk radius and will test our ability to predict detection probabilities for low luminosity debris disks. For young A-type stars, predictions for $p_{d,24}$ and $p_{d,70}$ agree well with observations; among older stars, we overpredict the detection rates by a factor of 2--3.  Models with $r_2 \\approx$ 5--10 $\\mu$m provide a better match to the data. Calculations which include the dynamical evolution of giant planets probably also change the detection probability at late times.  More sensitive observations of A-type stars at 70--160 $\\mu$m would enable better tests of models with our standard values for $r_2$ and $q$ and would discriminate among models with different choices for these two parameters and models with gas giant planets.  \\item In their comprehensive analytic study of planetesimal disks, \\cite{heng10} show that emission from long-lived disks of m-sized planetesimals can mimic debris disks produced by a collisional cascade. Imaging and multi-wavelength spectral energy distributions can distinguish between these two possibilities. Planetesimals at a distance $a$ (in AU) from the central star emit as blackbodies with equilibrium grain temperatures, $T_g \\approx$ 278 $L_\\star^{1/4} a^{-1/2}$. Grains in debris disks are hotter than blackbody grains and have an emissivity law with $q >$ 0 \\citep[e.g.,][]{bac93,heng10}. Recent {\\it Spitzer} data suggest debris disks have small grains \\citep[e.g.,][2008, 2009]{su05}. Data from {\\it Herschel} will yield direct measurements of $q$ and enable searches for long-lived planetesimal disks.  \\end{enumerate}  From the broad suite of calculations in KB08 and this paper, we conclude that debris disks are the inevitable outcome of icy planet formation in a disk of solid objects. The full set of models explains many of the general properties of debris disks, including the time evolution of IR excesses and colors and the frequency of debris disks as a function of stellar mass and time. Our results also lead to clear predictions which can be tested with new and existing observations."
    },
    "0911/0911.1010_arXiv.txt": {
        "abstract": " ",
        "introduction": "Dust grains dominate the opacity in protoplanetary disks whenever they are present. This implies that their radiation properties play a crucial role in determining the temperature and density structure of these disks. The initial population of (sub)micron-sized particles evolves over time towards planetesimals, eventually providing the building blocks for terrestrial planets.  The dust grains shield the interior of protoplanetary disks from energetic cosmic particles and stellar X-ray radiation and provide the surface for electron recombination. This regulates the ionization structure of disks, which is an important ingredient for the magneto-rotational instability to operate and to drive angular momentum transport. The dust particles are equally important for disk chemistry because chemistry on grain surfaces leads to the formation of molecular ices and, possibly, complex organic molecules, which will enter the gas phase when their evaporation temperature is reached.  Infrared spectroscopy is a powerful tool to characterize the properties of protoplanetary dust. With ground-based telescopes, the {\\it Infrared Space Observatory ISO} and the {\\it Spitzer Space Telescope}, an enormous amount of data has been obtained to characterize the mineralogy of disks around a variety of objects, ranging in luminosity from Herbig Ae/Be stars to T Tauri stars, and even brown dwarfs. These data allow us to put constraints on the chemical composition and amorphous/crystalline state of the dust particles, and to address questions such as radial distribution and mixing processes. Mid-infrared long-baseline interferometry is starting to contribute as well to our understanding of the structural properties of dust in the different radial zones of protoplanetary disks. ",
        "conclusions": "The last decade has seen huge progress in our knowledge of dust properties in protoplanetary disks, which was made possible mostly through observations with both the {\\it Infrared Space Observatory} and the {\\it Spitzer Space Telescope}. A first surprise was the vast variety in appearance of dust features around young stars. A more detailed analysis, however, showed that those dust features, although so different in appearance, could be related to a handful of dust species, with different sizes and structures.  The dust in protoplanetary disks bears the signature of processing compared to ISM dust, the grains are clearly larger and most young sources also show evidence for crystalline silicates.  Despite the large number of spectra now available to the community, it has proven difficult to pin down the relationships between the observed dust properties, on the one hand, and the stellar and disk properties on the other hand. The obviously expected relation of increasing grain size with age, or of increasing crystallinity with higher temperature of the central star has proven to be incorrect.  The problem lies in the fact that there are many parameters to take into account: the age, stellar luminosity, disk flaring angle and dust sedimentation, presence of a close companion, to name just a few. Furthermore, dust features cannot be directly related to a specific dust species and size: other parameters, such as the dust temperature or the shape of the dust particles also play an important role. In addition, good laboratory data for astronomical dust species over the full observable wavelength range remain a much needed ingredient. The same is true for reaction rates which determine the conversion between different grain species.  Therefore, we expect to make more progress in the coming years by comparing the wealth of data provided by {\\it Spitzer} in a thorough statistical study, eliminating as much as possible those parameters that can be determined, so that, e.g., when studying the size distribution, stars with a similar luminosity are being compared.  In the coming years, several new observatories and instruments will become available. With the launch of the {\\it Herschel Space Observatory} in 2009, covering the far IR wavelength region (PACS instrument 57-210~$\\mu$m) we will be able to  study the lattice vibrations of heavy ions or ion groups with low bond energies with a high signal-to-noise. In particular, studies of the forsterite band at 69~$\\mu$m will be promising in the context of determining the dust temperature and the composition of the olivines, while aqueous alteration can be studied through features of hydrous silicates at 100-110~$\\mu$m.  Further on the horizon lies the launch of the {\\it James Webb Space Telescope (JWST)}, where the unprecedented sensitivity will allow us to study the disks of brown dwarfs with the same ease as we now study T Tauri stars, opening up yet another region of the parameter space in dust processing studies.  Last but not least, a lot can be learned from comparing observations of protoplanetary disks and dust with the composition of primitive material in the solar system, provided by the analysis of meteoritic material and interplanetary dust particles of cometary origin as collected by the {\\it STARDUST} mission."
    },
    "0911/0911.4918_arXiv.txt": {
        "abstract": "{} {The circumstellar dust shells of intermediate initial-mass ($\\sim 1$ to 8 M$_{\\odot}$) evolved stars are generated by copious mass loss during the asymptotic giant branch phase. The density structure of their circumstellar shell is the direct evidence of mass loss processes, from which we can investigate the nature of mass loss.} {We used the {\\sl AKARI Infrared Astronomy Satellite} and the {\\sl Spitzer Space Telescope} to obtain the surface brightness maps of an evolved star R Cas at far-infrared wavelengths, since the temperature of dust decreases as the distance from the star increases and one needs to probe dust at lower temperatures, i.e., at longer wavelengths. The observed shell structure and the star's known proper motion suggest that the structure represents the interface regions between the dusty wind and the interstellar medium. The deconvolved structures are fitted with the analytic bow shock structure to determine the inclination angle of the bow shock cone.} {Our data show that (1) the bow shock cone of $1 - 5\\times 10^{-5}$ M$_{\\odot}$ dust mass is inclined at $68^{\\circ}$ with respect to the plane of the sky, and (2) the dust temperature in the bow shock cone is raised to more than 20 K by collisional shock interaction in addition to the ambient interstellar radiation field. By comparison between the apex vector of the bow shock and space motion vector of the star we infer that there is a flow of interstellar medium local to R Cas whose flow velocity is at least 55.6 km s$^{-1}$, consistent with an environment conducive to dust heating by shock interactions.} {} ",
        "introduction": "Numerous observations have elucidated the magnitude and ubiquity of mass loss across the upper-right side of the Hertzsprung-Russell diagram in the fifty years since \\citet{d56} discussed the existence of blue-shifted circumstellar cores in the spectrum of the red supergiant star $\\alpha$ Her, which constituted one of the first pieces of direct evidence for high rates of mass loss. Since such high rates of mass loss among asymptotic giant branch (AGB) stars rival evolutionary timescales and substantially affect stellar evolutionary tracks, careful investigations into mass loss from these stars are indeed necessary. Moreover, since mass loss defines the boundary conditions for stellar evolutionary tracks, the rate of mass loss from these stars is hardly time-invariant. These facts have had theorists perplexed, who are struggling with the basic challenge of how to lift so much mass away from the gravitational hold of the star \\citep[e.g.][]{gh04}.  Observations made with the {\\sl Infrared Astronomical Satellite} ({\\sl IRAS}\\/) during the 1980's in the far-infrared (far-IR) have demonstrated that extended shells of evolved stars  -- the anticipated effect of continuous dusty mass loss -- were present \\citep[e.g.][]{stencel88,y93a}. During the 1990s, the {\\sl Infrared Space Observatory} ({\\sl ISO}) and ground-based IR work began to refine those results \\cite[e.g.][]{izu97,meixner99}, indicating variations in the mass loss rate over time that resulted in multiple shells and axisymmetric structures. This decade, we are fortunate to have new instruments with higher resolution and sensitivity, such as the {\\sl AKARI Infrared Astronomy Satellite} \\cite[{\\sl AKARI}, formerly known as {\\sl ASTRO-F};][]{murakami07} and the {\\sl Spitzer Space Telescope} \\citep[{\\sl Spitzer};][]{w04} that can more carefully map out the mass loss history of evolved stars.  In parallel with investigations into the mass loss history of evolved stars, evidence of interactions between the circumstellar matter and interstellar medium (ISM) around AGB stars is growing, with new observations of bow shocks around R Hya \\citep{ueta06} and Mira \\citep{martin07,ueta08}, plus theoretical considerations of the phenomena \\citep[e.g.][]{villaver03,wareing07}. This report is one of the first of several studies of well-resolved extended circumstellar dust shells of AGB stars under {\\sl AKARI} and {\\sl Spitzer} observing programs labeled ``Excavating Mass Loss History in Extended Dust Shells of Evolved Stars'' (MLHES). In this paper, we explore the detection of an arcminute-sized dust shell around the oxygen-rich AGB star, R Cassiopeiae (HD 224490; hereafter R Cas), in context of interactions between the AGB wind from R Cas and the ISM. ",
        "conclusions": "We obtained far-IR images of an oxygen-rich Mira variable R Cas at 65, 70, 90, 140, and 160$\\mu$m using {\\sl AKARI} and {\\sl Spitzer} and revealed its very extended ($2\\arcmin$ to $3\\arcmin$ radius, corresponding to 0.1 pc at its adopted distance of 176 pc), slightly elliptical ($\\epsilon = 0.3$) dust shell, in which the central star is located offset from the geometric center of the shell in the direction of the measured proper motion of the central star. We recognize a positive gradient of the surface brightness along the direction of the ``shift'' of the central star, which apparently is caused by the surface brightness enhancement along the periphery of the shell in the east side seen in deconvolved images.  Fitting of the surface brightness using data in the 3 shortest bands suggests that the observed enhancement is caused by the temperature enhancement rather than the density enhancement, prompting a need to warm up dust grains primarily on the east side of the outer rim. Given the coincidence between the direction of the proper motion of the central star and the direction of the apex of the peripheric surface brightness/temperature enhancement in the shell, we infer that the observed shell structure represents the contact surface of the AGB wind-ISM interaction which is inclined to give an overall spherical shape instead of a typical parabolic bow shock structure. This AGB wind-ISM (collisional) interaction therefore warm up dust grains in the interface regions, causing the temperature enhancement towards the windward direction of the shell. Using the maps, we also estimated the total dust mass in the shell to be $1 - 5 \\times 10^{-5}$ M$_{\\odot}$.  The shape of the observed enhancement was fitted with the analytical function for the bow shock cone to derive the inclination angle of $68^{\\circ}$. The apex vector of the bow shock cone and the space motion vector of proper motion of the central star were compared to deduce that there is an ISM flow local to R Cas that has a flow velocity of at least 55.6 km s$^{-1}$. Then, the relative velocity of the ambient ISM flow with respect to the AGB wind-ISM interface regions is at least 37 km s$^{-1}$. Therefore, such weak shocks can play a role in heating dust grains in the outermost regions of these extended dust shells around evolved stars in addition to the interstellar radiation field that is originally expected to play a role in an environment where luminosity from the central source is not enough for required dust heating."
    },
    "0911/0911.3709_arXiv.txt": {
        "abstract": "Recent X-ray observations of hot gas in the galaxy cluster MS 0735.6+7421 reveal huge radio-bright, quasi-bipolar X-ray cavities having a total energy $\\sim 10^{62}$ ergs, the most energetic AGN outburst currently known. We investigate the evolution of this outburst with two-dimensional axisymmetric gasdynamical calculations in which the cavities are inflated by relativistic cosmic rays.  Many key observational features of the cavities and associated shocks are successfully reproduced.  The radial elongation of the cavities indicates that cosmic rays were injected into the cluster gas by a (jet) source moving out from the central AGN. AGN jets of this magnitude must be almost perfectly identically bipolar. The relativistic momentum of a single jet would cause a central AGN black hole of mass $10^9$ $M_{\\odot}$ to recoil at $\\sim 6000$ km s$^{-1}$, exceeding kick velocities during black hole mergers, and be ejected from the cluster-center galaxy. Observed deviations from bipolar symmetry in the radio cavities can be caused by subsonic flows in the ambient cluster gas, but reflection shocks between symmetric cavities are likely to be visible in deep X-ray images. When the cavity inflation is complete, $4PV$ underestimates the total energy received by the cluster gas. Deviations of the cluster gas from hydrostatic equilibrium are most pronounced during the early cavity evolution when the integrated cluster mass found from the observed gas pressure gradient can have systematic errors near the cavities of $\\sim$10-30\\%. The creation of the cavity with cosmic rays generates a long-lasting global cluster expansion that reduces the total gas thermal energy below that received from the cavity shock -- even this most energetic AGN event has a net cooling effect on cluster gas. One Gyr after this single outburst, a gas mass of $\\sim 6 \\times 10^{11}$ $M_{\\odot}$ is transported out beyond a cluster radius of 500 kpc. Such post-cavity outflows can naturally produce the discrepancy observed between the cluster gas mass fraction and the universal baryon fraction inferred from WMAP observations. ",
        "introduction": "\\label{section:intro}  The hot gas in galaxy clusters emits prolifically in X-rays, and has been extensively studied by X-ray telescopes {\\it Chandra} and {\\it XMM-Newton}, revealing complex structures such as X-ray cavities, thermal filaments, and weak shocks in the intracluster medium (ICM) as in the Perseus cluster (\\citealt{fabian06}). These interesting structures are presumed to be related to the active galactic nuclei (AGNs) located at the cluster centers. Although the ICM in many clusters losses energy by radiation in a timescale much shorter than the cluster age, high-resolution X-ray spectroscopy indicates that the gas does not cool to low temperatures (e.g., \\citealt{2001A&A...365L.104P, 2003ApJ...590..207P, 2001A&A...365L..87T}; for a review see \\citealt{2006PhR...427....1P}). The deficit of cool-phase gas and young stars in cluster cores, often known as the `cooling flow' problem, suggests that the ICM is prevented from the cooling catastrophe by one or more heating sources. Currently, heating of the ICM by central AGNs, as indicated by X-ray cavities and weak shocks, is regarded as the most successful energy source to balance radiative cooling (see \\citealt{2007ARA&A..45..117M} for a recent review). Recent observations show that the AGN energies associated with X-ray cavities ($\\sim 4PV$, where $P$ is the average pressure of the gas surrounding the cavity, and $V$ is the cavity's volume) are sufficient or nearly sufficient to stop the gas from cooling in most clusters containing detectable cavities \\citep{rafferty06}. Furthermore, the mean power of these AGN outbursts increases in proportion to the cluster cooling luminosity (\\citealt{birzan04}; \\citealt{rafferty06}). This correlation suggests that AGN outbursts are triggered by the gas cooling and operates as a self-regulating feedback mechanism, which is shown to be essential in suppressing global thermal instability and thus in maintaining the ICM in the cool core state \\citep{guo08b}.  Despite these supportive clues, the detailed mechanisms that transfer the AGN jet energy to the thermal energy of the ICM are still far from clear. Numerical simulations have been extensively used to understand this process. X-ray cavities are usually thought to form at the tips of AGN-induced jets. During cavity formation, weak shock waves emerge as the gas is displaced, and heat the ICM \\citep{brueggen07}. The buoyantly-rising cavities may induce sound waves, which may be viscously dissipated in the ICM and thus heat the gas (\\citealt{ruszkowski04}; but see \\citealt{mathews06}). Radio synchrotron emission has been observed from many X-ray cavities \\citep{birzan04}, suggesting that the cavities are formed by AGN-induced cosmic rays (CRs), which may leak into the ICM and heat the gas \\citep{2008MNRAS.384..251G}.  In most previous simulations, for simplicity, it is assumed that X-ray cavities are filled with ultrahot thermal gas, but CR-filled cavities have received much less attention. The only explicit studies of the formation and evolution of X-ray cavities filled with CRs are \\citet{mathews08a}, \\citet{mathews08}, and \\citet{mathews09}, who show several interesting effects of CR-formed X-ray cavities: ICM heating due to weak shock generation, ICM cooling associated with cluster expansion and outward mass transfer, and the dynamical relationship between post-cavity X-ray filaments and extended radio lobes. However, these studies focus on low-energy cavities ($\\sim 10^{58}-10^{59}$ erg) in the Virgo cluster. In this paper, we apply the same model to study the remarkably large X-ray supercavities observed in the cluster MS 0735.6$+$7421 (hereafter MS0735), created by the most energetic AGN outburst known \\citep{mcnamara05, gitti07, mcnamara2009}. In this cluster, two approximately symmetric cavities are observed in X-ray images, each associated with relativistic gas emitting radio synchrotron emission. The energy released by this outburst is estimated to be $6.4\\times 10^{61}$ erg, assuming an enthalpy of $4pV$ per cavity \\citep{rafferty06}. Weak shocks have also been detected in MS0735; a simple spherically-symmetric shock model suggests that the age and driving energy of the shock are $t_{\\rm s}=104$ Myr and $E _{\\rm s}=5.7\\times 10^{61}$ erg, respectively \\citep{mcnamara05}. MS0735 has been studied numerically by \\citet{soker09}, who assume that the cavities are formed by slow, massive, wide jets with a total injected energy of $7.2\\times 10^{61}$ erg and a very large injected mass $\\sim 10^{10}M_{\\sun}$.  One of the main objectives of the current paper is to investigate if AGN bubbles created by cosmic rays can reproduce the observed morphology of giant X-ray supercavities in X-ray maps of MS0735. We show that the cavity morphology contains important information about the process of CR injection: CRs are likely to have been injected into the ICM at a series of locations (or continuously) along the direction of a jet as it moves outward, instead of at a fixed location (\\S~\\ref{section:morphology}). We then follow the time evolution of the ICM energies during and after cavity creation and study the energetics of X-ray supercavities and their long-term influence on the cluster gas atmosphere. The investigation of this extremely powerful AGN outburst improves our understanding of how AGN outbursts affect the ICM, and their relevance in solving the cooling flow problem.  As the largest virialized structures in the universe, clusters of galaxies are useful cosmological probes, providing one of the current best constraints on the mean matter density $\\Omega_{\\rm m}$, dark energy density $\\Omega_{\\rm DE}$, and the dark energy equation of state parameter $w$ (see \\citealt{allen08} and references therein). These constraints are usually based on accurate measurements of the X-ray gas mass fraction profile $f_{\\rm gas}(r)$. In this paper, we study how the energetic AGN outburst in MS0735 affects the gas mass fraction by producing a buoyant outflow and global expansion of the cluster thermal gas. X-ray supercavities also have a significant effect on the assumptions of hydrostatic equilibrium and spherical symmetry, which are essential in measuring the cluster's total mass profile from X-ray observations.  The rest of the paper is organized as follows. In \\S~\\ref{section:equations}, we describe basic time-dependent equations and our numerical setup. Our results are presented in \\S~\\ref{section:results}. We summarize our main results in \\S~\\ref{section:conclusion} with a discussion of the implications. ",
        "conclusions": "\\label{section:conclusion}  By conducting a suite of two-dimensional axisymmetric hydrodynamical simulations, we investigated the formation and evolution of the X-ray supercavities recently observed in the cluster MS0735, which is the most energetic AGN outburst known. We assume that the X-ray supercavities are inflated by cosmic rays injected from AGN jets, and follow the co-evolution of the cluster gas and the CRs. We then study the thermal and hydrodynamic effects of this energetic AGN outburst on the cluster MS0735. Here we briefly summarize our main findings:  1. X-ray deficient cavities and weak shocks are successfully created by AGN-generated cosmic rays. In particular, run MS-2 reproduces very well many characteristic observational features of the cluster MS0735, such as the size and location of the cavities, the location and strength of the shock.  In this run, the age and total energy of this AGN outburst are $116$ Myr and $7.6\\times10^{61}$ erg, consistent with previous estimates \\citep{mcnamara05, mcnamara2009}.  2. Assuming that X-ray cavities are generated by relativistic AGN jets, the jet that produced the north cavity in run MS-2 had a momentum $\\sim 2 \\times 10^7 M_{\\odot}c$ which must be balanced by an almost identical southerly jet ejection. Otherwise, a one-sided jet in MS0735 would cause its AGN black hole to attain a recoil velocity of $\\sim 6000(M_{\\rm bh}/10^{9}M_{\\sun})^{-1}$ km s$^{-1}$, which is large enough to completely eject a massive black hole from the center of even the most massive elliptical galaxies and their surrounding group or cluster dark matter halos \\citep{merritt04}. The radio emission from the cavities in MS0735, while clearly bipolar, is not perfectly symmetric. This could be due to the production of somewhat different cavities, perhaps at different cluster radii,  by two jets that were identical when they left the AGN. In addition, we show that small subsonic motions in the hot gas of MS0735 are sufficient to explain the axisymmetric cavity radio and X-ray emission observed. However, if the cavities formed by the two opposing jets were initially approximately equal in power and formed at similar distances from the cluster center, the shocks expanding away from each cavity would collide at the symmetry plane, producing a second, somewhat weaker shock in each hemisphere as seen in our Figure \\ref{plot1}. Efforts should be made to detect these symmetry shocks in deep X-ray images since they contain information about the identical nature of the two jets and their cavities.  3. Run MS-1, where CRs are injected at a fixed location, produces an oblate, radially-flattened X-ray cavity unlike the observed cavity which is radially elongated. In contrast, in run MS-2 which fits observations very well, the CRs are injected continuously along the jet direction as it moves away from the cluster center. The radial elongation of the cavity may also be produced if the CRs are injected at a series of locations along the jet direction.  4. During the CR injection phase, more than half of the injected CR energy is converted into thermal and kinetic energy of the ICM. The CRs lose energy thorough $pdV$ work and shock generation as they displace the cluster gas. For instance, $\\sim 63\\%$ of the injected CR energy has been lost at $t=t_{\\rm agn} = 10^7$ yrs in run MS-2. The fraction of initial CR energy lost at time $t_{\\rm agn}$ increases with the CR luminosity during the injection phase.  5. The thermal energy of the ICM responds to CR-generated cavities in several complex ways. As the cavities are inflated in MS0735, shocks are driven into the surrounding gas, heating the local ICM. Mixing by CR buoyancy may  also increase the gas entropy near the post-cavity site as relatively low entropy, CR-enriched gas near the cluster core flows outward with the CRs and is replaced by inflowing gas of higher-entropy. On the other hand, the production of new cavities causes the cluster gas to readjust outwards in the cluster, increasing its global potential energy and decreasing its thermal energy (cooling) as its density decreases. Local heating by cavity shocks dominates during the early stage of cavity evolution, whereas cooling by global expansion dominates at later times, when the total cluster thermal energy decreases below the original cluster thermal energy before the outburst. Thus, from a global perspective, the cluster gas is eventually cooled by the cavity, as previously shown by \\citet{mathews08} for a low-energy outburst in the Virgo cluster. We have now confirmed this for the cluster MS0735 which is undergoing the most powerful AGN outburst known. This ultimate allocation of cavity energy is likely to be universal for all AGN outbursts.  6. The creation of X-ray supercavities in MS0735 with CRs produces a huge, long-lasting outward migration of cluster gas to large radii. This mass outflow arises because of the global cluster expansion due to the gas displaced as the cavity forms and heating by outward-propagating shock waves driven by the expanding cavity. The mass outflow following the single energetic outburst in MS073 is large. Repeated AGN outbursts may naturally explain why the cluster baryon mass fraction in MS0735 and other similar clusters is much less than the universal baryon fraction inferred from {\\it WMAP} data, and why it increases with the cluster temperature.  7. X-ray supercavities significantly disturb the cluster gas near them, locally invalidating the assumptions of hydrostatic equilibrium and spherical symmetry. This results in systematic errors (typically around $10-30\\%$ for MS0735) in measuring the cluster mass in these regions using spherically-averaged gas density and temperature profiles.  8. Assuming that X-ray cavities are filled with relativistic CRs, $4PV$ is often used to estimate the energy released by each AGN jet. When combined with a typical (buoyant) cavity lifetime $t_{\\rm life}$, the jet power $4PV/t_{\\rm life}$ can be estimated and compared with the bolometric X-ray luminosity of the cluster gas \\citep{birzan04}. Since the total CR energy contained inside the cavity is $E_{\\rm cr}V=3PV$, where $E_{\\rm cr}$ is the average CR energy density in the cavity, about $1/4$ of the jet energy ($4PV-E_{\\rm cr}V= PV$) is delivered to the ICM during the cavity formation in these approximations. In our simulations, we directly studied the energetics of the cavity, and find that more than half of the injected CR energy goes into the ICM, i.e., the CR energy lost during the cavity formation is actually more than $3PV$. This fraction increases with the CR luminosity. Although $4PV$ provides an approximate estimation for AGN jet energy, the actual AGN energy may be a few times larger (typically $\\sim 6PV-10PV$). This is consistent with previous simulations of cavities produced by (non-relativistic) thermal jets \\citep{binney07}, which, due to the additional inclusion of a large amount of kinetic energy in thermal jets, find that AGN energy is around $15 PV$, even larger than our estimation. Our calculations further show that $4PV$ begins to decrease after the cavity formation as CRs diffuse through the cavity walls \\citep{mathews08}. More generally, if $4PV/t_{\\rm life}$ underestimates the cluster jet power, it becomes easier to match this heating with the bolometric radiative losses $L_X$ as proposed by \\citet{birzan04} and \\citet{rafferty06}.  We ignore direct interactions of CRs with the ambient cluster gas, e.g., Coulomb interactions, hadronic collisions and interactions through the generation of hydromagnetic waves, which may become important especially when the cavities are disrupted and when the CRs are mixed with the ICM \\citep{2008MNRAS.384..251G}. These interactions may provide significant heating effects for the thermal gas. One byproduct of these interactions is the generation of $\\gamma \\text{-rays}$ through the decay of neutral pions, which may be studied by the recently-launched {\\it Fermi} telescope \\citep{ando08, mathews09}. Previous cluster observations with the High Energy Stereoscopic System (HESS) have not detected these $\\gamma \\text{-ray}$ signals in the VHE ($> 100$ GeV) band, providing upper limits on $\\gamma \\text{-ray}$ fluxes and ratios of cosmic ray pressure to thermal pressure in galaxy clusters (e.g., \\citealt{domainko09}).  The center of each cavity in the cluster MS0735 is located at a projected cluster-centric radius of around $170$ kpc, which is much larger than the average projected cluster radius of $20$ kpc for X-ray cavities typically observed in clusters \\citep{birzan04}. If strong AGN outbursts having energies similar to that of MS0735 create X-ray cavities much closer to the cluster center than those observed in MS0735, they produce much stronger heating in the cluster cool core and may even fully destroy it, transforming a cool core cluster to a non-cool core cluster \\citep{guo09}. It is of great interest to investigate if metallicity peaks at cluster centers are also destroyed in this process, producing much flatter metallicity profiles as observed in non-cool core systems (e.g., \\citealt{de-grandi01}; \\citealt{baldi07}). We will thoroughly investigate these models in a follow-up paper."
    },
    "0911/0911.3153_arXiv.txt": {
        "abstract": "{The Seyfert 1 galaxy Mrk~841 was observed five times between 2001 and 2005 by the  XMM-Newton X-ray observatory. The source is well known for showing  spectral complexity in the variable iron line and in the soft X-ray excess. } {The availability of multiple exposures obtained by  the reflection grating spectrometer (RGS) cameras  allows  thorough study of the complex absorption and emission spectral features in the soft X-ray band. This paper reports on the first study of Mrk~841 soft X-ray spectrum at high spectral resolution. } { The three combined exposures obtained in January 2001 and the two obtained in January and July 2005 were analysed with the {\\small SPEX} software. } {We detect a two-phase warm absorber. A medium ionisation component  (log$\\xi$$\\sim$1.5-2.2~ergs s cm$^{-1}$) is responsible for a deep absorption feature in the unresolved transition array of  the Fe M-shell  and for several absorption lines in the OVI-VIII band, and a  higher ionisation phase with log$\\xi$$\\sim$3~ergs s cm$^{-1}$ is required to fit absorption in the NeIX-X band.  The ionisation state and the column density of the gas present moderate variation from 2001 to 2005 for both phases.  The high ionisation component of the warm absorber has no effect on the Fe K band. No significant velocity shift of the absorption lines is measured in the RGS data. Remarkably, the 2005 spectra show  emission features consistent with photoionisation  in a high-density (n$_e$$\\ge$10$^{11}$~cm$^{-3}$) gas. A prominent  OVII line triplet was clearly observed  in January 2005  and  narrow radiative recombination continua (RRC) of OVII and CVI were observed in both 2005 data sets.  A broad Gaussian line  around 21.7~$\\AA$ was also required to fit  all the data sets.  The derived radial distance for the emission lines seems to suggest that  the photoionisation  takes place within the optical broad line region of the source. } {} ",
        "introduction": "The model that  explains active galactic nuclei (AGN) as being  powered by accretion on a supermassive black hole surrounded by gas (Rees, 1984) is nowadays widely accepted. In the past decade, the advent of high-resolution X-ray spectroscopy has allowed  detailed study of ionised material flowing out along the line of sight within the central region (e.g. Kaspi et al. 2001; Kaastra et al. 2002). Multi-wavelength simultaneous observations have shown that  X-ray outflows and UV winds are tightly connected and that the absorption must take place in a material with multiple ionisation layers (Gabel et al. 2005). More than 50\\% of the Seyfert~1 galaxies are characterised by this gas, which is mostly revealed in the X-ray band through detection of narrow absorption features from ionised elements, such as  carbon, oxygen, neon, iron (Creenshaw et al. 2003, Blustin et al. 2005).  The structure of the outflowing wind and its relation to the central engine where the AGN continuum radiation originates, are not entirely clear. The absorbing gas seems to be  distributed diversely in individual sources,  its location spanning a scale of light days in, e.g., NGC~4051 (Krongold et al. 2007), to a few pc in, e.g., NGC~3783 (Behar et al. 2003).  Variability in the ionised absorbers can often offer a key to studying the change in the opacity of the gas in relation to the continuum emission.  For example,  long Chandra observational campaigns on bright Seyfert Galaxies  have yielded a variety of results on the response of the warm absorber to the source luminosity fluctuations  (Netzer et al. 2002; Krongold et al. 2005).  More recent studies have shown that the presence of emission lines in the soft X-ray band brings additional information on the circumnuclear ionised gas. In some Seyfert 1s, emission lines  from highly ionised species can originate very close to or within the broad line region of the AGN  (Costantini et al. 2007, Longinotti et al. 2008), but only in sources with high signal-to-noise  has  it been possible to relate the X-ray emitter with the X-ray absorber (NGC~3783, Behar et al. 2003).   Mrk~841 is a bright Seyfert~1 galaxy at {\\itshape z} = 0.0365  (V{\\'e}ron-Cetty \\& V{\\'e}ron, 2001) that has been observed by several X-ray observatories in the past. Spectral variability was reported in the {\\it EXOSAT, Ginga} and {\\it  ROSAT} data by  George et al. (1993), and by Nandra et al. (1995). Evidence for a broad iron emission line was found in the {\\it ASCA} spectra reported by Nandra et al. (1997).  Weaver et al. (2001) have reported marginal evidence for line variability. The analysis of broad band {\\it BeppoSAX} data revealed a soft X-ray excess and a Compton reflection component (Bianchi et al. 2001, 2004), although these authors found that warm absorption provided an equally good description of the soft X-ray data. Evidence for a moderately deep OVII absorption edge was also reported in {\\it ASCA} data by Reynolds (1997). {\\it XMM-Newton} data at CCD resolution highlighted the presence of a variable and complex narrow Fe~K$\\alpha$ line (Petrucci et al. 2002, Longinotti et al. 2004,  Petrucci et al. 2007).   Besides the study of Fe~K band, the later paper focused on the analysis of the soft excess in Mrk~841 by testing two competing  models. On one hand, the soft excess could be modelled by reflection off a photoionised accretion disc whose spectrum is relativistically blurred and smeared  (Crummy et al. 2006). A valid alternative model was proposed  by Gierlinski \\& Done (2004) who explained the soft excess as arising from  ionised absorption in a relativistically smeared wind. According to Petrucci et al. (2007), both models were statistically indistinguishable in the EPIC data of Mrk~841.  This paper presents the very first analysis of the high-resolution soft X-ray spectra of Mrk~841, and it makes use of all the available {\\it XMM-Newton} data sets, i.e., the five observations previously studied by Petrucci et al. (2007).  \\begin{table} \\caption{\\label{tab:log} {\\it XMM-Newton} observation log  for Mrk~841} \\begin{tabular}{c c c c c} \\\\ \\hline\\hline Date          & OBSID &  Exp &  $^1$count rate   & $^2$F$_{0.3-2 keV}$  \\\\ (yyyy/mm/dd)  &        -    & (ks)   & Cts/s      & (10$^{-12}$ cgs)    \\\\ \\hline\\hline \\\\ 2001/01/13   & 0112910201 &  10 & 0.63$\\pm$0.01 &  26.6$\\pm$1.0  \\\\ 2001/01/13  & 0070740101 & 12 &   0.72$\\pm$0.01 & 26.6$\\pm$1.0    \\\\ 2001/01/14  & 0070740301 & 14 &    0.71$\\pm$0.01 & 26.6$\\pm$1.0   \\\\ \\\\ 2005/01/16  & 0205340201 &  45 &   0.19$\\pm$0.01 &   7.5$\\pm$0.5  \\\\ \\\\ 2005/07/17  & 0205340401 &  30   &  0.23$\\pm$0.01 & 9.1$\\pm$0.1   \\\\ \\hline \\end{tabular}  $^1$ In RGS 2\\\\ $^2$ Not corrected for absorption; the 2001 flux refers to the global fit of 3 data sets  \\end{table} ",
        "conclusions": "In the following we summarise the main results of the first analysis of Mrk~841 high-resolution X-ray spectra: \\begin{itemize}  \\item Warm absorber: a two-phase warm absorber with log$\\xi$$\\sim$1.5-2.2~ergs~s~cm$^{-1}$ and log$\\xi$$\\sim$3~ergs~s~cm$^{-1}$  was observed in the three spectra from 2001 through 2005,  a moderate decrease in the ionisation state was observed going from the brightest flux state (2001) to the dimmer state (2005),  no outflow was detected in the data and the high ionisation component has no effect on the Fe~K band.  \\item Emission features:  narrow OVII and NeIX  emission lines were detected in the Jan~2005 spectrum, the line ratio in the OVII triplet points to an origin in a gas at high electron density (n$_e$$\\ge$10$^{11}$~cm$^{-3}$), the presence of narrow RRC from OVII and CVI is consistent with photoionisation processes, if it is assumed that the emission component and the warm absorber originate within a gas at the same ionisation state, the inferred distance for the emission lines is constrained within 10$^{16}$~cm from the central source.  \\item Broad Gaussian: a broad Gaussian line centred at 21.7~$\\AA$ was interpreted as emission from OVII,  the FWHM is fully compatible with that of the optical broad emission lines in Mrk~841. \\end{itemize}"
    },
    "0911/0911.0798_arXiv.txt": {
        "abstract": "{\\bf globular clusters; stellar populations; galaxy formation} The ensemble of all star clusters in a galaxy constitutes its \\emph{star cluster system}. In this review, the focus of the discussion is on the ability of star clusters, particularly the systems of old massive globular clusters (GCSs), to mark the early evolutionary history of galaxies. I review current themes and key findings in GCS research, and highlight some of the outstanding questions that are emerging from recent work. ",
        "introduction": "The observational evidence available to date indicates that a major star-forming epoch in a galaxy's history will generate a new set of star clusters that accompanies its population of `field' stars. Thus, it is physically meaningful to think of a \\emph{subsystem} of star clusters as consisting of all clusters formed in a given starburst, and to treat the clusters as a proxy for the stellar subpopulation formed in the same burst (see de Grijs 2010 and Larsen 2010 for more extensive treatments of cluster formation in starburst environments). The huge advantage offered by star clusters is that they are easily bright enough to be measured individually within galaxies as distant as 100 Mpc and even beyond, and in giant galaxies particularly, they can be found in large numbers (see figure 1). We can then construct \\emph{distribution functions} of such key parameters as mass, age and heavy-element abundance (metallicity) for the clusters, instead of just the luminosity-weighted averages that we get from the unresolved field-star light.  The Milky Way star cluster system (our starting point for all such work and the `baseline' comparison system for all other galaxies) separates out rather cleanly into the two classic subsystems: the \\emph{open clusters} (found throughout the disc and spiral arms along low-eccentricity orbits) and the \\emph{globular clusters} (GCs, inhabiting the Galactic bulge and halo in a roughly isotropic distribution of orbits). In addition, the GCs are distinctly older than the open clusters (although with a small range of overlap around $\\sim 8$--10 Gyr), as well as more massive and less enriched in heavy elements, indicating that they belonged to a brief early stage of rapid star formation and chemical enrichment. The open clusters, like the general population of disc field stars, are found at any age but over a more restricted range of metallicities, marking the more gradual ongoing growth of the Galaxy's disc.  But in even the nearest external galaxies (the Magellanic Clouds, M31 and the other Local Group galaxies), this convenient dichotomy disappears. The Clouds, for example, contain small numbers of classically old, massive, metal-poor GCs as well as many analogues of open clusters, but we also find numerous examples of high-mass, \\emph{young} clusters that likely resemble GCs as they would have been closer to their formation time. Investigations of star cluster systems in other galaxies reveal still richer varieties, to the stage where every part of the star cluster age/mass/metallicity three-parameter space is occupied.  For thorough reviews and perspectives on the earlier literature up to $\\sim$2005, interested readers should see Harris \\& Racine (1979), Harris (1991, 2001), Ashman \\& Zepf (1998) and Brodie \\& Strader (2006). The present article will concentrate on recent developments in GC system (GCS) studies, leading up to a list of currently challenging questions. This short and biased discussion unfortunately cannot do justice to the diversity and richness of approaches now underway, and (happily) it will be doomed to be quickly superseded by the rapid advance of both theory and data. Perhaps the most important single implication of the work in this area, however, is that the old GCs represent a common thread in early galaxy evolution, dating back to the first star formation within the pregalactic gas clouds.  \\begin{figure} \\begin{center} \\vspace{-4.7cm} \\includegraphics[scale=0.55]{m87xy.eps} \\end{center} \\caption{A region near the centre of the Virgo giant elliptical galaxy M87, showing the brightest of its total population of about 14\\ 000 globular clusters (from Harris 2009$b$). {\\it (left)} Distribution of the blue, metal-poor GCs and {\\it (right)} the more centrally concentrated distribution of the red, metal-rich GCs. The circle in each panel has a radius of 100 kpc, projected at the distance of the Virgo cluster.} \\label{m87} \\end{figure} ",
        "conclusions": ""
    },
    "0911/0911.0421_arXiv.txt": {
        "abstract": "The QCD axion is the leading solution to the strong-CP problem, a dark matter candidate, and a possible result of string theory compactifications.  However, for axions produced before inflation, symmetry-breaking scales of $f_a \\gtrsim 10^{12}$ GeV (which are favored in string-theoretic axion models) are ruled out by cosmological constraints unless both the axion misalignment angle $\\theta_0$ and the inflationary Hubble scale $H_I$ are extremely fine-tuned.  We show that attempting to accommodate a high-$f_a$ axion in inflationary cosmology leads to a fine-tuning problem that is worse than the strong-CP problem the axion was originally invented to solve.  We also show that this problem remains unresolved by anthropic selection arguments commonly applied to the high-$f_a$ axion scenario. ",
        "introduction": "\\label{s:Introduction}  The axion lies at the intersection of particle physics, cosmology and string theory, potentially playing a crucial role in each.  If it exists, the axion solves the strong-CP problem of quantum chromodynamics (QCD), fits easily into a string theoretic framework, and appears cosmologically as a form of cold dark matter.  However, the axion scenario faces a challenge as astrophysical and cosmological observations place increasingly tight constraints on axion and inflationary parameters.  The axion solution to the strong-CP problem \\cite{tHooft:1976a,tHooft:1976b} involves the introduction of a new global $U(1)$ symmetry in the early universe, known as the Peccei-Quinn (PQ) symmetry \\cite{Peccei:1977hh}, which is broken at the symmetry-breaking scale $f_a$.  The resulting pseudo-Nambu-Goldstone boson is the axion \\cite{Wilczek:1978,Weinberg:1978}.  Through the effects of instantons, the axion field's potential develops a minimum at the CP-conserving point, thus naturally explaining the lack of observed CP violation in the strong interactions.  The current constraints on the axion allow two problematic scenarios. (1) The {\\it low-$f_a$} axion, with $f_a \\lesssim 10^{12}$ GeV, produced after cosmological inflation, exists in a narrow parameter space difficult to achieve in a string theory model and requires tuning of the underlying unified theory to avoid production of cosmologically catastrophic domain walls. (2) The {\\it high-$f_a$} axion, with $f_a \\gtrsim 10^{12}$ GeV, produced prior to inflation, exists at $f_a$ scales more easily accessible to string theory models but requires very careful tuning of the inflationary Hubble scale $H_I$ and the axion's initial condition $\\theta_0$ to evade cosmological constraints.  While some models have been designed to produce string-theoretic axions with $f_a \\sim 10^{12}$ GeV \\citep[e.g.,][]{Dasgupta:2008hb}, generic string theory models produce axions with $f_a \\sim 10^{16}$ GeV, near the string scale \\citep{Svrcek:2006yi}.  Some have suggested that in this case, tuning of $\\theta_0$ may come naturally from anthropic considerations \\citep{Linde:1991km,Wilczek:2004cr,Tegmark:2005dy,Hertzberg:2008wr,Freivogel:2008qc}.  Previous works considering a high-$f_a$ axion have addressed the relationship between $H_I$ and cosmological constraints on axions.  In \\cite{Fox:2004kb}, it was pointed out that if $H_I \\gtrsim 10^{13}$ GeV, which corresponds to an inflationary energy scale $E_I \\gtrsim 10^{16}$ GeV, primordial gravitational waves would be produced at a magnitude that would be detectable by the Planck satellite through its measurement of the cosmic microwave background (CMB) tensor-scalar ratio $r$ \\citep{Kamionkowski:2002mi,Cooray:2003da}.  Such high values of $H_I$ would also result in large inflationary fluctuations to the axion's misalignment angle, so even for very small values of $\\theta_0$, the axion field would overproduce isocurvature fluctuations in the CMB.  The conclusion of \\cite{Fox:2004kb} was that a detection of inflationary gravitational waves by Planck would confirm a high $H_I$ and therefore rule out a high-$f_a$ (string-theoretic) QCD axion.  A recent paper by Hertzberg, Tegmark and Wilczek \\cite{Hertzberg:2008wr} approached this issue in a converse way, stating that the assumption of the existence of a high-$f_a$ axion predicts a {\\it low} $H_I$ and therefore the {\\it non}-detection of inflationary gravitational waves.  They argue that anthropic selection on the string landscape supports both the existence of a high-$f_a$ axion at a non-negligible abundance and the very low $\\theta_0$ required to evade cosmological constraints.  In a companion paper \\citep{MackSteinhardt:2009}, we have looked at this issue as applied to multiple string-theoretic axion-like fields.  We pointed out that string theory models that include a QCD axion (generally, a high-$f_a$ axion) can also be expected to produce other axion-like fields at similar masses, and the cumulative impact of additional axion-like fields overproduces dark matter and/or isocurvature modes unless the axion model and $H_I$ are extremely fine-tuned.  In this paper, we argue that the high-$f_a$ axion scenario is too problematic to be considered an appealing solution, {\\it whether or not} (a) inflationary gravitational waves are detected by near-term satellite missions, (b) anthropic selection determines the density of dark matter, or (c) the QCD axion arises out of string theory.  We show that since both $\\theta_0$ and the inflationary Hubble scale must be restricted to extremely low values for the axion to evade cosmological constraints, a viable axion model requires incredibly careful fine tuning.  The axion's tuning problem is so extreme that it rivals or exceeds that of the strong-CP problem that the axion was originally invented to solve.  Anthropic selection arguments do not alleviate this tuning problem: the most stringent constraints on the axion come not from the (possibly anthropic) dark matter density, but from cosmological observables whose values are not anthropically selected.  While the tuning problem in QCD that inspired the axion's proposal was troubling enough to require revision of that otherwise very successful theory, a tuning problem of equal or worse magnitude in the axion itself -- a hypothetical particle for which there is no compelling observational evidence -- could negate the axion's purpose completely.  To quantify the level of fine tuning required for an axion model to evade cosmological constraints, we use a figure of merit \\citep{MackSteinhardt:2009} defined as the product of the inflationary slow-roll parameter $\\epsilon$ and the axion misalignment angle, $\\theta_0$, which measures the angular distance of the field from the CP-conserving $\\theta=0$ point at the onset of oscillation.  The angle $\\theta_0$ is a stochastic initial condition uniformly distributed in the interval $[0,\\pi]$.  With this definition, ${\\cal F}$ measures the volume of parameter space one is restricted to by observational constraints, taking into account tuning of both the initial condition of the axion and the level of tuning of the inflationary model.  An axion with $\\theta_0\\sim \\mathcal{O}(1)$ existing in a simple, minimally tuned single-field inflationary model has ${\\cal F} \\sim 10^{-2}$.  We show in \\S \\ref{s:tuning} that cosmological constraints strongly rule out such a model, and that for the parameter space considered, the best-case scenario still allowed by constraints has ${\\cal F} \\sim 10^{-11}$.  This indicates that the tuning problem caused by trying to accommodate the existence of a high-$f_a$ axion in inflationary cosmology at best rivals the tuning required to solve the strong-CP problem without an axion [$\\mathcal{O}(10^{-10})$], and, in most cases, far exceeds it.  The paper is organized as follows. In \\S \\ref{s:constraints}, we discuss the production of the high-$f_a$ QCD axion and the origin of its cosmological constraints.  In \\S \\ref{s:tuning} we discuss the high degree of fine tuning required for high-$f_a$ QCD axion models to remain consistent with inflationary cosmology and cosmological constraints.  In \\S \\ref{s:Bayes}, we show how Bayesian model comparison disfavors the high-$f_a$ axion model in the context of inflation.  In \\S \\ref{s:anthropic}, we explain how anthropic arguments have been used to justify the choice of initial conditions for the axion field, and in \\S \\ref{s:anthropicfail} we show that anthropic selection arguments are insufficient to excuse the degree of fine tuning required for the axion to evade cosmological constraints.  We then discuss the implications of these results in \\S \\ref{s:discussion}. ",
        "conclusions": "\\label{s:discussion}  We have shown that the fine tuning of axion and inflationary parameters required to accommodate a high-$f_a$ QCD axion in the context of a simple inflationary model is at best (in a small part of parameter space) comparable to the tuning that defines the strong-CP problem and much worse over the bulk of the parameter space considered.  In \\S \\ref{s:Bayes}, we also showed that a Bayesian model comparison disfavors the axion solution.  In \\S\\S \\ref{s:anthropic} and \\ref{s:anthropicfail}, we examined the possibility that anthropic selection effects could account for the apparent tuning of $\\theta_0$.  In order to be considered successful in this context, we require that: \\begin{enumerate} \\item Anthropic selection must provide a compelling argument for the smallness of the parameter ($\\theta_0$, via the dark matter density), and,  \\label{it:compelling} \\item The anthropic observable (dark matter density) must provide a stronger limitation on the parameter than non-anthropic observables (in this case, the CMB isocurvature mode fraction). \\label{it:dominant} \\end{enumerate} Item \\ref{it:dominant} must be satisfied if we live on the anthropic boundary, which some authors have suggested \\citep{Weinberg:1987dv,Hall:2007ja}.  We assume for this purpose that Item \\ref{it:compelling} is satisfied, but we have shown in \\S \\ref{s:anthropicfail} that Item \\ref{it:dominant} is not satisfied in general.  In regions of the parameter space for which Item \\ref{it:dominant} does hold, we re-examine the question of the fine tuning of the model, neglecting the (anthropically selected) smallness of $\\theta_0$.  We find that fine tuning of $H_I$ is necessary for models in that region to evade cosmological constraints.  For $f_a \\sim m_{Pl}$, the least-tuned inflationary models compatible with axions and cosmology have $\\epsilon \\lesssim 10^{-10}$, whereas for values of $f_a \\sim 10^{16}$ GeV (as predicted in the most generic string theory models), $\\epsilon \\lesssim 10^{-12}$. This would indicate that the axion solution to the strong-CP problem merely exacerbates the problem by moving it from particle physics to cosmology, all while requiring the assumption of an (as-yet) undetected particle.   We consider this to be a strong challenge to the high-$f_a$ axion scenario in the context of inflationary cosmology.  Since string-theoretic QCD axions generically appear at high $f_a$ scales, this is, by extension, a challenge to the realization of a QCD axion in string theory.  This challenge is made even stronger by the expectation of additional axion-like fields from string theory, as discussed in a companion paper \\cite{MackSteinhardt:2009}. We suggest that one of the following modifications of the axion scenario is required. \\begin{enumerate}[A.] \\item One of the components of this scenario -- generic models of inflation, string theory (in fact any theory that requires a high $f_a$ scale), or axions -- must be abandoned.  \\label{it:choosetwo} \\item An anthropic argument against CMB isocurvature fractions higher than those observed, or against large inflationary energy scales, must be presented. \\label{it:anthrowin} \\end{enumerate}  Option \\ref{it:choosetwo} might be realized in a number of different ways.  If the axion solution to the strong-CP problem is to be preserved, one could abandon the attempt to find a high-$f_a$ QCD axion and instead assume a low-$f_a$ QCD axion model (with $N=1$ to prevent the formation of domain walls), for which a non-generic string-theoretic model could potentially be found (such as has been discussed for warped throats \\citep{Dasgupta:2008hb}).  Alternatively, high-$f_a$ QCD axions are consistent with early-universe models in which inflation does not strongly excite the axion field.  Hybrid inflation, in which inflation is driven by evolution of two scalar fields, can result in very low energy scales of inflation, with $H_I$ values as low as $\\sim 10^4$ GeV \\citep{Linde:1993cn}.  Any inflationary model with $H_I \\lesssim 10^6$ GeV is unconstrained by isocurvature modes for $f_a \\lesssim m_{Pl}$, and would require tuning of only $\\theta_0$ to evade the dark matter density constraint.  In the cyclic model \\citep{Steinhardt:2004gk}, light fields are not excited and therefore isocurvature modes and an excess of axionic dark matter would not be produced.  Failing to detect signatures of primordial gravitational waves with satellite missions such as Planck \\citep{Fox:2004kb} would lend support to this option.  Finding a new solution to the strong-CP problem that does not involve an axion is another possibility, provided the axion-like fields naturally produced in string theory models do not themselves cause cosmological problems \\citep{MackSteinhardt:2009}.  If anthropic selection arguments can be made against CMB isocurvature modes or high inflationary energy scales (Option \\ref{it:anthrowin}), this could constitute another solution to the cosmological axion problem.  In the meantime, we argue that the scenario we present here -- a QCD axion that exists during single-field inflation -- does not constitute a consistent model without some degree of modification.  As stated in the Introduction, this conclusion holds {\\it even if}: $H_I$ is not high enough for near-term detectors to observe primordial gravitational waves, the dark matter density is indeed an anthropic parameter, and/or string theory is not the origin of the QCD axion.  Fine-tuning problems have long been considered ample motivation for the revision of a theory. In fact, the axion was hypothesized not for any observational or fundamental theoretical reason, but simply to solve a fine-tuning problem in QCD. If a theory's sole motivation is the solution of a fine-tuning problem, but it then produces an even more extreme tuning problem and an unseen particle, it is not a good theory. Unless an alteration of the axion model or early universe theory that does away with the tunings is adopted, the QCD axion should be considered an unhelpful complication of particle physics and an alternative solution to the strong-CP problem should be sought."
    },
    "0911/0911.2338_arXiv.txt": {
        "abstract": "Point source searches with the IceCube neutrino telescope have been restricted to one hemisphere, due to the exclusive selection of upward going events as a way of rejecting the atmospheric muon background. We show that the region above the horizon can be included by suppressing the background through energy-sensitive cuts. This approach improves the sensitivity above PeV energies, previously not accessible for declinations of more than a few degrees below the horizon due to the absorption of neutrinos in Earth. We present results based on data collected with 22 strings of IceCube, extending its field of view and energy reach for point source searches. No significant excess above the atmospheric background is observed in a sky scan and in tests of source candidates. Upper limits are reported, which for the first time cover point sources in the southern sky up to EeV energies. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0911/0911.2612_arXiv.txt": {
        "abstract": "The number of main-sequence stars for which we can observe solar-like oscillations is expected to increase considerably with the short-cadence high-precision photometric observations from the NASA Kepler satellite. Because of this increase in number of stars, automated tools are needed to analyse these data in a reasonable amount of time. In the framework of the asteroFLAG consortium, we present an automated pipeline which extracts frequencies and other parameters of solar-like oscillations in main-sequence and subgiant stars. The pipeline uses only the timeseries data as input and does not require any other input information. Tests on 353 artificial stars reveal that we can obtain accurate frequencies and oscillation parameters for about three quarters of the stars. We conclude that our methods are well suited for the analysis of main-sequence stars, which show mainly p-mode oscillations. ",
        "introduction": "Stars with sub-surface convection zones, like the Sun, display acoustic oscillations. The stochastic excitation mechanism limits the amplitudes of the oscillations to intrinsically weak values. However, it gives rise to a rich spectrum of oscillations. The excited pressure (p) modes probe different interior volumes, with the radial and other low angular-degree modes probing as deeply as the core. This differential penetration of the oscillations allows the internal structure and dynamics to be inferred as a function of depth. Seismic studies of the Sun have indeed proven to be very powerful in inferring its internal structure.  The fact that the solar-like oscillations have such small amplitudes has made observations of these oscillations in stars other than the Sun very challenging. Over the past few years asteroseismic observations of main-sequence stars up to evolved red-giant stars have been made using Doppler velocity measurements from ground-based spectrographs, e.g. ELODIE \\citep{baranne1996}, CORALIE, HARPS \\citep{queloz2001}, UCLES and UVES \\citep{odorico2000}, and photometric space-based instruments such as WIRE \\citep{buzasi2000}, MOST \\citep{matthews2000} and CoRoT \\citep{baglin2006}. This has led to detections of solar-like oscillations in more than ten main-sequence stars and of the order of one thousand red giants. For a recent, but pre-CoRoT, review of these results see \\citet{bedding2008}. CoRoT results for main-sequence stars are presented by e.g., \\citet{michel2008,appourchaux2008,garcia2009}, while \\citet{deridder2009} and \\citet{hekker2009} present first CoRoT results for red giants.  The NASA Kepler satellite was launched successfully into an Earth trailing orbit on March 7, 2009. The satellite contains a Schmidt telescope with a 0.95-m aperture and a 105-deg$^2$ field of view, equipped with a highly sensitive photometer with a spectral bandpass from 400 to 850~nm. It is designed to continuously and simultaneously monitor 100\\,000 stars brighter than 14$^{\\rm th}$ magnitude. Kepler will be pointed towards the constellations Cygnus and Lyra during the entire mission, which has a nominal length of 3.5 years. For most stars, data will be integrated over 30 minutes, while for approximately 512 stars at a time, data with a 1-minute cadence will be obtained. Although the driving goal for the development of Kepler is to observe transiting Earth-like exo-planets, these observations are very well suited for asteroseismology, including low-amplitude solar-like oscillations. Even so, asteroseismology will contribute to the exo-planet investigations by determining radii of the planet-hosting stars, which are needed to extract the planetary radii. The radii of stars exhibiting solar-like oscillations can be obtained using the difference in frequency between modes with consecutive radial orders, i.e., the large separation ($\\Delta \\nu$). This is a measure of the sound travel time across the star, which depends on the density of the star. See \\citet{stello2009} for an overview of current approaches to obtain radii. The asteroseismic potential of Kepler is described in more detail by \\citet{jcd2008c}.  Apart from the determination of radii, the detection of solar-like oscillations in stars at different epochs along stellar evolutionary life cycles offers the prospect to test theories of stellar evolution and stellar dynamos for many stars. The input data for probing stellar interiors are the mode parameters. Accurate mode parameters are a vital prerequisite for robust, accurate inference on the fundamental stellar parameters.  We expect in total of the order of a thousand solar-like stars to be observed by Kepler in short cadence. Short-cadence oscillation data are needed to observe solar-like oscillations in main-sequence stars and subgiants as these occur at frequencies of the order of a few-hundred up to several-thousand micro-Hertz. The short-cadence (1 minute) data have a Nyquist frequency of $\\sim$ 8300 $\\mu$Hz, while the Nyquist frequency of the long-cadence (30 minute) data is $\\sim$ 275 $\\mu$Hz.  In preparation for the Kepler mission, the asteroFLAG consortium has developed automated tools to analyse solar-like oscillations in main-sequence stars \\citep[e.g.][]{huber2009,mathur2009,mosser2009}, and tested these tools extensively, see e.g., \\citet{chaplin2008,stello2009}. Automated tools are needed to cope with the large number of stars we expect to be observed with Kepler. Here, we present an automated pipeline, built to determine oscillation parameters of solar-like oscillations in main-sequence stars and subgiants, which we describe in Section 3 (some mathematical details are deferred to Appendix A). We compare the results (Section 4) of the automated analysis with the input parameters used for simulations of realistic artificial data of a few hundred stars, prepared as if they were observed with Kepler in short-cadence (Section 2). % ",
        "conclusions": "With the methods described in Section 3, we could detect solar-like oscillations and their parameters in 260 out of 353 artificial main-sequence stars and subgiants and individual frequencies in 154 stars, not further discussed here. In general, we have tried to be very cautious to reduce the number of false detections. Special care is taken in the identification of the oscillation frequency interval as an incorrectly identified frequency range will imply miss-identifications for all oscillation parameters and the background fitting. Furthermore, parameters such as $\\nu_{\\rm max}$, $\\Delta \\nu$ and amplitude per radial mode are determined with two or more (independent) methods.   The input values for $\\nu_{\\rm max}$ and $\\Delta \\nu$ are reproduced by our analyses for the majority of stars. % Also, for the amplitudes per radial mode, our values are in agreement with the input values.    The widths and thus the life times of the modes can be determined with reasonable accuracy for Kepler stars brighter than 9$^{\\rm th}$ magnitude. Our method does not contain detailed fitting, nor does it take into account that lifetimes vary for modes of different degrees. Nevertheless, for the bright stars we do find values consistent with the input relation $\\Delta$~$\\sim$~$T_{\\rm eff}^4$ \\citep{chaplin2009a}.  An accurate determination of the value of the background at the oscillation frequencies is important for two reasons. First, the background level is taken into account in the extraction of oscillation parameters such as the amplitude per radial mode, and, secondly, it provides information on the power and time scales of atmospheric parameters. From the fact that we can obtain oscillation parameters which are consistent with the input parameters, we can on the one hand infer that our background estimate in the oscillation frequency range is accurate enough to determine oscillation parameters. On the other hand we still see scatter in the determined time scale and power of the granulation around the input values. The scatter in the granulation parameters does not mean per se that the background level in the frequency interval is uncertain, as we fit a function with 6 parameters, but the parameters should be treated with caution when using them for further investigations of activity and granulation.   In terms of our sensitivity to detect oscillations, we found clear evidence that the percentage of stars for which we can detect oscillations decreases for fainter stars. Also, we found evidence that it is harder to detect solar-like oscillations in cooler ($T_{\\rm eff}$ $<$ 5500 K) main-sequence stars than in the hotter ($T_{\\rm eff}$ $>$ 5500 K) ones. This is because in the simulations (as in real data) the pulsation amplitudes scale with luminosity, with a weaker dependence on temperature.   In conclusion, the analysis tools compiled into a pipeline presented here proved to identify oscillations for a large fraction of artificial main-sequence stars. For the majority of these stars we determine oscillation parameters within 3$\\sigma$ of the input values. The existence of such pipelines will be important to be able to perform an asteroseismic analyses on the many stars we expect to become available from Kepler."
    },
    "0911/0911.5081_arXiv.txt": {
        "abstract": "We study a sample of 16 microlensed Galactic bulge main sequence turnoff region stars for which high dispersion spectra have been obtained with detailed abundance analyses. We demonstrate that there is a very strong and highly statistically significant correlation between the maximum magnification of the microlensed bulge star and the value of the [Fe/H] deduced from the high resolution spectrum of each object. Physics demands that this correlation, assuming it to be real, be the result of some sample bias.  We suggest several possible explanations, but are forced to reject them all, and are left puzzled. To obtain a reliable metallicity distribution in the Galactic bulge based on microlensed dwarf stars it will be necessary to resolve this issue through the course of additional observations. ",
        "introduction": "}  Microlensing occurs when a ``lens'' (star, planet, black hole, etc) becomes closely aligned with a more distant ``source'' star, whose image it both magnifies and distorts.  Microlensing of stars in the Galactic bulge offers the possibility of studying in detail stars that are too faint for such even with the largest existing telescopes without a microlensing boost. The lens is usually  a star  the Galactic bulge, but sometimes in the disk of the Milky Way \\citep{dominik06}.  Bulge giants are bright enough that high-dispersion spectra can routinely be obtained with 8~m class telescopes. Extensive surveys of such giants at optical wavelengths \\citep[see, e.g.][]{fulbright06} have been carried out to construct the metallicity distribution function (MDF) for the bulge, many of them in Baade's Window ($b \\sim -4^{\\circ}$),  the innermost field of relatively low reddening. \\cite{zoccali08} have presented  results of a survey of [Fe/H] in the Galactic bulge from spectra  of about  800 stars with $\\lambda/\\Delta\\lambda = 20,000$. They find a radial gradient in [Fe/H] within the bulge with the mean value going from +0.03~dex at $b = -4^{\\circ}$ to $-0.12$~dex at $b = -6^{\\circ}$, and a sharp cutoff toward higher metallicities with less than 5\\% of the sample in Baade's Window having [Fe/H] $> 0.4$~dex. \\cite{rich07}, who reach into ($l,b) = (0^{\\circ},-1^{\\circ}$) with high dispersion in the near-IR, still find a sub-solar mean metallicity of $-0.22\\pm0.01$~dex.   The ability to obtain high resolution, high quality spectra of microlensed Galactic bulge dwarfs and to carry out a detailed abundance analysis offers an independent apparently unbiased way to determine the MDF of Galactic bulge stars, as well as their detailed chemical inventory. In principle the abundance analysis of a upper main sequence dwarf is much easier and less prone to error for spectra of a fixed signal-to-noise ratio than that of a much cooler but much brighter bulge giant with a very complex spectrum full of blends and strong molecular bands. ",
        "conclusions": ""
    },
    "0911/0911.5424_arXiv.txt": {
        "abstract": "Superluminous  Supernovae  (more  than a  100  times brighter than a typical supernova (SN); e.g.  SN2006gy, SN2005gj, SN2005ap, SN2008fz,  SN2003ma) have  been  a challenge  to  explain by  standard models. For  example, pair  instability SNe which  are luminous enough seem to  have too slow a rise, and  core collapse SNe do not seem  to be luminous  enough.  We present an  alternative scenario involving a quark-nova (QN), an explosive transition of the newly born neutron star  to a  quark star in  which a second  explosion (delayed) occurs  inside  the  already expanding ejecta  of  a  normal  SN.   The  reheated SN ejecta  can radiate at  higher levels for longer  periods of time primarily  due to reduced adiabatic expansion  losses, unlike the standard SN case. Our model is successfully applied to SN2006gy, SN2005gj, SN2005ap,  SN2008fz, SN2003ma  with encouraging fits  to the lightcurves.  There  are four predictions in our  model: (i) superluminous SNe optical lightcurves  should show  a double-hump  with  the SN  hump at  weaker magnitudes occurring days to weeks  before the QN;  (ii) Two shock breakouts should be observed vis-a-vis one for a normal SN. Depending on the time delay, this would manifest as two distinct spikes in the X-ray region or a broadening of the first spike for extremely short delays; (iii) The QN deposits heavy elements of  mass number $A> 130$ at the  base of the preceeding SN ejecta.  These QN r-processed elements should be clearly visible in the late spectrum (few days-weeks in case of strong ejecta mixing) of the superluminous SN and should be  absent in  the late spectrum  of normal  (single explosion) SNe; (iv) The QN yield will also contain lighter elements (Hydrogen and Helium). We expect the late spectra to include H$_{\\alpha}$ emission lines that should be distinct in their velocity signature  from standard H$_{\\alpha}$ emission usually attributed to preshocked circumstellar material. ",
        "introduction": "At high baryon density and vanishing pressure, the ground state of bulk matter may not be iron ($^{56}Fe$), but deconfined strange quark matter (SQM) made up of up, down, and strange (u, d, s) quarks (e.g. Bodmer 1971, Witten 1984)\\footnote{Introducing strange quarks into matter costs energy because  the strange quark mass is heavier than that of the up and down quarks. That is why the $\\Lambda^0$ baryon (also made up of an u-quark, a d-quark,  and an s-quark) in vacuum is heavier than the neutron/proton. It decays to them via weak interactions. But as  the baryon density is increased, up and down quarks Fermi levels go higher and higher. One needs 2 d-quarks for every u-quark to maintain neutrality. At some point, this becomes more expensive than introducing the s-quark because the mass of the s-quark is compensated by the fact that one needs only 1up, 1 down and 1 s-quark to make a zero charge unit. So the Fermi level of the d-quark comes down (in other words, the d-quark will decay to an s-quark as density becomes higher and higher). Eventually, we end up with SQM at high density. In general weak decays are prevented in bulk degenerate matter by filled Fermi sea. E.g., the neutron is unstable in vacuum but relatively more stable in a neutron star (the direct urca does occur but the lifetime of neutron is much longer).}. In that case, once the density for a transition to the (u, d, s) phase is reached in the core of a neutron star, it is likely that the entire star is contaminated and converted into a (u, d, s) star (e.g. Haensel et al. 1986; Alcock et al. 1986). However, the SQM hypothesis faces skepticism for lack of direct evidence of its existence and theoretical uncertainty regarding the threshold density (see Appendix A for a discussion on SQM and the disaster scenario).  In 2000, researchers at CERN evaluated results from seven separate experiments (involving collisions between large nuclei) to show that a state of matter in which quarks were not confined had been created (e.g. Margetis et al. 2000). Experiments  at the  relativistic heavy-ion collider (RHIC) at Brookhaven National Laboratory have since conclusively discovered a liquid-like  state of deconfined but interacting quarks and gluons termed the Quark-Gluon Plasma (sQGP) at high temperature (eg. Nagle \\& M\\\"uller 2006). Certain features of these experiments prove beyond doubt that for a short instant one has created, and in fact significantly exceeded, the conditions required for quark deconfinement. This result, along with the theoretical advance in color superconductivity  in cold and dense matter (eg. Rapp et al. 1998; Alford et al. 1998; Alford et al. 1999) revived the topic of SQM and  formation mechanisms of quark stars and opened doors for the exploration of the post-neutron-star-pre-black-hole stage in the evolution of compact stars.  Ouyed, Dey, \\& Dey (2002; hereafter ODD)  considered the intriguing possibility that the transition from a neutron star to a quark star might in fact be an explosive one  and involves two steps in a scenario called the Quark-Nova (QN). First, the neutron star core converts to (u,d) matter. Then, as shown in  Ker\\\"anen, Ouyed, \\& Jaikumar (2005; hereafter KOJ),  the dense (u, d) core collapses to the corresponding stable, more compact (u, d, s) configuration faster than the overlying material (neutron-rich hadronic envelope) can respond. In other words, the growing quark core collapses before it engulfs the overlaying layers of the parent neutron star. While this proposition of an explosive transition awaits more detailed theoretical and numerical investigations in order to be confirmed with certitude, it nevertheless found  many interesting applications in high energy astrophysics  phenomena such as Gamma-ray Bursters (Staff, Ouyed, \\& Bagchi 2007; Staff, Niebergal, \\& Ouyed 2008), Ultra-High Energy Cosmic Rays (Ouyed, Ker\\\"anen, \\& Maalampi  2005), Re-ionization (Ouyed, Pudritz, \\& Jaikumar 2009), and Cooling of magnetars (Niebergal et al. 2009). These studies  hint to the idea that QNe do go off in the universe.  In the QN  model, core deconfinement to (u,d) might occur during or after a SN explosion, so long as the central density of the protoneutron star is high enough to induce phase conversion. Theoretical and numerical work is underway on describing the details of the conversion of ordinary nuclear matter to quark matter inside neutron stars, in an attempt at improving initial efforts (eg. Olinto  1987;  Drago et al. 2007). This is intimately connected to recent advances in our understanding of the equation of state of high-density matter and color superconductivity.  Needless to say, the propagation of the burning front associated with this phase transition is crucial in determining the energetics of the explosion and ejecta dynamics.  Staff, Ouyed \\& Jaikumar (2006), using state-of-the-art simulations showed that massive neutron stars ($\\sim 1.5 M_{\\odot}$) born with core densities a few times nuclear density  and spinning at millisecond periods are most likely to reach the quark deconfinement density quickly. These heavy neutron stars (with progenitor mass $ > 25 M_{\\odot}$) would experience a QN episode generating a second explosion following the first (i.e. the SN). The delay between the first explosion and the second one varies from a few seconds to a few years depending on the birth parameters (mass, period and magnetic field) of the parent neutron star (see Staff, Ouyed, \\& Jaikumar 2006).  Here we describe the application of such a delay for superluminous SNe and briefly discuss the  general implication to  high energy astrophysics.    This proceedings note is organized as follows. In section 2, we describe the essential mechanism and features of the QN. In section 3, we describe in detail the efficacy of the dual-shock QN model in describing the outstanding features of a host of superluminous SNe, with emphasis on the successful interpretation of photometric and spectroscopic results. Section 4 contains an overview of QN dynamics, in particular, the explosive photon fireball that opens up the connection to high-energy astrophysical phenomena as well as r-process nucleosynthesis of the heavy elements in a unified way. ",
        "conclusions": "\\subsection{The photon fireball and explosive astrophysics} \\label{sec:fireball}  In core-collapse SNe, neutrinos being the lightest and most weakly interacting particle, carry away 99\\% of the star's binding energy  and drive the explosion. In QNe, neutrinos emitted from the quark core have long diffusion timescales, of order 10-100 ms for $T\\sim 20$ MeV (KOJ) and cannot escape before the entire star converts to $(u,d,s)$ matter. With a hadronic crust, the heating effect of the neutrinos implies that the mean free path in the crust is of the order of meters, and mass ejection is of the order of $10^{-5}M_{\\odot}$. Thus, neutrinos are not very efficient in depositing their energy in the outer layers of the star in a short enough time.  For a QN, the suitable agent of explosion are the photons, since the temperature of the quark core is large enough at the time of formation ($\\sim 20-50$ MeV) to sustain large photon emissivities. If quark matter is in a normal (i.e., non-superconducting) state ($T\\geq 50$ MeV), the photon emissivities are extremely large since $T>\\hbar\\omega_p/k_B$ where $w_p\\sim (20-25)$ MeV is the plasma frequency of normal quark matter (Usov 2001). The mean free path is small enough to thermalize these photons inside quark matter, and smaller still in the hadronic envelope ($\\omega_p({\\rm hadronic})\\sim 100$ MeV) so that energy deposition by photons is highly efficient.  If quark matter is superconducting and in the CFL phase, the result is similar. A huge build-up of thermal photons occurs because of medium effects. At $T=0$, a residual but exact $U(1)$ symmetry ensures that CFL photons travel freely in the quark medium. However, at $T\\sim (5-50)$ MeV, these photons thermalize due to electromagnetic interactions with charged light Goldstone mesons (Vogt, Rapp \\& Ouyed 2004; Jaikumar, Prakash \\& Sch\\\"afer 2002) and their mean free path is of order fermis. The resulting emissivity saturates the black body limit at $T\\sim 50$ MeV. Compared to the neutrino flux from hot quark matter, the photon flux is from 1-3 orders of magnitude higher for temperatures in MeV-tens of MeV. It follows that energy deposition in the crust is much more efficient for photons than neutrinos. Even a few percent of the photon energy, when deposited in the thin crust of the star, will impart a large momentum to it, leading to strong and ultra-relativistic mass ejection (Ouyed\\&Leahy 2008). Up to $10^{-2}M_{\\odot}$ can be ejected by the photon fireball, making the QN a highly explosive and luminous astrophysical phenomenon. As we have argued on the basis of observations as well, this lends itself naturally to an explanation of the superluminous SNe. Below, we discuss some astrophysical consequences and predictions that follow from the QN.  \\begin{figure}[t!] \\begin{center} \\epsfig{file=figure4.ps,height=5.0in,angle=-90.0,} \\caption{The r-process element abundance distribution as a function of mass number in a QN: parameters of the simulation are fixed to expansion timescale $\\tau\\sim 0.1$ms, 50\\% local energy deposition from nuclear reactions, and mass averaged electron fraction of ejecta $\\langle Y_e\\rangle$=0.03 and 0.12 as shown (from Jaikumar et al. (2007)). Note the peaks at Hydrogen and Helium. The solar heavy-element abundance is also shown for comparison.} \\label{fig:rprocess} \\end{center} \\end{figure}     \\subsection{The QCD-Astrophysics connection}  The discovery of a QN would have tremendous implications for Quark Matter and astrophysics. Among the immediate implications:  \\begin{itemize}  \\item It would  confirm that SQM (in a color superconducting, CFL-like, state) exists in a stable state in the universe.  \\item The  separation between the humps in the lightcurve (i.e. $t_{\\rm delay}+t_{\\rm prop.}$) is an estimate of the time delay between the two explosions, $t_{\\rm delay}$.  This delay is intimately linked to the value of the deconfinement density   (see Staff, Ouyed, \\& Jaikumar 2006) which is of relevance to QCD.  \\item The additional energy deposited by the QN ejecta is an estimate of the energy release during a QN explosion. Since the gravitational energy is transferred to internal and shock energy, superluminous SNe provide a window into the energetics of the phase transition in cool and dense matter.  \\item QN do not require the existence of extreme mass progenitors (see Appendix C). If the proposed observations reveal the expected sequence of spectroscopic features, our views  of the final states of massive stars have to be revised. Even a small fraction of progenitors with masses in the 25-40 $M_{\\odot}$ range undergoing a double explosion is enough to account for the rate of observed superluminous SNe.  \\item The first results on the production of heavy elements (with $A > 130$) in a QN described in Jaikumar et al. (2007; see also Charignon et al. 2009) were meaningfully compared to the striking r-process pattern seen in several old metal-poor stars. The observed delamination of a QN would make this connection more robust with immediate implications to our modeling of chemical evolution of galaxies and of the IGM. As shown in Ouyed, Pudritz, \\& Jaikumar (2009), with  QNe, a normal initial mass function (IMF) for the oldest stars can be reconciled with the mean metallicity of the early IGM post-reionization. It would be remarkable if the solution to the longstanding problem of heavy-element ($A >130$) nucleosynthesis  and the early IGM metallicity  enrichment is intimately linked to the conclusive spectral evolution of a dual-shock QN.  \\item Yasutake et  al. (2005) and Staff  et al. (2006)  have determined that the  evolutionary transition  from rapidly  rotating neutron  stars to quark  stars  due   to  spin-down  can  lead  to   an  event  rate  of $10^{-4}$-$10^{-6}$ per  year per  galaxy. Similar rates  were derived from studies of QNe contributions  to r-process material in the Galaxy by Jaikumar et al. (2007) who  estimated that 1 out every 1000 neutron stars might  have undergone  a QN.  Since  the Galaxy  likely contains about  $10^{8}$ neutron  stars this  suggests  an average  QN rate  of $10^{-5}$  per year  per galaxy.   Interestingly, the  fraction  of SN progenitors with  mass greater than $60M_{\\odot}$ can  be estimated as $\\sim 5\\times  10^{-3}$, using the Scalo (1986)  initial mass function for $M> 8 M_{\\odot}$.  Using a  SN rate of $\\sim 10^{-2}$ per year per galaxy,  we get $\\sim  5\\times 10^{-5}$  per year  per galaxy  for the explosion rate  of massive  star ($> 60  M_{\\odot}$). This  is, within uncertainties, the same as the QN rate.  \\end{itemize}   \\subsection{Four predictions}  \\begin{itemize}  \\item Two SN-like lightcurves are seen, with  the first a normal SN lightcurve, and  the second  brighter  one delayed  from  the first  by  by a  few days to a few weeks. The shape and peak  magnitude of the second lightcurve depend on the delay time (with generally lower peak for shorter delays) and also on the SN envelope mass and SN ejection velocity.  \\item Two shock breakouts should be observed for the QN vis-a-vis one for a normal SN. Depending on the time delay, this would manifest as two distinct spikes in the X-ray region or a broadening of the first spike for very short delays ($< 1$ day) between the two explosions. The double shock breakout specific to the QN model is currently being investigated in more details.   \\item Another tell-tale signature that would clearly distinguish a QN from a SN would be the detection of radioactive elements with lifetimes of order several days or few weeks. Such short-lived elements cannot realistically be detected in SN, since their decay times are too short compared to the gamma transparency of the SN ejecta. But the QN ejecta is relativistic and plows into the SN ejecta at very large speeds, producing the dual shock. If subsequent mixing effects are strong, there is the possibility that some of the short-lived heavy elements could be detected against the background of the late-time (decaying) light curve. Based on the r-process calculations of Jaikumar et al. (2007), we have identified some promising ``signature'' elements of the dual QN, including but not limited to:$^{156}$Eu$_{63}$, $^{166}$Dy$_{66}$, $^{175}$Yb$_{70}$, $^{183}$Ta$_{73}$, $^{191}$Os$_{76}$, $^{193}$Ir$_{77}$, $^{223}$Ra$_{88}$, $^{225}$Ac$_{89}$. (All of them are $\\gamma$-active with lifetimes of few days-few weeks): Specifically, gamma decay lines from these elements should be seen in late  spectra of superluminous SNe once the photosphere has receded deep into the ejecta.  \\item As seen from Figure \\ref{fig:rprocess}, the QN yield will also contain hydrogen and helium. We expect the spectra in dual-shock QNe (superluminous SNe) to include  H$_{\\alpha}$ emission lines  that should be distinct (delayed) from standard H$_{\\alpha}$ emission usually attributed to preshocked circumstellar material. There are a few caveats attached to this fourth prediction: (i)  As can be seen from Figure \\ref{fig:rprocess}, the hydrogen (and helium) abundance is no higher than the heavy element peaks. Since H and He are two elements whereas the band from A=100 to 250 has 150 elements, H would comprise only 1\\% by number (i.e. $ \\sim 10^{-4}$ by mass). Thus there will be no more than   $\\sim 10^{-4}M_{\\odot}$ of hydrogen in the QN ejecta. Once mixed with the preceding SN material the H might be hard to detect despite its   distinctive  velocity signature (with line-width $\\propto v_{\\rm QN}$); (ii)  If  the H is moving fast outward in a spherical shell, the line-of-sight Doppler velocity will range from $-v_{\\rm QN}$ to $+v_{\\rm QN}$, so that the line is spread over a wide frequency range and thus be very faint; (iii) if there is too much H$_{\\alpha}$ emission in the SN associated galaxy, the QN H$_{\\alpha}$ emission line  would not be detectable, even if it is at high velocity.   \\end{itemize}  Finally, it should be emphasized that there  are several active research fronts on understanding the possible implications of the  QN to other outstanding astrophysical phenomena (e.g.   gamma-ray bursters, ultra-high energy cosmic rays, re-ionization ect ...). Some of these  phenomena have been listed in the  ``Turner report\" which was entitled ``{\\it Connecting Quarks with the Cosmos: Eleven Science Questions for the New Century}\" (Committee On The Physics Of The Universe, 2003). It  would  be   remarkable  for  example  if  the   solution  to  some of these longstanding  problems  could find  answers in the discovery of stable quark matter in the universe via a QN. A Quark-Nova is a phenomenon that would naturally connect  quarks with the cosmos; the most powerful explosions in the universe signaling the birth of the tiniest of particles.  \\bigskip R.O. is grateful to  organizers of the CSQCDII conference. R. O. thanks the Canadian Institute for Theoretical Astrophysics (CITA) for  hospitality and support  during the completion of this work.  The research of R. O. and D. L. is supported by an operating grant from the National Science and Engineering Research Council of Canada (NSERC).   \\begin{appendix}"
    },
    "0911/0911.5612_arXiv.txt": {
        "abstract": "{Massive protostars have associated bipolar outflows with velocities of hundreds of km~s$^{-1}$. Such outflows can produce strong shocks when interact with the ambient medium leading to regions of non-thermal radio emission.} {We aim at exploring under which conditions relativistic particles are accelerated at the terminal shocks of the protostellar jets and can produce significant gamma-ray emission.} {We estimate the conditions necessary for particle acceleration up to very high energies and gamma-ray production in the non-thermal hot spots of jets associated with massive protostars embedded in dense molecular clouds.} {We show that relativistic Bremsstrahlung and proton-proton collisions can make molecular clouds with massive young stellar objects detectable by the {\\it Fermi}{} satellite at MeV-GeV energies and by Cherenkov telescope arrays in the GeV-TeV range.} {Gamma-ray astronomy can be used to probe the physical conditions in star forming regions and particle acceleration processes in the complex environment of massive molecular clouds.} ",
        "introduction": "Massive stars are formed in the dense cores of massive cold clouds (Garay \\& Lizano 1999, and references therein). The accumulation of gas in the core might proceed through previous stages of fragmentation and coalescence with the progressive result of a massive protostar that then accretes from the environment (e.g. Bonnell et al. 1997, Stahler et al. 2000)  or through direct accretion onto a central object of very high mass (e.g. Rodr\\'\\i guez et al. 2008 -RMF08-; see Shu et al. 1987 for the basic mechanism). In any case, the prestellar core is expected to have angular momentum, which would lead to the formation of an accretion disk. The strong magnetic fields inside the cloud that thread the disk should be pulled toward the protostar and twisted by the rotation giving rise to a magnetic tower, with the consequent outflows, as shown by numerical simulations (e.g. Banerjee \\& Pudritz 2006, 2007).  Evidence of molecular outflows is found through methanol masers, which are likely associated with shocks formed by the interaction with the external medium (e.g. Plambeck \\& Menten 1990). However, the most important evidence for outflows comes from the detection of thermal radio jets. These jets are observed to propagate through the cloud material along distances of a fraction of a parsec (e.g. Mart{\\'\\i} et al. 1993 -MRR93-). At the end point of the jets, hot spots due to the terminal shocks are observed in several sources. In a few cases, these hot spots are clearly non-thermal, indicating the presence of relativistic electrons that produce synchrotron radiation (e.g. Araudo et al. 2007 -ARA07-, 2008).  A population of relativistic electrons in the complex environment of the massive molecular cloud in which the protostar is being formed will produce high-energy radiation through a variety of processes: inverse Compton (IC) scattering of infrared (IR) photons from the cloud, relativistic Bremsstrahlung, and, if protons are accelerated at the shock as well, inelastic proton-proton ($pp$) collisions. If such radiation is detectable, gamma-ray astronomy can be used to shed light on the star forming process, the protostar environment, and cosmic ray acceleration inside molecular clouds.  This work is devoted to discuss under what conditions the terminal shocks of jets from massive protostars can efficiently accelerate particles, and produce gamma rays that may be detectable by the {\\it Fermi} satellite and Cherenkov telescopes in the near future. The model developed for the calculations is essentially different from the phenomelogic model presented by ARA07, since now the dynamics of the jet termination region is characterized, the shock power estimated, the conditions for particle acceleration analyzed, and the role of hydrodynamical instabilities for non-thermal radiation explored. In short, the acceleration and emission are consistently modeled together with the hydrodynamics in a more physical scenario.  Our study is partially based on early works on non-thermal emission in young stellar objects (YSO), as those by Crusius-Watzel (1990) and Henriksen et al. (1991), but we develop further some important aspects of the hydrodynamics-radiation relation, and focus on massive YSOs and the feasibility of their detection with the present observational facilities. ",
        "conclusions": "Romero (2008) pointed out that the detection of massive protostars at gamma-ray energies would open a new window to star formation studies. The detection of the cutoff in the SED would give important insights on the acceleration efficiency in the terminal shocks of the outflows. The SED can also shed light on the densities, magnetic fields, velocities, and diffusion coefficients in the shocked regions. Although we do not expect that massive protostars should be among the bright sources detected by {\\it Fermi} (Abdo et al. 2009), our calculations show that they could show up in further analysis of weaker sources after few years of observation. The emission levels above 100~GeV, close to 0.01~Crab, could be detectable by current and future Cherenkov telescopes for observation times moderately longer than 50~hours.  However, not only massive protostars, but also the regions in which they form, may be gamma-ray emitters. As mentioned in Sect.~\\ref{mod}, some amount of the highest energy particles may escape to the cloud far upstream the bow shock. It is hard to estimate the fraction of electrons and protons that would be released in the cloud, which depends strongly on the diffusion coefficient of the pre-shock cloud medium and the bow shock velocity and size. However, they might carry a non-negligible fraction of $L_{\\rm nt}^{e,p}$ if $\\Gamma\\sim 2$. In fact, in the case that $L_{\\rm nt}^{e,p}\\sim 0.01\\,L_{\\rm j}$ were in the cloud, massive YSOs may inject an amount of protons well above the average galactic level at several hundreds of GeV, and the radiation resulting from $pp$ may be detectable (for a general case, see Aharonian \\& Atoyan 1996), competing with that produced in the lobe itself. For leptons, the emission at high energies may be relevant for low magnetic fields, i.e. when the maximum energy is determined by diffusive escape and dominant relativistic Bremsstrahlung in the cloud. The spectrum of the gamma rays, generated by $pp$ collisions for protons and relativistic Bremsstrahlung for electrons, should be very hard since only the highest energy particles escape, peaking at $\\epsilon\\la E_{\\rm max}$. The cloud synchrotron emission should be quite diluted and dominated by the lobe.  A clustering of gamma-ray sources should be present in regions with large molecular clouds and star formation, as already inferred from EGRET data (e.g. Romero et al. 1999). The accumulation of cosmic rays accelerated in the radio lobes into the molecular cloud can produce extended gamma-ray sources. These radio lobes may be difficult to detect. Neither UV nor hard X-ray counterparts related to thermal Bremsstrahlung produced in the shock downstream regions are expected to be observed from these sources because of the large absorption and/or low emission levels. Deep inside the cloud, even radio emission may be missing due to strong free-free absorption, so the exact number of accelerators could be hard to estimate. Also, cosmic-ray re-acceleration inside the clouds due to magnetic turbulence (e.g. Dogiel et al. 2004) could result in stronger sources. Therefore, the combined effect of several protostars deeply embedded in giant clouds might be responsible for GeV-TeV sources found in star forming regions by EGRET, {\\it Fermi}, {\\it AGILE} and Cherenkov telescopes. We conclude that massive clouds with high IR luminosities and maser emission (tracers of massive star formation) deserve detailed study with {\\it Fermi} and ground-based Cherenkov telescopes."
    },
    "0911/0911.3728_arXiv.txt": {
        "abstract": "We investigate the fraction of starbursts, starburst-AGN composites, Seyferts, and LINERs as a function of infrared luminosity ($L_{\\rm IR}$) and merger progress for $\\sim$500 infrared-selected galaxies.  Using the new optical classifications afforded by the extremely large data set of  the Sloan Digital Sky Survey,  we find that the fraction of LINERs in IR-selected samples is rare ($< 5$\\%) compared with other spectral types. The lack of strong infrared emission in LINERs is consistent with recent optical studies suggesting that LINERs contain AGN with lower accretion rates than in Seyfert galaxies. Most previously classified infrared-luminous LINERs are classified as starburst-AGN composite galaxies in the new scheme. Starburst-AGN composites appear to ``bridge\" the spectral evolution from starburst to AGN in ULIRGs.  The relative strength of the AGN versus starburst activity shows a significant increase at high infrared luminosity. In ULIRGs ($L_{\\rm IR}>10^{12}L_{\\odot}$), starburst-AGN composite galaxies dominate at early $-$ intermediate stages of the merger, and AGN galaxies dominate during the final merger stages. Our results are consistent with models for IR-luminous galaxies where mergers of gas-rich spirals fuel both starburst and AGN, and where the AGN becomes increasingly dominant during the final merger stages of the most luminous infrared objects. ",
        "introduction": "Luminous Infrared Galaxies (LIRGs: $L_{\\rm IR}>10^{11}L_{\\odot}$\\footnote{$L_{\\rm IR} \\equiv L(8-1000 \\mu{\\rm m})$; \\citep[see][]{sanders96}}) were first discovered in small numbers $\\sim$4 decades ago \\citep{low68,kleinmann70a,kleinmann70b,becklin71,becklin72,rieke72}.  The importance of these objects to galaxy evolution was made more clear following the first all-sky survey carried out by the {\\it Infrared Astronomical Satellite} \\citep[{\\it IRAS}:][]{neugebauer84}. \\citet{soifer87} found that the space density of infrared (IR)-selected LIRGs in the local Universe ($z < 0.1$) rivaled that of the most powerful optically-selected starburst and Seyfert galaxies at similar bolometric luminosity, and that the most luminous objects $-$ ultraluminous infrared galaxies  (ULIRGs: $L_{\\rm IR}>10^{12}L_{\\odot}$) $-$ had similar space densities and bolometric luminosities as optically-selected quasi-stellar objects (QSOs).  There have been numerous studies related to the origin and evolution of U/LIRGs\\footnote{Previous studies of the properties of infrared galaxies versus log($L_{\\rm IR}/L_{\\odot}$) often divide the galaxy samples into decade luminosity bins and use the terms moderate, luminous, ultraluminous, and hyperluminous to refer to the decade ranges 10$-$10.99, 11$-$11.99, 12$-$12.99, and 13$-$13.99, respectively.  We follow this convention here, and use U/LIRGs when we wish to refer to all galaxies with log($L_{\\rm IR}/L_{\\odot}$)= 11$-$12.99 .}, and most now seem to agree that strong interactions and mergers of gas-rich galaxies are the trigger for the majority of the more luminous LIRGs \\citep[see the review by][]{sanders96}. The merger fraction increases with IR luminosity and approaches 100\\% for samples of ULIRGs \\citep[{\\it e.g.},][]{sanders88a, kim95a, clements96, farrah01}.  In the complete sample of IRAS 1{\\ts}Jy ULIRGs by \\citet{kim95a}, 117 out of 118 galaxies show strong signs of tidal interaction \\citep{veilleux02}.  There is less consensus on the nature of the power source of U/LIRGs.   It is clear that the IR luminosity in U/LIRGs can derive from dust reprocessing of either extreme starburst activity, Active Galactic Nuclei (AGN), or a combination of the two. Studies of moderate to large samples of U/LIRGs \\citep[{\\it e.g.},][]{kim95a,veilleux95,goldader95,genzel98} indicate that the dominant power source in lower luminosity LIRGs is an extended starburst, and that an AGN often makes an increasing contribution to the bolometric luminosity in the more luminous sources with obvious energetic point-like nuclei. However, different studies of the same objects disagree on the relative contributions of starburst and AGN activity to the bolometric luminosity, in particular for the ULIRGs where the dominant energy source powering their extremely luminous and compact nuclear cores continues to be the subject of intense debate \\citep[c.f.,][]{joseph99,sanders99}.  Although numerous studies at various wavelengths continue to be carried out to determine the energy source of ULIRGs \\citep[{\\it e.g.},][]{tran01,farrah03,farrah07,lip03,ptak03,ima07}, determining the relative contribution of starbursts and AGN within individual galaxies is still difficult.   One of the commonly proposed merger scenarios for ULIRGs \\citep[{\\it e.g.},][]{sanders88a,kim95b,farrah01,lip03,dasyra06} is based on the \\citet{toomre72} sequence in which two galaxies lose their mutual orbital energy and angular momentum to tidal features and/or an extended dark halo and coalesce into a single galaxy. Tidal interactions and associated shocks are thought to trigger star formation \\citep[{\\it e.g.},][]{bushouse87, kennicutt87, liu95, barnes04} which heats the surrounding dust, producing strong far-infrared (FIR) radiation. The FIR radiation rises to an ultra-luminous IR stage powered by starbursts and/or dust-shrouded AGN. As starburst activity subsides, the merger finally evolves into an optically bright QSO. In this scenario, ULIRGs plausibly represent a dust-shrouded transition stage that leads to the formation of optical QSOs \\citep[{\\it e.g.},][]{sanders88b,dasyra06,kawa06,zauderer07}.   A key element in testing the above scenario is to clarify the power source behind the strong IR emission, and the relationship between this power source and the evolutionary stage of the interaction. Comprehensive studies on large IR-selected samples are crucial to this analysis. Notable examples of such samples are the IRAS Bright Galaxy Survey \\citep[BGS: ][]{veilleux95}, the IRAS 1{\\ts}Jy ULIRGs sample \\citep[1{\\ts}Jy ULIRGs: ][]{kim95a, kim98, veilleux02}, and the Southern Warm Infrared Galaxy sample \\citep[SW01: ][]{kewley01a}.  Most previous studies use standard optical spectral diagnostic diagrams to classify the dominant power source in emission-line galaxies \\citep[]{baldwin81,veilleux87}. These diagrams are based on four optical emission line ratios that are sensitive to the hardness of the ionizing radiation field. More recently, the Sloan Digital Sky Survey (SDSS) has revolutionized these classification schemes by revealing clearly formed branches of star-forming galaxies, Seyferts, and Low Ionization Narrow Emission-line Region Galaxies (LINERs) on the diagnostic diagrams, for the first time \\citep[]{kewley06}. Kewley et al. shows that many galaxies previously classified as LINERs lie along a well-defined mixing branch from star-forming galaxies to Seyfert galaxies.  In light of this new classification scheme, we investigate the new spectral classification of IR galaxies as a function of IR luminosity and merger progress. We describe our sample selection and derived quantities in \\S~\\ref{sample}. The results are presented in \\S~\\ref{results}. We discuss the results in \\S~\\ref{discussion} and summarize in \\S~\\ref{summary}. For convenience of comparison with the old 1{\\ts}Jy ULIRG analysis by \\citet{veilleux99,veilleux02}, we adopt a Hubble constant of  H$_{0}$ = 75~ km~s$^{-1}$ Mpc$^{-1}$, and $\\Omega_{m} = 0.27$, $ \\Omega_{\\Lambda} = 0.73$ throughout the paper. ",
        "conclusions": "\\label{discussion}  \\subsection{IR-selected LINERs}\\label{liners}  We find that IR luminous LINERs are rare compared to other optical spectral types, even in the low-redshift BGS sample.  A majority of the previously classified IR-luminous LINERs are now reclassified as starburst-AGN composites in the new SDSS optical spectroscopic classification scheme. This result prompts us to discuss the following issues: \\ (1) The relationship between the new class of LINERs and starburst-AGN composite objects. \\ (2) The nature of IR-luminous LINERs compared with the LINERs from the SDSS. We discuss each of these issues separately below.  (1)  \\citet{kewley06} have shown that the host properties of composite galaxies are intermediate between those of AGN and high metallicity star-forming galaxies, and that LINERs form a unique, coherent class that is distinct from Seyferts, star-forming galaxies and composites. Specifically, LINERs are older, more massive, less dusty and less concentrated than Seyfert galaxies. These optical results highlight the intrinsically different nature of LINERs and composites.  Here we investigate whether similar differences exist in popular mid-IR diagnostics.  In Figure~\\ref{fig:sturm}, we investigate the positions of galaxies on the [FeIII]~$26.0 \\mu m$/[OIV]~$25.9 \\mu m$ versus [OIV]~$25.9 \\mu m$/[NeIII]~$12.8 \\mu m$ diagnostic diagram.  This diagram separates star-forming galaxies from those dominated by an AGN.   Previously, IR-luminous LINERs occupied a region in between starburst and Seyfert galaxies, whereas IR-faint LINERs occupied a region offset from the starburst-Seyfert sequence \\citep{sturm06}.    In our new classification scheme,  only 1 out of the 16 IR-luminous LINERs are classified as LINERs (the rest 15 objects are 8 composites, 1 Seyfert 2s, and 6 ambiguous classes between HII and Seyfert 2s), while 10 out of the 17 IR-faint LINERs remain LINERs (the rest 7 objects are ambiguous classes between LINERS and Seyfert 2s).    Figure~\\ref{fig:sturm} indicates that the majority (9/13) of the composites lie in along a mixing sequence connecting star-forming galaxies and Seyfert{\\it s}.  By contrast, most (8/10) of the LINERs lie offset from the mixing sequence. These results are consistent with our view of optically selected LINERs as a different class of objects from starburst-AGN composites.   \\begin{figure}[!ht] \\epsscale{1.56} \\plotone{f14.eps} \\caption{Mid-infrared diagnostic from \\citet{sturm06}, with objects re-classified using the Ke06 new scheme. The data were provided by Sturm, E.  The symbol meanings are labeled on the plot.  The bracket after the ambiguous class indicates between which classes the ambiguity lies. } \\label{fig:sturm} \\end{figure}     (2) There is growing evidence that a large fraction of optically selected LINERs are low luminosity AGN with low accretion rates \\citep[e.g.,][]{ho99, quataert01, barth02}.   The SDSS LINER population studied by \\citet{kewley06} supports this AGN nature. They show that at fixed $L$\\oiii/$\\sigma^4$ (an indicator for the black hole accretion rate), all differences between Seyfert and LINER host properties disappear.  With this interpretation of LINERs, it is not surprising that we find so few IR-luminous LINERs because ULIRGs tend to harbor AGN that favor high-accretion rates \\citep{weedman83, filipenko85, sanders88a}. Theory supports this interpretation; numerical simulations predict that the black hole accretion rate rises rapidly during the merger process when enormous quantities of gas flows into the central regions of the merging galaxies  \\citep[e.g.,][]{tani99, hopkins05}.  What about the few IR-luminous LINERs that do exist in our samples?  Do they belong to the same class as optical LINERs?   In Table~\\ref{tb9} we list all the possible LINERs found in our combined 1{\\ts}Jy ULIRG + BGS + SW01 sample. Altogether there are 9 objects in the three samples that fall into our LINER class,  however, only 4/9 of these LINERs (Arp 220 in the 1{\\ts}Jy ULIRG sample, NGC 6240 in the SW01/BGS samples and another 3 in the BGS sample) can be ``safely\" classified as LINERs.  The other 5/9 LINERs lie near the 0.1{\\ts}dex error region of the Seyfert-LINER boundaries, and may be Seyfert objects or intermediate between Seyfert and LINER types.   This can be clearly seen in Figure~\\ref{fig:linear1}, where all 9 LINERs are over-plotted with the SDSS data from \\citet{kewley06} on the \\sii/H$\\alpha$ versus \\oiii/H$\\beta$ and  \\oi/H$\\alpha$ versus \\oiii/H$\\beta$ diagnostic diagrams.  \\begin{figure}[!ht] \\epsscale{1.4} \\plotone{f15.eps} \\caption{ The positions of the 9 LINERs on the  \\sii/H$\\alpha$ versus \\oiii/H$\\beta$ and \\oi/H$\\alpha$ versus \\oiii/H$\\beta$ diagrams over-plotted on the SDSS data from \\citet{kewley06}. The 3 IR-luminous LINERs in the 1{\\ts}Jy ULIRG sample are shown as squares, the 3 IR-luminous LINERs in the SW01 sample are shown as diamonds, and the 3 crosses are the IR-luminous LINERs from the BGS sample. } \\label{fig:linear1} \\end{figure}  In Figure~\\ref{fig:linear1}, we indicate the positions of two well-studied IR-luminous LINERs: NGC 6240 and Arp 220. Recent high resolution X-ray and radio data for NGC 6240 \\citep{lira02, komossa03, iono07} and Arp 220 \\citep{clements02, ptak03, downes07} provide convincing evidence for the existence of AGN in these 2 objects.   However, the LINER emission may be excited by other ionization processes.  Both of these two LINERs show evidence for starburst-driven superwinds and/or shocks that may dominate the EUV emission of these galaxies \\citep{heckmanet87, tecza00,lira02, mc03, iwa05}.  Further investigation into the EUV power source of IR-luminous LINERs is required to draw robust conclusions about the nature of the few IR-luminous LINERs in our samples.  To summarize, the rarity of IR-luminous LINERs is consistent with the picture that most LINERs are (a) excited by low Eddington rate black holes and (b) reside in galaxies with an aged stellar population.   The special cases of Arp 220 and NGC 6240 may indicate a contribution from different LINER ionization sources: either starburst superwinds or shocks driven by galaxy collisions.  \\subsection{Are starburst-AGN composites playing the role of ``bridging\" in LIRGs and ULIRGs? }  In Fig.~\\ref{fig:sampleall}, we show 5 different galaxy samples on the \\nii/H$\\alpha$ versus \\oiii/H$\\beta$ diagnostic diagram. The new Ke06 classification boundaries are indicated in red.   We supplement the 1{\\ts}Jy ULIRG, BGS and SW01 samples with two optically selected samples: the NFGS field galaxy sample \\citep{jansen00} and the BGK00 galaxy close pair sample \\citep{barton00}.  The samples in Fig.~\\ref{fig:sampleall} are ordered from (a) to (e) by the level of interaction activity $-$ from no interaction (NFGS field galaxies), to non-IR selected galaxy pairs, to luminous IR-selected galaxies, and to the most luminous IR-selected galaxies (1{\\ts}Jy ULIRGs).  The fraction of starburst-AGN composite galaxies increases from the field galaxy sample (3.6\\%) $\\rightarrow$ galaxy pair sample (22.8\\% ) $\\rightarrow$ SW01 galaxies (29.6\\%) $\\rightarrow$ BGS galaxies (dominated by LIRGs) (37\\%) $\\rightarrow$ 1{\\ts}Jy ULIRGs (49.1\\%).  This plot indicates that the stronger the interaction between galaxies is, the more likely they are to contain starburst-AGN composite galaxies. \\begin{figure*}[!bthp] \\epsscale{1.14} \\plotone{f16.eps} \\caption{ Composite galaxies on the ``\\nii~ diagram\" for different types of galaxy samples. (a) Field galaxy sample: NFGS \\citep{jansen00}, (b) Galaxy pair sample: BGK00 \\citep{barton00}, (c) SW01 sample, (d) BGS sample, (e) 1{\\ts}Jy ULIRG sample. } \\label{fig:sampleall} \\end{figure*}  Our new results offer some hope for resolving previous disputes concerning the evolution of starburst and AGN activity in gas-rich major mergers, in particular for ULIRGs.  We have shown that a large fraction of ULIRGs are of composite starburst-AGN spectral type, with line ratios (or $D_{\\rm AGN}$) indicating an intermediate class between pure starbursts and pure AGN.   Spectral classification as a function of merger evolution does not change abruptly from ``pure\" starburst into ``pure\" AGN.  ULIRGs may spend a fairly large fraction of their merger history in a phase where starburst and AGN both make significant contributions to the EUV radiation.   In luminous IR mergers, starburst-AGN composites appear to ``bridge\" the spectral evolution from pure starburst to AGN-dominated activity as the merger progresses. For ULIRGs, we have shown that an initial apparent decrease in starburst activity from the wide binary to close binary stage is accompanied by a rise in the starburst-AGN composites.   Similarly, as the merger reaches its final stages, the fall in starburst-AGN composites is followed by a rise in Seyfert activity.  These effects can be seen in Fig.~\\ref{fig:moru}, where the starburst-AGN composite activity peaks in the diffuse merger stage.  We do not have sufficient data to determine whether a similar scenario occurs at the lower IR luminosities spanned by the BGS and SW01 samples ($L_{IR}<10^{12}L_{\\odot}$).     \\subsection{The Merger Scenario for LIRGs $\\rightarrow$ ULIRGs} There is now substantial agreement that ULIRGs are triggered by major mergers of gas-rich spirals \\citep[see the review by ][]{sanders96}.  In this merger scenario, the LIRG phase begins when tidal interactions between merging gas-rich disk galaxies initially trigger wide-spread starburst activity in one or both disks.  As the merger progresses, gas is funneled towards the merger nucleus and activates nuclear starbursts and AGN, further increasing the IR luminosity.  However, many questions still remain. The precise link between AGN fueling, the gas dynamics, and star formation either on a nuclear or a global scale is still poorly determined.  The relative strength of starburst and AGN activity, particularly in the ULIRGs is still debated, and the question of whether the majority of LIRG mergers become ULIRGs has not been clearly addressed.  Our new results shown in Figs. 9-16 provide new insight into how starburst and AGN activity evolve during interaction/merger of IR-luminous galaxies.  For lower luminosity IR-galaxies ($L_{\\rm IR} < 10^{12} L_{\\odot}$), Seyferts are rare ($D_{\\rm AGN} = 0.25-0.35$), and starbursts strongly contribute to the EUV radiation field either as ``pure\" starbursts or as composites.  The majority of lower luminosity LIRGs are wide or close pairs.  These pairs may either evolve into the single merger stage without becoming ULIRGs, or they may enter the ULIRG phase during or shortly after the close pair stage.  We consider both cases:  (1) Some lower luminosity IR galaxies ($L_{\\rm IR} < 10^{12} L_{\\odot}$) may reach the final single merger stage without becoming ULIRGs.  As these non-ULIRG galaxies evolve from the close pair to the single merger stage, the mean IR luminosity increases by $\\sim 2\\times$ (to $L_{\\rm IR} \\sim 10^11.5 L_{\\odot}$), with only a relatively small rise in the mean value of $D_{\\rm AGN}$($\\sim$0.4).  These results suggest that non-ULIRGs undergo a rise in starburst activity but only modest AGN growth as they evolve from the close pair to the single merger stage.  Some of these lower luminosity single mergers may be the product of minor mergers of objects with mass ratios larger than 5:1 \\citep{ishida04, ishida07}.   (2) The lower luminosity pairs that may become ULIRGs are likely to be major mergers with mass ratios closer to unity \\citep{ishida04, ishida07}.  In the ULIRG phase, the AGN contribution is elevated at all merger stages.  We observe clear changes in the optical classification of ULIRGs as a function of merger stage.  The diffuse merger stage (compared to the close binary stage) shows a dramatic increase in starburst-AGN composites (from $\\sim$45\\% to $\\sim$80\\%).  In the subsequent compact/old merger stages, the fraction of composite galaxies falls, and AGN dominate.  The mean value of $D_{\\rm AGN}$ rises (by $\\sim$0.22{\\ts}dex $-$ from 0.43 to 0.65), and the mean value of $L_{\\rm IR}$ rises ($\\sim 1.7\\times$), along with the emergence of a substantial Seyfert{\\ts}1 fraction ($\\sim$25\\%).   These results suggest that during the diffuse merger stage, the AGN becomes more powerful and increasingly visible via the optical emission-line ratios.   Once the merged nucleus begins forming a core (compact merger stage) starburst activity may be subsiding (and possible dust obscuration clears) as AGN activity becomes prominent.  We tentatively suggest that the dramatic increase in $D_{\\rm AGN}$ and the emergence of a substantial population of Seyfert{\\ts}1s (accompanied by a $\\sim 2\\times$ increase in $L_{\\rm IR}$) may signify a significant ``blowout\" phase, as ULIRGs transition to optical QSOs \\citep[e.g., ][]{hopkins05}.  A larger sample of ULIRGs is required to test this idea."
    },
    "0911/0911.4939_arXiv.txt": {
        "abstract": "Analytical and numerical results on equilibrium configurations of rotating fluid bodies within Einstein's theory of gravitation are reviewed. Particular emphasis is placed on continuous parametric transitions to black holes. In this connection the uniqueness of extremal Kerr black holes is discussed. ",
        "introduction": "The theory of figures of equilibrium of rotating fluid bodies under the influence of their own gravitational field is relevant for understanding the shape and the  structure of planets and stars. The problems to be solved are rather involved even within Newton's theory of gravitation, and analytic solutions are scarce. Basic references are the books by Lichtenstein\\cite{li} and Chandrasekhar\\cite{ch}. Within Einstein's theory of gravitation, which is relevant for describing compact objects like neutron stars, analytic solutions are only available in certain limiting cases. In general, a numerical treatment is necessary. For a detailed account of relativistic figures of equilibrium see Ref.~\\refcite{rfe} and the book by Friedman and Stergioulas\\cite{fs}. In the following, a brief review of some analytical and numerical results will be given. ",
        "conclusions": ""
    },
    "0911/0911.3034_arXiv.txt": {
        "abstract": "This paper briefly summarizes the status of the cosmic ray observations by EAS (Extended Air Shower) experiments with energy below $10^{16}$eV and the related studies of the hadronic interaction models. Based on the observed sharp knee structure and the irregularities of the cosmic ray spectrum around knee energy, plus the newly discovered electron and positron excess, the origin of the galactic cosmic rays and the single source model  interpretation are discussed, but  convincing evidence is not yet  available. High precision measurements of the mass composition of primary cosmic rays at knee energy will be very useful to disentangle the problem. To reach this goal, a better understanding of the hadronic interaction models is crucial. It is good to see that more dedicated accelerator and cosmic ray experiments will be conducted soon. As one EAS component, the muon distribution and muon charge ratio are important for testing the hadronic interaction models. In addition muons are an important background to neutrino experiments and all underground ultra-low background experiments. They are also a very useful tool for the meteorological studies. ",
        "introduction": "This paper is a writeup of the rapporteur talk at the 31st International Cosmic Ray Conference, July 15,2009 in Lodz/Poland.  The talk and this paper cover the works submitted to the HE session for cosmic ray energy below $10^{16}$ eV. The contributions include 26 papers in HE1.1 (observation and simulation at energies less than $10^{15}$ eV), 16 papers in HE1.2 (observation and simulation at energies of about $10^{15}$ - $10^{16}$ eV), 15 from HE1.5 (muons in EAS), 16 from HE1.6 (new experiments and instrumentation) and 20 in HE2.1 (particle interactions relevant for cosmic ray studies with energies less than $10^{16}$ eV).  Despite the large success of the observations of high energy $\\gamma $ rays during the last twenty years, direct evidence for the acceleration of VHE cosmic ray nuclei has not yet been convincingly obtained. Therefore, it has become consensus that from the point of view of gamma ray observations we must go to PeV energies ~\\cite{Aharonian} where the hadronic process should become visible. New activities following this approach are underway recently (e.g.,~\\cite{icrc0297},~\\cite{icrc0869}). Strictly speaking, gamma ray observations can only directly be related to the origin and acceleration mechanisms of the newly generated cosmic rays, observation on the cosmic ray particle themselves is the most fundamental and probably the only way to resolve the questions related to the acceleration of cosmic rays at ``early'' time. These cosmic rays compose the vast majority of the galactic cosmic rays as we can see them today. Many new observations, new understanding, new ideas and plans for the future were presented in this conference; it is impossible for me to cover all works and all details neither in the rapporteur talk nor in this paper. I do apology for the incompleteness of this summary paper. ",
        "conclusions": "More progress has been made in the study of the cosmic ray spectrum and more hints for a single source model are presented. The compelling evidence for the origin of the galactic cosmic rays is not yet found and this situation imposes a strong demand and great interest in precisely measuring the mass composition of cosmic rays around the knee energy which heavily relies on the progress in understanding the hadronic interaction models. With the ongoing and upcoming activities, we are quite confident that many of the problems will be resolved soon."
    },
    "0911/0911.1946_arXiv.txt": {
        "abstract": "In our ongoing multi-wavelength study of cluster AGN, we find $\\approx$75\\% of the spectroscopically identified cluster X-ray point sources (XPS) with L$_{0.3-8.0keV}>10^{42}$ ergs s$^{-1}$ and cluster radio galaxies with P$_{1.4 GHz}>3\\times10^{23}$ W Hz$^{-1}$ in 11 moderate redshift clusters (0.2$<$z$<$0.4) are located within 500 kpc from the cluster center. In addition, these sources are much more centrally concentrated than luminous cluster red sequence (CRS) galaxies.  With the exception of one luminous X-ray source, we find that cluster XPSs are hosted by passive red sequence galaxies, have X-ray colors consistent with an AGN power-law spectrum, and have little intrinsic obscuring columns in the X-ray (in agreement with previous studies).  Our cluster radio sources have properties similar to FR1s, but are not detected in X-ray probably because their predicted X-ray emission falls below our sensitivity limits. Based on the observational properties of our XPS population, we suggest that the cluster XPSs are low-luminosity BL Lac objects, and thus are beamed low-power FR 1s.  Extrapolating the X-ray luminosity function of BL Lacs and the Radio luminosity function of FR 1s down to fainter radio and X-ray limits, we estimate that a large fraction, perhaps all CRSs with L$>$L$^*$ possess relativistic jets which can inject energy into the ICM, potentially solving the uniform heating problem in the central region of clusters. ",
        "introduction": "Observations of galaxy clusters have provided important clues about the formation and evolution of stars, galaxies and AGN in these systems. And based upon the observed correlation between supermassive Black Hole mass and galaxy bulge mass, we expect that there should be numerous AGN in the bright ellipticals in rich clusters as well.  Radio-loud AGN have been directly implicated in cluster ICM heating because evacuated ``bubbles'' in the diffuse X-ray emitting gas have been observed that are spatially coincident with non-thermal radio emission. But what about other non-BCG radio-loud galaxies?  Is there a connection between cluster radio galaxies and the X-ray AGN population?  What is the spatial distribution of this combined population?  To answer these questions we present a multi-wavelength study of cluster AGN that differs from many previous studies (e.g. \\citep{1995AJ....109..853L,2007ApJ...664..761M}) that use X-ray or radio observations, but not both, to detect AGN in only one waveband. Most cluster AGN studies have either employed clusters in a flux-limited survey or simply selected clusters based upon availability.  However, since most flux-limited surveys tend to detect the most luminous clusters at higher-$z$, this selection naturally identifies more evolved structures at earlier times.  This mismatch can result in comparisons which can obscure any true evolution, since more massive clusters are placed at the beginning of the evolutionary sequence and less massive objects at the end.  We describe a unique cluster selection method below which avoids the difficulty just described and present a summary of our first results. ",
        "conclusions": ""
    },
    "0911/0911.3172_arXiv.txt": {
        "abstract": "{Is it significant that the intrinsic outputs of several BL Lacs are observed to level off at values of about \\({10}^{46}\\mbox{ erg}\\mbox{ s}^{-1}\\)? In searching for an answer, we compare \\(\\gamma\\)-ray observations by the \\emph{AGILE} satellite of the BL Lac S5 0716+714 with those of Mrk 421 and Mrk 501; the former are particularly marked by intense flares up to fluxes of \\(2\\times{10}^{-6}\\mbox{ photons}\\mbox{ cm}^{-2}\\mbox{ s}^{-1}\\) in the \\(0.1-10\\mbox{ GeV}\\) energy range. These ``dry\" BL Lacs show evidence of neither thermal disk emissions nor emission lines signaling any accreting or surrounding gas; the spectral distributions of their pure non-thermal radiations are effectively represented by the synchrotron self-Compton process. With source parameters correspondingly derived and tuned with simultaneous multiwavelength observations, we find for S5 0716+714 a total jet power of about \\(3\\times{10}^{45}\\mbox{ erg}\\mbox{ s}^{-1}\\), which makes it one of the brightest dry BL Lacs so far detected in \\(\\gamma\\) rays. We evaluate the mass of the associated Kerr hole to be around \\(5\\times{10}^8 M_{\\tiny{\\astrosun}}\\), implying that the source is significantly gauged in terms of the maximal power around \\(4\\times{10}^{45}\\mbox{ erg}\\mbox{ s}^{-1}\\) extractable via the Blandford-Znajek electrodynamical mechanism; other dry BL Lacs observed in \\(\\gamma\\) rays remain well below that threshold. These findings and those forthcoming from \\emph{Fermi}-LAT will provide a powerful test of electrodynamics in the surroundings of the hole, that are dominated by GR effects.} ",
        "introduction": "\\label{sezione1}  Blazars rank among the brightest active galactic nuclei on the basis of their inferred isotropic luminosities that may attain some \\(L_{iso}\\sim {10}^{48}\\mbox{ erg}\\mbox{ s}^{-1}\\). Actually, these sources radiate from a narrow relativistic jet closely aligned with the observer's line of sight. The jet emits highly beamed non-thermal radiations, with observed fluxes enhanced by aberration and Doppler effects of Special Relativity (Begelman, Blandford \\& Rees \\cite{bbr}; K\\\"{o}nigl \\cite{konigl86}; Urry \\& Padovani \\cite{urry}). So the luminosities \\(L_{iso}\\) greatly exceed the intrinsic outputs of the jets that easily level off at \\({10}^{-2} - {10}^{-3} L_{iso}\\).  The BL Lac objects (henceforth BL Lacs) in particular are blazars that show no or just weak and intermittent emission lines. Their spectra are represented well as a continuous spectral energy distribution (SED) \\(S_\\nu=\\nu F_\\nu\\) featuring two peaks: one at a lower frequency due to synchrotron emission by highly relativistic electrons; and a higher frequency counterpart due to inverse Compton upscattering by the same electron population of seed photons provided by the synchrotron emission itself (synchrotron self-Compton, SSC; see Jones, O'Dell \\& Stein \\cite{jones}, Marscher \\& Gear \\cite{marscher}; Maraschi, Ghisellini \\& Celotti \\cite{maraschi}), with possible additions from sources external to the beam (external Compton, EC; see Dermer \\& Schlickeiser \\cite{dermer93}; Sikora, Begelman \\& Rees \\cite{sikora}).  The BL Lacs also exhibit strong variability on timescales of days to minutes with substantial flux variations (flares) particularly at high energies as realized early on (see Setti \\& Woltjer \\cite{setti}, and references therein).  Here we focus on ``dry\" BL Lacs, that is, sources with no evidence of surrounding gas, such as emission lines or a big blue bump (see Peterson \\cite{peterson}, Kembhavi \\& Narlikar \\cite{kembhavi}) related to current accretion. They provide an appropriate testing ground for comparing their intrinsic outputs with maximal powers extractable from rotating supermassive black holes and from the dragged accretion disks by means of large-scale electromagnetic fields, via the intriguing, variously debated Blandford-Znajek electrodynamics (BZ, Blandford \\& Znajek \\cite{bz}; see also Ghosh \\& Abramowicz \\cite{ghosh}; Krolik \\cite{krolik}; Livio, Ogilvie \\& Pringle \\cite{livio}; Cavaliere \\& D'Elia \\cite{cavaliere}; McKinney \\cite{mckinney05}; Nemmen et al. \\cite{nemmen}; Tchekhovskoy, McKinney \\& Narayan \\cite{thcek}). The bare hole contribution can yield up to \\(L_K \\sim 2\\times{10}^{45}\\left({{M_{\\medbullet}}/{{10}^9\\,M_{\\tiny{\\astrosun}}}}\\right)\\mbox{ erg}\\mbox{ s}^{-1}\\), given the hole mass \\(M_{\\tiny{\\medbullet}}\\) in units of \\({10}^9M_{\\tiny{\\astrosun}}\\) and a magnetic field  \\(B\\sim{10}^4\\mbox{ G}\\) threading its horizon.  In the following, we adopt the standard, flat cosmology with \\(H_0 = 72 \\mbox{ km}\\mbox{ s}^{-1}\\mbox{ Mpc}^{-1}\\) and \\(\\Omega_\\Lambda = 0.74\\) (Dunkley et al. \\cite{dunkley}). ",
        "conclusions": "For dry BL Lacs with accretion rates \\(\\dot{m}<{10}^{-2}\\), the SSC radiation process provides a robust evaluation of the jet luminosities. Whence we conclude that \\(L_{BZ}\\) provides a significant \\emph{benchmark} for the output of the BL Lacs discussed here, and indeed  an upper limit to both their quiescent states and flares. In fact, during a recent flare S5 0716+714 was observed to be constrained by \\(L_{BZ}\\lesssim{10}^{46}\\mbox{ erg}\\mbox{ s}^{-1}\\), and a similar behavior was observed in 2008 for Mrk 421. Non-thermal, beamed powers in the range \\(L_K - L_{BZ}\\) also provide evidence of an accretion disk that is active mainly in launching and channeling the jets by means of large-scale fields.  Referring to Fig. \\ref{bs} and its caption, we note that during flares the sources move in the \\(L_{T}\\) - \\(\\gamma_{peak}\\) plane away from the envelope that is outlined by bright BL Lacs with increasing rates \\(\\dot{m}\\); the envelope ends up in the locus of the yet brighter Flat Spectrum Radio Quasars (FSRQs) with \\(\\dot{m}\\sim 1\\). The flares then move into a region of faster radiative cooling (Celotti \\& Ghisellini \\cite{celotti}, and references therein). This implies short-lived flares on timescales \\(\\la 1\\mbox{ day}\\), or requires shorter acceleration times \\(t_a\\sim \\gamma /E\\) with higher \\(E\\), as an alternative to structured sources such as decelerating (Georganopoulos \\& Kazanas \\cite{gk}) or spine-sheath jets with inner scale \\(R_1<R\\) (Tavecchio \\& Ghiselllini \\cite{tavecchio}). Faster acceleration and deviations from the envelope are consistent with flares caused by electron \\emph{boost} rather than episodes of increased accretion onto the disk.  \\begin{figure*} \\includegraphics[scale=0.40]{bs.eps} \\caption{Bright BL Lacs in their context, adapted from Celotti \\& Ghisellini (\\cite{celotti}) with historical data in terms of \\(L_T\\) and the energy \\(\\gamma_{peak}\\) related to the synchrotron peak (\\(\\gamma_{peak}=\\gamma_{p}\\times {10}^{1/2r}\\)). Blue circles indicate dry, while violet circles indicate wet BL Lacs. The lower-left region of the diagram corresponds to the source condition \\(t_c > R/c\\), and the upper-right to \\(t_c\\gtrsim t_a\\). Bright FSRQs lie at lower \\(\\gamma_{peak}\\) and higher powers, with increasing signs of current \\(\\dot{m}\\rightarrow 1\\); selection effects depopulate weaker sources in the lower region (see Padovani \\cite{padovani07}). The dashed oval highlights sources in the transition region from dry to wet BL Lacs, interesting to compare with \\(L_{BZ}\\) from Eq. \\ref{BZ} particularly when at \\(z\\gtrsim 0.3\\).} \\label{bs} \\end{figure*}  In this context, we recall (see Blandford \\cite{blandford}; Padovani et al. \\cite{padovani}) that sources lying along the envelope in Fig. \\ref{bs} at higher \\(L\\) and lower \\(\\gamma_{peak}\\) often exhibit stronger evidence of current accretion up to \\(\\dot{m}\\sim 1\\), such as thermal emissions and surrounding gas (big blue bump and broad emission lines), with a larger contribution from EC. In fact, the progression from dry BL Lacs to FSRQs is likely to involve an enhanced and extended disk contribution as described by Blandford \\& Payne (\\cite{bp}), starting with ``wet\" BL Lacs with \\(\\dot{m}\\sim{10}^{-1}\\); these feature larger EC contributions (Dermer et al. \\cite{dermer}) and looming evidence of gas, including some thermal disk emission and weak or intermittent lines (Celotti, Ghisellini \\& Fabian \\cite{celotti07}). The last step in this progression is constituted by the powerful FSRQs with extant broad lines, a big blue bump from disks accreting at full rates \\(\\dot m \\sim 1\\), and a dominant or towering EC (Maraschi \\& Tavecchio \\cite{maraschi01}).  We add that the outputs of even misaligned BL Lacs may be \\emph{calorimetrically} gauged from their feedback actions on the intra-cluster plasma surrounding their host galaxy when located in a cluster or a group of galaxies, as discussed by McNamara et al. (\\cite{mcnamara}). These authors evaluate average powers around \\({10}^{46}\\mbox{ erg}\\mbox{ s}^{-1}\\) injected into the cluster MS0735.6+7421, and possibly also in the cluster A2029 and the group AWM 4.  The whole of the above evidence provides observational support to the \\emph{relevance} of the electrodynamical BZ mechanism, and invites extended sampling of other interesting sources (see Fig. \\ref{bs}).  If in dry BL Lacs with \\(M_{\\tiny{\\medbullet}}< {10}^9 M_{\\tiny{\\astrosun}}\\) the \\(L_{BZ}\\) limit were found to be substantially exceeded by outputs \\(L_T> {10}^{46}\\mbox{ erg}\\mbox{ s}^{-1}\\), this would require \\(B > {10}^4\\mbox{ G}\\) at the Kerr horizon. These fields imply large dynamical stresses bounded only by \\(B^2/4 \\pi  \\leq\\rho c^2\\), associated with particle orbits plunging from the disk toward the hole horizon (Meier \\cite{meier2}) into a region fully controlled by strong gravity effects.  Thus, all such sources will provide powerful tests for the coupling of electrodynamics with General Relativity in full swing, and constitute an exciting arena for \\emph{AGILE} and \\emph{Fermi}-LAT data."
    },
    "0911/0911.1177_arXiv.txt": {
        "abstract": "Young binaries within dense molecular clouds are subject to dynamical friction from ambient gas. Consequently, their orbits decay, with both the separation and period decreasing in time. A simple analytic expression is derived for this braking torque. The derivation utilizes the fact that each binary acts as a quadrupolar source of acoustic waves. The acoustic disturbance has the morphology of a two-armed spiral and carries off angular momentum. From the expression for the braking torque, the binary orbital evolution is also determined analytically. This type of merger may help explain the origin of high-mass stars. If infrared dark clouds, with peak densities up to \\hbox{$10^7\\,\\,{\\rm cm}^{-3}$}, contain low-mass binaries, those with separations less than 100~AU merge within about $10^5$~yr. During the last few thousand years of the process, the rate of mechanical energy deposition in the gas exceeds the stars' radiative luminosity. Successive mergers may lead to the massive star formation believed to occur in these clouds. ",
        "introduction": "The youngest binaries are still embedded in molecular cloud gas. During this phase, the orbiting components experience dynamical friction with the surrounding medium. The braking torque causes the stars to spiral inward, so that their separation and period diminish with time. Such orbital decay may be significant in the very densest molecular clumps. In particular, infrared dark clouds, currently believed to be the birthsites of massive stars, have peak number densities as high as \\hbox{$10^7~{\\rm cm}^{-3}$} \\citep{r06,b07}. This impressive figure matches that in hot molecular cores, already known to contain luminous, high-mass objects \\citep{k00}. Binary mergers could be occurring in these environments, and could play a role in forming the massive stars.  There is still no direct evidence that either infrared dark clouds or hot molecular cores contain low-mass stars. However, observations show that, with few exceptions, massive stars originate in populous clusters \\citep{dw05}. Indeed, one account of high-mass star formation is that it proceeds through the coalescence of lower-mass cluster members \\citep{b98,s00}. At even the highest observed cluster densities, collisions between stars, either bare or with extended disks, occur at too low a frequency. However, the presence of dense, background gas should enhance the capture rate \\citep{bz05}. Specifically, the braking torque from dynamical friction might not only shrink initially wide binaries, but cause the component stars to merge. An accelerating succession of such mergers could then build up massive stars.  To assess this possibility, the present study calculates the torque on a binary embedded within an extensive gas cloud. Although the torque is created by dynamical friction, the usual analysis of that effect is difficult to apply in this context. Following \\citet{d64}, theorists have determined the frictional drag on a single mass traveling through ambient gas in a straight line \\citep[see also][]{rs71,rs80,o99}. Just as in the stellar dynamical case \\citep{c43}, the moving mass draws external matter into a trailing wake. It is the gravitational tug from this wake that provides the effective drag. The same essential mechanism is at work in the case of a binary. However, the trajectories of both perturbing masses are now closed orbits. Further complicating the analysis is the fact that each star draws in matter from the wake of its companion. Naive application of the standard dynamical friction formulae would grossly misestimate the torque.  The new approach introduced here concentrates on the fact that the binary as a whole must shed angular momentum to the external medium. The actual mechanism is that the orbiting stars create an oscillating gravitational potential that torques nearby gas. This fluctuating torque generates outgoing acoustic waves, which transport angular momentum. By calculating the total angular momentum efflux from the binary, the braking torque may be obtained without considering the complex star-gas interaction close to the stars themselves.  Section~2 below formulates the problem mathematically and derives the governing wave equation. The quadrupolar driving potential created by the rotating stars is established in Section~3, while Section~4 discusses the physical character of the generated waves. Section~5 presents the central result of this investigation. Here the angular momentum transported by the spiral wave is obtained, and thereby the torque (see eq.~(39)). It is also shown that the wave carries off all the mechanical energy released by the shrinking binary. The resulting temporal evolution of the system is considered in Section~6, along with the efficacy of braking in infrared dark clouds. Finally, Section~7 compares the results obtained here with the traditional analysis and discusses future extensions of this study. ",
        "conclusions": "We may compare, at least in a qualitative manner, our derived torque with that indicated by the traditional theory of dynamical friction. According to equation~(12) of \\citet{o99}, the retarding force on a mass $M$ moving at speed $V$ through a cloud of density $\\rho_0$ is \\begin{displaymath} F_{\\rm DF} \\,=\\, -{{4\\,\\pi\\,(G\\,M)^2\\,\\rho_0}\\over V^2}\\,\\,{\\cal I} \\,\\,. \\end{displaymath} The factor $\\cal I$, essentially a Coulomb logarithm, is a nondimensional function of the Mach number $V/c_s$ and the time since the mass first entered the cloud in question. Since $\\cal I$ will generally be of order unity, we may ignore it, along with other such factors, in the dimensional argument that follows.  To apply this formula to the binary problem, we interpret $V$ as the components' relative velocity $V_{\\rm rel}$, which is $\\omega\\,a_{\\rm tot}$. Then the torque is of order \\hbox{$a_{\\rm tot}\\,F_{\\rm DF}$}, so that \\begin{displaymath} \\Gamma_{\\rm DF} \\,\\sim\\, -{{\\omega^3\\,\\rho_0\\,G^2\\,I^2}\\over V_{\\rm rel}^5} \\,\\,. \\end{displaymath} Here, the moment of inertia has been approximated as \\hbox{$I\\,\\approx\\,M\\,a_{\\rm tot}^2$}. Comparison to equation~(39) shows that $\\Gamma_{\\rm DF}$ is the true $\\Gamma$ multiplied by a factor \\hbox{$(c_s/V_{\\rm rel})^5$}. This factor can be much smaller than unity for the hard binaries of interest. On the other hand, the nondimensional $\\cal I$ formally diverges for a velocity of $c_s$. The traditional theory is thus unreliable in this context.  However, it is not difficult to envision circumstances in which the present theory requires modification. From equation~(35a), the Mach number associated with the radial velocity amplitude is \\begin{eqnarray} {u_r\\over c_s} \\,&=&\\,{{G\\,I\\,\\omega^2}\\over{c_s^4\\,r}} \\\\ &\\sim&\\,\\left(V_{\\rm rel}\\over c_s\\right)^2 {r_{\\rm in}\\over r} \\,\\,. \\end{eqnarray} Since \\hbox{$r_{\\rm in}\\,\\ll\\,r$}, the induced velocity is normally subsonic. For binaries that are initially very hard, or late during the inspiral of any system, $u_r$ throughout the far field becomes supersonic. The disturbance then changes character from an acoustic wave to a tightly wound spiral shock. In the same regime, the density perturbation $\\rho_1$ is comparable to or even exceeds $\\rho_0$, so that a fully nonlinear treatment is necessary.  Additional modification of the theory would be required by the inclusion of a finite eccentricity in the binary orbits. While the system would still be periodic, the equivalent quadrupolar source would exhibit variation over a continuous distribution of frequencies. The same frequency distribution would then appear in the transmitted waves. It would be interesting to recalculate both the torque and energy loss under this more general condition and thereby follow the evolution of the eccentricity during orbital decay.  Returning to the theory's main astrophysical application, our result for the coalescence time is a reasonable one that adds credence to the underlying picture. The nearest and best-studied region of massive star formation is the Orion Nebula Cluster, whose general population formed 1-2~Myr ago \\citep{h97}. We do not know the age of the high-mass members, the Trapezium, with any precision. \\citet{ps01} have argued that they are relatively young, of order $10^5~{\\rm yr}$, based on the location of BM~Ori and its binary companion in the HR diagram. Assuming that the Trapezium stars, along with their companions, are coeval, their inferred age provides an upper bound to the formation time scale. We are thus encouraged by this matching of times. We stress, however, that our expression for $t_c$ in equation~(56) varies inversely with the imprecisely known ambient density $n$.  If massive stars indeed coalesce over such a period, perhaps they do so through an accelerating sequence of binary mergers. The binaries themselves might be created and disrupted rapidly out of the dense gas, leading to a statistically stable period distribution, as originally envisioned by \\citet{lb69}. Within this picture, one could in principle determine the mass distribution of the growing population of coalesced objects, thereby advancing the theory another significant step."
    },
    "0911/0911.3166_arXiv.txt": {
        "abstract": "{} {We present the main results of an imaging survey of possible young massive clusters (YMC) in M31 performed with the Wide Field and Planetary Camera 2 (WFPC2) on the Hubble Space Telescope (HST), with the aim of estimating their age and their mass. We obtained shallow (to B$\\sim 25$) photometry of individual stars in 19 clusters (of the 20 targets of the survey).  We  present the images and color magnitude diagrams (CMDs) of all of our targets.} {Point spread function fitting photometry of individual stars was obtained for all the WFPC2 images of the target clusters, and the completeness of the final samples was estimated using extensive sets of artificial stars experiments. The reddening, age, and metallicity of the clusters were estimated by comparing the observed CMDs and luminosity functions (LFs) with theoretical models. Stellar masses were estimated by comparison with theoretical models in the log(Age) $vs.$ absolute integrated magnitude plane, using ages estimated from our CMDs and integrated J, H, K magnitudes from 2MASS-6X.} {Nineteen of the twenty surveyed candidates were confirmed to be real star clusters, while one turned out to be a bright star. Three of the clusters were found not to be good YMC candidates from newly available integrated spectroscopy and were in fact found to be old from their CMD. Of the remaining sixteen clusters, fourteen have ages between 25 Myr and 280 Myr, two have older ages than 500 Myr (lower limits). By including ten other YMC with HST photometry from the literature, we assembled a sample of 25 clusters younger than 1 Gyr, with mass ranging from $0.6\\times 10^4 M_{\\sun}$ to $6\\times 10^4 M_{\\sun}$, with an average of $\\sim 3\\times 10^4 M_{\\sun}$. Our estimates of ages and masses well agree with recent independent studies based on integrated spectra.} {The clusters considered  here are confirmed to have masses significantly higher than Galactic open clusters (OC) in the same age range. Our analysis indicates that YMCs are relatively common in all the largest star-forming galaxies of the Local Group, while the lack of known YMC older than 20 Myr in the Milky Way may stem from selection effects.} ",
        "introduction": "Much of the star formation in the Milky Way is thought to have occurred within star clusters  (Lada et al. \\cite{lada91}, Carpenter et al. \\cite{carp}); therefore, understanding the formation and evolution of star clusters is an important piece of the galaxy formation puzzle. Our understanding of the star cluster systems of spiral galaxies largely comes from studies of the Milky Way. Star clusters in our Galaxy have traditionally been separated into two varieties, open and globular clusters  (OCs and GCs hereafter).  OCs are conventionally regarded as young ($<10^{10}$ yr), low-mass  ($<10^4 M_{\\sun}$), and metal-rich systems that reside in the Galactic disk. In contrast, GCs are characterized as old, massive systems. In the Milky Way, GCs can be broadly separated into two components: a metal-rich disk/bulge  subpopulation, and a spatially extended, metal-poor halo subsystem  (Kinman \\cite{tom}, Zinn \\cite{zinn}, see also Brodie \\& Strader \\cite{brodie}, Harris \\cite{h01}, for general reviews  of GCs).  However, the distinction between OCs and GCs has become increasingly blurred. For example, some OCs are luminous and old enought to be confused with GCs  (e.g., Phelps \\& Schick \\cite{phelps}). Similarly, some GCs are very low-luminosity systems (e.g., Koposov et al. \\cite{kopos}), and at least one has an age that is consistent  with the OC age distribution (Palomar 1, Sarajedini et al. \\cite{sara}). Moreover, a third category of star cluster, ``young massive clusters'' (YMCs) are observed to exist in both merging  (e.g., Whitmore \\& Schweizer \\cite{wischw}) and quiescent galaxies  (Larsen \\& Richtler \\cite{lari}). Indeed, YMCs have been known to exist in the Large Magellanic Cloud (LMC) for over half a century (Hodge \\cite{ho61}).  These objects are significantly more luminous than  OCs (M$_V\\la-8$ up to M$_V\\sim -15$), making them promising candidate young GCs. Once thought  to be absent in the Milky Way, recent observations suggest that their census may be quite incomplete, as some prominent cases have been found  recently in the Galaxy as well (Clark et al. \\cite{clark}, Figer \\cite{figer}, Messineo et al. \\cite{maria}).  Thus, a picture has emerged that, rather than being distinct groups, OCs, YMCs and GCs may represent regions within a continuum of cluster properties dependent upon local galaxy conditions (Larsen \\cite{lars}). The lifetime of a star cluster is  dependent upon its mass and environment. Most low-mass star clusters in disks are rapidly disrupted via interactions with giant molecular clouds (Lamers \\& Gieles \\cite{lamers}, Gieles et al. \\cite{gieles}). These disrupted  star clusters are thought to be the origin of much of the present field star populations (Lada \\& Lada \\cite{lala}). Surviving disk clusters may then be regarded as OCs or YMCs, depending upon their  mass. Star clusters in the halo may survive longer since they are subjected to  the more gradual dynamical processes of two-body relaxation and evaporation.  The clusters which survive for a Hubble time -- more likely to occur away from the  disk -- are termed GCs (see also Krienke \\& Hodge \\cite{krie1}). To date, no known {\\it thin} disk GCs have been identified in the Milky Way.  After the Milky Way, M31 is the prime target for expanding our knowledge of cluster systems in spirals. However, our present state of knowledge about the M31 cluster system is far from complete.  Similar to the Milky Way, M31 appears to have at least two GC  subpopulations, a metal-rich, spatially concentrated subpopulation of GCs  and a more metal-poor, spatially extended GC subpopulation (Huchra et al. \\cite{huchra}, Barmby et al. \\cite{barm00}). Also, again similar to the Milky Way GCs, the metal-rich GCs in M31 rotate and  show \"bulge-like\" kinematics (Perrett et al. \\cite{perrett}). However, unlike the case in  the Milky Way, the metal-poor GCs also show significant rotation (Huchra et al. \\cite{huchra}, Perrett et al. \\cite{perrett}, Lee et al.~\\cite{lee}). Using the Perrett et al. (\\cite{perrett}) data, Morrison et al. (\\cite{morri}) identified what  appeared to be a {\\it thin} disk population of GCs, constituting some  27\\% of the Perrett et al. (\\cite{perrett}) sample. Subsequently, it has been shown that at least a subset of these objects  are in fact young ($\\leq$ 1 Gyr), metal-rich star clusters rather than old ``classical''  GCs (Beasley et al. \\cite{beas}, Burstein et al. \\cite{burst}, Fusi Pecci et al.~\\cite{ffp}, Puzia et al. \\cite{puzia}, Caldwell et al. \\cite{C09}).  Fusi Pecci et al. (\\cite{ffp}, hereafter F05) presented a comprehensive study of bright young disk clusters in M31, selected from the Revised Bologna Catalog\\footnote{\\tt www.bo.astro.it/M31} (RBC, Galleti et al.~\\cite{rbc}) by color [$(B-V)_0\\le 0.45$] or by the strength of the $H\\beta$ line in their spectra ($H\\beta\\ge 3.5\\AA$). While these clusters have been noted since Vetesnik (\\cite{vetes}) and have been studied by various authors, a systematic study was lacking. F05 found that these clusters, that they termed -- to add to the growing menagerie of star cluster species -- ``blue luminous compact clusters'' (BLCCs), are fairly numerous in M31 (15\\% of the whole GC sample), they have positions and kinematics typical of {\\em thin disk} objects, and their colors and spectra strongly suggest that they have ages (significantly) less than 2 Gyr.  Since they are quite bright ($-6.5\\la M_V\\la -10.0$) and -- at least in some cases -- morphologically similar to old GCs (see Williams \\& Hodge \\cite{will}, hereafter WH01), BLCCs could be regarded as YMCs, that is to say, candidate  young GCs (see De Grijs \\cite{degris2}, for a recent review). In particular, F05 concluded that if most of the BLCCs have an age $\\ga 50-100$ Myr they are likely brighter than Galactic open clusters (OC) of similar ages, thus they should belong to a class of objects that is not present, in large numbers, in our own Galaxy. Unfortunately, the accuracy in the age estimates obtained from the integrated properties of the clusters is not sufficient to determine their actual nature on an individual basis, i.e., to compare their total luminosity with the luminosity distribution of OCs of similar age (see Bellazzini et al. \\cite{cefa}, hereafter B08, and references therein).  In addition to the question of the masses and ages of these BLCCs, it has become clear that the BLCC photometric and spectroscopic samples in M31 may suffer from significant contamination.  Cohen, Matthews \\& Cameron (\\cite{coh}, hereafter C06) presented NIRC2@KeckII Laser Guide Star Adaptive Optics (LGSAO) images of six candidate BLCCs. Their $K\\arcmin$ very-high spatial resolution images revealed that in the fields of four candidates there was no apparent cluster. This led C06 to the conclusion that some/many of the claimed BLCC may in fact be just {\\em asterisms}, i.e. chance groupings of stars in the dense disk of M31. The use of the near infrared $K\\arcmin$ band (required by the LGSAO technique) may be  largely insensitive  to very young clusters that are dominated by relatively few hot stars, which emit most of their light in the blue region of the spectrum. Hence, the imaging by C06 may be inappropriate to detect such young clusters (see, for example, the detailed discussion by Caldwell et al. \\cite{C09}). In any case, the study by C06 suggests that the true number of massive young clusters of M31 may have been overestimated.  \\begin{figure*} \\centering \\includegraphics[width=18cm]{cpos_lr.ps} \\caption{Location of the 20 targets of our survey (empty circles) projected against the body of M31. The $\\times$ symbols indicate the position of the additional ten Young Clusters we included in Sect.~\\ref{massageir}.} \\label{cpos} \\end{figure*}   Therefore, in order to ascertain the real nature of these BLCCs we have performed  a survey with the Hubble Space Telescope (HST) to image 20 BLCCs in the disk of M31 (program GO-10818, P.I.: J. Cohen).  The key aims of the survey are:  \\begin{enumerate}  \\item to check if the imaged targets are real clusters or asterisms, and to determine the fraction of contamination of BLCCs by asterisms,  \\item to obtain an estimate of the age of each  cluster in order to verify whether it is brighter than Galactic OCs of similar age. Ultimately the survey aims to provide firm conclusions on the existence of a significant population of BLCCs (YMCs) in M31, in addition to OCs (see Krienke \\& Hodge \\cite{krie1,krie2}, and references therein) and GCs.  \\end{enumerate}  \\begin{figure*} \\centering \\includegraphics[width=18cm]{cima_lr.ps} \\caption{F450W images of the 20 primary targets. Each image covers the central 10$^{\\arcsec}\\times$~10$^{\\arcsec}$ on the PC field (10$^{\\arcsec}$ = 38 pc at the assumed M31 distance modulus of 24.47). North is up and East to the left.} \\label{cima} \\end{figure*}     \\begin{table*} \\caption{Positional, photometric and spectroscopic parameters for the surveyed clusters.}  \\label{table:1} \\centering \\begin{tabular}{l c c c c c c c c c c c c}  \\hline\\hline Name & X$^a$ & Y$^a$ & R &  B &  V & (B-V)$_{0}$$^{F05}$ & (B-V)$_{0}$$^{(t.w.)}$ & H$_{\\beta}$$^{F05}$ & H$_{\\beta}$$^{G09}$ & ff$^b$ \\\\ &(arcmin)&(arcmin)&(arcmin)&&&&&(\\AA)&(\\AA)&\\\\  \\hline \\\\ B015D-D041 &  -19.27 &   9.22 &  21.36  & 19.11$\\pm$0.02 & 18.36$\\pm$0.03 & $\\dots$ &   0.15 & 7.32\t   & $\\dots$  & 1 \\\\ B040-G102  &  -35.40 & -11.92 &  37.35  & 17.54$\\pm$0.03 & 17.20$\\pm$0.04 & 0.18    &   0.11 & 7.41\t     & 7.58$\\pm$ 0.30\t & 1 \\\\ B043-G106  &  -33.62 & -11.37 &  35.49  & 17.04$\\pm$0.03 & 16.77$\\pm$0.04 & 0.17    &   0.04 & 5.53\t     & 5.70$\\pm$ 0.30\t & 1 \\\\ B066-G128  &  -29.55 & -13.17 &  32.35  & 17.56$\\pm$0.03 & 17.35$\\pm$0.04 & 0.25    &  -0.02 & 4.67\t     & 4.84$\\pm$ 0.30\t & 1 \\\\ B081-G142  &  -25.26 & -12.36 &  28.12  & 17.36$\\pm$0.02 & 16.86$\\pm$0.03 & 0.43    &    0.20 & 7.98\t     & 8.15$\\pm$ 0.30\t & 1 \\\\ B257D-D073 &   45.98 &   4.02 &  46.16  & 18.41$\\pm$0.02 & 18.00$\\pm$0.04 & $\\dots$ &   0.01 & 5.49\t   & 5.66$\\pm$ 0.30   & 1 \\\\ B318-G042  &  -52.14 &  -1.32 &  52.16  & 17.02$\\pm$0.03 & 16.82$\\pm$0.03 & 0.06    &   0.03 & $\\dots$\t& 5.49$\\pm$ 0.12    & 1 \\\\ B321-G046  &  -55.50 &  -7.41 &  55.99  & 17.82$\\pm$0.02 & 17.51$\\pm$0.03 & 0.11    &   0.06 & 6.29\t     & 6.85$\\pm$ 0.32\t & 1 \\\\ B327-G053  &  -47.67 &  -3.45 &  47.79  & 16.75$\\pm$0.03 & 16.58$\\pm$0.03 & 0.21    &  -0.03 & 4.09\t     & 3.78$\\pm$ 0.14\t & 1 \\\\ B376-G309  &   42.16 & -10.67 &  43.49  & 18.35$\\pm$0.02 & 17.97$\\pm$0.04 & 0.34    &   0.08 & $\\dots$\t& 6.40$\\pm$ 0.06    & 1 \\\\ B448-D035  &  -43.16 &  -2.97 &  43.26  & 18.01$\\pm$0.03 & 17.46$\\pm$0.04 & 0.50    &    0.20& 6.70\t     & 6.87$\\pm$ 0.30\t & 1 \\\\ B475-V128  &   45.00 &   4.06 &  45.18  & 17.55$\\pm$0.03 & 17.09$\\pm$0.04 & 0.20    &   0.11 & 5.96\t     & 6.13$\\pm$ 0.30\t & 1 \\\\ V031       &  -19.03 &   7.17 &  20.34  & 18.16$\\pm$0.03 & 17.62$\\pm$0.04 & 0.57    &   0.19 & 5.84\t     & 6.01$\\pm$ 0.30\t & 1 \\\\ B083-G146  &   19.83 &  22.08 &  29.68  & 17.85$^{d}$    & 17.09$^{d}$    & 0.65    &   0.56 & 3.75     & 1.75$\\pm$ 0.42    & 1 \\\\ B222-G277  &   10.22 & -16.16 &  19.12  & 18.00$\\pm$0.02 & 17.24$\\pm$0.03 & 0.57\t  &   0.56 & 8.47\t & 4.46$\\pm$ 0.31    & 1 \\\\ B347-G154  &   27.74 &  26.74 &  38.53  & 17.23$^{d}$    & 16.50$^{d}$    & 0.62    &   0.67 & $\\dots$  & 2.87$\\pm$ 0.17    & 2 \\\\ B374-G306  &   41.13 & -10.55 &  42.46  & 18.69$\\pm$0.03 & 18.23$\\pm$0.04 & 0.33\t  &   0.16 & 4.07\t   & 4.24$\\pm$ 0.30    & 1 \\\\ NB16       &    1.96 &   4.19 &   4.63  & 18.83$\\pm$0.04 & 17.59$\\pm$0.10 & 0.55\t  &   0.99 & $\\dots$\t& 3.34$\\pm$ 0.08    & 2 \\\\ VDB0       &  -47.16 &  -4.33 &  47.36  & 14.94$\\pm$0.09$^{c}$ & 14.67$\\pm$0.05$^{c}$ & 0.12\t  &   0.07 & 4.30   & 4.50$\\pm$ 0.07   & 1 \\\\ NB67-AU13  &    1.68 &   3.73 &   4.09  & 16.48$\\pm$0.02 & 15.92$\\pm$0.03 & 0.37\t  &   0.36 & $\\dots$\t & $\\dots$   & 1 \\\\ \\hline \\hline \\end{tabular} \\begin{list}{}{} \\item B and V magnitudes are from new aperture photometry performed on the CCD images of Massey et al.~(\\cite{massey}), except for B083 and B347 that are not included in the area covered by that survey. \\item[$^{\\mathrm{a}}$]  X and Y are projected coordinates in the direction along (increasing Eastward) and perpendicular to the major axis of M31, in arcmin. \\item[$^{\\mathrm{b}}$] ff is a flag indicating if the target has been selected from Table 1 or Table 2 of F05. \\item[$^{\\mathrm{c}}$] From Pap-I. \\item[$^{\\mathrm{d}}$] From the RBC. \\item[$^{\\mathrm{(t.w.)}}$] from this work: B and V from this table and E(B-V) as estimated in Sect.~3 from isochrone fitting. \\item[$^{\\mathrm{F05}}$] from Fusi Pecci et al. (2005): (B-V)$_{0}$ are calculated assuming a single value of E(B-V)=0.11 for all the clusters. \\item[$^{\\mathrm{G09}}$] from Galleti et al. (2009). \\end{list} \\end{table*}   In Perina et al. (\\cite{P09a}, hereafter Pap-I) we have described in detail the observational material coming from our survey, and the data reduction and methods of analysis that we homogeneously adopt for the whole survey. We did that by taking the brightest of our surveyed clusters (VdB0) as an example. In this contribution we apply the same process to the whole sample, obtaining metallicity, reddening and age estimates for all the targets of our survey. We incremented our final sample of candidate M31 YMC by including in the final analysis ten further  clusters having age estimates  available from the literature that are fully homogeneous with our own ones. In two companion papers, Hodge et al.~(\\cite{lessclu}, Pap-II, hereafter) identified and studied clusters of lower mass (with respect to those studied here) that were serendipitously imaged in our survey, while Barmby et al.~(\\cite{bar09}, Pap-III, hereafter) studied the structure of the clusters that are the main targets of the survey.   The paper is organized as follows. The sample is described in detail in Sect.~2, where we also summarize the data reduction procedure. In Sect.~3 we present the individual color magnitude diagrams (CMDs) and luminosity functions (LFs), we estimate ages, metallicities and reddening of each cluster. In Sect.~4 we derive the mass estimates for the clusters of our extended sample (including data from the literature), we compare our clusters with open and globular clusters of the Milky Way and we compare our estimates with those from the recent and extensive analysis of young M31 clusters by Caldwell et al. (\\cite{C09}, hereafter C09), that are based on integrated spectra. In Sect.~5 our main results are briefly summarized and discussed. Finally, in Appendix A we report on M31 clusters or candidate clusters listed in the RBC that have been serendipitously imaged within our survey, and, in Appendix B, we report on the nature of candidate BLCC=YMC M31 clusters that have an HST image in the archive, independent of this survey. ",
        "conclusions": "\\label{disc}  We presented the main results of a survey aimed at the determination of the nature  of a sample of 20 candidate YMC in the thin disk of M31 (one of which, VdB0, was studied in Pap-I). One of the targets surveyed  turned out to be a bright star projected onto the dense disk of M31, and thus erroneously classified as a possible cluster. All the other targets were revealed to be genuine star clusters and we were able to obtain  reliable CMDs for all of them. The main results from our own survey can be summarized as follows:  \\begin{enumerate}  \\item New integrated-light spectroscopy became available for many of our targets since the original selection was performed. Three of them (B083, NB16 and B347) were revealed by the new data to be not good YMC candidates as defined by F05. The CMDs obtained in this study confirms that they are likely old clusters.  \\item Among the remaining 17 targets, 16 are genuine clusters and one is in fact a star (NB67), as said above. Thus the fraction of spurious objects in our well-defined sample of BLCC=YMC is just 1/16 = 6.2\\%. Even excluding the two clusters considered at point 3., below, the incidence remains below 10\\%. The extended sample considered in Appendix~\\ref{appB} fully confirms these  results. We must conclude that M31 YMC are not especially plagued by contamination from spurious sources and most of the clusters considered in the original analysis by F05 should be real\\footnote{It may be useful to stress again that the clusters of our survey were selected among the class f=1 RBC entries, see Sect.~\\ref{sample} and Galleti et al.~(\\cite{svr}).}. In particular, {\\em asterisms}, suggested as a possible major contaminant of the sample by C06, are in fact found to be not a particular reason of concern, in this context (see also the discussion by C09).  \\item Two of the sixteen genuine clusters (B374 and B222) have integrated properties compatible with being YMCs but they do not show a detectable MS in the range of magnitudes sampled by our CMDs. We can provide only an upper limit to the age of these clusters ($\\ga 300$ Myr), but the available data suggest that they are good candidate intermediate-age clusters that indeed would merit follow-up with deeper HST photometry.  \\item The fourteen confirmed young clusters (including VdB0, studied in Pap-I) show a clear MS in the range of magnitudes sampled by our CMDs, hence we were able to obtain reliable estimates of their ages, reddenings and (an educated guess of) metallicities by comparison of the observed CMD and LF with theoretical models. Ten of them have ages in the range 25-100 Myr, the other four range between 140 Myr and 280 Myr. The adopted metallicities include $Z=0.004$ (one case), $Z=0.008$ (three cases), and $Z=0.019$ (solar metallicity, ten cases). The estimated reddenings range from E(B-V)=0.06 to E(B-V)=0.60, with E(B-V)=0.20-0.30 as most typical values.  \\end{enumerate}    \\begin{figure} \\centering \\includegraphics[width=9cm]{histo_mass.ps} \\caption{Upper panel: The mass distribution of the sample of YMC studied here (from Tab.~\\ref{table:3}, thick continuous line) is compared with the mass distribution of Galactic OCs (dotted line) and Galactic globular clusters (dashed lines). Masses of Galactic clusters are from B08. Lower panel: zoomed view of the distribution   of M31 YMC  compared with the distribution of the YMC of the Milky Way (data from M09).} \\label{histo} \\end{figure}    \\begin{figure*} \\centering \\includegraphics[width=16cm]{age_mass.ps} \\caption{Comparison between Galactic OCs (small filled circles), M31 YMC from the present study (big open squares), MW YMC from M09 (big open circles), M31's clusters from pap-II (small open squares), Magellanic Clouds clusters (grey open pentagons), and M33's clusters (grey crosses) in the log(age) $vs.$ log Mass plane. Masses of Galactic OCs are from B08, masses of Magellanic Clouds clusters are from \\cite{mclaug} and masses of M33 clusters are from \\cite{sanrom}. For M33 and the Magellanic Clouds  only clusters younger then 10 Gyr are shown.} \\label{agemass} \\end{figure*}  To increment our final sample of YMC we included ten further clusters for which the age was estimated from their CMDs (obtained from HST imaging) with methods  strictly homogeneous with those adopted here, from WH01 and P09b. In this way we assembled a final sample of 24 confirmed young clusters. For 22 of these we were able to obtain reliable estimates of the total stellar mass by coupling our age estimates with the integrated J,H,K magnitudes taken from the 2MASS-6X catalog. These clusters have masses ranging from $0.6\\times 10^4 M_{\\sun}$ to $6\\times 10^4 M_{\\sun}$, with an average of $\\sim 3\\times 10^4 M_{\\sun}$\\footnote{The remaining two clusters, that lack NIR photometry, also have masses lying in the same range, according to the estimates obtained using the integrated V magnitude instead of J,H,K ones.}. Our estimates of ages and masses are in good agreement with recent independent studies based on integrated light spectra (see also Pap-III for the comparison with the results by Pfalzner \\cite{pfalz}).  \\subsection{The nature of M31 YMC}  In the upper panel of Fig.~\\ref{histo} the mass distribution of our extended sample of M31 YMCs is compared with the distributions of Galactic OCs and GCs (masses from B08). The clusters considered here appear to lie in the middle of the two distributions, overlapping with the high-mass end of the OCs and with the low-mass end of GCs. This comparison provide a further confirmation that the YMCs (=BLCCs) of M31 are indeed more similar to the YMCs of the LMC than to classical OCs of the Milky Way, i.e. the original hypothesis advanced in F05. This is in full agreement with the main conclusions by C09, obtained with a completely independent method (less sensitive to age than ours) on a wider sample.  The lower panel of Fig.~\\ref{histo} compares our clusters with the YMCs seen toward the center of the Milky Way as listed by M09. The two samples have very similar mass distributions, suggesting that they are also similar in nature. An obvious difference between the two sets of clusters was already suggested in Pap-I and is confirmed here: the M31 YMCs of our sample are significantly older that the YMC discovered until now in the Galaxy ($\\ga 50$ Myr vs. $\\la 20$ Myr; see below for possible explanations). We confirm that the M31 YMCs studied here have larger sizes (half-light-radii) with  respect to their MW counterparts (see Pap-I and Pap-III); this seems in agreement with the age-size relations proposed by Pfalzner~(\\cite{pfalz}; see Pap-III for discussion).   A more thorough comparison between various samples of YMCs is presented in Fig.~\\ref{agemass}, where Galactic OCs and YMCs, YMCs from M33 (San Roman et al.~\\cite{sanrom}; for further discussion on M33's star clusters see Sarajedini \\& Mancone \\cite{sarajed}, Zloczewski et al.~\\cite{zlocz}, Park et al.~\\cite{park}), the LMC, the Small Magellanic Cloud (McLaughlin \\& van der Marel \\cite{mclaug}), and  M31 are plotted together in a log(age) $vs.$ log Mass diagram. Fig.~\\ref{agemass} is affected by a number of selection effects that deserve to be described in some detail.  \\begin{enumerate}  \\item The minimum mass threshold appears to increase with age (at least for age $\\ga$ 10 Myr, see the Galactic OCs if Fig.~\\ref{agemass}): this is due to the fact that the lower the mass of a cluster, the shorter is its dissolution time, as the cluster is less resilient to all the internal and external effects that may lead to its disruption (Gieles et al. \\cite{gieles}, Pap-III, and references therein). The minimum mass threshold  for samples in external galaxies is obviously due to the inherent magnitude limits.  \\item Also the maximum mass threshold increases with age in log~Age vs. log~Mass plots (Hunter et al. \\cite{hunter}; Gieles \\cite{giel09}; the effect is clearly evident in Fig.~\\ref{agemass} if one looks at the MW OCs, that cover the widest range in ages). This general behavior can be easily explained as a simple consequence of varying the sample size as a function of the age bin in the logarithmic scale. Assuming a power-law mass function and a constant Cluster Formation Rate (CFR) the number of cluster per logarithmic age bin increases with age. For an exponent of the power law mass function ($N(M)\\propto M^{-\\alpha}$) $\\alpha=2$, that is a reasonable approximation for most of the observed cluster systems, log~$M_{max} \\propto$ log~Age (see Gieles \\cite{giel09}, for detailed discussion and references).  \\item While the lack of massive ($M\\ga 10^4 M_{\\sun}$) clusters older than 400 Myr in the Milky Way is probably real, the typical limiting magnitude ($V\\sim 27$, Rich et al. \\cite{rich}) of available CMDs of M31 clusters prevent us from drawing firm general conclusions about objects in that age range in M31. The cases of B222 and B374, treated here, are excellent examples of clusters that may populate that region of the diagram but lack a reliable age estimate because the available photometry is too shallow (see Puzia et al.~\\cite{puzia}).  \\item The lack of massive (log~$(M/M_{\\sun})>3.6$) M31 clusters younger than 25-50 Myr may be due to the contribution of several biases. First, such young clusters may be hard to select from the RBC as there are no objects bluer than $(B-V)_0\\simeq 0.0$  in the list of confirmed clusters (see F05). This is not surprising as the RBC was intended to be a catalog of globular clusters. Second, for ages $\\la$ 8 Myr the $H_{\\beta}$ index is expected to fall below the threshold adopted to select YMC candidates (see, for example, Fig.~7 of F05), thus (possibly) preventing the selection of these objects for our survey. Third, very young objects should have their luminosity dominated by a few massive stars near their centers, thus leading to objects that may appear more like blended stars than like a star cluster at the distance of M31, even in HST images, thus preventing their inclusions in lists of candidate YMCs. Fourth, it can be hypothesized a positive correlation between the age of the clusters and their height above the disk plane, such that the youngest clusters are more deeply embedded in the thin dust layer of the M31 disc, out of our reach even from our privileged point of view, while most/some of the older clusters would be visible just because they lie above the densest part of that layer. There are indications that this kind of correlation actually holds in our own Galaxy (V.D. Ivanov, private communication).  \\item The lack of massive (log~$(M/M_{\\sun})>3.6$) MW clusters {\\em older} than 25-50 Myr may also be associated with an observational bias. Galactic YMC have been identified as clumps of bright stars in the near and mid IR and the youngest clusters, having the brightest RSG, are easier to detect in this way. Moreover the sample of Open/YM Galactic clusters is limited (essentially by the effect of interstellar extinction in the Galactic disc) to a volume of a few kpc around the Sun, while M31 (or M33) YMCs can be selected over the whole disk of their parent galaxy, thus introducing a bias that favors the detection of rarer cluster species (massive clusters) in the latter galaxies with respect to the MW.  \\item There seems to be a significantly under-dense region in Fig.~\\ref{agemass}, for masses $\\ga 10^3~M_{\\sun}$ and ages between $\\sim 15$ Myr and $\\sim 50$ Myr ($7.2\\la$~log~Age$\\la 7.7$). The same feature was noted by Whitmore et al. (\\cite{whitmore}) in their study of the cluster system of the Antennae and it was attributed by a degeneracy in age dating from broad band colors occurring in that age range due to the prompt onset of the RSG phase (see Whitmore et al. \\cite{whitmore}, for details, discussion and further references). Virtually all the clusters plotted in Fig.~\\ref{agemass} had their ages estimated from the CMD of their stars (instead of broad-band colors, see also Pap-II), hence our sample should not be affected by this bias, at least in principle. However the coincidence of the feature with that noted by Whitmore et al. (\\cite{whitmore}) suggests that the same kind of bias against ages in that interval may be at work also in Fig.~\\ref{agemass}.  \\item The samples of clusters from all the galaxies involved in Fig.~\\ref{agemass} have been selected according to different criteria, by color, magnitude, etc.  \\end{enumerate}  Given all the above considerations, it does not seem possible to draw any firm conclusion from the comparison shown in Fig.~\\ref{agemass}. The only straightforward conclusion is that {\\em YMCs in the age range 50-500 Myr are relatively common in all the most massive star-forming galaxies of the Local Group} (M31, M33, LMC and SMC). The only exception (the Milky Way) may be ascribable to observational biases, but it cannot be excluded that it is instead (at least partly) associated with intrinsic properties of the Milky Way, that appears peculiar under several aspects with respect to the typical spiral galaxies (and to M31, in particular see Hammer et al.~\\cite{hammer}, and  Yin et al.~\\cite{yin}). As the samples of M33 and M31 should be subject to the same kind of biases (as the distances are similar and the data have been collected with HST in both cases), the difference in the maximum mass limit between the two samples is likely real, and it can probably be ascribed to the difference in total mass between the discs of the two galaxies: larger discs should host more numerous populations of clusters, thus enhancing the probability of producing clusters with higher (maximum) masses (see Gieles \\cite{giel09}, and references therein).   \\subsection{Radial trends}  Given the wealth of data collected for our target clusters, it may be useful to look for correlations between their physical parameters, including their position within the M31 disc. Limiting the analysis to the young clusters (age $<$ 1 Gyr), that constitute a more homogeneous sample of bona-fide thin disk objects, it turns out that our sample is still too sparse for a thorough analysis of these correlations. In particular the covered ranges of age, mass and position are quite limited, thus not allowing us to reveal large scale trends, in most cases. Moreover, the adopted approach of CMD analysis provides just an educated guess of the metallicity of the clusters, aimed at obtaining the most reliable estimate of the clusters age, which was the main objective of our analysis. These limitations prevent the possibility of a meaningful study of the radial metallicity gradient with our data. It should also be recalled that the correlations bewteen the structural parameters of the clusters (mass, radius, density etc.) have already been discussed in Pap-III, hence here we consider only age, mass, de-projected galactocentric distance ($R_d$; assuming and inclination of $i=12.5\\degr$ of the disk with respect to the plane of the sky, see Simien et al. \\cite{simien} and Pritchet \\& van den Bergh \\cite{prit}), X, Y, and reddening.  Having checked all the combination of parameters, the only correlation  that appeared remarkable to us is presented in Fig.~\\ref{radtrend}.  It is a trend of decreasing age with galactocentric distance, that seems statistically significant if one consider the associated errors. Given the relatively limited range of galactocentric distance covered, in our view the observed distribution can be interpreted in two ways:  \\begin{itemize}  \\item as a part of a larger trend resulting from a inside-out wave of cluster formation. In this case the trend toward older mean ages should continue at lower radii and Fig.~\\ref{radtrend} shows the transition between a regime of decreasing age with galactocentric distance and an asymptotic regime of constant age in the outermost fringes of the disc;  \\item more likely, as a sharp transition in the epoch of the highest rate of star/cluster formation occurring at the onset of the $R_d\\sim 10$ kpc ``ring of fire''. This would be consistent with the well known burst of recent star formation that characterize this prominent structure of the M31 disc.  \\end{itemize}  While not especially conclusive or insightful, the result shown in Fig.~\\ref{radtrend} gives a clear idea of how useful YMCs can be as tracers of the structure and evolution of the disk itself, in particular if large and reliable samples can be assembled.   \\begin{figure} \\centering \\includegraphics[width=9cm]{rad_trend.ps} \\caption{Age as a function of the deprojected galactocentric distance for the young clusters (open squares with error bars). The cluster VdB0 has been labeled as it is by far, the youngest of the whole sample.} \\label{radtrend} \\end{figure}     \\subsection{Final remarks}  This research has demonstrated that the conspicuous population of bright disk objects studied by F05 consists of genuine YMC, similar to those found in the LMC, SMC and M33 galaxies. These clusters may open a new window to the study of the recent star formation history in the disk of M31. A systematic analysis over the whole extent of the M31 disk may provide the opportunity to study a rich system of young clusters using a sample much less affected by selection biases than in our own Galaxy, and to better constrain the models of dynamical evolution of clusters within the discs of spiral galaxies. M31 YMCs like those studied here provide also an excellent tracer of the disk kinematics in that galaxy, independent of (and in addition to) the HI gas. Recent wide-field surveys (Vansevicius et al.~\\cite{vanse}; see also Pap-II) suggest that a rich harvest of genuine YMCs await to be discovered in the disk of our next neighbor giant galaxy in Andromeda."
    },
    "0911/0911.3485_arXiv.txt": {
        "abstract": "We report the discovery with \\emph{Fermi}/LAT of $\\gamma-$ray emission from three radio-loud narrow-line Seyfert 1 galaxies: PKS~1502+036 ($z=0.409$), 1H~0323+342 ($z=0.061$) and PKS~2004-447 ($z=0.24$). In addition to PMN J0948+0022 ($z=0.585$), the first source of this type to be detected in $\\gamma$ rays, they may form an emerging new class of $\\gamma-$ray active galactic nuclei (AGN). These findings can have strong implications on our knowledge about relativistic jets and the unified model of AGN. ",
        "introduction": "The advent of the \\emph{Fermi Gamma-ray Space Telescope} (hereafter, \\emph{Fermi}), with its excellent performance, is changing our perception of the sky at high-energy $\\gamma$ rays. Specifically, the recent detection in the MeV-GeV energy range of the radio-loud narrow-line Seyfert 1 (RL-NLS1) quasar PMN J0948+0022 strongly supports the presence of a closely aligned relativistic jet in this peculiar system (Abdo et al. 2009a,c; Foschini et al. 2009a). This is quite surprising, since NLS1s are generally hosted in spiral galaxies and the presence of a fully developed relativistic jet is contrary to the well-known paradigm that this type of system is associated to ellipticals (e.g. see Marscher 2009 for a recent review).  Already in the past, several authors inferred from the multiwavelength properties of RL-NLS1s some parallelism between this type of source and blazars (e.g. Komossa et al. 2006, Yuan et al. 2008, Foschini et al. 2009b), suggesting similarities with different flavors of the blazar types (quasars or BL Lacs). What was missing in these previous studies is the $\\gamma-$ray detection, which is important to confirm the presence of a relativistic jet, to measure its power, and to study the characteristics of this type of sources by modeling their spectral energy distributions (SEDs). Indeed, blazar-like radio emission from radio-quiet AGNs has been already reported (e.g., Brunthaler et al. 2000, Brunthaler et al. 2005, Lister et al. 2009; see also the review by Ho 2008), but no $\\gamma$ rays have been yet detected from these sources. This suggests some evident differences between fully developed relativistic jets (i.e. emitting over the whole electromagnetic spectrum, from radio to $\\gamma$ rays) and jet-like (perhaps aborted?, cf Ghisellini et al. 2004) structures, whose emission at $\\gamma$ rays in the MeV-TeV energy range -- if any -- has not been yet reported.  Understanding these differences could give important insights into jet formation. Therefore, we started a larger program aiming at the detection of $\\gamma-$ray emission from other RL-NLS1s. Since no complete sample of RL-NLS1s is available in the literature, we have built a sample by merging all the sources of this type in the catalog of Zhou \\& Wang (2002) and in the lists of Komossa et al. (2006) and Yuan et al. (2008). We have adopted a threshold in the radio-loudness $R_L=f_{\\rm 5~GHz}/f_{\\rm 440~nm}=50$, in order to avoid doubtful sources at the border ($R_L=10$) with radio-quietness. The optical properties are those typical of NLS1s, i.e. narrow permitted lines (FWHM H$\\beta$$<2000$~km~s$^{-1}$), [OIII]/H$\\beta <3$ and the bump of FeII (see Pogge 2000 for a review). At radio frequencies, RL-NLS1s display strong and variable radio emission, with brightness temperature well above the inverse-Compton limit (see, e.g., Yuan et al. 2008, Table 3).  This selection resulted in a list of 29 sources, most of them selected from the \\emph{Sloan Digital Sky Survey} (SDSS). This low number should not be a surprise. According to Komossa et al. (2006), RL-NLS1s are more rare than quasars: 7\\% of NLS1s have $R_L>10$ and only 2.5\\% exceed $R_L=50$, while generally 10-20\\% of quasars are radio-loud. Although based on a small sample of objects (128), these results have been confirmed by studies on larger samples from the SDSS (Zhou et al. 2006, Whalen et al. 2006). It is worth noting that this sample is not complete, although the fact that it is based on the SDSS ensures a minimum degree of representativeness.  Here we report three new detections at $\\gamma$ rays of RL-NLS1s obtained with \\emph{Fermi}/LAT. In addition to the already known PMN J0948+0022 (reported by Abdo et al. 2009a,c; Foschini et al. 2009a), this increases the number of $\\gamma-$ray detections of NLS1s ($3+1=4$ out of 29 sources, $\\approx 14$\\% of our sample), suggesting that they form a new class of $\\gamma-$ray emitting AGNs.  Throughout this work, we adopted a $\\Lambda$CDM cosmology from the most recent \\emph{WMAP} results, which give the following values for the cosmological parameters: $h = 0.71$, $\\Omega_m = 0.27$, $\\Omega_\\Lambda = 0.73$ and with the Hubble-Lema\\^{i}tre constant $H_0=100h$ km s$^{-1}$ Mpc$^{-1}$ (Komatsu et al. 2009). ",
        "conclusions": "The analysis of the individual sources reveals some things worth noting. In 1H 0323+342 the dissipation region needs to be located very close to the central black hole ($\\sim 650$ times the Schwarzchild radius $R_{\\rm S}$) and hence a quite high magnetic field (up to $30$ gauss) is needed. We note that PKS 2004-447 seems to be different with respect to the three other RL-NLS1s, with optical/UV data sampling the synchrotron emission instead of the accretion disk. Indeed, the classification of this source is still open: while Oshlack et al. (2001) classified it as genuine RL-NLS1, Komossa et al. (2006) suggested it can be a narrow-line radio galaxy, on the basis that the bump of FeII is weak. Gallo et al. (2006) noted that there is no specific prescription on the value of the strength of the FeII bump and confirmed the NLS1 classification. The same authors noted that PKS 2004-447 is a compact steep-spectrum (CSS) radio source. On the other hand, this source is also in the CRATES catalog of flat spectrum radio sources (Healey et al. 2007), which includes also the three others NLS1s in this work. From the present study, the estimated jet power of PKS 2004-447 is well above the range typical of radio galaxies and, therefore, we favor the hypothesis of a genuine NLS1. However, in this case the classification remains uncertain and should be studied with further observations.  A synoptic analysis of the 4 RL-NLS1s in the present work shows that PKS 1502+036 carries a jet power comparable to PMN J0948+0022, while the remaining two sources are significantly less powerful -- even two order of magnitudes for proton powers (Table~\\ref{tab:outputmodel}). The comparison with the large sample of blazars reported in Celotti \\& Ghisellini (2008) and Ghisellini et al. (2009b), shows that the jet powers of PKS 1502+036 and PMN J0948+0022 are in the region of quasars, while 1H 0323+342 and PKS 2004-447 are in the range typical of BL Lac Objects. The $\\gamma-$ray emitting RL-NLS1s in our sample have small masses and high accretion rates. This is in contradiction to the larger masses and and smaller accretion rates expected from low-power blazars (HFSRQs) with strong emission lines, as suggested by Yuan et al. (2008). Anyway, when comparing the calculated powers with the distribution of powers in blazars (quasars plus BL Lacs; see Ghisellini et al. 2009b), RL-NLS1s are in the average range.  The main differences with respect to blazars are in the masses and accretion rates. The masses of the three newly discovered RL-NLS1s are around $10^{7}M_{\\odot}$, in agreement within one order of magnitude with the values obtained with other methods. In the case of 1H 0323+342, Zhou et al. (2007) found values of $1.8\\times 10^{7}M_{\\odot}$ with $H\\beta$ luminosity and $3\\times 10^{7}M_{\\odot}$ with the luminosity of the continuum at 5100\\AA. In the case of PKS 1502+036, Yuan et al. (2008) estimate the mass to be $4\\times 10^{6}M_{\\odot}$, by means of the virial method. These value are about 1-2 orders of magnitude lower than the typical blazar masses ($\\approx 10^{9}M_{\\odot}$, see Ghisellini et al. 2009b). The accretion rates can reach extreme values, up to 80 or even 90\\% the Eddington values in the cases of PKS 1502+036 and 1H 0323+342, respectively. These values are the most extreme ever found in any $\\gamma-$ray emitting AGNs, but usual for NLS1s.  What is at odds with this scenario is the type of host galaxy, which is elliptical in all blazars, while is likely to be spiral in RL-NLS1s (see, e.g. Zhou et al. 2006). In the case of 1H 0323+342, analysis of optical observations suggests two possibilities: Zhou et al. (2007) show the spiral arms of the host galaxy by means of observations with the \\emph{Hubble Space Telescope}, while Ant\\`on et al. (2008), on the basis of ground-based observations with Nordic Optical Telescope (NOT), suggest that these structures are the residual of a merging that occurred within the past $10^8$ years. This means that relativistic jets can form and develop independently of their host galaxies, with quasars, BL Lacs and, now, RL-NLS1s jointly characterized by the presence of a relativistic jet and with the differences in their observed SEDs mainly determined by their masses and accretion rates."
    },
    "0911/0911.1355_arXiv.txt": {
        "abstract": "We use 62,185 quasars from the Sloan Digital Sky Survey DR5 sample to explore the relationship between black hole mass and luminosity.  Black hole masses were estimated based on the widths of their H$\\beta$, Mg{\\small II}, and C{\\small IV} lines and adjacent continuum luminosities using standard virial mass estimate scaling laws.  We find that, over the range $0.2 < z < 4.0$, the most luminous low-mass quasars are at their Eddington luminosity, but the most luminous high-mass quasars in each redshift bin fall short of their Eddington luminosities, with the shortfall of order ten or more at $0.2 < z < 0.6$.  We examine several potential sources of measurement uncertainty or bias and show that none of them can account for this effect.  We also show the statistical uncertainty in virial mass estimation to have an upper bound of $\\sim 0.15$ dex, smaller than the 0.4 dex previously reported.  We also examine the highest-mass quasars in every redshift bin in an effort to learn more about quasars that are about to cease their luminous accretion.  We conclude that the quasar mass-luminosity locus contains a number of new puzzles that must be explained theoretically. ",
        "introduction": "\\label{sec:intro}  Supermassive black holes (SMBH), with masses between $\\sim 10^6 M_\\odot$ and $\\sim 10^9 M_\\odot$, are found at the center of nearly every galaxy where there have been sensitive searches.  While we suspect that the seeds for these SMBH might all have a common origin, their formation mechanism is not well understood.  Many galaxies at redshifts $z \\sim 2$ contain quasars, i.e., SMBH in the midst of luminous accretion.  The Soltan argument \\cite{Merritt2001} suggests black hole masses are largely accounted for via growth due to luminous accretion.  There are far fewer quasars at low redshift \\cite{Schmidt1983,Richards2006}, implying that at some point, SMBH cease their luminous accretion.  The quasar {\\em turnoff} mechanism is not well understood (Thacker et al. 2006\\nocite{turnoffreview}).  Finally, the black hole mass - stellar velocity ($M-\\sigma$) relation \\cite{msigma1,msigma2} suggests that SMBH are in some way co-evolving with their host galaxies, but the $M-\\sigma$ relation merely describes an end state and is not a theoretical explanation.  As we know of no rapid process by which SMBH can lose mass, the evolutionary tracks for SMBH consist of stages at progressively higher masses.  These stages must, in order, involve (1) a formation mechanism, (2) a period of luminous growth (the `quasar phase'), perhaps along with a period of nonluminous growth (`turnoff'), and (3) a period in which SMBH lie at the centers of galaxies without substantial growth, as we observe them today. Much current theoretical research concerns the origin and turnoff phases of quasar evolution, while the `quasar phase' appears to be relatively well understood.  In this paper we find a new feature of SMBH evolution during their quasar phase.  We examine the evolution of the quasar locus in mass-luminosity space as a function of redshift.  The $M-L$ locus is traditionally shown with all quasars on the same plot (as in Figure \\ref{fig:allml}).  The large size of the Sloan Digital Sky Survey (SDSS) DR5 quasar catalogue \\cite{SDSSDR5} allows a subdivision into several redshift bins, each containing thousands of quasars.  The Eddington limit produces an absolute upper bound on the luminosity of quasars which is proportional to the black hole mass, $L_{Edd}=1.3\\times 10^{46}(M/10^8M_\\odot)~{\\rm erg~s^{-1}}$ \\cite{Shapiro}.  While strictly applicable only for a spherical accretion flow of ionized gas, models for more realistic accretion configurations with rotation are still limited by a luminosity of this order (except under special circumstances, e.g. Begelman 2002)\\nocite{Begelman}.  Using a smaller data set ($N = 733$) than SDSS, Kollmeier et al. (2006)\\nocite{Kollmeier2005} appeared to confirm the applicability of $L_{Edd}$ to quasars, using an $M-L$ locus (similar to Figure \\ref{fig:allml}) to show that the most luminous quasars reach but do not exceed $L_{Ed}$.  For the lowest black-hole masses at every redshift, we confirm the conclusion of \\cite{Kollmeier2005}, but we show that the quasars with the highest SMBH mass at every redshift fall well short of $L_{Edd}$.  In {\\S~\\ref{sec:methods}}, we review the methods used to estimate masses and bolometric luminosities.  While we have added nothing original to this methodology, the remainder of our results are entirely dependent upon its accuracy.  In {\\S~\\ref{sec:lm}}, we subdivide the SDSS DR5 quasar catalogue by redshift and show that the quasar mass-luminosity distribution does not match what we should expect given our current theoretical understanding of quasar accretion.  In particular, we show that instead a sub-Eddington boundary (SEB) is present in each redshift bin.  We also use our limited statistics to make a first estimate for how the SEB evolves with redshift.  In {\\S~\\ref{sec:tests}}, we consider several alternative explanations for the discrepancies between the observed quasar locus and the Eddington luminosity, focusing on potential sources of measurement uncertainty or bias.  In {\\S~\\ref{sec:discussion}}, we comment on the potential implications of our new results. ",
        "conclusions": "\\end{table}  There is a small overlap in redshift between the H$\\beta$ mass and Mg{\\small II} mass samples as well as one between the Mg{\\small II} mass and C{\\small IV} mass samples (Table \\ref{table:vmesummary}).  Shen et al. (2008) show that the agreement between mass estimates for the same object taken using the H$\\beta$ and Mg{\\small II} relations (0.22 dex dispersion), is better than between estimates using the Mg{\\small II} and C{\\small IV} relations (0.34 dex, and correlated with the C{\\small IV}-Mg{\\small II} blueshift).  The C{\\small IV} mass estimates may be less accurate and we therefore focus here almost exclusively on masses obtained using the other two scaling relations.  We consider the implications of this disagreement further in \\S~{\\ref{sec:tests}}.    \\label{sec:discussion}  Quasar catalogues such as the Sloan Digital Sky Survey are now large enough to yield useful statistical conclusions after being subdivided.  In this work, we have explored just a few of the many ways in which the catalogue might be divided.  It is immediately apparent that when the quasar distribution is considered at any given redshift, the most massive central black holes do not accrete at their Eddington luminosity, but rather all fall well short of $L_E$.  While there are many measurements that contribute to these quasar distributions, each with its own set of assumptions, a close examination reveals no evidence that any these are responsible for massive quasars remaining sub-Eddington.  Rather, it appears that this `sub-Eddington boundary' is a new physical limit for quasars.  Thus, quasar accretion must be more complicated than had been previously thought.  The sub-Eddington boundary helps to recast the problem in a two-dimensional form.  This form emphasizes the luminosity as being a function of the mass and accretion rate, fundamental properties of the quasar, in a non-trivial way.  Quasar luminosity functions and quasar mass functions are both one-dimensional projections of the mass-luminosity plane.  The full two-dimensional quasar distribution contains complexities difficult to discern from either of these projections.  The sub-Eddington boundary presents a theoretical puzzle.  Not only must an improved theory explain the sub-Eddington boundary, but it must also explain its detailed shape and explain the evolution of that shape with redshift.  While the slope of the sub-Eddington boundary may not vary greatly with redshift, the mass scale at which the boundary becomes relevant is larger at higher redshift.  The results in this paper can be used to develop a series of tests which should be capable of discriminating between models.  Given the complications introduced by the sub-Eddington boundary, the case could be made that we really understand surprisingly little about the supermassive black holes that appear to be at the center of nearly every galaxy.  Not only is their seeding mechanism unknown, but as shown in this paper, the growth mechanism is quite poorly understood, contrary to expectations.  The Eddington limit is relevant at low black hole masses, but is only part of the story.  The sub-Eddington boundary developed in this work is the latest addition to a growing collection of puzzles regarding every phase of the evolution of galactic nuclei and surrounding regions.  The authors would like to thank Mihail Amarie, Forrest Collman, Doug Finkbeiner, Margaret Geller, Lars Hernquist, Gillian Knapp, Avi Loeb, Ramesh Narayan, Jerry Ostriker, and Michael Strauss for valuable comments.  This work was supported in part by Chandra grant number GO7-8136A (Chandra X-ray Center)."
    },
    "0911/0911.3216_arXiv.txt": {
        "abstract": "A new list of physical characteristics of 3914 astrometric radio sources, including all 717 ICRF-Ext.2 sources, observed during IVS and NRAO VCS sessions have been compiled. The list includes source type, redshift and visual magnitude (if available). In case of doubt detailed comment is provided. The list of sources with their positions was taken from the Goddard VLBI astrometric catalog with addition of two ICRF-Ext.2 sources.  At this stage the source characteristics were mainly taken from the NASA/IPAC Extragalactic Database (NED). 667 sources from our list are included into the IERS list. Comparison has shown a significant difference in characteristics for about half of these 667 common sources. We compiled a list of frequently observed sources without known physical characteristics for urgent optical identification and spectrophotometric observations with large optical telescopes. This presented list of physical characteristics can be used as a supplement material for the ICRF-2, as well as a database for kinematic studies of the Universe and other related works, including scheduling of dedicated IVS programs. ",
        "introduction": "Information on physical characteristics of the geodetic radio sources is important for planning of VLBI experiments and analysis of VLBI data to do a research in cosmology, kinematics of the Universe, etc.  In particular, the primary mainspring to this work was a support of the investigation of the systematic effects in apparent proper motion of geodetic radio sources. \\cite{MacMillan05,Titov2007}.  The official list of the physical characteristics of the ICRF radio sources is supported by the IERS ICRS Product Center (\\cite{Archival97}). The latest version of the IERS list is available in the Internet\\footnote{http://hpiers.obspm.fr/icrs-pc/info/car\\_physique\\_ext1}. However this list has some deficiencies: \\begin{itemize} \\itemsep=-0.3ex \\item Not all the sources observed in the framework of geodetic and astrometric experiments are included in the IERS list. \\item The characteristics of some sources in the IERS list are outdated or doubtful. \\end{itemize}  To overcome this problems, we performed a compilation of new list of the physical characteristics of geodetic radio sources using the latest information. In this paper we present our result of the first stage of this work.  The list of radio sources with their positions was taken from the Goddard VLBI astrometric catalogs\\footnote{http://vlbi.gsfc.nasa.gov/solutions/astro}, version 2007c, with removing duplicate source 1616+85A (L.Petrov, private communication) and addition of two ICRF-Ext.2 \\cite{Fey04} sources 1039-474 and 1329-665 not included in the Goddard catalog. It provides 3914 radio sources in total.  At this stage mainly the NASA/IPAC Extragalactic Database (NED)\\footnote{http://nedwww.ipac.caltech.edu/} was scoured. Some sources were checked with the CfA-Arizona Space Telescope LEns Survey (CASTLES)\\footnote{http://cfa-www.harvard.edu/glensdata/} and the HyperLeda\\footnote{http://leda.univ-lyon1.fr/} databases. In this list we have included only the optical characteristics of geodetic radio sources: source type, redshift and visual magnitude. The flux parameters are not included in our list because they are available from other centers. ",
        "conclusions": "A new list of the optical characteristics of geodetic radio sources has been compiled and available at  \\centerline{\\tt http://www.gao.spb.ru/english/as/ac\\_vlbi/sou\\_car.dat.}  This is only the first stage of our work. We are planning the following steps: \\begin{itemize} \\itemsep=-0.3ex \\item To continue searching for the missing and checking out the ambiguous characteristics through literature and astronomical databases. \\item To organize photometric and spectroscopy observations of geodetic radio sources with large optical telescopes.  In particular, such observational program has been included in the plan of the Pulkovo Observatory for 2008. The application for observation time on the Russian 6-meter BTA telescope for the second half of 2008 was handed over in February in cooperation with Pulkovo astrophysicists Kirill Maslennikov and Alexandra Boldycheva. \\end{itemize}  The authors would be happy to know whether this new list is useful either as a database for VLBI data analysts or as a supplement material for the ICRF-2 compiling.  We hope that this work will continue in cooperation with other interested groups."
    },
    "0911/0911.1213_arXiv.txt": {
        "abstract": " ",
        "introduction": "What kind of fundamental physics stands behind inflation in the very early Universe? When looking for the answer to the above question one often points towards the string theory. The intensive studies of inflation in the string theory began with the invention of stable de Sitter (dS) vacua in the famous KKLT paper \\cite{kklt}. Construction of such vacua is inevitable for inflationary model building because the inflaton (the field which drives inflation) should end its evolution in this kind of vacuum. The KKLT procedure consists of three steps. In the first step, the dilaton and the complex structure moduli are stabilized in a supersymmetry (SUSY) preserving minimum by turning on some non-trivial fluxes \\cite{Giddings}. In the second step, the K\\\"ahler moduli, including the volume modulus, are stabilized by some non-perturbative effects such as gaugino condensation or instantons.\\footnote{For the recent discussion on the reliability of the effective theory for the K\\\"ahler moduli see \\cite{addG,Gallego,Brizi}. } At this stage SUSY remains unbroken and the vacuum energy is negative. In the last step, the $\\ov{D3}$-branes are introduced to break supersymmetry and to uplift the minimum to a dS space. However, in the effective field theoretical description the $\\ov{D3}$-branes break supersymmetry explicitly. This is the main drawback of the KKLT model. The last step of the KKLT mechanism has been improved in the papers \\cite{matter}-\\cite{duplifting_lust} in which dS vacua were obtained due to spontaneous SUSY breaking.   Many models of inflation have been constructed within the KKLT framework \\cite{stringinf}. Especially interesting scenarios are those in which the volume modulus plays the role of the inflaton. Such models using a racetrack superpotential for the modulus, involving two non-perturbative terms, were proposed in \\cite{racetrack,lw}. With this kind of superpotential two different inflationary scenarios have been realized. In the first of them, inflation takes place in the vicinity of a saddle point of the potential with the axion, associated with the volume modulus, being the inflaton \\cite{racetrack}. In the second scenario, the real part of the volume modulus is the inflaton and inflation takes place in the vicinity of an inflection point of the potential \\cite{lw}. Unfortunately, in both scenarios SUSY breaking and uplifting are due to $\\ov{D3}$-branes which break supersymmetry explicitly.   Even though many forms of manifestly supersymmetric uplifting have been developed, a vast majority of them have not been applied to inflationary scenarios. To the best of our knowledge only two papers address the issue of spontaneous SUSY breaking in racetrack inflation models. The $D$-term uplifting \\cite{achucarro} of the saddle point racetrack inflation model \\cite{racetrack} was performed in \\cite{dupracetrack}. The string theory $\\alpha'$-corrections were used to realize saddle point racetrack inflation in a fully supersymmetric framework in \\cite{westphal}. In the latter approach the volume modulus $F$-term is responsible for uplifting and SUSY breaking.   In the present work we investigate $F$-term uplifting in both, saddle point \\cite{racetrack} and inflection point \\cite{lw}, racetrack inflation scenarios. We assume that the source of SUSY breaking is a hidden sector matter field. We consider two models of $F$-term SUSY breaking. The first one is the Polonyi model \\cite{polonyi}, while the second one is the quantum corrected O'Raifeartaigh model \\cite{O'Raifeartaigh,okklt}. In the case of the racetrack inflation model coupled to the Polonyi model, we show that both scenarios (saddle point and inflection point ones) significantly differ from the original ones with $\\ov{D3}$-brane uplifting. In our scenarios, the matter field dominates SUSY breaking as well as inflationary dynamics. In addition, fine-tuning of the parameters required for successful inflation is relaxed as compared to original models with non-SUSY uplifting, especially in the inflection point scenario. The situation is much different in the case of models coupled to the quantum corrected O'Raifeartaigh model. We find, rather unexpectedly, that the matter field responsible for SUSY breaking can be decoupled from the inflationary dynamics. This is a realization of ``pure'' $F$-term uplifting which does not alter significantly any feature of the original model except the fact that supersymmetry is broken spontaneously.   In most of the inflationary models based on the KKLT moduli stabilization, including those described above, the gravitino mass in the post-inflationary vacuum exceeds the Hubble scale during inflation \\cite{kl}. Typically, the scale of inflation is much larger than the electroweak scale while the gravitino mass sets a lower bound on the scale of SUSY breaking. This implies that such models are incompatible with a TeV-scale supersymmetry. The above mentioned correlation between the scale of inflation and the gravitino mass can be avoided if inflation ends in a SUSY (near) Minkowski minimum \\cite{kl}. However, it is not easy to construct a model realizing this idea. It was found in \\cite{bo,brustein,covi2} that the slow-roll parameter $\\eta$ is smaller than $-2/3$ for tree-level K\\\"ahler potential and arbitrary superpotential so the slow-roll condition $|\\eta|\\ll1$ is violated. This problem can be cured by adding string corrections to the K\\\"ahler potential. With such corrections a triple gaugino condensation model was constructed which accommodates arbitrary small gravitino mass and a high scale of inflation \\cite{bo}. It was later shown \\cite{bo2} that the number of the gaugino condensates can be reduced to two, as in the standard racetrack model, by admitting (effectively) positive exponents in the non-perturbative terms in the superpotential.\\footnote{The problem of disentangling the gravitino mass from the Hubble scale during inflation was investigated also in \\cite{lvgrav} (in the context of Large Volume Scenario) and  \\cite{openracetrack} (in the context of D-brane inflation).}  A part of the present work is devoted to the investigation of the impact of a matter field on models of volume modulus inflation ending in a SUSY (near) Minkowski minimum. In particular, we study a model with the racetrack superpotential with two positive exponents coupled to a matter field. The main issue we address is whether, and under what circumstances, inflation can be realized without inclusion of any string corrections to the K\\\"ahler potential. It occurs that the matter field always plays a dynamical role in inflation and cannot be decoupled.  We propose a model of inflection point inflation in which inflation is totally dominated by the matter field. The advantage of this model is that it can lead to inflation for any combination of the signs of the exponents in the racetrack superpotential. Matter field inflationary models have been constructed before in \\cite{matterturz,deCarlos}. However, our model is qualitatively different because it can accommodate an arbitrarily light gravitino.  The model that we propose resembles the one constructed many years ago in \\cite{Holman}. However, the model of \\cite{Holman} involves only the matter field so it is not coupled to the moduli sector. It is very important to study such coupling since the moduli fields are always present in supergravity models if the string theory is assumed to be the underlying theory. The impact of the moduli sector on the well established supersymmetric inflationary models such as chaotic \\cite{chaoticsugra} or hybrid inflation \\cite{hybrid} was studied in \\cite{Brax1,Brax2,Brax3,Davis1,Davis2}. In typical situations the presence of the moduli spoils inflation or at least significantly constrains the parameter space. Therefore, it is very encouraging that the inflection point model can be naturally coupled to the volume modulus.         The paper is organized as follows. In section \\ref{condsec} general constraints on the stability of dS vacua and slow-roll inflation originating from the K\\\"ahler potential are recalled. Section \\ref{racesec} contains a brief review of inflationary models based on the racetrack superpotential. Inflation in such models coupled to a matter field sector is investigated in section \\ref{racemattersec}. Two kinds of scenarios are considered. In the first, the Polonyi field $F$-term is used to uplift the racetrack potential. In the second, the uplifting is realized by the quantum corrected O'Raifeartaigh model. In section \\ref{scaleSUSYsec} the impact of the matter field on inflationary models admitting low scale of SUSY breaking is analysed. Section \\ref{conclsec} contains the summary of our results. ",
        "conclusions": "\\label{conclsec}    Our goal has been to construct supersymmetric models of inflation involving the volume modulus and a matter field. The primary purpose of including the matter field was to use its F-term to uplift the energy of the vacuum to a small positive value. In most of the models proposed so far, the uplifting was achieved by adding some terms which break SUSY explicitly (e.g.\\ terms associated with $\\overline{D3}$-branes). The uplifting can be realized in a fully supersymmetric way if the matter field is used. There is also another very important role of the matter field. Contributing to SUSY breaking, it can stabilize the directions perpendicular to the inflaton, some of which are usually tachyonic in supersymmetric models involving only moduli.   We concentrated on models in which the superpotential and the K\\\"ahler potential can be written as sums of terms depending on just one field, the volume modulus $T$ or the matter field $\\Phi$. For the volume modulus contributions we used the standard no-scale K\\\"ahler potential and the superpotential of the racetrack form. We investigated several models with different K\\\"ahler potentials and superpotentials for the matter field looking for such values of the parameters which can support a period of long enough slow roll inflation. Two scenarios have been considered. In one of them inflation takes place when the fields have values close to a saddle point of the potential. In the second scenario they are close to an inflection point.   We found that it is relatively easy to construct models of inflation if one does not insist on having a low scale of SUSY breaking. The superpotential for the matter field may have a simple form linear in $\\Phi$. We considered two K\\\"ahler potentials for the matter field: the canonical one and the canonical one with a correction proportional to $(\\Phi\\bar\\Phi)^2$. In these two cases the matter sector is the same as in the Polonyi model or in the quantum corrected O'Raifeartaigh model, respectively. For both cases we presented examples of saddle point and inflection point inflation. The models of inflection point inflation can be treated as effectively involving only two real fields because the imaginary components of the modulus and the matter field can be set to zero during and after inflation. All four real fields are to be taken into account in the case of the saddle point inflation models.   The main difference between models with different matter sectors is the role of the matter field in the inflationary dynamics. In models with the Polonyi uplifting, the inflaton is a mixture of the modulus and the matter field with a bigger contribution from the latter. For the case of the O'Raifeartaigh uplifting, the matter field is decoupled from the inflationary dynamics and the inflaton can be almost a pure modulus field.    In all models SUSY breaking at the inflationary region is dominated by the matter field but the strength of this dominance changes from case to case. It is much stronger for the inflection point models as compared to the saddle point ones. Within each type of models (inflection point or saddle point) the dominance of matter field SUSY breaking is stronger for the Polonyi uplifting. The reason is that the matter field submanifold has vanishing curvature for the Polonyi model while it is negative for the O'Raifeartaigh one. In the inflection point inflation models, the volume modulus contribution to SUSY breaking is at the level of one (a few) percent in the case of the Polonyi (O'Raifeartaigh) uplifting. In the saddle point models this contribution is typically bigger by one order of magnitude. For the saddle point model with the O'uplifting the contribution to SUSY breaking from the modulus may be even comparable to that of the matter field.   In the inflationary region there are three scalars with positive masses squared and one tachyonic state. In all the model investigated in this paper the inflaton and the lightest non-tachyonic state are dominated by two components of the same complex field (the modulus or the matter field). We found that the value of the inflationary Hubble constant is smaller than the masses of all scalar fields (other than the inflaton), so the isocurvature perturbations are suppressed. This suppression is stronger in models with O'uplifting.   The fine tuning of the superpotential parameters necessary to obtain a long enough period of inflation depends on the model. In the case of the saddle point models this fine tuning varies between a few times $10^{-4}$ (O'uplifting) and a few times $10^{-3}$ (Polonyi uplifting), so it is somewhat weaker than for analogous models with non-SUSY uplifting (typically at the level of $10^{-4}$). From the point of view of fine tuning the Polonyi uplifting seems to be better than the O'uplifting. This feature is more pronounce in the inflection point scenario. O'uplifting requires the constant term in the superpotential to be adjusted with the precision at the level $10^{-9}$ which is worse than $10^{-8}$ typical for inflection point models with non-SUSY uplifting. Using the Polonyi uplifting allows to decrease necessary fine tuning by four or even five orders of magnitude to the level $10^{-4}-10^{-3}$.    Another goal of the paper has been to construct models of inflation ending in  SUSY (near) Minkowski minimum. The main motivation for building such models is that they predict a low scale of SUSY breaking in contrast to the models with uplifting summarized above. This kind of models involving only the volume modulus have been constructed before \\cite{bo,bo2}. However, in those models the string corrections to the K\\\"ahler potential have to be taken into account in order to decrease the curvature of the K\\\"ahler manifold and fulfill the necessary condition for slow-roll inflation. In the present paper we addressed the question whether the contribution to SUSY breaking from the matter field could decrease the curvature of the K\\\"ahler manifold in such a way that inflation could be realized without inclusion of any string corrections to the K\\\"ahler potential. In particular, we investigated a model with the racetrack superpotential with two positive exponents. We assumed quadratic superpotential for the matter field which is the minimal choice allowing for the existence of SUSY Minkowski minimum. We found that inflation in such model can be realized if one parameter from the modulus sector is fine-tuned with the precision at the level of $10^{-4}$. However, the matter field dominates the inflaton field so ''pure'' volume modulus inflation cannot be achieved. Moreover, the parameter space of the matter field sector allowing for inflation is significantly constrained. In addition, the demand for the absence of significant isocurvature perturbations puts some extra constraints on the parameter space of the model because for many parameter sets the mass of the lightest non-tachyonic state during inflation is smaller than the Hubble scale.     We found that it is much easier to construct models of inflation ending in a SUSY (near) Minkowski minimum if the inflaton is totally dominated by the matter field. In such a case one only needs to fine-tune the vev of the matter field. Inflation can be realized for any combination of signs of the exponents in the racetrack superpotential including two negative ones which correspond to the KL model of moduli stabilization. However, in the case of KL model the volume modulus cannot be lighter than the matter field. Otherwise the inflaton would overshoot the SUSY Minkowski minimum and evolve towards the region of infinite volume. It is not necessary to restrict to the canonical K\\\"ahler potential for the matter fields because similar models can be also constructed for non-vanishing modular weight of the matter field. In the models with inflation dominated by the matter field masses of all non-tachyonic states are much bigger than the Hubble scale during inflation so the generation of the isocurvature perturbations is strongly suppressed.   The spectral index in all models discussed in this paper depends on the specific choice of parameters but cannot take arbitrary values. Its allowed values depend on whether inflation takes place in the vicinity of an inflection point or a saddle point. In the former case spectral index is bounded from below $n_s\\gtrsim0.93$, while in the latter case it is bounded from above $n_s\\lesssim0.95$. In any case the parameters can be chosen as to fit the WMAP5 data. The tensor to scalar perturbation ratio and the running of the spectral index predicted by all models under consideration are unmeasurably small.   In summary, our investigation shows that in models admitting a low scale of SUSY breaking it seems to be much more natural to consider the matter field, rather than the volume modulus, as the inflaton. On the other hand, if one does not insist on a low scale of SUSY breaking, the volume modulus is also a very good candidate for the inflaton. In particular, the volume modulus can dominate inflation in the racetrack models uplifted by the matter field $F$-term even though the matter field dominates SUSY breaking. Since the matter field sectors that we used in this work are described by superpotentials and K\\\"ahler potentials of very simple forms, we believe that in more realistic constructions, rigorously derived from the string theory, the $F$-term uplifting of racetrack inflation would be also accomplished. We hope that our analysis will be a useful guideline for the future work."
    },
    "0911/0911.4485_arXiv.txt": {
        "abstract": "{Certain types of globular clusters have the very important property that the predictions for their kinematics in the Newtonian and modified Newtonian dynamics (MOND) contexts are divergent.} {Here, we caution the recent claim that the stellar kinematics data (using 17 stars) of the globular cluster Palomar 14 are inconsistent with MOND.} {We compare the observations to the theoretical predictions using a Kolmogorov-Smirnov test, which is appropriate for small samples. } {We find that, with the currently available data, the MOND prediction for the velocity distribution can only be excluded with a very low confidence level, clearly insufficient to claim that MOND is falsified. } {} ",
        "introduction": "A plethora of observational data on various astronomical scales seem to support the idea that the amount of visible matter in the Universe is several times smaller than the total amount of matter (e.g. Hinshaw et al. 2009). The current paradigm of structure formation and evolution in the Universe is known as the $\\Lambda$ Cold Dark Matter ($\\Lambda$CDM) model. However, as long as the dark matter particle has not been discovered (and its cosmological abundance confirmed), it is worth considering alternative theories to explain the current data.  For instance, MOND (modified Newtonian dynamics) was proposed by Milgrom (1983) as an alternative to galactic dark matter. MOND stipulates that below a certain gravitational acceleration $a_0 \\sim 1.2 \\times 10^{-8}$ cm s$^{-2}$ the actual gravitational acceleration $g$ is stronger than expected in Newtonian gravity ($g_{\\rm N}$). Asymptotically, in MOND it reaches the value $\\sqrt{g_{\\rm N} a_0}$ for $g_{\\rm N} \\ll a_0$. This allows it to naturally explain various galaxy scaling relations (e.g., Faber \\& Jackson~1976, Tully \\& Fisher~1977, McGaugh et al.~2000, McGaugh~2004, Gentile~2008, Donato et al.~2009, Gentile et al.~2009).  The effects of MOND and dark matter are however often rather degenerate and model-dependent since the gravitational potential predicted by MOND can almost always be attributed  by a Newtonist to an {\\it ad hoc} dark matter distribution. Objects for which the predictions of the two theories are unambiguously different are unfortunately rare. We note that galaxy clusters are not good discriminant tests: indeed, in galaxy clusters, the acceleration predicted by MOND is not large enough (e.g., Sanders~1999), but some form of hot dark matter can be added within the MOND context to make the data consistent with the predictions (e.g., Angus et al.~2009). If there were systems where MOND predicted more gravity than observed, they would make a strong case against MOND.  One example of such objects that should be (almost) devoid of cold dark matter in the $\\Lambda$CDM cosmological model, and for which MOND predicts significantly stronger gravity than that attributable to visible matter, are tidal dwarf galaxies (TDGs). However, from the observations of three young TDGs around NGC~5291, Gentile et al. (2007) found that MOND remarkably fits the data with zero free parameters, whereas CDM fails to explain them.  Apart from TDGs, another type of object that has recently been put forward as a discriminant test for CDM and MOND are the globular clusters of our own Galaxy (Baumgardt et al.~2005). Of particular interest are those that are diffuse and distant from the Milky Way, to ensure that the internal acceleration probes the deep MOND regime (which is not always necessarily the case: see NGC~2419 in Baumgardt et al. 2009) and that, simultaneously, the gravitational acceleration due to the Milky Way (external field) is weak enough. In MOND the deviation from a Newtonian behaviour should start appearing around $a_0$, whereas in the ``Newtonian gravity plus CDM'' picture no discrepancy is expected since globular clusters should contain (almost) no cold dark matter.  In a recent paper, Jordi et al.~(2009) analysed the data of 17 stars from the globular cluster Palomar~14, and claimed that (within the assumption of Palomar~14 being on a circular orbit) MOND is inconsistent with the observed velocity dispersion: the MOND prediction is too high, whereas the data are consistent with the Newtonian prediction (with no dark matter). They discuss two separate cases, depending on the inclusion or not of a star (Star 15), which based on its line-of-sight velocity is not a definite member of the cluster. However, given such a small sample size, an appropriate test should be used to discern between the hypotheses. An example of such a test is the Kolmogorov-Smirnov (KS) test (e.g. Soong 2004). An illustrative example of the uncertain value of the velocity dispersion measured with a small number of stars is NGC~2419: the original analysis by Olszewski et al.~(1993) yielded a velocity dispersion for 12 stars of 2.7$\\pm$0.8~km$\\, {\\rm s}^{-1}$, which was re-evaluated by Baumgardt et al.~(2009) with 40 stars to be 4.1$\\pm$0.5~km$\\, {\\rm s}^{-1}$, i.e. about two sigma greater than the original. This gives a perfectly reasonable M/L in MOND, although the velocity dispersion profile, if confirmed, could be a problem for MOND (Sollima \\& Nipoti 2009). This will allow a very interesting test of MOND when data points at larger radii will be obtained in NGC~2419. In any case, this is an excellent example of an underestimation of the true velocity dispersion due to an originally too small sample size.  Therefore, the question we ask in the present research note is the following: do the Palomar~14 data really exclude MOND, given the small number of stars used in the analysis of Jordi et al.~(2009)? We then also investigate the minimum number of stars needed to exclude MOND in a globular cluster similar to Palomar~14. Or, equivalently, we ask how many globular clusters like Palomar~14 would be needed to exclude MOND.   \\begin{table} \\centering \\begin{scriptsize} \\centering \\caption[]{Velocities and cumulative distribution functions (observed and Gaussian, with and without Star 15). The star names are taken from Hilker (2006) when present, otherwise from Harris \\& van den Bergh (1984) and Holland \\& Harris (1992). } \\vspace{0.3cm} \\label{tab-cdf} \\begin{tabular} {l l l l l l} \\hline \\hline Name     &velocity       & obs cdf      & Gauss cdf    & obs cdf          & Gauss cdf      \\\\ & (km s$^{-1}$) & with Star 15 & with Star 15 & w/o Star 15      & w/o Star 15    \\\\ \\hline 15       &  69.99        & 0.059       & 0.034       & -                &  -                 \\\\ 8        &  71.38        & 0.118       & 0.234       & 0.063            &  0.209             \\\\ 3        &  71.75        & 0.177       & 0.332       & 0.125            &  0.302             \\\\ 14       &  71.80        & 0.235       & 0.347       & 0.188            &  0.316             \\\\ 12       &  71.83        & 0.294       & 0.356       & 0.250            &  0.324             \\\\ HH042    &  71.94        & 0.353       & 0.388       & 0.313            &  0.356             \\\\ 16       &  72.14        & 0.412       & 0.450       & 0.375            &  0.416             \\\\ 5        &  72.21        & 0.471       & 0.472       & 0.438            &  0.437             \\\\ 13       &  72.33        & 0.529       & 0.509       & 0.500            &  0.475             \\\\ 17       &  72.39        & 0.588       & 0.528       & 0.563            &  0.494             \\\\ 2        &  72.47        & 0.647       & 0.553       & 0.625            &  0.519             \\\\ 1        &  72.53        & 0.706       & 0.572       & 0.688            &  0.538             \\\\ 7        &  72.64        & 0.765       & 0.606       & 0.750            &  0.572             \\\\ 6        &  72.65        & 0.824       & 0.609       & 0.813            &  0.575             \\\\ HV004    &  73.23        & 0.882       & 0.768       & 0.875            &  0.740             \\\\ 9        &  73.50        & 0.941       & 0.828       & 0.938            &  0.805             \\\\ HV055    &  73.62        & 1.000       & 0.851       & 1.000            &  0.830             \\\\ \\hline  \\end{tabular} \\end{scriptsize} \\end{table} ",
        "conclusions": "Even assuming an isotropic MOND model with no rotation for a globular cluster on a circular orbit with M/L=2, we showed that, based on a KS test, which is the relevant statistical test for small samples, the currently available data are insufficient to discriminate between Newtonian gravity and MOND in Palomar~14, contrary to the claim of Jordi et al.~(2009). While the objects proposed by Baumgardt et al.~(2005) provide one of the best discriminating test between MOND and cold dark matter plus Newtonian dynamics, more observations would be needed to exclude the aforementioned MOND model if the observed velocity dispersion is representative of the true one: about 80 stars in Pal~14, or about 5 similarly problematic globular clusters for this model. In this respect, even though it might still not be conclusive, velocity data of 21 stars in Pal~3 and 24 stars in Pal~4 will be of prime interest for testing fundamental physics in our Galactic backyard.   \\begin{figure} \\centering \\includegraphics[width=8cm]{prob_plot.eps} \\caption{Mean P-value vs. number of stars in the sample for a number of Monte-Carlo realizations (10 per sample size) of distributions of stars following a Gaussian with a standard deviation of 0.85 km s$^{-1}$ (equivalent to the sample with Star 15). } \\label{nstars} \\end{figure}   \\begin{figure} \\centering \\includegraphics[width=8cm]{prob_plot_no15.eps} \\caption{Same as in Fig. \\ref{nstars}, but for a Gaussian with a standard deviation of 0.63 km s$^{-1}$ (equivalent to the sample without Star 15).} \\label{nstars_no15} \\end{figure}"
    },
    "0911/0911.4399_arXiv.txt": {
        "abstract": "The origin of supersonic infrared and radio recombination nebular lines often detected in  young and massive superstar clusters are discussed. We suggest that these arise from a collection of repressurizing shocks (RSs), acting effectively to re-establish pressure balance within the cluster volume and from the cluster wind  which leads to an even broader although much weaker component. The supersonic lines are here shown to occur in clusters that undergo a bimodal hydrodynamic solution (Tenorio-Tagle et al. 2007), that is within clusters that  are above the threshold line in the mechanical luminosity or cluster  mass  vs the size of the cluster (Silich et al. 2004). The plethora of repressurizing shocks is due to  frequent and recurrent thermal instabilities that take place within the matter reinserted by stellar winds and supernovae. We show that the maximum speed of the RSs and of the cluster wind, are both functions of the temperature reached at the stagnation radius. This temperature depends only on the cluster heating efficiency ($\\eta$). Based on our  two dimensional simulations (Wunsch et al. 2008) we calculate  the line profiles that result from several models and confirm our analytical predictions. From a comparison between the predicted and observed values of the half-width zero intensity of the two line components we conclude that the thermalization efficiency in SSC's above the threshold line must be lower than 20 \\%. ",
        "introduction": "Young super star clusters (SSCs) with ages  less than 10$^7$ years, stellar masses $\\sim 10^5$ to 10$^7$ M$_\\odot$, and sizes which span only a few parsecs have been found  in a large variety of galaxies. Many of these have turned out to be strong line emitters in the infrared and radio regimes (see, for example, Turner et al. 2000; Gilbert et al. 2000; Galliano \\& Alloin 2008). Some of them, as in the Antennae galaxy, Henize 2-10, NGC 5253 and IIZw 40 (see Gilbert \\& Graham 2007; Henry et al. 2007; Turner et al. 2003; Beck et al. 2002) are known to emit spatially extended Br$\\gamma$ lines from dense HII regions. Most of these present supersonic line widths (FWHM $\\sim$ 50 to 200 km s$^{-1}$) that may exceed the escape velocity of the clusters. As reported by Beck (2008), in some cases the emission lines can be fitted by a single gaussian component, as it is the case of NGC 5253, while others (as in He 2 - 10) require also of a low intensity, although much broader component (FWHM $\\sim$ 250 - 550 km s$^{-1}$), to fit the observed lines. A central issue pointed out by Turner et al. (2003) when dealing with the central supernebula in NGC 5253 which presents a mean radius smaller than 1 pc, is that if the supersonic line width is caused by the expansion of the  nebula, then its dynamical age is implausibly short. The supersonic infrared (Turner et al. 2003; Henry et al 2007) and radio recombination lines (Rodr\\'\\i guez-Rico et al. 2007) lack a full explanation. The core high intensity, broad component has been interpreted as arising from the photoionized gas leftover from the process of star formation, while the low intensity and even broader component may be the signature of an outflow, a cluster wind, that begins to disrupt the leftover cloud and unveil the central SSC. These ideas are supported by the high extinction and the estimated high  density gas ($n_{HII} \\sim 10^4 - 10^5$ cm$^{-3}$) in the observed HII regions, which may indicate that the star-forming episodes are young. Another possibility in the literature comes from an analogy to the broad lines presented by galactic compact HII regions. Henry et al (2007) and Beck (2008) have suggested that a plethora of massive stars blowing their winds may also contribute to the core line widths. Normal O stars however have winds with velocities $\\sim$ 10$^3$ km s$^{-1}$, but  in their view, a dense environment may reduce the speed to the observed values, as it is the case in compact HII regions in our galaxy (De Pree et al. 2000).  Here we take a completely different approach based on the recent theoretical contributions (Silich et al. 2004, 2007; Tenorio-Tagle et al. 2005, 2007; W\\\"unsch et al. 2007, 2008) that have led us to realize the physical conditions that prevail within the volume occupied by the large collection of massive stars expected within young and compact SSCs. The model assumes, as in the original contribution of Chevalier \\& Clegg (1985), that there is an even distribution of massive stars within the cluster volume, and considers their strong winds as well as their mass and energy released through supernova events. It considers also the direct interaction of all of these supersonic streams and thus the thermalization of the kinetic energy provided by massive stars. Our approach takes into consideration radiative cooling in the thermalized plasma and the thermalization efficiency $\\eta$ which accounts for the immediate  loss of energy resultant from the proximity of the sources and the large metallicities of the reinserted matter. Thus $\\eta$, the thermalization efficiency, measures the remaining energy available to the matter reinserted into the cluster volume. We neglect the radiation pressure as a possible wind driving mechanism because in a typical medium inside the young SSC's the dust grains are destructed by sputtering on a very short time scale. Moreover, the effect of radiation is small compared to the thermal pressure driving (see the Appendix).   Here we show that the detected broad nebular lines have indeed nothing to do with the expansion of the nebula. The supersonic broad emission lines  result only in clusters undergoing the  bimodal hydrodynamic solution (see Tenorio-Tagle et al. 2007),  those  found above the threshold line as defined in the  cluster mass or mechanical energy vs size plane by Silich et al. (2004).  The origin of the  most intense of the  broad lines is due to a plethora of repressurizing shocks (see, for instance, Zeldovich \\& Raizer, 1965), induced within the  dense thermally unstable reinserted gas as this strives to maintain pressure balance with the much hotter gaseous counterpart. The less intense, although much broader gaussian component detected  only in some cases, is here shown to be  caused by the cluster wind,  as it becomes photoionized and less dense upon its own expansion. Section 2 summarizes the hydrodynamics  of the matter reinserted within SSCs. Section 3 provides a handle  on the line-widths expected for clusters undergoing  the bimodal solution and shows that these are a measure of the heating efficiency attained by the matter reinserted within SSCs. Section 3 provides also some line profile examples calculated from our two dimensional hydrodynamic simulations (W\\\"unsch et al. 2008) and section 4 summarizes our conclusions. ",
        "conclusions": ""
    },
    "0911/0911.4925.txt": {
        "abstract": "The meaning of ``linear expansion'' is explained. Particularly accurate {\\em relative\\/} distances are compiled and homogenized a) for 246 SNe\\,Ia and 35 clusters with $v<30,000\\kms$, and b) for relatively nearby galaxies with 176 TRGB and 30 Cepheid distances. The 487 objects define a tight Hubble diagram from $300-30,000\\kms$ implying individual distance errors of $\\la7.5\\%$. Here the velocities are corrected for Virgocentric steaming (locally $220\\kms$) and -- if $v_{220}>3500\\kms$ -- for a $495\\kms$ motion of the Local Supercluster towards the warm CMB pole at $l=275$, $b=12$; local peculiar motions are averaged out by large numbers. A test for linear expansion shows that the corrected velocities increase with distance as predicted by a standard model with $q_{0}=-0.55$ [corresponding to $(\\Omega_{\\rm M},\\Omega_{\\Lambda})=(0.3,0.7)$], but the same holds -- due to the distance limitation of the present sample -- for a range of models with $q_{0}$ between $\\sim\\!0.00$ and $-1.00$. For these models $H_{0}$ does not vary systematically by more than $\\pm2.3\\%$ over the entire range. Local, distance-dependent variations are equally limited to 2.3\\% on average. In particular the proposed Hubble Bubble of Zehavi et~al. and Jha et~al. is rejected at the $4\\sigma$ level. -- Velocity residuals in function of the angle from the CMB pole yield a satisfactory apex velocity of $448\\pm73\\kms$ and a coherence radius of the Local Supercluster of $\\sim3500\\kms$ ($\\sim56\\;$Mpc), beyond which galaxies are seen on average at rest in co-moving coordinates with respect to the CMB. Since no obvious single accelerator of the Local Supercluster exists in the direction of the CMB dipole its motion must be due to the integral gravitational force of all surrounding structures. Most of the gravitational dipole comes probably from within $5000\\kms$. ",
        "introduction": "} \\label{sec:1} % % ****************************************************************** % 1.1. Preliminaries % ****************************************************************** \\subsection{Preliminaries} \\label{sec:1:1} % The basic prediction of all models of ideal universes that are homogeneous and isotropic is that, if the expansion is real, there must necessarily be a linear relation between redshift and distance. There are many proofs, but among the earliest are those by \\citet{Lemaitre:27,Lemaitre:31} and \\citet{Robertson:28} even before the observational announcement by \\citet{Hubble:29} in which he, in the final sentence of his paper, suggested linearity may be only local.  Modern theoretical proofs for global linearity are set out in many of the standard text books at various levels of mathematical sophistication. Popular among the simpler proofs is that a linear relation is the only one that preserves relative shapes of geometrical shapes in an expanding manifold, and which is the same from all vantage points. Deeper proofs, based on properties of the metric of homogeneous, isotropic models, follow \\citet{Robertson:29,Robertson:33} and \\citet{Walker:36}, and lead to the Robertson-Walker line element of the metric.  A unique property of a linear velocity field is that every vantage point in the field appears to be the center of the expansion and has the same ratio of velocity to distance over the entire field. Because of this important property, much effort over the past 80 years has been made by the observers to prove, or disprove, linearity, either locally or globally. The first comprehensive result was that by \\citet{Hubble:Humason:31}, enlarging the observational data available to Hubble in \\citeyear{Hubble:29}.  Many summaries of the theoretical expectation of linearity and observational verification for the standard model exist. From the theoretical side the classic text books include those by \\citet{Heckmann:42}, \\citet{Bondi:60}, \\citet{Robertson:Noonan:68}, \\citet{Harrison:81}, \\citet{Narlikar:83}, \\citet{Peebles:93}, \\citet{Peacock:99}. \\citet{Bowers:Deeming:84}, and \\citet{Carroll:Ostlie:96} are exemplary at the high end of the intermediate level. On the observational side, text book-like chapters by \\citet{Gunn:Oke:75} and \\citet{Sandage:75,Sandage:88,Sandage:95} are useful.  Deviations from a linear velocity-distance relation divide into two categories. (1) The relation is non-linear everywhere, i.e.\\ the deviation is global. (2) The deviation is local, going over into a linear relation at large distances.  The first category is the most fundamental because it denies the standard homogeneous model everywhere. The second contains models where either the local Hubble constant changes up to a certain distance such as increasing or decreasing outward, or where the local irregularities are due to streaming motions due, presumably, to a local inhomogeneous distribution of matter.  The most radical of the first category of global non-linear models can easily be disproved observationally, even as early as in \\citet{Hubble:Humason:31}, by noting that the slope of the local log(redshift) - magnitude relation is close to 0.2 as expected in the case of linear expansion. This conclusion was much strengthened by the tight Hubble diagram from brightest clusters galaxies in \\citet{Humason:etal:56} and then, from 1972 to 1975, by a series of papers on the velocity-distance relation based on new redshifts and apparent magnitudes of galaxy clusters measured at Mount Wilson and Palomar \\citep[][Table~1 for a summary]{Sandage:99}.  Non-linear models in the second category are more difficult to disprove because the deviations from the pure Hubble linear flow are much smaller than in the radical first group and can easily be mocked by systematic distance errors due to bias effects. The literature is large on deviations from a pure Hubble flow, either due to local streaming motions or to a variation of the Hubble constant outward, or to a local bulk motion relative to a globally significant distant kinematic frame. Some are easier to disprove than others.  One such non-linear model assumes that the local global distribution of matter is hierarchical. Following \\citet{Charlier:08,Charlier:22} and using a suggestion by \\citet{Carpenter:38} of a density-size relation for all objects in the universe, \\citet{deVaucouleurs:70,deVaucouleurs:71} postulated a universe made of hierarchies of decreasing mean density with increasing volume size up to some limiting distance. The velocity-distance relation in such a universe was formulated by \\citet{Haggerty:70}, and \\citet{Haggerty:Wertz:71}, following a prediction of \\citet{Wertz:70}, and was found to be nearly quadratic locally, but becoming linear at large distances. The model was shown, however, to be irreconcilable with observations \\citep{STH:72}.  It was early demonstrated \\citep[e.g.][]{% ST:75,Teerikorpi:75a,Teerikorpi:75b,Teerikorpi:84,Teerikorpi:87, Fall:Jones:76,Bottinelli:etal:86,Bottinelli:etal:87,Bottinelli:etal:88, Sandage:88,Sandage:94,FST:94} that the claimed changes of the Hubble constant outward were often caused by uncorrected observational selection bias \\citep[see also][Chapter~10]{Sandage:94,Sandage:95}.% % ****************************************************************** % Footnote 1 % ****************************************************************** \\footnote{It is important to emphasize that bias of this nature is {\\em not\\/} the famous Malmquist bias, but rather is an incompleteness bias that increases with distance. In contrast, the Malmquist bias is distance independent. It only concerns the error made in assigning a number to an absolute magnitude calibration of a distance indicator when using a sample that is magnitude, rather than, distance limited. On the other hand, an incompleteness bias causes an error that increases with distance, which, if uncorrected, makes the Hubble constant appear to increase outward with distance. The error of calling the incompleteness bias as Malmquist is wide spread in the literature.} % ****************************************************************** It must also be stated that in the early papers some of the claimed distortions of the velocity field were due to the omission of velocity corrections for Virgocentric infall and -- for the more distant galaxies -- for CMB dipole motion.  Hence, despite many papers and conferences on proposed streaming motions and searches for variations of the Hubble constant with distance, the problem is still open. What has been needed are distance indicators that have very small intrinsic dispersion so as to eliminate, or greatly reduce, the effect of distance-dependent incompleteness selection bias. At small distances the increasing number of Cepheid and TRGB distances becomes most helpful here, whereas for large distances Type Ia supernovae at maximum light have emerged as the best standard candles, just as \\citet{Zwicky:62} had proposed.    % ****************************************************************** % 1.2. The complications: the invisible world map must be %      inferred from the visible world picture % ****************************************************************** \\subsection{The Complications: The Invisible World Map Must be Inferred from the Visible World Picture} \\label{sec:1:2} % The concept of linearity of a velocity-distance velocity field, so simple to describe, is most complicated to define properly in an expanding universe. \\citet{Milne:35} proposed language that clarifies the problem. \\citet{Robertson:55} also uses the same language.  The problems are these.  (1) Because we do not observe the universe at large redshift at the same cosmic time as at low redshift, we must distinguish between what Milne called the world map and the world picture. The world map is the state of the universe at a given cosmic time. The world picture is what appears to us. We cut the series of world maps at different cosmic times as we look to different redshifts.  By a linear ``velocity''-``distance'' relation is meant that at a particular cosmic time there is a linear relation in the world map between the coordinate distance as measured by the metric, and redshift. But the metric distance, $R(t)r$, where $R(t)$ is the time dependent expansion scale factor and $r$ is the invariant (constant for all time) comoving radial metric ``distance''. It is not the same as a distance measured by an astronomer using observational data. There are many kinds of such ``astronomical'' distances, depending on the method by which they are measured. There is the distance when light left a galaxy, the distance when light is received, the luminosity distance by apparent magnitude, the distance by angular size, the round-trip distance by radar signals. All differ at large redshifts, and are the same only in the zero redshift limit. \\citet*{McVittie:74} discussion is useful here. Which distance do we use in formulating a velocity-``distance'' relation and proving that it is ``linear''?  There is also the difficulty with ``velocity''. What we measure is redshift, not velocity. The two are not the same except, again, in the zero redshift limit \\citep{Harrison:93}.  The problem is solved by not using the complicated concepts of velocity and distance but by transforming the world map into the observer's world picture {\\em using only observables}. Before 1958 this mapping was done by making a Taylor series expansion of the $R(t)$ expansion factor about the present cosmic time so as to sample the world map at the earlier cosmic times \\citep{Robertson:55}. This could be called a tangent mapping, and is not very useful for redshifts larger than those at moderately local distances. What was needed was a general transformation mapping that is valid for all redshifts.  \\citet{Mattig:58} solved the problem with his famous equation that relates the metric distance of a galaxy, $R_{0}r$, at the time of light reception with the redshift, no matter how large, in models with zero cosmological constant. Metric distance is further replaced by apparent magnitude by the \\citet{Robertson:38} equation that connects the two. In addition to the apparent magnitude and the redshift the equation contains the \\citet{Robertson:55} deceleration parameter, $q_{0}=-R_{0}\\ddot{R}/\\dot{R}^{2}$, which, in the simplest models with no cosmological constant, is determined by the mean matter density of the model. In that case the values of $H_{0}$ and $q_{0}$ observed at the present epoch are used to parametrize the past and the future of the world map.  The Mattig equation, or more complicated versions for models that incorporate the cosmological constant \\citep[e.g.][]{Mattig:59,Carroll:etal:92} which we use in the next section, provides the  connection between the world map and the world picture. It has properly been said that the Mattig equation is ``one of the single most useful equations in cosmology as far as observers are concerned'' \\citep[][p.~89]{Peacock:99}. Mattig's solution began the modern era of practical observational cosmology and has been used for many auxiliary and related problems concerning such observational data as galaxy counts, angular diameters, and others throughout the past 50 years.  Because the world map defined in the Robertson-Walker metric has a well defined linear relation between coordinate (metric) distance and redshift, and because a Mattig-like equation transforms the map in to the observed picture, the test for linearity becomes one of comparing the observed redshifts at a given apparent magnitude with the predicted redshifts from a model. The linearity test becomes, then, a search for residuals in the subtraction of the redshifts predicted {\\em from the model\\/} from the observed redshifts at a given photometric distance obtained from the observations as $(m-M)$. Streaming motions and/or variations of the Hubble constant outward are searched for as correlations (or not) of the residuals with direction and photometric distance.  The test for linearity is, then, model-dependent: we must compare the observed redshift-magnitude relation with the prediction for some adopted world model as it is transformed into the world picture by a Mattig-like equation of $(m-M)=f(z,q_{0},\\Lambda)$. We can only test linearity relative to the adopted model. We test for the sensitivity to the adopted model in \\S~\\ref{sec:3}.    % ****************************************************************** % 1.3. This paper % ****************************************************************** \\subsection{This Paper} \\label{sec:1:3} % The purpose of this paper is: (1) To test the linearity of the local expansion field within $z<0.1$. Beyond this limit the linearity of the expansion of space is well documented out to $z=0.4$ by the SNe\\,Ia of \\citet{Hicken:etal:09}, which have ample overlap with the present data, and which can be tightly fit to the SNe\\,Ia compiled by \\citet{Kessler:etal:09} extending to $z>1$. The high-$z$ data are essential to optimize the determination of the world model for which linear expansion is valid. But at smaller distances the character of the expansion field is still poorly known because so far the number of objects with sufficiently accurate distances has been small. Yet the test is important to understand the effect of the observed clumping of visible matter on the local dynamics. It is also important to find the minimum distance at which the cosmic value of $H_{0}$ can be found independent of peculiar velocities. Strong local deviations from linearity have in fact frequently been claimed up to the recent past by \\citet{Zehavi:etal:98} and \\citet{Riess:etal:09}. (2) To determine the size of the comoving volume that partakes of our observed velocity relative to the dipole of the CMB.  The plan of the paper is this. In \\S~\\ref{sec:2} the data are specified. In \\S~\\ref{sec:2:1} accurate {\\em relative\\/} distances of 246 SNe\\,Ia and 35 clusters within the adopted distance range are compiled from eight different sources. The distances have demonstrably rms errors of $\\le0.18\\mag$; only relative distances are needed in the present context. The distances on an arbitrary zero point, reduced to the barycenter of the Local Group, are listed in Table~\\ref{tab:ml} together with the appropriate velocities $v_{\\rm hel}$, $v_{220}$ (corrected for Virgocentric infall), and $v_{\\rm CMB}$ (corrected for motion with respect of the CMB dipole $A_{\\rm corr}$ if $v_{220}>3500\\kms$). In \\S~\\ref{sec:2:2} 30 Cepheid and 176 TRGB distances which are equally accurate as those of SNe\\,Ia are added in order to extend the sample down to $300\\kms$. They are reduced to the barycenter of the Local Group and corrected for the Virgocentric flow as in \\S~\\ref{sec:2:1}. We justify in \\S~\\ref{sec:2:3} why other distance indicators are not considered here.  In \\S~\\ref{sec:3} the problem of linear expansion is set out. The actual test in \\S~\\ref{sec:3:1} uses a Hubble diagram of {\\em all\\/} objects in the sample and analyzes the residuals from a Hubble line for a standard $\\Lambda$CDM model with $\\Omega_{\\rm M}=0.3$ and $\\Omega_{\\Lambda}=0.7$. (Note that this model implies a deceleration parameter of $q_{0}=-0.55$ because it holds that $q_{0}=0.5(\\Omega_{\\rm M}-2\\Omega_{\\Lambda})$ for all Friedmann models [e.g.\\ \\citealt{Sahni:Starobinsky:00}]). \\S~\\ref{sec:3:2} explores the dependence of linearity on the adopted model in terms of $q_{0}$.  An analysis of the velocity residuals of objects with $v_{220}<7000\\kms$ (the limit is set to avoid an overwhelming effect of distance errors) in function of direction is in \\S~\\ref{sec:4}, showing the size of the Local Supercluster and its motion toward $A_{\\rm corr}$, which is the direction of the warm pole of the CMB after correction for our Virgocentric velocity vector.  Results and conclusions are in \\S~\\ref{sec:5}.    % ****************************************************************** % 2. The data % ****************************************************************** ",
        "conclusions": "} \\label{sec:5} % Exceptionally accurate relative distances are compiled from eight sources in \\S~\\ref{sec:2:1}. They are based on SN\\,Ia magnitudes and on mean 21cm line widths and the Fundamental Plane in case of clusters. They are normalized to an arbitrary value of $H_{0}$ of 62.3 without loss of generality. Accurate TRGB and Cepheid distances were added in \\S~\\ref{sec:2:2}. The final Hubble diagram with 480 objects extends from 300 to $30,000\\kms$ and has a scatter of only $0.15\\mag$ (8\\% in linear distance) beyond $3000\\kms$; the increase of the scatter at shorter distances, due to peculiar motions, is compensated by large-number statistics.  The expansion of space is as linear as can be measured after allowance is made for a Virgocentric flow model and for the CMB motion of the Local Supercluster. Over the range of 300 to $30,000\\kms$ the value of $H_{0}$ does not change systematically by more than $\\pm2.3\\%$. At poorly populated distances $H_{0}$ may locally vary by 7\\%. But the so-called Hubble Bubble suggesting a sudden decrease of $H_{0}$ by 6.5\\% -- be it a local dip or a persistent feature beyond -- at the particularly well occupied distance around $7200\\kms$ is excluded by $4\\sigma$ or more.  The question at which distance the cosmic value of $H_{0}$ can be found is answered by Figure~\\ref{fig:06}: beyond $(m-M)=28.0$ ($\\sim\\!4\\;$Mpc) and out to at least $20,000\\kms$ any systematic deviation from linear expansion is limited to a few percent. At $20,000\\kms$ the expansion field is well tied to the equally linear large-scale expansion field \\citep[e.g.][]{Kessler:etal:09}. The cosmic value of $H_{0}$ can therefore be found quite locally if sufficient calibrators are available to compensate the locally important peculiar velocities. The available 78 TRGB distances outside $(m-M)=28.0$ yield a well determined value of $H_{0}=62.9\\pm1.6$ (statistical error) in good agreement with local Cepheids and SNe\\,Ia \\citep{TSR:08a}. Allowing for a $0.10\\mag$ error of the TRGB zero point, based on RR Lyr stars and other distance determinations, and for a generous variation of $H_{0}$ with distance of $0.10\\mag$ gives a compounded error of $H_{0}$ of $\\pm4.7$ (8\\%), including the statistical error. The discrepancy with values of $H_{0}>70$ \\citep[e.g.][]{Riess:etal:09} is not a subject of the present paper.  The Local Supercluster emerges as a comoving entity of radius $3500\\kms$, corresponding to a diameter of $\\sim\\!110\\;$Mpc, with the Virgo cluster at its center and several additional clusters (UMa, Fornax, Eridanus) and a large number of groups which must induce a network of small peculiar motions not considered in this paper \\citep[see][for a map]{Klypin:etal:03}. Inside the Supercluster the expansion is decelerated about inversely to the distance from the Virgo cluster. The galaxies of the Supercluster are concentrated toward the supergalactic plane \\citep{deVaucouleurs:56} which extends to at least $4500\\kms$ \\citep{Lahav:etal:00}. It is to be noted that the CMB apex $A_{\\rm corr}$ lies at supergalactic latitude $-39^{\\circ}$, which shows that objects near to the plane (like the Great Attractor and the Shapley Concentration) can contribute only a fraction of the acceleration of the Local Supercluster.  The Supercluster's reflex motion with respect to galaxies beyond $3500\\kms$ amounts to $448\\pm73\\kms$ in good agreement with the velocity of $495\\pm25\\kms$ toward an apex in the zone of avoidance ($l=275\\pm2$, $b=12\\pm4$) as inferred from the CMB dipole after correction for the local Virgocentric infall vector. The infall vector at the position of the LG amounts to $220\\pm30\\kms$ and diminishes for larger Virgocentric velocities according to the mass profile of the Supercluster. The question as to the exact shape of the Local Supercluster is very complex as evidenced for instance by the near component of the Centaurus cluster \\citep{Lucey:etal:91,Stein:etal:97}, and the Pegasus and Hydra clusters all three of which lie near to its outer boundary.  The Local Supercluster is the largest volume to our knowledge for which an unambiguous peculiar motion has been found. Claims for coherent streaming motions over still larger scales are either disproven or remain very doubtful in view of possible systematic errors of the applied distance indicators with large intrinsic scatter. In any case any large-scale streaming motion of which the Local Supercluster could partake is limited by the agreement between the absolute CMB dipole motion of $495\\pm25\\kms$ and the observed bulk motion of $448\\pm73\\kms$ with respect to galaxies beyond $3500\\kms$. The question as to the relation between volume size and corresponding peculiar velocity is decisive for the theory of structure evolution which predicts decreasing peculiar velocities with increasing structure size. $\\Lambda$CDM models with $\\Omega_{\\rm M}=0.3$, $\\Omega_{\\Lambda}=0.7$ predict quite generally that a bulk motion of $\\la500\\kms$ lies in the upper, yet permissible range for a volume of radius $3500\\kms$ (\\citealt{Dekel:00}, Fig.~1; \\citealt{Colless:etal:01}, Fig.~14).  The bulk motions of galaxy concentrations outside the Local Supercluster are still difficult to determine in view of the remaining distance errors. The one-dimensional bulk velocities in the inertial frame of the CMB of the 28 (super-) clusters and groups with $3500 < v_{\\rm CMB} < 10,000\\kms$ contained in Table~\\ref{tab:ml} are compiled in Table~\\ref{tab:cluster} (in order of RA). Columns (1)$-$(3) give the name of the cluster, its adopted distance modulus $\\mu^{0}$ from SNe\\,Ia, D$_{n}-\\sigma$ and/or fundamental plane distances, and the number of distance determinations, respectively. The observed mean cluster velocity, expressed as $v_{\\rm CMB}$ is in column~(4). The expected velocity $v_{\\rm model}$ in column~(5) is calculated for $H_{0}=62.3$, $\\Omega_{\\rm M}=0.3$, and $\\Omega_{\\Lambda}=0.7$ as throughout in this paper. The bulk velocity $v_{\\rm CMB}-v_{\\rm model}$ and its error follow in column~(6). The $1\\sigma$ error is estimated by assuming here a distance error of 7.5\\% for one distance determination and of 5\\% for two or three distance determinations. The average (absolute) bulk velocity of the 28 aggregates is $353\\kms$. However the signal of $353\\kms$ is smaller than the average error of $387\\kms$, questioning whether any of the bulk velocities are significant. If one considers instead only the 11 clusters and groups with $3500<v_{\\rm CMB}\\le5000\\kms$, for which the distance errors are less effective, one obtains an average bulk motion of $175\\pm250\\kms$, or $303\\pm432\\kms$ in three dimensions. In the sample of 11 only the Hydra cluster has a bulk radial velocity with a significance of $1.9\\sigma$, i.e.\\ $-349\\pm199\\kms$ toward the Local Supercluster. Peculiar velocities of large aggregates of some hundred $\\kms$ appear to be compatible with the data. The one-dimensional bulk velocity of the Local Supercluster of $480:\\sqrt{3}=280\\kms$, as seen from a random external observer, is therefore not exeptional.  % ******************************************************** %  Table 6: One-dimensional bulk velocities of (super-)clusters % ********************************************************  It is unphysical to seek for a single attractor accelerating the Local Supercluster as has been proposed in the case of the Great Attractor (see \\S~\\ref{sec:2:3:1}) or the Shapley Concentration \\citep{Mathewson:etal:92,Saunders:etal:00,Kocevski:Ebeling:06, Basilakos:Plionis:06}, because they are 53 and 34 degrees away from the apex direction $A_{\\rm corr}$. It is clear that instead the acceleration must be caused by the integrated force exerted by all surrounding mass concentrations and voids. A lower distance limit of the dominant accelerators is set by the radius of the Local Supercluster of $\\nobreak{\\sim\\!3500\\kms}$. A direct determination of the convergence depth is afforded by integrating the light of all-sky samples of galaxies up to the point where the light dipole  --  which is the same as the dipole of the gravitational pull if the mass-to-light ratio is assumed to be constant -- agrees with the CMB dipole \\citep[see][]{Yahil:etal:80a,Davis:Huchra:82,Lahav:87, Lynden-Bell:etal:89,Strauss:etal:92,Maller:etal:03}. Particularly suited for this purpose is the relatively absorption-free $2\\;\\mu$m all-sky 2MASS Redshift Survey from which \\citet{Erdogdu:etal:06a} have derived  --  by excluding a few nearby galaxies and integrating out to $3000\\kms$  -- an apex at $l=269\\pm9$, $b=37\\pm9$. This agrees well with the CMB apex at $l=276\\pm2$, $b=30\\pm2$ as seen from the Local Group (see \\S~\\ref{sec:4}).% % ****************************************************************** % Footnote 3 % ****************************************************************** \\footnote{In this case the dipole of the integrated light must be compared with the CMB apex as seen from the Local Group and not with the apex $A_{\\rm corr}$, i.e.\\ corrected for Virgocentric infall, because the integration starts at small velocities and includes the Virgo cluster.} % ****************************************************************** The agreement is only marginally improved if the integration is carried out to $5000\\kms$. The conclusion is that most of the acceleration of the Local Supercluster comes from a distance of $\\sim\\!4000\\pm1000\\kms$. A map of the dominant overdensities (in $2\\;\\mu$m) and voids within a shell centered at 4000 kms-1 is given by \\citet[][their Fig.~4]{Erdogdu:etal:06b}. The structures seen in this shell are probably the main players accelerating the Local Supercluster.    % *********************************************** %  Acknowledgments % ***********************************************"
    },
    "0911/0911.0435_arXiv.txt": {
        "abstract": "We report the discovery of the second accreting millisecond X-ray pulsar (AMXP) in the globular cluster NGC~6440. Pulsations with a frequency of 205.89~Hz were detected with the Rossi X-Ray Timing Explorer on August 30th, October 1st and October 28th, 2009, during the decays of $\\lesssim4$ day outbursts of a newly X-ray transient source in NGC~6440. By studying the Doppler shift of the pulsation frequency, we find that the system is an ultra-compact binary with an orbital period of 57.3 minutes and a projected semi-major axis of 6.22 light-milliseconds. Based on the mass function, we estimate a lower limit to the mass of the companion to be 0.0067 M$_{\\odot}$ (assuming a 1.4 M$_{\\odot}$ neutron star). This new pulsar shows the shortest outburst recurrence time among AMXPs ($\\sim1$ month). If this behavior does not cease, this AMXP has the potential to be one of the best sources in which to study how the binary system and the neutron star spin evolve. Furthermore, the characteristics of this new source indicate that there might exist a population of AMXPs undergoing weak outbursts which are undetected by current all-sky X-ray monitors. NGC~6440 is the only globular cluster to host two known AMXPs, while no AMXPs have been detected in any other globular cluster. ",
        "introduction": "\\label{sec:intro}  Accreting millisecond X-ray pulsars (AMXPs) are rapidly spinning neutron stars accreting matter from a low-mass stellar companion. These systems are thought to be the progenitors of the millisecond radio pulsars \\citep{Alpar82,Backer82}. Therefore, their study allows us to better understand how neutron stars in low-mass X-ray binaries (LMXB) evolve into radio millisecond pulsars. Furthermore, AMXPs are perfect laboratories to study the interaction between the neutron star magnetosphere and the accretion disk \\citep[see, e.g.,][and references therein]{Poutanen06}.    A total of ten AMXPs among the $\\sim100$ neutron star LMXBs were known by August 2009. Seven of these AMXPs showed outbursts lasting a few days to months, and recurrence times of at least ~2 years. In all seven cases, the pulsations were detected persistently during the whole outburst. The remaining three AMXPs showed pulsations only intermittently \\citep{Galloway07a,Gavriil07,Casella08,Altamirano08b,Patruno09}, bridging the gap between the small number of AMXPs and the large group of non-pulsating LMXBs \\citep{Altamirano08b}. The reason why only a small amount of neutron star LMXBs show pulsations is still under debate. One of the proposed scenarios for the non-pulsating systems is that the neutron-star magnetic field could be temporarily buried by the accreted matter \\citep{Cumming01}.  If the time-averaged mass accretion rate of the accretors is relatively high, the accreted matter can bury the field throughout the life of the X-ray binary so the neutron star does not pulsate. However, if the average accretion rate is low enough that the magnetic field can diffuse through the accreted matter faster than it is buried, then these systems will probably exhibit pulsations. Another of the proposed scenarios that explains why the majority of neutron stars in LMXBs do not pulsate is that of \\citet{Lamb09a}; if neutron stars undergo long periods of accretion (at high rates), then their magnetic poles are naturally forced to align to their spin axes as they spin up.  When the accretion rate decreases, the magnetic poles can move away from the rotation axis (due to magnetic dipole and other braking torques which cause neutron stars to spin down) and oscillations powered by accretion should become visible. Both \\citet{Cumming01} and \\citet{Lamb09a}  present clear predictions: we should find pulsations in many (if not all) of the systems which show low time-averaged mass accretion rates.  Very recently, a Chandra X-ray observation serendipitously discovered a new X-ray transient in the globular cluster NGC~6440 \\citep[identified as NGC~6440 X-2]{Heinke09a, Heinke09b, Heinke09c}. RXTE and Swift follow-up observations not only showed that this source has a very low time--average accretion rate \\citep[$\\lesssim2 \\cdot 10^{-12}$ M$_{\\odot}$ year$^{-1}$, see ][]{Heinke09d} but, as predicted by the models discussed above, also millisecond X-ray pulsations \\citep{Altamirano09}. One of the most interesting properties of this new AMXP is its very short recurrence time \\citep[$\\sim1$ month, see also][]{Heinke09d}. If its outbursts continue, it will be possible to study the evolution of spin and orbital parameters on timescales never before probed.  In this Letter we report the discovery of X-ray pulsations in NGC 6440 X-2. In a companion work \\citep{Heinke09d}, we report on a detailed description of the evolution of the source's outbursts based on our multiwavelength campaign on this source. Our results demonstrate that there may exist a population of AMXPs undergoing low-luminosity outbursts which are undetected by current all-sky X-ray monitors. ",
        "conclusions": "\\label{sec:discussion}  We have discovered the second accreting millisecond X-ray pulsar in the globular cluster NGC~6440. Pulsations were detected in three RXTE pointed observations performed during the decay phase of three outbursts, each of which lasted less than $\\sim3$ days \\citep{Heinke09d}. The frequency drifts we observe are consistent with an orbital period of 57.3 minutes and a semi-major axis ($a_x$ sin$i$) of 6.22 light-msec, which shows that this new AMXP is in an ultra-compact binary system.  The characteristics of this new pulsar are similar to those of other AMXPs, particularly SWIFT J1756.9--2508 \\citep{Krimm07,Patruno10}, which showed pulsations at a frequency of 182~Hz, an orbital period of 54.7 minutes, a projected semi-major axis of 5.94 lt-ms and a minimum companion mass of 0.0067 M$_{\\odot}$. As pointed out by \\citealt{Krimm07}, such a low mass for a companion in this type of system is probably inconsistent with brown dwarf models while white dwarf models suggest that the companion is probably a He-dominated donor.   The reason why only twelve\\footnote{See also the recent discovery of a 245~Hz AMXP by \\citet{Markwardt09}.} NS LMXBs exhibit coherent pulsations out of about 100 systems is still unclear.  There is a strong tendency for the detected AMXPs to have rather low time-averaged accretion rates \\citep[over many outbursts, see, e.g., ][]{Chakrabarty05,Galloway06a,Wijnands08}. The time-averaged accretion rate of NGC~6440 X-2 is at most similar to that of the AMXP SAX~J1808.4--3658, but could easily be much lower \\citep[see discussion by ][]{Heinke09d}. This is consistent with the general tendency and furthermore, with the theoretical predictions \\citep{Cumming01, Lamb09a} that a significant fraction (if not all) of the NS LMXB systems with low time-averaged accretion rates harbor accreting millisecond X-ray pulsars \\citep{Wijnands08}.  The outburst durations ($\\sim3$ days) of the newly discovered AMXP are short compared with those of the other AMXPs, although not uniquely so. For example, short-duration outbursts of the AMXP XTE J1751--305 have been detected in 2005, 2007 and 2009 \\citep{Grebenev05,Swank05, Markwardt07b, Markwardt07c,Linares07a, Markwardt09a}. Our discovery of NGC~6440 X-2 indicates that there might exist a population of AMXPs undergoing weak outbursts. Unfortunately, short low-luminosity outbursts cannot be detected with the current all-sky X-ray monitors (e.g., ASM aboard RXTE). More sensitive instruments like the ``Monitor of All-sky X-ray Image'' \\citep[MAXI, which should detect X-ray -- 2--30 keV -- sources of about 20 mCrab from 90 minutes observations, and sources of about 4.5 mCrab after a day observation, see][]{Matsuoka09} will probably allow us to discover many more potential low-luminosity AMXPs.  As discussed in \\citet{Heinke09d}, NGC~6440 X-2 has been in outburst at least five times in the last 250 days. If true, this is the AMXP with the shortest recurrence time. If NGC~6440 X-2 continues undergoing this type of short outbursts and RXTE is able to observe them, then NGC~6440 X-2 has the potential to be the best millisecond X-ray pulsar in which to study the evolution of the binary as well as that of the neutron star spin (by means of coherent timing analysis).   The globular cluster NGC~6440 is known for having at least 24 faint X-ray sources to a limiting luminosity of $\\sim2 \\cdot 10^{31}$ ergs s$^{-1}$ \\citep[0.5--2.5 keV, see ][]{Pooley02}. NGC~6440 is also known for hosting at least 6 millisecond radio pulsars, of which 3 are in binary systems \\citep{Freire08}. None of the known radio pulsars match the position (nor the characteristics) of NGC~6440 X-2. Although millisecond radio pulsars in binary systems have been found in other globular clusters \\citep[see, e.g.,][and references therein]{Ransom05, Freire08}, it is an intriguing question why NGC~6440 is the only globular cluster known today to host (two) AMXPs (SAX~J1748.9--2021 -- \\citealt{Gavriil07} and \\citealt{Altamirano08b} -- and NGC~6440 X-2 -- this paper), while no AMXPs have been detected in any other globular cluster. This might be only a selection effect, as current all-sky monitors cannot detect short low-luminosity outbursts. Monitoring X-ray observations of different globular clusters would be needed to further investigate this."
    },
    "0911/0911.1120.txt": {
        "abstract": "{We propose an alternate, calculable mechanism of dark matter genesis, ``thermal freeze-in,\" involving a Feebly Interacting Massive Particle (FIMP) interacting so feebly with the thermal bath that it never attains thermal equilibrium.   As with the conventional ``thermal freeze-out\" production mechanism, the relic abundance reflects a combination of initial thermal distributions together with particle masses and couplings that can be measured in the laboratory or astrophysically. The freeze-in yield is IR dominated by low temperatures near the FIMP mass and is independent of unknown UV physics, such as the reheat temperature after inflation. Moduli and modulinos of string theory compactifications that receive mass from weak-scale supersymmetry breaking provide implementations of the freeze-in mechanism, as do models that employ Dirac neutrino masses or GUT-scale-suppressed interactions. Experimental signals of freeze-in and FIMPs can be spectacular, including the production of new metastable coloured or charged particles at the LHC as well as the alteration of big bang nucleosynthesis. } \\end{center} \\end{titlepage}  \\tableofcontents   %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% ",
        "introduction": "%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% \\begin{figure} \\vspace{-2cm} \\centerline{\\includegraphics[width=14cm]{FI_vs_FO_figure2.pdf}} \\vspace{-1.5cm} \\caption{Log-Log plot of the evolution of the relic yields for conventional freeze-out (solid coloured) and freeze-in via a Yukawa interaction (dashed coloured) as a function of $x=m/T$.  The black solid line indicates the yield assuming equilibrium is maintained, while the arrows indicate the effect of increasing coupling strength for the two processes.  Note that the freeze-in yield is dominated by the epoch $x\\sim 2-5$, in contrast to freeze-out which only departs from equilibrium for $x\\sim 20-30$.\\label{fi_vs_fo}} \\end{figure} %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%  Many theories of Dark Matter (DM) genesis are based upon the mechanism of ``thermal freeze out\" \\cite{freezeout}.  In this process DM particles have a large initial thermal density which, as the temperature of the hot plasma of the early universe drops below the mass, dilutes away until the annihilation to lighter species becomes slower than the expansion rate of the universe and the comoving number density of DM particles becomes fixed. The larger this annihilation cross section the more the DM particles are able to annihilate and hence a thermal distribution with an exponential Boltzmann factor is maintained to a lower temperature, giving a lower final yield.  An attractive feature of the freeze-out mechanism is that for renormalisable couplings the yield is dominated by low temperatures with freeze-out typically occurring at a temperature a factor of $20-25$ below the DM mass, and so is independent of the uncertain early thermal history of the universe and possible new interactions at high scales.  Are there other possibilities, apart from freeze-out, where a relic abundance reflects a combination of initial thermal distributions together with particle masses and couplings that can be measured in the laboratory or astrophysically? In particular we seek cases, like the most attractive form of freeze-out, where production is IR dominated by low temperatures of order the DM mass, $m$, and is independent of unknown UV quantities, such as the reheat temperature after inflation.  In this paper we show that there is an alternate mechanism, ``freeze-in\", with these features.   Suppose that at temperature $T$ there is a set of bath particles that are in thermal equilibrium and some other long-lived particle $X$, having interactions with the bath that are so feeble that $X$ is thermally decoupled from the plasma.  We make the crucial assumption that the earlier history of the universe makes the abundance of $X$ negligibly small, whether by inflation or some other mechanism.  Although feeble, the interactions with the bath do lead to some $X$ production and, for renormalisable interactions, the dominant production of $X$ occurs as $T$ drops below the mass of $X$ (providing $X$ is heavier than the bath particles with which it interacts). The abundance of $X$ ``freezes-in\" with a yield that increases with the interaction strength of $X$ with the bath.  Freeze-in can be viewed as the opposite process to freeze-out. As the temperature drops below the mass of the relevant particle, the DM is either heading away from (freeze-out) or towards (freeze-in) thermal equilibrium.  Freeze-out begins with a full $T^3$ thermal number density of DM particles, and reducing the interaction strength helps to maintain this large abundance.  Freeze-in has a negligible initial DM abundance, but increasing the interaction strength increases the production from the thermal bath. These trends are illustrated in Figure~\\ref{fi_vs_fo}, which shows the evolution with temperature of the dark matter abundance according to, respectively, conventional freeze-out, and the freeze-in mechanism we study here.  In Section \\ref{idea}, as well as outlining the basic freeze-in mechanism and comparing its features with those of freeze-out, we also introduce the idea of a FIMP---a ``Feebly-Interacting-Massive-Particle\" (or alternatively ``Frozen-In-Massive-Particle\")---as distinct from a WIMP whose relic abundance is set by conventional freeze-out.  Two cases are considered: either the FIMP itself is the dark matter which is frozen-in, or the dominant contribution to the DM density arises from frozen-in FIMPs which then decay to a lighter DM particle.  For enhanced pedagogy the detailed calculation of the DM abundance in these cases is postponed until Section \\ref{calculation}.  We turn in Section \\ref{candidates} to the question of motivated DM candidates that have a relic abundance determined by the IR-dominated freeze-in mechanism, and show that the moduli and modulinos of compactified string theories with weak-scale supersymmetry breaking provide implementations of the freeze-in mechanism, as does any extra-dimensional extension of the Standard Model (SM) with some moduli stabilised at the weak scale.   Models following from Dirac neutrino masses, or involving kinetic mixing with a hidden sector also naturally accommodate FIMPs and the freeze-in mechanism. A striking difference with freeze-out production of DM, is that for freeze-in the final relic density is, in the simplest cases, automatically independent of the FIMP mass, allowing superheavy as well as weak-scale mass DM candidates, and we touch on a high-scale extra-dimensional realisation of such a scenario.  Theories utilising GUT-scale-suppressed non-renormalisable interactions involving SM or MSSM Higgs fields which become renormalisable when the Higgs gets its vacuum expectation value also naturally give rise to FIMPs.  However, in this case, there is also a UV contribution to the freeze-in yield which can dominate the IR contribution if the reheat temperature is sufficiently large, and we therefore postpone the discussion of such implementations until Section \\ref{comments}  where the effect of higher-dimension-operators on freeze-in yields is considered.  Experimental signals of freeze-in and FIMPs are briefly summarised in Section \\ref{signatures}, with a companion paper \\cite{companion} more extensively discussing the rich phenomenology that accompanies freeze-in to follow. The signals depend on the nature of the Lightest Observable Sector Particle (LOSP) that carries the conserved quantum number that stabilises DM.  (In theories with weak scale supersymmetry, the LOSP is the lightest superpartner in thermal equilibrium at the weak scale.)  There are two general possibilities:  The DM particle is the FIMP, and there are spectacular LHC signatures arising from the production of coloured or charged LOSPs, which stop in the detector and decay with a long lifetime. In addition, LOSPs decaying to FIMPs in the early Universe may lead to interesting modifications of Big-Bang Nucleosynthesis (BBN). Alternatively, the DM particle is the LOSP but its interactions are too strong to have a sufficiently large freeze-out relic density; instead its relic abundance results from freeze-in of FIMPs, with the FIMPs later decaying to the DM.  The second scenario can give signals via alteration of big-bang nucleosynthesis element abundances, and via increased DM pair annihilation relevant for indirect detection of DM.  Thermal relic abundances conventionally arise by decoupling of a species that was previously in thermal equilibrium, whether with or without a chemical potential.  Freeze-in provides the only possible alternative thermal production mechanism that is dominated by IR processes.  In Section \\ref{phasediagrams} we sketch ``abundance phase diagrams\" that show regions of mass and coupling where each mechanism dominates the production of the relic abundance. The topology of these phase diagrams, as well as the number of domains where different mechanisms dominate, depends on the form of the DM-bath interaction, as we illustrate with two examples arising from a quartic scalar interaction and a Yukawa interaction.  Furthermore, we argue that there exists a certain ``universality\" to the phase diagram behaviour at small coupling.  After presenting the detailed calculation of the relic abundance in various cases in Section~\\ref{calculation}, we discuss the physics of FIMP interactions mediated by higher-dimension-operators, as well as some variations of the basic freeze-in mechanism in Section \\ref{comments}.  Finally we conclude in Section \\ref{conclusions}.  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% ",
        "conclusions": "%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%   The nature and origin of the dark matter in the universe is unknown.  While there are many theories of DM, there are rather few cosmological production mechanisms, and even fewer that can be subjected to precision tests. For example, the non-thermal coherent oscillation of a scalar field, such as an axion, always involves an unknown initial field amplitude.  Thermal mechanisms that are independent of initial conditions are particularly interesting, and the decoupling of a heavy particle species from the thermal bath has received enormous attention.  Such decoupling can occur when the particle is relativistic, as with the neutrinos, or non-relativisitic, as with baryons. The former leads to hot dark matter, so that the latter case of thermal freeze-out has been widely studied for DM.  At first sight these appear to be the only two ways of thermally producing DM without any sensitivity to initial conditions.  If a massive particle of the thermal bath is to survive the early universe with a significant abundance it must decouple, and this can either happen when the particle is relativistic or non-relativistic.  We have argued that there is one other possibility:  the massive particle may have a negligible initial abundance and may be produced by collisions or decays of particles in the thermal bath. This is only a new possibility if the massive particle does not reach thermal equilibrium, and hence requires that it is a FIMP, interacting very feebly with the bath.  Since the FIMP abundance is heading towards thermal equilibrium we call this production mechanism ``freeze-in.\"  Freeze-in requires a very small coupling $\\lambda$ between the FIMP and the bath % \\begin{equation} \\lambda^2 \\sim \\frac{T_{eq}}{M_{Pl}} \\,\\, \\frac{m_{FI}}{m_{DM}} \\label{eq:lambda2} \\end{equation} % where $T_{eq}$ is the temperature of matter-radiation equality, $m_{FI}$ is the mass of the heaviest particle involved in the freeze-in reaction and $m_{DM}$ is the mass of the DM particle.  The mechanism works for a very wide range of FIMP masses: to be produced it must be lighter than the reheat temperature after inflation, and for DM to be cold it must be heavier than a keV.   Indeed the range of theories involving FIMP freeze-in is so large that in this paper we have concentrated on FIMPs with a mass of order the weak scale, such as moduli or right-handed sneutrinos.  It is not surprising that freeze-out is so popular: it presumably occurred with baryons (although with the added complication of a chemical potential) and it works with dimensionless coupling parameters of order unity.  Indeed, with couplings of order unity, if the only mass scale in the annihilation cross section is that of the DM particle, then this mass is predicted to be of order the weak scale.  On the other hand freeze-in of a FIMP requires a special situation.  The freeze-in reaction must involve a small dimensionless coupling parameter, eq.~(\\ref{eq:lambda2}); furthermore, UV sensitivity may reappear unless higher dimensional operators are suppressed or the reheat temperature after inflation is quite low.  However, small couplings for both renormalisable and higher dimensional operators occur very easily, for example by approximate symmetries or small wavefunction overlaps in higher dimensions.  Most importantly, precisely because FIMPs necessarily have very small couplings to the bath, eq~(\\ref{eq:lambda2}), they offer the prospect of exotic signals of dark matter generation.  If the FIMP is the lightest particle carrying the DM stabilising symmetry it is the DM, so that the lightest observable sector particle (LOSP) with this symmetry is unstable and decays to the FIMP with a rate suppressed by $\\lambda^2$.  Alternatively it may be the LOSP that is DM, and in this case it is the FIMP that has a very long lifetime.   The signals associated with the freeze-in mechanism all revolve around the long lifetime of decays between the FIMP and LOSP. In particular, we explored the signals that arise when the FIMP, LOSP and freeze-in masses are all of order the weak scale.  %If in addition the small coupling %takes the parametric form $\\lambda \\sim v/M_{Pl}$, where $v$ is the electroweak vev, as happens for weak-scale moduli, then the %temperature of matter-radiation equality from DM produced by the freeze-in mechanism has the %same parametric form as in the WIMP case: $T_{eq} \\sim v^2 / M_{Pl}$. In many theories a small coupling may arise from a parametric form that is linear in the weak scale $v$, $\\lambda \\sim v/M_{Pl}$.    In fact, in many models the Planck scale will be replaced by the string or compactification scale, as occurs for certain moduli, providing a better explanation for the order of magnitude of $\\lambda$.   The temperature of matter-radiation equality from DM produced by the freeze-in mechanism then has the same parametric form as in the WIMP case: $T_{eq} \\sim v^2 / M_{Pl}$. This not only explains why freeze-in can yield the observed abundance, but suggests that two mechanisms for DM production -- FIMP freeze-in and LOSP freeze-out -- yield broadly comparable abundances.  The cosmological signals of the FIMP arise from the late decays between the FIMP and the LOSP and, with masses of order the weak scale, the parametric form of the lifetime in the simplest theories is % \\begin{equation} \\tau \\, \\sim \\, \\frac{M_{Pl}}{T_{eq}\\, v} \\, \\sim \\, 1 \\, \\mbox{sec.} \\label{eq:tau} \\end{equation} % A precise computation shows that if the FIMP is frozen-in via two body decays or inverse decays that involve the LOSP, then the lifetime is closer to $10^{-2}$ seconds, as shown in eq.~(\\ref{eq:tauB1}). However, this should be viewed as a lower bound on the lifetime of decays between the FIMP and LOSP.  Further suppression of the decay rate occurs if the freeze-in process involves bath particles other than the LOSP or if the decays are to three or more particles.  Thus the lifetime may well be in the region where the decays can alter the nuclear abundances produced during BBN, possibly solving the $^7Li$ problem, and may even be long enough to allow a component of the DM to be warm.  If FIMP freeze-in dominates over LOSP freeze-out, then the LOSP annihilation cross section is larger than with conventional freeze-out.  This leads to an important astrophysical signal of LOSP DM: the DM annihilation in the galactic halo is boosted relative to the conventional case, leading to enhanced rates for indirect detection signals of photons, leptons and anti-protons.  Finally, FIMP dark matter yields important new collider signals.  The exotic particles of the LOSP sector can be produced at the LHC and will rapidly decay to the LOSP, which has a long lifetime, bounded from below by about $10^{-2}$ sec.  If the LOSP is charged or coloured it has a significant probability of being stopped in the detector, so that its decays can be observed.  Studying the production and decay of the LOSP at the LHC could directly probe the freeze-in reaction that created the FIMP dark matter in the early universe."
    },
    "0911/0911.2430_arXiv.txt": {
        "abstract": "{We analyse the radial and tangential velocity fields in the Galaxy as seen from the Sun by using as a first approximation a simple axisymmetric model, which we then compare with the corresponding fields in a barred $N$-body model of the Milky Way. This provides a global description of these quantities missing in the literature, showing where they take large values susceptible to be used in future observations even for sources well beyond the Galactic center.  Absolute largest proper motions occur at a distance slightly behind the Galactic Center, which are there 1.5 times larger than the highest local proper motions due to the Galactic differential rotation. Large proper motions well beyond the Galactic center are well within the current astrometric accuracy.} ",
        "introduction": "In view of the ongoing and forthcoming advances of astrometry applied to the Milky Way well beyond the Solar neighborhood, such as the measurements of the trigonometric parallaxes and proper motions of distant stellar masers (e.g., Reid et al.~\\cite{reid}), or the upcoming GAIA survey, we develop here a simple first-order model of a galactic disk making explicit where large proper motions and radial velocities due to differential Galactic rotation can be expected in the Galaxy.  It is indeed important to develop quantitative insight on these observables and separate the effects due to differential rotation from those due to, for example, a bar.  The level of analysis is elementary and could figure in astronomy textbooks. Despite literature search and discussions with astrometrists the points raised here do not seem to be commonly known. The probable reason for the lack of a previous similar discussion is that observing proper motions has been considered for a long time as difficult and limited to a few kpc.  Further, dust extinction has also limited observations in the Galaxy plane in the optical.  Therefore, a local analysis has been considered as sufficient (e.g. Mihalas~\\cite{mihalas}, Mihalas \\& Binney~\\cite{mihalasbinney}, Binney \\& Merrifield~\\cite{binneymerrifield}).  In radio astronomy, however, the need to develop a global vision of the Galaxy has been reached much sooner.  An analysis close to the present one concerning the Galaxy radial velocity field has been already done (see e.g. Burton~\\cite{burton}), since extinction in neutral hydrogen is weak and integrated radial velocity field can be directly measured over most of the Galaxy disk.  The most interesting aspect of the present analysis concerns the proper motion field at distances well beyond the Solar neighborhood, which shows a discontinuous behavior just beyond the Galactic center \\textit{allowing to measure easily stellar proper motions there and well beyond}.  The reason is easy to understand if we consider that the disk beyond the Galactic center is moving in first approximation with a velocity opposite to the Sun velocity.  The linear velocity difference of about $2 \\times 220\\,\\rm km\\, s^{-1}$ largely exceeds the inversely to distance decreasing proper motions, leading to proper motions up to about $440 \\,\\rm km\\, s^{-1} / 9\\,kpc \\approx 10\\,mas \\,yr^{-1}$, a well measurable quantity with current techniques.  Of course the circular velocity is not equal to the average azimuthal velocity, the velocity that can be observed by averaging proper motions, due to the non-negligible contribution of the velocity dispersion to radial support. Especially close to the Galactic center the velocity dispersion increases to values comparable to $220/\\sqrt 3\\,\\rm km\\, s^{-1} \\approx 130\\, km\\, s^{-1} $, therefore one should see, as regions beyond the Galactic bulge are measured at increasing distances, a progressive change from individual random motions at about $130 \\,\\rm km\\, s^{-1} / 8\\,kpc \\approx 3\\,mas \\,yr^{-1}$ superposed to a systematic proper motion with respect to an inertial frame due to the Sun motion of about $220 \\,\\rm km\\, s^{-1} / 8\\,kpc \\approx 5\\,mas \\,yr^{-1}$, progressively increasing to about $400 \\,\\rm km\\, s^{-1} / 10\\,kpc \\approx 8\\,mas \\,yr^{-1}$ as we explore regions 2\\,kpc beyond the Galactic center.  This insight was already exposed by Binney (1995) in a little read paper concerned primarily with the non-axisymmetric motion effects induced by a bar.  He already showed that large proper motions can be expected near the Galactic center even in the absence of a bar.  In any case, the systematics introduced by differential motion in the bar/bulge region must not be omitted, as pointed out by Zhao et al.~(\\cite{zhao}).  We also check that a fully consistent model of the Milky Way including the effects of a bar and velocity dispersions does not change the main conclusions reached with the simple model.  For this, we use one of the Milky Way $N$-body models of Fux~(\\cite{fux}) including a bar and spiral arms to quantify the differences existing between a fully consistent, self-gravitating, stable and barred model of the Milky Way, and the simple axisymmetric model.  The plan of the paper is as follows.  In Section 2 we give some definitions. In Section 3 we develop the first order axisymmetric model with analytic tools and a more elaborated rotation curve model. In Section 4 we analyze Fux's $N$-body model and compare it with the axisymmetric ones and with some observational results. In Section 5 we present our conclusions. ",
        "conclusions": "Starting from the analysis of simple axisymmetric models of the Milky Way, we have studied the differential velocity field, and in particular the proper motion field, providing a global description of observables.  Interestingly, there is a region behind the Galactic center extending to several kpc further out where the proper motions induced by differential rotation have the largest amplitudes in all the Milky Way, an apparently little known basic feature (with the exception of Binney\\cite{binney}) with direct observational consequences.  Here we clarify on a global scale the effects related to first order to the average circular rotation and to second order to a bar.  We have checked that these high proper motions are also found in self-consistent barred configurations, such as the Milky Way $N$-body model of Fux~(\\cite{fux}).  The proper motions at distances of 10\\,kpc to the observer are of the order of 7\\,$\\rm{mas \\,yr}^{-1}$ and thus are well within the accuracy of the forthcoming surveys such as GAIA. It is crucial to be able to separate the sources according to distance, at least to separate the sources closer to or beyond the Galactic center.  Provided that the position of the sources, such as stars or OH masers, are intense enough to be measured with respect to an absolute frame of reference, as well as their radial velocities, the global velocity field of the Milky Way including non-circular motions should become accessible.  This will be a first order information for constraining the strength of the bar and the mass distribution in the bulge, the keystone of the whole Galaxy."
    },
    "0911/0911.5356_arXiv.txt": {
        "abstract": "There is a growing body of evidence for the presence of multiple stellar populations in some globular clusters, including NGC~1851. For most of these peculiar globular clusters, however, the evidence for the multiple red giant-branches (RGBs) having different heavy elemental abundances as observed in $\\omega$~Centauri is hitherto lacking, although spreads in some lighter elements are reported. It is therefore not clear whether they also share the suggested dwarf galaxy origin of $\\omega$~Cen or not. Here we show from the CTIO 4m $UVI$ photometry of the globular cluster NGC~1851 that its RGB is clearly split into two in the $U - I$ color. The two distinct RGB populations are also clearly separated in the abundance of heavy elements as traced by Calcium, suggesting that the type II supernovae enrichment is also responsible, in addition to the pollutions of lighter elements by intermediate mass asymptotic giant branch stars or fast-rotating massive stars. The RGB split, however, is not shown in the $V - I$ color, as indicated by previous observations. Our stellar population models show that this and the presence of bimodal horizontal-branch distribution in NGC~1851 can be naturally reproduced if the metal-rich second generation stars are also enhanced in helium. ",
        "introduction": "The first photometric evidence for the presence of multiple stellar populations in globular clusters came from the discoveries of the discrete distributions of stars on the red giant-branch (RGB) and main-sequence (MS) of the most luminous globular cluster $\\omega$~Cen \\citep{lee99,bed04}. The fact that the distribution of RGB stars is not just showing a spread but is discrete is a compelling evidence for the heavy elements enrichment and the formation of successive metal-enhanced generations of stars in proto-$\\omega$~Cen. The similarity of this feature to that of the M54, suggested as a nuclear star cluster of the Sagittarius dwarf galaxy \\citep{lay97}, together with other peculiarities of $\\omega$~Cen, has led the community to conclude that $\\omega$~Cen was once part of a more massive dwarf galaxy that merged with the Milky way, as the Sagittarius dwarf galaxy is in the process of doing now \\citep{lee99,bek03}. In recent years, more and more evidence is reported for the presence of double or multiple stellar populations in other globular clusters, such as NGC~2808 \\citep{pio07}, NGC~1851 \\citep{mil08}, NGC~6388 \\citep{mor09}, and M22 \\citep{mar09,dac09}. For most of these peculiar globular clusters, however, the evidence for the discrete distribution of heavy elements as observed in the RGB of $\\omega$~Cen is hitherto lacking, although spreads in some lighter elements \\citep[][and references therein]{car08} and helium \\citep{lee05,dan05,pio07,yoon08} are reported. Therefore, the case of $\\omega$~Cen, and now that of M22 \\citep{mar09,dac09},  are still viewed as exceptional, and the presence of chemical inhomogeneity in other globular clusters is largely considered as due to the pollution mechanisms expected in normal globular clusters, such as the winds from the intermediate mass asymptotic giant branch (AGB) stars or fast-rotating massive stars \\citep{ven08,dec07}.  The purpose of this Letter is to report that one of these globular clusters, NGC~1851, is showing a clear split in the RGB from our CTIO 4m $UVI$ photometry. This is compared with the $Ca$-$by$ photometry \\citep{jleea09} to confirm that the two distinct RGB populations are different in the abundance of heavy elements. Stellar population models are then constructed to show that, when the metal-rich second generation stars are also enhanced in helium abundance, the split of the RGB discovered in $U - I$ color would not be detected in usual optical colors, such as $V - I$ used in the $HST/ACS$ survey. ",
        "conclusions": "We have shown that the RGB of NGC~1851 is split into two distinct subpopulations. The fact that the distribution of RGB stars does not just show a spread but is discrete can naturally eliminate the possibilities such as (1) star formation from not well mixed inhomogenous interstellar matter, (2) differential reddening (see also section 2), and (3) photometric errors, as all of these would produce a $spread$ in color rather than a discrete distribution. Consequently, the most plausible interpretation for the double RGBs is that the NGC 1851 underwent metal enrichment in its early stage of evolution and successively formed metal-enhanced second generation stars. Spectroscopic observations show star-to-star abundance variations of the lighter elements (elements lighter than Si, such as N, O, Na, \\& Al) in NGC~1851 \\citep{hes82,yon08,yon09}. This is generally interpreted as a result of pollution from winds of intermediate-mass AGB stars \\citep{ven08} or fast-rotating massive stars \\citep{dec07}. Based on the elemental abundances observed by \\citet{yon08} and \\citet{yon09} for eight bright RGB stars, we can estimate the differences in the lighter elements between the two RGB populations. Analysis of these data by \\citet[][see their Figure 3]{jleea09} indicates that, while N, Na, and Al are all enhanced in redder RGB ($\\Delta$[N/Fe] $\\approx$ 0.47, $\\Delta$[Na/Fe] $\\approx$ 0.53, $\\Delta$[Al/Fe] $\\approx$ 0.20), O is depleted ($\\Delta$[O/Fe] $\\approx$ $-$0.45), and no significant variation is shown in [C/Fe].  In order to investigate the effect of these elemental variations in the ($U - I$) color, in Figure 2, we have computed the differences in fluxes between two synthetic spectra using the spectral-synthesis program SPECTRUM \\citep{gra94}, one without and the other with these enhancements (and depletion for O) in lighter elements (panels a \\& b). Also shown in Figure 2 are for the two additional cases where (1) CNO and Na are enhanced by 0.3 dex while heavier elements are fixed (panels c \\& d), and (2) elements heavier than Al are enhanced by 0.3 dex while lighter elements are fixed (panels e \\& f). These simulations demonstrate that small variations either in lighter or heavier elements could cause significant change in ($U - I$) color. The observed variations in the lighter elements alone, however, would cause relatively small line blanketing equivalent to $\\Delta$($U - I$) $\\simeq$ 0.079 in terms of color difference. Despite uncertainties, taken at face value, this is only $\\sim$41\\% of the observed color difference [$\\Delta$($U - I$) $=$ 0.195 $\\pm$ 0.011] at given $I$ magnitude and suggests that other effects may also be responsible for the RGB split in ($U - I$) color. Indeed, recent $Ca$-$by$ photometry \\citep{jleea09} shows that besides lighter elements variations, RGB stars of NGC~1851 also show bimodal distribution in Ca, which can only be supplied by Type II supernovae \\citep[SNe II;][]{tim95}. In Figure 3, we show in our ($U$, $U - I$) CMD the ``Ca-normal'' and ``Ca-strong'' stars from \\citet{jleea09}. The Ca-strong stars lie well on the redder RGB sequence, whereas the Ca-normal stars are on the bluer RGB. According to \\citet{jleeb09}, the difference in Ca abundance is estimated to be $\\Delta$[Ca/H] $\\approx$ 0.15 dex, and other heavy elements, albeit small, are similarly enhanced in redder RGB population. We conclude therefore that the redder RGB population is richer than the bluer RGB population not only in lighter elements (N, Na, \\& Al) but also in heavy elements (such as Ca, Si, Ti, \\& Fe).  In order to better understand the origin of the RGB color difference in ($U - I$), and to place stronger constraints on the chemical combinations of the two subpopulations, we have constructed stellar population models based on the updated version of the Yonsei-Yale ($Y^2$) isochrones (Yi et al., in preparation) and HB evolutionary tracks (Han et al., in preparation). Readers are referred to \\citet{yoon08} and references therein for the details of our model construction. Figure 4 presents our synthetic CMDs for NGC~1851 in ($U$, $U - I$) and ($V$, $V - I$) planes. Our models are constructed under two different assumptions regarding the chemical enrichment in NGC~1851. First, we assumed that the second generation population is more enhanced in metallicity, but not in helium (Figure 4a; hereafter $\\Delta$Z-only model). Note that, while \\citet{cas08} and \\citet{sal08} suggested the difference in only CNO abundance, here we are assuming the difference in overall metallicity of 0.15 dex as discussed above. Second, in Figure 4b, we then assumed that both metal and helium abundances are enhanced (hereafter $\\Delta$Z+$\\Delta$Y model). The $U$ flux is very sensitively affected by CN and NH bands, and therefore by N abundance. In order to reflect this effect in the ($U$, $U - I$) CMD of the second generation population (redder RGB), we are also including the effects of the additional line absorptions from the enhancements in N and other lighter elements as discussed above, again by using SPECTRUM. More rigorous modeling should include the effects of these lighter elements enhancements in the construction of stellar evolutionary tracks, but we note that \\citet{dot07} found the enhancement in N has only negligible effect in the HR diagram morphology at globular cluster ages.  Comparison of the number ratio between the two subpopulations suggests that bluer RGB, brighter SGB, and redder HB are associated with one subpopulation, which we refer to as ``Pop-1'', while the redder RGB, fainter SGB, bluer HB are associated with the other subpopulation, which we refer to as ``Pop-2''. Pop-1 comprises about 75\\% of the total population, while Pop-2 takes about 25\\% of the whole population\\footnote{We have compared the ratio of the two subpopulations with samples selected in different manners (e.g. entire sample and a sample selected with various separation indices) and for all cases, the ratio comes out to be similar. This is also more or less consistent with the population ratio based on the SGB stars reported in \\citet{mil09}.}. The $\\Delta$Z-only model in Figure 4a ($\\Delta$Z = 0.0004 and $\\Delta$age = 0.1 Gyr) matches well with the observed CMD from the MS through the RGB in ($U$, $U - I$) CMD. Yet, the model fails to reproduce the blue HB and the narrow RGB in ($V$, $V - I$) CMD (Figure 4a, inset). This is because metal enhancement in Pop-2 moves HB to the opposite direction, i.e., makes HB morphology redder \\citep[see for example][]{lee94}. At the same time, the increase in metallicity, albeit small, will create a gap in ($V - I$) color between the two RGBs which is not shown in the observed CMD \\citep{mil08}. Therefore, the $\\Delta$Z-only model is in conflict with the observed CMDs of NGC~1851. Figure 4b shows that the $\\Delta$Y+$\\Delta$Z model ($\\Delta$Z $=$ 0.0004, $\\Delta$age $=$ 0.1Gyr, and $\\Delta$Y $=$ 0.05) is in good agreement with the observation from the MS to the HB. Note that, in our modeling, we are just employing the standard \\citet{rei77} mass-loss law, and the same mass-loss parameter $\\eta$ for both normal helium and enhanced helium populations. Input parameters adopted in our $\\Delta$Y+$\\Delta$Z model are listed in Table 1. The enhanced helium in Pop-2 easily overcomes the effect of metallicity on HB \\citep[see][]{lee05}, shifting HB morphology toward blue and reproducing the observed HB bimodality. At the same time, the increase in helium abundance in Pop-2 moves RGB slightly bluer and this effect cancels out the metallicity effect in ($V$, $V - I$) CMD, making apparently single and narrow RGB. Also, the model (Figure 4b, inset) is in better agreement with the observed SGB split\\citep{mil08}.  The $U$-band is much more sensitive to metal line blanketing, and therefore the increased helium has relatively small effect on $U - I$ color of RGB. One caveat to this helium enhanced scenario is that the blue HB in our model appears to be slightly brighter ($\\sim$0.05 mag) than the observation in $U$-band, although this could be due to the uncertainty in bolometric correction of the $U$-band.  The apparent helium enhancement in the second generation subpopulation in NGC~1851 is reminiscent of the cases of other globular clusters with multiple populations, including $\\omega$~Cen \\citep{nor04,lee05,pio05}, NGC~2808 \\citep{lee05,dan05}, NGC~6388 and NGC~6441 \\citep{calo07,yoon08}. Although the origin of this helium enhancement is currently not fully understood, general consensus is that it can be supplied either by SNe II \\citep{nor04,pio05} or by the intermediate mass AGB stars or fast-rotating massive stars \\citep{ven09,dec07}. In the case of NGC~1851, given the enhancements both in the lighter (such as N) and heavy (such as Ca) elements in the second generation population (Pop-2), all of the mechanisms seem to be responsible for the helium enhancement. A possible scenario is that, soon after the formation of the first generation stars (Pop-1), numerous SNe II explosions enriched both metal and helium of the leftover gas in the proto-NGC~1851. Winds and ejecta from the intermediate mass AGB stars may have added more helium and simultaneously enhance lighter elements. The second generation stars (Pop-2) would have then formed from the gas enriched in overall metallicity, helium, and lighter elements. We note that the present mass \\citep[$\\sim$$10^6$$M_\\odot$;][]{pry93} of NGC~1851 is too small to retain the ejecta from numerous SN explosions \\citep{bau08,dop86}. Therefore, as in $\\omega$~Cen \\citep{lee99,bek03}, we suggest that NGC~1851 is most likely the remaining core of more massive primeval dwarf galaxies that merged and disrupted to form the proto-Galaxy, as recently proposed by \\citet{lee07} for the origin of globular clusters with the extended HB. Our result presented here is calling a new survey of globular clusters employing $UV$ and/or Calcium filters in order to detect small difference in metallicity expected in the yet to be identified globular clusters with multiple populations. High resolution spectroscopy of stars in the two distinct groups are also needed to confirm the small difference in the abundance of heavy elements."
    },
    "0911/0911.0573_arXiv.txt": {
        "abstract": "Based on a set of cosmological N-body simulations we analyze properties of the dark matter haloes (DM) in a galaxy mass range ($10^{11} - 10^{13} h^{-1}M_{\\odot}$) in modified $\\lcdm$ cosmology with additional dynamically screened scalar interactions in DM sector. Our simulations show that scalar interactions support picture of the Island Universe. Rapid structure formation processes are shifted into higher redshifts resulting in a much smaller accretion and merging rates for galactic haloes at low redshifts. Finally, we present how this ``fifth'' force affects halo properties, like density profile, triaxiality, ellipticities and the spin parameter. ",
        "introduction": "Introduction}  In this paper we study the impact of the  long-range scalar interactions (LRSI) in the dark matter (DM) sector \\cite{GP1,GP2} on some properties of the dark matter haloes. LRSI was proposed to overcome some difficulties that $\\lcdm$ model faces at scales important to the galaxy formation. Large-scale clustering of LRSI cosmologies was studied in \\cite{NGP,LRSI1,LRSI2}. Here we will study  DM haloes density profiles, mass accretion histories, shape and spin parameters. ",
        "conclusions": ""
    },
    "0911/0911.4273_arXiv.txt": {
        "abstract": "\\par\\parskip=1.truecm  The proton flux  and the chemical composition of the cosmic radiation measured, respectively,  by the Kascade and Auger experiments entail radical changes in Cosmic Ray Physics.  A large discrepancy  emerges  by comparing the proton flux predicted by the dip model and that measured by Kascade in the critical energy interval $5\\times10^{16}$-$10^{17}$ eV. It is mentioned and substantiated that the  proton flux measurements of the Kascade experiment are  consistent with other pertinent empirical observations. It is shown that the chemical composition measured by Auger by two independent procedures,  using the mean depth reached by cosmic nuclei in giant air cascades, is incompatible with that predicted by the dip model. \\newline \\quad A notable consequence suggested here based on the failures of the dip model is that the spectral index softening of the primary cosmic radiation above $6\\times 10^{19}$ eV observed by HiRes and Auger experiments, is not due to the extragalactic cosmological protons suffering energy losses in the intergalactic space via the reactions,  $p$ $\\gamma$  $\\rightarrow$  $\\pi^0$ $p$, $\\pi^+$ $n$, but to some physical phenomena occurring in the cosmic vicinity. ",
        "introduction": "An abrupt progress has been recently occurred in Cosmic Ray Physics, which in many respects is a revolutionary change of the current notions of the discipline, due to some measurements  of the Auger and Kascade Collaborations.  The Auger experiment has determined the chemical composition of the cosmic radiation by measuring the mean depth reached by cosmic nuclei in giant terrestrial cascades in air \\cite{{augerloga},{sigmaauger2009}} or the equivalent variable, $X_{max}$. By two independent methods it has been established that in the energy interval $4\\times10^{17}$-$5\\times10^{19}$ eV the cosmic radiation attaining the solar cavity  consists predominantly of intermediate nuclei,  and not of pure protons. Figure 1 reports the $X_{max}$ versus energy measured by the Auger experiment (red dots) \\cite{augerloga}.  \\par The Auger apparatus determines  $X_{max}$  by fluorescent light released by nuclei penetrating the air. The longitudinal light profile recorded by the instrument is interpolated by an appropriate function, $f_{LP}$, which has a characteristic width denoted $\\sigma({X_{max}})$.  The measurement of $\\sigma({X_{max}})$ at a given energy for a number of atmospheric cascades determines the average value of the chemical composition of the cosmic radiation, which is a second method, besides $X_{max}$. Figure 2 shows $\\sigma({X_{max}})$ versus energy (black dots) measured by Auger \\cite{sigmaauger2009} along with its theoretical estimates of $\\sigma({X_{max}})$ for iron nuclei (turquoise lower band) and for protons (turquoise upper band).   In a series of measurements the Kascade experiment has determined that the proton flux in the energy interval $5\\times10^{16}$-$10^{17}$ eV  is $(1-2) \\times 10^{15}$ particles/($m^2$ sr s $eV^{1.5}$) \\cite{{kaskaqgsjet},{kaskasibyll}}. The proton flux has been measured by two methods hereafter referred to as QGSjet and Sibyll algorithms. Figure 3 reports the proton energy spectrum measured by Kascade using the QGSjet algorithm \\cite{kaskaqgsjet}. Data from the Sibyll algorithm \\cite{kaskasibyll} are equivalent for the purpose of this study and omitted for brevity.     \\begin{figure}[htb] \\vspace{-0.3cm} \\includegraphics [angle=0,width=8cm,height=8cm] {fig01aa.eps} \\vspace{-1.5cm} \\caption{ Measurements of  $X_{max}$ by Auger (red dots) compared with the corresponding $X_{max}$ derived from the dip model (turquoise band). The lower and upper theoretical values of $X_{max}$  of the dip model at a given energy correspond to different theoretical models of hadronic interactions in air.} \\label{fig:largenenough} \\vspace{-0.3cm} \\end{figure}    It is the aim of this $\\it {Letter}$ to show that the three independent quoted measurements  are incompatible with the corresponding predictions of the  dip model \\cite{{modelloafossa1},{fossa2}}.  This $\\it {Letter}$ benefits of  a critical examination  of the HiRes data on $X_{max}$ made in another study \\cite{codplochemical}.  Figure 4 shows that above  $10^{17}$ eV the HiRes data  \\cite{belzsalina} on $X_{max}$ converted into $<$$ln(A)$$>$ differ systematically from those of the Volcano Ranch, Yakutsk, Akeno, Agasa, Haverah Park, Fly's Eye and Auger experiment \\cite{codplochemical}. Figure 5 shows that the $\\sigma({X_{max}})$ profile measured by HiRes  \\cite{sokolskinorvegia}   also differs from that measured by Auger \\cite{sigmaauger2009} .   \\begin{figure}[htb] \\vspace{-0.3cm} \\includegraphics [angle=0,width=8cm,height=8cm] {fig02.eps} \\vspace{-1.5cm} \\caption{Measurements of  $\\sigma({X_{max}})$  by Auger (black dots) and the related theoretical estimates \\cite{sigmaauger2009} for a cosmic radiation consisting of protons only (upper colored band) and Fe nuclei only (lower colored band).} \\label{fig:largenenough} \\vspace{-0.4cm} \\end{figure}     \\begin{figure}[htb] \\vspace{-0.3cm} \\includegraphics [angle=0,width=8cm,height=8cm] {fig03aa.eps} \\vspace{-1.5cm} \\caption{Extragalactic proton flux at $10^{18}$ eV according to the dip model \\cite{aloisiofigure2} (black square) with the standard index  $\\gamma_s$= $2$  and its extrapolation (black curve) down to $10^{17}$ eV calculated in this study. The proton spectrum of the same model with $\\gamma_s$= $2.7$ (blue curve) normalized $\\it {ad}$ $\\it {hoc}$ to some experimental data \\cite{aloisiofigure2}, that measured by Kascade (red dots) \\cite{kaskaqgsjet}, and that predicted by the $\\it {Theory}$ $\\it {of}$ $\\it {Constant}$ $\\it {Spectral}$ $\\it {Index}$ (red curve) are also displayed for comparison \\cite{flussimisurati}.} \\label{fig:largenenough} \\vspace{-0.4cm} \\end{figure} ",
        "conclusions": "The fundamental tenet of the dip model  is that protons observed in the solar cavity above $4\\times 10^{17}$ eV have a cosmological origin and outnumber any other ion fractions. Since the dip model predicts unreal properties of the cosmic radiation, the $Logic$ dictates, out of a few alternatives, that spatial sources of the cosmic-ray protons are in the cosmic vicinity ($\\it {Galaxy}$ or $\\it{Local}$ $\\it{Group}$  or $\\it{Local}$ $\\it{Supercluster}$ $\\it{of}$ $\\it{Galaxies}$). In this circumstance: how could a softening of the spectral index from 2.6 to about 3.5 take place in the interval $5\\times10^{19}$-$3\\times10^{20}$ eV, if the extragalactic cosmological protons below $5\\times10^{19}$ eV do not reach the solar cavity ?  As the reaction $p$ $\\gamma$ $\\rightarrow$  $e^-$ $e^+$ $p$ does not alter any cosmic-ray spectrum in the solar cavity (first alternative), the companion reactions $p$ $\\gamma$ $\\rightarrow$ $\\pi^0$ $p$, $\\pi^+$ $n$, etc.  can not cause the index softening above $5\\times10^{19}$ eV just because the extragalactic cosmological protons are missing.  The $ad$ $hoc$ hypothesis that a cosmological extragalactic component would reach the solar cavity in the interval $6\\times10^{19}$-$3\\times10^{20}$ eV but not below this energy band (second alternative) confront with three barriers: (1) protons suffering the energy losses via $p$ $\\gamma$ $\\rightarrow$ $\\pi^0$ $p$ do not disappear but they accumulate below $6\\times10^{19}$ eV corrugating the spectrum; the dip model, as an example,  attempts to calculate this effect with a simplified calculation scheme. (2) Any extragalactic accelerators, operating in the specific interval $6\\times10^{19}$-$3\\times10^{20}$, and eventually above this maximum energy, would intrinsically  generate enough spill-over particles in the low energy edge, below $6\\times10^{19}$, which plausibly corrugate the galactic spectrum. (3) The rising dominance of heavy ions with energy discovered by the Auger Collaboration (see figures 1 and 2) in the band $10^{19}$-$4.12\\times10^{19}$ eV does not harmonize, for a number of reasons, with the aforementioned index softening.  The common value of 2.6-2.8 of the spectral index of the cosmic radiation close and below $6\\times10^{19}$ eV observed in all experiments (HiRes, Auger, Agasa, Yakutsk), disfavors and belittles the second alternative.   According to the present investigation, the softening of the spectral index of the cosmic radiation above $6\\times10^{19}$ eV is not due to any extragalactic cosmological protons.  This conclusion might represent the most notable result rooted to the reaction $p$ $\\gamma$ $\\rightarrow$  $e^-$ $e^+$ $p$ as conceived and framed in the dip model.     \\vspace{-0.2 cm}"
    },
    "0911/0911.4790_arXiv.txt": {
        "abstract": "In this paper we analyzed the behavior of the unusual dwarf nova EM Cyg using the data obtained in April-October, 2007 in Vyhorlat observatory (Slovak Republic) and in September, 2006 in Crimean Astrophysical Observatory (Ukraine). During our observations EM Cyg has shown outbursts in every 15-40 days. Because on the light curves of EM Cyg the partial eclipse of an accretion disc is observed we applied the eclipse mapping technique to reconstruct the temperature distribution in eclipsed parts of the disc. Calculations of the accretion rate in the system were made for the quiescent and the outburst states of activity for different distances. ",
        "introduction": "EM Cyg is an eclipsing dwarf nova which has relatively long orbital period ($P_{orb}= 6^h.98$). Dwarf novae are evolved binary stars, where the Roche lobe-filling red dwarf component accretes matter onto the white dwarf component. Loosing angular momentum, matter forms an accretion disc. Such systems produce outbursts, caused by instabilities in the disc.  EM Cyg was discovered by Hoffmeister in 1928. \\citet{mk69} investigated instabilities on the eclipse light curves and noted, that they reflect instabilities in accretion disc. \\citet{sti82} detected rapid oscillations on light curves with typical time scale of 14.6 and 16.5 seconds.  \\citet{br77}, using long-term photographic observations found outburst cycle to be 24-25 days and noted the seasonal changes of the light curve and the outburst cycle.  Because EM Cyg shows standstills \\citep{nor00}, it is classified as Z Cam type star. So EM Cyg is the only eclipsing dwarf nova among Z Cam type stars.  Using radial velocities, \\citet{rob74} estimated masses to be 0.7 $M_{\\odot}$ and 0.9 $M_{\\odot}$ for the white dwarf and the red dwarf respectively. These parameters corresponded to the thermal-timescale mass transfer in the system, and were doubtful. \\citet{sto81} found that the orbital inclination of the system is about $63^{\\circ}$ and the mass relation close to results obtained by \\citet{rob74}.  \\citet{nor00} found that the spectrum of EM Cyg in the range 6230-6650 $\\AA$ is contaminated by light from a K2-5V star (in addition to the K-type mass donor star). The K2-5V star contributes approximately 16 \\% of the light from the system in that band and, if not taken into account, has a considerable effect upon radial velocity measurements of the mass donor star. The revised value of the mass ratio, combined with the orbital inclination $i=67^{\\circ}$, leads to the masses of $0.99 M_{\\odot}$ and $1.12 M_{\\odot}$ for the mass donor and white dwarf respectively.  Recent more precise measurements of the third star light by Welsh et al. (2007) showed that masses to be $M_{wd}=1.0 M_{\\odot}$ and $M_{rd}=0.77 M_{\\odot}$.  There are several too different estimates of the distance to EM Cyg. From ellipsoidal variations in the infrared band, \\citet{jks81} found secondary spectral type to be K2V and the distance to the system to be 320 pc. \\citet{bay81} using K magnitudes of \\citet{jks81} determined the possible distance to be in the range 285-429 pc with most probable value of about 352 pc. But \\citet{sti82}, using private communication of Stover suggests 400 pc in his calculations.  Taking into account the light of the third star, \\citet{wel05} supposed the distance to be 450 - 500 pc.  Winter and Sion (2003) found 411 parsecs using the $M_v$ versus $P_{orb}$ relation \\citep{war95}, but they determined only 210 parsecs when they fitted the IUE spectra.  Latest paper of \\citet{go09} give the distances ranged from 200 to 380 pc for different WD+disc models of UV spectra.  Model fits for ultraviolet data \\citep{ws03, go09} give mass accretion rate estimates to be about $10^{-10} M_{\\odot}/yr$. This value lies below the upper limit ($2\\cdot10^{-9} M_{\\odot}/yr$), determined by \\citet{cz08} from the O-C diagram analysis.  According to \\citet{voi78}, there are no circular polarization above about 0.3\\% in EM Cyg.  Despite low Galactic latitude (+4.28$^\\circ$), the reddening estimate $E_{B-V}$ towards EM Cyg is only 0.03$\\div$0.05 \\citep{ver87, lad91, be94}.   \\begin{figure}[t] \\includegraphics[width=84mm]{halevin_fig1.EPS} \\caption{Long term light curve of EM Cyg, obtained in 2007 in Vyhorlat observatory  (Humenne, Slovak Republic).} \\end{figure} ",
        "conclusions": "In this paper, using large volume of minima timings we specified the orbital period of the system. Our R band light curves in the quiescent state show significant ellipsoidality effect. Firstly for this star we used the eclipse mapping technique of the outer parts of accretion disc. We estimated the mass accretion rate for different distances and showed, that the distance must be in the range between 300 and 400 pc. Using ellipsoidal variability of the secondary star we estimated the distance to be about 400 pc. Also we showed, that the magnetic braking effect can explain of $1.7\\cdot 10^{-8} M_{\\odot}/year$ mass transfer rate, which has been found for 400 pc distance."
    },
    "0911/0911.3665_arXiv.txt": {
        "abstract": "Radiative diffusion damps acoustic modes at large comoving wavenumber ($k$) before decoupling (``Silk damping''). In a simple WKB analysis, neglecting moments of the temperature distribution beyond the quadrupole (the tight-coupling limit), damping appears in the acoustic mode as a term of order $ik^2\\dot{\\tau}^{-1}$, where $\\dot{\\tau}$ is the scattering rate per unit conformal time.  Although the Jeans instability is stabilized on scales smaller than the adiabatic Jeans length, I show that the medium is linearly unstable to first order in $\\dot{\\tau}^{-1}$ to a slow diffusive mode.  At large comoving wavenumber, the characteristic growth rate becomes independent of spatial scale and constant: $(t_{\\rm KH}\\,a)^{-1}\\approx(128\\pi G/9\\kappa_Tc)(\\rho_m/\\rho_b)$, where $a$ is the scale factor, $\\rho_m$ and $\\rho_b$ are the matter and baryon energy density, respectively, and $\\kappa_T$ is the Thomson opacity. This is the characteristic timescale for a fluid parcel to radiate away its total thermal energy content at the Eddington limit, analogous to the Kelvin-Helmholz (KH) timescale for a radiation pressure-dominated massive star or the Salpeter timescale for black hole growth. Although this mode grows at all times prior to decoupling and on scales smaller than roughly the horizon, the growth time is long, about 100 times the age of the universe at decoupling.  Thus, it modifies the density and temperature perturbations on small scales only at the percent level.  The physics of this mode in the tight-coupling limit is already accounted for in the popular codes {\\tt CMBFAST} and {\\tt CAMB}, but is typically neglected in analytic studies of the growth of primordial perturbations.  The goal of this work is to clarify the physics of this diffusive instability in the epoch before decoupling, and to emphasize that the universe is formally unstable on scales below the horizon, even in the limit of very large $\\dot{\\tau}$. Analogous instabilities that might operate at yet earlier epochs are also mentioned. \\\\ ",
        "introduction": "\\label{section:introduction}  In standard analytic treatments of perturbation theory in the epoch near radiation-matter equality the tight coupling approximation is employed, which amounts to neglecting moments of the temperature distribution beyond the quadrupole.  In this limit, the scattering rate per unit conformal time $\\dot{\\tau}$ is very large and, neglecting gravity at large comoving wavenumber, the dispersion relation for acoustic modes is easily derived; the medium supports stable acoustic oscillations modified by radiative (``Silk'') damping on small scales (Silk 1967, 1968; Peebles \\& Yu 1970; Weinberg 1971; Ma \\& Bertschinger 1995; Hu \\& Sugiyama 1995ab, 1996, 1997; Dodelson 2003).  In the WKB approximation the damping term enters as a complex correction to the acoustic mode frequencies that is $\\propto k^2/\\dot{\\tau}$, becoming increasingly important at small spatial scales (see, e.g., Dodelson 2003, \\S8.4).  Most analytic calculations of radiative damping of acoustic modes neglect gravity at large comoving wavenumber, $k$, because the frequency of acoustic oscillations ($\\omega\\approx c_s k$, where $c_s$ is the sound speed) is large and the medium is Jeans stable. Here, I show that on scales {\\it smaller} than the Jeans length the medium is linearly unstable to a slow diffusive mode that is of order $\\dot{\\tau}^{-1}$.  Thus, on small scales where the medium is dynamically Jeans stable and supports stable (but radiatively damped) acoustic oscillations, it is unstable to an orthogonal diffusive mode.  This mode was earlier discussed by Yamamoto et al.~(1998). The physics of this slow diffusive mode is already included in CMB codes like {\\tt CMBFAST} and {\\tt CAMB}, which include gravity on small scales in the tight-coupling limit in their solution to the coupled equations for perturbations, and in earlier work employing full numerical solutions to the coupled perturbation equations. However, the physics of this mode has been often neglected in analytic studies of the growth of perturbations, and for this reason a more complete discussion is warranted; it may affect analytic estimates for the characteristic baryon-acoustic oscillation (BAO) scale, the physics of low-mass dark matter halos, and the transfer function.  In \\S\\ref{section:acoustic}, I review the physics of acoustic modes and radiative damping in the cosmological context, and I show that the unstable diffusive mode appears in a simple WKB treatment that includes gravity on small scales. The discussion is related to that of the ``terminal velocity'' mode in Yamamoto et al.~(1998).  I emphasize that the growth of this mode is not tied to the breakdown in the coupling between gas and radiation;  indeed, it appears at the same order in $\\dot{\\tau}$ as Silk damping, and a temperature difference between the radiation and the gas is not required for growth.  I note that analogous instabilities should operate in the epoch of neutrino decoupling, or whenever the medium is partially supported against self-gravity by a relativistic and diffusing particle species.  Section \\ref{section:background} presents a complimentary analysis in an expanding background closest in spirit to Silk (1967), and to the discussion of adiabatic modes in standard textbooks. In \\S\\ref{section:discussion}, I provide a brief summary. ",
        "conclusions": "\\label{section:discussion}  The original treatments of acoustic modes near the epoch of radiation-matter equality focused on the importance of radiative damping of acoustic modes at large comoving wave number.  The general equations derived by Silk (1967) (and Silk 1968; Peebles \\& Yu 1970; Weinberg 1971) contain the physics described above, but the slow diffusive mode these equations admit was not discussed because gravity was neglected when calculating the damping of acoustic modes on small scales.  This approximation eliminated the unstable diffusive mode. Similarly, the work of Hu \\& Sugiyama (1996) provides a concise analytic treatment of acoustic mode damping, but again neglects the importance of gravity at large comoving wavenumber. For the unstable mode I have highlighted, the $k$-dependence drops out in the high-$k$ limit because of the $k$-dependence of the diffusion timescale.  Thus, to first order in the limit of large $\\dot{\\tau}$ it is not consistent to neglect gravity at high-$k$. Yamamoto et al.~(1998) have discussed the same instability, but it is important to emphasize that the timescale for the mode's growth is a constant on small scales as long as the tight coupling approximation is valid (eq.~\\ref{unstable_scale}). In this way, the mode's existence is not a manifestation of the breakdown of the tight coupling approximation; even at very high $\\dot{\\tau}$, the instability exists with fixed $t_{\\rm KH}a$.  Most importantly, on all scales smaller than the Jeans length and at all times before decoupling, the medium is unstable, albeit on a long timescale.   Thus, it is not formally correct to think of the medium on scales below the horizon as stable.  Although it is dynamically stable, it is unstable on the Kelvin-Helmholtz timescale for a radiation-pressure supported self-gravitating medium ($\\sim\\kappa_T c/G$).  The instability should operate at all scales smaller than the Jeans length for which the medium is tightly coupled.   The physics of the mode is simply that self-gravitating radiation pressure supported media radiate their total internal energy content at the Eddington limit. It is qualitatively different than the mode described by Gamow (1949) and Field (1971).  The physics of this mode is already contained in the popular codes {\\tt CMBFAST} and {\\tt CAMB} for calculating the growth of perturbations in the tight coupling limit. However, precision analytic and numerical studies of primordial perturbations in density and temperature that neglect gravity on small scales may find systematic small differences with respect to more complete calculations that are attributable to the growth of this mode (e.g., in the transfer function [e.g., Yamamoto et al.~1998], or the BAO scale, or in the initial conditions for the formation of very small scale dark matter halos).  As noted in \\S\\ref{section:acoustic}, the instability identified in equation (\\ref{unstable_scale}) should operate at essentially all times before decoupling because it depends only on the ratio $(\\rho_m/\\rho_b)$ and the opacity $\\kappa_T$.  However, because the growth timescale is constant, it becomes increasingly long compared to the age of the universe at earlier times (see eq.~[\\ref{omegah}]). Thus, one expects it to be most important at the latest times for which the assumptions apply: the largest Jeans-stable modes at a time near decoupling.  Even here, the growth timescale is much longer than the age of the universe, as can be seen from equation (\\ref{omegah}).  More generally, the medium should be unstable to analogous modes any time it is supported against self-gravity by the radiation pressure provided by a diffusing particle species (\\S\\ref{section:dm})."
    },
    "0911/0911.0666_arXiv.txt": {
        "abstract": "{  The exploration of extragalactic objects with long-baseline interferometers in the near-infrared has been very limited. Here we report successful observations with the Keck interferometer at K-band (2.2 $\\mu$m) for four Type 1 AGNs, namely NGC4151, Mrk231, NGC4051, and the QSO IRAS13349+2438 at $z$=0.108.  For the latter three objects, these are the first long-baseline interferometric measurements in the infrared. We detect high visibilities ($V^2 \\sim$ 0.8$-$0.9) for all the four objects, including NGC4151 for which we confirm the high $V^2$ level measured by Swain et al.~(2003). We marginally detect a decrease of $V^2$ with increasing baseline lengths for NGC4151, although over a very limited range, where the decrease and absolute $V^2$ are well fitted with a ring model of radius 0.45$\\pm$0.04~mas (0.039$\\pm$0.003~pc).  Strikingly, this matches independent radius measurements from optical--infrared reverberations that are thought to be probing the dust sublimation radius.  We also show that the effective radius of the other objects, obtained from the same ring model, is either roughly equal to or slightly larger than the reverberation radius as a function of AGN luminosity. This suggests that we are indeed partially resolving the dust sublimation region. The ratio of the effective ring radius to the reverberation radius might also give us an approximate probe for the radial structure of the inner accreting material in each object. This should be scrutinized with further observations.  } ",
        "introduction": "The exploration of extragalactic objects, or in particular, Active Galactic Nuclei (AGNs), with long-baseline interferometers in the near-infrared (near-IR) has been very limited. While the brightest Type 1 AGN NGC4151 and Type 2 AGN NGC1068 have been observed by \\cite{Swain03} and \\cite{Wittkowski04}, respectively, further exploration has been hampered mainly by technical difficulties.  Here we report successful observations of four Type 1 AGNs with the Keck interferometer (KI) in the near-IR (K-band 2.2 $\\mu$m).  Type 1 AGNs are thought to give us a direct view of the innermost region of the putative dust torus as well as the central accretion disk, where the interesting effect of the latter should also be evaluated carefully. ",
        "conclusions": ""
    },
    "0911/0911.0720_arXiv.txt": {
        "abstract": "We describe the infrared properties of a large sample of early type galaxies, comparing data from the Spitzer archive with $Ks$-band emission from 2MASS. While most representations of this data result in correlations with large scatter, we find a remarkably tight relation among colors formed by ratios of luminosities in Spitzer-MIPS bands (24, 70 and 160$\\mu$m) and the $Ks$-band. Remarkably, this correlation among E and S0 galaxies follows that of nearby normal galaxies of all morphological types. In particular, the tight infrared color-color correlation for S0 galaxies alone follows that of the entire Hubble sequence of normal galaxies, roughly in order of galaxy type from ellipticals to spirals to irregulars. The specific star formation rate of S0 galaxies estimated from the 24$\\mu$m luminosity increases with decreasing $K$-band luminosity (or stellar mass) from essentially zero, as with most massive ellipticals, to rates typical of irregular galaxies. Moreover, the luminosities of the many infrared-luminous S0 galaxies can significantly exceed those of the most luminous (presumably post-merger) E galaxies. Star formation rates in the most infrared-luminous S0 galaxies approach 1--10 solar masses per year. Consistently with this picture we find that while most early-type galaxies populate an infrared red sequence, about 24\\% of the objects (mostly S0s) are in an infrared blue cloud together with late type galaxies. For those early-type galaxies also observed at radio frequencies we find that the far-infrared luminosities correlate with the mass of neutral and molecular hydrogen, but the scatter is large. This scatter suggests that the star formation may be intermittent or that similar S0 galaxies with cold gaseous disks of nearly equal mass can have varying radial column density distributions that alter the local and global star formation rates. ",
        "introduction": "If S0 galaxies did not exist it is possible that they would not have been missed. Late type, star-forming galaxies including spirals are expected in early galactic evolution and elliptical galaxies are thought to form by mergers among these ancient late type galaxies or, at a later stage, from dry mergers with each other. S0 or lenticular galaxies are differentiated from these more predictable and recognizable galaxies by being generally structureless with stellar bulges and rotating disks having light distributions that are somewhat less concentrated than in ellipticals. The distinction between E and S0 morphologies is not always easy, particularly when S0s are viewed face-on. This classification difficulty has resulted in the notion of E-S0 galaxies where the uncertain type may either be an intrinsic reality or a consequence of the limitations of the astronomical classifier. S0s are often a nuisance in large surveys since the S0-E distinction decreases with redshift and the two types of galaxies are often combined.  However, many local S0 galaxies are quite distinct from ellipticals, particularly as a result of their significant gas content, star formation rates, and optical colors. At infrared wavelengths many S0 galaxies exhibit a range of unusual properties not often found in ellipticals.  Recently we discussed the infrared properties of 34 S0 and E galaxies in the SAURON sample that have also been observed in the infrared with the Spitzer telescope (Temi, Brighenti \\& Mathews 2009; hereafter TBM09). This relatively small Spitzer-SAURON sample is useful since SAURON galaxies have been intensively observed at many other wavelengths. While the infrared emission from some S0 galaxies in the Spitzer-SAURON sample closely resembles gas-free elliptical galaxies, many S0 galaxies contain rotationally supported kpc-sized disks of star-forming cold gas (e.g. Combes, Young \\& Bureau 2007; Young. Bendo \\& Lucero 2009) and dust that emit more strongly in Spitzer bandpasses than ellipticals. The star formation rates in S0s in the Spitzer-SAURON sample, derived from their 24$\\mu$m emission, are small, $\\lta 0.2$ $M_{\\odot}$ yr$^{-1}$, but correlate with the magnitude of the H$\\beta$ spectral index produced by relatively young A-F stars (TBM09). However, in SAURON S0 galaxies we found that star formation rates estimated from Balmer emission line luminosities are often much less than star formation rates found from 24$\\mu$m emission. The insufficient emission line evidence for OB stars is curious since slightly less massive young stars are often visible. In any case, hydrogen line emission is easily absorbed by dust, so estimates of star formation rates in S0 galaxies are probably more reliably determined from 24$\\mu$m dust emission or by a combination of $H\\alpha$ and infrared luminosities (Kennicutt et al. 2009)   In this paper we discuss infrared Spitzer observations of a much larger, more comprehensive sample of E and S0 galaxies taken from the Spitzer archive. This extended sample reveals some remarkable attributes of S0 galaxies. Most important of these is a very narrow correlation between colors formed by ratios of mid and far-infrared Spitzer luminosities with $K$-band luminosities in which the dust properties differentiate many but not all S0s from gas-free E galaxies. Furthermore, it is remarkable that the color-color correlation for S0s is nearly identical to that of normal galaxies spanning all morphological types. Although ancillary data is less abundant for this extended sample of early-type galaxies than for the SAURON-Spitzer sample, the mass of neutral and molecular hydrogen is known for many of them and this provides several interesting comparisons with star formation rates based on their Spitzer 24$\\mu$m luminosities. ",
        "conclusions": "Our main conclusions are:  \\vskip.1in \\noindent(1){ S0 galaxies are sources of a broad range of mid and far infrared emission. At the gas-free extreme, their emission at 24$\\mu$m is circumstellar and emission at 70 and 160$\\mu$m is likely to result from interstellar dust heated by diffuse starlight and the hot virialized interstellar gas (TMB07). However, many S0 galaxies appear to contain substantial masses of cold gas and dust which is heated by star formation and radiates at all Spitzer MIPS wavelengths. Using estimates of the star formation rate from the 24$\\mu$m luminosity (Calzetti et al. 2007 and Eqn. 2), the most infrared-luminous S0 galaxies are currently forming stars at 1-10 $M_{\\odot}$ yr$^{-1}$. Many S0s have blue $U-V$ colors consistent with recent star formation. }  \\vskip.1in \\noindent(2){ For our large sample of early-type E and S0 galaxies we find that infrared color-color diagrams with ${\\rm Log}(L_{24}/L_{Ks})$ plotted against ${\\rm Log}(L_{70}/L_{Ks})$ or ${\\rm Log}(L_{160}/L_{Ks})$ result in unusually tight correlations. At low values of ${\\rm Log}(L_{70}/L_{Ks})$ or ${\\rm Log}(L_{160}/L_{Ks})$ these plots saturate at ${\\rm Log}(L_{24}/L_{Ks}) \\approx 30.1$, corresponding to the typical SED for gas-free galaxies where there is no evidence for star formation and all the 24$\\mu$m emission is circumstellar dust emission from old population giants. As the mass of cold dusty gas increases in S0 galaxies relative to that in the stars, ${\\rm Log}(L_{70}/L_{Ks})$ and ${\\rm Log}(L_{160}/L_{Ks})$ increase and the color-color correlations are elongated in the direction of decreasing $L_{Ks}$. }  \\vskip.1in \\noindent(3){ Color-color plots of ${\\rm Log}(L_{24}/L_{Ks})$ against ${\\rm Log}(L_{70}/L_{Ks})$ or ${\\rm Log}(L_{160}/L_{Ks})$ for E and S0 galaxies are coincident with the same plots for normal nearby (non-active) galaxies in the SINGS sample (Kennicutt et al. 2003). Furthermore, the SINGS galaxies are arranged along the infrared color-color correlations approximately following the Hubble morphological sequence, E $\\rightarrow$ Sa $\\rightarrow$ Sc $\\rightarrow$ Im. An important exception to this are the S0 galaxies which are spread along the entire color-color correlation. The specific star formation rates per unit $L_{K}$ (or stellar mass) for some S0 galaxies is comparable to those of spiral or irregular galaxies that lie near them in the infrared color-color plot (Fig. 2). The significant gas content in many S0 galaxies is likely to be relevant toward understanding the nature and origin of these galaxies. }  \\vskip.1in \\noindent(4){ Far-infrared luminosities correlate only weakly with the mass of neutral and molecular hydrogen for SINGS galaxies and S0 galaxies from our sample. This suggests that star formation may be intermittent or that disks with nearly equal masses of cold gas are distributed differently with galactic radius in similar S0 galaxies so that the variation of local cold gas (volume and column) densities alters the local and global star formation rates. }  \\vskip.1in \\noindent(5){ Early type galaxies do not have infrared properties that strictly adhere to the red sequence and blue cloud regions of the SDSS color magnitude plot. For example, in an IR color-magnitude plot quite a few galaxies in the SDSS red sequence occupy an infrared blue cloud region with larger Log$L_{24}/L_K$ than passively evolving ellipticals, indicating relatively high specific star formation rates. Conversely, a few galaxies occupying the green valley or blue cloud region of the SDSS color magnitude plot have infrared star formation rates that are indistinguishable from pure gas-free, passively-evolving old stellar populations that define the infrared red sequence. }    \\vskip.1in"
    },
    "0911/0911.4723_arXiv.txt": {
        "abstract": "The study of chemical abundances in stars with planets is an important ingredient for the models of formation and evolution of planetary systems. In order to determine accurate abundances, it is crucial to have a reliable set of atmospheric parameters. In this work, we describe the homogeneous determination of effective temperatures, surface gravities and iron abundances for a large sample of stars with planets as well as a control sample of stars without giant planets. Our results indicate that the metallicity distribution of the stars with planets is more metal rich by $\\sim$ 0.13 dex than the control sample stars. ",
        "introduction": "More than 400 stars with planets have been detected up to this date. One of the interesting properties of these stars concerns their metallicity distribution. Several studies (e.g., \\cite[Fischer \\& Valenti 2005]{fv05}) have confirmed the result first shown by \\cite{g97}: stars with giant planets are systematically metal-rich (by $\\sim$ 0.2 dex) relative to field FGK dwarfs not known to harbor planets. Two hypotheses have been proposed to explain this excess: primordial enrichment or pollution. Current results (e.g., \\cite[Fischer \\& Valenti 2005]{fv05}) show that the frequency of planets increases significantly for higher metallicities, thus giving strong support for the primordial hypotheses. Evidence for the occurence of pollution is still ambiguous (for a comprehensive review, see \\cite[Gonzalez 2006]{g06}). ",
        "conclusions": ""
    },
    "0911/0911.1089_arXiv.txt": {
        "abstract": "{Recent measurements of ultra-high energy cosmic rays and neutrinos are reviewed. With several new large scale observatories nearing completion or becoming fully operational only very recently, a large body of high quality and high statistics data is growing up now. Already these first data have started to open up a new window to the high energy Universe giving us first direct clues about the origin of the most energetic particles with energies of about $10^{20}$~eV as well as about their interactions from extragalactic sources to Earth. Also, for the first time full sky views of high energy neutrinos have become available with neutrino telescopes operating on either Hemisphere. While a ``smoking gun'' is still missing on galactic sources of cosmic rays, constraining upper limits to neutrino fluxes from various source candidates are reported. Thus, future neutrino telescopes, such as {\\sc KM3NeT} in the Mediterranean should aim at volumes significantly larger than one cubic kilometer. Besides seeking the sources of galactic and extragalactic cosmic rays, the new generation of cosmic ray and neutrino observatories touches a wide range of scientific issues and they have already provided important results on tests of fundamental physics.}  \\FullConference{European Physical Society Europhysics Conference on High Energy Physics, EPS-HEP 2009,\\\\ July 16 - 22 2009\\\\ Krakow, Poland}  \\begin{document} ",
        "introduction": " ",
        "conclusions": "Remarkable progress has been made in high energy cosmic ray and neutrino physics over the last two years. In case of UHECRs, the GZK-structure in the energy spectrum is observed with high statistical accuracy and anisotropies at the highest energies have emerged from the data. The coincidence of seeing anisotropies just above the GZK-threshold confirms the picture of the GZK-effect acting as a filter to nearby sources. Even though the directional correlation to nearby AGN is still under debate, the small scale of the angular correlations suggests light particles as primary particles. However, observations of the position of the shower maximum, $X_{\\rm max}$, in the atmosphere and the level of fluctuations of $X_{\\rm max}$ suggest mixed heavy component \\cite{M-ICRC09}. This puzzle is to be solved. It could be due to the poorly known hadronic interactions at cms energies 2-3 orders of magnitude higher than at the future LHC collider and/or due to a strongly increasing cross section at the highest energies, or due to weaker magnetic fields than generally assumed.  The sensitivity of the neutrino telescopes to point sources and diffuse fluxes has been improved over several orders of magnitude and over a wide range of energies during the last years. Still, no point sources or diffuse fluxes are observed. About 10 months of full IceCube equivalent data are at hand now and the chances are high to discover astrophysical neutrino sources within the next few years. However, based on the information from TeV gamma and UHECR observations, the neutrino fluxes from galactic and extragalactic sources seem to be lower than thought a few years ago. Hence, analyzing individual neutrino sources in some detail at TeV energies will require instruments much larger than a cubic kilometer. This should be considered when designing the future {\\sc KM3NeT} telescope. Similarly, due to the GZK break in the UHECR energy spectrum, much larger observatories than available at present are required to collect sufficient statistics for CR-astronomy of individual sources. Related R\\&D efforts, such as radio and acoustic particle detection methods, are ongoing to face these challenges.   \\subsection*{Acknowledgement} I would like to thank the organizers of the EPS high energy physics conference for giving me the opportunity to present these results at such an important meeting. Also, it is a pleasure to thank my colleagues from Auger and IceCube for stimulating discussions. The German Ministry for Research and Education (BMBF) is gratefully acknowledged for financial support."
    },
    "0911/0911.4504.txt": {
        "abstract": "We present a possible OVIII X-ray absorption line at $z=0.117 \\pm 0.001$ which, if confirmed, will be the first one associated with a broad HI Ly$\\beta$ (BLB: FWHM=$160^{+50}_{-30}$ km s$^{-1}$) absorber. The absorber lies along the line of sight to the nearby ($z=0.1372$) Seyfert 1 galaxy PKS~0558-504, consistent with being a WHIM filament. The X-ray absorber is marginally detected in two independent XMM-Newton spectra of PKS~0558-504, a long $\\sim 600$ ks Guest-Observer observation and a shorter, $\\sim 300$ ks total, calibration observation, with a combined single line statistical significance of 2.8$\\sigma$ (2.7$\\sigma$ and 1.2$\\sigma$ in the two spectra, respectively). When fitted with our self-consistent hybrid-photoionization WHIM models, the combined XMM-{\\em Newton} spectrum is consistent with the presence of OVIII K$\\alpha$ at $z=(0.117 \\pm 0.001)$. This model gives best fitting temperature and equivalent H column density of the absorber of log$T=6.56_{-0.17}^{+0.19}$ K, and logN$_H=(21.5 \\pm 0.3) (Z/Z_{0.01\\odot})^{-1}$ cm$^{-2}$, and predicts the marginal contribution of only two more lines within the XMM-{\\em Newton} RGS band pass, NeIX K$\\alpha$ ($\\lambda=13.45$ \\AA) and FeXVII L ($\\lambda=15.02$ \\AA), both with equivalent widths well within the 1$\\sigma$ sensitivity of the combined XMM-{\\em Newton} spectrum of PKS~0558-504 (EW$^{1\\sigma}<3$ m\\AA). The lack of detection of associated OVI in the archival FUSE spectrum of PKS~0558-504, allows us to infer a tighter lower limit on the temperature, of log$T>6.52$ K (at 1$\\sigma$).  The statistical sigificance of this single X-ray detection is increased by the detection of broad and complex HI Ly$\\beta$ absorption in archival FUSE spectra of PKS~0558-504, at redshifts $z=0.1183 \\pm 0.0001$ consistent with the best-fitting redshift of the X-ray absorber. The FUSE spectrum shows a broad (FWHM$=160^{+50}_{-30}$ km s$^{-1}$)  absorption complex, which we identify as HI Ly$\\beta$ $z_{BLB}=(0.1183 \\pm 0.0001)$. The single line statistical significance of this line is 4.1$\\sigma$ (3.7$\\sigma$ if systematics are considered). A possible HI Ly$\\alpha$ is marginally hinted in an archival low-resolution ($\\Delta\\lambda \\sim 6$ \\AA) IUE spectrum of PKS~0558-504, at a redshift of $z=(0.119 \\pm 0.001)$ and with single line significance of $1.7\\sigma$. Thus, the combined significance of the three (XMM-{\\em Newton}, FUSE, and IUE) independent tenative detections, is 5.2$\\sigma$ (5.0$\\sigma$ if the HI Ly$\\alpha$ is not considered, and 4.6$\\sigma$ if the systematics in FUSE are considered).  The detection of both metal and H lines at a consistent redshift, in this hot absorbing system, allows us to speculate on its metallicity. By associating the bulk of the X-ray absorber with the BLB line detected in the FUSE spectrum at $z_{BLB}=0.1183 \\pm 0.0001$, we obtain a metallicity of 1-4\\% Solar.  Although the absorber is only blueshifted by $\\sim -6000$ km s$^{-1}$ from the systemic redshift of PKS~0558-504, the identification of the absorbing gas with a high velocity nuclear ionized outflow, is unlikely. The physical, chemical and dynamical properties of the detected absorber are all quite different from those typically found in the Warm Absorber (WA) outflows, commonly detected in Seyferts and higher luminosity quasars. WA outflow velocities typically span a range of few hundreds to $\\sim 1-2$ thousands km s$^{-1}$; WA metallicities, when measured, are typically found to be at least Solar; high-ionization WAs are virtually always found to co-exist with lower-ionization X-ray and UV phases. All this strongly suggests that the absorber, if confirmed, is an intervening WHIM system. ",
        "introduction": "In the present epoch (z$<1-2$), over half (54 \\%, Fukugita, 2003) of the baryons are missing. They are predicted by Big-Bang Nucleosynthesis (e.g. Kirkman et al., 2003), inferred by density fluctuations of the Cosmic Microwawe Background (e.g. Bennet et al., 2003; Spergel et al., 2007), and seen at $z\\sim 3$ in the ``Ly$\\alpha$ Forest'' (e.g. Rauch, 1998; Weinberg et al., 1997), but by $z < 2$ they are unaccounted for in detected stars and gas (Fukugita, 2003). According to hydrodynamical simulations for the formation of structures in a $\\Lambda$-CDM Universe (e.g. Cen \\& Ostriker, 2006), most, if not all, of these missing baryons should be in a barely visible ``warm-hot'' phase of the filamentary intergalactic medium (the WHIM). The WHIM was shock-heated to temperatures of $10^5-10^7$ K during the continued process of collapse and structure formation, and enriched up to Z$^{WHIM}=0.1-1$ Z$_{\\odot}$ by galaxy super-winds (GSW, Cen \\& Ostriker, 2006). These GSWs efficiently removed metals from the cool and dense ISM phase in galaxies and spread them into the tenuous ($n_e\\sim 10^{-6}-10^{-4}$ cm$^{-3}$, i.e. overdensities $\\delta \\sim 5-500$ compared to the average density of the Universe) WHIM phase, enriching the gas to few percent, and up to tens of percent, of the solar metallicity values  (e.g. Cen \\& Ostriker, 2006). WHIM filaments are too tenuous to be detected through their bremsstrahlung and line emission with current X-ray instruments (e.g. Yoshikawa et al., 2003). However, integrated column densities of high-ionization metal ions along a random section of one of these filaments could reach values as high as $\\sim 10^{16}$ cm$^{-2}$ (OVII in the soft X-Rays) and $\\sim 10^{14}$ cm$^{-2}$ (OVI in the Far-Ultraviolet: FUV), imprinting metal absorption lines in the FUV and soft X-ray spectra of background sources with equivalent width ranging between EW $\\sim 1-20$ m\\AA\\ (the OVII K$\\alpha$ in the X-rays) and EW$\\sim 10-100$ m\\AA\\ (the $\\lambda\\simeq 1032$ \\AA\\ transition, inthe FUV).  The most intense of these absorption lines are the OVI$_{1s^22s\\rightarrow 1s^22p}$ doublet in the UV ($\\lambda=1031.926, 1037.617$ \\AA), and the CV K$\\alpha (r)$ ($\\lambda=40.268$ \\AA), OVII K$\\alpha (r)$ ($\\lambda = 21.602$ \\AA), CVI Ly$\\alpha$ ($\\lambda=33.74$ \\AA) and OVIII Ly$\\alpha$ ($\\lambda=18.97$ \\AA), in the X-ray band. Which of these lines dominates, depends on the ionization state of the gas, that is mainly on its temperature (and at a second order on the gas volume density). Dav\\'e and collaborators (2001) showed that the baryon temperature distribution in the intergalactic space, peaks at log$T\\sim 6.6$ K, and is strongly skewed toward low temperatures: 50\\% of the baryon in the WHIM are found between log$T=5.6$ K and log$T=6.7$ K at a weighted average temperature of log$T=5.9$ K, while the low (log$T=(5-5.6)$ K) and high (logT$=(6.6-7)$ K) temperature tails of the WHIM temperature distribution, contain 27\\% and 23\\% of the gaseous baryons in the Universe, respectively, around weighted average teperatures of log$T=5.4$ K and log$T=6.7$ (Conciatore et al., in preparation). So, while the vast majority (73\\%) of the baryons in the WHIM should absorb and emit in the X-rays at the wavelengths of the He-like (50\\%) and/or H-like (23\\%) ions of C and O, a substantial fraction of the WHIM (27\\%) should be detectable in the FUV,  through Li-like C and O transitions.  These theoretical predictions have been confirmed by FUV observations (e.g. Danforth \\& Shull, 2005, 2008; Tripp et al., 2008; Richther, 2006). Danforth \\& Shull (2005) analyzed FUSE data of 31 AGNs with $z<0.15$ to search for OVI absorption counterparts to 129 known Ly$\\alpha$ absorbers. They found 40 such systems, deriving a $dN_{OVI}/dz(\\ge EW_{Thresh})$ that agrees strikingly with the latest predictions by Cen \\& Fang (2006; see their Fig. 2). However, as pointed out recently by Tripp et al. (2008), the vast majority of the associated HI Ly$\\alpha$ absorption lines are too narrow to be produced in gas with typical WHIM temperatures. Danforth \\& Shull (2008) proposed that low-ionization metals and narrow HI Ly$\\alpha$ trace mildly photoionized medium (the low-$z$ Ly$\\alpha$ Forest: $\\sim 30$\\% of the baryons), while narrow-HI/OVI/NV associations trace a multiphase intergalactic medium, with a photoionized portion of the intervening filament imprinting the narrow-HI absorption and a shock-heated part (WHIM) imprinting the OVI/NV on the UV spectra ($\\sim 10$\\% of the baryons). Danforth \\& Shull (2008) estimate $\\Omega_{WHIM}^{OVI} = 0.34$ \\% (down to logN$_{OVI} > 13.4$ cm$^{-2}$). This is $\\sim 15$ \\% of the ``missing mass'' (strictly speaking, only a lower limit, given the large uncertainties in the ionization correction for the OVI-bearing gas), in good agreement with the WHIM baryon fraction predicted to reside in the low-temperature tail of the WHIM mass-temperature distribution.  Detecting the bulk of the 'Missing Baryons' in the ``X-Ray Forest'', instead, has proven to be extremely difficult. This is because of the unfortunate combination of (a) the still limited resolution ($R\\sim 400$ at $\\lambda = 21.6$ \\AA) and the low throughput ($A_{Eff} \\sim 20-40$ cm$^2$) of the current high-resolution X-ray spectrometers [the XMM-{\\em Newton} Reflection Grating Spectrometer (RGS, den Herder, et al. 2001) and the {\\em Chandra} Low Energy Transmission Grating (LETG: Brinkman et al., 2000)]; (b) the lack of bright ($f_{0.5-2 keV} \\ge 10^{-11}$ erg s$^{-1}$ cm$^{-2}$) extragalactic point-like targets (e.g. Conciatore et al., in preparation); and, (c) the dramatic steepening ($\\Delta\\alpha \\gs 1.5$) of the predicted number density of metal WHIM filaments per unit redshift at ion column densities $\\gs 10^{15}$ cm$^{-2}$.  Current evidence is still limited, and highly controversial: in 2005, Nicastro et al. (2005a,b) reported the first two $\\ge 3\\sigma$ detections of absorption systems identifiable with WHIM filaments, at $z=0.011$ and $z=0.027$, towards the blazar Mkn~421, which was in outburst. One of these two systems is at a redshift consistent with that of a known intervening HI Ly$\\alpha$ absorber (Shull, Stocke \\& Penton, 1996) that line, however, is too narrow, and so the absorbing gas too cold, to be physically associated with the X-ray metal filament. This claim has been questioned in two subsequent papers by Kaastra et al. (2006) and Rasmussen et al. (2007), based on: (a) apparent wavelength inconsistencies for some of the lines, which called into question their association with a given WHIM system; (b) statistical arguments for a significantly lower detection significance for the WHIM identifications than originally reported by Nicastro et al. (2005a,b); and (c) the lack of detection of these two systems in XMM-{\\em Newton} RGS calibration spectra of Mkn~421 (but see also Nicastro, Mathur \\& Elvis, 2008, for rebuttal arguments). Fang, Canizares \\& Yao (2007) proposed the detection of an OVIII WHIM filament along the line of sight to the blazar PKS~2155-304, once again near the redshift of a known narrow HI Ly$\\alpha$ absorber not due to the same gas as the putative X-ray filament, because of the line width. The above reported pieces of evidence for the existence of an ``X-ray Forest'' at low redshift, have been gathered by exploiting the technique of observing random lines of sight toward the brightest possible X-ray sources, with exposures long enough to reach extremely high S/N spectra. This technique has the advantage of not being biased toward the strongest WHIM systems, and so allows probes of the bulk of the WHIM mass distribution. However, the WHIM is predicted, and in the FUV has most likely been found (e.g. Stocke et al., 2006b), to correlate strongly with galaxy concentrations in the Universe: the densest of such concentrations are also supposed to host the nodes, and so the densest and hottest parts, of the WHIM network.  An alternative technique, therefore, is to target the WHIM search in the X-rays along those lines of sight toward bright background X-ray sources that intercept large scale structure concentrations. By exploiting this technique Buote et al. (2009), reported a $\\sim 3\\sigma$ detection of a high column density (N$_{OVII} \\ge 10^{16}$ cm$^{-2}$) OVII K$\\alpha$ absorption line, along the line of sight to the blazar H~2356-309, at the redshift of a very large concentration of galaxies in the Sculptor Wall. This technique has the disadvantage of probing only the extreme high-temperature/density and low-mass fraction ends ($\\sim 23$\\%) of the WHIM distribution.  All the above proposed WHIM detections so far, suffer the serious handicap of lacking the detection of a clearly associated HI counterpart. The detection of broad HI Lyman lines (mostly Ly$\\alpha$ at $\\lambda=1215.67$ \\AA\\ - BLA - and Ly$\\beta$ at $\\lambda=1025.72$ \\AA\\ - BLB) is vital for a proper assessment of the WHIM metallicity and mass, and require the use of both X-ray and FUV facilities. Without HI, metallicity cannot be determined and so a cosmological mass density $\\Omega_b^{WHIM}$ cannot be derived.  \\noindent The expected thermal FWHM of an absortion line imprinted from a gas of particles of mass $m$ at the equilibrium temperature $T$, is given by $2\\sqrt{2ln2}$ times the variance of the velocity distribution (the Doppler parameter $b$), that is FWHM$=2\\sqrt{2ln2} \\sqrt{kT/m}$. For HI, this implies a broadening of FWHM$\\sim 380$ km s$^{-1}$ at log$T=6.5$ % \\footnote{For heavier elements, the width scales with the inverse of the root square of the atomic weight. So, for example, for O, Ne and Fe, at log$T=6.5$, FWHM=95, 85 and 51 km s$^{-1}$.} % .  \\noindent The HI fraction, relative to the total H, in gas with temperatures in the log$T=(5.8-6.3)$ range (where 50\\% of the WHIM should be, e.g. Dav\\'e et al., 2001), spans the interval $10^{-6.2}-10^{-7.2}$. Thus only very shallow (EW(Ly$\\alpha \\simeq 8-80$ m\\AA) and broad HI absorption lines are expected to be imprinted onto FUV spectra by WHIM filaments, requiring spectra with S/N$\\simeq 10-100$ per resolution element to make them detectable.  Here we present the first tentative detection of a hot (X-ray) WHIM filament along the line of sight to the $z=0.1372$ Seyfert galaxy PKS~0558-504, that is physically associated with a BLA HI absorber. A full band analysis of the XMM-{\\em Newton} RGS spectrum of PKS~0558-504, is deferred to a companion paper (Papadakis et al., 2009, submitted). In \\S 2 we present the XMM-{\\em Newton}, FUSE and IUE data that we use in our analysis. \\S 3 and 4 are dedicated to the X-ray and FUV data reduction and analysis. In \\S 5 we critically discuss our findings, and in \\S 6 we summarize our conclusions. ",
        "conclusions": "We reported on the first combined and possibly associated Far-UV and X-ray tentative detection of a WHIM filament at a mean common redshift of $z=0.118 \\pm 0.001$, along the line of sight to the Seyfert PKS~0558-504. Our main findings are:  \\begin{itemize}  \\item{} OVIII Ly$\\alpha$ absorption is tentatively identified at $z=0.117 \\pm 0.001$ in two independent XMM-{\\em Newton} RGS1 spectra of the Seyfert galaxy PKS~0558-504: a high S/N 480 ks net GO spectrum taken in 2008, and the sum of a number of much lower S/N spectra taken between 2000-2001 as part of an XMM-{\\em Newton} calibration campaign, with a total net expsoure of 309 ks. The combined, single-line significance of these two detections is $2.8\\sigma$.  \\item{} When fitted with a model including proper parameterizations of the continua together with our hybrid (collisional ionization plus photoionization) WHIM absorption model, the two ful-band RGS1 and RGS2 spectra are consistent with the presence of a WHIM filament at $z=0.117 \\pm 0.001$ with logT$=6.56_{-0.17}^{+0.19}$ and N$_H = (3.2^{+3.1}_{-1,6}) \\times 10^{19} (Z/Z_{\\odot})^{-1}$ cm$^{-2}$. Based on X-ray data only, the parameters of this model are poorly constrained and both represent only lower limits at a $3\\sigma$ statistical confidence level.  \\item{} Archival FUSE data of PKS~0558-504 show the presence of a broad absorption complex at $\\lambda= 1147.1 \\pm 0.1$, with single-line statistical significance of $4.1\\sigma$ ($3.7\\sigma$ when systematics are included). This line can be identified as broad HI Ly$\\beta$ at $z_{BLB}=(0.1183 \\pm 0.0001)$, consistent with the redshift of the putative X-ray OVIII Ly$\\alpha$, within the large X-ray uncertainties. The width of this BLB absorber is marginally consistent, at a 2.5$\\sigma$ level, with the thermal width of HI ions in gas with logT$=6.56_{-0.17}^{+0.19}$ (the best-fit temperature derived from the X-ray data).  \\item{} Archival IUE-SWP, low-resolution, spectra of PKS~0558-504 taken in 1987 and 1989, also hint to the presence of an unresolved intervening HI Ly$\\alpha$ at $z_{BLA}=0.119 \\pm 0.001$, with a single-line statistical significance of only $1.7\\sigma$. The redshift of this absorber is consistent, within their relative 1$\\sigma$ uncertainties, with the redshift $z_{BLB}$ of the HI absorber detected with FUSE. The BLA (IUE) to BLB (FUSE) equivalent width ratio of these two lines, is consistent with the expected unsaturated value.  \\item{} We tentatively associate the OVIII Ly$\\alpha$ X-ray absorber at $z=0.117 \\pm 0.001$, with the far-UV BLB and BLA absorbers at $z_{BLB}=(0.1183 \\pm 0.0001)$ and $z_{BLA}=0.119 \\pm 0.001$, respectively, and identify them with an intervening WHIM filament at the mean common redshift of $z=0.118 \\pm 0.001$. The combined significance of this detection is $5.2\\sigma$ ($4.6\\sigma$ if systematics in the FUSE continuum modeling are taken into account and the IUE HI BLA line is not considered).  \\item{} The above identification allows us to estimate for the first time the equivalent hydrogen column density of the absorber and so its metallicity under the assumption that the bulk of the HI BLB and the OVIII absorbers are physically associated. These turn out to be: N$_H = (1.5_{-0.4}^{+0.7}) \\times 10^{21}$ cm$^{-2}$ and $Z=(1-4)$\\%$Z_{\\odot}$. If the HI BLB and the OVIII absorbers are not directly associated or are structured, the above estimate represent upper and lower limits, respectively.  \\item{} The non-detection of associated OVI absorption in the FUSE spectrum of PKS~0558-504, allows us to put a stringent lower limit on the temperature of the absorber: by combining X-ray and FUSE data we obtain: logT$=(6.56_{-0.04}^{+0.19})$ K.  \\item{} From the theoretical correlation between the temperature of WHIM filaments and their overdensity (relative to the average density in the Universe) and the expected 3D metallicity-temperature-overdensity relationship, we derive an overdensity of $\\delta\\sim 300$ for our system, and therefore a thickness along the line of sight of $5_{-1}^{+2}$ Mpc.  \\item{} Finally from this single detection we extrapolate the number density of OVIII Ly$\\alpha$ WHIM absorbers, and the cosmological mass density of WHIM. Both are consistent, within their large $1\\sigma$ uncertainties due to the low-number statistics, with the expected values, but their central-values are significantly higher due to the natural bias toward dense (i.e. high equivalent width) and rare absorbers associated with  serendipitous WHIM detections.  \\end{itemize}"
    },
    "0911/0911.1871_arXiv.txt": {
        "abstract": "{ We introduce new statistical methods for the study of cosmic voids, focusing on the statistics of largest size voids. We distinguish three different types of distributions of voids, namely, Poisson-like, lognormal-like and Pareto-like distributions. The last two distributions are connected with two types of fractal geometry of the matter distribution. Scaling voids with Pareto distribution appear in fractal distributions with box-counting dimension smaller than three (its maximum value), whereas the lognormal void distribution corresponds to multifractals with box-counting dimension equal to three.  Moreover, voids of the former type persist in the continuum limit, namely, as the number density of observable objects grows, giving rise to lacunar fractals, whereas voids of the latter type disappear in the continuum limit, giving rise to non-lacunar (multi)fractals.  We propose both lacunar and non-lacunar multifractal models of the cosmic web structure of the Universe.  A non-lacunar multifractal model is supported by current galaxy surveys as well as cosmological $N$-body simulations.  This model suggests, in particular, that small dark matter halos and, arguably, faint galaxies are present in cosmic voids.  } ",
        "introduction": "\\label{intro}  The large scale structure of the Universe is formed by matter clusters, filaments and sheets, and also cosmic voids.  Cosmic voids are the counterpart of matter structures.  They have arisen in observations of the galaxy distribution and have gradually become the subject of deep theoretical investigations.  Early studies of cosmic voids were conditioned by the small number of galaxies surveyed then.  For example, Otto et al.\\ \\cite{Otto} studied the significance of cosmic voids, to distinguish them from fluctuations of a homogeneous galaxy distribution.  They developed probabilistic methods for the study of voids but they actually concluded that cosmic voids were not really significant.  Betancort-Rijo \\cite{Bet} improved and generalized their methods and reached the opposite conclusion.  The r\\^ole of voids as a basic ingredient of the large scale structure is now well established \\cite{Einasto,Kauffmann,ElAd-Pir}.  However, the definition of what constitutes a void is still imprecise. Originally, voids were described as large regions devoid of galaxies and found by visual inspection. As this is hardly satisfactory, an objective way of identifying voids has been long sought.  To identify a void in a point distribution, one needs to decide if it is to be empty or it can have a few points inside and one further needs to determine its shape. The simplest option, preferred by Otto et al.\\ \\cite{Otto} and Einasto et al.\\ \\cite{Einasto}, is to define voids as empty spheres. Kauffmann and Fairall \\cite{Kauffmann} also demand that voids be empty but devise a void-finder that allows more general shapes, with the intention to find approximately ellipsoidal voids. ElAd and Piran \\cite{ElAd-Pir} devise an even more elaborate method: they assume that there is a set of highly clustered ``wall galaxies'' and implement a procedure to select them before the void-finding phase, which joins empty spheres with a ``thinness'' limitation.  Therefore, the found voids are not empty of galaxies and can have fairly complex shapes.  The preceding definitions of voids ignore that galaxy voids can contain dark matter.  Although there is no substantial observational knowledge of the geometry of voids in the dark matter distribution, we have information from cosmological $N$-body simulations.  For example, Gottl\\\"ober et al \\cite{Gott} re-simulate voids with higher resolution and find structures inside them, in a self-similar pattern.  This suggests that the definition of voids as empty regions of simple shape is not appropriate in this case.  Shandarin, Sheth \\& Sahni \\cite{Shan} and Sheth \\& van de Weygaert \\cite{Sheth-vdW} define voids as under-dense regions of a continuous density field, such that they are complementary to matter clusters. With this definition, voids contain matter and have very complex shapes.  Shandarin et al.\\ \\cite{Shan} and Sheth \\& van de Weygaert \\cite{Sheth-vdW} extend methods that have been used successfully for the study of clusters to the study of voids. In fact, under-densities and over-densities are symmetrical in a Gaussian random field.  Colberg et al \\cite{Col_etal} have presented an overview of various void definitions and have made a comparison of the results of applying the corresponding void finding algorithms to the same data set.  The choice of a simple definition of voids, in particular, their definition as empty spheres, is convenient for statistical studies of galaxy voids.  On the other hand, the geometrical aspects of dark-matter voids that arise in cosmological simulations are difficult to relate to the statistics of spherical voids.  The definition of dark-matter voids as under-densities in a Gaussian field is appealing but difficult to connect with any notion of empty voids.  In fact, a Gaussian density field is a valid description only in the linear regime of gravitational clustering, but empty or nearly empty voids are very nonlinear structures.  Therefore, a suitable description of the geometry of voids demands the use of nonlinear models of the cosmic structure.  We focus on fractal models, for they are well founded and provide a useful description of voids.  Fractal models of the large-scale structure of the Universe were introduced by Mandelbrot \\cite{Mandel} and have been well studied \\cite{Borga,Jones-RMP,Pietro_book}, but they are still controversial \\cite{Jones-RMP}.  The major controversy concerns the transition to homogeneity.  Mandelbrot \\cite{Mandel}, inspired by old hierarchical models of the Universe without transition to homogeneity, has indeed suggested that there might not be an outer cutoff to the fractal scaling range.  Assuming that there is a scale of homogeneity, scale invariance is nevertheless justified on smaller scales by observational and theoretical arguments.  The observation of a power-law two-point correlation function of galaxies in a range of scales \\cite{Pee1} was one of Mandelbrot's motivations and is still a strong argument. Furthermore, there are evidences of scaling in higher order correlation functions \\cite{Jones-RMP}.  Theoretical arguments for scale invariance are based on the absence of scales in the gravitational dynamics of collision-less cold dark matter (CDM). This nonlinear dynamics indeed gives rise to a hierarchical formation of structures on every scale up to the homogeneity scale, which grows with time.  Thus, we can reasonably expect that the resulting structure consists in a (multi)fractal attractor of the nonlinear dynamics.  Structure formation in CDM models is studied with $N$-body simulations, which support a multifractal model in a range of scales \\cite{Valda,Colom,Yepes,I4,I5}.  This range is restricted by the intrinsic limitations of these $N$-body simulations (even of state-of-the-art simulations).  The limitations are much less stringent for one-dimensional cosmological dynamics, which is simulated by Miller et al \\cite{Miller}, finding multifractal structure in a very long range of scales. One last but not least argument for scale invariance is that the cosmic web produced by the adhesion model \\cite{adhesion} is found to have multifractal features \\cite{V-Frisch,Bou-M-Parisi}.  Mandelbrot \\cite{Mandel} has also introduced the notion of fractal holes (``tremas'') and considered its application to the galaxy distribution.  In particular, he has introduced the notion of fractal {\\em lacunarity} as a measure of the size of voids in fractals of equal dimension. For any given dimension, one can construct a set of different fractals with progressively decreasing lacunarity, leading to the possibility of a {\\em non-lacunar} fractal. Mandelbrot \\cite{Mandel} indeed shows an example of non-lacunar fractal, the Besicovitch fractal, which in modern terminology belongs to the class of multinomial multifractals.  Mandelbrot \\cite{Mandel} is actually concerned about the low perceived lacunarity of the galaxy distribution and how to reconcile it with the power-law galaxy correlation function.  In our present study of cosmic voids, a non-lacunar fractal indeed appears as the most suitable model.  At any rate, most fractal models of the galaxy distribution proposed so far are lacunar and imply a self-similar distribution of cosmic voids.  The self-similarity of voids has been considered by Einasto et al \\cite{Einasto} as a probe for scale invariance in the large scale structure.  Following the ideas of Mandelbrot \\cite{Mandel}, we have established that the rank-ordering of fractal void sizes fulfills Zipf's power law and tried to confirm it with data from galaxy surveys \\cite{Gaite}.  Thus far, the scaling of galaxy voids remains moot: our analysis does not show any evidence \\cite{Gaite}, but analyses of recent surveys are more favourable \\cite{Tikho-Kara,Tikho,Tikho2}. However, it is questionable that these scalings hold in a sufficiently long range.  In a multifractal geometry, it is natural to define voids as the locations of mass depletions \\cite{I4}. This notion of voids is more general than the notion of voids as {\\em empty} holes and, actually, allows us to define voids in non-lacunar fractals.  We study here the geometry of multifractal voids, which is connected with the geometry of under-densities in a Gaussian field but is actually more complicated.  From the multifractal geometry of dark matter voids, we can derive the geometry of galaxy voids with a model of galaxy biasing.  We consider a simple biasing model, inspired by the ``peak theory'' of Gaussian fields \\cite{Kaiser}, but we substitute Gaussian peaks by multifractal mass concentrations, in accord with the multifractal model of dark matter halos that we have proposed in Refs.~\\cite{I,I4}.  Our model provides a new perspective on Peebles' ``void phenomenon'' \\cite{Pee2}, regarding the emptiness of voids.  We begin with the probabilistic analysis of voids in point distributions. We calculate the probability of spherical voids and the size of the largest void in a Poisson distribution (Sect.\\ \\ref{Poisson}).  We proceed to the general probability of voids in correlated point distributions (Sect.\\ \\ref{correl}) but we obtain less complete results.  We introduce the hierarchical Poisson model and the lognormal model as typical distributions and we connect them with fractal distributions.  To measure the statistics of spherical voids in samples, we devise an efficient void-finder in Sect.\\ \\ref{void-f}. In Sect.~\\ref{scaling}, we adopt a geometrical viewpoint centered on fractal distributions and, in particular, cut-out sets, which we consider as cosmic web models.  We generalize this construction to multifractals and thence we proceed to a general study of multifractal voids, which leads us to differentiate two types of voids according to their geometry (Sect.\\ \\ref{MF}). This general study of multifractal voids is combined in Sect.\\ref{sim-gal} with the results of Sect.~\\ref{correl} to characterize the voids in cosmological simulations and galaxy samples In particular, we deduce some properties of galaxy voids in Sect.\\ \\ref{bias}, assuming a multifractal model of galaxy biasing.  We summarize our results and present our conclusions in Sect.\\ \\ref{discuss}.  Finally, we include three appendices to deal with technical points. ",
        "conclusions": "\\label{discuss}  We have developed statistical methods for the study of cosmic voids and we have applied them to various distributions of observable objects.  The method of choice depends crucially on the number density $n$ of objects. If the density is sufficiently small for having a small number of objects in a volume in which the Universe is homogeneous, then the objects are weakly correlated and the analysis for Poisson voids in Sect.~\\ref{Poisson} is relevant. Of course, this case is the most amenable to analytical methods.  We have found the distribution of spherical voids and the volume of the largest void, and we have shown how to generalize these results to more general shapes of voids. Not surprisingly, to observe a really large void, namely, with volume $V \\gg 1/n$, one needs an exponentially large sample, with a number of objects $N_{\\rm t} \\sim e^{nV}$.  And this sample could only come from surveying volumes exponentially larger than the homogeneity volume.  The conclusion is that the voids observed in galaxy surveys, for example, are not due to Poissonian fluctuations. Indeed, galaxy samples have had for long time number densities such that there are many galaxies in a homogeneity volume.  The objects inside a homogeneity volume have correlations, which are necessary for the existence of voids but make the application of analytical methods more difficult.  For still moderate $n$, an expansion in powers of $n$ yields general results and, in particular, yields a formula for the distribution of voids, although the computations are hardly feasible. It is preferable to resort to significant models that are valid for any $n$. We have focused on a hierarchical Poisson model and on the lognormal model, which illustrate two different behaviors of the largest voids.  We define the hierarchical Poisson model by restricting the support of a Poisson distribution to random clusters. Thus, we construct a point distribution that saturates the general moment inequality ${\\bar\\mu}_k \\geq {{\\bar\\mu}_2}^{k-1}$. When ${\\bar\\mu}_2 \\gg 1$, the clusters are small and the distribution is strongly correlated, but then its large voids are just the voids in the random distribution of clusters. These voids are independent of the magnitude of the density of objects inside clusters. Therefore, the distribution of the volume of the largest void is of Fr\\'echet-Pareto type, in contrast with the pure Poisson case, in which the largest void vanishes when $n \\ra \\infty$ (Gumbel type).  It is natural to understand the hierarchical Poisson model as a real hierarchy of clusters of clusters in a self-similar pattern. Thus, it becomes a fractal model with hierarchical {\\em and} scaling correlation functions. For example, one-dimensional L\\'evy flights are constructed so as to have scaling voids that follow the Pareto distribution.  In contrast, the study of voids in higher-dimensional fractals demands a good deal of geometry. A convenient approach to fractal voids is the construction of cut-out fractals.  In particular, a fractal foam is a possible model of the sheet structure of the cosmic web.  Fractal foams have voids with simple shapes and scaling sizes. These properties of voids can be generalized to other monofractal distributions.  A monofractal distribution is characterized by its box-counting dimension $D_{\\rm b}$ and is essentially uniform throughout its support, like the hierarchical Poisson model. However, the adhesive gravitational dynamics gives rise to non-uniform distributions that contain filaments and nodes, in addition to sheets.  This motivates us to consider non-uniform fractal foam models and, ultimately, multifractal models of the cosmic web. A multifractal is characterized by a spectrum of dimensions, but its box-counting dimension still plays the most significant role in regard to voids. If $D_{\\rm b} < 3$, then the support of the multifractal has vanishing volume, namely, is itself a fractal set. Therefore, the distribution of voids is independent of the non-uniformity of the matter distribution inside its support. A more interesting type of multifractal voids arises in the case that $D_{\\rm b}=3$ and the support is the total volume.  Multifractals with support in a full volume have been called non-lacunar fractals by Mandelbrot \\cite{Mandel}.  Whether or not non-lacunar fractals have voids is a subtle question. Certainly, they have no voids with empty interior in which one can fit an empty sphere. However, they do have sets of mass depletions with large volumes in which the density vanishes. We have called them voids of second type. The lognormal model provides us with a good description of the statistics of this type of voids as they appear in a multifractal sample with number density $n$. We deduce, for example, that the volume of the largest spherical void is much larger than in a Poisson distribution but is still small in comparison with the homogeneity scale (its distribution is actually of Gumbel type).  On the other hand, these voids do not have scaling sizes.  A more geometrical picture of this type of voids is achieved by representing them as under-dense regions of a smoothed density field (Sect.~\\ref{excursion}). Thus, second type voids become geometrically similar to first type voids, although they necessarily contain mass. This mass is part of self-similar structures that can be observed as the smoothing length shrinks and the resolution increases.  Indeed, it has been observed by Gottl\\\"ober et al \\cite{Gott} in cosmological $N$-body simulations that a pattern of self-similar structure arises inside voids when the resolution increases.  Further statistical evidence supports that the voids in cosmological simulations belong to a non-lacunar fractal (Sect.~\\ref{sim}). Therefore, we can assert that dark matter voids are well described as second type voids in a non-lacunar fractal.  It is less certain that a similar description is appropriate for galaxy voids. Admitting that the distribution of galaxies is fractal in the relevant range of scales, we have examined the observational and theoretical arguments for lacunarity.  Scaling of galaxy voids implies lacunarity.  We have not found this scaling, neither in Ref.~\\cite{Gaite} nor in Sect.~\\ref{bias}, in spite of the favorable results of Tikhonov and Karachentsev \\cite{Tikho-Kara,Tikho,Tikho2}.  Another argument for lacunarity is the emptiness of voids.  Peebles argues that the formation of observable objects is suppressed in voids \\cite{Pee2}.  But it rather seems that the size of voids diminishes with the galaxy luminosities \\cite{Fur-Pir}. This is also predicted by a non-lacunar multifractal model of galaxy bias in which the mass (or luminosity) of galaxies is ruled by the local dimension as a measure of mass concentration (Sect.~\\ref{bias}). However, very weak concentrations may not form galaxies and remain inside voids as dark matter halos. This might produce a low lacunarity in the distribution of galaxies.  As regards the formation of voids, we have introduced multifractal models of the adhesion-like dynamics that gives rise to the cosmic web.  However, to obtain non-lacunar foams, we must modify slightly the construction introduced in Sect.~\\ref{MF_cutouts}; namely, we must require that the initial under-dense regions become depleted but {\\em not empty}.  This is necessary to obtain the proper cosmic web structure. However, the construction of a realistic model of the formation of voids along these lines is beyond the scope of this work.  In conclusion, a good description of cosmic voids and, in general, the structure of matter is provided by a non-lacunar multifractal, with dimension $D_{\\rm b} = 3$ (the dimension of non-fractal distributions) and without empty voids (first-type voids).  The distribution of galaxies is probably non-lacunar as well, but it could have a low lacunarity.  The peculiar structure of non-lacunar fractals can be partly responsible for the controversy about the application of fractal geometry to the distribution of galaxies. Indeed, that structure is very different from the structure of monofractals (or multifractals with $D_{\\rm b} < 3$) that has been normally assumed for the distribution of galaxies."
    },
    "0911/0911.4453_arXiv.txt": {
        "abstract": "We investigate the origin of the prompt and delayed emission observed in the short GRB 090510. We use the broad-band data to test whether the most popular theoretical models for gamma-ray burst emission can accommodate the observations for this burst. We first attempt to explain the soft-to-hard spectral evolution associated with the delayed onset of a GeV tail with the hypothesis that the prompt burst and the high energy tail both originate from a single process, namely synchrotron emission from internal shocks. Considerations on the compactness of the source imply that the high-energy tail should be produced in a late-emitted shell, characterized by a Lorentz factor greater than the one generating the prompt burst. However, in this hypothesis, the predicted evolution of the synchrotron peak frequency does not agree with the observed soft-to-hard evolution. Given the difficulties of a single-mechanism hypothesis, we test two alternative double-component scenarios. In the first, the prompt burst is explained as synchrotron radiation from internal shocks, and the high energy emission (up to about 1 s following the trigger) as internal shock synchrotron-self-Compton. In the second scenario, in view of its long duration ($\\sim 100$ s), the high energy tail is decoupled from the prompt burst and has an external shock origin. In this case, we show that a reasonable choice of parameters does indeed exist to accommodate the optical-to-GeV data, provided the Lorentz factor of the shocked shell is sufficiently high. Finally, we attempt to explain the chromatic break observed around $\\sim 10^3$ s with a structured jet model.  We find that this might be a viable explanation, and that it lowers the high value of the burst energy derived assuming isotropy, $\\sim 10^{53}$ erg, below $\\sim 10^{49}$ erg, more compatible with the energetics from a binary merger progenitor. ",
        "introduction": "Traditionally divided in ``long'' and ``short'' on the basis of their $\\gamma$-ray duration \\citep[longer or shorter than 2 s,][]{Kouvelioutou1993}, gamma-ray bursts (GRBs) are characterized by a prompt release of $\\gamma$- and X-ray photons, followed by a multi-wavelength afterglow \\citep{Costa1997} emission. The fireball model \\citep[e.g.][]{Meszaros1992,Sari1998} explains the GRB electromagnetic emission as a result of shock dissipation in a relativistic flow, or ``fireball'', taking place at distances greater than $10^{-5}$-$10^{-2}$ pc from the central source. Despite such a large distance, electromagnetic observations have revealed important clues about the progenitors, favoring two main models: the coalescence of a binary consisting of two neutron stars or a neutron star and a black hole for short GRBs; and the death of massive stars (collapsars) for long GRBs. Both of these are commonly believed to end in a BH-plus-torus system, where torus accretion powers the fireball jet.  In the internal-external shock scenario of the fireball model \\citep[see e.g.][]{Meszaros1993,Sari1998}, GRB prompt and afterglow emissions are thought to be produced by particles accelerated via shocks in an ultra-relativistic outflow released during the burst explosion. While the prompt emission is related to shocks developing in the ejecta (internal shocks, IS), the afterglow arises from the forward external shock (ES) propagating into the interstellar medium (ISM). Synchrotron and synchrotron-self-Compton (SSC) emission by the accelerated electrons are typically invoked as the main radiation mechanisms. The Fermi satellite\\footnote{http:$//$fermi.gsfc.nasa.gov} is currently bringing exciting new results, detecting high energy ($\\sim$ GeV) extended tails whose presence is particularly intriguing in the case of short GRBs \\citep[e.g.][]{Abdo2009,Giuliani2009,GCN8407,GCN9334}. Fermi observations of GRB 081024B and GRB 090510 clearly point to the existence of a longer-lasting high energy tail in the GeV range following the main event, motivating a deeper exploration and re-examination of the fireball physics and radiative mechanisms \\citep[e.g.][]{Asano2009,Corsi2009,Gao2009,Kumar2009,Zou2008}.  Here we study the conditions under which the high energy observations of GRB 090510 can be accommodated within the most popular theoretical models. This work is organized as follows. In Sec.~\\ref{sec1} we summarize the observations for GRB 090510. In Sec.~\\ref{sec2} we discuss in more detail the spectral properties of the emission during the first second, and their implications for the compactness of the source. We test whether the complex emission observed during the first 1 s can be attributed to a single component, specifically synchrotron emission from IS. After showing that a single mechanism does not offer a straightforward explanation, we test an IS synchrotron plus SSC scenario. In Sec.~\\ref{sec5} we examine separately the high energy emission in the context of the ES scenario. Finally, in Sec.~\\ref{sec6} we give our conclusions. Hereafter we adopt as $T_0$ the onset of the main GRB pulse which, as specified in the following section, is about $\\sim 0.5$ s after the precursor that triggered the Fermi/GBM. Also, hereafter $\\epsilon_e$ and $\\epsilon_B$ are the fraction of energy going into electrons and magnetic fields, respectively; $n$ is the ISM density in particles/cm$^3$;  $E_{iso}$ is the isotropic kinetic energy of the fireball; and $p$ is the power-law index of the electron energy distribution in the shock. ",
        "conclusions": "\\label{sec6} We have analyzed GRB 090510 in the context of the synchrotron IS and ES scenarios. We first attempted to explain the soft-to-hard spectral evolution associated to the delayed onset of a GeV tail with the hypothesis that both the prompt burst and the high-energy tail originate from synchrotron emission of electrons accelerated by IS. Considerations of the compactness of the source lead us to conclude that the high-energy tail should be produced in IS developing in a late-emitted shell, characterized by a Lorentz factor of the order of $\\Gamma \\sim 700$, greater than the one generating the prompt burst ($\\Gamma \\sim 200$). However, this condition on the Lorentz factor implies a hard-to-soft evolution of the peak frequency of the IS synchrotron component, which does not agree with the observed soft-to-hard evolution.  Given the difficulties of explaining the prompt and delayed high energy emission with a single mechanism (synchrotron emission from IS), we then tested two double-component scenarios. In the first, the emission observed during interval I is explained as synchrotron emission from IS, while the high-energy tail observed in interval II is explained as SSC emission from IS. In the second scenario, the high energy emission observed by the LAT is decoupled from the prompt burst, and has an external shock origin. This last scenario has the advantage of explaining in a simple way the smooth temporal behavior of the high-energy tail, up to 100 s after the burst, and the consistency of the broad-band SED observed at 100 s with a single spectral component. In the ES scenario, we show that a reasonable set of parameters does indeed exist to explain the optical-to-GeV observations of this burst, despite a high Lorentz factor being required to have the fireball entering the deceleration phase as early as $t_{obs}\\sim 0.3$ s (when the emission in the LAT is observed to peak).  The ES scenario thus seems to account more naturally for the observations, even if a more detailed modeling of the late-time chromatic break is required. We have suggested that a structured jet may indeed be a viable explanation of such a chromatic feature.  In conclusion, we stress that the high Lorentz factor implied by the ES scenario has some relevant consequences in relation to the physics of the central engine. The commonly accepted fireball model invokes a series of shells expanding outward with Lorentz factor of the order of a few hundred, where IS first generate the prompt emission, and then, with the merged shell continuing to expand outward toward the external medium, an ES generates the afterglow. If, on the other hand, the high-energy tail is attributed to the ES emission, this implies that the source emits first a very fast shell, which impacts on the external medium creating an early afterglow, plus a series of slower shells that catch up with each other generating the prompt $\\gamma$-ray emission. At the end of the IS phase, the merged, slower shell (with a more typical Lorentz factor of the order of a few hundred), would also decelerate and eventually generate an afterglow by interaction with the external medium. In this respect, we note that the additional component required to explain the shallow decay observed at late times in the optical band could be related to the ES generated by such a slower shell."
    },
    "0911/0911.1795_arXiv.txt": {
        "abstract": "We investigate the response of the star formation efficiency (SFE) to the main parameters of simulations of molecular cloud formation by the collision of warm diffuse medium (WNM) cylindrical streams, neglecting stellar feedback and magnetic fields. The parameters we vary are the Mach number of the inflow velocity of the streams, $\\Msinf$, the rms Mach number, $\\Mrms$ of the initial background turbulence in the WNM, and the total mass contained in the colliding gas streams, $\\Minf$. Because the SFE is a function of time, we define two estimators for it, the ``absolute'' SFE, measured at $t = 25$ Myr into the simulation's evolution (\\sfeabs), and the ``relative'' SFE, measured 5 Myr after the onset of star formation in each simulation (\\sferel). The latter is close to the ``star formation rate per free-fall time'' for gas at $ n = 100 \\pcc$. We find that both estimators decrease with increasing $\\Minf$, although by no more than a factor of 2 as $\\Msinf$ increases from 1.25 to 3.5. Increasing levels of background turbulence (injected at scales comparable to the streams' transverse radius) similarly reduce the SFE, because the turbulence disrupts the coherence of the colliding streams, fragmenting the cloud, and producing small-scale clumps scattered through the numerical box, which have low SFEs. Finally, the SFE is very sensitive to the mass of the inflows (at roughly constant density and temperature), with \\sferel\\ decreasing from $\\sim 0.4$ to $\\sim 0.04$ as the mass in the colliding streams decreases from $\\sim 2.3 \\times 10^4 \\Msun$ to $\\sim 600 \\Msun$ or, equivalently, the virial parameter $\\alpha$ increases from $\\sim 0.15$ to $\\sim 1.5$. This trend is in partial agreement with the prediction by \\citet{KM05}, since the latter lies within the same range as the observed efficiencies, but with a significantly shallower slope. We conclude that the observed variability of the SFE is a highly sensitive function of the parameters of the cloud formation process, and may be the cause of significant scatter in observational determinations. ",
        "introduction": "\\label{sec:Intro}  The control of the star formation efficiency (SFE) by turbulence is a central issue in our present understanding of star formation, and currently a topic of intense study \\citep[see, e.g., the reviews by][]{MK04, MO07}. In recent years, several groups have studied the SFE of molecular clouds (MCs) using numerical simulations of isothermal turbulence, in which the entire numerical box represents the interior of a molecular cloud \\citep[see, e.g., the reviews by][]{MK04, BP_etal07, VS07}. One confusing issue is that simulations of {\\it driven} turbulence seem to indicate that the SFE {\\it decreases} as the turbulent rms Mach number $\\Ms$ increases \\citep[e.g.,][]{KHM00, VBK03, VKB05}, while simulations of {\\it decaying} turbulence suggest that the SFE {\\it increases} with increasing $\\Ms$ \\citep{NL05}. This has prompted the question of how does turbulence actually originate and behave in real MCs. To answer this question, it has become necessary to investigate the entire evolutionary process of MCs.  The formation of MCs by collisions of warm neutral medium (WNM) streams has been intensely studied in recent years. A vast body of numerical simulations has shown that moderate, transonic compressions in the WNM can nonlinearly trigger a phase transition to the cold neutral medium (CNM) \\citep[e.g., ][]{HP99, KI00, KI02, WF00}, and that the dense gas produced by this mechanism is overpressured with respect to the mean WNM thermal pressure \\citep{VS_etal06} and turbulent, due to the combined action of Kelvin-Helmholz, thermal \\citep{Field65} and nonlinear thin-shell \\citep{Vishniac94} instabilities \\citep{Heitsch_etal05, Heitsch_etal06}. The turbulence produced by this mechanism is continually driven for as long as the compression lasts.  The physical scenario of MC evolution was outlined by \\citet{HBB01}, who estimated the column densities necessary for the cloud to become self-gravitating, molecular, and magnetically supercritical, finding them to be comparable. The one-dimensional physical conditions in the dense gas were calculated analytically by \\citet{HP99} and \\citet{VS_etal06}. A recent review of the subject has been presented by \\citet{HMV08}.  More recently, simulations including self-gravity and ``sink particles'', which represent gravitationally collapsed objects (stars or stellar clusters), and using finite-duration compressions in the WNM, although lacking stellar feedback and magnetic fields, have been used to study the evolution of the turbulent motions and of the SFE in a self-consistent manner \\citep[][hereafter Paper I]{VS_etal07}. Concerning the velocity dispersion in the clouds, these authors found that the turbulence is intermediate between driven and decaying, since what decays is the driving rate of the turbulence as the inflows weaken with time. However, they also found that the random motions are gradually replaced by global infall motions, as the cloud begins to contract gravitationally. Concerning the SFE, Paper I measured the masses of dense gas $\\Mdense$ and of the collapsed stellar objects $\\Msinks$ in the simulations, allowing a measurement of the SFE, defined as \\begin{equation} \\hbox{SFE} = \\frac{\\Msinks}{\\Mdense + \\Msinks}. \\label{eq:SFE_def} \\end{equation} The resulting SFE was however too high, reaching $\\sim 50$\\% roughly 6 Myr after the time at which star formation (SF) had begun (denoted $\\tsf$), although this excessive SFE can possibly be attributed to the neglect of stellar feedback in that simulation. Indeed, Paper I estimated, using a prescription by \\citet{FST94} and a standard IMF \\citep{Kroupa01}, that by 3 Myr after $\\tsf$, enough massive stars would have formed as to be able to destroy the cloud by ionization. At that point, the SFE was $\\sim 15$\\%, closer to the typical values $\\la 5$\\% reported observationally for full MC complexes \\citep{Myers_etal86}.  One obvious possible reason for the relatively high values of the SFE obtained in this type of simulations is the neglect of stellar feedback and magnetic fields. However, it is also of interest to perform a more systematic investigation of the degree of variability that the SFE could exhibit within the scenario of Paper I, simply by varying the parameters of the WNM compressions triggering the formation of the cloud complex. In this paper we undertake a first approach to such task, by varying three parameters of the WNM stream collisions modeled in the simulations. First, we consider the inflow speed $\\vinf$ and the velocity dispersion of the background turbulence initially present in the medium, both measured by their respective Mach numbers, $\\Msinf$ and $\\Mrms$, with respect to the unperturbed WNM. Subsequently, we consider the mass in the colliding streams $\\Minf$, as determined by their radius $\\Rinf$ and length $\\linf$. Since the parameter space covered by these three parameters is already quite large, in this work we do not consider variations in the collision angle of the streams, instead having them collide head-on in all cases.  The paper is organized as follows. In \\S \\ref{sec:the_model} we describe the numerical model and experiments. In \\S \\ref{sec:results} we present our results, and in \\S \\ref{sec:conclusions} we present a summary and our conclusions. ",
        "conclusions": "\\label{sec:conclusions}  In this paper, we have considered the scenario of molecular cloud formation by WNM stream collisions, and investigated the dependence of the SFE on three parameters of this scenario, namely the inflow speed, the rms Mach number of the background medium, and the total mass contained in the inflows. Since the SFE, defined as in eq.\\ (\\ref{eq:SFE_def}), is a time-dependent function because the cloud continues to accrete mass from the WNM while it forms stars, we have considered two estimators of its time integral, namely the absolute SFE after 25 Myr from the start of the simulation, \\sfeabs, and the ``relative'' SFE, 5 Myr after the onset of SF in the cloud, denoted \\sferel.  We have found a wide range of values of these estimators as we vary the parameters of the simulations. In general, the SFE decreases, although moderately, with increasing inflow velocity. In particular, \\sferel\\ decreases from $\\sim 0.4$ to $\\sim 0.2$ as $\\Msinf$ increases from 1.25 to 3.5. Note that the runs in this case have similar SFRs, and the decrease in the SFE and \\dSFE\\ is due to the faster increase in cloud mass rather than to a decrease in the SFR. That the SFR is similar in all three runs, in spite of the larger gas mass is probably due to the larger turbulent velocity dispersion in the dense gas caused by the larger inflow speed, which tends to inhibit the SFR, thus compensating the tendency to have a larger SFR due to the larger cloud masses.  Similarly, the SFE decreases with increasing background turbulence strength, as the latter progressively takes a dominant role in the production of the dense gas but, due to the relatively small scales at which the turbulence is excited, the clouds and clumps formed by it are significantly smaller than the cloud formed by the coherent stream collision. In this case, \\sferel\\ decreases from $\\sim 0.4$ to $\\sim 0.03$ as the rms Mach number of the background turbulence increases from $\\sim 0.02$ to $\\sim 0.3$.  Finally, the SFE in general decreases with decreasing mass of the inflows (at constant $\\vinf$). This may be a consequence of the fact that clouds formed by the collision of our inflows always have roughly the same density, temperature, and velocity dispersion, so smaller clouds are more weakly gravitationally bound, a condition known to decrease the SFE \\citep{Clark_etal05}. The end result is that \\sferel\\ decreases from $\\sim 0.4$ to $\\sim 0.03$ as the mass in the colliding streams decreases from $\\sim 2.3 \\times 10^4 \\Msun$ to $\\sim 600 \\Msun$. It is important to stress that this result is not at odds with the well known fact that the SFE {\\it increases} as the object mass decreases from the mass scale of a giant molecular cloud (${\\cal M} \\sim 10^4$ -- $10^6 \\Msun$, SFE $\\sim 0.02$; Myers et al.\\ 1986) to that of a cluster-forming core (${\\cal M} \\sim 10^3 \\Msun$; SFE $\\sim 0.3$ -- 0.5, Lada \\& Lada 2003), because in this case the cores' mean densities are much larger than those of the GMCs, while in our case the mean densities of the various clouds are always comparable. Thus, our clouds do not conform to Larson's (1981) density-size scaling.  The latter results, expressed in terms of the virial parameter $\\alpha$, exhibit partial agreement with the prediction by KM05 for the dependence of the SFE after a free-fall time (what those authors called the star formation rate per free-fall time, or \\sfrff), which is directly comparable to our \\sferel. Although their prediction, without any rescaling, lies in the same range of values as our observed efficiencies, it contains a much shallower dependence on $\\alpha$ than we observe. This may be due to the fact that those authors assumed that the clouds were supported by turbulent pressure, possibly provided by stellar feedback, while our simulations lack such support. However, the notion of turbulent support has been questioned recently by \\citet{VS_etal08}, and the low observed efficiency may be the result of the {\\it dispersal} (rather than support) of the parent cloud by its stellar products \\citep{HBB01}. More work is needed to determine the causes of the discrepancy. However, it is noteworthy that Run Mi1-Mb.06-Ma2e3, which has a value of $\\alpha$ closest to unity, has a reasonably realistic value of \\sferel$ \\sim 4$\\%.  We conclude that the SFE, even in the absence of further agents such as magnetic fields and stellar feedback, is a highly sensitive function of the parameters of the cloud formation process, and may be responsible for significant intrinsic scatter in observational determinations of the SFE."
    },
    "0911/0911.4386_arXiv.txt": {
        "abstract": "We present a study on the environments of the SDSS galaxies divided into fine classes based on their morphology, colour and spectral features. The SDSS galaxies are classified into early-type and late-type; red and blue; passive, H{\\protect\\scriptsize II}, Seyfert and LINER, which returns a total of 16 fine classes of galaxies. We estimate the local number density, target-excluded local luminosity density, local colour, close pair fraction and the luminosity and colour of the brightest neighbour, which are compared between the fine classes comprehensively. The morphology-colour class of galaxies strongly depends on the local density, with the approximate order of high-density preference: red early-type galaxies (REGs) -- red late-type galaxies (RLGs) -- blue early-type galaxies (BEGs) -- blue late-type galaxies (BLGs). We find that high-density environments (like cluster environments) seem to suppress AGN activity. The pair fraction of H{\\protect\\scriptsize II} REGs does not show statistically significant difference from that of passive REGs, while the pair fraction of H{\\protect\\scriptsize II} BLGs is smaller than that of non-H{\\protect\\scriptsize II} BLGs. H{\\protect\\scriptsize II} BLGs show obvious double (red + blue) peaks in the distribution of the brightest neighbour colour, while red galaxies show a single red peak. The brightest neighbours of Seyfert BLGs tend to be blue, while those of LINER BLGs tend to be red, which implies that the difference between Seyfert and LINER may be related to the pair interaction. Other various environments of the fine classes are investigated, and their implication on galaxy evolution is discussed. ",
        "introduction": "Galaxies in the Universe have a wide range of physical properties such as morphology, spectra, mass, and so on. One of the major goals of the extragalactic astronomy is to understand what determines such properties of galaxies, and the first step toward the goal is to understand exactly how different the various galaxies are. One efficient way to comprehend the nature of the diverse galaxies is to classify them using some criteria and to compare their statistical properties between the classes. In many previous studies, however, galaxies have been classified using simple schemes based on just one or two properties, which sometimes miss detailed aspects of galaxy evolution. For example, several unusual classes of galaxies such as blue early-type galaxies \\citep{fer05,lee06} and passive spiral galaxies \\citep{yam04,ish07}, or the fundamental parameter difference between finely-divided classes in the same morphology class \\citep{lee07,cho09a} are difficult to investigate using simple classification schemes. Recently, \\citet[][hereafter Paper I and Paper II, respectively]{lee08,lee09} presented studies on the optical and multi-wavelength properties of the Sloan Digital Sky Survey \\citep[SDSS;][]{yor00} galaxies divided into fine classes based on their morphology, colour and spectral features. Those finely-classified SDSS galaxies show distinguishable differences from each other in their photometric (both optical and multi-wavelength) and structural properties, which are hardly found in the studies using simple classification schemes. Based on the understanding of the basic properties of the various galaxy classes presented in Paper I and Paper II, now it is needed to approach the next question: what makes them have such diverse natures?  For a long time, the environments of galaxies have been regarded as important factors determining the nature of galaxies. Since close galaxies have influence on each other directly or indirectly, the local density about a galaxy is closely related to the properties of that galaxy \\citep{kue05,par07}. Environments in which galaxies have many chances to interact with each other, induce morphology transformation by the tidal interaction or merging of galaxies \\citep{spi51,mer83,coz07}. Such galaxy interactions convert the properties of galaxies, by triggering star formation \\citep{woo06} or by suppressing star formation with tidal gas stripping \\citep{mar03}. Moreover, the complete merging of two or more galaxies is a classical candidate of the early-type galaxy formation mechanism \\citep{too77,sea78}. In addition to the direct interaction between galaxies, dense environments affect galaxies with their strong gravitational potential \\citep{bek99}. Since large amount of hot gas is concentrated in high-density environments \\citep[e.g. galaxy clusters;][]{chu96}, the morphology transformation of galaxies happens in such environments by the \\emph{ram pressure stripping} \\citep{gun72,qui00,roe05}.  Due to the close relationships between environments and galaxy properties, many studies have been focused on the environmental effects in galaxy evolution. \\citet{bla05} inspected the environmental dependence of optical properties of the SDSS galaxies, finding that the colour is the most sensitive parameter to the local density, and that the other parameters are not sensitive to environments when the colour and luminosity are fixed. \\citet{kue05} found that the surface brightness profile type and axis ratio of galaxies are frequently correlated with the local density. \\citet{ber06} showed that the age, metallicity and $\\alpha$-element enhancement of galaxies are affected by the environments of galaxies. \\citet{mat07} argued that galaxy evolution is accelerated in dense environments, which took place mainly at high redshifts, based on the age -- density relations with respect to galaxy luminosity or mass. \\citet{rog07} investigated the star formation history of early-type galaxies, showing that a significant number of galaxies in high-density environments undergo low but detectable recent star formation. \\citet{owe07} inspected the environmental dependence of starburst galaxies, finding that the starburst in a galaxy is triggered when its nearby (within 20 kpc) neighbour has similar luminosity or mass. \\citet{wes07} presented that most AGN host galaxies are closer to the red sequence than the blue sequence, but they prefer lower density environments than red sequence galaxies. \\citet{par07} inspected the environmental dependence of physical properties of the SDSS galaxies, based on the local density estimated using an adaptive smoothing kernel, showing that variations of galaxy properties are almost entirely due to the environmental dependence of morphology and luminosity, and that other properties are almost independent of local density, when morphology and luminosity are fixed. \\citet{van08} found that the role of satellite galaxies in the transformation from blue sequence galaxies to red sequence galaxies is important, in the sense that the satellite galaxies may quench the star formation of their host galaxies. \\citet{par08} and \\citet{par09a} showed that the morphology of a galaxy depends sensitively on the distance and morphology of its nearest neighbour galaxy, and that galaxy properties in the general field do not directly depend much on the large-scale background density. They showed that the apparent dependence of galaxy properties on the background density environment can be largely explained by an accumulated result of the interactions between the nearest neighbour galaxies.  Since it is obvious that environments affect galaxy properties significantly, environments must have played an important role in determining the fine classes (i.e. morphology, colour, and spectral features) of galaxies as well. Thus, to understand the origin of the various fine galaxy classes, it is needed to compare their environments. This is the third in the series papers on the nature of galaxies in various classes based on the morphology, colour and spectral features. In this paper, we investigate the environments of the finely-divided galaxy classes, using several environmental parameters. The outline of this paper is as follows. Section 2 shows the data set we used, and \\S3 describes our classification scheme of the SDSS galaxies and the definition of the environmental parameters. In \\S4, we present the results, and their implication on galaxy evolution are discussed in \\S5. The conclusions in this study are summarised in \\S6. Throughout this paper, we adopt the cosmological parameters $h=0.7$, $\\Omega_{\\Lambda}=0.7$, and $\\Omega_{M}=0.3$. ",
        "conclusions": "Here, we summarise the new findings in this paper and their implication on galaxy evolution.  \\begin{itemize} \\item[(1)] The morphology-colour class of galaxies strongly depends on the local density. The approximate order of high-density preference is REGs -- RLGs -- BEGs -- BLGs. \\end{itemize}  \\begin{itemize} \\item[(2)] In low-density environments, the fraction of AGN host galaxies increases as the local density increases, but it decreases in high-density environments that approximately correspond to cluster environments. This indicates that high-density environments (like cluster environments) suppress AGN activity. \\end{itemize}  \\begin{itemize} \\item[(3)] The local luminosity density distribution of $h$REGs is almost the same as that of $p$REGs. It is an open question why the AGN-suppressing mechanisms at high-density environments do not suppress the star formation in $h$REGs. \\end{itemize}  \\begin{itemize} \\item[(4)] The non-AGN BEGs prefer intermediate-density environments, while the AGN BEGs prefer low-density environments like BLGs, although with low % significance. In the context of early-type formation, such preference for low-density environments of BEGs seems to reflect the later formation of early-type galaxies at lower-density environments. \\end{itemize}  \\begin{itemize} \\item[(5)] $p$RLGs that have properties very similar to those of REGs (Paper I and II) have environments similar to those of $p$REGs, which confirms the idea that $p$RLGs are intermediate objects between early-type galaxies and late-type galaxies as suggested in Paper I and II. \\end{itemize}  \\begin{itemize} \\item[(6)] The pair fraction of $h$REGs does not show statistically significant difference from that of $p$REGs ($1.4\\sigma$ significance), while the pair fraction of $h$BLGs is smaller than that of non-H{\\protect\\scriptsize II} BLGs (significant up to $2.4\\sigma$), indicating that the close pair interaction is less important for the star formation of BLGs compared to that of REGs. \\end{itemize}  \\begin{itemize} \\item[(7)] Red galaxies show a single red peak in the distribution of the brightest neighbour colour, whereas $h$BLGs and $s$BLGs show double (red + blue) peaks. The brightest neighbours of $s$BLGs tend to be blue, while those of $l$BLGs tend to be red, which shows the possible relationship between AGN types and close pair interactions. \\end{itemize}  \\begin{itemize} \\item[(8)] Some ROSAT-detected passive galaxies in high-density environments may be affected by galaxy cluster environments. However, about 28 per cent (15 of 54) of the ROSAT-detected passive galaxies reside in low-density environments, which are classified as X-ray Bright Optically Normal Galaxies (XBONGs). \\end{itemize}"
    },
    "0911/0911.0793_arXiv.txt": {
        "abstract": "{globular clusters: general, $N$-body simulations, stellar dynamics} Dynamical evolution plays a key role in shaping the current properties of star clusters and star cluster systems. A detailed understanding of the effects of evolutionary processes is essential to be able to disentangle the properties which result from dynamical evolution from those imprinted at the time of cluster formation.  In this review, we focus our attention on globular clusters and review the main physical ingredients driving their early and long-term evolution, describe the possible evolutionary routes and show how cluster structure and stellar content are affected by dynamical evolution. ",
        "introduction": "\\label{sec:intro}  Globular star clusters have long been considered the ideal astrophysical objects to explore many aspects of stellar dynamics and, in particular, to study the evolution of stellar systems as driven by the effects of two-body relaxation.  Only recently, however, the actual complexity of globular cluster dynamics has emerged from the wealth of new observational and theoretical studies that have clearly shown the close interplay between stellar dynamics, stellar evolution, the clusters' stellar content and the dynamics and properties of the host galaxy.  The notion of an `ecology of star clusters', introduced by Heggie (1992), nicely illustrates the close interplay among the different elements of the astrophysics of star clusters.  To use star clusters as tools to advance and guide our understanding of star formation in galaxies, it is essential to understand to what extent the current properties of star clusters and star cluster systems are the result of evolutionary processes. We also need to be able to discern the signatures of properties imprinted by formation processes from those determined by dynamical evolution. In this paper, we will review the main physical ingredients driving the early and long-term evolution of clusters and describe how cluster structure and stellar content is affected by dynamical evolution.  In \\S\\ref{sec:early}, we discuss the early evolution of clusters, i.e., the processes that can affect their stellar content early in their life and lead to their rapid dissolution.  In \\S\\ref{sec:longterm}, we review the long-term evolution of clusters as driven, primarily, by two-body relaxation, focusing our attention on the evolution of cluster structure, mass and stellar content. Conclusions are summarized in \\S\\ref{sec:concl}. ",
        "conclusions": "\\label{sec:concl}  The study of star cluster dynamics has a long history and significant progress has been made in understanding the physics of the main ingredients driving the dynamical evolution of star clusters.  Many fundamental problems, however, are still open and new observations continuously present us with new challenges and questions.  Understanding the role played by early evolutionary processes in the dissolution of star clusters and in shaping the properties of those which survive, how to identify a cluster's dynamical state, what energy sources support clusters in the PCC phase and what the role is of evolutionary processes in determining the current properties of globular cluster systems are just a few of the issues that are still awaiting firm solutions.  As the close interplay between dynamical evolution, the evolution of a cluster's stellar content and the role played by the host galaxy's environment has become clear, new and increasingly complex questions crossing the boundaries of different fields have arisen. Understanding the links between the formation and abundance of exotic objects (such as blue stragglers, X-ray sources and IMBHs) and a cluster's dynamical history, the relationship between a cluster's stellar mass function, its structure and its past dynamical evolution, the role played by the host galaxy's tidal field and the effect of its time variation during galaxy assembly in shaping the current properties of individual clusters and the global properties of cluster populations are some examples of unsolved problems that require increasingly sophisticated models. The observational evidence of the presence of multiple stellar populations in globular clusters is the latest major challenge and further illustrates the complexity of the formation and dynamical evolution of star clusters.  As clusters evolve, their structural properties and stellar content are modified by evolutionary processes. Clusters in different galaxies and at different galactocentric distances have different dynamical histories.  A better understanding of the effects of dynamical evolution is an essential step in our attempts to shed light on the relationship between current star cluster properties and those which were imprinted by star-formation processes."
    },
    "0911/0911.2796_arXiv.txt": {
        "abstract": "We present the instrumentation, the target selection method, the data analysis pipeline and the preliminary results of the \\textit{Thessaloniki Research for Transits} project (ThReT). \\textit{ThReT} is a new project aiming to discover Hot Jupiter Planets, orbiting Sun-like stars. In order to locate the promising spots for observations on the celestial sphere, we produced a sky-map of the Transit Detection Probability by employing data from the Tycho Catalogue and applying several astrophysical and empirical relationships. For the data reduction we used the ThReT pipeline, developed by our team for this specific purpose. ",
        "introduction": "ThReT is a new photometric survey initiated at the Aristotle University\u2018s facilities at Mt. Holomon, Chalkidiki. The latter offers sub-arcsecond seeing (mean value 0.82\u201d) conditions and limited light pollution skies. The employed instruments for the survey are a Celestron11\u201d f/6.3 telescope and a Fingerlakes PL6303E front illuminated CCD camera, mounted on a Synta EQ-6 german mount and guided through a smaller Konus 4\u201d reflector. The achieved pixel scale is 1.32 arcsec and the field of view is 58x33 arcmin. We use no optical filter for our observations. ",
        "conclusions": ""
    },
    "0911/0911.5455_arXiv.txt": {
        "abstract": "\\noindent  Recent measurements of the surface magnetic fields of classical T~Tauri stars (CTTSs) and magnetic cataclysmic variables show that their magnetic fields have a complex structure. The magnetic field associated with the octupole moment may dominate the magnetic field associated with other moments in some stars, such as the CTTS V2129~Oph. Previously, we studied disc accretion onto stars with magnetic fields described by a superposition of aligned or misaligned dipole and quadrupole moments. In this paper, we present results of the first simulations of disc accretion onto stars with an \\textit {octupole} field. As examples, we consider stars with a superposition of octupole and dipole fields of different strengths and investigate matter flow around them, the shapes of hot spots on their surfaces, and the light curves produced by their rotation. We investigate two possible mechanisms for producing phase shifts in the light curves of stars with complex fields: (1) change of the star's intrinsic magnetic field and (2) variation of the accretion rate, causing the disc to interact with the magnetic fields associated with different stellar magnetic moments. We also discuss the applicability of the potential approximation for extrapolation of the surface magnetic field to larger distances. ",
        "introduction": "Strong magnetic fields have been observed in a number of stars, including classical T Tauri stars (CTTSs) (Basri, Marcy \\& Valenti 1992; Johns-Krull 2007), magnetic white dwarfs (e.g., Warner 1995; Euchner et al.\\ 2002), and various types of neutron stars (e.g., Ghosh \\& Lamb 1978; Ghosh 2007). In most stars the structure of the field is unknown.  It is usually suggested that the magnetic fields of gaseous stars are generated and supported by some type of the dynamo mechanism that operates in the stellar interior or in the surface layers of the star. Numerical simulations of uniformly rotating, fully convective stars show the formation of complex, non-axisymmetric fields that can be represented only by a superposition of different multipoles (e.g., Chabrier \\& K\\\"uker 2006). However, in simulations of differentially rotating stars, an additional, ordered component appears (e.g., Dobler, Stix \\& Brandenburg 2006).  Measurements of the surface magnetic fields of CTTSs using different techniques indicate that they have a complex structure (Johns-Krull, Valenti \\& Koresko 1999; Johns-Krull 2007). Measurements of the magnetic fields of nearby low-mass stars using the Zeeman-Doppler technique show that their fields are often complex (Donati \\& Cameron 1997; Donati et al.\\ 1999; Jardine et al.\\ 2002). Recent observations of two CTTSs have shown that in one star (V2129 Oph) the surface magnetic field associated with the octupole moment dominates the fields associated with other moments (Donati et al.\\ 2007), whereas in the other star (BP Tau) the fields associated with the dipole and octupole moments are both significant and dominate the fields associated with other multipoles (Donati et al.\\ 2008).  Recently, the magnetic field structure of gaseous stars with complex fields has been studied analytically. It was found that the fraction of the flux through the stellar surface in open field lines is smaller in stars with complex fields than in those with purely dipolar fields (Gregory et al.\\ 2008; Mohanty \\& Shu 2008). The potential approximation is usually used to extrapolate the surface magnetic field to larger distances. Gregory et al.\\ (2006) calculated possible gas flows around stars with magnetic fields constructed from measurements using the potential approximation. Recently, we were able to compute the flow of gas around V2129~Oph and BP~Tau in global 3D MHD simulations in which the magnetic field is not assumed to be given by a potential but is instead calculated as part of the simulation (Long et al.\\ 2009).  Zeeman tomography of magnetic white dwarfs has shown that they also have complex magnetic fields (Euchner et al.\\ 2002). In some, such as HE~1045$-$0908, the magnetic field associated with the star's quadrupole moment dominates whereas the fields associated with its dipole and octupole moments are much weaker (Euchner et al.\\ 2005). In others, representation of the field requires inclusion of misaligned dipole, quadrupole, octupole, and other multipoles up to $n=4$ or $n=5$ (Euchner et al.\\ 2006; Beuermann et al.\\ 2007). These results are in accord with earlier indications of field complexity, such as asymmetric distribution of spots on the stellar surface (Meggitt \\& Wickramasinghe 1989; Piirola et al.\\ 1987; see Wickramasinghe \\& Ferrario 2000 for a review).  Neutron stars also have dynamically important magnetic fields. Soon after they have formed, the magnetic fields of some may be enhanced by a dynamo mechanism (e.g., Thompson \\& Duncan 1993) and may have a complex structure. There is evidence that at least some accretion-powered X-ray pulsars have complex magnetic field structures (Elsner \\& Lamb 1976; Gil et al.\\ 2002; Nishimura 2005) or that the dipole moment is off-center (e.g., Coburn 2001;  see also Ruderman 1991; Chen \\& Ruderman 1993). Accreting millisecond X-ray pulsars have relatively weak dipole magnetic fields (e.g. Psaltis \\& Chakrabarty 1999), but the details of their fields are unknown. They have almost sinusoidal light curves that are consistent with a dipole magnetic field with a small inclination to the spin axis (e.g., Gierli\\'nski \\& Poutanen 2005; Lamb et al.\\ 2008a,b). However, this does not exclude the presence of higher-order multipole fields near the star. An unusual phenomenon --- rapid, relatively large shifts in the phase of the light curve --- has been observed in many millisecond pulsars (e.g., Morgan et al.\\ 2003; Markwardt 2004; Burderi et al.\\ 2006; Hartman et al.\\ 2009; Patruno et al.\\ 2009). Several mechanisms have been proposed to explain this phenomenon.  Lamb et al.\\ (2008a,b) have explained several different properties of accreting millisecond pulsars, including the phase shifts of their light curves, with the ``nearly aligned moving spot model\". Motion of emitting areas located close to the spin axis alters both the shape and the arrival time of the pulse in a complex way. These movements could be caused by changes in the accretion rate. Earlier 3D MHD simulations of accretion onto stars with misaligned dipole magnetic fields (Romanova et al.\\ 2003, 2004) have shown that if the dipole is nearly aligned with the spin axis, the place where the accreting matter impacts the stellar surface can move prograde or retrograde relative to the surface, which favors the moving spot idea. In another model, the magnetic field of the neutron star is assumed to change in time (Burderi et al.\\ 2006). Here we consider stars with complex magnetic fields and investigate both of these mechanisms for producing phase shifts.  In our earlier studies, we were able to perform global 3D~MHD simulations of accretion onto stars with a dipole magnetic field (Koldoba et al.\\ 2002; Romanova et al.\\ 2003, 2004; Kulkarni \\& Romanova 2005) and onto stars with a combination of dipole and quadrupole fields (Long, Romanova \\& Lovelace 2007, 2008). 3D~simulations of accretion onto stars with multipolar fields are very challenging due to the very steep gradients of the magnetic field. In this paper, we report on the first simulations of accretion onto stars with an {\\it octupole} field component. Our methods allow us to investigate the general case of a magnetic field produced by superposition of dipole, quadrupole, and octupole fields oriented in different directions. However, even the simpler case of a magnetic field produced by superposition of misaligned dipole and quadrupole fields creates very complex field structures (e.g., Long et al.\\ 2008). Consequently, for clarity we omit quadrupole fields in the present work, restricting consideration to the simpler case of a superposition of octupole and dipole fields. We investigate matter flows onto stars with such fields, the shapes of the resulting hot spots, and the light curves produced by these hots spots. We also investigate mechanisms for producing phase shifts in the light curves of stars with complex magnetic fields. Finally, we discuss the validity of the potential approximation for computing the magnetic field outside such stars.  Section~2 describes the simulation method and the magnetic field configurations. Section~3 presents our results for different combinations of dipole and octupole fields. Section~4 shows examples of the phase shifts that are possible in the light curves of stars with complex fields. Section~5 summarizes and discusses our most important results. ",
        "conclusions": "We performed, for the first time, global 3D MHD simulations of accretion onto stars with predominantly octupolar magnetic fields, and for a combination of dipole and octupole fields. The calculation become possible due to enhancement of the grid resolution near the star. Simulations have shown that:  \\begin{enumerate}  \\item  If the disc interacts predominantly with the octupolar field, then matter flows into two octupolar rings on the surface of the star. In the case of an inclined octupole field, the hot spots are inclined and have a brightness amplification on one side of the ring.  \\item  At small misalignment angles of the octupole, $\\Theta=10^\\circ$, the light curves are quite sinusoidal and ``mimic'' the shape of the light curves of the pure dipole field, in particular at small inclination angles $i$. However, at higher $\\Theta$ (e.g., $\\Theta=20^\\circ$), the light curves strongly depart from sinusoidal \\textit{even for small $i$}.  \\item  If the dipole component strongly dominates, then matter flows in two funnel streams governed by the dipole field, and two round spots form in the vicinity of the dipole magnetic poles. If the dipole and octupole have comparable strength at the truncation radius, then both components channel matter to the star. Usually both octupolar rings and polar spots form on the surface of the star.   \\item  The potential field approximation is valid only in the region around the star where the magnetic stresses are higher than the matter stresses. At larger distances, the field is dragged by the disc and corona and strongly departs from being potential. The external magnetic field usually acquires a significant azimuthal component and inflates. A number of field lines wrap around the rotation axis, forming a magnetic tower which may propagate to larger distances due to magnetic force (e.g., Romanova et al.\\ 2009).  \\item  Accretion onto stars with multipolar fields may lead to phase shifts in the light curves. This may occur either (a) during dynamo-generated internal field reconstruction, or (b) during variation of the accretion rate when the complex magnetic field is fixed and has different phase angles $\\phi_n$ for each multipole component, but the disc interacts with multipoles of different orders at different accretion rates. The phase shifts would be expected to be accompanied by changes in the light curve pulse profile, because the shape and position of the hot spots is very different at different field configurations.  \\end{enumerate}   These new challenging 3D simulations helped us understand accretion onto stars with octupolar fields. Recent measurements by Donati et al.\\ (2007, 2008) of the two CTTSs V2129 Oph and BP Tau have shown that their fields have a significant octupolar component. In our next paper (Long et al.\\ 2009) we compare our numerical model with observations of these stars.   Magnetic fields in young stars of the CTTS type may have a complex structure and may vary with time due to internal dynamo processes. The recently measured magnetic fields of the accreting T Tauri stars CV Cha and CR Cha (Hussain et al.\\ 2009) show a complex structure consisting of a number of multipoles. In another CTTS, V2247 Oph, the magnetic field is complex and varies on the very short time-scale of a week (Donati et al.\\ 2009). Frequent phase shifts are expected in this star due to internal field reconstruction, as discussed in this paper.   In all these stars the phase shifts may also occur due to accretion rate variation.  A phase shift of 0.2 was recently observed in the millisecond pulsar SAX J1808.4-3658 on the 14th day of its outburst (Burderi et al.\\ 2006). This phase jump was observed only in the fundamental component of the almost sinusoidal light curve, and much less so in the 1st harmonic. Burderi et al. (2006) suggested that this phenomenon may be connected with fast spin-down of the millisecond pulsar or with some magnetic field reconstruction during a stage of enhanced accretion. We suggest that it raises the possibility that the magnetic field of the millisecond pulsar may be more complex than a dipole, and therefore during enhanced accretion the disc may interact with deeper layers of the magnetosphere where quadrupolar or other components might possibly influence the flow. We also suggest that the dipole may be slightly off-center (Ruderman 1991; Long et al.\\ 2008), and hence the disc matter accretes to both magnetic poles when the accretion rate is high, or to only one pole if it is low. Superposition of dipoles may also be a possible reason."
    },
    "0911/0911.0608_arXiv.txt": {
        "abstract": "{} {We investigate the emission mechanism and evolution of pulsars that are associated with supernova remnants.} {We used imaging techniques in both the optical and near infrared, using images with very good seeing ($\\leq$0\\myarcsec6) to study the immediate surroundings of the Crab pulsar. In the case of the infrared, we took two data sets with a time window of 75 days, to check for variability in the inner part of the Crab nebula. We also measure the spectral indices of all these wisps, the nearby knot, and the interwisp medium, using our optical and infrared data. We then compared the observational results with the existing theoretical models. } {We report variability in the three nearby wisps located to the northwest of the pulsar and also in a nearby anvil wisp in terms of their structure, position, and emissivity within the time window of 75 days. All the  wisps display red spectra with similar spectral indices ($\\alpha_\\nu = -0.58 \\pm 0.08$, $\\alpha_\\nu = -0.63 \\pm 0.07$, $\\alpha_\\nu = -0.53 \\pm 0.08$) for the northwest triplet. The anvil wisp (anvil wisp 1) has a spectral index of $\\alpha_\\nu = -0.62 \\pm 0.10$. Similarly, the interwisp medium regions also show red spectra similar to those of the wisps, with the spectral index being $\\alpha_\\nu = -0.61 \\pm 0.08$, $\\alpha_\\nu = -0.50 \\pm 0.10$, while the third interwisp region has a flatter spectrum with spectral $\\alpha_\\nu = -0.49 \\pm 0.10$. The inner knot has a spectral index of $\\alpha_\\nu = -0.63 \\pm 0.02$. Also, based on archival HST data and our IR data, we find that the inner knot remains stationary for a time period of 13.5 years. The projected average velocity relative to the pulsar for this period is $\\lsim 8 \\kms$.} {By comparing the spectral indices of the structures in the inner Crab with the current theoretical models, we find that the Del Zanna et al. (2006) model for the synchrotron emission fits our observations, although the spectral index is at the flatter end of their modelled spectra.} ",
        "introduction": "The Crab nebula is one of the most studied targets in the sky. It is a remnant from a supernova explosion that occurred in 1054. The Crab pulsar (m$_V$$\\sim$16), the first pulsar with optically detected pulses (Cocke et al. 1969), is located at the centre of the nebula, and is responsible for powering the nebula. The pulsar, together with the filaments and the continuum emitting pulsar-wind nebula (PWN), form the Crab nebula. The Crab nebula was the first detected astrophysical source of synchrotron radiation, as suggested by Shklovsky (1953) and was confirmed by polarisation observations a year later by Dombrovsky (1954). The inner part of the Crab nebula is a region consisting of jets, a torus of X-ray emission (Aschenbach \\& Brinkmann 1975), and complexes of sharp wisps (Scargle 1969). The structures give a unique opportunity to study the pulsar, along with its wind and the interaction it has with the rest of the remnant.  Evidence of activity in the Crab pulsar and its surroundings has existed for a few decades. Variability in the structures around the pulsar were mentioned by Lampland (1921) and by Oort \\& Walraven (1956). A very comprehensive study has been done by Scargle (1969), who found that the wisps may show relativistic motion, but conversely always seem to be more or less located at the same position. Using the WFPC2 onboard HST high resolution surveys were conducted by Hester et al. (1995; 2002). These data were complemented with X-ray data from ROSAT and Chandra to get a complete picture of the Crab pulsar and its immediate surroundings. The most outstanding discovery is the presence of two knots that are situated 0\\myarcsec65 and 3\\myarcsec8 southeast of the pulsar. The X-ray data have also shown many knots around the pulsar that show variability with time. For a complete look at the Crab pulsar and its nebula see Hester (2008).  Sollerman (2003) used ISAAC on the VLT, to conduct an IR survey on JHK${\\rm_s}$ bands in order to investigate the Crab pulsar wind nebula further. Using the VLT data, together with archive HST optical data he measured a spectral index for the knot equal to $\\alpha_{\\nu} = -0.8$. Melatos et al. (2005) went further into investigating the Crab PWN using IR images to check for short time-scale variations around the Crab pulsar. They found variations in the emission of the Crab wisps on a time scale of 1.2 ksec of $\\pm$24\\%  $\\pm$ 4\\% in the K band and $\\pm$14 $\\pm$5\\% in the J band. Apart from the variation estimations, Melatos et al. (2005) calculated the spectral indices from the structures near the pulsar, finding values for the spectral index that vary from $\\alpha_{\\nu} = -0.16$ for the rod (this the brightest feature in the shock located $\\sim$ 4\\myarcsecnodot in a southeastern orientation from the pulsar) down to $\\alpha_{\\nu} = - 0.76$ for the innermost knot.  Despite the amount of work that has been done on the Crab, there are still many unanswered questions regarding its nature and, in general, about the supernova explosion mechanism and how the pulsar emission mechanism operates and evolves. Also there is not any model that can fully explain the observed features of the emitted radiation, especially in the inner part of the PWN. In this article we present our near infrared and optical survey that we conducted using the Nordic Optical Telescope in La Palma, Spain, with a gap of two and a half months (in the case of the NIR) in order to check for variability in the enviroment around the Crab pulsar. In Sect. 2 we present our observing procedure followed by the data reduction process. In Sect. 3 we present our photometric results, and in Sect. 4 we present arguments for the nature of the wisps and the knots. We make our conclusions in Sect. 5.  \\begin{figure}[ht] \\includegraphics[width=1.0\\hsize]{Orient.ps} \\caption{I-band image of the Crab nebula taken with ALFOSC at the NOT on 7 December 2007. The pulsar has been removed with the PSF to reveal the nearby knot that lies at a distance of 0\\myarcsec65. The field of view is $\\sim$ 45$\\times$45 arcseconds, with north pointing upwards and east to the left. The areas that we used to measure the emission from the wisps and the interwisp regions are marked with circles.} \\end{figure} ",
        "conclusions": "We have presented data for the Crab pulsar in both the optical and infrared. Our infrared data suggest changes in the inner enviroment of the Crab for both dynamic structures and emissivity. Our estimated velocities for the wisps do not agree with currently known values (Hester et al. 2002). The reason for this difference can be the large time window we had. Ideally, the dynamic structure of the inner Crab should be studied using a daily based program. From our data we find that the inner knot does not show any variability in its emission and that it remains stationary over a period of 13.5 years (until 2007 December) relative to the pulsar; its average projected velocity relative to the pulsar within this time window is $\\lsim 8 \\kms$, corresponding to a relative space motion of $\\lsim 10 \\kms$, if the pulsar and knot are aligned with the jet axis.  The combination of the optical and NIR data allowed us to create spectra of the wisps and the knot and thus obtain hints to their nature. All the wisps (and the interwisp medium) display red spectra with similar indices ($\\alpha_{\\nu}$ that vary between $-0.49\\pm0.10$ and $-0.63\\pm0.07$), and the knot has $\\alpha_{\\nu}=-0.63\\pm0.02$. The latter seems to agree with the result of Sollerman (2003), who reported a spectral index for the knot equal to $\\alpha_{\\nu} = -0.8$ for a spectral range that covers both optical (U, V and I) from the HST and IR (J, H, K${\\rm_s}$) from the VLT. The spectral distributions we find for the wisps and the red knot are similar to what has been estimated for other PWNe, which suggests that these observed properties are the result of the same emission mechanism. The exception appears to be the PWN of PSR B0540-63.9, which for most of its PWN has a steeper spectrum than the other PWNe and which signals stronger cooling. A feature in the Crab PWN that sticks out is the red knot with its high degree of polarisation. It lies close to the pulsar and could be part of a shock structure shared with the sprite. Both the pulsar and the sprite have markedly flatter spectra than the knot though. On a somewhat larger scale, we find that our data are compatible with the model of Del Zanna et al. (2006), although our spectra are at the flatter end of the spectral index interval argued for by Del Zanna et al. The similarity of the spectral indices of the wisps with the interwisp medium can be an indication that the wisps are just enhancements of synchrotron emission, but synchrotron cooling instabilities should not be dismissed.  We highlighted that pulsar wind nebulae of Crab like remnants display similar red spectra, and observations have also revealed that in some cases they show similar structures to those of the Crab (e.g., Vela in X-rays). Further detailed comparisons are needed to constrain the physical processes giving rise to the observed properties. Because there are many open questions regarding the nature of the PWNe, it would be interesting to extend this comparison to related objects like Crab-like remnants with bow-shock pulsars such as the Guitar Nebula, in particular to study how age affects the properties of the pulsar wind nebulae."
    },
    "0911/0911.4392_arXiv.txt": {
        "abstract": "s{ The paper discusses the problem of the orientation of galaxies in groups in the Local Supercluster (LSC). The existence of the preferred orientation of galaxy group is shown. We found that the orientation of  galaxy groups in the Local Supercluster in the scale till about 20 Mpc is  strongly  correlated with the distribution of  neighbouring groups. The line joining the two brightest galaxies is in alignment  with  both the group major axes and the direction toward  the  centre  of  the  LSC, i.e. Virgo  cluster. These correlations suggest that two brightest galaxies were formed in filaments of matter directed towards  the  protosupercluster centre. Afterwards, the hierarchical clustering leads to aggregation of galaxies around these two galaxies. The groups are formed on the same or similarly oriented filaments. This picture is in agreement with the predictions of numerical simulations.} ",
        "introduction": "Starting from Binggeli paper \\cite{a1}, several authors studied the orientation of galaxy groups and clusters \\cite{a2,a3,a4,a5,a6,a7,a8,a9,a10}. On the basis of optical and X-ray data they found that structures exhibit a tendency to be orientated toward their neighbours. The interpretation of this effect has changed in the last thirty years, but the main idea that this should reflect  conditions during the structure formation is still very popular. Numerical simulations \\cite{a11,a12,a13,a14,a15,a16,a17,a18,a19,a20,a21} gave a better understanding of physical processes leading to structure formation. These simulations were performed in the framework of the cold dark matter(CDM) model, presently regarded as the correct description of the large scale structure formation.     Several  numerical simulations, using different approaches  and codes, led to the conclusion that the preferred orientation  of galaxy  clusters in the CDM model is a natural  consequence  of processes  leading to structure formation due to  gravitational interaction along filamentary structures. In order to  confirm, or  deny,  this  scheme of structure origin  we  carry  out  an analysis   of  the  Local  Supercluster  (LSC)  galaxy   groups alignment.  Groups  were  taken from the  Catalogue  of  Nearby Galaxies \\cite{a22}. We selected structures having at least 10 members. There are 61 such groups.  \\begin{figure} \\hskip 0.5cm \\psfig{figure=fig1a.eps,height=3.6in} \\psfig{figure=fig1b.eps,height=3.6in} \\caption{The distribution (from top to bottom) of the position angles $PA_g$, $PA_l$  $PA_V$ (left panel) and the distribution (from top to bottom) of the differences between position angles $PA_g-PA_V$, $PA_l-PA_V$, $PA_g-PA_l$ (right panel). \\label{fig1}} \\end{figure} ",
        "conclusions": "From the presented analysis of the orientation of galaxy groups in   the  Local  Supercluster  the  following  picture  of  the structure  formation appears. The two brightest  galaxies  were formed  first.  They  originated in the  filamentary  structure directed  towards the centre of the protocluster. This  is  the place where the Virgo cluster centre is located now.  Due  to gravitational clustering, the groups are formed in such a  manner that galaxies follow the line determined by  the  two brightest   objects.  Therefore,  the  alignment  of  structure position  angle  and  line joining two  brightest  galaxies  is observed.  The other groups are forming on the same  or  nearby filament.  The flatness of the LSC additionally contributes  to the  observed alignment of galaxy groups. The majority  of  the groups  lie  close  to  us. From the selection  effect  in  the catalogue  the  lack of groups further than the  Virgo  Cluster centres is observed.  This  picture is in agreement with predictions of  several  CDM models,  in  which structure formation is due  to  hierarchical clustering.   Moreover,  the  formation  is  occurring  on  the filamentary structure."
    },
    "0911/0911.1781_arXiv.txt": {
        "abstract": "We present radial velocity observations of four extremely low-mass ($0.2~M_\\odot$) white dwarfs. All four stars show peak-to-peak radial velocity variations of 540 -- 710 km s$^{-1}$ with 1.0 -- 5.9 hr periods. The optical photometry rules out main-sequence companions. In addition, no milli-second pulsar companions are detected in radio observations. Thus the invisible companions are most likely white dwarfs. Two of the systems are the shortest period binary white dwarfs yet discovered. Due to the loss of angular momentum through gravitational radiation, three of the systems will merge within 500 Myr. The remaining system will merge within a Hubble time. The mass functions for three of the systems imply companions more massive than $0.46~M_\\odot$; thus those are carbon/oxygen core white dwarfs. The unknown inclination angles prohibit a definitive conclusion about the future of these systems. However, the chance of a supernova Ia event is only 1\\% to 5\\%. These systems are likely to form single R Coronae Borealis stars, providing evidence for a white dwarf + white dwarf merger mechanism for these unusual objects. One of the systems, SDSS J105353.89+520031.0 has a 70\\% chance of having a low-mass white dwarf companion. This system will probably form a single helium-enriched subdwarf O star. All four white dwarf systems have unusal mass ratios of $\\leq0.2-0.8$ that may also lead to the formation of AM CVn systems. ",
        "introduction": "Mergers of binary white dwarfs (WDs) have been proposed to explain supernovae (SNe) Ia events, extreme helium stars including R Coronae Borealis (RCrB) stars, and single subdwarf B and O stars \\citep{iben84,webbink84,saio00,saio02,heber09}. However, radial velocity surveys of WDs have not revealed a large binary population that will merge within a Hubble time \\citep{marsh95,maxted00,napiwotzki01,napiwotzki02,karl03,nelemans05}. In addition to a few pre-WD + WD merger systems \\citep[e.g.][]{geier07, tovmassian10}, \\citet{napiwotzki04a} identify only eight merger candidates from the SN Ia Progenitor Survey (SPY) and the literature.  Radial velocity observations of extremely low-mass (ELM, $M\\sim0.2~M_\\odot$) WDs provide a new opportunity to find short period binaries. The Universe is not old enough to produce ELM WDs through single star evolution. These WDs must therefore undergo significant mass loss during their formation in binary systems. The majority of ELM WDs have been identified as companions to milli-second pulsars. However, not all ELM WDs have such companions \\citep{vanleeuwen07,agueros09a}. Radial velocity, radio, and x-ray observations of the lowest gravity WD found in the Sloan Digital Sky Survey (SDSS), J0917+4638, show that the companion is almost certainly another WD \\citep{kilic07a,kilic07b,agueros09b}.  Recently, \\citet{eisenstein06} identified a dozen ELM WDs in the SDSS Data Release 4 area. Here, we present radial velocity observations of four WDs from that sample; SDSS J082212.57+275307.4, SDSS J084910.13+044528.7, SDSS J105353.89+520031.0, and SDSS J143633.29+501026.8. Our observations are discussed in Section 2; an analysis of the spectroscopic data and the nature of the companions are discussed in Section 3. The future of these binary systems and the merger products are discussed in Section 4. ",
        "conclusions": "Almost all known double WD systems have mass ratios on the order of unity \\citep{nelemans05}. Recently, \\citet{kilic07b,kilic09} identified two WD binary systems with mass ratios $\\leq0.6$. However, these two systems (SDSS J0917 and LP400$-$22) will not merge within a Hubble time.  The four binary systems discussed in this paper will merge within a Hubble time. They have extreme mass ratios and are likely progenitors of RCrB stars, single He-enriched sdO stars, or AM CVn stars. These systems may even contribute to the SNe Ia population, but the probability of this event is small.  \\citet{liebert05} estimate that low-mass ($M<0.45 M_\\odot$) WDs make up about 10\\% of the local WD population, corresponding to a formation rate of 4 $\\times 10^{-14}$ pc$^{-3}$ yr$^{-1}$. \\citet{eisenstein06} discovered only 13 ELM WDs out of a sample of 9316 WDs found in a survey volume of $\\geq$ 4 kpc$^3$ in the SDSS Data Release 4 footprint. ELM WDs are rare, they make up about 0.14\\% of the local WD sample. Of course, the SDSS spectroscopic sample suffers from selection bias and incompleteness, but taken at face value, 0.14\\% corresponds to a formation rate of 6 $\\times 10^{-16}$ pc$^{-3}$ yr$^{-1}$. However, all four ELM WDs discussed in this paper are going to merge within a Hubble time, and some within the next few hundred Myr. Therefore, the formation rate of ELM WDs may be higher by an order of magnitude. For the SDSS DR4 survey volume of 4 kpc$^3$, the formation rate is $0.2-2 \\times 10^{-5}$ yr$^{-1}$, with the caveat that this number is highly uncertain due to selection bias present in the SDSS spectroscopic survey. \\citet{bogomazov09} estimate a CO+He WD merger rate of 6-8 $\\times 10^{-3}$ yr$^{-1}$. The ELM WD mergers discussed in this paper contribute only a small fraction to the overall CO+He WD merger rate in the Galaxy. \\citet{yungelson05} estimate a SNe Ia occurence rate of 10$^{-5}$ yr$^{-1}$ from CO+He WD progenitors. Our estimate of the formation rate of ELM WDs is similar to this result.  In this paper, we present radial velocity data for four stars. However, we are obtaining radial velocity observations for the remaining ELM WDs in the \\citet{eisenstein06} sample. To date, almost all of these ELM WDs show significant radial velocity variations indicating the presence of a companion star. We are continuing to follow-up these objects at the MMT to constrain their orbital periods accurately. The importance of these observations is that, we not only find short period binaries that will merge within a Hubble time, but also the majority of the binaries have extreme mass ratios. Understanding the prior history of these systems requires an understanding of the common envelope phase. Studying the mass ratios and physical characteristics of these systems will help in understanding common envelope evolution of close binary pairs."
    },
    "0911/0911.1865_arXiv.txt": {
        "abstract": "Stars are usually formed in clusters in the dense cores of molecular clouds. These embedded clusters show a wide variety of morphologies from hierarchical clusters with substructure to centrally condensed ones. Often they are elongated and surrounded by a low-density stellar halo. The structure of an embedded cluster, i.e.\\ the spatial distribution of its members, seems to be linked to the complex structure of the parental molecular cloud and holds important clues about the formation mechanism and the initial conditions, as well as about the subsequent evolution of the cluster. ",
        "introduction": "\\label{sec:intro}  Stars are born in the dense cores of turbulent molecular clouds that fragment and collapse under their own gravity. Once a gas clump becomes gravitationally unstable, it begins to collapse and the central density increases considerably until a proto\\-stellar core forms in the centre, which grows in mass by accretion from the infalling envelope. It is believed that the star formation process is controlled by the interplay between gravity and supersonic turbulence (Mac Low \\& Klessen 2004; Ballesteros-Paredes et al.\\ 2007; McKee \\& Ostriker 2007). Supersonic turbulence, which is observed ubiquitously in molecular clouds with typical Mach numbers in the range $1 \\lesssim \\mathcal{M} \\lesssim 10$, plays a dual r\\^ole. While on large scales it supports clouds against contraction, on small scales it can lead to local overdensity and collapse.  A large fraction of stars in the Milky Way form in clusters and aggregates of various size and mass (Pudritz 2002; Lada \\& Lada 2003; Allen et al.\\ 2007), appearing to form a hierarchy of structures. The interstellar medium shows a hierarchical structure (sometimes described as fractal) from the largest giant molecular cloud (GMC) scales down to individual cores and clusters, which are sometimes hierarchical themselves. There is no obvious change in morphology at the cluster boundaries, so clusters can be seen as the bottom parts of the hierarchy, where stars have had the chance to mix (Efremov \\& Elmegreen 1998; Elmegreen et al.\\ 2000; Elmegreen 2009). The complex hierarchical structure of turbulent molecular clouds provides a natural explanation for clustered star formation. Molecular clouds vary enormously in size and mass. In small, low-density clouds stars form with low efficiency, more or less in isolation or in distributed small groups of up to a few dozen members. Denser and more massive clouds may build up associations and clusters of several thousand stars. The latter is rare in the solar neighbourhood, but observed in more extreme environments like the Galactic centre or interacting galaxies.  Young clusters usually contain young stellar objects (YSOs) belonging to different evolutionary classes, from deeply embedded Class~0 and Class~1 protostars to relatively evolved pre-main-sequence stars (Class~2 and 3). The morphologies of embedded clusters show a wide variety. There are two basic types of clusters with regard to their structure (Lada \\& Lada 2003): (1) hierarchical clusters showing a stellar surface density distribution with multiple peaks and possible fractal substructure; (2) centrally condensed clusters exhibiting highly centrally concentrated stellar surface density distributions with relatively smooth radial profiles that can be approximated by simple power-law functions or King (1962) profiles.  The analysis of cluster structures requires a complete and unbiased sample of the cluster members. Since embedded stars are best visible in the infrared part of the spectrum, and since nearby GMCs can span several degrees on the sky, wide-field near and mid-infrared surveys are useful tools for identifying and mapping the distribution of young stars in GMCs. Consequently, recent infrared surveys such as 2MASS or several {\\em Spitzer Space Telescope} surveys have revealed a great deal about embedded clusters and their structures. On the other hand, numerical simulations are able to predict the formation of a stellar cluster and its properties with increasing reliability. This review attempts to summarise recent observational and theoretical findings about cluster structures and in particular their clues to star formation and early cluster evolution. ",
        "conclusions": "Embedded clusters show a wide variety of structures, common features are an elongated shape, substructure, and a low-density halo. Clusters seem to evolve from an initial hierarchical configuration consisting of several subclusters toward a centrally condensed state. Therefore the structures of embedded clusters hold important clues to their formation mechanism, and, in a later stage, reflect their dynamical evolution. Furthermore, they are an observational constraint to the results of numerical simulations and the knowledge of a cluster's structure can help identifying new cluster members and discriminating cluster members from field stars (albeit only in a statistical sense). While recent infrared surveys and hydrodynamical and N-body simulations have significantly improved our understanding of the structures of young clusters, there are still a lot of open questions, in particular when it comes to massive clusters or clusters in galaxies other than our Milky Way."
    },
    "0911/0911.3098_arXiv.txt": {
        "abstract": "The recent observations of the anomalous cosmic ray (ACR) energy spectrum as Voyagers 1 and 2 crossed the heliospheric termination shock have called into question the conventional shock source of these energetic particles. We suggest that the sectored heliospheric magnetic field, which results from the flapping of the heliospheric current sheet, piles up as it approaches the heliopause, narrowing the current sheets that separate the sectors and triggering the onset of collisionless magnetic reconnection. Particle-in-cell simulations reveal that most of the magnetic energy is released and most of this energy goes into energetic ions with significant but smaller amounts of energy going into electrons. The energy gain of the most energetic ions results from their reflection from the ends of contracting magnetic islands, a first order Fermi process. The energy gain of the ions in contracting islands increases their parallel (to the magnetic field ${\\bf B}$) pressure $p_\\parallel$ until the marginal firehose condition is reached, causing magnetic reconnection and associated particle acceleration to shut down. Thus, the feedback of the self-consistent development of the energetic ion pressure on reconnection is a crucial element of any reconnection-based, particle-acceleration model. The model calls into question the strong scattering assumption used to derive the Parker transport equation and therefore the absence of first order Fermi acceleration in incompressible flows. A simple 1-D model for particle energy gain and loss is presented in which the feedback of the energetic particles on the reconnection drive is included. The ACR differential energy spectrum takes the form of a power law with a spectral index slightly above $1.5$. The model has the potential to explain several key Voyager observations, including the similarities in the spectra of different ion species. ",
        "introduction": "\\label{intro} Anomalous Cosmic Rays (ACRs) are ions that have energies in the range of $5-100MeV/$nucleon, just below energies associated with galactic cosmic rays. It is known that the ACRs are produced from interstellar neutral atoms since their composition matches that of the local interstellar medium (LISM) \\citep{Cummings96,Cummings07}. In the standard model the LISM neutrals are ionized and picked up by the solar wind deep within the heliosphere and carried out to the heliospheric termination shock, where they undergo diffusive shock acceleration \\citep{Pesses81}. The termination shock (TS) marks the transition of the solar wind from supersonic to subsonic flow.  A major surprise when the Voyager spacecraft crossed the TS was that the number density of the ACRs did not peak at the shock \\citep{Stone05,Stone08}, indicating that their source was elsewhere. Since Voyager 1's crossing of the TS in December 2004 and Voyager 2's crossing in August-September 2007, the two spacecraft have seen increasing intensities of ACRs as they penetrate further into the heliosheath (HS), the subsonic plasma downstream from the shock. One possible explanation for the Voyager observations is that the ACRs peak along the flanks of the heliosphere where the spiral magnetic field from the sun has been in contact with the TS for a longer period of time \\citep{McComas06}. A compressional model of ACRs has also been proposed \\citep{Fisk06a}.  Here we suggest that it is not the TS that is the source of the ACRs. Instead it is the annihilation of the ``sectored'' magnetic field within the heliosheath as it is compressed on its approach to the heliopause (HP) that produces the ACRs. The dipole magnetic field emanating from the sun is quickly dragged in the azimuthal direction due to the sun's rotation. The resulting azimuthal field $B_\\phi$ reverses sign across the equator with the two polarities separated by the heliospheric current sheet. The difference between the rotation axis of the sun and the magnetic pole of the sun causes the heliospheric current sheet to flap in the vertical direction with a period of $26$ days. The resulting folded heliospheric current sheet creates a heliospheric field that displays a ``sector'' structure consisting of alternating signs of $B_\\phi$ \\citep{Wilcox65}. The ``sector-zone'' occupies a latitudinal extent that varies during the solar cycle, reaching nearly the poles when the fields from the sun are a maximum \\citep{Smith01}. The sectors remain a prominent feature of the heliospheric magnetic field out to the TS although the period of the reversals varies erratically from the nominal $13$ day value \\citep{Burlaga03,Burlaga05,Burlaga06}. Within the heliosheath the period of the reversals is even more variable, probably a consequence of the low flow speeds in this region and the time-varying location of the TS \\citep{Burlaga06}.  An obvious question is: Why does the sectored azimuthal field survive out to distances of order $90$AU? Why don't the reversed fields simply annihilate as a result of magnetic reconnection, releasing the stored magnetic energy? The heliosphere is, for all practical purposes, collisionless. It is well known that the rate of collisionless reconnection drops dramatically when the width of the current channel is greater than the ion inertial scale $d_i=c/\\omega_{pi}$ where $\\omega_{pi}$ is the ion plasma frequency \\citep{Yamada07,Cassak05}. At $1$AU the current sheet is on average around $10,000$km wide, which is in the range of $200d_i$ \\citep{Winterhalter94,Smith01}. The deviations from this mean width are substantial and magnetic reconnection is occasionally observed during crossings of the heliospheric current sheet \\citep{Gosling07a}. In any case, upstream of the TS the heliospheric current sheet is too wide to expect collisionless reconnection to dissipate the sector zone.  In the present paper we present MHD simulations of the global heliosphere showing the development of the sectored heliospheric field. The current sheet compresses and the spacing of the current sheets shrinks across the TS. As the plasma in the heliosheath approaches the heliopause its radial motion slows, causing the current sheet widths and sector spacing to further decrease. Current sheets upstream of the Earth's bow shock have been observed to compress by more than a factor of $30$ as they first compress across the shock and then pile up against the magnetopause \\citep{Phan07}. At the same time the magnetic field strength rises and the plasma density and pressure decrease as the plasma is squeezed away from the nose of the heliosphere along the magnetic field.  The resulting drop in the local plasma $\\beta$, the ratio of plasma to magnetic pressure, is similar to that seen just upstream of the Earth's magnetopause \\citep{Crooker79}.  When the width of the heliospheric current layers separating the sectored fields approach the ion inertial scale, the sectored field undergoes magnetic reconnection. In particle-in-cell (PIC) simulations of the sectored geometry, we demonstrate that essentially all of the magnetic energy in the ``sector zone'' is released into energetic particles as the sectors form a complex network of interacting magnetic islands.  A number of mechanisms have been proposed to explain ion acceleration during reconnection, mostly to address flare observations \\citep{Miller97}. Parallel electric fields are more important for electrons than ions \\citep{Litvinenko96,Pritchett08} and the localization of parallel electric fields in any case limits their overall importance \\citep{Drake06,Egedal09}. The resonance absorption of cascading MHD turbulence resulting from reconnection has also been proposed as an accelerator of ions \\citep{Miller98} although the mechanism for the efficient generation of this turbulence remains unspecified. The Petschek slow shocks bounding the reconnection outflow have also been proposed as sites of ion acceleration \\citep{Tsuneta96}. However, satellite crossings of these boundary layers indicate that ions are heated upon crossing into the reconnection outflows but do not reveal a significant energetic particle component \\citep{Gosling05}.  We find that the most energetic ions are accelerated through reflection in contracting islands, a first-order Fermi process that was studied earlier for electron acceleration \\citep{Drake06}. The rate of particle energy gain is given by \\begin{equation} \\frac{dE}{dt}=\\alpha\\langle\\frac{c_A}{L}\\rangle E, \\label{rate} \\end{equation} where $\\langle c_A/L\\rangle$ is a measure of the rate of contraction of magnetic islands and $\\alpha$ is a constant. Importantly, the rate of energy gain is proportional to the particle energy and is independent of the particle mass. The pickup particles gain the most energy and form the ACR spectrum because of their high seed energy. Equation (\\ref{rate}) implies that the ACR spectra of ions with different masses are similar because their seed energy and rate of energy gain are identical when expressed on a per nucleon basis. The energy spectra of all species assume a power law with an index of $(3+\\beta_0)/2$, where $\\beta_0$ is the initial pickup ion $\\beta$ where reconnection onsets. An ACR acceleration model based on the MHD description of turbulent reconnection of the sectored fields has also been recently proposed \\citep{Lazarian09} and its relationship with the present work is discussed in Sec.~\\ref{discussion}. ",
        "conclusions": "\\label{discussion}  The energy spectra of the ACRs have continued to unroll as the Voyager spacecraft have moved further into the HS and there is evidence that the source region of the higher energy particles has now been reached \\citep{Stone08}. The spectral indices of the ACR H and He spectra are around $1.75$. In the reconnection model presented here the power law index of the differential energy spectrum is $\\gamma/2=(3+\\beta_0)/2$. The MHD data shown in Fig.~\\ref{mhdB}(b) indicate that close to the HP the plasma $\\beta$ drops to $0.5$ so the estimated spectral index of the reconnection model approaches $1.75$, which is close to the value observed.  A central question, of course, is whether the rate of energy gain due to reconnection is sufficient to produce ACRs in the range of $100MeV/nuc$. The heating rate in Eq.~(\\ref{Edot}) depends on the typical Alfv\\'en transit time across a magnetic island. The characteristic scale length of the magnetic islands is the sector spacing. An upper limit on the sector spacing is given by the $13$ day sector periodicity times the local heliosheath velocity, which is around $100km/s$ \\citep{Richardson08}. Based on a local magnetic field of $0.15nT$ \\citep{Burlaga08} and a density of $0.002/cm^3$ \\citep{Richardson08}, we obtain an Alfv\\'en velocity of $74km/s$. The Alfv\\'en transit time across a typical magnetic island is therefore $18$ days. From Eq.~(\\ref{Edot}) the characteristic heating time $\\tau_h$ is around a year. The compression of the magnetic field and reduction of the density near the HP can increase $c_A$ by a factor of three, which reduces the heating time to $120$ days. The total e-folding time to increase the $10keV/nuc$ pickup ions to $100MeV/nuc$ therefore ranges from $3$ to $9$ years. The convective flow out of the sector field in the latitudinal direction very likely dominates the loss of energetic particles. For a sector-zone spanning $30^\\circ$ the characteristic width of the sector zone at a HP distance of $150AU$ is around $75AU$. A transverse velocity of $30km/s$ or $6AU/$year yields $\\tau_L=13$ years, which exceeds $\\tau_h$ as expected.  A key observation of ACRs is the similarity in the spectra of different species when expressed on a per nucleon basis. In the case of the shock acceleration model this result is because the rate of energy gain of a particle as it reflects back and forth across the shock is on average independent of mass when expressed on a per nucleon basis \\citep{Blandford87}. The contracting island mechanism for particle energy gain is analagous -- upon reflection from the end of a contracting island, the rate of energy gain is independent of mass when expressed on a per nucleon basis. In particular, the particle distribution functions of the various minor species are also given by the result in Eq.~(\\ref{F}), where $F_0$ is the distribution of seed particles of a given species but the spectral index $\\gamma$ is controlled by H with some contribution from He. Thus, at high energy the spectra of all of the minor species take the form of power laws with the same spectral indices as H.  A surprising observation is the nearly universal $f\\propto v^{-5}$ spectrum of super-Alfv\\'enic ions observed in the quiet-time solar wind throughout the heliosphere \\citep{Fisk06}. These distributions correspond to $F\\propto v^{-3}$ or $F\\propto E^{-1.5}$, the low $\\beta$ limit of Eq.~(\\ref{gamma}). An explanation of this universal spectral index has been offered based on ion acceleration in compressible turbulence \\citep{Fisk06}. We suggest that the non-pickup ions accelerated through reconnection within the heliosheath could be the source of this near universal spectrum of super-Alfv\\'enic particles. Since the pressure of the non-pickup particles is much less than that of the pickup particles, they can be treated as minor species just like the minor ACRs. Since they start from lower energy they never reach energies associated with ACRs, but like the minor species their spectral index is controlled by that of the pickup protons. In the limit of low $\\beta_0$, of course, the spectral index is $1.5$. The spectral index of $1.5$ arises from the approach to the marginal firehose condition of the ACR protons. An important question is whether such a spectrum of super-Alfv\\'enic ions can move upstream across the TS and into the inner heliosphere. An exploration of the transport of such particles is beyond the scope of the present manuscript but should be pursued.  Finally, we note that the reconnection of the sectored magnetic field in the heliosheath was independently proposed as the source of the ACRs by \\cite{Lazarian09}, denoted by LO in the following. Beyond the basic idea that the sectored field is a source for ACR acceleration and that the compression of the sectors on the approach to the heliopause would enhance the likelihood that reconnection would take place, the LO model is very different from that discussed here. Reconnection in the LO model is based on the turbulent reconnection model of \\cite{Lazarian99} although the explicit source of the turbulence is not identified. In the model presented in this manuscript reconnection is collisionless and fast reconnection is facilitated by the Hall term in Ohm's law \\citep{Birn01}. Recent scaling studies \\citep{Shay07} and observations \\citep{Phan07} confirm that Hall reconnection remains fast even in systems comparable in size to that defined by the sector spacing close to the HP. Further, in the model presented here the interaction of islands on adjacent current layers facilitates the rapid dissipation of the magnetic energy in the sectored region. While the mechanism for particle acceleration in both models is a first order Fermi process, the LO model is based on periodic reflection from the plasma flowing inward toward the x-line \\citep{Lazarian05} and curiously is taken in the strongly relativistic limit and so does not apply to ACRs. There is no evidence in even the largest scale PIC simulations carried out to date that such a mechanism is active during reconnection (see the particle distributions in Fig.~2 of \\cite{Drake09}). On the other hand it can't be ruled out that some additional source of turbulence and associated scattering could alter the particle dynamics allowing particle acceleration as proposed by \\cite{Lazarian05} to take place. One might make the case that the particles moving along the trajectories shown in Fig.~\\ref{energy_traj} are reflecting from the inflow of the reconnecting islands and the acceleration mechanism is therefore similar to that of the LO model. However, the spectra resulting from this acceleration process are not given by the \\cite{Lazarian05} model."
    },
    "0911/0911.4501_arXiv.txt": {
        "abstract": "Three years of optical monitoring of the low-mass X-ray binary (LMXB) 4U 1957+11 is presented. The source was observed in $V$, $R$ and $i$-bands using the Faulkes Telescopes North and South. The light curve is dominated by long-term variations which are correlated (at the $> 3 \\sigma$ level) with the soft X-ray flux from the All-Sky Monitor on board the Rossi X-ray Timing Explorer. The variations span one magnitude in all three filters. We find no evidence for periodicities in our light curves, contrary to a previous short-timescale optical study in which the flux varied on a 9.3-hour sinusoidal period by a smaller amplitude. The optical spectral energy distribution is blue and typical of LMXBs in outburst, as is the power law index of the correlation $\\beta = 0.5$, where $F_{\\rm \\nu,OPT}\\propto F_{\\rm X}^{\\beta}$. The discrete cross-correlation function reveals a peak at an X-ray lag of 2--14 days, which could be the viscous timescale. However, adopting the least squares method we find the strongest correlation at a lag of $0 \\pm 4$ days, consistent with X-ray reprocessing on the surface of the disc. We therefore constrain the optical lag behind X-ray to be between -14 and +4 days. In addition, we use the optical--X-ray luminosity diagram for LMXBs as a diagnostic tool to constrain the nature of the compact object in 4U 1957+11, since black hole and neutron star sources reside in different regions of this diagram. It is found that if the system contains a black hole (as is the currently favoured hypothesis), its distance must exceed $\\sim 20$ kpc for the optical and X-ray luminosities to be consistent with other soft state black hole systems. For distances $< 20$ kpc, the data lie in a region of the diagram populated only by neutron star sources (black hole systems are ten times optically brighter for this X-ray luminosity). 4U 1957+11 is unique: it is either the only black hole LMXB to exist in an apparent persistent soft state, or it is a neutron star LMXB which behaves like a black hole. ",
        "introduction": "4U 1957+11 (V1408 Aql) is a persistently active ($\\sim$ 1--10$\\times 10^{-10}$ erg cm$^{-2}$ s$^{-1}$; 1.5--12 keV) low-mass X-ray binary (LMXB). Its X-ray spectrum has always appeared soft since its discovery \\citep*{giacet74,wijnet02}. These two properties make this source unique; most LMXBs are transient or regularly perform transitions to harder X-ray states \\citep*[for more on X-ray states see e.g.][]{hasiva89,homabe05,mcclet06,fendet09}. The X-ray spectrum of 4U 1957+11 can be described by a blackbody plus a power law, with fluxes which vary on timescales of hours to months \\citep*{yaqoet93,singet94,nowawi99}. An apparent 117 day X-ray period \\citep{nowawi99} possibly due to a precessing accretion disc viewed almost edge-on was later shown to be much more complex; the long-term X-ray flux variations are likely to be driven by changes in the mass accretion rate \\citep{yaqoet93,wijnet02}.   \\begin{figure} \\centering \\includegraphics[width=7cm,angle=0]{finding.eps} \\caption{Finding chart for 4U 1957+11, indicating the two stars used for flux calibration. The image is a 100 sec exposure in i-band taken with the EM01 camera on FTN under good seeing conditions. The field-of-view is 60 $\\times$ 60 arcsec. The optical counterpart of 4U 1957+11 is marked in the centre. The magnitudes of stars 1 and 2 are given in Table 2. North is to the top and east is to the left. For a larger field, see \\citeauthor{marget78} (1978). } \\end{figure}    No estimates of the mass function \\citep[e.g.][]{casa07} have been established because the system has never faded to quiescence since its discovery, and type I X-ray bursts typical of neutron stars have not been detected -- as such it is uncertain whether the compact object in the 4U 1957+11 system is a black hole or a neutron star. It was argued \\citep{yaqoet93} that the X-ray luminosity is too low, and the inner disc radius too large for a candidate black hole X-ray binary (BHXB), but this may not be the case, and several X-ray spectral and timing properties are similar to other BHXBs \\citep{wijnet02,nowaet08}. The X-ray spectrum is typical of soft state BHXBs, and its behaviour shares many similarities with LMC X--3 \\citep[][and references therein]{wijnet02}, which contains a $> 5.8$ M$_{\\odot}$ black hole and is also in a persistent soft state \\citep[but unlike 4U 1957+11, LMC X--3 occasionally performs transitions to the hard state;][]{wilmet01}. Moreover, the luminosity is higher and consistent with soft state BHXBs if the distance to the source is further than the often assumed 7 kpc \\citep*[e.g.][]{marget78}. The X-ray spectrum also appears slightly harder at higher fluxes, which is not usually the case for neutron star X-ray binaries \\citep[NSXBs;][]{wijnet02}.  Recently, \\cite{nowaet08} fitted high-resolution X-ray spectra of 4U 1957+11 with an inclined disc blackbody with an interstellar absorption of n$_{\\rm H} = (1.5 \\pm 0.5) \\times 10^{21}$ cm$^{-2}$. The total absorption through the Galaxy in the direction of the source is $1.2 \\times 10^{21}$ cm$^{-2}$ \\citep{kalbet05}, so the measured absorption implies a large distance -- on the far side of the Galaxy or possibly in the halo \\citep[see also][]{yaoet08}. Indeed, \\cite{nowaet08} use \\emph{Chandra}, \\emph{XMM-Newton} and Rossi X-ray Timing Explorer (\\emph{RXTE}) data to infer a rapidly spinning black hole and a distance of $\\sim 10$--22 kpc. However, the authors do not explore the possibility of a neutron star using their data; instead they use the spectral and timing arguements above to favour a black hole.   \\begin{table*} \\caption{Faulkes Telescopes observations of 4U 1957+11 used in this work.} \\begin{tabular}{llllllll} \\hline Date      &MJD    &Camera  &Airmass&\\multicolumn{3}{c}{Exposure times (sec)}\\\\ &       &        &       &$V$\t     &$R$      &$i$\\\\ \\hline 2006-04-08&53833.6&DillCam &1.26   &120      &120      &120\\\\ 2006-05-10&53865.6&DillCam &1.08   &200      &200      &200\\\\ 2006-05-21&53876.6&DillCam &1.04   &200      &200      &200\\\\ 2006-05-27&53882.6&DillCam &1.02   &200      &200      &200\\\\ 2006-06-07&53893.6&DillCam &1.01   &200      &--       &--\\\\ 2006-11-14&54053.2&HawkCam &1.24   &200      &--       &200\\\\ 2007-04-22&54212.5&HawkCam1&1.56   &100      &--       &--\\\\ 2007-06-16&54267.5&HawkCam1&1.09   &100      &100      &100\\\\ 2007-07-10&54291.4&HawkCam1&1.20   &100      &100      &100\\\\ 2007-07-12&54293.5&HawkCam1&1.01   &--       &--       &100\\\\ 2007-08-06&54318.3&HawkCam1&1.12   &100      &100      &100\\\\ 2007-08-10&54322.3&HawkCam1&1.09   &100      &100      &100\\\\ 2007-08-19&54331.4&HawkCam1&1.03   &100      &100      &100\\\\ 2007-09-02&54345.3&HawkCam1&1.02   &100      &100      &100\\\\ 2007-10-05&54378.3&HawkCam1&1.08   &100      &100      &100\\\\ 2007-11-13&54417.3&HawkCam1&2.42   &100      &100      &100\\\\ 2008-06-13&54630.4&EM01    &1.70   &100      &100      &100\\\\ 2008-06-16&54633.5&EM01    &1.02   & 80      &110, 100 &2 $\\times$ 100 sec\\\\ 2008-06-20&54637.6&EA01    &1.12   &100      &100      &100\\\\ 2008-06-24&54641.5&EM01    &1.05   &--       &2 $\\times$ 100 sec&2 $\\times$ 100 sec\\\\ 2008-06-25&54642.5&EM01    &1.01   &--       &100      &2 $\\times$ 100 sec\\\\ 2008-06-30&54647.6&EA02 (FTS)&1.20--1.42&--  &--       &18 $\\times$ 100 sec\\\\ 2008-07-11&54658.5&EM01    &1.01--1.04&2 $\\times$ 100 sec&2 $\\times$ 100 sec&16 $\\times$ 100 sec\\\\ 2008-07-14&54661.5&EM01    &1.04   &100      &100      &100\\\\ 2008-07-14&54661.6&EA02 (FTS)&1.37   &2 $\\times$ 100 sec&2 $\\times$ 100 sec&4 $\\times$ 100 sec\\\\ 2008-07-16&54663.5&EM01    &1.19   &100      &100      &--\\\\ 2008-07-16&54663.6&EA02 (FTS)&1.37--1.42&2 $\\times$ 100 sec&100&3 $\\times$ 100 sec\\\\ 2008-07-30&54677.4&EA02 (FTS)&2.07   &200      &200      &200\\\\ 2008-08-09&54687.3&EM01    &1.05   &100      &100      &100\\\\ 2008-08-31&54709.5&EM01    &2.06   &100      &100      &100\\\\ 2008-09-10&54719.5&EM01    &2.16--2.41&2 $\\times$ 100 sec&100&2 $\\times$ 100 sec\\\\ 2008-09-15&54724.5&EM01    &2.13   &--       &100      &100\\\\ 2008-09-25&54734.4&EM01    &1.52   &100      &100      &100\\\\ 2008-11-04&54774.2&EM01    &1.18   &100      &100      &100\\\\ 2009-02-22&54884.6&EM01    &2.50   &100      &100      &100\\\\ 2009-05-17&54971.5&EM01    &1.12   &100      &100      &100\\\\ \\hline \\end{tabular} \\small \\\\ MJD = Modified Julian Day. The EA02 camera is on Faulkes Telescope South and all the other cameras are on Faulkes Telescope North. Data from 34 dates were used, totalling 144 images. \\normalsize \\end{table*}   An optical counterpart to 4U 1957+11 was discovered by \\cite{marget78}, which was found to vary sinusoidally on an orbital period of 9.33 hours \\citep{thor87}. However, a later photometric study revealed a more complex light curve of larger amplitude variability \\citep*{hakaet99}. In this latter study the authors establish a colour-dependency in the amplitude; the optical (UBVRI) spectral energy distribution (SED) is slightly bluer at lower fluxes. \\cite{hakaet99} modelled the light curve assuming it modulates on the orbital period, however the data cover just one 9-hour orbit. Since the optical behaviour is in contrast with the original sinusoidal modulation found \\citep{thor87}, the variability they measure may not be periodic. \\cite{shahet96} took a broad-band (4000--10,000\\AA) optical spectrum of 4U 1957+11, which displayed a mostly featureless continuum, with weak H$\\alpha$, H$\\beta$ and He \\small II \\normalsize emission lines. No absorption lines from the companion star were detected.  Here, we present the results of an optical monitoring campaign of 4U 1957+11 spanning 37 months during 2006--2009. Concurrent X-ray monitoring reveals a positive optical--X-ray correlation, and helps to constrain the origin of the optical variability and the nature of the compact object. The observations, reduction and flux calibration are described in Section 2. We analyze the results and give a comprehensive discussion, making comparisons to other sources, in Section 3. In Section 4 we summarize the results and conclusions. ",
        "conclusions": "We have monitored the optical flux of the persistent LMXB 4U 1957+11 in three wavebands over a period of three years with the two 2-m Faulkes Telescopes. Like the soft X-ray light curve, the optical continuum is dominated by long-term variations on timescales of days--months, likely originating in changes of the mass accretion rate. The amplitude is a factor $\\sim$ two in flux over these timescales; around half the amplitude of the X-ray variations. Contrary to previous shorter-timescale studies, no clear periodicities exist in the optical light curves. However we do not dispute the orbital period of 9.3 hours previously identified over a time span of $\\sim 10$ days by \\cite{thor87}. We suggest the X-ray flux (and so X-ray heating of the disc and star) during their observations may have been fairly stable (as is sometimes the case on those timescales; see Fig. 2) so that the low-amplitude sinusoidal variations from the heated face of the companion star were revealed (to our knowledge no X-ray data were acquired during that epoch). Under the assumption that the variability they see over one 9-hour period in 1996 was periodic, \\cite{hakaet99} model their complex light curve as an evolving accretion disc, and infer a mass ratio $q \\approx 0.4$--0.5 and an inclination angle of 70--75$^{\\circ}$. We remark that these estimates of the system parameters are invalid as the flux was probably not periodic at that time.  The optical SED is blue ($\\alpha \\approx +1.0$) and slightly redder at lower fluxes. This is indicative of an accretion disc (heated viscously or by irradiation) and rules out synchrotron emission (which has an even redder SED) as the dominating process. We establish a positive correlation between the optical and \\emph{RXTE} ASM 1.5--12 keV X-ray fluxes, of $F_{\\rm \\nu,OPT}\\propto F_{\\rm X}^{0.5}$, significant at up to the 4.5 $\\sigma$ level (conservatively, $> 3 \\sigma$) and evident from the three optical wavebands independently. The power law index of the correlation shows that the empirical correlation found over whole outbursts for LMXBs \\citep{russet06} can hold over much smaller ranges in luminosity.  Cross-correlating the optical and X-ray light curves with the least squares method reveals a lag of $0 \\pm 4$ days, consistent with X-ray reprocessing in which X-ray variations should precede the optical on light-travel timescales (seconds). However, the peak in the discrete CCF is at an X-ray lag of 2--14 days. This could be the viscous timescale for matter or waves of turbulence to travel from the optically-emitting outer disc to the X-ray-emitting inner disc. Combining these two methods we constrain the optical lag to be between -14 and +4 days. Optical and X-ray monitoring with more data and with higher S/N in X-ray is required to further constrain these lags.  BHXBs and NSXBs lie in different regions of the optical--X-ray luminosity diagram. We therefore overplot the data of 4U 1957+11, adopting a range of distances to the source, in order to constrain the currently disputed nature of the compact object. At distances $< 20$ kpc the data lie in a region of the diagram populated only by neutron star sources. If it were a BHXB at this distance, the optical emission would be one order of magnitude fainter than the other BHXBs, and it would be the faintest soft state BHXB. We therefore postulate that if its distance is $< 20$ kpc it is most likely a peculiar neutron star system.  However, taking into account the BHXB-like X-ray properties of 4U 1957+11, the most likely scenario is that the source is a distant BHXB. Its optical and X-ray luminosities are consistent with other soft state black hole systems if its distance exceeds $\\sim 20$ kpc. 4U 1957+11 may well be a Galactic halo BHXB, in analogy with the prototypical halo BHXB, XTE J1118+480. This would imply a high space velocity and a large supernova `kick' at its birth. Optical phase-resolved radial velocity studies during a faint period (when the star/disc ratio is highest) may be required to identify the compact object unambiguously.  If 4U 1957+11 indeed contains a black hole, it is a valuable addition to the very few persistently accreting black hole X-ray binaries known (e.g. LMC X--1; LMC X--3) and possibly only the second in our Galaxy, after Cygnus X--1. In addition, 4U 1957+11 would be the perfect candidate to test for the presence of a weak soft state radio jet. A steady, core radio jet has never been detected from a BHXB in the soft state \\citep[e.g.][]{fendet09} but it is speculated that faint jets (either intrinsically weak or with very low radiative efficiency) may be produced by soft state BHXBs. The only radio observations of 4U 1957+11 in the literature provide an upper limit of 2.1 mJy at 5 GHz \\citep{nelssp88}. X-ray binaries can be detected down to, for example 50 $\\mu$Jy or fainter with the Very Large Array \\citep{gallet06} and much fainter with future planned Square Kilometer Array Pathfinders."
    },
    "0911/0911.0819_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\label{sec:1}  How our planet was formed, how life came about, and whether life exists elsewhere in the universe are among some of the long-standing questions in human history. The latter, which has been the main drive behind many decades of searching for planets outside the solar system, is one of the most outstanding problems in planetary science and astrobiology. Although no Earth-like planet has yet been found, the success of observational techniques in identifying now more than 350  extrasolar planets has greatly contributed to addressing this question, and has extended the concept of habitability to billions of miles beyond the boundaries of our solar system. It is now certain that our planetary system is not unique and many terrestrial-size planets may exist throughout the universe.  The orbital and physical diversity of the currently known extrasolar planets play a crucial role in their habitability. In general, whether a planet can be habitable depends on its physical and dynamical properties, and the luminosity of its host star. The notion of habitability is normally defined based on the life as we know it, and uses the physical and orbital characteristics of Earth as an example of a habitable planet. In other words, a planet is habitable if it is Earth-like so it can develop and sustain Earthly life. This definition of habitability requires that a potentially habitable planet to maintain liquid water on its surface and in its atmosphere. The planet's capability in maintaining water is determined by its size and orbital motion, the luminosity of the central star, and the distribution of water in the circumstellar material from which the planet was formed.  How extrasolar habitable planets are formed is a widely addressed question that is still unresolved. While models of planetary accretion in the inner solar system present pathways (although in some cases incomplete) toward the formation of planets such as Earth and Venus, the orbital diversity of extrasolar planets present strong challenges to the applicability of these models to other planetary environments. For instance, systems with close-in giant planets may require massive protoplanetary disks to ensure that while planetesimals and protoplanets are scattered as giant planets migrate, terrestrial bodies can form and be stable. Systems with multiple planets also present a great challenge to terrestrial planet formation since the orbital architectures of such systems may limit the regions of the stability of smaller objects.  In a system with two stars, the situation is even more complicated. The interaction between one star and the protoplanetary disk around the other may inhibit planet formation by truncating the disk and removing circumstellar material \\cite{Artymowicz94}. This interaction may also  prevent the growth of km-size planetesimals to larger objects by increasing the relative velocities of these bodies and causing their collisions to result in fragmentation. Despite such difficulties, planets have, however, been detected in binary star systems (see \\ref{tab:1}) and observers have been able to identify three moderately close ($< 20$ AU) binaries, namely $\\gamma$ Cephei \\cite{Hat03}, GL 86 \\cite{Eggenberger01} , and HD 41004 \\cite{Zucker04}, whose primary stars are hosts to Jupiter-like planets.  The detection of planets in binary star systems is not a surprise. There is much observational evidence that indicates the most common outcome of the star formation process is a binary system \\cite{Math94,White01}. Also, as shown by Prato \\& Weinberger in chapter 1, there is substantial evidence for the existence of potentially planet-forming circumstellar disks in multiple star systems \\cite{Math94,Akeson98,Rodriguez98,White99,Silbert00,Math00,Trilling07}. These all point to the fact that planet formation in binaries is robust and many of these systems may harbor additional giant planets and/or terrestrial-size objects. This chapter is devoted to study the latter.  Whether binaries can harbor potentially habitable planets depends on several factors including the physical properties and the orbital characteristics of the binary system. While the former determines the location of the habitable zone (HZ), the latter affects the dynamics of the material from which terrestrial planets are formed (i.e., planetesimals and planetary embryos), and drives the final architecture of the planets assembly. In order for a habitable planet to form in a binary star system, these two factors have to work in harmony. That is, the orbital dynamics of the two stars and their interactions with the planet-forming material have to allow terrestrial planet formation in the habitable zone, and ensure that the orbit of a potentially habitable planet will be stable for long times. We organize this chapter with the same order in mind. We begin in section 2 by presenting a general discussion on the motion of planets in binary stars and their stability. Section 3, has to do with the stability of terrestrial planets, and in section 4, we discuss habitability and the formation of potentially habitable planets in a binary-planetary system\\footnote{The phrase \"binary-planetary system\" is used to identify binary star systems in which one of the stars is host to a giant planet.}.   \\begin{table} \\centering \\caption{Planets in double stars \\cite{Rag06}} \\label{tab:1} \\begin{tabular}{crrcr} \\hline\\noalign{\\smallskip} Star & \\hskip 10pt $a_{b}$ (AU) & \\hskip 10pt $a_{p}$ (AU) & \\hskip 10pt  $M_{p} \\sin i (M_{Jup})$ & \\hskip 10pt $e_{p}$\\\\ \\noalign{\\smallskip}\\hline\\noalign{\\smallskip} HD38529 & 12042 & 0.129 & 0.78 & 0.29 \\\\ & & 3.68 & 12.7 & 0.36 \\\\ HD40979 & 6394 & 0.811 & 3.32  & 0.23\\\\ HD222582 & 4746 & 1.35 & 5.11 & 0.76 \\\\ HD147513 & 4451 & 1.26 & 1.00 & 0.52 \\\\ HD213240 & 3909 & 2.03 & 4.5 & 0.45 \\\\ Gl 777 A & 2846 &0.128 & 0.057 & 0.1 \\\\ & & 3.92 &1.502 & 0.36\\\\ HD89744 & 2456 & 0.89 & 7.99 & 0.67 \\\\ GJ 893.2 & 2248 & 0.3 & 2.9 & -- \\\\ HD80606 & 1203 & 0.439 & 3.41 & 0.927 \\\\ 55 Cnc & 1050 & 0.038 & 0.045 & 0.174 \\\\ & & 0.115 & 0.784 & 0.02 \\\\ & & 0.24 & 0.217 & 0.44\\\\ & & 5.25 & 3.92 & 0.327\\\\ GJ 81.1 & 1010 & 0.229 &0.11 & 0.15 \\\\ & & 3.167 &0.7 & 0.3 \\\\ 16 Cyg B & 860 & 1.66 & 1.69 & 0.67\\\\ HD142022 & 794 & 2.8 & 4.4 & 0.57 \\\\ HD178911 & 789 & 0.32 & 6.292 & 0.124 \\\\ Ups And & 702 &0.059 & 0.69 & 0.012 \\\\ & & 0.83 & 1.89 &0.28 \\\\ & & 2.53 &3.75 & 0.27\\\\ HD188015 &684 & 1.19 & 1.26 & 0.15 \\\\ HD178911 & 640 & 0.32 & 6.29 & 0.124\\\\ HD75289 & 621 & 0.046 & 0.42 & 0.054 \\\\ GJ 429 & 515 & 0.119 & 0.122 & 0.05 \\\\ HD196050 & 510 & 2.5 & 3.00 & 0.28 \\\\ HD46375 & 314 & 0.041 & 0.249 & 0.04 \\\\ HD114729 & 282 & 2.08 & 0.82 & 0.31 \\\\ $\\epsilon$ Ret & 251 & 1.18 & 1.28 & 0.07 \\\\ HD142 & 138 &0.98 &1.00 & 0.38 \\\\ HD114762 & 132 & 0.3 & 11.02 & 0.25 \\\\ HD195019 & 131 & 0.14 & 3.43 & 0.05 \\\\ GJ 128 & 56 & 1.30 & 2.00 & 0.2 \\\\ HD120136 & 45 & 0.05 & 4.13 & 0.01 \\\\ \\hline $\\gamma$ Cep & 20.3 & 2.03 & 1.59 & 0.2\\\\ GL 86 & 21  & 0.11 & 4.01 &0.046  \\\\ HD41004 AB &  23 & 1.7 & 2.64 & 0.5 \\\\ \\hline \\noalign{\\smallskip}\\hline \\end{tabular} \\end{table} ",
        "conclusions": ""
    },
    "0911/0911.0502_arXiv.txt": {
        "abstract": "Using pulsewidth data for 872 isolated radio pulsars we test the hypothesis that pulsars evolve through a progressive narrowing of the emission cone combined with progressive alignment of the spin and magnetic axes.  The new data provide strong evidence for the alignment over a time-scale of about $1\\,$Myr with a log standard deviation of around 0.8 across the observed population.  This time-scale is shorter than the time-scale of about $10\\,$Myr found by previous authors, but the log standard deviation is larger.  The results are inconsistent with models based on magnetic field decay alone or monotonic counter-alignment to orthogonal rotation. The best fits are obtained for a braking index parameter, $n_{\\gamma}\\approx2.3$, consistent the mean of the six measured values, but based on a much larger sample of young pulsars. The least-squares fitted models are used to predict the mean inclination angle between the spin and magnetic axes as a function of log characteristic age. Comparing these predictions to existing estimates it is found that the model in which pulsars are born with a random angle of inclination gives the best fit to the data.   Plots of the mean beaming fraction as a function of characteristic age are presented using the best-fitting model parameters. ",
        "introduction": "One of the greatest of the many mysteries of radio pulsars is their evolutionary histories.  As pulsars age, powerful electromagnetic torques act to increase the rotation period, $P$, of the strongly magnetized neutron star, causing the spin-frequency, $\\nu=1/P$, to decrease over time, a phenomenon known as spin-down.  This can be described by the equation \\begin{equation} \\dot{\\nu}= -K \\nu^{n} \\mbox{~,} \\label{eq:SpinDown} \\end{equation} where $n$ is a parameter known as the braking index, with $n=3$ for a magnetic dipole rotating in a vacuum.  The spin-down torque acting on the star is proportional to $\\dot{\\nu}$, and it decays over time as $\\dot{\\nu}$ decreases.  There is also evidence that $K$ is not constant, but rather is also changing over time. The surface dipole magnetic field strength at the magnetic equator, assuming that the spin and magnetic axes are orthogonal, is conventionally given by $B_{\\mathrm{surf}}=3.2\\times10^{19}(P\\dot{P})^{1/2}\\,\\mbox{G}$; at the magnetic poles the field strength is $2 B_{\\mathrm{surf}}$ \\citep{ls06}. $B_{\\mathrm{surf}}$ tends to be bigger for younger pulsars than for older ones having a large spin-down age, \\(\\SpinDownAge{}\\LRdefined{}P/\\left[(n-1)\\dot{P}\\right]\\).  This in turn suggests that $K$ tends to reduce with increasing age, and much conjecture remains about the origin(s) of this evolution.  Many authors have argued that the main contributor to the $K$-evolution is the progressive alignment of the spin and magnetic axes: e.g. \\citet{cb83}, \\citet{cb86}, \\citet{lm88}, \\citet{xw91}, \\citet{kw92a}, \\citet{can93}, \\citet{pp96}, \\citet{tm98} and \\citet{wj08}.  But \\citet{mck93} and \\citet{gh96} have argued in favour of a random distribution of the inclination angles between the angular velocity and magnetic axes and against magnetic alignment.  Others have favoured magnetic field decay \\citep[e.g.][]{no90}, while there are also authors who argue for continuing counter-alignment to eventual orthogonality of the axes, via an electromagnetic torque exerted by magnetospheric electric currents flowing on open field lines \\citep{bgi84}.  According to \\citet{jon76}, pulsars initially counter-align (via a strongly temperature-dependent dissipative torque in the fluid interior) to reach the orthogonal rotator state after about $10^{3}\\,$yr, where they remain for $10^4$--$10^{5}\\,$yr (while the dissipative torque decays as the interior cools), after which alignment (via electromagnetic torque) occurs.  If $K$ were constant, the braking index $n$ would be equal to the \\emph{apparent} braking index, \\begin{equation} \\nAPP{app} \\LRdefined {\\nu\\ddot{\\nu}}/{\\dot{\\nu}^{2}} = 2 - {P \\ddot{P}}/{\\dot{P}^{2}} \\mbox{~.} \\label{eq:BrakingIndex:Definition} \\end{equation} If $\\ddot{\\nu}$ can be determined from observations of a pulsar, then $\\nAPP{app}$ can be computed.  However, if $K$ is varying, then in general $\\nAPP{app}\\neq{}n$.  Some authors \\citep[\\mdtyEG{}][]{jg99, tk01} have argued that, for pulsars of moderate age, $\\nAPP{app}$ exceeds the value of 3 that corresponds to magnetic dipole radiation.  This could be taken to imply that either magnetic alignment or field decay must be occurring.  However, it is likely that recovery from unseen pulsar glitches is the principal contribution to measured variations in $\\dot{P}$ \\citep{wmz+01} for these pulsars.  In this paper, we shall use the pulsar evolution models from \\citet[][hereafter CB83]{cb83}, \\citet[][hereafter CB86]{cb86} and \\citet[][hereafter C93]{can93}, which all invoke a progressive alignment of the spin and magnetic axes.  The CB83 model develops the pulsewidth-age relation on the assumption that all pulsars move along the same evolutionary track, differing only in age and orientation with respect to Earth.  The CB86 model relaxes this assumption and allows a distribution of alignment time-scales, as well as a time-varying relationship between characteristic and actual age.  The C93 models additionally allow a distribution of initial inclination angles.  All of these models predicted a minimum in the mean pulsewidth as a function of characteristic age and provided a good fit to the available data, consisting of 293 pulsewidth values from the \\citet{mt81} catalogue.    Here we use the ATNF (Australia Telescope National Facility) pulsar catalogue\\footnote{http://www.atnf.csiro.au/research/pulsar/psrcat/} \\citep{mhth05}, version 1.35, to investigate pulsar evolution through magnetic alignment using the Candy--Blair models.  This contemporary catalogue contains much larger data sets of 872 and 1420 isolated radio pulsars with 10~per~cent and 50~per~cent intensity pulsewidth values, $W_{10}$ and $W_{50}$.  It also contains many more old pulsars than was previously the case, which tend to have lower radio luminosities.  It thus allows a more precise analysis of pulsar evolution.  As a consistency check, we also examine separately those pulsars with pulsewidths measured during the Parkes multibeam pulsar survey \\citep{mlc+01,lfl+06}, giving $W_{10}$ and $W_{50}$ data sets of 377 and 934 pulsars respectively.   The paper is structured as follows. \\Secref{sec:MagneticAlignmentTheory} presents a summary of the theory of alignment between the spin and magnetic axes; \\secref{sec:PulsewidthEvolution} gives a mathematical description of pulsewidth evolution; in \\secref{sec:Data:Fitting} we least-squares fit the evolutionary pulsewidth models to the $W_{10}$ and $W_{50}$ data as functions of characteristic age; in \\secref{sec:analysis} we further analyse these results, in particular comparing them to the inclination angle estimates of \\citet{ran93b} and \\citet{gou94}; and conclusions are presented in \\secref{sec:Conclusions}. ",
        "conclusions": "\\label{sec:Conclusions}  From \\figref{fig:Models:123}, showing mean pulsewidths versus $\\log{t_c}$, we have seen that the pulsewidth data clearly favour pulsar evolution through magnetic alignment, and not magnetic field decay alone nor progressive counter-alignment. There is a clear minimum of mean pulsewidth at about $t_c \\approx 10^{6.9}\\,$yr. The alignment process causes a significant increase in the mean $W_{10}$  pulsewidth for all characteristic age bins greater than $3 \\times10^{7}\\,$yr. This is particularly significant in view of the fact that the ATNF catalogue contains numerous pulsars with $t_c > 10^8\\,$yr, which were not available to the earlier studies.  The three models of initial alignment that were specified in \\secref{sec:AlignmentModels} cannot be tested by fitting the Candy-Blair models to the pulsewidth data alone.  However, once a Candy-Blair model has been fitted to the pulsewidth data, it can be used to \\emph{predict} mean inclination angle evolution, $\\left<\\alpha\\right>\\!(\\log{t_c})$. Comparing these predictions to the datasets of \\citet{ran93b} and \\citet{gou94}, it was found that only Model~II (a random initial inclination angle) gives a good fit to the data, as shown in \\figref{fig:Alpha:Data}.   The alignment time-scale found from fitting Model~II to the ATNF~Cat~$W_{10}$ data (see \\tablerefs{table:Datasets} and \\ref{table:parameters:three:nGamma:min}) is $10^{6.1}\\,$yr, which is at least a factor of ten shorter than the timescales determined by a number of other authors (see \\tableref{table:alignment:timescales}). The fit requires that the log alignment time-scale varies with a standard deviation of about $0.83$ across the observed population.  The fit gives a limiting cone opening half-angle of $4\\fdg11$, and a braking index parameter $n_{\\gamma}$ of 2.30, consistent with the mean, $2.42\\pm{}0.05$, of the measured apparent braking indices (\\eqns{eq:BrakingIndex:Definition} and \\eqNo{eq:nApp:nGamma:tc}, and \\tableref{table:n}).  Consistent with empirical observations, the fit assumes the parameter $\\gamma=1/2$ (see \\eqn{eq:rho:t}).  For other values of $\\gamma$ the fitted parameters $n$ and $\\mu_{\\log}$ are determined from \\eqns{eq:parameters:gamma:n} and \\eqNo{eq:mu:gamma:log} respectively.  A plot of the mean beaming fraction for this fit is given in \\figref{fig:fb:mean:scaled}.   The initial predictions of pulsar alignment due to electromagnetic braking \\citep{dg70,mg70} predicted time-scales of around $10^{3}$--$10^{4}\\,$yr, which was rather small.  The time-scales will be increased if pulsars are born with their magnetic and rotation axes nearly orthogonal.  \\citeauthor{jon76} \\citeyearpar{jon76,jon76b,jon76a} found that the inclusion of a decreasing dissipative torque initially brought all pulsars close to counter-alignment, from which time the normal electromagnetic decay of $\\alpha$ would begin.  \\citeauthor{jon76} predicted that the alignment phase started at around $10^{3}$--$10^{4}\\,$yr, with a time-scale of around $10^{6}\\,$yr, which was believed to be more in accord with the observed population.  We have found alignment time-scales consistent with the predictions of \\citeauthor{jon76}, though with a significant spread in values, down to the time-scales consistent with the absence of the counter-alignment phase.  Contrary to the model of Jones, though, we do not find evidence for all pulsars starting their alignment phase with $\\alpha_{0}$ close to $90\\degr$ (i.e. Model~I).  This interesting counter-point to the model of Jones deserves further investigation."
    },
    "0911/0911.0444_arXiv.txt": {
        "abstract": "We have discovered a new transient low-mass X-ray binary, NGC 6440 X-2, with \\Chandra/ACIS, \\RXTE/PCA, and \\Swift/XRT observations of the globular cluster NGC 6440.  The discovery outburst (July 28-31, 2009) peaked at $L_X\\sim1.5\\times10^{36}$ \\ergss, and lasted for $<$4 days above $L_X=10^{35}$ \\ergss. Four other outbursts (May 29-June 4, Aug. 29-Sept. 1, Oct. 1-3,  and Oct. 28-31 2009) have been observed with \\RXTE/PCA (identifying millisecond pulsations, Altamirano et al. 2009a) and \\Swift/XRT (confirming a positional association with NGC 6440 X-2), with similar peak luminosities and decay times. Optical and infrared imaging did not detect a clear counterpart, with best limits of $V>21$, $B>22$ in quiescence from archival \\HST\\ imaging, $g'>22$ during the August outburst from Gemini-South GMOS imaging,  and $J\\gsim18.5$ and $K\\gsim17$ during the July outburst from CTIO 4-m ISPI imaging. Archival \\Chandra\\ X-ray images of the core do not detect the quiescent counterpart ($L_X<1-2\\times10^{31}$ \\ergss), and place a bolometric luminosity limit of $L_{NS}< 6\\times10^{31}$ \\ergss\\ (one of the lowest measured) for a hydrogen atmosphere neutron star.  A short \\Chandra\\ observation 10 days into quiescence found two photons at NGC 6440 X-2's position, suggesting enhanced quiescent emission at $L_X\\sim6\\times10^{31}$ \\ergss. NGC 6440 X-2 currently shows the shortest recurrence time ($\\sim$31 days) of any known X-ray transient, although regular outbursts were not visible in the bulge scans before early 2009.  Fast, low-luminosity transients like NGC 6440 X-2 may be easily missed by current X-ray monitoring. ",
        "introduction": "\\label{s:intro}  The dense cores of globular clusters are known to be efficient factories for dynamically producing tight binaries containing heavy stars, and thus X-ray binaries \\citep{Clark75,Hut91,Pooley03}.  Thirteen luminous ($L_X>10^{35}$ \\ergss) low-mass X-ray binaries (LMXBs) have previously been identified in Galactic globular clusters \\citep[see][]{Verbunt04}, concentrated in the densest and most massive clusters \\citep{Verbunt87,Verbunt02}.  Luminous LMXBs in globular clusters of other galaxies are clearly concentrated in the most massive, most metal-rich, and densest globular clusters \\citep{Kundu02,Sarazin03,Jordan04}.  A critical question for studies of globular cluster LMXBs is whether X-ray emission from a globular cluster is due to one LMXB or multiple LMXBs.  This affects the inferred nature of sources \\citep{Dotani90,diStefano02,Maccarone07} and luminosity functions \\citep{Sivakoff07}.  For example, an apparent contradiction in the qualities of the LMXB in M15 was resolved by the identification of two persistent X-ray sources in the cluster \\citep{White01}.  Identification with \\Chandra\\ of multiple quiescent LMXBs in several globular clusters \\citep{Grindlay01a,Pooley02b,Heinke03d} has suggested that transient LMXB outbursts from a cluster might arise from different sources, although some LMXBs possess distinctive characteristics \\citep[e.g. the Rapid Burster,][]{Homer01}. Here we report the discovery and outburst monitoring of the 14th luminous LMXB in a Galactic globular cluster, NGC 6440, the first cluster to show two luminous transient LMXBs (both of which show millisecond pulsations).  In a companion paper, \\citet{Altamirano09} identify this LMXB as a new ultracompact accreting millisecond X-ray pulsar (AMXP), after discovering coherent 206 Hz pulsations \\citep{Altamirano09b} and fitting the frequency drift to a 57.3 minute orbital period.  NGC 6440 is a globular cluster near the Galactic Center, at a distance of 8.5 kpc \\citep{Ortolani94}, with $N_H=5.9\\times10^{21}$ cm$^{-2}$ \\citep{Harris96}.  A luminous X-ray source (MX 1746-20) was identified in this cluster in 1971 \\citep{Markert75}, and the cluster was detected at a much lower flux level in 1980 \\citep{Hertz83}, identified with quiescent emission from the luminous X-ray source.  An X-ray source in the cluster (SAX J1748.9-2021) was observed in outburst again in 1998 \\citep{intZand99}, 2001 \\citep{intZand01}, and 2005 \\citep{Markwardt05}.  \\citet{Pooley02b} observed the cluster with \\Chandra\\ in 2000, identifying 24 sources within 2 core radii of the cluster.  One (CX1) was positionally identified with a blue variable optical counterpart during the 1998 outburst \\citep{Verbunt00}, and with the outbursting LMXB in 2001 \\citep{intZand01}.  \\citet{Altamirano07} identified intermittent 442 Hz pulsations in both the 2001 and 2005 outbursts (the latter also identified by \\citealt{Gavriil07}). Thus, CX1 can be confidently held responsible for the 1998, 2001, and 2005 outbursts.  We here identify a second transient luminous LMXB in NGC 6440, and confidently identify five outbursts from this transient during 2009 using its source position.  Preliminary results were presented in \\citet{Heinke09d,Heinke09f,Heinke09e}; these results supersede those.   \\begin{figure} \\figurenum{1} \\includegraphics[angle=0,scale=.8]{f1.eps} \\caption[6440-3images.ps]{ \\label{fig:image} \\Chandra/ACIS-S images of NGC 6440 during the July 2009 outburst of NGC 6440 X-2 (upper left, 0.5-2.5 keV) and 13 days later (upper right, 0.3-7 keV).  Below, we show a merged image from quiescent epochs in 2000 and 2003 (lower left, 0.3-7 keV), and an image from the 2001 outburst (lower right, 0.3-7 keV). The positions of NGC 6440 X-2 (red circle, 1'') and the cluster core (0.13') and half-mass (0.58') radii (large circles) are indicated.  North is up, east to the left. } \\end{figure} \\clearpage ",
        "conclusions": "\\label{s:discuss}  The discovery of millisecond pulsations \\citep{Altamirano09} clearly identifies NGC 6440 X-2 as a neutron star LMXB system, and provides the orbital period of 57.3 minutes. The ultracompact nature of NGC 6440 X-2 confirms the trend \\citep{Deutsch00} for globular clusters to have a higher fraction of ultracompact LMXBs than field systems, indicating a different formation mechanism \\citep{Ivanova05}.  However, it goes against the apparent trend \\citep{Zurek09} for ultracompacts in globular clusters to have shorter periods than ultracompacts in the field. The lack of optical detections is not surprising, when the ultracompact orbit is considered.  Using the relation between absolute magnitude, orbital period, and X-ray luminosity derived by \\citet{vanParadijs94}, we predict $M_V=3.8$ and thus $V=21.8$ at peak.  The scatter in this relation is about 1.5 magnitudes, and our optical observations all occurred during the outburst decays, so our limits on NGC 6440 X-2 are consistent with this relation.  NGC 6440 X-2's outbursts are unusual among globular cluster LMXBs, both for the faintness of the outbursts (peak 0.5-10 keV $L_X=1.5\\times10^{36}$ \\ergss\\ for the July outburst, $2.8\\times10^{36}$ \\ergss\\ for the August outburst), but more importantly for their brevity.  The July X-ray lightcurve indicates that the time spent above $10^{35}$ \\ergss\\ was no more than 4 days, and perhaps only 2.5 days (Fig. \\ref{fig:inset}), one of the shortest transient LMXB outbursts so far recorded \\citep[cf.][]{Natalucci00,intZand04,Wijnands07}. The August outburst lightcurve indicates $<3.5$ days spent above $10^{35}$ \\ergss, and the lightcurves from the other outbursts, while less constraining, are consistent with this timescale (Fig. \\ref{fig:inset}).  We note that another ultracompact AMXP, XTE J1751-305, has shown  similarly short and faint outbursts \\citep{Markwardt07,Linares07,Markwardt09}.  Why does NGC 6440 X-2 show such faint and brief outbursts?  The obvious drivers are accretion disk instabilities or magnetospheric instabilities.  However, magnetospheric instabilities occur on much faster timescales (e.g. the Rapid Burster, Lewin 1993), and would require that the disk remain viscous (and thus ionized) between outbursts, which seems unlikely.  Standard accretion disk instability models (e.g. Lasota 2001) predict quiescent periods 10 times longer, and brighter and longer outbursts, for systems with 57-minute orbital periods.  Naively, we expect longer intervals between outbursts for a 57-minute system than for the other known ultracompact systems, if NGC 6440 X-2 is indeed an evolutionary descendant of systems like them, as it will have a larger disk and lower mass transfer rate (e.g. Deloye \\& Bildsten 2003).   However, we are not aware of detailed modeling of outbursts of hydrogen-poor accretion disks with extreme mass ratios (and thus likely superhumps, Whitehurst 1988), suggesting an avenue for further study.  The X-ray record suggests that NGC 6440 X-2's activity has significantly increased in the past year.  It would be difficult to attribute this change to the disk (as it includes numerous outburst cycles), and we suggest that it represents a signal of mass-transfer variations from the donor.  Cyclical variations in the orbital period are well-established in longer-period cataclysmic variables \\citep{Borges08}, LMXBs \\citep{Wolff09}, and black widow pulsars \\citep{Arzoumanian94} and seem likely to be due to magnetic activity in the companion star.  Such activity has also been suggested to explain SAX J1808.4-3658's large current rate of orbital period increase \\citep{Hartman08}, and decades-long variations in mass transfer rates in LMXBs \\citep{Durant09}.  If such a mechanism is active here, it requires a partly nondegenerate, convective companion.  Alternatively, the increased activity could indicate a change in the orbital parameters, induced by a distant companion (making this a triple system; testable with monitoring of future outbursts by RXTE) or a recent close interaction with another star (its position outside the core suggests this is less likely).  It is difficult to estimate NGC 6440 X-2's mass transfer rate, due to its extreme faintness and the limitations of existing surveys.  We identify two limiting cases, one based on its recent outburst history, and one using the full bulge scan light curve. For the first case, we estimate the outbursts as lasting 3 days at an average $L_X\\sim1.5\\times10^{36}$ \\ergss, and occurring every 31 days. For canonical neutron star mass and radius estimates, this gives a time-averaged mass accretion rate of $3\\times10^{-11}$ \\Msun/year.  Although this represents the mass transfer rate over the past few months, it is clear that NGC 6440 X-2 has not shown such outbursts regularly over the entire bulge scan epoch, where only 5 outbursts have been identified.  Assuming (generously) that 2/3 of all outbursts have been missed (the October and November outbursts were missed by bulge scans, and it seems likely that outbursts in June/July and April/May were missed; Fig. \\ref{fig:lcurve}), and that the average outburst is like those seen so far, we estimate NGC 6440 X-2's mass transfer rate over the entire bulge scan epoch (10 years) as  $1.3\\times10^{-12}$ \\Msun/year.  This latter rate is consistent with an ultracompact binary of orbital period 57 minutes experiencing conservative mass transfer driven by general relativistic angular momentum loss \\citep{Deloye03}, though a higher rate is not inconsistent with a relatively high-entropy (partly nondegenerate) donor.  The tight upper limit on NGC 6440 X-2's quiescent emission is the third lowest for any neutron star LMXB, after the transients SAX J1808.4-3658 and 1H 1905+000 \\citep[][; Fig. \\ref{fig:yak_rev}]{Heinke09a,Jonker07}.  Deep \\Chandra\\ observations might substantially improve these limits (e.g. 100 ks could reduce the quiescent flux limit by a factor of 3). Long-term study of outbursts from this system will allow a better measure of the average mass accretion rate.  It will be of great interest to see if the outbursts continue to occur every $\\sim$31 days, turn off, or change their outburst frequency, as this system's behavior is extremely unusual.  \\begin{figure}[h] \\figurenum{8} \\includegraphics[angle=0,scale=.45]{f8.eps} \\caption[optical]{ \\label{fig:yak_rev} Measurements of, or limits on, the quiescent thermal luminosity of various NS transients, compared to estimates of, or upper limits on, their time-averaged mass accretion rates.  Data from compilations of \\citet{Heinke07a,Heinke09a}, with NGC 6440 X-2 added.  Predictions of standard cooling and several enhanced cooling mechanisms are plotted, following \\citet{Yakovlev04}.  Accreting millisecond pulsars are indicated separately (in red), while the effect of increasing the distance by a factor of 1.5 for any system is indicated with an arrow labeled ``D$\\times$1.5''. } \\end{figure}   This is the first globular cluster to show two transiently outbursting X-ray sources.  Many candidate quiescent LMXBs have been identified in globular clusters through their soft spectra, including eight in NGC 6440 \\citep{Grindlay01a,Rutledge02a,Pooley02b, Heinke03d}, although few have been observed to undergo outbursts.  Some of these quiescent LMXBs may be producing short, faint transient outbursts like NGC 6440 X-2's, which are at or near the noise level for existing surveys such as the \\RXTE/PCA bulge scans and All-Sky Monitor.  Even fainter X-ray transients have been studied in the Galactic Center with dedicated observations \\citep{Muno05,Wijnands06}.   \\Swift\\ could efficiently survey one or a few of the globular clusters richest in quiescent LMXBs for such small-scale outbursts."
    },
    "0911/0911.3392_arXiv.txt": {
        "abstract": "We determine an absolute calibration of the initial mass function (IMF) of early-type galaxies, by studying a sample of 56 gravitational lenses identified by the SLACS Survey. Under the assumption of standard Navarro, Frenk \\& White dark matter halos, a combination of lensing, dynamical, and stellar population synthesis models is used to disentangle the stellar and dark matter contribution for each lens. We define an ``IMF mismatch'' parameter $\\alpha\\equiv$\\mseld$/$\\msesed as the ratio of stellar mass inferred by a joint lensing and dynamical models (\\mseld) to the current stellar mass inferred from stellar populations synthesis models (\\msesed). We find that a Salpeter IMF provides stellar masses in agreement with those inferred by lensing and dynamical models ($\\langle \\log \\alpha\\rangle=0.00\\pm 0.03\\pm0.02$), while a Chabrier IMF underestimates them ($\\langle \\log \\alpha\\rangle=0.25 \\pm 0.03\\pm0.02$).  A tentative trend is found, in the sense that $\\alpha$ appears to increase with galaxy velocity dispersion. Taken at face value, this result would imply a non universal IMF, perhaps dependent on metallicity, age, or abundance ratios of the stellar populations.  Alternatively, the observed trend may imply non-universal dark matter halos with inner density slope increasing with velocity dispersion. While the degeneracy between the two interpretations cannot be broken without additional information, the data imply that massive early-type galaxies cannot have both a universal IMF and universal dark matter halos. ",
        "introduction": "The initial mass function (IMF) is a fundamental characteristic of a simple stellar population. Measuring the IMF from resolved stellar populations in the local universe has been a major astrophysical problem for decades \\citep[e.g.,][]{Sal55,Cha03}.  From the point of view of understanding distant unresolved stellar populations, the IMF holds the key to interpreting observables such as colors and their evolution in terms of star formation history and chemical enrichment history. Among the many astrophysical problems where the IMF plays a key role, in this paper we will focus on the determination of the stellar mass of galaxies from the comparison of stellar populations synthesis models \\citep[e.g.,][]{B+C03,Mar05} with broad band colors. Stellar masses derived in this manner are often used to construct stellar mass functions and to study the demographics of galaxies over cosmic time \\citep[e.g.,][]{B+E00,Fon++06,Bun++06}.  Unfortunately, the form of the IMF is degenerate with the derived stellar masses, since the luminosity is typically dominated by stars in a relatively narrow mass range. To facilitate comparison between results of different groups, it is therefore common practice to {\\it assume} a standard, universal, and non-evolving IMF. Popular choices are the so-called \\citet{Sal55} and \\citet{Cha03} IMFs.  However, this strategy leaves behind a systematic error that is believed to be of the order of a factor of two in stellar mass. Furthermore, the error need not correspond to a global error on the mean, but it could depend on the conditions of star formation and therefore for example on galaxy type, environment, metallicity and star formation epoch. Recent work, for example, has called into question the universality of the IMF and suggested it may be evolving with cosmic time \\citep[e.g.,][]{vDo08,Dav08}.  The goal of this paper is to determine the absolute normalization of stellar masses for early-type galaxies, its scatter from object to object, and its dependence on secondary parameters such as stellar velocity dispersion or stellar mass, by combining three independent probes of mass. Previous works have attempted to do this by comparing stellar masses determined from stellar populations synthesis models with those inferred from gravitational lensing \\citep[e.g.,][]{FSB08} or stellar kinematics \\citep[e.g.,][]{Pad++04,Cap++06}. However, the combination of the two latter techniques is particularly powerful as it allows one to reduce many of the degeneracies and disentangle the stellar and dark component \\citep[e.g.,][]{T+K04}.  To meet our goal we exploit the large and homogeneous sample of strong gravitational lenses discovered by the Sloan Lenses ACS Survey \\citep{Bol++06, Tre++06, Koo++06, Gav++07, Bol++08a, Gav++08, Bol++08b,Aug++09a}. These papers showed that the SLACS lenses are statistically indistinguishable within the current level of measurement errors from control samples in terms of size, luminosity, surface brightness, location on the fundamental plane, environment, stellar and halo mass. Thus our results can be generalized to the overall population of early-type galaxies.  Recently, \\citet{Gri++09} used ground based photometry to derive stellar masses for a subset of SLACS lenses. They compared the inferred stellar mass fractions with average lensing stellar mass fractions determined by \\citet{Koo++06} and \\citet{Gav++07} for subsamples of the SLACS lenses, finding a general agreement. The analysis presented here takes several steps forward: i) space-based HST photometry \\citep{Aug++09a} extending into the near infrared is used for stellar masses; ii) individual stellar fraction estimates from lensing and dynamical models are used for each galaxy; iii) a Bayesian framework is adopted to determine the IMF normalization and errors for each galaxies. This progress allows us to investigate for the first time the scatter of the IMF and its possible dependency on galaxy velocity dispersion and hence, e.g., redshift of formation of the stellar populations or metallicity.  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. Base-10 logarithms are used. ",
        "conclusions": "\\label{sec:conc}  We combined three independent diagnostics of mass (lensing, dynamics and stellar populations synthesis models) to determine an absolute normalization of the IMF for a sample of 56 early-type galaxies, spanning over a decade in stellar mass and a factor of two in velocity dispersion, under the assumption of a NFW dark matter density profile. On average, the absolute IMF normalization is found to be close to that of a Salpeter IMF and larger than that for a Chabrier IMF. Using the prescription outlined in this paper, stellar masses based on any stellar population synthesis models can be absolutely calibrated to better than 20\\%, a significant progress with respect to the range of a factor of $\\sim$2 spanned by standard choices of the IMF.  A tentative trend of IMF mismatch parameter $\\alpha$=\\mseld/\\msesed with galaxy velocity dispersion is found. Two possible explanations, not necessarily mutually exclusive, are suggested for the observed trend:  \\begin{itemize}  \\item The IMF is not universal, but rather depends on parameters such as metallicity, age, and abundance ratios of the stellar populations. In order to fully explain the observed trend, the normalization of the IMF and thus the abundance of low mass stars, must increase from Chabrier-like for $\\sigma~\\sim 200$~\\kms\\, to Salpeter-like for the most massive early-type galaxies.  \\item Dark matter halos are not universal. For a uniform Chabrier IMF, the observed trend of $\\alpha$ with $\\sigma$ could be explained if the inner slope of the dark matter halo were systematically steeper than NFW for the high velocity dispersion systems. For a uniform Salpeter IMF, the the dark matter halos would have to be NFW at the high mass end and flatter at lower masses.  \\end{itemize}  In conclusion, the data are inconsistent with both a universal IMF and universal NFW dark matter halos over the mass range probed by the SLACS sample. There is a fundamental degeneracy between the two interpretations that cannot be broken with the current dataset.  Further tests and more work are required to verify and extend our perhaps surprising results. Firstly, we need to extend our samples to cover a wider range of redshifts, masses, and morphological types and thus probe a larger variety of IMFs. Secondly, we need to use spectral stellar population diagnostics to obtain independent constraints on the stellar mass to light ratios as well as on the physical parameters that may correlate with IMF normalization. Thirdly, we need to improve the constraints on the inner slope of the dark matter halo as a function of velocity dispersion to break the current degeneracy. At the level of individual galaxies, some progress can be achieved by constructing more sophisticated lensing and dynamical models \\citep[e.g.,][]{Bar++09}. In fact, the models presented here only use a single measurement of stellar velocity dispersion from SDSS and the total mass enclosed by the Einstein Radius, while more radial information can be extracted from both diagnostics. At the level of joint analysis of sub samples of galaxies, we need to have enough objects so that they can be binned by velocity dispersions to perform a weak lensing analysis. The addition of weak lensing to the strong lensing, dynamics, and stellar populations diagnostics would allow us to probe systematic variations with velocity dispersion of the dark matter halo shape and of the stellar to virial mass to light ratio . If the amount of contraction is an important ingredient of the observed trend with velocity dispersion, we expect to see a parallel trend in the overall efficiency of converting baryons into stars, or perhaps in the spatial concentration of the stellar component relative to that of the halo.  \\medskip"
    },
    "0911/0911.1732_arXiv.txt": {
        "abstract": "The giant {\\sc Hii} region NGC\\,604 constitutes a complex and rich population to studying detail many aspects of massive star formation, such as their environments and physical conditions, the evolutionary processes involved, the initial mass function for massive stars and star-formation rates, among many others. Here, we present our first results of a near-infrared study of NGC\\,604 performed with NIRI images obtained with {\\it Gemini North}. Based on deep {\\it J\\,H\\,K} photometry, 164 sources showing infrared excess were detected, pointing to the places where we should look for star-formation processes currently taking place. In addition, the color-color diagram reveals a great number of objects that could be giant/supergiant stars or unresolved, small, tight clusters. A extinction map obtained based on narrow-band images is also shown.\\\\ ",
        "introduction": "NGC\\,604 is a giant {\\sc Hii} region (GHR) located in an outer spiral arm of M33, at a distance of 840 kpc. It is the second most luminous {\\sc Hii} region in the Local Group, after 30 Doradus in the LMC. Both are nearby examples of giant star-forming regions whose individual objects can be spatially resolve for further study. NGC\\,604 has been the target of many studies during the past few decades. A brief summary of known facts about NGC\\,604 includes the following. This GHR is ionized by a massive, young cluster, with at least 200 O stars (some as early as O3-O4). The cluster does not exhibit a central core distribution. Instead, the stars are widely spread over its projected area in a structure called 'scaled OB association' (SOBA; \\cite{Hunter_1996}; \\cite{Maiz_2004}; \\cite{Bruhweiler_2003}). Wolf-Rayet stars, a confirmed and many candidate red supergiant stars, a luminous blue variable and a supernova remnant are all part of NGC\\,604's stellar population (\\cite{Conti_1981}; \\cite{Dodorico_1981}; \\cite{Drissen_1993}; \\cite{Diaz_1996}; \\cite{Churchwell_1999}; \\cite{Terlevich_1996}; \\cite{Barba_2009}). The age of the central ionizing cluster has  been determined by different authors as between 3 and 5 Myr (\\cite{Gonzalez_2000}; \\cite{Bruhweiler_2003}; \\cite{Diaz_1996}; \\cite{Hunter_1996}; \\cite{Relano_2009}).\\\\ The interestellar medium reveals a complex structure with a high-excitation central region (made up of multiple two-dimensional structures), asymmetrically surrounded by a low-excitation halo. The whole region shows a very complex geometry of cavities, expanding shells and filaments, as well as dense molecular regions. All of these structures show different kinematic behaviour (\\cite{Maiz_2004}; \\cite{Tenorio-Tagle_2000}; \\cite{Sabalisck_1995}; \\cite{Relano_2009}). \\\\ Aiming to characterize the youngest stellar population and its environment, we performed near-infrared (NIR) photometry ({\\it J\\,H\\,K}) and analysed narrow-band images in Pa$\\beta$, Br$\\gamma$ and H$_2$(2-1). Taking into account that NIR observations are less affected by dust extinction characteristic in star-fomation environments, we can take a deep dive into NGC\\,604 to study those very young objects which are still immersed in their parental clouds at the sites of current star formation. ",
        "conclusions": "\\begin{figure}[] \\begin{center} \\includegraphics[width=0.49\\textwidth]{CCexcess_rad48_dist_1s.eps} \\includegraphics[width=0.49\\textwidth]{CMexcess_rad48_dist_1s_curvas.eps} \\caption{Color-color {\\it (left)} and color-magnitude {\\it (right)} diagrams for objects observed in a circular area centered on NGC\\,604. Red squares are objects that show an IR excess.} \\label{fig1} \\end{center} \\end{figure}   \\begin{figure}[b] \\begin{center} \\includegraphics[width=0.45\\textwidth]{astro_JHK_IRstars_radio.AAA.1.eps} \\includegraphics[width=0.53\\textwidth]{astro_IAUBr_Ha.ps} \\caption{{\\it (left)} Three-color (J, H and K$_s$) composite image of NGC\\,604 with 8.4 GHz radio-continuun countours overlaid (adapted from \\cite{Churchwell_1999}). Red circles are objects that show IR excess in our photometry. {\\it (right)} Extinction map  from Br$\\gamma$/H$\\alpha$, regions with higher extinction are in red/white.} \\label{fig2} \\end{center} \\end{figure}  The resulting color-magnitude (CM) and color-color (CC) diagrams are shown in Figures 1 and 2, respectively. We have included $\\sim$ 2000 selected objects, located within a radius of 48 arcsec ($\\sim$ 200 pc) and centered on  NGC\\,604, meeting a certain photometric quality level (magnitude error $\\leq$ median(error)+ 1.0 $\\times$ standard deviation(error)). Red symbols in both plots are objects that lie on the right side of the reddening line for a O6-O8 V star.For each of these objects, the error in (H-K$_s$) color is smaller than its distance to the reddening line, so that we can ensure that they undoubtedly show an IR excess.\\\\ Among the objects with intrinsic IR excesses in GHRs we expect to find Wolf-Rayet stars, early supergiants and Of stars, and massive young stellar objects (MYSOs). Many of these types of objects (or candidates) have already been found in NGC\\,604. Based on our deep {\\it J\\,H\\,K} photometry we found a total of 164 objects that show IR excesses. As can be expected, $\\sim$ 70\\% of these lie in a small area near the region's center and a large proportion are tightly grouped in regions coincident with the radio-continuum peaks at 8.4 GHz from \\cite{Churchwell_1999}, as shown in Figure 3 (left panel). These results are in complete agreement with the analysis of \\cite{Barba_2009} based on HST/NICMOS NIR images, where MYSO candidates also appeared aligned with the radio peak structures.\\\\ What we found supports the idea mentioned by many authors, that those areas may be embedded star-forming regions, also taking into account that regions coincident with the radio knots show conditions of high temperature and density (\\cite{Tosaki_2007}), and in two of them \\cite{Maiz_2004} identified compact {\\sc Hii} regions. Most objects with IR excesses will be the targets for further observing programs (making use of the integral-field spectroscopic facilities at {\\it Gemini North}) to elucidate their nature and derive accurate properties by means of spectroscopic study. Also, those objects that are located at the bright, massive end on the CM diagram deserve further observation and study.\\\\ On the basis of our narrow-band images we generated an extinction map, derived from the observed variations of the Br$\\gamma$ to H$\\alpha$ ratio (with the H$\\alpha$ image taken from \\cite{Bosch_2002}). The color scale used in the map shows regions with higher extinction in red/white. The present and future results of our NIR study will be placed  in the context of previous studies of NGC\\,604 to complete the picture of the overall formation and evolution scenario of this GHR.\\\\  {\\bf Acknowledgements}\\\\ RB acknowledges partial support from the Universidad de La Serena, project DIULS CD08102."
    },
    "0911/0911.1218_arXiv.txt": {
        "abstract": " ",
        "introduction": "Neutrinos from a core-collapse supernova can play an important role in probing both neutrino properties as well as throwing light on the supernova mechanism \\cite{Raff, Dighe:2008dq}. Neutrinos emitted during the explosion of core-collapse supernovae (SN) pass through very large density variation of matter and can undergo  MSW  resonant flavor conversion \\cite{msw}  which can give useful information on neutrino mass hierarchy and the third leptonic mixing angle $\\theta_{13}$. They are also influenced by the shock wave formed in the SN core and thus carry information about the explosion mechanism \\cite{fogli-lisi-mirizzi-montanino-0304056, dasgupta-dighe-0510219, choubey-harries-ross-0605255, fogli-lisi-mirizzi-montanino-0603033,schirito-Fuller-0205390,Tomas2004}.   Recently it was realized that a crucial feature in the study of SN neutrinos comes from the collective neutrino-neutrino interaction at very high densities of the core and this may change the emitted flux of different flavors substantially  \\cite{Pantaleone:1992eq}-\\cite{Fogli:2009rd}. The phenomenology of supernova neutrinos with collective effect including Earth matter effect on SN neutrinos \\cite{Dasgupta:2008my}, prompt SN \\cite{Duan:2007sh, Dasgupta:2008cd}, diffuse neutrino background from SN \\cite{Chakraborty:2008zp}, failed SN \\cite{Lunardini:2009ya} as well as CP violation in SN neutrinos \\cite{Gava:2008rp} has been studied. For a recent review we refer to \\cite{Duan:2009cd}.  It has been shown that effectively the collective evolution of a three-flavor ($\\nu_{e}, \\nu_{\\mu},\\nu_{\\tau} $) system can be treated like a two flavor ($\\nu_{e}, \\nu_{x} $) scenerio, where  $\\nu_{x}$ can be $\\nu_{\\mu}$ or $\\nu_{\\tau}$ or a linear combination of $\\nu_{\\mu}$ and $\\nu_{\\tau}$ \\cite{Dasgupta:2007ws, Fogli:2008fj}. The flavor evolution for this system is driven by the effective mass squared difference $\\Delta m^2$ and the mixing angle $\\theta_{13}$. This two flavor scenerio which has been extensively studied \\cite{Duan:2006an, Raffelt:2007cb, Raffelt:2007xt, Fogli:2007bk}  shows that for inverted hierarchy (IH, $\\Delta m^2 <0$), above a critical energy (split energy $E_{c}$),  the spectrum in both the electron neutrino ($\\nu_{e}$) and antineutrino ($\\bar\\nu_{e}$) sectors end up with a complete exchange or swap with $\\nu_{x}$ and  $\\bar\\nu_{x}$ respectively, this is referred as ``spectral swap''. Recently the studies in \\cite{Dasgupta:2009mg, Fogli:2009rd} analyzed the role of equipartition in energy and variation of luminosity and showed the interesting possibility of multiple splits in the supernova neutrino spectrum for IH. Single spectral split for Normal Hierarchy (NH, $\\Delta m^2 > 0$) was also reported for certain values of luminosities.  In this paper our goal is to study  (a) the impact of spectral splits on the electron fraction ($Y_e$) which is a diagnostic of successful r-process nucleosynthesis in supernova and (b) the inverse problem which is to see if it is possible to put any constraint on  the luminosities by demanding the neutron rich condition on $Y_e$ which is required for synthesis of heavy nuclei.  Though the site for r-process nucleosynthesis is not known definitely, supernovae are considered to be excellent candidates for it. One of the criteria for the rapid nucleosynthesis to take place is that it has to be in a neutron-rich region. With the two competing beta processes $n+\\nu_e \\rightarrow p+e^{-}$ and $p+\\bar\\nu_e \\rightarrow n+e^{+}$ occurring in the hot bubble and neutrino driven wind region, the minimal condition is that the electron fraction, $Y_e <0.5$. A more realistic constraint may be $Y_e <0.45$.  Effect of MSW neutrino oscillations on r-process nucleosynthesis were considered in \\cite{Qian:1993dg}, and subsequently in a large number of papers. Since r-process is expected to take place in the neutrino driven wind deep inside the supernova, the matter density in these regions obviously are high. Therefore, in order to have MSW effect one requires mass squared difference $\\Delta m^2 \\sim 1$ eV$^2$ or higher. Such large values of  $\\Delta m^2$ are possible only if one allows for sterile neutrinos. For only active neutrinos, the mass squared differences are of the order of $10^{-3}$ eV$^2$ (atmospheric) and $10^{-5}$ eV$^2$ (solar) only, and the MSW resonance for these are reached at distances far beyond the r-process region. Therefore with only active neutrinos, one expects no effect of MSW oscillations on r-process nucleosynthesis. However, one will have collective effects driven by these active neutrino mass squared differences and these collective flavor oscillations happen very close to the neutrinosphere -- close enough to impact r-process nucleosynthesis. The effect of the collective oscillations on the possibility of getting n-rich region in the hot bubble was studied in \\cite{Pastor:2002we}. The problem of r-process in the neutrino-wind region with different evolution scenarios was looked into in \\cite{Balantekin:2004ug}. With an improved understanding of the spectral splits in the collective oscillations \\cite{Dasgupta:2009mg, Fogli:2009rd}, we re-examine the problem of SN r-process  using the two flavor basis as mentioned above. In particular, recent papers have shown that (multiple) spectral splits could happen in either or both the neutrino and antineutrino channels, depending on the initial spectral conditions. Since spectral splits can change the neutrino and antineutrino spectra and hence the value of $Y_e$, and since these spectral split patterns depend on the initial conditions of the unoscillated spectra, r-process nucleosynthesis can in principle be used to put constraints on the initial spectral conditions. We perform this exercise as a proof of principle, in a simplified framework, in order to illustrate that such constraints indeed exist. We consider different models for the initial neutrino spectrum with different hierarchies among average energies and departures from energy equipartition among flavors. We demonstrate the occurrence  of spectral splits due to these variations on the probabilities and the fluxes for both hierarchies. We use the constraint that the environment should be neutrino rich for successful r-process and delineate the  resulting constraints that are obtained on the fractional luminosities.  The paper is organized as follows. In Section 2, we outline the neutrino-neutrino interaction in the context of the supernova problem. First the two flavor evolution equations are presented. Next we vary the fractional luminosities for different initial energy spectrum model and study the the effect of  collective oscillations, especially the effect of the spectral splits  on the emitted neutrino flux. In Section 3, we discuss how  the evolution of the electron fraction $Y_e$ is affected by collective oscillations. The value of $Y_e$ determines the possibility of having r-process nucleosynthesis in the supernova environment. We impose the constraint that the environment should be neutron rich and study the initial relative fluxes or relative luminosities which are allowed under this constraint. Finally Section 4 makes some concluding remarks. ",
        "conclusions": "Collective flavor oscillations driven by neutrino-neutrino interaction at the very high density region of core collapse supernovae control the emitted flux of neutrinos of different flavors. In the process one or more swaps of flavors for both neutrinos and antineutrinos take place depending on the initial neutrino flux and distributions. We study the phenomena of spectral splits and consequent flavor swaps for different models of neutrino spectrum, varying the relative luminosities of neutrinos and antineutrinos for both normal and inverted mass hierarchy. The effect of spectral splits is found to be more pronounced for inverted hierarchy and depending on the initial luminosities one can get single or dual splits in neutrinos and/or antineutrinos. For a specific choice of relative luminosity we also find a case for inverted hierarchy where the splits are not resolvable numerically and is akin to no spectral splits for all practical purposes. Single split patterns are also obtained for normal hierarchy for some choices of the luminosities. Next we consider the impact of the collective oscillations and the spectral splits on the electron fraction $Y_e$, which determines if the environment is neutron-rich and compatible with r-process nucleosynthesis or not. The minimal requirement for r-process is the electron-to-nucleon ratio $Y_{e}< $ 0.5, but a more favorable condition may be $Y_{e} <$ 0.45/0.40. We consider the flavor evolution reduced to an effective two flavor model with oscillation between $\\nu_{e}$ and $\\nu_{x}$ and their antiparticles.  The oscillation parameters are chosen as $\\Delta m^{2}=3\\times10^{-3}$ eV$^{2}$ and a small effective mixing angle $\\theta=10^{-5}$ in agreement with realistic $1-3$ mixing. For these values of parameters the ordinary MSW resonances take place beyond the r-process region and hence there will be no effect of neutrino conversions on $Y_e$ due to these. However, the inclusion of collective effects can affect the value of $Y_e$ even for these values of mass and mixing parameters.    The electron fraction ($Y_{e}$) as a function of the radius of the core is calculated and it shows an oscillatory behavior in the bipolar region due to collective effects, before saturating to a constant value which depends on the initial luminosities and the pattern of flavor swap. Different models of neutrino energy distributions are used. For each of the distributions initial fluxes of different flavors are varied and constraints on the initial neutrino fluxes consistent with successful r-process nucleosynthesis are shown in exclusion plots for these initial neutrino fluxes. While a detailed simulation of the r-process nucleosynthesis inside the supernova might bring some changes to the exclusion plots, this work illustrates the fact that such exclusion plots are possible to achieve.   The variation in the number of spectral splits with the variation in the luminosity give rise to different possibilities of neutrino and antineutrino spectrum at the detector. The constraints on luminosities obtained by ensuring r-process nucleosynthesis can provide additional inputs in narrowing down  the possible patterns.  \\textit{Note added}: While this paper was under peer review, a few papers \\cite{Friedland:2010, Dasgupta:2010b} have appeared which study the possible effects of a three flavor treatment of the collective oscillations of supernovae neutrinos. The differences of this with a two flavor treatment and their impact on the r-process nucleosynthesis need to be investigated. However from the above studies it is expected that the extra effect due to three flavors may be observable in IH but the changes in NH will be minor."
    },
    "0911/0911.2306_arXiv.txt": {
        "abstract": "The wide-band {\\it Suzaku} spectra of the black hole binary GX 339$-$4, acquired in 2007 February during the Very High state, were reanalyzed. Effects of event pileup (significant within $\\sim 3'$ of the image center) and telemetry saturation of the XIS data were carefully considered. The source was detected up to $\\sim 300$ keV, with an unabsorbed 0.5--200 keV luminosity of $3.8 \\times10^{38}$ erg\\,s$^{-1}$ at 8 kpc. The spectrum can be approximated by a power-law of photon index 2.7, with a mild  soft excess and a hard X-ray hump. When using the XIS data outside $2'$ of the image center, the Fe-K line appeared extremely broad, suggesting a high black hole spin as already reported by Miller et al. (2008) based on the same  {\\it Suzaku} data and other CCD data. When the XIS data accumulation is further limited to $>3'$ to avoid event pileup, the Fe-K profile becomes narrower, and there appears a marginally better solution that suggests the inner disk radius to be $5-14$ times the gravitational radius (1-sigma), though a maximally spinning black hole is still allowed by the data at the 90\\% confidence level. Consistently, the optically-thick accretion disk is inferred to be truncated at a radius $5-32$ times the gravitational radius. Thus, the {\\it Suzaku} data allow an alternative explanation without invoking a rapidly spinning black hole. This inference is further supported by the disk radius measured  previously in the High/Soft state. ",
        "introduction": "\\label{intro} One of the current  issues in black hole (BH) research is to measure the spin parameter, $a$. This is possible if we can estimate the radius of the innermost stable orbit  $R_{\\rm in}$, which decreases from $6 R_{\\rm g}$ ($R_{\\rm g}$ being the gravitational radius) down to 1.235  $R_{\\rm g}$  as $a$ increases from 0 to 1. A traditional method of measuring $R_{\\rm in}$ of black-hole binaries  (BHBs) is to parameterize the optically-thick disk emission component in their X-ray spectra in terms of the disk-blackbody ({\\tt diskBB}) model (Mitsuda et al. 1984). First applied successfully to GX~339-4 (Makishima et al. 1986) and  LMC X-3 (Ebisawa et al. 1993), this method has been calibrated using  BH masses estimated from companion star kinematics, and  confirmed to give reliable (e.g., within  $\\sim 30$\\%) estimates  of $R_{\\rm in}$  (Makishima et al. 2000) if  distance uncertainties can be neglected.  A more modern way to estimate $a$, applicable also to active galactic nuclei, is to utilize iron line profiles, which become broadened and skewed due to stronger relativistic effects as $a$ increases (e.g., Fabian et al. 1989). First found from the Seyfert galaxy MCG--6-30-15  by {\\it ASCA} (Tanaka et al. 1995), the broad Fe-K lines were later  reported in a number of Seyferts using, e.g., {\\it XMM-Newton}  (Nandra et al. 2007). The latest {\\it Suzaku} results on MCG--6-30-15 indicate a high spin parameter of $a > 0.917$ (Miniutti et al. 2007).  Similar high values of $a$ were derived from BHBs (Miller 2007, Miller et al. 2009), including  GX~339$-$4 observed with {\\it XMM-Newton} and {\\it Chandra} (Miller et al. 2004ab). Analyzing the {\\it Suzaku} data of this BHB acquired in 2007 February, Miller et al. (2008), hereafter MEA08, confirmed the broad Fe-K feature, and argued that the object is an extreme Kerr BH with $R_{\\rm in} \\sim R_{\\rm g}$.  The wide-band {\\it Suzaku} spectra of Cyg X-1 in the Low/Hard state (LHS), in contrast, gave $R_{\\rm in} \\sim 15 R_{\\rm g} $ (Makishima et al. 2008), via both the disk emission analysis and the Fe-K line modeling. Likewise, $R_{\\rm in} \\sim 8 R_{\\it g}$ was obtained from the {\\it Suzaku} data of GRO 1655-40 (Takahashi et al. 2008). Also, a  2--20 keV {\\it Tenma} observation of GX 339$-$4 (Makishima et al. 1986)  in the High/Soft state (HSS), with dominant disk emission, gave $R_{\\rm in}= 57\\,d_{8}$  km, or  $5.6\\,Q$ times  $R_{\\rm g}$; here $Q \\equiv d_8/m_7$, $d_8$ is the  distance in 8 kpc (Zdziarski et al. 2004), $m_7$ is the BH mass in units of a typical value of $7~M_\\odot$ (Hynes et al. 2004), and $R_{\\rm in}$ was recalculated applying a correction factor of 1.18 (Kubota et al. 1998; Makishima et al. 2000) and assuming an inclination of $i=30^\\circ$ (Gallo et al. 2004). These results imply  $a \\sim 0$.  To examine these apparent discrepancies on $a$, we reanalyzed the same {\\it Suzaku} data of GX 339-4 as MEA03, and found that the X-ray Imaging Spectrometer (XIS; Koyama et al. 2008) events suffer heavy pileup  and telemetry saturation. These effects, neglected by MEA08,  distort the continuum, and  {\\it indirectly} affect the Fe-K line shape. ",
        "conclusions": "\\label{discussion} We reanalyzed the  {\\it Suzaku}  XIS0 and HXD data of GX 339$-4$ acquired in the 2007 VHS. When the disk emission was modeled with {\\tt diskbb}, the ``blue\" XIS0 spectrum, using $r>2'$, gave a small  Fe-K  inner radius as $R_{\\rm Fe}< 3.6 R_{\\rm g}$ (\\S~\\ref{bluespec}). Although this  apparently reconfirms the high BH spin by MEA08, the Fe line EW is too large (\\S 3.3). In addition, a large-$R_{\\rm Fe}$ ($\\sim 10 R_{\\rm g}$) solution is also allowed when considering disk Comptonization that is usually the case in the VHS. Thus, the Fe line shapes degenerate, depending on the continuum modeling.  To further eliminate the XIS pileup effects, we utilized the ``red\" XIS0 spectrum from $r>3'$  (\\S 3.4). Then, the {\\tt diskbb} and {\\tt compbb} results became more consistent. Importantly, the large-$R_{\\rm Fe}$ solution (e.g., $5.0< R_{\\rm Fe}/R_{\\rm g} <14$ at $1\\sigma$), are preferred, although the small-$R_{\\rm Fe}$ solution still remains valid at the 90\\% confidence level. The implied Fe-K line EW and the center energy are consistent with the reflection parameters (\\S 3.4).   Assuming isotropic emission, the 0.5--200 keV luminosity is $3.8 \\times 10^{38}\\,d_8^2$ erg s$^{-1}$, or $\\sim 0.4\\,d_8^2m_7^{-1}$ times the Eddington limit ($L_{\\rm Ed}$). Since the VHS have so far been observed in other sources over a typical luminosity range of $ (0.2-1) L_{\\rm Ed}$ (Kubota \\& Makishima 2004), this indicates $ 0.5 < d_8^2/m_7 < 2.5$, or $ 0.7/\\sqrt{m_7} < Q < 1.3/\\sqrt{m_7}$. Considering extreme BH masses of $3\\;M_\\odot$ ($m_7= 0.43$) and  $15\\;M_\\odot$ ($m_7= 2.1$), we then obtain  $0.5 < Q < 2$.   The  {\\tt diskbb} radius $R_{\\rm in}/R_{\\rm g}=(5.6^{+1.0}_{-1.5}) \\; Q$, found in \\S~3.4, is in fact a lower limit, because a significant fraction of disk photons will be  Comptonized into the PL by hot electron clouds (Kubota \\& Makishima 2004). Then, the true disk area must be a sum of the {\\tt diskBB} area and that of  the seed photon source. Since the 0.5--200 keV photon number in the PL component is $\\sim5$ times larger than that contained in the 0.5--10 keV {\\tt diskBB} emission, the estimated radius will increase by a factor of  $\\sqrt{1+5}$, to $R_{\\rm in}/R_{\\rm g} =(10-16)\\; Q$. The above estimated uncertainty in $Q$ then yields $R_{\\rm in}/R_{\\rm g} =(5-32)$. Since we derived this range assuming $i = 30^{\\circ}$ (or $\\sqrt{\\cos i }$ = 0.93), uncertainties in $i$ can increase the upper bound, but would not affect the lower bound by more than $\\sim$ 7\\%.  The present  Fe-line and {\\tt diskbb} analyses consistently suggest $R_{\\rm in}/R_{\\rm g} \\gtrsim 5$, and hence $a<0.4$, in contrast to the value of $a = 0.89 \\pm 0.04$ by MEA08. If our interpretation is correct, GX 339-4 is inferred to be spinning only mildly (if at all), like Cyg X-1 (\\S~\\ref{intro}; Miller et al. 2009). In addition, the consistency between the two methods reconfirms that $Q$ is relatively close to unity.  We must admit that the assumption of a single PL continuum down to $\\sim 1$ keV might not be warranted. In that sense, the value of $R_{\\rm in}/R_{\\rm g}= 5.6\\,Q$ (\\S~\\ref{intro}) measured with {\\it Tenma} is more reliable, since it was obtained  in the  HSS where the spectral modeling is much less ambiguous, and by a non-CCD instrument that is free from any pileup. This, together with our estimates on $Q$, further argues against the large spin parameters.  So far, large values of $a$ were reported for GX~339$-$4 based on {\\it Chandra} data, and from  {\\it XMM-Newton} plus {\\it RXTE}  (Miller et al. 2004ab, and Reis et al. 2008). However, these measurements have not been examined for systematic effects due to the continuum choice employed therein. In addition, the former may be limited by the continuum bandwidth, and the latter could be still subject to CCD pileup.  \\smallskip The authors would like to express their thanks to the Suzaku team members and an anonymous referee who gave us valuable comments. The present work was supported by Grant-in-Aid for JSPS Fellows."
    },
    "0911/0911.2240_arXiv.txt": {
        "abstract": "Most black hole binaries show large changes in X-ray luminosity caused primarily by variations in mass accretion rate.  An important question for understanding black hole accretion and jet production is whether the inner edge of the accretion disk recedes at low accretion rate.  Measurements of the location of the inner edge ($R_{\\rm in}$) can be made using iron emission lines that arise due to fluorescence of iron in the disk, and these indicate that $R_{\\rm in}$ is very close to the black hole at high and moderate luminosities ($\\gsim$1\\% of the Eddington luminosity, $L_{\\rm Edd}$). Here, we report on X-ray observations of the black hole GX~339--4 in the hard state by {\\em Suzaku} and the {\\em Rossi X-ray Timing Explorer (RXTE)} that extend iron line studies to 0.14\\% $L_{\\rm Edd}$ and show that $R_{\\rm in}$ increases by a factor of $>$27 over the value found when GX~339--4 was bright.  The exact value of $R_{\\rm in}$ depends on the inclination of the inner disk ($i$), and we derive 90\\% confidence limits of $R_{\\rm in} > 35 R_{g}$ at $i = 0^{\\circ}$ and $R_{\\rm in} > 175 R_{g}$ at $i = 30^{\\circ}$.   This provides direct evidence that the inner portion of the disk is not present at low luminosity, allowing for the possibility that the inner disk is replaced by advection- or magnetically-dominated accretion flows. ",
        "introduction": "Accreting black holes exhibit a variety of spectral states \\citep{mr06}, most often showing strong thermal emission from an optically thick accretion disk \\citep{ss73} when they are at their brightest.  As luminosities drop below a few percent of the Eddington value ($L_{\\rm Edd}$), the level of thermal emission decreases, and the spectrum hardens, signaling a transition to the hard state \\citep{kalemci04}.  This state is also characterized by having a high level of X-ray variability and steady, self-absorbed compact jets, which are often observed in the radio band \\citep{fender01}.  In theory, it is predicted that the cool gas from the accretion disk will evaporate as the mass accretion rate drops \\citep{mlm00}.  The implied increase in the inner radius of the optically thick disk, $R_{\\rm in}$, is critical to the idea that a quasi-spherical advection-dominated accretion flow (ADAF) forms around the black hole \\citep{ny94}.  This model requires that $R_{\\rm in}$ increases to $\\sim$100--1000 $R_{g}$ \\citep{emn97}, where the gravitational radius is $R_{g} = GM/c^{2}$, $G$ and $c$ are constants, and $M$ is the black hole mass.  Recently, observational efforts to test the truncated disk picture have used two techniques:  Modeling the thermal emission from the disk and measuring the iron K$\\alpha$ emission line in the reflection component \\citep{lw88}.  The former technique is challenging because of uncertainties arising from electron scattering, irradiation of the disk, and the torque boundary condition at the disk's inner edge as well as the difficulty of modeling a soft X-ray component that is strongly affected by interstellar absorption.  Modeling the disk's thermal emission to estimate the inner radius ($R_{\\rm in}$) has led to mixed results with some in favor of an increase in $R_{\\rm in}$ \\citep{gdp08,cabanac09} and others in disagreement \\citep{rykoff07,rmf09}.  The iron line provides a measurement of $R_{\\rm in}$ because it is Doppler-broadened due to relativistic motion of the accretion disk material and gravitationally redshifted due to the black hole's gravitational field.  Broad iron lines have been seen from both stellar mass and supermassive black holes \\citep{tanaka95,fabian09}.  One challenge to the relativistic interpretation in the case of stellar mass black holes is that the iron line shape may be affected by scattering in a wind \\citep{lt07,tls09}.  However, arguments in favor of the relativistic interpretation are presented in \\cite{miller07}.  For stellar mass black holes at high luminosities, the iron line profiles have been used extensively, and values of $R_{\\rm in}$ less than the radius of the innermost stable circular orbit for a non-rotating black hole, $R_{\\rm in} < 6$~$R_{g}$, have been measured for several systems \\citep{miller02,miller04,miller08}, including GX~339--4.  The broad iron line has, somewhat surprisingly, persisted into the hard state \\citep{miller06a,tomsick08,hiemstra09}.  However, for GX~339--4, the previous detections of the iron line in the hard state were obtained at luminosities above 1\\% $L_{\\rm Edd}$, and it remains to be seen if there is evolution in the iron line profile at lower luminosities.  GX~339--4 is one of the most active black hole transients known, with 4 major outbursts in the past 7~yr.  The distance to GX~339--4 is $\\sim$8~kpc \\citep{hynes04,zdziarski04}, and the orbital period is 1.7~days \\citep{hynes03,lc06}.  The lower limit on the black hole mass is $>$6\\Msun~\\citep{hynes03,mcm08}, and here, we adopt a value of 10\\Msun, corresponding to $L_{\\rm Edd}\\sim 1.3\\times 10^{39}$ ergs~s$^{-1}$.  GX~339--4 has allowed for some of the most in-depth studies of compact jets \\citep{corbel00,coriat09}.  Here, we use {\\em Suzaku} and {\\em RXTE} observations of GX~339--4 to study the iron line profile at 0.14\\% $L_{\\rm Edd}$. ",
        "conclusions": "These results provide the most direct and quantitative evidence to date for the truncation of the accretion disk for stellar mass black holes in the hard state at low luminosities.  Figure~\\ref{fig:r_vs_l}a compares our constraint on $R_{\\rm in}$ at 0.14\\% $L_{\\rm Edd}$ to previous measurements at higher luminosities.  The data show that the inner edge of the disk moves sharply outward as the luminosity decreases from 1\\% to 0.1\\% $L_{\\rm Edd}$.  In addition, the drop in the iron line EW shown in Figure~\\ref{fig:r_vs_l}b is consistent with the increase in $R_{\\rm in}$ since one expects the line to become weaker if the disk is truncated.  The EW evolution and the fact that the line is well-described by a single component provide support for the interpretation that the narrow line comes from the disk.  Considering the broad lines seen in the brighter part of the hard state and the disk truncation that we see here, the overall evolution is very similar to the theoretical prediction that the disk evaporation will start in the middle of the disk and that the entire inner disk will evaporate when the luminosity reaches $\\sim$0.1\\% $L_{\\rm Edd}$ \\citep{taam08}.  While the study of stellar mass black holes, such as GX~339--4, provides a detailed look at the evolution of $R_{\\rm in}$, there is iron line evidence for truncated disks around supermassive black holes in Active Galactic Nuclei (AGN).  For NGC~4258, \\cite{reynolds09} measure a narrow line with an EW of 45~eV at a luminosity of $10^{-3}$\\% $L_{\\rm Edd}$ and infer that $R_{\\rm in} > 3\\times 10^{3} R_{g}$.  There is also evidence for a truncated disk in NGC~4593, and for this AGN, there is possible evidence for a change in $R_{\\rm in}$ over a period of 5~yr \\citep{mr09}.  Our finding for GX~339--4 is also important because a large increase in $R_{\\rm in}$ is required for the presence of an ADAF, and our result makes the ADAF model viable for the fainter part of the hard state.  However, it is known that the ADAF model does not give a complete physical description of the system since it does not incorporate the compact jet.  These jets have now been detected from GX~339--4 in the hard state both when the disk is truncated (this work) and when the disk is not truncated \\citep{miller06a,tomsick08}.  The possibility that disk truncation occurs at a different mass accretion rate than the transition to the hard state and the turn-on of the compact jet has not been considered in most theoretical work.  This finding strongly constrains models such as the magnetically-dominated accretion flow (MDAF) model \\citep{meier05} where jet production and properties depend on the accretion geometry.   \\begin{figure}[h] \\centerline{\\includegraphics[width=0.45\\textwidth]{fig4_r_vs_l.ps}} \\vspace{0.0cm} \\caption{Iron line measurements for GX~339--4 over a 0.5--100~keV luminosity range from 0.14\\% to 15\\% $L_{\\rm Edd}$. (a) The constraints on the inner radius of the accretion disk, including all GX~339--4 observations made with instruments with energy resolution better than 200~eV for which a determination of $R_{\\rm in}$ was obtained using relativistic reflection models.  The four measurements of $R_{\\rm in}$ in the range 2--5~$R_{g}$ between 1.3\\% $L_{\\rm Edd}$ and 15\\% $L_{\\rm Edd}$ use {\\em Suzaku} \\citep{miller08}, {\\em XMM-Newton} \\citep{miller04,miller06a,reis08}, and {\\em Swift} \\citep{tomsick08}. The grey 90\\% confidence upper limit is a {\\em Swift} measurement where broad features in the reflection component are detected, but the iron line is not clearly detected \\citep{tomsick08}.  The dashed line marks the central value obtained at 12--15\\% $L_{\\rm Edd}$.  The measurements indicate that a large increase in $R_{\\rm in}$ occurs between 0.5\\% $L_{\\rm Edd}$ and 0.14\\% $L_{\\rm Edd}$.  (b) The equivalent width of the iron line vs. luminosity.  The value we obtain at 0.14\\% $L_{\\rm Edd}$ is the lowest value, which is consistent with a drop in the strength of the reflection component due to truncation of the disk. \\label{fig:r_vs_l}} \\end{figure}"
    },
    "0911/0911.3841_arXiv.txt": {
        "abstract": "The Advanced Technology Large Aperture Space Telescope (ATLAST) is a set of mission concepts for the next generation UV-Optical-Near Infrared space telescope with an aperture size of 8 to 16 meters. ATLAST, using an internal coronagraph or an external occulter, can characterize the atmosphere and surface of an Earth-sized exoplanet in the Habitable Zone of long-lived stars at distances up to $\\sim$45 pc, including its rotation rate, climate, and habitability. ATLAST will also allow us to glean information on the nature of the dominant surface features, changes in cloud cover and climate, and, potentially, seasonal variations in surface vegetation. ATLAST will be able to visit up to 200 stars in 5 years, at least three times each, depending on the technique used for starlight suppression and the telescope aperture. More frequent visits can be made for interesting systems. ",
        "introduction": "We are at the brink of answering two paradigm-changing questions: Do Earth-sized planets exist in the Habitable Zones of their host stars? Do any of them harbor life? The tools for answering the first question already exist (e.g., Kepler, CoRoT); those that can address the second can be developed within the next 10-20 years (Kasting, Traub et al. 2009).  ATLAST is our best option for an extrasolar life-finding facility and is consistent with the long-range strategy for space-based initiatives recommended by the AAAC Exoplanet Task Force \\citep{Lunine2008}. ATLAST is a NASA Astrophysics Strategic Mission Concept for the next generation flagship UVOIR space observatory (wavelength coverage: 110 nm -- 2400 nm), designed to answer some of the most compelling astronomical questions, including ``Is there life elsewhere in the Galaxy?'' The ATLAST team investigated two different observatory architectures: a telescope with an 8-m monolithic primary mirror and two variations of a telescope with a large segmented primary mirror (9.2-m and 16.8-m). The two architectures span the range in viable technologies: e.g., monolithic vs. segmented apertures, Ares V launch vehicle vs. EELV, and passive vs. fully active wavefront control. This approach provides several pathways to realize the mission that would be narrowed to one as the needed technology development progresses and the availability of launch vehicles is clarified.  While ATLAST requires some technology development, all the observatory concepts take full advantage of heritage from previous NASA missions, as well as technologies available for missions in development.  The 8 m monolith architecture is similar to the Hubble Space Telescope (HST), although the optical design is different.  The 8 m mirror has an areal density exceeding that of the HST mirror, providing superb stiffness and thermal inertia.  The use of such a massive mirror provides excellent wavefront quality and is made possible given a launch vehicle with the capabilities of the proposed Ares V.  The 9.2 m and 16.8 m segmented mirror concepts rely heavily on design heritage from the James Webb Space Telescope (JWST), in development of lightweight segmented optics, the OTA deployment mechanics, and wavefront sensing and control.  The 9.2 m segmented concept can be launched on a Delta IV Heavy launch vehicle with a modified 6.5 m fairing; the 16.8 m concept requires an Ares V. The non-cryogenic nature of ATLAST makes the construction and testing of the observatory much simpler than for JWST.  We have also identified departures from existing NASA mission designs to capitalize on newer technologies, minimize complexity, and enable the required improvements in performance.  Four significant drivers dictate the need for a large space-based telescope if one wishes to conduct a successful search for biosignatures on exoplanets. First, and foremost, Earth-mass planets are faint -- an Earth twin at 10 pc, seen at maximum elongation around a G-dwarf solar star, will have V $\\sim$ 29.8 AB mag. Detecting a biosignature, such as the presence of molecular oxygen in the exoplanet's atmosphere, will require the ability to obtain direct low-resolution spectroscopy of such extremely faint sources. Second, the average projected angular radius of the Habitable Zone (HZ) around nearby F,G,K stars is less than 100 milli-arcseconds (mas). One thus needs an imaging system capable of angular resolutions of $\\sim$10 to 25 mas to adequately sample the HZ and isolate the exoplanet point source in the presence of an exo-zodiacal background. Third, direct detection of an Earth-sized planet in the HZ requires high contrast imaging, typically requiring starlight suppression factors of 10$^{-9}$ to 10$^{-10}$. Several techniques \\citep{Levine2009} are, in principle, capable of delivering such high contrast levels but all require levels of wavefront stability not possible with ground-based telescopes because the timescale of wavefront variations induced by the Earth's atmosphere is shorter than the time to measure the wavefront error to the required precision. A space-based platform is required to achieve the wavefront stability that is needed for such high contrast imaging. {\\footnote {Such stability is achievable regardless of the space telescope's aperture (within reasonable limits) -- so this driver is primarily for the space locale of the telescope rather than its aperture.}} Lastly, biosignature-bearing planets may well be rare, requiring one to search tens or even several hundred stars to find even a handful with compelling signs of life. The number of stars for which one can obtain an exoplanet's spectrum at a given SNR and in less than a given exposure time scales approximately as D$^3$, where D is the telescope aperture diameter.  This is demonstrated in Figure 1 where we have averaged over different simulations done using various starlight suppression options (internal coronagraphs of various kinds as well as an external occulter). To estimate the number of potentially habitable worlds detected, one must multiply the numbers in Figure 1 by the fraction of the stars that have an exoplanet with detectable biosignatures in their HZ ($\\eta_{\\oplus}$). The value of $\\eta_{\\oplus}$ is currently not constrained but it is not likely to be close to unity. One must conclude that to maximize the chance for a successful search for life in the solar neighborhood requires a space telescope with an aperture size of at least 8 meters.  \\setcounter{figure}{0} \\begin{figure} \\hskip 10pt \\includegraphics[height=1.6in]{o_postman_1_f1.jpg} \\caption{\\footnotesize The number of F,G,K stars as a function of telescope aperture where an R=70 SNR=10 spectrum could be obtained of an Earth-twin in the HZ in less than 500 ksec. These counts assume every star has an Earth twin.} \\end{figure}  All ATLAST concepts require many of the same key technologies. We believe these designs compose a robust set of options for achieving the next generation of UVOIR space observatory in the 2020 era. There are several fundamental features common to all our designs. All ATLAST concepts are designed to operate in a halo orbit at the Sun-Earth L2 point. The optical designs are diffraction limited at 500 nm (36 nm RMS WFE) and the optical telescope assembly (OTA) operates near room temperature (280 K -- 290 K). All OTAs employ two, simultaneously usable foci:  a three-mirror anastigmat channel for multiple, wide field of view instruments, and a Cassegrain channel for high-throughput UV instruments and instruments for imaging and spectroscopy of exoplanets (all designs have an RMS WFE of $<5$ nm at $<2$ arcsec radial offset from Cass optical axis). The ATLAST concept study is available on line and on the astro-ph archives \\citep{ATLAST1,ATLAST2}. ",
        "conclusions": ""
    },
    "0911/0911.1075_arXiv.txt": {
        "abstract": "We investigate to what extent the spin axes of stars in young open clusters are aligned. Assuming that the spin vectors lie uniformly within a conical section, with an opening half-angle between $\\lambda=0^{\\circ}$ (perfectly aligned) and $\\lambda=90^{\\circ}$ (completely random), we describe a Monte-Carlo modelling technique that returns a probability density for this opening angle given a set of measured $\\sin i$ values, where $i$ is the unknown inclination angle between a stellar spin vector and the line of sight. Using simulations we demonstrate that although azimuthal information is lost, it is easily possible to discriminate between strongly aligned spin axes and a random distribution, providing that the mean spin-axis inclination lies outside the range $45^{\\circ}$--$75 ^{\\circ}$.  We apply the technique to G- and K-type stars in the young Pleiades and Alpha Per clusters. The $\\sin i$ values are derived using rotation periods and projected equatorial velocities, combined with radii estimated from the cluster distances and a surface brightness/colour relationship. For both clusters we find no evidence for spin-axis alignment: $\\lambda=90^{\\circ}$ is the most probable model and $\\lambda > 40^{\\circ}$ with 90 per cent confidence. Assuming a random spin-axis alignment, we re-determine the distances to both clusters, obtaining $133\\pm 7$\\,pc for the Pleiades and $182\\pm 11$\\,pc for Alpha Per. If the assumption of random spin-axis alignment is discarded however, whilst the distance estimate remains unchanged, it has an additional $^{+18}_{-32}$ percent uncertainty. ",
        "introduction": "Most authors considering the statistics of orbital or rotational stellar motion assume, where relevant, that angular momentum vectors are randomly orientated.  It is possible however that the physical processes of star formation lead to a preferred axis of rotation over the scale of an individual star forming region (SFR). This might arise if the direction of average angular momentum of the molecular cloud giving rise to the SFR has a significant influence on the resulting angular momentum of individual stars -- for instance, if gas were constrained to collapse along strong, large-scale magnetic fields threading the cloud.  Historically, little has been discussed either theoretically or observationally about the possibility of spin-axis alignment during star formation. This would require a relatively undisturbed collapse along magnetic field lines with little disruption from turbulence or dynamical interactions (e.g. Shu, Adams \\& Lizano 1987).  Using circumstellar disc orientation as a proxy, some studies have suggested preferential spin alignment with the ambient magnetic field in SFRs (e.g. Tamura \\& Sato 1989; Vink et al. 2005), but others have found no evidence for disc axis alignment (M\\'enard \\& Duch\\^ene 2004).  A second important reason for assessing the degree of spin-axis alignment is that measurements of projected equatorial rotation velocities ($v \\sin i$, where $i$ is the unknown inclination of the spin-axis to the line of sight) and rotation periods can be combined to provide a powerful method to determine the distances (Hendry, O'Dell and Collier-Cameron 1993; Jeffries 2007a; Baxter et al. 2009), radii (Jackson, Jeffries \\& Maxted 2009) or star formation histories (Jeffries 2007b) of young clusters. Such statistical analyses {\\it must assume} that the orientation of spin axes are random.  If the spin axes in an individual cluster were in fact partially aligned, this would change the intrinsic $\\sin i$ distribution for the cluster, producing biased estimates of distances, radii and age spreads.  In this paper we use published rotation data for the young Pleiades and Alpha Per open clusters to investigate to what extent spin axes may be aligned once the star formation process has finished.  In Section 2 we discuss how well spin-axis orientation can be determined using measured rotation periods and projected equatorial velocities. In Section 3 we present a parameterised model for the observed $\\sin i$ distribution resulting from a group of stars with partially aligned spin axes and show how well such a model can be used to determine the degree of alignment from simulated datasets. In Section 4 we compare the models with measured $\\sin i$ distributions for G- and K-type stars in the Pleiades and Alpha Per clusters.  Section 5 presents and discusses the results of our analyses, including new, independent estimates of the distances to these clusters under the assumption that the spin axes {\\it are} randomly aligned, and in Section 6 we give our conclusions. ",
        "conclusions": "The results in this paper show that analysis of the measured $\\sin i$ values for groups of as few as 36 stars in a cluster can be used to investigate whether their spin axes are randomly oriented or well aligned in space. The information that can be gained is incomplete because (a) information is lost in the measurement technique, principally in the azimuthal direction of the spin axes and (b) there are significant uncertainties in the measurements of parameters used to determine $\\sin i$, which blur its distribution function.  The simulations in section 3.4 and Fig.~4 indicate that a cluster of stars with random spin-axis orientation can be identified from the strong increase in probability density with cone angle. In some cases it is easy to distinguish this from a well-aligned distribution of spin axes, but it depends on the average inclination, $\\alpha$, of the aligned spin axes. If $\\alpha \\leq 45^{\\circ}$ then a well aligned distribution produces a distinctive probability density that falls to zero for $\\lambda<90^{\\circ}$. However, if $45< \\alpha< 75^{\\circ}$ then the probability density of $\\lambda$ becomes quite flat, giving no clear indication of the underlying distribution of spin axes.  The simulations presented in section~3.4 also show that there are only modest gains to be made by observing very large samples of stars The exception would be a set of stars with a well-aligned distribution with a high mean inclination  ($\\alpha \\geq 75^{\\circ}$). In this case a relatively large sample ($\\approx 100$ stars) is required to clearly identify the high degree of alignment.  The first application of this technique using data for 36 stars in the Pleiades and Alpha Per is encouraging. For both clusters there is a clear increase in probability density with cone angle, entirely consistent with random spin-axis orientation.  However, we cannot rule out partial alignment of the spin axes, but place 90 per cent lower limits of about $\\lambda > 40^{\\circ}$ for both clusters. Values of $\\lambda$ that are much less than $90^{\\circ}$ could only be disguised in the data if the mean inclination were between $45^{\\circ}$ and $75^{\\circ}$ .  Our results also show that the analysis method is relatively insensitive to the exact levels of measurement uncertainty and small changes in the assumed mean distance to the cluster. These results show that for the method of analysis to be effective the average distance to the cluster taken from independent measurements needs to be well defined with a normalised uncertainty of a few per cent or better.   Should it turn out to be possible to assume that spin axes are always randomly directed within a cluster, then this allows measurements of average stellar distances, radii and age spreads. In this paper we have used this technique to derive new distance estimates for the Pleiades and Alpha Per of $133\\pm 7$\\,pc and $182\\pm 11$\\,pc respectively. These values are in good agreement with the main sequence fitting distances to these clusters. The Pleiades result is marginally higher than the Hipparcos parallax-based distance.  If larger samples of stars could be measured, then unlike the test for random spin alignment, there {\\it are} significant gains (improving by a factor of approximately $\\sqrt{N}$) to be made in the precision of the mean $\\sin i$ value, which in turn determines the statistical precision of the distance estimate. As there are far more $v \\sin i$ measurements than known periods in both the Pleiades and Alpha Per, it would be a fruitful project to search for more rotation periods in both these clusters. However, if spin-axis orientation cannot be assumed random, then although the estimated distance remains unbiased when analysed under the assumption of randomness, there are significant systematic errors of up to $^{+18}_{-32}$ percent, which render this technique ineffective as an independent means of estimating stellar radii or cluster distance.  Well defined distances should soon become available for many more open clusters following the launch of the GAIA satellite. In addition periods are now being measured systematically in nearby clusters (e.g. Aigrain et al. 2006). As these data become available it will be interesting to compare the distances derived from satellite based parallax measurement with those estimated from the measured period and rotational velocity. If all results agree within measurement errors then this would support the general assumption \\textit{that spin-axes are} randomly orientated in space over the scale length of a cluster. This is because it would become increasingly unlikely that the spins in all clusters were aligned, and had a mean inclination in the range $45^{\\circ}$--$75^{\\circ}$. Conversely, if the spin axes in some clusters are well aligned there would be a reasonable chance ($\\simeq 55$ per cent if we can assume that the mean inclination from cluster-to-cluster is a random variable) of observing a cluster with an average inclination of $\\leq 45^{\\circ}$ or $\\geq 75^{\\circ}$ in which case significant alignment could be identified using the method described in this paper."
    },
    "0911/0911.0009_arXiv.txt": {
        "abstract": "Kuiper belt object 136108 Haumea is one of the most fascinating bodies in our solar system. Approximately $2000\\times1600\\times1000$ km in size, it is one of the largest Kuiper belt objects (KBOs) and an unusually elongated one for its size. The shape of Haumea is the result of rotational deformation due to its extremely short 3.9-hour rotation period. Unlike other 1000 km-scale KBOs which are coated in methane ice the surface of Haumea is covered in almost pure H$_2$O-ice. The bulk density of Haumea, estimated around 2.6~g~cm$^{-3}$, suggests a more rocky interior composition, different from the H$_2$O-ice surface.  Recently, Haumea has become the second KBO after Pluto to show observable signs of surface features. A region darker and redder than the average surface of Haumea has been identified, the composition and origin of which remain unknown.  I discuss this recent finding and what it may tell us about Haumea. ",
        "introduction": "The Kuiper belt is currently the observational frontier of our solar system. Presumably the best kept remnants of the icy planetesimals that formed the outer planets, Kuiper belt objects (KBOs) have been the subjects of intense study in the past $\\sim$15 years. One intriguing KBO is 136108 Haumea (formerly 2003~EL$_{61}$). First famous for its super-fast rotation and elongated shape, Haumea went on to surprise us with a host of interesting properties.  Haumea's spin frequency of one rotation every $\\sim3.9$~hr is unparalleled for an object this large \\citep{2008ssbn.book..129S}. Its shape is rotationally deformed into a $2000\\times1600\\times1000$~km triaxial ellipsoid \\citep{2006ApJ...639.1238R} to balance gravitational and centripetal accelerations. To attain such a fast rotation, Haumea might have suffered a giant impact at the time when the Kuiper belt was massive enough to render such events likely.  Infrared spectroscopy has revealed a surface covered in almost pure H$_2$O ice \\citep{2007ApJ...655.1172T} which gives Haumea an optically blue colour \\citep{2007AJ....133..526T}. The surfaces of the remaining Pluto-sized KBOs (Eris, Pluto and Makemake) are covered in CH$_4$ ice instead, granting them the tag `Methanoids'.  Two satellites were discovered in orbit around Haumea \\citep{2006ApJ...639L..43B}, the largest of which is also coated in even purer H$_2$O ice \\citep{2006ApJ...640L..87B}. The two satellites have nearly coplanar orbits with fast-evolving, complex dynamics due mainly to tidal effects from the primary \\citep{2009AJ....137.4766R}.  Haumea's bulk density, derived assuming it is near hydrostatic equilibrium, is $\\rho\\sim2.6$~g~cm$^{-3}$ \\citep{2007AJ....133.1393L}.  The surface material has density $\\rho\\sim1$ in the same units implying that the interior must be differentiated and Haumea must have more rock-rich core.  A number of KBOs showing signs of H$_2$O ice in their surface spectra all lie close to Haumea in orbital space \\citep{2008ApJ...684L.107S}; this, plus the unusually fast spin, the differentiated inner structure and the two small satellites also covered in H$_2$O ice, all have been taken as evidence that Haumea is the largest remnant of a massive collision that occured $>1$~Gyr ago \\citep{2007AJ....134.2160R,2007Natur.446..294B}.  However, several potential members of the collisional family have been eliminated based on infrared photometry (Snodgrass \\etal, poster at this meeting).   \\begin{figure}[ht] \\begin{center} \\includegraphics[width=\\textwidth]{xxx20091029a_f1} \\caption{\\textbf{a)} Lightcurve of Haumea in two filters. Both $B$ and $R$ data were taken over 3 nights to ensure repeatability. The effect of the dark red spot is apparent at rotational phases  $0.7\\lesssim\\phi\\lesssim1.0$: the maximum and minimum that bracket that region appear darker and the $B$-band flux is consistently lower than the $R$-band flux indicating the spot is redder than elsewhere.  We measure a lightcurve period $P=3.9155\\pm0.0001$ hours and a photometric range $\\Delta m=0.29\\pm 0.02$~mag. The rotationally averaged colour is $B-R=0.972\\pm0.017$~mag. Best fit Jacobi ellipsoid models are overplotted: a thick solid grey line shows how the uniform surface model fails to fit the dark red spot region, and thinner lines show that a small ($S<<1$) and very dark ($\\chi<<100\\%$) spot or a large ($S\\sim1$) and not very dark ($\\chi\\sim100\\%$) spot fit the data equally well.  \\textbf{b)} Cartoon representation of the three spot models considered in a) showing the location of the spot on the surface of Haumea.  } \\label{fig1} \\end{center} \\end{figure} ",
        "conclusions": ""
    },
    "0911/0911.2183_arXiv.txt": {
        "abstract": "We present the simplest model that permits a largely analytical exploration of the $m=1$ counter--rotating instability in a ``hot'' nearly Keplerian disc of collisionless self--gravitating matter. The model consists of a two--component softened gravity disc, whose linear modes are analysed using WKB. The modes are {\\it slow} in the sense that their (complex) frequency is smaller than the Keplerian orbital frequency by a factor which is of order the ratio of the disc mass to the mass of the central object. Very simple analytical expressions are derived for the precession frequencies and growth rates of local modes; it is shown that a nearly Keplerian disc must be unrealistically hot to avoid an overstability. Global modes are constructed for the case of zero net rotation. ",
        "introduction": "Galactic nuclei are thought to harbour supermassive black holes and dense clusters of stars, whose structural and kinematical properties appear to be correlated with global galaxy properties \\citep{geb96, fm00, geb00}. The imprint of galaxy formation is expected to be recorded in the nature of stellar orbits. A remarkable case is that of our nearest large neighbouring galaxy M31, whose centre has a double--peaked distribution of stars \\citep{lds74, lau93, lau98, kb99}. \\citet{tre95} proposed that the off--centered peak marks the region, in a disc of stars, where lie the apoapses of many eccentric orbits. Self--gravitating models of such an eccentric disc have been proposed \\citep{bac01, ss01, ss02}. Of particular interest to the present investigation is the model proposed in \\citet{ss02}, because it included a few percent of stars on retrograde (i.e. counter--rotating) orbits. Here, it was proposed that the lopsidedness of the nuclear disc of M31 could have been excited by the presence of the retrograde stars, which were accreted to the centre of the galaxy in the form of a globular cluster that spiraled in due to dynamical friction. This proposal was motivated by the work of \\citet{tou02}, which suggested that even a small fraction of mass in retrograde orbits could excite a linear lopsided instability.  Counter--rotating streams of matter in a self--gravitating disc are known to be unstable to lopsided modes \\citep{zh78, ara87, saw88, ms90, pp90, sm94, ljh97}. The dynamics of galactic nuclei involve nearly Keplerian systems of stars or other collisionless matter \\citep{rt96, st99, sst99, tre01, ttk09}. \\citet{tou02} considered a {\\it softened gravity} version of Laplace--Lagrange theory of planetary motions, and showed that a nearly Keplerian axisymmetric disc is linearly unstable to a $m = 1$ mode when even a small fraction of the disc mass is counter--rotating. Softened gravity was introduced by \\citet{mil71} to simplify the analysis of the dynamics of stellar systems. In this form of interaction, the Newtonian $1/d$ gravitational potential is replaced by $1/\\sqrt{d^2 + b^2}$, where $b>0$ is called the {\\it softening length}. In the context of waves in discs, it is well--known that the softening length mimics the epicyclic radius of stars on nearly circular orbits. Therefore, a disc composed of cold collisionless matter interacting via softened gravity provides a surrogate for a ``hot'' collisionless disc.  The goal of this paper is to formulate and analyse the counter--rotating instability in the simplest model of a ``hot'' nearly Keplerian disc of collisionless self--gravitating matter.  To this end we make the following choices: (i) The discs are assumed to be made of matter whose self--interaction is through softened gravity; (ii) A Wentzel--Kramers--Brillouin (WKB) analysis is made of the linearised equations governing the perturbations. The unperturbed two--component nearly Keplerian disc is introduced in \\S~2 and the apse precession rate is defined. The equations governing the linearised perturbations and relevant potential theory for softened gravity is given in \\S~3.  The local (or WKB) approximation and dispersion relation for local modes is derived in \\S~4. This is used to discuss stability, instability and overstability. \\S~5 considers the construction of global  modes for the case of a Kuzmin disc with equal masses in the counter--rotating components, and we conclude in \\S~6. ",
        "conclusions": "The principal aim of this work is to present the simplest model that permits a largely analytical exploration of the $m=1$ counter--rotating instability in a ``hot'' nearly Keplerian disc of collisionless self--gravitating matter. To this end we have considered a two--component softened gravity disc, and performed a linearised WKB analysis of both local and global modes. We derive an analytical expression for local WKB waves for arbitrary $m$, which turns out to be quartic in the frequency $\\omega$. Specialising to $m=1$, we show that $\\omega$ is smaller than the (Keplerian) orbital frequency by the small quantitity $\\varepsilon = M_d/M$ (the ratio of the disc mass to the mass of the central object); in other words, the $m=1$ modes are {\\it slow modes}. The dispersion relation now reduces to a quadratic equation in $\\omega$. Hence the criteria for stability, instability and overstability can be readily derived in simple analytical forms. For a one--component disc (which does not have any counter--rotation), the $m=1$ modes are stable, consistent with the results of \\citet{tre01}. Equal mass in the two counter--rotating components corresponds to the case of not net rotation. In this case we find that the local modes are purely unstable (i.e. not overstable), consistent with \\citet{ara87, pp90, sm94, ljh97, tou02}. However the general case of arbitrary mass ratio in the two counter--rotating components corresponds to overstability, and we show analytically that the discs must be unrealistically hot to avoid an overstability. We finally contructed global  WKB modes, numerically,  for the case of a Kuzmin disc for the case of no net rotation, by using Bohr--Sommerfeld quantisation.         \\def\\etal{{\\it et~al.}} \\def\\apj{{Astroph.\\@ J. }} \\def\\mnras{{Mon.\\@ Not.\\@ Roy.\\@ Ast.\\@ Soc.}} \\def\\aap{{Astron.\\@ Astrophys.}} \\def\\aj{{Astron.\\@ J.}} \\def\\apjl{{Astrophysical.\\@ J. {\\rm Letters}}} \\def\\apss{{Astroph.\\@ Space \\@ Science}}"
    },
    "0911/0911.4888_arXiv.txt": {
        "abstract": "We present sensitive NIR (J, H and K) imaging observations toward four luminous massive star forming regions in the Norma Spiral Arm: G324.201+0.119, G328.307+0.432, G329.337+0.147 and G330.949-0.174. We identify three clusters of young stellar objects (YSO) based on surface density diagnostics.  We also find that sources detected only in the H and K-bands and with colors corresponding to spectral types earlier than B2, are likely YSOs. We analyze the spatial distribution of stars of different masses and find signatures  in two clusters of primordial mass segregation which  can't be explained as due to incompleteness effects.  We show that dynamic interactions of cluster members with the dense gas from the parent core can explain the observed mass segregation, indicating that the gas plays an important role in the dynamics of young clusters. ",
        "introduction": "While still deeply embedded, massive star clusters keep an imprint of the physical conditions in the natal molecular cloud. Massive stars form in molecular cores with typical radii of $\\sim 0.4$ pc, molecular number densities of $3\\times 10^5$ cm$^{-3}$ and masses of $5\\times 10^3$ M$_{\\odot}$ \\citep{plu97,fau04}.  Radio wavelength observations indicate that massive stars are formed in groups \\citep{ho81,gar93} and that they are generally located in the center of massive and dense cores \\citep{gar07};  an indication of mass segregation.    Thus, the study of young massive clusters is essential to fully understand the process of massive star formation.  Stellar mass segregation in a cluster can be either a consequence of dynamical evolution, or of biased massive star formation towards the core center. Dynamical mass segregation arises from kinetic energy equipartition between the cluster members. Massive stars will move slower than low-mass stars and sink to the center of the cluster in a process that takes place in a time comparable to the relaxation time (t$_r$) of the cluster \\citep{bin87}. On the other hand, primordial mass segregation takes place much faster, in a massive star accretion timescale. During this process, massive stars are created in the center of the molecular cloud, where the gas density is highest \\citep{bon01} and/or the mass accretion rate is high \\citep{kru06}.  Several studies have shown evidence of mass segregation in young open clusters having ages of at least one order of magnitude smaller than their relaxation time \\citep{gou04,bra05,che07}. This evidence suggests that mass segregation in those clusters is primordial. However, dynamical timescales are often calculated taking into account star-to-star interactions only, ignoring the gas. The later may have an impact in young clusters, where most of the mass is still in the form of gas. Recent studies have shown that the gaseous component may have an important effect in the dynamics of the stars by gravitational drag, making dynamic times even shorter \\citep[][in \\S~5.2]{esc03,esc04}. Thus, studies of mass segregation in young clusters should include the effects of dense gas in the cluster dynamics. On the other hand, Ascenso et al. (2009) have argued that sample incompleteness in observational data can mislead mass distribution analysis, producing mass segregation in non-segregated clusters. Therefore, to assess whether or not there is mass segregation in young clusters it is necessary to use robust mass segregation indicators.  In this paper we study the spatial distribution of stars in four luminous massive star forming regions located in the Norma spiral arm, at distances between 5.4-7.3 kpc from the Sun. The regions were observed in the near-infrared J, H, and K bands at the VLT.   Garay et al. (2009, hereafter Paper I) present dust mm-continuum observations towards these same regions.   Target selection and data reduction are presented in \\S~\\ref{section_observations1}. Results, including identification of young stars, analysis of their spatial distribution and extinction maps are presented in \\S~\\ref{section_results1}. In \\S~\\ref{section_discussion1} we study the radial distribution of cluster members of different masses. We also discuss the age of the clusters, the evidence of mass segregation and the presence of gaps in the color-magnitude diagrams. Conclusions are presented in \\S~\\ref{section_conclusions1}. ",
        "conclusions": "We used NIR data from the VLT-ISAAC instrument to identify the young population of stars in four regions of massive star formation: G324.201, G328.307, G329.337 and G330.949. These regions are located in the Norma spiral arm at distances between 5.4 and 7.3~kpc from the Sun. We used color criteria to identify sources with NIR-excess and found 236, 120, 186 and 228 NIR-excess sources in G324.201, G328.307, G329.337 and G330.949 respectively.  We identified clusters of YSOs in surface density maps of NIR-excess sources. We found new clusters towards G324.201, G328.307 and G329.337. The spatial distribution of sources detected in H and K, but not in the J-band (J-dropout), was investigated. We concluded that J-dropout sources with colors of massive stars are likely to correspond to cluster members.  We analyzed the spatial distribution of massive and intermediate-mass cluster members and found evidence of mass segregation in each of the clusters. The effects of data incompleteness on the detected mass segregation were assessed and concluded that they are not important.  The age of the clusters was compared with their relaxation time and migration timescales due to gravitational drag, in order to determine if mass segregation is primordial or dynamical. We found that the migration timescales for massive stars is comparable with the age of the cluster members. This evidence suggest that gravitational drag from the gaseous component may be a possible cause for mass segregation in clusters G324.201, G328.307 and G329.337.  Several main sequences in the color-magnitude diagram of regions G328.307, G329.337 and G330.949 were identified. We computed the H$_2$ column density derived from \\coa~for molecular components located in Scutum-Crux arm and compared those values with the column density derived from A$_K$ for the different main sequences. The H$_2$ column densities derived from both methods are in reasonable agreement, indicating that the reddening is mainly due to molecular clouds located along the line of sight."
    },
    "0911/0911.3699_arXiv.txt": {
        "abstract": " ",
        "introduction": " ",
        "conclusions": "In this brief paper, we have highlighted the importance of obtaining an accurate and precise measurement of the noble gas abundances in Saturn's atmosphere.  These can only be obtained by a Saturn entry probe.  For helium in Saturn and Jupiter, we will finally achieve a proper understanding of the extent and physical effects of helium phase separation, and a new generation of thermal evolution and structural models will yield a great leap forward in our understanding of these planets.  For the heavier noble gases, their abundances will constrain formation scenarios for these planets.  Jupiter and Saturn serve as the \\emph{calibrators} for our understanding of the formation and evolution of all giant planets, who now number over 300 in exoplanetary systems.  Accurate models are fundamental to understanding these planets.  In addition, Saturn's helium abundance, along with Jupiter's, are the only constraints yet possible on the physical interaction of H/He at megabar pressures, an important regime of condensed matter physics."
    },
    "0911/0911.4246_arXiv.txt": {
        "abstract": "The WiggleZ Dark Energy Survey is a survey of 240,000 emission line galaxies in the distant universe, measured with the AAOmega spectrograph on the 3.9-m Anglo-Australian Telescope (AAT). The primary aim of the survey is to precisely measure the scale of baryon acoustic oscillations (BAO) imprinted on the spatial distribution of these galaxies at look-back times of 4--8\\,Gyrs.  The target galaxies are selected using ultraviolet photometry from the GALEX satellite, with a flux limit of $NUV<22.8$ mag. We also require that the targets are detected at optical wavelengths, specifically in the range $20.0<r<22.5$\\,mag. We use the Lyman break method applied to the ultraviolet colours, with additional optical colour limits, to select high-redshift galaxies. The galaxies generally have strong emission lines, permitting reliable redshift measurements in relatively short exposure times on the AAT. The median redshift of the galaxies is $z_{med}=0.6$. The redshift range containing 90\\% of the galaxies is $0.2 < z < 1.0$.  The survey will sample a volume of $\\sim$1\\,Gpc$^3$ over a projected area on the sky of 1,000 degree$^2$, with an average target density of 350\\,degree$^{-2}$.  Detailed forecasts indicate the survey will measure the BAO scale to better than 2\\% and the tangential and radial acoustic wave scales to approximately 3\\% and 5\\%, respectively. Combining the WiggleZ constraints with existing cosmic microwave background measurements and the latest supernova data, the marginalized uncertainties in the cosmological model are expected to be $\\sigma(\\Omega_m)=0.02$ and $\\sigma(w)=0.07$ (for a constant $w$ model). The WiggleZ measurement of $w$ will constitute a robust, precise and independent test of dark energy models.  This paper provides a detailed description of the survey and its design, as well as the spectroscopic observations, data reduction, and redshift measurement techniques employed. It also presents an analysis of the properties of the target galaxies, including emission line diagnostics which show that they are mostly extreme starburst galaxies, and {\\em Hubble Space Telescope} images, which show they contain a high fraction of interacting or distorted systems. In conjunction with this paper, we make a public data release of data for the first 100,000 galaxies measured for the project. ",
        "introduction": "A major theme of cosmology for several decades has been the determination of the best cosmological model to describe our Universe. By the early 1990s, evidence was already accumulating from measurements of large-scale structure and the cosmic microwave background that if the Universe was flat, the matter density was well below the critical value, requiring an additional (large) contribution from a non-zero cosmological constant term \\citep{Efstathiou1990}. This ``$\\Lambda$CDM''cosmology, dominated by the cosmological constant ($\\Lambda$) and cold dark matter (CDM), although radical at the time, predicted an older age for the Universe than other models, so was more consistent with observational lower limits to the age \\citep{Ostriker1995}. Additional early support for a non-zero cosmological constant was provided by \\citet{Yoshii1991,Yoshii1995} in their analysis of faint near-infrared galaxy counts. The case for a $\\Lambda$CDM cosmology was strengthened dramatically in the late 1990s by the measurement from distant type Ia supernovae of the acceleration in the expansion rate of the Universe \\citep{Riess1998,Perlmutter1999}.  The goal of cosmology now is to test different models of the underlying physics responsible for this acceleration, dubbed ``dark energy''. This research has focused around the dark energy equation of state (parameterised by $w(z)$, so pressure equals $w$ times density) and how it changes with time. A central goal is to measure $w(z)$ at multiple redshifts to test various dark energy models \\citep{Frieman2008}. A powerful way to test the dark energy models is to measure the geometrical relations between distance and redshift. These can be measured by observing the baryon acoustic oscillation (BAO) scale imprinted on the distribution of baryonic matter at recombination \\citep{Cooray2001,Eisenstein2002,Blake2003,Seo2003,Hu2003,Linder2003,Glazebrook2005}. The BAO signal has been clearly measured in the cosmic microwave background \\citep[i.e.\\ at recombination;][]{Bennett2003}. More recently, the 2dF Galaxy Redshift Survey; \\citep[2dFGRS,][]{Colless2001} and the Sloan Digital Sky Survey \\citep[SDSS,][]{York2000} have permitted the detection of the BAO signal in the distribution of low-redshift ($z\\approx 0.2$--0.35) galaxies. The 2dFGRS detection was reported by \\citet{Cole2005} and \\citet{Percival2007}; the SDSS result was reported by \\citet{Eisenstein2005}, \\citet{Huesti2006} and \\citet{Percival2007}. These galaxy samples are too nearby to be sensitive to the cosmic effects of dark energy, but provide a superb validation of the concept.  Measurement of the BAO scale in the galaxy distribution at high redshifts provides a powerful test of the dark energy models. It gives a geometrical measurement of the Universe at different redshifts which delineates the cosmic expansion history with high accuracy and traces the effects of dark energy \\citep{Blake2003}. It also provides a unique and direct measurement of the expansion rate at high redshift \\citep{Glazebrook2005}. The most important advantage of the BAO measurement is that its measurement is completely independent of the supernova results and, in particular, is less susceptible to systematic errors. It is also complementary to the supernova results in that the measurement uncertainties are orthogonal to those of the supernova measurements \\citep{Blake2009a}.   The WiggleZ Dark Energy Survey is a major large-scale structure survey of UV-selected emission-line galaxies.  The survey is designed to map a cosmic volume large enough to measure the imprint of baryon oscillations in the clustering pattern \\citep[this requires $\\sim 1$ Gpc$^3$;][]{Blake2003} at a significantly higher redshift ($0.2<z<1.0$) than has been previously achieved.  Its core science goals are: \\vspace*{-2mm} \\begin{enumerate} \\parskip 0pt \\parsep 0pt \\parindent 1cm \\item To calculate the power spectrum of the galaxy distribution with sufficient precision to identify and measure the baryon acoustic oscillation scale to within 2 per cent; \\item To determine the value of the dark energy equation of state parameter, $w$, to within 10 per cent when combined with CMB data, and hence either independently validate or refute the currently-favoured ``cosmological constant'' model. \\end{enumerate}\\vspace*{-2mm} The redshift range of the WiggleZ Dark Energy Survey will make it the first BAO survey to sample the epoch when the universe makes the transition from matter-dominated to lambda-dominated expansion.  In this paper we present a detailed description of the survey and the key innovations that were developed to allow it to proceed. We start with a description of the survey design and its requirements in Section~\\ref{sec-design}.  This is followed in Section~\\ref{sec-target} with details of the selection of our target galaxies from UV data combined with optical data.  The UV selection of the target galaxies means they have strong emission lines, so their redshifts can be measured in relatively short exposures on a 4m-class telescope like the Anglo-Australian Telescope (AAT).  In Section~\\ref{sec-spectroscopy} we describe our AAT spectroscopic observations and data reduction. The processing of our data to obtain reliable spectroscopic redshifts is detailed in Section~\\ref{sec-redshifts}. This includes analysis of the reliability of our redshifts. In Section~\\ref{sec-results}, we present a summary of the initial results of the survey, notably the galaxy properties and the success of the survey design.  Following this, in Section~\\ref{sec-data} we describe the data products of the survey, as released in our first public data release accompanying this paper. We summarise this paper and outline the remaining stages of the survey in Section~\\ref{sec-summary}.  A standard cosmology of $\\Omega_m = 0.3$, $\\Omega_{\\Lambda} = 0.7$, $h = 0.7$ is adopted throughout this paper.  \\section[]{The Survey Design} \\label{sec-design}  \\subsection{Scientific goal}  The WiggleZ Dark Energy Survey is a major large-scale structure survey of intermediate-redshift UV-selected emission-line galaxies, focusing on the redshift range $0.2 < z < 1.0$.  The survey is designed to map a cosmic volume of about 1 Gpc$^3$, sufficient to measure the imprint of baryon oscillations in the clustering pattern at a significantly higher redshift than has been previously achieved.  The resulting measurements of cosmic distances and expansion rates, with uncertainties for each below $5\\%$, will constitute a robust test of the predictions of the ``cosmological constant'' model of dark energy.  Detailed forecasts for the survey are presented in \\citet{Blake2009a}.  \\subsection{Choice of target galaxies}  The cosmological measurements from a large-scale structure survey should be independent of the galaxy type used, given that the ``bias'' with which galaxies trace the underlying dark matter fluctuations is expected to be a simple linear function on large scales \\citep{Coles1993,Scherrer1998}. In this sense, the choice of the ``tracer population'' of galaxies can be determined by observational considerations such as telescope exposure times and the availability of input imaging data for target selection.  The WiggleZ Survey targets UV-selected bright star-forming galaxies, obtaining redshifts from the characteristic patterns of emission lines.  A series of colour cuts, described below, are used to boost the fraction of targets lying at high redshift.  The choice of blue galaxies was motivated by several considerations: (1) emission-line redshifts are generally obtainable in short 1-hour exposures at the AAT, even in poor weather, with no requirement to detect the galaxy continuum light; (2) the increasing star-formation rate density of the Universe with redshift results in a significant population of high-redshift $z > 0.5$ targets; (3) a large fraction of UV imaging data for target selection was already available as part of the Medium Imaging Survey conducted by the {Galaxy Evolution Explorer satellite \\citep[GALEX;][]{Martin2005}; (4) other large-scale structure surveys such as the Sloan Digital Sky Survey (SDSS) and the forthcoming Baryon Oscillation Spectroscopic Survey \\citep[BOSS;][]{Schlegel2009} have} preferred to observe Luminous Red Galaxies, allowing the WiggleZ Survey to constrain any systematic effects arising from the choice of tracer galaxy population; and (5) interesting secondary science is achievable from studying bright star-forming galaxies at high redshift.  The main disadvantage in the choice of WiggleZ Survey targets is that emission-line galaxies possess a significantly lower clustering amplitude than red galaxies, and hence larger galaxy number densities are required to minimize the shot noise error in the clustering measurement.  However, there is a benefit to the fact that WiggleZ galaxies avoid the densest environments: non-linear growth of structure---which erases linear-regime signatures in the power spectrum such as baryon oscillations---becomes less significant in the sample.  \\subsection{Target numbers}  The WiggleZ Survey goal is to cover $1{,}000$ deg$^2$ of the equatorial sky in order to map a sufficiently large cosmic volume to measure the imprint of baryon oscillations in the clustering power spectrum.  The survey was designed such that the average galaxy number density $n_g$ is related to the amplitude of the galaxy clustering power spectrum $P_g$ on the relevant baryon oscillation scales by $n_g \\sim 1/P_g$, implying that the contributions of sample variance and shot noise to the clustering error are equal.  This is the optimal survey strategy for a fixed number of galaxies (i.e.\\ fixed observing time).  The galaxy selection cuts described below result in an average target density of 340 deg$^{-2}$.  The overall survey redshift completeness, allowing for repeat observations of objects originally observed in the poorest conditions, is $70\\%$.  {The survey goal is hence to measure $\\sim 340{,}000$ galaxies over $1{,}000$ deg$^2$, of which $\\sim 240{,}000$ are expected to yield successful redshifts.}  We focus on the redshift range $0.2 < z < 1.0$ which we anticipate will yield our main baryon oscillation distance measurement.  {The fraction of redshifts lying in this range is 90 per cent,} therefore the average galaxy number density in the range $0.2 < z < 1.0$ is $n_g = 2.0 \\times 10^{-4} h^3 \\Mpc^{-3}$.  The matter power spectrum amplitude $P_m$ at the characteristic scale of baryon oscillations ($k = 0.15 h \\Mpc^{-1}$) is $3800 h^{-3} \\Mpc^3$ at $z=0$ (for $\\sigma_8 = 0.9$). Assuming a galaxy bias of $b=1.21$ \\citep{Blake2009a} and a growth factor $D=0.74$ (at $z=0.6$), the galaxy power spectrum amplitude at $z=0.6$ is $P_g = P_m D^2 b^2 = 3050 h^{-3} \\Mpc^3$. Converting to redshift-space (using the boost factor $1 + \\frac{2}{3} \\beta + \\frac{1}{5}\\beta^2$ where $\\beta=0.5$) we obtain $P_g = 4210 h^{-3} \\Mpc^3$.  The survey value of $n_g \\times P_g $ is therefore equal to 0.84, close to the optimal condition of $n_g P_g = 1$.  {The average redshift completeness of an individual observation (or ``pointing'') is $60\\%$ (we reach the value of 70\\% per galaxy mentioned above after repeating measurements made in poor conditions). Thus the total number of individual spectra required is $N = 240{,}000 / 0.6 = 400{,}000$.}  On average we can use 325 fibres per pointing and obtain 7.4 pointings per night, and therefore can measure the required $400{,}000$ spectra in 166 nights of clear weather.  This calculation is summarized in Table~\\ref{tab-design}.  \\subsection{Survey fields}  A final constraint on the survey design is the availability of suitable optical imaging catalogues in regions of sky overlapping the GALEX UV data.  We require the optical data to provide more accurate positions for fibre spectroscopy and to refine our target selection.  The optical data we use are from the fourth data release of the Sloan Digital Sky Survey \\citep[SDSS,][]{Adelman2006} in the NGP region and from the Canada-France-Hawaii Telescope (CFHT) Second Red-sequence Cluster Survey \\citep[RCS2,][]{Yee2007} in the SGP region.  The WiggleZ survey area, illustrated in Figure~\\ref{figsurvey}, is split into seven equatorial regions that sum to an area of 1000 deg$^2$, to facilitate year-round observing.  We require that each region should possess a minimum angular dimension of $\\sim 10$ deg, corresponding to a spatial co-moving scale that exceeds by at least a factor of two the standard ruler preferred scale [which projects to (8.5, 4.6, 3.2, 2.6) deg at $z = (0.25, 0.5, 0.75, 1.0)$].  The boundaries of the regions are listed in Table~\\ref{tab-survey}.  We have also defined high-priority sub-regions in each area (see Table~\\ref{tab-surveyhigh}) which will be observed first to ensure that we have the maximum contiguous area surveyed should we be unable to observe the complete sample.    \\begin{figure*} \\epsfig{file=fig_survey.ps,width=0.6\\linewidth,angle=-90} \\caption{The sky distribution of the seven WiggleZ survey regions compared to the coverage of the SDSS, RCS2 and GALEX data sets at the end of 2008.} \\label{figsurvey} \\end{figure*}    \\begin{table} \\caption{Design of the Observational Programme} \\label{tab-design} \\begin{tabular}{lll} \\hline Component     &  Value \\\\ \\hline Success rate (fraction $Q\\ge 3$)  &   60\\%  \\\\ Target fibres/field  &   325   \\\\ Fields/ clear night  &  7.4   \\\\ {\\bf Rate (successful targets/night)} &   1442  \\\\ \\hline Required number of targets  &  240,000  \\\\ Rate per night (from above)         &   1442  \\\\ Required clear nights  &   166   \\\\ Weather factor      &   1.33  \\\\ {\\bf Required total nights } &   221 \\\\ \\hline \\end{tabular} \\end{table}   \\begin{table} \\caption{Survey Regions: Full Extent} \\label{tab-survey} \\begin{tabular}{rrrrrr} \\hline Name & RA$_{min}$ & RA$_{max}$ & Dec$_{min}$ & Dec$_{max}$ & Area\\\\ & (deg)     & (deg)     & (deg)     & (deg)     & (deg$^2$) \\\\ \\hline 0-hr &  350.1 & 359.1  &  $-13.4$ & $+1.8$ &  135.7 \\\\ 1-hr &    7.5 &  20.6  &   $-3.7$ & $+5.3$ &  117.8 \\\\ 3-hr &   43.0 &  52.2  &  $-18.6$ & $-5.7$ &  115.8 \\\\ 9-hr &  133.7 & 148.8  &   $-1.0$ & $+8.0$ &  137.0 \\\\ 11-hr &  153.0 & 172.0  &   $-1.0$ & $+8.0$ &  170.5 \\\\ 15-hr &  210.0 & 230.0  &   $-3.0$ & $+7.0$ &  199.6 \\\\ 22-hr &  320.4 & 330.2  &   $-5.0$ & $+4.8$ &   95.9 \\\\ \\hline \\end{tabular} \\end{table}  \\begin{table} \\caption{Survey Regions: Priority Areas} \\label{tab-surveyhigh} \\begin{tabular}{rrrrr} \\hline Name & RA$_{min}$ & RA$_{max}$ & Dec$_{min}$ & Dec$_{max}$ \\\\ & (deg)     & (deg)     & (deg)     & (deg)     \\\\ \\hline 0-hr &  350.1  & 359.1  &  $-13.4$ &  $-4.4$    \\\\ 1-hr &    7.5  &  16.5  &   $-3.7$ &  $+5.3$    \\\\ 3-hr &   43.0  &  52.2  &  $-18.6$ &  $-9.6$    \\\\ 9-hr &  135.0  & 144.0  &   $-1.0$ &  $+8.0$    \\\\ 11-hr &  159.0  & 168.0  &   $-1.0$ &  $+8.0$    \\\\ 15-hr &  215.0  & 224.0  &   $-2.0$ &  $+7.0$    \\\\ 22-hr &  320.4  & 330.2  &   $-5.0$ &  $+4.8$    \\\\ \\hline \\end{tabular}  Note: the high-priority region of the 22-hr field is the full field. \\end{table}   \\section[]{Target Galaxy Sample} \\label{sec-target}  In this section we describe how we select our target galaxies from the GALEX ultraviolet data combined with additional optical imaging data. We stress that the target selection is motivated by our main science objective of measuring the BAO scale. For this we need to obtain as many high redshift galaxy redshifts as possible per hour of AAT observing without necessarily obtaining a sample with simple selection criteria in all the observing parameters. We nevertheless discuss the selection process in detail in this section to facilitate the use of the WiggleZ data for other projects that may need well-defined (sub) samples.  \\subsection{GALEX ultraviolet imaging data} \\label{sec-galexdata}  Our target galaxies are primarily selected using UV photometry from {the GALEX satellite \\citep{Martin2005}.} This satellite is carrying out multiple imaging surveys in two ultraviolet bands, in FUV from 135--175\\,nm, and in NUV from 175--275\\,nm. Specifically, the WiggleZ survey uses the photometry of the GALEX Medium Imaging Survey (MIS), which is an imaging survey of 1,000 square degrees of sky to a depth of 23\\,mag in both bands. The MIS consists of observations of 1.2\\,degree diameter circular tiles, with a nominal minimum exposure time of 1500\\,seconds. Several hundred extra GALEX observations will be dedicated to the WiggleZ project. These are in addition to the original MIS survey plan, but they are all taken with the standard MIS-depth exposures and cover more contiguous sky areas due to improved bright-star avoidance algorithms. Furthermore, GALEX observations in the original MIS plan that overlap the WiggleZ survey regions are being observed at a higher priority to maximise the fields available for the WiggleZ survey.  The MIS data we use for WiggleZ have a range of exposure times, as illustrated by the distribution shown in Figure~\\ref{fig-tiles}. The sharp peak at 1700\\,seconds arises because this is the maximum possible observing time in one satellite orbit. The figure also shows that there is a range in average Galactic dust extinction across the different GALEX observations. The effect of dust is discussed in Section~\\ref{sec-uniformity} below.  \\begin{figure} \\epsfig{file=fig_tiles.eps,width=\\linewidth,angle=0} \\caption{Properties of GALEX data used to select WiggleZ targets. Upper panel: distribution of exposure times. Lower panel: distribution of average dust extinction of each tile.} \\label{fig-tiles} \\end{figure}  All GALEX observations for the WiggleZ project are processed by the standard MIS data pipeline at Caltech to produce image catalogues \\citep{Morrissey2007}. The NUV images are processed in a standard way with objects detected according to a threshold criterion using SExtractor \\citep{Bertin1996}. For the NUV detections we use the ``NUV calibrated magnitude'' from the pipeline: this is the flux through an elliptical aperture scaled to twice the Kron radius of each source (termed ``MAG\\_AUTO'' by SExtractor). The detection limit in the raw NUV data is about 23 mag, as shown in Fig.~\\ref{fig-uvcount} (a). We used a different approach for the FUV photometry because relatively few objects at the faint limit of our survey ($NUV\\approx 22.8$) were independently detected in the FUV band in the relatively short MIS exposures. We therefore based the FUV photometry entirely on the NUV detections. For each object in the NUV catalogue, the corresponding FUV magnitude is obtained by measuring a fixed 12-arcsecond (8 pixel) diameter circular aperture (``aperture 4'' in the pipeline) in the FUV image at the NUV position. The full width half-maximum (FWHM) of the GALEX point spread function (PSF) is 4.3 arc seconds in the NUV and 5.3 arc seconds in the FUV.  Our use of total NUV magnitudes and aperture FUV magnitudes means that our $FUV-NUV>1$ colour selection could, in principle, be biased by colour gradients in resolved sources. Similarly we could suffer from source confusion with objects separated in the NUV data still being close enough to contaminate the FUV apertures, causing objects to incorrectly not satisfy the FUV-NUV colour limit (see Section~\\ref{sec-selection}).  In practice, the low FUV signal means that very few sources are rejected by this colour criterion.  The 1-$\\sigma$ calibration uncertainties in the GALEX photometry are 0.05 and 0.03 magnitudes in the FUV and NUV bands, respectively \\citep{Morrissey2007}.  The astrometric precision of the MIS catalogue positions is 0.5 arc second r.m.s., but this worsens with increasing distance from the centre of a tile \\citep{Morrissey2007}. The nominal FWHM of the PSF of 5 arc seconds also becomes poorer at increasing distance from the centre of a tile. Additionally, uncertainties in the photometric zero-points of the GALEX observations increase radially from the tile centre. For these reasons, the WiggleZ survey only utilises the inner 1.1 degree diameter region of each GALEX tile, where the astrometry, PSF and photometric zero-point uncertainties are negligibly different to the nominal values.  At UV wavelengths, it is especially important to correct the photometry for Galactic dust extinction. The standard GALEX pipeline records the $E(B-V)$ extinction \\citep[][hereafter SFD]{Schlegel98} for each detected source. We then apply the following corrections (including zero-point calibration to the AB system) to the NUV and FUV photometry \\citep{Wyder2007}: \\begin{eqnarray} FUV &=& FUV_{raw} + 18.82 - 8.2 \\times E(B-V) \\nonumber \\\\ & & +0.06\\times E(B-V)^2,\\\\ NUV &=& NUV_{raw} + 18.82 - 8.2\\times E(B-V) \\nonumber \\\\ & &  +0.67\\times E(B-V)^2. \\end{eqnarray} We note that the SFD dust maps have a resolution of 6.1\\,arc min and that the uncertainty in the reddening values is 16\\%. We also note that SFD determined that it was not reliable to calibrate the dust map normalisation from variations in the number counts of optically-selected galaxies: this over-estimates the dust amplitude because typical galaxy catalogues are both magnitude and surface brightness limited. We further discuss the sensitivity of our image catalogues to dust in Section~\\ref{sec-uniformity}.   \\subsection{Optical Data}  In addition to the GALEX photometry, we use optical photometry from the fourth data release of the Sloan Digital Sky Survey \\citep[SDSS,][]{Adelman2006} and from the CFHT Second Red-sequence Cluster Survey \\citep[RCS2,][]{Yee2007} to provide accurate astrometry and improved target selection.  The SDSS photometry covers 7,000 square degrees of sky in 5 bands from the near ultraviolet to the near infrared. The five bands are $u, g, r, i$ and $z$, with magnitude limits (95 per cent point source completeness) of 22.0, 22.2, 22.2, 21.3 and 20.5 mag respectively. The median PSF FWHM (in $r$) is 1.4 arc seconds.  The photometric calibration of the SDSS is nominally accurate to 0.02--0.03 magnitudes \\citep{Ivezic2004}, but the uncertainty at our faint detection limit of $r=22.5$ is more likely to be 0.2--0.3 magnitudes. We use the SDSS ``model'' magnitudes, calculated on the AB system. The location of objects detected in these bands is known to an rms astrometric accuracy of 0.1 arc seconds.  The RCS2 is an imaging survey of 1,000 square degrees of sky in 3 of the SDSS bands: $g, r$ and $z$. The magnitude limits (5-sigma point source limits) are 25.3, 24.8, and 22.5 respectively, significantly fainter than in the SDSS.  The typical seeing in the RCS2 imaging is ($0.6\\pm 0.1$) arc seconds for the best half of the data and ($0.8\\pm 0.1$) arc seconds for the remainder.  The internal photometric precision of the RCS2 is close to that of the CFHT Legacy Survey \\citep{Ilbert2006}, i.e.\\ 0.04 mag in each band; the magnitudes are on the AB system, calculated from a curve-of-growth analysis. The colours are measured in smaller (adaptive) apertures and then scaled to the total magnitude of the source in the band with the strongest detection (usually $r$). The RCS2 has an astrometric accuracy of 0.15 arc seconds, comparable to the astrometric precision of SDSS.  The precise astrometry of both these optical data sets is crucial to the WiggleZ survey, because the combination of the point-spread-function and astrometry of the GALEX source detections is too poor for the optical fibres used to feed light into the AAOmega spectrographs, which subtend an angle of 2 arc seconds on the sky.  In all our processing we use dust-corrected optical catalogues. For SDSS, the dust correction is provided as part of the public data and for RCS2 we apply a dust correction ourselves: in both cases this is the standard SFD correction.  \\subsection{Matching UV to optical samples}  The WiggleZ targets are selected from the GALEX UV photometry and then combined with either SDSS or RCS2 optical photometry. In the NGP region we combine the GALEX photometry with SDSS photometry, and in the SGP region we combine it with RCS2 photometry. The respective catalogues are combined by selecting the closest optical match to each GALEX source within a radius of 2.5 arc seconds. For the surface densities of the three data sets, this corresponds to 95 per cent confidence in a GALEX-SDSS match and 90 per cent confidence in a GALEX-RCS2 match. (The lower confidence for the RCS2 matches is mainly a consequence of the higher number density of sources in the RCS2 data.) We use the optical position for each matched source in our target catalogues as these are more precise than the GALEX positions.  Once an object is matched to the optical data, we can then apply further colour limits to refine our colour selection (see Section~\\ref{sec-selection}) noting that the $NUV-r$ colour is a total colour. However, the large size of the GALEX PSF (FWHM 4.3 arc seconds in the NUV) compared to the optical imaging means that the colours may be distorted if multiple objects are merged in the GALEX photometry, but are resolved in the optical imaging. In this case the GALEX photometry would be too bright compared to the optical measurement. As discussed in Section~\\ref{sec-properties}, this is possible given the evidence for multiple sources we present in Fig.~\\ref{fig-morphology}.  \\subsection{Uniformity of GALEX Data} \\label{sec-uniformity}  An important requirement for the survey is that no artificial structure be introduced by variations in the input catalogues. This is particularly true for the GALEX data as the tile diameter (1.1 degrees) is close to the BAO scale (see Section \\ref{sec-design}). The primary source of non-uniformity is likely to be foreground Galactic dust, although we also test the data for any calibration offsets. To demonstrate the importance of dust, we show the range of {\\em average} dust extinction across each tile, for a set of GALEX tiles in Fig.~\\ref{fig-tiles} (lower panel). This range ($0.03\\lesssim E(B-V) \\lesssim 0.07$; see Section \\ref{sec-galexdata} for the conversion to $A_{NUV,FUV}$) is sufficient to affect the detection rates as we show below.  In Fig.~\\ref{fig-dust} we present images of the SFD dust maps in each field. These clearly show structure on scales similar to and smaller than the BAO scale, so it is very important to quantify the effect of dust on our input catalogues and correct for it if necessary. We also note an additional effect: some of the regions of higher dust are associated with Galactic nebular emission. One such example is the Eridanus loop which coincides with the diagonal dust lane visible in the 3-hour field in Fig.~\\ref{fig-dust}. Spectra of targets in such regions can be contaminated by rest-wavelength nebular emission lines, confusing the target redshift measurement.  We tested the GALEX photometry in two ways: first by analysing the detected number counts and secondly by comparing the photometry of objects measured in overlapping GALEX tiles. The first test addresses all factors affecting image detection (notably dust) and the second tests the internal calibration of the photometry. The second tests did not reveal any evidence for significant differences in absolute calibration between overlapping tiles, so these will not be discussed in any more detail (small offsets were measured, but these were uncorrelated the numbers of images detected on the respective tiles).  \\begin{figure} \\epsfig{file=fig_alldust.ps,width=1.0\\linewidth,angle=0} \\caption{Distribution of dust in each of the survey fields. The grey scale shows the \\citet{Schlegel98} extinction values such that white is $E(B-V)=0.02$ and black is $E(B-V)=0.1$. The horizontal and vertical axes in each panel denote the Right Ascension and Declination, respectively, measured in degrees. } \\label{fig-dust} \\end{figure}  In Figure~\\ref{fig-uvcount}(a,b) we show the differential number counts of all sources detected in the NUV tiles in the 15-hour region. In the upper panel the counts are plotted as a function of raw observed magnitude and in the lower plot, as a function of dust-corrected magnitude. In each case the tiles are grouped according to the average dust extinction for each tile and the lines plotted give the average counts for each group of tiles as indicated by the key. We further limited the choice of tiles to those with exposures in the range 1600--1700 seconds to focus on the role of extinction. Our first observation from the plots is that the dust correction is working as expected: the uncorrected counts all show the same faint cut-off in raw magnitude, but they have different normalisations of their power-law slopes. In contrast, the dust-corrected counts show the same power-law normalisation, as desired.  \\begin{figure*} \\begin{minipage}{1.05\\linewidth} \\hspace*{-4mm} \\epsfig{file=fig_uvcount1.eps,width=0.52\\linewidth,angle=0} \\hspace*{-6mm} \\epsfig{file=fig_uvcount2.eps,width=0.52\\linewidth,angle=0} \\end{minipage}  \\caption{Number counts of GALEX NUV detections in tiles from the 15-hour region as a function of dust extinction. The NUV magnitude limit of the WiggleZ survey is indicated by a vertical dashed line in each plot. In the left hand panels (a, b) we plot all objects detected before and after the dust correction. The counts are binned by the average dust extinction in each tile, one curve for each bin (see key). The upper panel shows the raw counts; in the lower panel the dust correction has been applied to every individual object. The improved agreement of the power-law region of the number counts in panel b demonstrates the effectiveness of the dust correction. The right hand panels (c, d) are similar, but only count objects that satisfy the WiggleZ target selection criteria. Panel d shows that the target numbers at the NUV survey limit are more sensitive to dust than the total counts of detections shown in panel b.  } \\label{fig-uvcount} \\end{figure*}  Although we have shown that calibration variations in the GALEX data are minimal, the faint magnitude cut off in the counts has a strong effect on the number of WiggleZ targets selected as shown in Fig.~\\ref{fig-uvcount}(c,d). Here we show the number counts of WiggleZ targets, having applied the survey selection criteria (see Sec.~\\ref{sec-selection}) to the raw counts. This shows that, unlike the raw counts, the WiggleZ target counts at our survey limit ($NUV<22.8$) are strongly affected by dust.  The mean counts per tile vary by a factor of two over the full range of extinction, although 70 per cent of the tiles lie in a narrow range of extinction ($0.025<E(B-V)<0.045$) for which the counts vary by only 10 per cent. For the WiggleZ clustering measurements we will deal with this effect by scaling the random catalogues to follow the GALEX magnitude limits. We have tested this approach by measuring the angular power spectrum of the galaxies and comparing that to theoretical models. After scaling our selection function to allow for the combined effects of dust and exposure time on the GALEX counts we find there is no evidence for any large-scale systematic variation \\citep{Blake2009a}.  For measurements of the luminosity function of the WiggleZ sources, however, the correction needs to be with respect to the underlying power law and is therefore much larger. This will be discussed in more detail in later papers.  \\subsection{Galaxy Selection} \\label{sec-selection}  The main criteria used to select the WiggleZ survey targets from the GALEX and optical photometry are a $FUV-NUV$ colour cut to select high-redshift galaxies, and a $NUV-r$ colour cut to select emission line galaxies. These are shown in Fig.~\\ref{basicTselect}, a colour-colour plot of $FUV-NUV$ vs.\\ $NUV-r$. In the diagram we show the WiggleZ selection limits (dashed rectangle), data points from combined DEEP2 \\citep{Davis2003} and GALEX observations, and model tracks for several galaxy types over a range of redshifts.  The UV colour selection (specifically $FUV - NUV \\geq 1$ or $FUV$ undetected: a ``drop-out'') selects galaxies at redshifts $z\\geq 0.5$ because the Lyman break enters the FUV filter at a redshift of $z \\approx 0.5$. We note that in the MIS UV data used in our survey, for fainter galaxies (near our $NUV$ detection limit), the $FUV$ measurements have very low signal-to-noise, so the drop-out selection criterion is very crude. (The data points in Fig.~\\ref{basicTselect}, by contrast, were measured in much deeper GALEX data than used for our survey.)  The diagram also shows how we use the $NUV-r$ colour. The $NUV-r \\leq 2$ limit selects bluer star-forming galaxies, and the $-0.5 \\leq NUV-r $ limit helps remove spurious matches between the GALEX and optical data sets. These two colour cuts were chosen based on the colours of galaxy models as a function of redshift. The colours of these galaxy models and the colour cuts used for the WiggleZ target selection are shown in Fig.~\\ref{basicTselect}.  For the GALEX NUV photometry we impose a faint magnitude limit of 22.8 in NUV.  In addition, we also require that the NUV flux measurement has a signal-to-noise $\\geq 3$ to minimise the number of false detections. This requirement becomes increasingly important for dustier regions of the sky: in regions of low dust ($E(B-V)\\approx 0.03$) the signal-to-noise cut removes about 10 per cent of the raw $NUV \\leq 22.8$ sources, but in high dust regions ($E(B-V)\\approx 0.06$) it removes 25 per cent of the sources.  For the optical photometry (both SDSS and RCS2) we impose a magnitude range of $20<r<22.5$. The faint limit is set by the SDSS survey detection limit; the bright limit is to avoid low-redshift galaxies. As a consequence of this bright $r$ magnitude limit, the $NUV-r$ colour cut results in a bright magnitude limit in the NUV of 19.5, although this is not explicitly part of the WiggleZ target selection criteria. A further consequence of the $NUV-r$ colour cut (and indeed the average $NUV-r$ colours) is that the optical magnitudes of the selected galaxies do not peak at the faint $r=22.5$ limit, but instead peaks at $r\\approx 21.5$, about a magnitude brighter. Note that we do not apply any morphological selection to remove objects classified as ``stellar'' in the optical imaging because virtually no stars satisfy our photometric selection criteria. These basic selection criteria are listed in the first section of Table~\\ref{tab-selection}.  \\begin{figure} \\epsfig{file=fig_WGZTSelect.ps,width=\\figwidth\\linewidth,angle=0}  \\caption{The basic WiggleZ target selection criteria. The rectangle marked by the dashed line denotes the region in FUV-NUV,NUV-r colour space where WiggleZ targets lie. The solid lines are the model tracks for seven galaxy templates as a function of redshift for $z > 0.5$.  The first track, from left to right, is for a Lyman break galaxy template \\citep[Figure 1 in][]{Steidel1996}. The remaining tracks are two starburst templates \\citep{Kinney1996} and a series of (Im to E/S0) local galaxy templates \\citep{Coleman1980}, all extended in wavelength range by matching them with \\citet{Bruzual1993} models (Hsaio-Wen Chen, private communication).  The dashed lines at the end of the solid lines are these models for $z \\le 0.5$. Also shown as points are galaxies from deep GALEX imaging that were detected in all three bands and have published redshifts from the DEEP2 survey \\citep{Davis2003}. The triangles indicate objects with redshift $z > 0.5$.}  \\label{basicTselect} \\end{figure}  A further selection based on optical colour is used to reduce the number of targets selected with redshifts below $z=0.5$. These arise because the $FUV-NUV$ selection is imperfect due to photometric errors --- many of our targets are at the limits of the GALEX photometry in typical MIS exposures.  We do this by rejecting galaxies whose optical colours and magnitudes indicate they are very likely to be lying at lower redshifts (as shown in Fig~\\ref{fig-karl}).  Two different sets of low redshift rejection (LRR) criteria are used in the SDSS and RCS2 regions respectively, as listed in Table~\\ref{tab-selection}. In the SDSS data we are near the faint limit of the survey, so can only apply this to the brighter galaxies. However the brighter galaxies are more likely to be at lower redshift so this ameliorates this disadvantage. We therefore start by applying $g$ and $i$ magnitude limits in order to secure accurate colours.  Then we apply a ($g-r, r-i$) colour cut to select the 400 nm break in galaxies at redshifts $z<0.5$. In the RCS2 regions the optical data go much deeper, so no magnitude limits are required, but there are no $i$ data, so we have to use a $(g-r, r-z)$ colour cut which is less precise at discriminating at $z<0.5$. The success of these selection criteria is demonstrated in Fig.~\\ref{fig-selection} where we show redshift distributions of the objects observed with and without the LRR colour cuts.   \\begin{figure} \\epsfig{file=fig_karl.eps,width=0.8\\linewidth,angle=270} \\caption{The selection criteria applied to avoid low-redshift targets in the two survey regions. The points in each diagram are WiggleZ galaxies observed before we introduced the low redshift rejection cuts. Galaxies with redshifts below $z=0.6$ are plotted as blue points; the rest are plotted in red. Galaxies below the lines in each panel are rejected from the target catalogue as they have a high probability of being at low redshifts. Also shown as magenta lines are PEGASE-2 \\citep{Fioc1997} evolutionary tracks of four galaxy models ranging from E/SO (upper) to Irr (lower) types. Each model ranges from $z=0$ to $z=1$ with the high redshift ($z>0.5$) section dashed. } \\label{fig-karl} \\end{figure}   \\begin{table} \\caption{Photometric selection criteria for WiggleZ galaxies} \\label{tab-selection} \\begin{tabular}{ll} \\hline Criterion  &  Values     \\\\ \\hline \\multicolumn{2}{l}{\\em Select targets satisfying all these basic criteria:}\\\\ Magnitude &  $NUV<22.8$    \\\\ Magnitude &   $20 < r < 22.5$ \\\\ Colour    &  $(FUV-NUV) > 1$ or no FUV  \\\\ Colour    & $-0.5 < (NUV-r) < 2$  \\\\ Signal    &  $S/N_{NUV}>3$ \\\\ Optical Position  & matches within 2.5 arc seconds \\\\ \\hline \\multicolumn{2}{l}{\\em Then reject targets satisfying these:}\\\\ LRR$_{SDSS}$ & $g < 22.5$, $i < 21.5,$ \\\\ & $(r-i) < (g-r-0.1),  (r-i) < 0.4$ \\\\ LRR$_{RCS2}$ & $(g-r) > 0.6, (r-z) < 0.7(g-r)$ \\\\ \\hline \\end{tabular} \\end{table}   \\begin{figure} \\epsfig{file=fig_selection.eps,width=\\figwidth\\linewidth,angle=0} \\caption{The result of the optical colour cuts used to reduce the number of low-redshift targets. The top panel shows the effect in the Northern (SDSS) regions of the survey and the lower panel shows the Southern (RCS2) regions. In each panel the dashed histogram shows the redshift distribution of all galaxies observed before the dates on which the colour cuts were applied in each section of the survey (see dates in Table~\\ref{tab-dates}). In each panel the solid histogram shows the results of retrospectively applying the colour cuts to these same objects.  In both regions the cuts result in a significantly higher median redshift; we note that the deeper photometry in the RCS2 regions allows a more effective cut to be applied in the lower panel. } \\label{fig-selection} \\end{figure}  We also apply a prioritisation scheme to our target allocation as given in Table~\\ref{tab-priority}. The priorities given in the table are used when observing so that higher-priority objects are observed first. This is done for two reasons. First, there is a weak correlation between $r$ magnitude and redshift (shown in Fig.~\\ref{fig-fluxredshift}) so this also serves to select high-redshift objects. Secondly, this approach means that our later observations will be of brighter objects, allowing us more flexibility in the final stages of the observational campaign.  \\begin{table} \\caption{Priority selection scheme for WiggleZ observations} \\label{tab-priority} \\begin{tabular}{ll} \\hline Criterion  &  Priority   \\\\ \\hline repeat observations & 9 (highest) \\\\ white dwarf calibrators & 9 \\\\ {\\em New targets$^1$  by magnitude:} \\\\ $22.0 < r < 22.5$ & 8     \\\\ $21.5 < r < 22.0$ & 7     \\\\ $21.0 < r < 21.5$ & 6     \\\\ $20.5 < r < 21.0$ & 5     \\\\ $20.0 < r < 20.5$ & 4     \\\\ {\\em Other objects:} \\\\ Additional projects & 3 \\\\ Old good weather targets ($Q=$1,2) & 2 \\\\ Old good weather targets ($Q=$3--5) & 1 (lowest) \\\\ \\hline \\end{tabular}  Note 1: The ``new'' targets include objects that were previously observed in bad weather without obtaining a redshift. \\end{table}  \\begin{figure} \\epsfig{file=fig_fluxredshift.eps,width=\\figwidth\\linewidth,angle=0} \\caption{The correlation between redshift and $r$-band flux for the observed WiggleZ galaxies. As there is a weak correlation in both the SDSS (top panel) and RCS2 (lower panel) samples, we observe targets with fainter $r$ magnitudes at a higher priority in order to increase the median redshift of the sample. } \\label{fig-fluxredshift} \\end{figure}  Before observing we make a final check of all the target objects by visually inspecting them on SuperCOSMOS \\citep{Hambly2001} sky survey images to ensure none are artifacts associated with extremely bright ($r<12$) Galactic stars. This happens occasionally with automated data catalogues because bright objects can be incorrectly segmented into smaller, fainter objects, which are then matched to GALEX NUV detections. This is a problem, because the signal from these bright objects swamps adjacent spectra in such a way that we lose as many as 30 spectra from an observation.  The final result of the whole selection process is a target catalogue with an average surface density of 350 deg$^{-2}$ which is well matched to our survey goals. The WiggleZ target galaxies represent a small fraction of optical galaxies. Analysis of a sample of DEEP2 \\citep{Davis2003} galaxies for which deep optical and GALEX observations are available shows that in the SDSS regions, WiggleZ selects $2.2\\pm0.2$ per cent of the ($r<22.5$) optical galaxies. In the RCS2 regions, WiggleZ selects $2.6\\pm 0.2$ per cent of these same galaxies. These low percentages are mainly due to the WiggleZ NUV flux limit and NUV$-r$ colour cuts which only select galaxies with very high current star formation rates, but we also note that the median $r$ magnitude of the WiggleZ galaxies is brighter than $r=22.5$ (see Fig.~\\ref{fig-fluxredshift}).   \\section[]{Spectroscopic Observations} \\label{sec-spectroscopy}  The major component of this project is our large, 220-night observing campaign on the AAT. In this section we describe the management of this campaign, the procedures required to configure our AAOmega observations and the data processing.  \\subsection{AAOmega Spectrograph and Observing Setup}  The efficiency of the WiggleZ redshift survey is critically dependent on the high sensitivity and multiplexing gain made possible by the new AAOmega spectrograph and existing 2dF fibre positioner on the AAT \\citep{Smith2004,Saunders2004}. AAOmega is a bench-mounted, fibre-fed spectrograph consisting of blue and red arms split by a dichroic. Volume phase holographic (VPH) gratings are utilized resulting in an improved efficiency over reflection gratings. The bench for the spectrograph is mounted in a mechanically and thermally stable coud\\'e room. For more details about the performance characteristics of AAOmega, see \\citet{Sharp2006}.  In multi-object mode the combination of the 2dF positioner and AAOmega spectrograph allows observations with a total of 392 fibres over a circular field of view with a diameter of two degrees. The angular size of each fibre on the sky is 2 arc seconds. Our survey utilizes the 580V and 385R gratings in the blue and red arms, respectively, both providing a resolution of $R \\approx 1300$. This gives a dispersion of about 0.11 nm pixel$^{-1}$ in the blue arm and 0.16 nm pixel$^{-1}$ in the red arm.  The observable wavelength range of the system is 370--950 nm. The blue limit is determined by the fibre transmission characteristics (over the 38m-long fibres) and the red limit is due to the CCD response. In our early observations, we used the standard dichroic with the cut over centred at 570 nm. As we also required a small spectral overlap between the blue and red arms, this fixed the low-wavelength limit of the red spectrograph which could then only cover up to a maximum wavelength of 850 nm.  We alleviated this limit on observing redder wavelengths by purchasing a new dichroic for AAOmega, in partnership with the AAO, to extend our wavelength coverage.  The original AAOmega dichroic beam splitter cuts over at 570 nm while the new dichroic has a cut over at 670 nm. The new dichroic was specially designed for our project to allow improved spectral coverage out to the system limit of 950 nm, using the 385R grating, while still allowing continuous spectral coverage down to a new blue limit of 470nm (see Table~\\ref{tab-aaomega}). This redder wavelength range enables more reliable redshift identification of the emission-line galaxies in our sample. Specifically, an increase of 100 nm at the red end allows the detection of secondary emission-lines such as H$\\beta$/[OIII] lines at higher redshifts up to $z\\approx 0.95$. At redshifts above $z\\approx 0.95$ we often rely on just the [OII] (372.7, 372.9 nm) doublet for redshift identification, but we are aided by the fact that the doublet is often resolved at these redshifts. The new dichroic was first used for our 2007 August observing run (see Table~\\ref{tab-dates}). About 25 per cent of the 100,138 reliable redshifts measured up to the end of Semester 2008A on 2008 May 15 were taken with the original dichroic; the rest used the new dichroic which will be used for the remainder of the survey.   \\begin{table} \\caption{The two AAOmega setups used for our observations} \\label{tab-aaomega} \\begin{tabular}{lrrrr} \\hline Dichroic  &  $\\lambda_{C,Blue}$ & Blue Range  &  $\\lambda_{C,Red}$ & Red Range \\\\ \\hline 570 nm (old) &  477  &  370--580 & 725 & 560--850 \\\\ 670 nm (new) &  575  &  470--680 & 815 & 650--950 \\\\ \\hline \\end{tabular}  Note: for each dichroic the table lists the central wavelength setting and observable wavelength range for the blue and red arms of the spectrograph. For each VPH grating the blaze wavelength was set to be the same as these central wavelengths for maximum efficiency. All units are nm. \\end{table}   \\begin{table} \\caption{Significant Dates in our Observing Campaign} \\label{tab-dates} \\begin{tabular}{llr} \\hline Date  &  Activity & $N_z$\\\\ \\hline 2006 February & Pilot observations \\\\ 2006 August 19 & Start of survey (semester 2006B) \\\\ 2007 April 11 & Start of using LRR in SDSS fields&  12445\\\\ & Start of 0.7 deg radius observing& \\\\ 2007 June 8--21 & Observations without blue camera& 24740 \\\\ 2007 August 7 & Start of new dichroic use & 26293 \\\\ 2007 August 14 & End of 0.7 deg radius observing& 38540\\\\ 2007 October 4 & Start of using LRR in RCS2 fields& 42179 \\\\ 2008 May 13 & End of Semester 2008A and \\\\ & half-way point after 112 nights. & 100138\\\\ \\hline \\end{tabular}  Note: $N_z$ is the running total number of good redshifts ($Q\\ge3$) measured on each date. \\end{table}   \\subsection{Blank Sky and Guide Star Positions}  The 2dF/AAOmega system, like other multi-object fibre spectrographs requires fiducial (``guide'') stars to align the field accurately and blank sky positions where a subset of the target fibres can be placed to sample the sky background.  \\subsubsection{Selection of Blank Sky Positions}  The blank sky positions must satisfy two criteria: their number density on the sky should be high enough to ensure we can always configure at least 25 sky positions in any given 2dF observation, and each must be sufficiently far from any other object on the sky to avoid contamination. We chose a density of 40 sky positions per square degree, corresponding to 120 per 2dF observation. This was judged to be sufficient to allow 25 sky fibres to be configured after all the target and guide fibres were positioned.   We selected the sky positions automatically from the image catalogues for each region by setting down a square grid of the desired surface density. Any grid positions too close to a catalogue image were moved (randomly) within the corresponding grid cell as necessary until they satisfied a proximity test.  The proximity test for neighbouring objects was a function of their apparent $r$ (or $R$) magnitude: the sky position was moved if there were neighbours satisfying $r<12.5$, $r<16$, $r<17.5$, or $r\\ge 17.5$, within distances of 1.5, 1, 0.5, or 0.25 arc minutes respectively.   \\subsubsection{Selection of Guide Star Positions}  The guide stars for AAOmega must satisfy strict photometric and astrometric criteria. Most importantly, their positions must be accurate and on the same reference frame as the target objects. The new 2dF/AAOmega system has 8 fibre guide bundles on each plate so, in principle, observations are still possible if one or two of the guide stars are not suitable. However the diagnosis of the defective stars slows down the observing process, to the extent that they would amount to a significant time loss over the course of a large survey like WiggleZ.  To satisfy the astrometric criterion we selected guide stars from different sources in each of the two survey regions. In the northern SDSS survey regions we selected guide stars directly from the same SDSS catalogues as we used for our optical target photometry. We used the SDSS automated image classification flags to select stars from the catalogues (``mode''=1, ``type''=6, ``probPSF''=1).  In the southern (RCS2) regions we were unable to obtain guide star positions from our imaging data as they are much deeper and hence most guide stars are saturated, so they have unreliable positions. We therefore chose guide stars for these regions directly from the USNO-B astrometric catalogue \\citep{Monet2003} which was used to provide the astrometry for the RCS2 data (with fainter USNO-B stars that were not saturated in the RCS2 images). We undertook extensive testing of the USNO-B guide star astrometry, comparing the positions to both SDSS data (in the limited regions where they overlap) and also with the all-sky 2-MASS catalogue data \\citep{Skrutskie2006}. In the regions tested we found small but significant systematic position differences between 2-MASS and USNO-B; for the SDSS versus 2-MASS comparison, we found no systematic difference. In all cases the r.m.s.\\ scatter in the position differences (less than 0.2 arc seconds in each coordinate) was consistent with the published uncertainties in the respective catalogues. The offsets are tabulated in Table~\\ref{tab-offset}.  We have not adjusted the RCS2 or USNO-B astrometry for the small systematic offsets with respect to the SDSS/2-MASS astrometry, as this is not necessary for our spectroscopic observations.  There is no measurable offset between the astrometry of the RCS2 data (our targets) and the USNO-B catalogue (our guide stars), so AAOmega observations using these guide stars will ensure the target fibres are correctly positioned.   \\begin{table} \\caption{Astrometry Offsets in RCS2 Regions} \\label{tab-offset} \\begin{tabular}{rrr} \\hline Name & RA$_{USNO}-$RA$_{2MASS}$ & Dec$_{USNO}-$Dec$_{2MASS}$ \\\\ & (arc seconds)     & (arc seconds)  \\\\ \\hline 0-hr &  $-0.112\\pm 0.003$ &  $0.117\\pm 0.003$  \\\\ 1-hr &  $-0.133\\pm 0.004$ &  $0.102\\pm 0.003$  \\\\ 3-hr &  $-0.141\\pm 0.002$ & $-0.026\\pm 0.002$  \\\\ 22-hr &  $-0.031\\pm 0.003$ &  $0.137\\pm 0.003$  \\\\ \\hline \\end{tabular} \\end{table}  We did find that the USNO-B guide star positions, when compared to 2MASS, displayed a measurable tail (1--2\\% of the total) in the distribution of position differences, with them being greater than 0.5\\,arcsec, a size not seen with the SDSS data. Inspection of the images in this tail showed many of these objects were not single stars, but galaxies or binary stars. We therefore imposed a further check on the USNO-B guide stars, that their positions had to be consistent with the corresponding 2MASS data (after adding the systematic offset to put them on the USNO-B system) to within 0.5 arc seconds. We did not apply any image classification test to the USNO-B stars as the classifications were less reliable than for SDSS and, at these magnitudes, the majority of images are stellar.  The guide star magnitudes for 2dF must be in the range 12--14.5\\,mag, and must not differ in brightness by more than 0.5 mag to not exceed the dynamic range of the guide camera. Since the WiggleZ project is conducted largely in the dark of moon, we opted for stars in the faintest range $14<r<14.5$ to obtain as many as possible per field. In order to achieve good astrometry it is also important to avoid stars with high proper motion. We checked this using (USNO-B) catalogue data when available, but as a further check, we also avoided very red stars, as many of these are nearby M-dwarf stars which are more likely to have significant proper motions.  The guide star selection criteria for the two survey regions are summarised in Table~\\ref{tab-guides}.   \\subsubsection{Visual Inspection}  Finally, all the potential blank sky and guide star positions were inspected visually by team members, using SuperCOSMOS \\citep{Hambly2001} images at each position. This step is necessary because our input catalogues (like all automated image catalogues) suffer from a small rate of object mis-classification and even missing objects. These mistakes are not statistically significant, but could potentially ruin a complete observation.  For the sky positions, the visual check is simply to ensure there really is no object near the nominated position. The check is more demanding for the guide stars: we must ensure that they are really stellar, not multiple, and do not have any close neighbours (within about 30 arc seconds).  A total of some 50,000 different sky and guide star positions were defined across the different regions of the WiggleZ Survey, all of which were visually checked by a team member. To speed this process up, we designed a web-based interface that displayed large sets of SuperCOSMOS images, 1000 or more at a time, and allowed the team member to easily identify the small number of problematic positions.   \\begin{table} \\caption{The selection criteria for guide stars} \\label{tab-guides} \\begin{tabular}{lll} \\hline Criteria & SDSS  &  RCS2 (USNO-B)  \\\\ \\hline Magnitude & $14\\le r \\le 14.5$ & $13.95\\le R \\le 14.45$ \\\\ Colour    & $-0.5\\le (g-r) \\le 1.5$ & $-0.3\\le (B-R) \\le 1.7$ \\\\ Classification& stellar & any \\\\ Proper Motion& \\multicolumn{2}{c}{$p<0.015$ arc seconds yr$^{-1}$} \\\\ \\hline \\end{tabular} \\end{table}  \\subsection{2dF  field placement}  When the WiggleZ survey began taking data with AAOmega, the GALEX data set that formed the input catalogue for our survey was patchy and had a highly varying density of targets.  Although both of these factors have since improved tremendously, we do not use a simple grid of 2dF field placements as it invariably produces highly pathological fields that the 2dF robotic positioner is unable to configure in sufficient time (i.e.\\ less than the 1-hour typical observing time for each field).  Instead, we apply a simulated annealing algorithm to the problem of 2dF field placement.  The algorithm used is identical to that used by \\citet{Campbell2004} \\citep[see also][]{Metropolis1953} for the 6dF Galaxy Survey \\citep[e.g.][]{Jones2004}, but modified for use with the 2dF system.  Briefly, the algorithm generates an arrangement of 2dF placements that ensures that the smallest number of 2dF fields are used to give as large as possible coverage of the target set with as little variation in surface density as possible (see Fig.~\\ref{fig-anneal}).  \\begin{figure*} \\epsfig{file=fig_annealing.eps,width=\\figwidth\\linewidth,angle=0} \\vspace*{-8.5cm} \\caption{2dF field placement for the 15hr rectangle resulting from simulated annealing prepared for the 2008 April observing run. The target set is non-uniform with significant areas of GALEX targets yet to be observed (blank areas) and a number of targets within the available GALEX tiles having already been observed in previous runs (e.g.\\ note the apparent under-dense area at RA$\\sim$160 and Dec$\\sim$1).  The simulated annealing approach ensures that a minimum number of 2dF positions is used to observe the targets.} \\label{fig-anneal} \\end{figure*}   The annealing is performed prior to each observing run in order to have a fully efficient 2dF tile set to observe with.  In practice, one issue that regularly arises with this process is that the annealing algorithm simply aims to fill each 2dF field with $\\sim$350\\footnote{We enter 350 fibres as the typical number of target fibres available per 2dF field.  Combined with 8 guide star fibres, up to 25 sky fibres for sky subtraction and a variable number of broken fibres at any one time, 350 target fibres represents a good estimate as to the available number of fibres.}  fibres.  In regions where the target density is high such that there are many more targets than the 350 the annealing algorithm is asked to assign fibres to, the result can potentially be a 2dF field where there is a large target density gradient within that field.  This is especially true of areas that consist of a single, non-contiguous GALEX tile (Fig.~\\ref{fig-anneal})---the annealing algorithm does not know a priori to place the 2dF field such that it is centred on the GALEX tile's centre in order to make the fibre positioning run faster. Therefore we perform a small amount of manual intervention to optimize the resultant 2dF field placement by, for example, putting single GALEX fields in the middle of a 2dF circle.  Our later publications \\citep[e.g.][]{Blake2009b} will detail how we take account of the overlapping regions for the window function of our survey.  \\subsection{Observing priorities}  The WiggleZ survey has seven equatorial fields roughly distributed over the full range of R.A. thus there are target fields readily available throughout the year. Of course, this will change as the survey progresses since more time will have to be expended on observations of the more incompletely sampled target regions. The target selection within a given WiggleZ survey region is necessarily complex due to the rapid evolution of our data set and the complicated selection criteria (Section~\\ref{sec-selection}). However, choosing an individual field pointing for a given night is mostly based on minimizing its hour angle over the course of the observation. Other factors typically considered when prioritizing field centres to be observed, include: (1)\\,the field is part of a larger contiguous area in a WiggleZ region, (2)\\,it contains a high target density (i.e., it will use all of the fibres), and (3)\\,it does not lie on the edge of a GALEX tile or observed area (anticipating the GALEX data will be more contiguous in future observing runs).   \\subsection{Final target lists and calibration} \\label{sec-final}  Once the field pointings have been selected for a night of observing, combined target files for each field pointing are generated using custom software (developed by RJJ). These files combine the target galaxies (with suitable priorities assigned; see Table~\\ref{tab-priority}), guide stars, and blank sky positions. The software also reads the lists of objects observed earlier in the run to avoid repeat observations in overlapping regions. For the full 2-degree diameter field of view we generate a target file with an optimal number of 700--800 possible targets.  Our target catalogues are updated before each observing run to record which galaxies have already been observed. These are not normally re-observed with the exception that targets that were previously observed in poor conditions (typically with cloud cover more than 4/8 or seeing FWHM more than 3 arc seconds, resulting in a significantly reduced redshift completeness, less than 50 per cent) without obtaining an acceptable redshift are re-observed (see Table~\\ref{tab-priority}). Galaxies which were previously observed in good conditions are put in a reserve list and are occasionally re-observed at low priority ($P=2$ if no redshift was obtained; $P=1$ if a redshift was obtained).  A small number of additional calibration targets are included in the target file at high priority. First, where possible we include two standard stars: white dwarfs or hot sub dwarfs from the SDSS compilation by \\citet{Eisenstein2006}. These are used to assist in future spectrophotometric calibration. Secondly, in order to determine the reproducibility of our measurements, we include repeat observations of three previously allocated targets and three previously redshifted objects.  In regions where we are unable to use all the fibres on high-priority WiggleZ targets, we allocate a small number of targets for companion projects to the main survey: candidate cluster galaxies from the RCS2 project and candidate radio galaxy identifications from the FIRST \\citep{Becker1995} survey. These additional targets are placed at lower priority than all the new WiggleZ targets ($P=3$) and, in total, they account for less than 1 per cent of the observations made.  Finally, we note that a small number of the spectra from each plate at any time (up to 5 per cent) are affected by a fringing pattern superposed on the target spectrum. We are unable to measure redshifts for spectra strongly affected by the fringing.  The interference fringing is caused by small air gaps between the ends of the 2dF fibres and the 90 degree prisms that collect the target light on the focal plane.  The air gaps create etalons, introducing interference fringes which modulate the fibre transmission profile as a function of wavelength.  The air gaps are the result of thermal expansion between the steel fibre mounts and the glass fibres themselves.  The gap size is mechanically unstable, relaxing during each exposure after being disturbed during fibre placement, so the fringes are not removed by the flat field correction in the data reduction process. A hardware solution to the problem has been identified by the AAO, and will be applied in due course.  Allowing for the fringing fibres, as well as those devoted to background sky measurement and the various calibration measurements, the average number of fibres devoted to new targets per observation in the survey data is 325 (as quoted in Table~\\ref{tab-design}.   \\subsection{Field configuration}  The 2dF {\\sc configure} software \\citep{Shortridge2006} is then used to generate configuration files from the target lists. This optimally allocates all the fibres in use on the telescope to our targets. The configuration central wavelength is set to 670 nm due to our use of the new WiggleZ dichroic beam splitter. Within the {\\sc configure} program, we generally choose to weight the peripheral fiducial targets, use 25 sky fibres in a pointing and enforce a sky fibre quota.  A simulated annealing algorithm incorporated into recent versions (after v.7.3) of the {\\sc configure} code \\citep{Miszalski2006} has greatly improved the yield of high priority targets and the spatial uniformity.  Furthermore, high speed computers installed in the AAT control room in 2007 October, allow us to calculate optimal configurations for a typical field in less than half an hour.  The output of the {\\sc configure} software is a configuration file specifying the position of every fibre, which is then used by the 2dF positioner to place all the fibres on the one of the two field plates that is to be used for the observation. The efficiency and reliability of the positioner was improved significantly during the first year of our observations, leading to positioner configuration times of less than one hour for all the fibres.  These short configuration times provide greater observing flexibility, particularly during good observing conditions. During these clear periods of good seeing we can then expose for less than our conventional one-hour exposure, and still obtain the appropriate signal-to-noise ratio required for successful target redshifting.  With a large number of observations now completed, we can illustrate the success of this process by plotting the physical positions of the fibres from all our observations. This is shown for both plates combined in Figure~\\ref{fig-fibres1}. Note that we have not included data taken in the period 2007 April to 2007 August; during this period we restricted our observations to a maximum radius of 0.7\\,degrees from the centre of each 2dF field as there was a problem with the astrometric calibration of the 2dF positioner which reduced the accuracy of fibre placement at larger radii.  The fibre distribution in Figure~\\ref{fig-fibres1} reveals 8 areas of slightly lower density around the edge of the field, associated with the position of the 8 (equally spaced) fibre bundles used to measure the fiducial stars. A similarly weak effect was present in the 2dF Galaxy Redshift Survey (\\cite[2dFGRS;][]{Colless2001} but only 4 fiducial stars in that case). The 2SLAQ galaxy survey \\citep{Cannon2006}, by contrast, had a very marked discontinuity at the field edges as only one of the 2dF spectrographs was used (alternate banks of fibres were used for the companion quasar survey).  The fibre distribution in Figure~\\ref{fig-fibres1} also displays a strong radial gradient in surface density. The average fibre density decreases by a factor of about 3 from the centre to the edge of the field. This is a natural consequence of the smaller size of the GALEX field compared to the 2dF field: in many observations we do not fill the edges of the 2dF field with targets.  \\begin{figure} \\epsfig{file=fig_fibres1.eps,width=\\figwidth\\linewidth,angle=0} \\caption{The average distribution of fibres observed on the two 2dF field plates. Each dot represents one fibre observed in the first two years of the project excluding 2007 April--August when we only observed inside the central 0.7 degree radius. The smooth distribution (apart from very small regions at the edge due to the eight guide fibre bundles) demonstrates that no artificial spatial signal is introduced by the observing process. } \\label{fig-fibres1} \\end{figure}  \\subsection{Data reduction}  The AAOmega data are reduced during each observing run using the automated {\\sc 2dfdr} software developed at the AAO. We take standard calibration flat field tungsten and arc (CuAr, FeAr, Ne and He) exposures for each pointing throughout the night, whose integration times are typically 3\\,sec and 30\\,sec, respectively.  A two-dimensional bias frame is constructed for the blue arm data to correct for a small number of high bias columns on the blue CCD.  The data from the blue and red arms for a field pointing are reduced separately, and then spliced together resulting in a final spectrum used for target redshifting. The high-speed computers available at the AAT make it possible to reduce the data in real time, allowing the observers to have the majority of the data ready for redshift measurement by the end of the night. We note that, as in the case of most fibre-fed spectroscopy, we do not flux calibrate the spectra. This is mainly due to the inherent difficulty in calibrating the fraction of total light collected by each fibre from the corresponding target. In a later paper, we will present the results of producing approximate flux-calibrated spectra from the data using the white dwarf targets we include where possible in the SDSS fields (see Sec.~\\ref{sec-final}) and/or the optical photometry.  Significant refinement and further optimization of the {\\sc 2dfdr} code has been undertaken by one of us (SMC) and we have used these new versions of the software to have the most reliable, efficient reductions possible for our observing runs. Laplace filter cosmic ray rejection, Gaussian spectral extraction, 1D scattered light filtering and optimal sky subtraction have generally been used throughout the {\\sc 2dfdr} reductions \\citep[see][for more details]{Croom2005}. Once the reduced, spliced file is generated for a field, as a final reduction step we apply a principal component analysis (PCA) sky subtraction algorithm to the spectra. This provides for better removal of the OH sky line contamination at the red end of the spectra, thus making it easier to identify redshifts. The PCA code we use was adapted (by KG and EW) from software developed by \\citet{Wild2005}. Our standard approach for implementing the PCA sky subtraction is to generate an eigenspectra file from the first 8--10 fields obtained on an observing run. This eigenspectra file is then used to PCA sky subtract the rest of the spectra obtained for a given run. We have found the PCA sky subtraction is a useful step in that it makes the redshifting a less onerous task by cleaning up the ``cosmetics'' of the spectra, particularly for data taken in good observing conditions.  \\section[]{Redshift analysis} \\label{sec-redshifts}  In this section we describe the process of redshift measurement from our spectroscopic data, as well as various tests of the reliability of these measurements.   \\subsection{Measurement}  All redshifts in the WiggleZ survey are measured using an evolved version of {\\sc runz} --- the software used for both 2dFGRS \\citep{Colless2001} and 2SLAQ \\citep{Cannon2006}.  The original version of {\\sc runz} made use of both discrete emission line redshift determination and absorption line redshifts via template Fourier cross-correlation \\citep{Tonry1979}.  In the WiggleZ survey, it is rarely possible to make use of absorption lines---the survey is predicated on the ability to determine redshifts from emission lines alone. Therefore, {\\sc runz} has been modified significantly (by SMC), optimizing it to measure redshifts from emission lines. The emission line algorithm searches for sharp peaks in the spectrum, as candidate emission lines, and then tries to match sets of these to known combinations of strong lines at different redshifts.  The lines used are: [OII]$\\lambda$3727, H$\\beta$, [OIII]$\\lambda$4959,[OIII]$\\lambda$5007, H$\\alpha$, and [NII]$\\lambda$6583. We note that the adopted wavelength for the [OII]$\\lambda$3727 doublet is not the mid-point of the two transitions (which is 3727.42\\AA), but rather is $\\sim$0.4\\,\\AA\\ redder to allow for the average ratio found in our spectra.  For the cross-correlation measurement, {\\sc runz} uses a revised set of template spectra including high-redshift galaxies and QSOs, but the continuum level in typical spectra is so low that the cross-correlation redshifts are very rarely used.  Uncertainties in the redshift measurements are estimated in two different ways by {\\sc runz}. In the more common case when the redshifts are measured from emission lines, the redshifts are determined for each individual line by fitting Gaussian profiles, using the variance spectrum to determine the fitting errors.  The final emission line redshift (and error) is the variance-weighted mean of that derived from the individual lines.  For cross-correlation redshifts the error is derived from the width of the cross-correlation peak \\citep{Tonry1979}.   In practice, complete automation of the redshift measurements has proven to be problematic due to the noisy nature of many of the spectra and the presence of artifacts such as residuals from imperfect cosmic ray and sky removal and fibre fringing and cross-talk.  After a redshift has been assigned using emission lines or the method of \\citet{Tonry1979}, {\\sc runz} automatically generates an integer quality flag ($Q$) in the range 1--5 based on how well the template fits to the given spectra.  The user then inspects the given redshift. In many cases, the user manually re-fits the spectrum and assigns a new redshift and quality flag, in accordance with the criteria in Table~\\ref{tab-quality}.  In acknowledging that the assignment of $Q$ values varies between users, we note that in any future analysis only redshifts with values of $Q\\ge 3$ will be utilized.  Therefore, whilst there may be some debate as to what constitutes the difference between $Q=4$ and $Q=5$, the critical separation in quality occurs between $Q=2$ and $Q=3$.  This can be seen in Figure~\\ref{fig-quality} which displays redshift histograms for both $Q\\le 2$ and $Q=3$ redshifts for data collected up until 2008 March.  For the $Q=2$ redshift distribution, we see that there are multiple peaks that coincide with assigning [OII] to sky lines---such peaks are (with two exceptions) not present in the $Q=3$ distribution. Inspection of the two narrow peaks still visible in the $Q=3$ data shows that the first ($z=0.707$) corresponds to mis-matching [OII] to the weak sky line at 636.2 nm. The second peak ($z=0.761$) corresponds to mis-matching [OII] to rest-wavelength Galactic H$\\alpha$ emission at 656.3 nm which occurs in a few fields. The number of spectra in these two peaks is less than 0.1 per cent of the total sample.  \\begin{figure} \\epsfig{file=fig_quality.eps,width=\\figwidth\\linewidth,angle=0} \\caption{Redshift distribution as a function of the quality flag $Q$. The top two panels are for accepted measurements ($Q$ of 3 and above). These show that objects at higher redshifts (with the H-$\\beta$ line not visible) are generally not assigned the highest quality flags. The bottom panel is for rejected measurements ($Q\\le 2$): this has many sharp peaks in the distribution corresponding to misidentified sky lines. With two exceptions (see text), these sharp peaks are not evident in the good ($Q\\ge 3$) redshift distributions.} \\label{fig-quality} \\end{figure}   \\begin{table} \\caption{Redshift Quality Definitions} \\label{tab-quality} \\begin{tabular}{lrl} \\hline Q & \\% & Definition \\\\ \\hline 1 & 21 & No redshift was possible; highly noisy spectra.\\\\ 2 & 19 & An uncertain redshift was assigned.\\\\ 3 & 18 & A reasonably confident redshift; if based on [OII]  \\\\ &    & alone, then the doublet is resolved or partially resolved.\\\\ 4 & 33 & A redshift that has multiple (obvious) emission  \\\\ &    & lines all in agreement.\\\\ 5 &  9 & An excellent redshift with high S/N that may be  \\\\ &    & suitable as a template.\\\\ 6 &  0.5 & Reserved for Galactic stars used as calibration sources.\\\\ \\hline \\end{tabular}  Note: the second column lists the percentage of all observations (including repeats) in each category. \\end{table}  We present examples of a variety of WiggleZ spectra in Figure~\\ref{fig-spectraQ45} (high quality Q=4-5), and Figure~\\ref{fig-spectraQ3} (low quality Q=3).  The spectra include examples of those providing single [OII] doublet redshifts and standard multiple emission line detections. In Figure~\\ref{fig-spectraXX} we also present a random selection of QSO and star spectra from the survey. This demonstrates the relatively low redshifts of the few QSOs selected in the survey; many are only identified from the MgII emission line.   \\begin{figure*} \\epsfig{file=fig_spectraQ45.eps,width=0.9\\linewidth,angle=0} \\caption{Examples of randomly-selected high-quality (Q=4,5) spectra from the survey. For each spectrum, the upper panels (red) show the unsmoothed spectra zoomed in on the major emission lines, and the lower panel (blue) shows the whole spectrum, heavily smoothed in an optimal fashion. } \\label{fig-spectraQ45} \\end{figure*}   \\begin{figure*} \\epsfig{file=fig_spectraQ3.eps,width=0.9\\linewidth,angle=0} \\caption{Examples of randomly-selected low-quality (Q=3) spectra from the survey. For each spectrum, the upper panels (red) show the unsmoothed spectra zoomed in on the major emission lines, and the lower panel (blue) shows the whole spectrum, heavily smoothed in an optimal fashion.} \\label{fig-spectraQ3} \\end{figure*}  \\begin{figure*} \\epsfig{file=fig_spectraXX.eps,width=0.9\\linewidth,angle=0} \\caption{Examples of randomly-selected QSO and star spectra from the survey. For each object the whole spectrum is shown, heavily smoothed in an optimal fashion.} \\label{fig-spectraXX} \\end{figure*}    \\subsection{Redshift Reliability}  The one hour exposures that the WiggleZ survey uses are too short to allow any significant detection of the continuum, with a continuum S/N $\\sim$ 1 per (approx 0.14 nm) pixel. However, by design, we are able to reliably and accurately measure redshifts from the multiple bright emission lines in the WiggleZ galaxies. We commonly detect the [OII]$\\lambda$3727, H$\\beta$, and [OIII]$\\lambda\\lambda$4959, 5007 emission lines in the spectra. At low redshifts (z $< 0.26$), when the [OII] line moves outside of our wavelength range this is compensated by the ability to detect H$\\alpha$ in our spectra. At higher redshifts (z $>$ 0.95) we often rely on redshifts obtained from a single line; however, the spectral resolution is sufficient to see this line significantly broadened at these redshifts (the [OII] doublet is often resolved at z $\\geq$ 0.8) which allows us to confidently conclude this single line is [OII] and obtain the corresponding redshift. The reliability and accuracy of using the broadening to confirm a single line in a spectrum as the [OII] line of a high redshift emission line galaxy have already been demonstrated by \\citet{Breukelen2007}.  \\subsubsection{Internal Tests} \\label{sec-internal}  We have multiple $Q\\geq 3$ redshift measurements for over 6800 of the galaxies observed to date, for both the SDSS and RCS2 regions. These multiple ``correct'' redshift measurements can be used to estimate the accuracy of our redshift measurements, and assess the effectiveness of our redshift quality scheme. For the following analysis we define redshifts to be reliable (i.e.\\ the correct lines have been identified) if they agree within $|\\Delta z| < 0.002$. When converting to velocity differences, we calculate these according to $\\Delta v = c \\Delta z / (1 + z)$.  We start by assessing the internal agreement within each of the three acceptable quality categories, $Q=3$--5, by taking objects whose original and repeated measurements both have the same $Q$ value. There are 737 $Q=3$ repeats, 2263 $Q=4$ repeats, and 619 $Q=5$ repeats. The fractions of the redshift pairs that agree within each category are 68.5, 95.1, and 100 percent respectively. For the correctly identified redshifts, the mean difference is not significantly different from zero in any of the categories, and the standard deviations are $\\sigma_{\\Delta z Q=3}=0.0042$ (68.5 \\kms), $\\sigma_{\\Delta z Q=4}=0.0031$ (58.2 \\kms), and $\\sigma_{\\Delta z Q=5}=0.0022$ (41.6 \\kms) respectively. (These values are for the differences of two measurements so the uncertainty in an individual measurement would be a factor of $\\sqrt{2}$ times smaller.)  As expected, and as shown in Fig.~\\ref{fig-zdiff} the scatter in the redshift differences decrease with increasing $Q$.  Given the high reliability of the $Q=4,5$ measurements, we can now test a larger sample of $Q=3$ objects for which the repeated measurement has a higher quality ($Q\\ge4$) and can therefore be regarded as correct. There are 2250 such objects, of which 78.7 per cent are correct; the standard deviation in the redshift differences is $\\sigma_{z Q=3-4/5}=0.0039$ (67.5 \\kms). The fraction of reliable measurements is consistent with the internal value (68.5 per cent) reported above. Similarly, 1425 $Q=4$ objects are repeated by higher $Q=5$ measurements, of which 98.3 are correct and $\\sigma_{z Q=4-5}=0.0029$ (54.0 \\kms).  The average reliability and uncertainty values for the three quality categories are summarised in Table~\\ref{tab-redshift}. These values are consistent with earlier estimates made by \\citet{Blake2009a} for a subset of the current data.  \\begin{table} \\caption{Reliability and Uncertainty of Redshift Measurements} \\label{tab-redshift} \\begin{tabular}{crcc} \\hline Quality  &  Reliability & \\multicolumn{2}{c}{Uncertainty}  \\\\ $Q$  &              & $\\sigma_z$ & $\\sigma_v$ (\\kms)   \\\\ \\hline 3  &  78.7 \\%  &  0.00030  & 48.4 \\\\ 4  &  98.3 \\%  &  0.00022  & 41.2 \\\\ 5  & 100.0 \\%  &  0.00016  & 29.4 \\\\ \\hline \\end{tabular}  Notes: the reliability fractions were obtained by comparing with repeated measurements at higher quality ($Q$) values. The measurement uncertainties $\\sigma z, \\sigma v$ were obtained by scaling the internal standard deviations by $1/\\sqrt{2}$. The velocity differences for each measurement were calculated as $\\Delta v = c \\Delta z / (1 + z)$. \\end{table}  The comparisons of the repeat observations may also be used to check the internal redshift uncertainties $\\sigma z_{runz}$ for each measurement estimated by the {\\sc runz} code. Those internal uncertainties have a median value of $\\sigma z_{runz}\\approx 0.00015$ for all three quality bins combined, but the distribution is very non-uniform, with a long tail to higher values. Considering just this median estimate of $\\sigma_z =0.00015$, we find that it corresponds well to the $\\sigma_z = 0.00016$ derived for the $Q=5$ measurements in Table~\\ref{tab-redshift}. The empirical results for the $Q=3,4$ uncertainties in the table are larger than this, so in these cases the internal redshift uncertainty may be an underestimate.   \\begin{figure} \\epsfig{file=fig_zdiff.eps,width=\\figwidth\\linewidth,angle=0} \\caption{Redshift differences from repeat observations of WiggleZ targets. The three histograms show the distribution of redshift differences in pairs of repeated WiggleZ observations. The solid, dashed and dotted curves are for pairs in which the qualities of the both measurements are 3, 4, and 5 respectively.} \\label{fig-zdiff} \\end{figure}  We also examined the redshift differences as a function of redshift, for all galaxies with multiple $Q\\geq 3$ redshift measurements. We do not find any systematic variation with redshift of the mean or standard deviation of the differences. We conclude that the redshift uncertainties quoted above are valid over the entire redshift range of the WiggleZ survey. This is, however, a trend in the reliability of the $Q=3$ redshift measurements: the reliability decreases from about 80 per cent to about 40 per cent between redshifts of $z=0.8$ and $z=1.2$ \\citep[see Fig.~6 of][]{Blake2009a}.   \\subsubsection{External Tests} \\label{sec-external}  At the time of writing, there are very few other galaxy surveys in our redshift range that overlap with the WiggleZ survey, so there are few opportunities for external comparison. There is a small overlap between WiggleZ and Data Release 3 of the DEEP2 survey \\citep{Davis2003} in our 22-hour field. Although the DEEP2 survey (in that region) is aimed at $z>0.5$ targets, there are 53 sources in common with WiggleZ (positions matching to within 2.5 arc seconds). The estimated r.m.s.\\ uncertainty in the DEEP2 redshifts is 30 \\kms with less than 5 per cent incorrectly identified \\citep{Davis2003}.   We compare the WiggleZ and DEEP2 redshifts in Fig.~\\ref{fig-deep2}. Of the 34 sources that both surveys assign good redshifts to, only 3 disagree substantially (due to different line identifications).  Of the redshifts that disagree, all three are $Q=3$ WiggleZ measurements; no $Q=4$ or 5 WiggleZ measurements disagree with the DEEP2 measurements. On careful inspection of the WiggleZ spectra and the corresponding DEEP2 spectra we conclude that in two cases the DEEP2 value is correct; the third case is ambiguous so we cannot determine which is correct. This rate of misidentification among the $Q=3$ WiggleZ redshifts (20-30\\% of the $Q=3$ redshifts are wrong) is entirely consistent with the internal estimates in Section~\\ref{sec-internal} above.  Restricting the sample to those that agree within $|\\Delta z|<0.002$, there is no significant mean difference ($\\overline{\\Delta z}=-0.00003\\pm 0.00002$) and the rms difference is $\\sigma_{\\Delta z}=0.00047$ (80.5 \\kms). This scatter is entirely consistent with the estimates of the internal uncertainties in the respective surveys quoted above.  \\begin{figure} \\epsfig{file=fig_deep2.eps,width=\\figwidth\\linewidth,angle=0} \\caption{Comparison of WiggleZ redshifts with DEEP2 survey data. The left-hand panel compares the different redshifts with the symbols indicating the reliability status of the measurements in the respective surveys. The right-hand panel shows the histogram of redshift differences.} \\label{fig-deep2} \\end{figure} ",
        "conclusions": "\\label{sec-summary}  The WiggleZ Dark Energy Survey will be the first large-volume spectroscopic galaxy survey to measure the BAO scale in the key redshift range of $0.2<z< 1.0$. In this paper we have described the design and initial results of the WiggleZ Survey at the half-way point which coincides with our first public data release.  To emphasise the large size of the WiggleZ survey we note that, as of 2008 December, we have measured 119,000 galaxy redshifts, of which 73 000 have redshifts $z>0.5$. This means we have already more than doubled the number of known galaxy redshifts above $z=0.5$ from the two largest existing surveys: DEEP2 \\citep{Davis2003} with $N_{z>0.5}\\approx 28,000$ and the VIMOS VLT Deep Survey \\citep[VVDS,][]{Garilli2008} with $N_{z>0.5}\\approx 18,000$. At the time of writing, the total number of $0.5<z<1.5$ redshifts published from all surveys is 100,000 (as listed by NED). The WiggleZ survey, when completed, will more than double this quantity.  Our survey design relies on the use of GALEX UV satellite data to efficiently select high-redshift emission line galaxies. These objects, whilst fainter than giant red galaxies at the same redshift, can be observed very efficiently on a 4-metre class telescope because of their strong emission lines. We have added some optical colour selection to our UV colour selection to further improve the fraction of galaxies with high redshifts; the median redshift of the measured galaxies is now $z_{med}=0.63$. The high sensitivity of the new AAOmega spectrograph on the AAT and the short fibre configuration time of the 2dF positioner mean that we can observe a new field every 1.2 hours allowing us to make rapid progress on the survey.  In this paper we have also presented our first public data release, available from a dedicated web interface. The public data comprise photometry (optical and UV), redshifts and spectra of all objects observed up to the end of Semester 2008A.  We cannot make the BAO clustering measurement at this stage of the project. This is not only because we only have half the necessary sample but also because the window function is very irregular due to our highly variable coverage of the fields to date.  However, at the half-way stage we already have the largest ever sample of high redshift emission line galaxy spectra. Most importantly, we have measured the small-scale clustering of these UV-luminous galaxies and found it to be relatively strong which increases the accuracy with which we will be able to measure the BAO scale. At lower redshifts where we can measure the H$\\beta$ line we have shown that the WiggleZ galaxies are mostly extreme starburst-type galaxies as expected. We have compared the WiggleZ galaxies to optically-selected galaxies at the same redshift and confirmed that they are extremely blue with high star formation rates. There is some evidence from {\\em Hubble Space Telescope} imaging of a subset of the galaxies that the starburst activity is being driven by interactions."
    },
    "0911/0911.0171.txt": {
        "abstract": "{\\bf Abstract:} Aegaeon (Saturn LIII, S/2008 S1) is a small satellite of Saturn that orbits within a bright arc of material near the inner edge of Saturn's G ring. This object was observed in 21 images with Cassini's Narrow-Angle Camera between June 15 (DOY 166), 2007 and February 20 (DOY 51), 2009. If Aegaeon has similar surface scattering properties as other nearby small Saturnian satellites (Pallene, Methone and Anthe), then its diameter is approximately 500 m. Orbit models based on numerical integrations of the full equations of motion show that Aegaeon's orbital motion is strongly influenced by multiple resonances with Mimas. In particular, like the G-ring arc it inhabits, Aegaeon is trapped in the 7:6 corotation eccentricity resonance with Mimas. Aegaeon, Anthe and Methone therefore form a distinctive class of objects in the Saturn system: small moons in co-rotation eccentricity resonances with Mimas associated with arcs of debris. Comparisons among these different ring-arc systems reveal that Aegaeon's orbit is closer to the exact resonance than Anthe's and Methone's orbits are. This could indicate that Aegaeon has undergone significant orbital evolution via its interactions with the other objects in its arc, which would be consistent with the evidence that Aegaeon's mass is much smaller relative to the total mass in its arc than Anthe's and Methone's masses are. % ",
        "introduction": "Beginning in early 2004, images from the cameras onboard the Cassini spacecraft revealed the existence of several previously unknown small Saturnian satellites: Methone, Pallene, Polydeuces, Daphnis and Anthe \\citep{Porco05, Murray05, Spitale06, Cooper08}. Two of these moons --Anthe and Methone-- are in mean-motion resonances with Saturn's moon Mimas. Specifically, they occupy the 10:11 and 14:15 co-rotation eccentricity resonances, respectively \\citep{Spitale06, Cooper08, Hedman09}. Both of these moons are also embedded in very faint, longitudinally-confined  ring arcs \\citep{Roussos08, Hedman09}. This material probably represents debris that was knocked off the relevant moons at low velocities and thus remains trapped in the same co-rotation resonance as its source body.  Images from Cassini also demonstrated that a similar arc of material exists within Saturn's G ring, around 167,500 km from Saturn's center \\citep{Hedman07}. Images of this structure taken over the course of several years showed that it was also confined by a (7:6) corotation eccentricity resonance with Mimas. Furthermore, in-situ measurements of the plasma environment in the vicinity of the arc suggested that it contains a significant amount of mass in particles larger than the dust-sized grains that are the dominant source of scattered light observed in images \\citep{Hedman07}.   \\nocite{Porco09} In late 2008, during Cassini's Equinox Mission (2008-2010), images of the arc taken at lower phase angles and higher resolutions than previously possible revealed a small, discrete object.  Since the object was most visible at low phase angles and could be tracked over a period of roughly 600 days, it is almost certainly not a transient clump of dust but instead a tiny moonlet that represents the largest of the source bodies populating the arc.  The discovery of this object was therefore announced in an IAU circular, where it was designated  S2008/S1 (Porco {\\it et al.} 2009). More recently the International Astronomical Union has given it the name Saturn LIII/Aegaeon. As will be shown below, Aegaeon, like Anthe and Methone, occupies a corotation eccentricity resonance with Mimas, and all three of these small moonlets are associated with arcs of debris. These three objects therefore represent a distinct class of satellites and comparisons among the ring-moon systems have the potential to illuminate the connection between moons and rings.  Section 2 below describes the currently available images of Aegaeon and how they are processed to obtain estimates of the brightness and position of this object. Section 3 presents a preliminary analysis of the photometric data, which indicate that this object is approximately 500 m in diameter. Section 4 describes the orbital solutions to the astrometric data, which demonstrate that Aegaeon's orbit is indeed perturbed by the 7:6 corotation eccentricity resonance with Mimas. However,we also find that a number of other resonances, including the 7:6 Inner Lindblad Resonance, strongly influence Aegaeon's orbital motion. Finally, Section 5 compares the various resonantly-confined moon/ring-arc systems to one another in order to clarify the relationship between Aegaeon and the G ring. ",
        "conclusions": "Even though the currently available data on Aegaeon are sparse, they are sufficient to demonstrate that it is a interesting object worthy of further investigation. With a photometrically estimated diameter of less than a kilometer, Aegaeon is the smallest isolated moon of Saturn yet observed, and may be comparable in size to the largest particles in Saturn's main rings, which form the so-called ``giant-propellers\" in the A ring \\citep{Tiscareno09}. Aegaeon occupies a corotation eccentricity resonance with Mimas, like Anthe and Methone, and all three of these moons are associated with resonantly-confined arcs of debris. However, Aegaeon also appears to be a special case in terms of its orbital properties and its relationship with its arc. Its eccentricity and inclination are both extremely low, and the large number of resonant arguments on the boundary between circulation and libration lead to some interesting dynamical behavior. At the same time, the mass in the G-ring arc is probably a significant fraction of (and may even be comparable to) Aegaeon's mass, unlike the other arcs associated with small moons, opening up the possibility that interactions between the moon and the material in the arc could be responsible for some of Aegaeon's unusual orbital characteristics. Future analysis of this system could therefore provide insights into the orbital evolution of satellites coupled to disks of debris."
    },
    "0911/0911.1531_arXiv.txt": {
        "abstract": "Variations of frequency separations of low-degree p-modes are studied over the solar activity cycle. The separations studied are obtained from the frequencies of low-degree p-modes of the Global Oscillation Network Group (GONG). 10.7 cm radio flux is used as an index of solar activity. Small separations of the p-mode frequencies are considered to be mainly dependent on the conditions in stellar interiors. Thus they could be applied to diagnose the changes in the stellar interior. Our calculation results show that the magnitudes of variations of the mean large separations are less than 1 $\\sigma$ over the solar activity cycle. Small separations show different behaviors in the ascending and descending phases of activity. In the ascending phase, variations of the small separations are less than 1 $\\sigma$. However, the small separations have systematic shifts during 2004 - 2007. The shifts are roughly 1 $\\sigma$ or more. The variations of the ratios of the small to large separations with time are similar to the changes of the small separations. The effects of the changes in the large separations on the ratios are negligible. The variations of the separations may be a consequence of the influence from the surface activity or systematic errors in measurements or some processes taking place in the solar interior. ",
        "introduction": "It is well-known that the solar p-mode frequencies vary with the solar cycle \\citep{wood85, elsw90, libb90, elsw94, jime98, howe99, chap07}. It has been shown that the variations of the frequencies are significantly related to magnetic activities near the solar surface \\citep{wood91, jime04, tou05}. However, a number of attempts have been made to detect the relation between the frequency shifts and the variation of the solar interior. It has been proposed that the frequency shifts may relate to changes taking place in the base of the convection zone \\citep{jime98, sere05}, or in the deep core \\citep{jime98}.  Low-degree p-modes can be used to study the solar core because they can penetrate the deep interior \\citep{dzie97}. However, even these modes are primarily sensitive to the solar envelope structure. Thus frequency separations are proposed for a diagnostic tool. The frequency separations of low-degree p-modes in solar-like stars often proposed for diagnostic purposes are large separations $\\Delta_{l}(n)$ and small separations $d_{ll+2}(n)$. The large separations approximate to the characteristic frequency $\\nu_{0}$ of a star, defined by \\begin{equation} \\Delta _{l}(n)\\equiv \\nu_{n, l} - \\nu_{n-1, l} \\simeq \\nu_{0}, \\label{eqla} \\end{equation} where \\begin{equation} \\nu_{0} = (2\\int^{R}_{0}dr/c)^{-1}, \\end{equation} in which $c$ is sound speed at radius $r$ and $R$ is some fiducial radius of the star. The small separations are sensitive to the structure of stellar core \\citep{gou90a}, defined by \\begin{equation} \\begin{array}{lll} d_{l l+2}(n)& \\equiv & \\nu_{n, l} - \\nu_{n-1, l+2}\\\\ & =& \\frac{\\varphi_{l+2}-\\varphi_{l}}{\\pi}\\nu_{0}, \\end{array} \\label{eqsm} \\end{equation} where $\\varphi_{l}$ is the internal phase shift. For the low-degree p-modes, the $\\varphi_{l}$ is mainly dependent on conditions in the stellar core and is almost independent of those outside the core \\citep{rox03}. These frequency separations have been extensively investigated by many authors \\citep{chris84, ulr86, ulr88, gou87, gou90b, gou03, gou90a, rox03, aud94, rox05, flo05}. Additionally \\citet{rox93} defines differences $d_{01}(n)$ as \\begin{equation} \\begin{array}{lll} d_{01}(n)& \\equiv & \\frac{-\\nu_{n, 1} + 2\\nu_{n, 0}- \\nu_{n-1, 1}}{2} \\\\ & = & \\frac{\\varphi_{1}-\\varphi_{0}}{\\pi}\\nu_{0}, \\end{array} \\end{equation} and \\citet{yang07} define differences $\\sigma_{l-1l+1}(n)$ of low-degree p-modes as \\begin{equation} \\begin{array}{lll} \\sigma_{l-1 l+1}(n)& \\equiv & -\\nu_{n, l-1}+2\\nu_{n, l} - \\nu_{n, l+1} \\\\ & = & \\frac{\\varphi_{l+1}+ \\varphi_{l-1}-2\\varphi_{l}}{\\pi}\\nu_{0}. \\end{array} \\end{equation} The differences $\\sigma_{02}(n)$ are similar to the scaled small separations $d_{l l+2}(n)/(2l+3)$ in some cases and are mainly sensitive to conditions in stellar core. The difference $\\sigma_{02}$ averaged over $n$ is, however, more sensitive to changes in the central hydrogen abundance of solar-like stars than the scaled small separations \\citep{yang07}.  Moreover, the ratios $d_{l l+2}/\\Delta_{l}$ are considered to be essentially independent of the structure of the outer layers of a star and determined by interior structures \\citep{ulr86, rox03, rox05, flo05}. The ratios $d_{01}/\\Delta_{0}$ and $\\sigma_{02}/\\Delta_{0}$ also mainly rely on the interior structures \\citep{rox03, yang07}. Thus changes in stellar interior could be indicated by the variations in these separation ratios.  However frequency variations with solar activity are relative to the angular ($l$) degree, the azimuthal ($m$) degree, and the order $n$. The $m$-dependence is attributed to the solar near-surface activity \\citep{jime04, tou05} and could be removed by using the frequency centroids which incorporate all the $m$ components of a mode \\citep{chap05}. For the low-degree p-modes, centroid shifts can be regarded as being $l$-independent \\citep{chap05}. Thus separations of frequency centroids of low-degree p-modes should be insensitive to effects of surface activity. Using the p-modes of the Birmingham Solar-Oscillations Network (BiSON) collecting data using a Sun-as-a-star technique, \\citet{chap05} studied the impact of the solar activity cycle on the frequency separation ratios and found that the ratios change with the shifting level of global solar activity. \\citet{chap05} pointed out that some $m$ components are suppressed in the Sun-as-a-star observations and hence the ratios are expected to be affected by the surface activity. In this paper, we investigate the variations of frequency separations of low-degree p-modes as solar activity cycle proceeds using data collected by the resolved-Sun observations. In Section 2 we present our data and results. Then discussion and conclusion are represented in Section 3. ",
        "conclusions": "The characteristic frequency $\\nu_{0}$, i.e. the large separation, is an integral property of the entire star. Because the solar sound speed is a decreasing function of the solar radius, the contribution to $\\nu_{0}$ from the outer layers of the Sun is larger than that from the solar inner regions \\citep{gou90a}. Thus the large separation should be more sensitive to the changes in the outer layers than those in the inner regions and could be affected by the near surface activity. Moreover small separations are more sensitive to the conditions in stellar interior than those in outer layers. Thus the large separations should be more sensitive to the activity than the small separations. However, in fact, the magnitudes of variations of the large separations are about 0.07 $\\mu$Hz during the activity, which is less than 1 $\\sigma$. The magnitude of variations of $<d_{02}>$ is also less than 1 $\\sigma$ over the activity cycle. However, other small separations show a systematic shift during $\\sim$ 2004 - 2007, and the magnitudes of their variations are roughly 1 $\\sigma$ or more.  Small separations $<d_{ll+2}>$, $<\\sigma_{02}>$ and $<d_{01}>$ show a sudden change between 2001 and 2002. The fractional changes shown in Figure \\ref{fig2} of ratios in two GMFs (one in 1996 and one in 2001) are consistent with the prediction of \\citet{chap05}. However, if we use the data of 2002 instead of the data of 2001, the results would be different. The small separations show different behaviors in the ascending and descending phases of the cycle. In the ascending phase, variations of the small separations are less than 1 $\\sigma$. In the descending phase, the small separations show a systematic shift during 2004 - 2007. And the magnitudes of the changes in $<d_{13}>$, $<\\sigma_{02}>$ and $<d_{01}>$ are roughly 1 $\\sigma$ or more. It is not clear whether the variations are caused by the systematic errors in measurements.  The large separations increase slightly with increasing of the solar activity. But the magnitudes of temporal variations of the mean large separations of the low-degree p-modes are only about 0.07 $\\mu$Hz over the activity cycle, which is less than 1 $\\sigma$. Small separations $<d_{ll+2}>$, $<\\sigma_{02}>$ and $<d_{01}>$ show different behaviors in the ascending and descending phases of the cycle. In the ascending phase, variations of the small separations are less than 1 $\\sigma$. However, the small separations show a systematic shift during 2004 - 2007. The magnitudes of the changes in $<d_{01}>$, $<\\sigma_{02}>$ and $<d_{13}>$ in this period of time are about 1 $\\sigma$ or more. The variations of the ratios of the small to large separations with time are similar to the changes of the small separations. Changes in the large separations have a negligible impact on the variations in the ratios. The variations of the separations may be a consequence of the influence from the surface activity or systematic errors in measurements or some processes taking place in the solar interior. For further work, the data over cycle 22/23 would be required."
    },
    "0911/0911.4252_arXiv.txt": {
        "abstract": "We report the observations of PG\\,1553+113 during the first $\\sim$200\\,days of \\FermiGST science operations, from 4 August 2008 to 22 February 2009 (MJD 54682.7-54884.2). This is the first detailed study of PG\\,1553+113 in the GeV gamma-ray regime and it allows us to fill a gap of three decades in energy in its spectral energy distribution. We find PG\\,1553+113 to be a steady source with a hard spectrum that is best fit by a simple power-law in the \\Fermi energy band. We combine the \\Fermi data with archival radio, optical, X-ray and very high energy (VHE) gamma-ray data to model its broadband spectral energy distribution and find that a simple, one-zone synchrotron self-Compton model provides a reasonable fit. PG\\,1553+113 has the softest VHE spectrum of all sources detected in that regime and, out of those with significant detections across the \\Fermi energy bandpass so far, the hardest spectrum in that energy regime. Thus, it has the largest spectral break of any gamma-ray source studied to date, which could be due to the absorption of the intrinsic gamma-ray spectrum by the extragalactic background light (EBL). Assuming this to be the case, we selected a model with a low level of EBL and used it to absorb the power-law spectrum from PG\\,1553+113 measured with \\Fermi (200\\,MeV\\,-\\,157\\,GeV) to find the redshift which gave the best fit to the measured VHE data (90\\,GeV\\,-\\, 1.1\\,TeV) for this parameterisation of the EBL. We show that this redshift can be considered an upper limit on the distance to PG\\,1553+113. ",
        "introduction": "\\label{SEC:INTRO}  PG\\,1553+113 is a high-frequency peaked BL Lacertae object (HBL; \\citealp{Falomo:1994:SpecBlazars}; \\citealp{Beckmann:2002:BeppoSAXBLLacs}) whose redshift remains unknown despite continued efforts. Like all BL Lacs, we find its spectral energy distribution (SED) to have a double-peaked shape (in ${\\nu}F{\\nu}$ representation) that can thus be parameterised with four characteristic slopes. It has been detected from radio through hard X-rays and also in the very high energy (VHE; $E$\\,$\\gtrsim$\\,100\\,GeV) regime up to energies above 1\\,TeV (\\citealp{HESS:2006:PG1553Detection}; \\citealp{MAGIC:2007:PG1553Detection}). With these data three of the four components of its SED were sampled, namely, the rising and falling edges of the low-energy peak ($\\sim$10$^{-6}$\\,eV - 30\\,keV) and the falling edge of the high-energy peak ($\\sim$90\\,GeV - 1\\,TeV). We report here the first detailed analysis of the rising high-energy portion and, crucially, the high-energy peak, of the PG\\,1553+113 SED. These data, from observations made by the \\FermiGST Large Area Telescope (\\FermiLATc; \\citealp{Atwood:2009:LAT}) during its first $\\sim$200 days of operation, are combined with data from other wavebands to construct and model the broadband SED of PG\\,1553+113 in Section~\\ref{SEC:MODELING}.  Discovered and classified as a BL Lacertae object (BL Lac) by \\citet{Green:1986:PGCatalog}, PG\\,1553+113 is located at R.A. of $\\alpha_{J2000}$ = 15h55m43.04s and declination of $\\delta_{J2000}$ = +11d11m24.4s in the constellation of Serpens Caput. The logarithmic ratio of its 5\\,GHz radio flux, $F_{5GHz}$, to its 2\\,keV X-ray flux, $F_{2keV}$, has been found to range from log($F_{2keV}$/$F_{5GHz}$) = $-4.99$ to $-3.88$ (\\citealp{Osterman:2006:PG1553mwl}; \\citealp{Rector:2003:RadioHBL}). The high value of this ratio places PG\\,1553+113 at times among the most extreme of the HBLs; a BL Lac is classified as extreme when it has log($F_{2keV}$/$F_{5GHz}$) $\\geq$ -4.5 \\citep{Rector:2003:RadioHBL}. A number of the TeV blazars have, at one time or another, exhibited fluxes that place them in this extreme category (e.g., 1ES\\,0229+200, H\\,1426+428, 1ES\\,1959+650; \\citealp{Rector:2003:RadioHBL}).  In the radio band, PG\\,1553+113 has been detected at different mean flux levels. Its 4.8\\,GHz flux, for example, has ranged from 180 to 675\\,mJy (\\citealp{Bennett:1986:PG1553SEDned}; \\citealp{Gregory:1991:87GB}; \\citealp{Becker:1991:GHzSources}; \\citealp{Perlman:2005:XraySpectra}; \\citealp{Osterman:2006:PG1553mwl}). Its flux between 4.8 and 14.5\\,GHz was found to be variable on month timescales during the observations of \\citet{Perlman:2005:XraySpectra} and \\citet{Osterman:2006:PG1553mwl}. VLBA observations have resolved a jet extending at least 20\\,pc to the northeast of PG\\,1553+113 \\citep{Rector:2003:RadioHBL}. No evidence for superluminal motion has been reported in the literature to date; multi-epoch VLBA monitoring has commenced recently\\footnote{http://web.whittier.edu/gpiner/research/index.htm}.  PG\\,1553+113 is a bright optical source with $V$-band magnitude of $V_o\\,\\sim\\,14$ (\\citealp{Falomo:1990:PG1553}; \\citealp{Osterman:2006:PG1553mwl}). Observations taken between 1986 and 1991 with the ESO telescopes found its spectral index, $\\alpha_o$, to remain almost constant ($\\alpha_o\\,\\sim\\,-1$) and its magnitude to vary by ${\\Delta}V_o\\,=\\,1.4$ \\citep{Falomo:1994:SpecBlazars}. Low levels of optical variability were seen by \\citet{Osterman:2006:PG1553mwl} during their 2003 observing campaign. Limits on the magnitude of its host galaxy will be discussed later.  PG\\,1553+113 is also a bright X-ray source that has been observed by most X-ray missions (\\Einsteinc, \\ROSATc, \\ASCAc, \\BeppoSAXc, \\RXTEc, \\XMMNewtonc, \\Swift and \\Suzakuc). Although it has been detected at a number of different flux levels by these observatories, no evidence for strong or fast (sub-hour) flux variability has been observed at X-ray energies \\citep{Reimer:2008:PG1553Suzaku}. The \\Suzaku observations performed in July 2006 \\citep{Reimer:2008:PG1553Suzaku} provide the highest energy X-ray measurement so far obtained for PG\\,1553+113 at $\\sim 30$\\,keV with a 10-30\\,keV flux of $1.35$ x $10^{-11}$\\,erg\\,cm$^{-2}$\\,s$^{-1}$. No evidence for spectral hardening up to these energies was found in these data, indicating that all of the X-ray emission detected was due to synchrotron emission. The 2-10\\,keV fluxes measured by different X-ray missions are shown in Table~\\ref{TAB:X-RAY}, the highest being $6.9$ $\\times$ $10^{-11}$\\,erg\\,cm$^{-2}$\\,s$^{-1}$ in October 2006 (\\SwiftXRTc; \\citealp{Tramacere:2007:SwiftTeV}). The X-ray flux was found to double over a period of 10 days during the 3-week RXTE observing campaign of \\citet{Osterman:2006:PG1553mwl}. Despite the variations of a factor of approximately 20 in the 2-10\\,keV X-ray flux, the measured spectral properties of PG\\,1553+113 at these energies, also listed in Table~\\ref{TAB:X-RAY}, have not changed significantly over the course of the X-ray observations. It exhibits spectral curvature that can be well-described with either a broken power law or a log-parabolic shape. In the \\Suzaku data this curvature extends into the hard X-ray band ($\\leq 30$\\,keV) and steepening of the spectrum above $\\sim10$\\,keV beyond that predicted by either the broken power-law or log-parabolic model is observed \\citep{Reimer:2008:PG1553Suzaku}.  In the high energy (HE; $E$ $\\sim$ 100\\,MeV - 100\\,GeV) gamma-ray regime, PG\\,1553+113 was not detected by EGRET. An upper limit of $F{_{EGRET}}$ $<$ 9.97 $\\times$ 10$^{-8}$\\,cm$^{-2}$\\,s$^{-1}$ above 100\\,MeV was derived based on the summed exposure of cycles 1, 2, 3 and 4 \\citep{Hartman:1999:EGRET3EG}. At higher energies, in the VHE regime, PG\\,1553+113 is a confirmed gamma-ray emitter. First detected by H.E.S.S. \\citep{HESS:2006:PG1553Detection} and subsequently confirmed by MAGIC \\citep{MAGIC:2007:PG1553Detection}, PG\\,1553+113 has a flux that is approximately 3\\% that of the Crab Nebula at these energies. The combined 2005-2006 H.E.S.S. data follow a power law with photon index of $\\Gamma_{VHE}$=4.46\\,$\\pm$\\,0.34 between $\\sim$225\\,GeV and 1.3\\,TeV \\citep{HESS:2008:PG1553VLT}. The MAGIC 2005-2006 spectra are consistent with this having a power-law photon index of $\\Gamma_{VHE}$=4.21\\,$\\pm$\\,0.25 between $\\sim$90 and 500\\,GeV \\citep{MAGIC:2007:PG1553Detection}. No evidence for variability of the spectral index, to within the measurement uncertainties, is seen in the VHE measurements. The H.E.S.S. and MAGIC integral flux levels are consistent from April to August 2005 and from April to July 2006 with a mean flux of $I_{VHE}$ (E$>200$\\,GeV) = 2.0\\,$\\pm$\\,0.8 $\\times$ 10$^{-11}$\\,cm$^{-2}$\\,s$^{-1}$. Although the spectrum remained unchanged, the flux detected by MAGIC between January and April 2006 was lower than the preceding and proceeding measurements at $I_{VHE}$ ($E\\,>200$\\,GeV) = 0.6\\,$\\pm$\\,0.2 $\\times$ 10$^{-11}$\\,cm$^{-2}$\\,s$^{-1}$. When the systematic uncertainties on the fluxes are taken into account, this flux is marginally inconsistent with the VHE fluxes detected during the H.E.S.S. and other MAGIC observations suggesting that the PG\\,1553+113 flux varied by up to a factor of three on monthly timescales in 2006. PG\\,1553+113 has the steepest spectrum of all of the sources detected in the VHE regime, which makes it a promising target for the \\FermiLAT because extrapolating down to the \\Fermi energy range would make this an extremely bright source unless a dramatic spectral break occurs at energies below $\\sim$100\\,GeV. EGRET's non-detection could be interpreted in this context or as the result of PG\\,1553+113 being in a lower emission state at that time than during the VHE observations. PG\\,1553+113 is in the \\FermiLAT bright AGN source list (LBAS; \\citealp{Fermi:2009:LBAS}), the high-confidence AGN associations from the first three months of \\Fermi data. Listed as 0FGL\\,J1555.8+1110 with a flux of $I_{LBAS}$ ($E$\\,$>$\\,100\\,MeV)\\,=\\,8.0\\,$\\pm$\\,1.0 $\\times$ 10$^{-8}$\\,cm$^{-2}$\\,s$^{-1}$ and a photon index of $\\Gamma_{LBAS}$\\,=\\,1.70\\,$\\pm$\\,0.06, it has one of the hardest spectra of the 106 AGN that comprise this list. Indeed, if only the AGN with a significant detection across the entire \\Fermi bandpass are considered, its spectrum is the hardest of those in the LBAS. This combination of a very soft VHE spectrum and a very hard HE spectrum means that PG\\,1553+113 has a significant spectral break in the gamma-ray regime.  Despite significant efforts, the redshift of PG\\,1553+113 remains unknown. Its measurement is of great interest both for a better understanding of its SED, in particular since it is an extreme BL Lac and therefore has a very hard synchrotron emission spectrum, and also for studying EBL effects; some redshift estimates make it the most distant VHE source detected to date at $z<$0.78 \\citep{Sbarufatti:2005:Redshifts}, which, if shown to be the case, could imply an EBL density close to the minimum allowed by galaxy counts. The results of redshift studies undertaken to date are summarized in Section~\\ref{SEC:Z} and the archival gamma-ray data together with the \\Fermi data from PG\\,1553+113 are used to place constraints on its redshift. ",
        "conclusions": "\\label{SEC:CONCLUSION}  We have combined the \\Fermi data from PG\\,1553+113 with data from radio through VHE gamma rays to study its broadband emission. We demonstrated that a simple, one-zone SSC model provides a reasonably good fit to the observational results: It accounts for the different X-ray flux states observed while also allowing the gamma-ray data to remain approximately constant. More detailed theoretical modeling of the SED, which naturally might fit the data points better, is beyond the scope of this paper.  We have used gamma-ray data spanning four orders of magnitude to seek the EBL column density (as parameterized by \\citet{Franceschini:2008:EBL}) that best fits the measured spectrum of PG\\,1553+113. We find that an EBL integrated to a redshift of $z=0.75^{+0.04}_{-0.05}$ provides the best fit. Such a high value for the redshift would make PG\\,1553+113 the most distant source to be detected in the VHE regime, an attribute consistent with it being the VHE source with the largest spectral break in the gamma-ray regime ($\\Gamma_{VHE}$ - $\\Gamma_{HE}$ $\\sim$ 2.7). Three assumptions were made in our redshift calculation the first one being that the EBL behaves in a manner consistent with the model of \\citet{Franceschini:2008:EBL}. Of the many EBL models available in the literature (see \\citet{Finke:2009:EBL} for a recent comparison), that of \\citet{Franceschini:2008:EBL} predicts the lowest level of EBL across the energy range of interest here; it provides the minimum level of EBL photons that could exist based on known sources alone and therefore, in using it, the redshift derived for PG\\,1553+113 can be considered an upper limit. Had one of the other EBL models been used, the best-fit redshift would have been lower. Indeed, a more in-depth study of the absorption effects of the EBL on the spectrum of PG\\,1553+133 should consider all of the EBL models available in the literature. In this way, the systematic effects of using different realisations of the EBL could be determined; the data points provided in Table~\\ref{TAB:FermiSpectrum} allow for such a study to be undertaken.  Two additional assumptions were made in our redshift calculation, firstly, that the emission state of PG\\,1553+113 did not change significantly during the gamma-ray observations, which can be interpreted as an indication that it was in a low-flux state state at these energies throughout those observations. Indeed, the quiescent state is the state in which we are most likely to find a blazar \\citep{Vercellone:2004:DutyCycle}. Secondly, it was assumed that the power-law spectrum measured by \\Fermi does not suffer from significant intrinsic absorption up to 1\\,TeV. Should there be intrinsic absorption of VHE gamma rays at the source, a lower level of EBL absorption would be required to best fit the measured VHE spectrum. We note that, since the \\Fermi data show no evidence for spectral curvature, if it is significant intrinsic absorption that is mostly responsible for the sharp break in the gamma-ray spectrum, its effects would have to kick in at energies exactly above those accessible to \\Fermi and, at a significant level, in order to produce the sharp break observed in the data.  This is the first time that the measurement of the complete HE to VHE gamma-ray spectrum of a source has been used to constrain its redshift.  The value derived here is close to the limits derived by \\citet{HESS:2008:PG1553VLT} and \\citet{MAGIC:2007:PG1553Detection} and is higher than that derived by \\citet{Mazin:2007:PG1553}. Different parameterizations of the EBL were used in these analyses (those of \\citet{Kneiske:2004:EBL} and \\citet{Primack:2001:EBL}) so this is one of the factors contributing to the different results. Additionally, for these previous redshift constraints using gamma-ray data, it was necessary to make assumptions about the hardness of the intrinsic spectrum in the HE regime. The \\Fermi measurement of the HE spectrum from PG\\,1553+113 allows us to reduce the number of assumptions used in the constraining the redshift. Combining the \\Fermi measurement with the measured VHE spectra affords us complete coverage of the PG\\,1553+113 SED from 200\\,MeV to 1.1\\,TeV.  Its potentially large distance together with its consistently detectable gamma-ray flux make PG\\,1553+113 an excellent candidate with which to search for EBL cascading effects first proposed by \\citet{Protheroe:1986:EBLCascades} and, subsequently explored by many authors (e.g., \\citealp{Protheroe:1993:EBLCascades}; \\citealp{Aharonian:1994:EBLCascades}; \\citealp{Biller:1995:EBLCascades}; \\citealp{DAvezac:2007:EBLCascades}; \\citealp{Elyiv:2009:EBLCascades}). If present, pile-up from such cascades is predicted to occur below 100\\,GeV. Some authors predict that cascades produce a characteristic spectral bump, which could be detectable in the \\Fermi energy regime if the strength of the extragalactic magnetic field in the direction of the source being studied is less than $B$\\,=\\,10$^{-6}$\\,nG \\citep{DAvezac:2007:EBLCascades}. \\citet{Elyiv:2009:EBLCascades} predict that, for magnetic field strengths of $B$\\,$\\leq$\\,10$^{-7}$\\,nG, extended emission due to the cascading of the source photons should be detectable around extragalactic gamma-ray sources by \\Fermic.  The possibility that PG\\,1553+113 could be a particularly distant TeV source provides further impetus for IR/optical/UV astronomers to revisit the redshift measurement. A direct measurement would be very welcome, and would ultimately settle the question of the accuracy of estimates based on the EBL. The \\FermiLAT is continuing to accumulate data on PG\\,1553+113. Subsequent studies of these data will, therefore, allow us to measure its spectral shape with greater sensitivity at the highest energies accessible to \\Fermic.   \\begin{figure} \\resizebox*{0.9\\textwidth}{!}{\\includegraphics[draft=false]{fig1.eps}}\\\\% \\caption{\\label{FIG:FermiSpectrum}The differential spectrum from PG\\,1553+113 as measured by \\Fermi. The solid line shows the fit of a power-law to the overall spectrum derived from all of the data with energy $E>$ 200\\,MeV. The data-points (crosses) indicate the flux measured in each of the eight energy bins indicated by the extent of their horizontal lines, when the data in these energy ranges were analyzed with the photon index fixed at the value derived from the entire dataset. The grey shaded area shows the extent of the \\Fermi 68\\% confidence band.} \\end{figure}  \\begin{figure} \\resizebox*{0.9\\textwidth}{!}{\\includegraphics[draft=false]{fig2.eps}}\\\\% \\caption{\\label{FIG:FermiLightcurve}The lightcurve for PG\\,1553+113 measured by \\Fermi between MJD 54683 and 54883 (2008-08-05 - 2009-02-21). The data are binned in 2-day bins. The top panel shows the integral flux, $I_{2Day}$ ($E\\,>$\\,200\\,MeV), for each 2-day bin. The bottom panel shows the power-law photon index, $\\Gamma_{2Day}$, for each 2-day bin. The integral flux and photon index, with their uncertainties, calculated by \\texttt{gtlike} for the entire dataset are shown by the shaded horizontal bands. For twelve of the timebins, indicated by the triangles in the lower panel, the analysis failed to converge with reasonable error bars and therefore, these data were excluded from the variability analysis and the points are shown as ``$\\times$'' symbols.} \\end{figure}   \\begin{figure} \\resizebox*{0.9\\textwidth}{!}{\\includegraphics[draft=false]{fig3.eps}}\\\\% \\caption{\\label{FIG:SED-EBL}The spectral energy distribution for the gamma-ray data. The individual datasets are described in the text. The \\Fermi datapoints are shown as black crosses. The EGRET upper limit for emission above 100\\,MeV is shown as a black triangle. The H.E.S.S. data combined from May \\& August 2005 are shown as blue dots and those combined from April \\& July 2006 as blue open circles. The MAGIC data combined from April \\& May 2005 and from January to April 2006 are shown as cyan x's while those from July 2006 are shown as cyan open squares. The green solid line shows the power-law fit to the \\Fermi data, extended to higher energies with the level of EBL that best fitted the VHE data, which corresponds to the EBL integrated to a redshift of $z=0.75$. The green dotted line shows the extension of the \\Fermi power-law to higher energies without absorption. The upper and lower 68\\% uncertainty-contours for the \\Fermi data are shown as red dashed and blue dash-dotted lines, respectively. Each of them were also extended with the level of EBL that best fitted the VHE data, which corresponded to a redshift of $z=0.79$ for the upper contour and to a redshift of $z=0.70$ for the lower contour. Their unabsorbed extensions to higher energies are shown as red and blue dotted lines.} \\end{figure}    \\begin{figure} \\resizebox*{0.9\\textwidth}{!}{\\includegraphics[draft=false]{fig4.eps}}\\\\% \\caption{\\label{FIG:SSC}The spectral energy distribution for PG\\,1553+113 fit with an SSC model. The fitting procedure is described in the text. The yellow, green, blue and magenta lines are, respectively, the SSC model tuned to fit the \\Swift \\XRT data, the \\Suzaku data, the \\RXTE data and the \\Swift \\XRT data. The black dashed line is the SSC model when the highest energy electrons are omitted entirely from the fit. The parameters are described in the text and are summarised along with the dates of the observations in Table~\\ref{TAB:SED-PARAMS}. The X-ray data have been de-absorbed for a column density of $N_H$\\,=\\,3.67$\\times$10$^{-20}$\\,cm$^{-2}$. Apart from 7 hours of the H.E.S.S. data (taken in April 2006) the KVA (optical; \\citealp{Reimer:2008:PG1553Suzaku}), \\Suzaku \\citep{Reimer:2008:PG1553Suzaku}, MAGIC \\citep{MAGIC:2009:PG1553MWL} and H.E.S.S. \\citep{HESS:2008:PG1553VLT} observations were made in July 2006 and are thus quasi-simultaneous. They are shown as green filled circles. The \\SwiftXRT and \\UVOT data from October 2005 \\citep{Tramacere:2007:SwiftTeV} are shown as yellow triangles while the magenta open squares are \\SwiftXRT data (Swift observation ID 31368001) from March 2009. The \\RXTE data from \\citet{Osterman:2006:PG1553mwl} are shown as blue filled squares. The archival data (grey filled circles) come from: \\citet{Bennett:1986:PG1553SEDned}, \\citet{Becker:1991:GHzSources}, \\citet{Gregory:1991:87GB}, \\citet{Douglas:1996:TexasRadio}, \\citet{Gorshkov:2003:Radio} (\\radioc); \\citet{Falomo:1990:PG1553}, \\citet{Urry:2000:PG1553SEDned}, \\citet{Sbarufatti:2006:ESO-12Redshifts}, the Sloan Digital Sky Survey \\citep{SDSS:2009:SDSS7} via NED, \\citet{Fox:2006:HighVelClouds}, \\citet{Tramacere:2007:SwiftTeV}, (\\opticalUVc); \\citet{Donato:2005:BeppoSAXBlazarCat}, (\\Xrayc); H.E.S.S. \\citep{HESS:2008:PG1553VLT} and \\citet{MAGIC:2007:PG1553Detection}, (\\VHEgammarayc).} \\end{figure}  \\begin{figure} \\resizebox*{0.9\\textwidth}{!}{\\includegraphics[draft=false]{fig5.eps}}\\\\% \\caption{\\label{FIG:ZOOM}A zoom on the high-energy portion of the spectral energy distribution for PG\\,1553+113. The \\Fermi datapoints are shown by black filled circles; the 2006 H.E.S.S. and MAGIC data are shown as green solid squares while the H.E.S.S. and MAGIC data from 2005 are shown as grey solid circles. The EGRET upper limit is shown as a black triangle. The best SSC model fit for the \\Suzaku X-ray dataset, with EBL absorption applied, is shown as a green dashed line. The green dash-dotted line shows the SSC model before absorption for the EBL. The black solid line shows the Fermi power-law spectrum absorbed for $z$\\,=\\,0.75 using the model of \\cite{Franceschini:2008:EBL} while the dotted black line shows the unabsorbed Fermi spectrum. The shaded area shows the \\Fermi butterfly and the grey dashed lines, its unabsorbed extension to higher energies. The grey solid lines show the \\Fermi butterfly absorbed for the best fit redshift to each edge, as discussed in Section~\\ref{SEC:Z}.} \\end{figure}"
    },
    "0911/0911.5581_arXiv.txt": {
        "abstract": "To improve the planet detection efficiency, current planetary microlensing experiments are focused on high-magnification events searching for planetary signals near the peak of lensing light curves. However, it is known that central perturbations can also be produced by binary companions and thus it is important to distinguish planetary signals from those induced by binary companions. In this paper, we analyze the light curves of microlensing events OGLE-2007-BLG-137/MOA-2007-BLG-091, OGLE-2007-BLG-355/MOA-2007-BLG-278, and MOA-2007-BLG-199/OGLE-2007-BLG-419, for all of which exhibit short-term perturbations near the peaks of the light curves.  From detailed modeling of the light curves, we find that the perturbations of the events are caused by binary companions rather than planets.  From close examination of the light curves combined with the underlying physical geometry of the lens system obtained from modeling, we find that the short time-scale caustic-crossing feature occurring at a low or a moderate base magnification with an additional secondary perturbation is a typical feature of binary-lens events and thus can be used for the discrimination between the binary and planetary interpretations. ",
        "introduction": "The microlensing signal of a planet is characterized by a short-term perturbation to the smooth standard single-lens light curve of the primary-induced lensing event occurring on a background source star \\citep{mao91,gould92}.  The duration of the perturbation is several days for a gas giant and several hours for an Earth-mass planet, while the typical cadence of lensing surveys is roughly a day. As a result, it is difficult to detect planets from the survey observations alone.  To achieve the cadence required to detect short-term planetary signals, current planetary lensing searches are being conducted by combining survey and follow-up observations, in which the survey observations \\citep{udalski03,bond02} aim to maximize the number of detections of lensing events by monitoring a large area of sky and follow-up observations \\citep{albrow98,ghosh04} are focused on intensive monitoring of the events detected by the survey observations. However, the number of telescopes available for follow-up observations is far smaller than would be needed to monitor all the detected events from the survey observations.  As a result, follow-up observations are selectively conducted for events that can maximize the planet detection probability.  Currently, the highest priority is given to high-magnification events.  For these events, the planet detection efficiency is high because the source trajectory always passes close to the region of central perturbations.  In addition, the perturbation occurs near the peak of the light curve and thus follow-up observations can be prepared in advance.   Although the strategy to detect central perturbations has an important advantage of high sensitivity to planets, it also has a shortcoming of difficulty in the interpretation of observed signals.  One important cause of this difficulty is that the perturbation near the peak of a lensing light curve can be produced not only by a planet but also by a wide or a close binary companion with a mass roughly equal to that of the primary. Fortunately, the perturbation patterns produced by the planetary and the binary companion are intrinsically different \\citep{han08, han09a}, and thus it is in principle possible to discriminate between the planetary and binary interpretations.  However, this discrimination usually requires detailed modeling of the light curve, which demands a time-consuming search for a solution of the lensing parameters in the vast space of many parameters.  Therefore, simple diagnostics that can be used for the discrimination between the two interpretations will be useful not only for the immediate identification of the nature of the perturbation but also for the preparation of observational strategies for better characterization of planetary systems.    In this paper, we present the results of the analysis of three {\\it binary-lens} events with very strong short-term signals near the peaks of the light curves.  From the investigation of the light curves combined with the underlying physical geometry of the lens system, we find and present several features in lensing light curves that can be used as diagnostic tools for the discrimination between the planetary and binary interpretations of observed events.   The paper is organized as follows.  In section 2, we briefly describe the general features of central perturbations produced by a planet and a binary companion.  In section 3, we discuss the data used in our analysis.  In section 4, we explain the modeling procedure and present the best-fit lensing parameters of the individual lensing events obtained from the modeling.  In section 5, we present the features in the light curves that characterize the binary nature of the events and discuss the usability of the features in distinguishing the planetary and binary interpretations. We summarize the results and conclude in section 6. ",
        "conclusions": "We analyzed the light curves of three microlensing events with short-term perturbations near the peaks of the light curves. From detailed modeling of the light curves, we found that the perturbations of the events are caused by binary companions rather than planets. From close examination of the light curves combined with the underlying physical geometry of the lens system obtained from modeling, we found that the short time-scale caustic-crossing feature occurring at a low or a moderate base magnification with an additional secondary perturbation is a typical feature of binary-lens events and thus can be used for the discrimination between the binary and planetary interpretations."
    },
    "0911/0911.0647_arXiv.txt": {
        "abstract": "{} {Given that in most cases just thermal pressure is taken into account in the hydrostatic equilibrium equation to estimate galaxy cluster mass, the main purpose of this paper is to consider the contribution of all three non-thermal components to total mass measurements. The non-thermal pressure is composed by cosmic rays, turbulence and magnetic pressures.} {To estimate the thermal pressure we used public XMM-\\textit{Newton} archival data of five Abell clusters to derive temperature and density profiles. To describe the magnetic pressure, we assume a radial distribution for the magnetic field, $B(r) \\propto \\rho_{g}^{\\alpha}$, to seek generality we assume $\\alpha$ within the range of 0.5 to 0.9, as indicated by observations and numerical simulations. Turbulent motions and bulk velocities add a turbulent pressure, which is considered using an estimate from numerical simulations. For this component, we assume an isotropic pressure, $P_{\\rm turb} = \\frac{1}{3}\\rho_{\\rm g}(\\sigma_{r}^{2}+\\sigma_{t}^{2})$. We also consider the contribution of cosmic ray pressure, $P_{cr}\\propto r^{-0.5}$. Thus, besides the gas (thermal) pressure, we include these three non-thermal components in the magnetohydrostatic equilibrium equation and compare the total mass estimates with the values obtained without them.} {A consistent description for the non-thermal component could yield a variation in mass estimates that extends from 10\\% to $\\sim$30\\%. We verified that in the inner parts of cool core clusters the cosmic ray component is comparable to the magnetic pressure, while in non-cool core clusters the cosmic ray component is dominant. For cool core clusters the magnetic pressure is the dominant component, contributing more than 50\\% of the total mass variation due to non-thermal pressure components. However, for non-cool core clusters, the major influence comes from the cosmic ray pressure that accounts for more than 80\\% of the total mass variation due to non-thermal pressure effects. For our sample, the maximum influence of the turbulent component to the total mass variation can be almost 20\\%. Although all of the assumptions agree with previous works, it is important to notice that our results rely on the specific parametrization adopted in this work. We show that this analysis can be regarded as a starting point for a more detailed and refined exploration of the influence of non-thermal pressure in the intra-cluster medium (ICM).} {} ",
        "introduction": "\\label{intro} Clusters of galaxies are powerful tools for investigations of cosmological interests. The evolution of the mass function of a cluster is highly sensitive to cosmological models since the matter density controls the rate at which structures grow \\citep{Voit05}. In order to use clusters of galaxies as observational probes of dark energy in the Universe and to investigate the structure formation history including baryonic hydrodynamics, the non-thermal contribution must be well understood and quantified.  X-ray data are one of the most often methods used to determine the mass distribution of clusters of galaxies. To do so, hydrostatic equilibrium is usually  assumed, and the observed gas density and temperature profiles are used to compute the thermal pressure. In most  cases, only the gas (thermal) pressure is considered to evaluate the dynamical masses of galaxy clusters \\citep[e.g.,][]{David95,WF95,Finoguenov01,Reip02}. However, there is also a non-thermal pressure ($P_{\\rm NT}$), composed by the magnetic ($P_{\\rm B}$), turbulent ($P_{\\rm turb}$) and  cosmic ray ($P_{\\rm cr}$) components that is frequently assumed to be negligible and thus ignored. As a consequence of this, today the accuracy of the hydrostatic mass estimates is limited by the non-thermal pressure from these components.  Despite the difficulty to reliably calculate the small-scale properties of the magnetic field, the existence of intra-cluster magnetic fields is well established from the studies of the rotation measure of polarized radio frequencies and synchrotron emission from diffuse sources \\citep[e.g.,][]{Andernach88,giov93,Taylor94,Taylor02,Govoni04}. More recently, another indication that the intra-cluster medium (ICM) is permeated by a magnetic field came from the studies of X-ray cold fronts \\citep[sharp discontinuities in X-ray surface brightness profile and temperature,][]{Mark07}. In these cases, a parallel magnetic field can suppress transport processes in the ICM, making it difficult to mix different gas phases during a cluster merger. Even a very weak magnetic field can effectively inhibit transport processes such as thermal conduction and the settling of heavy ions \\citep{Sarazin86,SS90}.  Strong magnetic fields can make a significant contribution to the gas pressure support \\citep{loeb94}, contributing with a non-thermal component in the magnetohydrostatic equilibrium equation \\citep{D01b,Vogt05}. Indeed, magnetic fields as high as 10 - 100 $\\mu$G were found in Hydra A \\citep{Taylor93}, Cygnus A  \\citep{Dreher87} and in 3C 295 \\citep{Perley91}. Moreover, \\citet{Dolag99} performed numerical simulations and found that even clusters with an overall small magnetic field can be penetrated partially by regions of high magnetic fields.  Although on average the magnetic pressure in simulations is much smaller than the thermal pressure \\citep[$\\sim$ 5\\%,][]{DS00}, there are domains of high magnetic fields approaching or sometimes even exceeding equipartition with the thermal energy. Previous studies \\citep[][among others]{Dolag01a,Colafrancesco07} have analyzed the effects of the magnetic pressure in simulated galaxy clusters.  Magnetic fields and turbulence are possibly related to one another. It seems plausible that the turbulent motions in the ICM can maintain the magnetic field by converting kinetic energy into magnetic energy \\citep{Sanchez99}. The observed small-scale turbulence in the ICM can be due to bulk velocities and ongoing merger of substructures. Gas turbulence on small scales can also be driven directly by motions of galaxies, as for instance by jets and bubbles from the active galactic nuclei \\citep[AGN,][]{Churazov02}, although the latter may be confined to the inner regions of the cluster \\citep{Lau09}. The presence of random gas motion can also contribute to the pressure support in clusters of galaxies.  The chaotic nature of the ICM magnetic field makes it difficult for energetic particles to scape from the cluster, and thus cosmic-ray protons would be confined for timescales exceeding the Hubble time. The electron cosmic rays, on the other hand, have collisional and radiative lives much shorter. Thus, since the ICM is permeated by significant magnetic fields, one would expect the cosmic ray pressure to have some relevance in its support against gravity.  To consider deviations from the standard assumptions in computing cluster total mass, the main aim of this work is to analyze the effects of non-thermal pressure, that is to take into account magnetic, turbulent and cosmic ray components. Hydrostatic masses were derived using X-ray observational data for five Abell clusters: A496, A2050, A1689, A2667 and A2631. To do so, we use temperature and density profile fits from a previous work \\citep{Lagana08} and we introduce the $P_{\\rm NT}$ contribution in the magnetohydrostatic equilibrium equation. For these five clusters, we compare masses determined considering non-thermal pressure ($M_{\\rm NTP} (r)$) with their hydrostatic values ($M (r)$).  The  paper is organized as follows. We show the data sample in Sect.~\\ref{data}. The non-thermal components are described in Sect.~\\ref{NTC}. In this section we describe the structure of the intra-cluster magnetic field, the turbulence in the ICM and the cosmic ray component. In Sect.~\\ref{PNT}, we present the method of determining  the cluster mass, including the effects of the $P_{NT}$. Our results, as well as a discussion of them are presented in Sect.~\\ref{RD} and our conclusions in Sect.~\\ref{conc}. ",
        "conclusions": "\\label{conc} We have taken into account the effects of non-thermal pressure on the X-ray mass estimates for five Abell clusters (A496, A2050, A1689, A2667 and A2631). The masses derived considering just the thermal pressure were presented in a previous work by \\citet{Lagana08} and were used here for comparison. We summarize our main results below:  \\begin{itemize}  \\item{The inclusion of non-thermal pressure in the intra-cluster gas description is motivated by the increasing evidence for the presence of a magnetic field in clusters of galaxies. We assume a magnetic profile given by $B(r) \\propto B_{0} \\rho_{g}^{\\alpha}$, considering values for $B_{0}$ ranging from 5 up to 30 $\\mu$G for CC clusters, while for NCC clusters we assume 2$\\mu$G $< \\rm B_{0}< $ 8$\\mu$G. For each central value we let the shape parameter vary between 0.5 $ < \\alpha <$ 0.9. The magnetic pressure contributes with approximately 20\\% to the total mass variation.}  \\item{In order to take into account the influence of turbulent motion in the ICM and bulk velocities, we assume isotropic turbulent pressure based on the numerical simulation results of \\citet{Lau09}: $P_{\\rm turb}=1/3 \\rho_{\\rm g} (\\sigma_{r}^{2}+\\sigma_{\\rm t}^{2})$. The tangential ($\\sigma_{r}$) and the radial ($\\sigma_{\\rm t}$) dispersion velocity profiles were derived based in the same numerical simulations. This component can influence up to 5\\% of the total mass estimates.}  \\item{Energetic particles are confined by magnetic fields. Since ICM is permeated by magnetic fields, cosmic rays can also provide an important source of pressure. As the distribution of cosmic rays is poorly known, we considered the prescription of \\citet{Ando08} to derive the cosmic ray pressure. From our results we can see that this component can affect the cluster mass estimates by $\\sim$ 10\\%. In the inner parts of cool core clusters this component is comparable to the magnetic pressure.}  \\item{The plasma in clusters of galaxies may be affected by non-thermal processes which rise the pressure and hence cause an underestimate in X-ray measurements of the total mass. In this way the comparison between independent methods of mass estimates can constrain the contribution of non-thermal pressure. We compared weak-lensing results with X-ray mass measurements for A1689 to investigate this, and if the difference in mass estimates is in fact due to the non-thermal component, it can account for 2\\% to $\\sim$30\\%.}  \\item{We took into account the effects of non-thermal pressure on the total mass estimates. To derive the thermal pressure and compute the hydrostatic mass we used observed temperature and density profiles \\citep{Lagana08}. This is the first study that considers the influence of all three non-thermal pressure on the total mass estimates. Within the limits of our sample and with several reasonable assumptions to describe each of the non-thermal components, it is important to notice that our findings rely on the specific parametrization adopted in this work. This study indicates that further investigations are needed to make a detailed description of the influence of non-thermal components in ICM.}  \\item{Without the complete knowledge of the non-thermal contribution, we should be aware of the fact that clusters of galaxies have been used as observational tools for investigations of cosmological interest based only on the hydrostatic assumption. The non-thermal components are neglected. A thorough knowledge of non-thermal contribution will come from the combination of X-ray, radio and gamma ray data for a large sample of clusters. The next generation of cluster surveys will provide data to address these fundamental questions in cosmology\\footnote{http://wfxt.pha.jhu.edu/}.}  \\end{itemize}"
    },
    "0911/0911.2368_arXiv.txt": {
        "abstract": "Vilkovisky has claimed to have solved the black hole backreaction problem and finds that black holes lose only ten percent of their mass to Hawking radiation before evaporation ceases.  We examine the implications of this scenario for cold dark matter, assuming that primordial black holes are created during the reheating period after inflation.  The mass spectrum is expected to be dominated by 10-gram black holes.  Nucleosynthesis constraints and the requirement that the earth presently exist do not come close to ruling out such black holes as dark matter candidates.  They also evade the demand that the photon density produced by evaporating primordial black holes does not exceed the present cosmic radiation background by a factor of about one thousand.  \\vspace*{5mm} \\noindent PACS: 4.70.-s, 4.70.Dy,95.35.+d,98.80.-k.,98.80.Ft  \\\\ Keywords: Cosmology, primordial black holes, dark matter, Hawking radiation, evaporation. ",
        "introduction": "\\setcounter{equation}{0}\\label{sec1} \\baselineskip 8 mm  In a recent series of  papers, Vilkovisky \\cite{Vilk06a,Vilk06b,Vilk06c,Vilk08} has claimed to have completely solved the black hole backreaction problem, which is the major unresolved issue regarding black hole physics since Hawking's discovery of black hole radiation thirty years ago \\cite{Hawk75}.  Hawking's original calculation was semiclassical and ignored the effect of the black hole's evaporation on the spacetime metric itself; Hawking argued that any such backreaction should be negligible until the black hole approaches the Planck mass.  Vilkovisky maintains that if backreaction is properly taken into account, it exerts a negative feedback on the evaporation process, which ceases after the hole has lost only ten percent of its mass to thermal radiation. If the scenario is correct, then contrary to the accepted picture, primordial black holes of mass less than $10^{15}$ g will not have entirely disappeared by the present epoch of the universe and could conceivably provide the cold dark matter sought by cosmologists.  Vilkovisky's calculations are also semiclassical and do not assume any exotic physics.  Nevertheless, they have yet to receive full scrutiny by the physics community and it is premature to claim they are correct.  In this paper we merely propose to take his scenario seriously and examine the consequences.  We find that, indeed, black holes with mass less than $10^{15}$ grams can provide the $\\O \\approx$ .3 dark matter while safely evading basic observational constraints.  Such a conclusion may seem fairly obvious, since the black holes are evaporating only ten percent of their mass.  On the other hand, their evaporation profile does not at all resemble the usual one and it is worth verifying that this does not seriously alter the expected conclusions.  Furthermore, other phenomena, such as vortons\\cite{CA00}, behave similarly at least insofar as that they are stabilized against contraction by the current they generate.   We do also feel that it is worthwhile to stimulate discussion of Vilkovisky's work, which if correct, may have provided the answer to two outstanding problems.\\\\  Carr et al. have provided several reviews of the standard picture of primordial black hole (pbh) formation processes and observational constraints on their mass density; see for example \\cite{Carr05,Carr94,Carr75} and references therein.  Although the density of black holes with $M \\gtrsim 10^{15}$ grams is constrained by a variety of techniques, including microlensing and $\\gamma$-ray bursts, all constraints on holes with $M \\lesssim 10^{15}$ grams are imposed by assuming that they interfere with established cosmological processes through their evaporation, which in the Hawking picture ends in an explosive phase.  In the Vilkovisky picture, by contrast, the temperature profile of the emitted radiation differs sharply from the usual one, being nearly flat until it drops to zero, and evaporation never reaches an explosive phase.  The Hawking temperature is the usual one, \\be T = M_{pl}^2\\frac{\\k(u)}{2\\p} \\ee where $M_{pl}$ is the Planck mass, $\\k$ is the surface gravity and $u \\equiv t - r$ is the retarded time.  (We work in units with $c = G = k_{Boltzmann} = 1$.) However, in Vilkovisky's scenario the surface gravity is \\emph{not} $\\k = \\frac1{4M}$.  His considerations of the radiation at infinity lead him to describe the metric in terms of two ``Bondi charges.\"  The first is ${\\cal M}$, the Bondi mass, which is the mass remaining in a compact region after the escape of radiation (the ADM mass minus the mass of the radiation).  The second is a ``charge\" $\\cal Q$, which the black hole manufactures from the vacuum to protect itself from quantum evaporation.  We emphasize, however, that these two ``charges\" are {\\it not} merely assumed to exist but are a necessary consequence of the radiation metric.  They are in fact, the two leading coefficients of the metric expansion at infinity and are the only two that need to be specified.  Because in terms of these two coefficients, the asymptotic metric assumes the Reissner-Nordstr\\\"om form, the surface gravity becomes that for a Reissner-Nordstr\\\"om black hole: \\be \\k = \\frac{({\\cal M}^2-{\\cal Q}^2)^{1/2}}{[{\\cal M}+({\\cal M}^2-{\\cal Q}^2)^{1/2}]^2} \\ee  The evolution of $\\cal M$ and $\\cal Q$ are governed by two coupled differential equations: \\bea \\frac{d{\\cal M}}{du} &\\approx & - \\frac{M_{pl}^2}{48\\p}\\k^2\\\\ \\frac{d{\\cal Q}^2}{du} &=& \\frac{M_{pl}^2}{24\\p}\\k . \\eea  We plot $\\cal M,Q$ and $T$ in figure 1.  As one sees, $T$ is essentially flat but reaches a slight maximum of about \\be T_{max} = \\frac{M_{pl}^2}{8\\pi M}, \\label{TV} \\ee before dropping abruptly to zero, at which point evaporation stops.  Here $M$ is the initial mass of the black hole.  \\begin{figure}[htb] \\vbox{\\hfil\\scalebox{.7} {\\includegraphics{Mass.eps}\\includegraphics{Q.eps}\\includegraphics{Temp.eps}\\hfil} {\\caption{\\footnotesize{ ${\\cal M},  {\\cal Q}^2$ and $T$ plotted in terms of the initial mass vs. normalized retarded time $u$ for the Vilkovisky scenario.  Notice that the temperature is not proportional to $1/\\cal M$ but is reasonably flat until it drops to zero and that $\\cal M$ has dropped to about .9 of its initial value by the time evaporation ceases.} \\label{figs}}}} \\end{figure}   In what follows we shall assume that \\be T = T_{max} = constant = 4.8 \\times 10^{20} \\frac{M_{pl}}{M} \\ \\mathrm{MeV}\\label{TV2} \\ee or \\be T \\approx \\frac{10^{15}}{M} \\ \\mathrm{MeV},  \\label{TV3} \\ee where $M$ is in grams.  The retarded time taken for evaporation is \\be \\D u \\approx 96\\p \\frac{M^3}{M_{pl}^2}, \\ee where $u$ is evaluated at ${\\cal I}^+$.  We assume that asymptotically $u$ becomes the cosmic time in a flat Friedmann-Lema\\^ itre-Robertson-Walker universe, which in the matter-dominated case yields $\\D u = \\D t(1-3(t_0/\\D t)^{2/3})$ for initial time $t_0$. We see that for $\\D t >> t_0$ we may take $\\D u \\approx \\D t$.   Thus the cosmic time over which the hole is evaporating \\be \\t_{evap} \\approx 96\\p \\frac{M^3}{M_{pl}^2} \\approx 1.6 \\times 10^{-27} M^3 \\label{tevap} \\ee for $M$ in grams.\\footnote{As in Hawking's original work, Vilkovisky's calculation was carried out for a single scalar field.  If $N$ species of particles are emitted with equal probability, the evaporation time will decrease by a factor $1/N$; however Page's results\\cite{Page76}  show that emission is highly suppressed for particles with nonzero spin and thus only spin-zero need to be considered.  Unfortunately, we do not know the number of spin-0 species the energies of interest in this paper.  If $N \\lesssim 10$, the results to be discussed go through unchanged.  If $N \\gtrsim 10$, the limits to be discussed will necessarily be relaxed even further, because the integration time over pbh evaporation will be drastically reduced.\\label{foot1}}  Equating $\\t_{evap}$ with the age of the universe $t_u \\approx 4.5 \\times 10^{17}$ s, gives a  value, \\be M_{u} \\approx 6.6 \\times 10^{14} \\ \\mathrm{grams}. \\ee This is essentially the canonical $10^{15}$ gram black hole but in Vilkovisky's scenario it is not the mass of a hole that is evaporating during the current epoch; it is merely the mass of a hole whose evaporation is currently ceasing, by which time it will have lost only about ten percent of its mass, as can be seen in figure 1.  Thus, $M_{u}$ is the value above which black holes effectively do nothing whatsoever by the current age of the universe.  We expect that such a change should have drastic astrophysical consequences. ",
        "conclusions": ""
    },
    "0911/0911.5254_arXiv.txt": {
        "abstract": "{In this review, I present our current state of knowledge regarding both Classical Nova and Recurrent Nova systems. Two particular objects (V1974 Cyg and RS Oph) are chosen to illustrate the range of phenomena that may be associated with their outbursts. The wider importance of nova research is emphasised and some of the most pressing unsolved problems are summarised.} ",
        "introduction": "The outbursts of novae have been recorded for over 2000 years (for an historical review, see Duerbeck 2008). It was only in the 1920s during the ``Great Debate''  that it was realised that ``ordinary'' novae such as T Aur (1892 - often seen as the first well-studied nova outburst) were very distinct from ``supernovae'' such as S And (1885 - in M31). Later, Dwarf Novae (DNe) and Classical Novae (CNe) were in turn recognised as rather different beasts and certain of the Classical Novae were also subsequently reclassified as Recurrent Novae (RNe) when a second major outburst was recorded (the earliest example being T Pyx with the 1902 outburst repeating that first noted for this object in 1890).  The watershed in our understanding of the outbursts of CNe (and later RNe) came from a combination of the realisation that cataclysmic variables in general, and CNe in particular, had close binary central systems, together with the suggestion that explosive hydrogen burning was the root cause of the outburst (Kraft 1964). The relationship between DNe, CNe and RNe has long been established. The debate over the relationship of RNe to supernovae continues, and will be touched on here, and in several other contributions in these proceedings. ",
        "conclusions": "As has been noted above, we have made spectacular advances over the last 50 years in our understanding of the nova phenomenon. We have also realised the significance of their study for diverse branches of modern astrophysics. However, as might be expected, there are many important questions still to answer. The following list is ordered from pre-nova, through outburst, to wider issues. It is not exhaustive, and is inevitably subject to some personal biases, but it hopefully encompasses many of the most pressing problems we currently face:  \\begin{itemize}  \\item{What is the evolutionary track of a binary to the CN/RN phase? Conversely, can we use novae to help to trace the evolutionary history of binary stars in general? GK Per and V458 Vul may help with the former, and although inevitably rare, there may be more examples to be found. However, studies of novae in extragalactic systems, including the determination of the relationship of nova rates and sub-types to stellar populations, may turn out to be particularly important here.}  \\item{Is there a continuum of inter-outburst timescales (and other fundamental properties) from CNe through RNe (at least for the short period sub-types of RNe)? Again, studies of large samples of novae in extragalactic systems such as M31 are likely to help here.}  \\item{Do we need to revisit our TNR models particularly in the light of the Swift results? For example, what are the true durations of any super-Eddington and ``plateau'' phases of the bolometric evolution of novae? In this regard, and others, it is a great shame that there is no equivalent of the IUE satellite either currently in operation, or indeed planned for the longer term.}  \\item{Related to the previous point, what is the cause of the remarkable variability during the emergence of the SSS in RS Oph, and also, is a nuclear burning instability really the origin of the 35s modulation seen in the XRT data while the SSS was bright?}  \\item{There is still a significant (order of magnitude) discrepancy between the ejected mass found from models of the CN outburst compared to that derived from observations (although this seems less of a problem with the U Sco and RS Oph sub-classes of RNe). Consideration of clumpiness in the ejecta helps to a degree, but does not appear sufficient and it may mean we need to re-examine our models.}  \\item{Can studies of objects such as RS Oph help to determine more precisely the mechanism of the formation of jets in astrophysical sources? Although examples among the nova population may be rare, and perhaps confined only to the RNe, there is no doubting the wider importance of such work.}  \\item{Is there indeed a link between RNe and Type Ia SNe? Although it seems that for systems like RS Oph and U Sco, the mass of the WD is near the Chandrasekhar limit and increasing, there remain some fundamental questions. For example, what is the type of the WD (if ONe, then no SN explosion will occur); can the H in systems with red giant secondaries really be ``hidden'' at the time of any SN outburst; is the population of (appropriate sub-type) RNe sufficient to explain the observed SN Ia rate? A combination of detailed observations of individual Galactic RNe,  coupled with surveys of extragalactic novae, will help to answer these very important points.}  \\item{Can the MMRD be refined, possibly by choosing a suitable sub-set of CNe with particular characteristics, and/or by improving the homogeneity and cadence of observations of novae, both Galactic and extragalactic? Although their use as distance indicators for cosmology may have diminished in importance, they may still be very useful to determine the distances to particular extragalactic systems. In addition of course, the MMRD is often used to define the distance to newly discovered novae, and coupled with often poorly determined extinctions, distance estimates  can be very uncertain as a result. The advent of GAIA may prove a watershed here.}  \\end{itemize}"
    },
    "0911/0911.2704_arXiv.txt": {
        "abstract": "Protoplanetary disks with AU-scale inner clearings, often referred to as transitional disks, provide a unique sample for understanding disk dissipation mechanisms and possible connections to planet formation.  Observations of young stellar clusters with the {\\it Spitzer} Space Telescope have amassed mid-infrared spectral energy distributions for thousands of star-disk systems from which transition disks can be identified.  From a sample of 8 relatively nearby young regions ($d \\lesssim 400$ pc), we have identified about 20 such objects, which we term ``classical\" transition disks, spanning a wide range of stellar age and mass. We employed strict infrared continuum criteria to limit ambiguity: an 8 to 24 \\microns spectral slope limit ($\\alpha > 0$) to select for robust optically thick outer disks, and 3.6 to 5.8 \\microns spectral slope and 5.8 \\microns continuum excess limits to select for optically thin or zero continuum excess from the inner few AU of the disks. We also identified two additional categories representing more ambiguous cases: \"warm excess\" objects with transition-like spectral energy distributions but moderate excess at 5.8 \\micron, and \"weak excess\" objects with smaller 24 \\microns excess that may be optically thin or exhibit advanced dust grain growth and settling. From existing H$\\alpha$ emission measurements, we find evidence for different accretion activity among the three categories, with a majority of the classical and warm excess transition objects still accreting gas through their inner holes and onto the central stars, while a smaller fraction of the weak transition objects are accreting at detectable rates. We find a possible age dependence to the frequency of classical transition objects, with fractions relative to the total population of disks in a given region of a few percent at 1-2 Myr rising to 10-20\\% at 3-10 Myr.  The trend is even stronger if the weak and warm excess objects are included. This relationship may be due to a dependence of the outer disk clearing timescale with stellar age, suggesting a variety of clearing mechanisms working at different times, or it may reflect that a smaller fraction of all disks actually undergo an inner clearing phase at younger ages. Classical transition disks appear to be less common, and weak transition disks more common, around lower-mass stars ($M \\lesssim 0.3 \\msun$), which we suggest may be a further indicator of the stellar mass-dependent disk evolution that has been seen in previous studies. The difference in number statistics and accretion activity between the two classes further suggests that they are not connected but rather represent distinct evolutionary outcomes for disks. ",
        "introduction": "The lifetimes of circumstellar disks around young stars determine the relevant timescales available for planet formation to occur. Observational estimates are essential for constraining theories such as the core accretion (Pollack et al. 1996) and gravitational instability (Boss 1997) models.  A considerable amount of work has gone into estimating disk lifetimes primarily by examining infrared excess emission as a function of age.  Previous studies focused on near-infrared emission probing the innermost regions of disks have indicated typical disk lifetimes of roughly 3 Myr, with a wide dispersion from 1-10 Myr (Strom et al. 1989; Haisch et al. 2001; Hillenbrand 2005).  Constraints on the lifetime of disk material in the planet formation zone of 0.1-10 AU have been more difficult to achieve because of sensitivity and resolution limitations at the longer wavelengths which probe cooler dust.  Early results from ground-based and IRAS observations of the nearby Taurus star forming region by Skrutskie et al. (1990) first identified a small population of disks ``in transition\", lacking excess emission at $\\lambda \\lesssim 10$ \\microns but exhibiting considerable emission at longer wavelengths and thus implying a lack of dust (at least in small grains) in the inner few AU of the disks.  Given the age of Taurus, Skrutskie et al. estimated a typical transition timescale of $\\sim 0.3$ Myr.  A handful of such transition disks with inner ``holes\" have been studied in some detail in recent years \\footnote{The ``transition disk\" terminology is somewhat subjective and has a wide range of definitions in the literature.  For the purposes of this paper, we define a ``classical\" transition disk specifically as an optically thick protoplanetary disk with an $\\sim$ AU-scale inner region that is optically thin or completely evacuated of small dust grains.  We alternately refer to these as disks with inner holes.}. Calvet et al. (2002) modeled the spectral energy distribution (SED) of the 10 Myr-old classical T Tauri star (CTTS) TW Hya with an inner disk hole of size $\\sim 4$ AU, and hypothesized that a Jupiter-mass planet may have formed there. Similar conclusions have been drawn for the Taurus objects GM Aur (Rice et al. 2003; Bergin et al. 2004; Calvet et al. 2006), DM Tau (Bergin et al. 2004; Calvet et al. 2006) and CoKu Tau/4 (Forrest et al. 2004; Quillen et al. 2004; D'Alessio et al. 2005), though the latter has subsequently been found to harbor a stellar companion (Ireland \\& Kraus 2008).  Follow-up observations at high spatial resolution have directly confirmed the presence of holes or gaps in a few disks (Hughes et al. 2007; Ratzka et al. 2007; Brown et al. 2008; Hughes et al. 2009). However, the overall frequency of transition objects as a function of stellar age, mass, and environment has remained largely unconstrained.  With its superior mid-infrared sensitivity and spatial resolution, the {\\it Spitzer} Space Telescope has begun to address questions of disk evolution and lifetimes and the connection to planet formation in unprecedented detail.  {\\it Spitzer} observations taken with the IRAC and MIPS instruments are providing SEDs from 3.6 to 24 \\microns for large numbers of pre-main sequence stars in star forming regions and young clusters, allowing statistical studies of circumstellar disk evolution through the properties of dust emission as a function of stellar age, mass, and environment (Lada et al. 2006; Sicilia-Aguilar et al. 2006; Hern\\'andez et al. 2007, Balog et al. 2007; Kennedy \\& Kenyon 2009, among many others).  Here we report on the statistics of transition disks as identified from Spitzer SEDs for stars in nearby young regions, providing the first robust statistics of disks with inner holes around stars in young stellar clusters and associations with ages 1-10 Myr. ",
        "conclusions": "Understanding the properties of circumstellar disks with significant dust clearing in their inner regions may have profound implications for planet formation processes.  If indeed these transition disks are showing us the end-stages of protoplanetary disks, their statistics are indicative of disk dissipation mechanisms and timescales, which can then be used to constrain models of planetesimal growth and/or giant planet accretion. Some of the objects in our sample have ages of 1 Myr or less, which implies that the disk clearing mechanism(s) must operate very quickly. Our simple parametric model of disk dissipation indicates a characteristic ``transition\" time in the range $0.1 - 1\\times 10^6$ years, possibly dependent on the stellar age and/or mass. There have been competing claims in the recent literature regarding the transition timescale, with some advocating for a rapid transition of order $10^5$ years (e.g., Furlan et al. 2009) and others arguing for a much longer timescale comparable to the stellar age (e.g., Sicilia-Aguilar et al. 2008; Currie et al. 2009ab). Much of the discrepancy may lie in the wide range of adopted definitions of what constitutes a transitional disk; the studies that derive longer transition times tend to use much less restrictive selection criteria. Our results support a slower transition for disks only around stars older than $\\sim 3$ Myr.  As we mentioned in section 3.3.2, it is difficult to directly compare statistics for regions of different ages because they also appear to probe different stellar mass regimes. This uncertainty can only be addressed once the dependence of disk evolution on stellar mass has been better-delineated. All of the transition times were estimated assuming that every disk goes through the same inner hole phase.  However, our observation of a possible deficit of classical transition disks around late-type stars, as well as the apparent age dependence of the transition disk frequency, appears to contradict this assumption.  If classical transition disks do not represent a universal phase of disk evolution, then the clearing timescale may be longer than our estimates.  The presence of accretion, and hence gas within the dust hole, in a majority of the classical transition objects presents challenges to understanding how the inner disk is cleared of small grains while the outer disk (given the significant excesses at $\\lambda > 10$ \\micron) is not. One possibility is that significant grain growth to meter-size or larger bodies has occurred in the inner disk, as has been suggested by models of grain coagulation (Weidenschilling 1997; Dullemond \\& Dominik 2005). However, some other process is then needed to decouple the outer disk so that accretion from there cannot replenish the supply of small grains to the inner disk.  In this scenario, the inner disk should then drain onto the star in a viscous time, which for typical T Tauri parameters is of order a few $10^5$ years, similar to some of the clearing timescales estimated above. Several gap-creating mechanisms have been proposed for T Tauri disks, including photoevaporative mass loss induced by UV radiation from the central star (e.g., Clarke et al. 2001) and the dynamical influence of a giant planet (e.g., Bryden et al. 1999).  Najita et al. (2007) review these mechanisms and suggest that no single one operates in all cases. Our results support this view, particularly if the hole clearing timescale or incidence really does vary with stellar age and mass. The weak excess transition objects may represent a stronger case for the photoevaporation mechanism: the low or absent accretion activity and low disk masses as implied by the small fractional infrared excess are both properties predicted by the models at the time of inner hole formation (see also Cieza et al. 2008).  However, the large number of these objects, especially at older ages, may not be consistent with photoevaporation model predictions that the outer disk should dissipate very rapidly following inner hole formation (less than $10^5$ years; Alexander et al. 2006).  This may be resolved if photoevaporation is less efficient for M stars; current models have not yet fully explored the effect of stellar mass.  The inner disk clearing we see in the accreting transition objects (if not all of them) is most likely due to an advanced stage of dust evolution, either from growing planetesimals or possibly newly formed planets. Such processes occur across a large range of ages, from the youngest $\\sim 1$ Myr embedded star forming regions to some of the oldest known classical T Tauri stars such as the 10 Myr-old TW Hya, with some evidence for an increasing frequency with age. The cause of such disparate onset times for inner disk clearing is a key question that remains unanswered.  Initial conditions, such as the initial disk size and mass, may play a role. For example, disks of the same mass that are born with different initial sizes will viscously evolve at different rates, may form planets at different times, and may have different susceptibility to photoevaporation. All of these properties can influence the onset time of inner clearing and the amount of time required for the outer disk to dissipate. More massive stars may be born with more massive disks, which might have the best chance for forming giant planets.  Such a scenario could explain the young transition disks, which are preferentially associated with more massive stars. It may also explain the deficit of classical transition disks around the lowest mass stars, if in fact photoevaporation turns out not to be an efficient gas removal mechanism.  Alexander \\& Armitage (2009) have explored some of these parameters by calculating disk evolution models including the effects of giant planet formation, migration, and photoevaporation.  Interestingly, the transitional disk fraction they predict as a function of stellar age is qualitatively similar to our observed trend (although their values at 5-10 Myr are somewhat higher).  However, they used a solar-mass star for all models, so the stellar mass dependence remains to be considered.  We have shown that most classical transition disks exhibit trace amounts of excess emission at 5.8 \\micron, indicating that the inner holes are not totally devoid of small grains.  This is consistent with models of {\\it Spitzer}-IRS spectra of transition disks published to date, which can only explain the observed 10 \\microns silicate emission in most cases by placing optically thin dust inside the hole (e.g. Calvet et al. 2005). The origin of this dust is unclear.  It is possible that in objects that are still accreting, some dust is entrained in the accretion flow passing from the outer disk into the hole; numerical simulations by Rice et al. (2006) of a disk gap cleared by a giant planet show that pressure gradients at the outer edge of the gap can filter dust particles, allowing only small grains to pass through the gap along with the gas accretion flow. However, many of the non-accreting transitional systems also show hot dust excess, although some of these may be accreting at rates too low to produce observable diagnostics.  Some objects, particularly those with larger $f_{5.8}$, may still have optically thick material near the dust sublimation radius.  In that case, the emission may have been reduced from that of typical disks because of a lower scale height brought about by dust grain growth and settling.  An asymmetric inner disk geometry such as a warp might also produce lower excess emission at certain orientations; this may manifest itself through mid-infrared variability, which has been recently discovered in the transitional disk LRLL 31 (Muzerolle et al. 2009). Finally, there may also be {\\it in situ} dust production via colliding planetesimals, as Eisner et al. (2006) proposed for TW Hya, possibly tracing active terrestrial planet formation.  One obvious question is whether the three classes of evolved disks identified here are connected. We would argue that the classical and warm excess transition objects do likely represent an evolutionary sequence. The warm excess sources are akin to the pre-transitional disks described by Espaillat et al. (2007), and are most likely produced as a result of embedded giant planets carving out large gaps in the disks (Espaillat et al. 2008).  As the planet continues to grow via accretion from the outer disk, it should eventually stifle accretion through the gap and onto an optically thick inner disk, which may render the interior completely optically thin and result in a classical transition disk. However, the link between these and the weak excess disks is much more dubious.  The sheer numbers of the weak excess objects, particularly around older stars, are too great to simply be the result of further evolution of the outer disks of classical transition objects. The spectral type distributions are also distinctly different. Other studies have suggested that there may be separate disk evolutionary pathways (e.g. Lada et al. 2006; Cieza et al. 2007) of which the different disk classes may be representative. However, the weak excess category is less well-defined, and very likely contains a mixture of different disk types. The identification of true inner holes is not as robust for the weak excess disks around lower mass stars, as recently pointed out by Ercolano et al. (2009).  They may instead represent evolution without a cleared inner hole phase.  Similar objects have been dubbed ``homologously depleted\" (Wood et al. 2002; Currie et al. 2009b), suggesting that the dust opacity has been reduced through grain growth and mass loss across a wide range of disk radii.  We would like to point out that dust grain growth and settling acts to shrink the effective range of disk radii probed at a given wavelength; the models of D'Alessio et al. (2006) show that the majority of dust emission at 25 \\microns is located within $\\sim 1-2$ AU in a highly settled disk.  Thus, radial dust evolution in these disks cannot be definitively ruled out without measurements at far-infrared wavelengths. It is also very possible that some of the weak excess transition disks may in fact be recently-formed debris disks; further study of their gas content is required before this distinction can be made.  The recent discovery that the CoKu Tau/4 disk is circumbinary (Ireland \\& Kraus 2008) has fueled speculation as to whether all transitional disks are in fact the result of disk clearing by a stellar companion, and indeed whether such objects are truly transitional. We argue that regardless of hole formation mechanism, all objects showing these SED characteristics are likely to be in transition.  There are examples of close binaries whose SEDs betray no evidence for significant inner holes (DF Tau, separation $\\sim 13$ AU, and GW Ori, separation $\\sim 1$ AU, to name a few), which suggests that disks around close binary systems can and do evolve in a similar fashion.  Unfortunately, we lack binary statistics for our sample, particularly at close separations ($< 10$ AU). We can indirectly assess the importance of binaries by comparing the expected fraction of binaries with our observed transition disk statistics as a function of stellar age and mass. From the multiplicity study by Sterzik \\& Durisen (2004), we estimate a total multiplicity fraction of about 50\\% for solar-mass stars, and roughly 20\\% at $0.2 \\; \\msun$.  Since our transition disk selection is sensitive to maximum inner disk hole sizes of roughly 40 AU for solar type stars and 4 AU for low mass stars, any stellar companions must be within this range in order to produce the disk clearing. According to Sterzik \\& Durisen (2004), about 50\\% of the $1 \\; \\msun$ binaries have separations $<40$ AU while about 30\\% of the $0.2 \\; \\msun$ binaries will have separations within 4 AU, leading to combined fractions of 25\\% and 6\\%, respectively.  These estimates are inconsistent with the transition disk statistics in both magnitude and trend with mass, thus we suggest that stellar companions are unlikely to be a significant contributor. Nevertheless, sensitive searches for close stellar companions, both with AO imaging and radial velocity monitoring, need to be done before firm conclusions can be made.  Finally, the possibility of different disk decay rates at different stellar ages, as indicated in Figure~\\ref{holefrac}, is very intriguing in its own right.  It may be a further indicator that different disk dissipation mechanisms operate preferentially at different times.  The stellar mass dependence, where disks around less massive stars tend to last longer, may also be involved.  For example, if either photoevaporation or planet formation is less effective in mid- to late-M stars, their disks may take longer to completely erode.  Alternatively, stellar multiplicity may have an effect; evidence suggests that disks tend to disappear more quickly around binaries (e.g. Bouwman et al. 2006), and these may represent the faster decay rate at younger ages.  However, counter-examples (long-lived disks around binaries, short-lived disks around single stars) are plentiful.  More work needs to be done to fully elucidate the effects of these parameters, as well as others such as environment and initial conditions."
    },
    "0911/0911.0701_arXiv.txt": {
        "abstract": "Globular cluster systems in most large galaxies display bimodal color and metallicity distributions, which are frequently interpreted as indicating two distinct modes of cluster formation.  The metal-rich (red) and metal-poor (blue) clusters have systematically different locations and kinematics in their host galaxies.  However, the red and blue clusters have similar internal properties, such as the masses, sizes, and ages.  It is therefore interesting to explore whether both metal-rich and metal-poor clusters could form by a common mechanism and still be consistent with the bimodal distribution.  We show that if all globular clusters form only during mergers of massive gas-rich protogalactic disks, their metallicity distribution could be statistically consistent with that of the Galactic globulars.  We take galaxy assembly history from cosmological dark matter simulations and couple it with the observed scaling relations for the amount of cold gas available for star formation.  In the best-fit model, early mergers of smaller hosts create exclusively blue clusters, whereas subsequent mergers of progenitor galaxies with a range of masses create both red and blue clusters.  Thus bimodality arises naturally as the result of a small number of late massive merger events.  We calculate the cluster mass loss, including the effects of two-body scattering and stellar evolution, and find that more blue clusters than red clusters are disrupted by the present time, because of their smaller initial masses and larger ages.  The present-day mass function in the best-fit model is consistent with the Galactic distribution.  However, the spatial distribution of model clusters is much more extended than observed and is independent of the parameters of our model. ",
        "introduction": "A self-consistent description of the formation of globular clusters remains a challenge to theorists.  A particularly puzzling observation is the apparent bimodality, or even multimodality, of the color distribution of globular cluster systems in galaxies ranging from dwarf disks to giant ellipticals.  This color bimodality likely translates into a bimodal distribution of the abundances of heavy elements such as iron.  We know this to be the case in the Galaxy as well as in M31, where relatively accurate spectral measurements exist for a large fraction of the clusters.  The two most frequently encountered modes are commonly called {\\it blue} (metal-poor) and {\\it red} (metal-rich).  Bimodality in the globular cluster metallicity distribution of luminous elliptical galaxies was proposed by \\cite{zepf_ashman93}, following a theoretical model of \\cite{ashman_zepf92}.  The concept of cluster bimodality became universally accepted because the two populations also differ in other observed characteristics.  The system of red clusters have a significant rotation velocity similar to the disk stars whereas blue clusters have little rotational support, in the three disk galaxies observed in detail: Milky Way, M31, and M33. In elliptical galaxies, blue clusters have a higher velocity dispersion than red clusters, both due to lack of rotation and more extended spatial distribution.  Red clusters are usually more spatially concentrated than blue clusters (\\cite{brodie_strader06}). All of these differences, however, are in external properties (location and kinematics), which reflect {\\it where} the clusters formed, but not {\\it how}.  The internal properties of the red and blue clusters are similar: masses, sizes, and ages, with only slight differences.  Even the metallicities themselves differ typically by a factor of 10 between the two modes, not enough to affect the dynamics of molecular clouds from which these clusters formed.  Could it be then that both red and blue clusters form in a similar way on small scales, such as in giant molecular clouds, while the differences in their metallicity and spatial distribution reflect when and where such clouds assemble?  All scenarios proposed in the literature assumed different formation mechanisms for the red and blue clusters, and most scenarios envisioned the stellar population of one mode to be tightly linked to that of the host galaxy (e.g., \\cite{forbes_etal97}, \\cite{cote_etal98}, Strader et al.~2005).  The other mode is assumed to have formed differently, in some unspecified ``primordial'' way. This assumption only pushed the problem back in time but it did not solve it.  For example, \\cite{beasley_etal02} used a semi-analytical model of galaxy formation to study bimodality in luminous elliptical galaxies and needed two separate prescriptions for the blue and red clusters.  In their model, red clusters formed in gas-rich mergers with a fixed efficiency of 0.007 relative to field stars, while blue clusters formed in quiescent disks with a different efficiency of 0.002.  The formation of blue clusters also had to be artificially truncated at $z=5$.  \\cite{strader_etal05} and \\cite{rhode_etal05} suggested that the blue clusters could instead have formed in very small halos at $z > 10$, before cosmic reionization removed cold gas from such halos.  This scenario requires high efficiency of cluster formation in the small halos and also places stringent constraints on the age spread of blue clusters to be less than 0.5 Gyr. Unfortunately, even the most recent measurements of the relative cluster ages in the Galaxy (\\cite{deangeli_etal05, marin-franch_etal09}) cannot detect age differences smaller than 9\\%, or about 1 Gyr, and therefore cannot support or falsify the reionization scenario. ",
        "conclusions": ""
    },
    "0911/0911.1127_arXiv.txt": {
        "abstract": "We present the discovery of a brown dwarf or possible planet at a projected separation of 1.9$^{\\prime\\prime} =$ 29\\,AU around the star GJ~758, placing it between the separations at which substellar companions are expected to form by core accretion ($\\sim$5\\,AU) or direct gravitational collapse (typically $\\gtrsim$100\\,AU). The object was detected by direct imaging of its thermal glow with Subaru/HiCIAO. At 10--40 times the mass of Jupiter and a temperature of 550--640\\,K, GJ~758~B constitutes one of the few known T-type companions, and the coldest ever to be imaged in thermal light around a sun-like star.  Its orbit is likely eccentric and of a size comparable to Pluto's orbit, possibly as a result of gravitational scattering or outward migration.  A candidate second companion is detected at 1.2$^{\\prime\\prime}$ at one epoch. ",
        "introduction": "While the extrasolar planets and brown dwarf companions currently known from direct imaging mark significant discoveries, it is important to note that none of these systems present scenarios analogous to a Solar-like system.  Almost all imaged substellar companions either feature extreme orbital separations \\citep[hundreds of AU, e.g.][]{luhman2007,kalas2008}, star-like rather than planet-like temperatures \\citep[e.g.][]{marois2008}, or host stars at the extreme ends of the mass spectrum \\citep[A- and late M-type, e.g.][]{kalas2008,chauvin2005}. The closest match is probably Gl~229~B, a brown dwarf orbiting an M1 star at 40\\,AU \\citep{nakajima1995}. This shows how a remarkable gap still exists between recent direct-imaging discoveries and the exploration of Sun-like systems.  In this work, we present the discovery of a substellar companion with a mass of 10--40\\,$M_\\mathrm{Jup}$, a temperature of 550--640\\,K, and a projected separation of 1.9$^{\\prime\\prime} =$ 29\\,AU. The estimated semi-major axis range, overlapping with our own Solar System's outer planet orbits, along with the Sun-like parent star, make it a superior laboratory for advancing our conventional wisdom of planet or brown dwarf formation among solar-like systems. ",
        "conclusions": "We present the discovery of a substellar companion to the star GJ~758 at a separation of 1.9$^{\\prime\\prime}$ by angular differential imaging in H-band with Subaru/HiCIAO.  Common proper motion with its host is demonstrated.  We derive ranges of 10.3--39.6 $M_\\mathrm{Jup}$ for the mass and 549--637\\,K for the temperature of the companion, classifying it as either a methane-bearing high-mass planet or low-mass brown dwarf. A candidate second companion is detected at a separation of 1.2$^{\\prime\\prime}$, which could represent a migration partner of GJ~758~B.  This unique configuration of the GJ~758 system, its short dynamical timescale (observable motion in only 3 months), and its demonstrated accessibility by angular differential imaging, make it an excellent showcase target for the direct investigation of the formation and evolution mechanisms of planets and brown dwarfs."
    },
    "0911/0911.3408_arXiv.txt": {
        "abstract": "The classic example of a Type IIb supernova is SN 1993J, which had a cool extended progenitor surrounded by a dense wind.  There is evidence for another category of Type IIb supernova which has a more compact progenitor with a lower density, probably fast, wind.  Distinguishing features of the compact category are: weak optical emission from the shock heated envelope at early times; nonexistent or very weak H emission in the late nebular phase; rapidly evolving radio emission; rapid expansion of the radio shell; and expected nonthermal as opposed to thermal X-ray emission.  Type IIb supernovae that have one or more of these features include SNe  1996cb, 2001ig, 2003bg, 2008ax, and 2008bo.  All of these with sufficient radio data (the last four) show evidence for presupernova wind variability.  We estimate a progenitor envelope radius $\\sim1\\times 10^{11}$ cm for SN 2008ax, a value consistent with a compact Wolf-Rayet progenitor.  Supernovae in the SN 1993J extended category include SN 2001gd and probably the Cas A supernova.  We suggest that the compact Type IIb events be designated Type cIIb and the extended ones Type eIIb. The H envelope mass dividing these categories is $\\sim0.1\\Msun$. ",
        "introduction": "A Type IIb supernova has a spectrum with high velocity H lines near maximum, but in the late nebular phases more closely resembles a Type Ib supernova, with weak or absent H lines. The first supernova placed in this category was SN 1987K \\citep{filippenko88}. However, it was the nearby SN 1993J in M81 that has come to best define the properties of a SN IIb. In this case, the progenitor star was identified as a K supergiant \\citep{aldering94}, consistent with the interpretation of the early light curve which required a stellar radius of $\\sim 4\\times 10^{13}$ cm \\citep{woosley94}. Such a star is expected to have a slow, dense wind, and there is radio and X-ray evidence for a wind with mass loss rate $\\dot M\\approx 4\\times 10^{-5}\\ml$ for an assumed wind velocity of $v_w=10\\kms$ \\citep{fransson96,fransson98}. Radio VLBI observations provide strong support for the wind interaction picture \\citep{marcaide97,bartel02}. The interaction with the dense wind also produces some H$\\alpha$ emission in the late nebular phase \\citep{matheson00}.  A consistent scenario for SN 1993J is thus the explosion of an extended star that interacts with the dense wind from the progenitor star. However, there is also evidence for Type IIb supernovae arising from compact stars, and that is the topic addressed here. Some of that evidence has already been presented for SN 2001ig \\citep{ryder04}, SN 2003bg \\citep{soderberg06}, and SN 2008ax \\citep{pastorello08}. In \\S~2, we present the differences between SNe IIb that have extended or compact progenitors. The results are discussed in \\S~3. ",
        "conclusions": "One expectation of the two types of IIb is that the Type cIIb events should have a lower H mass envelope so that a large radial extent cannot be sustained. For SN 1993J, \\cite{woosley94} estimate a H mass of $0.2\\pm 0.05\\Msun$. For SN 2003bg, \\cite{mazzali09} estimate $0.05\\Msun$ of H, in keeping with expectations. However, \\cite{qiu99} suggest that SN 1996cb has more H than SN 1993J, although \\cite{deng01} estimate $0.1- 0.2\\Msun$ of H if the explosion energy is $10^{51}$ ergs. The implied dividing line between the types is $\\sim 0.1\\Msun$.  For lower amounts of H, we expect that the Type cIIb objects merge into the Type Ib's. A recent example of a Type Ib with H is SN 2007Y which showed high velocity ($\\sim 10,000\\kms$) H absorption near maximum light \\citep{stritz09}. The finding of H in the spectra of Type Ib events is common \\citep{branch02}. \\cite{branch02} estimate that a H mass $\\sim 0.01\\Msun$ is typically needed to produce the required optical depth in H$\\alpha$. There may  not be a clear separation of the Type cIIb and Type Ib events, but a gradual transition depending on remaining H mass.   The Cas A supernova was recently shown to be of Type IIb by the spectrum of its light echo \\citep{krause08}, which raises the question of whether the progenitor was compact or extended. Studies of the supernova remnant have not been conclusive; \\cite{chevalier03} argued that the progenitor was extended with the dense wind extending to the stellar surface, but there have been arguments for a blue phase before the supernova \\citep{hwang09}. The Type IIb identification was made on the basis of the close correspondence of the echo spectrum with the time integrated spectrum of SN 1993J.  The H$\\alpha$ line in the echo spectrum appears to be closer to the maximum light H$\\alpha$ in SN 1993J than in SN 2003bg \\citep{hamuy09}, but a more detailed comparison with the time integrated spectrum of a Type cIIb event would be useful.  The existing observations that allow a distinction between Type cIIb and eIIb supernovae, primarily at radio wavelengths, indicate a clear separation between the 2 types of events, but the numbers are small. Future time domain surveys should increase the discovery rate of Type IIb events so that this issue can be resolved. In addition, the early discovery of supernovae ($\\la 2$ days from explosion) will allow the classification of the events from their shock breakout and cooling envelope emission."
    },
    "0911/0911.3778_arXiv.txt": {
        "abstract": "Gas and star velocity dispersions have been derived for eight circumnuclear star-forming regions (CNSFRs) and the nucleus of the spiral galaxy NGC\\,3310 using high resolution spectroscopy in the blue and far red. Stellar velocity dispersions have been obtained from the Ca{\\sc ii} triplet in the near-IR, using cross-correlation techniques, while gas velocity dispersions have been measured by Gaussian fits to the H$\\beta$\\,$\\lambda$\\,4861\\,\\AA\\ and [\\OIII]\\,$\\lambda$\\,5007\\,\\AA\\ emission lines.   The CNSFRs stellar velocity dispersions range from 31 to 73\\,km\\,s$^{-1}$. These values, together with the sizes measured on archival HST images, yield upper limits to the dynamical masses for the individual star clusters between 1.8 and 7.1\\,$\\times$\\,10$^6$\\,M$_\\odot$, for the whole CNSFR between 2\\,$\\times$\\,10$^7$ and 1.4\\,$\\times$\\,10$^8$\\,M$_\\odot$, and 5.3\\,$\\times$\\,10$^7$\\,M$_\\odot$ for the nucleus inside the inner 14.2\\,pc. The masses of the ionizing stellar population responsible for the \\HII\\ region gaseous emission have been derived from their published H$\\alpha$ luminosities and are found to be between 8.7\\,$\\times$\\,10$^5$ and 2.1\\,$\\times$\\,10$^6$\\,M$_\\odot$ for the star-forming regions, and 2.1\\,$\\times$\\,10$^5$\\,M$_\\odot$ for the galaxy nucleus; they  therefore constitute between 1 and 7 per cent of the total dynamical mass.  The ionized gas kinematics is complex; two different kinematical components seem to be present as evidenced by different line widths and Doppler shifts. ",
        "introduction": "The gas flows in disc of spiral galaxies can be strongly perturbed by the presence of bars, although the total disc star formation rates (SFR) do not appear to be significantly affected by them \\citep{1998ARA&A..36..189K}. These perturbations of the gas flow trigger nuclear star formation in the bulges of some barred spiral galaxies. External environmental influences can have strong effects on the SFR, among them, the most important by far, is tidal interactions. The young extragalactic star clusters belonging to these systems have been the aim of different studies during the last decades \\citep[e.g.\\ ][]{1991MNRAS.253..245D,1993AJ....106.1354W,2005A&A...443...41M,2005A&A...431..905B,2006A&A...448..881B,2008A&A...489.1091M}. The enhancement of the SFR is highly variable depending on the star formation conditions, the degree of enhancement ranging from zero in gas-poor galaxies to around 10-100 times in extreme cases \\citep{1998ARA&A..36..189K}. Much larger enhancements are often seen in the circumnuclear regions of strongly interacting and merging systems \\citep{1998ARA&A..36..189K,1998ApJ...498..541K}.  Yet, the bulges of some nearby, non-interacting, spiral galaxies show intense star-forming regions located in a  roughly annular pattern around their nuclei. In the middle of last century, \\cite{1958PASP...70..364M} classified a sample of galaxies using as the principal classification criterion the degree of central concentration of light of each galaxy. An apparent fairly common phenomenon in some types of galaxies was pointed out by Morgan: their nuclear regions can consist of an extremely bright, small nucleus superposed on a considerably fainter background, or it may be made up of multiple ``hot-spots''. Almost a decade later, \\cite{1965PASP...77..287S} suggested a relationship between the existence of a bar and the presence of abnormal features in their nuclei for a survey of bright southern galaxies. These authors extended the survey to the whole sky \\citep{1967PASP...79..152S}, and found that $\\sim$\\,14\\,\\% of these galaxies presented peculiar nuclei. The distinctive nature (with respect to the more extended star formation in discs) of the luminous nuclear star-forming regions was fully revealed with the opening of the mid- and far-infrared (IR) spectral ranges \\citep[see e.g.\\ ][]{1972ApJ...176L..95R,1973ApJ...182L..89H,1978ApJ...220L..37R,1980ApJ...235..392T}.  In general, CNSFRs and giant \\HII\\ regions in the discs of galaxies are very much alike, although the former look more compact and show higher peak surface brightness \\citep{1989AJ.....97.1022K} than the latter. CNSFRs, with sizes going from a few tens to a few hundreds of parsecs \\citep[e.g.][]{2000MNRAS.311..120D} seem to be made of several \\HII\\ regions ionized by luminous compact stellar clusters whose sizes, as measured from high spatial resolution {\\it Hubble Space Telescope} (HST) images, are seen to be of only a few parsecs. Their large H$\\alpha$ luminosities, typically higher than 10$^{39}$\\,erg\\,s$^{-1}$, point to relatively massive star clusters as their ionization source. Although these \\HII\\ regions are very luminous (M$_v$ between -12 and -17) not much is known about their kinematics or dynamics for both the ionized gas and the stars. It could be said that the worst known parameter of these ionizing clusters is their mass. As derived with the use of population synthesis models their masses suggest that these  clusters are gravitationally bound and that they might evolve into globular cluster configurations \\citep{1996AJ....111.2248M}. {\\bf Further, deeper analysis as to whether or not such cluster would survive the hostile environment of the circumnuclear regions is extremely interesting but lies  outside the scope of the present work. } Classically it is assumed that the system is virialized hence the total mass inside a radius can be determined by applying the virial theorem to the observed velocity dispersion of the stars ($\\sigma_{\\ast}$). As pointed out by several authors \\citep[e.g.\\ ][]{1996ApJ...466L..83H}, at near-IR wavelengths (\\,8500\\,\\AA) the contamination due to nebular lines is much smaller  and since red supergiant stars, if present, dominate the light where the Ca{\\sc ii} $\\lambda\\lambda$\\,8498, 8542, 8662\\,\\AA\\ triplet (CaT) lines are found, these should be easily observable allowing the determination of $\\sigma_{\\ast}$ \\citep{1990MNRAS.242P..48T,1994A&A...288..396P}.  The equivalent width of the emission lines are lower than those shown by the disc \\HII\\ regions \\citep[see e.g.\\ ][]{1989AJ.....97.1022K,1997AJ....113..975B,1999ApJ...510..104B}. Combining GEMINI data and a grid of photo-ionization models \\cite{2008A&A...482...59D} conclude that the contamination of the continua of CNSFRs by underlying contributions from both old bulge stars and stars formed in the ring in previous episodes of star formation (10-20\\,Myr) yield the observed low equivalent widths.  This is the third paper of a series to study the peculiar conditions of star formation in circumnuclear regions of early type spiral galaxies, in particular the kinematics of the connected stars and gas. In this paper we present high-resolution  far-red spectra and  stellar velocity dispersion measurements ($\\sigma_{\\ast}$) along the line of sight for eight CNSFRs and the nucleus of the spiral galaxy NGC\\,3310.  NGC\\,3310 (UGC\\,5786, Arp217) is a starburst galaxy classified as an SAB(r)bc by \\cite{1991trcb.book.....D}, with an inclination of the galactic disc of about i\\,$\\sim$\\,40 \\citep{2000MNRAS.312....2S}. Its coordinates are $\\alpha_{\\rm 2000}$=10$^h$\\,38$^m$\\,45\\fs9, $\\delta_{\\rm 2000}$=+53$^{\\circ}$\\,30\\arcmin\\,12\\arcsec \\citep{1991trcb.book.....D}. These authors derived a distance to the galaxy equal to 15\\,Mpc, giving a linear scale of $\\sim$\\,73\\,pc\\,arcsec$^{-1}$. This galaxy is a good example of an overall low metallicity galaxy, with a high rate of star formation and very blue colours. This galaxy has a ring of star forming regions whose diameter ranges from 8\\,\\arcsec to 12\\,\\arcsec and shows two tightly wound spiral arms \\citep{2002AJ....123.1381E,1976A&A....48..373V} filled with giant \\HII\\ regions. These circumnuclear regions present low metal abundance \\citep[0.2-0.4\\,Z$_\\odot$][]{1993MNRAS.260..177P}, in contrast to what is generally found in this type of objects, which show high metallicities \\citep{2007MNRAS.382..251D}. In fact, in most cases, the [O{\\sc iii}]\\,$\\lambda$\\,5007\\,\\AA\\ forbidden line can barely be seen \\citep[see e.g.\\ ][hereafter Paper I and Paper II, respectively]{2007MNRAS.378..163H,2009MNRAS.396.2295H}.  The ages indicated by the colours and magnitudes of the star formation regions are lower than 10\\,Myr \\citep{2002AJ....123.1381E}. From near-IR J and K photometry these authors derived an average age of $\\sim$10$^7$\\,yr for the large scale ``hot-spots'' (star forming complexes). From the observed CaT line in the Jumbo \\HII\\ region \\cite{1990MNRAS.242P..48T} derived an age around 5 to 6\\,Myr. \\cite{2002AJ....123.1381E}, comparing their data with Starburst99 models \\citep{1999ApJS..123....3L}, estimated masses of the large ``hot-spots'' ranging from 10$^4$ to several times 10$^5$\\,M$_\\odot$. They found 17 candidate super star clusters (SSCs) with absolute magnitudes between M$_B$\\,=\\,-11 and -15\\,mag, and with colours similar to those measured for SSCs in other galaxies \\citep[see for example][]{1995AJ....110.1009B,2001ApJ...556..801L}.  We have measured the ionized gas velocity dispersions ($\\sigma_{g}$) from high-resolution blue spectra using Balmer H$\\beta$ and [O{\\sc iii}] emission lines. The comparison between $\\sigma_{\\ast}$ and $\\sigma_{g}$  {\\bf on an ample sample of objects} might throw some light on the yet unsolved issue about the validity of the gravitational hypothesis for the origin of the supersonic motions observed in the ionized gas in Giant \\HII\\ regions \\citep*{1999MNRAS.302..677M}. In Section 2 we describe the observations and data reduction. We present the results in Section 3, the dynamical mass derivation in Section 4 and the ionizing star cluster properties in Section 5. We discuss all our results in Section 6. Finally, the summary and conclusions are given in Section 7. ",
        "conclusions": "We have measured gas and stellar velocity dispersions in eight CNSFRs and the nucleus of the barred spiral galaxy NGC\\,3310. The stellar velocity dispersions have been measured from the CaT lines at $\\lambda\\lambda$\\,8494, 8542, 8662\\,\\AA, while the gas velocity dispersions have been measured by Gaussian fits to the H$\\beta$\\,$\\lambda$\\,4861\\,\\AA\\ and the [O{\\sc iii}]\\,$\\lambda$\\,5007\\AA\\ emission lines on high dispersion spectra.  Stellar velocity dispersions are between 31 and 73\\,km\\,s$^{-1}$. The stellar and gas velocity dispersion are in relatively good agreement, with the former being slightly larger. The velocity dispersions from [O{\\sc iii}]\\,5007\\,\\AA\\ behave similarly to those from H$\\beta$. However, the best Gaussian fits involve two different components for the gas: a ``broad component\" with a velocity dispersion larger than that measured for the stars by about 20\\,km\\,s$^{-1}$, and a ``narrow component\" with a velocity dispersion lower than the stellar one by about 30\\,km\\,s$^{-1}$. This last velocity component shows a relatively constant value for the two gas emission lines, close to 22\\,km\\,s$^{-1}$ and also close to that measured for CNSFRs in the previously studied NGC\\,3351 and NGC\\,2903. {\\bf The velocity dispersion for the broad component of H$\\beta$ is similar to that of [OIII], contrary to what we found in Papers I and II for NGC~2903 and NGC~3351.} We find a velocity shift between the narrow and broad components of the multi-Gaussian fits that vary between -25 and 20\\,km\\,s$^{-1}$ in radial velocity.  When plotted in a [O{\\sc iii}]/H$\\beta$ versus [N{\\sc ii}]/H$\\alpha$ diagnostic diagram, the broad and narrow  systems and those values derived using a single Gaussian fit show very similar values, lying in the region of low-metallicity \\HII-like objects. This is in contrast to what we found in Papers~I and II where the two systems are clearly segregated for the high-metallicity regions of NGC\\,3351 and NGC\\,2903, with the narrow component showing lower excitation and being among the lowest excitation line ratios detected within the SDSS data set of starburst systems.  The rotation curve of NGC\\,3310 shows a typical $S$ feature, with the presence of some perturbations, particularly near the location of the Jumbo region. The values derived from the gas emission lines and the stellar absorption features are in very good agreement. The position going through the nucleus shows maximum and minimum values more or less coincident with the location of the CNSFRs.  The upper limits to the dynamical masses estimated from the stellar velocity dispersion using the virial theorem for the CNSFRs are in the range between 2.1\\,$\\times\\,10^7$ and 1.4\\,$\\times\\,10^8$\\,M$_\\odot$. For the nuclear region inside the inner 14.2\\,pc this upper limit is 5.3\\,$\\times\\,10^7$\\,M$_\\odot$. The upper limits to the derived masses for the individual clusters are between 1.8 and 7.1\\,$\\times$\\,10$^6$\\,M$_\\odot$.  Masses of the ionizing stellar clusters of the CNSFRs have been derived from their H$\\alpha$ luminosities under the assumption that the regions are ionization bound and without taking into account any photon absorption by dust. These masses derived for the star forming complexes of NGC\\,3310 are between 8.7\\,$\\times\\,10^5$ and 2.1\\,$\\times\\,10^6$\\,M$_\\odot$, and is 3.5\\,$\\times\\,10^6$ for the nucleus (see table \\ref{parameters}). Therefore, the ratio of the ionizing stellar population to the total dynamical mass is between 0.01 and 0.07.  Derived masses for the ionized gas, also from their H$\\alpha$ luminosities, vary between 1.5 and 7.2\\,$\\times\\,10^5$\\,M$_\\odot$ for the CNSFRs. For the nucleus, the derived mass of ionized gas is 4.7\\,$\\times\\,10^4$\\,M$_\\odot$.  It is interesting to note that, according to our findings, the SSC in CNSFRs seem to contain composite stellar populations. Although the youngest one dominates the UV light and is responsible for the gas ionization, it constitutes only about 10 per cent of the total mass. This can explain the low EWs of emission lines measured in these regions.  This may well apply to studies of other SSC and therefore conclusions drawn from fits of single stellar population models should be taken with caution \\citep[e.g.\\ ][]{2003ApJ...596..240M,2004AJ....128.2295L}. Furthermore, the composite nature of the CNSFRs  means that star formation in the rings is a process that has taken place over time periods much longer than those implied by the properties of the ionized gas."
    },
    "0911/0911.3064_arXiv.txt": {
        "abstract": "On 2009 January 22 numerous strong bursts  were detected from the anomalous X-ray pulsar \\src. \\swi/XRT and \\xmm/EPIC observations carried out in the following two weeks led to the discovery of three X-ray rings centered on this source. The ring radii increased with time following the expansion law expected for a short impulse of X-rays scattered by three dust clouds. Assuming different models for the dust composition and grain size distribution, we fit the intensity decay of each ring as a function of time at different energies, obtaining tight constrains on the distance of the X-ray source. Although the distance strongly depends on the adopted dust model, we find that some models are incompatible with our X-ray data, restricting to 4--8 kpc the range of possible distances for \\src. The best-fitting dust model provides a source distance of $3.91\\pm0.07$ kpc, which is compatible with the proposed association with the supernova remnant G\\,327.24--0.13, and implies distances of 2.2 kpc, 2.6 kpc and 3.4 kpc for the dust clouds, in good agreement with the dust distribution inferred by CO line observations towards \\src. However, dust distances in agreement with CO data are also obtained for a set of similarly well-fitting models that imply a source distance of $\\sim$5 kpc. A distance of $\\sim$4--5 kpc is also favored by the fact that these dust models are already known to provide good fits to the dust-scattering halos of bright X-ray binaries. Assuming $N_{\\rm H}=10^{22}$ cm$^{-2}$ in the dust cloud responsible for the brightest ring and a bremsstrahlung spectrum with $kT=100$ keV, we estimate that the burst producing the X-ray ring released an energy of 10$^{44-45}$ erg in the 1--100 keV band, suggesting that this burst was the brightest flare without any long-lasting pulsating tail ever detected from a magnetar. ",
        "introduction": "Soft X-rays are efficiently scattered at small angles by interstellar dust grains. This effect, as predicted by \\citet{overbeck65} and observationally confirmed by \\citet{rolf83}, is responsible for the presence of X-ray scattering halos around several bright   X-ray sources. The X-ray scattering cross section depends on the grain properties, and the observed radial distribution of the photons scattered in the halo also depends on the positions of the scattering grains along the line of sight. Therefore, the study of the energy-dependent radial profile of the X-ray halos gives useful information on the dust grain size, composition, and spatial distribution   (see e.g. \\citealt{predehl95}).   Scattered X-rays are delayed with respect to direct ones, owing to their longer path length from the source to the observer. In the case of   variable sources, these delays give rise to time-dependent effects on the halos properties, that can be used, if the   dust spatial distribution is known, to constrain the source distance \\citep{ts73}. This method was applied for the first time to constrain the distance of the X-ray binary Cyg~X-3 \\citep{predehl00}.  An interesting situation occurs when the dust is concentrated only in a thin layer and the X-ray emission has a  short duration. In this case the halo appears as a X-ray ring with increasing radius,  because only the photons scattered at a well determined  angle are detected at a given time. Such ``expanding rings'' (or ``light echoes'') have been observed to date in a handful of gamma-ray bursts (GRBs) and by measuring their expansion rate it has been possible to derive accurate distances of Galactic dust clouds (\\citealt{vianello07} and references therein).    Here we report on the discovery of dust scattering  rings around the Galactic source \\src\\ \\citep{tiengo09gcn8848}. \\src\\ belongs to the small class of anomalous X-ray pulsars (AXPs), which, together with the soft gamma-ray repeaters (SGRs) are thought to be magnetars, i.e. isolated neutron stars powered by the energy stored in their extremely strong magnetic field (see \\citealt{mereghetti08} for a recent review). These sources are characterized by the sporadic emission of very bright, short bursts, reaching peak luminosities above $\\sim$10$^{46}$ erg s$^{-1}$ in the extreme cases known as giant flares. In 2008 October \\src\\ emitted several short bursts and its X-ray flux increased significantly \\citep{israel09sub}. No further bursts were reported until 2009 January 22, when the source started a new period of much stronger bursting activity \\citep{mereghetti09}. This prompted the follow-up X-ray observations that led to the discovery of the three dust scattering rings discussed here.  The paper is organized as follows. In Section 2 we briefly review some properties of the X-ray dust scattering process that will be used in our analysis of the \\swi\\ and \\xmm\\  observations described in Section 3. From the expansion rate of the three dust rings we could establish that they originate from the scattering of the same event in three dust layers at different distances (Section 4.1). While this result is independent on the dust properties, the estimate of the distance by modeling the rings intensity (Section 4.2) requires the knowledge of the X-ray scattering differential cross section, hence it depends on the assumed properties of the dust. We found that only a subset of the different dust models that we explored can give satisfactory fits (Section 5.1). The possible distances for \\src\\ from the best-fitting models span a large range of values (Section 5.2), but this can be restricted based on independent information on the gas distribution in the Galaxy (Section 5.3). ",
        "conclusions": "Three bright X-ray rings were discovered around the anomalous X-ray pulsar \\src\\ after its strong bursting activity of 2009 January 22. Their radii were seen to increase and their flux to decrease with time, as expected for a short burst of X-rays scattered by three thin dust layers. A similar phenomenon had been observed in a few GRBs, but on a much smaller scale: in the best case, the rings around GRB\\,031203 \\citep{vaughan04} were observed by \\xmm\\ from 5 to 20 hours after burst, but only $\\sim$3,000 ring photons could be collected. The $\\sim$65,000 ring photons we detected for \\src\\ allowed us to carry out a detailed analysis in which several subtle effects, that could have been neglected with a smaller statistics, had to be carefully taken into account.  Based only on simple, model-independent geometrical considerations, we found that the delayed X-rays seen in the rings were emitted by \\src\\ within the same time interval of the unscattered radiation detected between 6:10 and 6:56 UT of 2009 January 22 in hard X-rays. This short time period includes the phase of highest bursting activity ever observed from this source.   Interesting information on \\src\\ and on the dust clouds producing the X-ray rings could be derived from the analysis of the ring's spectra and intensity evolution. However, this analysis requires the knowledge of the interstellar dust properties affecting the X-ray scattering cross section. We therefore explored the consequences of using different dust models.  The energy released by the event that formed the X-ray scattering rings could be derived from the fit of the halo profiles, that provides at the same time its spectrum and fluence in the 2--5 keV band, and the source distance. Assuming a column density of 10$^{22}$ cm$^{-2}$ for the most distant cloud and a bremsstrahlung spectrum with $kT=30$--100 keV, a burst energy between 10$^{44}$ erg and 2$\\times$10$^{45}$ erg was derived (in the 1--100 keV range). These values are significantly larger than that  in the pulsating tail detected on 2009 January 22 at 6:48 UT ($E\\sim2.4\\times10^{43}d^2_{10~{\\rm kpc}}$ erg; \\citealt{mereghetti09}), that therefore cannot account for the rings. More likely candidates are two bright bursts, whose fluence could not be determined because they were so bright to saturate the detectors, and occurred at times compatible with the expansion rate measured for the rings. However, we cannot exclude that the event producing the scattering rings had a very steep spectrum and so dominated the emission at soft X-rays without showing up as an exceptional event at higher energies.    According to our estimate, the energy released by the event that generated the X-ray scattering rings is comparable to that emitted by the giant flares of SGR\\,0526--66 \\citep{mazets79} and SGR\\,1900+14 \\citep{hurley99}. These events were followed by  pulsating tails lasting several minutes, contrary to the case of \\src: the only pulsating tail detected by the continuous \\integral\\ ACS observation covering the time interval compatible with the origin of the X-ray rings lasted only 8 seconds \\citep{mereghetti09}. On the other hand, in case the X-ray scattering rings of \\src\\ were produced by a single short burst, it must have been exceptionally bright, since no other magnetar bursts without any pulsating tail and an emitted energy $>$10$^{44}$ erg have ever been detected. The intermediate flare emitted by SGR\\,1627--41 on 1998 June 18 \\citep{mazets99} was the former record holder, with $E\\sim10^{43}$ erg (assuming a distance of 11 kpc). We caution that large systematic uncertainties affect our estimate of the emitted energy: the total scattering cross-section and source distance depend on the dust model, the hydrogen column density and dust-to-gas ratio in the dust clouds are only poorly constrained, and the broad band spectrum is inferred from a narrow band X-ray spectrum and in analogy to similar events.   By using several dust models discussed in the literature we showed that only about half of them  provide an adequate fit to the three halo profiles at different energies. The best-fitting model (BARE-GR-B in \\citealt{zda04})  yields for \\src\\ a distance of 4 kpc, consistent with the value derived from  its association with the SNR  G\\,327.24--0.13 \\citep{gelfand07}. This model also gives distances for the three dust clouds that are consistent with CO observations in this direction.  Our best fitting dust model is the same that provided  the best fits to the halo of GX5--1 \\citep{smith06}, but also other models, that were compatible with the analysis of X-ray halos of other X-ray binaries (see, e.g., \\citealt{smith08}), gave good fits to our data. They resulted in distances for \\src\\ in the 4.8--5.4 kpc range. Other dust models giving relatively good fits imply distances of 6--8 kpc,  but require that the highest peak seen in the CO emission corresponds to the cloud responsible for the second X-ray ring. This is not the brightest ring, and thus these scenarios would imply  that the strength of the CO line is not proportional to the amount of dust that efficiently scatters X-rays.  We finally note  that, if the X-ray source and dust distance are not independently constrained, the quality of the fits of the X-ray halo profiles is mostly sensitive to the grain size distribution in the range 0.01--0.03 $\\mu$m, while the source distance is mainly determined by the largest grains. These grains, with size $>$0.1 $\\mu$m, are those less constrained by diagnostics at longer wavelengths. Therefore, we cannot exclude that the actual dust properties may be different from all the models we tested, providing a good fit to the X-ray rings, but for a significantly different distance of \\src."
    },
    "0911/0911.3587_arXiv.txt": {
        "abstract": "A feasibility study was carried out at the Astronomical Observatory of the Autonomous Region of the Aosta Valley demonstrating that it is a well-poised site to conduct an upcoming observing campaign aimed at detecting small-size (R$\\leq$R{\\tiny Neptune\\/}) transiting planets around nearby cool M dwarf stars. Three known transiting planet systems were monitored from May to August 2009 with a 25 cm f/3.8 Maksutov telescope. We reached seeing-independent, best-case photometric RMS less than 0.003 mag for stars with V$\\leq$13, with a median RMS of 0.006 mag for the whole observing period. ",
        "introduction": "The Astronomical Observatory of the Autonomous Region of the Aosta Valley (OAVdA; 45.7895\\deg  N, 7.47833\\deg E) has been identified as a potential site for hosting a photometric transit search for low-mass, small-size planets around nearby cool M dwarf stars. We carried out a study to gauge the near-term and long-term photometric precision achievable, observing known transiting systems in a range of transit depths under variable seeing and sky transparency conditions. ",
        "conclusions": "The feasibility study demonstrated that the OAVdA is a promising site for a long-term photometric survey aimed to detect transiting low-mass, small-size planets. We are planning to set up a system of five 40 cm identical telescopes for an high-precision observing campaign focused on several hundreds nearby M dwarfs, improving the longitudinal coverage of similar programs already ongoing, such as the Arizona-based MEarth project \\Citep{Dam-Nut08}."
    },
    "0911/0911.5197_arXiv.txt": {
        "abstract": "We describe the current status and the prospect for the development of monolithic Ge:Ga array detector for SAFARI. Our goal is to develop a $64\\times 64$ array for the 45 -- 110 $\\mu$m band, on the basis of existing technologies to make $3\\times 20$ monolithic arrays for the AKARI satellite. For the AKARI detector we have achieved a responsivity of 10 A/W and a read-out noise limited NEP (noise equivalent power) of $10^{-17}$ W/$\\sqrt{\\rm Hz}$. We plan to develop the detector for SAFARI with technical improvements; significantly reduced read-out noise with newly developed cold read-out electronics, mitigated spectral fringes as well as optical cross-talks with a multi-layer antireflection coat. Since most of the elemental technologies to fabricate the detector are flight-proven, high technical readiness levels (TRLs) should be achieved for fabricating the detector with the above mentioned technical demonstrations. We demonstrate some of these elemental technologies showing results of measurements for test coatings and prototype arrays. ",
        "introduction": "% \\label{doiy2_sec:monolithic} Ge:Ga photoconductor detector is a highly sensitive infrared detector that has been used for many astronomical applications including IRAS \\citep{doiy2:IRAS}, IRTS (\\citeauthor{doiy2:hiromoto92}, \\citeyear{doiy2:hiromoto92}; \\citeauthor{doiy2:shibai94}, \\citeyear{doiy2:shibai94}), ISO (\\citeauthor{doiy2:ISOPHOT}, \\citeyear{doiy2:ISOPHOT}; \\citeauthor{doiy2:ISOLWS}, \\citeyear{doiy2:ISOLWS}), AKARI (\\citeauthor{doiy2:doi02}, \\citeyear{doiy2:doi02}; \\citeauthor{doiy2:kawada07}, \\citeyear{doiy2:kawada07}; \\citeauthor{doiy2:shirahata09}, \\citeyear{doiy2:shirahata09}), Spitzer \\citep{doiy2:MIPS}, and Herschel \\citep{doiy2:PACS}.  For the AKARI satellite, we have developed the world's first two-dimensional monolithic Ge:Ga photoconductor array, achieving a large array format of $20\\times 3$ pixels with the pixel size of $500\\times 500$ $\\mu$m$^2$ (Figure~\\ref{doiy2_fig:AKARI_SW}). \\begin{figure}[tb] \\begin{center} \\includegraphics[width=9 cm]{doiy2_AKARI_SW} \\end{center} \\caption{Design of the $20\\times 3$ pixel Ge:Ga monolithic array module onboard AKARI. Two sets of array are used for two photometric bands. Top-left figure shows a sectional view of the detector (see text for the detailed descriptions).} \\label{doiy2_fig:AKARI_SW} \\end{figure} The Ge:Ga photoconductor is produced by doping a 0.5-mm-thick Ge wafer with Ga concentration of $1.6\\times 10^{14}$ cm$^{-3}$. For longitudinal configuration (top illumination), a transparent contact with a grid-shape metal electrode is fabricated on the top surface of the Ge:Ga wafer (also see $\\S$\\ref{doiy2_sec:transparent}). Ohmic contact of the electrode is formed by B$^+$ ion implantation on top and bottom surfaces of the wafer. Bias voltage is applied to the whole array in common from the top contact. The bottom contact of each array pixel is directly connected to each CRE element with a 100-$\\mu$m-diameter indium ball. The array pixels are electrically separated each other by 50-$\\mu$m-wide, 30-$\\mu$m-deep grid-shape ditches on the back surface.  The infrared light incident from the top surface travels between the top and bottom surfaces back and forth, due to high reflectivity of the non-AR-coated top Ge surface and the metal back surface. Thanks to the multiple reflections, long absorption length is obtained without any external optics like an optical cavity, and the resultant quantum efficiency is approximately 0.4, which corresponds to the responsivity at an optimum bias of $\\sim 20$ A/W. However, the multiple reflections cause interference fringes and result in the presence of periodic peaks in the spectral response, which is a serious drawback of this device especially for spectroscopic applications. Mitigating this fringe problem by the anti-reflection (AR) coating is an issue of new array development (see $\\S$\\ref{doiy2_sec:AR}).  The arrays are attached to a fan-out board made of quartz substrate, and the fan-out board is mounted on a rigid frame made of gold-coated KOVAR. These materials are chosen to minimize the thermal stress to the detector. The KOVAR frame is mounted on the aluminium housing with thermal isolation legs made of Vespel SP-1 and is thermally connected to the helium tank by a copper strap. The detector temperature is controlled to $\\sim 2.0$ K by self-heat dissipation of the CRE ($\\sim 1$ mW) and additional heater power. The total mass of the detector module is $\\sim 300$ g. The flight model detector module passes the vibration test for simulating the mechanical stress at the launch. ",
        "conclusions": ""
    },
    "0911/0911.0067_arXiv.txt": {
        "abstract": "The magnetorotational instability (MRI) plays a crucial role for cosmic structure formation by enabling turbulence in Keplerian disks which would be otherwise hydrodynamically stable. With particular focus on MRI experiments with liquid metals, which have small magnetic Prandtl numbers, it has been shown that the helical version of this instability (HMRI) has a scaling behaviour that is quite different from that of the standard MRI (SMRI). We discuss the relation of HMRI to SMRI  by exploring various parameter dependencies. We identify the mechanism of transfer of instability     % between modes through a spectral exceptional point that explains both the transition from a stationary instability (SMRI)              % to an unstable travelling wave (HMRI) and the excitation                 % of HMRI in the inductionless limit. For certain parameter regions we find new islands of the HMRI. ",
        "introduction": "The magnetorotational instability (MRI) \\citep{B09} is considered as the main candidate to solve the long-standing puzzle of how stars and black holes are fed by the accretion disks surrounding them. The central problem is that these accretion disks typically rotate according to Kepler's law, $\\Omega(r) \\sim r^{-3/2}$, which results in an angular momentum $r^2 \\Omega(r) \\sim r^{1/2}$. Hence, they fulfill Rayleigh's criterion stating that rotating flows with radially increasing angular momentum are hydrodynamically stable, at least in the linear sense. Such stable, non-turbulent disks would not allow the outward directed angular momentum transport that is necessary for the infalling disk matter to accrete into the central object.  In their seminal paper of 1991 \\cite{BH91} Balbus and Hawley had highlighted the key role of the magnetorotational instability (MRI) in explaining turbulence and angular momentum transport in accretion disks around stars and black holes. They had shown that a weak, externally applied magnetic field serves only as a trigger for the instability that actually taps into the rotational energy of the flow. This is quite in contrast to current-induced instabilities, e.g. the Tayler instability \\cite{T73}, which draw their energy (at least partly) from the electric currents in the fluid.  Soon after the paper by Balbus and Hawley it became clear that the principle mechanism of the MRI had already been revealed three decades earlier by Velikhov \\cite{V59} and Chandrasekhar \\cite{C60}. Actually, they had investigated the destabilizing action of an external magnetic field for the classical Taylor-Couette (TC) flow between two concentric, rotating cylinders rather than for Keplerian rotation profile. This is, however, not a crucial difference since a TC flow can be made very close to  a Keplerian one simply by adjusting the ratio of rotation rates of the inner and the outer cylinder.  The MRI in flows between rotating walls has attracted             % renewed interest during the last decade, mainly motivated by the increasing efforts to investigate MRI in the laboratory \\cite{AIP04,SGG08}.                                    % A first interesting experimental result was obtained in a spherical Couette flow of liquid sodium \\cite{SMTHDHAL04}. The authors                    % observed correlated modes of velocity and magnetic field perturbation in a parameter region which is quite typical for MRI. It                               % must be noted, however, that the                                  % background state in this spherical Couette experiments was already fully turbulent, so that the original goal that the MRI would destabilize an otherwise stable flow was not met. At Princeton University work is going on to identify MRI in a TC experiment with liquid gallium, and first encouraging results, including the observation of non-axisymmetric Magneto-Coriolis waves, have been obtained \\cite{N08, N09}.  Both experiments had been designed to investigate the standard version of MRI (SMRI) with only a vertical magnetic field being applied. In this case, the azimuthal magnetic field (which is an essential ingredient of the MRI mode) must be  produced from the vertical field by induction effects, which are proportional to the magnetic Reynolds number (${\\rm Rm}$) of the flow. ${\\rm Rm}$, in turn, is proportional to the hydrodynamic Reynolds number according to ${\\rm Rm}={\\rm Pm} {\\rm Re}$, where the magnetic Prandtl number ${\\rm Pm}=\\nu/\\eta$ is the ratio of viscosity $\\nu$ to magnetic diffusivity $\\eta=1/\\mu_0 \\sigma$. For liquid metals ${\\rm Pm}$ is typically in the range $10^{-6}...10^{-5}$. Therefore, in order to achieve ${\\rm Rm} \\sim 1$, we need ${\\rm Re}\\sim 10^5...10^6$, and wall-constrained flows (in contrast to wall-free Keplerian flows) with such high ${\\rm Re}$ are usually turbulent, whatever the linear stability analysis might tell (see, however, \\cite{JINATURE}).                                % This is the point which makes SMRI experiments, and their interpretation, so cumbersome.  One might ask, however, why not to substitute the induction of the necessary azimuthal magnetic field component of the MRI mode by simply externally applying this component               % as a part of the base configuration. Indeed, it was shown \\cite{HR05,RUEDIGER2005}                   % that the resulting ''helical MRI'' (HMRI), as we now call it, is then possible at                      % far smaller Reynolds numbers and magnetic field amplitudes than SMRI, making HMRI an ideal playground for liquid metal experiments.  First experimental evidence for HMRI was obtained in 2006 at the liquid metal facility PROMISE ({\\it P\\/}otsdam {\\it RO\\/}ssendorf {\\it M\\/}agnetic {\\it I\\/}n{\\it S\\/}tability {\\it E\\/}xperiment) which is basically  a Taylor-Couette (TC) cell made of concentric rotating copper walls, filled with  GaInSn (a eutectic which is liquid at room temperatures). In \\cite{SGGRSSH06,RHSGGR06,SGGRSH07,SGGSRH08} it was shown that the HMRI travelling wave appears only in the predicted finite window of the magnetic field intensity, with a frequency of the travelling wave that was also in good accordance with numerical simulations. Results of a significantly improved experiment (PROMISE 2) with strongly reduced Ekman pumping at the end-caps were published recently \\cite{SGGSRH09,SGGHPRS09}.  The connection of SMRI and HMRI is presently under intense debate \\cite{LGHJ06,RH07,PGG07,LV07,LGJ07,S07,RS08,L09,PG09}. The first essential point to note here is that HMRI and SMRI are               % connected. Indeed, Fig. 1 in \\cite{HR05} shows that there is a continuous and monotonic transition from HMRI to SMRI when ${\\rm Re}$ and the magnetic field strength are increased simultaneously.  A second remarkable property of HMRI for small ${\\rm Pm}$ (which has been coined ``inductionless MRI''), was clearly worked out in \\cite{PGG07}. It is the apparent paradox that a magnetic field is able to trigger an instability although the total energy dissipation of the system is larger than without this field.  The relevance of HMRI for Keplerian flows has               % been seriously put into question in \\cite{LGHJ06}. Using a local WKB analysis in the small-gap approximation, the authors had shown that the HMRI works only for comparably steep rotation profiles (i.e. slightly above the Rayleigh line) and disappears for profiles as flat as the Keplerian one. This result has been confirmed by Lakhin and Velikhov \\cite{LV07} and R\\\"udiger and Schultz \\cite{RS08}.  However, this disappointing result was soon relativized in \\cite{RH07} by solving the global eigenvalue equation for HMRI with electrical boundary conditions. It turned out that HMRI re-appears again for Keplerian flows provided that at least one radial boundary is    % highly conducting. A similar discrepancy between local and global results is well known for the so-called stratorotational instability (SRI) \\cite{DMNRHZ05} for which the existence of reflecting boundaries appears necessary for the instability to work \\cite{U06}. This artificial demand is of course a much stronger argument against the working of SRI than the necessity  of one conducting boundary is for the working of HMRI: considering, i.e., the colder outer parts of accretion disks, then the inner part can indeed be considered as a good conductor \\cite{BALBUS2008}.                                        %  Another argument that has been put forward against the relevance of HMRI for thin accretion disks is the necessity for a large ratio of toroidal to poloidal magnetic fields \\cite{L08}.  A further complication for applying HMRI to the real world     % is the fact that it appears in form of a travelling wave. The crucial point here is that monochromatic waves are typically not able                       % to fulfill the axial boundary conditions at the ends of the considered region. To fulfill them,  one has to consider wave packets. Only wave packets with vanishing group velocity will remain in the finite length system. Typically, the onset of this {\\it absolute instability}, characterized by a zero growth rate {\\it and} a zero group velocity, is harder to achieve  than the {\\it convective instability} of a monochromatic wave with zero growth rate. A comprehensive analysis of the relation of convective and        % absolute instability for HMRI can be found in \\cite{PG09}. From the extrapolation of the results of this paper it seems that Keplerian rotation profiles (with conducting boundaries)          % are indeed absolutely HMRI-unstable, but a final solution to this puzzle is still elusive.   In the present paper, we step back from those                 % important consideration of absolute and global instabilities and focus again on the local WKB method by considering the dispersion relation of MRI which had been                           % derived and analyzed in \\cite{LGHJ06,LV07,RS08}. In spite of these former investigations,                         % we feel that some points still need further clarification.      % This concerns a careful application of the                     % Bilharz stability criterion as well as some further parameter dependencies, in particular the dependence for small but finite magnetic Prandtl numbers. It also concerns the question in which sense the HMRI can be considered as a {\\it dissipation induced instability} which is quite common in many areas of physics \\cite{KM07,Ki07}.    To make the paper self-contained, we will start with a re-derivation of the dispersion relation in two forms which explicitly contain the relevant frequencies or the dimensionless parameters, respectively.  Then we will study the peculiar relation of SMRI and HMRI. As a main result of this paper we will describe in detail the mechanism of transition from SMRI to HMRI through a spectral exceptional point which appears at finite but small ${\\rm Pm}$. This provides  a natural explanation for the         % continuous and monotonic connection between SMRI          % (a destabilized slow magneto-Coriolis wave)                          % and HMRI (a weakly destabilized inertial oscillation).     % In addition to this, for high Reynolds numbers we will            % identify a second scenario for HMRI which leads to             % new islands of instability at small but finite values          % of $\\rm Pm$. ",
        "conclusions": "The helical magnetorotational instability is a more complicated phenomenon than the standard one. We found evidences that HMRI can be identified with the destabilization of an inertial                                                            % wave in contrast to SMRI that is a destabilized slow Magneto-Coriolis wave. We established two scenarios of transition from SMRI to HMRI: the first one is accompanied by the origination of a spectral exceptional point and a transfer of instability between modes while in the second scenario two independent eigenvalue branches become unstable. We distinguish between the essential HMRI that is characterized by small magnetic Prandtl numbers at which SMRI is not possible, smaller growth rates than SMRI, and by non-zero frequencies and the helically modified SMRI which is caused by a small perturbation of the unstable real eigenvalue branch and is thus characterized by high growth rates, small frequency and relatively high magnetic Prandtl numbers within the usual range of SMRI. With the use of the Bilharz stability criterion we established explicit expressions for the stability boundary and proved rigorously the bounds on the critical Rossby number for HMRI in the inductionless limit $(\\rm Pm=0)$. Nevertheless, we revealed that for $\\rm Pm \\ne 0$ these bounds can be easily exceeded---an indicator in favor of the HMRI for small negative Rossby numbers. Finally, we found that for small negative Rossby numbers the essential HMRI forms separated islands that can coexist simultaneously in the $({\\rm Pm},\\beta^*)$-plane."
    },
    "0911/0911.0298_arXiv.txt": {
        "abstract": "Astronomical and cosmological knowledge up to the dawn of modern science was profoundly embedded in myth, religion and superstition. Many of these inventions of the human mind remain today stored in different supports: medieval engravings, illuminated manuscripts, and of course also in old and rare books. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0911/0911.1125.txt": {
        "abstract": "We study the development of high energy ($E_{\\rm in} \\le 1$~TeV) cascades produced by a primary electron of energy $E_{\\rm in}$ injected into the intergalactic medium (IGM). To this aim we have developed the new code MEDEA (Monte Carlo Energy DEposition Analysis) which includes Bremsstrahlung and Inverse Compton (IC) processes, along with H/He collisional ionizations and excitations, and electron-electron collisions. The cascade energy partition into heating, excitations and ionizations depends primarily on the IGM ionized fraction, $x_{\\rm e}$, but also on redshift, $z$, due to IC on CMB photons. While Bremsstrahlung is unimportant under most conditions, IC becomes largely dominant at energies $E_{\\rm in} \\geq  1$ MeV. The main effect of IC at injection energies $E_{\\rm in} \\leq  100$ MeV is a significant boost of the fraction of energy converted into low energy photons ($h \\nu < 10.2$ eV) which do not further interact with the IGM. For energies $E_{\\rm in} \\geq  1$ GeV CMB photons are preferentially upscattered within the X-ray spectrum ($h \\nu > 10^4$ eV) and can free stream to the observer. Complete tables of the fractional energy depositions as a function of redshift, $E_{\\rm in}$ and ionized fraction are given. Our results can be used in many astrophysical contexts, with an obvious application related to the study of decaying/annihilating Dark Matter (DM) candidates in the high-$z$ Universe. ",
        "introduction": "High energy particles can be produced by several astrophysical and cosmological sources through different acceleration processes. The energy stored in relativistic particles might be a non-negligible fraction of the total energy of these systems and therefore an obvious question arises about how this energy is eventually thermalized and transferred to the surrounding environment. In spite of its importance, this question has received only a relatively limited attention. Previous works have considered non-relativistic initial energies of up to 10 keV (e.g. Shull 1979, Shull \\& van Steenberg 1985, Dalgarno et al. 1999; Vald\\'{e}s \\& Ferrara 2008, hereafter VF08; Furlanetto \\& Johnson Stoever 2009). In this energy range processes such as free-free emission with charged particles and Inverse Compton (IC) with a diffuse distribution of photons can be safely neglected.  However for many astrophysical applications it is necessary to deal with higher energy particles: extrapolating the results and fitting formulae presented in the aforementioned works can lead to substantially incorrect results. Particles can be accelerated to relativistic energies in a number of astrophysical systems, leading to the emission of synchrotron and Compton radiation observable in a broad range of energy bands. Active Galactic Nuclei, Stellar flares, Gamma Ray Bursts, Pulsar Wind Nebulae, Supernova Remnants are believed to house a population of shock accelerated electrons (see e.g. MacKinnon \\& Mallik 2009; Venter \\& de Jager 2008; Resmi \\& Bhattacharya 2008). It is then important to consider extensions of these works to compute at best the evolution of the energy cascade of relativistic electrons of energy $E_{\\rm in}$ into a partially ionized gas under realistic cosmological conditions including the presence of CMB photons.  Radio observations of galaxy clusters have revealed the existence of the so-called radio \"relics\" which can be explained with the presence in the intracluster medium of relativistic electrons accelerated by shocks due to cluster merging. This can give rise to X-ray emission through IC of CMB photons, as recent observations of the satellites Beppo-SAX and RXTE seem to indicate (see e.g. Rephaeli 1977; Fusco-Femiano, Landi, \\& Orlandini 2007; Rephaeli et al. 2008; Ferrari 2009).  Another astrophysical source of relativistic electrons could come from the decay or annihilation of DM particles. If indeed, as many theoretical models predict, this elusive matter component of the Universe injects relativistic electrons and positrons into the IGM and the consequent inverse Compton scattering with the CMB photons could generate a distortion of the black body spectrum by Sunyaev-Zeldovich (SZ) effect (see e.g. Lavalle, Boehm, \\& Barthes 2009; Colafrancesco \\& Mele 2001; Colafrancesco 2004; Colafrancesco, Profumo, \\& Ullio 2006). In the near future the recently launched satellite Planck (The Planck collaboration, 2006) and the ALMA array (Wootten \\& Thomson 2009) could detect the SZ deviation induced by DM decays/annihilations.  In the past few years a large number of works has investigated the effects and detectability of DM decays/annihilations into the high redshift IGM via observations of the redshifted 21~cm hyperfine triplet-singlet level transition of the ground state of neutral hydrogen (Shchekinov \\& Vasiliev 2006; Furlanetto, Oh, \\& Pierpaoli 2006;  Vald\\'{e}s et al. 2007; Natarajan \\& Schwarz 2009). The interest for such kind of studies is generated by the present or planned construction of large radio interferometers such as the Low frequency Array (LOFAR, http://www.lofar.org/), the 21 Centimeter Array (21CMA, http://21cma.bao.ac.cn/), the Murchison Widefield Array (MWA, http://www.mwatelescope.org/) and the Square Kilometer Array (SKA, http://www.skatelescope.org/). These instruments will in fact open a new observational window by the detection of the HI 21 cm line at $z>6$ (e.g. Peterson, Pen \\& Wu 2005; Bowman, Morales, \\& Hewitt 2005; Kassim et al. 2004; Wyithe, Loeb, \\& Barnes 2005).  To understand if observations of the redshifted HI 21 cm line can help constrain DM it is crucial to follow in detail the energy cascade from energetic primary photons or electrons up to energies much higher than previously studied (Shull \\& van Steenberg 1985, VF08) to include those DM candidates that can produce relativistic electrons. In our calculations we choose an initial electron energy 1 MeV $< E_{\\rm in} < 1$ TeV. The details of the energy degradation of electrons in this energy range are of great interest for many of the aforementioned astrophysical sources and in particular for the study of the effects of DM decays/annihilations, since recent data from PAMELA, ATIC, FERMI-LAT and HESS experiments have shown that there are electron-positron excesses in the cosmic ray energy spectrum which can be explained by annihilations of DM between 1 to 2 TeV (see e.g. Cirelli, Franceschini, \\& Strumia 2008; Liu et al. 2009; Berg et al. 2009; Hooper \\& Tait 2009; Chen, Takahashi \\& Yanagida 2009; He 2009 for a brief review with an extensive list of citations). Recently Slatyer et al. (2009) computed the energy injection rates for some annihilating DM candidates.  The rest of the paper is organized as follows. In Sec. 2 we describe the relevant physical process and their implementation into the multi-purpose code MEDEA (Monte Carlo Energy DEposition Analysis) we have developed to follow the energy cascade. In Sec. 3 we present the results and analyze their dependence on redshift, initial electron energy and redshift. In Sec. 4 we discuss the implications and draw some conclusions. The complete numerical outputs of MEDEA can be found in tabulated form in Appendix A. ",
        "conclusions": "We have introduced our code MEDEA, based on a Monte Carlo scheme that makes possible to follow the fate of electrons of energies up to 1 TeV in their secondary energy cascade. Our results represent a substantial generalization of previous works that considered exclusively non-relativistic electrons.  To perform our calculation we implemented in the code a large number of interactions such as collisional ionizations of H, He, HeI; collisional excitations of H, He; electron-electron Coulomb scattering; free-free interactions of electrons with protons; IC with CMB photons; direct collisional excitations to the $2s$ level of HI; indirect cascades from $n\\geq 3$ states of HI through the $2s$ level; recombinations. The energy range of the primary electron is $1$ MeV $<E_{\\rm in}<10$ MeV, the ionized fraction considered is $10^{-4}<x_{\\rm e}<0.99$ and redshift spans $10<z<50$.  The results presented here and summarized in numerical form in the Appendix Tables, can be used for many astrophysical applications such as cluster radio relics, DM decays/annihilations, Active Galactic Nuclei, Stellar flares, Gamma Ray Bursts, Pulsar Wind Nebulae, Supernova Remnants and, more generally, whenever it is necessary to deal with the interaction of energetic and/or relativistic particles with the surrounding thermal gas.  We will present new results and updates at the URL {\\small http://www.arcetri.astro.it/twiki/bin/view/DAVID/MedeaCode} periodically."
    },
    "0911/0911.5108_arXiv.txt": {
        "abstract": "As a test of the standard inflationary cosmology, which generically predicts nearly scale-invariant spectrum of primordial curvature fluctuations, we  perform Markov-Chain Monte-Carlo analysis to search for possible  modulations in the power spectrum and determine its shape together with the cosmological parameters using cosmic microwave background radiation data.  By incorporating various three-parameter features on the simple power-law spectrum, we find an oscillatory modulation localized around the comoving wavenumber $k\\simeq 0.009\\mathrm{Mpc}^{-1}$ at 99.995\\% confidence level which improves the log-likelihood as much as $-\\Delta 2\\ln {\\cal L}\\equiv \\Delta \\chi^2_{\\rm eff}=-22$. This feature can be detected even if we use only the cross correlation between the temperature and the E-mode polarization anisotropies. ",
        "introduction": "In the standard inflationary cosmology \\citep{guth,sato,staro,infreview} the observed large-scale structures and the anisotropies in the cosmic microwave background radiation (CMB) originate in tiny quantum fluctuations generated during the inflationary expansion stage in the very early Universe \\citep{yuragi,yuragi2,yuragi3}. Conventionally, the power spectrum of primordial curvature fluctuations, $P(k)\\equiv \\langle |\\calrk|^2\\rangle$, has been assumed to follow a simple functional form such as a power-law  $A(k)\\equiv \\frac{4\\pi k^3}{(2\\pi)^3}P(k)=A(k/k_0)^{n_s-1}$ which is quantified by the amplitude $A$ and the spectral index $n_s$, and the previous statistical analyses of the power spectrum have mostly focused on these two parameters (and at best the scale dependence of $n_s$, the running, which may be important to distinguish the inflation model\\citep{Kawasaki:2003zv}).  It is true that the simplest class of inflation models predicts a  power-law spectrum with $n_s$ close to unity \\citep{yuragi,yuragi2,yuragi3}, but we may not be able to identify the correct theory of the early Universe if we restrict our parameter space from the beginning.  From the viewpoint of observational cosmology, the shape of the primordial fluctuation spectrum should be determined purely from the observational data without any theoretical prejudices.  Along this line of thought, several Markov-Chain Monte-Carlo (MCMC) analyses have been performed to search for possible deviation from the power law, but none has detected statistically much significant modulations so far \\citep{4,4.2,5,7,WMAP3COSMO,2009PhRvD..80h3002I}. To be more quantitative, the presence of a nontrivial feature may be decided in terms of Akaike's Information Criterion (AIC) \\citep{akaike} which asserts that introduction of an additional fitting parameter is justified if and only if $\\chi^2_\\mathrm{eff}$ improves more than $-2$ with it. Most of the previous analyses resulted in the improvement of $\\Delta\\chi^2_\\mathrm{eff}$ much less than that required by AIC, and others reached it only at a marginal level, so that it was hardly possible to claim the presence of a feature. It should be emphasized, however, that those analyses did not have sufficient resolution to detect spectral fine structure,  because they tried to fit the primordial power spectrum in a broad range of scales in terms of a limited number of degrees of freedom using sparse sampling or some specific theoretical models. This lack of resolution  was  inevitable because exploring large dimensional parameter space is extremely time-consuming even with the help of MCMC analysis.   Recently, on the other hand, two of the present authors (RN and JY) performed reconstruction of primordial power spectrum from the angular power spectrum of the CMB temperature fluctuations, $C_\\ell^{TT}$, of the five-year WMAP data (hereafter WMAP5)\\citep{WMAP5a,WMAP5} using two different non-parametric methods, namely, the cosmic inversion method \\citep{MSY02,MSY03,KSY04,KSY05, KMSY04,2008PhRvD..78l3002N} and the maximum-likelihood reconstruction method\\citep{Tocchini,silk,2009PhRvD..79d3010N}.  They have probed the primordial power spectrum with finer resolution than aforementioned MCMC analyses and found an anomalously large % deviation from the best-fit power-law primordial spectrum of WMAP5 around the wavenumber $k\\simeq 0.009\\mathrm{Mpc}^{-1}$ or $kd\\simeq 124$ where $d=1.43\\times 10^4$Mpc is the distance to the last scattering surface. This scale corresponds to the multipole $\\ell \\simeq 120$ and the length scale $r\\simeq 710$Mpc.  In their reconstruction procedure, however, they had to fix values of the cosmological parameters and adopt an {\\it ad hoc} smoothing prior to ensure sensible reconstruction.  The latter tends to discard a large part of the information on the fine structure which may have affected the significance of the claimed spectral feature itself. Furthermore although the reconstructed power spectrum exhibited oscillatory features with both peaks and dips, the result of decomposition of the continuous reconstructed curve to statistically independent band-powers, which is necessary to discuss their statistical significance, revealed only a $3.3\\sigma$ peak but no dips \\citep{2009PhRvD..79d3010N}.  The purpose of this paper is to focus on the fine structures of the spectral shape.  We use a forward analysis implemented by MCMC simulation which circumvents the above-mentioned difficulties peculiar to the inverse mapping approach. We also overcome the lack of resolution, which was a very serious problem of previous MCMC analyses as mentioned above, by searching for and concentrating on the most prominent localized feature imprinted on the otherwise power-law spectrum. ",
        "conclusions": "In this paper we have found a spectral feature with a 4$\\sigma$ deviation from the best-fit power-law primordial power spectrum in a narrow range around $k \\simeq 0.009{\\rm Mpc}^{-1}$, whose existence is suggested separately by both TT data and TE data from five-year WMAP data.  This feature can be further tested and hopefully confirmed by forthcoming EE data from the Planck mission.  As emphasized above, the detected feature consists not only of a peak but also a dip.  Since we cannot lower the amplitude of the power spectrum by any sources uncorrelated with the preexistent fluctuations, this means that the feature was created together with the bulk of the power-law spectrum, suggesting some nontrivial phenomenon during inflation.  At the moment, while a number of inflation models have been proposed to produce spectral modulation, we are not aware of any model that can realize a spectrum with such a localized feature.  Hence efforts should be made toward construction of models that can account for the detected fine structures.  Currently inflation models are observationally constrained in terms of a small number of discrete parameters such as $n_s$, $dn_s/d\\ln k$, the tensor-to-scalar ratio $r$, and a nonlinearity parameter $f_{NL}$. It is difficult to determine the underlying particle-physics model with these parameters alone, because there are huge degeneracies. If our finding is further confirmed, it can certainly serve as an important clue to single out the correct model. On the other hand, whether this feature is a result of some nontrivial physics or a realization of an extremely rare event, we must take it into account in cosmological parameter estimation in the analysis of forthcoming higher precision data."
    },
    "0911/0911.2202_arXiv.txt": {
        "abstract": "Gamma-ray bursts have the potential to produce the particle energies (up to $10^{21}$\\,eV) and the energy budget ($10^{44}\\, \\rm{erg\\, yr^{-1}\\, Mpc^{-3}}$) to accommodate the spectrum of the highest energy cosmic rays; on the other hand, there is no observational evidence yet that they accelerate hadrons. Fermi recently observed two bursts that exhibit a power-law high-energy extension of the typical (Band) spectrum that extends to $\\sim 30$ GeV. On the basis of fireball phenomenology we argue that they, along with GRB941017 observed by EGRET in 1994, show indirect evidence for considerable baryon loading. Since the detection of neutrinos is the only unambiguous way to establish that GRBs accelerate cosmic rays, we use two methods to estimate the neutrino flux produced when the baryons interact with fireball photons to produce charged pions and neutrinos. While the number of events expected from the Fermi bursts is small, we conclude that an event like GRB941017 will be detected by IceCube if gamma-ray bursts are indeed the sources of the cosmic rays. ",
        "introduction": "}  The sources of the extragalactic cosmic rays with energies in excess of $\\sim$\\,3$\\times$10$^{18}$\\,eV remain a mystery, but one of the best motivated candidates is gamma-ray bursts (GRBs). Large cosmic-ray energies can be achieved in the prompt phase of the GRB fireball where internal shocks have the potential to accelerate charged particles up to $\\sim$10\\,$^{21}$\\,eV (\\citealt{vietri_1995}, \\citealt{waxman_1995}). Additionally, the total energy density in the Universe of cosmic rays must be matched by a sufficiently high hadronic energy density in the GRB with $\\rho_{\\mbox{\\tiny{had,\\,GRB}}}\\approx\\rho_{\\mbox{\\tiny{CR}}}$. GRB observations identify synchrotron photons produced by the electrons accelerated in the fireball with an energy $\\epsilon_e\\mbox{E}_{\\mbox{\\tiny{TOT}}}\\sim\\,$10$^{53}$\\,ergs, where E$_{\\mbox{\\tiny{TOT}}}$ is the total energy released by the burst. GRB fireballs also carry energy $\\epsilon_B\\mbox{E}_{\\mbox{\\tiny{TOT}}}$ in the form of magnetic fields, and, if they are the sources of cosmic rays, energy $\\epsilon_p\\mbox{E}_{\\mbox{\\tiny{TOT}}}$ in protons. Assuming equipartition, $\\epsilon_B\\simeq\\,\\epsilon_e$ and $\\epsilon_e+\\epsilon_p+\\epsilon_B=1$.  GRBs emerge as credible sources for the ultra high-energy cosmic rays because their observed flux can be accommodated with an energy density in protons that is similar to that in electrons, or $\\epsilon_p\\,\\simeq\\,\\epsilon_e$. Recent estimates of the local rate of GRBs yield a maximum of $\\dot{n}_{0}\\!\\!\\sim\\!\\! 1$~Gpc$^{-3}$~yr$^{-1}$ assuming that GRBs follow the star formation rate. For a stronger evolution with redshift, the local rate can be as low as $0.05$~Gpc$^{-3}$~yr$^{-1}$. Using this result we estimate the electromagnetic energy density from GRBs to be in the range $\\rho_{\\mbox{\\tiny{em,\\,GRB}}}\\approx\\dot{\\mbox{n}}_{0}\\,\\epsilon_e\\,\\mbox{E}_{\\mbox{\\tiny{TOT}}}=\\mbox{5}\\times\\mbox{10}^{42}\\mbox{\\,-\\,}10^{44}\\,\\,\\mbox{ergs\\,Mpc}^{-3}\\,\\mbox{yr}^{-1}$. In order for GRBs to be the sources of cosmic rays, their hadronic energy density $\\rho_{had,GRB}$ needs to produce the observed cosmic-ray energy density $\\rho_{\\mbox{\\tiny{had,\\,GRB}}}=\\dot{\\mbox{n}}_{0}\\,\\epsilon_p\\,E_{\\mbox{\\tiny{TOT}}}\\approx\\,10^{44}\\,\\mbox{ergs\\,Mpc}^{-3}\\,\\mbox{yr}^{-1}$. We therefore conclude that $\\epsilon_p/\\epsilon_e\\approx\\,1\\mbox{\\,-\\,}20$. The extragalactic spectrum is actually expected to extend to energies below the knee, but remains unobserved due to the larger contribution of galactic cosmic rays, see e.g. \\citet{ahlers}. If this is the case, the average fraction of proton to electron energy needs to be larger.  Fermi recently observed two bursts, GRB090510 and GRB090902b, that show a statistically significant deviation from the typical GRB spectrum described by the Band function\\,\\citep{090510,090902b}. A flux of high energy events is detected that extends to energies $\\sim$30\\,GeV following a power-law spectrum (Table~\\ref{tab:latparms}). This is similar to the potentially much more luminous burst observed by EGRET in 1994, GRB941017\\,\\citep{egret94}. In this paper we will investigate the possibility that the high energy component results from $\\pi^{0}$-decays and is therefore a signature of proton acceleration.  \\begin{table}[h!b!p!] \\caption{Spectral parameters of Fermi GRBs with power-law components. The total fluence in gamma rays $F_{\\gamma}^{\\mbox{\\tiny{TOT}}}$ is the sum of the Band fluence $F_{\\gamma}^{\\mbox{\\tiny{B}}}$ and the power-law fluence $F_{\\gamma}^{\\mbox{\\tiny{HE}}}$. The fluence for GRB900510 is calculated over the energy range 10 keV--30 GeV, and the fluence for GRB09092b is calculated over the range 10 keV--10 GeV. The parameters for GRB941017 are tabulated in~\\citet{egret94}.} \\begin{center} \\begin{tabular*}{1.0\\columnwidth}{c|cc} & \\small{GRB090510} & \\small{GRB090902b} \\\\[3pt] \\hline \\hline \\small{z}  & \\small{0.903} & \\small{1.822} \\\\ $T_{90}$   & \\small{2.1\\,s} & \\small{21.9\\,s}  \\\\ $F_{\\gamma}^{\\mbox{\\tiny{TOT}}}$ & \\small{$5.02\\times10^{-5}$\\,ergs\\,cm$^{-2}$} & \\small{$4.36\\times10^{-4}$\\,ergs\\,cm$^{-2}$}  \\\\ $F_{\\gamma}^{\\mbox{\\tiny{HE}}}$ & \\small{$1.84\\times10^{-5}$\\,ergs\\,cm$^{-2}$} & \\small{$1.05\\times10^{-4}$\\,ergs\\,cm$^{-2}$}  \\\\ $\\alpha_{\\gamma}$     & \\small{0.58} & \\small{0.61}  \\\\ $\\beta_{\\gamma}$      & \\small{2.83} & \\small{3.80}  \\\\ $E_{0}$  & \\small{2771\\,keV} & \\small{726\\,keV}  \\\\ $\\Gamma_{\\gamma}$     & \\small{1.62} & \\small{1.93}  \\\\ $E_{\\mbox{\\tiny{MAX}}}$ & \\small{30.53\\,GeV} & \\small{33.40\\,GeV} \\\\ $\\Gamma$     & \\small{1260} & \\small{1000}  \\\\ \\end{tabular*} \\end{center} \\label{tab:latparms} \\end{table}   The main contribution to the initial opacity of the fireball comes from the annihilation of photons into $e^{\\pm}$ pairs. The Fermi observation of a non-thermal spectrum up to an energy $E_{\\mbox{\\tiny{MAX}}}$ of tens of GeV can be used to constrain the minimum bulk Lorentz factor $\\Gamma_{\\mbox{\\tiny{MIN}}}$ required to make the source optically thin at the time of the gamma-ray display. For all photons with energy $E\\leq E_{\\mbox{\\tiny{MAX}}}$ the condition $\\tau_{\\gamma\\gamma}(E)<1$ must be fulfilled where $\\tau_{\\gamma\\gamma}$ is the opacity. The observation of photons with energies of tens of GeV requires highly relativistic outflows with $\\Gamma \\simeq 10^3$. Because, on the other hand, the observed energy flux of order $10^{-4}$\\,erg/s is typical of an average burst, the large boost factor implies that the photon density in the rest frame of the burst is low. This is a strong effect as the photon density is suppressed by $\\Gamma^{-4}$. From $\\tau_{\\gamma\\gamma}(E)=1$, we can determine $\\epsilon_{e}$ by finding the electromagnetic energy over the volume of the fireball as a fraction of the total GRB energy: \\begin{eqnarray} \\label{eq:epse} \\epsilon_{e}\\!\\!&\\approx&\\!\\!1-5 \\times 10^{-2} \\left( \\frac{\\Gamma}{300}\\right)^{4} \\left( \\frac{\\delta t}{10 \\,\\rm{ms}}\\right) \\left( \\frac{E_{\\gamma}}{1 \\,\\rm{MeV}}\\right)\\nonumber\\\\ & & \\times \\left( \\frac{T_{90}}{100\\, \\rm{s}}\\right) \\left( \\frac{10^{53}\\, \\rm{ergs}}{E_{\\mbox{\\tiny TOT}}}\\right) \\end{eqnarray} where $\\Gamma$ is the bulk Lorentz factor of the fireball, $\\delta t$ is the variability timescale, $E_{\\gamma}$ is the characteristic gamma-ray energy (which we take to be the peak energy of the event), $T_{90}$ is the duration of the burst, and $E_{\\mbox{\\tiny{TOT}}}$ is the total energy of the burst. We therefore estimate $\\epsilon_{e}$ for GRB090510 to be $\\sim0.05$ and for GRB090902b to be $\\sim0.02$. From the low values of $\\epsilon_e$ thus obtained we must conclude that protons dominate the fireball in order to accommodate the total energy of a solar mass generated by the primary engine.  In this letter we will first discuss the properties of the bursts. We subsequently compute the neutrino flux inevitably produced when the protons interact with fireball photons. Their observation would provide incontrovertible evidence for the pionic origin of the additional high energy component in the burst and support the speculation that GRB are the sources of the highest energy cosmic rays.  Is a kilometer-scale neutrino telescope such as IceCube sensitive enough to shed light on these questions? High energy neutrinos are produced in the fireball when protons produce pions in interactions with the low-energy photon field: \\begin{equation} p\\,\\gamma\\rightarrow \\Delta^{+}\\rightarrow \\left\\{ \\begin{array} {lll}n\\,\\pi^{+}&&\\mbox{1/3 of the cases}\\\\ p\\,\\pi^{0}&&\\mbox{2/3 of the cases} \\end{array} \\right.\\, \\end{equation} The neutral pions decay as $\\pi^{0}\\rightarrow 2\\gamma$, and the charged pions decay as $\\pi^{+}\\rightarrow\\mu^{+}\\,\\nu_{\\mu} \\rightarrow e^{+}\\,\\nu_{e}\\,{\\nu}_{\\mu}\\,\\nu_{\\mu}.$ Here, a single neutrino carries approximately $1/4$th of the $\\pi^{+}$ energy and a photon carries $1/2$ of the $\\pi^{0}$ energy. The calculation of the neutrino flux \\citep{WB} has been performed in detail for the EGRET bursts~\\citep{guetta, becker} with the following results: whereas an average burst produces only $\\sim 10^{-2}$ neutrinos, bursts that are unusually energetic or nearby produce an observable flux in a kilometer-scale neutrino telescope of order 10 events per year. We suggest that the power-law high-energy spectral feature identifies such bursts.  We will compute the neutrino fluxes expected in IceCube using two methods: the standard fireball model, and the bolometric method which relates the energy in neutrinos from the decay of charged pions to the observed photon energy assuming that it is of pionic origin. We will conclude that the neutrino rates from all 3 bursts are predicted to be larger than the average, but that a burst like GRB941017 extending to GeV energy will be observed by IceCube. We remind the reader that IceCube observes cosmic neutrinos in a background of neutrinos produced in the atmosphere. Given that neutrinos of GRB origin are relatively energetic and that the direction and time of the events can be correlated to satellite alerts, the atmospheric background is suppressed and a single neutrino may represent a conclusive observation. ",
        "conclusions": "Three gamma-ray bursts with statistically significant high-energy power-law spectral components have been detected thus far. The observed characteristics of the bursts indicate (Eq.~\\ref{eq:epse}) that it is likely that a large fraction of their total energy is hadronic rather than electromagnetic. As a result, we assume that these bursts are accelerating protons and calculate the associated fluxes of neutrinos. While the event rates of neutrinos derived from standard fireball phenomenology are small, we take the existence of the high-energy power-law components as indicative of the decay of $\\pi^{0}$-mesons. This allows us to calculate the magnitude of the neutrino flux from the related charged pions, and we find that a burst like GRB941017 will be observable in IceCube if its power-law component extends to energies in excess of $\\sim1$ GeV."
    },
    "0911/0911.2491_arXiv.txt": {
        "abstract": "We use analytic calculations of the post-recombination gravitational effects of cosmic strings to estimate the resulting CMB power spectrum, bispectrum and trispectrum.   We place a particular emphasis on multipole regimes relevant for forthcoming CMB experiments, notably the Planck satellite.    These calculations use a flat sky approximation, generalising previous work by integrating string contributions from last scattering to the present day, finding the dominant contributions to the correlators for multipoles $l> 50$.   We find a well-behaved shape for the string bispectrum (without divergences) which is easily distinguishable from inflationary bispectra which possess significant acoustic peaks.   We estimate that the nonlinearity parameter characterising the bispectrum is approximately $f_{NL} \\sim -20$ (given present string constraints from the CMB power spectrum. We also apply these unequal time correlator methods to calculate the trispectrum for parrallelogram configurations, again valid over a large range of angular scales relevant for WMAP and Planck, as well as on very small angular scales.   We find that, unlike the bispectrum which is suppressed by symmetry considerations, the trispectrum for cosmic strings is large.    Our current estimate for the trispectrum parameter is $\\tau_{NL} \\sim 10^4$, which may provide one of the strongest constraints on the string model as observational estimates for the trispectrum improve. ",
        "introduction": "Cosmic strings are a common feature in fundamental cosmology scenarios, such as brane inflation \\cite{Tye_or_review}, and they appear to be generic in realistic grand unified theories \\cite{Jeannerot}.  Cosmic strings leave a distinct `line-like' signature in the cosmic microwave background (CMB) \\cite{Gott,KAISERSTEBBINS} which offers one of the best prospects for their detection.   Strings create this imprint after recombination by perturbing photons through relativistic gravitational effects.  While inflationary fluctuations are believed to dominate the overall CMB power spectrum, current constraints suggest that up to 10\\% of the signal could be contributed by cosmic strings \\cite{Battye, Kunzetal}. A higher proportion is incompatible with WMAP because the cosmic string power spectrum around multipoles $l \\approx 200$ is dominated by metric fluctuations with a relatively featureless spectrum, rather than the pre-recombination acoustic peaks characterising inflation.  The gravitational strength of cosmic strings is determined by the parameter $G\\mu = (\\eta/m_{Pl})^2$ which is the ratio of the string tension $\\mu$ to the square of the Planck mass $m_{Pl}$.     The present WMAP limit on the string contribution to the CMB translates into a strong constraint on the parameter $G\\mu \\leq 2.5 \\times 10^{-7}$ \\cite{Battye}. \\par Inflationary CMB fluctuations in the standard picture are very nearly Gaussian, so higher order correlators offer the prospect of  further differentiation between these and competing signals from cosmic strings. The spherical harmonic transform of the three-point CMB correlator is the bispectrum $B_{l_1 l_2 l_3}$.   Usually discussions of the bispectrum are simplified further by focusing on the nonlinearity parameter $f_{NL}$ which is roughly the ratio of the three point correlator to the square of the two-point correlator, that is, $f_{NL}\\sim \\langle \\zeta\\zeta\\zeta \\rangle /\\langle\\zeta\\zeta\\rangle^2$ where the $\\zeta$ are the primordial curvature fluctuations (dimensionless) which seed the CMB anisotropies. For a perturbative theory like inflation, we expect the second-order fluctuations $\\zeta^{(2)}$ to be constructed from convolutions of the linear perturbations $\\zeta^{(2)} \\sim \\zeta^{(1)}* \\zeta^{(1)}$.  This implies that the leading order term in the three-point correlator can be expected to behave as $\\langle \\zeta^{(1)}\\zeta^{(1)}\\zeta^{(2)}\\rangle \\sim \\langle \\zeta^{(1)}\\zeta^{(1)} \\rangle^2$, that is, $f_{NL} \\sim 1$.   In fact, for standard single field inflation there is considerable further suppression from slow-roll and the primordial signal $f_{NL}\\approx {\\cal O}(0.01)$ \\cite{Maldacena}.  One can only obtain a significantly larger $f_{NL}$ in more exotic inflationary models with multiple scalar fields or non-canonical kinetic terms (see, for example, the reviews in \\cite{Chen, Fergusson}).  Similar arguments apply to the trispectrum or four-point correlator characterised by $\\tau_{NL} \\sim \\langle \\zeta\\zeta\\zeta\\zeta\\rangle/\\langle\\zeta\\zeta\\rangle^3$ with $\\tau_{NL} \\lesssim  1$ for standard single field inflation. \\par In contrast, the cosmic string signature has an inherently non-perturbative origin so we do not expect $f_{NL}$ to be small nor, more particularly, $\\tau_{NL}$.    For the higher order correlators from cosmic strings (measured relative to the string power spectrum $\\langle \\zeta\\zeta\\rangle_{cs}$), the relevant scaling with the parameter $G\\mu$ is $f_{NL} \\sim {\\cal A} (G\\mu)^{-1}$ \\cite{Hindmarsh2009} and, as we shall discuss here, $\\tau_{NL} \\sim {\\cal B}(G\\mu)^2$ with coefficients ${\\cal A}$ and ${\\cal B}$ determined by geometric and dynamical considerations. There is a suppressed amplitude ${\\cal A}$ for the bispectrum (due to symmetry cancellations \\cite{Hindmarsh2009}), and we contrast this here with one of the key results of this paper which is the much larger relative value of ${\\cal B}$ for the trispectrum.  Here, we shall compare the higher order correlators from strings we calculate with  the dominant inflationary power spectrum $\\langle \\zeta\\zeta\\rangle_{inf}$ and we will show that the trispectrum with $\\tau_{NL}\\ggg 1$ is  the more likely correlator to produce the strongest constraints on cosmic strings (or with which to identify them). \\par In this paper, we use analytic calculations of the post-recombination string signature to estimate the power spectrum, bispectrum and trispectrum relevant for forthcoming CMB experiments, particularly the Planck satellite.    We generalise previous analytic calculations of Hindmarsh and collaborators for the power spectrum \\cite{Hindmarsh} and bispectrum \\cite {Hindmarsh2009} based on a flat sky approximation.  While this was an important first step, the summation employed was essentially confined only to strings close to the surface of last scattering and it is only relevant for very small angular scales.   Here, we note that the dominant contribution on large angular scales (as well as very small angles) comes from the same gravitational GKS effects but from strings at late times well after last scattering. By integrating in time over the unequal time correlators, we obtain an approximate bispectrum valid for all $l \\gtrsim 50$ (where the flat sky approximation breaks down), i.e. useful for both WMAP and Planck.   We also apply these unequal time correlator methods to the first calculation of the trispectrum (for parralelogram configurations), which is again valid over a large range of angular scales.   Our bispectrum calculation also differs from ref.~\\cite{Hindmarsh2009} by eliminating the divergences they find in their Gaussian integrals for flattened triangles.  We note that these are actually cut-off by the behaviour of other terms in the integrand (due to momentum constraints), yielding a finite result in this regime.   The cross-sectional shape of the bispectrum we present is relatively featureless, except for a finite rise near the edges (for flattened triangles) and suppression towards the corners because of causality contraints (squeezed triangle limit).  This well-behaved shape is suitable for the bispectrum estimation methods being developed for Planck and is easily distinguishable  from bispectra predicted by inflation because of the absence of acoustic peaks \\cite{Fergusson}. Of course, these analytic calculations neglect important recombination effects which also provide significant contributions to the string bispectrum for $500\\lesssim l \\lesssim 2000$. However, determining the extent to which these contributions can confuse the non-Gaussian `line-like' signature of cosmic strings will be the subject of future study \\cite{FergLandetal}. \\par These results for the string bispectrum and trispectrum are important for CMB experiments and should be valid and dominant at both large and very small angular scales.    While the Planck satellite does not have the resolution to see individual string signatures, it should be possible to obtain statistically significant constraints on cosmic strings that compete with limits on $G\\mu$ from the power spectrum, especially for the trispectrum when these techniques are fully developed. On very small angular scales the string power spectrum begins to dominate over the inflationary signature ($l>3000$) because it is not influenced by exponentially decaying transfer functions.  Here direct detection of `line-like' signatures may prove possible provided that experiments can achieve $\\mu K$ sensitivities, as anticipated by AMI and ACT, for example. Again, searching in larger data sets for higher order correlators may provide statistically more significant results. ",
        "conclusions": "In this paper, we have endeavoured to analytically calculate the cosmic string power spectrum, bispectrum and trispectrum over a range of scales relevant for CMB experiments on both large and small angular scales.    We have been particularly focused on extending previous work to multipole ranges ($l<2000$) applicable for the Planck satellite,  an experiment which has the potential to dramatically improve constraints on all these correlators.   We have presented a relatively featureless shape for the bispectrum over the relevant range which should be easily distinguishable from the oscillatory peaks of inflationary bispectra \\cite{Fergusson} if there is a significant signal discovered.    Our preliminary estimates of $f_{NL}$ from cosmic strings indicate that Planck constraints from the bispectrum should be competitive with those from the power spectrum.      We note that we obtain a considerably smaller estimate than the $f_{NL} = -1000$ in ref.~\\cite{Hindmarsh} for several reasons, including the extension of  our analysis over lower multipole ranges relevant for WMAP, a more careful comparison with  $f_{NL}$ estimators used in the literature, and a lower normalisation of the string spectrum derived for Nambu strings \\cite{Battye}, rather than that obtained from field theory simulations \\cite{Kunzetal}.   Nevertheless, as we shall discuss, our estimate contains many uncertainties and more detailed and accurate forecasts are the subject of ongoing work \\cite{FergLandetal}.  We have also evaluated the CMB trispectrum (for parallelogram configurations) on the relevant multipole scales for WMAP and Planck. Again we find a relatively constant trispectrum, finite in all regimes and for which we do not expect dramatic features for other configurations.    Our results indicate a relatively larger signature  for the trispectrum,  in contrast to the amplitude of the bispectrum which is suppressed because it depends on the poor correlation between the string velocity and curvature.   Our preliminary estimate of $\\tau_{NL} $ a few times $10^4$ needs a more detailed analysis to characterise key uncertainties more carefully.   Nevertheless, the trispectrum deserves much closer scrutiny observationally and the prospect of constraining cosmic strings should motivate the development of suitable estimators.  These analytic calculations of the post-recombination gravitational effects of cosmic strings offer important physical insights into the CMB correlations they induce. However, we note that there are many directions in which they can be substantially improved.   Detailed numerical investigations of the CMB power spectrum created by cosmic strings and other topological defects already takes into account a far wider range of physical effects, notably recombination physics around decoupling and a better description of the evolving string network.    It should be possible to similarly develop the unequal time correlation methods presented for the late time GKS signatures here to calculate both the bispectrum and the trispectrum to high accuracy numerically.   In the meantime, however, it is a matter of comparing the present analytic results with CMB maps induced by cosmic string networks on both large (full sky) and small scales in order to get a more accurate normalisation and characterisation of the bispectrum and trispectrum.   Given the stark contrast between the string shapes with those predicted by inflation, there needs also to be a specific search for these signatures in present and forthcoming CMB data, a project which is being actively investigated \\cite{FergLandetal}.   There seem to be  good prospects of using forthcoming CMB data to obtain new insight into cosmic string scenarios using higher order correlators."
    },
    "0911/0911.2344_arXiv.txt": {
        "abstract": "We present an integral field spectroscopic study of the central $2\\times2$ kpc$^2$ of the blue compact dwarf galaxy Mrk~409, observed with the \\emph{Potsdam MultiAperture Spectrophotometer}. This study focuses on the morphology, two-dimensional chemical abundance pattern, excitation properties and kinematics of the ionized interstellar medium in the starburst component. We also investigate the nature of the extended ring of ionized gas emission surrounding the bright nuclear starburst region of Mrk~409. PMAS spectra of selected regions along the ring, interpreted with  evolutionary and population synthesis models, indicate that their ionized emission is mainly due to a young stellar population with a total mass of $\\sim 1.5 \\times 10^6$ \\msun, which started forming almost coevally $\\sim 10$ Myr ago. This stellar component is likely confined to the collisional interface of a spherically expanding, starburst-driven super-bubble with denser, swept-up ambient gas, $\\sim 600$ pc away from the central starburst nucleus. The spectroscopic properties of the latter imply a large extinction ($\\chbeta>0.9$), and the presence of an additional non-thermal ionization source, most likely a low-luminosity Active Galactic Nucleus. Mrk~409 shows a relatively large oxygen abundance ($12+\\log(\\mathrm{O/H})\\sim8.4$) and no chemical abundance gradients out to $R\\sim 600$ pc. The ionized gas kinematics displays an overall regular rotation on a northwest-southwest axis, with a maximum velocity of 60 \\kmsec; the total mass inside the star-forming ring is about $1.4 \\times 10^9 M_\\sun$. ",
        "introduction": "\\label{Sect:Introduction}   Blue Compact Dwarf (BCD) galaxies are systems undergoing violent bursts of star formation \\citep{Sargent1970}, that appear compact in the optical (starburst diameter $\\leq 1$ kpc), and have low intrinsic luminosities ($M_{B} \\geq -18$ mag) and low gas-phase metallicities ($7.1 \\leq 12+\\log\\mathrm{(O/H)} \\la8.3$). These characteristics make them excellent nearby laboratories for investigating some of the most outstanding questions in the contemporary extragalactic astronomy: i)~they are suitable sites to study the process of galaxy formation and evolution, as they are thought to be the local analogs of the building blocks from which larger systems were formed at high redshift \\citep{Kauffmann1993}; ii)~they allow the determination of the primordial helium abundance with a minimum of extrapolation to early conditions \\citep{Pagel1992,IzotovThuan2007}; and, iii)~being smaller and less massive than normal galaxies, BCDs neither can sustain spiral density waves, nor suffer from disk instabilities, hence they permit the investigation of the star-formation process in a relatively simple environment.  In order to properly characterize the BCD galaxy class, and to get insights into the above topics, detailed spectrophotometric analysis of individual objects, aimed to disentangle their different stellar populations and to elucidate their star-forming histories (SFH), are imperative \\citep[see, e.g.,][]{Cairos2002,Cairos2007}. However, very few such studies have been done so far \\citep{Papaderos2002,Papaderos2006,Noeske2000,GildePaz2000a,Fricke2001, Guseva2001,Cairos2002,Cairos2007}. These works, all based on traditional observing techniques (broad- and narrow-band imaging plus long-slit spectroscopy), have shown that combining spectroscopy and surface photometry is indeed essential for tackling the problem of BCD evolution. However, they suffer from several drawbacks, the most severe being the large amount of observing time that they require: acquiring images in several broad- and narrow-band filters, in addition to a series of long-slit exposures sweeping the region of interest, translate into prohibitively long observing times of two or more nights per galaxy. This usually means that the data are taken under varying atmospheric conditions and, potentially, slight differences in the instrumental setup. Clearly, uncertainties in the combination of such a data set will propagate throughout in the data analysis and interpretation. Long-slit spectroscopy has also the additional problem of the uncertainty on the exact location of the slit.  The relatively new observational method of Integral Field Spectroscopy (IFS) overcomes these problems. Detailed IFS studies have been carried out for SBS~0335-052\\,E with VLT/GIRAFFE \\citep{Izotov2006}, for the low-luminosity BCD II~Zw~70 with CAHA 3.5m/PMAS \\citep{Kehrig2008}, for five luminous BCDs with WHT/INTEGRAL \\citep{GarciaLorenzo2008}, for UM~408 with GMOS-IFU \\citep{Lagos2009} and for Mrk~1418 with PMAS \\citep{Cairos2009}.  Although the scope of these studies is somehow limited by a small field of view (FOV) or by low spectral resolution, they convincingly demonstrate that IFS is the method of choice for small, yet complex star-forming (SF) objects like BCDs. Each single exposure of an integral field spectrograph provides both spatial and spectral information, making the observations an order of magnitude more efficient than any  traditional observing technique. Furthermore, IFS provides simultaneously spectra for all spatial resolution elements, under the same instrumental and atmospheric conditions, which guarantees the homogeneity of the dataset.  Systems with a high degree of asymmetry in their SF component, such as BCDs, hosting two or more well-separated SF regions, can be studied in a comprehensive manner only with 2D spectroscopy. Long standing questions in the field, as the propagation of the star-formation, the generation and dispersal of heavy elements or the abundance gradients, can be approached efficiently only with IFS instruments.  While existing works indicate that the SF activity in BCDs proceeds largely in bursts \\citep[e.g.,][]{Krueger1994,MasHesse1999}, we are still lacking an understanding of the two-dimensional star-formation pattern. Established is solely that, in the majority (90\\%) of BCDs, the SF regions are not randomly scattered within the low surface brightness (LSB) host, but rather confined to its central part, out to $\\sim 2$ exponential scale lengths \\citep{Papaderos1996,Cairos2001}. This fact suggests a physical link between the mode of the SF activity in BCDs and the shape of the gravitational potential formed by the old, higher mass-to-light (M/L) ratio stellar LSB host.  Some of the most metal-poor BCDs, thought to be still in the process of their formation, such as SBS~0335-052\\,E \\citep{Thuanetal1997,Papaderos1998}  or SBS~1415+437 (\\citealt{Thuan1999}; \\citealt{Guseva2003}; see \\citealt{Papaderos2008} for a review), present clear signatures of uni-directional SF propagation.  This process can lead to a \\emph{cometary} morphology, with the ``comets head'' (the dominant SF region) located at the tip of an elongated LSB host delineating the trail of past propagating SF activities \\citep{Papaderos2008}. An irregular SF propagation pattern in form of a random-walk was reported  by a detailed analysis of individual stellar clusters in the luminous BCD ESO~338-IG04 \\citep{Ostlin2003}, while the potential role of SF propagation in Mrk~370 and Mrk~35 has been discussed in \\cite{Cairos2002} and \\cite{Cairos2007}, respectively.  In the two BCDs Mrk~86 and Mrk~409, the ring-morphology of the SF regions suggests that the SF activity could have been triggered by the shock wave produced by multiple SNs in a central starburst. In the case of Mrk~86, extensive studies \\citep{GildePaz2000b,GildePaz2002} support this hypothesis. However, because of the small spatial and spectral coverage of the available spectroscopic observations \\cite[][hereafterGdP03]{GildePaz2003a}, conclusive results for Mrk~409 can not be drawn.  In this paper we show the great potential of IFS when applied to BCDs by presenting a pilot study focused on the galaxy Mrk~409 (= NGC~3011).  Mrk~409 is a luminous ($M_{B} = -17.73$) BCD, belonging to the nE morphological class \\citep{LooseThuan1986}. Its intriguing gas ionized morphology --- a compact nuclear starburst region surrounded by two SF rings at projected galactocentric radii of 5\\arcsec\\ and 18\\arcsec, corresponding to 0.64 and 2.3 kpc --- makes it an excellent laboratory for investigating the star-formation pattern and abundance gradients in BCDs. We observed the galaxy with the PMAS spectrograph mounted on the 3.5m Calar Alto telescope. The PMAS field of view ($2\\times2$ kpc$^{2}$ at the Mrk~409 distance) allows us to map the nuclear region and the inner SF ring. The basic parameters of the galaxy are summarized in  Table~\\ref{Table:data}.  The paper is organized as follows: in Sect. \\ref{Sect:Data} we describe the observations, the data analysis and the method employed to derive the two dimensional maps. In Sect. \\ref{Sect:Results} we present continuum and emission line intensity, line ratio and velocity maps as well as results from the analysis of the integrated spectra of the SF regions identified in the galaxy. The results are discussed and summarized in Sect.~\\ref{Sect:Discussion} and Sect.~\\ref{Sect:Conclusions}, respectively.     \\begin{deluxetable}{lccccccc} \\footnotesize \\tablewidth{0pt} \\tablecaption{Basic data for Mrk~409} \\tablehead{ \\colhead{Parameter}  & \\colhead {Data} & \\colhead {Note} } \\startdata Other names                        &  NGC~3011, UGC~5259    \t   &      \\\\ RA~(2000)                          &  \\phs09 49 41\t    \t   &      \\\\ DEC~(2000)                         &  +32 13 16 \t    \t   &      \\\\ $v_\\mathrm{helio}$~(\\kmsec)        &  1527\t\t    \t   &      \\\\ $D$ (Mpc)                          &  26.3\t\t    \t   &\t  \\\\ $A_{B}$~(mag)                      &  0.071\t                   &  (1) \\\\ $M_{B}$~(mag)                      &  $-17.73$  \t           &      \\\\ $m_{B}$~(mag)                      &  $14.37\\pm0.03$\t           &  (2) \\\\ $m_{R}$~(mag)                      &  $13.33\\pm0.05$\t           &  (2) \\\\ $F(\\Ha)$~(ergs s$^{-1}$ cm$^{-2}$) &  $2.93\\pm0.17\\times10^{-13}$ &  (2) \\\\ $m_{J}$~(mag)               \t    &  $11.88\\pm0.02$\t\t   &  (3) \\\\ $m_{H}$~(mag)               \t    &  $11.34\\pm0.03$\t\t   &  (3) \\\\ $m_{K_{s}}$~(mag)           \t    &  $11.17\\pm0.04$\t\t   &  (3) \\\\ $M_{HI}$~($M_\\sun$)         \t    &  $0.31\\times10^{9}$\t   &  (4) \\\\ $M_\\mathrm{T}$~($M_\\sun$)   \t    &  $0.33\\times 10^{10}$\t   &  (4) \\\\ \\enddata \\label{Table:data} \\tablecomments{Coordinates, heliocentric velocity and distance, all taken from NED (http://nedwww.ipac.caltech.edu/). Distance calculated using a Hubble constant of 73 \\kmsec\\ Mpc$^{-1}$, and taking into account the influence of the Virgo Cluster, the Great Attractor and the Shapley supercluster. Absolute magnitude in the $B$ band, computed from the tabulated integrated $B$ magnitude and distance. (1) Absorption coefficient in the $B$ band, from \\cite{Schlegel1998}. (2) \\cite{GildePaz2003b}. (3) From 2MASS. (4) Neutral hydrogen mass $M_{HI}$ and total mass $M_\\mathrm{T}$ from \\cite{ThuanMartin1981}.} \\end{deluxetable} ",
        "conclusions": "\\label{Sect:Conclusions}   We presented PMAS integral field spectroscopy of the central $16\\arcsec\\times16\\arcsec$ of the BCD galaxy Mrk~409.  Mrk~409 is known to undergo intense SF  activity within the central part of an old,  elliptical host galaxy \\citep{GildePaz2005}. Previous work (GdP03) had revealed a compact (diameter $<5\\arcsec$) nuclear starburst region, surrounded by two nearly concentric rings of ionized gas emission with radii of 0.6 kpc and 2.8 kpc. The inner ring was proposed to be the result of in situ star formation in a dense shell of piled-up material on the surface of an expanding starburst-driven super-bubble.  The main results from the present study may be summarized as follows:  \\begin{itemize}  \\item The inner circumnuclear ring of Mrk~409 contains a young stellar population with a total mass of $\\sim 1.5\\times 10^6$ \\msun\\ , which likely started forming almost coevally $\\sim 10^7$ yr ago. This is strong evidence to the interpretation that circumnuclear SF activities in Mrk~409 were triggered through the collision of a spherically expanding, starburst-driven super-bubble with the ambient denser material. Although the details of this process remain unclear, our study suggests that sequential star formation in a spherical shell, as observed in Mrk~409, may contribute considerably to the build-up of the young stellar component in BCDs.  \\item While the circumnuclear SF ring presents properties typical of \\ion{H}{2} regions, diagnostic line ratios for the high-surface brightness starburst nucleus suggest an additional ionization source for this region, most probably a low-luminosity AGN.  \\item The nuclear and circumnuclear SF components in Mrk~409 also differ considerably in their extinction properties: SF regions in the ring show a moderate and relatively uniform intrinsic extinction ($0.2\\la\\chbeta\\la0.3$), while the nuclear emission is subject to strong dust obscuration ($\\chbeta\\approx0.9$). Consequently, corrections for extinction based on the luminosity-weighted spectrum of the central region of a BCD can lead, in the case of Mrk~409, to a severe over-estimation of the \\Ha\\ luminosity and star formation rate.  \\item  Mrk~409 shows a relatively high oxygen abundance $(12+\\log(\\mathrm{O/H})\\sim8.40$), without evidence for any significant chemical abundance gradient out to $R\\sim 0.6$ kpc.  \\item All the studied circumnuclear \\ion{H}{2} regions have electron densities in the low density limit. In the nuclear starburst region the electron density increases to a value as large as 500 cm$^{-3}$.  \\item The ionized gas component within the inner $\\sim 1$ kpc of Mrk~409 displays a smooth kinematical pattern which is likely dominated by rotation; the total mass inside the SF ring is estimated to be $\\simeq 1.4 \\times 10^9$ $M_\\sun$.  \\end{itemize}"
    },
    "0911/0911.1492_arXiv.txt": {
        "abstract": "The distribution of interstellar dust grains (ISDG) observed in the Solar System depends on the nature of the interstellar medium-solar wind interaction. The charge of the grains couples them to the interstellar magnetic field (ISMF) resulting in some fraction of grains being excluded from the heliosphere while grains on the larger end of the size distribution, with gyroradii comparable to the size of the heliosphere, penetrate the termination shock.  This results in a skewing the size distribution detected in the Solar System.  We present new calculations of grain trajectories and the resultant grain density distribution for small ISDGs propagating through the heliosphere.  We make use of detailed heliosphere model results, using three-dimensional (3-D) magnetohydrodynamic/kinetic models designed to match data on the shape of the termination shock and the relative deflection of interstellar \\HI\\ and \\HeI\\ flowing into the heliosphere.  We find that the necessary inclination of the ISMF relative to the inflow direction results in an asymmetry in the distribution of the larger grains (0.1 \\micron) that penetrate the heliopause. Smaller grains (0.01 \\micron) are completely excluded from the Solar System at the heliopause. ",
        "introduction": "Up to one percent of the mass of the interstellar cloud surrounding the heliosphere is carried by dust grains that interact with the heliosphere \\citep{SlavinFrisch:2008,Landgraf:2000}.  The flow of interstellar material past the Sun at 26.4 \\kms\\ drives large interstellar dust grains into the heliosphere, while small grains are diverted around the heliosphere by the interstellar magnetic field (ISMF).  The heliospheric trajectories of intermediate sized grains can be complicated and depend on the solar wind magnetic field and thus solar cycle phase.  Observations of ISDGs in the solar system by the Ulysses, Galileo and Cassini spacecraft show that the density of smallest grains are deficient in the inner heliosphere when compared to the nominal Mathis et al. ``MRN'' power law size distribution \\citep[e.g.][]{Frischetal:1999,MRN:1977}.  For instance, the density of grains of mass $10^{-14}$ g ($a \\sim 0.1$ \\micron) is reduced in the inner heliosphere by a factor of $\\sim 90$ below the MRN predictions and $10^{-15}$ g ($a \\sim 0.05$ \\micron) grains are deficient by three orders of magnitude. At the same time large ($a \\sim 1$ \\micron) grains are absent from the MRN distribution, but are abundant in the inflowing ISDGs.  While several alternative grain size distribution models exist \\citep[e.g.,][]{Kim_etal_1994,Weingartner+Draine_2001,Zubko_etal_2004}, the constraints of staying within the limits of cosmic abundances and explaining the interstellar extinction curve lead in all cases to the presence in the models of small grains that are absent from the observed distribution and a cutoff on the large grain size end that is below that of the observed distribution.  The deficiencies of small grains may result from small grains being deflected around the heliosphere because their high charge-to-mass ratios cause them to couple tightly to the ISMF. On the other hand, the circumheliospheric interstellar medium (CHISM) may simply be deficient in small grains (balancing the overabundance of large grains).  To evaluate which is the case requires calculations of the grain distribution as a function of grain size in the context of realistic heliosphere models. With recent data on the shape of the heliosphere from the crossing of the termination shock (TS) by Voyagers 1 and 2 \\citep{Stone_etal_2008}, and on the deflection of inflowing interstellar H$^0$ relative to He$^0$ \\citep{Lallement_etal_2005}, such models face tighter constraints than ever before.  The inferred asymmetry of the heliosphere can be explained by the orientation of the ISMF relative to the direction of the cloud-Sun relative motion.  The ISMF required to fit the heliosphere data, in turn, has direct implications for ISDG trajectories.  In this paper we present new calculations of grain trajectories and density distributions for small grains that are primarily excluded from the inner heliosphere.  Small grains experience enhanced charging rates due to the ejection of secondary electrons in the hot $10^5 - 10^6$ K plasma between the termination shock and heliopause \\citep{KimuraMann:1998}.  We calculate grain trajectories in 3-D, based on a self-consistent steady state MHD-kinetic heliosphere model with grain charge calculated at each location along the path according to the plasma conditions and radiation field. (We note that the recent Voyager 2 results \\citep{Richardson_etal_2008} on the lower than expected plasma temperature in the inner heliosheath are not incorporated in the heliosphere model we use.  We expect that this could result in somewhat lower dust charging in this region than we calculated, though the detailed effects of charging by thermal particles vs.\\ PUIs has yet to be fully explored.)  The model of plasma density and magnetic field in the heliosphere at each point of space properly accounts for the asymmetries that result from the angle between the ISMF and interstellar gas/dust inflow direction, and ISMF and ecliptic plane \\citep{PogorelovStoneetal:2007,Pogorelov_etal_2008}. When combined with recent models of grain charging \\citep{WeingartnerBarr:2006} and a realistic ultraviolet (UV) radiation field, the Lorentz force at each point in space can be used to calculate the grain trajectory.  Total 3-D grain densities are then calculated for a sample of interstellar grain sizes based on the total gas-to-dust mass ratio in the CHISM and ISDG models.  Many previous calculations have been made of grain interactions with the heliosphere \\citep[e.g.][]{Grun_etal_1994,Frischetal:1999,KimuraMann:1998, Landgraf:2000} but ours is the first to use a realistic heliosphere model, including distortion due to an ISMF orientation that is consistent with recent observational data. ",
        "conclusions": ""
    },
    "0911/0911.1806_arXiv.txt": {
        "abstract": "Establishing or ruling out, either through solid mass measurements or upper limits, the presence of intermediate-mass black holes (IMBHs; with masses of $10^2 - 10^5$ M$_{\\odot}$) at the centers of star clusters would profoundly impact our understanding of problems ranging from the formation and long-term dynamical evolution of stellar systems, to the nature of the seeds and the growth mechanisms of supermassive black holes.  While there are sound theoretical arguments both for and against their presence in today's clusters, observational studies have so far not yielded truly conclusive IMBH detections nor upper limits.  We argue that the most promising approach to solving this issue is provided by the combination of measurements of the proper motions of stars at the centers of Galactic globular clusters and dynamical models able to take full advantage of this type of data set.  We present a program based on {\\sl HST} observations and recently developed tools for dynamical analysis designed to do just that. ",
        "introduction": "There is solid observational evidence for the existence of black holes (BHs) in two very different regimes.  Stellar-mass BHs, with masses in the range $5-15$ M$_\\odot$, are the end products of normal high-mass stellar evolution. Solid mass estimates for these BHs come from measurements of the mass function in several tens of Galactic X-ray binaries. Supermassive BHs (SMBHs), at the other extreme, with masses between $10^{5.5}$ and $10^{9.5}$ M$_\\odot$, have been reliably weighed in galactic centers based on the motions of their surrounding stars and gas. A tight correlation between the masses of SMBHs and the velocity dispersion of their host spheroid (the $M_{\\rm BH}-\\sigma$ relation) evidences that an intimate link must exist between these BHs and the processes relevant to galaxy formation, and highlights the importance of understanding the seeds and growth mechanisms of SMBHs. IMBHs, with masses in between those of the above two regimes, are likely to play a key role in these questions but they are yet to be convincingly detected.  Existing theoretical work provides arguments for and against the presence of IMBHs at the centers of globular clusters (GCs).  Several possible channels of formation have been proposed (see van der Marel 2004 for a review), and, interestingly, all those mechanisms lead to IMBH masses consistent with an extrapolation of the $M_{\\rm BH}-\\sigma$ relation down to dispersions typical of GCs.  On the other hand, it has been argued that it may be difficult for GCs, given their relatively shallow potential wells, to retain growing massive black holes at their centers for too long (Baker et al. 2008; Holley--Bockelmann et al. 2008; Moody \\& Sigurdsson 2009).  This current debate, however, should not prevent astronomers to set out plans to look for the existence of these objects.  To the question of whether IMBHs are possible or not in GCs, the right attitude should be one that, although we do not know the answer at this moment, we will also act independently from theoretical prejudices and observationally probe for their presence anyway.  Any tight upper limits to the masses of central compact objects in a number of GCs are equally important for the field as any solid mass measurements eventually confirming their existence.  Certainly, any of these alternatives would clearly constitute an important improvement over today's situation. ",
        "conclusions": ""
    },
    "0911/0911.0611_arXiv.txt": {
        "abstract": "{ The today's energy density of the induced (second order) gravitational wave background in the frequency region $\\sim 10^{-3} - 10^3$ Hz is constrained using the existing limits on primordial black hole production in the early Universe. It is shown, in particular, that at frequencies near $\\sim 40$ Hz (which is the region explored by LIGO detector), the value of the induced part of $\\Omega_{GW}$ cannot exceed $(1-3)\\times 10^{-7}$. The spread of values of the bound is caused by the uncertainty in parameters of the gravitational collapse of black holes. }  \\PACS{98.80.-k, 04.30.Db}  \\begin{document} ",
        "introduction": "It is well known now that gravitational waves (GWs) can be effectively generated by density perturbations during the radiation dominated era. Tensor and scalar perturbations are decoupled at the first order, but it is not so in higher orders of cosmological perturbation theory. Namely, the primordial density perturbations and the associated scalar metric perturbations generate a cosmological background of GWs at second order through a coupling of modes \\cite{Matarrese:1993zf, Matarrese:1997ay, Carbone:2004iv}. In particular, a second order contribution to the tensor mode, $h_{ij}^{(2)}$, depends quadratically on the first order scalar metric perturbation, i.e., the observed scalar spectrum sources the generation of secondary tensor modes. By other words, the stochastic spectrum of second order GWs is induced by the first order scalar perturbations. Calculations of $\\Omega_{GW}$ at second order and discussions on perspectives of measurements of the second order GWs are contained in works \\cite{Mollerach:2003nq, Ananda:2006af, Baumann:2007zm, Saito:2008jc, Bugaev:2009zh}.  It is natural to conjecture that the detection of GWs from primordial density perturbations on small scales (not directly probed by observations) could be used to constrain overdensities on these scales. However, at the present time, gravitational wave background (GWB) is not yet detected. So, on the contrary, one can constrain GWB using existing limits on amplitudes of primordial density perturbations. Such limits are available, in particular, from studies of primordial black hole (PBH) production at the beginning of radiation era.  It is generally known that PBHs form from the density perturbations, induced by quantum vacuum fluctuations during inflationary expansion. The details of the PBH formation had been studied in \\cite{Carr:1975qj, Khlopov:1980mg}, the astrophysical and cosmological constraints on the PBH density had been obtained in many subsequent works (see, e.g., the recent review \\cite{Khlopov:2008qy}). The detailed constraints on amplitudes of density perturbation spectrum had been reviewed in recent works \\cite{Bugaev:2008gw, Josan:2009qn}.  Having the bounds on the primordial density spectrum, it is possible to determine the corresponding bounds on amplitudes of produced GWB.  The plan of the paper is as follows. In Sec. 2 we present the basic relations connecting the primordial power spectrum of density perturbations and induced GWB. The main result of the paper - the PBH bound on the energy density of GWB - is shown in Fig. \\ref{fig-gw-limit}. In Sec. 3 the mini-review of theoretical estimates for $\\Omega_{GW}$ is given. Sec. 4 contains the conclusions. ",
        "conclusions": "\\label{sec4}  We have presented in the previous section that the typical predictions of the energy density of GWB (characterized by the value of $\\Omega_{GW}$) do not exceed $10^{-8} - 10^{-9}$, although the examples of nonminimal variant of the PBB scenario and some variants of preheating models (see Eq. (\\ref{eq17})) show that larger intensities ($10^{-7}$ or even $10^{-6}$) are not, in general, excluded.  The result, obtained in this paper, can be formulated in the following way: the contribution to the total GWB, which is induced by primordial density fluctuations, is constrained, in the region of $\\sim 10^{-3} - 10^3$ Hz, as it is shown in Fig. \\ref{fig-gw-limit}.  This bound should be taken into account in elaborating inflationary models because, in principle, the induced part of GWB (produced as a result of inflation) can significantly exceed the first order part. The well known example is the running mass model \\cite{Bugaev:2009zh}.  From the experimental point of view, the bound of Fig. \\ref{fig-gw-limit} gives the maximum possible signal from second order GWs. It is important because if the experimentalists will be lucky do detect GWs above this bound, it will mean that the signal is not due to induced GWB (e.g., nonminimal variants of the PBB scenario mentioned above predict peak value of $\\Omega_{GW}$ close to $10^{-6}$).  The bound shown in Fig. \\ref{fig-gw-limit} is slightly more restrictive than the nucleosynthesis bound \\cite{Schwarztmann, Brustein:1996ut} (in LIGO region, at least). The latter bound depends on the effective number of neutrino species $N_\\nu$ while the former one is a straightforward consequence ot the PBH limit on the spectrum of primordial density perturbations (and Einstein's equations).   \\paragraph*{Acknowledgments.} We would like to thank Prof. A.A. Starobinsky for useful comments, and J.F. Dufaux and R. Easther for drawing our attention to their papers on the subject. The work was supported by FASI under state contract 02.740.11.5092."
    },
    "0911/0911.5172_arXiv.txt": {
        "abstract": "We  unify inflation and dark matter via a single scalar field $\\phi$. One of the main difficulties for this unification is that between inflation and dark matter one needs a successful reheating process and a long lasting period of radiation. Therefore the amount of energy density in the inflaton-dark matter field $\\phi$  must be fine tune after reheating to give dark matter. Here we show an alternative scheme in which the  inflaton decays completely, disappearing entirely from the spectrum. However, at low energies, before matter-radiation equality, the same interaction term that leads to the inflaton decay, regenerates $\\phi$. An essential feature is that the transition between the intermediate radiation dominated to the dark matter phase is related to a quantum generation of the scalar field $\\phi$ instead to purely classical dynamics. Thanks to this quantum transition the inflation-dark matter unification can take place naturally  without fine tuning. The unification scheme presented here has three parameters, the mass of the dark matter particle $m_o$, the inflation parameter $\\lm$ and the coupling $g$ for the inflaton interaction. Phenomenology  fixes the value for $\\lm$ and gives a constraint  between $g$ and $m_o$, leaving only the mass of the dark matter particle $m_o$ as a free parameter. These same three parameters $\\lm,m_o,g$  are present in models with inflation and a dark matter wimp particle but without unification. Therefore our unification scheme does not increase the number of parameters and it accomplishes the desired unification between inflaton and dark matter for free. ",
        "introduction": "Inflation has now become part of the standard model of the cosmology  \\cite{inflationarycosmology}. With new data coming soon, and in particular with the Planck satellite mission, the different inflationary models will need to pass a strong test. Inflation set up the initial perturbations from which gravity forms the large scale structures LSS in the universe. At the same time, a huge amount of evidence is gathering in structure formation via the large galaxy maps that are currently under way \\ci{sdss}. These surveys have detected not long ago the Baryon Acoustic Oscillation  BAO \\ci{sdss}, which have the same origin as the CMB \\ci{wmap} but are measured via galaxy distribution. The formation  structure requires the existence of  dark matter and therefore the nature of inflation and of dark matter are then essential building blocks to understand the observations.   Inflation is  associated with a scalar field, the \"inflaton\", and the energy scale at which inflation occurs is typically of the order of $E_{I}= 10^{16} GeV$ \\cite{inflationarycosmology}  but it is possible to have consistent inflationary models with $E_I$ as low as  $O(100) MeV$\\cite{lowinflation}.  On the other hand dark matter is described by an energy density which redshifts as $\\rdm\\sim a(t)^{-3}$ and is described by particles where its mass $m\\gg T$ with $T$ its temperature. These particles can be either fermions or scalar fields. In the case of scalar fields the scalar potential must be $V(\\phi)=m_o^2\\phi^2/2$ and independently on the value of the its mass  the classical equation of motion  ensures that $\\rp\\sim V \\sim a(t)^{-3}$.  Here, since we want to unify inflation with dark matter we will assume that DM is made out of the same field, a scalar field $\\phi$. \\ci{inf-dm}. A scalar field can easily give inflation and DM if the potential is flat at high  energies and at low energies the potential approaches the limit $V(\\phi)=m_o^2\\phi^2/2$. However, most of the time our universe was  dominated by radiation. Therefore, any realistic model must not only explain the two stages of inflation and dark matter but must also allow for a long period of radiation domination. Typically the inflaton decays while it oscillates around the minimum of its quadratic potential \\ci{inflationarycosmology}. If the inflaton decay is not complete then the remaining energy density of the inflaton after reheating must be fine tuned to give the correct amount dark matter.  In order to avoid a fine tuning problem we follow the quantum generation work presented in \\ci{fabios}. In \\ci{fabios} the quantum re-generation process was used to unify inflation with dark energy but here we use the same idea to unify inflation with dark matter. To reheat the universe  we couple $\\phi$ to a relativistic field $\\vp$ via an interaction term $\\lin$. This field $\\vp$ may be a standard  model \"SM\" particle, as for example neutrinos, but it could also be an extra relativistic particle not contained in the SM. An extra relativistic degree of freedom is fine with the cosmological data \\ci{rel.d}. After inflation the $\\phi$  decays into this field $\\vp$ at $E_{Dinf}$ and  to reheat the universe with \"SM\" particles we couple them to $\\vp$  at high energies, e.g. the same energy as the decay of $\\phi$. Thermal equilibrium \"TE\" between $\\vp$ and SM particles will be maintained as long as $\\vp$ and the SM particles remain  relativistic. At low energies we show that we can re-generate the dark matter field $\\phi$ at $E_{gen} $ using the same interaction term $\\lin$ as for the high energy inflaton $\\phi$ decay. The appearance of the dark matter field $\\phi$ at late times is then via a quantum transition  and not due to a classical evolution. A requirement on the mass $\\mmp\\equiv V''(\\phi)$ of the inflaton-dark matter field $\\phi$ is that $ \\mp(E_I)> E_{Dinf}\\gg E_{gen} > \\mp(E_{gen})$ so the simplest potential $V=\\mpo^2\\phi^2$ will not work.  The unification scheme presented here has three parameters, the mass of the dark matter particle $m_o$, the parameter in the inflation potential $\\lm$ and the coupling $g$ for the inflaton decay. Density perturbations normalized to COBE fixes the value for $\\lm$ and the correct amount of dark matter determines the cross section of dark matter at decoupling (wimp particle) gives a constraint between $g$ and $m_o$, leaving only the mass of the dark matter particle $m_o$ as a free parameter. We will show that the same coupling strength that gives the inflaton decay gives the dark matter re-generation at low scales and sets in combination with $m_o$ the wimp decoupling cross section. These same three parameters are present in models with inflation and a dark matter wimp particle \\ci{wimp} but without unification. This implies that our unification scheme does not increase the number of parameters and it accomplishes the desired unification between inflaton and dark matter for free. ",
        "conclusions": "We have presented a model where inflation and dark matter takes place via a single scalar field $\\phi$. This unification is realized  via a quantum re-generation of the $\\phi$ field at low energies. The late time  appearance of $\\phi$ solves the fine tuning problem of the initial condition of dark matter  density $\\Omega_{dmi}$ at high energies in models without quantum re-generation, and  allows for having a long lasting radiation dominated universe after reheating.   The unification scheme presented here has three parameters, the mass of the dark matter particle $m_o$, the inflation parameter $\\lm$ and the coupling $g$ for the inflaton decay. Phenomenology sets the values for $\\lm$ and gives a constraint  between $g$ and $m_o$, leaving only the mass of the dark matter particle $m_o$ as a free parameter. We have shown that the same coupling strength that gives the inflaton decay gives the dark matter re-generation at low scales and sets in combination with $m_o$ the wimp decoupling cross section. These same three parameters are present in models with inflation and a dark matter wimp particle but without unification. This implies that our unification scheme does not increase the number of parameters and it accomplishes the desired unification between inflaton and dark matter for free.        \\noindent{\\bf Acknowledgment} We would like to thank Luis A. Urena for useful discussions and comments. We thank for partial support Conacyt Project 80519, IAC-Conacyt Project.  \\appendix"
    },
    "0911/0911.4668_arXiv.txt": {
        "abstract": "We study the scattering of noncommutative vortices, based on the noncommutative field theory developed in \\cite{Balachandran:2006pi}, as a way to understand the interaction of cosmic strings. In the center-of-mass frame, the effects of noncommutativity vanish, and therefore the reconnection of cosmic strings occurs in an identical manner to the commutative case. However, when scattering occurs in a frame other than the center-of-mass frame, strings still reconnect but the well known 90$^{\\circ}$ scattering no longer need correspond to the head on collision of the strings, due to the breakdown of Lorentz invariance in the underlying noncommutative field theory. ",
        "introduction": "Topological defects such as magnetic monopoles, cosmic strings and domain walls, arise in a large class of spontaneously broken field theories. More recently, cosmic strings have also been shown to arise within string theory, providing a potential indirect way to search for observational signatures of the theory. The existence of defects often yields tight cosmological constraints, since they have the potential to overclose the universe, to yield nontrivial gravitational wave signatures, or to have nontrivial microphysical interactions. To balance this, there are a number of approaches to standard cosmological problems in which topological defects may play an important role.  Cosmic strings are of particular interest, since their self interactions allow a potentially catastrophic string network to lose energy in an orderly fashion, leading to a scaling solution which need not dominate the universe, and thus may contribute to cosmology in interesting ways. For example, while cosmic strings cannot play the central role in seeding structure formation in the universe, some contribution is still allowed~\\cite{Wyman:2005tu} by WMAP and SDSS data, as long as the defects account for no more than 14\\% of the temperature fluctuations in the cosmic microwave background radiation.  Central to an understanding of the cosmological implications of cosmic strings is therefore a detailed understanding of their self interactions. The evolution of cosmic string networks has been thoroughly investigated both numerically and analytically~\\cite{Vachaspati:1984dz,Kibble:1984hp, Bennett:1985qt, Bennett:1986zn, Allen:1990tv, Allen:1990mn}. The scattering of cosmic strings exhibits a crucial feature - they reconnect (intercommute or exchange end points) with a probability close to one, after they collide with each other. This property allows large cosmic strings to break down into smaller strings and loops of strings. The loops themselves are (assuming they are non-superconducting) entirely unstable, and shrink to zero size by emitting energy in the form of gravitational radiation and/or Goldstone bosons~\\cite{Hawking:1990tx, Fort:1993zb}.  In this paper we investigate the possibility of the reconnection of cosmic strings when the spacetime is noncommutative. It has been suggested that quantum gravity and string theory contain hints that spacetime may be noncommutative at a length scale close to the Planck scale. Given this possibility, it is natural to wonder whether it is possible for noncommutative cosmic strings to reconnect after they collide with each other.  There exists~\\cite{Gopakumar:2000zd, Bak:2000ac, Bak:2000im, Jatkar:2000ei, Bak:2000ym, Lechtenfeld:2001gf, Tong:2002xt, Lindstrom:2000kh} a variety of approaches to constructing and studying the properties of noncommutative solitons. In~\\cite{Gopakumar:2000zd} classical stable solitons were constructed for noncommutative scalar field theories, and noncommutative vortex solitons were constructed and studied in~\\cite{Bak:2000ac, Bak:2000im, Bak:2000ym, Jatkar:2000ei}. The moduli space dynamics of noncommutative vortices were analyzed in~\\cite{Tong:2002xt}, and the scattering of noncommutative solitons was studied in \\cite{Lechtenfeld:2001gf, Lindstrom:2000kh}.  In this paper we approach the question of the scattering, and hence reconnection, of cosmic strings by considering the noncommutative abelian Higgs model based on the twisted Poincar\\'e symmetry with deformed statistics developed in~\\cite{Balachandran:2006pi}. (See~\\cite{Chaichian:2004za, Aschieri:2005yw, Balachandran:2005eb, Balachandran:2007kv, Balachandran:2007vx, Akofor:2008ae, Balachandran:2009sq} for more details and developments.) We demonstrate that the nonlocal and Lorentz non-invariant nature of the noncommutative field theory plays a crucial role in the scattering of noncommutative vortices in $2+1$ dimensions, but do not find a significant modification of the behavior of the related cosmic strings in $3+1$ dimensions. The paper is organized as follows. In section~\\ref{comm-abelian-higgs} we briefly review the abelian Higgs model in the commutative case. In section~\\ref{noncomm-spacetime} we then review the particular formulation of noncommutative field theory that we study, providing a description that we hope will be useful to readers not familiar with this construction. In section~\\ref{noncomm-abelian-higgs} we construct the noncommutative abelian Higgs model, and in section~\\ref{sec:dynamics} we then discuss the low energy dynamics of noncommutative vortices and describe how their scattering is qualitatively and quantitatively different from that of their commutative counterparts, before concluding. Throughout this paper we use the mostly negative signature. ",
        "conclusions": "In this paper we have investigated the scattering of noncommutative vortices, and hence the interaction between noncommutative cosmic strings, with the goal of understanding how these may differ from their commutative counterparts. We have worked in the Moyal spacetime, have implemented the effects of noncommutativity by using the star product and by rewriting the ordinary fields as statistics deformed fields, and have focused on the noncommutative version of the abelian Higgs model. We have also used the geodesic approximation to probe the low energy dynamics of vortices, which allows us to express the relevant quantities in terms of the kinetic Lagrangian.  We have demonstrated several results, the first of which is that noncommutative cosmic strings reconnect after collision, just like their commutative relatives. The effects of noncommutativity in the Moyal spacetime can be captured through operators involving the total momentum operator. This allows us to show, within the geodesic approximation, in which we can phrase the relevant questions in terms of the kinetic Lagrangian, that in the center-of-mass frame the scattering of noncommutative cosmic strings is the same as in of the commutative case.  In non-center-of-mass frames, however, our formalism allows us to easily see that the scattering of noncommutative vortices can be somewhat different than in the commutative limit. While it is clear that cosmic strings will still reconnect after collision, unlike in the commutative case the well known 90$^{\\circ}$ scattering may not correspond to a zero impact parameter collision. Thus, the scattering of noncommutative vortices in $2+1$ dimensions can be seen to be quantitatively different from the commutative case, but the overall behavior of cosmic strings in $3+1$ dimensions remains essentially unchanged by the addition of noncommutativity."
    },
    "0911/0911.1798_arXiv.txt": {
        "abstract": "We report on the detection of the two shortest period non-interacting white dwarf binary systems. These systems, SDSS\\,J143633.29+501026.8 and SDSS\\,J105353.89+520031.0, were identified by searching for radial velocity variations in the individual exposures that make up the published spectra from the Sloan Digital Sky Survey. We followed up these systems with time series spectroscopy to measure the period and mass ratios of these systems. Although we only place a lower bound on the companion masses, we argue that they must also be white dwarf stars. With periods of approximately 1 hour, we estimate that the systems will merge in less than 100\\,Myr, but the merger product will likely not be massive enough to result in a Type 1a supernova. ",
        "introduction": "White dwarf stars (WDs) are the end point of stellar evolution for 98\\% of all stars \\citep{Weidemann00} and store the archaeological record of the Galaxy. WDs in binaries are particularly rich systems to study. Because of their intrinsically low luminosity the companions must also be faint, and are frequently rare or interesting objects. As an example, WDs are ideal targets for the direct detection of planets \\citep{Debes05cc2, Farihi08planet, Hogan09, Mullally09} and brown dwarf stars \\citep[e.g.][]{Farihi05bd2}.  The Sloan Digital Sky Survey \\citep[SDSS;][]{York00} has increased the number of spectroscopically identified WDs from two thousand to a few tens of thousands. This has, in turn, allowed follow-up surveys of specific types of WD systems, from pulsators \\citep[e.g.][]{Mullally05, Nitta09}, to extremely low mass WDs \\citep{Kilic07lowmass} to binaries involving main-sequence stars \\citep[e.g.][]{Silvestri06, Heller09}, and binaries with neutron stars \\citep{Agueros09}.  Binary WDs hold the solution to an enduring problem in astrophysics; the progenitors of Type Ia Supernova (SNIa). The origin of SNIa is of great interest given their role in galactic chemical evolution and determining the nature of dark energy. If a WD accretes enough material that its mass approaches the Chandrasekhar limit ($\\sim$1.4\\msolar), the star can no longer be supported by electron degeneracy pressure and explodes as a supernova. This scenario explains the lack of observed hydrogen in the spectra of SNIa, as well as the striking similarities in the lightcurves and spectra. However, the source of the accreted material, remains a subject of active debate.  In the double degenerate scenario for SNIa progenitors \\citep{Iben84, Webbink84}, two WDs in a tight binary in-spiral due to the emission of gravitational radiation. If the total mass is near the Chandrasekhar mass, the merger results in a supernova. While theoretically appealing, it is not clear that nature favors this method.  The SPY survey \\citep{Napiwotzki01}, a high precision radial velocity survey of over a thousand WDs, failed to find any candidates with periods short enough to merge within the lifetime of the Galaxy, and masses large enough to explode as SNIa \\citep{Nelemans05, Napiwotzki07}.  The SDSS offers an opportunity to build on the SPY survey with a much larger sample of stars. \\citet{Kleinman04} and \\citet{Eisenstein06} meticulously collected and classified the spectra of nearly 10,000 WD and sub-dwarf stars, of which approximately 8,000 were single stars with hydrogen or helium atmospheres (DAs and DBs respectively). The spectra were obtained as a series of 3 or more 15 minute exposures usually taken consecutively \\citep{Abazajian09}, which makes it possible to identify massive companions with orbital periods of a few hours or less and radial velocity amplitudes $\\gtrsim$170\\kms. The low luminosity of the WD means that any fainter companion must be a degenerate object (WD, brown dwarf, neutron star, etc.) or a very late M star; any other object would be more luminous than the white dwarf.  SWARMS \\citep[the Sloan White dwArf Radial velocity data Mining Survey,][]{Badenes09} exploits these individual exposures to mine the SDSS spectroscopic database for double degenerate white dwarf (DDWD) systems. Our survey is complementary to the SPY survey in that it has a lower radial velocity sensitivity, but is still sensitive to white dwarf companions for many thousands of objects. In this paper we present two binary systems, \\object{SDSS J143633.29+501026.8} and \\object{SDSS J105353.89+520031.0}, with periods of 1.1 and 0.96 hours respectively. These systems constitute the shortest period non-interacting double degenerate binaries yet found, and are significantly shorter than the 1.46 hour period of the previous record holder, \\object{WD0957$-$666} \\citep{Moran97}. As there are no visible absorption lines from the companions our mass estimates are only lower bounds, but in each case the companion is most likely another WD. ",
        "conclusions": "We report on the detection of the two closest non-contact  white dwarf binaries known. These systems were detected by mining the spectroscopic database of the SDSS, and followed-up with time resolved optical spectroscopy.  We argue that the companions must also be WDs: Main-sequence stars of the requisite mass are more luminous than the primary WDs, and a neutron star or black hole in such close proximity would have stripped off the thin outer layers of the primaries. With periods of about an hour, these systems will merge in less than 100\\,Myr and probably produce a high mass WD."
    },
    "0911/0911.4945_arXiv.txt": {
        "abstract": "We present a new, close relation between column densities of OH and CH molecules based on 16 translucent sightlines (six of them new) and confirm  the theoretical oscillator strengths of the OH A--X transitions at 3078 and 3082 \\AA\\ (0.00105, 0.000648) and CH B--X transitions at 3886 and 3890 \\AA, (0.00320, 0.00210), respectively. We also report no difference between observed and previously modelled abundances of the OH molecule. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0911/0911.3879_arXiv.txt": {
        "abstract": "We present H$\\alpha$ narrow-band imaging of 17 dwarf irregular galaxies (dIs) in the nearby Centaurus~A Group. Although all large galaxies of the group are or have been recently through a period of enhanced star formation, the dIs have normal star formation rates and do not contain a larger fraction of dwarf starbursts than other nearby groups such as the Sculptor Group or the Local Group. Most of the galaxies in the group now have fairly accurately known distances, which enables us to obtain relative distances between dIs and larger galaxies of the group. We find that the dI star formation rates do not depend on local environment, and in particular they do not show any correlation with the distance of the dI to the nearest large galaxy of the group. There is a clear morphology-density relation in the Centaurus~A Group, similarly to the Sculptor Group and Local Group, in the sense that dEs/dSphs tend to be at small distances from the more massive galaxies of the group, while dIs are on average at larger distances. We find four transition dwarfs in the Group, dwarfs that show characteristics of both dE/dSphs and dIs, and which contain cold gas but no current star formation. Interestingly the transition dwarfs have an average distance to the more massive galaxies which is intermediate between those of the dEs/dSphs and dIs, and which is quite large: 0.54$\\pm$ 0.31 Mpc. This large distance poses some difficulty for the most popular scenarios proposed for transforming a dI into a dE/dSph (ram-pressure with tidal stripping or galaxy harassment). If the observed transition dwarfs are indeed missing links between dIs and dE/dSphs, their relative isolation makes it less likely to have been produced by these mechanisms. An inhomogeneous IGM containing higher density clumps would be able to ram-pressure stripped the dIs at larger distances from the more massive galaxies of the group. ",
        "introduction": "The Centaurus~A Group of galaxies is one of the closest groups of galaxies outside the Local Group (at a mean distance of about 3.8 Mpc). Its composition is, however, much different from that of the Local Group, being a heterogeneous assembly of early to late-type galaxies, and, in fact, it has the largest dispersion of morphological types amongst all of the 55 nearest groups reported by \\citep{deV79}. Interestingly, it seems that all of its main galaxy members have been through, or are experiencing, a period of enhanced star formation. The most prominent member of the group is Centaurus~A itself (NGC~5128), a giant peculiar elliptical radio galaxy with powerful X-ray and radio jets, which most probably suffered a recent major merger. Both M83 and NGC~5253 are starburst galaxies \\citep{cal99}; NGC~4945 has a Seyfert nucleus \\citep{do96} and NGC~5102 is in a post-starburst phase \\citep{da08}. One is then led to wonder if the group's dwarf galaxy members also share this elevated level of activity. About 50 dwarfs are now known in the group \\citep{k07}, of which about two-thirds are gas-rich dwarfs (dIs), and it would be interesting to know if these dIs also have an elevated level of current star formation and a larger fraction of dwarf starbursts than the dwarf galaxy population of less ``active'' groups, such as the nearby Sculptor Group \\citep{scm03a}. In other words, is the star formation activity of the dwarfs influenced by their global group environment, or, does it only depend on local conditions and internal processes?  A reasonable estimate of star formation rates (SFRs) in dwarf galaxies can be obtained by H$\\alpha $ imaging, as it is known that in dwarf galaxies infrared and radio emission both can seriously underestimate the star formation rate \\citep[e.g.,][]{b03}. In the case of the Centaurus~A Group dwarfs, most of the dwarfs have relatively accurate distances \\citep[e.g.,][with distances determined via the tip of the red giant branch method with errors typically of $\\sim $0.5 Mpc]{k02}, and with optical diameters ranging from 1 to 10 arcmins they can have their full area imaged easily in one pointing. An H$\\alpha$ imaging survey can then produce much better star formation estimates, especially compared to many galaxy studies at higher redshifts for which SFR estimates are derived from longslit data, for example \\citep{bal97}, or even worse from limited aperture fibers such as SDSS or 2dF surveys. As is obvious from Figure 1 below, in some cases most of the star formation activity is found in the outside regions, and a slit passing through the nucleus would severely underestimate the true star formation rate. A detailed survey of HII regions in each dwarf is also useful in order to find individual high surface brightness HII regions required for follow-up spectroscopy for accurate abundance analysis \\citep[e.g.,][]{scm03a, scm03b}. The vast majority of nearby dwarfs have only been observed through longslits, not multi-object spectroscopy, so it is even more important to know exactly how to position the slit to align them on the brightest HII regions, from which one can reap the largest numbers of fainter emission lines.  Another useful outcome of a H$\\alpha$ survey of the dIs members of a group is that one can identify so-called transition dwarf galaxies, dwarfs that show characteristics in their morphology of both early-type dwarfs, dwarf ellipticals (dEs) or dwarf spheroidals (dSphs), and late-type dwarfs (dIs). These transition dwarfs have cold gas but no (or extremely low) active star formation so they will be mostly non-detections in the H$\\alpha$ survey. These dwarfs have been proposed as a ``missing link'' between the two classes of dwarfs \\citep[e.g.,][]{ggh03}. The question as to whether they represent a real evolutionary link between dEs/dSphs and dIs is still controversial. In the Local Group there is a clear morphology-density relation amongst the dwarfs, with dSphs being found predominantly at short distances from large galaxies, while dIs are more spread out in the group. Several scenarios have been proposed to efficiently transform a dI into a dSph in the vicinity of larger galaxies. Did dSphs end up as they are because of internal properties (from genetics), or are they dIs or pre-dIs that transformed into dSphs simply because they happen to be at the wrong place at the wrong time? The study of transition dwarfs in groups might bring clues to their nature, and, indirectly, to the nature of dSphs too.  In \\S 2 we will describe the H$\\alpha$ observations, and \\S 3 the SFRs for the Centaurus~A Group dwarfs will be derived. Section 4 will present the star formation trends for dIs, its dependence on environment, and will discuss the transition dwarfs in the Centaurus~A Group and other groups. ",
        "conclusions": "From our H$\\alpha$ imaging survey of dIs in the Centaurus~A Group, we find the following results:  The Centaurus~A Group dIs do not have enhanced star forming activities, and do not contain a larger fraction of dwarf starbursts when compared to other nearby groups such as the Sculptor Group or the Local Group. This is perhaps surprising given that all large galaxies in the group are exhibiting or have recently exhibited a period of enhanced star formation. The gas depletion timescales of the dIs are very large, and except for a few rare cases, are several times the present age of the universe.  We find that the SFRs of the Centaurus~A Group dIs do not depend on any global properties of the galaxies, such as magnitude, central surface brightness, HI mass-to-light ratio, or B-R color. They do not depend on local environment either, in particular there is no correlation with the distance of the dI to the nearest large galaxy of the group.  Nonetheless, there is a morphology-density relation in the Centaurus~A Group, similarly to the Sculptor Group and Local Group, in the sense that dEs/dSphs tend to be at small distances from the more massive galaxies of the group, while dIs are on average at larger distances. Interestingly, the four transition dwarfs of the Centaurus~A Group (ESO~269-G58, UGCA~365, ESO~384-G16 and UKS~1424-460) have an average distance intermediate between those of the dEs/dSphs and dIs. Together with the four transition dwarfs of the Sculptor Group and six of the Local Group, they have a quite large average distance of 0.54$\\pm$ 0.31 Mpc. This large distance poses some difficulty for the most popular scenarios proposed for transforming a dI into a dE/dSph. Both ram-pressure with tidal stripping, as well as galaxy harassment, need the dI to be much closer in its orbit to the large galaxy to have any effect. Also, after the first encounter, they produce transition objects that appear as strongly barred disturbed objects. If the observed transition dwarfs are indeed missing links between dIs and dE/dSphs they are unlikely to have been produced by these mechanisms.  A possible scenario is that the IGM in these groups of galaxies is inhomogeneous, with regions of higher densities $> 10^{-5}$ cm$^{-3}$, where clumps would be able to remove the HI from a dI by ram-pressure stripping. Several observations point indeed towards such a possibility. According to hydrodynamical simulations most low-redshift baryons are expected to be found in a warm-hot intergalactic medium (WHIM). This WHIM has already been detected in clusters of galaxies, e.g. Takei et al. (2007). In groups of galaxies too, OVI absorption lines studies in the UV \\citep{sem03} and X-ray absorption lines studies \\citep{ni02} detect warm temperature ionised gas pervasive to the groups, with densities possibly as high as 10$^{-4}$ cm$^{-3}$. This clumpy IGM could be associated loosely with galaxies (because of past mergers or disturbances, past bursts etc), or could be simply higher density regions of filaments expected from the cosmic web \\citep{dave2001}. This would also explain the existence of dSphs distant from any large galaxies such as Cetus and Tucana."
    },
    "0911/0911.3338_arXiv.txt": {
        "abstract": "We investigate the possibility of detecting light long-lived particle (LLP) produced by high energy cosmic ray colliding with atmosphere. The LLP may penetrate the atmosphere and decay into a pair of muons near/in the neutrino telescope. Such muons can be treated as the detectable signal for neutrino telescope. This study is motivated by recent cosmic electron/positron observations which suggest the existence of O(TeV) dark matter and new light O(GeV) particle. It indicates that dark sector may be complicated, and there may exist more than one light particles, for example the dark gauge boson $A'$ and associated dark Higgs boson $h'$. In this work, we discuss the scenario with $A'$ heavier than $h'$ and $h'$ is treated as LLP. Based on our numerical estimation, we find that the large volume neutrino telescope IceCube has the capacity to observe several tens of di-muon events per year for favorable parameters if the decay length of LLP can be comparable with the depth of atmosphere. The challenge here is how to suppress the muon backgrounds induced by cosmic rays and atmospheric neutrinos. ",
        "introduction": "The pursuit of new light boson (LB) with mass much lighter than W and Z bosons in the standard model(SM) has a long history. A well known example is a light gauge boson under an extra $U(1)$ gauge group beyond the SM gauge group $SU(3)_c\\otimes SU(2)_L\\otimes U(1)_Y$. However such kind of new LB is stringently constrained since the current precise measurements are in excellent agreement with the SM predictions. Obviously the interaction between the new LB and SM sector should be tiny to evade current constraints. With the same reason detecting LB experimentally is a very challenging task.  Recently the LB attracts more attention due to the new cosmic observations by PAMELA \\cite{Adriani:2008zq}, ATIC \\cite{:2008zz} and Fermi \\cite{Abdo:2009zk}. PAMELA collaboration reported excess in the positron fraction from 10 to about 100 GeV but absence of anti-proton excess \\cite{Adriani:2008zq}. It is still consistent with the results released by ATIC \\cite{:2008zz} or Fermi \\cite{Abdo:2009zk} experiments. The afterwards investigations suggested that these new observations need new source of electron/positron which may come from the dark matter (DM). The DM particles might annihilate or decay into electron/positron in the halo today. Moreover man began to realize the possible connection between the heavier DM at O(TeV) and the LB at O(GeV). If the annihilating DM produces the extra electron/positron, there is a mismatch between the DM annihilation cross section expected in the epoch of freeze-out and that required to account for recent new observations. Namely the present DM annihilation cross section is too small.  The new O(GeV) LB, via the so-called Sommerfeld enhancement, can fill the gap. Moreover, in order to be consistent with only electron/positron excess, LB decays preferably into charged leptons \\cite{ArkaniHamed:2008qn}.  Generally speaking such new boson may be scalar, pseudoscalar or gauge boson. The interactions between the LB and SM sector could arise from the mixing between the LB and photon or Higgs. It is quite interesting to investigate how to detect such kind of LB. Many authors have studied how to produce such LB and detect them at colliders, namely low-energy collider, large hadron collider or fix-target experiment. For the collider search, one often requires the life time of LB is short, as a consequence the charged leptons as the LB decay products could be observed at the detector. Such charged leptons could be clearly and easily identified. The construction of realistic model of dark sector showed that dark sector can be more complicated. The assumption that LB is short-lived might be incorrect. If dark sector contains an array of LBs besides a light gauge boson, some lighter ones may be long-lived due to the suppressed interaction with SM sector. Several recent works provided an interesting approach to search such long-lived particle (LLP) \\cite{Batell:2009zp,Schuster:2009au,Schuster:2009fc,Meade:2009mu}. The DM trapped inside the Sun/Earth would annihilate into LBs. If the LB has a long lifetime, it can travel through the Sun/Earth and decay into gamma rays or charged leptons which could be observed.  In this paper we point out another possible way to search such kind of LLP via the high energy cosmic rays. The high energy cosmic rays interact with atmospheric nucleons every second and can be treated as a natural and costless high energy hadron collider. If a proton with energy of $E\\sim 10^4$ GeV in cosmic ray collides with an atmosphere nucleon, the center-of-mass energy $\\sqrt{s}$ is approximately $\\sqrt{2m_N E}\\sim 10^2 $ GeV. Such mechanism can copiously produce LLPs through $pN$ collisions. If the lifetime of LLP is appropriate, it can penetrate the atmosphere and arrive at the neutrino telescope. The subsequently decay of LLP into a pair of muons can be observed by the telescope \\cite{Batell:2009yf,Bjorken:2009mm}.   If LLP decays near the detector, the difficult task is to distinguish the signal muon pair from the huge muon backgrounds which are produced by high-energy cosmic rays through hadron(for example $\\pi$) decays and/or QED process. Provided that the time information of muon is precisely recorded, one can identify two signal muons which are supposed to arrive at detector at the same time, not heavily polluted by two irrelevant coincident parallel muon events. In order to suppress the muon backgrounds, the analysis focused on quasi-horizontal events might be important \\cite{Illana:2006xg,Ahlers:2007js}. A more optimistic case is that LLP decays inside the detector with an 'obvious' decay vertex. The challenge here is how to distinguish almost parallel di-muon from the single muon. Another possible case is that for a single collision between cosmic ray particle and atmosphere nucleon, more than one LLP are produced. In this case the multi muon pairs appear in detector at the same time will be a significant signal.  In this work, we assume the LB has a long lifetime to penetrate through the atmosphere, some mechanisms which satisfy this requirement will be discussed. We consider the LB production by the cosmic rays and simulate di-muon events from LB decay. In order to detect such high energy muons, we focus on the large volume neutrino telescope IceCube \\cite{Klein:2008px} which will reach an effective detecting area in square kilometers. The detector are installed under the ice surface in a depth of 1.4 km in order to suppress muon backgrounds produced in atmosphere. Moreover, there is a extension of IceCube namely DeepCore which is still under construction \\cite{Resconi:2008fe}. DeepCore will be installed in the more deeper location which will suffer less above-mentioned muon backgrounds.  This paper is organized as following. In the  section II, we describe two classes of models which contain LLPs, discuss the main LLP production processes and calculate the LLP production cross section in $pN$ collision. We utilize PYTHIA to do the calculation and simulate the LLP events. In Section II, we also investigate the LLP production flux produced by primary cosmic ray. We find that for a LLP production process with O$(10)$ pb, the production rate of LLP can reach to O$(10^4)$ per square kilometer and per year. In Section III, we investigate the possibility to detect di-muon signal from LLP decay in the neutrino detector. The conclusions and discussions are given in the last section. ",
        "conclusions": "In this paper, we investigated the possibility of searching long-lived particle (LLP) produced by high energy cosmic ray colliding with atmosphere. The LLP may penetrate the atmosphere and decay into a pair of muons near/in the neutrino telescope. Such muons can be treated as the detectable signal. This study is motivated by recent cosmic electron/positron observations by PAMELA/ATIC/Fermi. The new data suggests new source of electron/positron which may come from O(TeV) dark matter. In order to understand the dark matter thermal history, new light O(GeV) particles have been proposed. It is quite natural to conjecture that dark sector is complicated. There are more than one light particles in dark sector, for example the dark gauge boson $A'$ and associated dark Higgs boson $h'$. In this paper, we studied the scenario with $A'$ heavier than $h'$ and $h'$ is treated as LLP.  We have studied the LLP production processes and found that the promising process is single $A'$ production $q\\bar{q}\\rightarrow A'$, and $A'$ subsequently decays into $h'$ rather than into SM particles. Our numerical calculations show that for $A'$ with mass of $1 GeV$, the production rate can reach $10^{4} km^{-2} s^{-1} sr^{-1}$ and the final di-muon rate from $h'$ is serval tens for some favorable parameter region. We have assumed the $\\kappa \\sim 10^{-3}$, $\\alpha'=\\alpha$ and the branching ratio of $A'\\rightarrow h' \\rightarrow \\mu^{+}\\mu^{-}$ is $1$. For different parameter selections, our results need to be multiplied by a factor of $(\\kappa/10^{-3})^2 Br_{A'\\rightarrow h'}Br_{h'\\rightarrow \\mu^{+}\\mu^{-}}$. Here it is worth remarking that a complete analysis for LLP production needs QCD correction to production process and simulations of cosmic ray shower included secondary hadrons and nucleons. The LLP could also be produced from secondary meson decay directly if allowed by kinematics \\cite{Batell:2009di}. These elements will increase LLP production rate efficiently.   We simulated the signal di-muon events and calculated the rate as functions of di-muon energy and the separation angle between two muons. Our numerical results showed that the large volume neutrino detector IceCube is suitable to detect LLP with lifetime about $10^{-8}\\sim 10^{-4}$ s. It is worth to mention that such parameter region could be compatible with the constraints from fixed-target experiments. For example, the Ref. \\cite{Schuster:2009au} has reported the constraints on LLP decay length $c \\tau$ between $1$ cm and $10^8$ cm by CHARM experiment result \\cite{Bergsma:1985qz}.  The scenario proposed in this paper could be cross-checked (tested) at the low energy $e^{+}e^{-}$ collider and/or large hadron collider. At the collider, the LLP will escape from the detector and act as the missing energy. For example, the well promising process is the $\\gamma A'$ associated production \\cite{Yin:2009mc}. The $\\gamma+ {E\\!\\!\\!\\! /\\,\\,}_T$ signal can be isolated from SM irreducible background $\\gamma Z \\rightarrow \\gamma \\nu \\bar{\\nu}$. If the LLP associated production with other light bosons which decay into charged leptons, the multi- $e^{+}e^{-}/\\mu^{+}\\mu^{-}+{E\\!\\!\\!\\! /\\,\\,}_T$ is a cleaner signal. Moreover, the interaction between SM and dark sector may be induced via the mixing of Higgs fields. It implies that SM Higgs may decay into LLP. Such possible invisible decay modes can even change our search strategies of SM Higgs boson."
    },
    "0911/0911.4803_arXiv.txt": {
        "abstract": "The snow line in a gas disk is defined as the distance from the star beyond which the water ice is stable against evaporation.  Since oxygen is the most abundant element after hydrogen and helium, the presence of ice grains can have important consequences for disk evolution.  However, determining the position of the snow line is not simple.  I discuss some of the important processes that affect the position of the snow line. ",
        "introduction": "Oxygen is the most abundant element after hydrogen and helium, and water is one of the most abundant oxygen-bearing molecules.  When water is a vapor, the mass fraction of solids in a solar composition disk is $\\sim 4.2\\times 10^{-3}$.  After water condenses, the solid mass fraction jumps to $\\sim 1.3\\times 10^{-2}$ \\cite[(Anders \\& Grevasse 1989)]{AndersGrevasse89}.  The properties of a protostellar disk will therefore change significantly when we cross the boundary between water vapor and ice grains.  These changes can have important consequences for other processes taking place in the disk.  One consequence is the change in opacity that accompanies the transition from a region where water is a vapor to a region where it is a solid.  This is illustrated in Fig.\\,\\ref{pollakopac} which is based on the Rosseland mean opacities of \\cite[Pollack et al. (1985)]{Polletal85}. \\begin{figure} \\vspace*{2.0 cm} \\begin{center} \\includegraphics[width=4in]{pollakopac.eps} \\vspace*{-2.0 cm} \\caption{Opacity of dust grains thought to have been present in the early solar nebula as a function of temperature.  The opacities are taken from \\cite[Pollack et al. (1985)]{Polletal85}.} \\label{pollakopac} \\end{center} \\end{figure} At temperatures above around 170 K, the opacity is due to mainly rocky grains, but just below this temperature the opacity jumps by nearly a factor of 5 because of the sudden condensation of water vapor to ice grains.  At still lower temperatures the opacity again drops because, in taking the Rosseland mean for low temperatures, the longer wavelengths are given more weight.  At these longer wavelengths the ice grains are inefficient scatterers and absorbers, so the opacity decreases.  Such changes in opacity can have an important effect on the evolution of the gas disk.  Another example of the possible significance of the snow line is the mechanism suggested by Stevenson and Lunine \\cite[(1988)]{StevLun88} for forming Jupiter.  Sunward of the snow line the disk temperature is high enough so that all the water is in the vapor phase.  Once you cross the snow line and water begins to condense, the mass fraction of water in the vapor phase decreases.  This causes a composition gradient across the snow line which drives diffusion of water vapor into the colder region and causes this region to become enhanced in solids.  Stevenson and Lunine \\cite[(1988)]{StevLun88} argued that this enhancement could lead to the formation of Jupiter.  The problem with this scenario is that this mechanism forms only one giant planet.  Some other mechanism would be required for Saturn.  But if it is correct, then the snow line had to have been somewhere in the vicinity of Jupiter, i.e. at around 5 AU.  Unfortunately, Jupiter appears to be sending us contradictory messages.  An analysis of Jupiter's composition by Lodders \\cite[(2004)]{Lodders04} finds that Jupiter is anomalously low in water, and concludes that the snow line was much further away from the sun; perhaps in the neighborhood of 10 AU.  A third reason for trying to determine the position of the snow line is that it tells us where we can expect to see icy bodies.  This has particular significance in view of the recent discovery of main-belt comets \\cite[(Hsieh and Jewitt 2008)]{HsieJewit08}.  These bodies have orbits that are dynamically similar to those of the main belt asteroids \\cite[(Jewitt et al. 2009)]{Jewitetal09} yet display cometary activity.  If main-belt comets were indeed formed in the asteroid belt, the snow line must have been closer to the sun than around 3 AU. ",
        "conclusions": "The snow line is an important concept that influences our understanding both of the composition of solar system bodies and the processes that lead to their formation.  In addition, the position of the snow line has consequences for the opacity of the disk and its evolution.  However, the position of the snow line is not merely a function of the temperature of the disk gas.  It also depends on the density of water vapor in the disk, and on the details of the radiation field, as well as on the size and composition of the ice grains.  As a result, it is not proper to talk about a ``snow line\" {\\it per se}, but of an ``ice stability region\" ({\\it ISR}).  The borders of this region will depend not only on distance from the central star, but also on the height above midplane.  The size distribution of ice grains in this {\\it ISR} may vary from place to place, and may even show a time dependence through the hysteresis effect described above.  Theoretical predictions of where the borders of this {\\it ISR} are located, as well as interpretations of observations of ice grains will need to base themselves on detailed modeling of the physics of such grains."
    },
    "0911/0911.1435_arXiv.txt": {
        "abstract": "We use observational data from  Type Ia Supernovae (SNIa), Baryon Acoustic Oscillations (BAO), and Cosmic Microwave Background (CMB), along with requirements of Big Bang Nucleosynthesis (BBN), to constrain the cosmological scenarios governed by Ho\\v{r}ava-Lifshitz gravity. We consider both the detailed and non-detailed balance versions of the gravitational sector, and we include the matter and radiation sectors. We conclude that the detailed-balance scenario cannot be ruled out from the observational point of view, however the corresponding likelihood contours impose tight constraints on the involved parameters. The scenario beyond detailed balance is compatible with observational data, and we present the corresponding stringent constraints and contour-plots of the parameters. Although this analysis indicates that Ho\\v{r}ava-Lifshitz cosmology can be compatible with observations, it does not enlighten the discussion about its possible conceptual and theoretical problems. ",
        "introduction": "Recently, a power-counting renormalizable, ultra-violet (UV) complete theory of gravity was proposed by Ho\\v{r}ava in \\cite{hor2,hor1,hor3,hor4}. Although presenting an infrared (IR) fixed point, namely General Relativity, in the  UV the theory possesses a fixed point with an anisotropic, Lifshitz scaling between time and space. Due to these novel features, there has been a large amount of effort in examining and extending the properties of the theory itself \\cite{Volovik:2009av,Cai:2009ar,Cai:2009dx,Orlando:2009en,Nishioka:2009iq,Konoplya:2009ig,Charmousis:2009tc,Li:2009bg,Visser:2009fg,Sotiriou:2009gy, Sotiriou:2009bx,Germani:2009yt,Chen:2009bu,Chen:2009ka,Shu:2009gc,Bogdanos:2009uj,Kluson:2009rk,Afshordi:2009tt,Myung:2009ur,Alexandre:2009sy, Blas:2009qj,Capasso:2009fh,Chen:2009vu,Kluson:2009xx}. Additionally, application of Ho\\v{r}ava-Lifshitz gravity as a cosmological framework gives rise to Ho\\v{r}ava-Lifshitz cosmology, which proves to lead to interesting behavior \\cite{Calcagni:2009ar,Kiritsis:2009sh}. In particular, one can examine specific solution subclasses \\cite{Lu:2009em,Nastase:2009nk,Colgain:2009fe,Ghodsi:2009rv,Minamitsuji:2009ii,Ghodsi:2009zi,Wu:2009ah,Cho:2009fc,Boehmer:2009yz,Momeni:2009au}, the phase-space behavior \\cite{Carloni:2009jc,Leon:2009rc}, the gravitational wave production \\cite{Mukohyama:2009zs,Takahashi:2009wc,Koh:2009cy,Park:2009gf,Park:2009hg,Myung:2009ug}, the perturbation spectrum \\cite{Mukohyama:2009gg,Piao:2009ax,Gao:2009bx,Chen:2009jr,Gao:2009ht,Cai:2009hc,Wang:2009yz,Kobayashi:2009hh,Wang:2009azb}, the matter bounce \\cite{Brandenberger:2009yt,Brandenberger:2009ic,Cai:2009in,Suyama:2009vy}, the black hole properties \\cite{Danielsson:2009gi,Cai:2009pe,Myung:2009dc,Kehagias:2009is,Mann:2009yx,Bertoldi:2009vn,Park:2009zra,Castillo:2009ci,BottaCantcheff:2009mp,Lee:2009rm,Varghese:2009xm,Kiritsis:2009rx}, the dark energy phenomenology \\cite{Saridakis:2009bv,Park:2009zr,Appignani:2009dy,Setare:2009vm}, the astrophysical phenomenology \\cite{Kim:2009dq,Harko:2009qr,Iorio:2009qx,Iorio:2009ek}, the thermodynamic properties \\cite{Wang:2009rw,Cai:2009qs,Cai:2009ph} etc. However, despite this extended research, there are still many ambiguities if Ho\\v{r}ava-Lifshitz gravity is reliable and capable of a successful description of the gravitational background of our world, as well as of the cosmological behavior of the universe \\cite{Charmousis:2009tc,Li:2009bg,Sotiriou:2009bx,Bogdanos:2009uj,Koyama:2009hc,Papazoglou:2009fj}.  Although the discussion about the foundations and the possible conceptual and phenomenological problems of Ho\\v{r}ava-Lifshitz gravity and cosmology is still open in the literature, it is worth investigating in a systematic way the constraints imposed by observations in a universe governed by Ho\\v{r}ava gravity. Thus, in the present work we use Big Bang Nucleosynthesis conditions, together  with Type Ia Supernovae (SNIa), Baryon Acoustic Oscillations (BAO) and Cosmic Microwave Background (CMB) data, in order to construct the corresponding probability contour-plots for the parameters of the theory. Furthermore, in order to be general and model-independent, we perform our analysis with and without the detailed-balance condition. As we will show, both the detailed-balance and beyond-detailed-balance formulations are compatible with observations, however under tight constraints  on the model parameters.   The paper is organized as follows: In section \\ref{model} we present the basic ingredients of Ho\\v{r}ava-Lifshitz cosmology, extracting the Friedmann equations, and describing the dark matter and dark energy dynamics. In section \\ref{Observational constraints} we constrain both the detailed-balance  and the beyond-detailed-balance formulations from the observational point of view. Finally,  section \\ref{conclusions} is devoted to the summary of the obtained results. ",
        "conclusions": "\\label{conclusions}  In this work we constrained Ho\\v{r}ava-Lifshitz cosmology using observational data. In particular, we considered scenarios where the gravitational sector is forced to satisfy the detailed-balance condition, and also  those where this condition is relaxed. Additionally, we have included the matter and radiation sectors following the usual effective fluid approach. These constructions, which cover the range of Ho\\v{r}ava-Lifshitz cosmology, were confronted with data from BBN, SNIa, CMB and BAO observations.  Our first result is that the detailed-balance formulation of Ho\\v{r}ava-Lifshitz gravity cannot fulfill observational requirements, without the analytic-continuation transformation. Under the analytic continuation we found that Ho\\v{r}ava-Lifshitz cosmology can be compatible with observations, and we presented the corresponding contour-plots on the model parameters. These likelihood-contours impose tight constraints on the model parameters, and the corresponding 1$\\sigma$-bounds are presented in Table \\ref{dblimits}. However, we mention that although analytic continuation is necessary for a realistic cosmology, it can fatally affect the gravitational theory itself, spoiling its initial stability and well-behaving nature. Therefore, the detailed-balance version of Ho\\v{r}ava-Lifshitz cosmology seems rather unlikely to be a robust description of nature.  The version of Ho\\v{r}ava-Lifshitz cosmology in which the detailed-balance condition has been abandoned, is also compatible with observations. We constructed the likelihood-contours for the two involved free parameters, namely the curvature and the dark-radiation coefficients. As we showed, observations lead to strong bounds in these parameters, and the corresponding 1$\\sigma$-allowed ranges are presented in Table \\ref{ndblimits}. This  feature was expected, since the data refer to redshifts in which the novel terms of Ho\\v{r}ava-Lifshitz cosmology are downgraded. However, these terms can have significant cosmological implications prior to nucleosynthesis, which could be probed by recent observations of high energy positrons and electrons by the PAMELA \\cite{pamela1,pamela2} and ATIC \\cite{atic} experiments.   Although the present analysis indicates that Ho\\v{r}ava-Lifshitz cosmology can be compatible with observations, it does not enlighten the discussion about possible conceptual problems and instabilities of Ho\\v{r}ava-Lifshitz gravity,  nor it can interfere with the questions concerning the validity of its theoretical background, which is the subject of interest of other studies. In particular, without a solid theoretical basis, it is not clear whether Ho\\v{r}ava-Lifshitz gravity is able to pass the basic parametrized post newtonian (PPN) tests that any physically interesting gravitational theory should \\cite{PPN,PPN1,PPN2,PPN3}. The present work just faces the problem from the phenomenological point of view, and thus its results can been taken into account only if Ho\\v{r}ava-Lifshitz  gravity passes successfully the aforementioned theoretical tests."
    },
    "0911/0911.4998_arXiv.txt": {
        "abstract": "Employing Eggleton's stellar evolution code with an optically thick wind assumption, we have systematically studied the WD + He star channel of Type Ia supernovae (SNe Ia), in which a carbon-oxygen WD accretes material from a He main-sequence star or a He subgiant to increase its mass to the Chandrasekhar mass. We mapped out the parameter spaces for producing SNe Ia. According to a detailed binary population synthesis approach, we find that the Galactic SN Ia birthrate from this channel is $\\sim$$0.3\\times 10^{-3}\\ {\\rm yr}^{-1}$, and that this channel can produce SNe Ia with short delay times ($\\sim$45$-$140\\,Myr). We also find that the surviving companion stars in this channel have a high spatial velocity ($>$400\\,km/s) after SN explosion, which could be an alternative origin for hypervelocity stars (HVSs), especially for HVSs such as US 708. ",
        "introduction": "Type Ia supernovae (SNe Ia) play an important role in the study of cosmic evolution, especially in cosmology. They have been applied successfully in determining cosmological parameters (e.g. $\\Omega$ and $\\Lambda$; Riess et al. 1998; Perlmutter et al. 1999). It is generally believed that SNe Ia are thermonuclear explosions of carbon-oxygen white dwarfs (CO WDs) in binaries. However, there is still no agreement on the nature of their progenitors (Hillebrandt and Niemeyer 2000; Podsiadlowski et al. 2008; Wang et al. 2008).  Over the past few decades, two families of SN Ia progenitor models have been proposed, i.e. the double-degenerate (DD) and single-degenerate (SD) models. Of these two models, the SD model is widely accepted at present. It is suggested that the DD model, which involves the merger of two CO WDs (Iben \\& Tutukov 1984; Webbink 1984; Han 1998), likely leads to an accretion-induced collapse rather than to an SN Ia (Nomoto and Iben 1985). For the SD model, the companion is probably a main-sequence (MS) star or a slightly evolved subgiant star (WD + MS channel), or a red-giant star (WD + RG channel) (Hachisu et al. 1996; Li and van den Heuvel 1997; Langer et al. 2000; Han and Podsiadlowski 2004, 2006; Chen and Li 2007, 2009; L\\\"{u} et al. 2009; Meng et al. 2009; Wang, Li and Han 2009; Meng and Yang 2009a). Meanwhile, a CO WD may also accrete material from a He star to increase its mass to the Chandrasekhar (Ch) mass, which is known as the WD + He star channel. Yoon and Langer (2003) followed the evolution of a CO WD + He star system, in which the WD can increase its mass to the Ch mass by accreting material from the He star. It is believed that WD + He star systems are generally originated from intermediate mass binaries, which may explain SNe Ia with short delay times implied by recent observations (Mannucci et al. 2006, Aubourg et al. 2008).  The purpose of this paper is to study SN Ia birthrates and delay times of the WD + He star channel and to explore the properties of the surviving companion stars after SN explosion. In Section 2, we describe the numerical code for the binary evolution calculations and the binary evolutionary results. We describe the binary population synthesis (BPS) method and results in Section 3. Finally, a discussion is given in Section 4. ",
        "conclusions": "The simulations give Galactic SN Ia birthrate of $\\sim$$0.3\\times 10^{-3}\\ {\\rm yr}^{-1}$, which is lower than that inferred observationally (i.e. $3 - 4\\times 10^{-3}\\ {\\rm yr}^{-1}$; Cappellaro and Turatto 1997). This implies that the WD + He star channel is only a subclass of SN Ia production, and there may be some other channels or mechanisms also contributing to SNe Ia, e.g. WD + MS channel, WD + RG channel or double-degenerate channel (see Wang, Li and Han 2009; Meng and Yang 2009).  In this paper, we assume that all stars are in binaries and about 50\\% of stellar systems have orbital periods less than 100\\,yr. If we adopt 46.3\\% of stellar systems have orbital periods below 100\\,yr by adjusting the parameter $a_{\\rm 1}$ in equation (7) of Wang et al. (2009b), the SN Ia birthrate from this channel will decrease to be $\\sim 0.28\\times10^{-3}\\,{\\rm yr}^{-1}$.  Hachisu et al. (2008) investigated new evolutionary models for SN Ia progenitors, introducing the mass-stripping effect on a MS or slightly evolved companion star by winds from a mass-accreting WD. The model can also provide a possible ways of producing young SNe Ia, but the model depends on the efficiency of the mass-stripping effect. We also find that the model produces very few young SNe Ia according to a detailed BPS approach. Thus, we consider the WD + He star channel as a main contribution to the formation of young SNe Ia.   US 708 is an extremely He-rich sdO star in the Galaxy halo, with a heliocentric radial velocity of +$708\\pm15$\\,${\\rm km/s}$ (Hirsch et al. 2005). We note that the local velocity relative to the Galatic center may lead to a higher observation velocity for the surviving companion stars, but this may also lead to a lower observation velocity. Considering the local velocity near the sun ($\\sim$220\\,km/s), we find that $\\sim$30\\% of the surviving companion stars may be observed to have velocity $V>700\\,{\\rm km/s}$ for a given SN ejecta velocity 13500\\,km/s. In addition, the asymmetric explosion of SNe Ia may also enhance the velocity of the surviving companions. Thus, a surviving companion star in the WD + He star channel may have a high velocity like US 708 (see also Justham et al. 2008).   For a single starburst, most of the SN explosions occur between $\\sim$45\\,Myr and $\\sim$ 140\\,Myr after the starburst, i.e. SNe Ia from the WD + He star channel will be absent in the old galaxies. In future investigations, we will employ the Large sky Area Multi-Object fiber Spectral Telescope (LAMOST) to search the HVSs originating from the surviving companions of SNe Ia."
    },
    "0911/0911.3740_arXiv.txt": {
        "abstract": "{} {We present a homogeneous and complete catalogue of optical groups identified in the purely flux limited ($17.5 \\leq I_{AB} \\leq 24.0$) VIMOS-VLT Deep redshift Survey (VVDS).} {We use mock catalogues extracted from the MILLENNIUM simulation, to correct for potential systematics that might affect the overall distribution as well as the individual properties of the identified systems. Simulated samples allow us to forecast the number and properties of groups that can be potentially  found in a survey with VVDS-like selection functions. We use them to correct for the expected incompleteness and also to asses how well galaxy redshifts trace the line-of-sight  velocity dispersion of the underlying mass overdensity. In particular, we train on these mock catalogues the adopted group-finding technique \\ie the Voronoi-Delaunay Method (VDM). The goal is to fine-tune its free parameters, recover in a robust and unbiased way the redshift and velocity dispersion distributions of groups ($n(z)$ and $n(\\sigma)$ respectively) and maximize, at the same time, the level of completeness and purity of the group catalogue.} {We identify 318 VVDS groups  with at least 2 members in the range $0.2 \\leq z \\leq 1.0$, among which 144 (/30) with at least 3 (/5) members. The sample has an overall completeness of $\\sim60$\\% and purity of $\\sim50$\\%. Nearly $45$\\% of the groups with at least 3 members are still recovered  if we run the algorithm with the particular parameter set which maximizes the purity ($\\sim75$\\%) of the resulting catalogue. We exploit the group sample to explore the redshift evolution of the fraction $f_b$ of blue galaxies ($U-B \\leq 1$) in the redshift range $0.2\\leq z \\leq1$. We find that the fraction of blue galaxies is significantly lower in groups than in the global population (\\ie in the whole ensemble of galaxies irrespectively of their environment). Both of these quantities increase with redshift, with the fraction of blue galaxies in groups showing a marginally significant steeper increase. We also investigate the dependence of $f_b$ on group richness: not only we confirm that, at any redshift, the blue fraction decreases in systems with increasing richness, but we extend towards fainter luminosities the magnitude range over which  this  result holds.} {} ",
        "introduction": "\\indent Galaxy groups and clusters are the largest and most massive gravitationally bound systems in the universe.  Because of this, they are very useful cosmological probes. For example, the evolution of their abundance or baryon fraction give insights into the value of fundamental cosmological parameters (\\eg \\citealp{borgani1999,newman2002, allen2002, ettori2003, zhang2006_baryon,ettori2009}), their mass and luminosity functions fix the amplitude of the power spectrum at cluster scales (\\eg \\citealp{rosati2002,finoguenov2010}), while their optical mass-to-light ratio allows to constrain the matter density parameter $\\Omega_m$ (\\eg \\citealp{girardi2000, marinoni2002_ML, sheldon2009_ML}).  Groups and clusters are also ideal laboratories for astrophysical studies. Several interesting physical processes are indeed triggered on scales characterized by such extreme density conditions. Their analysis is crucial in particular to understand the effects of local environment on galaxy formation and evolution (e.g. \\citealp{oemler1974, dressler1980, postman1984, dressler1997, garilli1999, treu2003, poggianti2006}).   \\subsection{The detection of galaxy groups and clusters}  A whole arsenal of algorithms allows to identify and reconstruct galaxy systems.  They range from the very first pioneer methods based on visual identification on photometric plates \\citep{abell1958,zwicky1968_clusters} to more recent techniques which exploit various physical properties of the systems as a guide for identification. For example, the thermal bremsstrahlung emission from the hot intracluster gas trapped inside the cluster gravitational potential allows to spot them by means of X-ray band observations. On the opposite side of the spectrum, in the centimetre regime, cluster detection is made possible thanks to the Sunyaev-Zeldovich effect (SZE, \\citealp{sunyaev_zeldovich1972,sunyaev_zeldovich1980}).  Indeed, the hot intracluster gas, by inverse-compton scattering the photons of the Cosmic Microwave Background (CMB), leaves a characteristic imprint in the CMB spectrum which can be exploited as a useful signature for identification.  A cluster potential well can also be detected through strong gravitational lensing or the cosmic shear induced by weak gravitational lensing \\citep{kneib2003,gavazzi2009,limousin2009,richard2010, limousin2010,morandi2010}. Clusters identification can be based also on the properties of the member galaxies. It has been observed that cluster cores host typically red galaxies, among which there are the brightest cluster galaxies (BCG). Thus, a cluster center can be identified as a {\\it R.A.-dec} concentration of galaxies with typical red colours (see for example the Red-Sequence Cluster Survey, \\citealp{gladders2000}, the first cluster survey based on this method), in some case adding also the constraint of a high luminosity (\\eg the {\\it maxBCG} method, \\citealp{hansen2005}, \\citealp{koester2007}).    An orthogonal approach, based on geometrical algorithms, consists in identifying systems from the 3D spatial distribution properties of their members. These algorithms vary from the earlier hierarchical method (\\citealp{materne1978}, \\citealp{tully1980h}) and the widely used `friend of friend' (FOF) method \\citep{huchra1982h}, to the more recent 3D adaptive matched filter method \\citep{kepner1999}, the `C4' method \\citep{miller2005} and the Voronoi-Delaunay Method \\citep[VDM,][]{marinoni2002}.  Finally, group-finding algorithms have also been developed which exploit information extracted from photometric redshifts (\\eg \\citealp{adami2005_CDFS,mazure2007}).   The availability of several identification protocols is not only useful in order to confirm clusters detection by an a-posteriori cross-correlation of various independent catalogues, but it is also crucial for spotting eventual systematics which might affect individual detection techniques.  For example, it has been shown by the first joint X-ray/optical survey \\citep{donahue2002_Xrayopt} that only $\\sim$ 20\\% of the optically selected clusters were also identified in X-rays, while $\\sim$60\\% of the X-ray clusters were detected in the optical sample. Understanding the possible selection effects hidden behind the different survey strategies is crucial in order to interpret this small overlap between the two different cluster catalogues (see for example \\citealp{ledlow2003_Xrayopt, gilbank2004_Xrayopt}). Moreover, using the RASS-SDSS galaxy cluster catalogue \\citet{popesso2004_I} show that a distinct class of `X-ray underluminous Abell clusters' does exist, with an X-ray luminosity $L_X$ which is one order of magnitude fainter than the one expected for their mass according to the typical $L_X$-mass relation \\citep{popesso2007_V}. This supports the concern of \\cite{donahue2002_Xrayopt} about the possible existence of biases in catalogues selected in different wavebands.  A major challenge we face is to extend cluster searches at high redshift.  Indeed, most of the methods described above suffer from major drawbacks when applied in this regime.  Both the X-ray apparent surface brightness and the gravitational lensing cross section of clusters decrease very rapidly with redshift. As a consequence, only very massive clusters can be detected at high $z$. On the contrary, the SZE detection efficiency does not depend on redshift, but large SZ survey are yet to be completed.  For what concerns cluster detection using the spatial distribution of members, we emphasize the difference between photometric and spectroscopic galaxy data sets. Several methods have been proposed to detect clusters with photometric data, mainly exploiting galaxy colours in different bands. On the one side, this method has been successfully used both for surveys (see for example the above-mentioned Red-Sequence Cluster Survey, \\citealp{gladders2000}) and single detections (\\eg the very recent work by \\citealp{andreon2008_z19}), but on the other hand the selection of red galaxies implies the selection of only the older structures, where galaxies lived enough time to be affected by the physical processes typical of the group environment (see for example the discussion in \\citealp{gerke2007_groupsblue}). Moreover, the depth required in photometric surveys to identify high-$z$ groups and clusters increases the number of foreground and background galaxies, due to the fact that objects surface number density is enhanced by the faint flux limit. This essentially limits the effectiveness of 2D identifications at high-$z$.  The third dimensions is thus imperative if we want to disentangle in an efficient way projection effects. Nonetheless, the uncertainty on the line-of-sight position of galaxies may be a concern when it is bigger (or even much bigger) than the typical velocity dispersion of group galaxies, as in the case of photometric redshifts.  \\subsection{This work and the existing groups and clusters samples}  To date, many local, optically selected group catalogues are available in literature. A review can be found in \\cite{eke2004}, where one of the largest catalogue of galaxy groups detected in redshift space from the Two Degree Field Galaxy Redshift Survey (2dFGRS) is presented. Similarly, several group catalogues have been extracted from the Sloan Digital Sky Survey data (\\eg \\citealp{miller2005, berlind2006, weinmann06}). Systematic searches of groups in redshift space have been undertaken also at intermediate redshift (\\eg within the CNOC2 survey, up to redshift $z=0.55$, \\citealp{carlberg2001}). The compilation of optically selected and complete samples of groups up to $z\\sim 1$ and beyond has been made possible only recently thanks to the completion of large and deep spectroscopic surveys, such as the DEEP2 Galaxy Redshift Survey \\citep{Davis2003}, the VIMOS-VLT Deep Survey \\citep{lefevre2005a}, and the zCOSMOS survey \\citep{lilly2007_zcosmos, lilly2009_zcosmos}.  \\cite{gerke2005_groups} present the first DEEP2 group catalogue: it contains 899 groups with two or more members identified in the redshift range $0.7\\leq z \\leq 1.4$ with the VDM method. The DEEP2 sample reaches a limiting magnitude of $R_{AB}=24.1$, and its galaxies are pre-selected in colour before being targeted for spectroscopic observations, in order to reduce the number of galaxies at $z \\lesssim 0.7$. The first zCOSMOS group catalogue \\citep{knobel2009_groups} comprises   $\\sim800$ groups with at least 2 members, covering the redshift range $0.1\\leq z \\leq 1.0$. The parent galaxy sample is purely flux limited ($15 \\leq I_{AB} \\leq 22.5$), and  groups are detected with the FOF method, combined with the VDM.   In this work, we make use of the VIMOS-VLT Deep Survey (VVDS, \\citealp{lefevre2005a}) to compile an homogeneous optically-selected group catalogue in the redshift range ($0.2<z<1.0$). We run the VDM code on a sample containing  more than 6000 flux limited galaxies ($17.5 \\leq I_{AB} \\leq 24.0$) for which reliable spectroscopic redshifts have been measured.  Particular attention has been devoted to optimally tune the parameters of the group-finding algorithm using VVDS-like mock catalogues. The selection function of the sample, essentially compensating only for the flux limited nature of the survey, is simple and mostly insensitive to possibly uncontrolled bias such as those which might affect colour selected samples. Moreover, the magnitude depth of the VVDS allows us to sample a galaxy population which is fainter in luminosity than that currently probed by other flux-limited surveys of the deep universe.   The paper is organized as follows: in \\S \\ref{data_and_mocks} the data sample and the mock catalogues are described.  The reliability of the virial line of sight velocity dispersion as estimated using galaxies is discussed in  \\S \\ref{preliminary_tests}. In \\S \\ref{the_algorithm} we review the basics of the VDM group-finding algorithm, while the strategy followed to fine tune  its parameters is presented in \\S \\ref{VDM_optimization}.  In \\S \\ref{real_VVDS_catalogue} we describe the properties of the VVDS group catalogue. The redshift evolution of the $U-B$ colour of group galaxies is analyzed in \\S \\ref{gal_group_properties}. Conclusions are drawn in \\S \\ref{Conclusions}.  We frame our analysis in the context of a $\\Lambda$ Cold Dark Matter model ($\\Lambda$CDM) specified by the following parameters: $\\Omega_m=0.3$, $\\Omega_{\\Lambda}=0.7$, $H_0=70$ km s$^{-1}$ Mpc$^{-1}$. Magnitudes are expressed in the AB system. ",
        "conclusions": "We have compiled a homogeneous catalogue of optical groups identified in the VVDS-02 field by means of the VDM algorithm, in the range $0.2 \\leq z \\leq 1.0$.  We used mock catalogues simulating the VVDS survey  to optimize the performances of  the group-finding algorithm (maximizing the completeness and the purity of the resulting group catalogue) as well as to minimize possible selection effects. Our main results are here summarized.  \\begin{itemize}  \\item[-] Using the mock catalogues, we verified that the VVDS-02h survey sampling rate allows us to recover at least $50$\\% of the groups (with a virial line of sight velocity dispersion $\\sigma_{vir} \\geq 350$ km/s) that are potentially present in the parent photometric catalogue up to $z=1$.  \\item[-] We tested how well $\\sigma_{vir}$ of the halo mass particles can be estimated using sparsely sampled galaxy velocities. We verified that with this method, given the characteristics of our survey (flux limit, sampling rate, redshift measurement error)  we are able to recover a sensible value of $\\sigma_{vir}$ for $\\sigma_{vir} \\geq 350$ km/s.  \\item[-] Applying the optimized algorithm to the VVDS real data set, we obtained a catalogue of 318 groups of galaxies (with at least two members) in the range $0.2 \\leq z \\leq 1.0$. Among these groups, 63 have a measured line of sight velocity dispersion greater than 350 km/s. The group catalogue is characterized by an overall completeness of $\\sim60$\\% and a purity of $\\sim50$\\%. Nearly $19$\\% of the total population of galaxies live in these systems.  \\item[-] the number density distribution as a function of both redshift ($n(z)$) and velocity dispersion ($n(\\sigma)$) of the VVDS groups with $\\sigma>350$ km/s scales in qualitative agreement with the analogous statistics recovered from the mock catalogues.  \\item[-] We studied the fraction $f_b$ of blue galaxies ($U-B \\leq 1$) in the range $0.2\\leq z \\leq 1$. We used a luminosity-limited subsample of galaxies extracted from our data ($M_B \\leq -18.9-1.1 z$), complete up to $z=1$. We found that $f_b$  is significantly lower in groups than in the global galaxy population. Moreover, $f_b$ increases as a function of  redshift irrespectively of the environment, with marginal evidence for a faster growth rate in groups.  We also analysed how $f_b$ varies as a function of group richness, finding that, at any  redshift explored, $f_b$ decreases in systems with increasing richness. \\end{itemize}  Further explorations of the properties of VVDS groups is left to future works. We only anticipate that the high degree of completeness of the catalogue can be potentially exploited for extracting cosmological information via, for example, cluster counts techniques. The high level of purity  makes the VVDS group sample  ideal also for astrophysical studies which aim at tracing various physical properties of galaxies as a function of local density and environment. We also mention that the cross-correlation studies of our optically-selected catalogue with samples inferred in the same field with independent techniques will help to gain insights not only on cluster selection biases but also on the physics at work within these extreme environments."
    },
    "0911/0911.5399_arXiv.txt": {
        "abstract": "It has been argued recently by Copi \\etals (2009) that the lack of large angular correlations of the CMB temperature field provides strong evidence against the standard, statistically isotropic, inflationary $\\Lambda$CDM cosmology. We compare various estimators of the temperature correlation function showing how they depend on assumptions of statistical isotropy and how they perform on the WMAP 5 year ILC maps with and without a sky cut. We show that the low multipole harmonics that determine the large-scale features of the temperature correlation function can be reconstructed accurately, independent of any assumptions concerning statistical isotropy, from the data that lie outside the sky cuts. The temperature correlation functions computed from our reconstructions are in good agreement with those computed from the whole sky. A Bayesian analysis of the large-scale correlations is presented which shows that the data cannot exclude the standard $\\Lambda$CDM model. We discuss the differences between our conclusions and those of Copi \\etal.   \\vskip 0.1 truein  \\noindent {\\bf Key words}: Methods: data analysis, statistical; Cosmology: cosmic microwave background, large-scale structure of Universe   \\vskip 0.3 truein ",
        "introduction": "Following the discovery by the COBE team of temperature anisotropies in the cosmic microwave background (CMB) radiation (Smoot \\etals 1992; Wright \\etals 1992), Hinshaw \\etals (1996) noticed that the temperature angular correlation function,  $C(\\theta)$,  measured from the  COBE maps was close to zero on large angular scales. This result attracted little attention until the publication of the first year results from WMAP (Bennett \\etals 2003; Spergel \\etals 2003, hereafter S03). The results from WMAP confirmed the lack of large-scale angular correlations in the temperature maps and led S03 to introduce the statistic \\begin{equation} S_{1/2} = \\int_{-1}^{1/2} [C(\\theta)]^2 d\\cos \\theta. \\label{C12} \\end{equation} The form of the statistic and the upper cut-off, $\\mu=\\cos \\theta = 1/2$, were chosen {\\it a posteriori} by S03 `in response' to the observed shape of the temperature correlation function computed using a particular estimator and sky cut (as described in further detail in Section 2). To assess the statistical significance of the lack of large-scale power, S03 computed a `p-value', {\\it i.e.} the fraction of models in their Monte Carlo Markov chains  which had a value of $S_{1/2}^{\\rm model} < S_{1/2}^{\\rm data}$, using the same estimator and sky cut that they applied to the data. For their standard six-parameter inflationary $\\Lambda$CDM cosmology, they found a p-value of $0.15\\%$, suggesting a significant discrepancy between the model and the data.   This problem was revisited by Efstathiou (2004a, herafter E04). The main focus of the E04 paper was to improve on the pseudo-harmonic power spectrum analysis used by the WMAP team (Hinshaw \\etals 2003) by using quadratic maximum likelihood (QML) estimates of the power spectrum, with particular emphasis on the statistical significance of the low amplitude of the quadrupole anisotropy. As an aside, E04 computed angular correlation functions from the QML power spectrum estimates and showed that they were insensitive to the presence of a sky cut. Similarly the $S_{1/2}$ statistic computed from these correlation functions was found to be insensitive to the size of the sky cut giving p-values of a few percent. E04 concluded that the correlation function and $S_{1/2}$ statistic offered no compelling evidence against the concordance inflationary $\\Lambda$CDM model. E04 did not explore in any detail the low p-value for the $S_{1/2}$ statistic reported by S03 (2003), but commented that it was probably simply an `unfortunate' consequence of the particular choice of statistic, estimator and sky cut chosen by these authors (in other words, a result of various {\\it a posteriori} choices).  The CMB temperature correlation function and $S$ statistic have been reanalysed in two recent papers (Copi \\etals 2007; Copi \\etals 2009, heafter CHSS09).  The arguments in the two papers are quite similar, and so for the most part we will refer to the later paper (since, as in this work, it analyses the 5 year WMAP temperature data, Hinshaw \\etals 2009). The Copi \\etals papers are largely motivated by evidence for a violation of statistical isotropy in the WMAP temperature maps, in particular evidence of alignments amongst the low order CMB multipoles ({\\it e.g.} Tegmark, de Oliveira-Costa and Hamilton 2003; Schwarz \\etals 2004; Land and Magueijo 2005a, b), although the statistical significance of these alignments has been questioned (de Oliveira-Costa \\etals 2004; Francis \\& Peacock 2009).  Copi \\etals make the valid point that statistical isotropy is often implicitly assumed in defining what is meant by the term `correlation function' and in defining estimators. They argue further that different estimators contain different information.  They then focus on pixel-based estimates of the correlation function applied to the WMAP data, including a sky cut, and find p-values for the $S_{1/2}$ statistic of $\\sim 0.025$--$0.04\\%$, depending on the choice of CMB map and sky cut. If no sky cut is applied, they find p-values of $\\sim 5\\%$ (similar to the p-values reported in E04). CHSS09 comment that the full-sky results are apparently inconsistent with the cut-sky analysis suggesting a violation of statistical isotropy.  Any analysis which claims to strongly rule out the simple inflationary $\\Lambda$CDM model deserves careful scrutiny, since a confirmed discordance would have profound consequences for our understanding of the early Universe. The purpose of this paper is to investigate carefully the analysis presented in CHSS09. In Section 2 we discuss estimators of the correlation function and relate the pixel-based estimator used by CHSS09 to the pseudo-power spectrum computed on a cut sky. In Section 3, we explicitly reconstruct the individual low order multipole coefficients $a_{\\ell m}$ from cut-sky maps using a technique first applied by de Oliveira-Costa and Tegmark (2006). This allows us to test the sensitivity of the large-angle correlation function to the presence of a sky cut, independent of any assumptions concerning statistical isotropy or Gaussianity. The results of this analysis are compared with the QML estimates of the correlation function used in E04.  Section 4 describes a Bayesian analysis of the $S_{1/2}$ statistic and contrasts it with the frequentist analysis applied by S03 and CHSS09.  Our conclusions are summarized in Section 5. ",
        "conclusions": "The low amplitude of the temperature autocorrelation function at large angular scales has led to some controversy since the publication of the first year results from WMAP. This paper has sought to clarify the following points:  \\smallskip  \\noindent [1] We have compared different estimators of the ACF showing: (a) how they depend on assumptions of statistical isotropy; (b) how they are interrelated; (c) how they perform on the WMAP 5 year ILC maps with and without a sky cut.   \\smallskip  \\noindent [2] The imposition of the KQ85 and KQ75 sky masks leads to little loss of information on the low multipoles that contribute to the large-scale angular correlation function. As demonstrated in Section 3, the low multipole harmonics can be reconstructed accurately, independent of any assumptions concerning statistical isotropy, from the data that lie outside the sky cuts. The ACFs computed from these reconstructions are in good agreement with the ACF computed from the whole sky and in good agreement with the maximum likelihood estimator (\\ref{C13}). There can be little doubt that the large-scale ACF for our realization of the sky is very close to the all-sky results shown in Figures 1, 3 and 6, independent of any assumptions regarding statistical isotropy.  \\smallskip  \\noindent [3] The Bayesian analysis presented in Section 4 shows that the posterior distribution of the $S^T_{1/2}$ is broad and cannot exclude the value $S^T_{1/2} \\sim 49000 \\; (\\mu {\\rm K})^4$ appropriate for the Komatsu \\etals (2009) best fitting inflationary $\\Lambda$CDM model.  The breadth of the posterior distribution $S^T_{1/2}$ distribution shows that it is fairly uninformative statistic and so is not a particularly good discriminator of theoretical models.  \\smallskip  \\noindent [4] Unusually low values of the $S_{1/2}$ statistic are found only if `sub-optimal' ACF estimators are applied to maps that include a Galactic mask. We have argued that the low $p$-values associated with these low values of $S_{1/2}$ are most plausibly a result of {\\it a posteriori} choices of statistic. This seems plausible because: (a)  $S$-type statistics are relatively uninformative and hence sensitive to {\\it a posteriori} choices; (b) the all-sky ACF (which is compatible with the concordance $\\Lambda$CDM model) can be recovered from the data outside the mask and so any physical model for the low $p$-values requires a fortuitous alignment of the temperature field with the Galaxy; (b) the analysis of the local ISW corrected maps presented in Section 4 suggests that any physical model of the low $p$-values requires a precise alignment of local structure with the large-scale potential fluctuations at the last scattering surface.   \\medskip  In summary, the results of this paper suggest that, irrespective of the imposition of Galactic sky cuts or assumptions of statistical isotropy, the large-scale correlations of the CMB temperature field provide unconvincing evidence against the concordance inflationary $\\Lambda$CDM cosmology.     \\vskip 0.1 truein  \\noindent {\\bf Acknowledgments:} The authors acknowledge use of the Healpix package and the Legacy Archive for Microwave Background Data Analysis (LAMBDA). Support for LAMBDA is provided by the NASA Office of Space Science. We are particular grateful to John Peacock and Caroline Francis for allowing us to use their ISW maps.  Yin-Zhe Ma thanks Trinity College Cambridge and Cambridge Overseas Trusts for support. Duncan Hanson is grateful for the support of a Gates scholarship."
    },
    "0911/0911.5350_arXiv.txt": {
        "abstract": "We present high-resolution (R $\\sim$ 40000), high-$S/N$ (20--90) spectra of an extremely metal-poor giant star {\\boos} in the ``ultra-faint'' dwarf spheroidal galaxy (dSph) {\\boo}, absolute magnitude M$_V \\sim -6.3$.  We derive an iron abundance of [Fe/H] = --3.7, making this the most metal-poor star as yet identified in an ultra-faint dSph. Our derived effective temperature and gravity are consistent with its identification as a red giant in \\boo.  Abundances for a further 15 elements have also been determined. Comparison of the relative abundances, [X/Fe], with those of the extremely metal-poor red giants of the Galactic halo shows that {\\boos} is ``normal'' with respect to C and N, the odd-Z elements Na and Al, the iron-peak elements, and the neutron-capture elements Sr and Ba, in comparison with the bulk of the Milky Way halo population having [Fe/H] $\\la$ --3.0.  The $\\alpha$-elements Mg, Si, Ca, and Ti are all higher by $\\Delta$[X/Fe]~$\\sim$~0.2 than the average halo values.  Monte-Carlo analysis indicates that $\\Delta$[$\\alpha$/Fe] values this large are expected with a probability $\\sim$ 0.02. The elemental abundance pattern in Boo--1137 suggests inhomogeneous chemical evolution, consistent with the wide internal spread in iron abundances we previously reported.  The similarity of most of the {\\boos} relative abundances with respect to halo values, and the fact that the $\\alpha$-elements are all offset by a similar small amount from the halo averages, points to the same underlying galaxy-scale stellar initial mass function, but that {\\boos} likely originated in a star-forming region where the abundances reflect either poor mixing of supernova ejecta, or poor sampling of the supernova progenitor mass range, or both. ",
        "introduction": "The discovery and analysis of extremely metal-poor stars (those with [Fe/H]\\footnote{[Fe/H] = log(N(Fe)/N(H))$_{\\rm{star}}$ -- log(N(Fe)/N(H))$_{\\odot}$} ~$<$~--3.0) in extremely low luminosity dwarf spheroidal galaxies is changing our perspective on the early chemical enrichment within these objects, the formation of the outer regions of the Milky Way halo, and the role of dwarf spheroidal galaxies (dSph) within the $\\Lambda$CDM paradigm of the manner in which the Milky Way formed.  Initial observations of the brighter dSph (Helmi et al.\\ 2006, and references therein) led to the conclusion that dSph contained no stars with [Fe/H] $<$ --3.0, while detailed studies of relative abundances (in particular [$\\alpha$/Fe]) showed patterns that were unlike those of Galactic halo stars in the solar neighborhood (Venn et al.\\ 2004, and references therein; see also Tolstoy, Hill, \\& Tosi 2009).  The paucity of brighter Milky Way dSph and their chemical relative abundances seem at odds with the CDM paradigm (Klypin et al.\\ 1999; Moore et al.\\ 1999) and the concept that these systems are the building blocks of at least part of the Galaxy's halo. The recent identification of several ``ultra-faint'' systems, with luminosities many orders of magnitude fainter than those of the classical dSph, through analyses of the imaging data from the Sloan Digital Sky Survey (e.g.~Belokurov et al.\\ 2006, 2007) has revitalized discussions of satellite luminosity and mass functions.  Further, these ultra-faint systems contain extremely low-metallicity stars (Kirby et al.\\ 2008; Norris et al.\\ 2008), perhaps redressing the apparent deficit in brighter dSph.  Arguably even more interesting than their relevance as building blocks is to understand the galaxies themselves. Are they the first objects?  Did they cause reionization?  What do they tell us of the first stars?  What was the stellar Initial Mass Function (IMF) at near zero metallicity?  How are the faintest dSph related to more luminous dSph and the Milky Way?  In a study of several ultra-faint dSph, Kirby et al. (2008) first reported stars with [Fe/H] $ < -3.0 $ (with metallicities as low as [Fe/H] $ = -3.3$), based on moderate-resolution spectra in the range 8300--8500\\,{\\AA}, while Norris et al. (2008) from multi-object spectroscopy of the {\\boo} dSph found a similar result, with the most metal-poor star having [Fe/H] $= -3.4$ based on the CaII K line (3933\\,{\\AA}) in moderate resolution blue spectra.  More recently, Frebel et al.\\ (2009) have obtained the first detailed elemental abundances, based on high-resolution, high $S/N$ spectra, of extremely metal-poor stars in the ultra-faint systems, with observations of Com Ber and U Ma II (M$_V \\sim -4, \\, -5.5$ respectively).  Of the six stars studied, one has [Fe/H] $= -3.2$, and another $-3.1$.  Perhaps the most interesting result of their investigation is that in these stars [$\\alpha$/Fe] is similar to that found in the bulk of Galactic halo stars.  That is to say, these low metallicity stars in these faint satellite galaxies were apparently enriched by SNe from a similar stellar IMF to that which enriched their Milky Way counterparts.  The most metal-poor -- and possibly oldest -- stars in the more luminous dSph also show the same level of enhancement of the $\\alpha$-elements (e.g. Koch et al.\\ 2008). These results stand in some contrast to results at higher abundance, in which regime the dSph member stars have significantly lower levels of [$\\alpha$/Fe] than seen in field halo stars of the same iron abundance (e.g.~Venn et al.~2004), plausibly reflecting the more extended star formation and self-enrichment in the dSph in contrast to the field halo (Unavane, Wyse \\& Gilmore 1996).  The ultra-faint systems have sufficiently low luminosities and low metallicities that one might observe the anticipated signatures of enrichment by single Type II supernovae, perhaps even by Population~III massive stars.  The {\\boo} system was discovered by Belokurov et al.~(2006) and is an ultra-faint dSph (M$_V \\sim -6.3$, luminosity $\\sim 3 \\times 10^4$~L$_\\odot$; Martin et al.~2008) at a distance of $\\sim 65$~kpc. The observed color-magnitude diagram is consistent with an old, metal-poor system.  We identified extremely metal-poor member stars in {\\boo}, together with a large internal dispersion in metallicity, through intermediate-resolution spectroscopy of the Ca II H and K lines (Norris et al.~2008).  We here report follow-up, high-resolution, high-S/N, observations of the particularly interesting star {\\boos} we identified which lies at $\\sim 2$~half-light radii from the center of {\\boo}, and for which we had derived [Fe/H] = --3.4.  In \\S2 we present data obtained for this star using the VLT/UVES system with resolving power R = 40000 and $S/N$ = 20--90.  In \\S3 we report chemical abundances from a model atmosphere analysis of this material for some 16 elements.  {\\boos} has [Fe/H] = --3.7, and relative abundances ([X/Fe]) that are very similar to those of Galactic halo stars of the same [Fe/H].  We discuss the implications of our results in \\S4. ",
        "conclusions": "{\\color{black} \\subsection {{\\boo} Membership}  {\\boos} lies 24{\\arcmin} from the center of {\\boo}, corresponding to 1.9 half-light radii (following Martin et al.\\ 2008).  Its radial velocity (V$_{\\rm r}$ = 99.1~{\\kms}) lies within 4~{\\kms} of the systemic velocity of the system (\\S2.3), while its ionization balances of both {\\ion{Fe}{2}}/{\\ion{Fe}{1}} and {\\ion{Ti}{2}}/{\\ion{Ti}{1}} are consistent with its being a red giant with [Fe/H] = --3.7 (\\S3.2). We refer the reader to the discussion by Norris et al.\\ (2008, Footnote 12), based on the discovery statistics of the HK and HES metal-poor star surveys (Beers et al.\\ 1992; Christlieb et al.\\ 2008), of the likelihood that such an extremely metal-poor giant belonging to the Galactic halo would lie in the direction of {\\boo}: they concluded that 0.02 such stars might be expected.  All of these facts confirm to us that {\\boos} is a member of the system.  \\subsection {Relative Abundances}  Figure~\\ref{Fig:XFe} presents [X/Fe] as a function of [Fe/H], in the range --4.5 $<$ [Fe/H] $<$ --2.55, for 12 representative elements in {\\boos} (the open red circle) and some $\\sim$~30 metal-poor Galactic red giants having high-resolution, high $S/N$, abundance analyses, together with data for dSph systems known to have member stars in the range [Fe/H] $\\la$ --3.1. For the halo stars, for reasons of homogeneity, we restrict the data to the results of the First Stars consortium (Cayrel et al.\\ 2004; Spite et al.\\ 2005; and Fran{\\c c}ois et al.\\ 2007), while for the ultra-faint dSph we plot the results of Frebel et al.\\ (2009; Com Ber and U Ma II, filled red circles) and for the more luminous Sextans dSph the data of Aoki et al.\\ (2009) (filled red triangles)\\footnote{We have modified the literature values to correct for differences in adopted solar abundances between them and the present work.  On the scale of Figure~\\ref{Fig:XFe}, however, this effect is small: for example, if one considers the abundances of Cayrel et al.\\ (2004) and Spite et al.\\ (2005) the mean difference in relative abundances caused by difference in the adopted solar values, over the elements in Table~2, is --0.02 dex.  The maximum absolute difference, for Na, is 0.11 dex.}.  It is important to recall that all of the data in the figure have been determined using 1D model atmospheres and the LTE approximation.  This should be borne in mind when comparison is made with predictions of stellar evolution and galactic chemical enrichment models.  For an appreciation of modifications that need to be made to the present abundances to take into account the role of more realistic 3D models and non-LTE effects we refer the reader to Asplund (2005), and references therein. That said, the question we are interested to address here is the similarity or otherwise between the most metal-poor dSph stars and those of the Galactic halo.  Insofar as we have established that our 1D, LTE techniques reproduce the abundances of Cayrel et al.\\ (see \\S3), Figure~\\ref{Fig:XFe} suffices for our needs.  As one moves from top to bottom in Figure~\\ref{Fig:XFe}, six pairs of related elements are plotted -- representing the CNO group, the light odd-Z elements, the $\\alpha$-elements, the Fe-peak below (Cr and Mn) and above (Co and Ni) iron, and the neutron-capture elements.  Initial inspection of the figure suggests an overall similarity between the relative abundances of {\\boos} and those of the Galactic halo at [Fe/H] $\\sim$ --3.5.  We make the following points: (1) For [C/Fe] and [N/Fe], given the large dispersion in the measured field star abundance ratios, and our large observational errors, the results for Boo-1137 are consistent with those for the Galactic halo ; (2) For [Cr/Fe], [Mn/Fe], [Co/Fe], and [Ni/Fe] (the Fe-peak elements for which we have data) the 1$\\sigma$ error bars all overlap the Cayrel et al.\\ (2004) regression lines (their Table 9), supporting the view that similar processes and enrichment occurred for the material from which {\\boos} and the Galactic halo formed. It also suggests that the present techniques produce results on the same system as those of Cayrel et al.\\ (2004), strengthening the similar conclusion reached in \\S3.2, based on our analysis of the Cayrel et al.~data; (3) For the light odd-Z elements, the {\\boos} [Na/Fe] and [Al/Fe] data overlap those of the halo at the 1$\\sigma$ level; (4) Of the two heavy neutron-capture elements, [Sr/Fe] lies within the values for halo stars, while [Ba/Fe] appears high. Given the complicated trend and scatter in [Ba/Fe] values seen in the figure, and the paucity of stars below [Fe/H] = --3.5, more data would be required to address this issue; and (5) For the representative $\\alpha$-elements, the [Mg/Fe] and [Ca/Fe] values of {\\boos} are both higher, by approximately twice their errors of measurement, than the Cayrel et al.\\ (2004) regression lines of the Galactic halo.  Given that we also have abundances for the $\\alpha$-elements Si and Ti, we examine the final point more closely in Figure~\\ref{Fig:Alphas}, which shows results for [Mg/Fe], [Si/Fe], [Ca/Fe], and [Ti/Fe], together with [$\\alpha$/Fe] (the average of [Mg/Fe], [Ca/Fe], and [Ti/Fe])\\footnote{Here we adopt [Ti/Fe] = ([Ti I/Fe] + [Ti II/Fe])/2, and exclude [Si/Fe] from the average because it is generally based on only one, relatively strong, line.  For each of the other elements five or more lines are available.} as a function of [Fe/H]. Also shown in Figure~\\ref{Fig:Alphas} are the regression lines of Cayrel et al.\\ (2004, their Table 9), supplemented by our linear least squares fits for Ti I, Ti II, and [$\\alpha$/Fe] (not given individually by Cayrel et al.).  There are interesting similarities in the panels of Figure~\\ref{Fig:Alphas}: the most obvious (and relevant for the present discussion) is that all of the {\\boos} relative abundances are larger that those of the Galactic halo at the [Fe/H] value of {\\boos}.  Is this significant?  To address this problem we proceed as follows.  Rows (1)--(6), columns (1)--(7), of Table 4 present the relevant input data: columns (1)--(3) contain the atomic species involved, the RMS scatter for material having [Fe/H] $<$ --3.0 about the regression lines in Figure~6 (from Cayrel et al.\\ (2004) and the present work), and the resulting values of the relative abundances of the Galactic halo at [Fe/H]$_{\\rm Boo-1137}$ = --3.66.  Columns (4)--(7) show for {\\boos} its relative abundance and error, the distance it falls above the halo line, and that distance expressed in units of the star's abundance error.  We then use Monte Carlo analysis to address the following question: if one draws putative stars at random from a gaussian distribution for a Galactic halo having the abundance dispersions in column (2), and superimposes on that a random gaussian error corresponding to the observational errors of {\\boos} in column (5), what fraction of the resulting abundances would be at least as large as the observed distance of {\\boos} above the halo lines (presented in column (6)).  The results are presented in the final column of Table 4.  One sees in rows (1)--(6) that, taken individually, the probabilities of {\\boos} lying above the Galactic halo values are in the range 0.03--0.14, broadly consistent with the results in column (7) of the table. The most stringent condition, as might be expected, comes from the average of the $\\alpha$-elements in the sixth row: the likelihood of finding the observed enhancement of the averaged $\\alpha$-elements is 0.017.  One might also ask the question : what is the probability of {\\it all} of the {\\boos} [Mg/Fe], [Si/Fe], [Ca/Fe], [Ti I/Fe], and [Ti II/Fe] values falling above their respective Galactic halo lines.  The answer, based on the Monte Carlo simulations, is that the fraction of relative abundances as large as seen for the five $\\alpha$ species is 3~$\\times$~10$^{-6}$. This assumes, however, that the abundances of each of the above species is independent of all of the others, which is not the case, given the observed propensity of several of them to often have correlated behaviour.  An example of such correlation is highlighted in Figure~\\ref{Fig:Alphas} where the red star represents the data of Cayrel et al.\\ (2004) for CS22968--014, which for all species fall below the least-squares lines of best fit to the Cayrel et al. data.  Had we performed the same test for CS22968--014 we would have concluded that the likelihood of all of the five species lying so far below the regression lines is 2~$\\times$~10$^{-4}$.  (We note for completeness that McWilliam et al.\\ (1995) also analyzed CS22968--014, and first reported the low relative abundances of its $\\alpha$-elements.)  The reader will see other examples of this effect in Figure~\\ref{Fig:Alphas}. Given such correlated behavior, we shall not consider further the test in this paragraph, which is inappropriate.  The reader may recall from \\S3.2 that we found small systematic differences between the relative abundances of Cayrel et al.\\ (2004) and those we obtained using our techniques.  Those relevant here are $\\Delta$[Mg I/Fe] = --0.049, $\\Delta$[Ca I/Fe] = --0.008, $\\Delta$[Ti I/Fe] = 0.033, and $\\Delta$[Ti II/Fe] = --0.019, in the sense [Cayrel et al.\\ -- present work].  If we adjust the results in column (4) for these four species and [$\\alpha$/Fe] to take the corrections into account, the fraction of Monte Carlo simulations having [$\\alpha$/Fe] lying at or above the observed {\\boos} values then increases to 0.024.  The above considerations depend on a knowledge of the dispersions in relative abundance in the Galactic halo.  In the range --4.1 $<$ [Fe/H] $<$ --3.1, Cayrel et al.\\ (2004) report dispersions for [Mg/Fe], [Si/Fe], [Ca/Fe], and [Ti/Fe] about their regressions against [Fe/H] of 0.11, 0.20, 0.11, and 0.09 dex, respectively.  Inspection of Figure~\\ref{Fig:Alphas} suggests that given the decreasing sample size as one moves to lower abundances one should proceed with caution.  For example, there are only five objects in the figure that have [Fe/H] $<$ --3.5.  More data are clearly needed before one regard these estimates as definitive.  With this caveat in mind, we note that these dispersions for [Fe/H] $<$ --3.1 are somewhat larger than reported for metal-poor dwarfs in the range --3.0 $\\la$ [Fe/H] $\\la$ --2.0: by Magain (1989) -- ``extremely small (if any)'', by Nissen et al.\\ (1994) -- $\\sigma$[Mg/Fe] = 0.06 dex, and by Arnone et al.\\ (2005) -- $\\sigma$[Mg/Fe] = 0.06 dex.  As discussed by these authors, and also by Argast et al.\\ (2002), the dispersion of these elements at lowest abundance places strong constraints on the yields of SNe, the IMF, and galactic chemical enrichment at the earliest times.  We shall return to the implications of this point below.}  We note in concluding this section that Feltzing et al.\\ (2009) have recently reported an anomalously large value of [Mg/Ca] = 0.73 (at [Fe/H] = --2.0) for one of seven stars they have observed in {\\boo} (all with [Fe/H] $>$ --3.0).  The value we obtain for {\\boos} is [Mg/Ca] = --0.05, which is not too dissimilar from the mean value of 0.11 $\\pm$ 0.06 that one obtains for their other six stars.  \\subsection{The Evolution and Chemical Enrichment of {\\boo}}  What are the implications from Figures~\\ref{Fig:XFe} and \\ref{Fig:Alphas} for the manner in which chemical enrichment occurred in {\\boo}?  As noted in \\S4.2 above, the initial impression of overall similarity between the abundances of {\\boos} and those of the Galactic halo in Figure~\\ref{Fig:XFe} is perhaps not too surprising. Intuitively, one might expect that for abundances as low as [Fe/H] = --3.7, and an old stellar population, it is more likely to find abundance patterns driven by enrichment from core-collapse supernovae from massive-star progenitors, without the later modification by Type SN Ia\\footnote{We implicitly ignore here possible abundance contamination effects that might result from mass transfer in a binary system as, for example, is believed to have occurred in the CEMP-s class of metal-poor stars (see Beers \\& Christlieb 2005).}.  This of course will depend on the rate and duration of star formation -- only the earliest stars will have enrichment from only core-collapse supernovae, and, should galactic chemical enrichment occur for long enough, one will see the SN~Ia signature downturn of [$\\alpha$/Fe] vs [Fe/H] to more solar-like values as one moves from lower to higher [Fe/H].  In the Sculptor and Draco dSph, for example, the onset of the decrease is evident already at [Fe/H] $\\sim$ --2.0, compared with --1.0 for the Galactic halo (see e.g. Tolstoy et al.\\ 2009, their Figure 11, and Cohen \\& Huang 2009, respectively).  The chemical evolution of {\\boo} probably involved a relatively short-lived epoch of star formation and self-enrichment in a dark matter dominated potential, terminated by catastrophic gas loss in (Type~II)-supernova-driven winds (Saito 1979; Wyse \\& Silk 1985; Dekel \\& Silk 1986) in which the bulk of the initial (gaseous) baryonic mass was lost. As noted earlier, the color-magnitude diagram of {\\boo} (Belokurov et al.\\ 2006) is consistent with an old, metal-poor population.  Forming stars would then have been chemically enriched by only core-collapse supernovae, resulting, for example, in enhanced [O/Fe] and [$\\alpha$/Fe] compared with the solar value.  The actual value of the relative enhancement depends on the mix of masses of the SNe progenitors, since model SNe yields show that more massive progenitors produce relatively more intermediate-mass elements than iron for many elements, so that an IMF biased towards the most massive stars will, with good sampling of the IMF and good mixing so that an IMF-average is achieved in star-forming regions, provide higher mean [O/Fe], [$\\alpha$/Fe], etc. (see e.g.~Wyse \\& Gilmore 1992; Argast et al.\\ 2002; Kobayashi et al.\\ 2006).  The low stellar-mass and low level of enrichment of {\\boo} mean that it would not be surprising if either the underlying IMF were not well-sampled in star-forming regions, or that they were not well-mixed, or both.  Our demonstration that the $\\alpha$-elements appear enhanced with respect to iron in Boo-1137, relative to the mean of the metal-poor halo stars, by $\\Delta[\\alpha$/Fe] $\\sim$ 0.2 is consistent with enrichment of the material from which {\\boos} formed being biased towards the ejecta of SNe with more massive progenitors than the bulk of the metal-poor Galactic halo. This could reflect a biased underlying galaxy-scale IMF, or incomplete mixing and/or poor sampling of an invariant IMF, or be due to the shorter lifetimes of the most massive core-collapse progenitors in an invariant IMF.  We also know that {\\boo} exhibits a large range in heavy element abundance -- from samples of 16 and seven members of {\\boo} Norris et al.\\ (2008) and Feltzing et al.\\ (2009) report ranges of 1.7 and 0.9 dex, respectively.  Assuming a normal underlying galaxy-scale mass function, for which the stellar mass at $\\sim 10$~Gyr after star formation burst is $\\sim 50$\\% of the stars formed, and for which there is one $\\sim 10$~M$_\\odot$ SN progenitor for every $\\sim 100$~M$_\\odot$ of stars formed, and one 25~M$_\\odot$ SN progenitor for every $\\sim 1000$~M$_\\odot$ formed, we can envisage that the early evolution of {\\boo} formed $\\sim 10^5~$M$_\\odot$ in stars and $\\sim 200$~Type II supernovae.  Nissen et al.\\ (1994) argued that the small dispersion, $\\sigma$[Mg/Fe] = 0.06, they obtained for the Galactic halo is consistent with evolution in a well-mixed region with enrichment from some 25 SNe in the mass range 13--40~M$_{\\odot}$ in a ``preceding generation of totally about 2 $\\times$ 10$^{4}$ stars''. {\\color{black} (From their Table 8, one would infer that the number of stars would be smaller by a factor of a few for the somewhat larger dispersions reported by Cayrel et al.\\ (2004) and discussed above.)}. Given the resultant small number of cells this would imply if applied to {\\boo}, it is not difficult to envisage incomplete mixing and poor sampling of the IMF across {\\boo}, to give the apparent higher values of [$\\alpha$/Fe] in {\\boos}.  This is the most conservative of the three suggested reasons behind the elemental ratio enhancements we gave above, and we advocate this conclusion. Individual star-forming regions may well be internally mixed, while still poorly sampling the core-collapse supernovae IMF. A prediction of this would be that as more data are collected, one will observe the signature of a small number of internally well-mixed cells having preferred abundances.  There remains the objection that these conclusions are based on results for just one star. Statistical chance in selecting one star from a possibly diverse parent population, or in enriching one new star from an inhomogeneous star forming region, may undermine our conclusions. The analyses of larger samples of extremely metal-poor stars in {\\boo} and in other dSph are awaited with much anticipation."
    },
    "0911/0911.2570_arXiv.txt": {
        "abstract": "{We have employed a penumbral model, that includes the Evershed flow and convective motions inside penumbral filaments, to reproduce the azimuthal variation of the net circular polarization (NCP) in sunspot penumbrae at different heliocentric angles for two different spectral lines. The theoretical net circular polarization fits the observations as satisfactorily as penumbral models based on flux-tubes. The reason for this is that the effect of convective motions on the NCP is very small compared to the effect of the Evershed flow. In addition, the NCP generated by convective upflows cancels out the NCP generated by the downflows. We have also found that, in order to fit the observed NCP, the strength of the magnetic field inside penumbral filaments must be very close to 1000 G. In particular, field-free or weak-field filaments fail to reproduce both the correct sign of the net circular polarization, as well as its dependence on the azimuthal and heliocentric angles.} ",
        "introduction": "%  Several investigations have proposed the presence of convective motions within the sunspot penumbra (Danielson 1961, Grosser 1989, M\\'arquez et al. 2006, Langhans 2006, S\\'anchez Almeida 2005, 2006, S\\'anchez Almeida et al. 2007), but only very recently have those motions been observationally pinpointed as occurring within penumbral filaments (Ichimoto et al. 2007; Rimmele 2008; Zakharov et al. 2008; cf. Bellot Rubio et al. 2005). Zakharov and co-workers have found that these convective flows appear similar to the upper part of convective rolls proposed by Danielson (1961), with an upflow at the filament's center that turns into downflowing lanes at its edges. The measured speed of these motions is about $1$\\kms. Superposed to this convective flow is the Evershed flow, with typical speeds of about $4-5$\\kms, although much larger values have been reported (del Toro Iniesta et al. 2001; Penn et al. 2003; Bellot Rubio et al. 2004; Borrero et al. 2005; S\\'anchez Almeida 2005). Recent 3D MHD simulations (Scharmer et al. 2008a; Rempel et al. 2009) suggest a relation between these two velocity fields, with the Evershed flow being formed by the deflection of the convective flow along the horizontal magnetic field inside penumbral filaments. \\\\  It is also well established that as the observer's line-of-sight penetrates through the penumbral ambient field and into the penumbral filament, the magnetic inclination and line-of-sight velocity undergo large variations. These are widely accepted as being responsible for creating the anomalous (i.e. asymmetric or even multi-lobed) polarization profiles observed in the penumbra (S\\'anchez Almeida \\& Lites 1992, Solanki \\& Montavon 1993; see Solanki 2003 for a review). Models incorporating such variations have successfully reproduced the azimuthal and Center-to-Limb variation of the net circular polarization (NCP) in visible and infrared Fe I lines (S\\'anchez Almeida 1996, 2005; Mart{\\'\\i}nez Pillet 2000, 2001; Schlichenmaier \\& Collados 2002; Schlichenmaier et al. 2002; M\\\"uller et al. 2002; Borrero et al. 2007). \\\\  However, the effect that the newly-discovered convective component of the velocity field inside penumbral filaments has on the net circular polarization (azimuthal and center-to-limb variation) has not been studied. The convective flow can potentially have important consequences for the NCP observed at disk center, or at all disk positions at locations perpendicular to  the line-of-symmetry of the sunspot. In both cases the Evershed flow is almost perpendicular to the line-of-sight, which should enhance the contribution of the convective flow. Furthermore, the NCP generated by the convective flow could be detected by spectropolarimeters operating at extremely high spatial resolution (Scharmer et al. 2008b) and it could be related to the non-zero NCP observed at the edges of penumbral filaments (Ichimoto et al. 2008).  In this paper we address this possibility and study the effect of the combined magnetic and convective flow pattern reported by Zakharov et al. inside penumbral filaments, on the azimuthal and center-to-limb variations of the net circular polarization in sunspot penumbrae. ",
        "conclusions": "%  We have developed a magnetohydrostatic model of a penumbral filament embedded in a surrounding potential field. The MHS equilibrium imposes a density, pressure and temperature structure inside the penumbral filaments such that the $\\tau=1$-level is formed inside the filament. Consequently, we do not need to specify its sub-surface structure,  which could be in the form of a flux-tube (filament with circular cross section) or in the form of a vertically elongated plume. Inside the filament we assume the presence of the Evershed flow along its axis and of a convective velocity field perpendicular to it. The filament's magnetic field is imposed to be homogeneous in its interior.  By means of Stokes radiative transfer calculations, we have shown that this model is able to reproduce the observed azimuthal variation of the net circular polarization $\\mathcal{N}(\\Psi)$, observed at different heliocentric angles for two different (visible and near-infrared) Fe I lines.  We have also studied the effect of the convective velocity field on the generated $\\mathcal{N}(\\Psi)$-curves. We have found that its effect is much smaller than the NCP generated by the Evershed flow. In addition, the NCP generated by the convective downflows ($\\mathcal{N}>0$) partly cancels with the NCP generated by the upflow at the filament's center ($\\mathcal{N}<0$).  Finally, we have employed our model to study the NCP generated by field-free gaps (Spruit \\& Scharmer 2006) and have found that this model does not reproduce satisfactorily the observed NCP. For that to happen, the magnetic field inside the filament should be around 1000 G, which is not compatible with the concept of a field-free gap.  Our results do not, by themselves, rule out the field-free gap model, since the model employed here is still rather simple, although it does account for the main features of the penumbral fine structure. A more elaborate model based on field-free gaps could still yield NCP curves closer to the observed ones.  In summary, our investigation confirms that the net circular polarization is produced mainly by the Evershed flow in filaments filled with a rather strong horizontal field, and embedded in an inclined magnetic field, as originally proposed by Solanki \\& Montavon (1993) and worked out in greater detail by Mart{\\'\\i}nez Pillet (2000), M\\\"uller et al. (2002), Schlichenmaier et al. (2002), Borrero et al. (2007) and others.  More elaborate models are already available thanks to recent 3D MHD simulations (Sch\\\"ussler \\& V\\\"ogler 2006, Heinemann et al. 2007, Rempel et al. 2008, 2009). These simulations, reveal a complex picture that shares similarities and differences with both the flux-tube and the gappy penumbral model (see Borrero 2009, Schlichenmaier 2009). In a next step it is important to introduce non-grey radiative energy transfer into such simulations, so that similar analyses as carried out here can be performed."
    },
    "0911/0911.3817_arXiv.txt": {
        "abstract": "A gap in the total solar irradiance (TSI) measurements between ACRIM-1 and ACRIM-2 led to the ongoing debate on the presence or not of a secular trend between the minima preceding cycles 22 (in 1986) and 23 (1996). It was recently proposed to use the SATIRE model of solar irradiance variations to bridge this gap. When doing this, it is important to use the appropriate SATIRE-based reconstruction, which we do here, employing a reconstruction based on magnetograms. The accuracy of this model on months to years timescales is significantly higher than that of a model developed for long-term reconstructions used by the ACRIM team for such an analysis. The constructed `mixed' ACRIM~--- SATIRE composite shows no increase in the TSI from 1986 to 1996, in contrast to the ACRIM TSI composite. ",
        "introduction": "\\label{intro}  Solar total irradiance has been measured by a number of space-borne instruments without interruptions since 1978 \\citep{froehlich-2006}. However, no single instrument remained operational over the whole period, and a cross-calibration and construction of a composite of measurements by several radiometers is not trouble-free. A detailed description of individual data sets, their problems and corrections is given by \\citet{froehlich-2006}, cf. \\citet{scafetta-willson-2009}. As a consequence, three different composites of the total solar irradiance (TSI) have been constructed: PMOD \\citep{froehlich-2006}, ACRIM \\citep{willson-mordvinov-2003}, and IRMB \\citep{dewitte-et-al-2004}.  The main disagreement concerns the TSI levels during the minima preceding cycles 22 (which took place in 1986) and 23 (in 1996) and thus the presence of a long-term trend during this period. The origin of the discrepancy lies in the fact that the launch of the ACRIM-2 experiment on UARS, originally planned to overlap with ACRIM-1 on SMM, had to be postponed. Thus for more than 2 years (July 1989 to October 1991; the so-called ACRIM gap) only the data from HF/Nimbus7 and ERBS/ERBSE are available. The ERBS irradiance measurements were made only approximately every 2 weeks. In September 1989, the HF radiometer was switched off for several days due to saturation of its output signal. When it returned to normal operating mode, it apparently showed an abrupt increase in irradiance by about 0.4~W\\,m$^{-2}$ \\citep[for more details, see][]{lee-et-al-1995,chapman-et-al-96,lockwood-froehlich-2008}. There is also an indication for another either abrupt or gradual increase of a similar magnitude \\citep{lee-et-al-1995,froehlich-2006}. This change in the HF sensitivity was allowed for in the PMOD composite but not in the ACRIM one, which led to the disagreement in the long-term behaviour mentioned above. The ACRIM composite shows an increase in the TSI of 0.037\\% (about 0.5~W\\,m$^{-2}$) between the minima of 1986 and 1996 \\citep{willson-mordvinov-2003}, whereas the PMOD composite gives nearly the same TSI values (difference less than 0.06~W\\,m$^{-2}$ or 0.004\\%) for both.  Recently, \\citet{scafetta-willson-2009} have proposed to use the SATIRE model of TSI variations, in order to bridge the ACRIM gap and thus estimate the long term change independently. The SATIRE models \\citep{solanki-et-al-2005a,krivova-solanki-2008a} calculate variations of solar irradiance from the solar surface magnetic field evolution. Several versions of the model have been developed for different applications. Each version is optimised for a different time scale. For reconstructions restricted to the period after 1974, direct full-disc measurements of the solar photospheric magnetic field, i.e. magnetograms, can be employed together with continuum images to identify magnetic features (i.e. sunspots, faculae and the network) \\citep{krivova-et-al-2003a,wenzler-et-al-2004a,wenzler-et-al-2005a,% wenzler-et-al-2006a}. The brightnesses of the individual types of features are time-independent and are calculated from semi-empirical model atmospheres \\citep{unruh-et-al-99}. Variations of TSI result from the emergence, evolution and decay of magnetic features, which is obtained from the magnetograms and continuum images. Thus variability on all considered time scales is modelled in a self-consistent way, with no additional assumptions on a long-term trend. In the following we refer to this version of the model as SATIRE-S, where the S stands for `Satellite era'. This model is indeed well suited for a self-contained test of the irradiance trends around the ACRIM gap.  In contrast, the model SATIRE-T (where the T stands for Telescope era) by \\citet{krivova-et-al-2007a} that has been employed by \\citet{scafetta-willson-2009} is not suited for such an analysis. It has been developed in order to provide an insight into irradiance changes on decadal to centennial time-scales and is, by design, significantly less accurate, particularly on time scales of months. The main challenge of the longer-term models is that the only direct proxy of solar activity going back to the Maunder minimum is the sunspot number. The evolution of the bright magnetic elements (faculae, network) has to be derived somehow from the sunspot record. For this, \\citet{krivova-et-al-2007a} use the coarse physical model developed by \\citet{solanki-et-al-2000,solanki-et-al-2002}, which allows a reconstruction of the solar photospheric magnetic field from the sunspot number. Unfortunately, the evolution of the facular and network components cannot be recovered on a daily basis using such a model: only averages over several months can be considered. This is clear from the Eqs.~(1--7) of that work and the corresponding values of the involved time scales for the decay and flux transfer in active and ephemeral regions (about 3~months to years). Note that variations of the irradiance on time scales of days to weeks is still well reproduced by this model, since they are driven by the evolution of sunspots, which are adequately represented by the daily sunspot number record.  Thus, by its very conception, the model by \\citet{krivova-et-al-2007a} employed by \\citet{scafetta-willson-2009} is expected to be relatively accurate on time scales of days to the solar rotation and from the solar cycle to centuries, whereas its accuracy on time scales of several months to years is limited. This is exactly opposite to what \\citet{scafetta-willson-2009} assumed. Here we repeat their analysis employing the more accurate SATIRE-S model by \\citet{wenzler-et-al-2006a,wenzler-et-al-2009a} based on quasi-daily full-disc magnetograms and continuum images of the Sun recorded by the National Solar Observatory, Kitt Peak (NSO/KP) between 1974 and 2003. We show that employing a more appropriate model gives a rather different result than obtained by  \\citet{scafetta-willson-2009}. We stress that it is not the aim of this paper to produce a new TSI composite, but rather to show that the approach taken by \\citet{scafetta-willson-2009} is completely consistent with a deeper minimum preceding cycle 23 than that preceding cycle 22, if an appropriate model is used. ",
        "conclusions": "We have compared the SATIRE-S model of the TSI variations with the measurements by the ACRIM-1 and ACRIM-2 experiments, in order to bridge the so-called ACRIM gap (July 1989 to October 1991), as proposed by \\citet{scafetta-willson-2009}. This gap is the source of the on-going debate about the presence of the secular variation in solar irradiance between the minima of cycle~22 and 23. The SATIRE-S model calculates the TSI variations from the continuously evolving distribution of the solar surface magnetic field obtained from NSO/KP magnetograms and continuum images \\citep{wenzler-et-al-2009a} covering the period 1974--2002. In contrast to the SATIRE-T model by \\citet{krivova-et-al-2007a} employed by \\citet{scafetta-willson-2009}, in this model variations on all covered time scales are modelled in a self-consistent way, with no additional assumptions regarding the long term trend. Also the accuracy of the SATIRE-S model is significantly higher than that of the SATIRE-T model since it uses direct measurements of the solar photospheric magnetic flux rather than its modelled evolution. Thus it is best suited for such a test.  The constructed `mixed' ACRIM-1~-- WSK09~-- ACRIM-2 composite does not show an increase in the TSI from 1986 to 1996, in contrast to the ACRIM composite. Independently of the value of the model's free parameter, a slight decrease is found. The magnitude of this decrease cannot be estimated very accurately from such an analysis (and therefore such a `mixed' composite should not be considered as a replacement of real measurements), but it lies between approximately 0.15 and 0.7~W\\,m$^{-2}$ (0.011--0.05\\%) for different values of the model's single free parameter. Note that irradiance changes due to non-magnetic effects, if any, cannot be revealed by either SATIRE-S used here nor by SATIRE-T employed by \\citet{scafetta-willson-2009}."
    },
    "0911/0911.1812_arXiv.txt": {
        "abstract": "Using archival, high-resolution far-ultraviolet \\textit{HST}/STIS spectra of 34 Galactic O and B stars, we measure \\ion{C}{1} column densities and compare them with measurements from the literature of CO and \\htwo\\ with regard to understanding the presence of translucent clouds along the line-of-sight.  We find that the CO/\\htwo\\ and CO/\\ci\\ ratios provide good discriminators for the presence of translucent material, and both increase as a function of molecular fraction, $f^N = 2N(H_2)/N(H)$.  We suggest that sightlines with values below CO/\\htwo\\ $\\approx10^{-6}$ and CO/\\ci\\ $\\approx1$ contain mostly diffuse molecular clouds, while those with values above sample clouds in the transition region between diffuse and dark. These discriminating values are also consistent with the change in slope of the CO v. \\htwo\\ correlation near the column density at which CO shielding becomes important, as evidenced by the change in photochemistry regime studied by Sheffer et al. (2008).  Based on the lack of correlation of the presence of translucent material with traditional measures of extinction we recommend defining 'translucent clouds' based on the molecular content rather than line-of-sight extinction properties. ",
        "introduction": "In the traditional three phase interstellar medium (ISM), the cold neutral medium (CNM) contains the bulk of the interstellar ``clouds''.  These clouds range from the dense dark clouds, the centers of which are the birthplaces of the next generation of stars, typically observed by CO radio emission, to the tenuous diffuse clouds, which can be probed by absorption spectroscopy of background stars.  In the literature, these are typically defined by their extinction properties, with the diffuse ISM having \\av$\\lesssim1$ and the dense clouds having \\av$\\gtrsim5$ \\citep{vDB88}.  This distinction by extinction indicates the importance of the ultraviolet portion of the interstellar radiation field on the photochemistry in the clouds.  The diffuse ISM is fully exposed, and therefore most species exist in atomic or ionized forms.  In dense clouds the ultraviolet light is sufficiently shielded such that molecules are not photodissociated and virtually all the hydrogen exists in molecular form (\\htwo), with carbon monoxide (CO) the second most abundant molecule.  Sightlines with intermediate extinction tend to be referred to as ``translucent''.  A translucent sightline may simply be the concatenation of multiple diffuse clouds, resulting in a higher value of \\av\\ than is typical of a sightline that might pass through a single diffuse cloud.  One may also sample the transition region between the diffuse and dense clouds, where species make the transition between their atomic and molecular forms -- what could be considered a translucent ``cloud''.  In these regions the chemistry is sensitive to the physical conditions, such as density, temperature and radiation field, which may be changing rapidly from one location to the next.  With increased depth into a cloud comes an increased fraction of molecular content, both in \\htwo\\, because of self-shielding, and CO, because of shielding by \\htwo, as well as the attenuation of photodissociating flux by dust \\citep{vDB88}.  In the case of hydrogen, the transition from atomic to molecular form is quite sharp, with the local molecular fraction ($f^n = 2n(H_2)/n_H$, $n\\equiv$ space density) reaching essentially a value of one at very short depths into a cloud.  Observations of the integrated molecular fraction ($f^N = 2N(H_2)/N(H)$, $N\\equiv$ column density) show a sharp jump from values less than $\\sim10^{-4}$ to values larger than $\\sim0.01$ consistent with the presence of this transition \\citep{Spitzer75,Gillmon06}, and as a consequence \\htwo\\ is ubiquitous in the diffuse ISM \\citep{Shull00}. The gas-phase carbon makes a transition from ionized form (\\cii) to molecular (CO).  Although the calculated depth at which this transition happens and the relative abundances of all the carbon-bearing species depends on the physical parameters and chemical networks used, this qualitative description appears in all models \\citep{Rollig07}.  The far-ultraviolet (far-UV: $\\lambda\\lesssim2000$\\AA) portion of the electromagnetic spectrum allows for the direct observation of both \\htwo\\ and CO in the ISM through absorption spectroscopy towards distant O and B stars \\citep{Savage77,Rachford02,Federman80,Burgh07,Sheffer08}.  In particular, the study of \\citet{Rachford02} investigates the molecular content and physical properties along lines-of-sight considered translucent by their extinction properties, reaching the conclusion that they are likely the projections of multiple diffuse clouds with low extinction rather than a single translucent cloud with higher extinction. \\citet{Burgh07} suggested that the ratio of CO to \\htwo\\ would serve as a better discriminator of diffuse and translucent clouds and they  measure a transition from low to high values of CO/\\htwo\\ with increased molecular fraction.  The CO abundance is more sensitive to the effects of geometry, dust shielding and fragmentation of clouds, and thus if a given sightline were simply a collection of diffuse clouds, there would not be an increase in the measured CO/\\htwo\\ \\citep{Kopp00}.  If the term ``translucent'' should refer to the transition region between the physical states of ``diffuse'' and ``dense'' then these results demonstrate a weakness in defining translucent clouds based solely on line-of-sight extinction or \\htwo\\ properties.  \\citet{Snow06} recommend a different definition for ``translucent'', based upon the carbon content, specifically the transition from ionized to molecular form.  This suggests that observations of neutral carbon (\\ci), also available in the far-UV \\citep{Jenkins79,Jenkins01}, may give further insight as to the presence of translucent clouds along a given sightline.  Steady-state PDR models predict that \\ci\\ is formed through recombination of \\cii\\ and photodissociation of CO, and its abundance should peak in the region between the ionic and molecular parts of the cloud.  \\citet{Snow06} then further break down cloud structure into the following categories (see their Table 1): diffuse atomic, where the low molecular fractions mentioned earlier are observed; diffuse molecular, where the hydrogen is primarily molecular, but carbon still in ionized form; translucent, where the carbon makes the transition to molecular; and dense, where both the hydrogen and carbon are fully molecular.  We agree with this type of categorization, eschewing the use of extinction as the sole discriminator.  In this work, we expand upon the study of \\citet{Burgh07} by including measurements of the \\ci\\ lines observed in the \\textit{HST}/STIS E140H mode.  We compare these measurements to CO and \\htwo\\ in an attempt to better isolate the transition from diffuse to translucent and better understand the conditions under which this transition occurs. ",
        "conclusions": "Based upon the results reported here, we suggest that the definition of ``translucent'' based on measures of extinction be replaced with one based on molecular content, primarily CO.  In the range of \\av\\ sampled in this study, extinction properties give no indication of the physical or chemical conditions in the clouds sampled along the line of sight.  We believe this to result primarily from the facts that extinction correlates better with total hydrogen rather than molecular hydrogen and that the penetration of ultraviolet radiation is very sensitive to cloud structure \\citep{Kopp00}.  These observations provide good constraints for models of the interstellar medium.  Significant column densities of CO, and high values of CO/\\htwo, even at relatively low extinction, suggest that the interstellar medium cannot be too fragmented and porous; i.e., translucent material that is sufficiently shielded from photodissociating far-UV radiation exists in increasing abundance along sightlines with high molecular fraction.  These sightlines cannot simply be concatenations of multiple diffuse clouds.  Finally, we point out that, like CO/\\htwo, CO/\\ci\\ can be a good measure of the presence of translucent material along a given line-of-sight. Although the Cosmic Origins Spectrograph aboard the \\textit{Hubble Space Telescope} has some ability to measure \\htwo\\ along some sightlines \\citep{McCandliss09}, both it and the Space Telescope Imaging Spectrograph are optimized for performance in the 1150--1550\\AA\\ range, where CO and \\ci\\ both have many absorption features and may be measured simultaneously."
    },
    "0911/0911.1165_arXiv.txt": {
        "abstract": "We analyze the temporal variation of the diurnal anisotropy of sub-TeV cosmic ray intensity observed with the Matsushiro (Japan) underground muon detector over two full solar activity cycles in 1985-2008. We find an anisotropy component in the solar diurnal anisotropy superimposed on the Compton-Getting anisotropy due to the earth's orbital motion around the sun. The phase of this additional anisotropy is almost constant at $\\sim$15:00 local solar time corresponding to the direction perpendicular to the average interplanetary magnetic field at the earth's orbit, while the amplitude varies between a maximum (0.043$\\pm$0.002 $\\%$) and minimum ($\\sim$0.008$\\pm$0.002 $\\%$) in a clear correlation with the solar activity. We find a significant time lag between the temporal variations of the amplitude and the sunspot number and obtain the best correlation coefficient of +0.74 with the sunspot number delayed for 26 months. We suggest that this anisotropy might be interpreted in terms of the energy change due to the solar-wind induced electric field expected for GCRs crossing the wavy neutral sheet. The average amplitude of the sidereal diurnal variation over the entire period is 0.034$\\pm$0.003 $\\%$, which is roughly one third of the amplitude reported from AS and deep-underground muon experiments monitoring multi-TeV GCR intensity suggesting a significant attenuation of the anisotropy due to the solar modulation. We find, on the other hand, only a weak positive correlation between the sidereal diurnal anisotropy and the solar activity cycle, in which the amplitude in the ``active'' solar activity epoch is about twice the amplitude in the ``quiet'' solar activity epoch. This implies that only one fourth of the total attenuation varies in correlation with the solar activity cycle and/or the solar magnetic cycle. We finally examine the temporal variation of the ``single-band valley depth'' (SBVD) quoted by the Milagro experiment and, by contrast with recent Milagro's report, we find no steady increase in the Matsushiro observations in a 7-year period between 2000 and 2007. We suggest, therefore, that the steady increase of the SBVD reported by the Milagro experiment is not caused by the decreasing solar modulation in the declining phase of the 23rd solar activity cycle. ",
        "introduction": "The sidereal anisotropy of high-energy galactic cosmic ray (GCR) intensity provides us with unique information about the magnetic structure of the heliosphere and/or the local interstellar space surrounding the heliosphere, through which GCRs propagate to the earth. Examining the temporal variation of the anisotropy is also of interest, since such a variation, if any, may reflect the change in the global magnetic structure in response to the solar  activity- and/or magnetic-cycle variations. While recent observations with air shower arrays and underground muon detectors have reported a consistent global structure of the anisotropy, conclusive reports on its temporal variation are still limited \\citep[e.g., see][]{Nagashima89,Nagashima04,Amenomori05,Guillian07,Abdo09,Amenomori09}. The galactic anisotropy is measured as a sidereal daily variation (SDV) of cosmic-ray intensity recorded in a fixed directional channel of the detector on the spinning earth. The temporal variation of the SDV in connection with the solar magnetic cycle was investigated from a continuous observation of the SDV of $\\sim$10 TeV cosmic-ray intensity over 12 years between 1973 and 1987 with an air-shower (AS) detector at Mt. Norikura in Japan. While the SDV in each year was statistically significant, its temporal variation was quite small with a statistical significance of only two sigmas \\citep{Nagashima89}. Comparing data divided into two 5-year periods, 1997-2001 including the solar maximum and 2001-2005 approaching the solar minimum of the 23rd solar activity cycle, the Tibet III air shower experiment also concluded that the variation of 3 TeV cosmic ray anisotropy during the solar activity cycle was insignificant \\citep{Amenomori06}. In the present paper, we analyze the long-term variation of the SDV by analyzing 0.6 TeV cosmic ray data observed with a Japanese underground muon detector over 24 years between 1985 and 2008. This observation period is long enough to examine the variation of both the 11-year solar activity- and 22-year solar magnetic-cycles.  Recently, the Milagro experiment utilizing a large water Cherenkov detector surrounded by an air shower array, on the other hand, has reported on a steady increase in the magnitude of the anisotropy of 6 TeV cosmic ray intensity over 7 years between 2000 and 2007 \\citep{Abdo09}. This period corresponds to the declining phase of the 23rd solar activity cycle when the yearly mean sunspot number decreased from $\\sim$100 to $\\sim$10 toward a long-lasting activity minimum epoch. The increase reported by the Milagro experiment was so large that the anisotropy amplitude more than doubled between 2000 and 2007. If this increase did indeed, result from the decreasing solar activity, it is reasonable to expect that the effect should be observable also in the sub-TeV region, because the physical processes responsible for the solar modulation, such as adiabatic deceleration, pitch angle scattering by the magnetic irregularities and the solar wind convection effects in the heliosphere, are all expected to be more effective for lower energy cosmic rays. It has actually been established that the amplitude of the sidereal anisotropy decreases with decreasing energy between $\\sim$0.1 and $\\sim$1 TeV due to a larger attenuation by the solar modulation effect, while the amplitude is more constant for multi-TeV cosmic rays with energies over $\\sim$1 TeV, implying insignificant influence of the solar modulation on TeV GCRs \\citep{Munakata97,Nagashima98,Amenomori05,Amenomori09a}. In the present paper, we also examine whether the effect reported by Milagro is observed in the sub-TeV cosmic ray anisotropy.  The outline of this paper is as follows. In \\S 2, we briefly describe the data and analysis method and present analysis results. We first show in \\S 2.1 how the diurnal anisotropy in solar time changes in a significant correlation with the solar activity, indicating that solar modulation effects act on sub-TeV GCRs. We then show in \\S 2.2 the long-term variation in the sidereal anisotropy. Our results are discussed in \\S 3. ",
        "conclusions": "A clear $\\sim$11-year variation was found in the solar diurnal anisotropy. The amplitude of the residual anisotropy, after subtracting the contribution from the Compton-Getting anisotropy due to the earth's orbital motion around the sun, varies in an $\\sim$11-year cycle between a maximum amplitude of 0.043$\\pm$0.002 $\\%$ and a minimum of 0.008$\\pm$0.002 $\\%$, while its phase remains fairly constant around 15:10$\\pm$00:25 LT. A possibility is that the anisotropy is connected with the temporal variation of the IMF which, on large scales, is influenced by the neutral sheet tilt angle (TA). \\citet{Erdos80} first suggested that GCRs are expected to experience the energy change when they cross the neutral sheet on their orbital motion through the heliosphere to the earth. They computed the energy change along individual trajectories and deduced the directional anisotropy expected at the earth. Preliminary results of recent numerical model simulations using different tilt angles \\citep{Kota07} showed some qualitative similarities to the Matsushiro observations for solar anisotropies: high tilt angles in the maximum period produced larger anisotropies while the phases turned out fairly stable typically in the 3rd quadrant between 12:00 and 18:00 LT. We tentatively note that the $\\sim$15:00 phase corresponds to arrival direction of particles that have the best chance to gain energy from the interacting with the corotating magnetic sectors divided by a wavy current sheet. It is most interesting, on the other hand, that the variation of the observed amplitude of the ``additional'' anisotropy, exhibits a significant time lag of +26 months relative to the SSN. This time lag corresponds to the time for the solar wind with an average speed of 400 km/s to propagate through $\\sim$180 AU, which is larger than the heliocentric distance of the solar wind termination shock (TS) in the nose direction \\citep{Decker05,Stone08}. Even with the slowing down of the subsonic solar wind, this time-lag corresponds to distances deep in the heliosheath, and possibly close to the heliopause \\citep{Washimi96}. The cyclic variation induced by the sun propagates to the outer heliosphere with the average solar wind speed. The global distribution of the solar wind plasma and the magnetic field, which governs the TA and the spatial distribution of GCRs, is reorganized when the variation propagates throughout the entire region filled with the solar wind. This implies, therefore, that the intensity of 0.6 TeV GCRs is still a subject to the solar modulation operating over an entire region within the TS and possibly also in the subsonic region beyond the TS. The possible role of heliosheath outside the TS is largely unknown at present and remains to be explored.  We discuss the sidereal anisotropy, which shows no clear correlation with either the solar activity- or magnetic-cycles. Furthermore, there is even a factor of three attenuation of the amplitude relative to air shower observations due to the solar modulation. A possible explanation for this may be found in the region where the attenuation takes place. The solar diurnal anisotropy of 0.6 TeV GCRs at the earth's orbit is mainly produced by the global distribution of the GCR density inside the TS, while the sidereal anisotropy originates from the GCR anisotropy in interstellar space outside the heliosphere and it is attenuated as cosmic rays propagate to the earth through the heliosphere losing their original directional information. If the attenuation occurs predominantly inside the TS, it is reasonable to expect significant solar cycle variation of the anisotropy at the earth's orbit, as we saw in the solar diurnal anisotropy. The situation may change, however, if the attenuation occurs in the region outside the TS. Although our knowledge about the heliosheath between the TS and the heliopause is still limited, the Magneto-Hydrodynamic (MHD) simulations of the interaction between the solar wind and the interstellar plasma suggest the existence of a long heliotail extending over thousand of AU downstream the interstellar plasma flow \\citep[e.g., see][]{Baranov93,Washimi96,Zank96,Linde98,Pogorelov04}. Because of both the slower solar wind in the heliosheath and the larger extent of the heliotail, the propagation time of the solar cycle variation through the entire heliotail is expected to take much longer than inside the TS, probably even longer than 11 or 22 years of the solar activity- or magnetic-cycles. In this case, the plasma regions originating from the active and quiet sun are both expected to exist alternatively in the heliotail, as confirmed by a recent MHD simulation \\citep{Washimi07}(Dr. H. Washimi, 2009, private communication). If the observed attenuation of the sidereal anisotropy occurs predominantly in the heliotail, therefore, it is possible to expect the attenuation magnitude showing no strong and direct correlation with the solar activity- and/or magnetic-cycles, as we saw in the observation with Matsushiro. This is just one of possible interpretations and we need to confirm it by analyzing the GCR propagation in the model heliosphere including the heliotail predicted by MHD simulations.  Finally we note that the magnetic structure of the heliosheath is largely unknown. The possible role of the heliosheath is particularly interesting in the light of the recent finding of the IBEX spacecraft, mapping the global heliosphere by detecting energetic neutral atoms (ENAs) \\citep{McComas09}. The first results of IBEX showed that ENAs are sensitive to, and give information on the structure of the large-scale magnetic field surrounding the heliosphere \\citep{Heerik10}. TeV cosmic rays also sense the global structure of the magnetic field within the whole heliosphere as well as the field surrounding the heliosphere, and can provide an additional tool to explore the global structure of these fields."
    },
    "0911/0911.0355_arXiv.txt": {
        "abstract": "We observed two fields near M32 with the \\textsl{Advanced Camera for Surveys/High Resolution Channel} (ACS/HRC) on board the \\textsl{Hubble Space Telescope} (HST). The main field, F1, is 1\\farcm8 from the center of M32; the second field, F2, constrains the M31 background, and is 5\\farcm4 distant.  Each field was observed for 16-orbits in each of the $F435W$ (narrow $B$) and $F555W$ (narrow $V$) filters.  The duration of the observations allowed RR Lyrae stars and other short-period variables to be detected.  A population of RR Lyrae stars determined to belong to M32 would prove the existence of an ancient population in that galaxy, a subject of some debate.  We detected 17 RR Lyrae variables in F1 and 14 in F2.  A $1\\sigma$ upper limit of 6 RR Lyrae variables belonging to M32 is inferred from these two fields alone.  Use of our two ACS/WFC parallel fields provides better constraints on the M31 background, however, and implies that $7_{-3}^{+4}$ (68 \\% confidence interval) RR Lyrae variables in F1 belong to M32.  We have therefore found evidence for an ancient population in M32.  It seems to be nearly indistinguishable from the ancient population of M31.  The RR Lyrae stars in the F1 and F2 fields have indistinguishable mean $V$-band magnitudes, mean periods, distributions in the Bailey diagram and ratios of RRc to RR$_{\\mathrm{total}}$ types.  However, the color distributions in the two fields are different, with a population of red RRab variables in F1 not seen in F2.  We suggest that these might be identified with the detected M32 RR Lyrae population, but the small number of stars rules out a definitive claim. ",
        "introduction": "Messier 32 (M32) is the only elliptical galaxy close enough to possibly allow direct observation of its stars down to the main-sequence turn-off (MSTO).  It is a vital laboratory for deciphering the stellar populations of all other elliptical galaxies, which can only be studied by the spectra of their integrated light, given their greater distances. Major questions about M32's star formation history remain unanswered. M32 appears to have had one or more relatively recent episodes of star formation \\citep[within the last 3 Gyr: e.g.,][]{OConnell80,Rose85,Rose94,G93,T00b,Coelho09}, which also appears to be true for many elliptical galaxies \\citep[e.g.,][]{G93,T00b,TMBO05}.  These conclusions rest on painstaking and controversial spectral analysis of their integrated light. In contrast, the most direct information about a stellar population comes from applying stellar evolution theory to color--magnitude diagrams (CMDs).  Little however is known about M32's ancient population \\citep[see, e.g.,][]{Brown00,Coelho09}.   With our \\textsl{Advanced Camera for Surveys/High Resolution Channel} (ACS/HRC) data (Cycle 14, Program GO-10572, PI: T. Lauer) we have obtained the deepest CMD of M32 to date.  A comprehensive analysis of this CMD is discussed in a companion paper \\citet[hereafter M09]{M09} and we refer to it for further details. However here we want to stress that, due to the severe crowding in our fields, even with the high spatial resolution of HRC it is not possible to reach the MSTO with sufficient precision to claim the presence of a very old population.  RR Lyrae variables are low-mass stars burning He in their cores. They are excellent tracers of ancient stellar populations, completely independent of the MSTO, and knowledge of their properties provides important information on their parent stellar populations. Because they are located on the horizontal branch (HB) in a CMD, they are at least 3 magnitudes brighter than MSTO dwarfs and therefore detectable to relatively large distances. RR Lyrae are also very easy to characterize, with ab-type RR Lyrae (RRab) pulsating in the fundamental mode, rising rapidly to maximum light and slowly declining to minimum light, and c-type RR Lyrae (RRc) pulsating in the first harmonic mode, with their luminosities varying roughly sinusoidally. Most importantly for our purpose, the mere presence of RR Lyrae stars among a population of stars suggests an ancient origin, as ages older than $\\sim 10$ Gyr are required to produce RR Lyrae variables. Thus the detection of RR Lyrae stars in M32 is presently the only way to confirm the existence of an ancient stellar population in this galaxy.  \\citet{AlonsoGarcia04} were the first to attempt to directly detect RR Lyrae stars in fields near M32. They imaged a field $\\sim3\\farcm5$ with WFPC2 to the east of M32 and compared it with a control field well away from M32 that should sample the M31 field stars. They identified $12\\pm8$ variable stars claimed to be RR Lyrae stars belonging to M32 and therefore suggested that M32 possesses a population that is older than $\\sim10$ Gyr.  They were however unable to classify these RR Lyrae variables and could not derive periods and amplitudes for them.  Very recently, \\citet[hereafter S09]{s09} used ACS/WFC parallel imaging from our present data set to find RR Lyrae variables in two fields close to M32.  They found 681 RR Lyrae variables, with excellent photometric and temporal completeness (Sec.~\\ref{sec:detection}).  These RR Lyrae stars were roughly equally distributed between the two fields, with the same mean average magnitudes, metallicities, and Oosterhoff types in each field.  It was therefore impossible for them to separate the variables into M31 and M32 populations.  \\emph{It is still therefore an open question as to the precise nature or even presence of RR Lyrae variables in M32.}  In this paper, we present newly-detected RR Lyrae variables observed with ACS/HRC and also a detailed analysis of the fields near M32 where RR Lyrae stars have been found with HST. The paper is organized as follows.  In Section 2 we describe our observations and the data reduction we performed. We move on to describe the technique used to identify and characterize the RR Lyrae variable stars in Section 3, where we present their periods and light curves.  In Section 4 we show that we have clearly detected RR Lyrae variables in M32, as long as we include the results from our ACS/WFC parallel fields. In Section 5 we discuss the properties of the RR Lyrae stars, such as the location of their instability strip, reddenings, mean periods and Oosterhoff types, as well as pulsational relations such as period--metallicity--amplitude.  In this section we also derive estimates of the distance moduli to and metallicities of our fields. We summarize our findings and present our final conclusions in Section 6. ",
        "conclusions": ""
    },
    "0911/0911.0739_arXiv.txt": {
        "abstract": "{The source of emission from \\sgra, the supermassive \\bh\\ at the Galactic Center, is still unknown. Flares and data from multiwavelength campaigns provide important clues about the nature of \\sgra\\ itself. } {Here we attempt to constrain the physical origin of the broadband emission and the radio flares from \\sgra.} {We developed a time-dependent jet model, which for the first time allows one to compare the model predictions with flare data from \\sgra. Taking into account relevant cooling mechanisms, we calculate the frequency-dependent time lags and photosphere size expected in the jet model. The predicted lags and sizes are then compared with recent observations. } {Both the observed time lags and size-frequency relationships are reproduced well by the model.  The combined timing and structural information strongly constrain the speed of the outflow to be mildly relativistic, and the radio flares are likely to be caused by a transient increase in the matter channelled into the jets.  The model also predicts light curves and structural information at other wavelengths which could be tested by observations in the near future. } {We show that a time-dependent relativistic jet model can successfully reproduce: (1) the quiescent broadband spectral energy distribution of \\sgra, (2) the observed 22 and 43 GHz light curve morphologies and time lags, and (3) the frequency-size relationship. The results suggest that the observed emission at radio frequencies from \\sgra\\ is most easily explained by a stratified, optically thick, mildly relativistic jet outflow. Frequency-dependent measurements of time-lags and intrinsic source size provide strong constraints on the bulk motion of the jet plasma. } ",
        "introduction": "\\label{s:intro} The compact radio source Sagittarius A* (hereafter \\sgra), at an estimated distance of $8.4\\pm0.6$ kpc \\citep{ReidMentenZheng2009} and mass of $\\sim 4\\times10^6\\;\\msun$ \\citep{GhezSalimWeinberg2008, GillessenEisenhauerTrippe2009}, is our closest supermassive \\bh. When compared to other active galactic nuclei (AGN), \\sgra\\ is extremely under-luminous at all wavelengths, suggesting either that the mass accretion rate (\\mdot) onto the \\bh, or the radiative efficiency, is very low.  Polarization measurements suggest that $\\mdot\\leq 10^{-7}M_\\odot/{\\rm yr}$ \\citep{MarroneMoranZhao2007}. Recent multiwavelength campaigns that have monitored \\sgra\\ simultaneously from radio to X-rays have found that even though the luminosity of \\sgra\\ is low relative to its AGN counterparts, its emission is quite variable at all wavelengths.  Flaring activity at both sub-mm and higher frequencies has been observed on timescales ranging from minutes to hours, suggesting that the emission at these frequencies is produced very close to the supermassive \\bh, probably within a few tens of gravitational radii ($R_g$=$GM/c^2$). Although flaring appears to be nearly simultaneous at these frequencies \\citep{MarroneBaganoffMorris2008, EckartSchoedelGarcia2008,Dodds-EdenPorquetTrap2009}, the flares seem to have an increasingly long time-delay at longer wavelengths. For example, \\citet{Yusef-ZadehWardleHeinke2008} found that the 43 GHz flares precede the 22 GHz flares by $20\\pm10$ minutes, and that 350 GHz flares precede those of 43 GHz by $30-60$ minutes.  The physical origin of the observed electromagnetic radiation from \\sgra\\ has been under active debate for more than a decade. Attempts to model the emission from \\sgra\\ can be broadly classified into two categories: (1) accretion inflow models \\citep{Melia1992, NarayanMahadevanGrindlay1998,LiuMelia2002, YuanShenHuang2006}, and (2) relativistic outflow/jet models \\citep{FalckeBiermann1999,FalckeMarkoff2000,YuanMarkoffFalcke2002}. The radio, NIR, and X-ray flare light curves were modeled assuming expansion of a spherical plasma blob \\citep[see e.g.,][]{Yusef-ZadehWardleHeinke2008, EckartBaganoffMorris2009}. Recently \\citet{FalckeMarkoffBower2009} suggested that the radio time-lag data and measurements of intrinsic size at various frequencies \\citep[e.g.,][]{BowerFalckeHerrnstein2004,ShenLoLiang2005, DoelemanWeintroubRogers2008} can be used to resolve the degeneracy between inflow and outflow models.  It has long been suspected that radio flares such as those reported by \\citet{Yusef-ZadehWardleHeinke2008} could be ``smoking gun evidence'' of jets in sources such as \\sgra\\ \\citep{Falcke1999b}.  While jets have not been directly imaged in \\sgra, several observations strongly suggest the presence of jets in this source. The flat/slightly inverted radio spectral energy distribution (SED) \\citep[signature of a compact self-absorbed jet in the radio; see e.g.,][] {BlandfordKonigl1979} observed from \\sgra\\ are similar to those of other low-luminosity AGNs.  M81 and \\sgra\\ have very similar spectra and polarization properties in the cm-radio band, and M81 has been found to have weak jets \\citep{NagarFalckeWilson2005}.  Radio observations of the stellar X-ray binary system A0620$-$00 \\citep{GalloFenderMiller-Jones2006} at almost near quiescent luminosity suggest that compact jets are present at bolometric luminosities as low as $\\sim$10$^{-7}$ L$_{\\rm Edd}$. If jet physics and accretion flow scale with the mass of the \\bh, then \\sgra\\ (L$_{\\rm bol}$$\\sim$10$^{-9}$ L$_{\\rm Edd}$) is also expected to harbor a faint, compact jet.  Steady-state jet models have previously been used to model the broadband quiescent spectra of \\sgra\\ successfully \\citep{FalckeMarkoff2000, YuanMarkoffFalcke2002}; steady-state models, however, cannot predict the flaring activity.  In this paper we develop a proper treatment of the time-dependent cooling processes of the particle distribution, while keeping the basic hydrodynamic outflow model the same as in \\citet{Falcke1996}, allowing a semi-analytical treatment of the problem.  Once dominant cooling mechanisms are taken into account, the time-dependent jet model naturally accounts for all the observational constraints, namely, (a) the broadband SED, (b) the observed light-curve morphologies and time-lags, and (c) the observed frequency--size relationship. ",
        "conclusions": "\\label{s:summary} In this paper we have shown that a time-dependent relativistic jet model can explain most of the observational features seen in \\sgra. The main conclusions of this work are: \\begin{itemize} \\item The model presented here can describe the quiescent broad band SED of \\sgra, as well as its long-term radio variability. \\item Assuming that the radio flares are caused by a transient density enhancement at the base of the jet, which then propagates outward, the model can also fit the 22 and 43 GHz light curve morphology of the flares seen on 2006 July 17 by \\citet{Yusef-ZadehWardleHeinke2008}.  We predict light curve morphology and delays at longer wavelengths, which can be used to test the model with future observations. \\item The model predicted frequency--size relationship also matches quite well with the observed radio data. \\end{itemize}  The radio emission in our model originates in the self-absorbed, optically thick part of the outflow, where adiabatic losses dominate over radiative cooling. Therefore variability in the radio light curves largely traces the expansion of the jet plasma. Assuming that the flares have the same frequency--size relationship as the quiescent emission, the time delay measurements combined with frequency--size measurements strongly constrain the bulk speed of the jet plasma in this model and rule out subrelativistic expansion speeds.  An alternate model used by \\citet{Yusef-ZadehWardleHeinke2008} to fit the radio light curves assumes adiabatic, subrelativistic expansion of a spherical blob of plasma, i.e. the classic \\citet{vanderLaan1966} model.  Our modeling, at the very least, shows that the van der Laan model is not unique and in fact the jet model, in contrast to the van der Laan model, not only correctly predicts the light curves and the overall variability amplitudes, but also fits the frequency--size relation without violating any other observational constraints.  It must be kept in mind that the jet model does not attempt to explain the NIR or X-ray flares.  The NIR/X-ray flares trace particle (re)energization and cooling very close to the \\bh, while only marginally affecting the optically thick radio flux \\citep{MarkoffFalckeYuan2001}, and no tight correlation between radio and X-ray flares have been found so far.  A full general relativistic MHD model including radiative cooling \\citep[see e.g.][]{FragileMeier2009,MoscibrodzkaGammieDolence2009} would be required to model the jet launching region.  However the simple model presented here captures the important physics and shows that the jet model explains the radio flaring properties naturally, due to simple adiabatic expansion of overdensities in the outflow (likely linked to variations in the accretion rate).  Future coordinated multiwavelength campaigns, especially measurements of time-lags, sizes and positional offset at other wavelengths, will enable a better understanding of the velocity profile of the jet from \\sgra."
    },
    "0911/0911.4188_arXiv.txt": {
        "abstract": "\\noindent Long-term monitoring of Low Mass X-ray Binaries (LMXBs) by the All Sky Monitor on board the Rossi X-ray Timing Explorer now covers $\\sim$13 yrs and shows that certain LMXB types display very long-term ($\\sim$ several to tens of years) quasi-periodic modulations. These timescales are much longer than any ``super-orbital'' periods reported hitherto and likely have a different origin. We suggest here that they are due to long-term variations in the mass-transfer rate from the donor, which are a consequence of solar-like magnetic cycles that lead to $P_{orb}$ changes (as proposed by Richman, Applegate \\& Patterson 1994 for similar long-term variations in CVs). Atoll sources display much larger amplitude modulations than Z sources over these timescales, presumably because Z sources are Eddington limited and hence unable to respond as readily as Atoll sources to fluctuations in the mass-transfer rate from the donor. ",
        "introduction": "Low Mass X-ray Binaries (LMXBs) contain a neutron star (NS) or black hole (BH) primary onto which material is transferred from a low-mass ($M \\leq 1 M_{\\odot}$) late-type, essentially normal main sequence (MS) star (spectral type $\\sim$A-M), with longer period systems containing a sub-giant. There are $\\sim$190 luminous ($\\lesssim10^{38}$ erg\\space s$^{-1}$) LMXBs known in our galaxy (Liu, van Paradijs \\& van den Heuvel 2007). Cataclysmic Variables (CVs) are very similar, but have white dwarf (WD) primaries at a $\\sim10^3$ factor reduction in luminosity and have orbital periods that range from hours to $\\sim10$ d. While the flux from LMXBs is predominantly in X-rays which originate from the inner accretion disc and the NS surface (where applicable), the flux from CVs is predominantly in the optical and originates from the entire accretion disc, the hot spot(s) on the WD surface and the disc-mass-transfer stream impact region. (For details see e.g. Frank, King \\& Raine 1992).  Long-term, non-orbital, quasi-periodic variations on timescales of $\\sim$tens-hundreds of days in LMXBs were discovered early in X-ray astronomy (see Charles et al. 2008 for a recent review). These ``super-orbital'' periods are thought to be related to the properties of the accretion disc, and there are two basic mechanisms actively under consideration: radiation-induced warping (Ogilvie \\& Dubus 2001) and precession (Whitehurst \\& King 1991) of the accretion disc. Radiation pressure from the intense X-rays arising near the compact object causes the warping of the accretion disc, while tidal forces in high mass-ratio binaries lead to the precession of the accretion disc (an effect first observed in CVs, see e.g. Warner 1995). Either of these effects can lead to the periodic obscuration of the compact object in X-ray binaries, causing a super-orbital modulation in the X-ray flux (Clarkson et al. 2003). The stability criteria established by Ogilvie \\& Dubus (2001) suggest that this effect should not arise in most LMXBs, however it is observed, for example, in the long-period LMXB CygX-2 (Clarkson et al. 2003).  The presence of even longer-term quasi-periodic variations (on the order of decades) has been predicted in CVs as a consequence of the modulation in the mass transfer-rate due to a solar-type magnetic-activity cycle occurring in the donor star (Applegate \\& Patterson 1987, Warner 1988). The donor stars in short-period LMXBs and CVs are tidally locked and therefore rotate at the orbital period of the binary system, which is $\\lesssim1$ d for most systems. This rapid rotation is expected to generate much stronger magnetic fields (Schrijver \\& Zwaan 1992). Long-term variability changes in the surface activity of stars (such as starspot activity) have been detected on timescales of decades (Baliunas 1985), similar to the $\\sim$11 yr magnetic-activity cycle of our Sun.  Applegate (1992) suggested that magnetically active donors become more oblate as their outer layers are spun up, due to angular momentum distribution changes brought about by their magnetic activity. As a result, the volume of the Roche lobe changes during the magnetic cycle, while the volume of the donor remains unchanged. The magnetic activity cycle therefore governs the Roche-lobe volume and the structure of the donor, which will be varying in oblateness as the cycle progresses. Such a variation will modulate the amount by which the donor overfills the Roche lobe and therefore the mass transfer rate will also vary on the magnetic-activity time-scale (as discussed by Richman, Applegate \\& Patterson 1994 for CVs, hereafter R94).  Considering the similarities between LMXBs and CVs, it is not unreasonable to expect similar longer-term modulations in LMXBs for the same reason. Online archival datasets now allow for a detailed investigation of long-term variations in the lightcurves of X-ray sources. Here we report on an analysis of the $\\sim$13-year X-ray history of LMXBs contained in the RXTE/ASM datasets.  Shortly before completing this paper, we became aware of Durant et al (2009). While they mention similar long-term timescales of the X-ray flux modulations in the 16 brightest persistent LMXBs contained in the RXTE ASM datasets, we propose here a mechanism which would explain the observed dichotomy between Atoll and Z sources, and for the origin of these long time-scales. ",
        "conclusions": "Long-term variation in the mass transfer rate due to the magnetic activity cycle of the donor, should therefore translate into long-term quasi-periodic modulations of the X-ray flux for sources in which the additional material can be accreted onto the NS without violating the Eddington limit, such as Atoll sources and bursters. However, in Z sources very little (if any) of the additional material will be accreted and much lower amplitude (if any) long-term X-ray flux modulations are expected as a result of the magnetic activity cycle of the donor.  Therefore, we conclude that RXTE ASM lightcurve data may now provide the first evidence for very long-term quasi-periodic modulations of the X-ray flux as a result of the modulation in the mass transfer-rate due to a solar-type magnetic-activity cycle in the donor star, similar to those proposed for CVs."
    },
    "0911/0911.1078_arXiv.txt": {
        "abstract": "We investigate the precision with which the parameters describing the characteristics and location of nonspinning black hole binaries can be measured with the Laser Interferometer Space Antenna (LISA). By using complete waveforms including the inspiral, merger and ringdown portions of the signals, we find that LISA will have far greater precision than previous estimates for nonspinning mergers that ignored the merger and ringdown.  Our analysis covers nonspinning waveforms with moderate mass ratios, $\\,q \\geq 1/10$, and total masses $10^5\\lesssim M/\\MSun\\lesssim 10^7$. We compare the parameter uncertainties using the Fisher matrix formalism, and establish the significance of mass asymmetry and higher-order content to the predicted parameter uncertainties resulting from inclusion of the merger. In real-time observations, the later parts of the signal lead to significant improvements in sky-position precision in the last hours and even the final minutes of observation. For comparable-mass systems with total mass  $M/\\MSun\\sim10^6$, we find that the increased precision resulting from including the merger is comparable to the increase in signal-to-noise ratio. For the most precise systems under investigation, half can be localized to within $\\cal O$(10 arcmin), and 10\\% can be localized to within $\\cal O$(1 arcmin). ",
        "introduction": "Gravitational waves carry a tremendous amount of information through the universe. It is the goal of the emerging field of gravitational wave astronomy to access that information and bring it to bear on the problems of astrophysics and cosmology. The current generation of gravitational wave detectors, such as the Laser Interferometer Gravitational-wave Observatory (LIGO) \\cite{Abramovici:1992ah}, are focused on the detection of gravitational waves from isolated astrophysical systems. LIGO and its contemporaries will also provide some minimal information on the parameters of these systems such as their mass, luminosity distance, and approximate sky position. The quality of this information will be limited by the relatively low signal-to-noise ratios (SNR) expected for LIGO events \\cite{Ajith:2009fz}.  The Laser Interferometer Space Antenna (LISA) \\cite{LISA1}, a space-based detector of gravitational waves in the milliHertz band, will detect the inspiral and merger of supermassive black hole binaries (BHBs) with very large SNRs ($100\\sim 10^{4}$)  out to redshifts of $z\\sim10$ or greater \\cite{Flanagan:1997sx, Baker:2006kr}. These large SNRs make it possible to extract a large amount of information from each event including mass, mass ratio, spins, orientation, luminosity distance, and sky position. Because sources of gravitational waves are strongly dominated by gravitational dynamics, and because the waves are expected to propagate through intervening matter with little interaction, these observations may provide an unusually clean and direct measurement of the source system parameters. Of particular interest are the distance and sky position parameters, which will drive LISA's ability to narrow the set of candidate source galaxies or clusters for merger events, potentially opening up a range of multi-messenger astronomy opportunities.  For instance, the coincident measurement of gravitational and electromagnetic signatures from a single source galaxy (\\ie standard sirens \\cite{Holz:2005df}) will allow a direct measurement of the redshift-luminosity relationship, thereby constraining the dark-energy equation of state. While cosmological models predict that the dark-energy-dominated era began fairly recently at $z\\sim1$, measurements for much larger redshifts are not currently possible by other methods, and the standard siren method is limited only by the achievable range of coincident electromagnetic observation.  Extracting this information requires a waveform model that can provide templates with sufficient fidelity to distinguish between signals with different parameters in the presence of instrumental noise. For BHBs, the complete waveform signal is traditionally divided into three different regimes: the inspiral, which can be described using post-Newtonian (PN) orbital dynamics; the ring-down, which can be treated using black-hole perturbation theory; and the merger, which bridges the two and can be predicted using numerical relativity.  Ideally, an estimate of LISA's ability to measure the parameters of observed BHBs would include the information contained in the complete waveform. However, the difficulty associated with modeling the merger has led the majority of such studies (the exceptions being \\cite{McWilliams_PhD}, \\cite{Babak:2008bu}, and \\cite{Thorpe:2008wh}) to include only the inspiral portion of the waveform. Until recent advances in numerical relativity \\cite{Pretorius:2005gq,Baker:2005vv,Campanelli:2005dd} opened the door to a complete understanding of General Relativity's predictions for these signals, it was not clear whether theoretical knowledge about the final strongest moments of the signal would be available for system parameter measurements. A naive guess at the consequences of omitting the merger would be that the loss in parameter precision would be proportional to the loss in SNR.  This assumes that the two portions of the signal have equal density of information per unit SNR and that that information is independent.  There are two reasons that are sometimes cited for expecting that the effect of the merger on parameter precision will be less than that on SNR. The first is that the merger encompasses very few GW cycles compared with the observed portion of the inspiral and it is expected that information content correlates with the number of cycles.  The second is that, in the low-frequency limit, the sensitivity of LISA to parameters such as sky position is entirely generated by the orbital motion of the LISA constellation and the merger is too short in duration to experience a significant orbital modulation.  In this work we investigate the significance of the merger to LISA parameter estimation using quasi-analytic waveforms that are tuned to match the results of numerical relativity.  We restrict ourselves to non-spinning binaries with moderate mass ratios ($q\\leq1/10$) and explore the parameter space around a candidate system with total redshifted mass of $1.33\\times 10^6 \\MSun$, mass ratio $q=1/2$, and redshift $z=1$.  Astrophysically, we expect black holes to have spin. While including spin \\cite{Vecchio:2003tn, Lang:2006bz} and mergers each separately improve parameter estimation, it is not known how these effects will combine.  Therefore, including both spin effects and mergers will be an important followup to this investigation.  In section \\ref{sec:meth}, we discuss the methods employed for generating models of the complete waveform signals and the instrument, and for estimating parameter measurement precision. In section \\ref{sec:res}, we examine LISA's ability to measure binary black-hole system parameters.  The primary novelty of our results is that we assume theoretical knowledge of the complete signal is applied in the observational analysis.  We examine the impact of including the merger and higher harmonic content (\\ref{subsec:harm}) for comparable mass systems near $10^6 \\MSun$, the variation of the results across a range of masses (\\ref{subsec:mtot}) and mass ratios (\\ref{subsec:mrat}). In  (\\ref{subsec:timeEvo}) we study how sky position information accumulates in time, which will impact how LISA ultimately interacts with other astronomical instruments. We summarize our key results in section \\ref{sec:conc}. ",
        "conclusions": "\\label{sec:conc}  We have investigated the precision with which black hole binary system parameters can be measured from LISA observations including merger waveforms in the analysis of nonspinning binaries with moderate mass ratios ($q\\geq 1/10$). We have further studied how the expected performance depends on mass ratio and total system mass, and the impact of including or neglecting the merger signal and higher harmonics. The luminosity distance and the polarization phase depend on both the inclusion of the merger and the presence of higher harmonics, although the improvements from including these two elements are not independent. The inclination, the orbital phase constant, the ecliptic latitude and ecliptic longitude also depend on both the merger and the harmonic content. For these parameters the improvements resulting from including both the merger and higher harmonics are essentially independent.  For comparable-mass systems near $10^6\\MSun$, ignoring the merger reduces the SNR by a factor of $\\sim 3$ and results in a similar loss of median precision in parameter estimation, even more so for the sky position estimates. For sky position, ignoring the merger results in a more significant loss of precision than ignoring higher harmonics. Parameter estimates are roughly independent of mass ratio through $1>q>1/4$, for sky position in particular, though for smaller mass ratios $q\\lesssim1/10$ the precision begins to decrease.  For  $q=1/2$, the best parameter estimates are obtained for systems near $10^6\\MSun$, which merge in the middle of LISA's sensitivity band. Decreasing the mass by an order of magnitude to $\\sim 10^5\\MSun$ results in a precision loss of roughly a factor of 5 with a diminished relative contribution from the merger.  Increasing the mass to $\\sim 10^7\\MSun$ results in a similar loss. Though we have left out the effects of spin, including precession, our median sky position precision estimates are similar to those obtained with precession, but ignoring the merger \\cite{Lang:2006bz}. Each method locates the systems in the sky within ${\\cal O}(10$ arcmin). Our best cases ($\\sim$ top 10\\%) are localized at the level of ${\\cal O}(1$ arcmin). We estimate that LISA will usually be able to locate larger mass systems (near $10^7\\MSun$) quite well, in some cases better than systems with masses near $10^5\\MSun$, and far better than earlier estimates based on inspirals alone. Our results for these more massive systems do not, however, reproduce the preliminary (but widely discussed) extraordinarily precise sky localization results found in \\cite{Babak:2008bu}, though we do achieve such high precision for the $\\sim10^6\\MSun$ systems, where both the inspiral and merger are in-band and can contribute. For equal-mass systems near $10^6\\MSun$ the sky angle estimates improve over the last several hours up to a few minutes before merger in rough proportion to the time remaining before merger."
    },
    "0911/0911.0554_arXiv.txt": {
        "abstract": "First spectroscopic and new photometric observations of the eclipsing binary FM~Leo are presented. The main aims were to determine orbital and stellar parameters of two components and their evolutionary stage. First spectroscopic observations of the system were obtained with DDO and PST spectrographs. The results of the orbital solution from radial velocity curves are combined with those derived from the light-curve analysis (\\textit{V}-band ASAS-3 photometry and supplementary observations of eclipses with 1~m and 0.35~m telescopes) to derive orbital and stellar parameters. \\textit{JKTEBOP}, Wilson-Devinney binary modelling codes and a two-dimensional cross-correlation (TODCOR) method were applied for the analysis. We find the masses to be $M_{1}$ = 1.318 $\\pm$ 0.007 and $M_{2}$ = 1.287 $\\pm$ 0.007 $\\rmn M_{\\sun}$, the radii to be $R_{1}$ = 1.648  $\\pm$ 0.043 and $R_{2}$ = 1.511  $\\pm$ 0.049 $\\rmn R_{\\sun}$ for primary and secondary stars, respectively. The evolutionary stage of the system is briefly discussed by comparing physical parameters with current stellar evolution models. We find the components are located at the main sequence, with an age of about 3~Gyr. ",
        "introduction": "\\label{sec1}   Binary stars are the main source of fundamental data on stellar masses and radii. High quality and quantity of binaries' observations can help in testing of stellar evolution theory. Accurate masses and radii (determined better than 3~$\\%$) derived from detached binary systems' analysis have led to a number of new and interesting results on the properties and evolution of normal stars \\citep{tor}. Such data can be used to improve the treatment of convection, diffusion, and other non-classical effects which are essential to our understanding of stellar structure.  FM~Leo, also known as a HD~97422, HIC~54766 is an EA-type eclipsing binary. Its extensive photometry was listed by ASAS \\citep{pojm} under the code ASAS~111245+0020.9. Its maximum magnitude comes to $V_{\\rm max} = 8.47$~mag. An analysis of the light-curve reveals two eclipses of nearly identical depth and a flat maximum proving that FM~Leo is a detached binary consisting of two nearly identical spherical components.  An analysis of ASAS-3 \\citep{pojm} and Hipparcos \\citep{pery} databases yields an orbital period of $6 \\fd 72863$ \\citep{otero}. The star is classified as F8 \\citep{pery}. The Hipparcos parallax of the binary is $\\pi = 8.35 {\\pm}$ 1.17~mas.  In this paper we present results of the exact determination of stellar parameters of FM Leo. Since there was no radial velocity (RV) curve for this star in literature, we resorted to our own observations. We also present new photometric observations of eclipses of the binary. Our data and reduction procedures are described in Sect.~\\ref{sec2}. In Sect.~\\ref{sec3} we analyse observations using \\textit{JKTEBOP} (photometry) \\citep{south04a,south04b} and our own procedure (spectroscopy) to find the solution for the system and present absolute properties of FM~Leo, while Sect.~\\ref{sec4} is devoted to discussion on the evolutionary status and the age of the system. ",
        "conclusions": "\\label{sec5}  First spectroscopic and new photometric observations of the eclipsing binary FM~Leo combined with data from literature (ASAS photometry) have allowed us to derive definitive orbital parameters and physical properties of the component stars. The orbital and physical analysis is presented for the first time. With the mass and radius of each component of FM~Leo determined better than 1~$\\%$ and 4~$\\%$ respectively, we compared the observations with current stellar evolution models interpolated for this system. Using several methods with different input parameters and plotting them on various planes, we find the components are located at the main sequence, with an age of about 3~Gyr. The metallicity of the system is comparable with the solar one.  To confirm the results we calculated the distance to the system from determined radius, temperature and apparent magnitude and compared it with the value from Hipparcos ($\\pi~=~8.35~\\pm~1.17~$~mas, $d~=~119.8~\\pm~19.5~$~pc) \\citep{pery}. We estimated the parallax to the system of $\\pi~=~8.55~\\pm~2.05$~mas and the distance $d~=~130.4~\\pm~25.5$~pc. The estimation of reddening was performed by using maps of dust infrared emission by \\citet{schlegel}. We received the color excess of $E(B-V)~=~0.043$. The test confirmed our estimations are acceptable.  To make radii estimations more precise new photometric data of times of eclipses are needed. It would open a potential application of fully characterized FM~Leo for tests of evolutionary tracks and isochrones."
    },
    "0911/0911.3651.txt": {
        "abstract": "Using available astrometric and radial velocity data, the space velocities of cataclysmic variables (CVs) with respect to Sun were computed and kinematical properties of various sub-groups of CVs were investigated. Although observational errors of systemic velocities ($\\gamma$) are high, propagated errors are usually less than computed dispersions. According to the analysis of propagated uncertainties on the computed space velocities, available sample is refined by removing the systems with the largest propagated uncertainties so that the reliability of the space velocity dispersions was improved. Having a dispersion of $51\\pm7$ km~s$^{-1}$ for the space velocities, CVs in the current refined sample (159 systems) are found to have $5\\pm1$ Gyr mean kinematical age. After removing magnetic systems from the sample, it is found that non-magnetic CVs (134 systems) have a mean kinematical age of $4\\pm1$ Gyr. According to $5\\pm1$ and $4\\pm1$ Gyr kinematical ages implied by $52\\pm8$ and $45\\pm7$ km~s$^{-1}$ dispersions for non-magnetic systems below and above the period gap, CVs below the period gap are older than systems above the gap, which is a result in agreement with the standard evolution theory of CVs. Age difference between the systems below and above the gap is smaller than that expected from the standard theory, indicating a similarity of the angular momentum loss time scales in systems with low-mass and high-mass secondary stars. Assuming an isotropic distribution, $\\gamma$ velocity dispersions of non-magnetic CVs below and above the period gap are calculated $\\sigma_\\gamma=30\\pm5$ km s$^{-1}$ and $\\sigma_\\gamma=26\\pm4$ km s$^{-1}$. Small difference of $\\gamma$ velocity dispersions between the systems below and above the gap may imply that magnetic braking does not operate in the detached phase, during which the system evolves from the post-common envelope orbit into contact. ",
        "introduction": "Cataclysmic variables (hereafter CVs) are short-period interacting binary stars consisting of a white dwarf primary and low-mass secondary which overflows its Roche lobe and transfers matter to the white dwarf usually via a gas stream and an accretion disc. A bright spot is formed in the location where the matter stream impacts on the accretion disc. However disc formation is prevented in systems with strongly magnetic primaries. Nevertheless mass transfer continues through channels as accretion flows. For comprehensive reviews of the properties of CVs, see \\citet{War95} and \\citet{Hel01}.  According to standard CV formation and evolution scenario, CVs are formed from much wider -detached- main-sequence binaries with orbital periods typically 10-1000 d of the pairs with large mass ratios; secondaries being less than a solar mass, but the primaries are more massive ($M_{1}\\simeq1-9M_{\\odot}$). The primary star evolves in a nuclear time scale of roughly $t_{nuc}\\simeq10^{10}(M_{1}/M_{\\odot})^{-3}$ yr \\citep{KS96,Hel01}. When the primary fills its Roche lobe, dynamically unstable mass transfer leads to a common envelope (CE) phase during which dynamical friction extracts orbital angular momentum to eject the envelope of the giant. The time spent in common envelope phase is very short compared to other phases of the evolution, of the order of $10^{3}-10^{4}$ yr. After the CE phase, system finds itself in a new detached phase called post common envelope phase or pre-CV state; the secondary component is still a main-sequence star, while the primary component has been changed to a white dwarf. This detached system approaches a semi-detached state by either the nuclear expansion of the secondary, or by shrinking orbit due to the loss of orbital angular momentum by gravitational radiation and magnetic braking. During the post-CE phase, the orbit shrinks from the post-CE period of about 0.1-10~d to the orbital period $P\\simeq9(M_{2}/M_{\\odot})$~h at the onset of mass transfer as a CV. The post-CE phase can be very short if the system emerges from the CE phase at almost semi-detached configuration \\citep{KS96}.  The standard scenario proposed for the evolution of CVs after the onset of the mass transfer is sensitive to orbital angular momentum loss mechanisms too as in the post CE phase of the evolution. It has been proposed two main mechanisms: magnetic braking and gravitational radiation. For CVs with orbital periods ($P$) above the period gap, the secondary star is partially convective and, thus, it has a magnetic field. It has been proposed that the magnetic field in the secondary star causes orbital angular momentum loss by magnetic braking \\citep{VZ81,Retal82,Retal83,PS83,SR83,K88} until the orbital period decreases to about 3~h at which the secondary star becomes fully convective and magnetic braking ceases. Consequently, the secondary star relaxes to its thermal equilibrium and shrinks inside its Roche lobe. At this point, the mass transfer shuts off completely and the gravitational radiation becomes dominant mechanism for the angular momentum loss \\citep{Pac67}. Since the mass transfer ceased, the system cannot be observed as a cataclysmic variable and it evolves towards shorter orbital periods through emission of gravitational radiation. At the orbital period of about 2~h, the secondary star starts to fill its Roche lobe again and re-starts the mass transfer at a much lower rate than in the mode of long-period CVs. When the orbital period decreases to the observed minimum of about 80 min, the secondary star becomes a degenerate (brown-dwarf like) object and, thus, further mass transfer causes the orbital period to increase \\citep{Pat98}, creating systems called period bouncers. For a comprehensive review of the formation and evolution of CVs, see \\citet{Ritter08}.  Though the above model successfully explains the period gap, i.e. drop in the number of known CVs in between orbital periods in the range about $2<P$(h) $<3$, some predictions of it are contradictory with observations. Since CVs should spend most of their lifetime near the period minimum, a significant accumulation of CVs near the period minimum is expected. Although such an accumulation has not been observed for many years \\citep{KB99}, there is now a claimed discovery of a spike at the period minimum \\citep{Southworthetal08,Gaensickeetal09}. According to population synthesis \\citep{dK92,dKR93,Kol93,Pol96,KB99,Howetal97,KS02}, the vast majority of CV population should have orbital periods shorter than about 2~h . However, such an accumulation has not been seen in the distribution of orbital periods of CVs. Recently, \\citet{Litt08} found seven systems with brown dwarf donors and proposed that the missing population of post-period minimum CVs has finally been identified, although their masses and radii are inconsistent with model predictions. In addition, the predicted space density of CVs is about a few $10^{-5}$ to a few $10^{-4}$ pc$^{-3}$ \\citep{dK92,Pol96}, whereas the space density derived from the observations is in agreement only with the lower limit of these predictions \\citep{Pretetal07a,Schetal02,Aketal08}. Although a number of alternative suggestions \\citep{LP94,KK95,Clemetal98,Koletal98} and some alternative angular momentum loss mechanisms \\citep{Andetal03,TS01,Scheetal02,Willetal2005,Willetal2007} has been introduced as possible solutions to these problems, none of them could solve the problem with a complete success.  The observational data sets, which can be used to test the above predictions of the standard evolution model of CVs, are strongly biased by the selection effects. Brightness dependent selection effects are the strongest since the known CV sample is not approximately magnitude-limited, even if a limiting magnitude as bright as $V=13$ is adopted \\citep{Pretetal07b}. Nevertheless, kinematical age of a cataclysmic variable, which is a time span since formation of component stars, can be used to test the predictions of the model. The age of a CV does not affect its mass transfer rate at a given orbital period. Hence the age distribution of CVs is not biased by brightness-selection \\citep{Kol01}. Nevertheless, shorter the orbital period, it is harder to measure radial velocity variation for determining orbital parameters.  \\cite{KS96}, who determined the age structure of a modeled population for galactic CVs by applying standard models for the formation and evolution of CVs, predicted that systems above the period gap ($P\\gtrsim$ 3~h) must have an average age of 1~Gyr, with most of them being younger than 1.5~Gyr, while systems below the gap ($P\\lesssim$ 2~h) should display a wide range of all ages above about 1~Gyr, with a mean of 3-4~Gyr \\cite[see also][]{RB86}. Using the empirical relation between age $t$ and space velocity dispersion $\\sigma(t)$ of field stars in the solar neighbourhood \\citep{Wie77,Wieetal92}, $\\sigma(t)= k_{1}+k_{2}t^{1/2}$ ($k_{1}$ and $k_{2}$ are constants), an intrinsic dispersion of the $\\gamma$ velocities $\\sigma_{\\gamma} \\simeq 15$ km~s$^{-1}$ for the systems above the period gap, and $\\sigma_{\\gamma} \\simeq 30$ km~s$^{-1}$ below the gap were predicted by \\cite{KS96}. \\cite{Kol01} states that the difference between ages of systems above and below the period gap is mainly due to the time spent evolving from the post-common envelope (post-CE) orbit into contact. However, \\cite{vPaetal96} (hereafter referred to as vPAS96), who collected the observed $\\gamma$ velocities for a sample of CVs from published radial velocity studies and analysed this CV sample statistically, could not detect such a difference between the velocity dispersions for the systems above and below the period gap. Using precise $\\gamma$ velocity measurements of only four dwarf novae above the period gap, \\cite{Noretal02} suggest a velocity dispersion of $\\sim$8~km~s$^{-1}$. It should be noted that \\cite{Kol01} also expresses that if magnetic braking does not operate in the detached phase, the $\\gamma$ velocity dispersions of CVs are $\\sigma_{\\gamma} \\simeq 27$ km~s$^{-1}$ above the period gap versus $\\sigma_{\\gamma} \\simeq 32$ km~s$^{-1}$ below the gap.  Aim of this paper is to derive the observed $\\gamma$ velocity dispersions to test predictions. Kinematical age profiles for CVs according to different orbital period regimes are investigated in order to understand orbital period evolution of CV orbits. In order to do this, we have collected the measured $\\gamma$ velocities of CVs from the published radial velocity studies since the work of vPAS96 and combined our $\\gamma$ velocity collection with their sample. ",
        "conclusions": "The $\\gamma$ velocities of CVs are taken from a variety of sources from the literature. For all but 59 systems, published $\\gamma$ velocities are by-products of measurements of the components' semi-amplitude. In addition, measurements based on emission lines remain problematic. As a consequence, systematic errors affect specifically the determination of $\\gamma$ velocity and error values are high.  After analysing available kinematical data of CVs, it is concluded that there is not considerable kinematical difference between dwarf novae and non-magnetic nova-like stars. Magnetic and non-magnetic systems display different kinematical properties. However, kinematics based on $\\gamma$ velocity measurements of magnetic systems remain problematic as their $\\gamma$ velocities may be significantly contaminated by the flow velocities of the magnetically channelled plasma.  Kinematics of the present sample show that systems above the orbital pediod gap are younger than systems below the gap. This result is in agreement with the standard theory of the CV evolution. Kinematics of the present sample of CVs also roughly confirm the prediction of \\cite{KS96} who predicted about 2 Gyr age difference between the CV groups of below and above the period gap. Smaller age difference found in this study shows a similarity of the angular momentum loss time scales in systems with low-mass and high-mass secondary stars \\citep{Kol01}.  Assuming an isotropic distribution, observational $\\gamma$ velocity dispersions of CVs are in agreement with the theoretical predictions of \\cite{Kol01} who suggested that the difference of $\\gamma$ velocity dispersions of the systems below and above the period gap is about 5 km~s$^{-1}$ with systems above the gap having a $\\gamma$ velocity dispersion of $\\simeq 27$ km~s$^{-1}$. This agreement between the observations and the theory implies that magnetic braking does not operate in the detached phase."
    },
    "0911/0911.4627_arXiv.txt": {
        "abstract": "We address the existence of globally neutral neutron star configurations in contrast with the traditional ones constructed by imposing local neutrality. The equilibrium equations describing this system are the Einstein-Maxwell equations which must be solved self-consistently with the general relativistic Thomas-Fermi equation and $\\beta$-equilibrium condition. To illustrate the application of this novel approach we adopt the Baym, Bethe, and Pethick (1971) strong interaction model of the baryonic matter in the core and of the white-dwarf-like material of the crust. We illustrate the crucial role played by the boundary conditions satisfied by the leptonic component  of the matter at the interface between the core and the crust. For every central density an entire new family of equilibrium configurations exists for selected values of the Fermi energy of the electrons at the surface of the core. Each such configuration fulfills global charge neutrality and is characterized by a non-trivial electrodynamical structure. The electric field extends over a thin shell of thickness $\\sim \\hbar/(m_e c)$ between the core and the crust and becomes largely overcritical in the limit of decreasing values of the crust mass. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0911/0911.1833_arXiv.txt": {
        "abstract": "We investigate the physical meaning of the prominent blue shifts of Ne~{\\sc{viii}}, which is observed to be associated with quiet-Sun network junctions (boundary intersections), through data analyses combining force-free-field extrapolations with EUV spectroscopic observations. For a middle-latitude region, we reconstruct the magnetic funnel structure in a sub-region showing faint emission in EIT-Fe 195. This funnel appears to consist of several smaller funnels that originate from network lanes, expand with height and finally merge into a single wide open-field region. However, the large blue shifts of Ne~{\\sc{viii}} are generally not associated with open fields, but seem to be associated with the legs of closed magnetic loops. Moreover, in most cases significant upflows are found in both of the funnel-shaped loop legs. These quasi-steady upflows are regarded as signatures of mass supply to the coronal loops rather than the solar wind. Our observational result also reveals that in many cases the upflows in the upper transition region (TR) and the downflows in the middle TR are not fully cospatial. Based on these new observational results, we suggest different TR structures in coronal holes and in the quiet Sun. ",
        "introduction": "It has well been established that the Ne~{\\sc{viii}}~(770.4~{\\AA}) line is on average blue shifted in coronal holes and the quiet Sun. This line is formed in the upper transition region (TR) and lower corona, and thus its Doppler shift is very useful to study TR dynamics and the solar wind origin. In coronal holes (CHs), the blue shift of Ne~{\\sc{viii}} is believed to be an indicator of solar wind outflow \\citep[e.g.,][]{Hassler1999,Wilhelm2000,Tu2005a}. Also in the quiet Sun, significant blue shifts of Ne~{\\sc{viii}} were found at the network junctions and interpreted \\citep{Hassler1999} to be possible sources of the solar wind. This idea challenges the conventional notion that the solar wind originates from coronal holes. To help clarifying this issue we have, through data analyses combining force-free-field extrapolation with EUV spectroscopy, further studied the prominent blue shifts of Ne~{\\sc{viii}} in the quiet Sun and investigated its physical meaning. ",
        "conclusions": ""
    },
    "0911/0911.4411_arXiv.txt": {
        "abstract": "We update a previous investigation of cosmological effects of a non-standard interaction between neutrinos and dark matter. Parameterizing the elastic-scattering cross section between the two species as a function of the temperature of the universe, the resulting neutrino-dark matter fluid has a non-zero pressure, which determines diffusion-damped oscillations in the matter power spectrum similar to the acoustic oscillations generated by the photon-baryon fluid. Using cosmic microwave background data in combination with large scale structure experiment results, we then put constraints on the fraction of the interacting dark matter component as well as on the corresponding opacity. ",
        "introduction": "Both the existence and the energy content of dark matter in the universe have been firmly estabilished by several observations (see \\cite{Bertone2004} for a review), yet, after decades of research its nature is still unknown. The favored candidate for dark matter has three main feature namely, it seems to be a long-lived, cold and collisionless species. However, it is possible to relax some of these assumptions and in fact, several different dark matter candidates have been proposed through years (\\cite{Jungman:1995df,axion}) and it is not even clear if only one species is responsible for the whole dark matter in the universe.\\\\ An interacting component has been proposed by many authors; couplings with a light scalar or pseudoscalar boson, as in the Majoron model has been investigated in \\cite{hannestad1,hannestad2,pierpaoli,Dodelson2004} and interaction with neutrinos and more generally with the electromagnetic plasma have been considered for example, in \\cite{boehm2002, boehm2003a, boehm2003b, Hooper2004, boehm2003c}.\\\\ In this paper we study the scenario where at least a fraction of dark matter is coupled to neutrinos. We parameterize the neutrino-dark matter interaction in terms of the opacity $P\\equiv\\langle\\sigma_{dm-\\nu}|v|\\rangle/m_{dm}$, the ratio of the thermal averaged dark-matter-neutrino scattering cross section to the mass of the dark matter particle, where $P=Q \\, a^{-2}$ with $Q$ a constant and $a$ the scale factor. We will briefly review the underlying theoretical basis of this choice in the next Section.\\\\ The main observable consequence of such a coupling is a pattern of oscillations in the matter power spectrum, which can in principle mimic the baryon acoustic oscillations observed in the large scale clustering of galaxies \\cite{Eisenstein2005,Percival2009}. The cosmic microwave background (CMB) radiation is  also affected by this interaction. Indeed, for sufficiently small scales, the temporal growth of a dark matter perturbation is lowered because of the coupling with neutrinos, and this determines a small enhancement of the small scale peaks  at a level which depends on both the strength of the interaction and the percentage of interacting dark matter. This effect can be also described in terms of a lowering of the neutrino \"viscosity parameter\" $c^2_{vis}$ \\cite{Hu1998} from its standard value $c^2_{vis}=1/3$.\\\\ \\begin{figure}[t] \\begin{center} \\begin{tabular}{ccc} \\resizebox{80mm}{!}{\\includegraphics{./fig1.eps}}\\\\ \\resizebox{80mm}{!}{\\includegraphics{./fig2.eps}}\\\\ \\resizebox{80mm}{!}{\\includegraphics{./fig3.eps}}\\\\ \\end{tabular} \\caption{A perturbation of interacting dark matter (red line) of wavenumber $k=0.81\\,h\\,$Mpc$^{-1}$ is plotted against a perturbation of non-interacting dark matter (black dotted line) in the upper panel. The effects of this interaction are clearly seen in the angular power spectum of the CMB (middle panel) and on the matter power spectum (lower panel). We have chosen a very large value of Q to magnify the effect on the main cosmological observables.} \\end{center} \\end{figure} This paper updates and improves the results of a previous work \\cite{Mangano2006}. In fact, we present the results of a likelihood analysis based on a Monte Carlo Markov Chain sampling of the full cosmological parameter space, which allows us to take into account for possible degeneracies with the new parameters introduced in the analysis. To this aim, we use the latest CMB datasets available from the WMAP experiment \\cite{Komatsu2009} in combination with large scale structure data from SDSS.  The paper is organized as follows. In Section II we briefly discuss possible theoretical models for the dark matter-neutrino interaction and corresponding opacity due to scatterings. We refer to \\cite{Mangano2006} for a more detailed discussion of these topics. In Section III we explain the method considered in the analysis and the data used to constrain both the strength of the interaction and the fraction of interacting dark matter. Finally we give our results and conclusions. ",
        "conclusions": "In this paper we have investigated the effects of a non-standard interaction between neutrinos and dark matter using the to date available data from WMAP and SDSS Collaborations on cosmic microwave background and large scale structures. These effects are the presence of diffusion -- damped oscillations in the matter power spectrum and a small enhancement of the peaks in the cosmic microwave background angular power spectrum. Parameterizing the opacity parameter as $Q \\,a^{-2}$ with $Q$ a constant, which holds in a large class of models for neutrino -- dark matter interactions, we found that  the bound on $Q$ becomes stronger as the percentage $\\Omega_X h^2$ of interacting dark matter component grows. Typical upper bound on $Q$ is of the order of $10^{-42} \\div 10^{-40}$ cm$^2$ MeV$^{-1}$, while the bound becomes weaker for very low values of $\\Omega_X$; we note that cosmological constraints on neutrino masses can be affected because of possible degeneracies.  \\smallskip This research was funded by NSF CAREER AST-0605427 and by LANL IGPP-08-505."
    },
    "0911/0911.1002_arXiv.txt": {
        "abstract": "A self-consistent global fitting method based on the Markov Chain Monte Carlo technique to study the dark matter (DM) property associated with the cosmic ray electron/positron excesses was developed in our previous work. In this work we further improve the previous study to include the hadronic branching ratio of DM annihilation/decay. The PAMELA $\\bar{p}/p$ data are employed to constrain the hadronic branching ratio. We find that the $95\\%$ ($2\\sigma$) upper limits of the quark branching ratio allowed by the PAMELA $\\bar{p}/p$ data is $\\sim 0.032$ for DM annihilation and $\\sim 0.044$ for DM decay respectively. This result shows that the DM coupling to pure leptons is indeed favored by the current data. Based on the global fitting results, we further study the neutrino emission from DM in the Galactic center. Our predicted neutrino flux is some smaller than previous works since the constraint from $\\gamma$-rays is involved. However, it is still capable to be detected by the forth-coming neutrino detector such as IceCube. The improved points of the present study compared with previous works include: 1) the DM parameters, both the particle physical ones and astrophysical ones, are derived in a global fitting way, 2) constraints from various species of data sets, including $\\gamma$-rays and antiprotons are included, and 3) the expectation of neutrino emission is fully self-consistent. ",
        "introduction": "\\label{Int}  In a previous work (\\cite{Liu:2009sq}, Paper I) we have developed a Markov Chain Monte Carlo (MCMC) code to fit the parameters of dark matter (DM) models, which are proposed to explain the recent reported abnormal excesses of cosmic ray (CR) positrons and electrons by PAMELA \\cite{Adriani:2008zr}, ATIC \\cite{Chang:2008zzr}, HESS \\cite{Aharonian:2008aa,Aharonian:2009ah} and Fermi-LAT \\cite{Abdo:2009zk}. One assumption adopted in that work is that the DM particles only couple with leptons. To extent the discussion including hadronic channels will be natural and necessary. Actually the non-excess of PAMELA $\\bar{p}/p$ data \\cite{Adriani:2008zq} have been studied in many works to set constraints on the hadronic couplings of DM particles (e.g., \\cite{Cirelli:2008pk,Yin:2008bs, Donato:2008jk}). However, these studies can only give some illustration constraints on the model parameters instead of a full scan of the possible parameter space. Our MCMC fitting scheme makes it possible to fully scan the high-dimensional parameter space and derive the model-independent constraints from the data directly.  Shortly after the proposal of using DM annihilation to account for the data, the accompanied $\\gamma$-ray and radio emission are investigated as constraint and/or future probe of the current models (e.g., \\cite{Bertone:2008xr,Zhang:2008tb,Bergstrom:2008ag}). Besides the photon emission, neutrino can also be regarded as an another probe to test the DM interpretation of the data \\cite{Liu:2008ci, Spolyar:2009kx,Buckley:2009kw}. In Ref. \\cite{Liu:2008ci} it is shown that Antares and IceCube are promising to detect the neutrino signals of DM annihilation in the Galactic center (GC) or the massive subhalo for models to explain PAMELA and ATIC data, assuming an NFW profile \\cite{Navarro:1996gj} of DM spatial distribution. It is also claimed that for annihilation DM scenario the IceCube/DeepCore detector can explore much of the parameter space to explain the PAMELA and Fermi-LAT $e^{\\pm}$ data, still for an NFW profile \\cite{Spolyar:2009kx}. For decaying DM scenario IceCube/DeepCore also has the potential to exclude some of the parameter space, depending on the final states and DM profile \\cite{Buckley:2009kw}. In all of these works, the expectation of neutrino signals will strongly depend on the annihilation/decay final states and DM profile, and the constraints from various kinds of data are not taken into account simultaneously.  Based on the above points, in this work we improve our MCMC code to include the hadronic branching ratio, and then we re-examine the neutrino signals according to the parameter sets derived in  MCMC calculations. We  include the positron fraction data from PAMELA \\cite{Adriani:2008zr}, the electron + positron data from Fermi-LAT and HESS \\cite{Abdo:2009zk, Aharonian:2008aa,Aharonian:2009ah}, the diffuse $\\gamma$-ray data of the GC ridge from HESS \\cite{Aharonian:2006au}, and the antiproton-proton ratio data from PAMELA \\cite{Adriani:2008zq} in this MCMC study.  The paper is organized as follows. In Sec. II we introduce the production and propagation of the CR electrons/positrons and antiprotons. In Sec. III we give the MCMC fitting results of DM model parameters, for both annihilation and decaying scenarios. The neutrino signals from GC are discussed in Sec. IV. Finally, Sec. V is the summary. ",
        "conclusions": "\\label{Sum}  Based on the MCMC code developed in Paper I, we further study the properties of DM models which may connect with the recent CR lepton excesses, after incorporating the hadronic branching ratio and the constraint from PAMELA $\\bar{p}/p$ data. Our global fitting results show that DM with annihilation or decay channels dominantly to the combination of $\\tau^+\\tau^-$ and $\\mu^+\\mu^-$ can well describe the electron and/or positron data measured by PAMELA, Fermi-LAT and HESS. The $\\bar{p}/p$ data from PAMELA limit the branching ratio to quark pairs to be less than $5\\%$. This conclusion is qualitatively consistent with the previous studies \\cite{Cirelli:2008pk,Yin:2008bs, Donato:2008jk}, but should be regarded as the first quantitative result based on the global fitting method.  We then calculate the neutrino fluxes from the DM annihilation or decay according to the MCMC fitting parameters.Due to the  $\\gamma$-ray constraint on the DM profile, Our expected neutrino flux for DM annihilation scenario is smaller than that in the previous works where the NFW profile is priorly adopted. However, as we have shown in Ref. \\cite{Bi:2009am} and Paper I, the NFW profile will be strongly constrained by the diffuse $\\gamma$-rays from the GC ridge observed by HESS \\cite{Aharonian:2006au} for DM annihilation scenario. The derived $2\\sigma$ upper limit of the slope of DM central profile is $\\sim 0.5$ in the MCMC study, which is much smaller than the NFW profile with $\\gamma=1$. Although smaller, the neutrino flux still could be detected by the forth-coming detector such as IceCube. However, a careful analysis about the angular distribution and spectral distribution is necessary to separate the DM-induced signal from the high atmospheric background. For DM decay scenario the constraint from $\\gamma$-rays is weaker, and the differences of the neutrino flux between this work and other works are smaller.  \\appendix*"
    },
    "0911/0911.5598_arXiv.txt": {
        "abstract": "{} {We investigate the nature and properties of far-infrared (FIR) sources from the AKARI Deep Field South (ADF-S).} {We performed an extensive search for the counterparts of $1000$ ADF-S objects brighter than 0.0301 Jy in the {{\\it WIDE-S}} (90 $\\mu$m) AKARI band in the public databases (NED and SIMBAD). We analyzed the properties of the resulting sample: statistic of the identified objects, quality of position determination of the ADF-S sources, their number counts, redshift distribution and comparison of morphological types, when the corresponding information was available. We also made a crude analysis of the clustering properties of the ADF-S sources and we constructed spectral energy distributions (SEDs) of selected objects with the best photometry, using three different models. } {Among 1000 investigated ADF-S sources, 545 were identified at other wavelengths using the public databases. From them, 518 are known galaxies, and 343 of these galaxies were not known previously as infra-red sources. Among the remaining sources there are two quasars and infrared and radio sources of unknown origin. Among six stellar identifications, at least five is probably an effect of contamination. We found redshifts of 48 extragalactic objects and morphological types of 77 galaxies. We present models of SEDs of 47 sources with sufficiently good photometric data.  } {We conclude that the bright FIR point sources observed in the ADF-S are mostly nearby galaxies. Their properties are very similar to properties of the local population of optically bright galaxies, except an unusually high ratio of peculiar or interacting objects and a smaller percentage of elliptical galaxies. The percentage of lenticular galaxies is the same as in the optically bright population which suggests that galaxies of this type may frequently contain a significant amount of cool dust. It is possible that source confusion plays a significant role in more than 34~\\% of measurements. The SEDs display a variety of galaxy types, from very actively star forming to very quiescent. Thanks to the AKARI long wavelength bands it was revealed for the first time that these galaxies form a population of objects with very cool dust.} ",
        "introduction": "\\label{sec:introduction}  Active star formation (SF) is followed by heavy element production from the birth and death of stars. Several of the heavy elements produced by stars end up substantially depleted into dust grains. These dust grains in galaxies tend to absorb ultraviolet (UV) light, emitted by young stars, and re-emit it in the far infrared (FIR). Indeed, there is an extreme category of galaxies which have large amount of dust and are extremely luminous in the FIR and submillimeter (submm) wavelengths. Heavily hidden SF is suggested in these galaxies.  By examining the luminosity functions (LFs) at UV and FIR from GALEX and IRAS/Spitzer, \\citet{takeuchi05a} proved that the FIR LF shows much stronger evolution than that of UV, though both evolve very strongly. This indicates that the fraction of hidden SF rapidly increases toward higher redshifts up to $z < 1$. There is another important observable closely related to the dust emission from galaxies: the cosmic IR background (CIB). Recently, \\citet{takeuchi06} constructed the IR spectral energy distribution (SED) of the Local Universe. The energy emitted in the IR is 25--30~\\% of the total energy budget. In contrast, the IR (from near/mid-IR to millimeter) contribution is roughly (or even more than) a half in the CIB spectrum \\citep[e.g.,][]{dole06}. This also suggests a strong evolution of the IR contribution to the cosmic SED in the Universe.  Thus, understanding the radiative physics of dust is a fundamental task to have an unbiased view of the cosmic SF history. Especially, exploring the evolution of SEDs of galaxies at each epoch is a key to have a unified view of the SF history. First step to this is to know the properties of Local galaxies. Vast amount of new knowledge about the IR universe was provided by IRAS \\citep[e.g.,][]{soifer87}, followed by MSX \\citep[e.g.,][]{egan03}, ISO \\citep[e.g.,][]{genzel00,verma05} and Spitzer \\citep[e.g.,][]{soifer08}.  Recently, after IRAS, a Japanese IR satellite AKARI has performed an all-sky survey \\citep{murakami} and various smaller but deeper surveys at different IR wavelengths. In particular, with the aid of the Far-Infrared Surveyor \\cite[FIS:][]{kawada}, observations in four FIR bands were possible. Among % the observed fields, the lowest Galactic cirrus emission density region near the South Ecliptic Pole was selected % for observations since it can provide the best FIR extragalactic image of the Universe. This field is referred to as the AKARI Deep Field South (ADF-S). This survey is unique in having continuous wavelength coverage with four photometric bands (65, 90, 140 and 160 $\\mu$m) mapped over a% wide area ($\\sim$12 square degrees). 2268 infrared sources were detected, down to $\\sim 20$~mJy at 90~$\\mu$m, and infrared colors for about 400 of these were measured.  By the advent of AKARI surveys, a new generation of large database of the Local Universe will be available. In this paper, we present first result of our cross identification (cross-ID) of the sources in the ADF-S. After showing the process of identification of the ADF-S objects, we discuss their statistical properties. Further, we show the UV-optical-FIR SEDs of selected nearby star-forming galaxies from our sample with the best photometric data. We have made a analysis of these SEDs by fitting a few simple model of dust emission from galaxies.  The paper is organized as follows: in Section~\\ref{sec:data}, we present the data. Section~\\ref{sec:basic} is devoted to the basic analysis of the properties of identified sources, their distribution on the sky, and the quality of the catalog and possible biases, e.g. related to the source confusion. Then we discuss the statistical properties of the obtained sample, e.g., the number counts, redshift distribution, galaxy morphologies, and other properties. In Section~\\ref{sec:sed} we present SEDs of the identified galaxies and we attempt to model the dust emission from them. We present our conclusions in Section~\\ref{sec:conclusion}. ",
        "conclusions": "\\label{sec:conclusion}  \\begin{enumerate} \\item In this paper, we presented the catalog of counterparts of the ADF-S 90~$\\mu$m sources, detected at the $>6 \\sigma$ level. We found counterparts for 545 among 1000 sources from the analyzed catalog. We discussed the properties of these sources and tried to derive conclusions about the average properties of the sample. \\item The point source ADF-S catalog itself appeared to be quite reliable in the position determination, with most of the counterparts located at the angular distance significantly lower than the nominal resolution of the FIS detector. A small number of counterparts detected at higher angular distances can be possibly an effect of contamination. In the position determination we observe a small (of an order of 4'', but systematic bias, which depends on the location in the ADF-S. This small systematics should be taken into account in the future works on these data. \\item  We conclude that FIR sources detected in the ADF-S are mostly nearby galaxies. This conclusion is supported, first of all, by source number counts in the FIR, but also a large number of bright optical counterparts and the redshift distribution of counterparts. It should be noted, however, that the completeness of the identified sample, close to 100\\% for sources brighter than 0.1 Jy, falls down rather rapidly for fainter sources, to 55~\\% for the whole 10$\\sigma$ catalog. \\item The population of identified galaxies appears surprisingly \"normal\", similar to one expected for local optically bright galaxies. The main differences are: \\\\ {\\it a)} a significantly lower percentage of elliptical galaxies, which can be explained by the fact they are less dusty than other galaxies, \\\\{\\it b)} much higher percentage of peculiar galaxies, which can be attributed to strong star-forming activity  of these objects, related to intergalactic interactions, and hence stronger radiation of cold dust in them). \\\\ c) We detected also a fraction of lenticular galaxies (all from the cluster of galaxies Abell~S0463) which is practically the same as expected in the optically bright galaxy population. It suggests that these galaxies contain a significant amount of cold dust and supports more complex models of their formation than a simple secular evolution. \\item The estimated source confusion is on the level higher than 34~\\% of the number of identified sources. It suggests that to remove the effect of the source confusion from the IR flux measurements and for a proper estimation of local environment of the FIR-bright galaxies dedicated observations of their closest neighborhood will be required. \\item The SEDs of the identified sources display a variety of properties. In the first approach, the examined galaxies seem to be either extremely quiescent or very active in star formation. The lenticular galaxies usually belong to the % actively star forming group. The analysis suggests that to explain well the FIR properties of otherwise \"normal\" galaxies from our sample new updated models should be developed. \\end{enumerate}  \\begin{figure}[tb] \\centering \\includegraphics[width=8.5cm]{alpha_v1.eps} \\caption{The histogram of the parameter $\\alpha$ of the model by \\cite{dale2002}. }\\label{dale_alpha} \\end{figure}  \\begin{figure}[tb] \\centering \\includegraphics[width=8.5cm]{PAH.eps} \\caption{The histogram of the amount of polycyclic aromatic hydrocarbons (PAH) in the dust of the analyzed galaxies, according to the model by \\cite{li2001}.}\\label{pah} \\end{figure}"
    },
    "0911/0911.2371_arXiv.txt": {
        "abstract": "A new method is explored to detect extensive air showers: the measurement of radio waves emitted during the propagation of the electromagnetic shower component in the magnetic field of the Earth.  Recent results of the pioneering experiment LOPES are discussed.  It registers radio signals in the frequency range between 40 and 80 MHz.  The intensity of the measured radio emission is investigated as a function of different shower parameters, such as shower energy, angle of incidence, and distance to shower axis.  In addition, new antenna types are developed in the framework of \\Lstar and new methods are explored to realize a radio self-trigger algorithm in real time. ",
        "introduction": "An intense branch of astroparticle physics is the study of high-energy cosmic rays to reveal their origin, as well as their acceleration and propagation mechanisms \\cite{behreview,cospar06,wuerzburg}. At energies exceeding $10^{14}$~eV cosmic rays are usually studied by indirect measurements --- the investigation of extensive air showers initiated by cosmic particles in the atmosphere. Different techniques are applied, like the measurements of particle densities and energies at ground level, or the observation of \\Cerenkov and fluorescence light. An alternative technique has been recently revitalized --- the detection of radio emission from extensive air showers at energies exceeding $10^{16}$~eV.  \\begin{figure}[t] \\epsfig{file=lopes-lat.eps,width=0.96\\columnwidth} \\caption{Measured field strength as a function of the distance to shower axis for an individual shower \\cite{icrc09-nehls}.} \\label{lat} \\end{figure}  Radio emission from air showers was experimentally discovered in 1965 at a frequency of 44 MHz \\cite{jelleynature}. The early activities in the 1960s and 1970s are summarized in \\cite{allanrev}.  Only recently, fast analog-to-digital converters and modern computer technology made a clear detection of radio emission from air showers possible.  LOPES, a LOFAR Prototype Station had shown that radio emission from air showers can be detected even in an environment with relatively strong radio frequency interference (RFI) \\cite{radionature}. Further investigations of the radio emission followed with LOPES \\cite{badearadio,petrovicinclined,buitinkthunder,niglfreq,nigldirect} and the CODALEMA experiment \\cite{codalema,Ardouin:2006nb}, paving the way for this new detection technique, see also \\cite{Falcke:2008qk}.  The LOPES experiment registers radio signals in the frequency range from 40 to 80~MHz \\cite{lopesspie}. In this band are few strong man made radio transmitters only, the emission from air showers is still strong (it decreases with frequency), and background emission from the Galactic plane is still low. An active short dipole has been chosen as antenna.  An inverted V-shaped dipole is positioned about 1/4 of the shortest wavelength above an aluminum ground plate. In this way a broad directional beam pattern is obtained.  LOPES comprises 30 antennas \\cite{nehlspune} located on site of the KASCADE-Grande experiment \\cite{kascadenim,grande} one of the best air shower experiments operating in the energy range between $10^{14}$ and $10^{18}$~eV. The LOPES data acquisition is triggered by large air showers registered with KASCADE-Grande. The latter measures the showers simultaneously to LOPES and delivers precise information on the shower parameters, such as shower energy as well as position and inclination of the shower axis.  All antennas, including the complete analog electronics chain, have been individually calibrated with a reference radio source \\cite{Nehls:2008ix}.  Most likely, the dominant emission mechanism of the radio waves in the atmosphere is radiation due to the deflection of charged particles in the Earth's magnetic field (geosynchrotron radiation) \\cite{huege2003,huege2005,huegefalcke}. In the frequency range of interest the wavelength of the radiation is large compared to the size of the emission region: the typical thickness of the air shower disc is about 1 to 2~m only. Thus, coherent emission is expected which yields relatively strong signals at ground level. ",
        "conclusions": ""
    },
    "0911/0911.2830_arXiv.txt": {
        "abstract": "We have carried out long-term (14 years) V and R photometric monitoring of 12 carbon-rich proto-planetary nebulae.  The light and color curves display variability in all of them. The light curves are complex and suggest multiple periods, changing periods, and/or changing amplitudes, which are attributed to pulsation. A dominant period has been determined for each and found to be in the range of $\\sim$150 d for the coolest (G8) to 35$-$40 d for the warmest (F3). A clear, linear inverse relationship has been found in the sample between the pulsation period and the effective temperature and also an inverse linear relationship between the amplitude of light variation and the effective temperature. These are consistent with the expectation for a pulsating post-AGB star evolving toward higher temperature at constant luminosity. The published spectral energy distributions and mid-infrared images show these objects to have cool (200 K), detached dust shells and published models imply that intensive mass loss ended a few thousand years ago. The detection of periods as long as 150 d in these requires a revision in the published post-AGB evolution models that couple the pulsation period to the mass loss rate and that assume that intensive mass loss ended when the pulsation period had decreased to 100 d.  This revision will have the effect of extending the time scale for the early phases of post-AGB evolution. It appears that real time evolution in the pulsation periods of individual objects  may be detectable on the time scale of two decades. ",
        "introduction": "Proto-planetary nebulae (PPNs) are objects in transition between the asymptotic giant branch (AGB) and planetary nebula (PN) phases of stellar evolution. In this phase, the high AGB mass loss has ended and the star is surrounded by a detached, expanding shell of gas and dust. Prior to the {\\it Infrared Astronomical Satellite} ({\\it IRAS}) satellite in 1983, this phase remained as an observational ``missing link'' in the post-main sequence evolution of stars of intermediate (1$-$8 M$_\\sun$) mass. However, the use of the {\\it IRAS} database has allowed the identification of PPNs on the basis of their infrared excesses and the colors of their circumstellar dust shells \\citep{partha86, veen89, hri89}. This has led to the identification of $\\sim$50 PPNs and an equal number of additional candidates \\citep{hri94, meixner99, ueta00, garlar97}. They show a double-peaked spectral energy distribution (SED), with a peak in the mid-infrared from re-radiated emission from the circumstellar dust and a peak in the visible$-$near-infrared from the reddened photosphere. The spectral types range from B0 to K0. Reviews of the properties of PPNs have been published by \\citet{kwok00} and \\citet{hri03} and they are included in the more general review of post-AGB stars by \\citet{vanwin03}.  Millimeter radio observations of CO, OH, and HCN have reveled the oxygen-rich or carbon-rich chemistry of the circumstellar shells \\citep{lik89,lik91,omont93}. High-resolution visible spectroscopy has shown the objects to be iron-poor and led to the determination of the chemistry of the photospheres \\citep{vanwin03}. High-resolution imaging with the {\\it Hubble Space Telescope} ({\\it HST}) has revealed the morphology of the nebulae, which are generally elliptical or bipolar with sizes up to a few arcseconds in radius \\citep{ueta00,su01,sahai07,siod08}.  We are carrying out a long-term study of the variability of PPNs. It initially began with radial velocity monitoring to search for possible binary companions.  This was motivated in part by the idea that a binary companion could cause the dusty torus seen in bipolar PPNs, which apparently restricts the outflow and results in the bipolar morphology \\citep{livio88}. The results of this initial study did not reveal any binaries, but did reveal periodic velocity variations \\citep{hri00a,hri09e}.  The radial velocity monitoring was followed-up by the initiation of a study of the light variations in PPNs, which has been ongoing for the past 15 years. We identified an initial sample of $\\sim$40 PPN candidates and have monitored them for light variability. All of the PPNs are found to be light variables, and the preliminary results have been presented \\citep{hri00a,hri02,hri07a}. Arkhipova and collaborators \\citep[e.g.,][]{ark00} have also presented the results of their light curve studies of several PPNs .  The results of our ongoing studies will be published in a series of papers.  In this initial publication, we focus on 12 carbon-rich PPNs. We present in detail our observing program and discuss the observed light curves, which we investigated for periodicity. The results are then discussed in the broad context of post-AGB variability and evolution. In two companion papers, radial velocity data are presented for four of the brightest of these 12 PPNs, and together they are used in the interpretation of the discovered pulsational variability \\citep{hri09a, hri09b}. ",
        "conclusions": "In this study, we carried out a photometric V and R monitoring survey of 12 carbon-rich PPNs over an interval of 14 years. This is the first systematic, long-term ($>$10 yr)  light curve study of a homogeneous group of PPNs. The targets were all of spectral type F and G. Below are the primary results of this study.  1. All of these PPNs varied in brightness and color, with the general result that they are redder when fainter.  2. Periods of the variations were found for all of the targets (except perhaps AFGL 2688). These range from $\\sim$38 to 153 days, with a maximum V variation (peak-to-peak within a season) ranging from 0.13 to 0.67 mag. The objects with the later spectral types, mid- to late-G, have the longer periods and the larger variations.  3. The objects show evidence for multiple periods and/or changing periods and/or changing amplitudes. Two of them have light curves that suggest a longer beat period (IRAS 22272+5435 and 22223+4327) and two of them have light curves with a suggestion of intervals of alternating deeper and shallower minimum (IRAS 20000+3239 and Z02229+6208). Nevertheless, all except the two with the shortest periods (IRAS 07134+1005 and 19500$-$1709) show a clearly dominant period in the data set.  These results are consistent with the variation being due to pulsation and not binary interactions.  4. An approximately linear correlation was found between the period of variation and the effective temperature over the temperature range observed in this study (5250$-$8000 K), with the shorter periods associated with higher temperatures. The range of light variation is also seen to decrease approximately monotonically with temperature.  5. The presence of pulsation periods as long as 150 d in PPNs with detached shells implies that the intensive AGB mass loss had ended before the period had decreased to 150 d. This has significant implications for modeling post-AGB evolution rates, since the rates depend upon the  level of post-AGB mass loss.  Since some post-AGB evolution codes tie the pulsation period to a mass loss prescription (i.e., Bl\\\"{o}cker) that assumes that intensive mass loss has ended by the time that the pulsation period has reduced to 100 d, these models will need to be revised. This would have the result of significantly increasing the initial stages of post-AGB evolution.  6. An average rate of change of the pulsation period with time can be determined by combining the empirically-determined rate of variation of period with temperature and the model-dependent rate of change of temperature with time.  These lead to an approximate rate of period change of $-$0.2 d/yr, a rate that is potentially detectable within a few decades.  Our observations of individual PPNs over this 14 year interval are complicated by the evidence for multiple periods, but give some tentative results that are not inconsistent with this value.  As can be seen, the study of the light curves of PPNs and the determination of their periods has the potential to constrain the models of post-AGB stellar evolution and to reveal evolution in real time.  In addition, these light curves together with radial velocity curves, when compared with appropriate pulsation models, have the potential to determine fundamental properties of the stars, such as mass and luminosity.  These will be explored further in subsequent papers. We are continuing to monitor the light curves of many of this PPNs to extend the time baseline."
    },
    "0911/0911.0004_arXiv.txt": {
        "abstract": "One of the most fundamental problems in the study of Kuiper belt objects (KBOs) is to know their true physical size. Without knowledge of their albedos we are not able to distinguish large and dark from small and bright KBOs. \\emph{Spitzer} produced rough estimates of the sizes and albedos of about 20 KBOs, and the \\emph{Herschel} space telescope will improve on those initial measurements by extending the sample to the $\\sim$150 brightest KBOs. \\emph{SPICA}'s higher sensitivity instruments should allow us not only to broaden the sample to smaller KBOs but also to achieve a statistically significant sample of KBO thermal light curves (\\emph{Herschel} will measure only six objects). A large sample covering a broad range of sizes will be key to identify meaningful correlations between size and other physical and surface properties that constrain the processes of formation and evolution of the solar system. ",
        "introduction": "In this paper I discuss the importance of the upcoming \\emph{`Space Infrared Telescope for Cosmology \\& Astrophysics'} (hereafter, \\emph{SPICA}) for the study of the icy small bodies of the outer solar system. Here, I focus on the study of Kuiper belt objects (KBOs) but the same ideas can be applied to any atmosphereless bodies. I will mainly discuss how \\emph{SPICA} can help us to measure the sizes of KBOs, and how that leads to more accurate estimates of the size distribution and total mass of KBOs. I will also mention how this new infrared space telescope might probe the rotational properties, chemical composition and thermophysical parameters of those icy bodies. Other uses of a space-based infrared telescope for solar system studies are detailed elsewhere, in the context of \\emph{SPICA}'s precursor observatory, \\emph{Herschel} \\citep{2009EAS....34..133L}. In \\S2 to \\S4 I begin by summarising what we know about KBOs and why their study is interesting and important.  Although many accounts exist elsewhere I believe these proceedings should include a broad overview of the subject. In \\S5 and \\S6 I discuss how \\emph{SPICA} may contribute to the study of KBOs. ",
        "conclusions": "The enhanced sensitivity of the upcoming \\emph{`Space Infrared Telescope for Cosmology \\& Astrophysics'} (\\emph{SPICA}) will play a key role in the study of Kuiper belt objects.  Multi-wavelength thermal observations will enable us to measure the albedos and sizes of several hundreds of bodies and to investigate their surface composition and thermophysical properties. The high sensitivity -- two orders of magnitude better than \\emph{Herschel} -- will allow the detection of KBOs as small as a few tens of kilometres across, very close to the sizes of cometary nuclei. Sizes measurements of binary Kuiper belt objects will lead to estimates of their bulk density, a first step into the interior structure of these bodies. Time resolved thermal observations of a statistical ensemble of KBOs will help break the degeneracy between albedo variability and shape as the cause for observed light curves, which is important for assessing the shape and angular momentum distributions of Kuiper belt objects.  By extending the far-infrared domain to potentially hundreds of Kuiper belt objects and associated families, \\emph{SPICA} should allow us to identify meaningful correlations and patterns between orbital, physical and chemical properties and thus help us probe the formation and evolutionary processes that operated at the early stages of our solar system."
    },
    "0911/0911.5621_arXiv.txt": {
        "abstract": "Characterizing the surfaces of rocky exoplanets via their scattered light will be an essential challenge in investigating their habitability and the possible existence of life on their surfaces. We present a reconstruction method for fractional areas of different surface types from the colors of an Earth-like exoplanet. We create mock light curves for Earth without clouds using empirical data. These light curves are fitted to an isotropic scattering model consisting of four surface types: ocean, soil, snow, and vegetation. In an idealized situation where the photometric errors are only photon shot noise, we are able to reproduce the fractional areas of those components fairly well. The results offer some hope for detection of vegetation via the distinct spectral feature of photosynthesis on the Earth, known as the red edge. In our reconstruction method, Rayleigh scattering due to the atmosphere plays an important role, and for terrestrial exoplanets with an atmosphere similar to our Earth, it is possible to estimate the presence of oceans and an atmosphere simultaneously. ",
        "introduction": "\\label{s:intro}  A major milestone in exoplanet research will be the discovery of Earth-like planets. Indeed planets a few times as massive as the Earth, {\\it super-earths} have already been discovered using the radial velocity method; for instance, $M_{\\rm p} \\sin i$ of GJ581e is estimated to $\\sim 2M_{\\oplus}$ where $M_{\\rm p}$ is the planetary mass and $i$ is the inclination of its orbit \\citep{mayor2009}. Microlensing and transit methods are well-suited for the detection of such low-mass planets; MOA-2007-BLG-192-Lb \\citep{bennett2008} and CoRoT-7b \\citep{queloz2009} are reported to have masses of $3.3 {+4.9 \\atop -1.6} M_{\\oplus}$ and $4.8 \\pm 0.8 M_{\\oplus}$, respectively.  The {\\it Kepler} mission, launched on 2009 March 6, aims to detect a number of terrestrial planets with masses on the order of $M_{\\oplus}$. Such low-mass planets are very likely to be rocky and may have bodies of liquid water on their surfaces if they orbit in Habitable Zone \\citep[HZ, e.g.,][]{kasting1993,kasting2003} of the primary star. Their discovery will definitely trigger ever more serious investigations of techniques to search for signatures of life, or {\\it biomarkers}.  The most conventional and extensively studied biomarkers are based on spectroscopic identification of molecular species in the planetary atmospheres, such as ${\\rm O}_2$ and ${\\rm O}_3$ \\citep[e.g.,][]{leger1993, desmarais2002}.  For the most favorable known exoplanetary systems, detection of such atmospheric absorption features is already within reach during the transit or the secondary eclipse; absorption features of ${\\rm Na}$, ${\\rm C}$, ${\\rm O}$, and ${\\rm H}$ in a hot Jupiter HD209458b \\citep{charbonneau2002, vidal2003, vidal2004, ballester2007} and ${\\rm H}_2{\\rm O}, {\\rm CH}_4$, ${\\rm CO}$, and ${\\rm CO}_2$ in HD189733b \\citep{tinetti2007, swain2008, swain2009} have been reported. Spectroscopy of non-transiting terrestrial planets orbiting in the HZ will likely require major and specialized observatories in space, but such facilities are being actively studied and evaluated by both NASA and ESA at present \\citep[e.g.,][]{tpf-c, tpf-i, darwin}.  An even more ambitious and direct approach to the search for life may be possible via multi-band photometry of terrestrial exoplanets that is observationally less demanding.  In particular, the light scattered by the surface of a planet carries important information on properties of the surface.  \\citet{ford2001} computed for the first time the diurnal variation of scattered light by the Earth observed at a distance of 10 pc.  They showed that the fractional variation of light curves of the Earth is 10 \\% -20\\%, which indeed agrees with the result of Earthshine observations \\citep[e.g.,][]{goode2001}.  More importantly, these model light curves exhibited an increase in brightness due to the presence of vegetation around at $\\lambda = $750 nm.  This sharp increase of the albedo is known as the {\\it red edge}. It is a fairly generic feature of vegetation on the Earth and is due to bio-pigments associated with photosynthesis. The red edge feature has been detected by Earthshine observations \\citep{woolf2002,arnold2002,rodriguez2006}. It was subsequently discussed as a possible biomarker in extrasolar terrestrial planets by \\citet{seager2005} and \\citet{kiang2007a,kiang2007b}. \\citet{tinetti2006a, tinetti2006b} and \\citet{rodriguez2006} performed more accurate simulations of the scattered radiation spectra of the Earth. They paid particular attention to the red edge, and discussed the detectability of its spectral feature.  \\citet{palle2008} focused on the determination of the rotation period of an extrasolar terrestrial planet from the time variation of its scattered light. They concluded that the period can be reliably estimated in the presence of realistic partial cloud coverage of the planetary surface. They even found that the global circulation of the cloud pattern systematically modulates the estimate of the spin rotation period of the planet.  More recently, \\citet{cowan2009} approached the problem in a different and interesting manner. They performed a principal component analysis (PCA) on the real light curves of the Earth observed by the Extrasolar Planet Observation and Deep Impact Extended Investigation (EPOXI) mission. They identified two major eigenspectra that roughly correspond to ocean and land on the Earth and extracted a rough distribution of these components as a function of longitude.  This was the first attempt to solve the inverse problem and to extract the surface properties from the observed data in a model-independent fashion.  \\citet{williams2008} and \\citet{oakley2009} paid attention to the variation of scattered light according to the orbital motion of the planet. They pointed out that the reflectivity of ocean dramatically increases at a crescent phase and that the existence of the liquid water may be observationally inferred through the effect.  \\citet{oakley2009} also suggested a possibility of reconstructing the land/ocean distribution along longitude using the gap between the reflectivity of ocean and that of land.  In this paper, we develop still another method to estimate the fractional areas of different surface types on extrasolar planets from multi-band photometry. We are particularly interested in terrestrial exoplanets in habitable zone and consider the detectability of vegetation using the red-edge feature. Our attempt is to reproduce the scattered light curves as a sum of the four surface types, i.e., ocean, soil, vegetation, and snow.  While PCA attempts to extract only orthogonal eigenspectra in a model-independent manner, the correspondence with real surface types is not clear. Indeed, we would like to answer an ambitious, but well-defined, scientific question: ``If we discover an Earth-like exoplanet in the near future, to what extent can we reconstruct the surface information, in particular the presence of vegetation observationally?''. For that purpose, we intentionally and specifically adopt the major surface types of the Earth, including vegetation, into the analysis even though our method is thus inevitably model dependent.  We discuss the feasibility of our method by applying it to mock scattered light curves of the {\\it cloudless} Earth. We hope that our current results are relevant for a future TPF mission \\citep{tpf-c} such as the Occulting Ozone Observatory which aims at multi-band photometry of Earth-like planets \\citep{kasdin2010}.  The rest of this paper is organized as follows. Section \\ref{s:simulation} describes the generation of mock light curves of the Earth. The methodology, assumptions, and results of our inversion model to reconstruct the fractional areas are presented in Section \\ref{s:inverse}.  We discuss the importance of the photometric band selection and comparison with PCA in Section \\ref{s:dis}.  Finally we summarize the conclusion and discuss future directions of investigation in Section \\ref{s:sum}. ",
        "conclusions": "\\label{s:dis}    \\subsection{Band Selection} \\label{s:band}  Our inversion method relies entirely on the difference of the wavelength-dependence of the scattering spectra among ocean, soil, vegetation, and snow, or their colors in short. Therefore the selection of the observed bands is crucial.  First we repeat the same analysis performed in Section \\ref{s:inverse} but with four bands.  In the case of five bands, we can fit up to five unknowns; we selected fractional areas for the four types and $\\tau$ as the five fitting parameters. In the case of 4 bands, however, we cannot fit $\\tau$ simultaneously and thus fix $\\tau=\\tau_0$. The result is shown in Figure \\ref{fig:comparison2}.  The left and right panels correspond to the bluer bands ($j=$ 1, 2, 3, and 4) and redder bands ($j=$ 4, 5, 6, and 7), respectively. The difference can be easily understood; the bluer bands are more sensitive to the effect of Rayleigh scattering and the fractional area of ocean is recovered fairly well. Also the red edge feature between bands 3 and 4 still carries the vegetation signature, while the fractional area of soil becomes more uncertain because it is brighter in the redder bands. The result with the redder bands shows consistently opposite generic features; the lack of the blue bands and the red edge makes it difficult to detect the signature of ocean and vegetation, respectively, while the soil is more easily reconstructed.   \\begin{figure}[!tbh] \\centerline{\\includegraphics[width=16.0cm] {fig14.eps}} \\caption{Reconstructed fractional areas using simulated light curves in two different sets of four bands; Left: bluer four bands ($j=$ 1, 2, 3, and 4), Right: redder four bands ($j=$ 4, 5, 6, and 7).} \\label{fig:comparison2} \\end{figure}  A closer look at Figure \\ref{fig:comparison2} indicates some degeneracy between reflection spectra of selected surface types (Figure \\ref{fig:effalbedo}). This is more clearly illustrated in Figure \\ref{fig:band3}, where we use three bands ($j=2$, 3, and 4) only. In this case we cannot determine four surface types, and we choose ocean, soil, and vegetation in the left panel, and ocean, vegetation, and snow in the right panel, while we neglect the remaining surface type from the fit. The result naturally is degraded compared with the 4-band and 5-band cases. The vegetation signature is still there because we chose bands 3 and 4 that bracket the red edge, and the fractional areas of soil and snow compensate for each other. The fact that the spectra of snow and soil are similar in shorter wavelength bands (except for their amplitude) partly explains the behavior of the lower panels in Figure \\ref{fig:band3}. While this similarity may come largely from our single scattering approximation as described in Section \\ref{ss:inversion}, we obtained a similar result even when we use the snow effective albedo of $\\tau=0$. Therefore the degeneracy between snow and soil is fairly generic as long as we use the short wavelength bands in the reconstruction.  \\begin{figure}[!tbh] \\centerline{\\includegraphics[width=16.0cm] {fig15.eps}} \\caption{Reconstructed fractional areas using simulated light curves in three bands; Left: ocean, soil, and vegetation are considered. Right: ocean, vegetation, and snow are considered.} \\label{fig:band3} \\end{figure}  \\subsection{Principal Component Analysis (PCA)} \\label{s:pca}  \\citet{cowan2009} performed a PCA of the multi-band light curves of the Earth as observed by the EPOXI mission.  They extracted two major eigenspectra in a model-independent manner. We also performed PCA on the same data-set, and obtained results consistent with those of \\citet{cowan2009}. We then performed PCA of our mock light curves and again extracted two major eigenspectra; these are shown in Figure \\ref{fig:eigenvector5band}.  In order to compare our current methodology with PCA, we decompose these two eigenspectra into a combination of the effective albedo of the four surface types with $\\tau = \\tau_0$ (solid lines in Figure \\ref{fig:effalbedo}). We find that the first eigenspectrum roughly corresponds to (soil+vegetation-ocean), and the second one roughly corresponds to (vegetation-(soil+snow+ocean)) as displayed in Figure \\ref{fig:pcadecompose}.  This result indicates that the extracted eigenspectra do not necessarily correspond to any single surface type.  This is not surprising, of course. While PCA extracts orthogonal eigenspectra by definition, the wavelength dependence of albedos of real materials is not orthogonal in general. Moreover, they are likely to be degenerate in PCA because the fractional areas are complementary; for example, the fractional area of ocean decreases when fractional area of land increases. In other words, the time dependences of these components are necessarily correlated. Therefore, it seems natural that the first eigenspectrum is roughly (land-ocean). Our method is quite model dependent, but it allows decomposition of the light curves into physical components that can be interpreted in a simple way. While the model independence of PCA is a great advantage, the final interpretation is not straightforward. Further comparison with PCA is beyond the scope of this paper, but clearly these two methodologies are very complementary.  \\begin{figure}[htbp] \\begin{center} \\includegraphics[width=12cm]{fig16.eps} \\caption{ Left: the eigenspectra extracted by PCA of our fiducial mock light curves displayed in Figure \\ref{fig:scatter2}. The contribution of the first eigenspectrum (solid) is 94.3 \\% and that of the second one (dashed) is 5.7 \\%. Right: the time variation of the first eigenspectrum (solid) and the second eigenspectrum (dashed). } \\label{fig:eigenvector5band} \\end{center} \\end{figure}  \\begin{figure}[htbp] \\begin{center} \\includegraphics[width=12cm]{fig17.eps} \\caption{ Left: decomposition of the first eigenspectrum of Figure \\ref{fig:eigenvector5band} where the coefficients are ocean, -2.19; snow, 0.03; vegetation, 0.86; soil, 1.07. Effective albedos with $\\tau = \\tau _0 $ are used. Right: decomposition of the second eigenspectrum of Figure \\ref{fig:eigenvector5band} where the coefficients are ocean, -1.83; snow, -0.58; vegetation, 2.36; soil, -1.54. Effective albedos with $\\tau = \\tau _0 $ are used. } \\label{fig:pcadecompose} \\end{center} \\end{figure}"
    },
    "0911/0911.2709_arXiv.txt": {
        "abstract": "{} {Our aim is to study the very high energy (VHE; $E>100$ GeV) $\\gamma$-ray emission from BL\\,Lac objects and the evolution in time of their broad-band spectral energy distribution (SED).} {VHE observations of the high-frequency peaked BL\\,Lac object PKS\\,2005$-$489 were made with the High Energy Stereoscopic System (HESS) from 2004 through 2007. Three simultaneous multi-wavelength campaigns at lower energies were performed during the HESS data taking, consisting of several individual pointings with the XMM-Newton and RXTE satellites.} {A strong VHE signal, $\\sim$17$\\sigma$ total, from PKS\\,2005$-$489 was detected during the four years of HESS observations (90.3 hrs live time). The integral flux above the average analysis threshold of 400 GeV is $\\sim$3\\% of the flux observed from the Crab Nebula and varies weakly on time scales from days to years.  The average VHE spectrum measured from $\\sim$300 GeV to $\\sim$5 TeV is characterized by a power law with a photon index, $\\Gamma = 3.20\\pm0.16_{\\rm stat}\\pm0.10_{\\rm syst}$. At X-ray energies the flux is observed to vary by more than an order of magnitude between 2004 and 2005. Strong changes in the X-ray spectrum ($\\Delta$$\\Gamma_{\\mathrm X} \\approx 0.7$) are also observed, which appear to be mirrored in the VHE band.} {The SED of PKS\\,2005$-$489, constructed for the first time with contemporaneous data on both humps,  shows significant evolution. The large flux variations in the X-ray band, coupled with weak or no variations in the VHE band and a similar spectral behavior, suggest the emergence of a new, separate, harder emission component in September 2005.} ",
        "introduction": "PKS 2005$-$489 is one of the brightest BL Lac objects, at all wavelengths, in the Southern Hemisphere. It was initially discovered as a strong radio source in the Parkes 2.7\\,GHz survey \\citep{discovery_paper} and later identified as a BL Lac object \\citep{BL_id}. It belongs to the complete 1-Jy radio catalog \\citep{stickel}, and is one of the few extragalactic objects detected in the EUV band \\citep{euv}. Its redshift, $z=0.071$ \\citep{redshift}, is determined from weak, narrow emission lines observed during a low optical state.  PKS 2005$-$489 is classified as a High-frequency peaked BL Lac object \\citep[HBL;][]{giommipadovani94}, because of its high X-ray-to-radio flux ratio \\citep{Sambruna_paper} and because its broad-band spectral energy distribution (SED) peaks in the UV--soft X-ray band. As is typical of HBLs, the X-ray spectrum is dominated by the synchrotron emission of high-energy electrons. The second SED hump is expected to peak in the GeV$-$TeV $\\gamma$-ray band, and is commonly believed to be produced by the same electrons up-scattering via the inverse Compton mechanism seed photons of lower energy.  In the X-ray band, PKS\\,2005$-$489 has been studied extensively, showing an extreme flux and spectral variability. Large flux variations with correlated spectral hardening of the generally steep spectrum (photon index\\footnote{The power-law spectrum is described as $N(E)=N_0\\,E^{-\\Gamma}$} $\\Gamma=2.7-3.1$) were observed during five EXOSAT observations \\citep{giommi90,sambruna94}. Two ROSAT observations in 1992 confirmed the EXOSAT results and similarly show a soft spectrum \\citep[$\\Gamma\\simeq3$;][]{Sambruna_paper}. A harder spectrum ($\\Gamma=2.3$ from 2 to 10 keV) was observed in September 1996, during \\sax observations of a brighter X-ray state \\citep{padovani}. In October$-$November 1998  PKS\\,2005$-$489 underwent a period of exceptional activity, with several strong X-ray flares.  RXTE monitoring observations \\citep{perlman} were performed during the entire epoch, and the 2$-$10 keV flux reached $\\sim$3\\ergs{-10}, approximately 30 times higher than the ROSAT values.  These RXTE observations yielded a detection up to 40 keV and showed variations in the photon index between $\\Gamma=2.3$ and 2.8.  An X-ray flare alert also triggered \\sax observations on November 1$-$2, 1998.  The X-ray spectrum was measured between 0.1 to 200 keV and was characterized by a curved shape and harder photon indices, in both the soft ($\\Gamma_1=2.0$) and hard ($\\Gamma_2=2.2$) X-ray bands \\citep[$<$ and $>$2 keV;][]{gt}. More recently, observations  with the \\swift satellite have generally shown PKS\\,2005$-$489 in a low-flux, steep-spectrum state \\citep[$\\Gamma\\simeq3$,][]{massaro08}.   Strong correlations between X-ray and $\\gamma$-ray emission have been observed in many  HBL \\citep[e.g.][]{pian98,maraschi99,1959,fossati08,2155chandra}, and are typically expected in a synchrotron-Compton scenario. Therefore the very bright flux and large flux/spectral variability at X-ray energies make PKS\\,2005$-$489 one of the most promising targets for observing a similar behavior in the $\\gamma$-ray domain.  In the VHE ($E>100$ GeV) band, PKS\\,2005$-$489 was detected neither during observations made between 1993 and 2000 by either the CANGAROO or Durham groups \\citep{cangaroo1,cangaroo2,cangaroo3,durham}, nor by HESS with a partial array during its commissioning in 2003. In 2004, HESS discovered VHE $\\gamma$-ray emission from PKS\\,2005$-$489 with a significance of 6.7$\\sigma$, at a flux of a few percent of the Crab Nebula  \\citep{hess_discovery}. The measured spectrum was soft ($\\Gamma=4.0$). In the MeV-GeV band, PKS\\,2005$-$489 is one of the few HBL detected by  EGRET. % However, the observed significance is marginal: 3.7$\\sigma$ above 100 MeV ~\\citep{egret1} and 4.1$\\sigma$ above 1 GeV \\citep{egret2}. It was instead detected by {\\it Fermi} with high significance ($>$10$\\sigma$), during the first three months of operation \\citep[2008, Aug.--Oct.,][]{lbas}.  Because of the high potential for strong $\\gamma$-ray activity, PKS\\,2005$-$489 has been monitored at VHE by HESS every year since its detection in 2004. Several campaigns of coordinated observations with the X-ray satellites \\xmm~ and RXTE were also performed. These simultaneous  observations, sampling the same particle distribution through two different emission processes, represent a powerful diagnostic tool for probing the conditions of the inner blazar jet, especially during flaring events \\citep{coppi}. \\begin{figure} \\centering \\includegraphics[width=0.5\\textwidth]{figures/PKS2005_thsq.eps} \\caption{Distribution of $\\theta^2$ for on-source events (points with statistical errors) and normalized off-source events (shaded) from observations of PKS\\,2005$-$489.  The curve represents the $\\theta^2$ distribution expected from simulations of a $\\gamma$-ray point source (photon index $\\Gamma=3.20$) at a zenith angle of $40^{\\circ}$. The dashed line represents the cut on $\\theta^2$ applied to the data.} \\label{thtsq_plot} \\end{figure} \\begin{figure} \\centering \\includegraphics[width=0.5\\textwidth]{figures/PKS2005_monthly_lc.eps} \\caption{Integral flux, I($>$400 GeV), measured by HESS from PKS\\,2005$-$489 during each dark period of observations (i.e. moonless night time within a month). Only the statistical errors are shown. The horizontal line represents the average flux for all the HESS observations. The four horizontal line segments represent the average annual flux observed during the corresponding years (see Table \\ref{results}).} \\label{monthly_plot} \\end{figure}  In this article the results of all HESS observations taken from 2004 through 2007 are presented, together with the results of the multi-wavelength observations. These campaigns characterize the SED of PKS\\,2005$-$489 during different states and, for the first time, sample both humps of the SED simultaneously. A re-analysis of the 2004 HESS data is also provided, which benefits from an improved calibration of the absolute energy scale with respect to the previously published result. \\begin{figure*} \\centering \\includegraphics[width=0.6\\textwidth]{figures/PKS2005_2004lc.eps}\\\\ \\includegraphics[width=0.6\\textwidth]{figures/PKS2005_2005lc.eps}\\\\ \\includegraphics[width=0.6\\textwidth]{figures/PKS2005_2006lc.eps}\\\\ \\includegraphics[width=0.6\\textwidth]{figures/PKS2005_2007lc.eps}\\\\ \\caption{Integral flux, I($>$400 GeV), measured by HESS from PKS\\,2005$-$489 during each night of observations.  The individual plots represent each of the four years (2004$-$2007) of data taking. Only the statistical errors are shown.  The horizontal lines represent the average annual flux observed during the respective year.  The vertical lines at MJD 53282, 53640, and 53642 represent the nights of \\xmm~ observations. The epoch of the RXTE observations is between the vertical lines at MJD 53609 and 53623.} \\label{nightly_plots} \\end{figure*}    \\begin{table*}[t] \\begin{minipage}[t]{2.0\\columnwidth} \\caption{Results from long-term HESS observations of PKS\\,2005$-$489.} \\label{results} \\centering \\begin{tabular}{c c c c c c c c c c c c c} \\hline\\hline \\noalign{\\smallskip} Dark & MJD & MJD & Time & On & Off & $\\alpha$ & Excess & Sig & I($>$400 GeV)\\footnote{The integral flux above 400 GeV is calculated using the excess of events above an energy threshold, which differs slightly from the observed excess.  The quoted error is statistical only and the 20\\% systematic error on the observed flux is not shown.} & Crab\\footnote{The percentage is calculated relative to the HESS Crab Nebula flux above 400 GeV \\citep{HESS_crab}.} & $\\chi^2$\\,\\,(NDF\\footnote{The $\\chi^2$, degrees of freedom (NDF), and corresponding $\\chi^2$ probability P($\\chi^2$) are given for fits of a constant to I($>$400 GeV) binned nightly within a dark period, or monthly within a year, or yearly within the total.}) & P($\\chi^2$)$^c$\\\\ Period & First & Last & [hrs] & & & & &  [$\\sigma$] & [10$^{-12}$\\,cm$^{-2}$\\,s$^{-1}$] & \\% & & \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} 06/2004 & 53171 & 53185 & 8.1  & 678 & 5877  & 0.0919 & 138 & 5.4 & $2.64\\pm0.50$ & 3.0 & 7.2\\,\\,(10) & 0.71\\\\ 07/2004 & 53199 & 53205 & 1.5  & 105 & 985   & 0.0909 & 15  & 1.5 & $1.06\\pm1.18$ & 1.2 & 1.4\\,\\,(1)  & 0.75\\\\ 09/2004 & 53255 & 53268 & 5.6  & 342 & 3171  & 0.0916 & 52  & 2.8 & $1.89\\pm0.78$ & 2.1 & 2.1\\,\\,(5)  & 0.84\\\\ 10/2004 & 53285 & 53292 & 9.0  & 569 & 5065  & 0.0916 & 105 & 4.5 & $2.58\\pm0.62$ & 2.9 & 1.5\\,\\,(4)  & 0.83\\\\ 07/2005 & 53582 & 53595 & 9.4  & 573 & 4504  & 0.0908 & 164 & 7.3 & $3.56\\pm0.67$ & 4.0 & 9.2\\,\\,(9)  & 0.42\\\\ 08/2005 & 53609 & 53618 & 5.1  & 286 & 2159  & 0.0989 & 73  & 4.5 & $2.86\\pm0.93$ & 3.2 & 4.6\\,\\,(4)  & 0.33\\\\ 09/2005 & 53639 & 53646 & 18.1 & 1072 & 9125 & 0.0928 & 225 & 7.1 & $2.93\\pm0.46$ & 3.3 & 13.1\\,\\,(6) & 0.041\\\\ 06/2006 & 53908 & 53909 & 1.3  & 127 & 824   & 0.0938 & 50  & 4.9 & $4.72\\pm1.32$ & 5.3 & 0.0\\,\\,(1)  & 0.97\\\\ 07/2006 & 53938 & 53940 & 4.4  & 397 & 2927  & 0.0901 & 133 & 7.3 & $4.33\\pm0.70$ & 4.8 & 1.3\\,\\,(2)  & 0.52\\\\ 08/2006 & 53967 & 53977 & 8.4  & 500 & 4261  & 0.0913 & 111 & 5.1 & $2.74\\pm0.57$ & 3.1 & 15.1\\,\\,(7) & 0.035\\\\ 09/2006 & 53995 & 54002 & 7.4  & 405 & 4116  & 0.0920 & 26  & 1.3 & $1.31\\pm0.58$ & 1.5 & 6.0\\,\\,(5)  & 0.30\\\\ 06/2007 & 54264 & 54270 & 5.3 & 333 & 3006 & 0.0924 & 55 & 3.1 & $1.02\\pm0.60$ & 1.1 & 1.9\\,\\,(6) & 0.93\\\\ 07/2007 & 54291 & 54304 & 4.4 & 309 & 2412 & 0.0964 & 76 & 4.5 & $2.51\\pm0.70$ & 2.8 & 15.2\\,\\,(8) & 0.056\\\\ 08/2007 & 54321 & 54329 & 1.8 & 93 & 849 & 0.0977 & 10 & 1.0 & $0.52\\pm0.92$ & 0.6 & 2.9\\,\\,(3) & 0.41\\\\ 09/2007 & 54345 & 54345 & 0.5 & 11 & 114 & 0.1000 & 0 & $-0.1$ & $<6.60$\\footnote{The upper limit is calculated at a 99.9\\% confidence level \\citep{UL_tech}.} & $<$7.4 & $-$ & $-$ \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} 2004 & 53171 & 53292 & 24.2 & 1694 & 15098 & 0.0917 & 310 & 7.7 & $2.37\\pm0.33$ & 2.6 & 2.0\\,\\,(3) & 0.57\\\\ 2005 & 53582 & 53646 & 32.6 & 1931 & 15785 & 0.0930 & 462 & 11.0 & $3.09\\pm0.35$ & 3.5 & 0.6\\,\\,(2) & 0.73\\\\ 2006 & 53908 & 54002 & 21.5 & 1429 & 12128 & 0.0914 & 320 & 8.8 & $2.84\\pm0.34$ & 3.2 & 13.5\\,\\,(3) & 0.0037 \\\\ 2007 & 54264 & 54345 & 12.0 & 746 & 6381 & 0.0947 & 141 & 5.3 & $1.50\\pm0.40$ & 1.7 &  3.8\\,\\,(3) & 0.28\\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} Total & 53171 & 54345 & 90.3 & 5800 & 49392 & 0.0924 & 1233 & 16.7 & $2.57\\pm0.18$ & 2.9 & 10.2\\,\\,(3) & 0.016 \\\\ \\noalign{\\smallskip} \\hline \\end{tabular} \\end{minipage} \\end{table*} ",
        "conclusions": "PKS\\,2005$-$489 is detected at VHE in each of the four years it was being observed by HESS (2004--2007).  The 2005--2007 data clearly confirm the VHE discovery reported by HESS \\citep{hess_discovery} and quadruple the statistics of the initial spectrum measurement.   Re-analysis of the previously published 2004 HESS data, using the improved calibration of the detector's energy scale, results in a $\\sim$3 times higher flux and a similar photon index.  The measured VHE $\\gamma$-ray flux is low ($\\sim$3\\% Crab) and only shows weak variations on time scales ranging from days to years. The observed time-averaged (2004--2007) VHE spectrum is soft, with a photon index $\\Gamma = 3.20\\pm0.16$. Although evidence of VHE spectral variations is marginal by itself, the VHE spectrum seems to track the X-ray slope variations when multi-wavelength coverage is available. Observations performed with \\xmm~ and RXTE in 2004 and 2005 reveal remarkable changes in the X-ray spectrum, but without shifting the location of the synchrotron peak with respect to historical observations.  Interpreting these measurements along with the HESS data suggests the emergence of a new jet component in the SED that is characterized by a harder electron spectrum.  This component must be separate: its particles cannot interact with the synchrotron photons of the observed SED peak, otherwise higher VHE fluxes than observed would be implied.  PKS\\,2005$-$489 is found overall in a very low state, in both the X-ray and VHE bands, during the observations presented in this article. PKS\\,2005$-$489 has historically demonstrated a 100$\\times$ dynamical range in the X-ray band. Thus, dramatically higher VHE fluxes ($10^2-10^4\\times$) can be expected in the future, unless such an increase in the X-ray flux is counter-balanced by a strong ($>$10$\\times$) and simultaneous increase in the blazar's magnetic field. Further monitoring of this object is highly encouraged, as it is one of the few HBL easily detected at VHE during low states and has the potential for extreme brightness and variability.  The results presented here confirm the strong diagnostic potential of coordinated optical--X-ray--VHE observations. Future studies can be significantly improved by incorporating data from the recently launched Fermi $\\gamma$-ray satellite.  The Fermi data will provide information on the lower energy side of the inverse-Compton peak, enabling contemporaneous measurements of both sides of each blazar hump."
    },
    "0911/0911.2223_arXiv.txt": {
        "abstract": "The Cold Spot on the Cosmic Microwave Background could arise due to a supervoid at low redshift through the integrated Sachs-Wolfe effect. We imaged the region with MegaCam on the Canada-France-Hawai'i Telescope and present galaxy counts in photometric redshift bins.  We rule out the existence of a 100Mpc radius spherical supervoid with underdensity $\\delta=-0.3$ at $0.5<z<0.9$ at high significance.  The data are consistent with an underdensity at low redshift, but the fluctuations are within the range of cosmic variance and the low density areas are not contiguous on the sky.  Thus, we find no strong evidence for a supervoid.  We cannot resolve voids smaller than 50Mpc radius; however, these can only make a minor contribution to the CMB temperature decrement. ",
        "introduction": "The nature of the Cold Spot on the Cosmic Microwave Background (CMB) has been the source of broad speculation.  The 10\\degr~diameter region identified in \\emph{Wilkinson Microwave Anisotropy Probe} (WMAP) temperature maps \\citep{WMAP1} is curious due to its pronounced temperature and morphology. Extensive study of the region was first motivated by the Spot's non-Gaussian properties \\citep{Vielva04,Cruz05}.  Under the standard cosmological model, the primordial fluctuations on the CMB are homogeneous, isotropic and described by a Gaussian random field.  In this scenario, the existence of the Spot is unlikely at the $\\sim$0.5\\% level \\citep{Cruz06}, although see \\citet{Zhang09} for a critical view.  The mean temperature decrement within a 5\\degr~radius is $\\Delta T=-100$\\muK, and it shows no systematic variation with frequency in WMAP data, making it inconsistent with contamination from synchrotron or dust emission \\citep{Cruz06,rudnick07}.  Many sources of secondary anisotropies have been proposed to explain the Cold Spot, including the integrated Sachs-Wolfe (ISW) and Sunyaev-Zeldovich (SZ) effect, as well as more exotic physics including textures arising in the early Universe.  Evidently there is no sufficiently massive cluster in the local universe to produce an SZ temperature decrement over such a large angular scale; however, the other hypotheses have not been directly addressed through observations.  \\citet{Cruz08} review these possibilities.  \\begin{figure}[th] \\begin{center} \\includegraphics{figs/fieldloc.pdf} \\end{center} \\caption{The 7 Cold Spot fields overlaying the WMAP linear combination (ILC) temperature map.  The dashed circle indicates the angular scale of 100Mpc radius at z=0.5. \\label{fig:fields}} \\end{figure}  In this work, we address the question of whether a supervoid exists along the line-of-sight.  A significant CMB decrement can be induced by a time-varying gravitational potential through the integrated Sachs-Wolfe effect (ISW) \\citep{sachswolfe}.  On linear scales, CMB photons traversing a large-scale void lose energy as the potential decays under the accelerated cosmological expansion.  Non-linear structure growth can also contribute to a cold imprint on the CMB through the Rees-Sciama effect \\citep{RS68}.  This scenario garnered interest due to the recent direct measurement of voids imprinted on the CMB \\citep{supervoids}.  The 10\\muK~signal was measured on 4\\degr~scales and has been speculated to arise from 100Mpc scale underdensities.  Could a similar structure be responsible for the Cold Spot?  Furthermore, \\citet{rudnick07} found a coincidental depression in source counts in the NRAO VLA Sky Survey (NVSS), although the significance of this alignment has been debated \\citep{Smith08}.  Through an independent methodology, \\citet{McEwen07} found that the Cold Spot region contributes strongly to the positive NVSS-WMAP cross correlation.  These findings motivate further investigation of the void hypothesis.  Survey data covering the Cold Spot are limited.  The ISW signal arising from the local Universe has been studied by \\citet{maturi07} who consider structures within 100Mpc.  They find no significant signal at the location of the Cold Spot.  \\citet{Francis09a} investigate the ISW signal traced by 2MASS at $z<0.3$.  They detect an underdensity at the Cold Spot, but find that it contributes to the temperature decrement by only 5\\%.  The Cold Spot is most prominent in a compensated filter. \\citet{Cruz06} use a Spherical Mexican Hat Wavelet with a scale radius of 4.17\\degr~and find a temperature decrement of $-16.09$\\muK~in this filter with a 1$\\sigma$ range due to cosmic variance of 3.55\\muK. Assuming that a void contributes only within the positive range of the wavelet, we need a $\\sim$10\\muK~contribution to reproduce the Spot on top of a $1-2\\sigma$ primary fluctuation on the CMB.  A 10\\muK~temperature decrement requires a $\\sim$100Mpc scale underdensity.  A spherical underdensity with $\\delta=-0.3$ must have a radius of 200Mpc to produce the effect \\citep{inouesilk07,Sakai08}. For the purposes of this work, we compute the expected ISW signal using the order-of-magnitude derivation presented by \\citet{rudnick07} which is in agreement with this result at low redshift.  To test the existence of a 100Mpc supervoid at $z<1$, we carried out an imaging survey of the Cold Spot region with MegaCam on the Canada-France-Hawai'i Telescope (CFHT).  Using galaxy counts, we estimate the large-scale density distribution in photometric redshift bins to $z=0.9$ and we compare our measurement to the expected distribution from mean counts across the sky.  We assume a standard WMAP5 \\LCDM~cosmology with $\\Omega_m=0.31$, $H_0=0.72~{\\rm km~s^{-1}~Mpc^{-1}}$ and $\\sigma_8=0.80$.   \\begin{table*}[ht] \\begin{center} \\caption{Survey fields\\label{table:fields}} \\begin{tabular}{|c|c|c|c|cccc|} \\tableline &       &           & Area      &\\multicolumn{4}{|c|}{95\\% Completeness (AB mag)}   \\\\ ID& R.A.  &  Dec.     & (sqr deg) & $g$& $r$& $i$& $z$     \\\\ \\tableline A & 03:05:26.5 & -23d15m21s &0.80 &22.7 & 21.6 & 22.4 & 20.9 \\\\ B & 03:05:45.3 & -19d48m07s &0.82 &22.8 & 22.8 & 22.5 & 21.0 \\\\ C & 03:08:45.8 & -17d01m35s &0.81 &23.2 & 23.4 & 22.4 & 20.8 \\\\ D & 03:14:40.1 & -20d45m16s &0.81 &22.6 & 23.2 & 22.4 & 20.5 \\\\ E & 03:18:53.2 & -20d30m25s &0.77 &22.4 & 23.4 & 22.4 & 21.1 \\\\ F & 03:21:20.0 & -18d15m23s &0.88 &23.9 & 23.3 & 22.1 & 21.5 \\\\ G & 03:28:45.3 & -20d29m45s &0.79 &23.1 & 23.0 & 22.1 & 21.1 \\\\ \\tableline \\end{tabular} \\end{center} \\end{table*} ",
        "conclusions": "We investigated the distribution of galaxies at $z<0.9$ in the direction of the CMB Cold Spot.  We detect an underdensity at low redshift, $0.1<z<0.3$, that is consistent with the density expected from a supervoid.  This underdensity appears to be present in 2MASS as well \\citep{Francis09a}.  However, due to the limited sky coverage of our survey, we cannot draw a definite conclusion regarding the existance of a coherent supervoid structure.  Our measurements disfavor a supervoid with $\\delta=-0.3$ at z$=0.3-0.5$, and we can rule out a supervoid at $z=0.5-0.9$.  We must be cautious in interpreting the density measurements due to possible systematic shifts in the selection function, especially in the near and far redshift bins.  In \\citet{supervoids,Granett09} we found the imprint of supervoids in SDSS to be stronger than predicted from linear ISW. This suggests that a more modest supervoid with $\\delta\\gtrsim -0.3$ could be the origin of the Cold Spot decrement.  Significant progress will be made in understanding the Cold Spot with wide-area large-scale structure surveys including Pan-STARRS-1 \\citep{panstarrs}.  Additionally, polarization data from the Planck mission \\citep{Planck} may provide additional information on the intrinsic CMB fluctuations in this area."
    },
    "0911/0911.2824_arXiv.txt": {
        "abstract": "Over the last decade observations at submillimetre (submm) and millimetre (mm) wavelengths, with their unique ability to trace molecular gas and dust, have attained a central role in our exploration of galaxies at all redshifts.  Due to the limited sensitivities and angular resolutions of current submm/mm telescopes, however, only the most luminous objects have been uncovered at high redshifts, with interferometric follow-up observations succeeding in resolving the dust and gas reservoirs in only a handful of cases.  The coming years will witness a drastic improvement in the current situation, thanks to the arrival of a new suite of powerful submm observatories (single-dish and interferometers) with an order of magnitude improvement in sensitivity and resolution. In this overview I outline a few of what I expect to be the major advances in the field of galaxy formation and evolution that these new ground-breaking facilities will facilitate. ",
        "introduction": "Key to understanding the origins of galaxies and their supermassive black holes (SMBHs) is that benchmark period of cosmic history called the epoch of reionization (EoR) in which the radiation fields from the first stars and accreting black holes reionized the surrounding gas -- an epoch of time that lasted from a couple of hundred million years to about one billion years after the Big Bang. While the first glimpses into the EoR have now been made owing to the discovery of a handful of quasars, Lyman-alpha emitters (LAEs) and gamma-ray bursts (GRBs) at $z\\gs 6$ (e.g.\\ Fan et al.\\ 2003; Kashikawa et al.\\ 2006; Kawai et al.\\ 2006), including the recently discovered GRB at $z\\sim 8.3$ (corresponding to about 625 million years after the Big Bang -- Tanvir et al.\\ 2009), it remains {\\it terra incognita} and a top priority of extragalactic astronomy in the years to come.  Another major challenge for extragalactic astronomy will be to make sense of the many different types of galaxies found at $z \\ls 6 $ in the context of galaxy evolution after the EoR. It is during this $\\sim 12.7\\,\\rm{Gyr}$ epoch -- i.e.\\ from the first one billion years to the present day -- that the bulk growth and evolution of galaxies and SMBHs took place. High-resolution radio/near-IR/submm observations have shown that accurate measurements of the sizes, morphologies, and kinematics of the gas and stars in high-$z$ galaxies are crucial steps in delineating their evolutionary paths and linking them to galaxy populations at later cosmic epoch. Yet our understanding of the 'evolutionary tree' of galaxies is still rudimentary, with even seemingly simple questions such as 'how do galaxies attain their gas?', and 'what is the co-evolution of SMBHs and galaxies?' lacking definite answers.  While answering these questions will require a 'pan-chromatic' effort, I shall here focus on the unique contributions we can expect from the next generation of large submm observatories such as a European 25m aperture submm telescope at Dome C, the Cornell Caltech Atacama Telescope (CCAT), the Atacama Large Millimeter Array (ALMA), and the Northern Extended Millimeter Array (NOEMA). The dramatic improvements in sensitivity, angular resolution, frequency coverage, and calibration capabilities of these facilities represent a quantum leap in our ability to map the gas and dust reservoirs in galaxies at all cosmic epochs, heralding in a new golden era in extragalactic astronomy. ",
        "conclusions": ""
    },
    "0911/0911.2015_arXiv.txt": {
        "abstract": "We describe a flasher system designed for use in monitoring the gains of the photomultiplier tubes used in the VERITAS gamma-ray telescopes. This system uses blue light-emitting diodes (LEDs) so it can be operated at much higher rates than a traditional laser-based system. Calibration information can be obtained with better statistical precision with reduced loss of observing time. The LEDs are also much less expensive than a laser. The design features of the new system are presented, along with measurements made with a prototype mounted on one of the VERITAS telescopes. ",
        "introduction": "This article concerns a light flasher system developed for the Very Energetic Radiation Imaging Telescope Array System (VERITAS)~\\cite{weekes02, holder06}. VERITAS is an array of four imaging atmospheric Cherenkov telescopes (IACTs) located at the Fred Lawrence Whipple Observatory on Mount Hopkins in southern Arizona. The VERITAS telescopes measure the Cherenkov light caused by relativistic electrons in air showers initiated by high energy astrophysical gamma rays. Each telescope consists of a 12-m-diameter reflector viewed by a `camera' comprising 499 29-mm-diameter photomultiplier tubes (PMTs). The PMTs are read out using FADCs operating at rate of 500 megasamples per second. A flasher system is used to monitor the gains and timing characteristics of the PMTs.  The present VERITAS flasher system~\\cite{hanna07} uses a nitrogen laser which produces short pulses of ultraviolet ($\\lambda$ = 330 nm) light that are distributed through quartz optical fibres. Each telescope is supplied by one fibre, which ends approximately four metres from the front of the camera in a termination box equipped with an opal diffuser. This system provides uniform and simultaneous illumination to all PMTs in the camera. It delivers light pulses which are similar in shape and wavelength to the Cherenkov-light pulses seen in normal observing and is therefore well-suited for testing the PMTs in their nominal region of operation.  This system has a few disadvantages. It cannot be pulsed at a rate much higher than 10 Hz which means that calibration runs can be long, especially when one is obtaining high-statistics data sets. Such runs are usually performed under dark-sky conditions and are therefore at the expense of scientific observing time. Another negative feature of the laser-based system is the cost and lifetime of the laser which uses a factory-sealed cartridge capable of producing a large but finite number of pulses. Replacing the cartridge costs several thousand dollars and can cause considerable down time if a spare is not on hand.  To overcome these disadvantages, we have developed a flasher system based on blue light-emitting diodes (LEDs) which are bright enough and fast enough to illuminate the VERITAS PMTs with pulses of duration and intensity that are very similar to those coming from the laser-based system and yet cost less than a dollar each. The LED flasher can be used to test the linearity of the PMT-FADC chain and can be used in conjunction with a neutral-density filter to determine the single-photoelectron response in all the channels. Moreover, because the system can be flashed at rates limited only by the VERITAS data acquisition system, runs which require very high statistics can be carried out in only a few minutes. ",
        "conclusions": "We have developed an LED-based flasher system that can be used for all of the tasks currently assigned to the VERITAS laser system. In addition it can be run at much higher rates so calibration runs can be carried out more quickly or can benefit from improved statistics for the same length of run.  The LED flasher is also much less expensive than a laser and the LED system is inherently a distributed one. A separate light source is installed on each telescope of the array. This could make it an attractive option for future projects such as AGIS~\\cite{agis} and CTA~\\cite{cta} where a large number of telescopes distributed over a large area would make a system based on a central laser sending light through optical fibres impractical."
    },
    "0911/0911.3155_arXiv.txt": {
        "abstract": "We use 62,185 quasars from the Sloan Digital Sky Survey DR5 sample and standard virial mass scaling laws based on the widths of H$\\beta$, Mg{\\small II}, and C{\\small IV} lines and adjacent continuum luminosities to explore the maximum mass of quasars as a function of redshift, which we find to be sharp and evolving.  This evolution is in the sense that high-mass black holes cease their luminous accretion at higher redshift than lower-mass black holes.  Further, turnoff for quasars at any given mass is more highly synchronized than would be expected given the dynamics of their host galaxies.  We investigate potential signatures of the quasar turnoff mechanism, including a dearth of high-mass quasars at low Eddington ratio.  These new results allow a closer examination of several common assumptions used in modeling quasar accretion and turnoff. ",
        "introduction": "\\label{sec:intro}  Supermassive black holes (SMBH) are found today at the centers of nearly every galaxy where there have been sensitive searches.  These SMBH are the remnants of bright quasars which are found in many galaxies at redshifts $z \\sim 2$.  The mechanisms involved in the quasar {\\em turnoff} process are not well understood (cf. Thacker et al. 2006\\nocite{turnoffreview}).  In this work, we combine the large quasar sample provided by the Sloan Digital Sky Survey (SDSS) \\cite{Schneider2007} with recent advances in virial mass estimation \\cite{Vestergaard2006,Shen2008} to develop new constraints on quasar turnoff.  It had been previously thought that most quasars were within a factor of $\\sim 10$ of their Eddington luminosity \\cite{Kollmeier2006}, and quasars have been commonly modeled as `light-bulbs', operating either near Eddington or without substantial luminous accretion, although this is likely an oversimplification \\cite{Hopkins2008}.  In Paper I \\cite{Steinhardt2009}, we used the quasar distribution in the mass-luminosity plane as a function of redshift.  Had the relationship between quasar mass and luminosity been trivial (as implied by the the light-bulb approximation) the mass-luminosity plane would not present an improvement over the quasar mass function or luminosity function.  Instead, we found a more complex relationship between mass and luminosity than had been previously assumed.  Several new puzzles have emerged from this view of the quasar distribution, including a sub-Eddington boundary that restricts luminosities of the highest-mass quasars at every redshift to lie below their Eddington limit.  We also noted an upper mass limit varying with redshift.  In this work, we use this new, multi-dimensional view of the quasar distribution to investigate this high-mass `turnoff' in detail.  We use these results to evaluate three basic assumptions: (1) that the dynamics of quasar turnoff are closely linked to the dynamics of the host galaxy, (2) that a segregation of quasars by mass is equivalent to a segregation of quasars by luminosity, (a corollary of the light-bulb approximation), and (3) that SMBH undergo no fundamental change during turnoff, but rather run out of fuel (which implies that an individual SMBH might have many luminous episodes during its lifetime and would undergo another one today given sufficient fuel).   In \\S~{\\ref{sec:revml}}, we briefly review the SDSS sample and exhibit its behavior in the $M-L$ plane.  A full discussion of this sample and its uncertainties can be found in Schneider et al. (2003) and Paper I\\nocite{Schneider2003,Steinhardt2009}.  In \\S~\\ref{sec:permanent}, we establish that massive quasars are indeed turning off, and doing so earlier than lower-mass quasars.  In \\S~\\ref{sec:signatures}, we use this to select individual quasars near turnoff and search for signatures of their turnoff mechanisms.  In \\S~\\ref{sec:synch}, we show that the synchronization of turnoff of different quasars at a given mass is closer than that of the dynamics of their host galaxies.  We also show that the synchronization of quasars at fixed mass is characteristically different from that of quasars at fixed luminosity as reported in Amarie et al. (2009)\\nocite{Amarie2009}.  In \\S~\\ref{sec:discussion}, we evaluate the three basic assumptions above in light of this new evidence. ",
        "conclusions": "\\label{sec:discussion}  Here and in Paper I \\nocite{Steinhardt2009}, we have considered what the quasar locus in the mass-luminosity plane can teach us about quasar accretion and turnoff.  Quasars have been modeled as very simple objects, in part because there has not yet been sufficient data to require more complex assumptions.   In Paper I, we falsified the assumption that quasars at all combination of mass and redshift were near their Eddington luminosity while accreting luminously.  With a combination of the large quasar catalog from the Sloan Digital Sky Survey (SDSS) and virial mass estimation, we can now evaluate several other common assumptions.  It is widely assumed that the dynamics of quasar turnoff are closely linked to the dynamics of the host galaxy.  However as shown in \\S~\\ref{sec:synch}, accretion rate and turnoff are highly synchronized amongst quasars of a common mass.  Either galaxies must be much more highly synchronized than currently believed or this synchronization must have an alternate source.  Galaxies continue to merge and virialize at later times than their quasars are predominantly turning off, and the decline in galaxy merger rates is believed to be much more gradual than that observed for quasars.  So, why are quasars so well synchronized?  One possibility is that the seeding mechanism comes not from the dynamics of their host galaxy but rather from some mechanism with a cosmic clock.  The $\\sim 700$ Myr timescale for quasar turnoff at $M_{BH} \\sim 10^{10} M_\\odot$ corresponds to the Hubble time at $z \\sim 8$, and extrapolating the slope of the e-folding time as a function of mass (\\S~\\ref{sec:permanent}) to the highest-mass quasars would yield the Hubble time at $z \\sim 20$.  It is plausible that the seeds of these massive quasars turn on during the epoch of reionization and, if the seeding mechanism is one that yields a nearly-fixed quasar life span, turnoff would be synchronized to roughly the observed precision.  This would require a seeding mechanism only weakly linked to the dynamics surrounding galaxy, but rather triggered by cosmic dynamics.  It is also widely assumed that most quasars are near Eddington while accreting luminously.  The presence of a sub-Eddington boundary at which high-mass quasars cannot reach their full Eddington luminosity directly contradicts this claim and slows the maximum allowed growth rate.  However the boundaries of the quasar locus are very sharp, and quasars of a given mass are highly synchronized in both the maximum luminosity they can attain and in their turnoff time.  In the absence of techniques for mass estimation of high-redshift quasars, it was widely assumed that a segregation of quasars by mass was equivalent to a segregation by luminosity.  Since individual quasars are variable, this cannot be entirely true, but there is certainly a correlation between quasar mass and luminosity as evidenced by the quasar distribution in the mass-luminosity plane (Figure \\ref{fig:allzcontour}).  Further, the high-mass, low-luminosity boundary suggests that even if high-mass quasars cannot reach $L_{Edd}$, luminous accretion is rare below a few percent of Eddington as well.  This also implies that the large Milliquas sample \\cite{Milliquas} of $\\sim 10^6$ new SDSS quasars selected photometrically but not bright enough for SDSS spectroscopy consists predominantly of quasars at lower $M_{BH}$ than those in the SDSS spectroscopic quasar catalog rather than a mixture of quasars at low $M_{BH}$ and quasars at higher $M_{BH}$ and very low Eddington ratio.  Finally, it is widely believed that a typical SMBH will have many periods of luminous accretion separated by dormant periods.  In part, this is because the Salpeter time for black hole accretion at $L_{Edd}$ is only $\\sim 10^8$ yr \\cite{Natarajan1998}, and a quasar could not grow at $L_{Edd}$ for several Gyr.  However, the sub-Eddington boundary restricts more massive quasars at every redshift to $0.2~L_{Edd}$ or below (with the peak number density below $0.1~L_{Edd}$), raising the possibility that rather than accreting luminously for less than 1 Gyr, quasars might accrete luminously, but sub-Eddington, in one long, continuous burst for several Gyr from when they first turn on until turnoff.  Certainly the sub-Eddington boundary requires that quasars accrete luminously far longer than previously believed.  When quasars of a given mass no longer reach a given Eddington ratio, the sub-Eddington boundary never moves back to a higher luminosity.  When quasars of a given mass begin to decline in number density, their turnoff is synchronized and sharp, and the population does not reappear at lower redshift.  While this is insufficient evidence to conclude that turnoff is a permanent process for any individual quasar, a quasar lifecycle consisting of one long period of accretion followed by permanent turnoff is consistent with all of these observations.  As shown in \\S~\\ref{sec:signatures}, the sharpness of the high-mass, low-luminosity boundary (HMLLB) appears to require either that virial mass estimation is wrong near the HMLLB or that quasars at the HMLLB undergo a rapid transition to an object with a very different spectral energy distribution.  One possibility would be transition to RIAF accretion \\cite{Narayan2008}.  However, the HMLLB does not occur at a constant Eddington ratio as accretion disk theory would suggest, but rather has a sub-linear dependence on mass.  The HMLLB is also evolving with redshift:  the HMLLB for high-mass quasars is at a higher redshift than lower-mass quasars.  In this paper, we have shown several new puzzles presented by the quasar mass-luminosity plane in addition to the sub-Eddington boundary discussed in Paper I.  These new puzzles which cannot be explained by selection effects are further evidence that supermassive black hole accretion and turnoff is more complex than previously believed.  Rather, many of the common assumptions made in quasar toy models are poor matches to the quasar locus in the $M-L$ plane.  These new boundaries indicate both that a more nuanced quasar model is required and that the quasar mass-luminosity plane now contains enough detail that it might be able to provide strong tests for such new models.  The authors would like to thank Mihail Amarie, Doug Finkbeiner, Lars Hernquist, John Huchra, and Michael Strauss for valuable comments.  This work was supported in part by Chandra grant number 607-8136A."
    },
    "0911/0911.4891_arXiv.txt": {
        "abstract": "Dynamical friction arises from the interaction of a perturber and the gravitational wake it excites in the ambient medium. We study the effects of the presence of a boundary on dynamical friction by studying analytically the interaction of  perturber with uniform rectilinear motion in a uniform homogeneous medium with a reflecting planar boundary. Wake reflection at a medium's boundary may occur at the edges of truncated disks perturbed by planetary or stellar companions as well as in numerical simulations of planet-disk interaction with no-outflow boundary conditions. In this paper, we show that the presence of the boundary modifies the behaviour of dynamical friction significantly. We find that perturbers are invariably pushed away from the boundary and reach a terminal subsonic velocity near Mach 0.37 regardless of  initial velocity. Dynamical friction may even be reversed for Mach numbers less than 0.37 thereby accelerating instead of decelerating the perturber. Perturbers moving  parallel to the boundary feel additional friction orthogonal to the direction of motion that is much stronger than the standard friction along the direction of motion. These results indicate that the common use of the standard Chandrasekhar formula as a short hand estimate of dynamical friction may be inadequate as observed in various numerical simulations. ",
        "introduction": "Dynamical friction, the reaction force a perturber feels from the gravitational wake it excites in the ambient medium,   is one of the fundamental processes in the formation of astronomical structures. Applications of dynamical friction to collisionless systems include mass segregation in star clusters \\citep{r1,r2}, galaxy distribution in hierarchical cluster formation \\citep{r3,r4}, the decay of satellites in galactic halos and galaxy mergers \\citep{r5} (and references therein).  Dynamical friction also occurs in collisional particle systems and was applied to the study of the growth of planetesimals \\citep{r6,r7}, the eccentricity excitation of planetary embryo orbits \\citep{r8,r9,r10}, and the confinement of planetary rings \\citep{r11}. For gaseous systems, applications include the orbital decay of compact stars orbiting supermassive black holes \\citep{r12}, the heating of intercluster gas by decelerating supersonic galaxies \\citep{r13,r14,r15}, the characterization of the differences in X-ray emissions in galaxy clusters between the standard cold dark matter model and modified newtonian dynamics \\citep{r16}.  Boundaries appear in dynamical friction because both the perturber and the ambient medium that surrounds it have finite sizes. In the common analyses of dynamical friction, these two boundaries enter the force through the Coulomb logarithm that requires cutoffs at small and large distances from the perturber to avoid divergence. In the context of collisionless systems,  the classical derivation of dynamical friction using a perturber in uniform rectilinear motion \\citep{r17} leaves the choice of the two cutoffs somewhat indeterminate and gives the force as: \\begin{equation} {\\bm F}=-16\\pi^2(GM)^2m\\log \\left(\\frac{\\rmx}{r_{\\rm min}} \\right)\\left[\\int_0^V{\\rm d}u\\ u^2f(u)\\right]\\frac{\\bm V}{V^3}, \\label{chandra} \\end{equation} where $G$ is the gravitational constant, $M$ is the perturber's mass, $m$ is the individual mass of a background medium star,  $V$ is the perturber's velocity, $f$  is the medium's velocity distribution function, $r_{\\rm min}$ and $\\rmx$  are the small and large distance cutoffs respectively. The small distance cutoff, $r_{\\rm min}$,  is believed to be the larger of the perturber's radius and the distance related to the typical velocity in the ambient medium whereas the larger cutoff radius,  $\\rmx$ ,  is usually taken as the size of the system (e.g. see \\citep{r5}). A more accurate analysis of the gravitational scattering of a star in rectilinear motion in a collisionless infinite system \\citep{r18}, shows that there is a natural large distance cutoff that equals the distance travelled by the perturber, $Vt$, where $t$ denotes the time the perturber takes crossing the medium. This cutoff arises for durations smaller than the crossing time because the far stars do not have enough time to completely deflect the perturber by the standard (asymptotic) angle that depends only on the distance to the perturber and its velocity but not on time. The farther the star, the more incomplete the deflection, leading to a scattering angle that depends on time. Equivalently, it can be shown that only stars with impact parameters smaller than $Vt$  substantially affect the perturber's velocity. This analysis was applied to accurately estimate the relaxation time in an infinite dilute star system \\citep{r19}. When steady state is considered, the distance travelled by the perturber becomes comparable to the system's size and the upper cutoff, $\\rmx$, is taken as the system's size. Similar prescriptions are applied to circular motion in spherically symmetric systems where it is shown that only the vicinity of resonances within the system contribute to dynamical friction \\citep{r20} including the standard long and short distance cutoffs.  For a perturber moving in a gaseous medium, the small distance cutoff is taken as the perturber's accretion radius or physical size.  The large distance cutoff depends on the perturbation regime.  In steady state,  it is taken as the system's size \\citep{r21,r22,r23} whereas in the time-dependent regime, in which the perturbation's duration is much smaller than the crossing time, it is given by the distance travelled by the sonic shockwave that the perturber excites in the medium \\citep{r24,r25,r26,r27,r38} and \\citep{r29} hereafter Paper I. Although the nature of the large distant cutoffs used to truncate the density perturbation is of different physical origin for collisionless and collisional media, the corresponding friction forces are similar \\citep{r26}.  In both collisionless and collisional media, dynamical friction was determined in the two limiting cases of steady state and the time-dependent regime. For the latter, the independence of dynamical friction of the medium's size is straightforward as the perturbed region that acts back on the perturber is much smaller than the ambient medium. The situation is much less clear for steady state in which the density enhancement raised by the perturber has reached the medium's boundary. Suppose for instance that a perturber is set in motion near the boundary of a gaseous medium such that the distance the sonic shockwave has to travel to reach the boundary is much smaller than the medium's size. Before the shockwave reaches the boundary, dynamical friction is in the time-dependent regime. However as the boundary is reached, the medium's full size cannot serve as a large distance cutoff. Instead, the proximity of the perturbed region to the medium's boundary will affect the way the density enhancement acts back on the perturber depending on how the boundary responds to an incoming shockwave. It is the object of this paper to find out how dynamical friction may be changed as the perturber's shockwave reaches the medium's boundary.  As such a task is quite difficult, we will make a number of simplifying assumptions aimed at illustrating that the classical behavior of dynamical friction may be changed drastically because of the presence of a boundary.  We assume that the perturber moves in a homogeneous gaseous medium along a rectilinear trajectory with uniform velocity near the medium's boundary that is taken to be planar and reflecting. The boundary's curvature is neglected as the medium's size is supposed to be large compared to the perturber's distance to the boundary. Wake reflection at a medium's boundary may occur at the edges of truncated disks perturbed by planetary or stellar companions \\citep{r30,r31,r32,r42} as well as in numerical simulations of planet-disk interaction with no-outflow boundary conditions \\citep{r33,r34,r35}. Interestingly, we shall show that most of the significant changes on dynamical friction come not from the reflected component but from the truncation of the shockwave incident on the boundary. To derive the friction force for this set-up, we shall use the general expression of dynamical friction derived in Paper I that allowed us to study the influence of a perturber's acceleration on dynamical friction.  The paper is organized as follows: in section 2, we recall the general expression of dynamical friction derived in Paper I and characterize how the medium's size affects dynamical friction in general.  In section 3, we derive the dynamical friction force for a perturber moving normal to the medium's boundary and show that perturbers are pushed away from the boundary and reach a finite terminal velocity.  In section 4, we examine dynamical friction for motion parallel to the boundary and show the appearance of an additional force component normal the trajectory  that pushes the perturber away from the boundary.  Section 5 includes a summary of main results and possible further refinements.  In Appendix A, we characterize the gas flow forced by the perturbers  in the configurations studied in  sections 3 and 4  by deriving the corresponding velocity field and density perturbation and linking them to the known results of the Bondi-Hoyle-Lyttleton accretion. ",
        "conclusions": "We set out to examine the effect of the presence of a boundary on dynamical friction. We showed that the expressions of the force in  infinite (\\ref{ostriker}) and finite  (\\ref{steady1})  media may not be used in the presence of a boundary such as the medium's outer and inner limits or any other internal edge.  In particular, dynamical friction along the direction of motion may be reversed and a new drag component perpendicular to the boundary  is present and forces the perturber to move away from the medium's edge.  These results are of particular relevance for the numerical simulation of dynamical friction once the perturber's wake reaches the boundaries of the physical system or the simulation box as for instance  in problems related to disk-planet interactions and  galaxy merging where timescales are observed to differ significantly from the prediction of the Chandrasekhar formula \\citep{r43}.  We have considered separately motion normal and parallel to the boundary to isolate the effects of the instability of zero-velocity perturbers, the reversal of dynamical friction, and  normal dynamical friction. Whereas qualitatively accurate for a perturber on a trajectory inclined by an arbitrary angle with respect to the boundary, these effects may not be added quantitatively linearly because dynamical friction depends on the phase vectors ${\\bm \\Delta}$ and ${\\bm \\Sigma}$. This is why it is not straightforward to carry over the first two effects to rotational motion in bounded systems (such as disk-planet interactions) as they would require the perturber's orbit to have some eccentricity and hence a velocity that makes a time-dependent angle with the boundary. Normal dynamical friction however readily applies to circular motion in a spherical system. The dependence of the normal force on the distance to the boundary is translated for a  perturber in circular motion around an attracting center as a radial force that depends on the orbital radius and the location of  the medium's edge for subsonic perturbers, and additionally on the velocity for supersonic perturbers. This time-dependent force affects the perturber's energy but not its angular momentum and is most efficient where motion is resonant with the sound speed.  In our derivation of dynamical friction, we neglected the boundary's curvature. It is straightforward that for a perturber on a normal trajectory that moves away from a convex boundary, the friction force will be further reduced with respect to  that given by (\\ref{incidentsub},\\ref{incidentsuper}) as the wake behind the perturber is truncated further (as in Figure 1).  The current analysis may also serve as a guide to examine the effect of the ambient medium's boundary in collisionless systems. As mentioned earlier, the derivation of dynamical friction from the velocity impulse caused by a test star on the perturber is centered physically around the latter \\citep{r5}. When steady state is considered, the large distance cutoff is applied in the perturber-centered frame which is equivalent to assuming a comoving symmetric boundary.  Although they differ in magnitude, the  friction forces in a gaseous medium and a collisionless star system share a strong similarity in the way they depend on the Mach number \\citep{r26}. Such similarity suggests that the conclusions we found  here will  hold for collisionless systems."
    },
    "0911/0911.4858_arXiv.txt": {
        "abstract": "The dark energy problem has led to speculation that not only may $\\Lambda$CDM be wrong, but that the FLRW models themselves may not even provide the correct family of background models. We discuss how direct measurements of $H(z)$ can be used to formulate tests of the standard paradigm in cosmology. On their own, such measurements can be used to test for deviations from flat $\\Lambda$CDM. When combined with supernovae distances, Hubble rate measurements provide a test of the Copernican principle and the homogeneity assumption of the standard model, which is independent of dark energy or metric based theory of gravity. A modification of this test also provides a model independent observable for flatness which decorrelates curvature determination from dark energy. We investigate these tests using Hubble rate measurements from age data, as well as from a Hubble rate inferred from recent measurements of the baryon acoustic oscillations. While the current data is too weak to say anything significant, these tests are exciting prospects for the future. ",
        "introduction": "There is a growing realisation in cosmology that the dark energy problem means that we have to test the fundamentals of our cosmological model as rigorously as we can. Despite the fact that the standard flat $\\Lambda$CDM paradigm works so well on many fronts, the unknowns at the heart of it have led us into the territory of trying to invert observations to ascertain directly any temporal evolution of dark energy~-- whether it is in fact modified gravity, global inhomogeneity such as a void model, quintessence and so on. Normally in physics, we would hope to work the other way round: a model based on known physics can predict outcomes of an experiment which can then potentially falsify the model (or theory). Unfortunately, once we move away from the comfort of $\\Lambda$CDM we have little physics to guide us at all; worse, many suggestions appear degenerate or even unfalsifiable. Generalised gravity theories, quintessence models, void models, etc. all have free \\emph{functional} degrees of freedom, and not just free parameters.  So dark energy reconstruction tries to deduce functional degrees of freedom directly from data sets we have.  In particular, if Hubble-scale inhomogeneity is responsible for the apparent acceleration, then all bets are off as to what the effective equation of state\\footnote{That is, the equation of state which reproduces some particular family of observations, such as the distance modulus, if we utilise a FLRW model.} could be. In the simplest void models which describe this~\\cite{voids,FLSC}, a violation of the Copernican principle~-- that we are not located at a special place in the universe~-- is required, because they place us very near the centre of symmetry. Clearly such a radical suggestion warrants much more study; for example, it is not yet understood how perturbations evolve, although the generalised Bardeen equation is known~\\cite{CCF}.   One important aspect of the reconstruction approach will rely on formulating consistency tests of the standard paradigm. Ideally these tests would be constructed so that signals of deviations from (flat) $\\Lambda$CDM are easy to spot and quantify. If it's possible, we would like these tests to be independent of the standard paradigm, and rely as little as possible on assumptions that we normally take for granted. For example, consistency between the background evolution and the evolution of perturbations can be used to rule out whole classes of quintessence models~\\cite{MHH}. Using perturbations in this way is only possible within given classes of dark energy, however, because of assumptions about the sound speed which must be made, amongst other things. Tests which may be formulated purely from background observables would not suffer this difficulty.    Here our principle aim is to investigate the curvature test, introduced in~\\cite{CBL}. This is a test which uses background observables only, and tests for consistency of the FLRW paradigm itself. In FLRW models, the luminosity distance may be written as \\begin{equation}\\label{d_L} d_{L}(z)=\\frac{c(1+z)}{H_0 \\sqrt{-\\Omega_k}}\\sin{\\left( \\sqrt{-\\Omega_k}\\int_0^z{\\mathrm{d}z'\\frac{H_0}{H(z')}}\\right)}, \\end{equation} where $\\Omega_k$ is the curvature parameter today, and the expansion rate $H(z)$ takes the value $H_0=H(0)$ today. (Note that this is actually valid for either sign of curvature; for $\\Omega_k=0$ we can take the limit $\\Omega_k\\to0$.) The area distance is defined using $d_L=(1+z)^2d_A$, and another distance measure we will use is $D=(1+z)d_A/(c/H_0)$, which is the dimensionless comoving distance. We may rearrange Eq.~(\\ref{d_L}) to give an expression for the curvature parameter in terms of $H(z)$ and $D(z)$~\\cite{CBL}: \\begin{equation}\\label{OK} \\Omega_k=\\frac{\\left[h(z)D'(z)\\right]^2-1}{[D(z)]^2}=\\mathscr{O}_k(z), \\end{equation} where $h(z)=H(z)/H_0$ and $'=\\d/\\d z$. This tells us how to measure the value of the curvature parameter today from distance and Hubble rate observations, independently of any other model parameters or dark energy model or field equations. Remarkably this tells us the curvature today from these measurements at any single redshift, provided that $D(z)$ and $H(z)$ come from unrelated families of observations. If the underlying geometry of the background cosmology is FLRW, $\\mathscr{O}_k(z)$ will be measured as constant and so be independent of the redshift of measurement. As a consistency test of the standard paradigm it has all the right ingredients. It is a straightforward comparison of different observables, both of which can be formulated fairly independently of contamination from the standard model.  At the simplest level, then, we can use this to see if our distance measurements are consistent with our $H(z)$ measurements, if we choose to stick within the standard paradigm. But we can get more information than that.  For measurements of $\\mathscr{O}_k(z)$  not to be constant as a function of redshift we are looking at pretty radical changes to the standard model~-- an alteration to the underlying homogeneous and isotropic FLRW models themselves. One important situation where this will happen is in void or Hubble-bubble models. These models fit the SNIa distance modulus by having us near the centre of a large Hubble-scale depression in the energy density and accompanying curvature.  These models violate the Copernican principle which lies at the heart of cosmology. If isotropy about us were determined at all redshifts, this curvature test may then be used as a direct test of the Copernican principle.  However, deviations from $\\mathscr{O}_k(z)=$const. will be signalled in any non-FLRW cosmology; in particular, in Swiss-cheese models which satisfy the Copernican principle on very large scales will have an $\\mathscr{O}_k(z)$ which oscillates with redshift. Furthermore, backreaction and averaging will induce non-constant $\\mathscr{O}_k(z)$~\\cite{wilt}.  (Note that in all these models it won't necessarily be saying anything specific about the curvature.)  A similar background test which is specifically to test any deviations the flat $\\Lambda$CDM scenario was introduced independently in~\\cite{om,litmus}. The $Om$ diagnostic defined by~\\cite{om}: \\be \\label{eq:om} Om(z)=\\frac{h^2(z)-1}{(1+z)^3-1}=\\frac{1-D'(z)^2}{[(1+z)^3-1]D'(z)^2}, \\ee is a constant at different redshifts in a spatially flat $\\Lambda$CDM model, equal to today's value of the matter density parameter $\\Omega_m$. In a similar spirit to $\\Omega_k$ we can measure the variables on the rhs, and expect our measurements to yield the same number regardless of redshift. Fig.~\\ref{fig:om} shows the kind of behaviour $Om(z)$ exhibits from different FLRW models. \\begin{figure}[htbp] \\begin{center} \\includegraphics[width=0.95\\columnwidth]{om.jpg} \\caption{A plot of $Om(z)$ in different models, obtained using distances. The figure shows the range of behaviour from curvature in each fan of curves (key, top left). The three grey fans show how differing $\\Omega_m$ values interact with curvature for $\\Lambda$CDM (key, top right). The effect of changing $w$ by a constant is illustrated in the red and brown fans. Similar curves are found if we use $h$ rather than $D$ to form $Om(z)$, but without the cross-over at high redshift. It is clear that, although $Om(z)$ deviates from a constant because of curvature this deviation is not strong for small curvature values; therefore, $Om(z)$'s utility lies in picking up non-$\\Lambda$ behaviour.} \\label{fig:om} \\end{center} \\end{figure} From this we can also derive~\\cite{litmus}: \\ba \\label{eq:L} \\mathscr{L}(z)\\!\\!&=&\\!\\!2[(1+z)^3-1] D''(z)\\nonumber\\\\&& +3(1+z)^2D'(z)[1-D'(z)^2]\\nonumber\\\\ &=&\\!\\!0~\\mathrm{for~all~flat}~\\Lambda\\mathrm{CDM~models}. \\ea Testing for deviations from $\\mathscr{L}(z)=0$ provides an alternate test of flat $\\Lambda$CDM.  The main issue in using the curvature test lies in observing $H(z)$ independently of distance measurements. For example, $H(z)$ is often reconstructed from  supernovae data~\\cite{varun_alexei_rev06}, but this implicitly uses Eq.~(\\ref{OK}) and so can't be used here. There are other methods, two of which we use here. Measurements of the Baryon Acoustic Oscillations give the `volume distance' which is~ \\be\\label{DV} D_V(z)^3=\\left(\\frac{c}{H_0}\\right)^3\\frac{z D(z)^2}{h(z)}, \\ee from which we can derive $H(z)$ if we have an another method to determine $D(z)$. Alternatively, we can reformulate the curvature test directly in terms of $D_V(z)$ and $D(z)$. This method is not entirely independent of cosmological model, however, because it assumes that there is no scale dependence in the evolution of perturbations~-- this is not the case in inhomogeneous models where the curvature is varying spatially~\\cite{CCF}.  Spectroscopic age dating of passively evolving galaxies give another measure of $H(z)$~\\cite{jimenez_verde}, which we use here. This method measures \\be \\frac{\\d t}{\\d z}=\\frac{1}{(1+z)H(z)}. \\ee Other possible methods include measurements of the time drift of redshifts~\\cite{zdot} and the dipole in the SN observations~\\cite{ruth}, but these are not considered further here. For distance measurements we shall use the latest complete SNIa data set~\\cite{Hicken09}. By using the smoothing method presented in~\\cite{Shafieloo05,Shafieloo07,Shafieloo09b} we can turn this data into $D(z)$ and $D'(z)$ without relying on any fiducial cosmological model or parameterisation. We use $\\mathscr{O}_k(z) = \\Omega_k$ in Eq.~(\\ref{OK}) and $Om(z)$  in Eq.~(\\ref{eq:om}) as  two diagnostics which we attempt to reconstruct using these independent datasets. Any deviations from constant will signify problems with the FLRW models themselves in the case of $\\mathscr{O}_k(z)$, and problems with flat $\\Lambda$CDM in the case of $Om(z)$. Finally we make use of a third test: $\\mathscr{O}_k(z)D(z)^2$ may be used as a model independent test for flatness if we look for deviations from zero. ",
        "conclusions": "We have investigated some model independent consistency tests of the FLRW models using a mixture of the latest complete SNIa data and some Hubble rate data. The SNIa data is turned into the dimensionless distance function $D(z)$ using a model independent smoothing method.   One source of the $H(z)$ data is from the BAO measurements of $D_V(z)$, which relies on assumptions about how perturbations evolve in the standard FLRW paradigm. The other is from the relative ages of passively evolving galaxies, which has far fewer assumptions about the underlying cosmological model built in to it.  We have considered three tests: \\begin{itemize} \\item Tests for flat $\\Lambda$CDM: $Om(z)$ and $\\mathscr{L}(z)$. The first of these was investigated using both Hubble rate data and distance data independently. When using Hubble rate data, everything is consistent with flat $\\Lambda$CDM. Using distance data, there seems to be slight evidence for deviation from flat $\\Lambda$CDM. We should note that a larger $Om(z)$ at low red shifts, as we see in Fig.~\\ref{fig:OM}, is indicative of a larger $w(z)$ for dark energy (considering the effect to be from dark energy, and not due to curvature) and suggest that some decaying models of dark energy (to dark matter or something else) must have a good fit to the data. This is in agreement with \\cite{Shafieloo09a}. However larger reconstructed $Om(z)$ at high redshifts (despite of the fact that the quality of the data is poor in this range) seems to be related to the apparent extra brightness of the supernovae data (in comparison with $\\Lambda$CDM model) at $z > 1$ as reported and studied in~\\cite{arman_leandros}. The second test $\\mathscr{L}(z)$ relies on second derivatives of distance data, but is surprisingly useful even with present data; here, deviations from flat $\\Lambda$CDM are seen with mild significance. It is interesting to note the close resemblance of the curves we have reconstructed here in a model independent way, to those same functions in void models~\\cite{FLSC}.  \\item Test for the sign of curvature: $\\mathscr{O}_k(z)D(z)^2$ is a test for flatness, which is independent of dark energy. It requires two independent data sets, and seems most useful at the moment for $z\\lesssim1$. For $z\\gtrsim1$ the errors on SNIa distances dominate. There is no evidence for deviation from flatness, though this is not surprising because it relies on Hubble rate data.  \\item Copernican test for homogeneity: $\\mathscr{O}_k(z)$ is a model independent test for the FLRW models themselves, which is independent of dark energy or theory of gravity. It requires two independent data sets, and is very promising for $z\\gtrsim1$. Again, because it relies on Hubble rate data, there is no evidence for deviations from FLRW, nor indeed from flatness. Note that for $0.8\\lesssim z\\lesssim1.6$ the data shows a mild trend towards closed models. This is likely of no significance. \\end{itemize}    {It is early days for tests as sensitive as these because they require non-parametric construction, and we are still dealing with only a few hundred data points. We can easily envisage reaching thousands of data points in the next few years at which point these tests will come into their own as foundational tests for the concordance model. If the concordance model is correct then these tests will provide a useful way to quantify how confident we are in a more absolute way than we can at present. If, on the other hand, we find deviations for any of these tests, then they will help point us in other directions. In particular, if the Copernican tests $\\mathscr{O}_k(z)$ and $\\mathscr{C}(z)$ fail then we have a serious problem for cosmology, and the standard FLRW models. But with that would come a strong indication that dark energy cannot be due to modified field equations on either side; rather, we would know that there is something wrong with the spacetime metric we are using or that we are not describing the universe in the mathematically correct way.}"
    },
    "0911/0911.4405_arXiv.txt": {
        "abstract": "We study the well pronounced Moreton wave that occurred in association with the X17.2 flare/CME event of October 28, 2003. This Moreton wave is striking for its global propagation and two separate wave centers, which implies that two waves were launched simultaneously. The mean velocity of the Moreton wave, tracked within different sectors of propagation direction, lies in the range of $v\\approx900-1100$~km~s$^{-1}$ with two sectors showing wave deceleration. The perturbation profile analysis of the wave indicates amplitude growth followed by amplitude weakening and broadening of the perturbation profile, which is consistent with a disturbance first driven and then evolving into a freely propagating wave. The EIT wavefront is found to lie on the same kinematical curve as the Moreton wavefronts indicating that both are different signatures of the same physical process. Bipolar coronal dimmings are observed on the same opposite East-West edges of the active region as the Moreton wave ignition centers. The radio type II source, which is co-spatially located with the first wave front, indicates that the wave was launched from an extended source region ($\\gtrsim 60$~Mm). These findings suggest that the Moreton wave is initiated by the CME expanding flanks. ",
        "introduction": "Large-scale, large-amplitude disturbances propagating in the solar atmosphere were first recorded by \\cite{Moreton:1960b} and \\cite{Moreton:1960} in the chromospheric H${\\alpha}$ spectral line; therefore called ``Moreton waves''. Typical velocities are in the range $\\approx500-1000$~km~s$^{-1}$, and the angular extents are $\\approx90-130\\degr$ \\citep[e.g.][]{Warmuth:2004,Veronig:2006}. Since there is no chromospheric wave mode which can propagate at such high speeds, Moreton waves were interpreted as the intersection line of an expanding, coronal fast-mode shock wave and the chromosphere which is compressed and pushed downward by the increased pressure behind the coronal shockfront \\citep{Uchida:1968}. They occur in association with major flare/CME events and type II bursts \\citep{Smith:1971}, the latter being a direct signature of coronal shock waves. Typically, the first wavefront appears at distances of $\\ge$100~Mm from the source site and shows a circular curvature. The fronts are seen in emission in the center and in the blue wing of the H$\\alpha$ spectral line, whereas in the red wing they appear in absorption. This behavior was interpreted as a downward motion of the chromospheric plasma with typical velocities of $\\approx 5-10$~km~s$^{-1}$ \\citep{Svetska:1976}. Sometimes the trailing segment of the wave shows the upward relaxation of the material, i.e., the chromosphere executes a down-up swing \\citep{Warmuth:2004b}. In the beginning of the wave propagation, the leading edge is rather sharp and intense. As the disturbance propagates, the perturbation becomes more irregular and diffuse and its profile broadens \\citep{Warmuth:2001, Warmuth:2004}. Thus, the perturbation amplitude decreases and the wavefronts get fainter until they can no longer be traced at $\\approx300-500$~Mm from the source active region \\citep{Smith:1971, Warmuth:2004}.  A decade ago, large-scale waves were for the first time directly imaged in the corona by EIT \\citep[Extreme-Ultraviolet Imaging Telescope;][]{Delaboudiniere:1995} aboard the SoHO (Solar and Heliospheric Observatory) spacecraft, so-called ``EIT waves'' \\citep{Thompson:1997}. Similarities in the propagation characteristics led to the assumption that in at least a fraction of the events are the EIT waves the coronal counterpart of the chromospheric Moreton waves \\citep[e.g.][]{Thompson:1997, Warmuth:2001, Vrsnak:2002, Vrsnak:2006, Veronig:2006}. Basic questions regarding the nature of Moreton and EIT waves are whether they are caused by the same or by different disturbances, and whether they are initiated by the associated flare or the CME. For recent reviews, we refer to \\cite{Chen:2005}, \\cite{Mann:2005}, \\cite{Warmuth:2007} and \\cite{Vrsnak:2008}.  Here, we study the fast and globally propagating Moreton wave that occurred in association with the powerful X17.2/4B flare and fast CME event from the NOAA AR10486 (S16$\\degr$, E08$\\degr$) on October 28, 2003. Due to its extreme powerfulness and geo-effectiveness, diverse aspects of this flare/CME event have been analyzed in a number of studies \\citep[e.g.][]{Gopalswamy:2005, Klassen:2005, Aurass:2006, Kiener:2006, Hurford:2006, Mandrini:2007}. The relationship of the Moreton wave to radio observations has already been studied in \\cite{Pick:2005}. In this paper, we focus on the kinematical analysis of the Moreton wave, its relationship to the flare, the CME, coronal dimmings and type II radio bursts, in order to get insight into the wave characteristics and its initiating agent. ",
        "conclusions": "We analyzed the intense Morton wave of October 28, 2003 together with its associated phenomena: the X17.2/4B flare event, the fast halo CME ($v\\approx$~2500~km~s$^{-1}$), the EIT wave, the radio type II burst, and coronal dimmings, in order to study the wave characteristics and its driver. The main findings of the study and our interpretation are summarized in the following:  \\begin{enumerate} \\item The Moreton wave propagates in almost all directions over the whole solar disk, and can be observed for a period as long as 11~min (during 11:02--11:13~UT). These are both unusual properties of a Moreton wave. But we note that a wave with similar characteristics (global propagation, $v\\approx$~1100$-$1200~km~s$^{-1}$) was launched from the same AR on 2003 October 29 in assocation with an X10 flare/CME event \\citep[][]{Balasubramaniam:2007}. We studied the wave kinematics in detail for four different propagation directions finding mean velocities in the range 900--1100~km~s$^{-1}$. Typical values  of the sound speed and the Alfv$\\acute{\\hbox{e}}$n speed in the solar corona are 180~km~s$^{-1}$ and 300--500~km~s$^{-1}$ (Mann et al.\\ 1999), respectively, which corresponds to a magnetosonic Mach number of $\\approx$2--3 of the Moreton wave under study, if interpreted as a propagating fast-mode shock. The band-split of the associated type II burst was determined to 0.15--0.25, from which we inferred magnetosonic Mach number $M_s=1.23-1.45$. \\item We find two radiant points (source centers) for the Moreton wave on opposite East-West edges of the source active region (AR~10486). This has not been reported before and indicates that either two independent waves were launched simultaneously, or that the wave was caused by the expansion of an extended source with a preferred axis. \\item The associated coronal wave (``EIT wave'') was observed by the EIT instrument in one frame. The determined EIT wavefront lies on the extrapolated kinematical curve of the Moreton wave, identifying both the EIT wave and the H$\\alpha$ Moreton wave as different signatures of the same disturbance, according to Uchida's theory. \\item In two propagation directions, a deceleration of $-2960$~m~s$^{-2}$ and $-1420$~m~s$^{-2}$ was observed in the distance-time diagrams of the H$\\alpha$ Moreton wavefronts. In the initial phase, the perturbation profiles show amplitude growth and steepening of the disturbance. Later on, they display a broadening, amplitude decrease, and signatures of chromospheric relaxation appear at the trailing segment of the disturbance. This indicates that the coronal perturbation is a freely propagating large-amplitude simple-wave. The width of the disturbance is in the range of 50--150~Mm. \\item The extrapolated start time of the wave fits roughly with the extrapolated start time of the CME observed in LASCO (but we note that the uncertainty of the CME start time is of the order of 10~min) but occurs slightly before (2~min) the first HXR peak of the flare impulsive phase (cf.\\ Fig.~\\ref{img:plot_HXR}). \\item The associated bipolar coronal dimmings, which are generally interpreted as footprints of the expanding CME and associated coronal mass depletion, appear on the same opposite East-West edges of AR~10486 as the Moreton wave ignition centers do. The bipolar dimmings show a fast increase during the impulsive phase, followed by a much slower evolution, which indicates a characteristic restricted growth. The expansion velocity of the dimming region is $\\approx$100~km~s$^{-1}$ at the time, where the Moreton wave is launched. \\item The position of the type II burst source is co-spatial with the Moreton wavefront in direction 4$-$6. The back-extrapolated start time for the type II burst is $\\approx$11:02:30~UT, i.e., somewhat later than the start time derived for the Moreton wave with respect to its radiant point. This suggests that the perturbation was not launched from a point-like source but from an extended source whose surface was at a distance $\\gtrsim$60~Mm from the back-extrapolated radiant point. The launch time of the type II burst, which indicates the shock front formation, and the Moreton wave perturbation profile steepening coincide within the measurement errors. \\end{enumerate}   From several case studies \\citep[e.g.][]{vourlidas03,Mancuso:2003,yan06,Cho:2008,ontiveros09} it was derived that radio as well as white light signatures in coronagraphs provide evidence that the lateral expansion of a CME is able to produce shocks in the corona. The results of the study, presented in this paper, support a scenario in which the launch of the observed Moreton and EIT wave is initiated by the lateral expansion of the CME. In such a case we expect that the perturbation is driven only over a short time/distance, as the CME lateral-expansion range is limited \\citep[for the theoretical background see, e.g.][]{zic08,temmer09}. Thus, the disturbance would evolve into a freely propagating wave, which is consistent with the observed perturbation characteristics. The behavior of CME expansion (radial and lateral direction) at low coronal heights might be an important condition for the formation of coronal shock waves."
    },
    "0911/0911.4319_arXiv.txt": {
        "abstract": "We present ultraviolet integrated and azimuthally-averaged surface photometric properties of a sample of 44 dIm, BCD, and Sm galaxies measured from archival NUV and FUV images obtained with {\\it GALEX}. We compare the UV to \\ha\\ and $V$-band properties and convert FUV, \\ha, and $V$-band luminosities into star formation rates (SFRs). We also model the star formation history from colors and compare the integrated SFRs and SFR profiles with radius for these methods. In most galaxies, the UV photometry extends beyond \\ha\\ in radius, providing a better measure of the star formation activity in the outer disks. The H$\\alpha$ appears to be lacking in the outer disk because of faintness in low density gas. The FUV and $V$-band profiles are continuous with radius, although they sometimes have a kink from a double exponential disk. There is no obvious difference in star formation properties between the inner and outer disks. No disk edges have been observed, even to stellar surface densities as low as $0.1\\;M_\\odot$ pc$^{-2}$ and star formation rates as low as $10^{-4}\\;M_\\odot$ yr$^{-1}$ kpc$^{-2}$. Galaxies with low HI to luminosity ratios have relatively low FUV compared to $V$-band emission in the outer parts, suggesting a cessation of star formation there. Galaxies with relatively high HI apparently have fluctuating star formation with a Gyr timescale. ",
        "introduction": "\\label{sec-intro}  Outer edges of dwarf irregular (dIm) galaxies present an extreme environment for star formation. Dwarf galaxies already challenge models of star formation because of their low gas densities even in the central regions (Hunter \\& Plummer 1996, Meurer \\et\\ 1996, van Zee \\et\\ 1997, Hunter \\et\\ 1998, Rafikov 2001). Outer parts of dwarfs, where the gas density is even lower, therefore, present a particularly difficult test of our understanding of the cloud/star formation process. Yet, from broad-band images we see that stars have formed in the outer parts to very low surface brightness levels (for example, D.\\, Hunter \\et\\ 2009, in preparation).  We are using UV images obtained with the {\\it Galaxy Evolution Explorer} satellite ({\\it GALEX}; Martin \\et\\ 05) to trace and characterize star formation from the centers out into the outer disks of dwarf galaxies where \\ha\\ may not be an effective tracer of recent star formation (see, for example, Thilker \\et\\ 2005, 2007a; Boissier \\et\\ 2007). Fortunately, the UV---ultraviolet light coming directly from OB stars---can easily trace young stars in the outer galaxy, and the UV has the advantage in a patchy star-forming environment of integrating over a longer timescale ($\\leq$100--200 Myrs) than does \\ha\\ ($\\leq$10 Myrs).  In this paper we compare the UV surface brightness profiles to those of \\ha\\ and to those of the older stars as seen in broad-band images. We compute integrated and azimuthally-averaged profiles of several measures of star formation rates (SFRs): derived from FUV, \\ha, and $V$-band luminosities and model stellar population fits to broad-band UV, optical, and near-infrared colors. In a companion paper, we examine characteristics of individual star-forming regions identified on the {\\it GALEX} images and compare those in the outer disk with those in the inner disk (Melena \\et\\ 2009). The galaxies discussed in these two studies are also part of a larger, multi-wavelength survey of dIm, Blue Compact Dwarf (BCD), and Sm galaxies that has been assembled in order to examine the drivers of star formation in tiny galaxies (Hunter \\& Elmegreen 2004, 2006). ",
        "conclusions": "\\label{sec-concl}  We have presented UV integrated and azimuthally-averaged surface photometric properties of a sample of 44 dIm, BCD, and Sm galaxies. The UV measurements come from analysis of archival NUV and FUV images obtained with {\\it GALEX}, and we compare the UV to \\ha\\ and $V$-band properties. We convert FUV, \\ha, and $V$-band luminosities into SFRs. We also fit model stellar populations to colors for an alternate SFR measure for 7 different assumed star formation histories. We compare integrated SFRs and SFR profiles with radius in these four measures.  In most of the galaxies, the measures of SFRs track each other with radius. However, the UV profile often extends further in radius than the \\ha\\ profile, providing a better measure of the star formation activity in outer disks. Most of the dIm galaxies have constant or slightly increasing UV color with radius, as do the BCDs, while the Sm galaxies, which are also larger than the others, sometimes get bluer with radius. Star formation often extends so far into the outer disk that it would appear to be sub-threshold, according to conventional ideas. Such star formation may result from local triggering or from gravitational instabilities with angular momentum removal during cloud formation (Melena \\et\\ 2009).  We find that integrated SFRs determined from \\ha\\  are lower than SFRs determined from FUV for all but the highest SFR systems. The discrepancies are unlikely due to under-estimated extinction, plausible alternative conversion factors, or a larger contribution from older stars where the SFR is lower. On the other hand, the SFR determined from $M_V$ is the same on average as the FUV-based SFR for all but the lowest SFR systems, but with a large scatter about the equivalency relationship. This is consistent with a generally constant SFR that varies by factors of a few over long times. In addition, the SFR determined from modeling colors is always higher than the other 3 SFR measures, perhaps because the star formation rate has decayed over time, although it almost always has the same radial profile shape as $SFR_{\\rm V}$.  Galaxies with the largest discrepancies between FUV-observed star formation rates and $V$-band observed star formation rates also have the lowest relative HI masses. We suggest that the outer disks of these galaxies have very low gas column densities, causing a cessation in star formation, while in galaxies with relatively high HI masses, the star formation in the outer disk can be active or not with fluctuations on a Gyr timescale.  The lack of H$\\alpha$ in the outer disk is most likely the result of faint emission measures, rather than the result of radially varying or peculiar IMFs.  This conclusion follows quantitatively from a correlation between the ratio of star formation rates in FUV and H$\\alpha$ and the absolute H$\\alpha$ star formation rate. In our interpretation, a significant amount of H$\\alpha$ flux is missing from the outer disk star-forming regions, and so the star formation rate determined from H$\\alpha$ is too low.  We fit the observed correlation to a model with a constant disk thickness and a density-dependent star formation law.  The stellar surface densities in the outer parts of our galaxies reach values as low as $\\sim0.1\\;M_\\odot$ pc$^{-2}$, and the star formation rates get as low as $\\sim10^{-4}\\;M_\\odot$ yr$^{-1}$ kpc$^{-2}$ for Gyr periods.  Most likely the star formation in these regions is sub-threshold in the conventional sense. There is no evident break in the FUV radial profiles from the inner disks that might correspond to a star formation threshold, aside from the occasional presence of a kink in the exponential disk that seems unrelated to color changes. The kink is often far inside the furthest measured radius anyway, and thus apparently unrelated to a physical disk edge. No such edges have been observed yet."
    },
    "0911/0911.0748.txt": {
        "abstract": "{It is well established that the atomic interstellar hydrogen is  filling  the  galaxies  and  constitutes the  building  blocks  of molecular clouds.}  {To understand  the formation and the evolution of molecular clouds,  it is  necessary to investigate  the dynamics  of turbulent and  thermally bistable as well as barotropic flows.} {We  perform high resolution 3-dimensional  hydrodynamical simulations  of 2-phase,  isothermal and polytropic  flows.}   {We compare  the  density  PDF  and Mach  number density relation in the various simulations and conclude that 2-phase flows  behave  rather  differently  than polytropic  flows.   We  also extract the clumps and study their statistical properties such as mass spectrum,  mass-size relation  and internal  velocity  dispersion.  In each case,  it is found  that the behaviour  is well represented  by a simple  powerlaw.   While  the  various exponents  inferred  are  very similar  for   the  2-phase,  isothermal  and   polytropic  flows,  we nevertheless find significant differences, such as for example the internal velocity  dispersion  which  is  smaller  for  2-phase  flows.}   {The structures statistics are very similar  to what has been inferred from observations, in particular the  mass spectrum, the mass-size relation and  the velocity  dispersion-size  relation are  all powerlaws  whose indices well agree with the observed values.  Our results suggest that in  spite  of  various   statistics  being  similar  for  2-phase  and polytropic flows,  they nevertheless present  significant differences, stressing the  necessity to consider  the proper thermal  structure of the interstellar atomic hydrogen for computing its dynamics as well as the  formation  of   molecular  clouds.} ",
        "introduction": "Understanding the  interstellar turbulence  is of great  importance in the  context of  molecular cloud  and  star formation.  As such,  many theoretical  studies  and numerical  simulations  have been  performed during the  last decades (see e.g.   the reviews by  MacLow \\& Klessen 2004, Scalo  \\& Elmegreen 2004 and  Elmegreen \\& Scalo  2004). So far, most of the  studies  have considered  isothermal flows.  Although this constitutes  a reasonable  assumption  for the  densest  parts of  the molecular  clouds,  it  is  not  an appropriated  assumption  for  the description of  the interstellar atomic  hydrogen which is  2-phase in nature (e.g., Dickey \\& Lockman 1990, Field et al. 1969,  Wolfire et al.  1995) and therefore for  the   formation  of  molecular  clouds.    Indeed,  recent  works (Hennebelle \\& Inutsuka 2006, V\\'azquez-Semadeni et al. 2007, Heitsch et al.  2008a,  Hennebelle  et  al.  2008,  Banerjee  et  al.  2009)  have investigated  the possibility  that molecular  clouds  are multi-phase objects. It seems therefore important to understand the dynamical properties of a turbulent  2-phase medium as  the interstellar atomic  gas.  Various studies have been already performed  along this line. This includes 1D calculations (Hennebelle  \\& P\\'erault 1999,  Koyama \\& Inutsuka 2000,  S\\'anchez-Salcedo  et al.  2002, Inoue et al. 2006), 2D  calculations (Koyama  \\& Inutsuka 2002, Audit  \\& Hennebelle 2005, Heitsch et  al. 2005, 2006) and 3D calculations  (Kritsuk  \\&  Norman  2002, Gazol et al. 2005, V\\'azquez-Semadeni et al. 2006). The influence of the magnetic field on the dynamics of a thermally bistable flow has been investigated by Hennebelle \\& P\\'erault (2000), Piontek \\& Ostriker (2004, 2005), de Avillez \\& Breitschwerdt (2005), Hennebelle \\& Passot (2006), and more recently by Inoue \\& Inutsuka (2008), Hennebelle et al. (2009), Heitsch et al. (2009), Inoue et al. (2009) and Gazol et al. (2009).   In  a  previous study,  Hennebelle  \\&  Audit  2007, (hereafter  HA07)  and Hennebelle et al. (2007), we have investigated the dynamics of 2-phase flows  by  the  means   of  bidimensional  high  resolution  numerical simulations.   In particular, we  have studied  the properties  of the dense  and  cold  clumps  formed  out  of  the  warm  gas  by  thermal instability, showing that they present many similarities with observed clumps.  In the present paper, we investigate the properties of clumps formed  in more realistic  3D simulations.   To better  understand the influence  of the  2-phase  physics and since the nature of the inter-clump gas within molecular clouds is not well known (see \\S 4.1 for a discussion), we  also  perform isothermal  and polytropic simulations having identical setups as the 2-phase ones but for the thermal  properties. We then compare the  results obtained for the various types of simulations.   The  plan of  the paper  is the  following, in  the second  section we describe the numerical experiments we perform, in the third section we present the  global properties of the  numerical simulations including the density  PDF and the Mach number-density  distribution. The fourth section  is  devoted to  the  clump  properties  including their  mass spectrum, velocity  dispersion and  mass-size relation.  In  the fifth section we summarize our results and conclude the paper.    \\begin{figure*} %\\includegraphics[width=8cm,angle=0]{FIG/NEW/coupe_A1_2.ps} %\\includegraphics[width=8cm,angle=0]{FIG/NEW/coupe_isotherme2.ps} \\includegraphics[width=8cm,angle=0]{fig3_lr.ps} \\includegraphics[width=8cm,angle=0]{fig4_lr.ps} \\caption{Density   cut  through   the  simulations.   The   left  plot corresponds to the  2-phase run and the right  one to the isothermal run.  Bright  regions correspond to  a density between  20 cm$^{-3}$ and 100  cm$^{-3}$ while  the pervasive grey  area corresponds  to a density of a few cm$^{-3}$.} %Dark-blue, blue, green and red correspond respectively % to  densities of the order of 1 cm$^{-3}$, 3 cm$^{-3}$, 20 cm$^{-3}$ and 100 cm$^{-3}$. \\label{coupes} \\end{figure*} ",
        "conclusions": "In this paper, we perform  3D high resolution numerical simulations of converging flows  using either a standard  interstellar atomic cooling function,  an isothermal  or a  polytropic equation  of state  with an adiabatic index of $\\gamma=0.7$.  We  investigate  the density  PDF  and  Mach number-density  relation, showing  that, as expected,  the thermal  behaviour of  the gas  has a drastic  influence.   While as  previous  authors,  we  find that  the isothermal runs tend to  produce lognormal density distributions, when the adiabatic  exponent is smaller  than one, namely  $\\gamma=0.7$, we find  that  the  density  distribution  follow  a  powerlaw  for  high densities  confirming  the  result  of  Passot  \\&  V\\'azquez-Semadeni (1998).  The 2-phase  case is very different thought  the high density part of the distribution could  reasonably be described by a lognormal distribution.   We stress  that higher  numerical resolution  may well change this conclusion.  However,  the distribution of the low density (say  in the range  10-300 cm$^{-3}$)  cold gas  departs significantly from it, making unclear the  use of the lognormal density distribution an adequate choice for molecular clouds.  We  compute  the mass  spectrum  of  the  clumps for  various  density thresholds,  as  well  as  the  mass-size relation  and  the  internal velocity dispersion-size  relation.  We find that  the 3 distributions agree  with  the  statistics  inferred  for  the  interstellar  clouds (e.g.  Heithausen et  al. 1998)  reasonably well. In particular, the 3 distributions are well fitted by a powerlaw whose exponent values are close to the observed ones. While the  first 2 appear to be  reasonably similar for isothermal and  2-phase flows, we find  that the  internal velocity  dispersion is  about 2  to  3 times larger for the  clumps of the isothermal simulation  than for the ones of  the   2-phase  case,  suggesting  that  energy is less efficiently injected into the  dense clumps in 2-phase   flows  than it is in isothermal ones.  Altogether, these results confirm the  claim made in HA07 that 2-phase flows behave differently than isothermal flows. Although it appears to us that  higher resolution simulations should be  performed to further assess this  conclusion, the present  study suggests that  the 2-phase nature  of the  flow may  have significant  implications even  for the physics of high density gas."
    },
    "0911/0911.4633_arXiv.txt": {
        "abstract": "\\par\\parskip=1.truecm Some recent measurements of the chemical composition of the cosmic radiation indicate that at the energy of $3\\times10^{18}$  eV,  around the ankle, light cosmic ions dominate the spectrum as it occurs in the preknee energy region. \\newline Taking advantage of a recent theory of cosmic radiation which provides a quantitative explanation of the knee, the second knee and the ankle, the chemical composition of cosmic radiation is explicitly calculated giving individual ion spectra and  ion fractions  from 10$^{12}$ eV to $5\\times10^{19}$ eV. The calculation assumes two components of the cosmic radiation feeding the ion flux at Earth: one originated in the disc volume and another one,  called $\\it {extradisc}$ $\\it {component}$,  which from the disc boundaries traverses the Galaxy reaching the solar system.   Data above 10$^{17}$ collected during half century of experimentation by Auger, HiRes, Agasa, Akeno, Fly' s Eye, Yakutsk, Haverah Park and Volcano Ranch experiments are reviewed, examined and compared with the theoretical $<$$ln(A)$$>$. \\newline The comparison between computed and measured $<$$ln(A)$$>$ exhibits a good global accord up to $2\\times10^{19}$ eV except with the HiRes experiment and an excellent agreement in the range $10^{15}$-$10^{17}$ eV with Kascade, Eas-top, Tunka and other experiments. The accord requires a flux of the extradisc component of $1.8\\times10^{14}$ particles/($m^2$ sr s $eV^{1.5}$) at $10^{19}$ eV, twice that generated by disc sources. ",
        "introduction": "The spectrum \\footnote{ This paper is an english translation of an INFN Report to be published in italian: $\\it {Composizione }$ $\\it {Chimica}$ $\\it {della}$ $\\it {Radiazione}$ $\\it {Cosmica}$ $\\it {intorno}$ $\\it {alla}$ $\\it {Caviglia}$ $\\it {e}$ $\\it {Indici}$ $\\it {Spettrali}$ } of the cosmic radiation between the knee and the ankle has been recently calculated according to the $\\it {Theory}$ $\\it {of}$ $\\it {Constant}$ $\\it {Spectral}$ $\\it {Indices}$ [1-4]. The galactic mechanisms and the parameters generating the knee are the same at the origin of the ankle: this circumstance constitutes a major characteristic of the solution of the knee and ankle problem. Another notable aspect of this solution is the presence of a structure at $(5-7)\\times10^{17}$ eV which is called the second knee, just a direct consequence of the theory \\cite{seconginoc}. Figure 1 shows an example of excellent quantitative accord between computed and measured cosmic-ray spectrum [4].   \\begin{figure}[htb] \\vspace{-0.3cm} \\includegraphics [angle=0,width=8cm,height=8cm] {fig01.eps} \\vspace{-2cm} \\caption{ The differential cosmic-ray spectrum (black thick line) resulting from the theory normalized to $3.79 \\times 10^{17}$ particles/($m^2$ sr s $eV^{1.5}$) at $10^{12}$ eV (green square). The data are from the Kascade $\\cite{kaskaflussotot}$ and Haverah Park \\cite{haverahtot}  experiments.} \\label{fig:largenenough} \\vspace{-0.3cm} \\end{figure}    \\begin{figure}[htb] \\vspace{-0.3cm} \\includegraphics [angle=0,width=8cm,height=8cm] {fig02ING.eps} \\vspace{-1.5cm} \\caption{Energy spectra of individual ions and group of ions of Eas-top (red squares) $\\cite{eastop1}$ and Kascade (blue small circles) $\\cite{kascaionspec}$ experiments compared with the theory of constant spectral indices (black curve). Fluxes are normalized at $10^{12}$ eV. Ion grouping and flux of the Eas-top experiment shown in this figure  are examined elsewhere $\\cite{nota2flussidipro}$.} \\label{fig:largenenough} \\vspace{-0.4cm} \\end{figure}     The $\\it {Theory}$ $\\it {of}$ $\\it {Constant}$ $\\it {Spectral}$ $\\it {Indices}$ determines the energy spectra of individual ions, or group of ions, in the interval $10^{11}$-$5\\times10^{19}$ eV with only one normalization point for the particle flux. As a consequence, once the spectral indices and the abundances at a given energy are assigned, the chemical composition of the cosmic radiation is known at any energy.  Fig. 2 shows that individual ion spectra in the knee energy region are also in good agreement with the theory. The simplicity of the theory and  its anchorage to the empirical parameters governing the cosmic-ion motion in the Galaxy are not minor aspects besides the excellent quantitative agreement with ion spectra exhibited in figures 1 and 2. \\par The chemical composition expressed by $<$$ln(A)$$>$ has been also calculated with all cosmic-ray sources placed in the disc \\cite{salina}, thus ignoring any extragalactic component. Figure 3 reports the results of the calculation.  The profiles of $<$$ln(A)$$>$ measured by the Kaskade and Tunka experiments along with that extracted with the measurements of $X_{max}$ of the Auger Collaboration are shown in figure 3. There is an evident disagreement above $10^{17}$ eV between theory and data. All other experiments above $10^{17}$ eV have been purposely omitted from figure 3 to render the disagreement between theory and data more evident. According to the above quoted measurements the theoretical chemical composition,  as generated by the galactic cosmic-ray sources above $4\\times$$10^{17}$ eV,  is incompatible with the data, primarily because it is too heavy, and secondly, because  above $2\\times$$10^{19}$ it has a flat trend, contrary to the measured $<$$ln(A)$$>$,  which increases with energy.    \\begin{figure}[htb] \\vspace{-0.35cm} \\includegraphics [angle=0,width=8.5cm] {fig03.eps} \\vspace{-1.5cm} \\caption{Comparison of the $<$$ln(A)$$>$ derived from the $\\it {Theory}$ $\\it {of}$ $\\it {Constant}$ $\\it {Indices}$ with the Tunka \\cite{prosintunkaloga},  Kascade \\cite{kaskalogaorand} and Auger \\cite{augerloga}  data. A significant discrepancy between this calculation and the Auger data appears above $4 \\times 10^{17}$ eV. The theoretical evaluation of $<$$ln(A)$$>$ (blue line) assumes that all cosmic-ray sources are placed in the disc volume, with no extragalactic component.} \\label{fig:largenenough} \\vspace{0.8cm} \\end{figure}   Notice that the Auger Collaboration disposes today (2009) of a powerful and redundant  apparatus  and a number of recorded atmospheric cascades superior to that of any other experiment.  The purpose of this work is to calculate the chemical composition of the cosmic radiation at Earth,  not only assuming cosmic-ray sources in the disc,  but also postulating a component that penetrates the disc volume from the exterior reaching the Earth. This cosmic-ray component is called here $\\it {extradisc}$ being a particular component of the more vast class of the extragalactic cosmic rays.  Both the extradisc cosmic rays and their flux will be denoted by the same symbol $\\it I_{ed}$ ($\\it ed$ for extradisc). Let us anticipate here that $\\it {extradisc}$ $\\it {component}$ reconciles the results of the theory with the Auger data on $<$$ln(A)$$>$ shown in figure 3,  still preserving a quantitative accord with the knees, the ankles and the second knee.   \\begin{table}[htb] \\begin{center} \\caption{ Parameters of the theory denoted Low Energy ($LE$) and High Energy ($HE$) ion blends. Fluxes are multiplied by $E^{2.5}$ and expressed in units  of $(m^{-2} sr^{-1} s^{-1} eV^{1.5})$.}  \\begin{tabular}{lrrrr} \\hline \\vspace{-0.05cm} Blend.        & LE              &          & HE              &         \\\\ \\hline \\vspace{0.2cm} & $10^{12}$ eV    &          & $10^{14}$ eV    &         \\\\ \\hline \\vspace{0.2cm} & \\%              & $\\gamma$ & \\%              & $\\gamma$ \\\\ \\hline \\vspace{0.2cm} H             & 42.4            & 2.77     & 33.5            & 2.67     \\\\ He            & 26.5            & 2.64     & 27.3            & 2.64     \\\\ CNO           & 11.9            & 2.68     & 13.6            & 2.65     \\\\ Ne-S          &  9.2            & 2.67     & 10.9            & 2.63     \\\\ Ca(17-20)     &  1.2            & 2.67     &  2.8            & 2.63     \\\\ Fe(21-28)     &  8.7            & 2.59     & 12.0            & 2.62     \\\\ \\hline & Flux            & $\\gamma$ & Flux            & $\\gamma$ \\\\ \\hline H             & $1.15\\ 10^{17}$ & 2.77     & $3.93\\ 10^{16}$ & 2.67     \\\\ He            & $7.19\\ 10^{16}$ & 2.64     & $3.20\\ 10^{16}$ & 2.64     \\\\ CNO           & $3.24\\ 10^{16}$ & 2.68     & $1.60\\ 10^{16}$ & 2.65     \\\\ Ne-S          & $2.50\\ 10^{16}$ & 2.67     & $1.28\\ 10^{16}$ & 2.63     \\\\ Ca(17-20)     & $3.14\\ 10^{15}$ & 2.67     & $3.27\\ 10^{15}$ & 2.63     \\\\ Fe(21-28)     & $2.36\\ 10^{16}$ & 2.59     & $1.41\\ 10^{16}$ & 2.62     \\\\ \\hline Total         & $2.71\\ 10^{17}$ & 2.70     & $1.18\\ 10^{17}$ & 2.64     \\\\ \\hline \\end{tabular} \\end{center} \\vspace{-0.5 cm} \\end{table}  The structure of this paper preassumes the existence of the $\\it {Theory}$ $\\it {of}$ $\\it {Constant}$ $\\it {Spectral}$ $\\it {Indices}$ [1-4] and the companion paper \\cite{salina} where the chemical composition of the cosmic radiation is evaluated under the restricted assumption that all cosmic-ray sources feeding the particle flux at Earth are placed in the disc volume. Section 2 presents a brief survey of the bases of the calculation and a flash of the cultural background of the theory. Section 3 is devoted to a survey on cosmic-ion energy spectra useful for the normalization of the theory at $10^{12}$ eV. Section 4 summarizes the calculation of the chemical composition obtained with the galactic sources only. Section 5 reports a comparison between data and theory under the assumption that all cosmic ray sources are placed in the disc and not in its exterior. Section 6 reports many details of the calculation of the ion abundances of the $\\it {extradisc}$ $\\it {component}$ and the attenuation of its intensity while penetrating the disc volume from its periphery up to the solar system.  The results of the calculation reported  in Section 6 are utilized in Section 7,  where the $<$$ln(A)$$>$  for the disc and extradisc components are given. The comparison of the predicted $<$$ln(A)$$>$ with the data for the two components $\\it I_{ed}$ and $\\it I_{d}$ (disc component)  are given in Sections 8 and 9. Section 10 marks the disagreement on the chemical composition between the HiRes experiment and others.  Conclusions are in the last Section 11.    \\begin{figure}[htb] \\vspace{1.35cm} \\includegraphics [angle=0,width=8cm]{fig04.eps} \\vspace{0.1cm} \\caption{ Energy spectra of $11$ individual ions measured by balloon and satellite detectors at energies below $10^{15}$eV suggesting and founding the  postulate of the constant spectral indices ( see ref. \\cite{cod-merid} and references therein).} \\label{fig:largenenough} \\vspace{0.1cm} \\end{figure}    \\begin{figure}[htb] \\vspace{0.1cm} \\includegraphics [angle=0,width=8cm]{fig05.eps} \\vspace{0.5cm} \\caption{Spectral indices of cosmic ions from Hydrogen to Nickel measured by balloon-borne and satellite detectors at $10^{12}$ eV (filled red dots) and $10^{14}$ eV (open blue dots). All indices extracted from the spectra around $10^{14}$ eV fall between $2.6$ and $2.7$. } \\label{fig:toosmall} \\vspace{0.1cm} \\end{figure}    \\begin{figure}[htb] \\vspace{-0.35cm} \\includegraphics [angle=0,width=8cm]{fig06.eps} \\vspace{-1.5cm} \\caption{Measured fluxes of individual ions of the Wiebel-Sooth Compilation ($LE$ blend) (red dots) and $HE4$ Compilation (open blue dots).} \\label{fig:largenenough} \\vspace{-0.7cm} \\end{figure}    \\begin{figure}[htb] \\vspace{-0.3cm} \\includegraphics [angle=0,width=8cm] {fig07.eps} \\vspace{-1.5cm} \\caption{  Theoretical spectra of 6 ions for the $LE$ blend and the resulting $\\it{complete}$ $\\it{spectrum}$ (thick green upper curve).} \\label{fig:largenenough} \\vspace{-0.4cm} \\end{figure}     \\begin{figure}[htb] \\vspace{-0.3cm} \\includegraphics [angle=0,width=8cm]{fig08.eps} \\vspace{-2.1cm} \\caption{ Theoretical spectra of 6 ions for the $HE4$ ion blend and the resulting $\\it{complete}$ $\\it{spectrum}$ (thick upper green curve) normalized at $10^{12}$ eV with a flux of $3.89\\times10^{17}$ particles/$m^2$ sr s $eV^{1.5}$. } \\label{fig:largenenough} \\vspace{-0.5cm} \\end{figure} ",
        "conclusions": "The comparison of the chemical composition of the cosmic radiation expressed by $<$$ln(A)$$>$ derived from the $\\it {Theory \\quad of \\quad Constant \\quad Indices }$ with the observed $<$$ln(A)$$>$  extracted from the Auger data on $X_{max}$ dictates the necessity of introducing an extradisc component of the cosmic radiation above $10^{17}$ eV. In Sections 1 and 2 it is explained why this component is better termed $\\it {extradisc}$ $\\it {component}$,  $I_{ed}$. The variable $r$ is a free parameter of the theory defined as the extradisc-to-disc flux ratio in the solar cavity at $10^{19}$ eV, e.g. $r$=$I_{ed}$/$I_{d}$.  From the ensemble of the profiles  of $<$$ln(A)$$>$ extracted from the measurements of $X_{max}$ of the Auger, Yakutsk, Fly' s Eye, Agasa, Akeno, Haverah Park and Volcano Range experiments it emerges a global accord with the theory.   In the energy interval $10^{15}$-$10^{17}$ eV there is an excellent agreement with the observations of Eas-top, Kascade, Tunka (see fig. 3) and other experiments. Below $10^{17}$ eV  the extradisc component is a negligible fraction of the disc component $I_d$, and as a consequence, the theoretical and empirical analysis reported in the companion paper [8] remains valid.  Because of the unknowns entailed in the conversion of $X_{max}$ into $<$$ln(A)$$>$ above $10^{17}$ eV,   the accord between the theoretical profile and data in figures 27, 28, 30, 31, 32 and 33 is more than satisfactory. Values of $r$= $I_{ed}$/$I_{d}$ higher than 2 would alleviate the gap between the empirical profile of $<$$ln(A)$$>$ and the theory, as shown in figure 27. Presently, values of $r$ higher than 2 are neither convincing nor necessary, due to large systematic errors in the experimental data. The last two figures 35 and 36 intend to illustrate this position.   \\par Figure 35 depicts in a flash the present experimental uncertainties in flux measurements quoting arbitrarily data from Kascade, Yakutsk and Auger experiments. Fluxes from other experiments do not agree as well. Such uncertainties impede to identify  a more precise and reliable value of the ratio $r$ in figure 26 on an empirical basis,  being $r$=$2$ the present adequate global estimate.   Let be $X^{cal}_{max}$ (H) the theoretical atmospheric depth for protons, $X^{cal}_{max}$ (Fe) that for Fe nuclei and $D_{max}$ the difference between the two depths i.e. $D_{max}$ = $X^{cal}_{max}$ (H)- $X^{cal}_{max}$ (Fe) at a given energy $E$. Error sources generating differences in  $<$$ln(A)$$>$ in theories and in experiments might arise both from $X^{cal}_{max}$ (H) and  $D_{max}$ as well. Ideally, different hadronic models may result in a combination of $X^{cal}_{max}$ (H) and $D_{max}$ having the same $<$$ln(A)$$>$. Figure 36 reports the $D_{max}$   versus energy in some current hadronic models and, at the arbitrary energy of $10^{17.5}$ eV, the $X^{cal}_{max}$ (H). Figures 36 and 12 show that the $X^{cal}_{max}$ (H)  and $D_{max}$ profiles used in this paper for the conversion of  $X_{max}$ into  $<$$ln(A)$$>$ are quite similar to those adopted by others (for example, Auger). \\par Notice that the hadronic code called QGSjet model-03 \\cite{orandellogdea}, having a larger $D_{max}$ of 120-130 $g$/$cm^2$ in the relevant range $10^{17}$-$5\\times 10^{19}$ eV and a deeper $X^{cal}_{max}$ (H) of 739 $g$/$cm^2$  at $10^{17.5}$ eV (see figure 36), it would shift upward the simulated $X_{max}$ profiles in figure 12 by 9 $g$/$cm^2$ at $10^{17.5}$ eV. The same correction applied to the Auger data in figure 3 (or figure 14) would shift upward by 0.57 units of $<$$ln(A)$$>$ (9 $g$/$cm^2$ for $r$=$0$ or 11 $g$/$cm^2$ for $r$=$2$) at $4.12\\times10^{17}$ eV alleviating the gap between Auger data and theory. As far as uncertainties in hadronic models above $10^{17}$ eV used in the analysis of giant atmospheric cascades remain large, the determination of a  value of $r$ higher than 2 based on the $<$$ln(A)$$>$ data would appear premature.  This same conclusion is suggested by the flux measurements in figure 35.  Let us finally mention the systematic disagreement between the $<$$ln(A)$$>$ extracted from the $X_{max}$ measured by the HiRes experiment and  theory,  amounting to about one unit in the entire interval $10^{17}$-$5\\times 10^{19}$ eV (see fig. 29). The examination of the HiRes data on $<$$ln(A)$$>$ made in Section 9 suggests a number of inconsistencies with the data of numerous experiments and with different versions of the HiRes experiment itself (HiRes Prototype, HiRes Stereo, data revision \\cite{belzsalina2008}).  \\vspace{-0.2 cm}"
    },
    "0911/0911.2403_arXiv.txt": {
        "abstract": "We consider the effect of radial fluid injection and suction on Taylor-Couette flow.  Injection at the outer cylinder and suction at the inner cylinder generally results in a linearly unstable steady spiralling flow, even for cylindrical shears that are linearly stable in the absence of a radial flux.  We study nonlinear aspects of the unstable motions with the energy stability method. Our results, though specialized, may have implications for drag reduction by suction, accretion in astrophysical disks, and perhaps even in the flow in the earth's polar vortex. ",
        "introduction": "There are many contexts in which rotating flows arise and their instabilities are naturally of interest.  For an axisymmetric, incompressible, inviscid flow with no variation along the axis of symmetry, Rayleigh (1916) provided a necessary and sufficient condition for stability against axisymmetric perturbations (Chandrasekhar, 1961). Rayleigh also showed that a necessary criterion for the instability of such flows to non-axisymmetric perturbations is an analogue to his classic inflection point criterion for plane flows.  He expressed this criterion in words  but the correct mathematical expression is easily worked out following Rayleigh's lead (Spiegel and Zahn, 1970). For the study of the effect of viscosity on these results, the Taylor-Couette configuration of flow between coaxial cylinders provides a good proving ground (Chandrasekhar, 1961), although shear and rotation come together in other contexts as well (Yecko and Rossi, 2004 and references therein).  When criteria for the simplest two-dimensional circular shear flows indicate stability, the question naturally arises whether some simple extraneous effects may be destabilizing. In the case of the linearly stable plane shear flow, magnetic fields (Stern, 1963 and Chen and Morrison, 1991), Rossby waves (Lovelace {\\it et al.}, 1999) and suction (Hocking, 1974; Doering {\\it et al.}, 2000) have been found to be destabilizing mechanisms, given suitable conditions. In the case of  a stable Taylor-Couette flow, it had been seen even earlier that a magnetic field may be destabilizing (e.g. Chandrasekhar, 1961; Balbus and Hawley, 1998). In an attempt to add to this compendium of instabilities, we here investigate the instability that arises when a formerly circular flow is turned into a spiralling flow by the imposition of continous suction on the inner cylindrical boundary with a compensatory injection of fluid on the outer boundary. For brevity, we shall hereinafter refer to the induced radial inflow as  {\\it accretion} and the mechanism driving it as {\\it suction}.  A practical example of the conversion of differential circular flow to spiraling motion, or {\\it swirl}, arises in the study of drag reduction on a cylindrical airfoil (Batchelor, 1967). In that case,  when suction is imposed on the central cylindrical airfoil, the drag is reduced because boundary layer separation is inhibited. Spiralling flows are also central to the structure and dynamics of accretion disks that form in a great variety of astrophysical situations such as binary stars (which are often X-ray sources), flows around massive black holes at the centers of galaxies, and in the disks around newly born stars where planetary systems are thought to form.  Of course, accretion disks are more complicated than the simple model that we consider here.  In most astrophysical cases there is ionization so hydromagnetic effects  become important and can produce the magneto-rotational instability (see Balbus and Hawley, 1998).  However, protoplanetary nebulae such as the primitive solar nebula are believed to have been cool so that other mechanisms may need to be considered to get the full picture (for at least some phases) of the turbulent dynamics of planet formation. The feature of the present study that may have relevance to the fluid dynamics of astrophysical disks is the radial inflow, or accretion that, in some astrophysical examples, is fed by inflow from an external mass source. A similar combination of motions occurs also in the terrestrial polar vortex that may play a role in the transport of ozone in the upper atmosphere (McIntyre, 1995).  With these motivational examples in mind, we turn to an idealized pilot problem to learn what may happen when rotation and accretion combine to produce a spiralling motion as in Taylor-Couette flow with suction. Although we cannot draw immediate conclusions about the applications of our results to the motivational examples just described, especially to accretion disks, we regard it as suggestive that the combination of rotation and accretion (i.e., radial injection and suction) in a flow leads to instability. As we shall see in what follows, taking into account both accretion and rotation leads to very different stability properties of the flow than when only one of these effects alone is included. For example, in the classical problem of the linear and nonlinear stability of Taylor-Couette flow between concentric cylinders,  the system is linearly stable when the outer cylinder is rotating and the inner one is stationary. Likewise, simple radial accretion without rotation is linearly stable. However, we find that, at sufficiently large rotational Reynolds numbers in a stable TC flow, small accretion rates lead to linear instability. Moreover we find that the rotational motion, rather than simple plane shear, plays a major role in the stability characteristics.  An initially plane-parallel shear layer perturbed by transverse suction may also exhibit a small-suction instability (Hocking, 1974), while larger suction rates absolutely stabilize the corresponding laminar flow for arbitrarily large Reynolds numbers (Doering {\\it et al.}, 2000). This strong stabilizing effect of suction for even very high Reynolds numbers in the plane case does {\\it not} carry over to the rotational case; in the cylindrical geometry transient growth of perturbations is generally encouraged by strong suction.  The following discussion of these issues is organized as follows. In section \\ref{secform} we describe Taylor-Couette flow with suction. Then we treat the linear stability of this flow in section \\ref{seclin} where we find linear instabilities in the simple flow of the model problem. In section \\ref{secnarrow} the asymptotic behavior of the onset of linear instability is computed for different limits in the narrow-gap approximation. We next go on to consider some nonlinear aspects of the instability. First, in section \\ref{secenergy} we describe the energy stability method and then derive some rigorous bounds on the energy stability limit of the flow in section \\ref{secbounds}. That  limit is computed numerically in section \\ref{secnumc}. The concluding section \\ref{conclusion} summarizes the results and mentions some possible implications. In an appendix we note that physically relevant bounds on the mean energy dissipation rate for non-steady flows (including statistically stationary turbulence) remain unknown in the Taylor-Couette geometry with suction. ",
        "conclusions": "\\label{conclusion} It is not unusual to suspect that a given flow may be unstable only to encounter mathematical arguments to the contrary. A familiar example for fluid dynamicists is provided by kinematic dynamo theory. In a working kinematic dynamo, a magnetic field grows exponentially from a small magnetic perturbation on a suitable flow.  To many this prospect once seemed to be forbidden by Cowling's  antidynamo theorem. The way out of the dilemma was to (finally) appreciate that Cowling's theorem posited an axisymmetric flow and to focus on flows that did not respect this constraint. In the present problem we confront a similar situation: many flows with differential rotation may seem to have enough energy to excite the growth of disturbances to the flow yet Rayleigh's criteria appear to forbid a purely fluid dynamical instability. In this instance, we have contravened Rayleigh's proscription by simply violating one of the conditions of Rayleigh's demonstrations. In both of these examples, the value of the relevant (anti-)theorem has been to point the way to making the (seemingly) forbidden event happen by vitiating one of the theorem's premises.  As we mentioned at the outset, there are several interesting problems where differential rotation plays an important role. We have not hidden our particular interest in the situation of accretion disks where matter rotates differentially around a central gravitating body. These disks are often quite luminous and the source of their emissions has been thought to be in the dissipation of the turbulence that they had long been assumed to sustain. But these objects typically have Keplerian rotation laws where their circular velocities vary as $1/\\sqrt{r}$, where $r$ is the distance from the central object. Hence there was no generally accepted purely fluid dynamical process to produce this turbulence, though the matter was steeped in controversy. Moreover, turbulence in the disk, if it occurred, was supposed to transport angular momentum out of the disk. This would allow the purely circular motion to begin to turn inward so that the gas in the disk could fall onto the central object and be accreted. The resulting spiral motion, or swirl, is then a central feature of the accretion disks yet its dynamical role has been largely glossed over in the discussions of disk instabilities.  For us, the intriguing feature of the instability of swirling (a.k.a. spiraling) flows in disks is that their instabilities may promote accretion that may in turn enable the inward flow that was the source of instability. This is positive feedback {\\bf par excellence} and we may reasonably anticipate that the instability treated here is likely to cause subcritical bifurcation in disks, though we have not yet shown that. Nevertheless, this physical argument leads us to suppose that rather than thinking of a competition among instability mechanisms, we may imagine a cooperation amongst them. Thus in the linear study of the magneto-rotational instability Kersale {\\it et al.} (2004) find that the resulting fluid stresses drive accretion for suitable boundary conditions. This self-induced accretion means that a swirl is produced that is likely to contribute to the instability we consider here.  This interaction of instabilities in the fully nonlinear regime bids fair to give rise to many of the processes that have been sought to explain the observed wealth of behavior in swirling flows.  However, the work reported here does not treat disks explicitly. We have chosen instead to consider a fluid flow field that is well known to be stable and to show that by converting it from purely circular to swirling motion by the agency of accretion we could render it linearly unstable. This required only a simple change in the boundary conditions to vitiate the applicability of Rayleigh's inflection point theorem for circular motion. It remains to learn whether this destabilization will apply to Keplerian flows, though the mechanism seems generic and we do think the instability is inherent to physically realistic, spiraling flows. As to real disks, there we are in the domain of shallow gas theory and that will lead to a more difficult investigation, but in the beginning it will be a linear one. In the  meantime, it could well be that other instabilities will be uncovered as other variations on disk flows are examined.  These will also join the magnetorotational instability that now leads the pack and whose consequences are currently being vigorously explored by many investigators. \\break \\bigskip  \\noindent {\\it Acknowledgment.} This work was initiated, and much of it carried out, during the 2007 Geophysical Fluid Dynamics Program at Woods Hole Oceanographic Institution. We are grateful to the US National Science Foundation and the US Office of Naval Research for their ongoing support of this program. This work was also supported in part by NSF Awards PHY-0555324 and PHY-0855335.  \\appendix"
    },
    "0911/0911.0892_arXiv.txt": {
        "abstract": "In this paper we present a dust emission map of the starless core TMC-1C taken at 2100~\\micron.  Along with maps at 160, 450, 850 and 1200~\\micron, we study the dust emissivity spectral index from the (sub)millimeter spectral energy distribution, and find that it is close to the typically assumed value of $\\beta = 2$.  We also map the dust temperature and column density in TMC-1C, and find that at the position of the dust peak ($A_V \\sim 50$), the line-of-sight-averaged temperature is $\\sim$7~K.  Employing simple Monte Carlo modeling, we show that the data are consistent with a constant value for the emissivity spectral index over the whole map of TMC-1C. ",
        "introduction": "Starless cores are often identified through maps of their dust emission at (sub)millimeter wavelengths, and recent surveys of nearby molecular clouds have increased the number of such objects considerably \\citep[e.g.][]{Hatchell07, Enoch08, Kauffmann08, Simpson08}.  Analysis of the properties of starless cores is often restricted by sparse sampling of their spectral energy distribution (SED), uncertain dust emissivities and assumptions about their geometries.  As a result, the measured (sub)millimeter fluxes are often used in conjunction with assumed values for the dust temperature and/or emissivity to determine core masses, though the systematic errors can be substantial (a factor of a few).  With sufficient spectral coverage and knowlege of source geometry, it should be possible to map the column density, temperature and dust emissivity spectral index in a starless core, though this has yet to be accomplished.  One of the best-studied starless cores is TMC-1C, located in the Taurus molecular cloud at an approximate distance of 140 pc \\citep{Torres09}.  Previous studies of the dust emission from TMC-1C at 450, 850 and 1200 \\micron\\ have determined that it is cold ($T_d \\sim$6 K) and dense ($n \\sim$10$^6$ cm$^{-3}$) at its center, and becomes progressively less dense and warmer at larger radii \\citep{Schnee07a, Schnee05}.  The dust-derived mass of TMC-1C is a factor of a few larger than its virial mass, and self-absorbed \\nthp\\ spectra show evidence for sub-sonic infall \\citep{Schnee07b, Schnee05}.  The chemistry of TMC-1C is characterized by depletion of CO and its isotopologues, as well as depletion of CS and \\nthp\\ \\citep{Schnee07b}.  Previously reported mass and temperature profiles for TMC-1C are estimated assuming a constant value for the emissivity spectral index ($\\beta = 1.8$) \\citep{Schnee07a}.  However, a wide range of values for $\\beta$ in other regions have been measured or calculated on theoretical grounds.  For instance, in the diffuse interstellar medium (ISM), an emissivity spectral index of $\\beta = 2$ fits the typical far-infrared spectrum well \\citep{Draine07, Draine84}, while grain growth in protostellar disks results in values of $\\beta \\simeq 1$ \\citep{Beckwith91}.  Since starless cores are in an intermediate stage between the diffuse ISM and protostars, one might expect that the dust in a starless core would have an emissivity spectral index in the range $1 \\le \\beta \\le 2$, although the timescale of grain growth at densities $n_H \\sim 10^6$ cm$^{-3}$ is several hundred times the free-fall time of a typical starless core \\citep{Chakrabarti05}.  On the other hand, the growth of icy mantles on dust grains could steepen the slope of the dust SED, and this has been cited as the reason for high values of the emissivity spectral index in parts of the Orion Ridge ($\\beta \\sim 2.5$) \\citep{Lis98} and in Sgr B2(N) ($\\beta \\sim 3.7$) \\citep{Kuan96}.  In the cold and dense evironment in TMC-1C, the dust grains could form similar icy mantles, which might drive the emissivity spectral index to values $\\beta > 2$.  In addition to changes in $\\beta$, grain growth by either coagulation or accretion of an icy mantle will change the opacity of the dust grains, $\\kappa$.  For instance, evidence that dust opacities are higher in dense cores than in the diffuse ISM can be found by comparing gas temperature and dust emission profiles \\citep{Keto08} or by comparing dust and gas emission profiles \\citep{Keto04}.  From models of dust coagulation for conditions that are plausible for a starless core, \\citet{Ossenkopf94} show that the dust opacity at 1.3~mm can be a factor of five higher than in a diffuse environment. Using 160~\\micron\\ data for Taurus, \\citet{Terebey09} measure an opacity of $\\kappa_{160~\\micron} = 0.23 \\pm 0.046$ cm$^2$g$^{-1}$ in the cold cloud ($A_V$ ranging from 0.4 to 4.0).  The opacity is 2.6 times greater than the diffuse ISM opacity \\citep{Weingartner03}, and provides evidence for higher opacity values in dense regions \\citep[c.f.][]{Ossenkopf94}.  Laboratory studies of possible dust grain materials have measured the dust emission and absorption properties at far-infrared and longer wavelengths.  For instance, \\citet{Aannestad75} measured the absorption of silicate grains and found that $\\beta \\simeq 2$ for grains without an ice mantle, and $\\beta \\simeq 3$ for grains with an ice mantle.  The steepest spectral indices calculated by \\citet{Aannestad75} were for fused quartz and olivine with ice mantles, both of which had $\\beta = 3.5$, close to the extremely steep $\\beta = 3.7$ reported by \\citet{Kuan96}.  In a study of the millimeter-wave absorption of silicate grains in the temperature range 1.2 - 30~K, \\citet{Agladze96} found that the emissivity spectral index depends on the dust temperature and for some materials can vary between $1.5 \\le \\beta \\le 2.5$, with the largest value for $\\beta$ at a temperature $T_d \\sim 10$~K.  A temperature-dependent emissivity spectral index was also found by \\citet{Mennella98} for cosmic dust analog grains in the temperature range $24 \\le T_d \\le 295$~K and a wavelength range $20~\\micron \\le \\lambda \\le 2$~mm, with a steeper SED at lower temperatures.  The range of emissivity spectral indices is $1.2 \\le \\beta \\le 2.4$ for the various grain types considered by \\citet{Mennella98}.  In a study of amorphous silicates in the temperature range $10 \\le T_d \\le 300$~K and wavelength range $0.1 \\le \\lambda \\le 2$~mm, \\citet{Boudet05} also report an anticorrelation between $T_d$ and $\\beta$, with values of the emissivity spectral index between $1.5 \\le \\beta \\le 2.5$.  Based on these laboratory measurements, one might expect that in starless cores like TMC-1C, which have temperatures close to 10~K \\citep{Schnee09}, the emissivity spectral index would be close to $\\beta \\simeq 2.5$.  Observations of nearby star-forming regions suggest an anticorrelation between $T_d$ and $\\beta$ similar to that measured in the laboratory. For instance, observations of the Orion, $\\rho$ Ophiuchi and Taurus molecular clouds at far-infrared and (sub)millimeter wavelengths by \\citet{Dupac03} show that the emissivity spectral index may vary from $2.5 \\ge \\beta \\ge 1.0$ as the dust temperature rises from $10 \\le T_d \\le 80$~K.  Similarly, in a study of the far-infrared and (sub)millimeter SEDs of luminous infrared galaxies, \\citet{Yang07} report that the emissivity spectral index varies from $2.4 \\ge \\beta \\ge 0.9$ as the dust temperature varies from $30 \\le T_d \\le 60$~K. In a study of the starless core L1498, \\citet{Shirley05} find an emissivity spectral index of $\\beta = 2.44 \\pm 0.62$ for an assumed dust temperature of $T_d = 10.5 \\pm 0.5$~K, consistent with the hypothesis that the cold dust in a starless core has a steep spectral index.  However, variations in the dust properties along the line-of-sight and noise in the observed maps can mimic an anticorrelation between $T_d$ and $\\beta$ \\citep{Shetty09a, Shetty09b}.  Furthermore, at high densities and low temperatures, grain coagulation will decrease $\\beta$ while the possible $T_d$--$\\beta$ relation will increase the emissivity spectral index.  Therefore, it is not at all clear that one should expect the emissivity spectral index to vary significantly in a starless core.  Observations of dust emission and extinction in the dark cloud B68 show no evidence for variations in the dust properties \\citep{Bianchi03}, suggesting that to some extent the various pressures on $\\beta$ may cancel out.  In this paper we present maps of the starless core TMC-1C at 160, 450, 850, 1200 and 2100~\\micron.  By spatially resolving the emission from TMC-1C over a wide range of wavelengths, we then map the column density, temperature and emissivity spectral index of the dust, which has typically been difficult to do with previous datasets. ",
        "conclusions": "\\label{DISCUSSION}  In Section \\ref{VARIABLEBETA} we found that there is an anti-correlation between the dust temperature and emissivity spectral index in the two dimensional maps of TMC-1C.  Although this phenomenon has been reported in other star-forming regions \\citep{Dupac03} and in other galaxies \\citep{Yang07}, this is not a known feature of starless cores.  The assumption that the temperature along each line-of-sight is constant, even though the temperature profile in TMC-1C is very likely to be colder at higher densities, can lead to the mistaken conclusion that there is an anti-correlation between $T_d$ and $\\beta$ \\citep{Shetty09a, Shetty09b}.  Noise in the dust emission maps can also drive a spurious $T_d - \\beta$ anti-correlation.  Here we investigate whether the dust temperature and emissivity spectral index are truely anti-correlated, and conclude that this is not necessarily the case.  We then determine what constant value of $\\beta$ best characterizes TMC-1C.  \\subsection{Testing the $T_d - \\beta$ Anti-correlation}  To determine the effects of temperature variations along the line of sight and noise on the dust properties derived from flux maps at 160, 450, 850, 1200 and 2100~\\micron, we consider a spherical model of a starless core with plausible properties.  The density in the core is given by \\begin{equation} \\label{DENSITYEQ} n(r) = \\frac{n_0}{1 + (r/r_0)^{2.5}} \\end{equation} where $r$ is the radius, $n_0 = 10^6$~cm$^{-3}$ and $r_0 = 0.005$~pc. The temperature in the core is given by \\begin{equation} \\label{TEMPEQ} T_d(r) = T_0 \\left(1 + r/r_1 \\right) \\end{equation} where $T_0 = 6.4$~K and $r_1 = 0.15$~pc.  The emissivity spectral index is constant in the core at a value $\\beta = 2.2$.  This model core is colder and denser at the center, as expected for an externally heated starless core.  We do not claim that these are the temperature and density profiles for TMC-1C, but there is a general agreement between the shape of the model SED and that of TMC-1C.  The density and temperature profiles are shown in Figure \\ref{MODELSPHEREPLOT}, as is the total emergent SED.  For each line of sight passing through the model core, we determine the volume of intersection between a cylinder and each spherical shell of the core.  Given the density and temperature distributions from Equations \\ref{DENSITYEQ} and \\ref{TEMPEQ}, the flux from each layer is given by \\begin{equation} \\label{SHELLEQ} S_\\nu = \\frac{B_\\nu(T_d) \\kappa_\\nu n_r V_r}{d^2} \\end{equation} where $V_r$ is the volume of intersection and $d$ is the distance to the core.  The two-dimensional flux distribution is calculated by summing the flux from each layer along each line of sight.  For the sake of convenience, the total flux density of the model core, summed over all radii, is scaled by a common factor such that the peak of the SED has a value of unity.  We ``observe'' the model core by deriving two-dimensional flux maps at 160, 450, 850, 1200 and 2100~\\micron\\ and adding Gaussian random noise equal to 5\\% of the peak flux in each map.  Using equations \\ref{FLUXEQ} -- \\ref{KAPPAEQ}, we then derive two-dimensional maps of the column density, dust temperature and emissivity spectral index, assuming that these quantities are constant along the line of sight (as was done in Section \\ref{VARIABLEBETA}).  The derived column density map peaks at the center of the core, where the temperature is seen to decrease, as one would expect from the input density and temperature profiles (see Figure \\ref{MODELSPHEREPLOT}).  The derived emissivity spectral index map covers a wide range of values, mostly between $1.0 \\le \\beta \\le 3.5$, even though the input emissivity spectral index was held constant.  The median derived emissivity spectral index is \\MODELBETA\\ $\\pm$ \\MODELBETARMS, which is close to the input value of $\\beta = 2.2$.  Furthermore, the derived temperature and emissivity spectral index are seen to be anti-correlated.  The combination of noise and temperature variations along the line of sight have been shown to drive an anti-correlation between these two quantities \\citep{Shetty09a, Shetty09b}, so it is not surprising that this effect is seen here. Interestingly, the portion of the $T_d - \\beta$ parameter space covered by the model starless core overlaps significantly with that of TMC-1C (see Figures \\ref{VARIABLEBETA} and \\ref{MODELSPHEREPLOT}). From this we conclude that there is no definitive evidence that the emissivity spectral index in TMC-1C varies with temperature, and that the more reliable estimates of column density and temperature will come from fitting the observed fluxes while holding the emissivity spectral index constant.  \\subsection{Average Emissivity Spectral Index in TMC-1C}  Estimates of the emissivity spectral index in TMC-1C derived from its average SED (see Section \\ref{MAP1D}) and from two dimensional maps using the 450, 850 and 1200~\\micron\\ maps to predict the lower resolution 160 and 2100~\\micron\\ maps (see Section \\ref{CONSTANTBETA}), $\\beta = \\SINGLESEDBETA \\pm \\SINGLESEDBETARMS$ and $\\beta = \\MAPSINGLEBETA \\pm \\MAPSINGLEBETARMS$, are in good agreement with each other.  Given that an analysis of the flux maps from an idealized core, with line-of-sight temperature variations and noise, gives the correct {\\it average} emissivity spectral index (within errors), we suggest that the value of the emissivity spectral index in TMC-1C is well-represented by $\\beta \\simeq 2.2 \\pm 0.5$. This is consistent with the emissivity spectral index in the diffuse ISM of $\\beta = 2$ \\citep{Draine07, Draine84} and with the emissivity spectral index measured in the starless core L1498 \\citep[$\\beta = 2.44 \\pm 0.62$,][]{Shirley05}.  However, due to the large uncertainty in $\\beta$ in this study, we cannot rule out the possibility that the emissivity spectral index in TMC-1C could also be intermediate between the shallower SEDs seen in circumstellar disks and the SED of the ISM.  When deriving the dust temperature and mass of prestellar cores from sparsely-sampled SEDs, it is common to assume that the emissivity spectral index of the dust is equal to 2 \\citep[e.g.][]{Kirk07, Simpson08}.  Observations of L1498 \\citep{Shirley05} and our observations of TMC-1C are consistent with $\\beta = 2$, though the best-fit values of the emissivity spectral index in both cores are slightly steeper, which would result in lower temperatures and higher masses.  However, with so few measurements of $\\beta$ in starless cores, there is not sufficient information with which to re-examine the dust properties in previously published studies.  Upcoming surveys of the Gould Belt molecular clouds using {\\it Herschel}\\footnote[1]{http://starformation-herschel.iap.fr/gouldbelt} and {\\it SCUBA-2}\\footnote[2]{http://www.jach.hawaii.edu/JCMT/surveys/gb/} promise to make statistically significant measurements of the emissivity spectral index in starless cores."
    },
    "0911/0911.4363_arXiv.txt": {
        "abstract": "{We   present   new  long-baseline   spectro-interferometric observations of the Herbig~Ae star HD~163296 (MWC~275) obtained in the  $H$ and $K$  bands with  the AMBER  instrument at the VLTI.  The observations cover a range  of spatial resolutions between $\\sim$3 and  $\\sim$12~milliarcseconds,  with  a  spectral  resolution  of $\\sim$30. With a total of 1481 visibilities and 432 closure phases, they represent the most comprehensive $(u,v)$ coverage achieved so far for a  young star.  The circumstellar material  is resolved at the sub-AU spatial scale and closure  phase measurements  indicate a small  but significant deviation  from point-symmetry.  We  discuss the  results assuming that the near-infrared excess in  HD~163296 is dominated by the emission of a circumstellar disk.  A  successful fit to the spectral energy distribution, near-infrared visibilities and closure phases is found with a model in which a dominant contribution to the $H$ and $K$ band emission originates in an optically thin, smooth and point-symmetric  region extending  from  about 0.1  to  0.45~AU. At  a distance of 0.45~AU from the star, silicates condense, the disk becomes optically thick and develops a puffed-up rim, whose skewed emission can account for the non-zero closure phases. We discuss the  source of the inner disk  emission and tentatively exclude dense molecular gas as well as optically thin atomic or ionized gas as its possible origin. We propose  instead that the  smooth inner emission is  produced by very refractory grains in a partially cleared region, extending to at least  $\\sim$0.5~AU.  If  so, we may  be observing the  disk of HD~163296 just before it reaches the transition disk phase. However, we  note that the nature of  the refractory grains or, in fact, even the possibility of any grain surviving at the very high  temperatures we require  ($\\sim 2100-2300$ K  at 0.1~AU from the star) is unclear and should be investigated further.} ",
        "introduction": "Herbig~AeBe stars (HAeBe) are intermediate-mass young stars, surrounded by large amounts of dust and  gas. The distribution of this circumstellar material  remains actively  debated.  Various  types  of models  can reproduce the spectral energy distribution (SED) by considering material in geometrically thin accretion disks \\citep{hillenbrand92},       in       a       spherical       envelope \\citep{miroshnichenko97},       a    puffed-up   inner    disk   rim \\citep{dullemond01,    isella05}   or       a    disk    plus a  halo \\citep{vinkovic06}. Fitting the SED alone is therefore highly ambiguous.  Near-infrared (NIR) long baseline interferometry has allowed us to directly probe the properties of matter within the  innermost   astronomical  unit  (AU),   where  key quantities for the star-disk-protoplanets interactions are set.  The first interferometric  studies of  HAeBe  showed that  the NIR  characteristic sizes were larger than expected by classical accretion disk models \\citep{millangabet01}, and were found to be correlated with the stellar luminosity \\citep{monnier02}.  This supports the idea that the NIR emission is dominated  by  the thermal  emission  of  hot  dust   heated  by  stellar radiation. \\cite{natta01} suggested that an inner, optically thin cavity produced by dust sublimation exists inside the disk. At the edge of this region, where dust condensates, the disk is expected to puff up because of the direct illumination from the star \\citep{dullemond01,isella05},  explaining  the  size-luminosity  law derived for Herbig~Ae (and late Be) stars by \\cite{monnier02}. Based  on  a  small  number of  interferometric  observations,  simple geometrical models  were proposed to explain the  global morphology of these    regions   \\citep{millangabet01,eisner04,monnier05,monnier06}. However, when  larger sets of  data became available, it  became clear that the  regions probed by  NIR interferometry are much  more complex and  that   a  deeper   understanding  requires  the   combination  of photometric and multi~wavelength interferometric measurements at the milli-arcsecond resolution and more sophisticated models.  In this  study, we present an  analysis of the  inner disk surrounding the HAe star HD~163296  (MWC275). This isolated Herbig~Ae star is described well by a spectral type of A1, a $\\sim$30~L$_\\odot$ luminosity,        and        a       $\\sim$2.3~M$_\\odot$        mass \\citep{vandenancker98,natta04,montesinos09}.  It is is located at 122$^{+17}_{-13}$~pc and  exhibits a  NIR excess interpreted  as the emission  from   a  circumstellar  disk   \\citep{hillenbrand92}.  A large-scale disk was detected in scattered light \\citep{grady00}, as well as at millimeter  wavelengths \\citep{mannings97}.   This inclined  disk, traced  out to  540~AU,  was found  to  be in  Keplerian rotation,  and probably evolving towards a debris disk phase \\citep{isella07}.  In addition, it  also exhibits  an  asymmetric outflow  on  large scales  ($\\geq$27'') perpendicular  to the  disk, with  a  chain of  six Herbig-Haro  knots (HH409)     that     traces    the     history     of    mass     loss \\citep{devine00,wassell06}. The spectrophotometric  observations that probe  the intermediate and small spatial scales are also compatible with the presence of a disk. \\cite{doucet06} studied  the warm  dust emitting in  the mid-infrared, located in the surface layers  of the intermediate regions of the disk (30-100~AU) and concluded that the emission was consistent with a disk that has little flaring.  This conclusion is consistent with the classification of HD~163296 by  \\cite{meeus01} in their Group~II, whose SED can be explained by assuming that the inner part of the disk shields the outer part  from stellar radiation. In the  innermost regions, the far-UV  emission lines  have  been attributed  to  optically thin  gas accreting onto the stellar surface, a magnetically confined wind, or shocks at the base of the jet \\citep{deleuil05,swartz05}.  Weak X-ray emission (L$_x$/L$_\\star \\sim 5 10^{-6}$) was detected on large scales  and attributed to  the jet \\citep{gunther09}.   In the NIR, the SED time variability was  interpreted to be due to changes in the inner disk  structure, on timescales similar to  the generation of the HH objects \\citep{sitko08}.  With  a disk,  signs  of  accretion, and  a  bipolar outflow,  HD~163296 provides an excellent case study to understand how  circumstellar material is distributed on the sub-AU scale. Its NIR disk was resolved by IOTA, PTI, and Keck-Interferometer \\citep{millangabet01,eisner09, monnier05}, and at milli-arcsecond (mas) resolution with the long CHARA baselines \\citep{tannirkulam} (T08, hereafter). T08 found  that their observations  could not  be  reproduced using models where the majority of the $K$ band emission originates in a dust rim, but that an additional NIR emission inside the dust sublimation  radius could explain the visibilities and the SED.  They interpreted this additional emission as being produced by gas,     as     suggested      for     other     Herbig~AeBe     stars \\citep{eisner07,isella08,kraus08}.  The advent of spectro-interferometry, as provided by the AMBER instrument at VLTI, allows us to simultaneously measure the emission at various  NIR  wavelengths  and  consequently,  to  derive  temperature profiles  for the  emission,  bringing additional  constraints on  its nature.   In this  paper, we  present  an observational  study of  the circumstellar disk around HD~163296 at  the sub-AU scale, using the largest interferometric dataset obtained  for a young star so far.  The paper is organized    as    follows:     in    Sect.~2,    we    describe    the spectro-interferometric observations obtained at AMBER/VLTI and the data processing;  in  Sect.~3,  we  present the  obtained  visibilities  and closure phases.  In Sect.~4, we outline a successful disk model that reproduces all observables and we discuss its physical origin.  We summarize our results in Sect.~5.   \\begin{figure}[t] \\centering \\includegraphics[width=0.4\\textwidth]{./UV_spatialfrequency.ps} \\caption{\\label{uv} $(u,v)$ plane coverage of the observations in spatial frequencies. The observing nights are plotted with different symbols and the  corresponding telescope configurations are  reported in the figure.} \\end{figure}    \\begin{table*}[t] \\centering \\caption{\\label{obs}       Log       of      the       interferometric observations.  The seeing and $\\tau_{0}$ were measured at 650~nm.} % \\begin{tabular}{c c c c c c c c c c} \\hline \\hline Date & Spectral window & Telescope & Baseline & Projected & Position  & Calibrator & Seeing & $\\tau_{0}$ & FINITO \\\\ &     [$\\mu$m]     &    configuration&   name&   length  B$_{\\rm{L}}$~[m]   & angle B$_{\\rm{PA}}$~[$^\\circ$] & name & [''] & [ms] & \\\\ \\hline \\hline 24-05-2008 &[1.67-1.80]-[2.06-2.42] & A0-D0-H0 & A0-D0 & 26.7 & 81.6 & HD156897 & 1.0 & 3.9 & Y \\\\ & & & D0-H0 & 53.5 & 81.6 & & & &  \\\\ & & & A0-H0 & 80.2 & 81.6 & & & &  \\\\ 26-05-2008  & [1.62-1.79]-[2.05-2.31]  & D0-G1-H0  & D0-G1  &  69.7 & 138.2& HD156897& 0.8& 4.4 & Y \\\\ 04-06-2008 & [1.60-1.82]-[2.03-2.33] &  E0-G0-H0 & E0-G0 & 14.3 &56.5 & HD156897 & 0.6& 3.9 & N\\\\ & & & H0-G0 & 28.7 & 56.5& & & & \\\\ & & & E0-H0 & 43.0 & 56.5 & & & &\\\\ 05-06-2008 & [1.61-1.82]-[2.02-2.29] &  E0-G0-H0 & E0-G0 & 15.9 &71.1 & HD156897& 0.8&4.1 & N\\\\ & & & H0-G0 & 31.7 & 71.1& HD163955& & & \\\\ & & & E0-H0 & 47.7 & 71.1& & & & \\\\ 24-06-2008 &[1.65-1.80]-[2.02-2.33]  & U1-U2-U4 & U1-U2  & 55.8 & 33.7 & HD156897 & 1.1& 1.4 & N\\\\ & & & U2-U4 & 82.6 & 89.8 & & & &\\\\ & & & U1-U4 & 122.8 & 67.6 & & & & \\\\ 06-07-2008 & [1.63-1.79]-[2.01-2.45] & D0-G1-H0 & D0-H0 & 60.2 & 70.4& HD160915 & 1.0& 2.3 & N\\\\ & & & D0-G1 & 69.2 & 137.1& & & & \\\\ & & & G1-H0 & 71.4 & 7.6 & & & & \\\\ 08-07-2008  &  [1.60-1.82]-[2.01-2.44]& A0-K0-G1  &  A0-G1  & 81.4  & 121.6& HD160915& 0.8& 2.9 & N\\\\ & & & G1-K0 & 86.5 & 29.9& & & & \\\\ & & & A0-K0 & 108.4 & 74.9& & & & \\\\ 19-07-2007  & [1.61-1.80]-[2.06-2.42]  & A0-K0-G1  & A0-G1  &  82.4 & 124.6& HD156897& 1.0 & 1.6 & N \\\\ & & & G1-K0 & 88.6 & 32.7& & & &\\\\ & & & A0-K0 & 115.2 & 77.6& & & &\\\\ \\hline \\hline \\end{tabular} \\end{table*} ",
        "conclusions": "This paper had discussed the largest set of NIR interferometric data collected so far for a young star.  HD~163296 has a well studied disk at large spatial scale, which motivated our interpretation of the  NIR  interferometry  using  a  star +  inner  disk  model.   Both interferometric  and  photometric data  can  be  accounted  for by  an inclined  disk with  a low  density inner  region.  The  NIR continuum emission is then not dominated by the thermal emission from  the  dusty  disk  rim  located  at  the  sublimation  radius  of astronomical silicates, but by the additional optically thin component located inside. This component emits about 32\\% of the stellar luminosity, but 54\\% (50\\%) of the observed radiation in $H$ ($K$) band.  Because of the silicate condensation  and the strong  increase in opacity,  the disk rim forms in this low-density  region at $\\sim 0.45$~AU from the star and  emits  about 16\\%  (36\\%)  of the  $H$  ($K$)  band flux.   The combination  of  the unresolved  stellar  radiation,  the smooth  and point-symmetric emission of the inner disk region, and the skewed disk rim emission  can successfully explain the  visibilities and non-zero closure phases measured in both bands.  The nature of the emission in the inner disk remains a matter of discussion.  We argue against gas being  mainly responsible for this continuum emission.  A dense, cold disk, as expected for viscous accretion models (Muzerolle et al. 2004), would produce strong molecular lines  that are not seen in high-resolution spectra. Non-LTE tenuous gas layers, in an atomic state if only heating by the  star is assumed, or  fully ionized by  additional sources of energy, cannot account  for the observed  properties of the  NIR continuum. However, we  emphasize that  self-consistent models  of dust-free gaseous disks are  not currently  available, but are  needed to  exploit the full potential of the interferometric observations.  We suggest instead that a small fraction of refractory grains survive very close to the star.  We propose models for the optically thin  emission  of the  innermost  region,  using various kinds of refractory grains,  distributed from  0.10~AU  to 0.45~AU. The dust  surface density provides  only a lower limit  to the gas surface density, as we do know neither the exact nature of the grains nor  their abundance.   However, we find  that a  low density region is consistent with the location and properties of the rim, as condensation of silicates will occur,  and with the lack of molecular features in the  spectrum of HD~163296.  We expect the gas in the inner disk indeed to be mostly atomic, in non-LTE, and although its continuum emission will  be weak, hydrogen  lines can be strong.   The models used in Sect.~4.5.2 to argue  for the presence of refractory grains, make  a  number  of  crude  assumptions,  and  improved  models  that self-consistently compute the grain temperature and emission in  the thin  disk as  well as  the properties  of the  rim  are being developed.   Our  study  indicates   that  the  inner  region  of HD~163296 is quite empty, with a  very low surface density that  is inconsistent with a dense  accretion  disk.   For   comparison,  a  surface  density  of 1~g/cm$^2$ at $\\sim$0.10~AU corresponds, in a standard accretion disk ($\\alpha$=0.01),     to     an     accretion     rate     of     $\\sim 10^{-11}$~M$_{\\odot}$yr$^{-1}$, much lower  than typical values for Herbig~Ae stars  \\citep{garcialopez06}.  With  our data, we  do not constrain the outer radius of this low density region, which can be larger than  0.5~AU. However,  we know that  at large radii  the HD 163296  disk is  massive  and  dense, as  shown  by the  millimeter interferometric observations of \\cite{isella07}. It seems likely that HD~163296 has a dense disk with an inner cavity, and that  we observe it  just before it  reaches the transition disk phase, as suggested by \\cite{sitko08}.  \\vskip 0.1cm  It  is  fair  to  emphasize  that  we  make  no  claim  that  our interpretation is unique. As stated at the beginning of Sect.~4, we assumed that the NIR emission of HD 163296, at spatial scales of less than 0.5~AU, is dominated by the emission of a circumstellar disk. Moreover, we interpreted the non-zero values of the CP as evidence of the  asymmetric emission  of a disk  rim.  While  the properties that we derived for the  smooth inner emission are probably robust, the existence  of the  rim is  less so. In  particular, the  lack of visibility bounces  at large baselines argues  against the presence of a rim.   In this case, the observed  closure phases may possibly be caused by any asymmetric brightness distribution, such as a symmetric flared disk with a stellar contribution that is off-centered by a few percent of the inner disk radius with respect to the disk \\citep{malbet01}, by a hot spot on the disk, or by a density  discontinuity.   Only  a larger  \\textit{(u,v)}  coverage, providing access  to more details  of the morphology could  solve this ambiguity.\\\\  Our interpretation  of the smooth,  inner emission as  originating in refractory grains requires their survival at very high temperatures ($\\sim 2100-2300$ K), much higher than expected, even for the most refractory  grains, at the pressure of  the low-density inner disk \\citep{pollack94,  kama09}.  However, clearing  the innermost regions and optically thin  emission from left-over refractory grains could be  a common phenomenon  in Herbig~Ae stars. Their  presence within the first  few tenths of an  AU is a promising  interpretation of the observed depletions in refractory dust species - such as iron - in   jets   of   young    stars   that   are   launched   from   this region~\\citep{nisini05,  podio06}.  Similar  interferometric studies, with  a large  number  of measurements  in  various wavelength  bands simultaneously, should be performed for  a large sample of stars.  A higher level of complexity in models is  also needed to account for both the dust and  the gas  emission in a  self-consistent way.   Finally, the advent  of  the next-generation of imaging  instruments will  hopefully provide unambiguous constraints on these complex environments."
    },
    "0911/0911.3319_arXiv.txt": {
        "abstract": "{% The surface of the Sun provides us with a unique and very detailed view of turbulent stellar convection. Studying its dynamics can therefore help us make significant progress in stellar convection modelling. Many features of solar surface turbulence like the supergranulation are still poorly understood. }{% The aim of this work is to give new observational constraints on these flows by determining the horizontal scale dependence of the velocity and intensity fields, as represented by their power spectra, and to offer some theoretical guidelines to interpret these spectra. }{% We use long time-series of images taken by the Solar Optical Telescope (SOT) on board the Hinode satellite; we reconstruct both horizontal (by granule tracking) and vertical (by Doppler effect) velocity fields in a field-of-view of $\\sim 75\\times75$~Mm$^2$. The dynamics in the subgranulation range can be investigated with unprecedented precision thanks to the absence of seeing effects and the use of the modulation transfer function of SOT for correcting the spectra. }{% At small subgranulation scales down to 0.4~Mm the spectral density of kinetic energy associated with vertical motions exhibits a $k^{-10/3}$-like power law, while the intensity fluctuation spectrum follows either a $k^{-17/3}$ or a $k^{-3}$-like power law at the two continuum levels investigated (525 and 450~nm respectively). We discuss the possible physical origin of these scalings and interpret the combined presence of $k^{-17/3}$ and  $k^{-10/3}$ power laws for the intensity and vertical velocity as a signature of buoyancy-driven turbulent dynamics in a strongly thermally diffusive regime. In the mesogranulation range and up to a scale of 25~Mm, we find that the amplitude of the vertical velocity field decreases like $\\lambda^{-3/2}$ with the horizontal scale $\\lambda$. This behaviour corresponds to a $k^2$ spectral power law. Still in the 2.5--10~Mm mesoscale range, we find that intensity fluctuations in the blue continuum also follow a $k^2$ power law. In passing we show that granule tracking cannot sample scales below 2.5~Mm.  We finally further confirm the presence of a significant supergranulation energy peak at 30~Mm in the horizontal velocity power spectrum and show that the emergence of a pore erases this spectral peak. We tentatively estimate the scale height of the vertical velocity field in the supergranulation range and find 1~Mm; this value suggests that supergranulation flows are shallow.  }{% } ",
        "introduction": "A complete understanding of thermal convection in stars remains one of the main challenges of present day astrophysics. The Sun offers an unsurpassed detailed view to address this question.  Solar convection is highly non-linear -- typical Reynolds numbers are over $10^{10}$ -- making fully detailed direct numerical simulations down to the viscous dissipation scales unaffordable. In spite of this difficulty, some observed features of solar convection like for instance granulation are rather well understood \\citep[e.g.][]{SN00}. Dynamical features at larger scales like the supergranulation pattern are much less understood. Supergranulation is characterised by a horizontal velocity field spanning scales from 15~Mm to 75~Mm according to the recent work of \\cite{RMRRBP08}.  Despite several attempts, supergranulation has not been identified in large-eddy or direct numerical simulations \\citep[e.g.][]{RLRNS02,RLR05,SGSNB09}. Its origin remains unknown, although some scenarios have been proposed. The classical picture \\citep{SL64} is a linear thermal convection scenario which associates supergranulation cells with  the second ionisation of helium. Other scenarios put forward the idea that supergranulation is a surface phenomenon, as actually indicated by the recent results of local helioseismology \\citep{GB05}. \\cite{RRMR00} suggested that it could be the result of a large-scale instability of surface granular convection. \\cite{RR03b} proposed that some fixed flux boundary condition is imposed by the granulation to the layers just below: with these boundary conditions, the buoyancy destabilises the largest available scales. The associated convective motions have a fairly low intensity contrast \\citep{vdb74}, very much like supergranulation cells \\citep{MRT07}. \\cite{RR03b} further showed that an effective finite but large horizontal convection scale can be obtained in this framework by considering the dynamical effects of a mean magnetic field on the flow. Collective plume interactions \\citep{R03a,CCT07} and travelling-wave sheared \\citep{GK06} convection or magnetoconvection \\citep{GK07} have also been suggested as a possible origin of supergranulation.  Overall the constraints on the large-scale dynamics of solar convection imposed by models and theories are still very loose. Guidance from observations is thus much desired. Particularly if a dynamical connection exists between granulation and supergranulation, some hints of such a scale interaction should be found in the dynamics of intermediate horizontal scales. Flows in the 3--10~Mm range are little known. This range of scales is traditionally referred to as the mesogranulation range after the work of \\cite{NTGS81}, but the very existence of genuine enhanced convective motions at these scales has been much debated \\citep{SDF92,SB97,RRMR00,SSH00}. As pointed out in \\cite{RRMR00}, the main problem with mesogranulation was its way of detection, namely through the measurement of the horizontal divergence of the flows, which is highly sensitive to the way data are reduced. The very nature of the flows at these scales is nonetheless important to answer the question of the origin of supergranulation.  In order to make progress, it is clear that the large-scale side of the granulation spectral peak should be investigated in more detail. Simple models like the mixing-length theory or plume dynamics \\citep{RZ97} do not predict any special spectral feature when the scale (horizontal or vertical) grows. Hydrodynamic numerical simulations like the recent ones of \\cite{SGSNB09} seem to go in the same direction and do not exhibit any spectral feature reminiscent of supergranulation \\cite[see also][]{NSA09}.  However, kinetic energy spectra derived from the radial velocity measured by SOHO/MDI (i.e. from dopplergrams) do show a rise of the spectral kinetic energy density at a scale of 11~Mm and a peak at 36~Mm \\citep{HBBBKPHR00}.  These results are clearly confirmed by the recent measurement of the horizontal surface flows at disc centre with the wide-field high resolution camera CALAS at Pic-du-Midi \\citep{RMRRBP08}. These observations are therefore clearly at odds with the most advanced numerical models of solar surface convection.  \\begin{figure}[t] \\centerline{\\includegraphics[width=\\linewidth]{fig/vel_field29_45mn_m.ps}} \\caption[]{Velocity field derived from the motion of granules during a 45~min window, interpolated by a Daubechies wavelet of 4*953~km. Contours mark the horizontal divergence of the velocity.} \\label{velo} \\end{figure}  \\begin{figure}[t] \\centerline{\\includegraphics[width=\\linewidth]{fig/fluct_spectres_abs.eps}} \\caption[]{Mean spectra and fluctuations: eleven independent horizontal kinetic energy spectra (dotted lines) computed with a three-hours time window, along with their average (thick solid line) are represented. Note that the fluctuations may reach an amplitude of 40\\% compared to the mean value. The straight line on the right indicates the power law $E(k)\\sim k$ characteristic of decorrelated noise. The vertical straight line marks the 2.3~Mm scale. Here $N_x=N_y=80$ and $N_p=1024$. } \\label{mean_spec} \\end{figure}  In this paper we wish to provide accurate observational facts to guide modelling efforts and impose more constraints for theories. For this purpose we use new data sets collected by the Hinode mission and especially by the SOT instrument. Because of their high spatial resolution and their very low noise, these data are appropriate to either refine previous measurements or investigate new properties of the flows. In particular the subgranulation range can now be investigated as never before. We shall focus our attention on the power spectra of the flows because the information that can be obtained from such representations is usually the most robust aspect of these highly turbulent flows.  The paper is organised as follows: after a brief presentation of the data and the techniques to process them (Sect.~\\ref{sectiontwo}), we detail the spectral properties of the measured velocity fields (Sect.~\\ref{sectionthree}) and the intensity field (Sect.~\\ref{sectionfour}).  A discussion and physical interpretation of the results follows (Sect.~\\ref{discussion}). The main results and conclusions are summarised at the end of the paper (Sect.~\\ref{conc}).  \\begin{table*}[t] \\caption[]{The position and amplitude of the supergranulation peak for various data sets.} \\centering \\vspace{1mm} \\begin{tabular}{ccccccc} \\hline Date & Instrument & FOV (Mm$^2$)& Filter & Time slot & $\\lambda_{\\rm SG}$ (Mm) & $E_{\\rm SG}$ (km$^3/s^2$) \\\\ \\multicolumn{7}{c}{ Wide-field data} \\\\ 1996 May-Jul & SOHO/MDI & global & Ni I, 676.8~nm & 62d & 36.4 & (200)\\\\ 2000 Apr 22& TRACE & 225x250 & WL & 6x3 h &  35.6 & 200 \\\\ 2007 Apr 24& TRACE & 238x275 & WL & 2x3.4 h &  29.8 & 490 \\\\ 2007 Mar 13& CALAS & 290x216 & 575~nm & 2x3 h & 36.4 & 350 \\\\ \\multicolumn{7}{c}{Small-field data} \\\\ 2007 Mar 8& Hinode/SOT& 75.9x76.8  & G-band &  4h   & no peak   & 220 \\\\ 2007 Mar 15& Hinode/SOT& 77.3x33.9 & G-band &  3x3h & 57/17-27  & 240/350  \\\\ 2007 Aug 29-30& Hinode/SOT & 76.6x76.6 & blue & 11x3h & 28.7 & 600 \\\\ 2007 Oct 1 & Hinode/SOT & 2x(76x76) & blue & 2x3.8h & 27/31 & 460/545 \\\\ \\hline \\end{tabular} \\label{SG_table} \\end{table*}   \\begin{figure*}[ht] \\centerline{\\includegraphics[width=0.5\\linewidth]{fig/many_spectres_abs256.eps} \\includegraphics[width=0.5\\linewidth]{fig/scale_to.eps}} \\caption[]{Left: kinetic energy spectra of the horizontal velocity field for various time windows. The vertical line marks the 2.3~Mm scale below which noise dominates. The black dots mark the scale at which the spectrum of a given time average disconnects from the 30-min one. Right: relationship between the length scale and the correlation time scale for the horizontal velocity field. Here $N_x=N_y=80$ and $N_p=256$.} \\label{KES} \\end{figure*} ",
        "conclusions": "}  The power spectrum of surface flows at the Sun's surface has been investigated using mainly image series from the Hinode/SOT instrument. The results may be summarised as follows.  \\begin{enumerate} \\item The smallest scales, i.e. the subgranulation range, can be investigated for the first time without a seeing effect and using corrections of the MTF of the instrument. On the one hand, the spectral cut-off of intensity fluctuations is similar to a $k^{-17/3}$ power law for images taken with a narrow filter (width $\\sim9$~pm) in the continuum on the red side of the 525~nm FeI line and to $k^{-3}$ for images taken with the broadband filter (width $\\sim220$~pm) around 450.45~nm.  On the other hand, the vertical velocity horizontal spectrum measured 80~km or 190~km above the continuum at 557.6~nm shows a steep cut-off in $k^{-10/3}$, which can be related to the  $k^{-17/3}$ power law by invoking a balance between buoyancy and inertia. The values of the two scaling exponents suggest that the temperature fluctuations impact the dynamics at a scale well below the one determined by the unit P\\'eclet number.  Finally, we also note the decrease of the granulation peak amplitude with altitude, possibly betraying the deceleration of the upflows.  \\item The granulation spectral peak at 1.7~Mm is conspicuous in the vertical velocity spectrum. The location of the intensity spectrum peak depends on the wavelength: it is close to 1.5~Mm in the blue continuum at 450~nm, 2~Mm in the continuum near 525~nm, and to 2.5~Mm in the core of the FeI line at 525.0~nm, which probes the dynamics in a layer 200~km above the continuum.  \\item Moving to the large scales, we note that the spectral density of vertical kinetic energy decreases as $k^2$ with decreasing $k$ (equivalently, the amplitude of $v_z$ decreases as $\\lambda^{-3/2}$ with the horizontal scale $\\lambda$). This $k^2$ power law also emerges from the relationship between the scale and the correlation time of the horizontal flows.  Furthermore we note that the spectral density of intensity fluctuations in the blue continuum also scales like $k^2$ in the 2.5--10~Mm range.  \\item The power spectrum of horizontal flows has been explored in the 2.5--40~Mm range.  It clearly reveals the supergranulation peak as well as the large-scale side of the granulation peak. These observations confirm in passing the result of the numerical simulations of \\cite{RRLNS01} that granule tracking cannot determine plasma flows at scales below $\\sim$2.5~Mm. The energy density minimum is found at a scale of 12~Mm. These results confirm with a completely different technique the spherical harmonics spectra obtained by \\cite{HBBBKPHR00} from dopplergrams. Focusing on the supergranulation peak, we could estimate its amplitude to be in the range of 200--700~km$^3/$s$^2$ and show its strong variability. From Figs.~\\ref{KES} and \\ref{pore} we may add that the visible part of the granulation peak points to an amplitude in the range of 200--400~km$^3/$s$^2$. Interestingly enough, we discovered that the emergence of a pair of pores is able to erase the supergranulation peak completely, which points out the sensitivity of the supergranulation amplitude flow to the ambient magnetic field. Finally, combining all the data, we tentatively estimated the scale-height of the vertical velocity field at the supergranulation horizontal scale. We found a scale height of approximately 1~Mm, indicating that supergranules are likely surface flows. However, this estimate should be verified for a number of supergranules to obtain a statistically robust value.  \\end{enumerate}  These observational results provide a set of landmarks that should be recovered by future numerical simulations. Future observational work will focus on the large-scale side of the supergranulation peak to determine what scaling law, if any, controls the very large-scale dynamics of the Sun between supergranulation and differential rotation. Another point deserving a detailed investigation is the dynamical influence of the magnetic fields on the surface flows and their power spectrum."
    },
    "0911/0911.3313_arXiv.txt": {
        "abstract": "We present 2D high-resolution hydrodynamic simulations of the relativistic outflows of long-duration gamma-ray burst progenitors. We analyze the properties of the outflows at wide off-axis angles, produced by the expansion of the hot cocoon that surrounds the jet inside the progenitor star. We find that the cocoon emission at wide angles may have properties similar to those of the subclass of short-duration gamma-ray bursts with persistent X-ray emission. We compute the predicted duration distribution, redshift distribution, and afterglow brightness and we find that they are all in agreement with the observed properties of short GRBs with persistent emission. We suggest that a SN component, the properties of the host galaxies, and late afterglow observations can be used as a crucial test to verify this model. ",
        "introduction": "Gamma-ray bursts (GRBs) are divided in two classes based on their duration and on the hardness of their prompt emission spectra (Mazets et al. 1981; Kouveliotou et al. 1993). Long-duration bursts (LGRBs) were the first class for which afterglow observations were possible (Costa et al. 1997). Afterglow observations have allowed us to study them in great detail. We now know that long GRBs are associated with the deaths of massive compact progenitor stars (Woosley 1993; Stanek et al. 2003; Hjorth et al. 2003), take place in compact star-forming galaxies, and can be seen up to redshift $z=8.2$ (GRB~090423, Tanvir et al. 2009; Salvaterra et al. 2009).  Much less is known about the class of short-duration bursts (SGRBs) and their progenitors. It has been suspected for a long time that SGRB progenitors are binary systems of two compact objects that merge after their orbital energy has been radiated, mostly in gravitational waves (Eichler et al. 1989; Lee \\& Ramirez-Ruiz 2007; Ramirez-Ruiz \\& Lee 2009). With the advent of HETE2 and Swift, afterglow observations of SGRBs have become possible (Fox et al. 2005; Gehrels et al. 2005; Hjorth et al. 2005; Villasenor et al. 2005). The localization of SGRBs has confirmed that some are associated with early type galaxies and explode at relatively large distances from the center of the host galaxy, confirming their origin as due to the merger of compact objects (Nakar et al. 2007; Berger 2009). However, not all SGRBs fall within such a simple scheme. Some short bursts, especially those showing a plateau of persistent prompt emission in the X-rays (Lazzati et al. 2001; Villasenor et al. 2005) are associated with small, star forming galaxies, and explode close to the center of their hosts (Troja et al. 2008). This led to the speculation that short bursts have a more diverse family of progenitors than long bursts, possibly involving different compact objects (neutron stars vs. black holes, Belczynski et al. 2006; Troja et al. 2008) or transient proto-magnetars (Metzger et al. 2008). Virgili et al. (2009) studied GRB populations based on a new Type I/Type II classification (see Zhang et al. 2007, 2009) and came to the conclusion that most SGRBs closely follow star formation and cannot therefore originate from a traditional merger path. They suggest that a large fraction of SGRBs may be associated with massive stars, analogously to LGRBs. A close association between the progenitors of SGRBs and LGRBs was theoretically explored within the multi-jet scenario (Toma et al. 2005ab) and in the cannonball scenario (Dado et al. 2009).  In this letter, we show that LGRB progenitors may produce SGRBs at large off-axis angles. This letter is organized as follows: in \\S~2 we present our numerical setup, in \\S~3 we describe the properties of the relativistic outflow at large off-axis angles, and in \\S~4 we discuss some tests that we have performed to assure that this model is consistent with observations of SGRB rates and their afterglows. Our results are summarized and discussed in \\S~5. ",
        "conclusions": "We presented the results of numerical simulations of long GRB outflows from massive progenitor stars within the collapsar scenario. In this letter, we focused on the outflow at large off-axis angles, showing that it has a very thin structure and that it can lead to emission with the characteristics of a SGRB. The presence of slower material behind the leading edge of the outflow suggests that the emission at lower frequencies may last longer than the prompt $\\gamma$-ray emission, a characteristic that has been observed in the subclass of SGRBs with persistent X-ray emission. We propose that at least some SGRBs with persistent emission are due to collapsars seen at wide angles, $40\\degr$ to $50\\degr$ away from the axis along which a LGRB is released. We check this possibility by computing the number, the duration distribution, and the redshift distribution of such a short GRB population, finding that the predictions are in good qualitative agreement with observed quantities. A better agreement is prevented by the fact that our simulations do not incorporate the diversity in progenitor and engine populations of LGRBs. Diversity in the properties of the GRB progenitor may also cause a small fraction of short GRBs due to on-axis collapsars with a short-duration engine (Janiuk \\& Proga 2008; Proga et l. 2009). There are several ways in which the validity of this model can be checked against observations. The smoking gun of a SGRB from an off-axis collapsar would be the detection of an accompanying SN explosion in the late afterglow, as observed in LGRBs (Stanek et al. 2003; Hjorth et al. 2003). Such accompanying SN may however be fairly different from those accompanying LGRBs, due to the difference in the observing angle. In particular, SNe associated with off-axis collapsars should be characterized by slower expansion velocities towards the observer, a less luminous peak (due to the fact that the $^{56}$Ni is ejected preferentially along the axis, Maeda et al. 2006), and a longer time to reach maximum. All these characteristic make the SN detection more difficult, and could explain the unsuccessful searches of SN components in several SGRBs\\footnote{The other explanation being the fact that these bursts are not from off-axis collapsars, given their low redshift.} (GRB~060505 and GRB~060614, Della Valle et al. 2006; Fynbo et al. 2006; Gal-Yam et al. 2006). Second, if our model is correct, the host galaxies of some SGRBs should have properties fairly similar to those of LGRBs. The comparison of host galaxies of SGRBs and LGRBs shows marked differences (Berger 2009; Fong et al. 2009), especially at low redshift. At $z\\sim1$, the host galaxies of the two populations become more similar (see, e.g., Figure 3 in Berger 2009), as our model would predict. Given the small number of galaxies involved, it is premature to give any firm conclusion, and only further studies with more extensive samples can give better constraints. Along similar lines, Nysewander et al. (2009) found that the density of the ambient medium of LGRB and SGRB afterglows are similar. Finally, SGRBs from off-axis collapsars may be identified through their afterglow properties, especially if they explode in a uniform environment (Figure~\\ref{fig:aglow}).  The simulation that we have presented is just one case of a big family of possibilities, with diverse progenitors and diverse engine properties. The properties of the SGRBs at large off-axis angles are likely to depend mostly on the luminosity of the jet, with more energetic events from more luminous jets. The volume of the cocoon inside the progenitor star affects instead the temperature and Lorentz factor of the cocoon outflow. Harder SGRBs are therefore expected from more compact stars. Lacking  a definite radiative model, however, it is difficult to clearly connect the Lorentz factor of the outflow to a peak frequency of the spectrum. A definitive measurement of a high Lorentz factor ($\\Gamma>100$) for a SGRB would be inconsistent with this model, since the cocoon shell can hardly have Lorentz factors $\\Gamma>20$. Finally, short bursts from off-axis collapsars should have little internal variability, unless some localized dissipation process, such as magnetic reconnection (Lyutikov \\& Blandford 2003) of relativistic turbulence (Narayan \\& Kumar 2009), provides the internal energy for the radiation of photons.  In conclusion, we would like to stress that what we presented is a set of necessary conditions for the production of SGRBs. These are by no means sufficient and at least two important assumptions have to be made to successfully predict SGRBs from our simulation. First, we have to assume that all the radiation is released at a fairly constant and small radius. Should this fail, the light curve would be smeared on a longer time and the burst would not be classified as a SGRB. Second, we have to postulate that the unknown radiation mechanism is able to produce photons in the MeV range, even if it expands at a much slower Lorentz factor compared to the on-axis jet. Should this fail, the cocoon emission would likely produce an x-ray transient rather than a SGRB."
    },
    "0911/0911.3914_arXiv.txt": {
        "abstract": "{} {We present a catalog of 213 type--2 AGN selected from the zCOSMOS survey. The selected sample covers a wide redshift range (0.15$<z<$0.92) and is deeper than any other previous study, encompassing the luminosity range $10^{5.5}$L$_{\\odot}<$L$_{\\rm [OIII]}<$ 10$^{9.1}$ L$_{\\odot}$. We explore the intrinsic properties of these AGN and the relation to their X-ray emission (derived from the XMM-COSMOS observations). We study their evolution by computing the \\lOIIIb\\ line luminosity function (LF) and we constrain the fraction of obscured AGN as a function of luminosity and redshift.} {The sample was selected on the basis of  the optical emission line ratios, after applying a cut to the signal-to-noise ratio (S/N) of the relevant lines. We used the standard diagnostic diagrams (\\OIIIb/\\Hb\\ versus \\NII/\\Ha\\ and \\OIIIb/\\Hb\\ versus \\SII/\\Ha) to isolate AGN in the redshift range 0.15$<z<$0.45 and the diagnostic diagram \\OIIIb/\\Hb\\ versus \\OII/\\Hb\\ to extend the selection to higher redshift (0.5$<z<$0.92).} {Combining our sample with one drawn from SDSS, we found that the best description of the evolution of type--2 AGN is a luminosity-dependent density evolution model. Moreover, using the type--1 AGN LF we were able to constrain the fraction of type--2 AGN to the total (type--1 + type--2) AGN population. We found that the type--2 fraction decreases with luminosity, in agreement with the most recent results, and shows signs of a slight increase with redshift. However, the trend with luminosity is visible only after combining the SDSS+zCOSMOS samples. From the COSMOS data points alone, the type--2 fraction seems to be quite constant with luminosity.} {} ",
        "introduction": "According to the standard unified model \\citep[e.g.;][]{Antonucci1993}, AGN can be broadly classified into two categories depending on whether the central black hole and its associated continuum and broad emission-line region are viewed directly (type--1 AGN) or are obscured by a dusty circumnuclear medium (type--2 AGN). Type--1 AGN are characterized by power-law continuum emission, broad permitted emission lines ($\\gsim$1000 km s$^{-1}$) and are thus easily recognizable from their spectra. In contrast, type--2 AGN have narrow permitted and forbidden lines ($\\lsim$1000 km s$^{-1}$) and their stellar continuum, often dominated by stellar emission, is similar to normal star-forming galaxies (SFGs). The main difference between AGN and SFGs is the ionizing source responsible for their emission lines: non-thermal continuum from an accretion disc around a black hole for AGN or photoionization by hot massive stars for normal SFGs.  To identify type--2 AGN, we thus need to determine the ionizing source. \\citet{Baldwin1981} demonstrated how this is possible by considering the intensity ratios of two pairs of relatively strong emission lines. In particular, they proposed a number of diagnostic diagrams (hereafter BPT diagrams), which were further refined by \\citet{Veilleux1987}, based on \\lOIIIb, \\lOI, \\lNII, \\lSII, \\lHa\\ and \\lHb\\ emission lines, where \\Ha\\ and \\Hb\\ refer only to the narrow component of the line. The main virtues of this technique, illustrated in Fig. \\ref{fig:red}, are: 1) the lines are relatively strong, 2) the line ratios are relatively insensitive to reddening corrections because of their close separation, and 3) at least at low redshift (z $\\lesssim$ 0.5) the lines are accessible using ground-based optical telescopes. Several samples have been selected in the past using the BPT diagrams and the method select AGN reliably with high completeness \\citep{Zavadsky2000,Zakamska2003,Hao2005ss}.   At high redshift, however, the involved lines are redshifted out of the observed optical range and the classical BPT diagrams can no longer be used. In these circumstances, it is thus desirable to devise a classification system that is based only on the blue part of the spectrum.  For this reason, \\citet{Rola1997}, \\citet{Lamareille2004}, and \\citet{PerezMontero2007} proposed alternative diagrams based on the strong lines \\OII, \\NeIII, \\Hb, and \\OIIIb, which provide moderately effective discrimination between starbursts and AGN. Since this technique is more recent than classical BPT diagrams, it has been used by fewer studies in the literature. We also note that, the use of the ratio of two lines that are not close to each other in wavelength (\\lHb\\  and \\lOII) makes this diagram sensitive to reddening effects which, due to differential extinction of the emission lines and the stellar continuum \\citep{Calzetti1994}, also affect the EW measurements.  An important issue to address in AGN studies is their evolution. The overall optical luminosity function of AGN, as well as that of different types of AGN, holds important clues about the demographics of the AGN population, which in turn provides strong constraints on physical models and theories of AGN and galaxy co-evolution.  Many studies have been conducted and many results obtained in the past few years to constrain the optical luminosity function of  type--1 AGN at both low \\citep{Boyle1988,Hewett1991,Pei1995,Boyle2000,Croom2004} and high redshift \\citep{Warren1994, Kennefick1995,Schmidt1995,Fan2001,Wolf2003,Hunt2004,Bongiorno2007,Croom2009}. In contrast, there are not many type--2 AGN samples available in the literature and consequently very few studies of their evolution have been conducted.  In the local Universe, \\citet{Huchra1992} selected 25 Seyfert--1 and 23 Seyfert--2 galaxies from the CfA redshift survey \\citep{Huchra1983} and used these AGN to measure their luminosity function. \\citet{Ulvestad2001} also computed the local luminosity function of a sample selected from the Revised Shapley-Ames Catalog \\citep{Sandage1981}, and using the BPT diagrams, \\citet{Hao2005lf} derived the luminosity function of a sample selected from the SDSS at z $<0.13$.\\\\ The only sample that spans a relatively wide redshift range, from the local Universe up to z$\\sim$0.83, is that selected by \\citet[][hereafter R08]{Reyes2008} from the SDSS sample, which is however limited to bright objects ($10^{8.3}$ L$_{\\odot} <$ L$_{\\rm [OIII]} <$ 10$^{10} L_{\\odot}$). Thus, a sample of type--2 AGN encompassing a wide redshift interval and including lower luminosity objects is highly desirable.   The zCOSMOS survey \\citep{Lilly2007,Lilly2009} is a large redshift survey in the COSMOS field. From this sample, using the standard BPT diagrams at low redshift and the diagram from \\citet{Lamareille2004} at high redshift, we selected a sample of 213 type--2 AGN in a wide redshift range (0.15$<z<$0.92) and luminosity range ($10^{5.5} \\rm L_{\\odot} < L_{\\rm [OIII]} < 10^{9.1} \\rm L_{\\odot}$). Here we present the main properties of this sample, their \\OIIIb\\ line luminosity function, and the derived type--2 AGN fraction as a function of luminosity and redshift.  The paper is organized as follows: Sect. \\ref{sec:cosmos} presents a brief overview of the COSMOS project and in particular of the zCOSMOS sample, while in Sect. \\ref{sec:sample_selection} we describe in detail the adopted method to select the sample. Sections \\ref{sec:Opt_comp} and \\ref{sec:Xcomp} compare our sample with both other optical samples and with the X-ray selected sample in the same field (XMM-COSMOS; \\citealt{Hasinger2007,Cappelluti2009,Brusa2007}; Brusa et al., in prep) respectively. Finally, in Sect. \\ref{sec:lf_method}  we derive our emission-line AGN luminosity function, and in Sect. \\ref{sec:lf_result}, we compare the results with those in previous works and the derived evolutionary model, as well as the type--2 AGN fraction as a function of luminosity and redshift. Finally, Sect. \\ref{sec:conclusion} summarizes our work.  Throughout this paper, we use AB magnitudes and assume a cosmology with $\\Omega_{\\rm m}$ = 0.3, $\\Omega_{\\Lambda}$ = 0.7 and H$_{0}$ = 70 km s$^{-1}$ Mpc$^{-1}$. ",
        "conclusions": "\\label{sec:conclusion}  We have presented the faintest optically selected sample of type--2 AGN to date. The sample, selected from the zCOSMOS survey, consists of 213 sources in the redshift range 0.15 $<z<$0.92, spanning the \\OIIIb\\ luminosity range $10^{5.5}$L$_{\\odot}<$L$_{\\rm \\OIIIb}<$10$^{9.1}$ L$_{\\odot}$. The only other sample at z$>$ 0.15 is one selected in a similar way from the SDSS by \\citet{Zakamska2003}, which, however, covers significantly brighter luminosities ($10^{7.3}$L$_{\\odot}<$L$_{\\rm \\OIIIb}<$10$^{10.1}$L$_{\\odot}$).  Our sample has been selected using the first 10,000 spectra of the zCOSMOS dataset on the basis of their emission line properties. In particular, we used the standard BPT diagrams (\\OIIIb/\\Hb\\ versus \\NII/\\Ha\\ and \\OIIIb/\\Hb\\ versus \\SII/\\Ha) to isolate AGN in the redshift range 0.15$<z<$0.45 and the more recent diagnostic diagram \\OIIIb/\\Hb\\ versus \\OII/\\Hb\\ to extend the selection to higher redshift (0.5$<z<$0.92), after applying a cut to the S/N ratio of the involved lines.  Cross-checking the zCOSMOS emission-line sample with the XMM-COSMOS catalog, we found a significant incompleteness in the \\blue\\ diagnostic diagram used to select high redshift type--2 AGN (z$>$0.5). Our hypothesis is that LINERs as well as composite AGN/SF sources can be misclassified by this diagram.  For the selected type-2 AGN sample, we computed the \\lOIIIb\\ line luminosity function using the 1/V$_{\\rm max}$ method. The selection function takes into account (and corrects for) the sources that were not spectroscopically observed (but were in the photometric catalog) and those for which a secure spectroscopic identification has not been obtained. The correction is performed using a statistical weight associated with each galaxy that has secure redshift measurement. Since the sample was selected from a magnitude-limited sample, applying a criterion based on the line fluxes, the maximum volume V$_{\\rm max}$ for each object has been estimated as the minimum between the volume $V_{max}(m_i)$ associated with the maximum apparent magnitude and the volumes $V_{max}(f_{i,l})$ associated with the minimum flux of the used lines. We have extended the \\lOIIIb\\ LF to luminosities about 2 orders of magnitude fainter than those previously available (down to L$_{\\rm \\OIIIb}$=10$^{5.5} L_{\\odot}$).  To enlarge the luminosity range, we combined our faint zCOSMOS sample with the sample of bright type--2 AGN from the SDSS (R08) and found that the evolutionary model that represents the combined luminosity functions most accurately is an LDDE model. By fixing the evolutionary parameters (p1, p2, $\\alpha$, z$_{\\rm c,0}$, L$_{a}$) using the results obtained by DC08, we obtained as best-fit model parameters $\\gamma1$=0.56, $\\gamma2$=2.42, and L$^*$=2.7$\\times$10$^{41}$erg s$^{-1}$ with the normalization $\\Phi^*_L$=1.08$\\times$10$^{-5}$Mpc$^{-3}$.     Finally, by comparing the LF for type--2 and type--1 AGN (obtained by converting the broad-band LF from the VVDS into an \\OIIIb\\ LF), we constrained the type--2 quasar fraction as a function of luminosity. We found that the fraction of type--2 AGN is high at low \\OIIIb\\ luminosities and then decreases at higher luminosities, in agreement with that found from several studies in the X-ray band \\citep{Ueda2003,Barger2005,Treister2006a,Sazonov2007,Gilli2007}. The same decreasing trend with luminosity is found in all individual redshift bins, and we found on average a slightly higher fraction towards higher redshift (consistent however with being constant inside the errors). In particular, we found that at 0.15$<z<$0.3 the fraction of type--2 AGN decreases with luminosity from $\\sim$65\\% to $\\sim$50\\%, while at 0.3$<z<$0.45, and 0.5$<z<$0.92 the type--2 fraction ranges from $\\sim$80\\% to $\\sim$25\\%. However, analysis of a sample uniformly selected across a wider luminosity range is needed to confirm these results, which are derived by combining two different samples (zCOSMOS and SDSS). We can not exclude that the observed trend could still be an artifact produced by the different selection functions in the two samples. From the COSMOS data points alone, the type--2 fraction seems to be quite constant with luminosity."
    },
    "0911/0911.3852_arXiv.txt": {
        "abstract": "Cosmological observables are used to construct cosmological models.  Since cosmological observations are limited to the light cone, a fixed number of observables (even measured to arbitrary accuracy)  may not uniquely determine a cosmological model without additional assumptions or considerations.  A prescription for constructing a spherically symmetric, inhomogeneous cosmological model that exactly reproduces the luminosity-distance as a function of redshift and the light-cone mass density as a function of redshift of a $\\Lambda$CDM model is employed to gain insight into how an inhomogeneous cosmological model might mimic dark energy models. ",
        "introduction": "}  In the last decade or so remarkable progress has been made in measuring cosmological parameters to unprecedented accuracy. Parameters such as the Hubble constant, $H_0=72\\pm 8 \\textrm{ km s}^{-1}\\textrm{ Mpc}^{-1}$ \\cite{Hubblekey}, the temperature of the cosmic background radiation (CBR), $T_0=2.728\\pm 0.004\\ \\textrm{K}$ \\cite{COBET}, and many other parameters, are now known to impressive precision.  Of course, we don't invest so much time and effort determining cosmological parameters because of an interest in numerology, but rather, because the parameters are necessary inputs for the task of constructing a standard cosmological model.  In turn, we do not construct cosmological models simply to make predictions to compare with observations, but rather, to guide us to, in the words of Einstein \\cite{Einstein}, ``a deeper and more consistent comprehension of the physical and astronomical facts.''  The precision cosmological measurements have lead to the latest cosmological model, usually called the \\textit{standard cosmological model,} the \\textit{concordance model}, or simply $\\Lambda$CDM, where $\\Lambda$ indicates the inclusion of  Einstein's cosmological constant (or more generally, dark energy), and CDM stands for cold dark matter. The model is but one possible realization of homogeneous and isotropic cosmological solutions to Einstein's equations. (We will refer to homogeneous/isotropic models as Friedmann--Lema\\\"{i}tre--Robertson--Walker (FLRW) models.)  The question yet to be answered is whether today's standard model will lead us to a deeper and more consistent comprehension of the physical and astronomical facts.  A troubling feature of the $\\Lambda$CDM model is the present composition of the universe: $95\\%$ is dark!  Of the dark components, roughly $25\\%$ of the total mass-energy density is dark matter, associated with galaxies, clusters of galaxies, and other bound structures.  The bulk of the mass-energy of the universe, about $70\\%$ in the standard $\\Lambda$CDM model, is attributed to dark energy, which is usually said to drive an accelerated expansion of the Universe.  The observations and phenomena that lead to the assumptions of dark matter and dark energy are real, but that does not necessarily imply that dark matter and dark energy as usually envisioned are real.  Dark matter and dark energy are good names for the phenomena.  But naming is not explaining!  In some sense, an even more remarkable feature of the standard cosmological model is that it seems capable of accounting for all cosmological observations; \\textit{i.e.,} it seems to work!  Here, we address the issue of whether the agreement with observations prejudices our judgment regarding the likelihood that the $\\Lambda$CDM model is the final answer. In particular, we will discuss the possibility that dark energy, \\textit{per se}, does not exist.  The most important point regarding dark energy (and also, for that matter, the acceleration of the expansion of the universe), which is central to the motivation for this paper, is that all evidence for dark energy is \\textit{indirect.} The only effect of dark energy is on the expansion history of the Universe.  The expansion rate of the Universe is perhaps the most fundamental quantity in cosmology.  In the standard cosmological model the expansion rate is determined from the $00$-component of the Einstein equations, $G_{00}=\\kappa T_{00}$, where $\\kappa = 8\\pi G$.  In FLRW models, $G_{00}=3H^2 + k/a^2$, where $H=\\dot{a}/a$ is the expansion rate (here, $a$ is the Robertson-Walker scale factor and $k=\\pm1$ or $0$, depending on the geometry).  With the assumption of a perfect-fluid stress tensor, $T_{00}=\\rho$, where $\\rho$ is the mass density.  The stress tensor is assumed to consist of several components with different equations of state.  For any component $i$, the equation of state for that component is defined as $w_i=p_i/\\rho_i$.  If $w_i$ is constant, the energy density in component $i$ evolves with redshift $z\\equiv a_0/a-1$  as $\\rho_i(z) = \\rho_i(0)\\left(1+z\\right)^{3(1+w_i)}$, where $a_0$ is the present value of the scale factor.   Then the expansion rate as a function of redshift can be expressed as \\begin{equation} H^2(z)=H_0^2\\left\\{ \\Omega_\\Lambda \\exp\\left[\\int_0^z\\frac{dz_1}{1+z_1} 3 \\left[1+w_\\Lambda(z_1)\\right] \\right] + \\Omega_k(1+z)^2 + \\Omega_M(1+z)^3 + \\Omega_R(1+z)^4  \\right\\} . \\label{hsquared} \\end{equation} Here, $\\Omega_i$ is the ratio of the present energy density in component $i$ compared to the critical density $\\rho_C=3H_0^2/\\kappa$.  The subscript ``$k$'' indicates a spatial-curvature contribution ($w_k=-1/3)$, ``$M$'' indicates a matter component ($w_M=0$), ``$R$'' indicates a radiation component ($w_R=1/3$),  and ``$\\Lambda$'' represents dark energy.    In the event that the function  $w_\\Lambda(z)$ in the last term is independent of $z$, the last term in $H^2(z)$ becomes $\\Omega_\\Lambda(1+z)^{3(1+w_\\Lambda)}$.  The equation of state for a cosmological constant is $w_\\Lambda=-1$.  It appears that the expansion rate as a function of $z$ depends on four constants $\\{\\Omega_k,\\Omega_{M},\\Omega_R,\\Omega_\\Lambda\\}$ and one function $w_\\Lambda(z)$.  In reality, the situation is much more predictive.  First, since $H^2(0)=H_0^2$, it follows that $\\sum_i\\Omega_i=1$, removing one constant.  The value of $\\Omega_k$ is well determined by WMAP to be small, $\\Omega_k=-0.0026^{+0.0066}_{-0.0064}$ at the $68\\%$ confidence level \\cite{wmap5}, effectively removing yet another constant. The value of $\\Omega_R$ is also well determined by CBR measurements ($5\\times 10^{-5}$) removing a further constant.  One may then make the less secure assumption that $\\Omega_M$ is ``known,'' or more profitably, use observational constraints to marginalize over the $\\Omega_i$'s, leaving only the function $w_\\Lambda(z)$ undetermined. The function must be determined through its effect on the expansion history of the Universe, which is manifest through various cosmological observables.  Many cosmological observables depend on the expansion history of the Universe.  They often depend on the expansion history through the coordinate distance $r$ of a source of redshift $z$.  Recall that the Robertson-Walker metric can be written in the form \\begin{equation} ds^2 = dt^2 - a^2(t)\\left[\\frac{dr^2}{1-kr^2}+r^2d\\Omega^2\\right] , \\label{rw} \\end{equation} where here $d\\Omega$ is the angular differential and $r$ is the comoving ``radial'' coordinate.  The radial coordinate of a source at redshift $z$ is determined by integrating the null geodesic equation ($ds^2=0$) to obtain \\begin{equation} r(z) = \\left. \\begin{array}{l} \\sin \\\\ 1 \\\\ \\sinh \\end{array} \\right\\} \\left[ \\int_0^z \\frac{dz'}{H(z')} \\right] , \\label{rofz} \\end{equation} where $\\sin$, $1$, $\\sinh$ obtains for $k=+1,\\ 0, \\-1$, respectively.  Observables such as the luminosity distance, the angular-diameter distance, the volume element, and the age of the Universe, all depend on the time evolution of the expansion rate, hence on the properties of the stress tensor, hence on the dark energy fluid characterized by $w_\\Lambda$.  At present, a constant $w_\\Lambda(z)$ seems to fit the data, and furthermore $w_\\Lambda=-1$, the value of $w_\\Lambda$ for a cosmological constant, is not disallowed.  It is also true that since there is no apparent reason to exclude vacuum energy from the stress tensor, there is no reason \\textit{not} to imagine it enters the determination of the expansion rate.  The problem is that the magnitude of the  expected vacuum energy is very many orders of magnitude larger than expected. If the cosmological constant is part of nature, then either some principle must be found to explain the much smaller-than-expected magnitude of $\\Omega_\\Lambda$, or one must resort to anthropic arguments.  This situation has led many to consider alternate explanations of the observations: modified gravity, new long-range forces, new ultra-light scalar fields, extra dimensions, branes, bulk, Lorentz-invariance violation, \\textit{etc.}   Some of these approaches are, in many respects, more drastic than a small cosmological constant.  For instance, theories of modified gravity typically do not respect the principle of general covariance.  Since we should not abandon our principles without good cause, it is tempting to retreat to the position that since a cosmological constant is adequate to account for the data, just assume that it exists and call it a day?  We resist the retreat because our goal in cosmology is more than just explaining the observations.  The ingredients of a cosmological model must be deeply grounded in fundamental physics.  Dark matter, dark energy, modified gravity, mysterious new forces and particles, {\\em etc.,} unless part of an overarching model of nature, should not be part of a cosmological model.  We may plant new ideas and features into cosmological models, but they will ultimately prove barren unless grounded in fundamental physics.  No matter how well a concordance model fits the data, its ingredients must not be discordant with fundamental physics.  In this paper we explore a less traveled path, and consider the possibility that the Friedmann equation is not an adequate description of our inhomogeneous Universe \\cite{Celerier:1999hp,Buchert:1999mc,Schwarz:2002ba,Rasanen:2003fy,Kolb:2004am,Rasanen:2004js,Kolb:2005me,Rasanen:2005zy,Barausse:2005nf,Notari:2005xk,Rasanen:2006kp}. In this approach there is no need for dark energy, so no need for new long-range forces, modifications of general relativity, new ultralight particles, or anthropic reasoning.  The phenomenon of inhomogeneities mimicking dark energy is sometimes referred to as ``backreaction.''  In this paper we follow the prescription of Mustapha, Hellaby, and Ellis \\cite{Mustapha:1998jb} [MHE], recently applied to $\\Lambda$CDM models by C\\'{e}l\\'{e}rier, Bolejko, Krasinski, and Hellaby \\cite{Celerier:2009sv},\\footnote{Similar formalism was developed by Chung and Romano \\cite{Chung:2006xh} to fit only the luminosity-distance, and by Yoo, Kai, and Nakao \\cite{Yoo:2008su} who simultaneous fit the luminosity distance and require no spatial variations of the age of the Universe.} and explore the implications of the fact that one can construct a spherically symmetric inhomogeneous model that \\textit{exactly} reproduces the redshift dependence of the luminosity distance and mass density of any given $\\Lambda$CDM model.  We show that in general one expects the cosmological solution thus obtained to have a ``mixmaster'' approach to the initial singularity, with the radial scale factor diverging and the angular scale factor vanishing.  We then show that there is another cosmological solution with the same density and curvature functions as the singular model  but with regular initial conditions where both scale factors vanish at the singularity, and that this regular solution well \\textit{approximates} another $\\Lambda$CDM model.  We then discuss the utility of different averaging prescriptions for understanding the results.  In the next section we briefly review spherically symmetric cosmological solutions.  In Sect.\\ \\ref{recon} we review the construction of such an inhomogeneous model that reproduces key observational features of $\\Lambda$CDM. In Sect.\\ \\ref{bang} we evolve the model off the light cone, paying special attention to its behavior near the bang.  In Sect.\\ \\ref{average} we calculate some spatially averaged features of the model.  Section \\ref{conclusions} presents the conclusions and discusses their implications. ",
        "conclusions": "}  In the discussion of what we have learned in this investigation, it is useful again to refer the reader to Fig.\\ \\ref{rhat} and to reiterate that all direct cosmological information we have is limited to the light cone, illustrated by the solid line in the figure.  Any statement about the Universe off of the light cone is unsupported by direct observations, and hence requires a cosmological model for the evolution of the Universe.  An interesting lesson from the consideration of this paper is that if one is just restricted to the two cosmological observables employed in this paper, $\\hat{d}_L(z)$ and $\\widehat{\\rho}(z)$, it is impossible to  determine \\textit{uniquely} a cosmological model \\cite{Mustapha:1998jb,Celerier:2009sv}. In our case, the fiducial $\\Lambda$CDM model and the singular LTB model result in identical observations for $\\hat{d}_L(z)$ and $\\widehat{\\rho}(z)$.  That result is valid even if one imagines perfect astronomical observations, since by construction the two cosmological observables are degenerate in the two models.  There are two ways to break the degeneracy, one observational and the other theoretical.   One promising avenue to break the degeneracy with observational data might be to take advantage of the fact that in LTB models there are two different expansion rates, a radial expansion rate $H_r$, and an angular expansion rate $H_\\perp$ (they were defined in Sect.\\ \\ref{LTB}). Figure \\ref{h} shows the differences between the expansion rates in the singular model.  Note that the fact that $H_r=H_{\\Lambda CDM}$ is guaranteed by the input requirements used to construct the model. Quartin and Amendola \\cite{Quartin:2009xr} have recently discussed using cosmic parallax and redshift drift to distinguish between void models and dark energy. It is possible to pile on additional observational constraints to break the degeneracy.  \\begin{figure} \\begin{center} \\includegraphics[width=12cm]{h.eps}  \\caption{The expansion rates as a function of redshift compared to the result for the fiducial $\\Lambda$CDM model. The expansion rates $H_r$ and $H_\\perp$ were defined in Sect.\\ \\ref{LTB}.}  \\label{h}  \\end{center} \\end{figure}  While this can be done, it misses the important point that if one is restricted to light-cone observations, it is unclear whether any number of observations alone, even of arbitrary precision and with no systematic uncertainties, can uniquely determine a cosmological model, even if it is possible to rule out any LTB model.  An FLRW model is a zero-dimensional (dimensions of space) model, and the LTB solution is a one-dimensional model. Intuition from other fields lead us to believe that the range of solutions in two-dimensional, or fully inhomogeneous three-dimensional models, will be richer.  So it is unclear whether any number of cosmological observations, even to arbitrary accuracy, can uniquely specify a cosmological model.  This point was discussed by Kolb, Marra, and Matarrese (KMM) \\cite{Kolb:2009rp}.  In that work they introduce the concept of cosmological background solutions.  This concept will be useful in elucidating several points, so we briefly review some concepts from KMM.  A cosmological background solution is defined as a mean-field geometry of suitably averaged Einstein equations, in which the average expansion is described by a single scale factor. A cosmological background solution depends upon the spatial curvature and the mass-energy content. Associated with the mass-energy content is a stress-energy tensor that describes a fluid (or fluids) with equation(s) of state that satisfy local energy conditions.  KMM then define several types of background solutions; of relevance here is the \\textit{Phenomenological Background Solution (PBS)}: a background solution that effectively describes the observations, that is, the data on the past light cone of an observer. The energy content and curvature of the PBS are not necessarily related to the spatial averages of the energy content and curvature of the observable universe. The equation of state of the stress-energy tensor of the PBS need not satisfy any of the local energy conditions of the mass-energy content of the universe.  In this nomenclature, one might imagine that the true background solution is the LTB solution, or more generally some inhomogeneous model, and the fiducial $\\Lambda$CDM solution is merely a phenomenological background solution. It is unclear whether it is possible, in principle, to exclude this possibility just on the basis of light-cone observations.  Observations alone may not be able to determine the true cosmological model, but as remarked in the introduction, we expect the cosmological model to be part of our deeper comprehension of the physical and astronomical facts; it must fit into the larger framework of particle physics, general relativity, astrophysics, and cosmology.  As an example of larger considerations, consider the theory of structure formation.  In the fiducial $\\Lambda$CDM model, structure grows in a background that is close to homogeneous and isotropic, but with small seed perturbations initially generated by inflation. Even in this zero-dimensional model, structure formation is a complicated affair (at least in the nonlinear regime). Structure formation in LTB models is an even more complicated affair \\cite{Biswas:2006ub,book}: each spherical shell evolves as an independent FLRW model.  One would expect the evolution of structure formation in LTB models to differ from the fiducial $\\Lambda$CDM model. However, since structure formation in a three-dimensional model is more complicated still, it is difficult to see how structure formation might be conclusive.  The same argument might apply to other cosmological tests like the integrated Sachs--Wolfe effect.  Now, we return to the argument that purports to prove that inhomogeneous models cannot mimic the effects of dark energy (see, \\textit{e.g.,} Refs.\\ \\cite{lamuros,Siegel:2005xu,Ishibashi:2005sj}).  The argument might be paraphrased as follows: Perform an expansion of the FLRW metric to  Newtonian order in potential and velocity but take into account the fully nonlinear density inhomogeneities.  In the conformal Newtonian gauge, the metric will be of the form \\begin{equation} ds^2 = a^2(\\eta) [-(1 + 2 \\psi) d\\eta^2 + (1 - 2 \\phi) d\\bf{x}^2] , \\end{equation} where $\\eta$ is conformal time and $\\phi \\simeq \\psi$.  The volume average of the equation then yields \\cite{Siegel:2005xu} \\begin{equation} \\left( \\frac{\\dot{a}}{a} \\right)^2 = \\frac{\\kappa}{3} \\langle\\rho\\rangle \\left( 1 - 5W + 2K \\right) , \\end{equation} where $W$ and $K$ are the Newtonian potential and kinetic energy (per unit mass), \\begin{equation} \\label{wandk} W = \\frac{1}{2} \\langle (1 + \\delta) \\phi \\rangle , \\qquad K = \\frac{1}{2} \\langle (1 + \\delta) v^2 \\rangle . \\end{equation} Here, $\\delta = \\delta \\rho / \\langle\\rho\\rangle$ is the density perturbation in the matter rest frame and $v$ is the peculiar velocity.   The argument is that observationally both $W$ and $K$ are small, hence, inhomogeneities cannot mimic dark energy.  However the model constructed in Sec.\\ \\ref{recon} seems to belie this simple and compelling argument.  If one interprets the observations of $\\widehat{\\rho}(z)$ of the LTB model in terms of the fiducial $\\Lambda$CDM phenomenological background solution, one would conclude that the perturbations are very small (\\textit{viz.}\\ zero), hence $W$ and $K$ zero, and inhomogeneities could not mimic dark energy. But, at least for $\\hat{d}_L(z)$ and $\\widehat{\\rho}(z)$, inhomogeneities do exactly that!  We point out that the very notion of peculiar velocities requires a background solution (a norm) about which to define ``peculiarity'' \\cite{Kolb:2008bn}. Also, Enqvist, Mattsson, and Rigopoulos \\cite{Enqvist:2009hn} explicitly demonstrate how spherically symmetric perturbations of a dust Universe can lead to acceleration while the perturbations of the gravitational potential remain small on all scales.  Perhaps the general lesson is that there are highly non-Gaussian inhomogeneities in the late Universe, and the coherence of structures causes small deviations in observables to sum to a large deviation.  We can venture a different way to frame the argument: If the observed density fluctuations, $\\delta\\widehat{\\rho}(z)$, originated from the growth of perturbations in an FLRW cosmology, and the nonlinear perturbations on one scale did not affect the growth of structure or the evolution of the background solution on any other scale, and inhomogeneities do not show coherence or non-Gaussianity, and $\\phi\\simeq\\psi$, then one can say that the interpretation of observations would yield a value of $W$ and $K$ that are too small to mimic dark energy.  This way of framing the argument makes manifest the fact that there are dynamical assumptions that enter, and it cannot simply be a matter that small peculiar velocities rule out the backreaction proposal.  If any of the dynamical assumptions fail, the argument fails.  Another lesson that can be extracted from the exercise of this paper concerns using averaging procedures to quantify the effect of inhomogeneities on observables. This was the subject of the previous section. The conclusion was that while the two averaging prescriptions discussed give an indication that the LTB model would have observables that might be interpreted as arising from dark energy, the effective equation of state parameters do not resemble that of the fiducial $\\Lambda$CDM model.  This might be because the averaging approach is not very useful, or it might imply something quite troubling. Perhaps the equation of state parameter $w$, which has received so much attention, has a physical interpretation only in FLRW models, and in an inhomogeneous Universe, where there is no natural time slicing, it is an unphysical construct.  This paper has developed and discussed issues related to the interpretation of cosmological observables in constructing cosmological models and theories. Observations of $\\hat{d}_L(z)$ and $\\widehat{\\rho}(z)$ do not, by and of themselves, prove that there is dark energy or that the Universe is accelerating in the usual sense; they do so only if one assumes that the Universe is homogeneous and isotropic and the dynamics of the Universe are governed by general relativity. This paper discussed aspects of the program to develop understanding of how an inhomogeneous Universe might mimic dark energy.  While consideration of the results of this investigation might lead one to conclude that there are no arguments that conclusively exclude the backreaction proposal, we are no closer to developing an inhomogeneous cosmology alternative to dark energy.  One attempt to construct inhomogeneous cosmological models involves studying ``Swiss-cheese'' models comprised of low density voids embedded in an Einstein--de Sitter background \\cite{Sugiura:1999fm,Kozaki:2002ka,Brouzakis:2006dj,Marra,Bolejko:2008xh,Vanderveld:2008vi,Brouzakis,Ghassemi:2009ug,Clifton:2009nv}. If the void regions are packed together the regions between them form a network of high-density filamentary, or spaghetti-like, structures.  Further refinements of this approach have included spherical mass concentrations referred to as ``meatballs'' \\cite{Kainulainen:2009sx}.  A Swiss-cheese, meatball, and spaghetti cosmological model may include ingredients that are desirable on their own, or work well in combination with one other ingredient, but combining them all may lead to a model that is not easily digestible.  Whatever direction the study of inhomogeneous cosmologies might take, consideration of the results of this investigation should be important."
    },
    "0911/0911.3063_arXiv.txt": {
        "abstract": "We report the detection of pulsed $\\gamma$-rays for PSRs J0631+1036, J0659+1414, J0742--2822, J1420--6048, J1509--5850 and J1718--3825 using the Large Area Telescope (LAT) on board the {\\em Fermi Gamma-ray Space Telescope} (formerly known as GLAST). Although these six pulsars are diverse in terms of their spin parameters, they share an important feature: their $\\gamma$-ray light curves are (at least given the current count statistics) single peaked. For two pulsars there are hints for a double-peaked structure in the light curves. The shapes of the observed light curves of this group of pulsars are discussed in the light of models for which the emission originates from high up in the magnetosphere.  The observed phases of the $\\gamma$-ray light curves are, in general, consistent with those predicted by high-altitude models, although we speculate that the $\\gamma$-ray emission of PSR J0659+1414, possibly featuring the softest spectrum of all {\\em Fermi} pulsars coupled with a very low efficiency, arises from relatively low down in the magnetosphere. High-quality radio polarization data are available showing that all but one have a high degree of linear polarization. This allows us to place some constraints on the viewing geometry and aids the comparison of the $\\gamma$-ray light curves with high-energy beam models. ",
        "introduction": "\\newlength{\\figwidth} \\setlength{\\figwidth}{0.95\\hsize}  \\begin{table*}[t] \\caption{\\label{RefTableSpin}Rotational and derived parameters for six pulsars. } \\begin{center} \\begin{tabular}{llccccc} \\hline \\hline \\multicolumn{1}{c}{PSR} & \\multicolumn{1}{c}{PSR}  & $P$ & $\\dot{E}$ & $\\tau_\\mathrm{c}$\\tablenotemark{a} & $B_\\mathrm{S}$\\tablenotemark{b} & $B_\\mathrm{LC}$\\tablenotemark{c}\\\\ \\multicolumn{1}{c}{(J2000)} & \\multicolumn{1}{c}{(B1950)} & (sec) & ($10^{35}$ erg\\,s$^{-1}$) & ($10^{3}$ yr) & ($10^{12}$ G) & ($10^{3}$ G)\\\\ \\hline J0631+1036  &           &  0.288 &   1.73 &  43.6 & $5.55$ & $2.18$\\\\ J0659+1414  & B0656+14  &  0.385 &   0.38 & 111   & $4.66$ & $0.77$\\\\ J0742--2822 & B0740--28 &  0.167 &   1.43 & 157   & $1.69$ & $3.43$\\\\ J1420--6048 &           &  0.068 & 104    &  13.0 & $2.41$ & $71.3$\\\\ J1509--5850 &           &  0.089 & 5.15   & 154   & $9.14$ & $12.2$\\\\ J1718--3825 &           &  0.075 & 12.5   &  89.5 & $1.01$ & $22.6$\\\\ \\hline \\end{tabular} \\end{center} \\tablenotetext{a}{Spin-down age $\\tau_\\mathrm{c}=P/(2\\dot{P})$.} \\tablenotetext{b}{Magnetic field strength at the surface of the star in Gauss $B_\\mathrm{S}=3.2\\times10^{19}(P\\dot{P})^{1/2}$.} \\tablenotetext{c}{Magnetic field strength at the light cylinder in Gauss $B_\\mathrm{LC}=3.0\\times10^{8}(\\dot{P}/P^5)^{1/2}$.} \\end{table*}  The {\\em Fermi Gamma-ray Space Telescope} (formerly known as GLAST) was successfully launched on 2008 June 11. The study and discovery of $\\gamma$-ray pulsars is one of the major goals of this mission. Studying pulsars at these high energies is important, because a large fraction of the total available spin-down energy loss rate ($\\dot{E}=4\\pi^2I\\dot{P}P^{-3}$) is emitted in $\\gamma$-rays. Here $I$ is the moment of inertia of the star (generally taken to be $10^{45}$~g$\\,$cm$^{2}$), $P$ its spin period and $\\dot{P}$ its spin down rate. By studying individual {\\em Fermi} detections as well as the population of $\\gamma$-ray pulsars as a whole, models for the high-energy emission can be constrained (e.g. \\citealt{hgg07,wrwj09}).  The models can be divided into three different families that place the emitting regions at different locations in the pulsar magnetosphere. In the so-called polar-cap model (e.g. \\citealt{dh96}) the $\\gamma$-ray photons are produced close to the neutron star surface (within a few stellar radii) near the magnetic axis.  At the other extreme are the outer-gap models (e.g. \\citealt{mor83,chr86a,ry95}), which place the emitting region near the light cylinder. Finally, in slot-gap models (e.g. \\citealt{mh04}) the particle acceleration occurs in a region bordering the last open field lines at a large range of emission altitudes. The two-pole caustic model (\\citealt{dr03}) is a geometrical realization of the slot-gap model.  It has recently been demonstrated that $\\gamma$-ray pulsars can be discovered via blind searches in the {\\em Fermi} data \\citep{aaa+09c}. Nevertheless the detection threshold for pulsed $\\gamma$-rays is lower when accurate positions and spin frequencies (including their unpredictable timing irregularities) are already known. Therefore a set of pulsars with $\\dot{E}>10^{34}$ erg\\,s$^{-1}$ is being timed in the radio band, allowing {\\em Fermi} to search for pulsations with the highest possible sensitivity \\citep{sgc+08}. In addition these radio observations allow us to determine the difference in arrival time of the $\\gamma$-ray pulses with respect to the radio pulses, an important parameter to distinguish between different high-energy models.  In this paper we report the detection of pulsed $\\gamma$-rays for six pulsars that were found by folding the {\\em Fermi} data on the ephemerides obtained from radio observations (e.g. \\citealt{wjm09}). These {\\em Fermi} detections will be included (although with more limited count statistics) in the LAT pulsar catalog paper \\citep{aaa+09e}. The six pulsars of this paper have moderate to large spin-down luminosities ($\\dot{E} > 10^{34.5}$ erg\\,s$^{-1}$, see Table \\ref{RefTableSpin}) and have (at least given the current count statistics) $\\gamma$-ray light curves which are consistent with a single peak. Five of the six pulsars have a strong degree of linear polarization in the radio band, which can be used to constrain the emission geometry. Therefore the combination of radio and $\\gamma$-ray data for these objects make them valuable for tests of the underlying beaming geometry.  The paper is organised such that we will start with describing the {\\em Fermi} observations in Section \\ref{SctObs}. This is followed by a description of the methods used to constrain the emission geometry using radio data in Section \\ref{SctMethods}.  The results, including the $\\gamma$-ray light curves and spectral parameters obtained with {\\em Fermi}, are presented in Section \\ref{SctResults} and discussed in Section \\ref{SctDiscussion}. Finally the results and conclusions are summarised in Section \\ref{SctConclusions}. ",
        "conclusions": "\\label{SctConclusions}  We report here on the detection of pulsed $\\gamma$-rays by {\\em Fermi} for PSRs J0631+1036, J0659+1414, J0742--2822, J1420--6048, J1509--5850 and J1718--3825.  These six pulsars are young to middle-aged and, except for PSR J1420--6048, have relatively small values of $\\dot{E}$ compared to other known $\\gamma$-ray pulsars. In all cases the $\\gamma$-ray light curves appear single peaked (at least with the present count statistics), but there is a hint that at least two of them have closely-spaced double peaks. As {\\em Fermi} continues its all-sky survey, the quality of the light curves will increase, helping to resolve this issue.  We present high quality radio polarization profiles for these pulsars and discuss their geometries in the context of RVM fitting and simple beam modelling. Unfortunately, the narrow phase range of the radio emission generally leaves a strong degeneracy between the fit $\\alpha$ and $\\beta$ values. This, in combination with the limited $\\gamma$-ray count statistics makes it difficult to distinguish between single-peaked and double-peaked light curves and hence to make the comparison with the model predictions.  We show that models where the $\\gamma$-ray emission occurs at relatively high altitudes in the pulsar magnetosphere, such as the outer-gap or two-pole caustic models, do a good job in predicting the phase of the $\\gamma$-ray emission relative to the radio emission.  However, the shape of the $\\gamma$-ray light curve is less well modelled, and additional inputs to the model are needed to explain the strong bridge emission in some pulsars and the strong single $\\gamma$-ray component in PSR~J0742--2822 in particular. Finally, PSR~J0659+1414 warrants further attention.  It has a peculiar light curve and phase offset with the radio profile, its $\\gamma$-ray efficiency is low and its $\\gamma$-ray spectrum is extremely soft. It may be an aligned rotator with $\\gamma$-ray emission arising relatively low down in the magnetosphere. Higher signal-to-noise ratio $\\gamma$-ray light curves in combination with possible additional geometrical constraints, such as from PWN imaging, will result in stronger constraints on the models."
    },
    "0911/0911.2912_arXiv.txt": {
        "abstract": "Wave properties and instabilities in a magnetized, anisotropic, collisionless, rarefied hot plasma in fluid approximation are studied, using the 16-moments set of the transport equations obtained from the Vlasov equations. These equations differ from the CGL-MHD fluid model (single fluid equations by Chew, Goldberger, and Low \\cite{Chew,Baran}) by including two anisotropic heat flux evolution equations, where the fluxes invalidate the double polytropic CGL laws. We derived the general dispersion relation for linear compressible wave modes. Besides the classic incompressible fire hose modes there appear four types of compressible wave modes: two fast and slow mirror modes -- strongly modified compared to the CGL model -- and two thermal modes. In the presence of initial heat fluxes along the magnetic field the wave properties become different for the waves running forward and backward with respect to the magnetic field. The well known discrepancies between the results of the CGL-MHD fluid model and the kinetic theory are now removed: i) The mirror slow mode instability criterion is now the same as that in the kinetic theory. ii) Similarly, in kinetic studies there appear two kinds of fire hose instabilities - incompressible and compressible ones. These two instabilities can arise for the same plasma parameters, and the instability of the new compressible oblique fire hose modes can become dominant. The compressible fire hose instability is the result of the resonance coupling of three retrograde modes - two thermal modes and a fast mirror mode. The results can be applied to the theory of solar and stellar coronal and wind models. ",
        "introduction": "Frequent particle collisions turn the plasma distribution function into an isotropic one, and thus the thermal pressure is isotropic as well. If collisions rarely occur the presence of a magnetic field will maintain a ``non-mixed'' state of the energies of the longitudinal and transverse motions of particles. Thus, the transverse and longitudinal kinetic particle temperatures will differ from each other, $T_\\bot\\ne T_\\|$ . Typical examples of such plasmas are the coronal and solar wind plasmas which are very anisotropic and inhomogeneous, in cross-field direction in particular \\cite{Asc05}. For the observational motivation of our modeling see the references in our previous paper \\cite{Dzhal08}. So, due to the anisotropy of the kinetic temperatures of the particles (especially protons and heavy ions) the corresponding partial pressures become anisotropic in this way. This makes the total thermal pressure anisotropic too, $p_\\bot\\ne p_\\|$.  In the corona the electron and ion gyroradii $r_{B}$ and gyrotimes $\\tau_{B}$ become smaller than any of the particle collisional mean free paths and times and smaller than any of the typical scales of variations of macroscopic thermodynamical quantities. Thus, the condition of a strongly magnetized plasma is well satisfied. That means, particles gyrating around the magnetic field lines are localized across the field at a distance of the Larmor radius which plays the role of a free path length of particles. Thus, the dynamical motion of a collisionless plasma with characteristic scales of $L\\gg r_{B}$ and $\\tau\\gg \\tau_{B} $ behaves across the magnetic field as a fluid. However, under such circumstances a traditional hydrodynamical description of the plasma is hardly possible \\cite{Marsch06}. The isotropic MHD equations are applicable only if the plasma is collision-dominated and the distribution functions are close to Maxwellian's. In the opposite limiting case, when the plasma is collisionless at all, the local distribution function strongly differs from the Maxwell function. To describe the plasma in the fluid approximation in this case usually the single fluid CGL-MHD equations by Chew et al. \\cite{Chew,Baran} are applied. Instead of the energy equation of the isotropic MHD these equations include double-polytropic laws in the form $p_\\bot/\\rho B = $ const and $p_\\| B^2/\\rho^3=$ const. Many studies of wave instability problems are based on these equations. Similar to the low-frequency kinetic considerations there appear two kinds of instabilities: incompressible fire hose and compressible mirror instabilities \\cite{veden58,chandra58,parker0,barnes66,Hasegawa69,Gary98}. In comparison to the results based on the kinetic theory the CGL equations provide the correct instability criterion for the classic fire hose instability. However, there exists a discrepancy in the criterion for the slow-mode mirror instability. Moreover, there appear basic differences between the nonlinear stages of these instabilities compared with the kinetic results. It has been shown \\cite{Hellinger00,Hellinger01} by hybrid kinetic simulations that a new type of fire hose instability may arise for oblique propagation due to the proton temperature anisotropy. Unlike the classical fire hose instability this instability is compressible and has a maximum growth rate at oblique propagation. The growth rate of the new instability is comparable to (or in some parameters ranges even larger than) the maximum of the standard fire hose instability growth rate. Both fire hose instabilities may occur at the same time for the same plasma parameters. A similar second type of fire hose instability driven by the electron temperature anisotropy has also been found \\cite{Hollweg70,Pae99,Li00}. The CGL theory cannot give an analogy of this second type of fire hose instability.  To remove such discrepancies in the CGL-MHD model, generalized polytropic laws were used \\cite{Abraham73,Hau93,Wang03} introducing some artificial polytropic indices such as  $p_\\bot/\\rho B^{\\gamma_1-1} = $ const and $p_\\| B^{\\gamma_2-1}/\\rho^{\\gamma_2} =$ const. With a suitable choice of these free polytropic indices, $\\gamma_1$ and $\\gamma_2$, it is in principle possible to remove some of the discrepancies. However, the origin of these new indices is not clear as they do not follow directly from the kinetic equations when the fluid equations are derived. Later it has been shown that in the experimental data and in the particle simulations these two adiabatic invariants become invalid for realistic plasmas \\cite{Marsch87,Quest96}. The conservation of the two CGL adiabatic invariants in an ideal collisionless plasma leads to a strong pressure anisotropy $p_\\perp < p_\\|$ which is much larger than the observed values \\cite{Grooker77,Hill95}. However, due to non-ideal effects such as heat flux these invariants are broken \\cite{Hau96}, and this leads to properties which are quiet different from those predicted by the CGL equations. Deriving the CGL-MHD equations the third moments of the distribution function, hence the heat fluxes have been ignored without any proof \\cite{Baran}, which is the main shortage of these equations.  In the present paper we study the linear wave instability problem on the base of more correct equations -- the 16-moments transport equations, which are derived from the Vlasov collisionless magnetized plasma kinetic equations by the fast gyromotion ordering technique \\cite{Oraev68,Oraev85,Ramos03}. These equations include additionally two dynamic evolution equations of the heat fluxes and no polytropic laws are possible. This allows us to resolve the main discrepancies of the CGL fluid theory. We consider the wave peculiarities which can appear in the anisotropic compressible plasma. We have already considered the incompressible wave instability on the base of these equations \\cite{Dzhal08}. In Section 2 we formulate the basic equations, which are the integrated moment equations of the kinetic Vlasov equations. In Section 3 the linear compressible wave equation and the general dispersion relation are derived. To compare the results with the CGL-MHD theory the CGL dispersion equation is deduced from the new dispersion equation as a special case in Section 4. The solutions and analysis of the new dispersion equation are the topic of Section 5. In Section 6 we show that similar to the the kinetic theory the existence of two kinds of fire hose instabilities is possible in the fluid approximation too. The mass density fluctuations due to compressible wave modes and their instability is discussed in the Section 7. A discussion and some conclusions are presented in Section 8. ",
        "conclusions": "\\subsection{Model equations} Compared with a full kinetic model the fluid description of a plasma has the mathematical advantage of a smaller number of dimensions. Moreover, many observed dynamic phenomena in space plasmas are large-scale structures --- already averaged over both temporal and spatial scales. This suggests to find ways for a description of a plasma as a fluid. It is easily possible in the case of a collision-dominated plasma which is described by a Maxwellian distribution function. In this case the usual isotropic MHD equations are received. The situation becomes much more complicated for a smaller frequency of the collisions between the particles of a hot magnetized plasma (such as space plasmas in most cases). Due to the magnetic field the collisionless plasma becomes anisotropic with respect of its local direction. To describe in this case the plasma as a fluid the transport model equations are deduced from the kinetic equations for the moments of the distribution function. These moments are such quantities as plasma mass density, fluid speed of the plasma, anisotropic thermal pressure, anisotropic thermal flux, etc. Basically, the number of these moments is infinite, and the equations of these moments are coupled among each other. However, if the conditions $ \\omega / \\Omega_{B} \\ll 1$ and $k\\, r_{B} \\ll 1$ are satisfied ($ \\omega $ -- frequency of perturbations, $ \\Omega_{B} $ -- gyration frequency, $k $ -- wave number, $r_{B} $ -- gyration radius), it is possible to break off the chain of these equations, if some additional conditions related to the given exact analytical form of the particle distribution function are fulfilled. This method is called the standard method of MHD ordering of the kinetic equations, and the resulting new equations are called ``transport equations''. However, the application of this ordering method has some subtleties -- it is not trivial \\cite{Ramos03}. Depending on these and on additionally chosen conditions, the obtained transport equations can be different. By including higher order moments it is possible to increase the accuracy of the transport equations. With increasing order of the ordering the accuracy of the model equations also increases, but they become more complicated for the analysis. Even though these equations can never be complete without supposing additional conditions, they describe well such phenomena as Alfv\\'{e}nic and acoustic (electronic and ionic) waves, they can include such an important kinetic effect as Landau damping. However, depending on the order of ordering such kinetic effects as drift-waves and other micro-instabilities can be lost. Besides, to deduce the transport equations an exact analytical type of the function of particle distribution with anisotropic temperatures, e.g. bi-Maxwellian  or $k$-distributions, is required. That means, the fine structure of the realistic distribution functions and the related microphysics are ignored.  Among these model equations the CGL equation are most simple with respect to the included moments and the order of ordering. The heat flux tensor in these equations is ignored at all. Strictly speaking, i.e. the phase speed of the perturbations should be much larger than the thermal velocity of the particles. Such condition are met for Alfv\\'{e}nic modes, but for acoustic modes this condition is impracticable. The inclusion of a heat flux tensor to the equations was carried by many authors \\cite{Barak,Snyder,Mahajan}. Results of our studies in the present paper (types of instabilities, conditions of their existence, thresholds and values of growing rates of instabilities, comparisons of these with the corresponding results for mirror and ion-acoustic instabilities in the low-frequency kinetic approach) let us come to the conclusion that the used equations derived by Ramos \\cite{Ramos03} are more correct.  To sum it up it can be said that we have used more complete fluid transport equations describing the macroscopic behavior of a magnetized anisotropic collisionless plasma. In particular, we include heat fluxes and their evolution which, contrary to the CGL-MHD model, are a basic feature and cannot be ignored. We consider only two parallel heat fluxes corresponding to parallel and perpendicular thermal motions of particles. Perpendicular heat fluxes have been neglected.  \\subsection{Mirror instability} In strongly magnetized and weakly collisional turbulent plasmas the anisotropy of the pressure develops in a spontaneous way \\cite{Schekoch,Sharma}. In a high-beta plasma that triggers a number of instabilities, above all firehose and mirror \\cite{chandra58,parker0,veden58,barnes66,Hasegawa69}. The nonlinear evolution of the instabilities should have the tendency to compensate on the average the pressure anisotropies generated by the turbulence. Thus the development of the instability of the modes further changes the distribution function, tending to make it more isotropic, or to strengthen the magnetic structurization. If $p_\\bot > p_ \\| $, there appear two kinds of electromagnetic instability: ion-cyclotron instability at frequencies $ \\omega < \\Omega _ {Bi} $ ($ \\Omega _ {Bi} $ -- ion cyclotron frequency) and magneto-mirror instability at a very low frequencies, $ \\omega\\approx 0$. Because there exist numerous observations of low-frequency turbulence in magneto-active plasmas, for instance in magnetosheaths, in solar wind, and in cometary comas (see references in \\cite{Gedal01}) the mirror instability was studied theoretically in detail. It was shown that in ionic high-beta plasma the mirror modes become unstable if an anisotropy indicator $p_\\bot / p_ \\| - 1 $ exceeds some critical value \\cite{chandra58,veden58}. This instability causing a local deformation of the magnetic field makes the plasma spatially inhomogeneous.  It occurs because a part of the particles captured in ``weak mirror traps'' subdivides the distribution of particles into passing and trapped species \\cite{Kivel}.  The fluid analogy of the kinetic mirror instability has a similar simple description \\cite{Hasegawa69,Thompson}. Basically, the instability occurs because at low frequencies the changes of the perpendicular pressure of the plasma and of the magnetic field occur in opposite phases. Really, as follows from Eq. (\\ref{per}), for $\\omega\\to 0$ neglecting small density perturbations, $p_\\bot' \\sim - p_\\bot (\\frac{p_\\bot}{p_ \\|} - 1) \\frac {B'}{B}$. Hence, in those places of the plasma where $p_\\bot >  p_ \\| $ a decrease of the plasma pressure increases locally the intensity and the pressure of the magnetic field. Meeting the condition $ \\frac {p_\\bot}{p_ \\|} -1 > \\frac{p_{\\mathrm m}}{p_\\bot}$ the force (caused by the total pressure) in perpendicular direction decreases. Thus, the increase in intensity of the magnetic field locally reduces the total pressure which, in its turn, pushes together the magnetic field lines even more, i.e. leads to a further growth of the magnetic field. It causes instability. From the invariance of the first magnetic moment of the plasma follows that the energy of the particles will simultaneously grow in the direction perpendicular to the magnetic field. From the conservation of total energy follows that the parallel energy should decrease accordingly. Under these conditions (when the magnetic moment and the total energy are conserved) the plasma will naturally flow from an area with high magnetic field intensity to an area with a weak magnetic field. That means there is a conversion of perpendicular to parallel energy. It looks like an acceleration of particles along the magnetic field caused by a certain force. The name of this force is magneto-mirror. But this force is a pseudo-force as there is only a swapping of energy from a perpendicular into a longitudinal direction, the total energy doesn't change. If the instability criterion is not fulfilled the fluid mirror modes show an oscillatory behavior.  Contrary to the fluid description of the mirror instability in the kinetic description not all particles equally react to the magnetic field changes \\cite{Tajiri}. Particles with small parallel speed ``do not feel'' the mirror force to the same degree as the particles with larger parallel speed. In this way with changing magnetic field its energy is not conserved, and the perpendicular pressure changes synchronously with the magnetic field change. Unlike the fluid approach in the kinetic treatment the mirror modes become non-oscillating (exponentially damped) if the instability criterion is not satisfied. This is because of the resonant origin of the kinetic mirror instability \\cite{South}.  We can compare the growing rate of the mirror instability in our fluid description Eq. (\\ref{gromir}) with the different kinetic estimates. There is good agreement. Note that the more exact numerical calculations of the kinetic growing rate for bi-Maxwellian plasma \\cite{Gary93} show that a maximum of the growth rate occurs at $k_\\bot r _ {Bi} \\sim 1$, well above the threshold, and $ \\Gamma_{\\mathrm max} \\sim k_\\| v_ {Ti \\|} $.  We focused our interest on the properties of the mirror modes because of the high probability of their realization in practice. The properties of the other modes, fire hose and thermal modes, are well known. The first one is the prototype of Alfv\\'{e}nic waves and the second one is analogous to the ionic magnetic sound in the kinetic approach. At places of a crossing of wave branches we have a mixing of modes, i.e. there is a resonant interaction of wave modes. As usually in hydrodynamics if we consider the anisotropy of the plasma as well as its spatial inhomogeneity, these points of mode intersection will introduce singularities into the wave equations. Unlike the kinetic approach where wave modes can grow or fade as a result of a resonance of particles and fluctuations, in the fluid description it occurs as the result of resonant interaction of different modes.  In the present paper the linear instability problem is investigated in a homogenous, unlimited plasma. The classic incompressible fire hose modes are not influenced by the heat fluxes, and their instability criterion is the same as that in kinetic theory. However, the two CGL mirror modes are strongly modified by the heat fluxes, and there appear two additional thermal branches. These thermal modes result from including the two dynamic evolution equations of the thermal fluxes. In the initial state nonzero heat fluxes are supposed, $\\gamma \\ne 0$. However, even for zero initial heat fluxes ($\\gamma = 0$) the thermal modes appear. The deduced 8-th order polynomial dispersion equation describes the interaction of all of the 4 types of compressible modes and their stability. If the initial heat fluxes along the magnetic field ($\\gamma \\ne 0$) are included, the mode dispersion behavior and the instability criteria are different for the same type of modes running along and backward with respect to the magnetic field.  To sum it up it can be said that the main shortages of the CGL fluid theory have been removed. It is shown that the fluid slow mirror instability criterion is the same as that in the kinetic theory: $p_\\bot^2/p_\\| = p_\\bot + p_{\\mathrm m}$. We have shown that in some selected ranges of the parameters $\\alpha$ (parameter of pressure anisotropy),  $\\beta$ (parameter of magnetic field), and $\\gamma$ (parameter of initial heat flux) in the plasma there exists such a propagation angle ($l = \\cos^2\\theta$) in which at the same time two kinds of fire hose instability can develop. The discovered new instability is compressible, slightly periodical (${\\mathrm Re}(\\omega) \\ne 0$), and it has a larger growing rate than the incompressible fire hose instability for parallel propagation. It seems that these modes are analogous to the two kinds of fire hose instabilities found in kinetic theory \\cite{Hollweg70,Pae99,Li00,Hellinger00}. This new instability develops when the three retrograde modes (two thermal and fast mirror) interact resonantly.  We found a strong dependence of the growth rate on the parameters $\\alpha, \\beta, \\gamma$ and $l$. There appear different unstable and stable wave branches simultaneously within the given parameter ranges. Only the mode with the highest growth rate will dominate, and after some exponential growth the nonlinear stage of the instability should be considered. It is of basic importance for the found instabilities that in the collisionless plasma there is a plasma pressure anisotropy that is kinetically supported. The origin of this pressure imbalance is not important for our fluid approximation; it is the background of large-scale flows only. In principle many kinds of kinetic wave turbulence can support such a pressure anisotropy, the existence of which in the considered plasma situation is shown by observations.  \\subsection{Solar wind} We think that the discovered wave modes and their instabilities in the anisotropic fluid approximation are interesting for those plasmas for which the approximation for a magnetized hot plasma with rare collisions can be applied. Important candidates for such conditions are the solar wind and the solar corona plasma. Macroscopic turbulence observed in the solar wind \\cite{Marsch06} and in the stable coronal turbulent background (appearing in the nonthermal broadening of coronal emission line profiles \\cite{Asc08}) may be a consequence of these instabilities. Moreover, it is now generally accepted that the observed large ion temperature anisotropies are related to the physical mechanism by which the solar corona and solar wind are heated \\cite{Hollweg02,Marsch06}.  Near the Sun heavy ions are stronger heated than protons and electrons. These findings have strengthened the arguments in favour of the kinetic ion-cyclotron model of heating and acceleration of particles in the solar wind. However, this mechanism has a number of shortages. For example, the observed properties of low-frequency wave turbulence are close to those of Alfv\\'{e}nic modes and their power spectrum has a maximum around one--two hours. In order to realize the ion-cyclotron resonance, Hollweg \\cite{Hollweg08} assumed that low-frequency waves by the nonlinear cascade should finally turn to high-frequency modes. Within the frames of our fluid model such observed low-frequency modes can easily be explained. Indeed, in the ideal case (without a heat flux, $\\gamma=0$) near the instability threshold of the Alfv\\'{e}nic fire hose modes it is possible to receive very low frequencies. For the nonideal case (with a heat flux, $ \\gamma\\ne 0$) the mirror modes have every chance to explain the observable low frequencies. The new observed facts which can essentially modify our ideas about the physical nature of the solar wind are presented in a recent paper \\cite{Brovsky}. It appears that the solar wind consists of sets of magnetic filaments. We think that the nonlinear evolution of the large-scale mirror modes are capable of creating such structures.  \\subsection{Identification of wave modes} In the preceding Section 7 we presented the relations between the fluctuations of the plasma mass density and the magnetic field as well as the two components of the fluid speed. These formulae can be used for identifying the modes; in particular the simple analytic estimates of the asymptotic limiting cases are helpful for recognizing easily the modes in the observed data. Modern space and ground-based observations with spectral and imaging methods allow to detect a reach spectrum of different kinds of wave motions in the corona \\cite{Asc05,Asc08,Marsh09,Wang09}. From such observed data we get information mainly on the amplitudes of density perturbations and the fluid velocity component along the line-of-sight, the wave frequency, the phase speed, the mode life time, and on the spatial orientation of the magnetic loops along which the waves are running. These data allow us to compare the observed wave motions with theoretical predictions \\cite{Nakar05}. So far all interpretations are based on the well developed isotropic MHD wave theory which is based on the collision-dominated plasma description. However, in such theoretical interpretations we should be more careful, especially concerning the outer corona. For example, even in the lower corona close to the transition region the simulated electron heat flux is not described correctly by the collision-dominated Spitzer law \\cite{Landi01}. The observed fluxes are closer to the collisionless theoretical estimates. The appearance of anisotropic wave modes is an important evidence: if there exists strong enough heat fluxes, then the phase velocities along the magnetic field will differ from those in the opposite direction. Observed life times can be compared with the growing times of the instabilities.  Our theory should be further improved. In the 16-moments transport equations the next order terms should be retained. Finite particle gyroradii should result in a $k$-dependence of the maximum of the growing instability rates such as shown in the kinetic theory \\cite{Lang02}. An inclusion of small collisional terms in the basic equations would be the best way to describe the coronal plasma.   \\begin{acknowledgement}  The present work has been supported by the German Science Foundation (DFG) under grant No. 436 RUS 113/931/0-1 (R) and by the Russian Foundation for Basic Research (RFFI) under project No. 09-02-00494 which is gratefully acknowledged. \\end{acknowledgement}"
    },
    "0911/0911.5229_arXiv.txt": {
        "abstract": "We discovered multiple high-velocity (ranging from -900 to -650 km s${}^{-1}$) and  narrow (FWHM $\\sim$ 15 km s${}^{-1}$) absorption components corresponding to both the D2 and the D1 lines of Na~{\\sc i} on a high dispersion spectrum of V1280 Sco observed on 2009 May 9 (UT), 814 d after the $\\it V$ band maximum.  Subsequent observations carried out on 2009 June  and July  confirmed at least 11 distinct absorption components in both systems. Some components had deepened during the two months period while their HWHMs and wavelengths remained nearly constant. We suggest these high velocity components originate in cool clumpy gas clouds moving on the line of sight, produced in interactions between pre-existing cool circumstellar gas and high velocity gas ejected in the nova explosion. The optical region spectrum of V1280 Sco in 2009 is dominated by the continuum radiation and exhibits no forbidden line characterizing the nebular phase of typical novae. Permitted Fe~{\\sc ii} lines show doubly peaked emission profiles and some  strong  Fe~{\\sc ii} lines are accompanied by a  blue shifted ($\\sim$ -255 km s${}^{-1}$) absorption component. However, no high-velocity and narrow components corresponding to those of Na~{\\sc i} could be detected in Fe~{\\sc ii} lines nor in the Balmer lines. The 255 km s${}^{-1}$ low velocity absorption component is most probably originating in the wind from the nova. ",
        "introduction": "The bright nova V1280 Sco (Nova Sco 2007 $\\sharp$1, \\timeform{16h57m40s.91}, \\timeform{-32D20'36''.4}, equinox 2000.0)  was discovered on 2007 February 4 by two Japanese observers (Y. Sakurai and Y. Nakamura) independently, as  communicated by \\citet{yamaoka2007}. \\citet{naito2007} obtained a low dispersion spectrum of the object on Feb. 5.87 and found that it had a smooth continuum together with the Balmer and Fe~{\\sc ii}  lines showing P Cygni profiles. The nova had brightened greatly during the first 10 days and reached the maximum brightness $\\it V$ = 3.8 on Feb. 16 \\citep{munari2007a}. The outburst amplitude ($\\it A$) is 15 mag or larger because \\citet{das2008} noted that no star was visible down to $\\it B$ and $\\it R$ magnitudes of 20.3 and 19.3, respectively, on pre-discovery plates. \\citet{munari2007a} reported their photometric and spectroscopic monitoring of the nova and noted that on Feb. 20.24 the nova showed a typical classical nova spectrum of the Fe~{\\sc ii} type characterized by a rich forest of strong permitted emission lines of Fe~{\\sc ii} displaying deep P Cyg profiles.  \\citet{das2007} observed the nova in the near IR region on 2007 March 4.95 and found that the continuum in the 1.08 - 2.35 $\\mu$ region had risen sharply indicating the  dust formation in the nova ejecta. \\citet{puetter2007} reported spectroscopic observations in the visual - near IR regions carried out in 2007 May and found that the nova was in a very low excitation state showing strong C~{\\sc i} lines and no discernible He~{\\sc i} emission. They estimated the reddening, which is in part due to the dust shell, from the O~{\\sc i} lines and derived E($\\it B - V$) = 1.7 mag. Photometric behaviors of the nova in the optical  region  until 2007 October 9 are summarized by \\citet{munari2007b}. They pointed out a re-brightening in 2007 May and noticed an unexpected large re-brightening in 2007 September.  High spatial resolution interferometric observations in the near and mid IR regions were carried out using the Very Large Telescope Interferometer (VLTI). \\citet{chesneau2008} observed spectra and visibilities of the nova from 2007 February 28 to June 30 and determined an apparent linear expansion rate for the dust shell. They  pointed out that the approximate time of the mass ejection during which the dust shell had formed was close to the date of the maximum brightness.  They used the observed expansion velocity and the linear expansion rate to derive an upper limit of the distance to be  $\\sim$ 1.9 kpc.  We have continued photometric observations in $\\it B, V, Rc, Ic$ and $\\it y$ bands at Osaka Kyoiku University and in $\\it J, H$ and $\\it Ks$ bands at Higashi Hiroshima Observatoty until 2009 September. Low resolution spectroscopic observations were conducted  at Nish-Harima Astronomical Observatory. Part of our photometric and spectroscopic observations has been reported in \\citet{naito2009}. A long-term (spanning three years) light curve of V1280 Sco reveals unique features among classical novae. Figure 1 shows long-term light curves in the $\\it V$ and the $\\it y$ bands including the latest observations obtained in the summer of 2009. After the initial rapid decline in the $\\it V$ band, the nova showed a temporary re-brightening in 2007 May. Then it declined and reached a minimum ($\\sim$ 16 mag in the $\\it V$ band) in 2007 August. After 2007 September, the nova recovered its brightness in all bands and a bright plateau ($\\sim$ 10 mag in the $\\it V$ band)  had continued throughout 2008. Although the nova showed a slight decline starting from 2009 May, it was still brighter than 11 mag in the $\\it V$ band in 2009 August, more than 900 d after the  $\\it V$ band maximum. One of the notable feature of the optical light curve is a close resemblance between the $\\it y$, which is free from emission lines, and the $\\it V$ magnitudes. Even in the summer of 2009, its $\\it y$ magnitude is only slightly fainter than the $\\it V$ magnitude. This implies that the optical region spectrum is dominated by the continuum radiation  even 900 d after the $\\it V$ band maximum. A detailed discussion of the light and color curves will be presented in a separate paper. A low resolution optical spectrum of V1280 Sco obtained at Higashi Hiroshima Observatoty on 2009 February shows no hint of the [O~{\\sc iii}] forbidden lines. Thus, the nova has not entered the Nebular phase yet. \\begin{figure}[t] \\begin{center} \\FigureFile(75mm,70mm){fig1.eps} \\end{center} \\caption{Light curves in the $\\it V$ (dots) and the $\\it y$ (open triangles) bands obtained at Osaka Kyoiku University. Data of photometric starndard stars given by \\citet{henden2007} are used. Epochs of Subaru observations in 2009 are indicated by arrows. } \\end{figure} We have carried out high resolution optical region spectroscopic observations of the nova in 2009 in order to clarify the physical conditions in the continuum radiating source. Here we focus on  a remarkable finding concerning the highly blue-shifted narrow absorption components associated with the Na~{\\sc i} yellow doublet lines. ",
        "conclusions": "We have discovered multiple high-velocity absorption components corresponding to both the D2 and the D1 lines of Na~{\\sc i} on  high resolution spectra of V1280 Sco observed three years after the explosion. There are at least 11 sharp (FWHM $\\sim$ 15 km${}^{-1}$) components and they are highly blue-shifted, ranging from -900 to -650 km s${}^{-1}$, on the helio-centric velocity scale. These absorptions are not associated with Fe~{\\sc ii} or the Balmer emission lines. This implies that the   high-velocity absorptions  are produced in a cool environment where iron atoms are in neutral state and most hydrogen  atoms stay in the ground state. This must be a very rare finding because we know no previous similar observation in the literature.  \\citet{williams2008} reported observations of short-lived blue-shifted metallic absorption systems, including the Na~{\\sc i} D lines, near the maximum light of various novae. These absorption lines have expansion velocities from 400 to 1000 km s${}^{-1}$ and velocity dispersions between 35 and 350 km s${}^{-1}$. They are usually accelerated outward and progressively weaken and disappear over timescales of weeks (within 100 d). They proposed a spiral ring model to interpret their observations.    They suggest that some material ejected by the secondary star before the nova outburst is spiraling around the binary system. After the nova outburst, a rapidly expanding luminous photosphere is produced and it collides with the pre-existing gas stream.  Our observations of narrow and high-velocity absorption components associated with the Na~{\\sc i} D lines appear hard to be interpreted in this scheme. Our data were obtained more than 800 d after the explosion, while the transient absorptions reported in \\citet{williams2008} usually disappear over timescales of weeks. Their analysis of absorption line systems observed in LMC 2005 shows the excitation temperature to be around 10${}^{4}$ K, which is too high to explain the absence of high-velocity absorption components associated with Fe~{\\sc ii} lines. Furthermore, the observed line widths of the high-velocity absorption components in V1280 Sco are narrower than the widths noted in  \\citet{williams2008}.  Finally, the large number (at least 11) of the absorption components looks difficult to be interpreted in this scenario.  We propose that the observed high velocity absorption components originate in many cool clumpy gas clouds moving toward the observer on the line of sight. These gas clouds are produced after the nova explosion in interactions between the pre-existing cool circumstellar gas and the high velocity gas ($\\sim$ 2000 km s${}^{-1}$, Naito et al. 2009)  ejected in the nova explosion. \\citet{naito2009} suggested that the 2000 km s${}^{-1}$ gas might be associated with the second mass ejection episode corresponding to the temporal re-brightening observed in 2009 May. In the case of V1280 Sco, the binary system had been embedded in a relatively dense circumstellar cloud, and the pre-existing circumstellar gas had not necessarily been supplied from the secondary star as suggested in \\citet{williams2008}. The prompt formation of dust in this nova \\citep{chesneau2008} is likely to be triggered by the interaction between the surrounding clouds and the expanding ejecta. Many small gas clouds might have been produced during the turbulent interaction in the shock and expelled outward with velocities ranging from -1000 to -500 km s${}^{-1}$. We suppose that some of the clouds are moving toward us on the line-of-sight and produce multiple absorption lines. The observed increase in line depths for several components may be the result of rapid cooling of these clouds, increasing the number density of neutral Na capable of absorbing the D line photons. This picture is hinted by a direct image of the old nova GK Per observed with the WIYN 3.5 m telescope \\citep{slavin1995}. Confirmation of this scenario could come from future high spatial resolution imaging observations.  \\vskip 3mm We thank Dr. M. Kato for providing a $\\it y$ band filter and Dr. I. Hachisu for helpful comments and suggestions. Co-operations of students at Osaka Kyoiku University are gratefully acknowledged. This research was partly supported by Grants-in-Aid from the Ministry of Education, Culture, Sports, Science and Technology (No. 19540240 KS)."
    },
    "0911/0911.5119.txt": {
        "abstract": "{}{}{}{}{} %5 {} token are mandatory  \\abstract % context heading (optional) % {} leave it empty if necessary {Low-luminosity   galaxies   are  known   to  outnumber  the   bright   galaxy population    in   poor  groups  and   clusters of   galaxies. Yet,   the  investigation    of   low-luminosity galaxy   populations outside     the  Local     Group remains    rare   and     the dependence      on different     group environments   is  still     poorly  understood.   Previous  investigations   revealed photometric  scaling relations for  early-type dwarfs and a  strong dependence of morphology with environment.} % aims heading  (mandatory) {The present   study  aims   to  analyse   the photometric    and   spectroscopic  properties  of    the low-luminosity  galaxy population   in the nearby, well-evolved   and    early-type   dominated NGC~5846    group  of    galaxies.  It  is     the     third     most    massive aggregate  of early-type   galaxies      after   the          Virgo     and         Fornax     clusters            in        the           local        universe. Photometric scaling relations  and   the  distribution    of   morphological  types  as well as  the characteristics of  emission-line galaxies are investigated.} % methods heading (mandatory) {Spectroscopically selected low-luminosity  group  members from   the  Sloan  Digital  Sky  Survey with   $cz<3000$ km   s$^{ -1}$   within  a radius  of  $2\\degr=0.91$ Mpc    around  NGC~5846  are   analysed.   Surface    brightness  profiles    of  early-type    galaxies are   fit  by    a S\\'{e}rsic  model $\\propto$  $r^{1/n}$. Star  formation  rates, oxygen abundances  and emission characteristics  are  determined for  emission-line galaxies.} % results heading (mandatory) {Seven new group members    showing     no     entry    in     previous  catalogues are  identified    in     the outer ($>$80     arcmin) parts of the  system. Several   photometric scaling  relations for  dEs as  well  as  the morphology-density  relation for  dwarf galaxies  are  reproduced. Moreover, the correlation between host  and  satellite morphologies   in poor groups  of  galaxies is   confirmed. Nucleated dwarfs  are  found  to be located in  the   vicinity  to the  brightest  ellipticals  in  the group.  Only two faint  galaxies show  fine structure. Emission-line  dwarfs show no interaction induced activity.} % conclusions heading (optional), leave it empty if necessary {} ",
        "introduction": "Redshift surveys   have  shown    that  most    galaxies in     the  nearby    universe are       located in   poor   groups      \\citep{tullyfisher}. These   aggregates     are   not only   made  up  of   the  ordinary   bright   Hubble  types  but  consist  also  of  a low-luminosity  dwarf  galaxy population outnumbering their brighter counterparts by far.  Previous  observations studying  the  distribution and  morphological properties of  such low-luminosity    galaxy     populations   have  mainly           focused            on        large-scale      structures,  e.g.   the          Virgo and Fornax       clusters  \\citep{sandagebinggeli,    ferguson}  representative of    high-density environments.   Complementary, poor   groups  serve as ideal  laboratory   to    study   the   environmental   dependence  of    low-luminosity    galaxy   populations   in   a  low-density  environment with galaxy  interactions,     mergers,  and  the    coalescence    of  individual   galaxies      as   the      driving physical processes.  However, detailed photometry and spectroscopy of low-luminosity galaxies is challenging. Indeed, only     the    Local      Group and      a     few      other nearby  galaxy  aggregates   have       been   investigated      in     detail to   analyse        the   connection   between    low-luminosity galaxies     to their environment \\citep{karachentsev,cote,jerjen,gruetzbauch}. Because   $\\Lambda$CDM  cosmology    suggests a   hierarchical  growth   of galaxies   in the   course  of   mergers, it is  obvious to  classify poor groups  based on their                 optical         morphologies \\citep{zabludoff}.    Spiral     dominated   aggregates   represent      unevolved systems  in  this scenario  while     aggregates   with     tidal    signatures   and   the  formation     of   tidal     dwarf  galaxies    mark   an intermediate    stage  \\citep{temporin}.  Groups     dominated   by   early-type    morphologies  are in      an    advanced   phase  of  coalescence. Dynamically,  these well-evolved aggregates     are  more    likely  to      show   virialised systems compared   to   groups with a   higher   spiral fraction.   Moreover, the  dwarf-to-giant     ratio   (DGR) is      expected   to   increase      with  time as        bright,  massive      galaxies are more  affected  by    mergers than    their  faint            counterparts  \\citep{zabludoffmulchaey}.     The number   of   faint   galaxies   in groups   and clusters    is   puzzling,  however. $\\Lambda$CDM  cosmology   suggests  more dwarfs   than observed, an  issue widely   referred to   as the   missing satellite  problem  \\citep{klypin}. \\\\  This paper focuses  on  the low-luminosity  galaxy population of  the  X-ray  bright  galaxy group  around  the   giant elliptical    NGC~5846      in the     Local  Supercluster.     The    group   was    first       catalogued  in       \\citet{devaucouleurs}.        Since  then, the          system has   appeared   in    many       similar  catalogues   \\citep{gellerhuchra,garcia}.  \\citet{zabludoffmulchaey}   identified   13  dwarfs    in    the central            region     of             the  system.      \\citet{mahdavi},     hereafter       [MTT05]   state        $251\\pm10$            group members   including           83 spectroscopically        confirmed     ones.   \\citet{carrasco}    have discovered  16   additional     low   surface brightness    dwarfs. The NGC~5846   group  of galaxies    is  located   at   a   distance of    26.1 Mpc   in   the  Virgo III   Cloud of   galaxies, a major   component of    the Local  Supercluster  \\citep{tully82}.  Being   an elongated  structure   ranging from   the   Virgo  cluster   out    of the    Supergalactic Plane,     this  cloud    contains   about   40     luminous   galaxies  with      the  NGC~5846    system     as   the   biggest aggregate in   the   region. The   NGC~5846     group    is    the   most   massive   of    only     three  dense  groups  (NGC~4274, NGC~5846     and M96)    in the Local Supercluster   that   are   dominated      by  elliptical  and    lenticular    galaxies,   being  the     third    most  massive aggregate  of  early-type     galaxies   (after the     Virgo   and       Fornax clusters)        in      the       local     universe.      Moreover, with           an  average          galaxy       density\\footnote{Densities       taken      from       the      Nearby        Galaxies        Catalog \\citep{tullyfisher},     see \\citet{tully88}   for  details.}  of $\\rho=0.78\\pm0.08$  Mpc$^{-3}$  compared        to $\\rho=0.49\\pm0.06$     Mpc$^{-3}$ for the  Local  Group       and $\\rho\\sim5$  Mpc$^{-3}$ for   the   central regions   of  the     Virgo  cluster, the    NGC~5846     system  presents also  one    of   the  densest galaxy     environments  found  for     poor   groups.     The  group   is  dominated   by   early-type   galaxies with a  massive  central elliptical  surrounded  by a   symmetric     X-ray halo.   The large   early-type  fraction,  the prevalence   of  dwarf  galaxies together  with  the strong   and extendend   diffuse  X-ray     emission marks    the   NGC~5846 group   as an    evolved   system.   The    group  is composed  of    two smaller  aggregates    around    the  two  brightest ellipticals       NGC~5846     and    NGC~5813. [MTT05]   determine       the group\u00b4s    virial      mass  to       be     8.3         $\\pm$   0.3  $\\times$10$^{13}$        M$_{\\odot}$         via  the  median      virial   mass estimator      proposed     by  \\citet{heisler}  and  derive a   velocity dispersion of  322    km s$^{-1}$.  The   goals  of this  work  are  to     study   new    faint  members  using  the  {\\bf   S}loan   {\\bf   D}igital  {\\bf  S}ky   {\\bf  S}urvey   (DR4) \\citep{adelman} to  extend  the  list  of  known  group members   and to   study  photometric  scaling relations, the distribution  of   morphological types,   as  well  as the  spectroscopic   properties of the   low-luminosity  galaxy population.   The  paper is    organized as  follows. Section  2 gives    a   brief    overview    on  the     sample   selection   and    data   analysis procedures.    Section 3 focuses     on  the  low-luminosity galaxy     population   and   presents   the    photometric  and   spectroscopic    data   of all    individual objects.  Finally,  the   results  are summarized  and discussed  in Section     4. A  group distance of 26.1 Mpc    based  on  the  work of    [MTT05] is     used    throughout  this paper corresponding to a distance  modulus  of $m-M=32.08$. The  systemic velocity   of the group is  assumed   to be represented    by NGC~5846    with   a value of   $v_{r}=1710$    km    s$^{-1}$.      Concerning             the              separation             between  dwarf and  bright    galaxies, an absolute  magnitude  of      $M_{B}=-16$    is    used \\citep{fergusonbinggeli}   except   otherwise  stated. Total magnitudes  presented  in  this work are   SDSS  model magnitudes \\footnote{SDSS    model magnitudes  are  based   on a   matched galaxy  model  as optimal measure   of the flux of a galaxy. The  code  fits    a  deVaucouleurs profile as     well as  an  exponential   profile to   the  two-dimensional  image. The best-fit model  is stored as  the model magnitude.}.  \\section {Sample Selection and Data  Analysis} In order to extend   the   existing sample  of  low-luminosity  galaxies   in the  NGC~5846  system,  we   have queried the  SDSS    for  all galaxies with $cz<3000$  km  s$^{-1}$   within  a   radius of  $2\\degr=0.91$   Mpc,   similar   to     the second turnaround    radius  $r_{2t}=1.85$\\degr   as presented by [MTT05]\\footnote{Bound  group members uncouple from  the   Hubble flow  at first  turnaround, i.e.   the zero-velocity  surface    around the  group  \\citep     {sandage},    begin to   collapse,  and    again   expand    to   a      radius of  second   turnaround,   finally  oscillating and exchanging kinetic    energy with   other     group  members  \\citep {bertschinger}.}.      Galaxies   beyond   these   limits   are  very  likely to  be    characterised as field  or    background       objects.    The  query  resulted    in   a   total    sample of   74   galaxies:   19  bright objects  and  55 dwarfs (see  Table~\\ref  {faint} and Figure~\\ref{chart})    with seven   objects found in     the outer   parts ($>1.33\\degr$)     of the  system  that show     no    entry in  previous  catalogues.  Though being   slightly   fainter    than  $M_{B}=-16$,  UGC 9751  and UGC 9760  are listed  as bright members  due to their classification as Scd and Sd in NED\\footnote{\\tt  http://nedwww.ipac.caltech.edu/}.  Corrected   FITS   frames   (bias    subtracted,   flat-fielded   and   cleaned    of  bright   stars)   have    been   extracted   from   the    SDSS for   these   55   dwarfs    in the $g^{\\prime}$,  $r^{\\prime}$ and   $i^{\\prime}$ bands\\footnote{See \\citet {fukugita} for a description of the  SDSS photometric system.}.   SDSS  model magnitudes have  been   adopted  from   the   database  for each  individual   galaxy  and transformed   to    the Johnson-Morgan-Cousins    $B$   band       via    transformation     equations   from    \\citet{smith}.    Absolute          magnitudes  were  derived assuming an average    extinction  value of $A_{\\rm{B}}=0.22$ from  \\citet{schlegel} for   the whole   galaxy population. An  isophote    analysis was carried out for    the  dE   subsample using    the  \\texttt     {ellipse}      task  within    the      \\texttt   {IRAF}   \\texttt   {stsdas} package to     fit  ellipses       to the  galaxy images      and   measure   the      deviations from   purely     elliptical  isophote shapes.   A  detailed description    of   the procedure   is  given  by    \\citet{jedrzejewski}.    The  analysis   resulted    in   surface  brightness profiles   and  the harmonic content    of     the isophotes. A   \\citet{sersic}   law   $\\mu   (r)  =   \\mu  _e  +  1.086~b_n~({\\left(  {r/r_e }   \\right)^{1/n}   - 1})$ with an   effective      radius  $r_{e}$, an effective     surface     brightness $\\mu_{e}$   and a shape   parameter       $n$   as   free parameters was   fit     to     each    object.   The quantity $b_{n}$   is   a   function   of  $n$   and  chosen  so  that   half     the galaxy     luminosity is     located    within     the    effective  radius    \\footnote{The  exact    relation     between  $n$     and    $b_{n}$    is   derived       by $\\Gamma(2n)=\\gamma(2n,b_{n})$,   where  $\\Gamma(a)$   and $\\gamma(a,x)$    are the complete  and  incomplete gamma   functions, respectively. A   good approximation    is $b_{n}=1.9908n-0.3118$ for $0.5\\le  n \\le 2$, which    has  been used  in  this work.}.    The     fitting was      carried    out using          a   Levenberg-Marquardt    algorithm  until  a      minimum      in     $\\chi^{2}$    was achieved,       yielding the      parameters $n$,  $r_{e}$,       and       $\\mu_{e}$   with      the corresponding   uncertainties.       The innermost   parts    of the   surface     brightness profiles      ($\\le1.4$       arcsec)     have   not  been  taken     into           account     for  the           fitting   procedure     due     to seeing.  Central        surface            brightnesses $\\mu_{0}\\hspace{-1pt}_{_{B}}$    have  been  derived   from     the fit.       The     results of     the       isophote    analysis         are presented      in    Table~\\ref{surface}. S\\'{e}rsic   models    have  been subtracted    from   the original       images      yielding     residual frames    which are   investigated      for    photometric   substructures.  In addition  to     the photometric   data,    wavelength    and      flux calibrated  spectra  (sky     subtracted  and     corrected   for  telluric absorption)  have  been examined    for  all       sample  galaxies.     The     spectra cover    a    wavelength   range    of  $\\lambda\\lambda3800-9200$  \\AA~and have  been checked    for    spectral  lines    in  both   emission and    absorption.   Identified  spectral    lines     were     fit  by    Gaussians yielding central  wavelengths     and   full      widths     at half maximum    (FWHM).    The  derived     central  wavelengths    were  subsequently  used to verify SDSS redshifts.  \\section {The Low-Luminosity Galaxy Population}  \\begin{figure} \\centering \\includegraphics[width=\\columnwidth]{chart.eps} \\caption{\\label{chart}Projected spatial   distribution  of    NGC~5846    group  members     as  studied   in this     work. Filled  triangles indicate bright  members, while open  ones  represent    the low-luminosity  galaxies.   Numbers refer  to  the  galaxy identification   of  Table~\\ref {faint}. The circle shows the 2\\degr (0.91 Mpc) query radius.} \\end{figure}  \\begin{figure} \\centering \\includegraphics[width=\\columnwidth]{cmd.eps} \\caption{\\label{cmd}Colour-magnitude diagram of  NGC~5846   group    members. Numbers  as    in Figure~\\ref   {chart}.  The    dashed  line  indicates the red sequence   linear fit to   bright galaxies and  early-type  dwarfs: $r^{\\prime}$$-$$i^{\\prime}=0.75(\\pm  0.05)-0.027(\\pm0.003)$. The  vertical line separates dwarfs from bright galaxies.} \\end{figure}  \\begin{figure} \\centering \\includegraphics[width=\\columnwidth]{n-MB.eps} \\caption{\\label {sersic}S\\'{e}rsic shape   parameter  versus  absolute  blue   magnitude of   early-type  galaxies  in  the   NGC~5846   group.   Open symbols   represent  dwarf  ellipticals  while   filled  triangles  show  bright  ellipticals.  The  horizontal  lines  indicate exponential  and de Vaucouleurs r$^{1/4}$ laws.} \\end{figure}   Table~\\ref   {faint}  lists    the  NGC~5846  galaxy   group   sample  as    studied   in        this work.   Morphological   types   for dwarfs  have been  classified through visual  inspection   of  SDSS   images  and spectra. Objects  with  a  smooth   light  distribution  and exponential profiles were   classified   as dEs.  The   dE ellipticity     was   determined     from  the      isophote  analysis. Dwarfs showing a blue patchy  appearance and  emission lines   were classified  as dIrrs.      Morphological  types     for  bright  galaxies  have      been  taken   from NED.        Dwarfs are        identified       with   increasing     projected     radial    distance  from NGC~5846.  We  extend      previous group  member lists  with      seven    new   objects     in   the     outer    parts        ($>1.33\\degr$)   of   the system showing       no   entry         in previous catalogues.     Despite       exhibiting   slightly  brighter      values       than    $M_{B}=-16$, some galaxies  have     been classified   as  dwarfs   due to       their    nearly     exponential   surface   brightness     profiles     or        their compactness.   35   dwarfs  ($\\sim$63\\%)   have     been    identified            as         dEs          according                to            their optical   appearance       and  red   colours  $(B-V)_{0}\\simeq0.85$.     18   objects    ($\\sim$32\\%)    are      dwarf     irregulars      showing a patchy,   blue    optical  appearance      with emission   lines  indicating      ongoing       star   formation. Three   galaxies  ($\\sim$1\\%) reveal photometric fine    structure and  are in  the transition  to  bright     galaxies.           11   dwarfs    ($\\sim$31\\%  of        the      dEs)  are found to be nucleated on basis of the analysis of the surface brightness profiles.  Our  list   of  nucleated  dwarfs  is   nearly  identical to  the  list  obtained by  [MTT05]  with   exception of  only   three   objects classified differently. Taking      the        Schechter  luminosity    function derived by [MTT05],  we   conclude   that  with    an      absolute    magnitude of   $M_{B}=  -13.32$     for  our faintest     object,     the  sample of  our work comprises  at  most 34\\%    of  the   total number  of   galaxies expected in   the  system down to  a limiting magnitude of  $M_{B}=-10$. Our sample yields  a dwarf-to-giant-ratio (DGR)  of  $\\sim$3 as  lower limit. Following  the photometric dwarf  classification of [MTT05],  this value could rise  to $\\sim$12, confirming that the NGC 5846 group is indeed one  of the richest groups in the Local Supercluster.  \\addtocounter{table}{1}  The  projected   spatial   distribution   of     all    investigated   galaxies  is  illustrated     in   Figure~\\ref  {chart}.   The     distribution indicates     two   substructures     of    the  system around     the    two brightest    ellipticals   NGC~5846   and    NGC~5813.  A   third, however    much  smaller    subgroup   can   be  found around UGC   9760,  which   in contrast    to the    NGC~5846  and NGC~5813    aggregates shows no    noteworthy     X-ray      counterpart.  We present  on Figure~\\ref    {cmd}    a   colour-magnitude     diagram   for      the    studied  group galaxy     population.  The  diagram   indicates     the   prevalence  of   an       early-type dwarf morphology     and  the  continuation    of    a red sequence similar   to    that  found  in    galaxy  clusters \\citep{gladders}   into  the dwarf regime. Figure~\\ref  {sersic} shows the  $n-M_{B}$ relation  consisting of  our dE sample  as  well as  the bright  galaxies classified  as   ellipticals   in    NED.   S\\'{e}rsic  shape     parameters of     the   low-luminosity    galaxy    population      show  a mean   value of  $n=1.19$  with  a scatter of  $\\sigma_{n}=0.23$  representing nearly exponential       surface     brightness    profiles      as      expected            for dwarf    ellipticals.   Three dwarfs (N5846\\_10,  N5846\\_15, N5846\\_43) show comparatively large error bars of the shape parameter due  to their faintness or contamination of the surface brightness profile by a central nucleus.  Our   data    confirm the  trend  of   more  luminous        galaxies showing    greater  shape      parameters,  thus    a  steeper surface  brightness     profile  than    the  less  luminous  objects.  Interestingly     there are   no   objects     with   intermediate  S\\'{e}rsic parameters   between   $2<n<3$.   NGC~5813 shows   a   comparatively high shape parameter  around $n\\sim8$.   \\subsection       {Scaling relations} When  addressing   the  question    of    the influence of the environment   on   the   evolution    of   galaxies,    scaling        relations     of the photometric properties      of         early-type   galaxies     are      a   useful        tool.  However   the      environment   does  not seem to affect  all early-type  systems   likewise.  The fundamental   plane   of  bright ellipticals  is  found  as     much  independent   of environment as possible, while dwarf   ellipticals show a  stronger    variation  \\citep{nieto,petersoncaldwell}. Elliptical galaxies and   bulges     are   known to show     a   tight   correlation      in      the    $\\mu-M_{B}$     plane   with      higher     surface    brightnesses    indicating     smaller total luminosities    \\citep{kormendy}.  Dwarf     galaxies    exhibit        an   opposite          trend,  being        a        distinct      class of  objects     with      possibly  different   formation      and        evolution   histories.   \\citet{binggelijerjen},  hereafter  [BJ98]     have investigated      the    $\\mu-B$  plane     for      Virgo   cluster   early-type     dwarfs      and      confirm   this  trend.          Figure~\\ref {scaling}(a)      shows         the  $\\mu_{0}\\hspace{-1pt}_{_{B}}-M_{B}$ plane       for our       dEs and     the    relationship   taken from   the [BJ98](dash-dotted line) sample   of Virgo    dwarf ellipticals (assuming  a  Virgo distance  of 16.5   Mpc  corresponding  to a  distance modulus  of $\\sim$31.09).  Remaining relations  \\ref     {scaling}(b)-(f)  focus   on        $M_{B}$,  $\\log    n$,  $\\mu_{0}\\hspace{-1pt}_{_{B}}$,   $\\log r_{0}$  and     ellipticity $\\epsilon$.     S\\'{e}rsic     shape   parameters   used    here    should  not     be    mixed     up    with       those of   [BJ98]: $n=1/n_{\\rm{BJ98}}$.  The  parameter  $r_{0}$ is   the  scale    radius of     the S\\'{e}rsic    model: $I(r)=I_{0}\\exp[-r/r_{0}]^{1/n}$. Table~\\ref    {surface}  presents     the corresponding  values  of    the  isophote  analysis. Objects    with    larger  errors  than     the   data itself   are not listed.  Our data  in  the  $\\mu_{0}\\hspace{-1pt}_{_{B}}-M_{B}$ plane        are in  agreement   with  the work     of [BJ98].   One   object  falling  outside the $2\\sigma$   limit   is     N5846\\_01   with   a  comparatively   too    high   surface  brightness   with    respect   to   its   luminosity  (see section \\ref{individual}).  In the   $\\log n-M_{B}$   plane,   our  data show   a larger   scatter $\\sigma_{\\rm{log n}}=0.19$. Compared to  the  dwarf sample  of the Virgo cluster, our dwarfs match  the Virgo dwarf  relation, however. Our sample also confirms the  trend  between shape  parameter  $n$ and  scale radius $r_{0}$ with smaller   objects exhibiting  higher shape  parameters. Again N5846\\_01 is  the  only object in  our   sample  to  fall off    the trend.  The     scale radius    $r_{0}$ shows    also  a    strong correlation     with the   central surface    brightness as    seen   in the $\\mu_{0}\\hspace{-1pt}_{_{B}}-\\log  r_{0}$  plane. The  NGC~5846  dwarfs  match with   the  relation obtained  for the Virgo   dwarfs with a  lower scatter in magnitude   ($\\sigma=0.83$mag  arcsec$^{-2}$)  than   [BJ98]  who  derive   a  value  of   $\\sigma=1.25$mag  arcsec$^{-2}$.   Similarly  to [BJ98]    we  also derived  the   best    fitting linear  combination  between     shape parameter  $n$, central   surface  brightness   $\\mu_{0}$ and absolute $B$   band magnitude with a  scatter of  $\\sigma=0.59$mag. Finally,   we  also  checked for   a  correlation   between  shape   parameter $n$ and  ellipticity. There is a  slight trend  of  flattened   objects showing shallower surface brightness profiles.  \\begin            {figure*} \\centering \\includegraphics  [width=450pt]{relations1sigma.eps} \\caption          {\\label  {scaling}  Photometric  scaling relations   of the   low-luminosity  galaxy  population in   the  NGC~5846 group.  (a)-(d): dash-dotted   lines are  the   relations    of the Virgo   cluster   early-type     dwarfs  from [BJ98]  assuming  a   distance to  the  Virgo cluster of  16.5 Mpc.  The  scatter of   our data  with  respect  to the   [BJ98] relations  is  given  in the   lower left  corner  of each panel.  $1\\sigma$ deviations are    indicated as   dashed lines.   (a):  $\\mu_{0}\\hspace{-1pt}_{_{B}}-M_{B}$    plane. N5846\\_01,  classified as  ultra-compact dwarf by [MTT05] is  located  outside   $2\\sigma$. (e):  best-fitting linear     combination    of    shape parameter    $n$ and   central   surface brightness $\\mu_{0}$  with  respect      to   absolute    magnitude   $M_{B}$.    N5846\\_01    was  excluded    from   the     fit   and    is  not  shown in the diagram. The dotted line  indicates  identity.  (f):   S\\'{e}rsic  shape  parameter   versus  ellipticity  . Objects      with  high  ellipticity tend to   have shallower surface    brightness profiles. The dotted  line indicates an exponential profile.} \\end              {figure*}  The     structural    properties     of    early-type      systems    can     furthermore    be     investigated     in     the  \\citet{hamabekormendy} relation,   a photometric    projection  of     the fundamental     plane   of     galaxies relating     the  logarithm     of  effective radii    and effective   surface   brightnesses  linearly  \\citep{ziegler,diserego}.   This relation   divides    early-type    systems     in   \\emph{regular} and \\emph{bright}   classes    separated     at   the effective   radius    of    $r_{e}\\simeq  3$   kpc \\citep{capaccioli}.   In  hierarchical  evolution scenarios,  regular   galaxies  are   thought  to     be  the    progenitors   of    bright ones    evolving    through    succesive   mergers   along the Hamabe-Kormendy    relation   \\citep{capelato}. However,  recent    numerical   simulations  of   \\citet{evstigneeva}  suggest   that    low  mass systems like  dwarf galaxies    can  only     follow this  scenario  with     a  large   amout    of    dissipation  involved.   Figure~\\ref  {hamabe} shows  the Hamabe-Kormendy    relation  for  the   faint   galaxy     population       of       our    work.      Open      squares     are     dwarfs from  X-ray  dim  and   X-ray     bright       groups   studied   by  \\citet{khosroshahi}.  All   faint      galaxies  of    our    sample   represent regular  systems   and  show  properties   similar to   the dwarfs  from  \\citet{khosroshahi}   suggesting no  strong  difference  of the   structural properties  of the  early-type dwarfs  from the NGC~5846   group compared  to  other group environments.  Again, N5846\\_01 is the  only dwarf  clearly separated from the main  cloud of dEs.  \\begin{figure} \\centering \\includegraphics[width=230pt]{hamabe-kormendy.eps} \\caption{\\label{hamabe}$\\mu_{e}-\\log(r_{e})$   plane for   early-type   galaxies.   Effective  surface  brightnesses   are    shown in  the    Johnson $B$  band.    The vertical   solid   line   separates  bright   and   regular    ellipticals      while   the   dash-dotted   line    represents   the Hamabe-Kormendy relation   for elliptical  galaxies and bulges.  Black triangles   indicate  the  NGC~5846  faint  galaxy sample.  Open squares  refer to the  groups studied by \\citet{khosroshahi} for comparison.} \\end{figure}  \\begin{table*} \\begin{minipage}[t]{\\textwidth} \\caption{\\label{surface}Surface  photometry    of  the    low-luminosity  galaxy     sample.  Surface    brightnesses  are    shown  in     magnitudes arcsec$^{-2}$, effective radii  in arcseconds.  Numbers  in parantheses indicate errors  of the   last significant digit.  The quality of the fits  is shown in  the last three columns.} \\centering \\renewcommand{\\footnoterule}{} \\begin {tabular}{lcccccccccc} \\hline\\hline \\multicolumn {1}{c}{galaxy}   &  $n_{i^{\\prime}}$ &   $n_{r^{\\prime}}$ &   $n_{g^{\\prime}}$  &  $\\mu_{0}\\hspace{-1pt}_{_{B}}$           &   $r_{e_{i^{\\prime}}}$  &  $r_{e_{r^{\\prime}}}$ &  $r_{e_{g^{\\prime}}}$  &  $\\chi_{\\nu,i^{\\prime}}^{2}$  &  $\\chi_{\\nu,r^{\\prime}}^{2}$  &  $\\chi_{\\nu,g^{\\prime}}^{2}$    \\\\ \\hline N5846\\_01                     &      1.84(4)      &      1.82(3)       &       1.64(3)       &                 17.1(2)                  &         1.6(4)          &        1.7(3)         &         1.6(3)         &             0.00043             &             0.00025             &              0.00055              \\\\ N5846\\_02                     &      1.12(3)      &      1.06(2)       &       1.10(2)       &                 22.6(1)                  &         6.9(8)          &        6.9(6)         &         7.0(5)         &             0.00528             &             0.00317             &              0.00229              \\\\ N5846\\_03                     &      1.08(1)      &      1.14(2)       &       1.11(2)       &                 21.8(1)                  &         5.1(2)          &        5.7(4)         &         5.9(5)         &             0.0018              &             0.00325             &              0.00221              \\\\ N5846\\_04                     &      1.43(2)      &      1.53(2)       &       1.30(2)       &                 21.63(9)                 &         8.4(9)          &        9.5(9)         &         7.6(5)         &             0.00159             &             0.00064             &              0.00099              \\\\ N5846\\_05                     &      1.31(2)      &      1.27(1)       &       1.14(2)       &                 21.79(9)                 &         6.1(5)          &        6.3(3)         &         5.7(3)         &             0.00288             &             0.00109             &              0.00217              \\\\ N5846\\_09                     &      0.91(9)      &      1.09(9)       &       1.3(1)        &                 22.3(6)                  &          9(3)           &         10(4)         &         12(7)          &             0.00681             &             0.00957             &              0.0078               \\\\ N5846\\_10                     &      1.4(3)       &       2.0(7)       &       1.9(8)        &                  23(4)                   &            -            &           -           &           -            &             0.00468             &             0.00335             &              0.00513              \\\\ N5846\\_11                     &      1.01(3)      &      1.15(3)       &       1.34(5)       &                 20.7(2)                  &         9.0(9)          &         10(1)         &         10(2)          &             0.00824             &             0.00925             &              0.01286              \\\\ N5846\\_13                     &      1.22(2)      &      1.33(1)       &       1.21(1)       &                 21.52(6)                 &         11.3(8)         &        12.3(5)        &        10.9(7)         &             0.00286             &             0.00065             &              0.00066              \\\\ N5846\\_15                     &      1.8(6)       &      1.20(8)       &       1.5(1)        &                 22.2(6)                  &            -            &         10(3)         &         13(7)          &             0.0038              &             0.00293             &              0.00142              \\\\ N5846\\_16                     &      1.42(2)      &      1.45(2)       &       1.43(2)       &                 21.8(1)                  &         10.0(9)         &        10.1(8)        &         10(1)          &              0.001              &             0.00069             &              0.00106              \\\\ N5846\\_17                     &      0.97(6)      &      1.11(6)       &       0.95(3)       &                 21.9(2)                  &          6(1)           &         7(2)          &         5.6(7)         &             0.00375             &             0.00238             &              0.00143              \\\\ N5846\\_21                     &      1.37(2)      &      1.21(1)       &       1.35(1)       &                 20.79(6)                 &          11(1)          &        9.7(5)         &        10.8(6)         &             0.00332             &             0.00085             &              0.0015               \\\\ N5846\\_22                     &      1.23(2)      &      1.23(2)       &       1.18(2)       &                 21.8(1)                  &         5.2(4)          &        5.3(3)         &         5.3(4)         &             0.00354             &             0.00212             &              0.00356              \\\\ N5846\\_23                     &      1.41(3)      &      1.31(3)       &       1.22(2)       &                 21.9(1)                  &          9(1)           &         8(1)          &         8.7(7)         &             0.00163             &             0.00327             &              0.00156              \\\\ N5846\\_24                     &      0.98(2)      &      0.93(1)       &       0.90(2)       &                 21.72(9)                 &         8.0(5)          &        7.7(4)         &         7.7(5)         &             0.00107             &             0.00118             &              0.0026               \\\\ N5846\\_25                     &      1.09(1)      &      1.02(1)       &       0.97(1)       &                 21.18(5)                 &         6.5(3)          &        6.3(3)         &         6.4(2)         &             0.00508             &             0.00374             &              0.0028               \\\\ N5846\\_27                     &      0.91(3)      &      0.81(2)       &       0.84(4)       &                 22.4(2)                  &         6.9(7)          &        6.0(5)         &         6.3(9)         &             0.00099             &             0.00075             &              0.0024               \\\\ N5846\\_28                     &      0.80(1)      &      0.81(1)       &       0.89(2)       &                 21.05(9)                 &         6.5(4)          &        6.0(3)         &         5.8(4)         &             0.00345             &             0.00289             &              0.00406              \\\\ N5846\\_31                     &      1.02(6)      &      1.10(5)       &       1.07(8)       &                 22.2(4)                  &          8(2)           &         9(2)          &          8(2)          &             0.00173             &             0.0018              &              0.00227              \\\\ N5846\\_32                     &      0.99(2)      &      0.99(1)       &       1.05(2)       &                 21.83(9)                 &         10.9(9)         &        11.3(7)        &        11.9(9)         &             0.0016              &             0.00131             &              0.00168              \\\\ N5846\\_33                     &      0.78(2)      &      0.79(2)       &       0.83(2)       &                 23.3(1)                  &          11(1)          &         12(1)         &         13(1)          &             0.00362             &             0.00173             &              0.00329              \\\\ N5846\\_34                     &      1.20(2)      &      1.23(3)       &       1.13(2)       &                 22.5(1)                  &         9.5(8)          &         10(1)         &         9.4(9)         &             0.00161             &             0.00255             &              0.00225              \\\\ N5846\\_35                     &      1.28(4)      &      1.26(3)       &       1.20(2)       &                 22.2(1)                  &          7(1)           &        7.1(9)         &         7.3(7)         &             0.00541             &             0.00358             &              0.00206              \\\\ N5846\\_37                     &      1.3(1)       &       1.1(1)       &       1.2(1)        &                 22.9(6)                  &          11(6)          &         7(2)          &         10(6)          &             0.00111             &             0.00376             &              0.00245              \\\\ N5846\\_38                     &      1.08(5)      &      1.02(4)       &       1.21(5)       &                 21.8(2)                  &         4.0(6)          &        3.8(5)         &         4.4(7)         &             0.00575             &             0.00515             &              0.00399              \\\\ N5846\\_39                     &      0.74(6)      &      0.99(3)       &       1.19(9)       &                 22.1(4)                  &          6(1)           &        6.1(6)         &          7(2)          &             0.0024              &             0.0019              &              0.00173              \\\\ N5846\\_40                     &      1.15(7)      &      1.38(5)       &       1.18(6)       &                 22.5(3)                  &          10(3)          &         13(3)         &          9(2)          &             0.00661             &             0.0016              &              0.00416              \\\\ N5846\\_43                     &      1.3(2)       &       1.8(4)       &       1.2(2)        &                  23(1)                   &          6(5)           &           -           &          4(3)          &             0.00621             &             0.00222             &              0.00715              \\\\ N5846\\_46                     &      1.36(4)      &      1.36(5)       &       1.45(5)       &                 19.6(2)                  &         3.4(6)          &        3.3(7)         &         3.3(8)         &             0.0148              &             0.0148              &              0.00764              \\\\ N5846\\_47                     &      0.87(2)      &      0.82(2)       &       0.85(2)       &                 22.3(1)                  &          11(1)          &        10.6(9)        &         11(1)          &             0.00337             &             0.00217             &              0.00371              \\\\ N5846\\_50                     &      1.17(4)      &      1.27(3)       &       1.35(4)       &                 21.6(2)                  &         5.5(8)          &        5.8(8)         &          6(1)          &             0.00782             &             0.00516             &              0.00784              \\\\ N5846\\_51                     &      1.05(2)      &      1.21(2)       &       1.19(2)       &                 20.9(1)                  &         4.9(3)          &        5.5(5)         &         5.4(4)         &             0.00177             &             0.0034              &              0.00283              \\\\ N5846\\_52                     &      1.14(4)      &      1.78(9)       &       1.41(5)       &                 21.7(3)                  &          8(1)           &         14(8)         &         10(3)          &             0.00397             &             0.00245             &              0.00322              \\\\ N5846\\_56                     &      1.10(2)      &      1.09(2)       &       1.10(2)       &                 22.12(9)                 &          15(1)          &         15(1)         &         15(1)          &             0.00268             &             0.00264             &              0.00186              \\\\ \\hline \\end{tabular} \\end{minipage} \\end{table*}  \\begin{table*} \\begin{minipage}[t]{\\textwidth} \\caption{\\label{emission}Spectral properties of emission-line dwarfs. Fluxes present the integrals of the fitted  Gaussians and are given in units  of 10$^{-17}$ ergs s$^{-1}$ cm$^{-2}$. H$\\alpha$ and H$\\beta$  fluxes have  been corrected for extinction. Numbers in parentheses indicate errors  of the last  significant digit.} \\centering \\begin{tabular}{lcccccccccc} \\hline\\hline &    H$\\beta$     & $[$O{\\scriptsize~III}$]$ & $[$O{\\scriptsize~III}$]$ & $[$N{\\scriptsize~II}$]$ &    H$\\alpha$    & $[$N{\\scriptsize~II}$]$ & $[$S{\\scriptsize~II}$]$ &  $[$S{\\scriptsize~II}$]$  &           SFR           &  $12+\\log$O/H  \\\\ \\multicolumn{1}{c}{\\raisebox{1.5ex}[-1.5ex]{galaxy}}  & $\\lambda$4861.3 &     $\\lambda$4958.9      &     $\\lambda$5006.8      &     $\\lambda$6548.1     & $\\lambda$6562.8 &     $\\lambda$6583.4     &     $\\lambda$6717.0     &      $\\lambda$6731.3      &  M$_{~\\odot}$yr$^{-1}$  &     [dex]      \\\\ \\hline N5846\\_06                                             &      348.1      &           45.8           &          128.4           &          14.0           &      992.0      &          51.7           &          66.6           &           47.7            &        0.0064(3)        &     8.2(1)     \\\\ N5846\\_07                                             &      619.6      &            -             &           4.2            &          62.5           &     1765.9      &          173.7          &          107.2          &           81.3            &        0.0114(4)        &     8.4(1)     \\\\ N5846\\_08                                             &        -        &            -             &            -             &           7.7           &      51.0       &           7.2           &          12.5           &            8.0            &            -            &     8.5(1)     \\\\ N5846\\_09                                             &      79.0       &            -             &           18.9           &           3.1           &      225.1      &           8.5           &          21.4           &           16.5            &            -            &     8.1(1)     \\\\ N5846\\_12                                             &      300.1      &           93.9           &          270.4           &           7.9           &      855.4      &          16.6           &          47.5           &           26.7            &        0.0055(3)        &     7.8(2)     \\\\ N5846\\_14                                             &      409.9      &           71.9           &          212.0           &          10.1           &     1168.3      &          47.8           &          89.3           &           64.0            &        0.0075(4)        &     8.1(2)     \\\\ N5846\\_15                                             &        -        &            -             &            -             &            -            &      24.2       &            -            &            -            &             -             &            -            &       -        \\\\ N5846\\_17                                             &        -        &            -             &            -             &            -            &      46.7       &           9.1           &           6.5           &            6.8            &            -            &     8.6(1)     \\\\ N5846\\_18                                             &        -        &            -             &            -             &            -            &      30.7       &            -            &           7.7           &            7.4            &            -            &       -        \\\\ N5846\\_19                                             &      43.7       &            -             &           19.3           &           8.6           &      124.6      &          14.0           &          21.3           &           12.3            &            -            &     8.4(1)     \\\\ N5846\\_20                                             &      171.1      &           5.4            &           39.9           &          17.5           &      487.5      &          42.4           &          51.4           &           37.5            &            -            &     8.4(1)     \\\\ N5846\\_27                                             &        -        &            -             &            -             &            -            &      38.8       &            -            &           7.9           &             -             &            -            &       -        \\\\ N5846\\_28                                             &      298.8      &           15.6           &           50.1           &          21.2           &      851.5      &          67.1           &          64.5           &           52.6            &        0.0055(3)        &     8.3(1)     \\\\ N5846\\_33                                             &        -        &            -             &            -             &            -            &      29.7       &            -            &            -            &             -             &            -            &       -        \\\\ N5846\\_38                                             &      103.8      &           12.1           &           40.8           &           8.6           &      295.8      &          20.5           &          16.2           &           14.4            &            -            &     8.3(1)     \\\\ N5846\\_39                                             &      34.0       &            -             &           10.4           &            -            &      96.8       &           9.3           &          19.2           &           12.7            &            -            &     8.4(1)     \\\\ N5846\\_40                                             &        -        &            -             &            -             &            -            &      42.8       &            -            &           9.3           &            5.9            &            -            &       -        \\\\ N5846\\_41                                             &     2576.2      &          1222.0          &          3664.0          &          20.6           &     7342.2      &          68.1           &          182.0          &           127.8           &        0.0473(9)        &     7.6(2)     \\\\ N5846\\_42                                             &      914.9      &          475.5           &          1385.0          &          10.9           &     2607.5      &          31.4           &          76.4           &           48.9            &        0.0168(5)        &     7.7(2)     \\\\ N5846\\_46                                             &      223.7      &           15.3           &           43.8           &          16.6           &      637.6      &          51.7           &          53.9           &           43.1            &        0.0041(3)        &     8.3(1)     \\\\ N5846\\_49                                             &      28.9       &            -             &           8.7            &            -            &      82.3       &           2.4           &           5.3           &            7.7            &            -            &     8.0(1)     \\\\ N5846\\_50                                             &        -        &            -             &            -             &            -            &      32.2       &            -            &            -            &             -             &            -            &       -        \\\\ N5846\\_51                                             &      62.6       &           8.9            &           38.8           &           8.6           &      178.4      &          15.8           &          24.5           &           20.9            &            -            &     8.4(1)     \\\\ N5846\\_54                                             &      290.9      &            -             &           45.4           &          31.4           &      829.2      &          89.8           &          75.4           &           57.6            &        0.0053(3)        &     8.4(1)     \\\\ N5846\\_55                                             &      461.1      &            -             &           17.9           &          47.0           &     1314.2      &          140.1          &          81.2           &           61.9            &        0.0085(4)        &     8.4(1)     \\\\ N5846\\_56                                             &      40.6       &            -             &           15.6           &          12.0           &      115.7      &          10.5           &          14.1           &           10.2            &            -            &     8.3(1)     \\\\ \\hline \\end{tabular} \\end{minipage} \\end{table*}  \\subsection       {\\label{individual}Comments on individual objects} Most of  the  investigated  dwarfs do not  show any  photometric    or  spectroscopic peculiarities. A    few galaxies  in  the   sample  are    worth a   closer  look,   however. Figure~\\ref {fine}   presents  all    objects  which evidently exhibit  fine structure in  their optical appearance  with $r^{\\prime}$ band images and  residual  frames  shown  for each  galaxy.  These objects  lie  furthermore in the   transition   to bright galaxies due to   their relatively large angular diameter.\\\\  {\\bf N5846\\_01}  is the  closest dwarf elliptical to NGC~5846  with a   much higher   central  surface brightness than the  rest   of our dwarfs.   It is furthermore the most reddish galaxy  from our  low-luminosity  galaxy sample.  With    an average   effective  radius   of   only  $r_{e}\\sim$  200 pc  in   all   studied passbands,  this     object   is also the   most   compact     of    our  sample.   Due    to   the     visual  appearance  and the     close proximity    to NGC~5846,    [MTT05] suggested   that N5846\\_01  has very     likely   been  affected  by    tidal  stripping   and  can therefore rather    be classified   as ultra  compact   dwarf.  This   idea     is also       supported by     the   fact    that the    nearby object NGC~5846A shows  a similar   morphological appearance.  In fact, NGC~5846A is   one of the   extremely  rare   \"compact   ellipticals\" like M32,   the prototypical  object   of  this    class. Thus,  NGC5846\\_01   fits  perfectly into   the picture   that  the    structural differences    of  compact ellipticals and    ordinary dEs     result   from   the fact   that  the compacts  formed  within  the potential  well    of another  massive   galaxy whereas  dEs  evolved  as  isolated   systems \\citep{burkert}.  {\\bf N5846\\_07} is  classified   as S0/a  galaxy by  [MTT05]   and  originally considered    as almost certainly  background  object in  their  sample due to its   optical appearance. Spectroscopy confirmed  the group membership  however. Assuming this distance, the   galaxy has a projected  diameter of  $D_{25}\\sim7$ kpc which    locates   the  galaxy  in  the   transition between  bright and    dwarf galaxies. The spectrum   of  N5846\\_07 reveals strong H$\\alpha$ emission  together with the  weaker H$\\beta$, [NII]   and  [SII] features    indicating ongoing star  formation.   Two spiral    arms are  identified  in   the   optical  image which  could    also be interpreted as  signatures of ongoing interaction  with the nearby irregular  dwarf N5846\\_08,  only 21kpc away.  {\\bf  N5846\\_41/42}  are    two   large  HII    regions   of    the    irregular  dwarf  KKR15     \\citep{huchtmeier, karachentseva}    classified  as individual  galaxies in  SDSS.  The host  galaxy shows an  extremely blue  colour,  the  majority of  light contained in  the two   HII  regions  with the brighter  one, N5846\\_41,  being the  bluest object  of our   sample.  Due   to the high concentration  of light, N5846\\_41/42 present  by far the highest S/N ratio of  all spectra  in   our sample. Besides  the usual H$\\alpha$,   H$\\beta$,  [OIII], [NII]  and [SII] features  both  spectra   also show    the  weaker  Balmer lines   H$\\gamma$,     H$\\delta$,  H$\\xi$    along  with    the    more   uncommon   HeI     $\\lambda$5875.6,  [Ne    III] $\\lambda$3868.7   and  ArIII $\\lambda$7135.8  transitions. Of    all emission-line  galaxies  in   our sample,   N5846\\_41/42 show  the  lowest oxygen abundance   with an average value of 12+log(O/H)=7.68 (see section \\ref{emissiondwarfs}).  {\\bf N5846\\_54} and  {\\bf N5846\\_55}  were   both classified as  S0  galaxies in    the [MTT05] sample   and are the   only objects  of  our work that show photometric substructures. With    a projected linear  extent  of  $D_{25}\\sim4$ and  7  kpc  these objects    can,  similarly  to  N5846\\_07, be placed in  the  transition  to dwarf   galaxies.  The  surface brightness  profile  of N5846\\_55 exhibits a  rather  flat gradient  up  to  14 arcsec, wherefrom   the  light    distribution  falls    much  steeper,   indicating    a  possible  bar   feature  \\citep{gadotti}.  Spectroscopically, both galaxies show ongoing  star formation.  \\begin            {figure*} \\centering \\includegraphics  [width=420pt]{fine.eps} \\caption          {\\label   {fine} Modeling   of   faint   galaxies with     significant fine     structure. Figures    show SDSS    $r^{\\prime}$ band images (upper    panel)  and     residual  images     (lower panel) of N5846\\_07, N5846\\_54  and N5846\\_55.  Since  no S\\'{e}rsic model  was  obtained for N5846\\_07, gaussian smoothing  was applied to the galaxy.} \\end              {figure*}  \\begin{figure} \\centering \\includegraphics[width=\\columnwidth]{dichte.eps} \\caption{\\label{clustering}Clustering  properties for  both  early  and late-type  dwarfs of    the NGC~5846 system.  Contours  are   linearly spaced and   show    the projected     number    density  of   the    distinct morphological   populations.   Triangles    refer  to  bright  group  members. The values give the number of galaxies  per 0.25 square degree.} \\end{figure}  \\begin{figure} \\centering \\includegraphics[width=\\columnwidth]{dEN.eps} \\caption{\\label{nuclei}Distribution of nucleated dwarfs  in  the  NGC~5846   group. Black triangles  are   bright group  members, open  triangles  dEs without nucleus. Diamonds show nucleated  dwarfs, concentrated around NGC~5846 and NGC~5813 (circles indicate a radius of 260kpc).} \\end{figure}  \\subsection{Spatial distribution of morphological types} We have studied  the morphology distribution  for the  low-luminosity  galaxy population in  the NGC~5846  system  for both  the early- and  late-type populations. Figure~\\ref {clustering}   illustrates the     clustering for     both     the  dE     and  dIrr  samples    with   the projected  number density   shown   separately   for    each    morphological  type.   The   diagram    displays   the       same  field   of   view   as   presented in Figure~\\ref   {chart}.   Density   contours   have      been   created     using    a        grid   with a    spatial     resolution   of     0.5\\degr $\\times$ 0.5\\degr. The  segregation between   early- and   late-type low-luminosity  glaxies   in  the NGC~5846 system  is  evident: dwarf  ellipticals are found predominantly in  the vicinity of bright   galaxies while irregular   systems  are clearly  detached  from  this distribution  and diffusely arranged  in the  outer regions of the  group.  Morphology  shows also a  strong  correlation  with the projected number  density. Early-type   dwarfs populate  the high-density  regions    around  NGC~5846   and  NGC~5813,    while   irregulars show   comparatively   low  agglomeration.    Moreover, the   satellite morphologies  resemble the  morphology  of their   elliptical hosts  (NGC~5846 and   NGC~5813).  Similar  results   have  been   found by \\citet{weinmann} who,  on the   basis of   a  large sample  of  SDSS  groups,   inferred that  the  satellite  morphologies in   groups   correlate strongly with the   morphology of  the central   host galaxy.   Regarding the   radial  morphology  distributions,  Figure~\\ref{clustering}  indicates that   the projected  number    density  of  the  early-type  population    shows    a    much  higher    radial   gradient    than   the    irregular subsample,  merely presenting  any   radial  variation.  In addition  to  the  optical picture, the diffuse X-ray  component correlates also  with the early-type dwarf distribution  (see [MTT05]), indicating   the  dark matter  potentials of the  group.  We have   also checked  the  distribution   of nucleated  dwarfs  within   the  NGC~5846  system    shown  in   Figure~\\ref{nuclei}. With the exceptions  of two objects, all nucleated   dwarfs  are found    in  the vicinity ($\\sim$ 260kpc)  of NGC~5846  and  NGC~5813. Research on  the   distribution of   nucleated dwarfs  in    the  Virgo cluster has also  shown that   nuclei  are predominantly found  in the  central regions    and not  in the  outskirts   \\citep{binggelicameron}. \\citet   {oh} argued   that  this  could   be  explained    assuming   nuclei   to   be   the result     of orbital   decay   of   globular  clusters.    Through  a series   of  numerical simulations    they   showed   that  for   dwarf   galaxies   exposed    to  little   external   tidal  perturbation  dynamical friction  can   lead   to significant orbital  decays  of globular clusters   and  the formation   of  compact nuclei  within a   Hubble-time.   Thus, a   larger    fraction   of nucleated   dwarfs   is   expected   in    the   centres   of   galaxy    clusters   where   the   extragalactic   tidal perturbation  tends to preserve  the integrity of  dwarf galaxies unlike in the outskirts where this perturbation tends to be disruptive.   \\subsection{\\label{emissiondwarfs}Emission-line dwarfs} Line                fluxes              from          H$\\alpha$,           H$\\beta$,              $[$N\\begin{scriptsize}         II\\end{scriptsize}$]$ $\\lambda\\lambda$6548.1,                 6583.4                                     \\AA,                                         $[$O\\begin{scriptsize} III\\end{scriptsize}$]$  $\\lambda\\lambda$4958.9,   5006.8    \\AA,               and           $[$S\\begin{scriptsize}              II\\end{scriptsize}$]$ $\\lambda\\lambda$6717.0, 6731.3 \\AA~features   together    with     star formation   rates       and    oxygen     abundances  have  been measured  for dwarfs showing  emission  lines  in their  spectra. The  data   are     presented   in     Table~\\ref     {emission}. Star formation    rates   (SFRs) were   derived    via   H$\\alpha$ fluxes    using     the      \\citet{kennicutt98} relation:   $\\textrm {SFR}(M_{\\odot}  \\textrm      {\\space  year$^{ -1}$})=7.9      \\cdot  10^{      -42}L(\\textrm     {H}\\alpha)(\\textrm       {ergs       s$^{-1}$)}$.      H$\\alpha$  extinction        was       taken into  account        using  an    average       extinction           value            of     $A(\\textrm{H}\\alpha)=0.95$         mag    based        on \\citet{kennicutt83}          and  \\citet{niklas}.  Additionally,      galactic    extinction      $A_{R}=0.15$     mag     was       also   taken into account  so that      a      total   extinction value    of   1.10   mag     for  H$\\alpha$         was  allowed.    Errors    for    star   formation rates  were estimated assuming    H$\\alpha$      photon  noise as    the   dominant    error   source.   Star     formation   rates      of   galaxies exhibiting   H$\\alpha$       fluxes     with        S/N$<$15   are  not            listed          in          Table~\\ref{emission}.         Emission characteristics    have         been     analysed                                                     calculating                                  the line                                                                                                                                              flux ratios                                                                          $\\textrm{[N\\begin{tiny}II\\end{tiny}]}~\\lambda6583/\\textrm{H$\\alpha$}$, $\\textrm{[O\\begin{tiny}III\\end{tiny}]}~\\lambda5007/\\textrm{H$\\beta$}$,                                                                             and $\\textrm{[S\\begin{tiny}II\\end{tiny}]}~(\\lambda6716+\\lambda6731)/\\textrm{H$\\alpha$}$                            according                            to \\citet{veilleuxosterbrock}.   Oxygen   abundances        have                been                                         estimated              using the    $N_{2}$      index              ($\\textrm{[N\\begin{tiny}II\\end{tiny}]}~\\lambda6583/\\textrm{H$\\alpha$}$)                 method               as defined  by \\citet{denicolo}:  $12+\\log(\\textrm{O/H})=9.12(\\pm   0.05)+0.73(\\pm    0.10)\\times    N_{2}$.    Extinction   and    reddening     effects together  with   specific    absorption   components   underlying      H$\\alpha$    and   H$\\beta$   can       considerably   affect       line   flux measurements. A  correction was  applied    to     the   emission-line      galaxies    by      taking      H$\\alpha$    extinction into account   and assuming    a   Balmer decrement    of H$\\alpha$/H$\\beta=2.85$    for    HII  region-like galaxies \\citep{veilleuxosterbrock}  to  consider   H$\\beta$ extincion too.  \\begin            {figure*} \\centering \\includegraphics  [width=450pt]{osterbrock.eps} \\caption          {\\label{osterbrock}Classification of emission-line galaxies   in the NGC~5846  group   according to Veilleux \\& Osterbrock   (1987). The line   separates HII  region-like galaxies from AGNs. All measured objects show flux ratios explained by simple photoionization.} \\end              {figure*}  \\begin            {figure} \\centering \\includegraphics  [width=\\columnwidth]{oxygen.eps} \\caption          {\\label{oxygen}Oxygen  abundances versus  abolute B  magnitudes of  NGC~5846  irrgeluar dwarfs.  The abundances have been determined indirectly via   the $N_{2}$  method as  proposed by  Denic\u00f3lo  et  al. (2002).  The dash-dotted  line shows the metallicity-luminosity  relation from \\citet{richer}. 1$\\sigma$ deviations are shown as dashed lines.} \\end              {figure}  Figure \\ref{osterbrock} distinguishes  HII  region-like galaxies from  AGNs  such  as  Seyferts  or    LINERS, according to the  diagnosis proposed by \\citet{veilleuxosterbrock}.  All  emission-line     galaxies  can    be explained     by          photoionization.  Only few  bright galaxies     have SDSS spectra.  Furthermore   their  H$\\beta$  flux  is   strongly  contaminated by  the   old  stellar population,  which makes   a proper measurement of the  line ratio problematic.  In addition,   the  [OIII]  feature at   $\\lambda5007$\\AA~   is    not    seen    in    most  of     the investigated bright              galaxies.        However,                just              from                the              measurements              of   the $\\textrm{[N\\begin{tiny}II\\end{tiny}]}~\\lambda6583/\\textrm{H$\\alpha$}$                                                                              and $\\textrm{[S\\begin{tiny}II\\end{tiny}]}~(\\lambda6716+\\lambda6731)/\\textrm{H$\\alpha$}$  ratios,   all   these   objects  can    be  classified   as   HII region-like  objects. There  is no   evident correlation   between  star    formation activity  and location    of emission-line  dwarfs  within   the group. \\citet{lequeux}   suggested  that     the  oxygen    abundances   correlate    with  total     galaxy  mass    for   irregular    galaxies with more massive galaxies  exhibiting  higher metal   content.  Since  the  galaxy   mass  is    a  poorly  known  parameter,   the metallicity-luminosity relation instead   of the  mass-metallicity  relation is   usually considered.   Figure~\\ref{oxygen}   shows the   oxygen abundance  of  emission-line dwarfs measured  via the  $N_{2}$ method versus absolute  blue   magnitude.  Most  of  our   dwarfs lie   above the   metallicity-luminosity  relation from \\citet{richer}. Only the   two  HII  regions N5846\\_41/42    of  KKR15   \\citep{huchtmeier} (see  Section  \\ref{individual}) and  N5846\\_12  fall below this  sequence. The error    bars reflect  merely  the  uncertainty     of  the   $N_{2}$  flux    ratio   and   the  error   of    the   linear least squares   fit   of \\citet{denicolo}. ",
        "conclusions": "We have undertaken an   analysis  of the   photometric  and  spectroscopic   properties  of  the low-luminosity  galaxy    population  of the NGC~5846 group  of galaxies     located   in   the  Virgo III   Cloud of    galaxies in    the Local  Supercluster.  The    group    is of particular  interest since it is   the   most   massive   of    only     three  dense  groups  (NGC~4274, NGC~5846     and    M96)    in the Local Supercluster  that   are dominated     by  elliptical and    lenticular    galaxies,  being  the    third    most  massive    aggregate  of early-type     galaxies  (after the Virgo   and      Fornax clusters)       in      the       local      universe. Of the  19 bright galaxies in  our sample, E and S0 galaxies amount  to $\\sim58$\\%. Seven  new   group   members in   the  outer   regions   ($>$80   arcmin)  of    the group    have   been   identified    via   SDSS   and NED redshifts  complementing    existing     member     catalogues. \\\\  The  dwarf-to-giant-ratio  (DGR) of the   NGC~5846 system  is large:  our   sample of spectroscopically determined group  members  yields a  DGR    of $\\sim$3.  Taking into   account   the photometric classification   of   group members down to  $M_{R}=-10$   by [MTT05] this   value could rise up  to $\\sim$12. \\citet{fergusonsandage} have  shown that  the  DGR increases with  the  richness of  a  group and   constitute an  early-type-dwarf-to-giant-ratio (EDGR),  (dE+dE,N+dS0 over  E+S0) of   5.77 for  the  Virgo  cluster  and  3.83 for the Fornax  cluster down   to $M_{B}=-13.5$. For  the  Coma cluster \\citet{seckerharris} state  an EDGR of $5.80 \\pm 1.33$  down to $M_{B}=-13.5$. Our data yield  an  EDGR of 2.69 for the   NGC~5846 system down to our faintest dwarf  ($M_{B}=-13.32$). This value  shows  that the NGC~5846  system has less early-type dwarfs per giant  early-type galaxy compared to  clusters but   still a   much  larger  EDGR than   other  groups  (Leo:  1.48,  Dorado: 1.44,  NGC~1400:  2.08; all   down  to  $M_{B}=-13.5$, see \\citet{fergusonsandage}),   indicating  that  NGC~5846  is  indeed  one   of  the  more  massive aggregates in    the Local  Supercluster. \\\\  A    colour-magnitude diagram    of   the     investigated   group    members    reveals      a    red sequence displaying   the     well-known  trend of   early-type    galaxies  exhibiting    redder       colours     (higher    metallicities)     for     higher luminosities    \\citep{caldwell}.  In contrast, irregular   dwarfs  do not  form a   sequence  but   show a   large spread  in   colours.   There is  no evident segregation between  bright galaxies and dwarfs. \\\\  Photometric    scaling relations  have  been  studied   for  early-type dwarfs  and  compared with   the  relations   derived  by [BJ98]    for  Virgo cluster    dwarf    ellipicals.   If      bright    ellipticals       and      early-type     dwarfs share    a    similar   evolution, a    continous sequence   in    scaling  relations  with   respect    to    the   galaxy     luminosity    would  be    expected.  It   is      known that   in the surface-brightness-luminosity    ($\\mu-M$)       plane   \\citep{kormendy},           early-type   dwarfs      and     bright   ellipticals    follow different   trends,  marking dwarf         galaxies a   distinct    class   of objects      that have      undergone   a   different evolutionary path compared to   ordinary ellipticals.    Our    data    match  the     trends  from   Virgo      cluster   dEs suggesting     a  similar  structure  and origin    of     the dwarfs  in both      the    Virgo  cluster     and  the   NGC~5846   group. One  special  object    falling    off the       main cloud   of  early-type  dwarfs is   NGC5846\\_01,  however.  It   is   the  closest   dwarf     to  NGC~5846 and shows  a comparatively   high  surface brightness and  compactness  with respect  to   the  other  dwarfs    investigated.   In  the       $\\mu_{0}\\hspace{  -1pt}_{_{B}}-M_{B}$   plane  the object   lies        outside   the  $2\\sigma$   limit       compared  to       the   relation      of [BJ98].        Due to     the proximity       to NGC~5846,    the    galaxy  has      very likely        been   affected     by   tidal stripping   and   can  therefore  rather be  classified      as ultra      compact      dwarf. This   idea is also     supported by    the   fact   that the    nearby object  NGC~5846A,   the  elliptical  companion to NGC 5846,  shows   a similar morphological appearance,    though being   too  bright   to   be   considered as     dwarf. In  fact,  NGC~5846A   is one  of  the    extremely   rare     \"compact  ellipticals\"  like     M32,  the      prototype    of     this   class.    NGC5846\\_01   confirms   the idea    that  the   structural differences   of  compact ellipticals and     ordinary dEs    result    from   the fact    that  the  compacts   formed within  the   potential well   of another    massive galaxy whereas   dEs   evolved  as   isolated    systems \\citep{burkert}.    Our  data also  show that early-type     dwarfs  with   high  ellipticity tend      to  have    shallower   surface     brightness  profiles.   We    also    checked   the location  of  the  NGC~5846  low-luminosity   galaxy population     in   the  $\\mu_{e} -\\log(r_{e})$      plane  with   respect to        the Hamabe-Kormendy   relation,    a     photometric projection     of the      fundamental   plane    of    galaxies        relating     the          logarithm of  effective    radii          and    effective  surface  brightnesses      linearly.  The      NGC~5846   dwarfs   represent    regular     galaxies in     the  classification        of    \\citet{capaccioli}  indicating   that    these    systems   are   the     building      blocks     of     more massive    galaxies,      in  agreement     with  $\\Lambda$CDM  cosmology.  Moreover,     the   dwarfs      are  comparable      with  dwarfs     from X-ray   dim         and   X-ray  bright  groups  studied   by \\citet{khosroshahi}. \\\\  Scaling relations  for dwarf    ellipticals are  not  the  only  way    to enlighten  their  origin  and evolution. Morphological  segregation  within clusters and    groups   also   gives   the   opportunity    to     test   the   formation  and    environmental  dependence      of    the      faint galaxy population.   The   well-known      morphology-density     relation      \\citep{dressler}  is     not   only   restricted     to    the  normal Hubble types    but    is  also    found  to      apply    to     dwarf      galaxies  \\citep{binggeli}.    This   relation  has  been investigated in detail  in our    immediate vicinity, the Local  Group \\citep{grebel}.   We    confirm      this morphology-density         relation     for     dwarf galaxies in  the   NGC~5846   system.    While      early-type        dwarfs  are   concentrated        around  the         two   massive  ellipticals NGC~5846 and    NGC~5813, objects     of          irregular    type  are       more    randomly   distributed       and  found  predominantly       in the outer  regions.  This morphological   segregation  could    be  explained    by  gas    stripping   due   to   the    infall  of the    dwarfs  in the group    centre.   This   scenario     was    also   investigated        by  \\citet     {weinmann}  and    \\citet{park},    who    inferred   that satellite morphologies  in  groups  tend to    be     similar     to   those    of hosts.    With     the   comparatively     high   density   of  the NGC~5846  system   and       the       hot   and       dense        halo    gas       present  around       NGC~5846     and         NGC~5813,     the dependence of the  satellite  morphology   on host  morphology   can be      interpreted     by the hydrodynamic  and   radiative    interaction    of the  hot  X-ray gas of   the  two host    galaxies with their   surrounding  satellites.  We have  also   checked  the   distribution   of   nucleated dwarfs  in   the NGC~5846   system with    respect   to  non-nucleated  ones.    Interestingly, nearly   all    dwarfs    that     show    a   nucleus superimposed      on the    underlying   smooth   surface  brightness   profile    are  found    in   the  vicinity   of  NGC~5846    and     NGC~5813 within a  radius    of $\\sim$260kpc.   This       is     in agreement     with      the   work          of \\citet{binggelicameron}      on  the  Virgo cluster, who   show  that  dwarfs  located    near the  centre  of   the   cluster   are   mostly   nucleated   while those  in the    outskirts   are non-nucleated.  Assuming    the formation      of   the    nuclei   due     to the     orbital  decay      of globular   clusters, \\citet{oh}   argued that this  could  be   explained    by  extragalactic  tidal  perturbation  which   tends to  preserve   the  integrity  of dwarf  galaxies at cluster centres but tends to disrupt dwarf galaxies in their outskirts. \\\\  Only two objects  of  the low-luminosity  galaxy population show   photometric fine structure.  These galaxies exhibit   ongoing star  formation   and diameters in        the   transition  to     bright  galaxies.  This    lack of  photometric  peculiarities  in the low-luminosity  galaxy  population emphasises  the fact  that the NGC~5846  system is  an  old,   well-evolved aggregate where no   recent   interactions between  individual     members have  occured.    Emission characteristics      and star    formation rates     of   the     irregular         dwarfs        show typical  values  for low-luminosity  galaxies,   additionally  supporting   the  idea   of  a    system   without   recent   activity   enhancing   star   formation. Oxygen  abundances    have    been    derived    indirectly       via         the         $N_{2}$          index         method     proposed        by \\citet{denicolo}.   Our abundances    show a   large  scatter     ($\\sigma=0.36$ dex) with    respect to    the  metallicity-luminosity    relation  of \\citet{richer} due    to the  intrinsical  scatter in  the  $N_{2}$   method  itself.   Nevertheless,   all  objects  outside   1$\\sigma$   deviations lie above   the  metallicity-luminosity  relation,  indicating a comparatively  high oxygen abundance for irregular dwarfs in the NGC~5846  group. \\\\  The study  of       the  low-luminosity   galaxy      population   of  the       NGC~5846  group    has revealed  that   the  dwarf galaxies  in  this massive group  match  the   trends observed  for  dwarf  galaxies  in other aggregates   indicating a similar  formation  and evolution  scenario. The spatial distribution of morphological types suggests that the structural properties of dwarf galaxies are clearly dependent on the location within the group. The structural properties of the dwarfs confirm the evolved state of the system. The group  remains   special    due to      the     dominating early-type   morphology,   the   strong      X-ray emission   and  the    large dwarf-to-giant-ratio compared to other groups."
    },
    "0911/0911.4125_arXiv.txt": {
        "abstract": "We present results of a gravitational-lensing and optical study of MACS\\,J1423.8+2404 ($z{=}0.545$, MACS\\,J1423), the most relaxed cluster in the high-redshift subsample of clusters discovered in the MAssive Cluster Survey (MACS). Our analysis uses high-resolution images taken with the {\\it{Hubble Space Telescope}} in the F555W and F814W passbands, ground based imaging in eight optical and near-infrared filters obtained with Subaru and CFHT, as well as extensive spectroscopic data gathered with the Keck telescopes. At optical wavelengths the cluster exhibits no sign of substructure and is dominated by a cD galaxy that is 2.1 magnitudes (K-band) brighter than the second brightest cluster member, suggesting that MACS\\,J1423 is close to be fully virialized. Analysis of the redshift distribution of 140 cluster members reveals a Gaussian distribution, mildly disturbed by the presence of a loose galaxy group that may be falling into the cluster along the line of sight. Combining strong-lensing constraints from two spectroscopically confirmed multiple-image systems near the cluster core with a weak-lensing measurement of the gravitational shear on larger scales, we derive a parametric mass model for the mass distribution. All constraints can be satisfied by a uni-modal mass distribution centred on the cD galaxy and exhibiting very little substructure. The derived projected mass of $M(<65\\arcsec {\\rm [415\\,kpc]} )=(4.3\\pm0.6)\\times 10^{14}$ M$_{\\sun}$ is about 30\\% higher than the one derived from X-ray analyses assuming spherical symmetry, suggesting a slightly prolate mass distribution consistent with the optical indication of residual line-of-sight structure. The similarity in shape and excellent alignment of the centroids of the total mass, K-band light, and intra-cluster gas distributions add to the picture of a highly evolved system. The existence of a massive cluster like MACS\\,J1423, nearly fully virialized only $\\sim$ 7 Gyr after the Big Bang, may have important implications for models of structure formation and evolution on cosmological time scales. ",
        "introduction": "The MAssive Cluster Survey \\citep[MACS,][]{macs} has compiled a complete sample of 12 very X-ray luminous galaxy clusters at $z{>}0.5$ \\citep{highzmacs}, providing a unique opportunity for comprehensive studies of the densest regions of the cosmic web at intermediate redshifts. As noted by \\citet{macslargescale}, MACS\\,J1423 stands out among these 12 extreme systems as the most dynamically relaxed; it is, indeed, the most massive cool-core cluster known at these redshifts. The system has been used in several statistical studies \\citep{laroque03,1423cosmo,bonamente06,1423MT} where its relaxed dynamical state has been particularly important for cosmological work using the baryon fraction \\citep{allen04,schmidtAllen1423,allen08,ettori09}. Since the abundance of massive relaxed clusters at high redshift has important implications for theoretical and numerical models of structure formation and evolution, we here attempt to further test the system's relaxation state using gravitational lensing. All our results use the $\\Lambda$CDM concordance cosmology with $\\Omega_{\\rm{M}} = 0.3, \\Omega_\\Lambda = 0.7$, and a Hubble constant \\textsc{H}$_0 = 70$ km\\,s$^{-1}$ Mpc$^{-1}$. Magnitudes are quoted in the AB system. ",
        "conclusions": "Using optical and lensing observations of MACS\\,J1423 we find overwhelming evidence for a uni-modal mass distribution whose centre, ellipticity, and orientation are in excellent agreement with the corresponding parameters of the cluster light and intra-cluster gas distributions (Fig.~\\ref{fig2}). We quantify the dominance of a single, central mass component by computing the substructure fraction, $f_{\\rm sub}$, defined as the mass within a given cluster-centric radius that can be attributed to substructure, divided by the total mass within the same radius \\citep{smithtaylor}. In the case of MACS\\,J1423, only cluster galaxies (the cD galaxy is by definition excluded) contribute to the substructure fraction. We find $f_{\\rm sub}(<250\\,{\\rm kpc}){=}0.048 \\pm$0.025, comparable to values measured for the most relaxed clusters at $z\\sim0.2$ \\citep{locuss}. Within the high-redshift MACS subsample, the substructure fraction of $f_{\\rm sub}(<500\\,{\\rm kpc}){=}0.045\\pm$0.025 of MACS\\,J1423 is in stark contrast to the value of  0.25$\\pm$0.12 measured for MACS\\,J1149.5+2223, the most complex cluster lens studied to date \\citep{1149}.  The radial mass profile of MACS\\,J1423 as derived from our lensing analysis is well constrained across the full ACS field of view, i.e., out to a radius of 100\\arcsec (640\\,kpc). Previous results for the mass within 65\\arcsec (415\\,kpc) of 5.0$^{+3.1}_{-0.9}\\times 10^{14}$ M$_{\\sun}$, \\citep{laroque03}\\footnote{\\citet{laroque03} quote the mass within a three-dimensional radius, which we convert to a projected mass assumed a $\\beta$-model with $b=2/3$.} using observations of the Sun\\-yaev-Zel`dovich effect, and of 3.1$^{+0.9}_{-0.7} \\times 10^{14}$ M$_{\\sun}$ \\citep[3$\\sigma$ confidence level,][]{schmidtAllen1423} from an analysis of the {\\em Chandra} X-ray data, agree within the errors with our measurement of $4.35\\pm0.6 \\times 10^{14}$ M$_{\\sun}$ (3$\\sigma$). We note, however, that the existing X-ray analyses of MACS\\,J1423 assume spherical symmetry which will bias the resulting mass estimates low for a prolate matter distribution. The difference of $\\sim$ 30\\% between the lensing and X-ray mass may thus be due to an elongation of the system along the line of sight, consistent with our finding of residual accretion onto the cluster core (Section~\\ref{velspec}).  \\citet{locuss} investigated the relationship between the effective Einstein radius and the strong lensing mass (measured within 250 kpc) for a sample of 20 strong lensing clusters at $z\\sim0.2$ drawn from the Local Cluster Substructure Survey (LoCuSS). According to their statistical analysis, one would expect a mass of 2.55$^{+0.90}_{-0.65}$ 10$^{14}$ M$_{\\sun}$ for an \\emph{undisturbed} cluster with an Einstein radius of 21$\\arcsec$. The excellent agreement with our measurement of 2.46$\\pm$0.2 10$^{14}$ M$_{\\sun}$ underlines again the relaxed dynamical state of MACS\\,J1423.  Although our analysis of the radial velocity distribution of cluster members suggests that MACS\\,J1423 evolved through mergers along a line-of-sight filament (Section~\\ref{velspec}), the rate of current or recent infall is low and does not result in a noticeable perturbation of the cluster core. In fact, the large luminosity gap of K$_{\\rm gap}{=}2.06$ mag between the cD and the second brightest cluster galaxy, as well as the excellent match of the dark-matter, intra-cluster gas, and cluster light distributions, represent strong evidence for a dynamically highly evolved cluster core \\citep{jones03,Lgapcluster}, undisturbed by major mergers for at least 1 Gyr, the characteristic time scale for relaxation via dynamical friction. The existence of  MACS\\,J1423 only 7 Gyr after the Big Bang thus adds an important observational constraint to N-body simulations of structure formation and evolution in the current cosmological model."
    },
    "0911/0911.3369_arXiv.txt": {
        "abstract": "We used 3.6, 8.0, 70, 160~$\\mu$m {\\it Spitzer} Space Telescope data, James Clerk Maxwell Telescope HARP-B CO~$J$=(3-2) data, National Radio Astronomy Observatory 12 meter telescope CO~$J$=(1-0) data, and Very Large Array H{\\small I} data to investigate the relations among PAHs, cold ($\\sim20$~K) dust, molecular gas, and atomic gas within NGC~2403, an SABcd galaxy at a distance of 3.13~Mpc.  The dust surface density is mainly a function of the total (atomic and molecular) gas surface density and galactocentric radius.  The gas-to-dust ratio monotonically increases with radius, varying from $\\sim100$ in the nucleus to $\\sim400$ at 5.5~kpc.  The slope of the gas-to-dust ratio is close to that of the oxygen abundance, suggesting that metallicity strongly affects the gas-to-dust ratio within this galaxy.  The exponential scale length of the radial profile for the CO $J$=(3-2) emission is statistically identical to the scale length for the stellar continuum-subtracted 8~$\\mu$m (PAH 8~$\\mu$m) emission. However, CO $J$=(3-2) and PAH 8~$\\mu$m surface brightnesses appear uncorrelated when examining sub-kpc sized regions. ",
        "introduction": "The James Clerk Maxwell Telescope (JCMT) Nearby Galaxies Legacy Survey (NGLS) is an 850~$\\mu$m and CO~$J$=(3-2) survey with the goals of studying molecular gas and dust in a representative sample of galaxies within 25~Mpc.  The sample contains three subsets: a subset of 31 galaxies from the {\\it Spitzer} Infrared Nearby Galaxies Survey \\citep[SINGS;][]{kabetal03} sample; an H{\\small I} flux-limited sample of 36 galaxies in the Virgo Cluster; and an H{\\small I} flux-limited sample of 72 field galaxies.  At this point in time, CO~$J$=(3-2) observations of the SINGS galaxies and Virgo Cluster galaxies have been completed.  Aside from the 3.6-160~$\\mu$m data from SINGS that resolve structures less than a kiloparsec in size in the closest galaxies, H{\\small I} Nearby Galaxies Survey \\citep[THINGS;][]{wbbbktl08} are also available for many of these galaxies.  Together, these data are a powerful dataset for studying the properties of the interstellar medium (ISM) on sub-kiloparsec scales within nearby galaxies such as variations in gas-to-dust ratios on sub-kiloparsec scales.  Gas-to-dust ratios within nearby galaxies have been a major topic of research since the launch of the Infrared Astronomical Satellite (IRAS).  Debate on this subject began when dust masses derived from IRAS 60 and 100~$\\mu$m data seemed a factor of 10 too low to match what was expected either from the depletion of metals from the ISM or from the ratio of gas column density to dust extinction \\citep[e.g.][]{yskl86, yxkr89, dy90}.  More recent calculations of global gas and dust masses using {\\it Spitzer} Space Telescope \\citep{wetal04} and James Clerk Maxwell Telescope data for spiral galaxies produce gas-to-dust mass ratios of $\\sim150$ \\citep[e.g][]{rtbetal04, betal06, ddbetal07}.  Although the measured ratios are reliant upon CO to H$_2$ mass conversion factors that do not vary with metallicity or star formation activity, the ratios that were measured are close to what is expected from the ratio for the Milky Way derived from the comparison of gas column densities to optical extinction \\citep[e.g.][]{k03, w03} or from the depletion of metals from the ISM \\citep[e.g.][]{w03, l04, k08}.  Therefore, the natural expectation would be that dust emission from kiloparsec-sized regions should be correlated with the sum of atomic and molecular gas emission on those spatial scales.  However, since the gas-to-dust ratio depends on metallicity and since metallicity varies with radius within most galaxies \\citep[e.g.][]{ve92, zkh94, vetal98}, the gas-to-dust ratio should decrease with radius.  For the most part, these assumptions are only beginning to be tested, mainly because the combination of H{\\small I}, CO, and far-infrared dust emission data that is resolved on kiloparsec scales has not been available for nearby galaxies.  Related to this topic is the association between polycyclic aromatic hydrocarbon (PAH) emission in the mid-infrared and other constituents of the cold ISM.  Several observational studies using data from either the Infrared Space Observatory \\citep[ISO; ][]{ketal96} or {\\it Spitzer} have shown that, in nearby spiral galaxies, PAH emission is strongly correlated with cold ($\\sim20$ K) dust emission in the far-infrared, at least on scales $\\gtrsim1$~kpc \\citep[e.g.][]{mtl99, hkb02, betal06, betal08, zwcl08}.  Additionally, \\citep{retal06} found that PAH emission is correlated with CO spectral line emission at submillimetre and millimetre wavelengths within nearby spiral galaxies, although these results are based on using radial averages. These results strongly suggested that PAH emission is a tracer of both dust and gas in the cold interstellar medium (ISM).  Some of the SINGS/THINGS/JCMT NGLS galaxies have already been used in comparisons of H{\\small I} and CO emission with star formation, including measures of star formation that incorporate 24~$\\mu$m emission \\citep[e.g.][]{ketal07, blwbdmt08, wetal09}.  These studies have found that star formation is correlated with molecular gas surface density or with a sum of molecular and atomic gas surface density; no correlation exists between star formation and H{\\small I} emission by itself.  Since 24~$\\mu$m dust emission is included in these comparisons, the results imply that dust emission should be correlated with molecular or total gas surface densities but not with atomic gas surface densities, although a direct comparison between dust, atomic gas, and molecular gas is needed.  As a first look into these topics with the combined SINGS, THINGS, and JCMT NGLS data, we compare PAH, cold dust, CO~$J$=(3-2), and H{\\small I} on sub-kiloparsec scales for NGC~2403.  This SABcd spiral galaxy \\citep{ddcbpf91} was selected because it is relatively nearby \\citep[the distance to the galaxy is $3.13 \\pm 0.14$~Mpc; ]{fetal01}, the galaxy and the dust emission is very extended \\citep[the optical disc is 21.9~arcmin $\\times$ 12.3~arcmin; ][]{ddcbpf91}. While the inclination is $62.9^\\circ$ \\citep{dwbtok08}, the galaxy is close enough to face-on that it is possible to distinguish spiral structures and individual star-forming regions.  As is found in similar late-type spiral galaxies, the regions with the strongest star formation are located outside the centre of of the galaxy \\citep[e.g][]{drms99, betal08}, while the cold dust emission peaks in the centre \\citep{betal08}.  Consequently, emission related to star formation can be easily disentangled from emission related to dust surface density. ",
        "conclusions": "We find that the dust surface density depends on the total (atomic and molecular) gas surface density and on radius.  The gas-to-dust ratio varies from $\\sim100$ in the nucleus to $\\sim400$ at 5.5~kpc, which is the limit at which the dust surface densities can be detected.  The logarithm of the slope for the gas-to-dust ratio ($d\\log(\\sigma_{dust}/\\sigma_{gas})/dr$)), which is $0.097\\pm0.002$~dex~kpc$^{-1}$, is very close to the slope in 12+log(O/H), which is usually reported as between -0.07 to -0.09~dex~kpc$^{-1}$.  This implies that the gas-to-dust ratio is strongly dependent on metallicity.  Additional abundance variations, such as variations in the C/O ratio, could explain the small disparity between $d\\log(\\sigma_{dust}/\\sigma_{gas})/dr$ and 12+log(O/H).  Aside from these strong radial variations in the gas-to-dust ratio, the ratio does not appear to vary within the galaxy.  We also find that the PAH 8~$\\mu$m and CO~$J$=(3-2) radial profiles have statistically identical scale lengths.  On spatial scales of 220~pc, however, PAH 8~$\\mu$m and CO~$J$=(3-2) emission are uncorrelated.  We propose two mechanisms that could link the PAH emission with star formation.  First, the CO emission may appear associated with PAH emission if molecular cloud formation is triggered by stellar potential wells.  The stars in these potential wells would enhance the radiation field that heats the diffuse ISM, which would enhance PAH emission.  Thus, the PAH emission would appear associated with CO emission on large scales but not on small scales.  Second, the two wave bands may be related through different excitation mechanisms that are linked through star formation and that affect the ISM on large areas.  While PAH emission does not trace star formation on sub-kpc scales, it is expected to be enhanced by increases in the interstellar radiation field that accompany enhanced star formation. Meanwhile, CO emission could be enhanced by the increase in cosmic rays that accompanies enhanced star formation.  Thus, radially-averaged PAH and CO emission would both be linked through star formation, even though PAH and CO emission may not be spatially correlated on sub-kpc scales.  However, we also do not rule out the possibility that the may just coincidentally have radial profiles with similar scale lengths.  This paper is only a first look at the gas-to-dust ratio and the relation between PAH and CO emission within the JCMT NGLS sample.  The SINGS, THINGS, and JCMT NGLS samples were all selected to contain many of the same galaxies.  The subset of galaxies found in all three samples can be used for an extended comparison of dust, molecular gas, and atomic gas surface densities and for comparisons of PAH to CO emission.  Moreover, the gradients in the gas-to-dust ratio could then be compared to abundance gradients to examine the relation between the two further."
    },
    "0911/0911.4255_arXiv.txt": {
        "abstract": "{ Ultra-cool dwarfs as wide companions to subgiants, giants, white dwarfs and main sequence stars can be very good \\emph{benchmark} objects, for which we can infer physical properties with minimal reference to theoretical models, through association with the primary stars. We have searched for benchmark ultra-cool dwarfs in widely separated binary systems using SDSS, UKIDSS, and 2MASS. We then estimate spectral types using SDSS spectroscopy and multi-band colors, place constraints on distance, and perform proper motions calculations for all candidates which have sufficient epoch baseline coverage. Analysis of the proper motion and distance constraints show that eight of our ultra-cool dwarfs are members of widely separated binary systems. Another L3.5 dwarf, SDSS 0832, is shown to be a companion to the bright K3 giant $\\eta$ Cancri. Such primaries can provide age and metallicity constraints for any companion objects, yielding excellent benchmark objects. This is the first wide ultra-cool dwarf + giant binary system identified. } % ",
        "introduction": "\\label{intro} Current brown dwarf (BD) models have difficulty in accurately reproducing observations, and benchmark objects are thus needed to calibrate both atmospheric and evolutionary models. In order for a brown dwarf to be considered a benchmark it must have one or more properties (e.g. age, mass, distance, metallicity) that can be constrained relatively independently of models. BDs as members of binaries are very useful for this purpose, where the host star can provide age, distance and in some cases metallicity constraints. How useful these systems are, is dependent on how well we understand the nature and physics of the primary stars. Evolved stars, e.g. subgiants, giants and white dwarfs can constraint accurate ages, as well as metallicity in the case of subgiants and early phase giant stars, via robust spectroscopic models. These types of binary system can populate the full age range of the disk up to 10 Gyrs\\cite{pin}. In this paper we report the discovery of nine common proper motion (PM) binary systems with ultra-cool dwarf (UCD) components. One is a benchmark L dwarf companion to the giant star $\\eta$ Cancri, and the other eight are late M and  early L dwarf companions to K-M dwarf stars. ",
        "conclusions": ""
    },
    "0911/0911.4063_arXiv.txt": {
        "abstract": "A digitalized temperature map is recovered from the first light sky survey image published by the Planck team, from which an angular power spectrum is derived. The amplitudes of the low multipoles measured from the preliminary Planck power spectrum are significantly lower than that reported by the WMAP team. Possible systematical effects are far from enough to explain the observed low-$l$ differences. ",
        "introduction": "The first light survey image covering about 10\\% of the sky has been published by the Planck team$^{[1]}$. Recently, through a visual comparison of this first light sky survey with the WMAP and COBE results by using greyscale images, Cover$^{[2]}$  found that there are substantial differences at large scale (low-$l$) between the WMAP and Planck sky maps. In this work, we extract the temperature intensity information from the published Planck image and located their corresponding pixels on the sky to produce a sky map. We then derive a CMB power spectrum from the map and find the amplitudes of the low multipoles measured from the spectrum being significantly lower than WMAP release, which is consistent with what is seen by Cover$^{[2]}$ and also consistent with our findings from an improved spectrum by using the WMAP data and our own software$^{[3]}$. ",
        "conclusions": "\\label{sec:conclusion} From the Planck first light sky survey image we see that the most prominent feature of CMB anisotropy, revealed by WMAP with its first acoustic peak, is certainly real and well confirmed by Planck. On the other hand, for the low-$l$ components (see Table~1), significant differences between the Planck preliminary spectrum and WMAP release exist. It is obvious that the published Planck first light image is possibly contaminated by systematic instrumental effects and might be unsuitable for precise cosmology study, e.g. for cosmological parameters estimation. However, as discussed in \\S~\\ref{sec:Effect of the systematics}, for the low-$l$ components, the fact that they are significantly lower than WMAP is quite beyond doubt.  Recently, we found a remarkable inconsistency between the input data and reconstructed temperatures in WMAP team's map-making and observed serious ecliptic contamination in published temperature maps$^{[3]}$. That indicates there must be a serious problem in the map making routine of the WMAP team. We then independently developed a map-making software which can completely overcome the problem of inconsistency. We used the improved software to produce new WMAP temperature maps in safe mode with stronger constraints on data selection for planet avoidance. The power spectrum derived from our safe mode maps has nearly zero quadrupole component, which is consistent with what revealed by the Planck first light survey.  Tegmark et al.$^{[6]}$ found in the released first year WMAP maps both the CMB quadrupole ($l=2$) and octopole ($l=3$) having power along a particular spatial axis, and more works$^{[7-10]}$ found that the preferred directions of these two low-$l$ components are highly aligned and close to the ecliptic plane. Such a strange orientation of large-scale patterns of CMB maps in respect to the ecliptic frame has long puzzled cosmologists. Our work$^{[3]}$ pointed out that the CMB quadrupole component apparent in the published WMAP maps is not of cosmological origin, but mainly came from  ecliptic contamination induced by inappropriate map-making. From the comparison between Planck and WMAP for greyscale maps by Cover$^{[2]}$ and for power spectra in our present work, the Planck first light survey image strongly supports the artificial origin of quadrupole observed in official WMAP maps and the real CMB quadrupole is most possibly near zero.  Table~1 shows that not only the quadrupole, but also other low-$l$ components, from $l=3$ (octopole) to $l=5$, are very weak. We have found two kinds of systematic errors in official WMAP maps. One is the pixels in the WMAP scan ring of a hot pixel being systematically cooled$^{[11,12]}$.  Another one is the significant correlation between the temperature and the observation number of a sky pixel$^{[11,13]}$. These two kinds of systematic effects also remain in our safe mode maps. The sky coverage of WMAP mission is nonuniform: the number of observations is greatest near the ecliptic poles and in rings $\\sim 141\\degree$ from each pole, and least in the ecliptic plane. The non-uniform exposure pattern should produce artificial power component at low-$l$ through the systematic effects mentioned above. It is expected that after properly correcting systematics low-$l$ spectrum of WMAP safe-mode maps can be more consistent with the Planck survey result."
    },
    "0911/0911.2473_arXiv.txt": {
        "abstract": "We present multi-epoch Very Long Baseline Array (VLBA) imaging of the $^{28}$SiO $v$=1 and $v$=2, $J$=1-0 maser emission toward the massive young stellar object (YSO) Orion \\SI. Both SiO transitions were observed simultaneously with an angular resolution of $\\sim$0.5~mas ($\\sim$0.2~AU for $d$=414~pc) and a spectral resolution of $\\sim$0.2~\\kms. Here we explore the global properties and kinematics of the emission through two 19-epoch animated movies spanning 21 months (2001 March 19 to 2002 December 10). These movies provide the most detailed view to date of the dynamics and temporal evolution of molecular material within $\\sim$20-100~AU of a massive ($\\gsim8M_{\\odot}$) YSO. As in previous studies, we find that the bulk of the SiO masers surrounding \\SI\\ lie in an {\\sf X}-shaped locus; the emission in the South and East arms is predominantly blueshifted and emission in the North and West is predominantly redshifted. In addition, bridges of intermediate-velocity emission are observed connecting the red and blue sides of the emission distribution. We have measured proper motions of over 1000 individual maser features and find that these motions are characterized by a combination of radially outward migrations along the four main maser-emitting arms and motions tangent to the intermediate-velocity bridges. We interpret the SiO masers as arising from a wide-angle bipolar wind emanating from a rotating, edge-on disk. The detection of maser features along extended, curved filaments suggests that magnetic fields may play a role in launching and/or shaping the wind. Our observations appear to support a picture in which stars with masses as high as at least $8M_{\\odot}$ form via disk-mediated accretion. However, we cannot yet rule out that the \\SI\\ disk may have been formed or altered following a recent close encounter. ",
        "introduction": "At a distance of $\\sim414$~pc (Menten et al. 2007; Kim et al. 2008), radio ``\\SI'' in the Kleinmann-Low (KL) nebula of Orion is believed to be the nearest example of a massive young stellar object (YSO). \\SI\\ is highly embedded and has no optical or infrared counterpart;  at wavelengths of 8 and 22$\\mu$m Greenhill et al. (2004a) estimated continuum optical depths of $>$300. However, high-resolution 7-mm radio continuum observations have revealed an extended ($\\sim$70~AU across) source that Reid et al. (2007) proposed may be an ionized accretion disk surrounding an early B-type star.  In addition to its proximity, Orion~KL has the distinction of being one of only three star-forming regions known to exhibit SiO maser emission (Hasegawa et al. 1986; Zapata et al. 2009), and these masers were definitively linked with \\SI\\ by Menten \\& Reid (1995). SiO masers offer key advantages for studying stellar sources in that they arise at small radii ($T_{\\rm ex}>10^{3}$~K, $n({\\rm H}_{2})\\sim10^{10\\pm1}$~cm$^{-3}$), are unaffected by extinction from dust and neutral gas, and can sample gas kinematics via emission line Doppler shifts and proper motions. Moreover, their high brightness temperature makes them observable with extraordinarily high angular resolution using Very Long Baseline Interferometry (VLBI). Indeed, previous VLBI observations have established that the SiO masers lie at projected distances of $\\sim$20-100~AU from \\SI\\ (Greenhill et al. 1998, 2004b; Doeleman et al. 1999, 2004; Kim et al. 2008). These scales are of considerable interest, since they correspond to regions where accretion is expected to occur and where winds and outflows from YSOs are expected to be launched and collimated. At present, little is known about the dynamics of these regions in higher mass YSOs ($M_{\\star}\\gsim8M_{\\odot}$) owing to the dearth of direct observations with the requisite resolution. Consequently, persistent gaps in our knowledge of the formation process for OB stars have remained (e.g., Bally \\& Zinnecker 2005; Zinnecker \\& Yorke 2007; Tan 2008).  We have used the National Radio Astronomy Observatory's\\footnote{The National Radio Astronomy Observatory is a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc.} Very Long Baseline Array (VLBA) to monitor the SiO maser emission surrounding \\SI\\ at monthly intervals over several years (see also Greenhill et al. 2004b). Our images have higher dynamic range and higher cadence than any observations of the \\SI\\ masers to date. Here we showcase our results in the form of 19-epoch movies spanning 21 months. These movies chronicle for the first time the kinematics and evolution of the molecular material $\\sim$20-100~AU from a massive YSO. This is part of a series of papers examining \\SI\\ on 10-1000~AU scales based on VLBA monitoring and complementary observations of other maser and thermal lines using the Very Large Array (VLA) and the Green Bank Telescope (e.g., Goddi et al. 2009a,b and in preparation; Greenhill et al., in preparation; Matthews et al., in preparation). Together, these data are allowing us to forge a comprehensive new picture of \\SI\\ and its role in shaping the Orion~KL region. ",
        "conclusions": "} \\subsection{A New Model for the SiO Maser Kinematics\\protect\\label{model}} Earlier VLBI observations of SiO masers around \\SI\\ established the {\\sf X}-shaped distribution of emission as well as the spatial separation between the red- and blue-shifted arms. Based on these observations, Greenhill et al. (1998) and Doeleman et al. (1999) proposed that the SiO masers arise from limb-brightened edges of a biconical outflow, oriented along a southeast/northwest direction. However, our new, more sensitive observations are better explained by a new model---namely that the masers are associated with an edge-on disk whose rotation axis is oriented along the northeast/southwest direction (see also Greenhill et al. 2004b; Kim et al. 2008). Support for this model comes from the radial velocity gradients along each of the four arms (indicative of differential rotation), as well as the presence of the bridge emission and its associated velocity gradients (\\S~\\ref{globvel}). Furthermore, the canting of the maser arms (Fig.~\\ref{fig:nest}) suggests a reflexive symmetry about a plane whose position angle (PA$\\approx141^{\\circ}\\pm1^{\\circ}$) closely matches the  ridge of 7-mm continuum emission observed by Reid et al. (2007) on 2000 November~10 (PA$\\approx142^{\\circ}\\pm3^{\\circ}$). The presence of the well-defined dark band in Fig.~\\ref{fig:nest} (see \\S~3.1.1) also suggests that the conditions favorable for excitation of the SiO masers set in rather abruptly at a fixed scale height above a flattened, disk-like structure.  Underlying the new kinematic model for \\SI\\ is the assumption that the SiO masers trace real, physical motions of gas clumps rather than, e.g., illumination patterns or shocks transversing a fixed medium. The contrast between the pattern of motions observed in the Western bridge versus the arms (\\S~\\ref{proper}) strongly supports this interpretation, as shocked material would not be expected to have two apparent kinematic components (both tangential and radially outward).  Moreover, owing to gradients in temperature and density between regions close to the disk plane and the outer reaches of the wind, shock conditions are unlikely to be similar over the scales of tens of AU from \\SI\\ where maser motions are observed.  Finally, the multitude of linear maser trails seen in Fig.~\\ref{fig:nest} also supports a kinematic interpretation for the SiO masers; features that move many times their characteristic sizes without significantly changing morphology would not be expected to arise as a shock front transverses a clumpy, inhomogeneous medium and are a strong indicator that we are tracing motions of individual clumps.  \\subsection{The Mass of Source I\\protect\\label{mass}} One of the longstanding controversies surrounding \\SI\\ has been the mass of the central star. While our present data can offer important new constraints on this quantity, obtaining a precise mass estimate is complicated by the likelihood that the maser-emitting material is not in purely Keplerian rotation (see also Kim et al. 2008). For example, some degree of turbulence is almost certain present in the gas (e.g., as evidenced by small-scale complexities in the radial velocity fields in the arm regions), and the outward motions of the maser clumps (\\S~\\ref{proper}) imply forces acting on the maser gas opposite to those of gravity (e.g., radiative and/or magnetic forces).  Such effects can lead to underestimates of the enclosed mass (e.g., K\\\"onigl \\& Pudritz 2000; Pi\\'etu et al. 2005; Bujarrabal et al. 2005). Moreover, the disk itself might contain up to a few solar masses of material (Reid et al. 2007) and thus have a non-negligible mass relative to the central source.  One means of estimating of the mass of \\SI\\ comes from the observed transverse motions of maser features in the Western bridge region (Fig.~\\ref{fig:PM}). We assume that the bridge features lie at $r\\sim35$~AU in a circularly rotating disk (i.e., just outside the edge of ionized inner disk measured by Reid et al. 2007) and that the \\SI\\ disk has a nearly edge-on inclination to our line-of-sight ($i\\sim85^{\\circ}$). Taking the mean space motion of the $v$=2 bridge features ($V_{\\rm 3D}\\sim$13.5~\\kms) then implies $M_{\\star}\\gsim7M_{\\odot}$.  A second mass estimate can be derived by assuming that the material within the four arms is part of an outflowing wind (see \\S~\\ref{wind}) and therefore must be moving at or near escape velocity.  Taking the mean space velocities of the $v$=2 masers within the four arms ($V_{\\rm 3D}\\sim16.0$~\\kms) and a fiducial radius $r\\sim$25~AU from \\SI\\  (roughly equal to the radius of the midpoint of the base of each arm from \\SI) also implies a central mass  of $M_{\\star}\\gsim7M_{\\odot}$. These kinematically-derived masses are somewhat smaller than the estimate of Reid et al.  (2007) based on the 7-mm radio continuum luminosity ($M_{\\star}\\approx10M_{\\odot}$). Since a star with $M_{\\star}\\lsim7M_{\\odot}$ would be unable to produce an \\HII\\ region and would not have sufficient luminosity to power the SiO maser emission ($L_{\\star}\\gsim10^{4}~L_{\\odot}$; Menten \\& Reid 1995), it therefore seems probable that \\SI\\ is  somewhat higher than the above kinematically determined values---i.e., $M_{\\star}\\sim8$-$10M_{\\odot}$.  If magnetic fields are threading the \\SI\\ disk and are responsible for powering the wind (see \\S~\\ref{MHD}), then the models described by K\\\"onigl \\& Pudritz (2000) predict that  only the material at the {\\sl base} of the wind is expected to exhibit Keplerian rotation. We therefore have measured the locations of the peak rotational velocity of the SiO emission along the northeastern edge of the dark band that runs parallel to the disk midplane (see Fig.~\\ref{fig:cartoon}). Taking a mean from the 19 epochs of data and from the red- and blue-shifted sides of the disk, we find $|V_{\\rm max}|\\approx$19~\\kms\\ at $r\\approx20$~AU. This implies $M_{\\star}\\gsim8M_{\\odot}$.  If dust is mixed with the SiO maser-emitting gas as proposed by Elitzur (1982), then radiation pressure on the grains could also influence the  gas kinematics, assuming the grains can efficiently transfer some of their momentum to the gas. This would again lead to observed velocities that are smaller than predicted by Kepler's Law for a given mass.  However, estimating the magnitude of this effect on the observed gas velocities will depend on the intrinsic stellar mass and luminosity, as well as the detailed properties of the grains---all of which are uncertain (e.g., Kwok 1975). Moreover, it is unclear whether dust could survive at the temperatures expected near the base of the \\SI\\ wind ($T\\gsim$2000~K; Goddi et al. 2009a; see also below).  We emphasize that the interpretation of \\SI\\ as a single, luminous YSO with $M_{\\star}\\sim8$-$10M_{\\odot}$ seems to provide the simplest explanation for {\\em both} the radio continuum observations and the SiO maser kinematics. A binary of equivalent mass, or a less massive, less luminous star, would not be able to produce the observed radio continuum emission nor provide the necessary luminosity to power the SiO masers. On the other hand, a significantly more massive star would be difficult to reconcile with the observed SiO kinematics (i.e., the transverse motions in the bridges, which we interpret as material orbiting the central mass, and the systematically outward motions in the arms, which we interpret as material traveling at or near escape velocity). Our results therefore seem to be  in contradiction with the scenario recently proposed by G\\'omez et al. (2008), in which \\SI\\ was ejected from a multiple stellar system $\\sim$500~years ago and now comprises a tight binary with a mass in the range $12~M_{\\odot}<M_{\\star}<19~M_{\\odot}$.  In addition to the problem of the discrepancy between our derived mass for Source~I and the value proposed by G\\'omez et al. (2008), accounting for the properties of Source~I's disk in such a picture (e.g., its size, density, and symmetry) might also be problematic. The passage of another star within $\\sim230\\pm70$~AU of \\SI\\ (see G\\'omez et al.) and the subsequent formation of a tight binary most likely would have  disrupted any previously existing disk around \\SI\\ (e.g., Moeckel \\& Bally 2006), requiring formation of the existing disk structure within the past 500~yr. Nonetheless, under certain conditions, such rapid disk regrowth might be possible via either Bondi-Hoyle accretion or the tidal shredding of the interloper.  In the case of Bondi-Hoyle accretion, the Bondi radius of an accreting object of mass $M$ is defined as $r_{\\rm B}=GM/v^{2}_{i}$, where $G$ is the gravitational constant and $v_{i}$ is a characteristic velocity, which may be taken to be the motion of the object with respect to the surrounding gas (e.g., Krumholz et al. 2006). From Goddi et al. (in prep.), the motion of \\SI\\ with respect to the ambient medium is $v_{i}\\sim$12~\\kms, implying $r_{\\rm B}\\approx$50~AU for $M=8M_{\\odot}$. Such a radius is comparable to the observed size of the \\SI\\ disk.  While the mass of the disk surrounding \\SI\\ is uncertain (see Reid et al. 2007), we can roughly estimate this quantity, $M_{d}$, by assuming the disk shape is a flattened cylinder with $r\\sim r_{B}=$50~AU and $h$=14~AU (\\S~\\ref{morph}) and that the mean particle density is comparable to the value required to explain the SiO maser emission ($n\\approx10^{10}$~cm$^{-3}$). Assuming the bulk of the disk material is ionized, we adopt a mean molecular weight per particle of $\\mu$=0.6, implying $M_{d}\\approx0.002M_{\\odot}$. Given a time frame of 500~yr, this implies a minimum required mass flux ${\\dot M}_{d}\\sim 4\\times10^{-6}~M_{\\odot}$~yr$^{-1}$ (comparable to the mass {\\em outflow} rate estimated from the SiO $v$=0 emission; Greenhill et al., in prep.). Using Equation~1 of Krumholz et al. (2006), we can now estimate the required ambient density to support this mass accretion rate as $\\rho_{a}={\\dot M}_{d}v^{3}_{i}[4\\pi G^{2}M^{2}_{\\star}]^{-1}\\approx 3.0\\times10^{-17}$~g~cm$^{-3}$. Assuming the ambient material is purely molecular ($\\mu$=2.3) then implies an ambient particle density $n_{\\rm H_{2}}\\approx 8\\times10^{6}$. Although the latter value is rather high, it is comparable to values previously measured for the Orion hot core region (e.g., Masson \\& Mundy 1988 and references therein) and thus may be roughly consistent with plausible values for this Orion~KL region. We conclude that in the absence of more sophisticated calculations, we  cannot rule out disk augmentation or rebuilding via Bondi-Hoyle accretion may have occurred during the past 500~yr. However, we note that one additional caveat is that disks formed in this manner are predicted to have rather chaotic and asymmetric structures (e.g., Krumholz et al.), in contrast to the \\SI\\ case.  One additional scenario that might reconcile our new \\SI\\ measurements with the findings of G\\'omez et al. (2008) is the possibility that \\SI\\ has recently formed or altered its disk by tidally shredding a lower mass interloper (see Davies et al. 2006). If this interloper had a mass as high as $M_{\\star}\\approx 3M_{\\odot}$, this would bring the total mass of the \\SI\\ system in marginal agreement with values proposed by G\\'omez et al. Furthermore, such an event might offer a natural explanation for the maser emission and outflows associated with \\SI\\ (see Bally \\& Zinnecker 2005). Nonetheless, this picture would likely require a considerably closer passage between the two stars than estimated by G\\'omez et al. (i.e., as close as a few tens of stellar radii; see Davies et al.). In addition, the relatively large resulting mass of the disk may in turn require an implausibly luminous central star in order to account for the observed radio continuum emission (see Reid et al. 2007).  In summary, while our latest observations of \\SI\\ present a compelling case for disk-mediated accretion in a massive YSO, it is clear that further modeling will be required to explore possible interaction scenarios and to better constrain whether such an event is likely to have influenced its present disk and outflow properties.   New N-body simulations as well as further discussion of the interaction history of \\SI\\ will be presented by Goddi et al. (in preparation).  \\subsection{What Drives the SiO Maser Emission from Source I?\\protect\\label{wind}} Because SiO maser emission is extremely rare around YSOs, its origin in the case of \\SI\\ has been another longstanding puzzle. Cunningham et al. (2005)  proposed that the \\SI\\ masers arise from a shear layer along the walls of a cavity that has been evacuated as a bipolar wind expands into a rotating, collapsing envelope. This model is able to reproduce the magnitudes and directions of the proper motions in the arms as well as the frothy appearance of the SiO-emitting material. However, the observed breadth of the arms appears to be greater than expected for an interface region. Moreover, this model cannot readily explain several other features of the SiO masers, including the transverse motions observed in the inter-arm bridges, the presence of groups of maser features beyond the four main arms (the ``isolated'' features in Fig.~\\ref{fig:cartoon}), or the linear trajectories of features that persist unperturbed over many months (Fig.~\\ref{fig:nest}, \\ref{fig:streamer}, \\& \\ref{fig:twizzler}). The Cunningham et al. model also predicts higher densities along the outer edges of the arms (their Fig.~1), which appears to be inconsistent with our observation that the $v$=2 masers (which preferentially occur in higher density gas) lie on average closer to \\SI\\ than the $v$=1 masers (\\S~\\ref{morph}; Fig.~\\ref{fig:histo}).  Wright et al. (1995) proposed a slightly different scenario---namely that the SiO masers arise from material ablated from the surface of an accretion disk by a wind or outflow. Because our current observations provide evidence for the presence of an accretion disk, we now favor some variant of this ``boiling disk'' picture. Constraining the driving mechanism for this wind will require detailed modeling, but here we comment briefly on the likely applicability of various classes of disk wind models to the \\SI\\ case.   \\subsubsection{Disk Photoionization} As discussed by Hollenbach et al. (1994), YSOs hot enough to produce an \\HII\\ region are capable of mass-loss via photoevaporation of their disks as material is heated to temperatures in excess of the local escape temperature. In the ``weak wind'' case, an ionized flow is predicted to set in beyond the disk radius, $r_{g}$, where the sound speed is roughly equal to the escape speed. For \\SI, if we take the sound speed as $\\sim$11~\\kms\\ (assuming $T=8000$~K on the surface of the ionized disk; Reid et al. 2007), a mean mass per particle of 1.13$\\times10^{-24}$~g within the ionized disk (Hollenbach et al. 1994), and $M_{\\star}\\approx8M_{\\odot}$, this predicts $r_{g}\\approx1\\times10^{15}$~cm--- roughly a factor of three larger than what is observed.  In addition to this discrepancy, the outflowing material is predicted to be mostly ionized, raising the problem of how to maintain sufficient quantities of dense, molecular material in the disk wind and how to account for the SiO maser emission in the bridge regions.  \\subsubsection{Line-Driven Winds} For hot stars, radiation pressure mediated by ultraviolet absorption line opacity offers another means of powering a wind. Classically, such winds tend to have velocities too high ($\\gsim$400~\\kms) and densities too low ($\\rho<< 10^{-14}$ g cm$^{-3}$) to readily account for the maser-emitting gas surrounding \\SI\\ (e.g., Lamers et al. 1995). Drew et al. (1998) showed that the presence of a rotating, circumstellar disk helps to produce a slower, denser  wind along the equatorial direction. However, for $M_{\\star}\\approx10M_{\\odot}$, the characteristic wind speeds at more oblique angles (i.e., angles consistent with the outward motions along the arms of \\SI) still are an order of magnitude too high to match those of the \\SI\\ SiO masers (see Drew et al.). Furthermore, the high poloidal velocities predicted by this model are inconsistent with outflow speeds inferred from other line tracers, such as SiO $v$=0 emission (Goddi et al. 2009a; Greenhill et al., in preparation). It thus appears that line-mediated radiation pressure is unlikely to play a significant role in powering the outward migrations of the SiO maser-emitting gas around \\SI.  \\subsubsection{Radiatively-Driven (Dust-Mediated) Winds} Elitzur (1982) has suggested that gas and dust may be mixed within the zone from which the SiO maser emission arises around \\SI. This would be in contrast to the situation in the extended atmospheres of evolved stars, where the SiO maser emission arises from a region just inside the dust-formation radius (e.g., Reid \\& Menten 1997). Elitzur further showed that in such a case, dust-mediated radiation pressure  alone might be able to account for the outflow of material from \\SI.  Although we cannot yet exclude this general class of model for driving the \\SI\\ wind, several of the previous assumptions made by Elitzur (1982) require revision as a result of more recent, high-resolution observations. For example, Elitzur assumed a luminosity for the \\SI\\ of $10^{5}L_{\\odot}$, which is likely an order of magnitude too high (Reid et al. 2007) and a mass-loss rate of $10^{-3}M_{\\odot}$~yr$^{-1}$---two orders of magnitude larger than the value recently derived by Greenhill et al. (in preparation) based on SiO $v$=0 measurements. Elitzur's model also depends upon energy input by turbulent motions of $\\sim$7~\\kms\\ in order to maintain the gas at a higher temperature than the dust (as otherwise the gas would rapidly thermalize to the dust blackbody temperature and population inversion could not be maintained). Such large turbulent motions are inconsistent with the low mean line-of-sight velocity dispersion of the maser-emitting material seen in our present data (Fig.~\\ref{fig:veldisp}) as well as with the ordered velocity field of the SiO masers (Fig.~\\ref{fig:mom1}). Lastly, as already noted above, it would likely be difficult for grains to survive within $\\lsim$100~AU from \\SI, particularly near the base of the wind, where temperatures may be $\\gsim$2000~K (Goddi et al. 2009a). Given these new developments, the dust-driven wind scenario for \\SI\\ now faces a number of challenges.  \\subsubsection{Magnetohydrodynamic Winds\\protect\\label{MHD}} Apparent curved and helical trajectories of certain SiO maser features (\\S\\ref{morph} \\& \\ref{proper}), strongly hint that magnetic fields may play a role in shaping the dynamics of the \\SI\\ region. For example, it would seem difficult to explain features such as the streamer emulating from the Western bridge (Fig.~\\ref{fig:streamer}) in the absence of magnetic fields. While pressure gradients within the disk might act to bend the paths of outflowing material, the observations of pronounced curvature at large vertical displacements from the plane would require that the pressure scale height of the disk is comparable to the observed vertical extent of the masers.  Moreover, we see numerous proper motion trajectories that appear to be linear, including along the outer edges of the individual arms, contrary to what would be expected if pressure gradients were important. We therefore suggest that the streamers are more likely to be comprised of gas clumps constrained to move along a magnetic field lines like ``beads on a wire'' (see e.g., Blandford \\& Payne 1982).  Previous evidence for a magnetic field associated with \\SI\\ was provided by polarization measurements of the SiO $v=0$, $J$=1-0 and $J$=2-1 emission (Tsuboi et al. 1996; Plambeck et al. 2003). After correcting for Faraday rotation, the position angle for the polarization vectors derived by Plambeck et al. (57$^{\\circ}$) is consistent with magnetic field lines threading roughly perpendicular to the disk defined by the SiO masers and 7-mm continuum emission and parallel to the outflow direction traced by the SiO $v=0$ emission and the H$_{2}$O masers (see Goddi et al. 2009a; Greenhill et al., in preparation).  Assuming a magnetic field is present, a magnetocentrifugal wind (e.g., Blandford \\& Payne 1982; K\\\"onigl \\& Pudritz 2000) would be a natural candidate for powering a disk wind from \\SI. In this scenario, gas clumps fragment from the disk, are swept outward along field lines by hydromagnetic forces (e.g., Emmering et al. 1992), and are induced to excite maser emission when irradiated, shocked, or heated by collisions. Magnetic phenomena may also provide an explanation for the elongated emission fronts visible in three of the arms. For example, tangled magnetic field lines (e.g., Kigure \\& Shibata 2005; Banerjee \\& Pudritz 2006) or instabilities within a magnetically-driven flow (e.g., Kim \\& Ostriker 2000) might account for the elevated velocity dispersions and rope-like morphologies of these features (see \\S~\\ref{morph}). Doeleman et al. (1999) originally proposed that these elongated structures might arise at the shocked interfaces of an outflow. However, the observation that the emission fronts are not preferentially oriented perpendicular to the outflow direction (as expected for shock fronts) seems to argue against this interpretation.  It is believed that a necessary condition for launching a magnetically-powered wind is that the vertical magnetic field is close to equipartition (e.g., Ferreira 2007)---i.e., $\\frac{1}{2}nV^{2}=\\frac{1}{2}B^{2}\\mu^{-1}_{0}$, where $n$ is the gas number density, $V$ is the mean velocity of a gas molecule, $B$ is the magnetic field strength, and $\\mu_{0}$ is the permeability of free space. Assuming $n=10^{10}$~cm$^{-3}$ (\\S~3.1.1) and $V$=14.0~\\kms\\ (\\S~\\ref{proper}), the implied magnetic field strength for \\SI\\ is $\\sim$0.3~G. One possible source for this field might be the original interstellar magnetic field threading the molecular cloud out of which \\SI\\ was born. Based on  measurements of OH 1665-MHz masers across a $\\sim10^{4}$~AU region surrounding \\SI, Cohen et al. (2006) derived field strengths of 1.8 to 16.3~mG. Since the OH masers are believed to arise from material with molecular hydrogen density $n_{\\rm H_{2}}\\sim 7\\times10^{6}$~cm$^{-3}$ (Gray et al. 1992), and magnetic field strength is expected to scale as the square root of the gas density, it is plausible that field strengths of order the  value predicted by assuming equipartition might now be present within the denser, SiO-emitting gas.  Direct measurements of the magnetic field strength in the disk of \\SI\\ would be of considerable interest, both for understanding its role during the accretion/outflow process and for providing new clues on the nature of magnetic fields in B-type stars during later evolutionary stages. Little is presently known about the range of magnetic field strengths present in  early-type B stars on the main sequence owing to the difficulty of measuring fields $\\lsim$3~kG in rapidly rotating stars (e.g., Landstreet 1992; Schnerr et al. 2008 and references therein). Moreover, it has been suggested that magnetic fields with strengths as low as a few Gauss and originating at the time of the star's formation might account for the origin of magnetized neutron stars (Ferrario \\& Wickramasinghe 2005).  We note that even if magnetohydrodynamic forces are not the main driving mechanism for the \\SI\\ wind, magnetic forces may still play a key role in shaping the outflow, similar to what has been proposed for planetary nebulae (e.g., Blackman 2008 and references therein). In addition, magnetic fields may provide a mechanism to enhance the coupling between gas and dust grains, thereby increasing the efficiency of a possible dust-driven wind (Hartquist \\& Havnes 1992)."
    },
    "0911/0911.2703_arXiv.txt": {
        "abstract": "I present a previously unpublished method for calculating and modeling multiple lens microlensing events that is based on the image centered ray shooting approach of \\citet{em_planet}. It has been used to model all a wide variety of binary and triple lens systems, but it is designed to efficiently model high-magnification planetary microlensing events, because these high-magnification events are, by far, the most challenging events to model. It is designed to be efficient enough to handle complicated microlensing events, which include more than two lens masses and lens orbital motion. This method uses a polar coordinate integration grid with a smaller grid spacing in the radial direction than in the angular direction, and it employs an integration scheme specifically designed to handle limb darkened sources. I present tests that show that these features achieve second order accuracy for the light curves of a number of high-magnification planetary events. They improve the precision of the calculations by a factor of $>100$ compared to first order integration schemes with the same grid spacing in both directions (for a fixed number of grid points). This method also includes a $\\chi^2$ minimization method, based on the Metropolis algorithm, that allows the jump function to vary in a way that allows quick convergence to $\\chi^2$ minima. Finally, I introduce a global parameter space search strategy that allows a blind search of parameter space for light curve models without requiring $\\chi^2$ minimization over a large grid of fixed parameters. Instead, the parameter space is explored on a grid of initial conditions for a set of $\\chi^2$ minimizations using the full parameter space. While this method may be somewhat faster than methods that find the $\\chi^2$ minima over a large grid of parameters, I argue that the main strength of this method is for events with the signals of multiple planets, where a much higher dimensional parameter space must be explored to find the correct light curve model. ",
        "introduction": "\\label{sec-intro} Gravitational microlensing has opened a new window on the study of extrasolar planets as it is the only method that is currently able to detect low-mass planets in orbits beyond $1\\,$AU. In fact, six of the ten published microlensing planet discoveries have been of planets of less than Saturn's mass \\citep{ogle169,gaudi-ogle109,bennett-ogle109,sumi-ogle368,moa310} and two of these have masses below $10\\mearth$ \\citep{ogle390,bennett-moa192}. The range of planetary separations probed by microlensing is particularly relevant to tests of the core accretion model for planet formation, as microlensing is particularly sensitive to planets just beyond the ``snow-line\" \\citep{ida_lin,kennedy-searth} where core accretion predicts that the most massive planets should form.  Seven of these ten published planetary microlensing events \\citep{ogle71,ogle169,gaudi-ogle109,bennett-moa192, dong-moa400,moa310} have been found in high-magnification events, which although rare, have a much higher planet detection probability \\citep{griest_saf}. A corollary of this is that high-magnification events have significant sensitivity to events with signals from multiple planets \\citep{gaudi_nab_sack} and to planets in stellar binary systems. This point was demonstrated with the discovery of the Jupiter-Saturn analog system OGLE-2006-BLG-109Lb,c \\citep{gaudi-ogle109,bennett-ogle109}. Furthermore, this event also demonstrated that microlensing can detect the orbital motion of a planet when the caustic structures are sufficiently large as can occur for a massive planet \\citep{dong_ogle71} or a ``resonant\" caustic. (When a planet is close to the Einstein ring, the planetary and central caustics merge to form a ``resonant\" caustic.)  This sensitivity to orbital motion and to systems with more than two lens masses makes the modeling of these microlensing events significantly more challenging than other events. In fact, there is currently a backlog of planetary microlensing events that appear to have three or more lens masses. There are three such events discovered through the end of the 2008 Galactic bulge observing season, but the analysis is complete for only one of these \\citep{gaudi-ogle109,bennett-ogle109}. In contrast, the analysis is complete for 8 of the 9 planetary microlensing events discovered through the end of 2008 that can be modeled with a single planet and host star.  In this paper, we present a general method for light curve modeling that is designed to be able to model the most complicated and difficult high-magnification microlensing events. This method has gradually evolved from the first general method for calculating finite source light curves for binary lens events \\citep{em_planet}. Versions of this modeling method have been used to analyze possible planetary events observed in the 1990's \\citep{rhie_ben96,bennett-macho_planets,mps-98blg35}, to make the first successful real-time predictions of caustic crossings toward the Galactic bulge and Magellanic Clouds \\citep{iauc96blg3,macho-98smc1}, and to analyze binary lensing events seen by the MACHO collaboration \\citep{macho-lmc9,macho-binaries}. More recent versions of this code have been used to model all of the known planetary microlensing events, and this code has played a major role in modeling 7 of the 10 published planetary microlensing signals \\citep{bond-moa53,ogle390,ogle169,gaudi-ogle109,bennett-moa192, bennett-ogle109,sumi-ogle368}.  This method has several unique features. In Section~\\ref{sec-int}, I present a numerical integration scheme  with two new features that improve its precision by more than a factor of 100 (for a fixed number of integration grid points). The first feature is an integration scheme that is specifically designed to handle the singular derivative at the limb of a limb darkened source profile. This method also features a polar coordinate integration grid with a much larger grid spacing in the angular than in the radial direction to take advantage of the lensing distortion of the images. In Section~\\ref{sec-lc_calc}, I present tests of this method using a number of previously analyzed planetary microlensing events. These tests confirm the dramatic improvement in precision for high-magnification events.  Next, in Section~\\ref{sec-markov}, I present an adaptive $\\chi^2$ minimization recipe based on the \\citet{metrop} algorithm that is designed to rapidly descend to a minimum of a complicated $\\chi^2$ surface. This method, along with the integration scheme described in Section~\\ref{sec-int}, was critical for the analysis of the one double-planet microlensing event that has been successfully modeled \\citep{gaudi-ogle109,bennett-ogle109}. This event was unusually time consuming to model, due to the important effect of the orbital motion of one of the planets, but the optimizations described in Section~\\ref{sec-int} and \\ref{sec-markov} made the modeling of this event tractable.  A global fit strategy, designed to find all the competitive $\\chi^2$ minima for a microlensing event is presented in Section~\\ref{sec-global}, and in Section~\\ref{sec-global-ex}, I present examples of this method for a number of single-planet events. Finally, in Section~\\ref{sec-conclude}, I discuss ways in which this method can be improved and reach some conclusions. ",
        "conclusions": "\\label{sec-conclude}  I have presented a previously unpublished general method for modeling multiple mass microlensing events that has been optimized for high-magnification events. This method has been developed over a number of years from the first general method for calculating binary lens light curves, the image centered ray shooting method of \\citet{em_planet}. This method is specifically designed to be computationally efficient for the most demanding high magnification events, \\ie, those with more than two lens masses and/or orbital motion. There are three aspects to this method: an efficient method for numerical calculation of the integrals that are needed to calculation microlensing magnification, an adaptive version of the Metropolis algorithm to quickly find a $\\chi^2$ minimum in parameter space, and a global search strategy that can find all the important local $\\chi^2$ minima even in parameter spaces with many dimensions. The computational efficiency of the first two elements of this method was critical for finding the solution for the solution of the only triple-lens microlensing system yet to be published \\citep{gaudi-ogle109}, as well as the study of the lens masses and orbits that are consistent with the light curve data for this event \\citep{bennett-ogle109}.  This method has also been successfully tested on all eight published single planet microlensing events \\citep{bond-moa53,ogle71,ogle390,ogle169,bennett-moa192, dong-moa400,moa310,sumi-ogle368}, as well as four planetary events from the 2009 observing season. For two of these 2009 events, this method found the correct solution before the planetary signal was completed, using data that spanned less than 25\\% of the planetary deviation (although these events did have characteristics that were relatively easy to model). The details of several of these test calculations were presented in Section~\\ref{sec-global-ex}.  In Section~\\ref{sec-lc_calc}, I demonstrated the dramatic improvement in high-magnification light curve calculation precision given by my limb-darkening optimized integration method and polar coordinate grid with a much larger angular than radial grid spacing. For the high-magnification events tested, these features improve the light curve calculation precision by a factor of $\\simgt 100$. Nevertheless, it should still be possible to make significant additional improvements. One improvement would be to implement the hexadecapole approximation \\citep{pej_hey,gould-hex} to do finite source calculations, where the point source approximation does not quite work. This will reduce the number of lens magnification calculations that require the full finite source integrals. For events without caustic crossings, this can dramatically improve calculation efficiency \\citep{dong_ogle71}, but for most caustic crossing events, the improvement is likely to be only a factor of two or so.  Further modifications to the integration grid scheme I present here could provide a more dramatic improvement in computational efficiency. I have used a 16:1 grid spacing ratio for the most efficient high-magnification light curve calculations presented here. However, this is a compromise value. The integration of the bright images that are generated by the primary lens star could probably benefit from a larger ratio, but the smaller images (or parts of images) that are directly perturbed by the planet are done more efficiently with a more modest grid-spacing ratio. This compromise could be avoided by going to an adaptive grid scheme in which the dimensions of the grid are adjusted to match the distortion of the images over the entire image plane. Of course, it would be quite time consuming directly calculate the lens distortions over the entire lens plane, so this would require a simpler prescription to estimate the lens distortions. It might be that a direct calculation over a coarse grid would work.  Of course, the main goal of this method is to be able to model complicated microlensing events that have yet to be successfully modeled. So, the best demonstration of the strength of this method would be the successful modeling of a number of these events."
    },
    "0911/0911.2716.txt": {
        "abstract": "{}{We present a new Bayesian non-parametric deprojection algorithm DOPING (Deprojection of Observed Photometry using and INverse Gambit), that is designed to extract 3-D luminosity density distributions $\\rho$ from observed surface brightness maps $I$, in generalised geometries, while taking into account changes in intrinsic shape with radius, using a penalised likelihood approach and an MCMC optimiser.}{We provide the most likely solution to the integral equation that represents deprojection of the measured $I$ to $\\rho$. In order to keep the solution modular, we choose to express $\\rho$ as a function of the line-of-sight (LOS) coordinate $z$. We calculate the extent of the system along the ${\\bf z}$-axis, for a given point on the image that lies within an identified isophotal annulus. The extent along the LOS is binned and density is held a constant over each such $z$-bin. The code begins with a seed density and at the beginning of an iterative step, the trial $\\rho$ is updated. Comparison of the projection of the current choice of $\\rho$ and the observed $I$ defines the likelihood function (which is supplemented by Laplacian regularisation), the maximal region of which is sought by the optimiser (Metropolis Hastings).}{The algorithm is successfully tested on a set of test galaxies, the morphology of which ranges from an elliptical galaxy with varying eccentricity to an infinitesimally thin disk galaxy marked by an abruptly varying eccentricity profile. Applications are made to faint dwarf elliptical galaxy Ic~3019 and another dwarf elliptical that is characterised by a central spheroidal nuclear component superimposed upon a more extended flattened component. The result of deprojection of the X-ray image of cluster A1413 - assumed triaxial - the axial ratios and inclination of which are taken from the literature, is also presented.}{} ",
        "introduction": "\\label{sec:intro} \\noindent An integral step in the construction of dynamical models of galaxies is the recovery of the intrinsic luminosity density from the surface brightness that is observed projected on the plane of the sky \\citep{magog98, kronawitter00, krajnovic04}. Such deprojection is non-trivial and indeed offers no unique solutions except for very specific configurations of geometry and inclination. As demonstrated by \\cite{rybicki87}, the deprojection is degenerate for axisymmetric systems viewed at inclination angles, $i$, other than edge-on ($i=90$\\deg). This is a consequence of the fact that the observed surface brightness cannot yield any information on a density term whose Fourier transform is non-zero only within a cone of half angle $90^{\\circ} -i$ (the ``cone of ignorance'').  \\cite{binneygerhard} report a family of analytical konus densities. \\cite{kochanekrybicki} found a family of konus densities that have arbitrary densities in the equatorial plane. \\cite{vandenbosch} finds that konus densities contribute at most a few percent of the total galactic mass to the centre of elliptical galaxies with nuclear cusps, implying that their dynamical influence is minimal. \\cite{magorrian} suggests that nearly face-on disk-like konus densities can be recognised via the unique signature they imprint on the line-of-sight (LOS) velocity profiles.  These problems notwithstanding, a great deal of effort has been put into the development of methodologies aimed at deprojecting two-dimensional photometric information. These include parametric formalisms designed by \\cite{palmer}, \\cite{bendinelli} and \\cite{cappellari}, as well as non-parametric methods, such as the Richardson-Lucy Inversion scheme \\citep{richardson, lucy} and a method by \\cite{romkoch} (hereafter, RK).  The parametric methods work by making series expansions of the density (or brightness) and fit the brightness (or density) to the coefficient of the expansions; convergence is defined at a preset accuracy level. In \\cite{bendinelli}, the density is derived via a Gaussian expansion of the surface brightness profile, an idea further developed by \\cite{cappellari}. In \\cite{palmer}, the density is expanded in terms of angular polynomials and the projections of these are then fit to the surface brightness. These methods suffer from the basic drawback that the answer depends on the choice of the basis functions. Thus, the solution is forced to conform to a subset of all possible solutions. Even more worrisome is the fact that the validity of the goodness of fit measures, or ``$\\chi^2$ quantities'' that are employed in these schemes to identify acceptable fits, is questionable, particularly in the presence of inhomogeneous noise \\citep{bissantz_01}. On the whole, the fitting of non-linear functions, to what is usually noise-ridden incomplete data, over large dynamical ranges, is worrisome.  Deprojection of surface brightness profiles has also been attempted with the Abel integral equation, under the assumption of sphericity or axisymmetry with an edge-on inclination \\citep{merrittmeylanmayor, merritttremblay, gebhardt96}.  An example of a non-parametric inversion scheme is the Richardson-Lucy algorithm \\citep{richardson, lucy}, which has been widely used in the stellar dynamical context. It is a simple deprojection scheme that works by iterating toward increasingly better approximations to the density that fits the data. However, absolute convergence is not sought in this framework --- rather, the iterations are stopped when the density is judged to be a good fit to the observations. In lieu of this imposed clause in the code, progressive iterations would produce increasingly more unphysical densities. This lack of a robust convergence criterion is cause for dissatisfaction with the Richardson Lucy scheme. Moreover, implementations of the same, within Astronomy, have not incorporated either radial variations of intrinsic shape or deviations from sphericity.  The shortcomings of Lucy's algorithm are overcome by the analysis of RK in their incorporation of monotonicity and positivity into the sought solution and by their more satisfying convergence criterion. RK construct their density profile as a series of stacked blocks in the space of one quadrant of the meridional plane of their axisymmetric (by assumption) galaxy. The density values at radially adjacent locations are connected through linear interpolation. The summation of the projections of each of the density blocks along the LOS, give the surface brightness at a given location in the plane of the sky (where a brightness measurement is reported). This estimated brightness is then compared to the observed brightness; the algorithm attempts to minimise this statistic while imposing the smoothing condition through a bias function, thus providing a more satisfying convergence criterion than included in Lucy's algorithm. However, this scheme too fails to allow for deprojection in general triaxial geometries; specifically, it is designed to reproduce axisymmetric systems. Moreover, the validity of interpolation, in systems where local gradients can be steep, is also worrisome.  \\cite{magorrian} also advances a scheme similar to RK's expect that he implements a penalty function in his definition of likelihood. The imposed penalty is achieves nearest-neighbour smoothing. The fundamental shortcomings of this scheme are the same as what plagues RK's method. There is no apriori reason to belive that the galaxy under consideration is axisymmetric; triaxiality is a much more general model. Moreover, the requirement in this work, for density to behave like a power-law on local scales, implies that the method will fail in the presence of even moderate gradients in density; in this sense, the recovered answer could be sensitive to the binning details of the 2D grid on which the density structure is placed.  % with radially independent %axial ratio (eccentricity) and major axis position angle. The latter %assumptions are often unsatisfactory (e.g. Merritt, Meylan \\& Mayor %1997), and are indeed violated by most galaxies, for which isophotal %twists and varying projected eccentricity are often observed %(e.g. Ferrarese et al. 2006).  In this paper, we present a new, Bayesian non-parametric algorithm that implements an MCMC optimiser, in order to tackle the deprojection of observed photometry of galaxies. Completely free-form solutions for the 3D density are provided with the constraint of positivity imposed by hand. The scheme is a penalised likelihood procedure. We refer to this algorithm as DOPING - an acronym for Deprojection of Observed Photometry using an INverse Gambit. The algorithm is easy to implement and each run typically takes about a a couple of minutes on a state-of-the art personal computer.  The most distinguishing feature of DOPING is that it can tackle deprojection in virtually any geometry, as long as we can express the intrinsic shape parameters (such as eccentricities) in analytical relations with the projected shape parameters. DOPING can be applied to deproject surface brightness maps of elliptical as well as disk galaxies. This is possible, while taking intrinsic shape variation into account.  \\footnotetext{Such a broad class of solid shapes are perhaps an overkill for astrophysical applications, but these shapes do indeed show up in nature, for example, the wide range of shapes of images of deposited nanoparticles (taken with Transmission Electron Microscope) indicate that these would require such a description \\citep{yong_06}.}   %Thus, DOPING is able to perform deprojection of faceted %3-D nano-crystallites deposited on a substrate, in the %realization of an electronic device (Chakrabarty $\\&$ Paul, in %preparation). However, it is possible that unlike the nano-islands, in %other irregular systems we would not have a handle on the inclination %and/or the extent of different points in the system along the %line-of-sight (as in irregular galaxies). For such systems, %deprojection is only possible under assumptions.  The first major application of this algorithm to galaxies (Chakrabarty $\\&$ McCall, 2009) is the study of the deprojected luminosity profiles of 100 early type galaxies observed as part of the ACS Virgo Cluster Survey \\citep{coteacs04}. As these systems do not exhibit significant variations in position angle, the preliminary version of DOPING that is presented here, considers position angle to be a constant. However the code can account for changes in position angle and the skeletal scheme for the inclusion of a radially varying position angle is presented later in \\S5.6.  The following section begins with a discussion of the broad framework of our algorithm DOPING, a moves to an exposition of the technical details of the code in \\S2. In \\S3, we talk about the application of DOPING to a test case and corroborate the robustness of the algorithm. Application to the real ACS photometry of the galaxy vcc9 is also discussed in \\S4. \\S3 explores the effects of varying ambient conditions such as inclination and the assumed geometry. \\S5 is devoted to discussions and conclusions that are to be drawn from the work. In the appendix, a discussion of the details of various aspects of DOPING is presented. ",
        "conclusions": "\\noindent Thus, DOPING is advanced as a simple but powerful deprojection algorithm, that can be treated as a black-box by the user, is fast and offers 3-D density distributions in general geometries, without resorting to making unconstrained approximations to the form of the density or blindly accepting validity of commonly used goodness of fit measures in light of the inhomogeneous errors of measurement \\citep{bissantz_01}.  The greatest novelty borne by DOPING is its applicability to general systems. Even though in the above examples, triaxial systems were investigated - ranging from razor thin discs to 2-component or bulge+disk galaxies - even when inclination is unknown DOPING can deal with all systems that offer an analytical relation between the intrinsic and measured projected shape parameters. The class of geometries that bear an $m$-fold symmetry allows for this. 3-D modelling of images of systems with even more general geometries can also be performed with DOPING, as long as the system can be imaged at various known inclinations. That deprojection into the third dimension is possible in such general geometries, in a non-parametric way, is due to the fact that the representation of the sought density is modular, i.e. not dependent on a characteristic of the system geometry.  The recovery of the 3-D density is performed iteratively, by searching for the most likely 3-D density structure that projects to the observed 2-D image. This search is robustly undertaken by an MCMC optimiser. Since the choice of the 3-D density is constrained via its projection, distinct 3-D density distributions will project to the same 2-D image. To lift this degeneracy, in DOPING, the system geometry and orientation are specified completely. That axisymmetry is an invalid assumption - at least in real disc-like galaxies - was shown above. Consequently, the description of galaxy geometry as triaxial is a suitable alternative. However, triaxiality entails two axial ratios and inclinations, not all of which can be specified, given the constricted level of achievable observed information. When observed information is available, DOPING can perform deprojection under triaxiality without invoking assumptions, while in lieu of the same, assumptions are invoked. The former case is demonstrated above via the example of deprojection of a galaxy cluster. A benefit of the speed of the algorithm is that a suite of 3-D density models, corresponding to a range of assumed values, is achieved quickly.  Thus, the strengths of DOPING include generalised applicability, ability to incorporate substructure and non-parametric density recovery, along with logistical advantages such as high speed of runs and user-friendliness. Such characteristics render DOPING a very useful tool in three dimensional modelling. In particular, the all important estimation of galactic masses will be aided by a tool such as DOPING."
    },
    "0911/0911.0654_arXiv.txt": {
        "abstract": "We derive a compact, semi-algebraic expression for the cold quark matter equation of state (EoS) in a covariant model that exhibits coincident deconfinement and chiral symmetry restoring transitions in-medium. Along the way we obtain algebraic expressions for: the number- and scalar-density distributions in both the confining Nambu and deconfined Wigner phases; and the vacuum-pressure difference between these phases, which defines a bag constant. The confining interaction materially alters the distribution functions from those of a Fermi gas and consequently has a significant impact on the model's thermodynamic properties, which is apparent in the EoS. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0911/0911.2698_arXiv.txt": {
        "abstract": "By identifying the recently introduced Barbero-Immirzi field with the QCD axion, the strong {\\it CP} problem can be solved through the Peccei-Quinn mechanism. A specific energy scale for the Peccei-Quinn symmetry breaking is naturally predicted by this model. This provides a complete dynamical setting to evaluate the contribution of such an axion to the cold dark matter content of the Universe. Furthermore, a tight upper bound on the tensor-to-scalar ratio production of primordial gravitational waves can be fixed, representing a strong experimental test for this model. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0911/0911.5247_arXiv.txt": {
        "abstract": "{Contrary to theoretical expectation, surprisingly low concentrations of molecular oxygen, \\molo, have been found  in the interstellar medium by means of orbiting telescopes.} {Observations of the $(N_J=1_1-1_0)$ ground state transition of \\molo\\ with the Odin satellite resulted in a \\gapprox\\,$5 \\sigma$ detection toward the dense core \\roa. At the frequency of the line, 119\\,GHz, the Odin telescope has a beam width of 10\\amin, larger than the size of the dense core, so that the precise nature of the emitting source and its exact location and extent are unknown. The current investigation is  intended to remedy this.} {Telluric absorption makes ground based \\molo\\ observations essentially impossible and observations had to be done from space. Millimetre-wave telescopes on space platforms were necessarily small, which resulted in large, several arcminutes wide, beam patterns. Although the Earth's atmosphere is entirely opaque to low-lying \\molo\\ transitions, it  allows ground based observations of the much rarer \\moloisoiso\\ in favourable conditions and at much higher angular resolution with larger telescopes. In addition, \\roa\\ exhibits both multiple radial velocity systems and considerable velocity gradients. Extensive mapping of the region in the proxy \\coiso\\ $(J=3-2)$ line can be expected to help identify the \\molo\\ source on the basis of its line shape and Doppler velocity. Line opacities were determined from observations of optically thin \\coisoiso\\ $(J=3-2)$ at selected positions.} {During several observing periods, two \\coiso\\ intensity maxima in \\roa\\ were searched for in the \\moloisotwo\\ line at 234\\,GHz with the 12\\,m APEX telescope. These positions are associated with peaks in the mm-continuum emission from dust.  Our observations resulted in an upper limit on the integrated \\moloiso\\ intensity of \\tadv\\,$ < 0.01$\\,K\\,\\kms\\ ($3 \\sigma$) into the \\asecdot{26}{5} beam.} {Examining the evidence, which is based primarily on observations in lines of \\moloiso\\ and \\coiso, leads us to conclude that the source of observed \\molo\\ emission is most likely confined to the central regions of the \\roac. In this limited area, implied \\molo\\ abundances could thus be higher than previously reported, by up to two orders of magnitude.} ",
        "introduction": "Oxygen is the most abundant of the astronomical metals \\citep[e.g.,][and~references~therein]{asplund2009}. Consequently, in its molecular form, it was also expected to be very abundant in the UV-shielded regions inside molecular clouds \\citep[e.g.,][]{black1984,bergin2000,charnley2001,roberts2002,spaans2001,viti2001,willacy2002,quan2008} and to contribute significantly to the cooling, hence the energy balance, of dense clouds \\citep{goldsmith1978}.  Because of the high \\molo\\ content in the Earth's atmosphere, astronomical \\molo\\ sources cannot be observed from the ground. Dedicated space missions\\footnote{ SWAS in 1998, see {\\ttfamily http://cfa-www.harvard.edu/swas/}, and Odin in 2001, see {\\ttfamily http://www.snsb.se/eng\\_odin\\_intro.shtml} } came into operation near the beginning of the new millenium. Their unsuccessful searches \\citep{goldsmith2000,pagani2003} were highly disappointing and it was hard to understand that, in the interstellar medium (ISM), \\molo\\ is an elusive species (see references cited above).  Eventually, after more than 20 days of Odin-observing during three different runs, came a real break-through: for the very first time, \\molo\\ was finally detected in the ISM \\citep{larsson2007}. The \\molo\\ emitting object, \\roa, is a dense clump \\citep{loren1990} in a region of active star formation (L\\,1688). On the basis of theoretical model calculations, the detection on the basis of theoretical model calculationsof this kind of source had earlier been predicted by \\citet{black1984} and \\citet{marechal1997a}, where the latter authors made their specific prediction with regard to Odin.  Odin carries a 1.1\\,m telescope which is designed for observations in the submillimetre regime, between roughly 480 and 580\\,GHz ($0.5-0.6$\\,mm). However, the \\molo\\ discovery was made with a dedicated 119\\,GHz (2.5\\,mm) receiver aboard Odin, fix-tuned to the frequency of the ground state O$_2$ ($N_J = 1_1 - 1_0$) transition at 118\\,750.343\\,MHz. At this frequency, the telescope beam size is 10\\amin, larger than the angular dimension of the dense \\roac, which is about 4\\amin\\ \\citep[FWHM of devonvolved CS core,][]{liseau1995}.  It follows that the true \\molo\\ source is likely under-resolved, the consequence of which directly affects estimates of the abundance of \\molo, i.e. $N$(\\molo)/$N$(\\molh): depending on the adopted model, the Odin observations imply an abundance which is currently uncertain by two orders of magnitude \\citep{liseau2005}.  In Fig.\\,2 of \\citet{larsson2007}, the Odin-\\molo\\ line is compared to transitions of other molecular species in \\roa. Whereas lines of \\water\\ and CO are optically very thick over large parts of the cloud and have self-absorbed profiles, the optically thin \\molo\\ line displays a simple, Gaussian shape. This line shape is similar to that of a carbon recombination line, displayed at the top of the figure and which most likely originates in the \\ro-PDR. If also the main source of \\molo\\ emission, the abundance would indeed be very low.  However, the \\molo\\ line shape is also similar to that of the C$^{18}$O (3-2) line, also shown in the figure. This suggests that the C$^{18}$O line can be used as a tracer of the molecular oxygen emission and we set out to map the 10\\amin\\ Odin beam in the (3-2) transition of C$^{18}$O with the APEX beam of size 19\\asec.  It was expected that a detailed comparison of the line centre velocity with that of the \\molo\\ line would help to narrow down the exact location of the \\molo\\ emission, since two distinct velocity components are known to be present in \\roa. This information is needed to understand, where, i.e. in what physical conditions, the majority of the \\molo\\ molecules is excited: in the cold and dense dark cores \\citep{difrancesco2004}, in the extended warm Photon Dominated Region \\citep[PDR;][]{hollenbach2009} or in the hot shocked gas of the outflow from VLA\\,1623 \\citep{liseau2009}? With the \\coiso\\ proxy for \\molo\\ emission, probable emission regions were identified, which were then observed for \\moloisotwo.  There exists earlier work for this line and the \\roc. \\citet{goldsmith1985} observed \\roa\\ in the same transition and with comparable beam size (26\\asec), albeit at an offset 11\\asec\\,E and 61\\asec\\,N relative to the position of SM\\,1N. They obtained  \\trs$ <120$\\,mK ($1 \\sigma$) over 0.34\\,\\kms. At similar channel resolution (0.32\\,\\kms) and toward essentially the same position, \\citet{liszt1985} obtained an rms-noise value of \\trs\\,$<17.5$\\,mK with the 12\\,m NRAO telescope (34\\asec). These papers also present energy level diagrams. Observations made with the 10\\,m telescope of the Caltech Submillimeter Observatory (CSO) in July 1991 and for the position \\rahms{16}{23}{25}, \\decdms{$-24$}{15}{49} (B1950) \\footnote{This corresponds to \\radot{16}{26}{26}{4}, \\decdms{$-24$}{22}{33} in J2000 coordinates and is at (+25\\asec, +80\\asec) relative to the origin of the \\coiso\\ map (Fig.\\,\\ref{c18o32_map}).}  resulted in an rms noise temperature of 16\\,mK in a 0.25\\,\\kms\\ velocity bin and of 12\\,mK after binning to 0.50\\,\\kms\\ (E. van Dishoeck, J. Keene \\& T. Phillips, private communication).  The derivation of molecular abundances requires knowledge of the \\molh\\ column density. One of the widely exploited techniques to estimate $N$(\\molh) is to use observations of C$^{18}$O, the transitions of which in many cases can be shown to be optically thin. We discovered, however, that in the dense core regions of \\roa, this not to be the case everywhere and that appropriate opacity corrections using the $^{13}$C$^{18}$O line needed to be made.  This paper is organised as follows: In Sect.\\,2, our APEX observations of the \\roca\\ in transitions of \\moloiso, \\coiso\\ and \\coisoiso\\ are described. Sect.\\,3 presents our results, which are discussed in Sect.\\,4. Finally, in Sect.\\,5 our main conclusions are briefly summarised. ",
        "conclusions": "Summarising, we briefly conclude the following:  \\begin{itemize} \\item[$\\bullet$] \\coiso\\,(3-2) mapping observations with APEX of a $10^{\\prime} \\times 5^{\\prime}$ region in \\roa\\ have revealed a complex radial velocity field. The central $200^{\\prime \\prime} \\times 200^{\\prime \\prime}$ have been spatially sampled at the Nyquist frequency. \\item[$\\bullet$] The \\vlsr\\ of the \\molo\\,119\\,GHz line appears confined to a particular region (SM\\,1), which is also a prime emitter in \\coiso\\ and N$_2$H$^+$. \\item[$\\bullet$] The observation of \\moloiso\\ toward SM\\,1 (P\\,3) and SM\\,1N (P\\,2) resulted in upper limits. Combined with the \\coiso\\ data, this leads to a ratio of $N$(\\coiso) to $N$(\\moloiso) much larger than unity. \\item[$\\bullet$] From the \\molo\\ and \\moloiso\\ observations we infer an \\molo\\ abundance $5 \\times 10^{-7} < {\\rm O_2} \\le 2.5 \\times 10^{-6}$. \\item[$\\bullet$] The \\molo\\ source is likely relatively compact, within a factor of a few of 1 arcminute, and should become readily detectable by upcoming Herschel HIFI observations. \\end{itemize}  \\acknowledgement{We wish to thank Cathy Horellou and Daniel Johansson for making part of the APEX observations.}"
    },
    "0911/0911.0181_arXiv.txt": {
        "abstract": "We report the discovery of a new Wolf-Rayet star in Aquila via detection of its circumstellar nebula (reminiscent of ring nebulae associated with late WN stars) using the {\\it Spitzer Space Telescope} archival data. Our spectroscopic follow-up of the central point source associated with the nebula showed that it is a WN7h star (we named it WR\\,121b). We analyzed the spectrum of WR\\,121b by using the Potsdam Wolf-Rayet (PoWR) model atmospheres, obtaining a stellar temperature of $\\simeq 50$\\,kK. The stellar wind composition is dominated by helium with $\\sim 20$ per cent of hydrogen. The stellar spectrum is highly reddened ($E_{B-V} = 2.85$\\,mag). Adopting an absolute magnitude of $M_v = -5.7$, the star has a luminosity of $\\log L/L_{\\odot} = 5.75$ and a mass-loss rate of $10^{-4.7} \\, \\msun \\, {\\rm yr}^{-1}$, and resides at a distance of 6.3 kpc. We searched for a possible parent cluster of WR\\,121b and found that this star is located at $\\simeq 1$ degree from the young star cluster embedded in the giant H\\,{\\sc ii} region W43 (containing a WN7+a/OB? star -- WR\\,121a). We also discovered a bow shock around the O9.5III star ALS\\,9956, located at $\\simeq 0.5$ degree from the cluster. We discuss the possibility that WR\\,121b and ALS\\,9956 are runaway stars ejected from the cluster in W43. ",
        "introduction": "\\label{sec:intro}   Wolf-Rayet (WR) stars are evolved massive stars possessing a strong stellar wind. Interaction of the WR wind with the material lost during the preceding evolutionary phase results in the origin of a compact circumstellar nebula (e.g. Chevalier \\& Imamura 1983; Robberto et al. 1993; Brighenti \\& D'Ercole 1995). The majority of WR circumstellar nebulae is associated with late WN-type (WNL) stars (Gvaramadze et al. 2009 and references therein), i.e. with stars that only recently entered the WR phase (e.g. Moffat, Drissen \\& Robert 1989). These nebulae could be sources of mid- and far-infrared (IR) emission (e.g. Mathis et al. 1992; Hutsemekers 1997; Marston et al. 1999) and can be detected with the modern IR surveys. Barniske et al.\\ (2008) found circumstellar emission from dust and gas around two extremely luminous WN stars near the Galactic center.  The detection of compact nebulae reminiscent of nebulae associated with known WNL and Luminous Blue Variable (LBV) stars might suggest that their central stars are evolved massive stars as well and could be used for selection of candidate WNL, LBV or related stars for spectroscopic follow-ups. The importance of this channel for search for evolved massive stars was pointed out by Gvaramadze et al. (2009), who reported the discovery of a ring nebula in Cygnus (using the archival data from the Cygnus-X Spitzer Legacy Survey; Hora et al. 2009) and the results of spectroscopic follow-up of its central star, showing that the star belongs to the WN8-9h subtype. Inspired by this discovery, we searched for more ring nebulae in the archival data of the {\\it Spitzer Space Telescope} and found dozens of objects (Gvaramadze, Kniazev \\& Fabrika 2009; cf.\\ Carey et al.\\ 2009).  In this paper, we present the results of study of one of the stars selected via detection of their circumstellar nebulae. Our spectroscopic follow-up of this star showed that it is a WR star of the WN7h subtype, which provides a further proof of the effectiveness of our method for revealing massive post--main-sequence stars. ",
        "conclusions": "We have discovered a new WNL (WN7h) star via detection of a circular IR nebula (typical of evolved massive stars) and spectroscopic follow-up of its central star. This discovery provides further evidence that the circumstellar nebulae associated with WR stars are inherent almost exclusively to WR stars of the WNL subclass. It also shows that the IR imaging could be an effective tool for revealing WNL and related evolved massive stars. The new WNL star is one of many dozens of (candidate) evolved massive stars revealed via their association with compact IR nebulae, detected in the MIPS $24 \\, \\mu$m data from the {\\it Spitzer Space Telescope} archive. The spectroscopic study of these stars is currently underway and its results will be presented in the forthcoming papers."
    },
    "0911/0911.3657_arXiv.txt": {
        "abstract": "The surface magnetic field strength of white dwarfs is observed to vary from very little to around $10^9\\,$G. Here we examine the proposal that the strongest fields are generated by dynamo action during the common envelope phase of strongly interacting stars that leads to binary systems containing at least one white dwarf. The resulting magnetic field depends strongly on the electrical conductivity of the white dwarf, the lifetime of the convective envelope and the variability of the magnetic dynamo. We assess the various energy sources available and estimate necessary lifetimes of the common envelope. In the case of a dynamo that leads a randomly oriented magnetic field we find that the induced field is confined to a thin boundary layer at the surface of the white dwarf. This then decays away rapidly upon dispersal of the common envelope. The residual field is typically less than $10^{-8}$ times the strength of the external field. Only in the case where there is some preferential direction to the dynamo-generated field can an induced field, that avoids rapid decay, be produced. We show that a surface field of magnitude a few\\,per cent of the external field may be produced after a few Myr. In this case the residual field strength is roughly proportional to the lifetime of the dynamo activity. ",
        "introduction": "\\label{intro}  Surveys of the galactic white dwarf (WD) population have discovered magnetic field strengths ranging up to about $10^9\\,\\rm{G}$ \\citep{schmidt2003}. Typically WDs fall into two categories, those with field strengths of the order $10^6\\,\\rm{G}$ or higher and those with fields weaker than around $10^5\\,\\rm{G}$. We focus here on highly magnetic WDs (hereinafter MWDs) with field strengths greater than $10^6\\,\\rm{G}$. \\citet{landstreet1971} proposed a fossil field mechanism based on the evolution of Ap/Bp stars but it is difficult to argue that such a field can survive stellar evolution in which all parts of the star have passed through convective phases. Here we build on the proposal of \\citet{tout2008} that the origin of such fields lies in the interaction between the WD and its companion star in a binary system.  This assertion is based on observations from the SDSS that approximately $10$\\,per cent of isolated WDs are highly magnetic \\citep{liebert2005} as are $25$\\,per cent in cataclysmic variables \\citep{wickramasinghe2000} while there are none to be found in wide detached binary systems.  Of the 1\\,253 binary systems comprising a WD and non-degenerate M-dwarf star surveyed in the SDSS Data Release Five \\citep{silvestri2007} none have been identified with magnetic fields greater than the detection limit of about $3\\,\\rm{MG}$. The relatively high occurrence of MWDs in strongly interacting binaries compared with elsewhere suggests that the generation of the strong magnetic fields is likely the result of the interaction and subsequent evolution of the binary system. The MWDs observed in isolated systems may be explained by either the total disruption of the companion star during unstable mass transfer or the coalescence of the MWD and the core of its companion following loss of sufficient orbital energy to the common envelope (hereinafter CE) or via gravitational radiation.  When a giant with a degenerate core expands beyond its Roche Lobe mass transfer may proceed on a dynamical time scale.  A dense, typically main-sequence star, companion cannot accrete the overflowing material fast enough so that it instead swells up to form a giant CE. As a result of energy transfer and angular momentum to the CE during orbital decay of the dense companion and the remnant core, strong differential rotation is established within the envelope. Also, owing to its size and thermal characteristics together with the nuclear energy source at the core, the CE is expected to be largely convective. In the mechanism proposed by Tout et al. this is expected to drive strong dynamo action giving rise to powerful magnetic fields \\citep{regos1995,tout1992}. If sufficiently strong dynamo action occurs in the CE then comparable magnetic fields may be induced in the degenerate core (DC) that will evolve into a WD once the envelope has been removed.  We show here that strong surface fields can result from CE evolution. The strength of such fields is highly dependent on the electrical conductivity of DC, the lifetime of the CE and the variability of the magnetic dynamo.  In section 2 we outline the various sources of energy in the binary/CE system and their relation to the energy requirements of the magnetic dynamo. In section 3 we discuss the governing equations of the system and derive the form of the magnetic field for a general spatially and temporally varying magnetic diffusivity $\\eta({\\textbf{\\em{r}}},t)$ via the static induction equation. Then in section 4 we give an overview of the numerical methods we have used to solve the various stages of the problem. In section 5 we present our results and we discuss these and conclude in section 6. ",
        "conclusions": "The statistical occurrence of magnetic WDs in close interacting binary systems suggests that the origin of their strong fields is a product of some feature of their binary evolution. Indeed, the argument that the fields originate through flux conservation during the collapse of Ap/Bp stars does not explain the higher frequency in interacting binaries. By building on the proposal that the magnetic field is generated by dynamo action within a common convective envelope we have shown that strong fields may be transferred to the DC.  These fields can then be preserved upon dissipation of the envelope.  Following the dissipation of the CE, we have found that the system rapidly relaxes on a timescale of about $7\\,\\rm{Myr}$ to the state we would anticipate given no external field. Although there is no significant dissipation of magnetic energy during this time, the redistribution of magnetic flux towards the interior of the WD results in significant decay of the surface field by around a factor of at least $10-100$. Further decay is prevented by the increasing electrical conductivity towards the interior of the WD. This prevents field diffusing beyond around $r=0.9r_{\\rm{c}}$. Once the field has relaxed to its new configuration it continues to decay but on a timescale much longer than the time scale for cooling of the WD so we expect any further loss of field strength to be minimal.  The maximal rate of transfer of field energy from the CE to the DC occurs when the field is kept fixed. In this case the residual strength of the field following dissipation of the CE is linearly related to the lifetime of the CE. If we combine equations (\\ref{me}) and (\\ref{res}) we find  \\begin{equation} M=5.7\\times10^{35}\\left(\\frac{B_{\\rm{res}}}{10^7\\,\\rm{G}}\\right)^2\\left(\\frac{a}{0.01\\,\\rm{au}}\\right)\\left(\\frac{r_{\\rm{I}}}{0.01R_{\\odot}}\\right)^2\\left(\\frac{t_{\\rm{CE}}}{2\\times10^{15}\\,\\rm{s}}\\right)^{-2}\\,\\rm{J}. \\end{equation}  \\noindent If we assume that all of the energy released through orbital decay is transformed into magnetic energy then this requires an envelope lifetime of $1.1\\times10^{12}\\,$s or $3.6\\times10^4\\,$yr to produce a $10^7\\,$G magnetic field. This lifetime is extended if there is variation in the direction of the field produced by the dynamo or there are energy sinks other than the magnetic field. Because the required envelope lifetime scales linearly with $B_{\\rm{res}}$, weaker fields could be produced on a much shorter time scale. Energetically there is nothing to prevent formation of a $10^9\\,$G MWD. As shown, the final field strength scales proportionally with the CE lifetime. If the dynamo is confined to a smaller region or the CE lifetime is extended then this strength of field could be produced. Scenarios this extreme seem unlikely and this is complemented by the rarity of MWDs with such strong magnetic fields. We do not rule out the possibility that the strongest fields may be produced by the merger of a white dwarf and the core of a giant star which have both been strongly magnetized during CE phases.  We have also shown that the production of sufficiently strong fields is almost certainly dependent on some preferred orientation of the magnetic dynamo. In the case where there is no preferred direction the DC magnetic field is confined to a layer of thickness about $(\\frac{\\alpha}{\\eta})^{-1/2}$, where $\\alpha$ is the frequency for field variation, and is almost entirely meridional. Despite the generation of strong surface fields during the CE phase in this case, the total magnetic energy of the DC field is constrained by the depth of the magnetic layer. Consequently once the CE disperses and the field of the exposed WD begins to diffuse inwards, the majority of the field strength is lost. This result is confirmed by the numerical simulations. Typically a field strength of $10^{-5}$ relative to the average magnetic field strength of the dynamo is the largest possible residual field. This may be sufficient to produce WDs with weaker fields but the formation of the strongest WD magnetic fields would require far more energy than the system could provide. In the case where the field has some preferred orientation, the proportion of field retained is increased up to a few per cent depending on the lifetime of the CE dynamo.  Whilst the arguments we have presented here suggest that the origin of the fields of MWDs may be the result of dynamo action in CEs of closely interacting binaries, there are still many uncertainties. Most of these are the result of insufficient understanding of the formation and evolution of CEs. The field structures we have used to construct our models are simple but complex field geometries are observed in WDs such as Feige~7 and KPD 0253+5052. If CE dynamos are responsible for the origins of MWDs, they should support complex geometries. Further progress depends on a better understanding of the physical processes and energy transfer within the CE."
    },
    "0911/0911.4131_arXiv.txt": {
        "abstract": "Type III solar radio storms, observed at frequencies below $\\sim$ 16 MHz by space borne radio experiments, correspond to the quasi-continuous, bursty emission of electron beams onto open field lines above active regions. The mechanisms by which a storm can persist in some cases for more than a solar rotation whilst exhibiting considerable radio activity are poorly understood. To address this issue, the statistical properties of a type III storm observed by the STEREO/WAVES radio experiment are presented, examining both the brightness distribution and (for the first time) the waiting-time distribution. Single power law behavior is observed in the number distribution as a function of brightness; the power law index is $\\sim$ 2.1 and is largely independent of frequency. The waiting-time distribution is found to be consistent with a piecewise-constant Poisson process. This indicates that during the storm individual type III bursts occur independently and suggests that the storm dynamics are consistent with avalanche type behavior in the underlying active region. ",
        "introduction": "Type III solar radio storms consist of many individual type III radio bursts produced quasi-continuously over several days \\citep{fain70a,fain70b,fain71}. Such storms can persist for long intervals -- they have been observed for up to half a solar rotation, and in some cases over a whole rotation  \\citep{fain70a}. It has been shown that they are associated with both active regions and metric type I storms \\citep{kays87,gopa04}. The type III storm emission is produced by electron beams on open field lines \\citep{boug84a, boug84b}, and so type III storm activity may shed light on the dynamics of the open/closed field line boundaries associated with the underlying active region. The dynamics of type III storms are still not well understood, particularly the mechanisms by which they are able to persist for such long times whilst at the same time generating considerable numbers of radio bursts. Furthermore, type III storms have also been observed as precursors to coronal mass ejections \\citep{rein01}, and so a better understanding of their properties may help to elucidate the dynamics of active regions prior to eruptive events.  One important statistical property is the distribution of power $S$ in the individual bursts. In their initial type III storm study, \\citet{fain70a} noted that there were more faint bursts than bright ones. More recently, \\citet{mori07} studied the statistics of type III bursts using Geotail (24Hz - 800 kHz) and Akebono (20kHz - 5MHz) radio data. Using observations covering 4 Bartels rotations and containing both ordinary type III emission and intervals of type III storm activity, they found that the overall occurrence of type III bursts followed a double power law; for example, at a frequency of 860 kHz and below an intensity of $\\sim 10^{-15}$ V$^{2}$m$^{2}$Hz$^{-1}$ ($\\sim$ 2.7 $\\times$ 10$^{-18}$ Wm$^{-2}$Hz$^{-1}$), the number distribution (simply the number of bursts at a given brightness) followed a $S^{-3.6}$ power law whereas above this limit, a $S^{-0.52}$ power law was observed. \\citet{mori07} named the lower flux distribution `micro-type III bursts', and concluded that although they were associated with the same active regions as the `ordinary' type III bursts, they were produced by independent, coexisting processes. This is consistent with the association of the Type III storms with meter-wave Type I `noise storms', which do not have the close association with flares that characterizes ordinary type III bursts \\citep[e.g.][]{gopa04}. In comparison, \\citet{fitz76} found in a study of ordinary type III bursts observed over a period of three months that the number/unit flux density fell off with increasing flux following a $S^{-\\beta}$ power law with $\\beta$ varying from 1.33 at 110 kHz to 1.69 at 4.9 MHz. This in fact appears compatible with the \\citet{mori07} study, since a $S^{-1.69}$ power law in the number/unit flux density corresponds to a $S^{-0.69}$ power law in the number distribution. For comparison, at higher frequencies (164 MHz, 237 MHz and 327 MHz), the number distribution of type I bursts in radio noise storms has been found to follow a similarly steep $S^{-3}$ power law \\citep{merc97}.  A second statistical property is the Waiting-Time Distribution (WTD), that is the distribution of times between individual type III events within the storm. WTD analysis has been applied to soft X-ray flares \\citep{whea00, whea02} and Coronal Mass Ejections \\citep{whea03} and in both cases it has been shown that the data are consistent with a time-dependent Poission process, in which events occur independently. The WTD provides models with extra constraints; for example, flare avalanche models \\citep[e.g.][]{lu1993} predict that flares occur independently, and so the WTD is an important observable.  To our knowledge the WTD has not been used to examine type III storms, particularly at the low radio frequencies typically measured by spacecraft (below $\\sim$ 16 MHz). In this Letter we analyze both the brightness distribution and the WTD of a type III storm observed by the S/WAVES radio experiment \\citep{boug08} on board the STEREO spacecraft \\citep{kais08}, in order to examine in more detail the underlying processes responsible for type III storm emission. ",
        "conclusions": "As mentioned in the introduction, type III emission is thought to arise from non-thermal electron beams on open field lines. Emission may occur at the fundamental or its harmonic; type III bursts whose trajectory intersects the spacecraft show that the emission is at the fundamental mode \\citep{lebl98,dulk2000}, but the possibility of harmonic emission perhaps should not be discounted \\citep{kell80, boug84b}. Since the local electron plasma frequency depends on density, that is $f_{p} [\\rm{Hz}] = 8.98\\times10^{3} \\times (\\it{n}_{\\rm{e}} [\\rm{cm}^{-3}])^{1/2}$, a density model is required to determine the height of the radio emission. Using the model developed by \\citet[]{lebl98}, in this example 2 MHz corresponds to $\\sim$ 3 $R_{S}$ and 9 MHz corresponds to $\\sim$ 1.8 $R_{S}$. The model is scaled to the solar wind density at 1 AU; between 17 - 20 November 2006 the average solar wind density measured by the ACE spacecraft at 1 AU is approximately 4 $\\rm{cm}^{-3}$. If the type III emission is harmonic, then then the local density is reduced, and the emission comes from higher in the corona - between approximately 2.5 and 9 $R_{S}$ for 9 MHz and 2 MHz.  It is generally accepted that the radio emission arises from the non-linear conversion of Langmuir waves excited at the local plasma frequency by non-thermal electron beams. As the electrons move along the magnetic field, the faster electrons outrun the slower electrons leading to a bump-on-tail distribution and instability \\citep{robi00}. The growth of Langmuir waves causes the distribution to relax to marginal stability; in stochastic growth theory, fluctuations in the ambient density are thought to cause fluctuations about marginal stability which allows wave growth without rapid destruction of the electron beam \\citep{robi92, robi93}. This suggests that the electron beams are being emitted from closer to the Sun. The more general association of type III storms (open field) and type I storms (closed field) suggests that the storm activity is related to reconfiguration of the boundary between open and closed magnetic fields in the underlying active region. Analysis of Potential Field Source Surface maps \\citep{schr03} shows that a small region of open field lines was persistently present adjacent and to the west of the active region during the storm. The interface between the open and closed field regions rotated behind the limb before the active region itself, which may help to explain the observed decline in storm activity prior to the active region rotating past the limb.  In this type III storm, a $S^{-2.1}$ power law was found in the number distribution, corresponding to a $S^{-3.1}$ power law in the number/unit flux density. The power law index was found to be largely independent of frequency. Although the exact relationship between the primary energy release, the energy and properties of the electron beam, and the energy converted to radio emission is not well understood, it is likely that this reflects power-law behavior in the underlying dynamics. In a previous study, \\citet{mori07} observed a $S^{-3.6}$ power-law in the number distribution of the storm (`micro') burst population, but this was averaged over several storm intervals and used data binning. In this analysis, maximum likelihood was used, although this technique should be used with care if there are departures from the power-law behavior due to the method of event selection, background subtraction, etc.  It is necessary to study other events to understand whether this difference represents variability from storm to storm, or is due to some other effect. The rate at which bursts are observed is found to change through the storm. Using a Bayesian block decomposition, we find the observed WTD to be consistent with a piecewise-constant Poisson process, which indicates that individual type III bursts occur independently of one another. It is interesting to note that the Bayesian block structure has no correlation with the occurrence of flares in the active region. A GOES C3.3 event occurred on 12 November 2009 at 2155 UT, for example, in the middle of a long Bayesian block.  These two properties are predicted by avalanche models of driven systems. In the solar context, such models have been developed to understand the properties of solar flares \\citep{lu1993}. This suggests that in the active region under consideration, magnetic energy and flux, slowly injected from below and being driven e.g. by twisting due to foot point motion, is being released in a series of independent events which result in the emission of electron beams onto open field lines. In particular, the long lifetime of type III storms indicates that the active region can remain in such a dynamically balanced state for a prolonged period of time even as the burst rate slowly changes. This stability was not disturbed by the $\\sim$ 21 GOES flares in this time interval from AR 10923, including 5 C-class events before 14 November 2006."
    },
    "0911/0911.0418_arXiv.txt": {
        "abstract": "Incorporating the QCD axion and simultaneously satisfying current constraints on the dark matter density and isocurvature fluctuations requires non-minimal fine-tuning of inflationary parameters or the axion misalignment angle (or both) for Peccei-Quinn symmetry-breaking scales $f_a > 10^{12}$ GeV.  To gauge the degree of tuning in models with many axion-like fields at similar symmetry-breaking scales and masses, as may occur in string theoretic models that include a QCD axion, we introduce a figure of merit ${\\cal F}$ that measures the fractional volume of allowed parameter space: the product of the slow roll parameter $\\epsilon$ and each of the axion misalignment angles, $\\theta_0$. For a single axion, $\\mathcal{F} \\lesssim 10^{-11}$ is needed to avoid conflict with observations. We show that the fine tuning of $\\mathcal{F}$ becomes exponentially more extreme in the case of numerous axion-like fields.  Anthropic arguments are insufficient to explain the fine tuning because the bulk of the anthropically allowed parameter space is observationally ruled out by limits on the cosmic microwave background isocurvature modes.  Therefore, this tuning presents a challenge to the compatibility of string-theoretic models with light axions and inflationary cosmology. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0911/0911.5698_arXiv.txt": {
        "abstract": "{The Gas Pixel Detector, recently developed and continuously improved by Pisa INFN in collaboration with IASF-Roma of INAF, can visualize the tracks produced within a low Z  gas by photoelectrons of few keV. By reconstructing the impact point and the original direction of the photoelectrons, the GPD can measure the linear polarization of X-rays, while preserving the information on the absorption point, the energy and the time of individual photons. Applied to X-ray Astrophysics, in the focus of grazing incidence telescopes, it can perform angular resolved polarimetry with a huge improvement of sensitivity, when compared with the conventional techniques of Bragg diffraction at 45$^\\circ$ and Compton scattering around 90$^\\circ$. This configuration is the basis of  POLARIX and HXMT, two pathfinder missions, and is included in the baseline design of IXO, the very large X-ray telescope under study by NASA, ESA and JAXA.} ",
        "introduction": "Polarimetry is the last unexplored probe in High Energy Astrophysics, being particularly interesting in X-rays. In this energy domain the nonthermal emission is typically predominant and consequently a high degree of polarization is expected  \\citep{Novick1975,Rees1975}. On the basis of a literature developed since the very beginning of X-ray astronomy, polarimetry can easily discriminate among competitive models \\citep{Meszaros1988,Dyks2004,Bucciantini2005}. Moreover the relatively high fluxes potentially allow for collecting a sufficient number of photons to reach minimum detectable polarization at the level of a few percent and below. Indeed, even for an instrument without any systematic effect, a partial degree of polarization is always measured because of statistical fluctuations which can be reduced only with a larger amount of data.  Gas detectors able to image the track of photoelectrons are today the most valuable alternative to the instruments based on classical techniques (Bragg diffraction at 45$^\\circ$ and Thomson/Compton scattering), which have succeeded only in the measurement of the polarization of one of the brightest object in the X-ray sky, the Crab Nebula \\citep{Weisskopf1978,Dean2008}. In the following we describe the Gas Pixel Detector and briefly summarize the opportunity of missions currently discussed for this instrument. ",
        "conclusions": "The Gas Pixel Detector is one of the most advanced instruments to image the tracks of photoelectrons in a gas, both for performances and readiness level. It allows to measure the linear polarization in the energy range $\\sim$1-10 keV, reconstructing also the impact point, the energy and the time of the events. The unique possibility to join the polarimetric informations to timing, spectral and imaging capabilities is particularly interesting in X-ray Astrophysics. A large literature, unfortunately still without a significant experimental feedback, describes polarimetry as a fundamental tool to distinguish among competitive models, often equivalent on the basis of spectral or timing informations alone. The GPD provides very concrete possibilities to perform measurements at the level of 1\\% and below on board small missions, like POLARIX or HXMT, to be launched in a few years. In the future, the GPD will be also part of the International X-ray Observatory which will definitely elevate X-ray polarimetry as a systematic tool for a deeper understanding of many different classes of astrophysical sources."
    },
    "0911/0911.2840_arXiv.txt": {
        "abstract": "{With their period-luminosity relation, ``Classical Cepheids'' (CC) are the most common primary distance indicators within the Local Group, also providing an absolute calibration of important secondary distance indicators. However, the predicted position of these pulsators in the HR diagram, along the so called {\\it blue loop}, that is the expected distribution of Cepheids within the instability strip is affected by several model inputs, reflecting upon the predicted $PL$ relation.} {The aim of this work is to quantitatively evaluate the effects on the theoretical $PL$ relation of current uncertainties on the chemical abundances of Cepheids in the Large Magellanic Cloud (LMC) and on several physical assumptions adopted in the evolutionary models.  We will separately analyse how the different factors influence the evolutionary and pulsational observables and the resulting $PL$ relation.} {To achieve this goal we computed new sets of updated evolutionary and pulsational models.} {As a result, we find that present uncertainties on the most relevant H and He burning reaction rates do not influence in a relevant way the loop extension in temperature. On the contrary, current uncertainties on the LMC chemical composition significantly affect the loop extension and also reflect in the morphology of the instability strip; however their influence on the predicted pulsational parameters is negligible. We also discussed how overshooting and mass loss, sometimes suggested as possible solutions for the long-standing problem of the Cepheid mass discrepancy,  influence the $ML$ relation and the pulsational parameters.} {In summary, the present uncertainties on the physical inputs adopted in the evolutionary codes and in the LMC chemical composition are negligible for the prediction of the main pulsational properties. On the other hand, the inclusion of overshooting in the previous hydrogen burning phase and/or of mass loss is expected to significantly change the resulting theoretical pulsational scenario for Cepheids, as well as the calibration of their distance scale. These systematic effects are expected to influence the theoretical Cepheid calibration of the secondary distance indicators and in turn the resulting evaluation of the Hubble constant.} ",
        "introduction": "\\label{sec:intro}  During the central helium burning phase, intermediate mass stars show an excursion toward higher effective temperature, in the Hertzprung-Russell (HR) diagram, subsequently coming back toward the asymptotic giant branch ({\\em blue loop}). During this phase the stars cross the instability strip, becoming Classical Cepheids.  These variable stars, thanks to their characteristic Period-Luminosity ($PL$) relation, are the most common primary distance indicators within the Local Group. Moreover, from space, they are observable also in external galaxies, up to distances of 20-30 Mpc \\citep[see e.g.][]{free01}. On this basis they provide the absolute calibration of important secondary distance indicators, such as the maximum luminosity of Supernovae Ia, the Tully-Fisher relation, surface brightness fluctuations, and the planetary nebulae luminosity function, that are the basis for measurement of the Hubble constant \\citep{free01,sa01}.  Any systematic error affecting the Cepheid $PL$ relations affects the extragalactic distance scale and the estimate of the Hubble constant. Therefore, even though the Cepheid $PL$ relation is assumed to be universal and calibrated through the sample in the Large Magellanic Cloud \\citep[LMC, see e.g.][]{mf91,u99}, the possibility that the Cepheid properties and $PL$ relation depend on the chemical composition of the parent galaxies has been actively debated during the last decade \\citep[see e.g.][]{S04,M06,r05,r08,M05,b08}. On the other hand the theoretical prediction for the existence of a $PL$ relation for Classical Cepheids relies on the assumption that intermediate mass stars undergoing central helium burning are characterized by a mass-luminosity ($ML$) relation, as predicted by stellar evolution models. On very general grounds the pulsation period is expected to be tightly correlated with the luminosity, the mass and the effective temperature of the star, for each given chemical composition. It is only thanks to the $ML$ relation that the period can be simply related to the luminosity and the effective temperature, allowing a theoretical prediction for the period-luminosity-color relation \\citep[$PLC$, see][]{f84,f91,f95,cc86,ls86,s88,cl91,s97}.  However the theoretical $ML$ relation depends on the several physical processes included in stellar evolutionary codes as well as on the adopted chemical composition.  By averaging the $PLC$ relation on the color extension of the pulsation instability strip one obtains the $PL$ relation. This implies that the $PL$ relation has a statistical nature and observationally holds only when a statistically significant sample of Cepheids (at the same distance) is available. Thus the predicted temperature extension of the blue loop for a given mass also reflects upon the expected distribution of Cepheids within the instability strip and in turn upon the theoretical $PL$ relation.  This is a very critical point because the predicted extension of the blue loop in the HR diagram is affected by several model inputs (chemical composition, efficiency of external convection, nuclear burning cross sections, etc..) producing effects that are not yet completely understood \\citep[see e.g.][]{Xu2004a,stot94,bru90}. In this context, a detailed comparison of evolutionary and pulsational models with the observations can provide insight into the physics of stellar evolution and pulsation.  The aim of this work is to quantitatively evaluate the effects on the predicted $PL$ relation of current uncertainties on the chemical abundances of Cepheids in the LMC and of several physical assumptions adopted in the evolutionary models.  In particular, we will separately analyse how the different factors influence the evolutionary and pulsational observables (extension in effective temperature of the blue loop, Mass-Luminosity relation, morphology of the instability strip, period distribution) and the resulting $PL$ relation. For the chemical composition we choose that of the LMC with its associated uncertainty.  All the relations based on the newly computed sets of evolutionary and pulsation models will be made available to the community.  In a forthcoming paper (Valle et al. 2009 in preparation) we will focus on the comparison of these new predictions with the observations.  The organization of the paper is the following: in Sect.~\\ref{sec:2} we present the theoretical scenario both from the evolutionary and the pulsational point of view; Sect.~\\ref{sec:blueloop} contains a general discussion about the problems in calculating blue loop models; in Sect~\\ref{sec:rates} the dependence of the predicted extension of the {\\it blue loop} on the uncertainties affecting nuclear reaction rates is investigated; in Sect.~\\ref{sec:deltaz} we focus on the effects of chemical abundance uncertainties on the most relevant evolutionary and pulsational features, whereas in Sect.~\\ref{oversh} we take into account possible noncanonical effects in the evolutionary scenario. The conclusions, reported in Sect.~\\ref{sec:concl}, close the paper. ",
        "conclusions": "\\label{sec:concl}  On the basis of an updated set of evolutionary and pulsation models for Classical Cepheids in the LMC, we have investigated the effect of the uncertainties on the chemical composition and on the physical assumptions adopted in the codes on the most relevant pulsation observables and thus on the theoretical calibration of the Cepheid distance scale. We found that present uncertainties on relevant nuclear reaction rates have only a negligible effect on evolutionary and pulsational theoretical prediction. Moreover the uncertainties in the metal and helium abundances affect the results of the evolutionary computations but do not significantly changethe pulsation scenario. On the other hand, the  still present uncertainties on the efficiency of the overshooting phenomenon in the previous H burning phase and on the mass-loss rates are found to be the most important source of uncertainty in the theoretical Cepheid mass-luminosity and period-luminosity relations. In particular, using a theoretical $PL$ relation that relies on the assumption of mildly overshooting evolutionary models, one infers distances that are significantly shorter than the values obtained when a canonical theoretical $PL$ is used, especially in the optical bands. Therefore we conclude that the uncertainty on the Cepheid $ML$ relation is expected to significantly affect the Cepheid calibration of the extragalactic distance scale and in turn the evaluation of the Hubble constant. The application of the theoretical analysis performed in this paper to LMC Cepheid data will be addressed in a forthcoming paper (Valle et al. 2009 in preparation)."
    },
    "0911/0911.3594_arXiv.txt": {
        "abstract": "{} { Radio halos are faint radio sources usually located at the center of merging clusters of galaxies. These diffuse radio sources are rare, having so far been found only in about 30 clusters of galaxies, suggesting that particular conditions are needed to form and maintain them. It is interesting to investigate the presence of radio halos in close pairs of interacting clusters in order to possibly clarify their origin in relation to the evolutionary state of the merger. In this work, we study the case of the close pair of galaxy clusters A399 and A401. } {A401 is already known to contain a faint radio halo, while a hint of diffuse emission in A399 has been suggested based on the NVSS. To confirm this possibility, we analyzed deeper Very Large Array observations at 1.4 GHz of this cluster.} {We find that the central region of A399 is permeated by a diffuse low-surface brightness radio emission that we classify as a radio halo with a linear size of about 570 kpc and a central brightness of 0.3 $\\mu$Jy/arcsec$^2$. Indeed, given their comparatively small projected distance of $\\sim 3$ Mpc, the pair of galaxy clusters A401 and A399 can be considered as the first example of double radio halo system. The discovery of this double halo is extraordinary given the rarity of these radio sources in general and given that current X-ray data seem to suggest that the two clusters are still in a pre-merger state. Therefore, the origin of the double radio halo is likely to be attributed to the individual merging histories of each cluster separately, rather than to the result of a close encounter between the two systems.} {} ",
        "introduction": "An ever increasing number of galaxy clusters exhibit Mpc-scale synchrotron radio halos associated with the intracluster medium. These elusive radio sources are located at the cluster center and are characterized by a regular shape and extremely low surface brightness at levels of $\\sim 1 \\mu$Jy/arcsec$^2$ at 1.4 GHz. To date, about 30 radio halos are known (see Giovannini et al. 2009 for recent compilation), and most have been found in clusters that show significant evidence of an ongoing merger (e.g. Buote 2001). The rarity of radio halos seems to suggest that particular conditions are required for their formation and maintenance. Merger events may play a strong role in the re-acceleration of the radio-emitting relativistic particles, thus providing the energy that powers these extended sources (e.g. Brunetti et al. 2009). In this context, it is interesting to investigate radio halos in close pairs of interacting clusters of galaxies. These systems may share the same recent dynamical history, so they are unique laboratories for studying the formation of radio halos and possibly clarifying the origin of these radio sources in relation to the evolutionary state of the merger. \\begin{figure*}[t] \\centering \\includegraphics[width=18 cm]{fig1.ps} \\caption{Left: total intensity radio contours of A399 at 1.4 GHz with the VLA in D configuration. The image has an FWHM of 45$''\\times 45''$. The first contour level is drawn at 120 $\\mu$Jy/beam and the rest are spaced by a factor of $\\sqrt{2}$. The sensitivity (1$\\sigma$) is 40 $\\mu$Jy/beam. The total intensity radio contours overlaid on the XMM X-ray image in the 0.2$-$12 keV band. The X-ray image has been convolved with a Gaussian of $\\sigma=12\\arcsec$. Right: zoom of the total intensity radio contours in the center of A399 at 1.4 GHz with the VLA in C array. The image has an FWHM of 14$''\\times 13''$. The first contour level is drawn at 120 $\\mu$Jy/beam and the others are spaced by a factor of $\\sqrt{2}$. The sensitivity (1$\\sigma$) is 40 $\\mu$Jy/beam. The contours of the radio intensity are overlaid on the optical DSS2 image.} \\label{fig1} \\end{figure*}  In this work we analyze the close pair of galaxy clusters A399 and A401, which are located at a redshift of $z=0.071806$ and $z=0.073664$ (Oegerle \\& Hill 2001). This is an exceptional pair, since the two clusters are separated in projection by an angular distance of 36\\arcmin, corresponding to a linear separation of 3 Mpc in our chosen cosmology (see below). A399 and A401 are two rich clusters of galaxies with a similar average gas temperature of $kT\\simeq 7$ and $kT\\simeq 8$ keV, respectively. The X-ray luminosities in the 0.1$-$2.4 keV band for the two clusters are $L_{X}=3.8$ and $6.5 \\times10^{44}$ erg/s (Reiprich \\& B{\\\"o}hringer 2002). The X-ray excess and the slight temperature increase in the region between the two clusters indicate a physical link between this pair of clusters (e.g. Fujita et al. 1996, Fabian et al. 1997, Markevitch et al. 1998). Recent XMM X-ray analyses (Sakelliou \\& Ponman 2004, Bourdin \\& Mazzotta 2008) show that neither of the two clusters contains a cooling core. However, the reasonably relaxed morphology of the clusters and the absence of major temperature anomalies argue against models in which A399 and A401 have already experienced a close encounter. The link region between A399 and A401 has been studied with the Suzaku satellite by Fujita et al. (2008). The metallicity of the hot intergalactic medium in this region is found to be comparable to those in the inner regions of the clusters. The authors conclude that the metal enrichment in this region is likely caused by strong winds from galaxies before the cluster formation rather than by cluster collision. These analyses indicate that the clusters are just starting to mildly interact and that the sub-features found in their inner regions are related to the individual merging histories of each cluster separately, rather than to the remnant of a previous merger of the two systems. Indeed, in the light of all these indications, it seems likely that A399 and A401 are two merger remnants, just before they merge together to form a single rich cluster of galaxies.  A401 is already known to contain a faint radio halo (Harris et al. 1980, Roland et al. 1981, Giovannini et al. 1999, Bacchi et al. 2003), while a hint for a diffuse emission in A399 has been suggested by Sakelliou \\& Ponman (2004) based on the NRAO VLA Sky Survey (NVSS) images. In this letter, we report on the results of deeper observations performed at 1.4 GHz with Very Large Array (VLA) in C and D configurations.  Throughout this paper we assume a $\\Lambda$CDM cosmology with $H_0$ = 71 km s$^{-1}$Mpc$^{-1}$, $\\Omega_m$ = 0.27, and $\\Omega_{\\Lambda}$ = 0.73. At the distance of A399, 1\\arcsec~ corresponds to 1.35 kpc. ",
        "conclusions": "In summary, we find that the central region of A399 is permeated by a diffuse low-surface brightness radio emission which is classified as a radio halo with a length-scale of about 190 kpc and a central brightness of 0.3 $\\mu$Jy/arcsec$^2$. Indeed, given their projected distance of $\\sim 3$ Mpc, the pair of galaxy clusters A401 and A399 can be considered as the first example of double radio halo system.  The origin of the double radio halo is likely attributed to the individual merging histories of each cluster separately, rather than to the result of a close encounter between the two systems. The two clusters must have experienced a similar sequence of mergers which prevented the formation of the cool cores and left significant residual kinetic energy in the gas to power the radio halos.  It seems likely that the two small radio halos in A399 and A401 are in proximity to merge together eventually forming a single giant radio halo or a large scale diffuse emission identified with a filamentary supercluster-like structure. This offers new clues to how giant radio halos could form, since it is now clear that magnetic fields and relativistic particles can already be present in the intracluster medium of two clusters well {\\it before} a major merger event between them."
    },
    "0911/0911.3900_arXiv.txt": {
        "abstract": "We map the stellar structure of the Galactic thick disk and halo by applying color-magnitude diagram (CMD) fitting to photometric data from the SEGUE survey.  The SEGUE imaging scans allow, for the first time, a comprehensive analysis of Milky Way structure at both high and low latitudes using uniform SDSS photometry.  Incorporating photometry of all relevant stars simultaneously, CMD fitting bypasses the need to choose single tracer populations.  Using old stellar populations of differing metallicities as templates we obtain a sparse three-dimensional map of the stellar mass distribution at $|Z|>$1 kpc.  Fitting a smooth Milky Way model comprising exponential thin and thick disks and an axisymmetric power-law halo allows us to constrain the structural parameters of the thick disk and halo. The thick-disk scale height and length are well constrained at $0.75\\pm0.07$ kpc and $4.1\\pm0.4$ kpc, respectively. We find a stellar halo flattening within $\\sim$25 kpc of $c/a=0.88\\pm0.03$ and a power-law index of $2.75\\pm0.07$ (for 7 kpc $\\lesssim R_{GC} \\lesssim$~30 kpc). The model fits yield thick-disk and stellar halo densities at the solar location of $\\rho_{thick,\\sun}=10^{-2.3\\pm0.1}$ $M_\\sun$pc$^{-3}$ and $\\rho_{halo,\\sun}=10^{-4.20\\pm0.05}$ $M_\\sun$pc$^{-3}$, averaging over any substructures. Our analysis provides the first clear in situ evidence for a radial metallicity gradient in the Milky Way's stellar halo: within $R\\lesssim$15 kpc the stellar halo has a mean metallicity of $[Fe/H]\\simeq-1.6$, which shifts to $[Fe/H]\\simeq-2.2$ at larger radii, in line with the two-component halo deduced by \\cite{carollo07} from a local kinematic analysis. Subtraction of the best-fit smooth and symmetric model from the overall density maps reveals a wealth of substructures at all latitudes, some attributable to known streams and overdensities, and some new. A simple warp cannot account for the low latitude substructure, as overdensities occur simultaneously above and below the Galactic plane. ",
        "introduction": "\\label{sec:intro}  Studying the stellar structure of the Milky Way has a long history, as early models based on star counts were already constructed by the likes of \\cite{herschel1785} and \\cite{kapteyn1922}. Today we know that four `components' make for a sensible approximate description of the Milky Way stellar body, each with different structural, kinematic and population characteristics: the bulge, the thin and the thick disk, and the stellar halo. Accurate determination of the properties of these components is a difficult undertaking, as it requires surveys with (1) sufficient sky coverage to assess the overall geometry, (2) sufficient depth for mapping stars to $\\sim$10 to 30 kpc, and (3) sufficient information (e.g., color-based luminosity estimates) to obtain reasonable distance estimates for these stars. Furthermore, the presence of dust in the plane of the disk clouds our view of most of the Galaxy at wavelengths shorter than a few microns. For this reason, most previous efforts to study Milky Way structure at low latitudes have used (near-)infrared data, such as the work by \\cite{kent91} based on data from the Spacelab infrared telescope, or more recent work using 2MASS data \\citep[e.g.][]{momany06,reyle09} or the GLIMPSE survey \\citep{benjamin05}. While near-infrared surveys provide the only way of studying the low latitude regions, because they suffer much less from dust extinction, they are currently less sensitive than their optical counterparts and often rely on tracer populations with less accurate distances.  Over the past few years, data from 2MASS \\citep{2mass} and the Sloan Digital Sky Survey \\citep[SDSS;][]{york00,dr7} have played a key role in revolutionizing our empirical picture of the stellar components of the Milky Way. As the SDSS survey avoided the Galactic plane, much of the work using SDSS data has focused away from the bulk of the Milky Way's stars to high galactic latitudes. Nonetheless, \\cite{juric08} have been able to constrain the structural parameters of the stellar halo and the thin and thick disk.  Beyond the global parameters, these surveys have also led to the discovery of substructure in both the stellar disk \\citep{newberg02,ibata03,martin04,bellazzini06} and halo \\citep{newberg02,fieldofstreams,grillmair06,hercaqui,juric08,bell08}. This detailed substructure provides information on how the Milky Way formed and evolved; as galaxies assemble through mergers and accretions of smaller systems, these events leave identifiable signatures in their structure and kinematics.  This also has cosmological implications, as the structure of the dark halo and the interactions with dark subhalos can significantly influence the appearance of a galaxy \\citep[e.g.][]{kazantzidis08,younger08}.  In this paper we derive a new map of the stellar distribution in the Milky Way, in order to study the large-scale structure of the thick disk and stellar halo. Though covering only the Northern celestial hemisphere, and that only sparsely, this map covers both high and low latitudes and uses rigorously derived stellar distances. This map is based on one of the constituent projects of the extended SDSS survey (SDSS II), SEGUE \\citep[Sloan Extension for Galactic Understanding and Exploration; ][]{segue}, which is an imaging and spectroscopic survey aimed at the study of the Milky Way and its stellar populations. Whereas the main SDSS survey avoided low Galactic latitudes, SEGUE imaging scans go through the Galactic plane (see Figure \\ref{fig:dataoverview}), allowing a `picket-fence' view of the Galaxy at these low latitudes. It is this deep, uniform data set that enables a systematic analysis of stellar structure as function of Galactic latitude and longitude, which we base on color-magnitude diagram (CMD) fitting techniques developed by \\cite{sdssmatch}. A crucial advantage of an analysis based on CMDs is that they bypass the often ad-hoc and sometimes problematic choices of stellar `tracers' (such as K-dwarfs or F-stars) to delineate the structure. Instead, this procedure results immediately in a mass-density map for stellar population components of, say, different metallicity.  To avoid having to fit the complex stellar population make-up of the thin disk, we in the end consider only stars that are more than 1 kpc away from the Galactic mid-plane, and treat the thin disk as a small perturbation to our results. As SEGUE is a northern hemisphere survey, it stays far away from the Galactic bulge, which can be safely ignored. On the other hand, the resulting maps of stellar mass density in the thick disk and halo allow us to measure the structural parameters of these components. In addition, we show that subtracting a smooth Galactic model reveals a wealth of substructure, especially at low latitudes, which we will study in detail in a forthcoming publication.  The outline of this paper is as follows. Section \\ref{sec:data} describes the data used for our analysis, and \\S\\ref{sec:cmdfitting} describes the CMD-fitting techniques we utilize. The resulting stellar distributions are presented in \\S\\ref{sec:results}, where we also fit smooth Galactic models to them, derive structural parameters, and look at deviations from the best fitting model. Finally, we summarize and discuss our results in \\S\\ref{sec:discussion}. ",
        "conclusions": "\\label{sec:discussion}  We have applied CMD-fitting techniques to SEGUE photometric data to study the thick disk and stellar halo of the Galaxy at both high and low latitudes. Three template stellar populations, all with an age of $\\sim$14 Gyr and metallicities of [Fe/H]=$-$0.7, $-$1.3 and $-$2.2, are fit to the data, yielding a three-dimensional census of stellar mass. Since the thick disk and stellar halo consist of predominantly very old stars, differences in turn-off color are indicative of differences in metallicity. A change of turn-off color is seen in the halo at distances of 15 to 20 kpc, which is therefore interpreted as a change in metallicity. At distances smaller than $\\sim$10 kpc our results indicate a mean metallicity of [Fe/H]$\\sim-$1.6, similar to the work by \\cite{ivezic08}, who find a value of [Fe/H]$\\sim-$1.5 out to $\\sim$9 kpc. At larger distances, D$\\gtrsim$15 kpc, our results indicate that the mean metallicity of the stellar halo is [Fe/H]$\\sim-$2.2, however. This transition is consistent with the inference from local samples of stars (D$\\lesssim$4 kpc) by \\cite{carollo07} and preliminary results from \\cite{beers09}, but this is the first time this change in population properties is measured {\\it in situ}.  Structural parameters of the Galaxy can be derived from the resulting stellar-density maps through the comparison with models. Our fits of models with an exponential thin and thick disk and a power-law halo (Eq. \\ref{eq:galacticmodel}) yield constraints on the thick disk local density, $\\rho_{thick,\\sun}$, scale height, $H_{thick}$, and scale length, $L_{thick}$, and the halo local density, $\\rho_{halo,\\sun}$, flattening, $q_{halo}$ and power-law index, $n_{halo}$.  As pointed out by \\cite{fieldofstreams}, \\cite{bell08}, and \\cite{juric08}, the SDSS data of the stellar halo contains clear evidence of the presence of substructure with respect to smooth halo models. \\cite{bell08} also demonstrated that the larger known substructures in the halo significantly influence their model fits. To test the sensitivity of our model fits to known substructure, we perform our fits to the full data set, as well as to a data set with the Sagittarius stream and Monoceros overdensity removed. The results are tabulated in Table \\ref{tab:modelfits}, and are not strongly influenced by the presence of these structures. In either case, the reduced $\\chi^2$ of the best fit, 4.2 and 3.9, shows that the smooth model is a poor fit to the data.  The values we find for the local density of the thick disk, $\\rho_{thick,\\sun}=10^{-2.3\\pm0.1}$ $M_\\sun$pc$^{-3}$ and the local halo density, $\\rho_{halo,\\sun}=10^{-4.20\\pm0.05}$ $M_\\sun$pc$^{-3}$, agree well with star-count studies. In addition, our results for the thick-disk structural parameters, $H_{thick}=0.75\\pm0.07$ kpc and $L_{thick}=4.1\\pm0.4$, are well within the range of parameters found in previous studies \\citep{siegel02} and consistent with the results of \\cite{juric08}. For the flattening of the stellar halo, where we find $q_{halo}=0.88\\pm0.04$, our analysis seems to diverge from the work by \\cite{bell08} and \\cite{juric08}, who use SDSS data to study the halo out to similar distances. On the other hand, the power-law index we recover, $n_{halo}=-2.75\\pm0.07$, is in excellent agreement.  Subtracting the best model from the stellar-density maps unveils abundant substructure (Fig. \\ref{fig:subtracted}). Most of the overdensities seen can be attributed to know structures in the halo and outer disk of the Galaxy: the Sagittarius stream, Virgo overdensity, and the Monoceros overdensity or Low Latitude Stream. Whereas the origin of the Sagittarius stream is known to be the result of the disruption of the Sagittarius dwarf galaxy, the nature of the other two entities is currently under debate, in particular the interpretation of the overdensities at low Galactic latitudes remains contentious.  A detailed analysis of all substructures is outside the scope of this paper, but an in-depth study of the low latitude substructure is planned to be presented in a future publication."
    },
    "0911/0911.0958_arXiv.txt": {
        "abstract": "{Radiation driven winds are the likely origin of AGN outflows, and are believed to be a fundamental component of the inner structure of AGNs. Several hydrodynamical models have been developed, showing that these winds can be effectively launched from AGN accretion discs.} {Here we want to study the acceleration phase of line-driven winds in AGNs, in order to examine the physical conditions for the existence of such winds for a wide variety of initial conditions. } {We built a simple and fast non-hydrodynamic model QWIND, where we assume that a wind is launched from the accretion disc at supersonic velocities of the order of a few 10$^2$~km/s and we concentrate on the subsequent supersonic phase, when the wind is accelerated to final velocities up to 10$^4$~km/s.  } {We show that, with a set of initial parameters in agreement with observations in AGNs, this model can produce a wind with terminal velocities of the order of $10^4$~km/s. There are three zones in the wind, only the middle one of which can launch a wind: in the inner zone the wind is too ionized and so experiences only the Compton radiation force which is not effective in accelerating gas. This inner \u2018failed wind\u2019 however plays an important role in shielding the next zone, lowering the ionization parameter there. In the middle zone the lower ionization of the gas leads to a much larger radiation force and the gas achieves escape velocity This middle zone is quite thin (about 100 gravitational radii). The outer, third, zone is shielded from the UV radiation by the central wind zone and so does not achieve a high enough acceleration to reach escape velocity. We also describe a simple analytic approximation of our model, based on neglecting the effects of gravity during the acceleration phase. This analytic approach is in agreement with the results of the numerical code, and is a powerful way to check whether a radiation driven wind can be accelerated with a given set of initial parameters.} {Our analytical analysis and the fast QWIND model are in agreement with more complex hydrodynamical models, and allow an exploration of the dependence of the wind properties for a wide set of initial parameters: black hole mass, Eddington ratio, initial density profile, X-ray to UV ratio. } ",
        "introduction": "Outflowing winds are now believed to be common, and quite possibly ubiquitous, in the inner parts of Active Galactic Nuclei (AGNs) and quasars. The most striking evidence comes from the 10\\% of Broad Absorption Line (BAL) quasars which show broad absorption lines blueshifts spanning 10-20 thousands km~s$^{-1}$.  Less spectacular, but more common, evidence for outflows comes from the 50\\% of AGN with Narrow Absorption Lines (NALs) both UV and X-ray (Reynolds 1997, George et al. 1998, Crenshaw et al. 1999, Krongold et al. 2003, Vestergaard 2003, Piconcelli et al. 2005, Ganguly \\& Brotherton~2007).  Several attempts have been made to explain the origin of such winds. Two main scenarios have been suggested: magnetically driven winds (Blandford \\& Payne 1982, Konigl \\& Payne 1994, Konigl \\& Kartje~1994, Everett \\& Murray~2007), or radiation driven winds.  One of the most promising explanations is through radiation line-driven winds arising from the accretion disc. The physics of radiation line-driven winds based on the work of Sobolev (1960) was developed by Castor, Abbott \\& Klein (1975, hereafter CAK), and Abbott (1982, 1986) for winds from hot stars. CAK showed that resonance line absorption in an accelerating flow could be hundreds of times more effective than pure electron scattering. More recently the same theory has been applied to accretion discs in compact binaries ( Proga, Drew \\& Stone 1998) and to AGN discs (Murray et al.~1995, Proga et al. 2000, hereafter P00, Proga 2003). These authors use the powerful hydrodynamical code ZEUS2D (Stone \\& Norman 1992) to solve the wind equations, assuming a Shakura-Sunyaev (1972, hereafter SS) $\\alpha$-disc, and plausible initial conditions. The main results of these works are (1) the demonstration that a wind can be launched and accelerated up to velocities of the order of 10$^4$~km~s$^{-1}$, and (2) the determination of the wind geometry and physical state for the given starting conditions. These models have been tested with several different choices of the initial conditions (Proga \\& Kallman~2004), showing that a wind can arise for a wide range of black hole mass and accretion rates.  Further recent improvements in this field are the modeling of Schurch \\& Done~(2007), where the interaction between the X-ray radiation and the outflowing wind is throughly analyzed, and the analysis of possible effects of radiation winds at larger distances from the accretion disc, such as on a parsec-scale X-ray heated torus (Dorodnitsyn et al.~2008), and on large-scale AGN outflows (Kurosawa \\& Proga~2009).   One key aspect of the physics of radiation driven winds, clearly emerging from the models mentioned above, is that regardless of the details of the initial launching phase, it is impossible to avoid the strong, probably dominant, effect of radiation pressure in the subsequent acceleration phase, where the wind gains more than 99\\% of its kinetic energy. This is easily estimated from the comparison between the amount of momentum absorbed by the gas and its final momentum. For example, Hamann (1998) estimates that at least $\\sim25$\\% of the UV radiation emitted by the Broad Absorption Line Quasar PG~1254+047 is absorbed by a gas with column density of the order of 10$^{23}$ cm$^{-2}$. It is sufficient that the luminosity of this source is of the order of 10\\% of the Eddington luminosity to conclude that the momentum in the outflowing wind is of the order of that absorbed in the UV wavelength range.  The distinction between the {\\em launching} and {\\em acceleration} phases of an accretion disc wind, which we made above, allows the modelling of the problem to be separated into these two parts. {\\bf The {\\em launching} phase has been investigated in several numerical simulations, both in hydrodynamical (Ohsuga et al.~2005) and magneto-hydrodynamical (Hawley \\& Krolik~2006) regimes. Recently, the creation of outflows from accretion discs has been investigated through MHD simulations including the effects of radiation pressures (Ohsuga et al.~2009). }  The main aim of the work presented here is to explore a wide set of initial conditions for the {\\em acceleration} phase of a wind in AGN  using a deliberately simplified approach in the hope that this can produce an intuitive understanding of quasar winds and so guide future detailed simulations. We will make use both of analytic approximations and of a simplified numerical code for quasar winds (which we named {\\em QWIND}). {\\em QWIND} treats in detail the radiation force mechanism, but not the internal gas pressure in the wind. As the wind velocity in the acceleration phase is always many times the thermal velocity of the wind gas, this is a reasonable approach.  Our study is motivated by the observations of fast outflows in Broad absorption Line (BAL) quasars (Weymann 1997), by photometric evidence for a highly flattened structure in BAL winds (Ogle et al.~1999) and by results suggesting an axially symmetric, but not spherical, spatial distribution of the Broad Emission Line (BEL) gas (Wills \\& Browne 1986, Brotherton 1996, Maiolino et al. 2001, Rokaki et al. 2003). Elvis (2000) has proposed a specific quasar unification model based on this kind of structure.  This model predicts that a geometrically thin, funnel-shaped outflow is present in all AGNs and quasars, which allows the observational properties of the different classes of sources to be explained through orientation effects. This model provided our initial motivation, as we suspected that a radiation driven wind might naturally create a thin, funnel-shaped structure.  We therefore developed the {\\em QWIND} model in order to easily and quickly explore the acceleration of a radiation-driven wind for the huge range of physically possible initial conditions, in terms of black hole mass, accretion rate, relative strength of the X-ray radiation, initial density and temperature of the outflowing gas. Our approach does not add more physical insights for the single wind solutions, than the already available hydrodynamical models mentioned above, but allows a much wider exploration of the initial parameters space.  The purpose of this paper is to present the model code, {\\em QWIND}, and to demonstrate that {\\em QWIND} produces results in agreement with both an analytic treatment and with the more complex and detailed approach of P00. The structure of this paper is the following: in Section 2 we briefly review previous results on the topic of radiation-driven winds.  In Section 3 we present the {\\em QWIND} code and we show examples of its possible applications. In Section 4 we present an analytical treatment of the wind equations, which provides interesting constraints on the existence of such winds, as a function of the initial parameters. In Section~5 we briefly compare our results with those obtained by more complex hydrodynamical codes. Finally, in Section~6 we present our conclusions and outline future work. ",
        "conclusions": "We have shown that a radiation-driven wind can be accelerated to velocities of about 10$^4$~km~s$^{-1}$ from a Shakura-Sunyaev accretion disc at a distance from the central black hole of the order of that of the BEL clouds, for densities and temperatures in the range observed in quasars warm absorbers.  Thanks to our simple approach, we were able to make an analytic estimate of the parameter space in which a wind solution is allowed, and to build a fast code, {\\em QWIND}, which can be used to better investigate the dependence on the several unknown initial parameters.  The main limitation of our study is the assumption of a substantial initial velocity (a few 100~km~s$^{-1}$) which makes the wind supersonic from the beginning.  We did not discuss the physical mechanism through which the disc provides the initial kinetic energy to the gas.  Our main results are: \\begin{enumerate} \\item Our model reproduces the global results of more complex hydrodynamical codes.  This makes us confident that our approach can be used to explore the huge space of initial conditions, in order to understand in which cases a wind can be launched.  The results of our analytical study are in agreement with those of the numerical code {\\em QWIND}, in the sense that parameters space regions excluded by the analytical analysis are also excluded with the more precise numerical approach.  \\item For initial parameters typical of AGN warm absorbers, a wind can arise only for relatively high accretion efficiencies ($\\epsilon_{EDD}>0.3$) and in a {\\em narrow range of disc radii}. This is due to overionization at too small radii, and to a too small UV radiation and initial velocity at too large radii, and is an inherent feature of radiation driven disc winds in AGNs.  \\item The inner part of the gas, which is too overionized to be effectively accelerated, forms a ``failed wind'' which removes the X-ray radiation from directions close to the disc plane, allowing a fast decrease of the ionization parameter with the distance form the central source, and so allowing the line driven wind acceleration at larger radii. This component is analogous to the ``hitchhiking gas'' of Murray \\& Chiang~(1995).  \\item We have shown how QWIND can be used to gain a simple physical understanding of the dependence of the wind properties on the initial parameters, such as the X-ray to UV ratio, the Eddington ratio, the initial density and temperature, the central black hole mass. Some of the results are at first counter-intuitive, as, for example, the non-monotonic dependence of the wind final velocity on the Eddington ratio: in some cases more luminous sources can have slower winds, with the other physical parameters being held constant.  \\item The terminal velocity of the wind is typically of the order of $1-2~10^4~$km~s$^{-1}$, and the cone covering factor is a few percent of the lines of sight. This is in reasonable agreement with observations of BAL quasars.  \\item The wind angle above the disc is substantial ($\\sim20$deg), and so is neither equatorial nor polar. Some $\\sim35$\\% of the viewing directions will see the X-ray and UV continuum source through the wind.   \\end{enumerate}  Our findings could have deep implications for wind-embedded BEL models, like that of Elvis (2000).  Black holes accreting in conditions outside the narrow allowed ranges would not have winds, and so neither absorption lines, nor the BELs would be observed. Such 'naked quasars' would be picked up from their non-stellar colors in the Sloan Digital Sky Survey or by their high X-ray flux in X-ray surveys. A few ``lineless'' quasars have been found but they are rare (e.g. McDowell et al.~1995). If AGN winds are radiation-driven, then some self-regulation mechanism must be involved in the accretion process, to keep the AGNs at the conditions needed for a radiation driven wind.  Our main aims at this stage were to check the reliability of our approach, and to give significant examples of its use. The next major step, which will be presented in a forthcoming paper, is a systematic survey of initial parameters, to give a more general view of the physical conditions needed to accelerate winds by line driving. In particular, we will investigate how the results are affected by changes in the wind inner radius and in the initial density and velocity profiles."
    },
    "0911/0911.5099_arXiv.txt": {
        "abstract": "{The INFN and INAF Italian research institutes developed a space-borne X-Ray polarimetry experiment based on a X-Ray telescope, focussing the radiation on a Gas Pixel Detector (GPD). The instrument obtains the polarization angle of the absorbed photons from the direction of emission of the photoelectrons as visualized in the GPD. Here we will show how we compute the angular resolution of such an instrument.} ",
        "introduction": "The X-ray telescopes are based on the grazing angle principle. The radiation is reflected with small incidence angles on the surfaces of hyperboloid and paraboloid mirrors and then is focused. The GPD is a gas detector which is able to image the photoelectron tracks. The polarization is measured using the dependence of the photoelectric cross section from the photon polarization direction \\cite[Costa 2001]{Costa0} \\cite[Bellazzini 2007]{Bellazzini}. The photoelectron is emitted with more probability in the direction of the electric field of the photon. The track created by the photoelectron path, is drifted and amplified by the Gas Electron Multiplier (GEM) and collected on a fine sub-divided pixel detector. Using different mixtures of gas it is possible to properly select the energy band of the instrument in the range of about $1-30$ keV. This GPD has the capability to preserve the imaging while reaching a good sensitivity in polarization as well as in spectroscopic and timing measurements. \\begin{table} \\begin{tabular}{|l|ccc} \\hline {\\textbf{Characteristic}} & {\\textbf{Value}} & {\\textbf{Unit}}\\\\ \\hline optics energy band & 0.1-10 & keV\\\\ GPD energy band & 1-30 & keV\\\\ GPD Area  & $15 \\cdot 15$ & $mm^{2}$\\\\ GPD height & 10  & mm\\\\ GPD transistors & $16.5 \\cdot 10^{6}$ & n. \\\\ GPD pixels & $105600$ & n. \\\\ GPD pixel matrix & $300 \\cdot 352$ & n. \\\\ \\hline \\end{tabular} \\label{tab1} \\caption{GPD characteristics} \\end{table} ",
        "conclusions": ""
    },
    "0911/0911.0991_arXiv.txt": {
        "abstract": "Formation-flying studies to date have required continuous and minute corrections of the orbital elements and attitudes of the spacecraft. This increases the complexity, and associated risk, of controlling the formation, which often makes formation-flying studies infeasible for technological and economic reasons. Passive formation-flying is a novel space-flight concept, which offers a remedy to those problems. Spacecraft in a passive formation are allowed to drift and rotate slowly, but by using advanced metrology and statistical modelling methods, their relative positions, velocities, and orientations are determined with very high accuracy. The metrology data is used directly by the payloads to compensate for spacecraft motions in software. The normally very stringent spacecraft control requirements are thereby relaxed, which significantly reduces mission complexity and cost. Space-borne low-frequency radio astronomy has been identified as a key science application for a conceptual pathfinder mission using this novel approach. The mission, called FIRST (\\textbf{F}ormation{}-flying sub{}-\\textbf{I}onospheric \\textbf{R}adio astronomy \\textbf{S}cience and \\textbf{T}echnology) Explorer, is currently under study by the European Space Agency (ESA). Its objective is to demonstrate passive formation-flying and at the same time perform unique world class science with a very high serendipity factor, by opening a new frequency window to astronomy. ",
        "introduction": "In passive formation-flying, the spacecraft are allowed to drift slowly so that no expensive and complex position control systems are required to maintain the spacecraft in predefined positions. Instead, in the passive formation-flying paradigm, the evolving orbits and the orientations of the spacecraft are constantly monitored. By using advanced metrology and statistical modelling methods, the relative positions of the spacecraft can be determined with high accuracy, even if simple ranging sensors are used. This knowledge enables the continuous phase reconstruction of a high-performance radio telescope aperture to be performed, while the individual constellation satellites rotate and drift over time.  The FIRST Explorer constellation consists of six daughter spacecraft with radio astronomy antennas, and a mother spacecraft for science and metrology data processing and communications. The location of the constellation at the second Lagrange point (L2) allows for a stable, low-drift orbit that is sufficiently far away from Earth to avoid severe radio frequency interference (RFI), while at the same being close enough to maintain operations using standard telemetry systems. A novel use of credit card sized mini solar sails on the spacecraft enables the constellation to stay within an overall mission envelope \\cite{JENAM2009}.  The main science objective of this pathfinder mission is to provide an all-sky survey at very low radio frequencies, which cannot be observed from the ground; because the ionosphere effectively acts like a shield for frequencies below $\\sim10$ MHz. Such observations have not been performed since the Radio Astronomy Explorer (RAE 1 and 2) missions in the 1970's, as illustrated by the all-sky map \\cite{Novaco78} in Fig.~\\ref{fig:rae2}, they were very limited in terms of sensitivity and angular resolution. For all practical purposes, the low-frequency Universe is therefore uncharted. A mission that could fill this gap, and produce astronomical data of high quality, has therefore been on the wish list of many astronomers for more than three decades. The major science goal is therefore an all-sky survey below 10 MHz with better than 1 degree resolution.  \\begin{figure}[b] \\includegraphics[width=\\columnwidth]{figures/rae2.eps} \\caption{ RAE-2 all-sky map of the galaxy at 4.70 MHz by J.~C.~Novaco \\& L.~W.~Brown (1978) \\cite{Novaco78}. Repreduced by permission of J.~C.~Novaco and the American Astronomical Society (AAS).} \\label{fig:rae2} \\end{figure}  In addition, the FIRST Explorer can be used to image the low-frequency Sun and the evolution and propagation of coronal mass ejections (CME), which is important for our understanding of space weather and its sometimes severe consequences for vital infrastructure on Earth. It can also perform long-term studies of all the radio-planets by providing dynamic spectra of the intense low frequency radio bursts that originate in the planets' auroral zones. Due to the ionospheric cut-off, those planetary radio emissions are, with the exception of the Jovian decametric radiation, not accessible from Earth. Similar radio emissions are most likely produced also by extra-solar planets with magnetic fields and atmospheres. With an intensity that can exceed their host star and a frequency content directly proportional to the planet's magnetic field strength, those burst radio emissions hold a great promise as beacons for future extra-solar planet searches.  Larger space-observatories based on the enabling passive formation-flying technology, but consisting of several small satellite constellations, are being considered as the next step. These observatories will have enhanced sensitivity and resolving power to address many fundamental science objectives in radio astronomy, such as detection and observation of extra-solar planets as well as comprehensive studies of the dark ages and the epoch of re-ionisation, by observations of high red-shift $21$ cm line emissions from neutral Hydrogen, which is the only radiation believed to have survived from the period before the Universe became transparent \\cite{Loeb2004}. With suitable long integration times, FIRST Explorer can in itself gather high red-shift $21$ cm line spectra that may be able to prove the dark ages hypothesis. A no-result is also very useful since it will put limits on the maximum strength of these emissions at the lowest frequencies, which are not accessible from Earth.  The FIRST Explorer mission concept study was commissioned by the European Space Agency (ESA) \\cite{FIRST} and based on an idea from one of the authors, Dr Constantinos Stavrinidis, Head of the Department of Mechanical Engineering at ESA/ESTEC, to reduce formation flying mission costs. The study objective put forward by Dr Stavrinidis was to:  \\emph{ ``Assess the potential of an innovative and valuable science mission to achieve with low cost, a `passive' formation flying instrument, by allowing the formation to drift, and overcoming the need \\ for demanding and expensive technologies for position control systems to constantly servo the various spacecraft into tightly defined positions''.}  The cost of current and previous formation flying mission designs are driven up by the need to \\ employ sophisticated, often optical and thereto expensive, range sensors to monitor the position and orientation of separate spacecraft to high precision; and the need to use complex active multi-axis micro thrusters on each spacecraft, to constantly servo the constellation into a defined position and orientation.  This is a challenge for a formation involving only two or three spacecraft, but as the number increases the complexities of the servo control algorithm grows exponentially and then the growing time-delays between command and control make the challenge even more difficult.  The passive formation flying concept optimises the ``knowledge'' of the relative positions of the spacecraft -- rather than optimising control of the position of the spacecraft. More sophisticated computational modelling is used as a trade-off against less sophisticated conventional ACOS (Attitude and Orbit Control System) instrumentation using active multi-axis thrusters.  This paper examines one possible mission concept using example design criteria, and based on existing low-cost technology that has been demonstrated in space or on the ground. It is not a full mission design study and many of the system parameters require further detailed design trade-off studies, but the paper illustrates what could be achieved as an alternative to conventional science and technology demonstrator missions.  By adopting a different paradigm, where the instrument adapts to a changing environment rather than the other way round and by postulating a mission scenario in a low-drift environment, we have designed a fresh (holistic) approach to the problem. The science goals have been iterated with the mission architecture approach, to develop a mission concept that both tests the low cost passive formation flying idea, and yet still targets valuable new astronomical science results that cannot be achieved from the ground or with conventional spacecraft designs. The resulting mission design is known as the FIRST Explorer, as it explores both the novel mission technology and control concepts and points the way to larger formations that can address some of the grand science challenges open to future space-borne low frequency radio astronomy missions. ",
        "conclusions": ""
    },
    "0911/0911.0675_arXiv.txt": {
        "abstract": "We present a study of the radial velocity offsets between narrow emission lines and host galaxy lines (stellar absorption and H~I 21-cm emission) in Seyfert galaxies with observed redshifts less than 0.043. We find that 35\\% of the Seyferts in the sample show [O~III] emission lines with blueshifts with respect to their host galaxies exceeding 50 km s$^{-1}$, whereas only 6\\% show redshifts this large, in qualitative agreement with most previous studies. We also find that a greater percentage of Seyfert 1 galaxies show blueshifts than Seyfert 2 galaxies. Using {\\it HST}/STIS spatially-resolved spectra of the Seyfert 2 galaxy NGC~1068 and the Seyfert 1 galaxy NGC~4151, we generate geometric models of their narrow-line regions (NLRs) and inner galactic disks, and show how these models can explain the blueshifted [O~III] emission lines in collapsed STIS spectra of these two Seyferts. We conclude that the combination of mass outflow of ionized gas in the NLR and extinction by dust in the inner disk (primarily in the form of dust spirals) is primarily responsible for the velocity offsets in Seyfert galaxies. More exotic explanations are not needed. We discuss the implications of this result for the velocity offsets found in higher redshift AGN. ",
        "introduction": "Seyfert galaxies, quasars, and other active galactic nuclei (AGN) often show radial velocity offsets between their broad emission lines, narrow emission lines, and host galaxy lines (stellar absorption and/or H~I 21-cm emission), and there is a long history of investigations into the nature of these offsets, beginning with Penston (1977), Netzer (1977), Osterbrock \\& Cohen (1979), and Gaskell (1982). In addition, many studies have concentrated on differences in line profile, asymmetry, width (e.g., full-width at half maximum [$FWHM$]), peak, and/or velocity centroid as a function of ionization potential or critical density for lines from the same region, such as the narrow-line region (NLR; Osterbrock 1981; Pelat, Alloin, \\& Fosbury 1981; De Robertis \\& Osterbrock 1986; De Robertis \\& Shaw 1990; Busco \\& Steiner 1992; Moore, Cohen, \\& Marcy 1996; Komossa et al. 2008a) or broad-line region (BLR; Wilkes 1984; Crenshaw 1986; Osterbrock \\& Mathews 1986; Corbin 1990; Tytler \\& Fan 1992; Sulentic, Marziani, \\& Dultzin-Hacyan 2000; Richards, et al. 2002; Snedden \\& Gaskell 2007). Even for the same emission line from a low-redshift AGN, such as the commonly used [O~III] $\\lambda$5007 line, a perusal of the NASA Extragalactic Database (NED) demonstrates that the heliocentric redshift in terms of radial velocity (cz) can differ substantially in excess of the quoted errors from one study to the next, due to wavelength calibration uncertainties, different measurement techniques, or aperture effects, for example. Thus, some care must be taken when determining velocity offsets between different lines, but there is no doubt from the above references that the offsets in many sources are real and significant.  Recent investigations have sparked renewed interest in the study of velocity offsets, because they have been used to indicate the possible existence of multiple AGN in a galaxy, or an AGN that is displaced with respect to the gravitational center of a galaxy. The most convincing evidence comes from galaxies that show active nuclei that are spatially offset from their optical centers, as well as emission lines that are offset in radial velocity from the systemic velocities of the host galaxies. For example, the evidence for multiple AGN in NGC~3341 ($z =$ 0.027) appears to be incontrovertible (Barth et al. 2008). Three separate nuclei can be identified in a Sloan Digital Sky Survey (SDSS) image of this galaxy. The central nucleus has a LINER/H~II composite spectrum, an offset nucleus at a distance of 5.1 kpc from the center has a Seyfert 2 spectrum at a velocity offset of $-$200 km s$^{-1}$, and another offset nucleus at a distance of 8.4 kpc has a possible LINER spectrum with a similar velocity offset. Barth et al. suggest that the offset nuclei are due to the fueling of the supermassive black holes (SMBHs) in two dwarf galaxies that are merging with the primary disk galaxy.  Another AGN that has generated a lot of interest and speculation is the quasar SDSS J092712.65+294344.0 ($z =$ 0.71), which shows two sets of emission lines (one set of narrow lines and one set of broad plus narrow lines) separated by 2650 km s$^{-1}$ (Komossa, Zhou, \\& Lu 2008b). In this case, the SDSS image shows only an unresolved quasar -- multiple nuclei have not been detected in groundbased images. Most of the suggested explanations for the double set of emission lines are based on mergers or interactions of galaxies and, possibly, their SMBHs. Komossa et al. (2008b) suggest that the system with broad lines is a recoiling SMBH carrying its BLR and high-ionization NLR with it and leaving behind the bulk of the (lower-ionization) NLR. The recoil is due to coalescence of two black holes which, under the right conditions, can result in anisotropic gravitational radiation that carries away linear momentum, ejecting the merged SMBH at velocities up to $\\sim$4000 km s$^{-1}$ in the opposite direction (Campanelli et al. 2007; Baker et al. 2008). Another possible explanation is that this SDSS quasar contains a binary SMBH, in which the BLR and high-ionization NLR are due to accretion onto the smaller SMBH, and the low-ionization NLR is an envelope surrounding the pair of SMBHs (Dotti et al. 2008; Bogdanovi\\'{c}, Eracleous, \\& Sigurdsson 2009). Heckman et al. (2009) suggest a third possibility, which has been used to explain the properties of the nearby AGN NGC 1275: the narrow redshifted lines are associated with a small galaxy falling into the center of a rich cluster of galaxies where it encounters a large galaxy with an AGN that shows both broad and narrow lines. Finally, Shields, Bonning, and Salviander (2009) suggest that the double set of emission lines in SDSS J092712.65+294344.0 may be due to a chance alignment of two AGN, possibly within a massive cluster.  Although two sets of emission lines that are offset in radial velocity are rare, many AGN show velocity offsets between the emission lines from their NLRs and the stellar absorption lines or H~I 21-cm emission from their host galaxies. Comerford et al. (2009) present a study in which 32 of 91 Seyfert 2 galaxies with red host galaxies at 0.34 $< z <$ 0.92 show velocity offsets between their [O~III] emission lines and stellar absorption lines that lie in the range $\\sim$50 km s$^{-1}$ to $\\sim$300 km s$^{-1}$ (two of these AGN show double peaked emission lines as well). Comerford et al. suggest that the velocity offsets are likely due to displaced AGN that are moving with respect to their host galaxies. The same explanations for the double set of emission lines in SDSS J092712.65+294344.0 can then be invoked to explain the velocity offsets between the NLR and host galaxy, if one of the AGN in these scenarios is dispensed with or replaced by an inactive SMBH. However, Comerford et al. (2009) suggest that the most plausible explanation is inspiralling SMBHs after a galaxy merger, which move to the gravitational center of the merger remnant via dynamical friction. They leave open the possibility, however, that the velocity centroids of the emission lines are offset from the systemic velocity of the host galaxy due to a combination of outflows and dust extinction in the NLR, although they say this is an unlikely explanation for their sample.  In this paper, we show that in low-redshift Seyferts, velocity offsets of the narrow emission lines with respect to the host galaxies are indeed due to mass outflows in the NLR and partial extinction of the outflows by circumnuclear dust. This is not a new idea, as a number of previous studies suggested that the combination of dust, either inside or between the NLR clouds, and radial motion, in the form of inflows or outflows, could be the cause of asymmetries and velocity shifts in the narrow lines (Capriotti, Foltz, \\& Byard 1979; Heckman et al. 1981; Whittle 1985; Dahari \\& De Robertis 1988; De Robertis \\& Shaw 1990; Veilleux 1991a). Subsequently, a number of studies based on long-slit observations with ground-based telescopes suggested that outflows and dust obscurations could produce the observed profile asymmetries (Storchi-Bergmann, Wilson, \\& Baldwin 1992;  Arribas, Mediavilla, \\& Garc\\'{i}a-Lorenzo 1996; Christopoulou et al. 1997; Rodr\\'{i}guez-Ardila et al. 2006). We can now investigate this issue in more detail using spatially-resolved optical spectra of the NLR ($\\sim$0\\arcsecpoint1 resolution) from the Space Telescope Imaging Spectrograph (STIS) on board the {\\it Hubble Space Telescope} ({\\it HST}). ",
        "conclusions": "We find that $\\sim$40\\% of Seyfert galaxies show radial velocity offsets $|\\Delta v| \\geq$ 50 km s$^{-1}$ between their narrow emission lines and host galaxy lines. The distribution of offsets peaks near zero km s$^{-1}$, but has a strong blue tail extending up to $\\sim$ $-$250 km s$^{-1}$; only $\\sim$6\\% of the Seyferts in the sample show significantly redshifted emission lines, up to 100 km s$^{-1}$. The Seyfert 1 distribution is shifted to higher blueshifts on average.  To investigate the nature of the velocity offsets, we relied on our previous investigations into the reddening, physical conditions, and kinematics of the NLRs in NGC~1068 (Seyfert 2) and NGC~4151 (Seyfert 1), which were based on spatially resolved {\\it HST}/STIS spectra and phototionization and kinematic models. The parameters from these studies allowed us to generate geometric models of the biconical outflow in the NLR and extinction by dust in the inner galactic disk. We also collapsed the high-resolution {\\it HST}/STIS spectra obtained at multiple slit locations to simulate ground-based observations. The resulting [O~III] emission-line profiles are asymmetric to the blue and have blue-shifted centroids, which can be attributed to preferential extinction of the redshifted clouds in the geometric models of these two Seyferts.  It is clear from the above examples that models that incorporate biconical outflow and extinction by the inner galactic disk are more likely to produce blueshifted velocity centroids than redshifted ones, because the blueshifted clouds are closer to the observer and less likely to be extincted in most cases. However, we have shown that redshifted velocity centroids are possible in certain cases, for example when the disk is highly inclined and covers most of the near cone. Furthermore, one would expect from these models that more Seyfert 1s would show blueshifts than Seyfert 2s, because the near cone would be pointed more directly at us and therefore less extincted on average. All of these features of the geometric models are in agreement with the observed velocity offsets. However, we note that the [O~III] emission in NGC~1068 shows a higher blueshift than that in NGC~4151, which indicates that individual variations in velocities and extinctions of the NLR clouds amongst AGN are important in determining the overall distribution in velocity offsets.  We conclude that the velocity offsets of emission lines in low-$z$ Seyfert galaxies are primarily due to a combination of radial outflow in their NLRs and extinction by dust on the same size scale as their NLRs. More exotic explanations are not required. Furthermore, we find that the dust responsible for the extinction must be primarily exterior to the ionized gas. The dust features seen in the inner kpc of Seyfert galaxies, mostly in the form of dust spirals (Martini et al. 2003, and references therein), are the likely culprits.  As we have already noted, the bumps and peaks in the emission-line profiles are due to large bright knots of emission-line gas traveling at their own peculiar velocities, on top of the general outflow pattern. Thus, another contributor to the velocity offsets could be an uneven distribution of knots in velocity space in an AGN, resulting in more (or brighter) blueshifted knots than redshifted ones, or vice-versa. This effect alone, however, will not lead to an asymmetric distribution of velocity offsets.  It is instructive to compare the distributions of velocity offsets between narrow emission and stellar absorption lines in the low $z$ ($<$ 0.043, Nelson \\& Whittle 1995) and moderate $z$ (0.34 $< z <$ 0.92, Comerford et al. 2009) samples. The magnitudes of the velocity offsets in the two distributions are essentially the same, and they have a similar percentage of ``significant'' ($\\gtsim$~50 km s$^{-1}$) offsets (41\\% at low $z$, 35\\% at moderate $z$). However, the distribution for the moderate $z$ sample is roughly symmetric around zero km s$^{-1}$, which is surprising, because these AGN are likely to have outflows in their NLRs as well. In fact, Comerford et al. use the symmetric distribution as a reason to disfavor the outflows plus dust explanation and favor the inspiralling black hole scenario. This, in turn, leads them to a merger fraction of $\\sim$30\\% for red galaxies at these redshifts, which is a surprising large number, particularly since these galaxies show no evidence for star formation.  The nature of the Comerford et al. (2009) sample provides a couple of possible explanations for a symmetric distribution of velocity offsets in the context of mass outflows. Their sample contains only Seyfert 2 galaxies that have been selected to have red host galaxies, to specifically avoid contamination from ionized gas around star-forming regions. The lack of significant star formation suggests a lack of dusty gas, and therefore a lack of extinction, in their NLRs. The velocity offsets would therefore be dominated by uneven distributions of emission-line knots in at least some of the AGN, with no preference for blueshifted profiles. Another possible explanation is that the red colors are due to highly inclined disk galaxies, which, as shown in Figure 7, will lead to a significant number of redshifts and a more symmetric distribution of velocity offsets. Nevertheless, there is a local example of velocity (and spatial) offsets due to multiple AGN in the nearby galaxy NGC~3341 (Barth et al. 2008), and one might expect more of these at higher redshifts from galaxy mergers. However, the emission-line ratios in NGC~3341 indicate significant star formation, as might be expected from mergers, and if it was at the appropriate redshift, NGC~3341 would have not been included in the Comerford et al. sample. Both the low-$z$ and moderate-$z$ samples are rather small; larger sample would be helpful in clarifying the role that mergers might play in producing velocity offsets.  Finally, we note that ground-based spectra of local Seyfert galaxies, particularly those obtained at high spectral resolution (Vrtilek \\& Carleton 1985; Veilleux 1991b; Nelson \\& Whittle 1995), show a wide variety of narrow emission-line profiles. The profiles are often very irregular in appearance, with bumps or shelves in their red or blue wings, and a number of them show multiple peaks, like those in NGC~1068 (Figure 4), due to the contributions of distinct emission-line knots. Of particular interest is the [O~III] profile of Mrk~78 (Vrtilek \\& Carleton 1985; Whittle et al. 1988; Nelson \\& Whittle 1995), which shows two large well-defined peaks separated by $\\sim$800 km s$^{-1}$, and yet there is no evidence for displaced or double SMBHs in Mrk~78. This velocity separation is actually larger than those in the two ``dual AGN'' of Comerford et al. (2009), which show velocity separations between double-peak profiles of 630 km s$^{-1}$ and 440 km s$^{-1}$. Thus, even these double-peaked AGN may not require the presence of double SMBHs."
    },
    "0911/0911.1359_arXiv.txt": {
        "abstract": "Calculations of cosmological hydrogen recombination are vital for the extraction of cosmological parameters from cosmic microwave background (CMB) observations, and for imposing constraints to inflation and reionization. The \\textit{Planck} mission and future experiments will make high precision measurements of CMB anisotropies at angular scales as small as $\\ell\\sim 2500$, necessitating a calculation of recombination with fractional accuracy of $\\approx 10^{-3}$. Recent work on recombination includes two-photon transitions from high excitation states and many radiative transfer effects. Modern recombination calculations separately follow angular momentum sublevels of the hydrogen atom to accurately treat nonequilibrium effects at late times ($z<900$). The inclusion of extremely high-n ($n\\gtrsim100$) states of hydrogen is then computationally challenging, preventing until now a determination of the maximum $n$ needed to predict CMB anisotropy spectra with sufficient accuracy for \\textit{Planck}. Here, results from a new multi-level-atom code (\\textsc{RecSparse}) are presented. For the first time, `forbidden' quadrupole transitions of hydrogen are included, but shown to be negligible. \\textsc{RecSparse} is designed to quickly calculate recombination histories including extremely high-$n$ states in hydrogen. Histories for a sequence of values as high as $n_{\\rm max}=250$ are computed, keeping track of all angular momentum sublevels and energy shells of the hydrogen atom separately. Use of an insufficiently high $n_{\\rm max}$ value (e.g., $n_{\\rm max}=64$) leads to errors (e.g., $1.8\\sigma$ for \\textit{Planck}) in the predicted CMB power spectrum. Extrapolating errors, the resulting CMB anisotropy spectra are converged to $\\sim0.5\\sigma$ at Fisher-matrix level for $n_{\\rm max}=128$, in the purely radiative case. ",
        "introduction": "Measurements of cosmic microwave background (CMB) temperature anisotropies by the Wilkinson Microwave Anisotropy Probe (WMAP) have ushered in the era of precision cosmology, confirming that the Universe is spatially flat, with a matter budget dominated by dark matter and a baryonic mass fraction $\\Omega_{b}h^{2}$ \\cite{wmap_5} in agreement with the measured ratio of deuterium-hydrogen abundances (D/H) \\cite{steigmanreview}. WMAP measurements of large-scale CMB polarization also yield the optical depth $\\tau$ to the surface of last scattering (SLS), meaningfully constraining cosmological reionization. Together with surveys of supernovae \\cite{perlmutter_supernovae,reiss_supernovae}, galaxies \\cite{sdss_pspec,sdss_lrg,bao_discovery_cole}, and galaxy clusters \\cite{cluster_cosmo}, WMAP measurements build the case that the Universe's expansion is accelerating, due to ``dark energy\" or modifications of general relativity \\cite{caldwell_darkenergy,carroll_1r}, and constrain other physical parameters (such as the sum of neutrino masses $\\sum_{i}m_{\\nu_{i}}$ \\cite{neutrino,neutrinob,mabert} and the effective number of massless neutrino species $N_{\\nu}$).  CMB temperature observations (WMAP, BOOMERANG \\cite{boomerang_b}, CBI \\cite{cbi} and ACBAR \\cite{acbar_a}) probe properties of the primordial density field, such as the amplitude $A_{s}$, slope $n_{s}$, and running $\\alpha_{s}$ of its power spectrum. These observations constrain deviations from the adiabatic, nearly scale free and Gaussian spectrum of perturbations predicted by the simplest models of inflation, but also offer controversial hints of deviations from these models (see Refs. \\cite{wmap_5,acbar_c} and references therein).  Experimental upper limits to  B-mode polarization anisotropies (e.g. DASI \\cite{dasi_a} and BICEP \\cite{bicep}) impose constraints to the energy density of relic primordial gravitational waves \\cite{seljak_pol,kamion_pol}.  The \\textit{Planck} satellite, launched in May 2009, will obtain extremely precise measurements of the CMB temperature anisotropy power spectrum ($C_{\\ell}^{\\rm TT}$) up to $\\ell\\sim 2500$ and the E-mode polarization anisotropy power spectrum ($C_{\\ell}^{\\rm EE}$) up to $\\ell\\sim 1500$ \\cite{planck}. Robust measurements of the acoustic horizon and distance to the SLS will break degeneracies in dark energy surveys \\cite{planck,detf_b,bao_discovery_eisenstein,bao_discovery_cole}. Polarization measurements will yield the optical depth $\\tau$ to the SLS \\cite{planck}, further constraining models of reionization and breaking the degeneracy between $n_{s}$ and $\\tau$ \\cite{planck}. Cosmological parameters will be determined with much greater precision. More precise values of $n_{s}$ and $\\alpha_{s}$ will be obtained from CMB data alone, helping to robustly constrain inflationary models and alternatives to inflation \\cite{planck}. The advent of \\textit{Planck}, ongoing (SPT \\cite{spt} and ACT \\cite{act}) experiments at small scales, and a future space based polarization experiment like CMBPol \\cite{cmbpol_c,cmbpol_d} all require predictions of primary anisotropy multipole moments $C_{\\ell}$ with ${\\cal O}(10^{-3})$ accuracy.  During atomic hydrogen (H) recombination, the Thomson scattering opacity drops, decoupling the baryon-photon plasma and freezing in acoustic oscillations. The phases of acoustic modes are set by the peak location of the visibility function \\cite{bond_acoustic,peebles_acoustic}, while damping scales \\cite{silk_damp,hu_sugiyama_silk_acoustic} and the amplitude of polarization \\cite{bond_pol,polnarev_pol_a} are set by its width. Small-scale CMB anisotropies are also smeared out by free electrons along the line of sight, suppressing power on small scales so that $C_{\\ell}\\to C_{\\ell}e^{-2\\tau}$, where $\\tau$ is the total optical depth of this w \\cite{scottsmear}. An accurate prediction of the time-dependent free-electron fraction $x_{e}(z)$ from cosmological recombination is thus essential to accurately predict CMB anisotropies.  Recent work has highlighted corrections of $\\Delta x_{e}(z) /x_{e}(z)\\gsim 0.1\\%$ to the standard recombination history computed by \\textsc{RecFast} \\cite{seager_recfast}. These corrections will propagate through to predictions of anisotropies, and neglecting them would lead to biases and errors in \\textit{Planck} measurements of cosmological parameters \\cite{lewis_rec_cmb_param,wong_better_rec_b}. The use of the CMB as a probe of the first ionizing sources and of physics at energy scales greater than $10^{16}~{\\rm GeV}$ thus requires an accurate treatment of  the $\\sim {\\rm eV}$ atomic physics of recombination \\cite{improving_wong}.  Direct recombination to the hydrogen ground state is ineffective because of the high optical depth to photoionization \\cite{peebles_rec,zeldovich_rec}. Recombination proceeds indirectly, first through recombination to a $n\\geq2$ state of H, and then by cascades to the ground state. Because of the optical thickness of the Lyman-$n$ (Ly$n$) lines, the resulting radiation may be immediately absorbed, exciting atoms into easily ionized states.  There are two ways around this bottleneck \\cite{peebles_rec,zeldovich_rec}.  In the first, the sequence of decays from excited H levels ends with a two-photon decay (usually $2s\\to 1s$). The emitted photons may have a continuous range of energies, allowing escape off resonance and a net recombination. In the second, photons emitted in the $np\\to 1s$ transition redshift off resonance due to cosmological expansion, preventing re-excitation and yielding some net recombination. The dominant escape channel is from the $2p\\to 1s$ Lyman-$\\alpha$ line. These resonant transitions give off line radiation and distort the CMB \\cite{rybicki_specdist_a,dubrovich_a}.  Peebles, Sunyaev, Kurt, and Zel'dovich modeled recombination assuming that all net recombination resulted from escaping the $n=2$ bottleneck \\cite{peebles_rec,zeldovich_rec}. This three-level-atom (TLA) treatment included recombinations to excited states, under the assumption of equilibrium between energy levels $n$ and angular momentum sublevels $l$ for all $n\\geq 2$ (note the use of $l$ for atomic angular momentum and $\\ell$ for CMB multipole number). This sufficed until the multi-level-atom (MLA) model of Seager et al. \\cite{seager_phys_recfast}, which included hydrogen (H) and helium (He), separately evolved excited states assuming equilibrium between different $l$, accurately tracked the matter/radiation temperatures $T_{\\rm M}$/$T_{\\rm R}$ \\cite{weymann_tmtr,sunyaev_tmtr}, accounted for line emission using the Sobolev approximation \\cite{sobolev_sobolev_a}, and included ${\\rm H}_{2}$ chemistry. This treatment underlies the \\textsc{RecFast} module used by most CMB anisotropy codes, including those used for WMAP data analysis \\cite{seager_recfast}.  The higher precision of \\textit{Planck} requires new physical effects to be considered, among them two-photon transitions from higher excited states in H and He \\cite{kholupenko_twophoton,chluba_twophoton_a,chluba_twophoton_b, hirata_twophoton,chluba_lymana_b}, other forbidden and semiforbidden transitions in He \\cite{wong_scott_forbidden,hirata_helium_b,helium_forbidden}, feedback from Ly$n$ lines \\cite{chluba_lyman_feedback}, and corrections to the Sobolev approximation due to a host of radiative transfer effects in H and He resonance lines \\cite{kholupenko_helium_a,sunyaev_helium_a,chluba_lymana_a,chluba_lymana_b}. Most recent work on recombination has focused in one way or another on the radiative transfer problem. Here we direct our attention to the populations of very high-$n$ states.  One important effect is the breakdown of statistical equilibrium between states with the same value of the principal number $n$ but different angular momenta $l$. This effect is dramatic at late times. When $l$ sublevels of a level $n$ are resolved, increases in $x_{e}(z)$ of $\\sim 1\\%$ at late times result \\cite{chluba_highn_a,chluba_highn_b}. This changes predicted $C_{\\ell}$'s at a statistically significant level for \\textit{Planck}. Highly excited states in hydrogen also change the recombination history at a level significant for \\textit{Planck}. While levels as high as $n=300$ were included in the treatment of Ref.~\\cite{seager_phys_recfast} underlying \\textsc{RecFast}, $l$ sublevels were not resolved. It is thus important to update cosmological recombination histories to include high-n states of H while resolving $l$ sublevels, in order to predict the $C_{\\ell}$'s as well as CMB spectral distortions from recombination.  Simultaneously including very high $n$ and resolving the $l$ sublevels is computationally expensive, taking nearly a week on a standard workstation for $n_{\\rm max}=100$ \\cite{chluba_highn_b}, using a conventional multilevel-atom recombination code. This becomes prohibitively expensive for higher values of $n_{\\rm max}$, unless considerable resources are devoted to the problem. To date, this has prevented a determination of how $x_{e}(z)$ converges with $n_{\\rm max}$ and how high $n_{\\rm max}$ must be to predict $C_{\\ell}$'s for \\textit{Planck}. The existence of electric dipole selection rules $\\Delta l=\\pm 1$ means the relevant rate matrices are sparse, and we have used this fact to develop a fast code, \\textsc{RecSparse}, to explore convergence with $n_{\\rm max}$. While the computation time $t_{\\rm comp}$ for standard $l$-resolving recombination codes scales as $t_{\\rm comp} \\propto n_{\\rm max}^{6}$, with \\textsc{RecSparse} the scaling is $t_{\\rm comp}\\propto n_{\\rm max}^{\\alpha}$, where $2<\\alpha<3$. With \\textsc{RecSparse}, we can calculate recombination histories for $n_{\\rm max}=200$ in $~4$ days on a standard work-station; this would likely take weeks using a conventional code. For the first time, we have calculated recombination histories for $n_{\\rm max}$ as high as $250$ \\textit{with $l$ sublevels resolved}.  While previous computations have included some forbidden transitions, none have included optically thick electric quadrupole (E2) transitions in atomic hydrogen. We include E2 transitions, and find that the resulting correction to CMB anisotropies is negligible.  We find that the correction to CMB $C_{\\ell}$'s due to extremely excited levels is $0.5\\sigma$ or less if $n_{\\rm max}\\geq 128$, in the purely radiative case. This paper is not the final word on recombination; atomic collisions must be properly included and the effect of levels with $n>n_{\\rm max}$ must be included to conclusively demonstrate absolute convergence. The end goal of the present recombination research program is to include all important effects in a replacement for \\textsc{RecFast}, as the interplay of different effects is subtle.  In Sec. \\ref{mla}, we review the formalism of the multilevel atom (MLA), and follow by explaining how we extend the MLA to include very high-n states (Sec. \\ref{highnsec}) and electric quadrupole transitions (Sec. \\ref{quadtheorysec}). State populations, recombination histories, and effects on the $C_{\\ell}$'s are presented in Sec. \\ref{res_highn}. We conclude in Sec. \\ref{concl}.  We use the same fiducial cosmology as in Ref.~\\cite{hirata_lymana}: total matter density parameter $\\Omega_{m}h^{2}=0.13$, $\\Omega_{b}h^{2}=0.022$, $T_{\\rm CMB}=2.728~{\\rm K}$, $N_{\\nu}=3.04$, and helium mass fraction $Y_{\\rm He}=0.24$. ",
        "conclusions": "\\label{concl} We have developed a new recombination code, \\textsc{RecSparse}, optimized for tracking the populations of many energy shells of the hydrogen atom while resolving angular momentum sublevels. The code runs more quickly than would be anticipated using simple scaling arguments, which would yield the the scaling $t_{\\rm comp}\\propto n_{\\rm max}^{6}$. Using \\textsc{RecSparse}, we find empirically that for the range of $n_{\\rm max}$ values used, computation time scales as $t_{\\rm comp}\\propto n_{\\rm max}^{\\alpha}$, where $2<\\alpha<3$. With this code, we have computed cosmological hydrogen recombination histories for a series of $n_{\\rm max}$ values going as high as $n_{\\rm max}=250$ and explored the highly nonequilibrium state of the resulting atomic hydrogen.  The resulting correction $\\Delta x_{e}(z)$ satisfies $\\Delta x_{e}(z)/x_{e}(z)<0.01$ for $z>200$ when $n_{\\rm max}=250$ and converges with $\\Delta x_{e}(z)/x_{e}(z)\\propto n_{\\rm max}^{-1.9}$. The correction to the $C_{\\ell}$'s becomes of order the cosmic variance when $n_{\\rm max}=250$. In light of realistic error estimates for \\textit{Planck}, the resulting CMB anisotropy spectra $C_{\\ell}^{XX}$ are converged to $0.5\\sigma$ at Fisher-matrix level for $n_{\\rm max}=128$ in the purely radiative case, assuming error extrapolations may be trusted.  To definitively answer the question of absolute convergence, collisions must be included to speed the approach to Saha equilibrium at high $n$, allowing a conclusive treatment of states beyond the truncation limit, with $n>n_{\\rm max}$. Future work should also properly account for the overlap of the Lyman resonance line series at high $n$. It will also be interesting to determine if there is coherent stimulated emission between excited states, given its relevance for the detectability of faint CMB spectral distortions from the epoch of recombination. Finally, the sparse-matrix methods applied here or similar techniques could be profitably applied in the development of fast recombination codes for CMB data analysis, even at early times in recombination, when only lower values of $n_{\\rm max}$ are relevant."
    },
    "0911/0911.1997_arXiv.txt": {
        "abstract": "Inelastic dark matter reconciles the DAMA anomaly with other null direct detection experiments and points to a non-minimal structure in the dark matter sector. In addition to the dominant inelastic interaction, dark matter scattering may have a subdominant elastic component.  If these elastic interactions are suppressed at low momentum transfer, they will have similar nuclear recoil spectra to inelastic scattering events.  While upcoming direct detection experiments will see strong signals from such models, they may not be able to unambiguously determine the presence of the subdominant elastic scattering from the recoil spectra alone.  We show that directional detection experiments can separate elastic and inelastic scattering events and discover the underlying dynamics of dark matter models. ",
        "introduction": "The annual modulation anomaly from the DAMA experiment is an intriguing hint of the identity of dark matter \\cite{Bernabei:2005hj,Bernabei:2008yi}.  Direct detection experiments such as DAMA measure the energy of nuclear recoils from incident dark matter particles.  The signal experiences an annual modulation because of the variation of the Earth's relative motion with respect to the dark matter in the halo.  While the DAMA experiment has measured an $8.2\\sigma$ modulation with the correct phase, its results are in conflict with those from  CDMS  \\cite{Akerib:2005za, Akerib:2005kh, Ahmed:2008eu}, XENON10 \\cite{XENON10, Collaboration:2009xb}, CRESST \\cite{Angloher:2004tr, CRESSTII}, ZEPLIN-II \\cite{Alner:2007ja}, and ZEPLIN-III \\cite{ZEPLINIII}, which measure the total, unmodulated scattering rate.  All current direct detection experiments are optimized to look for elastic scattering, which has an exponentially falling recoil energy spectrum \\cite{Lewin:1995rx}.  A distinguishing feature of DAMA's measured modulation spectrum is that it is suppressed at energies below $\\sim$ 25 keVnr, where elastic events should dominate.  DAMA's results, taken together with the null results from all other direct detection experiments, can be  elegantly explained if the dark matter scatters inelastically off of nuclei to an adjacent state that is $\\OO(100 \\keV)$ more massive than the ground state \\cite{TuckerSmith:2001hy,Chang:2008gd, Cui:2009xq,SchmidtHoberg:2009}.  In the case of inelastic dark matter (iDM), a minimum velocity is needed to upscatter and the recoil spectrum is suppressed below this kinematic threshold.  Inelastic dark matter has three features that lead to the consistency of all current experiments.  First, iDM has a much larger dependence on the Earth's velocity than elastic dark matter and thus the annual modulation is often 25\\% or larger, in comparison to $2.5\\%$ for elastic scattering.  This decreases DAMA's unmodulated cross section by a factor of 10 or more and therefore reduces the expected signal by the same factor in all experiments looking for unmodulated scattering. Second, iDM dominantly scatters off heavier nuclei and lighter nuclei may not have enough kinetic energy to excite the transition.  As a result, CDMS' sensitivity is strongly suppressed due to inelastic kinematics.  Lastly, inelastic scattering events have higher recoil energies than elastic events.  Recent experiments such as XENON10 and ZEPLIN-III have shrunk their recoil energy window to eliminate higher energy scattering events, reducing the acceptance for iDM. However, XENON10 has recently reanalyzed their data over a larger signal window to be sensitive to iDM and their latest results are used in this article \\cite{Collaboration:2009xb}.   A dynamical alternative to the inelastic hypothesis was recently discussed in \\cite{Chang:2009yt,Feldstein:2009tr}; in this scenario, the interaction between the dark matter and Standard Model (SM) is elastic, but is suppressed by a form factor \\begin{equation} F_{\\dm}(q^2) = c_0 + c_1 q^2 + c_2 q^4 + \\cdots. \\end{equation} If the first term in this expansion vanishes, the elastic scattering rate is multiplied by additional factors of $q^2 = 2 m_N E_R$ and goes to zero as $E_R \\rightarrow 0$.  Form factor-suppressed scattering can arise from multipole, polarization, and charge-radius interactions between the dark sector and the Standard Model \\cite{Pospelov:2000bq}.  The recoil spectrum for form factor elastic dark matter (FFeDM) is suppressed at low energies and more closely resembles the spectrum for inelastic events rather than standard elastic events, which peaks at low energies.  However, it is more challenging to reconcile DAMA and the null experiments using only FFeDM because the form factor depends on the product $m_N E_R$ and any suppression for light nuclei can be compensated by looking at larger energies.  In this paper, we consider a scenario where inelastic scattering is complemented by a form factor-suppressed elastic component. These scenarios arise, for example, in composite dark matter models \\cite{Lisanti:2009} and yield new challenges for direct detection phenomenology. As is illustrated in Fig.~\\ref{Fig: AnnMod}, the modulation amplitude for inelastic scattering with subdominant charge-radius scattering (dashed line) resembles that for inelastic scattering (solid line) and is markedly different from the regular elastic spectrum.  Distinguishing the elastic and inelastic contributions is critical for understanding the underlying structure of the theory.  This article addresses how directional detection experiments, which can measure the direction of the nuclear recoil in addition to its energy \\cite{Sciolla:2008vp,Finkbeiner:2009ug}, can differentiate the contributions to the scattering rate.  In the following section, we review the direct detection phenomenology.  In Sect.~\\ref{Sec: Directional Detection}, we introduce directional detection experiments and show how they can distinguish the dynamics in the dark sector.  We conclude in Sec.~\\ref{Sec: Discussion}. \\begin{figure}[b] % \\centering \\includegraphics[width=3.4in]{DAMAdata.pdf} \\caption{Comparison of the modulation amplitude for two scattering scenarios: completely inelastic scattering (solid) and inelastic scattering with a subdominant elastic charge-radius component (dashed).  The points show the modulation amplitude measured by DAMA and DAMA/LIBRA \\cite{Bernabei:2005hj}.  The electron equivalent energy (keVee) is the recoil energy rescaled by the quenching factor for iodine ($q_I = 0.085$). } \\label{Fig: AnnMod} \\end{figure} ",
        "conclusions": "\\label{Sec: Discussion}  We have shown that directional detection experiments can distinguish mixed inelastic-elastic scattering scenarios that would otherwise be difficult to discover unambiguously, even with next-generation experiments such as LUX and XENON100.  In particular, directional detection experiments can detect elastic wide angle scatters that are kinematically forbidden in inelastic transitions. To evade current null experiments, a relatively large exposure of several $\\text{kg-yrs}$ is needed, but this is well within reach of current detectors.  However, these experiments are currently optimized to look for spin-dependent elastic scattering and need to use heavier nuclei, such as iodine or xenon, to have sensitivity to inelastic recoils.  Distinguishing different scattering mechanisms is critical for understanding the underlying symmetry structure of dark matter theories.  Many types of inelastic models can give small elastic scattering contributions.  For instance, nontrivial scattering mechanisms arise in models where the dark matter has a finite size, such as composite \\cite{Alves:2009nf,Nussinov:1985xr, Khlopov:2008ki, Ryttov:2008xe}, atomic \\cite{Kaplan:2009de}, mirror  \\cite{Mohapatra:2001sx}, and quirky \\cite{Kribs:2009fy} dark matter.  All these models have several scattering channels with a hierarchy of scales given by a series of higher dimensional operators.  Both elastic and inelastic operators may be allowed, and approximate discrete symmetries may cause the elastic scattering rate to be subdominant to the inelastic rate \\cite{Lisanti:2009}.  Another class of models that leads to form factor-suppressed elastic scattering is characterized by a pseudo-Goldstone mediator between the dark sector and the SM \\cite{Chang:2009yt}.  In this case, additional dimension six operators in the Lagrangian may no longer be negligible in comparison to the standard spin-independent and spin-dependent operators.  These higher dimension operators lead to momentum suppression in the scattering rate.  There has been much interest recently in models with light mediators \\cite{ArkaniHamed:2008qn, Morrissey:2009ur, Baumgart:2009tn, Cholis:2008qq, Finkbeiner:2009mi, Cheung:2009qd}, given the results from PAMELA \\cite{Adriani:2008zr}, Fermi \\cite{Abdo:2009zk}, ATIC \\cite{:2008zzr}, and HESS \\cite{Aharonian:2009ah}.  If the dark sector does indeed have non-minimal structure, it may give rise to several types of scattering mechanisms.  Distinguishing these different scattering events is of fundamental importance for understanding the symmetry structure and unraveling the dynamics of the dark sector."
    },
    "0911/0911.1503_arXiv.txt": {
        "abstract": "The gamma-ray burst (GRB) 080503 was classified as a short GRB with extended emission (Perley et al. 2009). The origin of such extended emission (found in about a quarter of Swift short GRBs) is still unclear and may provide some clues to the identity of the elusive progenitors of short GRBs. The extended emission from GRB\\,080503 is followed by a rapid decay phase (RDP) that is detected over an unusually large dynamical range (one decade in time and $\\sim 3.5$ decades in flux), making it ideal for studying the nature of the extended emission from short GRBs. We model the broad envelope of extended emission and the subsequent RDP using a physical model for the prompt GRB emission and its high latitude emission tail (Genet \\& Granot 2009), in which the prompt emission (and its tail) is the sum of its individual pulses (and their tails). For GRB\\,080503, a single pulse fit is found to be unacceptable, even when ignoring short timescale variability.  The RDP displays very strong spectral evolution and shows some evidence for the presence of two spectral components with different temporal behaviour, likely arising from distinct physical regions. A two pulse fit (a first pulse accounting for the gamma-ray extended emission and decay phase, and the second pulse accounting mostly for the X-ray decay phase) provides a much better (though not perfect) fit to the data. The shallow gamma-ray and steep hard X-ray decays are hard to account for simultaneously, and require the second pulse to deviate from the simplest version of the model we use. Therefore, while high latitude emission is a viable explanation for the RDP in GRB\\,080503, it does not pass our tests with flying colors, and it is quite plausible that another mechanism is at work here. Finally, we note that the properties of the RDP following the extended emission of short GRBs (keeping in mind the very small number of well studied cases so far) appear to have different properties than that following the prompt emission of long GRBs. However, a larger sample of short GRBs with extended emission is required before any strong conclusion can be drawn. ",
        "introduction": "GRB080503 was detected by the {\\it Swift} Burst Alert Telescope (BAT) on 2008 may 3 \\citep{peretal09}. Its prompt gamma-ray emission presents a short ($\\sim 0.32\\;$s) intense initial spike followed by an extended emission lasting several minutes, for a total duration of $T_{\\rm 90} \\approx 232\\;$s. The first short spike duration in the $15-150\\;$keV band is $0.32 \\pm 0.07\\;$s and its peak flux is $(1.2\\pm0.2)\\times 10^{-7}\\;$erg$\\;$cm$^{-2}\\;$s$^{-1}$. The count rate hardness ratio between the $50-100\\;$keV and the $25-50\\;$keV of the initial spike is $1.2\\pm0.3$, consistent with other short {\\it Swift} bursts, but also with some long bursts. The fluence of the extended emission (measured between $5\\;$s and $140\\;$s) in the $15-150\\;$keV band is $(1.86\\pm0.14)\\times 10^{-6}\\;$erg$\\;$cm$^{-2}$, about thirty times that of the initial spike, higher than for any other short {\\it Swift} bursts, but within the range of such ratio measured for BATSE short bursts. \\citet{peretal09} found a spectral lag between the $50-100\\;$keV and the $25-50\\;$keV bands consistent with zero. All this led them to associate GRB080503 with the ``short'' \\citep{kouvetal93} class.  The {\\it Swift} X-Ray Telescope (XRT) started observing GRB080503 about $82\\;$s after the burst, detecting a bright early afterglow that rapidly decayed ($\\alpha=2 - 4$ with $F_{\\nu} \\propto t^{-\\alpha}$) to below the detection threshold during the first orbit, which makes a record overall decline of $\\sim 6.5$ decades, with a steep decay clearly observed for $\\sim 3.5$ decades (see figure 6 of Perley et al. 2009). Chandra detections at $\\sim 3\\;$days after the GRB indicate the presence of a separate X-ray afterglow component $\\sim 10^6$ times fainter than the peak emission from the X-ray tail of the prompt extended emission. The gamma-ray lightcurve of the extended emission (which can be found with the X-ray light-curve of the extended emission in figure \\ref{fig_fits}) presents two relatively well defined spikes at $\\sim 28\\;$s and $\\sim 36\\;$s, followed by some less defined variability with several narrow spikes (see also figure 1 of Perley et al. 2009). The X-ray RDP lightcurves present a first break at $\\sim 180\\;$s seen in both soft and hard bands; later on, the hard X-ray band clearly shows a steepening (to a temporal slope of about $-4.2$) that is only marginally seen in the soft X-rays. The simultaneous gamma-ray decay is much shallower, with a temporal slope only about $-2$.  The rapid decay phase (RDP) of GRB080503 is a smooth temporal and spectral continuation of the extended emission in the gamma-rays (Perley et al., 2009). This suggests that it is the tail of the prompt emission (O'Brien et al., 2006; Butler \\& Kocevski, 2007). The most popular model to explain the RDP following the prompt emission of long GRBs is High Latitude Emission (HLE; Kumar and Panaitescu 2000). In this model, after the prompt emission stops, photons from increasingly larger angles relative to the line of sight still reach the observer, due to the curvature of the (assumed to be quasi-spherical) emitting surface. The blueshift of such photons decreases with the angle from the line of sight, and therefore with their arrival time. This results in a simple relation between the temporal and spectral indexes of the flux, $\\alpha = 2+ \\beta$, where $F_{\\nu} \\propto t^{-\\alpha} \\nu^{-\\beta}$, that holds at late times (when $t-t_0 \\gg \\Delta t$, where $t_0$ is the onset time of the pulse and $\\Delta t$ its width) for each pulse of the prompt emission.  The RDP following the extended emission of the short GRB\\,080503, however, shows strong spectral evolution that is different than that typically seen in the RDP following the prompt emission of long GRBs. Whereas in the latter case the RDP usually shows a monotonic softening of the spectrum with time (see Zhang, Liang \\& Zhang 2007 and references therein), the RDP in GRB\\,080503 shows a similar softening in the XRT energy range together with a hardening in the spectral slope between the XRT and BAT energy ranges. This behavior suggests the presence of two distinct spectral components.  Such extended emission is observed in about a quarter of Swift short bursts (Norris and Gehrels 2009; Norris et al. 2009). It usually lasts tens of seconds, and has been shown to be always softer than the initial short spike, however showing negligible lags as the initial spike \\citep{nobo06}. The same authors also found that the extended emission is softer than the initial spike, making it closer to long bursts, which are softer than short bursts. However, the initial spike of short burst with and without EE show similar properties \\citep{noge09}. Many authors recently turned to environmental constraints to explain differences between short GRBs with and without EE, but with results at best uncertain, sometime contradictory (Rhoads 2008; Troja et al. 2008; Fong, Berger \\& Fox 2009; Sakamoto \\& Gehrels 2009; Nysewander, Fruchter \\& Pe'er 2009). There is as of now no strong explanation for the mechanism behind extended emission.  It is therefore very important to test whether HLE can indeed explain the RDP of GRB080503, as this would allow some comparison with results for the RDP following the prompt emission without extended emission (Willingale et al., 2009). This will stress similarities and differences between the extended emission and the prompt emission itself, ultimately leading to a better understanding of the origin of extended emission.  Genet \\& Granot (2009; hereafter GG09) have developed a simple yet physical and self-consistent model for the prompt and high latitude emission. The very large dynamical range over which the RDP of GRB080503 is observed allows us to test its behaviour at late times, which is rarely possible in other {\\it Swift} bursts. In order to test whether the extended emission is coming from a mechanism similar to the prompt emission itself, we apply the GG09 model to GRB080503 data by folding it through the response matrices generated from the four energy bands of data (0.3-1.3 keV and 1.3-10.0 keV for the XRT, 15-50 keV and 50-150 keV for the BAT). Section \\S$\\;$\\ref{sec_modeldescription} summarizes the theoretical model used to fit the burst lightcurves. Section \\S$\\;$\\ref{sec_datafit} describes how the data are reduced and fitted and the results of our fits, and in section \\S$\\;$\\ref{sec_discussion} we discuss the implications of these results. ",
        "conclusions": ""
    },
    "0911/0911.1029_arXiv.txt": {
        "abstract": "Active galactic nuclei (AGN) found at the centers of clusters of galaxies are a possible source for weak cluster-wide magnetic fields. To evaluate this scenario, we present 3D adaptive mesh refinement MHD simulations of a cool-core cluster that include injection of kinetic, thermal, and magnetic energy via an AGN-powered jet. Using the MHD solver in FLASH 2, we compare several sub-resolution approaches that link the estimated accretion rate as measured on the simulation mesh to the accretion rate onto the central black hole and the resulting feedback. We examine the effects of magnetized outflows on the accretion history of the black hole and discuss the ability of these models to magnetize the cluster medium. ",
        "introduction": "Faraday rotation measures of galaxy clusters suggest they contain magnetic fields with average strengths of $\\sim 1$~$\\mu$G (e.g., \\citep{Feretti}), but the origin of these fields is unknown. Possible seed fields, such as those generated by the Biermann battery mechanism, require amplification \\citep{Widrow}. Accretion disks around active galactic nuclei (AGN) are a promising candidate amplification source, since AGN-launched jets require magnetic fields and the energetics of AGN-blown bubbles suggest that AGN can amplify the seed fields to the observed strengths.  We would like to directly test this scenario in simulation, but since AGN accretion disks are typically smaller then $\\sim 100$ AU, and cosmological simulations have resolutions on the order of $\\sim 1$ kpc, we must include the amplification and jet launching as a subgrid model. We investigate two commonly-used AGN feedback models: one based on jets \\citep{Cattaneo}, the other on already-inflated bubbles \\cite{Sijacki}. We link the feedback energies of these models to the strength of an injected magnetic field. We inject magnetic fields with the form described by \\cite{Li}, in which the field has both toroidal and poloidal components. ",
        "conclusions": ""
    },
    "0911/0911.3506_arXiv.txt": {
        "abstract": "We model a coronal loop as a bundle of seven separate strands or filaments. Each of the loop strands used in this model can independently be heated (near their left footpoints) by Alfv\\'en/ion-cyclotron waves via wave-particle interactions. The Alfv\\'en waves are assumed to penetrate the strands from their footpoints, at which we consider different wave energy inputs. As a result, the loop strands can have different heating profiles, and the differential heating can lead to a varying cross-field temperature in the total coronal loop. The simulation of TRACE observations by means of this loop model implies two uniform temperatures along the loop length, one inferred from the 171:195 filter ratio and the other from the 171:284 ratio. The reproduced flat temperature profiles are consistent with those inferred from the observed EUV coronal loops. According to our model, the flat temperature profile is a consequence of the coronal loop consisting of filaments, which have different temperatures but almost similar emission measures in the cross-field direction. Furthermore, when we assume certain errors in the simulated loop emissions (e.g., due to photometric uncertainties in the TRACE filters) and use the triple-filter analysis, our simulated loop conditions become consistent with those of an isothermal plasma. This implies that the use of TRACE/EIT triple filters for observation of a warm coronal loop may not help in determining whether the cross-field isothermal assumption is satisfied or not. ",
        "introduction": "In our recent papers \\citet{Bourouaine2008a,Bourouaine2008b} (hereafter referred to as BR2008a; BR2008b) we modelled a coronal loop as one single flux tube (or monolithic loop) in which the confined plasma is heated by ion-cyclotron/Alfv\\'en waves via the resonant wave absorption mechanism. When considering Coulomb collisions in that model, it turned out that they are not strong enough to maintain local thermal equilibrium, and it was found that the electron temperature profile is different from the ion profile. Furthermore, it was shown that the warm coronal loop (e.g., $T\\leq 1.5$ MK), which has a quasi-uniform electron temperature profile and enhanced density \\citep[relative to the static model by][]{Serio1981}, is the consequence of a uniform cross-section of the flux tube forming the loop. Modelling of this type of loop profile with a uniform temperature was motivated by the many observations of EUV (Extreme Ultraviolet) loop emissions (using the high spatial resolution imaging provided by TRACE \\citep{Lenz1999,Aschwanden1999,Aschwanden2000}.  Under the assumption that the loops have a single cross-field temperature, techniques based on filter-ratio analysis have been employed to infer the electron temperature along the EUV coronal loops \\citep{Noglik2007,Noglik2008,Aschwanden2008,Schmelz2009}. Most of the obtained results indicate that the warm loops have roughly constant temperatures along much of their lengths. Furthermore, it was found that their densities are roughly uniform and much larger than those predicted by static models with uniform heating \\citep{Aschwanden2001,Winebarger2003}. However, it was concluded that static loop models with footpoint heating may produce flat temperature profiles and enhanced apex densities, but in the case of long loops the theoretical enhancement in density was found to be not large enough as compared to the observed values.  It is believed that the difficulty of reproducing the observed properties of TRACE loops with static models lies in that they are not in thermal equilibrium. Therefore, it was suggested that impulsive heating caused by nanoflares could explain the observed loop characteristics. However, in such process, numerical simulations showed that the single-loop cooling would appear simultaneously in both the TRACE 171~{\\AA} and 195~{\\AA} filters for a few seconds \\citep{Reeves2002}. This cooling time is much shorter than the lifetime of the active-region loops observed by TRACE, which generally live for several hours while maintaining their high densities and flat temperature profiles.  \\begin{figure}[t] \\includegraphics[width=14cm,height=5.5cm]{f1.eps} \\caption{Coronal model loop that consists of a bundle of many magnetic strands} \\label{fig.1} \\end{figure}  Therefore, it has been suggested that the relevance of the impulsive heating heavily relies on the assumption of a multi-strand or multi-thread nature of coronal loops, as it is illustrated in Fig.~\\ref{fig.1}, and these strands are at different stages of heating and cooling exhibiting a broad differential emission measure (DEM) distribution.  The concept of multi-stranded loops is supported by many observations which exploit the high spatial resolution of TRACE and indicate that an observed \"fat\" single loop could be composed of many small-scale filaments having most probably different temperatures across the magnetic field \\citep{Schmelz2001,Schmelz2003,Schmelz2005,Martens2002}. Hydrodynamic models based on multi-stranded loops have been suggested and used to calculate numerically the flare emission \\citep{Warren2006, Reeves2007}. Other multi-strand coronal loop models involving nanoflare heating have been proposed to simulate EUV coronal loops \\citep[e.g.,][]{Cargill1997, Warren2002, Sarkar2009}.  The lack of observations of very fine-scale structures at scales below the TRACE resolution still permits the possibility that the observed loops may be composed of very thin \"strands or filaments\" which have different temperature and density profiles. As a result, the observed loop intensity could be caused by an ensemble of emissions from many unresolved strands.  \\citet{Reale2000} produced the observed features of warm coronal loops by assuming them to consist of a bundle of static, uniformly heated strands. Their model, however, implements dense and hot strand ($\\sim5$~MK) to reproduce the flat TRACE filter ratios. This seems to contradict previous observations which indicate that, generally, hot loops are not co-spatial with relatively cool loops \\citep{Sheeley1980,Habbal1985, Brooks2007}.  Following the concept of a multi-strand structure of coronal loops, we will also model the observed TRACE loop as a bundle of seven isolated strands. However, concerning loop heating we propose here that the strands are heated through the dissipation of parallel propagating ion-cyclotron/Alfv\\'{e}n waves via resonant wave absorption. The same kinetic model that has been used in the papers BR2008a and BR2008b to describe the heating of a monolithic loop is now applied to the heating of an individual small-scale strand. Thus various fine structures can have different temperatures and density profiles. Therefore, an observation of a neighbouring set of filaments with a low-resolution instrument could lead to an apparently fat loop, which yet is simply composed of a finite number of strands. Since in the corona the cross-field heat transport is very weak, and as the plasma is dominated by the magnetic pressure, we can assume that separate loop strands are thermally isolated.  Once we have determined the different plasma profiles of single loop strands, we can synthesize the total emission of the composite model loop, thereby assuming its observation as a monolithic loop in the TRACE/EIT filters, and thus we can extract its temperature by adopting the filter-ratio technique. As a result we find that, if the coronal loop is composed of unresolved filaments having different temperatures and roughly identical emission measures across the coronal loop, a quasi-uniform temperature is inferred along the loop length.  The paper is organized as follows: in Sec.~2 we model a TRACE/EIT coronal loop as a bundle of seven separate filaments heated through the dissipation of high-frequency Alfv\\'en waves. We assume different wave energy inputs at the footpoints of the different strand. Then in Sec.~3, we synthesize the emission of the modelled coronal loop in the three TRACE passbands, and thus we can derive the loop temperature from two filter ratios 171:195 and 171:284. Finally, we discuss the obtained results and conclude in Sec.~4. ",
        "conclusions": "In this work we synthesized the emission of seven strands which together are assumed to constitute a coronal loop as seen in a low-resolution TRACE or SOHO/EIT observation. Each of the loop strands used in the model is heated by Alfv\\'en/ion-cyclotron waves via wave-particle interactions. This process leads to proton heating in the electron-proton plasma, and due to electron-proton collisions the electrons can then be heated as well up to a certain temperature which is below or equal to the proton temperature. It turns out that the plasma, in case of hot loop strands (where the proton temperature exceeds $>1$~MK) is not in local thermal equilibrium and the electrons are cooler than the protons.  In our model, the Alfv\\'en waves, which are assumed to penetrate the strands from their footpoints, are there generated with different intensities. Consequently, different heating profiles occur within each strand due to the wave absorption or heating process. Therefore, this differential heating leads to a varying cross-field temperature in the total coronal loop. Moreover, due to the ion acceleration process caused by wave-particle interactions, the model produces roughly similar electron density values (or emission measure since the strands have the same widths) across the loop.  The simulated TRACE observation of the modelled loop implies two different quasi-uniform temperature profiles along the loop length, one derived from the filter-ratio 171:195 and the other for the 171:284. This flatness behaviour in the temperature profiles is consistent with results based on the filter ratio analysis applied to TRACE/EIT coronal loop. According to our model results, such flat temperature profiles can occur when the loop consists of a small number of threads (forming a discrete loop like-structure) which have different temperature profiles but rough similar emission measure, $EM$, across the total loop.  It is worth noting that, due to the week binary collisions between the plasma species, the loop plasma can be far from local thermal equilibrium if the ion temperature exceeds one mega Kelvin. Therefore, it is ideally suited for applying the kinetic description of a plasma rather than using the single fluid approach to model the coronal loop heating.  Furthermore, most of the previous hydrodynamic models invoked arbitrary heating functions in the energy equation to model the loop heating. However, our model adopts a physical mechanism for loop heating. The considered heating process relies on wave-particle interactions described by the quasi-linear theory of the kinetic Vlasov equation. This process does not only lead to plasma heating (or ion heating) but also accelerates the ions in the loop strands due to the wave pressure, and thus, can fill these strands with hot plasma. This loop-filling mechanism, which is different from the evaporation process, is naturally involved in our model, and it turns to be important in order to reproduce the observed features of coronal EUV loops."
    },
    "0911/0911.5087.txt": {
        "abstract": "{ We report measurements of the thermal emission of the young and massive planet CoRoT-2b at 4.5 and 8 $\\mu$m with the  $Spitzer$ Infrared Array Camera (IRAC). Our measured occultation depths are $0.510\\pm0.042$ \\% and $0.41\\pm0.11$ \\% at 4.5 and 8 $\\mu$m, respectively. In addition to the CoRoT optical measurements, these planet/star flux ratios indicate a poor heat distribution to the night side of the planet and are in better agreement with an atmosphere free of temperature inversion layer. Still, the presence of such an inversion is not definitely ruled out by the observations and a larger wavelength coverage is required to remove the current ambiguity. Our global analysis of CoRoT, $Spitzer$ and ground-based data confirms the large mass and size of the planet with slightly revised values ($M_p=3.47\\pm0.22$~$M_J$, $R_p=1.466\\pm0.044$~$R_J$).  We find a small but significant offset in the timing of the occultation when compared to a purely circular orbital solution, leading to $e \\cos{\\omega} =-0.00291\\pm0.00063$ where $e$ is the orbital eccentricity and  $\\omega$ is the argument of periastron. Constraining the age of the system to be at most of a few hundreds of Myr and assuming that the non-zero orbital eccentricity is not due to a third undetected body, we model the coupled orbital-tidal evolution of the  system with various tidal $Q$ values, core sizes  and initial orbital parameters. For $Q_s' = 10^5 - 10^6$, our modelling is able to explain the large radius of CoRoT-2b if $Q_p' \\le 10^{5.5}$ through a transient tidal circularization and corresponding planet tidal heating event. Under this model, the planet will reach its Roche limit within 20 Myr at most. ",
        "introduction": "Transiting planets are key objects for our understanding of the atmospheric properties of exoplanets. Indeed, their special geometrical configuration gives us the opportunity not only to deduce their density but also to study directly their atmospheres without the challenging need to spatially resolve their light from that of their host star. In particular, their emergent flux can be directly measured during their occultation (secondary eclipse) when they are hidden by their host star, as was demonstrated by Charbonneau et al. (2005) and  Deming et al. (2005). Since 2005, many exoplanet occultation  measurements have been gathered, the bulk of them by  the $Spitzer$ $Space$ $Telescope$ (see, e.g., Deming 2009), the few others being due to the $Hubble$ $Space$ $Telescope$ (Swain et al. 2009), CoRoT (Alonso et al. 2009b, 2009c; Snellen et al. 2009a), Kepler (Borucki et al. 2009) and  ground-based telescopes (Sing \\& L\\'opez-Morales 2009, de Mooij \\& Snellen 2009, Gillon et al. 2009b). Combining the photometric measurements at different wavelengths allows us to map the  Spectral Energy Distribution (SED) of the planet and to constrain its chemical composition, its thermal distribution efficiency and the presence of a possible stratospheric thermal inversion (see, e.g., Charbonneau et al. 2008, Knutson et al. 2008). Such inversions have been detected for the highly irradiated planets HD 209458b (Burrows et al. 2007b, Knutson et al. 2008), TrES-2b (O'Donovan et al. 2009), TrES-4b (Knutson et al. 2009), XO-1b (Machalek et al 2008) and XO-2b (Machalek et al. 2009). These results are in rather good agreement with the theoretical division of hot Jupiters into two classes based on their level of irradiation (Hubeny et al. 2003, Burrows et al. 2007b, Harrington et al. 2007, Burrows et al. 2008, Fortney et al. 2008). Under this division, the planets warmer than required for condensation of high-opacity gaseous molecules like TiO/VO, tholins or polyacetylenes should show a stratospheric temperature inversion due to the absorption in their upper-atmosphere of a significant fraction of the large incident flux by these  compounds. The less irradiated planets would lack these gazeous compounds and the resulting temperature inversion. Still, this simple division was recently challenged by the absence of thermal inversion reported for the strongly irradiated planet TrES-3b by Fressin et al. (2009). This result indicates that, in addition to the  irradiation amplitude, other effects like chemical composition, surface gravity and the stellar spectrum have probably an impact on the temperature profile of hot Jupiters.  With an irradiation $\\sim1.3\\times10^9$ erg~s$^{-1}$~cm$^{-2}$, the planet CoRoT-2b could be expected to show such a temperature inversion, according to the theoretical division mentionned above. This planet is the second one discovered by the CoRoT transit survey mission (Alonso et al. 2008, hereafter A08). Spectral analysis and evolution modelling of the host star leads to a solar-type dwarf with a mass $M_\\ast=0.97\\pm0.06$~$M_\\odot$ and an effective temperature \\teff$=5625\\pm120$~K (Bouchy et al. 2008, hereafter B08).  A08 derived for the planet a radius of $1.465\\pm0.029$~$R_{J}$ and a mass of $3.31\\pm0.16$~$M_{J}$, leading to a density of $1.31\\pm0.04$~g cm$^{-3}$, very close to the value for Jupiter. This density is surprising because the radius of massive planets is expected to approach Jupiter's asymptotically. In this context, it is worth noticing the probably young age of the system. Indeed, the presence in the stellar spectrum of the Li I absorption line and the strong emission in the Ca II H and K line cores (B08) suggest that  the star is still close to the Zero-Age Main-Sequence (ZAMS) and is thus younger than 0.5 Gyr (B08), in full agreement with the short rotational period of $\\sim4.5$ days deduced from CoRoT photometry (A08). Still, the youngness of CoRoT-2b is not enough to prevent it from falling into the sub-group of planets with a radius larger than predicted by basic models of irradiated planets (Burrows et al. 2007a, Fortney et al. 2007). Most of these planets show an orbital eccentricity compatible with zero. Nevertheless, these planets could still have undergone during their evolution a tidal heating large enough to explain their low density (Jackson et al. 2008b; Ibgui \\& Burrows 2009), and it is thus important to measure very precisely their present eccentricity to constrain their tidal and thermal history. The precise measurement of a planet's occultation provides strong constraints on the orbital eccentricity, especially on the parameter $e \\cos{\\omega}$, where $e$ is the eccentricity and $\\omega$ the argument of periastron (see e.g. Charbonneau et al. 2005, Knutson et al. 2009). In the case of CoRoT-2b, the dynamical interest of such occultation observations is reinforced by the large jitter noise of the young host star (B08), which makes the precise determination of a tiny eccentricity very challenging for the radial velocity (RV) method alone.  With the goals of better characterizing the atmospheric properties of CoRoT-2b (SED, inversion) and improve our understanding of its low density (tidal heating), we observed the occultation of this planet at 4.5 and 8 $\\mu$m with $Spitzer$/IRAC (DDT program 486). A partial transit was also observed with VLT/FORS2 to bring one more constraint on the orbital parameters.  We report here the results of the analysis of these new data. Section 2 presents our IRAC and VLT observations and their reduction. We analyzed this new photometry in combination with CoRoT transit photometry and published RVs. This combined analysis is presented in Sec. 3. We present and discuss our results in Sec. 4 and give our conclusions in Sec. 5. ",
        "conclusions": "Using $Spitzer$ and its IRAC camera, we measured an occultation of the young and massive planet CoRoT-2b at 4.5 $\\mu$m and 8 $\\mu$m. In addition, we observed a partial transit of the planet with the Very Large Telescope and its FORS2 camera.  Our global analysis of CoRoT, $Spitzer$ and ground-based (FORS2 photometry + RVs) data confirms the low density of the planet ($\\rho_p=1.105\\pm0.060$~$\\rho_{J}$ with $M_p =3.47\\pm0.22$~$M_J$ and $R_p=1.466\\pm0.044$~$R_J$). Constraining the system to be of at most a few hundreds of Myrs  old and the present orbit to be slightly eccentric ($e \\cos{\\omega} = -0.00291\\pm0.00063$) and using coupled tidal-orbital evolution modelling, we find a self consistent thermal \\& tidal evolution history that may explain the radius through a transient tidal circularization \\& corresponding tidal heating inside the interior of the planet. Under this scenario, the planet will be tidally disrupted within 20 Myr at most.  The  occultation depths that we measured at 4.5 $\\mu$m and 8 $\\mu$m are, respectively, $0.510\\pm0.042$ \\% and $0.41\\pm0.11$ \\%. In addition to the optical measurements reported by Alonso et al. (2009b) and Snellen et al. (2009b), these values favor a poor heat distribution to the night side of the planet and the absence of thermal inversion, but measurements at other wavelengths are needed to confirm this latter point."
    },
    "0911/0911.0749_arXiv.txt": {
        "abstract": "Using the Yale stellar evolution code, models of \\astrobjh{} based on asteroseismic measurements are constructed. A $\\chi^{2}$ minimization is performed to approach the best modeling parameters which reproduce the observations within their errors. By combining all non-asteroseismic constraints with asteroseismic measurements, we find that the observational constraints favour a model with \\textbf{a mass of 1.00$^{+ 0.01}_{- 0.02}$ $M_{\\odot}$, an age t = 6.433 $\\pm$ 0.04 Gyr, a mixing-length parameter $\\alpha$ = 1.75 $\\pm$ 0.25 , an initial hydrogen abundance $X_{i}$ = 0.605$^{+ 0.01}_{- 0.005}$ and metal abundance $Z_{i}$ = 0.0275$^{+ 0.002}_{- 0.001}$}. \\astrobjh{} is in post-main sequence phase of evolution. The modes of $l$ = 1 show up the characteristics of avoided crossings, which may be applied to test the internal structure of this type stars. Asteroseismic measurements can be used as a complementary constraint on the modeling parameters. \\textbf{The models with mass 1.00 - 1.10 $M_{\\odot}$ can reproduce the observational constraints. Existing observed data of \\astrobjh{} do not rule out these models.} ",
        "introduction": "Solar-like oscillations have been confirmed for several main-sequence and subgiant stars, such as \\astrobjca{} \\citep{bedd04}, \\astrobjcb{} \\citep{kjel05}, \\astrobjp{} \\citep{egge04b}, \\astrobjb{} \\citep{carr05}, etc. The large and small frequency separations of p-modes can provide a good estimate of the mean density and age of the stars \\citep{ulr86, ulr88}.  It has been proven that asteroseismology is a powerful tool for determining the fundamental parameters of the stars \\citep{egge05, egge06}.   \\astrobjh{ (HR 6623, HD 161797, HIP 86974)} is a G5 IV subgiant star, in the neighbourhood of the Sun. It is considered to be a solar-type star with mass 1.14 $M_{\\odot}$, effective temperature $T_{e}$ = 5596 $\\pm$ 80 K, and radius R = 1.77 $\\pm$ 0.07 $R_{\\odot}$ \\citep{fuhr98}. Recently, \\citet{bona08} detected solar-like oscillations on $\\mu$ Her and identified individual frequencies in the range of 900 - 1600 $\\mu$Hz. These seismological data will provide a constraint on the fundamental parameters of \\astrobjh{}. In this work, we try to determine modeling parameters of \\astrobjh{} based on asteroseismic constraints using the Yale Rotation Evolution Code (YREC7) in its non-rotating configuration.  The observational constraints available for \\astrobjh{} are summarized in Sect. 2, while the details of the evolutionary models and the computational method are given in Sect. 3. The results are presented in Sect. 4 and the conclusion is given in Sect. 5. ",
        "conclusions": "The mass of 1.00 $M_{\\odot}$ for the M2 and M2b models is less than the value in the previous literatures \\citep{fuhr98, take05}, which may result from that the results of \\cite{fuhr98} and \\cite{take05} were obtained without the asteroseismic constraints. The initial hydrogen mass fraction is 0.605 and 0.650 for the M2 and M1 model, respectively. Although there is a difference of 0.0015 in Z between the M1 and M2 model, the increase of the hydrogen or decrease of helium is mainly compensated by an increase of the mass in order to reach the same location in the H-R diagram: a helium-mass degeneracy \\citep{lebr93}. The M1 and M2 models have the same value of $\\chi_{osci}^{2}$ and the almost same asteroseismic features. This implies that the helium-mass degeneracy is difficult to be removed, even if the asteroseismic constraints are taken into account.  Comparing with the observed frequencies, the theoretical frequencies have a systematic shift. \\textbf{The mode of $l$ = 1 near 1238 $\\mu$Hz deviate from its expected asymptotic value in Fig. \\ref{fig6}. It may undergo an avoided crossing and be a mixed-mode. The mixed-mode could have higher inertia than other p-modes. The shift of this mode may be different from that of other p-modes. Thus we divide the $\\chi_{osci}^{2}$ into two parts: one $\\chi_{oscm}^{2}$ for mixed-modes and one $\\chi_{oscp}^{2}$ for all modes except for mixed-modes. Assuming only the $l=1$ mode near 1238 $\\mu$Hz is a mixed-mode, we obtained the value of this new $\\chi_{osci}^{2}$ is 1.07, 1.10, 3.61 and 0.72 for M1, M2, M2a and M2b, respectively. Under this assumption, M2b has a smallest $\\chi_{osci}^{2}$. }  We confirmed the results that an analysis of the H-R diagram does not allow us to determine the mixing length parameter for an evolved solar-type star \\citep{fern03} but the observed oscillation frequencies could provide a constraint on this parameter. The internal structure of the evolved solar-type stars is sensitive to the mixing length parameter at given ($T_{eff}, L$).  In this work we constructed the models for the \\astrobjh{} using the Yale stellar evolution code. By combining the non-asteroseismic constraints with the asteroseismic observations, we find that a model for $\\mu$ Her can reproduce the all non-asteroseismic and asteroseismic constraints well: \\textbf{the model with a mass of 1.00$^{+ 0.01}_{- 0.02}$ $M_{\\odot}$, an age t = 6.433 $\\pm$ 0.04 Gyr, a mixing-length parameter $\\alpha$ = 1.75 $\\pm$ 0.25 , an initial hydrogen abundance $X_{i}$ = 0.605$^{+ 0.01}_{- 0.005}$ and metal abundance $Z_{i}$ = 0.0275$^{+ 0.002}_{- 0.001}$. However, the models with mass 1 - 1.1 $M_{\\odot}$ and age $6.2 - 6.7$ Gyr also can reproduce the non-asteroseismic and asteroseismic constraints. Existing observational constraints do not rule out those models.}  The modes of $l$ = 1 show up the characteristic avoided crossings, which may be applied to test the internal structure of an evolved solar-type star. The asteroseismic observations put important constraints on the models for \\astrobjh{}, but they are not enough to really test the differences in the models. More accurate oscillation frequencies, especially the modes of $l$ = 1, are need to investigate the internal structure of this type star."
    },
    "0911/0911.0780_arXiv.txt": {
        "abstract": "{{\\bf hydrodynamics; stars: formation; stars: mass function; open clusters and associations: general}} We review progress in numerical simulations of star cluster formation.  These simulations involve the bottom-up assembly of clusters through hierarchical mergers, which produces a fractal stellar distribution at young ($\\sim 0.5$\\,Myr) ages. The resulting clusters are predicted to be mildly aspherical and highly mass-segregated, except in the immediate aftermath of mergers. The upper initial mass function within individual clusters is generally somewhat flatter than for the aggregate population. Recent work has begun to clarify the factors that control the mean stellar mass in a star-forming cloud and also the efficiency of star formation. The former is sensitive to the thermal properties of the gas while the latter depends both on the magnetic field and the initial degree of gravitational boundedness of the natal cloud. Unmagnetized clouds that are initially bound undergo rapid collapse, which is difficult to reverse by ionization feedback or stellar winds. ",
        "introduction": "Hydrodynamic simulations of star cluster formation are in their infancy. This is hardly surprising, since only in the last decade has it been possible to move beyond modelling the formation of single stars and binaries. Even with current high-performance-computing facilities, it is still unfeasible to undertake cluster simulations that can follow the collapse of individual stars down to stellar densities (cf. Bate 1998).  Thus, `star formation' in the simulations actually means the accumulation of gravitationally bound gas within `sink particles' of specified radius.  In general, the most ambitious simulations (in terms of numbers of stars formed and, hence, the ability to track hierarchical cluster assembly) are those that necessarily use rather large sink radii. This suppresses the proper modelling of close-binary formation (see Goodwin 2010) and compromises the modelling of stellar dynamical interactions. Other shortcomings include the fact that most simulations to date omit magnetic fields and model the gas using a simple barotropic equation of state.  These shortcomings should not detract from the achievement of being able to model the hydrodynamics of cluster formation at all. First, it should be stressed that there are a number of ongoing efforts to add new physical ingredients to the simulations, including magnetic fields, radiative transfer and the effects of thermal and mechanical feedback. Second, however, it should be borne in mind that even in their simplest, `vanilla' forms, these simulations represent an enormous advance on attempts to model young clusters as an $N$-body system plus background potential (e.g., Geyer \\& Burkert 2001; Boily \\& Kroupa 2003).  Although such simulations can be valuable in studying the later stages of cluster dispersal, once star formation has ceased (Goodwin \\& Bastian 2006), they are no substitute during the early, embedded phase of cluster evolution for hydrodynamical simulations that can capture the complex interplay between fragmentation, accretion, infall and stellar dynamics. This review concentrates on predictions of those `turbulent-fragmentation' simulations that have been most thoroughly analysed to date. The discussion will, therefore, focus on observables (e.g., mass segregation, mass functions, clustering statistics, cluster morphology) and should therefore be read in conjunction with Lada (2010), who reports on the current observational status in these areas. This is followed by a preliminary survey of the results to date from simulations that include a more complete suite of input physics. Finally, we try and link the insights gained from such simulations to issues of star formation on a Galactic scale.  Before embarking on this description, it is worth stressing that the simulations described here relate to the very earliest (deeply embedded) phases of cluster formation, typically covering about the first 0.5 Myr following collapse of the parent cloud. Clustering statistics are still evolving at the end of these simulations. Consequently, questions of ultimate cluster survival have to be followed over considerably longer timescales [see the discussion of cluster `infant mortality' by de Grijs (2010)]. ",
        "conclusions": "\\subsection{Observable properties of young star clusters predicted by simulations}  There are generic similarities running through all simulations described above, which make clear predictions about the observational properties of star clusters at an age of $\\sim 0.5$ Myr.  The assembly of clusters through scale-free hierarchical merging imparts the resulting clustering with a fractal character, in the sense that the structures have no preferred mass scale and the pattern of nested substructures looks similar at all scales.  These clusters are generally mass segregated, except if they are observed shortly after a cluster merger event. The time between such events is, however, longer than the internal mass-segregation timescale, so that mass segregation is the predicted observational norm.  The clusters are typically elliptical, their relatively mild flattening (typical axis ratio $<2:1$) being a trade-off between extension in the plane of the most recent merger and ongoing sphericalization by relaxation effects.  \\subsection{Implications of cluster-formation simulations for galactic-scale star formation}  \\subsubsection{The IMF in the field versus clusters}  Most of the clusters formed in these simulations are not long-lived and are expected to contribute to the galactic-field population.  The largest clusters are $\\sim 1000$ M$_\\odot$ in mass and thus, depending on the subsequent history of gas removal from such objects, they might later be identified observationally as open clusters.  Thus, simulations can start to shed light on the formation of both the `field' and `cluster' population in the Galaxy.  In the bottom-up scenario outlined here, essentially all stars are in clusters at some point.  Since it is the most populous clusters that are most likely to survive, the ultimate destiny of a particular star as a field or cluster member is largely determined by how many mergers it undergoes, which depends at least partly on the the degree of gravitational boundedness of the parent cloud locally.  Since `cluster' and `field' stars follow similar early evolutionary histories, one does not expect major differences between `field' and `cluster' stars. For example, the closest dynamical interactions (which may determine the formation of binary systems and the possible destruction of protoplanetary discs) often occur in the early formation stages when essentially {\\it all} stars (regardless of their ultimate designation as `field' or `cluster' objects) are in small-$N$ groupings. The lower end of the IMF (as measured by the ratio of stars to brown dwarfs, for example) is thus expected to be similar in rich clusters and what will ultimately become the field.  In broad terms, successive cluster mergers do not greatly affect the lowest-mass stars, since these tend not to undergo significant accretion in cluster cores.  The upper end of the IMF {\\it does}, however, depend on the history of successive cluster mergers, since mergers furnish the opportunity for further accretional growth by the most massive stars.  Two observational consequences of this scenario are therefore (i) that the slope of the IMF is somewhat flatter in populous clusters than in the total population and (ii) that the IMF in a cluster is truncated at the upper end at a value that depends on cluster richness. Weidner \\& Kroupa (2005; see also Pflamm--Altenburg \\ea 2009) have set out a number of potential observational consequences of a scenario in which the integrated galactic IMF is, in fact, steeper than the IMF in individual clusters and may vary with galactic properties [see Hoversten \\& Glazebrook (2008) for possible observational corroboration of this effect].  The simulations support such a distinction, but in the cases studied so far the difference is rather modest (see figure 5).  Finally, it should be stressed that the above is derived from simulations that omit stellar feedback. Pilot studies involving both ionizing radiation and feedback from stellar winds suggest that the latter may be important in limiting individual stellar masses at the high-mass end and may thus introduce an additional truncation to the stellar IMF that is independent of cluster environment.  \\subsubsection{The efficiency of star formation}  \\begin{figure} \\parbox[t]{0.48\\columnwidth} {\\includegraphics[width=0.45\\columnwidth]{price09.ps}} \\parbox[t]{0.52\\columnwidth} {\\includegraphics[width=0.435\\columnwidth]{clark08.ps}} \\caption{Effect of magnetic-field strength (increasing left to right in the left-hand panel; Price \\& Bate 2009) and initial ratio of kinetic to gravitational energy (increasing left to right in the right-hand panel; Clark \\ea 2008) on star-formation efficiency.} \\end{figure}  The fraction of a cloud that is incorporated in stars per free-fall time is an important quantity since it controls the relationship between the total molecular gas in the Galaxy and the Galactic star-formation rate. Current observational estimates of these quantities imply a star-formation efficiency of $<10$\\% (Evans \\ea 2009). It is found that, in the absence of magnetic fields, simulations starting from gravitationally bound initial conditions form stars much more rapidly than this. In fact, pilot simulations (Dale \\ea 2005; Dale \\& Bonnell 2008) suggest that, in this case, the collapse is so rapid that by the stage that feedback from massive stars (photoionization or winds) switches on, it can be hard to disrupt the dense accretion flows and quench the star-formation rate significantly.  Instead, it would seem that low star-formation efficiencies can be achieved either in simulations that start from globally unbound clouds (Clark \\& Bonnell 2004; Clark \\ea 2008) or where clouds are threaded by (moderately supercritical) magnetic fields (Price \\& Bate 2009): see figure 10.  Either of these conditions helps to stop runaway collapse in the early stages of cloud evolution and thus may make it easier for stellar feedback to quench star formation at a later stage. This is, however, speculation: currently, simulations have just reached the stage where they have started to implement these effects separately.  Further investigation is required to understand how these processes interact in the formation of stars and clusters from molecular clouds."
    },
    "0911/0911.1786_arXiv.txt": {
        "abstract": "This paper presents theoretical star formation and chemical enrichment histories for the stellar halo of the Milky Way based on new chemodynamical modeling. The goal of this study is to assess the extent to which metal-poor stars in the halo reflect the star formation conditions that occurred in halo progenitor galaxies at high redshift, before and during the epoch of reionization. Simple prescriptions that translate dark-matter halo mass into baryonic gas budgets and star formation histories yield models that resemble the observed Milky Way halo in its total stellar mass, metallicity distribution, and the luminosity function and chemical enrichment of dwarf satellite galaxies. These model halos in turn allow an exploration of how the  populations of interest for probing the epoch of reionization are distributed in physical and phase space, and of how they are related to lower-redshift populations of the same metallicity. The fraction of stars dating from before a particular time or redshift depends strongly on radius within the galaxy, reflecting the ``inside-out\" growth of cold-dark-matter halos, and on metallicity, reflecting the general trend toward higher metallicity at later times. These results suggest that efforts to discover stars from $z > 6 - 10$ should select for stars with [Fe/H] $\\lesssim -3$ and favor stars on more tightly bound orbits in the stellar halo, where the majority are from $z > 10$ and $15-40$\\% are from $z > 15$. The oldest, most metal-poor stars -- those most likely to reveal the chemical abundances of the first stars -- are most common in the very center of the Galaxy's halo: they are in the bulge, but not of the bulge. These models have several implications for the larger project of constraining the properties of the first stars and galaxies using data from the local Universe. ",
        "introduction": "How and when the first stars and galaxies formed is a central and abiding question in modern astrophysics. The high-redshift frontier has steadily advanced into the end of reionization, $z \\sim 6 - 8$, but the truly first stars lie still beyond our view. It is now recognized, on the one hand, that if the first stars are massive and isolated in their own small dark matter halos \\citep{Abel:02:93, Bromm:02:23} they will be very faint and difficult to detect directly at high redshift. On the other hand, chemical abundances in the most metal-poor stars may preserve a record of early chemical evolution dating all the way back to the first stars \\citep{Freeman:02:487a, Beers:05:531}. This effort, sometimes termed ``Galactic Archaeology'', has the strategic goal of using the most metal-poor stars in the Milky Way and Local Group to address open questions about the first stars and galaxies as a complement to direct study at high redshift.  The goal of this series of papers is to advance the theoretical basis on which these tests can be performed. A large portion of this theoretical foundation consists in demonstrating that the oldest stars in the Milky Way can be used in such a fashion. Stated another way, it would be fatal to the whole enterprise if it turned out that the Milky Way contained no stars from $z > 6$, or from the Epoch of Reionization. As yet we have no purely empirical basis for claiming that it does; all such statements are still necessarily model-dependent. Dark-matter only numerical simulations \\citep{White:00:327, Scannapieco:06:285, Diemand:07:859} have already demonstrated that Milky-Way sized DM halos should contain significant structures that originated before $z > 6$ (roughly the current observational frontier). This is well-established and seems necessarily true in the CDM paradigm. Yet while the dark matter holds some interest in its own right, it is still only the dark scaffolding on which the Galaxy's stellar populations are built. When we consider the stars, a host of other questions arise: How many low-metallicity stars should we expect in the Milky Way? How many of these should date from $z > 6$, if any? What are the most advantageous places to look for these true survivors, and what should we expect to find there? To what extent are these early extremely metal-poor (EMP) stars obscured by later populations of similar metallicity but of less relevance to the first stars? All these questions must be answered to accomplish the overarching goal of ``Galactic Archaeology'', but they have no easy answers.  That metal-poor stars\\footnote{This paper adopts the \\cite{Beers:05:531} terminology for labeling metal-poor stars, e.g. EMPs, or extremely metal-poor stars, have [Fe/H] $<-3$, and so on.} address the first stars has been taken for granted as a working hypothesis by many observational programs \\citep{Cayrel:04:1117, Christlieb:02:904, Frebel:05:871}. Yet this claim rests on no firm empirical proof: though some metal-poor stars are known to have ages consistent with 13 Gyr ago, these ages cannot yet be determined precisely or for a large sample. The observational approach rests rather on the somewhat different claim that these stars are related to the first stars by way of the ``chemical clock''; they are considered related to the first star formation in their region of space by virtue of their low metallicity, not by their age. Yet not all [Fe/H] $\\sim -3$ stars are necessarily created equal, and some may arise much later and be chronologically unrelated to the epoch of first star and galaxy formation in which we are interested. The relationship between the two clocks is likely to be a complex one. Further, even if we take the relationship between EMP stars and the first stars at face value, this assumed relationship does not tell us directly how much and what kind of information EMPs contain about the first stars, or how many stars are needed to answer the open questions.  These issues of interpretation reduce to a core problem that must be solved for the larger enterprise to advance. That is: the readily observable quantities in the present-day Galaxy -- the luminosity, temperature, chemical abundances, and orbits of individual stars and their trends in stellar populations -- are only indirectly related to the properties of stars and gas in early galaxies that preceded the Milky Way, and to the physical processes that governed their formation. For example, nucleosynthetic yields from supernovae enter newly formed stars only after some degree of dispersal and mixing in interstellar gas. The degree of mixing may vary with time and place. For another example, long-lived stars formed in the shallow potential wells of early galaxies -- the exact populations we now study in the MW halo -- have been stripped from their parent halo and had their present orbits in the Galaxy set by a long and stochastic series of mergers and interactions that cannot be reconstructed exactly, if at all. These two examples are likely essential ingredients in the chain of events leading from the stars we want to understand to the stars we actually observe. There appears to be little alternative to an attempt at a single model that includes these processes in building quantitative links between theoretical ideas about early galaxies and observations in the present day.  The observational frontier is being advanced by large surveys that detect and measure galactic structure and substructure, and by mining these surveys for metal-poor stars and then measuring their detailed abundances in intensive spectroscopic campaigns. As examples of the former, there is now a proliferating number of known stellar streams \\citep{Yanny:00:825} and faint dwarf galaxies (Belokurov et al. 2007) discovered in the Sloan Digital Sky Survey. Using SDSS, \\cite{Carollo:07:1020} confirmed two distinct chemical and kinematic components of the stellar halo near the Sun, using a much larger sample than available before SDSS. As examples of the latter, the HK  \\citep{Beers:92:1987} and HES \\citep{Cohen:04:1107} surveys have been thoroughly mined for metal-poor stars, leading to discovery of the most metal-poor stars known \\citep{Christlieb:02:904, Frebel:05:871}, to the discovery of a class of carbon-enhanced metal-poor stars (Lucatello et al. 2005), to increasingly precise measurements of the stellar metallicity distribution \\citep{Schoerck:08:1172}, and to intensive studies of the abundances and distribution of the heaviest chemical elements \\citep{Barklem:05:129}. These discoveries are important motivations for this theoretical work. They provide crucial information about the structure of the Milky Way's halo for which any full synthetic model must ultimately account. Each of these individual studies provides data on a key piece of the overall puzzle; collectively, they suggest that a very large quantity of information may lie in stellar halo populations. Yet observational studies have a strong incentive to carefully select their target populations, so it is difficult to draw general inferences about the behavior of the full Milky Way system without a model that calculates its predictions at a level of detail commensurate with the observations. Given the complexity of the datasets and the large dynamic range between the extremes of the Milky Way on the one hand and abundances in individual stars on the other, building such model presents a significant challenge.  Paper I in this series \\citep{Tumlinson:06:1} made an attempt at such a model in the form of a new stochastic, hierarchical model of Galactic chemical evolution, based purely on the ``Extended Press-Schecter'' merger-tree formalism \\citep{Somerville:99:1}. That model favored close tracking of chemical abundances over a realistic dynamical model of halo formation, with the goal of constraining the IMF of the first stars by way of their contributions to chemical abundances at low metallicity. That technique worked adequately for the tests performed in that paper, which tracked gas and chemical enrichment budgets with a high level of detail but which could not calculate the spatial or kinematic distributions of model stellar populations. By itself, it yields only a part of a model. Using a similar semi-analytic merger-tree technique, \\cite{Salvadori:07:647} studied the chemical evolution of the halo in response to the first stars, but with the similar limitation that no detailed treatment of the dynamics was included. At the same time, \\citeauthor{Bullock:05:931} and their collaborators built dynamical models of halo formation based in non-cosmological N-body simulations and applied simple chemical evolution prescriptions to assess the relationships between kinematics and chemical abundances \\citep{Robertson:05:872, Font:06:585, Font:06:886}. Their work has been compared extensively to surveys of the Milky Way halo and has been shown to provide a realistic description of the halo's basic properties - mass, degree of substructure, and trends in chemical abundances. While each was successful in their own way, none of these studies excelled at understanding the formation of the first stars and galaxies by locating their descendants in the Milky Way. Paper I lacked any treatment of dynamics, and the \\citeauthor{Bullock:05:931} halo models did not simulate the Milky Way in a cosmological ``live halo'' and did not follow the merger histories of the dwarf galaxies or -- most importantly -- the formation history of the Milky Way host. Because of these limitations neither the Paper I nor the \\citeauthor{Bullock:05:931}  models were well-suited for studying the present-day properties and distribution of stellar populations formed at very high redshift.  The need for a strong synthetic and flexible Milky Way model and the limitations of the previous efforts inspire the present study, which attempts to retain the strengths of the previous work while rectifying some of its major weaknesses. With respect to Paper I, the present paper takes a step backwards to reassess a number of key assumptions that underlie the basic approach to probing the early Universe with local signatures; the spirit of the new approach is to develop a workable model with a minimum degree of complexity. These reassessments are possible because the present paper deals with a much more sophisticated realization of the Galaxy's dark matter halo, and they are necessary because the earlier study made a number of simplifying assumptions that are now testable in detail. The goals of the present study are:  \\begin{itemize} \\item[1.] To assess the fundamental basis for the ``Galactic Archaeology'' effort \\citep{Beers:05:531} by demonstrating that the Milky Way should contain stellar populations that date from the Dark Ages and the Epoch of Reionization, by extending previous DM-only modeling efforts to account for how the Galaxy's progenitors acquire and process their gas and build up their chemical abundances. \\item[2.] To identify the most important global and local physical influences on the chemodynamical assembly of the Milky Way's stellar halo, to assess the sensitivity of the resulting stellar populations to these influences, and to help guide future in-depth studies. \\item[3.] To lay out a framework for the calculation of key observational quantities related to Galactic chemical evolution, in absolute physical units where possible, and their sensitivity to the key influences; for example, how does cosmic reionization affect the radial density profile of stellar populations in the halo? \\item[4.] To suggest methods by which stellar populations that originate at high redshift can be isolated from later populations with which they may be easily confused in random samples; for example, when looking at any random star with [Fe/H] $< -2$, what is the chance that it reflects conditions at $z > 6$? \\end{itemize}  It is with all these ideas in mind that this paper is organized as follows:  Section~2 describes the N-body simulations of Milky-Way-like DM halos and \\S~\\ref{section-results1} describes the mass assembly history of the inner portions of the MW halo, considering dark matter only. Section~\\ref{section-methods2} describes the chemical evolution models that run within the dark-matter halo merger trees. Section~\\ref{section-results2} shows how the models provide a realistic description for the gross properties of the real Milky Way halo, as an observational check on the model construction and parameter choices. Section~\\ref{section-results3} shows the metal abundances of stars from the first MW progenitors and maps out broad trends for their variation within the MW halo at the present time. Section~\\ref{section-results4} discusses the general results of the models and draws some implications for observational studies of metal-poor stars and star formation in the early Universe. ",
        "conclusions": ""
    },
    "0911/0911.1579_arXiv.txt": {
        "abstract": "{ Existing theory and models suggest that a Type I (merger) GRB should have a larger jet beaming angle than a Type II (collapsar) GRB, but so far no statistical evidence is available to support this suggestion. In this paper, we obtain a sample of 37 beaming angles and calculate the probability that this is true. A correction is also devised to account for the scarcity of Type I GRBs in our sample. The probability is calculated to be 83\\% without the correction and 71\\% with it. ",
        "introduction": "% \\label{sect:intro}   There are two intrinsically different phenomena that give rise to gamma-ray bursts (GRBs). One is the merging of two compact objects, such as that of a neutron star and a black hole; the other is the core collapse of a massive star ``collapsar\"), such as the birth of a Type Ib/c Supernova. According to a classification scheme developed recently (Zhang, 2007 and Zhang et al, 2009), a GRB obtained from the former channel is defined as a Type I GRB, while one from the latter would be classified as a Type II GRB. This classification scheme, though more intrinsic than the classical short/hard vs long/soft categories, is not easy to carry out at the present stage, since many different criteria need to be applied in order to determine the progenitor of a GRB, few of which are decisive. One of these criteria is that Type I GRBs are usually short/hard ( \\begin{math} T_{90}\\leq2\\,{\\rm s} \\end{math} ), while Type II GRBs are usually long/soft ( \\begin{math} T_{90}>2\\,{\\rm s} \\end{math} ). This criteria, though supported by many individual cases, is not decisive, as any short GRB could theoretically have been a long GRB had it occurred at a high enough redshift. Of the criteria that are decisive, such as whether or not a GRB emanates gravitational radiation, most are impractical and the remainder is not applicable to a great majority of GRBs at the present stage.  GRBs eject their energy in the form of jets. These jets have already been modeled, and their beaming angles (also called opening angles) can be calculated from observed physical data (Sari et al, 1999). The size of the beaming angle, according to the models, is highly dependent on the nature of the GRB progenitor: theoretically, Type II GRBs should have smaller beaming angles than Type I GRBs, due to the collimation effect of the stellar wind hugging collapsars. However, this dependency has not yet been supported statistically, possibly due to the aggravatingly small sample of Type I GRB jet beaming angles available, hence this paper, which specifically seeks to find statistical evidence of this  theoretical relationship.  This paper is divided into 5 sections, of which this introduction is the first. In the second, we will list the relevant observational data and use it to obtain the jet beaming angles of 37 GRBs. The third section will cover data reduction and statistical analysis. The results and discussion of the analysis will be presented in the fourth section, and the entire paper will be summarized in the fifth section, which is the conclusion. ",
        "conclusions": "\\label{sect:analysis}  In this paper, we have obtained a Sample I of 3 jet beaming angles and a Sample II of 34 jet beaming angles. These two samples are expected to be representative of Type I and Type II GRBs respectively. After that, normal (Gaussian) distributions have been fitted to the samples, resulting in Fit I and Fit II. Taking these fits to be representative of Type I and Type II GRBs, we then proceed to calculate the probability that a random variable taken from Fit I is larger than another taken from Fit II, thereby deriving the probability that a Type I GRB has a larger opening angle than a Type II GRB. Taking into account the uncertainty caused by the very small Sample I, we then devise a correction for the probability stated above, and reach a probability that is much lower but still significantly high nevertheless.  If we take Sample I to be a sample perfectly representative of Type I GRBs, then it could be said, with an 83\\% degree of certainty, that Type I GRBs have larger jet beaming angles than Type II GRBs. However, this 83\\% drops to 71\\% once we take into account the uncertainty caused by our small sample in Sample I. In either case, it could be justifiably concluded, from our current samples, that Type I GRBs generally have a larger beaming angle in comparison to Type II GRBs.  Our results are supportive of current models and theory. Type I GRBs in general have larger beaming angles in comparison to Type II GRBs, with a fairly high degree of certainty, though by far not high enough to become a decisive criterion as to whether a GRB is of Type I or Type II origin (so far).  It is possible that with a larger Sample I, higher levels of certainty could be reached. Therefore, further observations that yield data concerning Sample I jet beaming angles (preferably Type I GRB jet beaming angles) are required for progress. It is also possible that with a large enough Sample I, the correction methods used in this paper (i.e. the correction for Sample I) will become obsolete, however, given the current rate at which Sample I opening angles are being derived, it seems unlikely that this would be the case in the near future. Lastly, it would be desirable to find a parametrical fit which could be supported theoretically, and which yields a higher level of acceptance than the Gaussian.        \\normalem"
    },
    "0911/0911.1609_arXiv.txt": {
        "abstract": "The mean-square current quadrupole moment associated with vorticity fluctuations in high-Reynolds-number turbulence in a differentially rotating neutron star is calculated analytically, as are the amplitude and decoherence time of the resulting, stochastic gravitational wave signal. The calculation resolves the subtle question of whether the signal is dominated by the smallest or largest turbulent eddies: for the Kolmogorov-like power spectrum observed in superfluid spherical Couette simulations, the wave strain is controlled by the largest eddies, and the decoherence time approximately equals the maximum eddy turnover time. For a neutron star with spin frequency $\\nu_{\\rm s}$ and Rossby number ${\\rm Ro}$, at a distance $d$ from Earth, the root-mean-square wave strain reaches $h_{\\rm RMS} \\approx 3\\times 10^{-24}\\, {\\rm Ro}^3 (\\nu_{\\rm s} / 30\\,{\\rm Hz})^3 (d/{\\rm 1\\,kpc})^{-1}$. Ordinary rotation-powered pulsars ($\\nu_{\\rm s} \\lesssim 30\\,{\\rm Hz}$, ${\\rm Ro} \\lesssim 10^{-4}$) are too dim to be detected by the current generation of long-baseline interferometers. Millisecond pulsars are brighter; for example, an object born recently in a Galactic supernova or accreting near the Eddington rate can have $\\nu_{\\rm s} \\sim 1\\,{\\rm kHz}$, ${\\rm Ro}\\gtrsim 0.2$, and hence $h_{\\rm RMS} \\sim 10^{-21}$. A cross-correlation search can detect such a source in principle, because the signal decoheres over the time-scale $\\tau_{\\rm c} \\approx 1 \\times 10^{-3} \\, {\\rm Ro}^{-1} (\\nu_{\\rm s} / 30\\,{\\rm Hz})^{-1} \\, {\\rm s}$, which is adequately sampled by existing long-baseline interferometers. Hence hydrodynamic turbulence imposes a fundamental noise floor on gravitational wave observations of neutron stars, although its polluting effect may be muted by partial decoherence in the hectohertz band, where current continuous-wave searches are concentrated, for the highest frequency (and hence most powerful) sources. This outcome is contingent on the exact shape of the turbulent power spectrum, which is modified by buoyancy and anisotropic global structures, like stratified boundary layers, in a way that is understood incompletely even in laboratory situations. ",
        "introduction": "} Shortly after the discovery of radio pulsars, speculation arose that the superfluid interior of a differentially rotating neutron star is turbulent \\citep{gre70}. Since then, the theme has resurfaced intermittently during the quest to understand pulsar rotational irregularities, like glitches and timing noise \\citep{and78,tsa80}. Neutron star turbulence can be hydrodynamic, taking the form of a Kolmogorov-like cascade of {\\em macroscopic} eddies at high Reynolds numbers \\citep{per05,per06a,mel07,per08}. Complicated vorticity patterns of this sort are observed in terrestrial experiments on spherical Couette flow, which undergo transitions to nonaxisymmetric flow states at high Reynolds numbers, e.g.\\ spiral, shear herringbone, or Taylor-G\\\"{o}rtler vortices \\citep{bel79,buh90,jun00,sha01,nak02a,nak02b,nak05a,per06b,per08}, relaminarization \\citep{nak05b}, or Stewartson layer disruption \\citep{hol03,hol04,hol06,wei08}. Neutron star turbulence can also be quantum mechanical, comprising a self-sustaining tangle of quantized {\\em microscopic} vortices \\citep{gla74,jou04,jou06}, excited by bulk two-stream instabilities \\citep{and07}, interfacial two-stream instabilities \\citep{bla02,mas05}, or meridional circulation \\citep{per05,per06a,mel07}. In general, macroscopic and microscopic superfluid turbulence appear to trigger each other; it is an unsolved, chicken-or-egg question as to which comes first \\citep{bar01,tsu09}.  Turbulence powered by differential rotation is axisymmetric when averaged over long times but nonaxisymmetric instantaneously. Turbulent flows therefore emit stochastic gravitational waves. In an incompressible fluid, the waves arise mainly from current quadrupole (and higher multipole) source terms. In a compressible fluid, the mass multipoles also contribute; indeed, they can dominate, e.g.\\ during post-glitch Ekman pumping in a neutron star \\citep{van08}. Recently, \\citet{per06b} pointed out that there exists a fundamental theoretical uncertainty regarding the shape and strength of the gravitational wave signal emitted by hydrodynamic turbulence. The mechanical stress-energy in Kolmogorov-like turbulence is contained mostly in large eddies near the stirring scale. Naively, therefore, one might expect the gravitational wave signal to look like a `dirty sinusoid', which reflects circulation on the largest scales and decoheres in approximately one rotation period. However, the instantaneous wave strain is proportional to the second time-derivative of the stress-energy tensor, and this quantity is greatest for small eddies near the dissipation scale, which turn over most quickly. If the latter effect dominates, one might expect the signal to resemble white noise. Of course, large eddies match better to low-order multipoles than small eddies, and low-order multipoles typically dominate the gravitational wave strain far from the source \\citep{was09}. A careful calculation is therefore required to select between the various possibilities and reliably estimate the detectability of the signal.  In this paper, we undertake such a calculation by combining the formalism of \\citet{kos02} and \\citet{gog07}, developed to calculate the gravitational radiation from a turbulent, first-order phase transition in the early Universe, together with the formalism of \\citet{was09}, developed to calculate the gravitational radiation from nonaxisymmetric vorticity fluctuations in neutron stars, e.g.\\ due to clusters of quantized superfluid vortices. In \\S\\ref{sec:tur2}, we analyze global hydrodynamic simulations of incompressible, shear-driven neutron star turbulence to extract the vorticity correlation function, which feeds into the statistics of stress-energy fluctuations in the source. We then calculate analytically the current multipole moments, root-mean-square (RMS) wave strain, and decoherence time of the resulting, stochastic gravitational wave signal in \\S\\ref{sec:tur3}. The results are applied in \\S\\ref{sec:tur4} to estimate the detectability of the signal, e.g.\\ with long-baseline interferometers like the Laser Interferometer Gravitational-Wave Observatory (LIGO), and its polluting effect on continuous-wave searches currently under consideration. Hydrodynamic turbulence imposes a fundamental, quantifiable noise floor on gravitational wave observations of neutron stars. Astrophysical implications, including the rate of gravitational wave braking in different types of neutron stars, are briefly canvassed. ",
        "conclusions": ""
    },
    "0911/0911.1101_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\setcounter{equation}{0}  Magnetic fields are found to be associated with nearly all structures in the universe. They are observed on stellar upto possible supercluster scales \\cite{mag}. In some processes such as star formation they can play an important role. Equally the physics of cosmic rays indicate the existence of a large scale galactic magnetic field. Furthermore, ultra high energy cosmic rays which are most likely of extra galactic origin could be used to study the properties of galactic and extra galactic magnetic fields \\cite{cr}.   As to the origin of the observed magnetic fields there seem to be generally speaking two broad classes of mechanisms \\cite{mag}. On the one hand there are battery-type mechanisms which work on the basis of charge separation which leads to a current and finally induces a magnetic field. On the other hand there are dynamo type mechanisms which amplify an initial seed magnetic field. The former have a coherence scale of the order of the domain associated with the battery-mechanism, which in a cosmological context is always smaller than the horizon at the epoch of creation. The latter involves magnetic fields whose correlation length is of the order of the region of interest, which in the case of a galactic magnetic field, would be a proto galactic scale of the order  of  1 Mpc today which requires a mechanism to create magnetic fields with rather large coherence scales. A natural mechanism for this is provided by the amplification of perturbations of the electromagnetic field during inflation. Since quantum perturbations are stretched beyond the horizon during inflation and upon leaving the causal domain becoming classical, there is no problem with the coherence length. However, in general, in a background geometry with flat spatial section the resulting field strength of the primordial magnetic field after inflation in standard electrodynamics is far too small to, for example, serve as a seed field for a potential galactic dynamo explaining the observed  galactic magnetic field of  the order of  $10^{-6}$ G. Starting with \\cite{tw} this motivated an intensive study of alternative models of some kind of coupling of linear electrodynamics to either other fields in the theory, such as scalar fields \\cite{mag-sc} or gravity \\cite{tw,mgrav} including extra dimensions \\cite{mag-ex}. Other models break explicitly Lorentz invariance considering a non zero photon mass \\cite{tw,BLI}. Recently there has also been interest in magnetic field generation within models of nonlinear electrodynamics which naturally occur when quantum corrections and self couplings of the electromagnetic field are taken into account \\cite{nled}. Furthermore, there are models within the standard model or its supersymmetric extensions \\cite{other}. In the case of models with curved spatial sections it was shown that even in the minimally coupled model of linear electrodynamics strong enough magnetic fields can be created \\cite{tk}.  Here we are returning to a model where the electromagnetic field is gravitationally coupled which first has been proposed in the context of the generation of primordial magnetic fields during inflation in \\cite{tw}. In particular the model under consideration derives form the one loop effective action of vacuum polarization in QED in a gravitational background \\cite{dh}. As shown in \\cite{dh}  this leads to birefringence of the electromagnetic wave where the photons have  velocities depending on their polarization which can exceed the speed of light. This last observation leads to an interesting causal structure of these space-times \\cite{gs}. ",
        "conclusions": "\\setcounter{equation}{0} Primordial magnetic field generation has been investigated in a model where  the electromagnetic field is coupled to various curvature terms which is motivated by the form of the one loop effective action of vacuum polarization in QED in a gravitational background \\cite{dh,tw}. Here the resulting magnetic field spectrum has been calculated explicitly  by matching a stage of power law inflation directly to the standard radiation dominated phase, solving the mode equation and finding the appropriate Bogoliubov coefficient. The ratio of magnetic field energy density over background radiation energy density $r$ calculated at a galactic scale of 1 Mpc has been employed to compare the model with the minimal required values necessary to seed the galactic dynamic field either directly or relying on the mechanism of a galactic dynamo. For simplicity it was assumed that the coefficients of all additional curvature terms are of the same order. A constraint on the parameters, which is part of the approximation used to derive the mode equation on superhorizon scale, leads to a maximum value of the magnetic field strength or equivalently the ratio $r$ which is a function of $\\beta$ and $\\frac{H_1}{M_{\\rm P}}$. Here $\\beta$ is the exponent characterizing inflation and $H_1$ is the value of the Hubble parameter at the beginning of the radiation dominated era. As can be seen from figure \\ref{fig1} there is a region in parameter space where the resulting  magnetic field strength is  larger than $B_s>10^{-20}$ G corresponding to $r>10^{-37}$, namely, $\\beta<-2.4$, $\\frac{H_1}{M_{\\rm P}}>10^{-18}$ bounded by the corresponding contour line. De Sitter inflation corresponds to $\\beta=-1$. In this case there is no significant magnetic field generation since the mode functions during inflation as well as during the radiation dominated stage are plane waves.  Finally, it has been checked that the resulting maximum magnetic field strength satisfies the bound due to gravitational wave production \\cite{cd}. Imposing this bound leads to a maximal value of $\\frac{H_1}{M_{\\rm P}}$ which, however, is always much larger than the maximal value considered here, as can be seen from figure \\ref{fig2}. In summary, there is a range of parameters $\\beta$ and $\\frac{H_1}{M_{\\rm P}}$ for which magnetic fields are generated that are strong enough to explain the galactic magnetic field."
    },
    "0911/0911.4512_arXiv.txt": {
        "abstract": "{The Keck telescope's High Resolution Spectrograph (HIRES) has previously provided evidence for a smaller fine-structure constant, $\\alpha$, compared to the current laboratory value, in a sample of 143 quasar absorption systems: $\\da=(-0.57\\pm0.11)\\times10^{-5}$. The analysis was based on a variety of metal-ion transitions which, if $\\alpha$ varies, experience different relative velocity shifts. This result is yet to be robustly contradicted, or confirmed, by measurements on other telescopes and spectrographs; it remains crucial to do so. It is also important to consider new possible instrumental systematic effects which may explain the Keck/HIRES results. \\citet[][arXiv:0904.4725v1]{GriestK_09a} recently identified distortions in the echelle order wavelength scales of HIRES with typical amplitudes $\\pm$250\\,\\ms. Here we investigate the effect such distortions may have had on the Keck/HIRES varying $\\alpha$ results. Using a simple model of these intra-order distortions, we demonstrate that they cause a random effect on \\da\\ from absorber to absorber because the systems are at different redshifts, placing the relevant absorption lines at different positions in different echelle orders. The typical magnitude of the effect on \\da\\ is $\\sim$$0.4\\times10^{-5}$ for individual absorbers which, compared to the median error on \\da\\ in the sample, $\\sim$$1.9\\times10^{-5}$, is relatively small.  Consequently, the weighted mean value changes by less than $0.05\\times10^{-5}$ if the corrections we calculate are applied.  Unsurprisingly, with corrections this small, we do not find direct evidence that applying them is actually warranted. Nevertheless, we urge caution, particularly for analyses aiming to achieve high precision \\da\\ measurements on individual systems or small samples, that a much more detailed understanding of such intra-order distortions and their dependence on observational parameters is important if they are to be avoided or modelled reliably. ",
        "introduction": "\\label{sec:intro}  The Standard Model of particle physics is parametrized by several dimensionless `fundamental constants', such as coupling constants and mass ratios, whose values are not predicted by the Model itself. Instead their values and, indeed, their constancy must be established experimentally. If found to vary in time or space, understanding their dynamics may require a more fundamental theory, perhaps one unifying the four known physical interactions. One such parameter whose constancy can be tested to high precision is the fine-structure constant, $\\alpha\\equiv e^2/\\hbar c$, characterising the strength of electromagnetism. Earth-bound laboratory experiments, conducted over several-year time-scales, which use ultra-stable lasers to compare different atomic clocks based on different atoms/ions \\citep[e.g.~Cs, Hg$^+$, Al$^+$, Yb$^+$, Sr, Dy; e.g.][]{PrestageJ_95a}, have limited the time-derivative of $\\alpha$ to $\\dot{\\alpha}/\\alpha=(-1.6\\pm2.3)\\times10^{-17}{\\rm \\,yr}^{-1}$ \\citep{RosenbandT_08a}  \\begin{figure*}[t!] \\begin{center} \\includegraphics[width=0.70\\textwidth]{SAIT_MURPHY_FIG1.eps} \\caption{\\footnotesize Transitions used in MM analyses versus their rest-frame wavelength and their sensitivity to variations in $\\alpha$. Note the very different signature of a varying $\\alpha$ in absorption systems where only the Mg and Fe{\\sc \\,ii} transitions redwards of 2300\\,\\AA\\ are detected and fitted, compared with the more complicated signature for systems containing a subset of the bluer transitions. Note that the relative strengths of the different transitions typically observed in quasar absorption systems is not portrayed here.} \\label{fig:shifts} \\end{center} \\end{figure*}  Important probes of variations over cosmological space- and time-scales are narrow absorption lines imprinted on the spectra of distant, background quasars by gas clouds associated with foreground galaxies \\citep{BahcallJ_67b}. In particular, electronic resonance transitions from the ground states of metallic atoms and ions are useful indicators of $\\alpha$ variation for redshifts up to $\\sim$4 -- when the universe was $\\sim$10\\,\\% of its current age -- where the transitions are easily accessed from ground-based optical telescopes. The many-multiplet (MM) method \\citep{DzubaV_99a,WebbJ_99a} utilises the relative wavelength shifts expected from different transitions from different multiplets of various atom/ions to measure $\\alpha$ from quasar absorption spectra.  The velocity shift, $\\Delta v_i$, of transition $i$ due to a small relative variation in $\\alpha$, $\\Delta\\alpha/\\alpha\\ll 1$, is determined by the $q$-coefficient for that transition, \\begin{equation}\\label{eq:da1} \\omega_{z,i} \\equiv \\omega_{0,i} + q_i\\left[\\left(\\alpha_z/\\alpha_0\\right)^2-1\\right]\\,~{\\rm or} \\end{equation} \\begin{equation}\\label{eq:da2} \\frac{\\Delta v_i}{c} \\approx -2\\frac{\\Delta\\alpha}{\\alpha}\\frac{q_i}{\\omega_{0,i}}\\,, \\end{equation} where $\\omega_{0,i}$ \\& $\\omega_{z,i}$ are the rest-frequencies in the lab and at redshift $z$, $\\alpha_0$ is the lab value of $\\alpha$ and $\\alpha_z$ is the shifted value measured from an absorber at $z$. The MM method is the comparison of measured velocity shifts from several transitions (with different $q$-coefficients) to compute the best-fit $\\Delta\\alpha/\\alpha$. Figure \\ref{fig:shifts} illustrates the wavelength shifts experienced by the transitions typically utilized in MM analyses.  Some evidence for $\\alpha$-variation has emerged over the last decade from quasar spectra observed with HIRES \\citep{VogtS_94a} on the Keck I 10-m telescope in Hawaii. The first tentative evidence in \\citet{WebbJ_99a} was subsequently strengthened with larger samples of absorbers (\\citealt{MurphyM_01a}; \\citealt[hereafter \\citetalias{MurphyM_03a}]{MurphyM_03a}). MM analysis of 143 absorption spectra, all from the Keck/HIRES instrument, currently indicate a smaller $\\alpha$ in the clouds at the fractional level $\\da=(-0.57\\pm0.11)\\times10^{-5}$ over the redshift range $0.2<\\zabs<4.2$ \\citep[][hereafter \\citetalias{MurphyM_04a}]{MurphyM_04a}.  Obviously, confirmation of such a potentially fundamental result must be made with many other telescopes and spectrographs. Similar MM studies using the Ultraviolet and Visual Echelle Spectrograph on the ESO Very Large Telescope in Chile are also beginning to yield constraints on \\da, but none yet rule out (or confirm) the Keck/HIRES results \\citep{MurphyM_08a}. At the same time, further searches for subtle systematic errors which, despite extensive searches \\citep[e.g.][]{MurphyM_01b}, have evaded detection so far, must be considered in detail. Of particular importance are systematic errors in the wavelength calibration of the quasar spectra, which is usually established using exposures of a standard thorium-argon (ThAr) hollow-cathode emission-line lamp taken immediately before and/or after the quasar exposure. After the wavelength--pixel mapping is derived from the ThAr exposure, that solution is simply applied to the quasar exposure. Since the quasar and ThAr light illuminate the spectrograph slit differently and, in general, take slightly different paths through the spectrograph to the CCD, wavelength calibration errors may ensue.  \\citet{GriestK_09a} (hereafter \\citetalias{GriestK_09a}) recently observed a quasar absorption system using Keck/HIRES with an iodine gas absorption cell for more direct wavelength calibration, thereby allowing comparison with the ThAr wavelength scale. For varying-$\\alpha$ studies, the most important and robust result from \\citetalias{GriestK_09a} is the identification of distortions in the ThAr wavelength scale across each echelle order (hereafter `intra-order distortions'). See their figures 4 and 5. These may result from differential and variable vignetting of quasar light compared to ThAr light, where only the former enters HIRES from the telescope, the optical axes of which differ slightly, causing the quasar light path to rotate as the telescope tracks the quasar (\\citealt{SuzukiN_03a}; \\citetalias{GriestK_09a}). \\citetalias{GriestK_09a} also discuss velocity offsets between the ThAr and I$_2$ wavelength scales. If those results are robust, they are less important for varying $\\alpha$ studies because, as equation (\\ref{eq:da2}) makes clear, such velocity offsets will not directly influence a \\da\\ measurement\\footnote{This is not strictly true when many quasar exposures are combined, as is generally the case for most of the Keck/HIRES sample. If different velocity shifts apply to the different exposures of the same quasar, and if the relative weights of the exposures vary with wavelength when forming the final, combined spectrum, then small relative velocity shifts will be measured between transitions at different wavelengths. Nevertheless, the effect on \\da\\ of overall velocity shifts is of secondary importance to the intra-order distortions considered in detail here.}.  This paper aims to assess the impact these intra-order distortions may have had on the Keck/HIRES results for varying $\\alpha$. The following section describes the general effect the distortions will have and provides a crude calculation of their expected magnitude. Section \\ref{sec:calc} details a more refined calculation of the correction to \\da\\ for each absorber in the Keck/HIRES sample and Section \\ref{sec:res} discusses the results. In section \\ref{sec:model} we search for direct evidence for the need to apply the corrections and consider the effect of model errors on our calculations. We conclude in Section \\ref{sec:conc}. ",
        "conclusions": "\\label{sec:conc}  Exploring systematic effects which may explain the Keck/HIRES evidence for a varying $\\alpha$ is clearly an important problem. However, various properties of the Keck/HIRES results make the task difficult; the consistency between the average \\da\\ values in the low-$z$ Mg/Fe{\\sc \\,ii} and the more diverse high-$z$ systems being an important one. We have demonstrated here that if we model the intra-order distortions identified by \\citetalias{GriestK_09a} in a simple way, and extrapolate the model to all echelle orders (not just those within the I$_2$-cell calibration range) the effect on the overall evidence for varying $\\alpha$ is very small. As expected, the distortions affect individual values of \\da\\ randomly in sign and magnitude, because the differing redshifts place the transitions at varying positions with respect to echelle order edges where the distortions are worst. Indeed, the effect for a typical absorber is to shift \\da\\ by $0.4\\times10^{-5}$ which, compared to the median error on \\da, $\\sim$$1.9\\times10^{-5}$, is too small for us even to find direct evidence of the need to apply the corrections we calculate.  Despite the above conclusions, it is important to emphasise that we do not yet fully understand the origin of the intra-order distortions identified by \\citetalias{GriestK_09a}, how they depend on various observational parameters (e.g.~telescope pointing direction, temperature, time etc.) and therefore how they may differently affect spectra of different quasars in the Keck/HIRES sample. Indeed, \\citetalias{GriestK_09a} show that separate exposures taken through the I$_2$ cell seem to have somewhat different intra-order distortion patterns. And while we have made a simple attempt to address this problem of model errors, it is not enough to completely rule out intra-order distortions as an important systematic error for the Keck/HIRES results. If future Keck/HIRES observations allow much smaller statistical errors on \\da\\ in individual absorbers or small samples, intra-order distortions like those identified by \\citetalias{GriestK_09a} must either be eliminated or carefully modelled."
    },
    "0911/0911.1337_arXiv.txt": {
        "abstract": "Planets have been observed in tight binary systems with separations less than $20$ AU. A likely formation scenario for such systems involves a dynamical capture, after which high relative inclinations are likely and may lead to Kozai oscillations. We numerically investigate the fate of an initially coplanar double-planet system in a class of binaries with separation ranging between $12 - 20 $ AU.  Dynamical integrations of representative four-body systems are performed, each including a hot Jupiter and a second planet on a wider orbit. We find that,  although such systems can remain stable at low relative inclinations ($\\lesssim 40^\\circ$), high relative inclinations are likely to lead to instabilities. This can be avoided if the planets are placed in a  \\emph{Kozai-stable zone} within which mutual gravitational perturbations can suppress the Kozai mechanism. We investigate the possibility of inducing Kozai oscillations in the inner orbit by a weak coupling mechanism between the planets in which the coplanarity is broken due to a differential nodal precession.   Propagating perturbations from the stellar companion through a planetary system in this manner can have dramatic effects on the dynamical evolution of planetary systems, especially in tight binaries and can offer a reasonable explanation for eccentricity trends among planets observed in binary systems.  We find that inducing such oscillations into the orbit of a hot Jupiter is more likely in tight binaries and an upper limit can be set on the binary separation above which these oscillations are not observed. ",
        "introduction": "Among all extrasolar planets discovered to date, $20 \\%$  are in multiple stellar systems \\citep{des07,egg07b}. Binary stellar systems with circumstellar and circumbinary debris disks have been observed with an overall occurrence rate higher than that for debris disks around single stars \\citep{tri06}. Numerical simulations suggest that protoplanetary disks embedded in binary star systems should be able to form both terrestrial and giant planets \\citep{bar02,qui02,bos06}.  Most binary and multiple stellar systems found to be harboring planets are wide, with separations larger than $100$ AU \\citep{egg04,mug04}. However, several systems with separations as low as  $\\sim 20$ AU have been shown to harbor giant  planets. These are the binary systems HD196885, $\\gamma$ Cephei, and Gliese 86, and the higher order system HD41004 \\citep{cor08,hat03,els01,zuc04}.  There are two important things to note about these systems. First, since radial-velocity surveys  have always been in biased against close binaries \\citep{egg07b}, this sample probably underrepresents the frequency of planets in such binaries.  In particular,  the lower limit for the separation of binaries that can harbor planets may be smaller than $\\sim 20$ AU.  Second, such systems constitute a unique data set which can allow us to test theoretical models for planet formation and evolution. This is because the presence of a stellar companion at such proximity can greatly influence these processes. A stellar companion as far as $\\sim 50$ AU can weaken the chances for the formation of a giant planet by stirring, heating or truncating  protoplanetary disks \\citep{jan07,bos06,may05,kle01,nel00a,nel00b}.  Therefore, the fact that systems as tight as $20$ AU are harboring giant planets is puzzling and  does seem to require further investigation.  The most likely formation scenario for such systems may involve dynamical encounters  in dense stellar systems where the binary companion is captured or moved inward from a wider orbit after the planet formation has taken place. A stellar companion can also affect the dynamical evolution of a planetary system, and therefore its long term survival,  through secular perturbations.   In the case of a large relative inclination between the planetary orbit and that of the stellar companion, Kozai oscillations can take place where angular momentum exchange between the orbits leads to large-amplitude synchronous oscillations in the eccentricity and inclination of the planet \\citep{koz62,holm97,inn97,for00,tak05}. This has the potential to disrupt the system if it results in close encounters, orbital crossings, or strong planet-planet interactions for systems with more than one planet.  A renewed interest in planetary dynamics in general and in the stability of planets in binary systems particularly has been triggered by the discovery of planets in binary and multiple stellar systems \\citep{dvo03,bar04,ray05,ray06,riv07}. Looking into the dynamical history of planetary systems is key to understanding the processes involved in their formation and subsequent evolution.  In fact, this problem has received special interest from dynamicists even before extrasolar planets were discovered, with many studies applying general results of the three body problem to the special case of a binary stellar system harboring a planet \\citep[e.g,][hereafter HW99]{dvo93,inn97,hol99}. By performing numerical simulations on planetary orbits in binary systems, HW99 derived empirical expressions for the maximum semi-major axis of the planet as a function of both binary mass fraction and eccentricity, above which the system would become unstable. Their study considered the stability of binary systems harboring planets in the coplanar case only. Numerical simulations by \\citet{dav03} of Earth-like planets in binary systems with different initial configurations led to the estimate that 50\\% of binary systems allow an Earth like planet to remain stable for 4.6 Gyr. The stability of Earth-like planets in binaries was also studied by \\citet{fat06}.  With numerical simulations they calculated the survival time of an Earth-like planet orbiting a Sun-like star in the presence of a stellar companion. They explored a range of binary configurations including relative inclinations between the binary orbit and the orbit of the planet. \\citet{hag06} performed numerical simulations to model the binary system $\\gamma$ Cephei. This is a tight binary with separation $18.5$ AU and an eccentricity of $0.36$.  A jovian planet with a minimum mass of $1.7 M_{J}$ was shown to be orbiting the primary of this system at $2.1$ AU \\citep{hat03}. The question of the presence of an Earth-like planet within the habitable zone (HZ) of the primary was addressed and a range of semi-major axes and eccentricities of the binary as well as orbital inclination of the Jupiter-like planet was explored.   This study concludes that, within timescales in the range  $10-100$ Myr,  such a planet can exist on a stable orbit but not within the HZ of the primary.  In this study, we investigate the stability of double-planet systems in tight binaries. The tightest stellar system in which a planet was ever claimed is HD188753.  This is a hierarchical triple star system in which the primary is  a Sun-like star with mass $1.06$ $M_\\odot$. The binary  companion is a double-star system with total mass $1.63$ $M_\\odot$ orbiting the primary at separation  $\\sim{12.3}$ AU  and eccentricity 0.5 \\citep{gri77}.  A hot Jupiter of mass $1.14$ $M_{J}$ was claimed to be orbiting the primary star with a period of 3.35 d  \\citep{kon05}. Because of the proximity of the companion, the presence of a planet in this system is suggestive of a dynamical history in which the binary companion was either moved inward to a tighter orbit or captured  after the claimed planet had already formed and reached its current position \\citep{por05,pfa05,bos06,jan07}. Although the observation of this planet was refuted by  \\citet{egg07a},  we find this system to be a good representative of the lower limit (in separation) of a class of tight binary systems harboring giant planets.  We test in this study the capacity of such systems to harbor two planets including a hot Jupiter. This class of binaries with separation $\\lesssim 20$ AU is of special interest since it is likely that such systems may have had active dynamical histories and are expected to induce large perturbative forces to planetary systems with dramatic results.  Studying such systems  in the presence of more than one planet, can uncover the role played by binary companions in shaping planetary systems. This  is especially important since multi-planet systems seem to be common in nature \\citep{fis03}. The dynamics of double-planet systems in wide binaries was studied analytically by \\citet{tak08}, where three different dynamical classes were identified as possible outcomes: (1) Decoupled systems in which planetary orbits experience independent Kozai cycles due to their weak mutual interaction compared to perturbations from the companion; (2) Weakly coupled systems in which mutual gravitational interactions become more dominant and relative nodal precession can cause mutual inclinations to grow and induce Kozai oscillations in the inner orbit; (3) Dynamically rigid systems in which the orbital elements of both planets oscillate in concert. These possible outcomes depend on the initial conditions and therefore the relative strength of the different perturbative forces. Here we expand on this study by looking into systems with closer companions and test the possibility of inducing eccentricity oscillations in a hot Jupiter orbit through a weakly coupled system.  Since planets with orbital radii $a \\lesssim 0.1$ AU are expected to be tidally circularized, the recent observations of a few hot Jupiters on considerably eccentric orbits (such as XO-3, HAT-P-2, and HD185269 \\citep{joh08,loe08,joh06}) does raise questions.  Why are the orbits of these planets not circularized when the orbits of most hot Jupiters are?  \\citet{mat08} show that close-in planets on eccentric orbits can be explained by constraining their tidal $Q$ factors under the assumption that these orbits are in the process of being circularized.   They also explore the possibility that these eccentricities are induced by a perturber on a wider orbit for the system GJ436 and show that such a planet, if present, would cause a radial-velocity amplitude above the current detection limit, and therefore, would have been observed.  Although such eccentricities can, in principle, be explained by the presence of a companion, whether stellar or sub-stellar, one has to keep in mind that a planet on a close-in orbit experiences strong GR forces. Therefore, a stellar companion as close as $12 - 20$ AU cannot induce such eccentricities directly into a hot Jupiter orbit. We test here the possibility of inducing such eccentricities to a hot Jupiter through a second planet in the presence of a binary companion. We  also investigate the dynamics and stability criteria of a double-planet system orbiting the primary of a tight binary system with separation in the range $ a_{b} = 12 - 20$ AU. All systems we investigate in this study contain a hot Jupiter of mass $1 M_{J}$ and separation $a_{1} = 0.05$ AU.  The stability region for a second planet, placed on a wider orbit, that is initially circular and coplanar  with a variable semi-major axis $a_{2}$, was explored for two different binary configurations.   In the first, the binary companion orbits at  $ a_{b} = 12$ AU  with an eccentricity $ e_{b} = 0.5$, and in the second, $ a_{b} = 20$ AU and $ e_{b} = 0$.  For the tightest binary separation  ($ a_{b} = 12$ AU), two different planetary masses $m_{2}=M_{J}$  and $m_{2} =M_{\\oplus}$ were investigated. In this class of systems, planets experience mutual gravitational perturbations in addition to perturbations from the stellar companion. A hot Jupiter orbit is strongly dominated by GR forces. Because of the tightness of these systems, the interplay between these different forces is expected to have a dramatic effect on the planetary system.  Throughout this paper, the subscripts 0, 1, 2, and b will be used to refer to the primary star, the inner planet, the outer planet, and the binary companion, respectively. We use a semi-analytical approach to predict the stability region of the outer planet $m_{2}$. Our numerical simulations incorporate a wide range of parameters including mutual inclinations of the orbits and perturbative forces to the Keplerian motion of each planet including planet-planet interactions and GR effects. Such numerical simulations can reveal to us the likelihood of a compact stellar system to maintain multiple planets and can be a useful guide for future observational efforts such as HARPS and the Kepler Mission.   An overview of secular perturbations affecting this class of systems is presented in \\S 2 . In \\S 3 we discuss the analytical background for our predicted stability region for the outer planet in a stellar system resembling HD188753.  The numerical techniques used in this study are presented in \\S 4, results and discussion in \\S 5, and finally,  \\S 6 summarizes our conclusions and suggestions for future work.\\\\ ",
        "conclusions": "We test the stability region within a tight binary for a double-planet system composed of a Jupiter mass planet at $0.05$ AU and a second planet on a wider orbit. Two binary separations are investigated, $12$ AU and $20$ AU.   The stability region is tested for a Jupiter mass planet in the outer orbit for both binary separations.  In the case of  $12$ AU separation, the stability region for a potentially habitable  planet is also tested where $m_{2}=M_{\\oplus}$.  \\subsection{ A $12$ AU Binary: The Case of HD188753}  \\subsubsection{Equal-Mass Planets} The stability region of the double-planet system described above in a binary system resembling HD188753 is investigated. Figure 5 summarizes  the results of all our simulations of this system spanning the whole expected region of stability for a second planet of Jupiter mass. It shows that as long as the orbits are highly inclined ($\\gtrsim 40^\\circ$), Kozai oscillations will always lead to instabilities in this system. \\begin{figure} \\epsscale{0.85} \\plotone{f5.eps} \\caption{The stability region for a second planet in HD188753 of mass $m_{2}=m_{1} = M_{J}$. Empty triangles represent stable systems while filled triangles represent systems which experience instabilities. All systems with the binary orbital plane highly inclined relative to the planetary plane are unstable, except for systems within the Kozai-stable zone.} \\end{figure} For small relative inclinations $i_{init} \\leq 40^\\circ$, the simulations give a \\emph{lower limit} for the semi-major axis of this planet in the range $0.063$ AU $< a_{2, \\rm min} < 0.066$ AU, which is consistent with the G93 predicted value of $0.065$ AU. The \\emph{upper limit} lies in the range $1.3$ AU $ < a_{2, \\rm max} < 1.4$ AU for the coplanar case, larger than the value predicted by HW99.  These results are obtained using full 4-body simulations which include planetary interactions with all other particles in the system and therefore, are expected to set more realistic limits on stability boundaries than do studies that treat planets as massless point particles.  Our results also confirm the existence of a \\emph{Kozai-stable zone} within which stability is maintained even in the presence of large relative inclinations. This is seen in the region  $a_{2} = 0.1 - 0.3$ AU.  At $a_{2} = 0.3$ AU, we can see a transitional region between stability and instability as the system approaches the edge of the Kozai-stable zone. We compare in Figure 6 the time evolution of both planets in two different cases. \\begin{figure} \\epsscale{0.85} \\plotone{f6.eps} \\caption{ Orbital elements of both planets (inner (dotted) and outer (solid)) as they evolve with time for the following systems: - Left panel:  ($m_{1} = m_{2}= M_{J}$),  $a_{1} = 0.05 $ AU, $a_{2} = 0.2 $ AU (within the \\emph{Kozai-stable zone}), and $i_{init}  = 70^{\\circ}$. - Right panel: Same masses,  $a_{2} = 1.0$ AU , and $i_{init}  = 50^{\\circ}$. The figure shows the distance at pericenter  ($ r_{p}$) and the eccentricity ($e$) of the planets. } \\end{figure} In the first the outer planet is within the Kozai-stable zone at $a_{2} = 0.2 $ AU. Although its orbit is highly inclined relative to the orbit of the companion, we see stable orbits with no sign of growth. The second case represents the outer planet outside the Kozai-stable zone at $a_{2} = 1.0 $ AU , where mutual gravitational perturbations become too weak to suppress Kozai oscillations. This case leads to the ejection of the planet within less than $80,000$ years.  Figure 5 also shows a region of the explored space  ($0.35 < a_{2} < 0.8$ AU, and $i_{ini} \\geq 40^{\\circ}$) within which the planets are weakly coupled and therefore Kozai oscillations are induced in the inner orbit by interaction with the outer planet.  Large eccentricities are gained during these oscillations which eventually bring the inner planet to the surface of the star.  We discuss this effect in detail in \\S 5.3.  \\subsubsection{A Terrestrial Planet}  There has been great interest lately in the question of habitability of exoplanets \\citep[e.g.,][]{dav03,jia05,fat06,hag06,hag07}.  Although observations of Earth mass planets within the Habitable Zones (HZ) of Sun-like stars may not become reality in the very near future, numerical simulations to understand the dynamics and stability limits of these planets in different environments is necessary to guide future efforts.  It has been shown that  Earth-like planets can survive the migration of  a giant planet  \\citep{man07, raym06}. Therefore, we look into the stability of  a double-planet system with a habitable planet  and a hot Jupiter in HD188753.   The central star in this system with mass $1.06$  $M_{\\odot}$ is of G-type with a habitable zone (HZ) between $0.95$ AU and $1.37$ AU \\citep{kas93}. Therefore, our predicted stability region (Fig. 5) does leave space for a habitable planet in this system, but only for small values of the initial relative inclinations of the orbits.  Figure 7 shows the results of our simulations performed for a second planet of mass $m_{2}=M_{\\oplus}$ in HD188753.  Since an Earth mass planet is three orders of magnitude less massive than a Jupiter mass, perturbations from the companion are expected to become less dominant leaving more space for stability. Therefore, a maximum integration time of  $6.6 \\times 10^{6}$ yr was used for all systems shown to account for any delayed instabilities. \\begin{figure} \\epsscale{0.85} \\plotone{f7.eps} \\caption{The results of all numerical simulations performed for a system similar to HD188753 with a second planet of mass $ = 1 M_{\\oplus}$ . Empty triangles represent stable systems while filled triangles represent unstable systems. The stable region includes the HZ of this star, but only for low relative inclinations among the orbits. } \\end{figure} Our results show that the upper limit for stability of an Earth mass planet at low relative inclinations ($i_{init} \\leq 40^\\circ$) is larger than that for a Jupiter mass planet  and lies in the range $1.4$ AU $ < (a_{2})_{max} < 1.5$ AU, confirming the dependence of perturbations, and therefore stability limits,  on the planetary-masses. The HZ of the central star in HD188753 lies completely within the stable region at low inclinations.  On the other hand, the \\emph{Kozai-stable zone} is located outside the HZ, leaving no space for a habitable planet when high relative inclinations are present. At the lower limit of $a_{2}$,   although an Earth mass planet can survive down to a region between $0.06$ AU and $0.063$ AU such a planet will be extremely hot and inhabitable.  In addition to dynamical stability, an important condition for habitability is climate stability which can be comprimised if the planet's orbit becomes eccentric.  This has the potential to drive the planet's orbit outside the HZ at apastron or/and periastron, or may lead to extreme heat variations on the surface of the planet. We show in Figure 8 the orbital elements of an Earth-mass planet in this system with low relative inclination and initially on a circular orbit at $a_{2} \\sim 1.0 $ AU. It experiences mild eccentricity oscillations with an amplitude of $e \\sim 0.1$, with no sign of growth. \\begin{figure} \\epsscale{0.85} \\plotone{f8.eps} \\caption{Orbital elements of the outer planet in a double planet system in HD188753 with the outer planet in the Habitable Zone of the primary. System elements are : $a_{1}=0.05$ AU; $a_{2}=1.0$ AU; $m_{1}=M_{J}$;  $m_{2} =M_{\\oplus}$ and $i_{init}  = 0^{\\circ}$ . The figure shows from top to bottom:  semi-major axis, eccentricity and argument of pericenter of the outer planet. } \\end{figure} Still, habitability  of this planet can become questionable since such an eccentricity can drive the planet near the edges of the HZ at periastron and apastron.  Therefore, the presence of  a perturber to the Keplerian motion of a potentially habitable planet can alter the positions of the outer and the inner edges of the HZ through their effect on the orbit.  In this case, the outer edge of the HZ is expected to move inward to  $ a_{2} \\sim 1.24$ AU from the primary while the inner edge move from $a_{2}\\sim 0.95$ AU to $\\sim 1.0$ AU, altering the width of the HZ by $\\sim 40\\%$.   \\subsection{ A $20$ AU Equal mass Binary} It has been shown by \\citet{jan07} that a stellar system as tight as HD188753  is unlikely to support a disk massive enough to form a Jovian planet \\emph{in situ}. Therefore,  the presence of a planet in such a system would be indicative of  a dynamical history.  On the contrary, the truncated disk around a binary with slightly larger separation such as $\\gamma$ Cephei (separation $20$ AU and eccentricity $0.4$) has been shown  \\citep{jan08} to have sufficient  mass to form such a planet. Such wider binaries may host giant planets and still have no active dynamical histories, and therefore, maintain  dynamical stability with multiple planets which are coplanar with the binary orbit.  Therefore, a  binary with separation $20$ AU, represents an upper limit in this study to a class of binaries with  low separation in which the presence of a hot Jupiter is questionable.  We show in this section, our results for the stability region of a double planet system in a binary with two stars of solar mass orbiting at $20$ AU.  The planetary system used in this test is the same system used in \\S5.1.1 in which two coplanar, equal mass planets of Jupiter size are place on circular orbits.  The inner planet is a hot Jupiter at $0.05$ AU, while the semi-major axis of the second planet, $a_{2}$, is varied along the expected range. The stability region is also tested here in the two dimensional space [$i_{init}, a_{2}$] where $i_{init}$ is the initial relative inclination between the binary orbit and the common planetary plane.  We apply the same criteria as in \\S 3 to predict the stability region for the outer planet when the inner orbit is fixed. These results are shown in Figure 9 where one can see that the stability region for this system is similar to that of HD188753 in its structure but varies in scale.  The larger binary separation has allowed more space for the outer planet to survive.  In some cases the outer planet can experience Kozai oscillations (with moderate eccentricities) for as long as $1$ Myr without disrupting the system as long as it does not reach the upper limit for stability at apastron. The Kozai stable zone expands with the binary separation in general since mutual planetary interactions become more dominant relative to the Kozai effect.  In this system the edge of the Kozai stable zone is at $a_{2} = 0.5$ AU. We also see that the two planets are placed in the \"weakly coupled\" dynamical class in a small region of the shown parameter space  ($ 0.5 \\gtrsim a_{2} \\lesssim 0.8$ AU and $i_{init} = 50^\\circ$) within which the inner planet experiences Kozai oscillations induced by the outer planet. The width of this region is narrower than that observed in the case $a_{b} = 12$ AU  (see \\S 5.3 for a more detailed discussion). \\begin{figure} \\epsscale{0.85} \\plotone{f9.eps} \\caption{The stability region for a double planet system in a binary with parameters: Two stars of Solar mass, $a_{b} = 20$ AU, and $e_{b} =0$.   Planetary masses are $m_{2}=m_{1} = M_{J}$ initially on circular coplanar orbits. The inner planet is a hot Jupiter at $a_{2} = 0.05$ AU. Empty triangles represent stable systems, while filled triangles represent the unstable systems.} \\end{figure}    \\subsection{Induced Eccentricity Oscillations in a hot Jupiter}  We test the possibility of inducing Kozai oscillations in the hot Jupiter orbit in our systems by propagating perturbations from the stellar companion through a second planet of equal mass.  We find that  inducing large eccentricities in the hot Jupiter orbit is more likely in binaries with smaller separation.  In the case of a $12$ AU binary separation this effect is observed within the region of the explored phase space defined by  ($0.35 \\leq a_{2} \\geq 0.8$ AU, and $i_{ini} \\geq 40^{\\circ}$). This behavior is not observed within the  \\emph{Kozai-stable zone}   ($a_{2} < 0.3$) AU, and shows its strongest behavior in the region  ($0.4$ AU $\\leq a_{2} \\leq 0.5$ AU).  As the distance between the planets increases, this effect gets weaker until it disappears beyond $a_{2} > 0.8$ AU.   Figure 10 shows the time evolution of the orbital elements of a system representative of this behavior with the outer planet at $a_{2}=0.5$ AU and $i_{init}=50^\\circ$. \\begin{figure} \\epsscale{0.85} \\plotone{f10.eps} \\caption{A case representative of the systems which experience induced eccentricity in the inner orbit, showing orbital elements of both planets, inner (dotted line) and outer (solid line). The initial system parameters are: $m_{1} = m_{2}= M_{J}$,  $a_{1} = 0.05 $ AU, $a_{2} = 0.5 $ AU, $e_{1}=e_{2}=0$, and $i_{init}=50^\\circ$. The outer planet is experiencing Kozai oscillations induced by the stellar companion. Due to fast nodal precession of the outer planet, the mutual inclination between the planets $I_{12}$ grows, and oscillates with the same period as $\\Omega_{2}$.   The amplitude of these oscillations grows to a value twice the common initial inclination relative to the binary plane,  $i_{init}$, causing the inner planet's eccentricity to grow chaotically. } \\end{figure} The outer planet is experiencing Kozai oscillations induced by the stellar companion. The nodal precession of this planet creates a split in the planetary orbits and the mutual inclination between the planets, $I_{12}$, grows and oscillates with the same period as $\\Omega_{2}$.   The amplitude of these oscillations is twice the common initial inclination relative to the binary plane $i_{init}$.  Angular momentum exchanged between the planetary orbits, propagating the angular momentum transfer from the binary leads to mutual planetary Kozai oscillations.  Since the angular momentum of the perturber ($m_{2}$ in this case) is changing, the eccentricity oscillations induced in the inner orbit lead to chaotic growth. As the planet approaches the central star, it should experience strong tidal dissipation  and may be circularized into a tighter orbit. Our simulations do not include tidal effects and therefore, we cannot predict the final outcome of these events. In general,  tidal torques are not significant in the evolution of a circular orbit at $\\sim 0.05$ AU where GR is the most dominant force.  Once a planet gains significant eccentricities driving it very close to the surface of the star at periastron, tidal forces may become a dominant factor but do act on longer timescales when compared to Kozai precession timescales.  To demonstrate the interplay between the different forces acting on the inner planet, we show in Figure 11 \\begin{figure} \\epsscale{0.85} \\plotone{f11.eps} \\caption{ Eccentricity of both planets (outer in (solid) and inner in (dotted)) and mutual inclination between planets $I_{12}$ shown for the same system in Figure 10, but with GR forces absent in the top two panels and and $m_{2}=M_{\\oplus}$ in the lower two panels.  } \\end{figure} the relative inclination between the planetary orbits and the eccentricities of both planets in a system similar to that shown in Figure 10, but once with the GR forces removed from the calculation (top two panels) and second using an Earth-mass planet in the outer orbit (lower two panels).  The former shows the faster growth of eccentricity when compared to Figure 10 and therefore demonstrates the goal played by GR forces in stabilizing the inner orbit against other perturbations, while the later confirms the dependency of this mutual planetary Kozai effect on the planetary masses.  Inducing such oscillations into the orbit of a hot Jupiter requires specific initial conditions to be satisfied (see \\S 2.4).  We find that the region in parameter space that allows these conditions to be satisfied is narrower for wider binaries. We show in Figure 12 the range in semi-major axis of the outer planet that will allow induced Kozai oscillations in the inner orbit as a function of binary separation.   This is shown for an inner planet on an initially circular orbit at $a_{1}=0.05$ AU, and a binary orbit with inclination $i_{init} = 50^{\\circ}$ relative to the common planetary plane. All binaries in this figure share the following parameters ($m_{0} = m_{b}= M_{\\odot}, m_{1}=m_{2}=M_{J}$ , and $e_{b} = 0$). \\begin{figure} \\epsscale{0.85} \\plotone{f12.eps} \\caption{ The region in parameter space which allows Kozai oscillations to be induced in a hot Jupiter orbiting at $0.05$ AU.  On the x-axis, $a_{b}$ represents the variable binary separation. Binary parameters for all binaries are:  $m_{0} = m_{b}= M_{\\odot}, m_{1}=m_{2}=M_{J},  e_{b} = 0$, and $i_{init} = 50^{\\circ}$.  The variable semi-major axis of the outer planet $a_{2} $ is plotted on the y-axis.  In order to overcome GR precession and induce Kozai oscillations in the inner orbit, the outer planet must lie below the line $a_{2} = 0.8$ AU. The dashed line represents the limit of the \\emph{ Kozai-stable zone} as a function of binary separation.  Inside the dashed region between the two limits,  both planets experience Kozai oscillations (outer planet induced by the binary and the inner planet induced by the outer planet). All circles represent results of numerical simulations.  The large filled circles are systems in which the outer planet experiences Kozai oscillations and always lead to eccentric hot Jupiters. Small filled circles are systems in which the outer planet is shielded from the  Kozai effect, but in  less than $50\\%$ of the simulations we observe induced eccentricities to the hot Jupiter.  Empty circles represent systems which did not show any sign of this effect.} \\end{figure} These results show that tighter binaries are more likely to produce this effect and that  a cut-off binary separation $a_{b, \\rm cut}$ exists beyond which it is not observed.  For a hot Jupiter of mass $M_{J}$ and semi-major axis $a_{2} =0.05$ AU, this cut-off binary separation lies in the range $ 30 < a_{b, \\rm cut} < 40$ AU. As the inner planet's orbit is made wider, mutual planetary interactions grow larger which should allow a wider binary to induce eccentricities through weak coupling between the planets into the inner orbit."
    },
    "0911/0911.3332_arXiv.txt": {
        "abstract": "We present a detailed analysis of the Astrophysical Research Consortium 3.5m telescope spectrum of QSO SDSS~J0838+2955.  The object shows three broad absorption line (BAL) systems at 22,000, 13,000, and 4900 \\kms\\ blueshifted from the systemic redshift of z=2.043.  Of particular interest is the lowest velocity system that displays absorption from low-ionization species such as \\mgii, \\alii, \\siII, \\siII*, \\feii\\ and \\feii*.  Accurate column densities were measured for all transitions in this lowest velocity BAL using an inhomogeneous absorber model. The ratio of column densities of \\siII* and \\feii* with respect to their ground states gave an electron number density of log $n_e$ (cm$^{-3}$) = 3.75 $\\pm$ 0.22 for the outflow.  Photoionization modeling with careful regards to chemical abundances and the incident spectral energy distribution predicts an ionization parameter of log $U_H$ = -1.93 $\\pm$ 0.21 and a hydrogen column density of log $N_H$ (cm$^{-2}$) = 20.80 $\\pm$ 0.28.  This places the outflow at 3.3$^{+1.5}_{-1.0}$ kpc from the central active galactic nucleus (AGN).  Assuming that the fraction of solid angle subtended by the outflow is 0.2, these values yield a kinetic luminosity of $(4.5^{+3.1}_{-1.8})\\times10^{45}$ erg~s$^{-1}$, which is (1.4$^{+1.1}_{-0.6}$)\\% the bolometric luminosity of the QSO itself. Such large kinetic luminosity suggests that QSO outflows are a major contributor to AGN feedback mechanisms. ",
        "introduction": "\\label{sec:introduction}  Quasar outflows have been increasingly invoked as a primary feedback mechanism to explain the formation and evolution of supermassive black holes, their host galaxies, the surrounding intergalactic medium (IGM), and cluster cooling flows \\citep{Elvis2006}.  Several of these feedback mechanisms require that the kinetic luminosity ($\\dot{E}_k$) of the outflowing material would be $\\sim$5\\% the bolometric luminosity of the QSO itself assuming the active galactic nucleus (AGN) is accreting near the Eddington limit \\citep{Scannapieco2004}. The kinetic luminosity of absorption outflows (observed as blueshifted troughs in the quasar spectrum) is directly proportional to the distance $R$ of the outflow from the central AGN, the total hydrogen column density $N_H$, and the fraction $\\Omega$ of the solid angle subtended by the wind (see Section 5).  These parameters must therefore be accurately determined in order to infer the significance of such outflows for AGN feedback.  Previous attempts to measure $\\dot{E}$$_k$ \\citep{Wampler1995, deKool2001, deKool2002, deKool2002b, Hamann2001} had large statistical uncertainties \\citep[typically more than an order of magnitude, see ][]{Dunn2009} and suffered from unquantified systematic errors due to the use of inadequate absorption models. The latter lead to unreliable measurements of column densities in the observed troughs ($N_{ion}$), which are crucial for determining almost every physical aspect of the outflows. To determine $N_H$ we use photoionization modeling to find the ionization equilibrium that can reproduce the measured $N_{ion}$. To find $R$, in addition to the ionization equilibrium, we also need the number density of the gas which can be determined by comparing the ionic column densities of excited metastable states to their ground states \\citep{Wampler1995,deKool2001,Korista2008}.  We assume $\\Omega$ = 0.2, as is the case for broad absorption line (BAL) quasars in general \\citep{Hewett2003}.  We caution that $\\Omega$ is not directly constrained by our observations, and therefore express the kinetic luminosity in terms of $\\Omega_{0.2}$ (see Section~5).  In order to determine accurate  $\\dot{E}_k$ in quasar outflows, we launched a multistage research program. We first developed analysis techniques that allow us to determine reliable trough column densities \\citep[][and references therein]{Arav2008}. We then performed an exhaustive search through $\\sim$50,000 QSOs with redshift $z>$0.25 and Sloan Digital Sky Survey (SDSS) magnitude $r<$19.3 in the SDSS Data Release 6 \\citep{Adelman2008}; $\\sim$100 objects were found that exhibited absorption systems containing transitions from excited metastable states of \\hei*, \\siII*, \\feii*, \\feiii* and/or \\niII*.  Of these objects, approximately one third exhibited only \\hei* (E=159,856 cm$^{-1}$) without transitions from other energy levels nor other ions in metastable states that could serve as density diagnostics.  Of the remaining $\\sim$70 QSOs, $\\sim$30 objects displayed the intrinsic absorption system containing the excited states at velocities $\\gtrsim$1,000 \\kms\\ with respect to the systemic restframe of the QSO and continuum flux levels $\\gtrsim$10$^{-16}$ \\ergsAcm\\ in the observed frame.  Here we investigate one of the brightest objects from this list with wide enough absorption troughs to allow sufficiently accurate extraction of column densities with a moderate resolution spectrum.  QSO SDSS~J083817.00+295526.5 (hereafter, SDSS~J0838+2955) exhibits three systems of BALs with the lowest velocity system at -4,900 \\kms\\ containing absorption from \\siII\\ and \\siII* $\\lambda\\lambda$1527,1533 in the SDSS spectrum.  With an SDSS r magnitude of 17.85 and a redshift z=2.043, this object was ideal for follow-up observations because the optical portion of the spectrum bracketed both high ionization species of \\civ\\ and \\siIV\\ as well as the low ionization species of \\cii, \\mgii, \\alii, \\aliii, \\siII, and \\feii.  The plan of this paper is as follows.  In Section~\\ref{observations}, we detail the observations and data reduction of SDSS~J0838+2955.  Gas column and number densities are measured in Section~\\ref{column} after determining the inhomogeneity of the absorber across the continuum source.  Photoionization modeling of the outflow with careful regards to the incident spectral energy distribution (SED) is presented in Section~\\ref{photoion}.  We discuss the results for distance and kinetic luminosity in Section~\\ref{results}.  For this analysis, we adopt a standard $\\Lambda$CDM cosmology with H$_0$=70 \\kms\\ Mpc$^{-1}$, $\\Omega_m$=0.3, and $\\Omega_{\\Lambda}$=0.7 \\citep{Spergel2003}.  We list all transitions by their rounded vacuum wavelengths in the text, but use the unrounded vacuum wavelengths during analysis. ",
        "conclusions": "\\label{results}  With estimates of $n_H$, $N_H$, $U_H$, and $Q$ in hand, the large-scale physical properties of the outflow can be calculated.  From the definition of the ionization parameter \\citep[][Equation (14.7)]{Osterbrock2006}, the distance of the outflow from the central AGN is  \\begin{equation} R = \\left( {{Q} \\over {4 \\pi c U_H n_H}}\\right)^{{1} \\over {2}} = 3.3^{+1.5}_{-1.0}~ \\mathrm{ kpc}. \\end{equation}  \\noindent This large distance with a velocity v=4900 \\kms\\ suggests that this outflow will continue to expand into the IGM, where its mass and kinetic energy will be dissipated as heat into the intracluster medium.  Assuming the outflow is in the form of a thin partial shell, its mass is  \\begin{equation} M = 4 \\pi \\mu m_p \\Omega R^2 N_H = (1.9^{+1.8}_{-0.9}) \\times 10^8~\\Omega_{0.2}~ \\mathrm{\\Msun} \\end{equation}  \\noindent where $\\mu \\approx$ 1.43 is the mean atomic mass for solar abundances, $m_p$ is the mass of the proton, and $\\Omega$ is the portion of the solid angle subtended by the shell. The adopted parameter space for $N_H$ and $U_H$ is used to determine the expected uncertainty in the mass.  A full discussion about what $\\Omega$ value is appropriate for these outflows is given in \\citet{Dunn2009}.  To summarize their argument: there are good reasons to assume that the $\\Omega$ of these low-ionization BAL outflows with excited metastable states is similar to the canonical value for \\civ\\ BALs \\citep[$\\Omega \\sim$ 0.2; see][]{Hewett2003}, and that their scarcity is due to the selection effect that low-ionization BALs are rare in general. See also \\citet{Hall2003} for further discussion regarding this point.  In order to accommodate for the uncertainty in $\\Omega$, we leave the final answer as a function of $\\Omega_{0.2} \\equiv \\Omega / 0.2$. The mass flux is then given by  \\begin{equation} \\dot{M} = 8 \\pi \\mu m_p \\Omega R N_H v = 590^{+410}_{-240}~ \\Omega_{0.2}~ \\mathrm{\\Msun\\ yr}^{-1} \\end{equation}  \\noindent where the adopted parameter space for $N_H$ and $U_H$ is used.  If the solution to $N_H$ and $U_H$ were assumed to be uncorrelated, then the error in  $\\dot{M}$ due to photoionization modeling would be 0.30 dex instead of the actual 0.19 dex.  Finally, the kinetic luminosity of the outflow is  \\begin{equation} \\dot{E}_k = \\frac{\\dot{M} v^2}{2} = 4 \\pi \\mu m_p \\Omega R N_H v^3 = (4.5^{+3.1}_{-1.8}) \\times 10^{45}~ \\Omega_{0.2}~ \\mathrm{erg~s}^{-1} \\end{equation}  \\noindent which is (1.4$^{+1.1}_{-0.6}$ $\\Omega_{0.2}$)\\% the bolometric luminosity of the QSO itself.  The kinetic luminosity of this component is sufficiently large to affect its host environment and contribute to AGN feedback mechanisms.  Indeed, this value of $\\sim$1.4\\% is more likely a lower limit considering that the other outflow components $a$ and $b$ could further contribute to AGN feedback \\citep[see][]{Arav2009}. Moreover, in Seyfert galaxies the hot phase of the outflow seen in X-rays can carry 70\\%-99\\% of the kinetic luminosity of the outflow \\citep{Gabel2005b,Arav2007}.  We point out that a few earlier investigations have found outflows with similar physical properties.  Based on the \\citet{Hamann2001} analysis of QSO 3C~191,  the kinetic luminosity of the outflow observed in this object is $\\sim$9$ \\times 10^{43}~ \\Omega_{0.2}$ erg s$^{-1}$ and its mass flux is $\\sim$310$~ \\Omega_{0.2}$ \\Msun\\ yr$^{-1}$  for an outflow situated at $\\sim$28 kpc.  In QSO~FBQS~1044+3656, \\citet{deKool2001} found the outflow to be $\\sim$1 kpc from the AGN with a kinetic luminosity of $\\sim$9$ \\times 10^{44}~ \\Omega_{0.2}$ erg s$^{-1}$ and mass flux of $\\sim$150$~ \\Omega_{0.2}$ \\Msun\\ yr$^{-1}$.  Given the bolometric luminosity and assuming that the black hole is accreting near the Eddington limit, we find that the amount of mass being accreted is \\citep{Salpeter1964}:  \\begin{equation} \\dot{M}_{acc} = {{L_{Bol}} \\over {\\epsilon_r c^2}} \\approx 60~\\mathrm{\\Msun\\ yr}^{-1} \\end{equation}  \\noindent assuming an accretion efficiency $\\epsilon_r$ of 0.1. The mass flux of the outflow is $\\sim$10 times greater than the accretion rate, which is physically plausible assuming the outflow entrained a substantial amount of material on its journey to $\\sim$3 kpc away from the AGN.  Furthermore, SDSS~J0838+2955 is an example of a QSO with outflows at different distance scales considering that the outflow contributing to component $a$ is most likely much closer than component $c$.  The argument is as follows.  The significant variation in optical depth of component $a$ is most likely due to a change in flux of the ionizing source or transverse velocity of the outflow out of the line of site to the central AGN \\citep{Barlow1994}.  Although the spectrum itself cannot be used to determine variations in luminosity (because the ARC spectrum was flux calibrated to the SDSS spectrum), several slit guider images were used to conduct differential photometry between the target and foreground stars of similar color in the field.  The change in $r$ magnitude (the color corresponding to the peak efficiency of the guider CCD) was determined to be less than 0.09 mag at the 2$\\sigma$ level and thus the variation in absorption of the outflow was most likely due to transverse motion.  An upper limit on the distance of this component $a$ of the outflow can be estimated assuming that the transverse velocity cannot exceed the Keplerian velocity.  The Keplerian velocity is determined from the mass of the black hole $M_{BH}$ which in turn is estimated from a mass-luminosity relation.   In more detail, the luminosity of the QSO is expected to be $\\nu L_{\\nu} \\approx 3 \\times 10^{46}$ erg s$^{-1}$ near the restframe wavelength of 5100 \\AA.  From Equation (9) in \\citet{Peterson2004} and assuming L/L$_{Eddington} \\approx$0.1, the mass of the central black hole is expected to be $M_{BH} \\approx 7 \\times 10^{9}$ \\Msun, although the mass--luminosity relation is not well constrained in this high-mass regime.  From this black hole mass, the radius of the UV continuum region arising from the accretion disk can be estimated to be $\\sim$30 light-days \\citep[Equation (3.20) in][]{Peterson1997} assuming the accretion rate is near the Eddington limit.  We will choose a lower limit on the line of sight covering factor of component $a$ to be $\\sim$0.5 because the \\civ\\ absorption trough had an initial residual intensity of $\\sim$0.43. The transverse velocity of the outflow across half of the $\\sim$60 light-day wide continuum source in the restframe time interval of 4.3 yrs / 1+z = 1.4 yrs must be v $\\gtrsim$ 30 light-days / 1.4 yrs $\\approx$ 18,000 \\kms\\ in the restframe of the QSO.  Finally, assuming the tangential velocity of the outflow does not exceed the Keplerian velocity, then the distance of outflow $a$ from the central AGN must be R $<$ $G M_{BH}/v^2 \\approx$  0.1 pc."
    },
    "0911/0911.3042_arXiv.txt": {
        "abstract": "The cause of enhanced  acoustic power surrounding active regions, the acoustic halo, is not as yet understood. We explore the properties of the enhanced acoustic power observed near disk center from 21 to 27 January 2002, including AR 9787. We find that (i) there exists a strong correlation of the enhanced high frequency power with magnetic-field inclination, with greater power in more horizontal fields, (ii) the frequency of  the maximum enhancement increases along with magnetic field strength, and (iii) the oscillations contributing to the halos show modal ridges which are shifted to higher wavenumber at constant frequency in comparison to the ridges of modes in the quiet-Sun. ",
        "introduction": "\\label{introduction}  The application of local helioseimology to probe the nature of solar active regions is critically important in understanding their origin and evolution. The influence of magnetic fields on incident acoustic waves is complex, but is coming under increased scrutiny due to advances in numerical simulations of wave propagation in magnetised plasmas (e.g. \\opencite{KC06}; \\opencite{SET06}; \\opencite{PK07}; \\opencite{CGD08}; \\opencite{RSK09}). A successful model should reproduce all of the observed features, including changes in the surface amplitudes of waves in the vicinity of active regions.   The acoustic halo is an observed enhancement of high frequency acoustic power surrounding regions of strong magnetic field. This enhancement was discovered independently by \\inlinecite{BLFJ92}, \\inlinecite{BBLT92} and \\inlinecite{TL93}. Observations made using Ca II K-line images showed these enhancements extending tens of Mm beyond active regions \\cite{BLFJ92, TL93}. \\inlinecite{BBLT92} showed patchier and more confined halos using Doppler velocity measurements taken in the Fe I $\\lambda$5576 line, and with higher spatial resolution. They also showed a peak in the power at 5.5~mHz. \\inlinecite{Hindman:1998p398} and later \\inlinecite{JH02} demonstrated that halos seen in SOHO/MDI \\cite{MDI95} Dopplergrams tend to be prominent in intermediate magnetic field strengths of 50-250 G and are absent in continuum intensity observations. Acoustic halos  have been shown to be more prominent in Ca II K-line power maps and less visible in the Doppler Ni I line shifts \\cite{Ladenkov:2002p369}. \\inlinecite{Finsterle2004} found that the acoustic halo increases in size with increasing height into the canopy over active regions. On the other hand, \\inlinecite{M05} found that high-frequency oscillations in TRACE UV filtergrams show deficits in magnetic fields, rather than enchancements. There is no established explanation for the enhanced power, although several ideas have been suggested (see Section~\\ref{discussion}).   In this paper we present an analysis revealing several new properties of acoustic halos and of surface acoustic power in general. We use data obtained by SOHO/MDI over 7 days containing several large active regions. In Section~\\ref{one} we demonstrate (i) a strong association of the enhanced high frequency power with nearly horizontal field, and (ii) an increase in the temporal frequency of the maximum enhancement with field strength. In Section~\\ref{two} we show that (iii) the oscillations contributing to the halos show modal ridges which are shifted to higher wavenumber at constant frequency than the ridges due to waves in the quiet-Sun. In the Discussion we list possible theories put forward, to date, that attempt to explain the existence of the acoustic halos, noting that none of them have been shown to explain all of the observed properties successfully. ",
        "conclusions": "\\label{discussion} We find properties of the acoustic halo that confirm prior analyses, but also discover several heretofore unknown properties. The halos are locations of acoustic power increase up to 140\\% of the quiet-Sun value at higher frequencies and are strongest at intermediate magnetic field strength (150 to 350~G).  These results are in general agreement with previous studies \\cite{JH02,Hindman:1998p398}. We have discovered a clear association with near horizontal ($\\pm 30^\\circ$) magnetic field. The dependence of acoustic amplitude on field inclination occurs over frequencies between 2.5 and 6.5~mHz, and for field strengths between 100 and 800 gauss, but is particularly strong between 5 and 6.5~mHz and at the aforementioned intermediate field strengths.  In addition, we show that the frequency of the peak enhancement depends on the field strength, with the peak frequency increasing slightly with field strength, and that in these regions of enhanced power the ridges of the  acoustic power spectrum are shifted towards higher $k$. Qualitatively there appears to be a larger shift for higher $k$ indicating that this is a shallow surface effect. For constant $k$ the phase speed of the ridge is reduced, indicating an increase in travel times, that could be caused by either a slower wave speed  or a longer path distance.  Any theory attempting to explain the acoustic halo or fully explain the interacation of waves with the magnetic field needs to incorporate these properties. A quantitative study of changes in the power spectrum requires a much more in depth analysis than our rudimentary study discussed here since, amongst other things, we have made no attempt to remove or account for distortions due to foreshortening and projection effects from the power spectra. On the other hand, we are confident through the multiple comparisons performed here with different masks and solar regions that the ridge shifts are qualitatively real and are clearly associated with the halo regions.  Various authors have pointed to a few different mechanisms to explain the halo. Some of the original discovery papers suggested enhanced acoustic emission \\cite{BDLJHP92, BBLT92}. More recently, \\inlinecite{JKWM08} have studied the effects of vertical magnetic fields on solar convection. The convection occurs on smaller scales and higher frequencies than in quiet-Sun regions, and shifts the oscillation power to higher frequencies. It is not as yet understood whether these properties are dependent upon the inclination of the field. However, the lack of halos in continuum observations \\cite{Hindman:1998p398} and the relative lack of association with emission sites, as identified with seismic holography analyses sensitive to propagating waves \\cite{Donea:2000p120} casts doubt on this interpretation.  \\inlinecite{Hanasoge:2008p816} claims that halos may be more likely related to scattering effects. Roughly in agreement with the findings in this paper, he also sees the halo as being confined to magnetic field inclinations $|\\gamma| < 40^\\circ$ and an increase of energy in the higher wavenumber regime. \\inlinecite{KC2009}  propose that high-frequency fast mode waves refracted in the higher atmosphere due to a rapid increase in the Alfv\\'en speed are responsible for the seismic halos. They get a 40-50\\% increase in power of high frequency vertical velocity in regions of intermediate field strength, although, like \\inlinecite{Hanasoge:2008p816}, they find that the horizontal velocity component of the halo is stronger. However, it is difficult to completely reconcile the observations with these simulations, since the latter show only a weak halo being generated in the vertical velocity component (up to a maximum of 5\\%) and the strongest enhancement (up to 35\\%) being generated  in the horizontal velocity component.  If the halo were strongest in horizontal velocities there would be a line-of-sight dependence with the halo power being stronger at larger heliocentric distance. This has not as yet been observed.  \\inlinecite{Kuridze:2008p578}  suggests this enhanced power is caused by high frequency waves trapped  beneath the bipolar magnetic canopy. The low frequency waves are reflected below the photosphere and so are not trapped. In regions of open field, the high frequency waves are permitted to propagate upwards and are lost to the atmosphere.  These field free cavities have granular dimensions ($\\approx 0.5$~Mm), but with SOHO/MDI we must be observing \\textit{many} of these cavities within the resolution which is three times this size, and these cavities must be widely distributed to cover the area of the halo regions. A somewhat similar idea is put forth by  \\inlinecite{MHS05}. They suggest that the lack of haloes in the TRACE observations \\cite{M05} caused by reflections (and mode conversion) due to an overlying inclined magnetic field.  \\inlinecite{Simoniello:2009} have studied Sun-as-a-star oscillations and observe enhanced power in high frequency (5.7-6.3~mHz) bands, up to 18\\%, at solar maximum. We believe that this is due to a large surface of the Sun being covered by plage and the associated enhanced power, as presented in this paper, associated with the increased magnetic activity.  Understanding the properties of surface acoustic power is important in modeling helioseismic signatures due to subsurface perturbations. It has also been suggested that acoustic power be used as a predictive tool to forecast the emergence of active regions \\cite{HKZM09}. The results shown here, namely that the surface acoustic power is a strong function of the local magnetic field properties, may be useful in the removal of surface contributions from acoustic power maps and as a part of a correction to local helioseismic methods (e.g. \\inlinecite{LB05},\\inlinecite{CR07}).  At this point, we cannot confidently suggest a mechanism explaining all the properties of the acoustic halo. However, a more detailed analyses of the characteristics of the acoustic halo will likely help to pin-point the mechanism behind the acoustic halos. The recently launched Solar Dynamics Observatory hosts the Helioseismic and Magnetic Imager instrument which has a resolution four times better than SOHO/MDI and, along side improvements in numerical modeling, may be able to help resolve the nature of acoustic halos."
    },
    "0911/0911.3873_arXiv.txt": {
        "abstract": "The core accretion theory of planet formation has at least two fundamental problems explaining the origins of Uranus and Neptune: (1) dynamical times in the trans-Saturnian solar nebula are so long that core growth can take $> 15$~Myr, and (2) the onset of runaway gas accretion that begins when cores reach $\\sim 10 M_{\\oplus}$ necessitates a sudden gas accretion cutoff just as Uranus and Neptune's cores reach critical mass. Both problems may be resolved by allowing the ice giants to migrate outward after their formation in solid-rich feeding zones with planetesimal surface densities well above the minimum-mass solar nebula. We present new simulations of the formation of Uranus and Neptune in the solid-rich disk of \\cite{icelines} using the initial semimajor axis distribution of the Nice model (Gomes et al.\\ 2005; Morbidelli et al.\\ 2005; Tsiganis et al.\\ 2005), with one ice giant forming at 12~AU and the other at 15~AU. The innermost ice giant reaches its present mass after 3.8-4.0~Myr and the outermost after 5.3-6~Myr, a considerable time decrease from previous one-dimensional simulations (e.g.\\ Pollack et al.\\ 1996). The core masses stay subcritical, eliminating the need for a sudden gas accretion cutoff.  Our calculated carbon mass fractions of $22\\%$ are in excellent agreement with the ice giant interior models of \\cite{podolak95} and \\cite{marley95}. Based on the requirement that the ice giant-forming planetesimals contain $> 10\\%$ mass fractions of methane ice, we can reject any solar system formation model that initially places Uranus and Neptune inside of Saturn's orbit. We also demonstrate that a large population of planetesimals must be present in both ice giant feeding zones throughout the lifetime of the gaseous nebula. This research marks a substantial step forward in connecting both the dynamical and chemical aspects of planet formation. Although we cannot say that the solid-rich solar nebula model of Dodson-Robinson et al.\\ (2009) gives {\\it exactly} the appropriate initial conditions for planet formation, rigorous chemical and dynamical tests have at least revealed it to be a {\\it viable} model of the early solar system. ",
        "introduction": "\\label{introduction}  The canonical core accretion theory of planet formation, in which planetesimals collide to form solid cores which then destabilize the surrounding gas to accrete an atmosphere (Safronov 1969; Pollack et al.\\ 1996), has at least two fundamental problems explaining the origins of Uranus and Neptune.  First, dynamical times in the trans-Saturnian solar nebula are so long and solid surface densities $\\Sigma$ are so low ($< 1$~g~cm$^{-2}$) according to the assumed $\\Sigma \\propto R^{-2}$ mass distribution (Pollack et al.\\ 1996) that planet growth takes $> 15$~Myr, far longer than both observed and theoretical protostellar disk lifetimes (Haisch et al.\\ 2001; Alexander et al.\\ 2006). Second, runaway gas accretion begins when solid cores reach 10 to 15~$M_{\\oplus}$, requiring a sudden and complete gas accretion cutoff just as Uranus and Neptune reach their current masses. \\cite{pollack96} pointed out these problems in their seminal paper on the viability of the core accretion theory. More recently, Benvenuto et al.\\ (2009) showed that Uranus and Neptune could grow within a few Myr in a population of planetesimals with a distribution of radii between 30 and 100~km. However, planetesimals as small as 30~km are not consistent with the prevailing theory of planetesimal formation, based on the streaming instability, which produces planetesimals around 100 km and in some cases up to the radius of Ceres (457 km; Johansen et al.\\ 2007).   Uranus and Neptune's total masses, 14.5 and 17.2~$M_{\\oplus}$ respectively, place them squarely in the predicted critical mass range for nucleating an instability in the surrounding gas and accreting Jupiter's mass or more in under 1000 years (Mizuno 1980; Papaloizou and Nelson 2005). {\\it The first challenge for theorists is to find a combination of the parameters that control core accretion---feeding zone location, ice inventory and planetesimal surface density---that leads to solid planet cores of} $> 14 M_{\\oplus}$ {\\it that form within observed protostellar disk lifetimes and are subcritical with respect to the surrounding gas density.} An ice giant formation theory should also account for the planets' bulk composition, particularly their 20--50 $\\times$ solar tropospheric C/H ratios (Encrenaz 2005). Treating feeding zone location as a free parameter creates the further challenge of moving Uranus and Neptune into their current orbits.  Two previous theories attempted to explain both the timely formation and subsequent orbital evolution of the ice giants. Thommes et al.\\ (1999, 2002) proposed that Uranus and Neptune are failed gas giants that formed between Jupiter and Saturn. Jupiter scattered the ice giants into orbits with semimajor axes $a > 15$~AU once it reached runaway gas accretion, while interactions with planetesimals further forced the ice giants slowly outward. The ``collisional damping scenario'' was put forth by Goldreich et al.\\ (2004a, 2004b). According to Goldreich et al., Uranus and Neptune formed {\\it in situ} from a dynamically cold planetesimal disk that also produced three other proto-ice giants. The protoplanets formed quickly despite long dynamical times because the planetesimal disk scale height fit within the Hill sphere (the protoplanet's zone of gravitational dominance), leading to high solid accretion rates. Dynamical friction could no longer damp the eccentricities of the $\\sim 5$ trans-Saturnian oligarchs once they attained a surface density comparable to the surrounding planetesimal disk. The oligarchs suffered close encounters and the resulting instability ejected all proto-ice giants but Uranus and Neptune.   The assumptions underlying the order-of-magnitude analysis in Goldreich et al. (2004a, 2004b) have ultimately proven unreliable. \\cite{levison07} demonstrated that the collisional damping scenario cannot reproduce the current Solar System: rather than ejecting three of five ice giants, the trans-Saturnian protoplanets simply spread out and all planets were retained. Furthermore, the collisional damping scenario requires that oligarchs grow while planetesimals fragment to sizes $\\ll 1$~km. Since low-velocity particles ($v < 10$~cm~s$^{-1}$) in the COLLIDE-2 microgravity experiment burrowed into the target material without producing ejecta (Colwell 2003), there is no reason planetesimals should fragment in the dynamically cold planetesimal disk required to produce Uranus and Neptune {\\it in situ}.   The Thommes et al.\\ (1999, 2002) ``failed gas giant'' model has substantial success reproducing the current Solar System and does not require finely tuned planetesimal behavior. Studies of planet formation in the 5-10~AU region demonstrate the efficiency of growing ice giant-sized cores between Jupiter and Saturn [\\cite{hubickyj05}, \\cite{saturn}]. However, the compositions of Uranus and Neptune strongly indicate an origin in the trans-Saturnian solar nebula. Tropospheric abundances of methane show carbon enrichments of 20--50 times solar (Encrenaz 2005), and interior models find methane mass fractions of $\\sim 20$\\% (Marley et al.\\ 1995; Podolak et al.\\ 1995). The combined dynamical and chemical model of the solar nebula calculated by \\cite{icelines} shows that the methane condensation front is beyond Saturn's orbit during the first $5 \\times 10^5$ of solar nebula evolution. Without methane ice present during the planetesimal-building epoch---which lasts only $5 \\times 10^4$ years according to Johansen et al.\\ 2007---neither planet could obtain its methane-rich composition.   The Nice model of planetary dynamics \\cite{tsiganis05}, \\cite{gomes05}, \\cite{morbidelli05} uses initial conditions that place Uranus and Neptune initially in the methane ice-rich regions beyond 10~AU. In the Nice model, Neptune and Uranus assume initial semimajor axes of $\\sim 12$ and $\\sim 15$~to~17~AU. When planetesimal perturbations pull Jupiter and Saturn across their 1:2 mean motion resonance (MMR), their eccentricities suddenly increase, forcing close encounters between all possible pairs of giant planets except Jupiter and Saturn. In about half of the simulations, Neptune is scattered across Uranus' orbit, leapfrogging to $\\sim 23$~AU within a few $10^5$ years. Slow outward migration due to interaction with a planetesimal disk pulls Uranus and Neptune into their current orbits over the course of $\\sim 40$~Myr.    Although the Nice model explains the current orbits of the giant planets, it is incomplete without an assessment of the planets' ability to form in their predicted initial orbits. With high solid surface- density, methane-rich planetesimals, the protostellar disk model of \\cite{icelines} contains promising initial conditions for ice giant formation between 12 and 15~AU. None of the three dynamical theories discussed---the Nice model, the failed gas giant theory and the collisional damping scenario---treats the growth of the ice giants' envelopes, a gap in the literature that this work is partly intended to fill. Verifying that $\\sim 10-15 M_{\\oplus}$ solid cores can form is an important step---one which N-body simulations show is extremely difficult even in the inner nebula (McNeil et al.\\ 2005; Chambers 2008)---but one also has to verify that the ice giant atmospheres stay under $10\\%$ of the total planet mass and do not experience runaway growth.  In this paper, we demonstrate that Uranus and Neptune can form by core accretion, in the feeding zone approximation, in the trans-Saturnian solar nebula using the Nice model initial semimajor axis distribution. In \\S \\ref{simulation} we describe the numerical methods used in our experiments. In \\S \\ref{results} we discuss the results of our core accretion simulations, focusing on formation timescale, accretion efficiency and solid/gas ratio. In \\S \\ref{discussion} we discuss the strengths and weaknesses of our model as a realistic descriptor of Uranus and Neptune's formation. We present our conclusions in \\S \\ref{conclusions}. ",
        "conclusions": "\\label{conclusions}  We have shown that it is possible to assemble Uranus and Neptune from $1 M_{\\oplus}$ seed cores in a $0.12 M_{\\odot}$ solar nebula, with initial semimajor axes of 12~AU and 15~AU as predicted by the Nice model (Gomes et al.\\ 2005; Morbidelli et al.\\ 2005; Tsiganis et al.\\ 2005).  The resulting growth timescales are 3.8~Myr for the innermost ice giant and 5.3~Myr for the outermost ice giant.  However, numerous details still remain to be fully resolved:  \\begin{enumerate}  \\item What is the migration history of the planets after formation and can we explain the current orbits?  \\item If Planets I and O do not change places, how can Planet I's lack of a massive gaseous atmosphere be explained?  \\item If Planets I (Neptune) is scattered to a wide orbit while the disk is still gas-rich, what happens to Planet O (Uranus)?  \\item How can the Late Heavy Bombardment be explained if the giant planets' orbital evolution must take place soon after they finish forming?  \\end{enumerate}  Because we chose the initial semimajor axes of the Nice model, both ice giants' post-formation orbital evolution may be explained according to that model. If Uranus and Neptune form between 12 and 15~AU in the feeding zone approximation, they end up with either different formation timescales or different masses. If the proto-ice giants begin limiting their own solid accretion rates by ejecting planetesimals, the outermost planet core can continue to grow after the innermost core has assembled the bulk of its mass, but the inner planet may evolve into a gas giant. One way to form Uranus and Neptune at 12~AU and 15~AU but keep their near-equal masses may be for Neptune (the innermost planet in this scenario) to get scattered out to a wide orbit upon reaching its present mass---but in this scenario we cannot be assured of Uranus' survival and continued growth.  One way our results differ from the scenario described by the Nice model suite of papers is that the orbital evolution of the giant planets must take place shortly after Uranus and Neptune finish forming. No disk where Uranus and Neptune fall in a gap in the planetesimal distribution {\\it at any time during their growth} can form planets with $\\sim 90 \\%$ solids by mass. As stated in \\S \\ref{discussion}, we still lack a mechanism for checking Planet I's growth while Planet O finishes forming. If Neptune is the innermost ice giant, Jupiter and Saturn's 1:2 MMR crossing might perform such a function by suddenly sending Neptune into an extremely wide orbit.  The formation timescale of the ice giant forming at 15~AU is quite close to both the upper limit of observed protostellar disk lifetimes (Haisch et al.\\ 2001; Currie et al.\\ 2009) and the predicted onset of photoevaporation for the solar nebula (Alexander et al.\\ 2006), indicating that quick ice giant formation in the trans-Saturnian nebula is only marginally possible even in the optimistic feeding zone approximation. Since the compositions of Uranus and Neptune strongly indicate that they originated in the trans-Saturnian solar nebula, however, we encourage more detailed N-body investigations of embryo growth rates between 12 and 15~AU.  Without knowing for sure whether Uranus and Neptune swapped orbits, we cannot pair each planet with its unique formation zone. However, the chemically evolving solar nebula simulations of \\cite{icelines} show that the CO and N$_2$ ice lines fall between 12 and 15~AU. Detailed measurements of the carbon and nitrogen abundances of both planets and their satellites could resolve the question of which planet formed outermost in the solar nebula: the planet that formed near 15~AU should be slightly more carbon- and nitrogen-rich than the planet that formed near 12~AU. Such detailed measurements of Uranus and Neptune's atmospheric compositions await another space mission to the outer planets, but ground-based spectroscopy of their satellites (e.g.\\ Stern \\& McKinnon 2000; Bauer et al.\\ 2002) may make some progress on this question. Our formation model reproduces the ice giants' solid/gas ratios and carbon abundances inferred from gravitational moment measurements and spectroscopy.  Our results indicate that planet formation in the outer nebula is a fundamentally inefficient process. Even in the feeding zone approximation, where we calculate what is best viewed as an upper limit to accretion rate, building ice giants within the lifetime of the solar nebula requires substantially more solid and gas mass than the planet can ever accrete. The dynamical times in the massive \\cite{icelines} disk (total mass $0.12 M_{\\odot}$) beyond 15~AU are still too long to allow gas or ice giant formation. From the combined constraints of the methane ice line location and the upper limits to ice giant growth rates, we can constrain the ice giant formation zone to between 10 and 15~AU.   The fact that the solar nebula model of \\cite{icelines} can form all the giant planets at locations at or near those predicted by the Nice model, as well as reproduce important features of the compositions of the planets and their satellites---such as ammonia enrichment in the Saturnian system (Dodson-Robinson et al.\\ 2008)---indicates that substantial progress has been made at modeling the ice inventory and chemical evolution of the solar nebula. Here we have created the first successful progression (to the best of our knowledge) from primordial solar system ice inventory to detailed planet composition. This research marks a substantial step forward in connecting both the dynamical and chemical aspects of planet formation. Although we cannot say that the solid-rich solar nebula model (Dodson-Robinson et al.\\ 2009) gives {\\it exactly} the appropriate initial conditions for planet formation, rigorous chemical and dynamical tests have at least revealed it to be a {\\it viable} model of the early solar system.    Support for S.D.R.'s work was provided by NASA through the Spitzer Space Telescope Fellowship Program. P.B. received support from the NASA Origins Grant NNX08AH82G. We thank both referees for comments that improved the clarity and focus of the manuscript."
    },
    "0911/0911.5178_arXiv.txt": {
        "abstract": "We present a detailed study of the collapse of a spherical perturbation in DGP braneworld gravity for the purpose of modeling simulation results for the halo mass function, bias and matter power spectrum.  The presence of evolving modifications to the gravitational force in form of the scalar brane-bending mode lead to qualitative differences to the collapse in ordinary gravity.  In particular, differences in the energetics of the collapse necessitate a new, generalized method for defining the virial radius which does not rely on strict energy conservation.  These differences and techniques apply to smooth dark energy models with $w\\neq -1$ as well. We also discuss the impact of the exterior of the perturbation on collapse quantities due to the lack of a Birkhoff theorem in DGP. The resulting predictions for the mass function, halo bias and power spectrum are in good overall agreement with DGP N-body simulations on both the self-accelerating and normal branch. In particular, the impact of the Vainshtein mechanism as measured in the full simulations is matched well.  The model and techniques introduced here can serve as practical tools for placing consistent constraints on braneworld models using observations of large scale structure. ",
        "introduction": "\\label{sec:intro}   Modified gravity models have attracted a great deal of interest recently as an alternative explanation of the observed accelerated expansion of the Universe \\cite{RiessEtal,KowalskiEtal,WMAP5,Giannantonio,Pietrobon}.  In order for this scenario to work, gravity must be significantly modified from General Relativity (GR) on cosmological scales, but has to reduce to GR locally in order to satisfy stringent Solar System constraints at a few AU. Thus, a working modified gravity model has to include a non-linear mechanism to restore GR in high-density environments, which can have a noticeable impact on the formation of large-scale structure on intermediate scales of a few to tens of Mpc \\cite{LueEtal04,HPMpaperII,HPMhalopaper,DGPMpaper}.  One popular modified gravity scenario is the DGP braneworld model \\cite{DGP1}. Here, a four-dimensional Friedmann-Robertson-Walker universe is imbedded as a brane in five-dimensional Minkowski space.  In this model, gravity is 5-dimensional on the largest scales, and becomes four-dimensional at scales below the crossover scale $r_c$, a fundamental parameter of the model.  The modification of the Friedmann equation depends on the choice of embedding for the brane \\cite{Deffayet01} which yields two branches of the model: the \\emph{self-accelerating} branch, on which accelerated expansion occurs at late times without any cosmological constant or dark energy, and the \\emph{normal} branch where no such acceleration occurs, and a brane tension or other form of stress-energy with negative pressure has to be added on the brane.  On scales much smaller than the horizon and crossover scales, DGP gravity can be described by an effective scalar-tensor theory \\cite{NicRat,KoyamaMaartens}, where the additional scalar degree of freedom, the brane-bending mode $\\ph$ is associated with displacements of the brane from its background position. The brane-bending mode yields an additional gravitational force which influences dynamics of non-relativistic particles.  On the normal branch of DGP this force is attractive, while it is repulsive in the self-accelerating branch. Non-linear interactions of $\\ph$, via the so-called Vainshtein mechanism, ensure that it has tiny effects within the Solar System. For values of $r_c$ of order the current horizon, these interactions become important as soon as the density contrast becomes of order unity.  Hence, it is crucial to consistently follow the full brane-bending mode interactions in a cosmological simulation, as has recently been done \\cite{DGPMpaper,ScII}. While we study the behavior of $\\ph$ in specific DGP models, the form of the non-linear interactions is expected to be generic to braneworld models with large extra dimensions \\cite{deGrav,galileon}.  Given the considerable computational expense of these simulations, it is worthwhile to develop a model which captures the main modified gravity effects, enabling forecasts and constraints properly marginalized over cosmological parameters. In the halo model of large structure \\cite{CooraySheth}, which assumes that all matter is in bound dark matter halos, the abundance and clustering of halos are determined by the linear power spectrum once characteristic quantities of the collapse of a spherical perturbation are determined, namely, the linear collapse threshold and the virial radius or overdensity.  Spherical collapse is well studied for General Relativity with a cosmological constant (e.g., \\cite{Peebles80,LahavEtal91}), and has also been explored for quintessence-type dark energy \\cite{WS98,HB05,BasilakosEtal}.  A first study of spherical perturbations in the context of DGP was undertaken in \\cite{LueEtal04}. We extend previous studies by deriving the full $\\ph$ field profile of an isolated mass in closed form, and by carefully defining the interior and exterior forces during collapse along with the energetics implied by the profile.  In particular, the potential energy required for the virial theorem and for the total Newtonian energy differ, with the latter not being strictly conserved during collapse.\\footnote{The ``total energy'' here is defined for a Newtonian cosmology and its non-conservation does not imply violation of covariant energy-momentum conservation.} This lack of a conservation law also applies to dark energy models with an equation of state $w\\neq -1$.  We show how this problem can be circumvented by properly defining the condition for virial equilibrium based on forces.  We discuss the general properties and parameterization of DGP models in \\refsec{models}, and the spherical collapse calculation in \\refsec{rev}.  The halo model calculations are outlined in \\refsec{HM}, and the results and comparisons with simulations are given in \\refsec{res}.  We conclude in \\refsec{concl}. The Appendix contains derivations of the $\\ph$ profile and other quantities needed in the collapse calculation, as well as a discussion of the potential energy and virial theorem in DGP. ",
        "conclusions": "\\label{sec:concl}  By studying the collapse of a spherical perturbation under DGP braneworld gravity, we have shown how simulation results on the mass function, halo bias and power spectrum can be understood semi-analytically. In DGP gravity, force modifications are carried by the scalar brane-bending mode.  The global properties of the response of the brane-bending mode to matter control how the Vainshtein mechanism modifies force and energy conditions during collapse.  These conditions are important for the calculation of virial equilibrium (see the Appendix for detailed discussions of these results).  In particular, the presence of evolving modifications to the gravitational force either through the brane bending mode or through the background expansion violate conservation of Newtonian total energy for traditional definitions of the potential energy contribution.  This violation applies to smooth dark energy models with $w\\neq -1$ as well.   We introduce a new, general technique for defining the virial radius which does not rely on strict energy conservation.  Under the halo model, these spherical collapse predictions give rise to predictions for the mass function, halo bias and power spectrum.  We have shown that these predictions are in good qualitative agreement with DGP N-body simulations on both the self-accelerating and normal branch. In particular, the use of spherical collapse for the mass function always provides slightly conservative limits on mass function deviations when compared with the simulations.  Hence the semi-analytic techniques introduced here can be used as a practical tool for extending simulation results for the purpose of studying parameter constraints on braneworld models from observations of the mass function.  While the absolute power spectrum agreement is not quite as good, these techniques can still be useful in combination with \\emph{linearized} DGP simulations. Since our spherical collapse predictions appear to capture accurately the impact of the Vainshtein mechanism on the power spectrum, they could provide an effective way of taking non-linear interactions into account in results obtained from linearized DGP simulations. Such simulations are very easy to implement and an order of magnitude cheaper computationally than the full DGP simulations. Likewise they are more readily extendable to higher resolution and can be used to cover a wider range in parameter space."
    },
    "0911/0911.0870_arXiv.txt": {
        "abstract": "Verified and updated calibrated absolute solar flux in the He~II 30.4~nm spectral band-pass as measured by the Solar EUV Monitor (SEM) allows us to study variations of the solar EUV irradiance near the minima of Solar Cycles 22/23 and 23/24. Based on eight (1996 to 2007) NASA sounding rocket flights, a comparison of SEM data with the measurements from three independent EUV instruments was performed to verify and confirm the accuracy of the published SEM data. SEM calibrated data were analyzed to determine and compare minima for solar cycles 22/23 and 23/24. The minima points were calculated using SEM first order daily averaged flux smoothed by a running mean (RM) filter with the window of averaging equal to 365 days. These minima occurred on June 2, 1996 (22/23) and November 28, 2008 (23/24). The 23/24 minimum showed about 15\\% lower EUV flux in the 30.4 nm band-pass than the 22/23 minimum. The 365-day RM curve around the 23/24 minimum has significant asymmetry (fast decrease of the EUV flux to the minimum and a long, near-horizontal profile after the minimum). This profile is quite different from the much faster and symmetrical change of the flux around the 22/23 minimum. SEM flux was compared with both high spectral resolution (0.1~nm) Mg~II index calculated from the Solar Radiation and Climate Experiment (SORCE) using the Solar Stellar Irradiance Comparison Experiment (SOLSTICE) data and with the NOAA composite Mg~II index spectrum. ",
        "introduction": "SOHO/CELIAS SEM \\citep{Hovestadt95} is a highly stable and accurate solar EUV spectrometer. The stability of the spectrometer is provided by a transmission diffraction grating based on thin gold bars \\citep{Schattenburg90}. The two (plus and minus) first order bands are centered about the He II (30.4~nm) EUV spectral line. The high accuracy of SEM EUV measurements is achieved by correcting for changes of the flight SEM sensitivity to the solar EUV irradiance based on comparisons of the flight data with the data from a number of NASA sounding rocket calibration flights. The sounding rocket flights have provided solar measurements to correct for these changes in sensitivity, which are due to minor degradation of SOHO SEM's thin film Al filters \\citep{Judge98}. A model of such time dependent degradation was created based on the data from the three earliest SOHO sounding rocket under-flights (1996 to 2000). The model was confirmed by the following four under-flights (2001 to 2006). This comparison of the SOHO SEM solar flux measured during the sounding rocket under-flights showed that the absolute solar flux is accurate \\citep{Judge08} to within $\\pm 5\\%$ of the measured flux.  The long, (more than 13.8 years) and practically uninterrupted absolute EUV flux in the mean (plus and minus) first order band around the He II spectral line is ideally suited for an analysis of the 22/23 and 23/24 solar minima. ",
        "conclusions": "The SEM EUV 30.4 nm scaled flux matches the Mg II index for the 22/23 and 23/24 minima. SEM flux shows the minimum for 2008 is lower ($15 \\pm 6\\%$) than the minimum for 1996, $0.97\\times 10^{10} {ph/cm^{2}/s}$ and $1.12\\times 10^{10} {ph/cm^{2}/s}$, respectively. This is consistent with recent 04/14/08 SDO/EVE sounding rocket measurements using both EVE/ESP and EVE/MEGS irradiance, and with the result shown by \\citet{Lockwood07}. The 2008 minimum shows a sharp decrease of irradiance and a steady, slow increase. The 1996 minimum has symmetric wings. The 23/24 minima determined from SEM and Mg II data with a RM window of 365-d practically coincide. They occurred on 28 Nov 08 (SEM) and 22 Nov 08 (Mg II)."
    },
    "0911/0911.0041_arXiv.txt": {
        "abstract": "We present new analytical estimates of the large scale bias of neutral hydrogen ($\\HI$) We use a simple, non-parametric model which monotonically relates the total mass of a halo $\\Mtot$ with its $\\HI$ mass $\\Mhi$ at zero redshift; for earlier times we assume limiting models for the $\\Omega_{\\HI}$  evolution consistent with the data presently available, as well as two main scenarios for the evolution of our $\\Mhi$ - $\\Mtot$ relation. We find that both the linear and the first nonlinear bias terms exhibit a strong evolution with redshift, regardless of the specific limiting model assumed for the $\\HI$ density over time. These analytical predictions are then shown to be consistent with measurements performed on the Millennium Simulation. Additionally,  we show that this strong bias evolution does not sensibly affect the measurement of the $\\HI$ power spectrum. ",
        "introduction": "The 21cm emission line of neutral hydrogen has been the workhorse of astronomy for over 50 years. Its use for galactic astronomy cannot be overestimated, and it now promises to revolutionize cosmology as well \\citep[cf.][and references therein]{21cm:fob06,21cm:bl07,21cm:pl08,Mao:08,Visbal:08}.  Compared to present probes of large-scale structure, the 21 cm intensity mapping technique \\citep{21cm:cppm08,21cm:wl08} offers three advantages. First, it allows to reconstruct the spatial distribution of $\\HI$ in the universe over volumes much larger than the ones currently probed by galaxy redshift surveys. Since the precision with which the power spectrum of density fluctuations can be measured is approximately proportional to the number of independent Fourier modes that fit into the survey volume, a larger survey volume leads to higher precision in the measured power spectrum and hence cosmological parameters \\citep{Loeb:2008hg}. Second, 21 cm intensity mapping allows to probe a very wide range of redshifts and to map $\\HI$ distribution in space and time not only for the epoch after reionization ($z\\lesssim 6$) but also during reionization and even in the so-called dark ages, before the first galaxies were formed ($z\\gtrsim 20$). Third, using radio telescopes is one of the most economical ways to probe the dark energy evolution among all current and planned experiments, although several crucial advances in calibration and foreground removal must be made for the intensity mapping approach to reach its stated goals.  As a dark energy probe, 21 cm intensity mapping is primarily useful for measuring the baryon acoustic oscillations (BAO) in the matter power spectrum \\citep{Wyithe:2007rq,Chang_et_al:08,21cm:wl08,Ansari:08,Seo_et_al:09}, which provides important information on constraining cosmological models and parameters. Measurements of the scale of the BAO in the matter power spectrum as a function of redshift, when combined with the true physical scale, i.e., the sound horizon scale measured from the cosmic microwave background (CMB), probe the angular diameter distances and Hubble parameters to various redshifts, and therefore dark energy properties \\citep[see, for instance, ][]{hu_white:96, eisenstein_etal:98, eisenstein:02, blake_glazebrook:03, linder:03,hu_haiman:03, seo_eisenstein:03}. The distinct oscillatory feature of the BAO allows us to isolate and measure the BAO information independent of the overall broadband shape of the HI power spectrum. Therefore in principle BAO measurements do not require a knowledge of the HI bias, i.e., the relation between the neutral hydrogen and the total matter clustering, which affects the overall shape of the HI power spectrum.  However, understanding the HI bias is important for the 21cm BAO surveys in the following ways. First the prediction of the signal to noise of the BAO measurements is sensitive to the knowledge of the clustering bias of HI \\citep{Chang_et_al:08,Seo_et_al:09}. Second, understanding HI bias will help us to extract cosmological information from the broadband shape of the power spectrum in addition to the BAO. Third, the nonlinear effects on the BAO, such as the degradation and the shift of the feature, may differ for different biased tracers \\citep[e.g.,][]{padmanabhan_white:09, metha_etal:10}. Knowing the properties of HI bias will allow us to approximately estimate the degree of the expected nonlinear effect on the BAO in the HI distribution.   In this work, we focus on two critical aspects of the large-scale $\\HI$ bias: its dependence on the relation between $\\HI$ mass and halo mass and its evolution with redshift. To study the first aspect,  we use a non-parametric model where we assume a one-to-one correspondence between a  $\\HI$ mass present in a dark matter halo $\\Mhi$ and the total mass $\\Mtot$ of this halo. Using this $\\Mhi -\\Mtot$ relation we estimate the large-scale bias of the $\\HI$ using  a model containing elements on the halo occupation distribution (HOD) formalism and the $\\HI$ mass function measured at $z=0$ \\citep{igm:zmsw05}. We then use this relation to paint $\\HI$ on Millennium Simulation halos \\citep{springel_etal:05} and to measure the power spectrum of neutral hydrogen in order to test our model assumptions.   With respect to the redshift evolution of the bias, it is absolutely critical to know how  the $\\HI$ mass function evolves with time. However, to the best of our knowledge, little is known about the redshift evolution of $\\HI$ both on the theoretical and the observational sides \\citep[but see][]{wyithe_etal:09}. A detailed study of the bias evolution requires two extra pieces of information which are currently poorly constrained by observational data: the evolution of the total HI mass in the universe and the evolution of the HI mass function. For both of these, we empirically assume limiting evolution models consistent with the few available data that should bracket the actual evolution. We then proceed to predict the bias evolution in these scenarios, being aware of the fact that the actual evolution will lie somewhere in between these limiting cases.  The paper is organized as follows. In \\S\\ref{app:biashod} we introduce the basics of our model, we derive the relation $M_{\\rm HI}-Mtot$ at $z=0$, we introduce the limiting cases for its evolution and we define the bias. In \\S\\ref{section:results} we present our main results: analytic predictions for the evolution of HI bias, measurement of its scale dependence on the Millennium Simulation and how it would affects measurements of the power spectrum. Finally, we conclude in \\S\\ref{section:conclusions}. To compare analytical and numerical results, throughout the paper we use  a flat $\\Lambda$CDM cosmology with $\\Omega_M=0.25$,  $\\sigma_8=0.9$, Hubble parameter $h=0.73$ and initial power spectrum index of $n_s=1$, consistent with the cosmological parameters used in the Millennium Simulation. ",
        "conclusions": "\\label{section:conclusions}  In this paper we have estimated analytically and in the Millennium Simulation the large-scale bias of neutral hydrogen from a non-parametric model which establish a one-to-one relation between the total mass  and the $\\HI$ mass of galactic halos. Although only an approximation, this relation is expected to hold for a large range of masses.  The bias parameters have been calculated using a model which uses elements of the HOD formalism, in a way that can be considered complementary to the galaxy-HOD model of $\\HI$ galaxies carried out by \\cite{wyithe_etal:09}; we do not deal in our case with the individual galaxies inside dark matter halos but with the total neutral hydrogen inside them. In terms of large-scale clustering, these approaches are completely equivalent.  Using this  approach together with different models for the redshift evolution of the $\\HI$ content in the universe and the relation between neutral hydrogen and total matter contents of galactic halos, we analyze the evolution of bias with redshift. We find that both linear and nonlinear bias increase significantly with redshift. This contrasts with the finding by \\citet{Wyithe:2007rq}, based on an extrapolation to lower redshifts from a reionization model of \\citet{Wyithe:2007gz}. Since \\citet{Wyithe:2007rq} do not account for clustering of galactic halos that host all neutral hydrogen after reionization, they underestimate the $\\HI$ bias at $z \\lesssim 4$. On the other hand, our results are in  agreement with the high value of the linear large-scale bias at redshift $z\\sim 3$ measured in the model by \\cite{bagla_khandai:09} where they populate $\\HI$ in high-resolution $N-$body simulations.  We also find that, given the limitations of our model at high masses, the estimation of the bias does not change significantly when the very high mass halos are neglected.  Our main conclusion is that, despite a wide range of possible models for the evolution of $\\HI$ mass function that we consider, in all models the bias evolution is similar --the neutral hydrogen distribution is mildly antibiased at $z=0$ but becomes strongly biased ($b_\\HI \\sim 2$) by $z \\sim 4$. This result is encouraging for the planned radio intensity mapping experiments, since large bias implies a stronger 21 cm signal. Nevertheless, this strong redshift evolution does not significantly compromise the measurements of the neutral hydrogen power spectrum along the radial direction."
    },
    "0911/0911.5208_arXiv.txt": {
        "abstract": "Recent observations of the black hole (BH) - bulge scaling relations usually report positive redshift evolution, with higher redshift galaxies harboring more massive BHs than expected from the local relations. All of these studies focus on broad line quasars with BH mass estimated from virial estimators based on single-epoch spectra. Since the sample selection is largely based on quasar luminosity, the cosmic scatter in the BH-bulge relation introduces a statistical bias leading to on average more massive BHs given galaxy properties at high redshift \\citep{Lauer_etal_2007b}. We here emphasize a previously under-appreciated statistical bias resulting from the uncertainty of single-epoch virial BH mass estimators and the shape of the underlying (true) BH mass function, which leads to on average overestimation of the true BH masses at the high-mass end \\citep{Shen_etal_2008b}. We demonstrate that the latter virial mass bias can contribute a substantial amount to the observed excess in BH mass at fixed bulge properties, comparable to the Lauer \\etal\\ bias. The virial mass bias is independent of the Lauer et al.\\ bias, hence if both biases are at work, they can largely (or even fully) account for the observed BH mass excess at high redshift. ",
        "introduction": "\\label{sec:intro}  The redshift evolution of the local scaling relations between galaxy bulge properties and the mass of the central supermassive black hole (SMBH) has important clues to the establishment of these tight relations across cosmic time. There is currently a huge effort in measuring this evolution, either in terms of the BH mass--bulge velocity dispersion ($M_\\bullet-\\sigma$) relation, or the BH mass--bulge stellar mass/luminosity ($M_\\bullet-M_{\\rm bulge}$, $M_\\bullet-L_{\\rm bulge}$) relation \\citep[e.g.,][]{Shields_etal_2003,Peng_etal_2006a,Peng_etal_2006b,Salviander_etal_2007,ShenJ_etal_2008, Woo_etal_2006, Woo_etal_2008, Treu_etal_2007,Jahnke_etal_2009, Greene_etal_2009, Merloni_etal_2009,Decarli_etal_2009b,Bennert_etal_2009}. With a few exceptions, most of these studies report a strong positive evolution in BH mass for fixed bulge properties, which reaches as high as $\\sim 0.6$ dex offset in BH mass from the local relation at the highest redshift probed ($z\\sim 2$).  This observed strong evolution is difficult to understand in two aspects. For the $M_\\bullet-\\sigma$ relation, numerical simulations based on self-regulated BH growth \\citep[e.g.,][]{Robertson_etal_2006} and other theoretical arguments \\citep[e.g.,][]{Shankar_etal_2009c} favor a mild evolution in the normalization of the $M_\\bullet-\\sigma$ relation, in conflict with the otherwise claimed strong evolution. For the $M_\\bullet-M_{\\rm bulge}$ or $M_\\bullet-L_{\\rm bulge}$ relation, many of the observed hosts at earlier epoches are already bulge-dominated\\footnote{We note that there are also some observed hosts which are likely to be late-type \\citep[e.g.,][]{Sanchez_etal_2004,Letawe_etal_2007,Merloni_etal_2009}. }, hence unless their bulges continue to grow a considerable amount over time, it is difficult to understand how these systems would have migrated to the local relation.  An alternative explanation for this observed strong evolution is that it is caused, at least partly, by the systematics involved in deriving both host properties and BH masses. All these studies mentioned above focus on broad line quasar samples, with host properties measured from the quasar-light subtracted images or spectra. The systematics in subtracting the quasar light, as well as in converting observables (such as photometric colors or spectral properties) to physical quantities (such as bulge stellar mass) is a potential contamination to the final results. More importantly, the samples are predominately selected in quasar luminosity and hence are {\\em not} unbiased samples for studies of the BH scaling relations.  One statistical bias resulting from using quasar-luminosity selected samples was discussed in detail in \\citet{Lauer_etal_2007b}. Because there is cosmic scatter in the BH scaling relations and because the galaxy luminosity (or stellar mass, velocity dispersion, etc) function is bottom heavy, there are more smaller galaxies hosting the same mass BHs than the more massive galaxies. Hence when selecting samples based on quasar luminosities, low-mass BHs are under-represented, leading to an apparent excess in BH mass at fixed host properties. If the cosmic scatter in BH scaling relations increases with redshift and reaches $\\ga 0.5$ dex, it can in principle account for all the excess in BH mass observed for the most luminous quasars \\citep[e.g.,][]{Lauer_etal_2007b,Merloni_etal_2009}.  A second statistical bias, resulting from the uncertainty in the BH mass estimates, was pointed out in Shen \\etal\\ (2008b, also see Kelly \\etal\\ 2009a). In all these studies, the BH masses are estimated using the so-called virial method based on single-epoch spectra. In this method, one assumes that the broad line region (BLR) is virialized, and the BH mass can be estimated as $M_{\\rm vir}\\approx G^{-1}RV^2$, where $R$ is the BLR radius and $V$ is the virial velocity; one further estimates the BLR radius using a correlation between $R$ and the continuum luminosity $L$, i.e., $R\\propto L^{C_1}$, found in local reverberation mapping (RM) samples \\citep[e.g.,][]{Kaspi_etal_2005,Bentz_etal_2006,Bentz_etal_2009}, and estimates the virial velocity using the width of the broad lines. In this way, one can estimate a virial BH mass using single-epoch spectra: $\\log M_{\\rm vir}= \\log R + 2\\log({\\rm FWHM})+ {\\rm const}= C_1\\log L+2\\log({\\rm FWHM})+{\\rm const}$. The coefficients are calibrated empirically using RM samples or inter-calibrations between various lines \\citep[e.g.,][]{McLure_Jarvis_2002,Vestergaard_Peterson_2006,Mcgill_etal_2008}.  Even though the virial method is widely used, these BH mass estimates based on several lines (usually \\halpha, \\hbeta, \\MgII\\ and \\CIV) have a non-negligible uncertainty $\\sigma_{\\rm vir}\\ga 0.4$ dex, when compared to RM masses or masses derived from the $M_\\bullet-\\sigma$ relation \\citep[e.g.,][]{Vestergaard_Peterson_2006,Onken_etal_2004}. This uncertainty must come from the imperfectness of using luminosity and line width as proxies for the BLR radius and virial velocity, i.e., there are substantial {\\em uncorrelated} variations $\\sigma_{L}$ in luminosity and variations $\\sigma_{\\rm FWHM}$ in line width, which together contribute to the overall uncertainty $\\sigma_{\\rm vir}$, i.e., \\begin{equation}\\label{eqn:sig_vir} \\sigma_{\\rm vir}=\\sqrt{(C_1\\sigma_{L})^2 + (2\\sigma_{\\rm FWHM})^2}\\ . \\end{equation} One can then imagine a statistical bias will arise if the underlying active black hole mass function (BHMF) is bottom heavy. In particular, the variations (uncompensated by the variations in FWHM) of luminosity at fixed true BH mass, $\\sigma_L$, will scatter more lighter BHs into a luminosity bin than heavier BHs, and bias the mean BH mass in that bin. This is the Malmquist-type bias (or Eddington bias) emphasized in \\citet[][sec 4.4]{Shen_etal_2008b}, which has received little attention in the studies on the evolution of BH scaling relations to date.  In this paper we examine the impact of this mass bias on the observed evolution in BH scaling relations with realistic models for the underlying true BHMF and quasar luminosity function (LF). In \\S\\ref{sec:mass_bias} we review the statistical mass bias discussed in \\citet[][]{Shen_etal_2008b} and demonstrate its effects with simple models; in \\S\\ref{sec:sim} we consider more realistic intrinsic BHMF and quasar LF, and estimate the mass bias as functions of luminosity and redshift; we discuss the impact of this mass bias and conclude in \\S\\ref{sec:disc}. Throughout the paper we use cosmological parameters $\\Omega_0=0.3$, $\\Omega_\\Lambda=0.7$ and $H_0=70\\ {\\rm km\\,s^{-1}\\,{Mpc}^{-1}}$. Luminosities are in units of ${\\rm erg\\,s^{-1}}$ and BH masses are in units of $M_\\odot$. Unless otherwise specified, ``luminosity'' refers to the bolometric luminosity, and we are only concerned with the active BH population. ",
        "conclusions": "\\label{sec:disc}  The virial mass bias discussed here depends on the amplitude of the scatter $\\sigma_L$. This is the variation in luminosity at fixed true BH mass which is {\\em not compensated} by the variation in FWHM. How large is $\\sigma_L$? For the best studied local RM samples and for \\hbeta\\ only \\citep[e.g.,][]{Kaspi_etal_2005,Bentz_etal_2006,Bentz_etal_2009}, the scatter of luminosity around fixed BLR size is on the order of $\\la 0.2$ dex. However, this is for the best cases. We generally expect larger $\\sigma_L$ for statistical quasar samples with single-epoch spectra for the following reasons: 1) single-epoch spectra have larger scatter than averaged spectra in RM samples; 2) these samples usually cover a luminosity range that is poorly sampled by current RM data; 3) the UV regime of the spectrum (in particular for \\CIV) has more peculiarities and larger scatter in the $R-L$ relation than the \\hbeta\\ regime, but is usually used at high redshift due to the drop off of \\hbeta\\ in the optical spectrum; 4) even if some variations in luminosity truthfully trace the variations in BLR radius $R$, they may still be uncompensated by FWHM if, for instance, there is a non-virial component in the broad line which does not respond to BLR radius variations \\citep[especially for \\CIV, see discussions in ][]{Shen_etal_2008b}; 5) in many statistical quasar samples, the spectra have low signal-to-noise ratios, which lead to increased uncorrelated scatter between the measured luminosity and FWHM, and bias the virial mass estimates more. Hence it is very plausible that $\\sigma_L\\ga 0.4$ in most cases, which then accounts for $\\ga 0.2$ dex in $\\sigma_{\\rm vir}$ (e.g., Eqn.\\ \\ref{eqn:sig_vir}).  Therefore we conclude that the mass bias discussed here contributes at least $0.2-0.3$ dex at $L_{\\rm bol}\\ga 10^{46}\\ {\\rm erg\\,s^{-1}}$, comparable to the statistical bias resulting from the cosmic scatter in the BH scaling relations \\citep[e.g.,][]{Lauer_etal_2007b}. In Fig.\\ \\ref{fig:bias} we also show the ``observed'' BH mass offset (excess at fixed galaxy properties) from the literature. We note that the amount of the uncorrelated scatter $\\sigma_L$ might increase with redshift both due to the switch from the \\hbeta\\ line to the more problematic UV lines and due to the often decreased spectral quality at high redshift. Since the virial mass bias is independent of the Lauer \\etal\\ bias, the two statistical biases together can account for a large (or even full) portion of the BH mass excess seen in the data, hence no exotic scenarios are needed to explain this strong redshift evolution. It is possible, however, that there is still a mild evolution in these BH scaling relations, which is inherent to the cosmic co-formation of SMBHs and their hosts. But given these statistical biases, and given other systematics with single-epoch virial estimators \\citep[not discussed here, but see, e.g.,][]{Denney_etal_2009}, it is premature to claim a strong positive evolution in the BH scaling relations."
    },
    "0911/0911.2428_arXiv.txt": {
        "abstract": "{There is a lack of state-of-the-art information on very young open clusters, with implications for determining the structure of the Galaxy.} {Our main objective is to study the timing and location of the star formation processes which yielded the generation of the giant HII region Sh 2-284. This includes the determination of different physical variables of the stars, such as distance, reddening, age and evolutionary stage, including pre-main-sequence (PMS) stars.} {The analysis is based on $UBVR_CI_C$ CCD measurements of a field of 6.5'$\\times$6.5' containing the cluster, and JHK$_s$ photometry in the 3.5'$\\times$3.5' subfield, centered in the apparent higher condensation. The determination of cluster distance, reddening and age is carried out through comparison with ZAMS, post-MS and PMS isochrones. The reference lines used are obtained from theoretical post-MS and PMS isochrones from the Geneva and Yale groups, for metalicity Z=0.004, in agreement with the spectroscopic metallicity determination published for several cluster members.} {The results amount to E(B-V)=0.78$\\pm$0.02, DM=12.8$\\pm$0.2 (3.6~kpc), LogAge(yr)=6.51$\\pm$0.07 (3.2 Myr). The distance result critically depends on the use of low metallicity ZAMS and isochrones. A PMS member sequence is proposed, with an age of LogAge(yr)=6.7$\\pm$0.2 (4.7 Myr) which is therefore coeval within the errors with the post-MS cluster age. The mass function for this population in the mass range above 1.3-3.5 M$_\\odot$ is well fitted by a Salpeter mass function. The presence of a different star generation in the cluster with a distinctly older age, around 40 Myr, is suggested. On the other hand, the NIR photometry results indicate a large number of sources with (H-K$_s$) excess, practically distinct from the optical PMS candidate members. } {The analysis of our deep $UBVRIJHK$ photometry of Dolidze 25 therefore reveals a young cluster with coeval MS and PMS populations of age 3.2-5 Myr. In addition, a distinctly older cluster member population of age 40 Myr is suggested. The distance determined for the cluster from quantitative fits to ZAMS and isochrones is distinctly lower than previously published values. This result originates in the consistent use of low metallicity models for ZAMS fitting, applying published metallicity values for  the cluster} ",
        "introduction": "The open cluster Dolidze~25 (06h 45m 06s, +00$^{\\circ}$ 18' 00'', Epoch 2000) is one of the targets in our long term project on the search for and characterization of PMS members in young open clusters \\citep[and references therein, DAY-I in the following]{del07}. The cluster was included in their series of photoelectric studies by \\citet{mof75} and \\citet{mof79}. A UBV CCD photometric study was more recently published \\citep{tur93}.  In addition to the general objectives of this project (DAY-I), Dolidze~25 presents in principle two additional features of interest. It is located in the Galactic anticenter direction, and has been claimed to have a metallicity distinctly lower than the solar value \\citep{len90,fit92}. The cluster is therefore of special interest in our project, because of the expected finding of PMS stars with low metallicity and possibly at large Galactocentric distances.  The low metallicity of the cluster would contradict the value expected from the published values of the Galactic metallicity gradient. Any of the published values obtained from B-type stars, between -0.05 to -0.07~dex/kpc \\citep{fu09}, would require too large Galactocentric distances for our cluster to reach the \\cite{len90} metallicity value. On the other hand, recently published results \\citep{prza08,prz08}, with an improved method of spectroscopic analysis of B-type stars, suggest a clearly tighter result for abundances of B-type stars in the solar neighbourhood than those found in previous works. The accurate abundances could be considered to define the zero point of B-type stars metallicity at the solar Galactocentric distance, which would imply necessary corrections of previous values towards higher metallicities. We note however, that local metallicity variations with respect to an average gradient have been reported several times \\citep{fu09,ped09}, and the possible presence of discontinuities and slope changes, especially for large Galactocentric distances are still open. In this context, the analysis of cluster parameters, and the eventual improvement of their values is of importance. The cluster was selected as target for one of the COROT\\footnote{http://earth-sciences.cnes.fr/COROT/} Additional Programs \\citep{rip06}. These authors performed VIMOS observations at VLT which revealed the presence of a large number of emission line objects in the region (partial report ftp://ftp.na.astro.it/pub/astrows06/presetazione-FCUSANO.ppt at http://www.na.astro.it/; http://earth-sciences.cnes.fr/COROT/A-corot-week.htm).  Finally, the cluster is located in the center of the giant H-alpha bubble which encompasses the HII region Sh 2-284 (shortened to S284 in the following). Thus the analysis of the stellar content of this field could help to better understand the connection between the star formation processes and the physics of the ionized gas. The S284 region has been the object of a mid-infrared study with Spitzer IRAC bands \\citep{pug07}. The analysis of these observations \\citep[P09 in the following]{pug09} includes Dolidze~25, and addresses the properties of massive star formation in the region.  \\defcitealias{del07}{DAY-I} \\defcitealias{pug09}{P09}  In the present study we report UBVRIJHK observations centered in the Dolidze~25 cluster. In the next section we give an overview of the region. In Sect.~\\ref{obs} we describe the different data sets, with reduction procedures. Section~\\ref{ubvri-results} presents the optical results and a discussion of the determination of cluster parameters and cluster membership. Section~\\ref{jhk-results} presents the near-IR results and IR-excess sources. Section~\\ref{disc} contains the discussion of the results, and finally the last section resumes the main conclusions. ",
        "conclusions": "\\label{sum}  \\begin{itemize}  \\item We obtained deep UBVRIJHK photometry in the central 6.5' part of the Dolidze~25 cluster, reaching the 10$\\sigma$ limiting magnitudes of $V$ = 23.3, $J$ = 19.7, and $K_S$ = 18.6. Magnitudes, colours, and positions of all sources, as well as possible membership assignments, are published in Table~\\ref{tab-1}, available on-line.  \\item The membership analysis revealed 214 candidate members, of which 35 are main-sequence or post-MS stars, and the rest are PMS candidates. The PMS candidates comprise optically selected sources as well as sources with IR excess.  \\item The use of low metallicity ZAMS and isochrones to estimate distance, gives a new distance for the cluster to be 3.6~kpc.  \\item The ages of the optically selected PMS stars are likely around 5 Myr and those of the IR-excess sources possibly slightly younger.  \\item We find a possible existence of two generations of member stars, one with MS and PMS stars of age below 5 Myr, and another older population of age around 40 Myr. The spatial distribution of the two potential generations is found to be distinct, giving independent support to this suggestion.  \\item Dolidze~25 seems to be located in the centre of an ionized bubble which probably originated from the older generation of stars. We cannot exclude that the young cluster generation is located in the near side of the bubble shell, though.  \\end{itemize}  In the previous presentation of results we have described the presence of post-MS cluster members at similar luminosity but different evolutionary states, as inferred from their color indices. At the same time, we detect the presence of MS candidate members with masses below the presently PMS stars. This evidence adds to the indication of the presence in the field of a wide range in distances, with foreground and background groupings of stars (see Fig.~\\ref{fig-cm}) producing the general view of a region where star formation has been going on for a long time, both at different places of the H{\\sc ii} region, and in the same place at different epochs. At present, we observe this mixture of populations, with a star forming process which takes place in an already formed cluster, giving rise to the presence of two star generations of different age and being members of the same cluster. The evolved state of some of the most luminous cluster members, together with the presence of low mass MS member stars, can indeed be interpreted in this context. We can not exclude, though, that the younger population is seen in projection in front of the older one, possibly being formed in the shell of the bubble. Finally, we recall that the interpretation of two generations of member stars is furthermore favoured from inspection of the spatial distribution."
    },
    "0911/0911.5514_arXiv.txt": {
        "abstract": "From photometric observations of $\\sim$ 47,000 stars and spectroscopy of $\\sim$ 11,000 stars, we describe the first extensive study of the stellar population of the famous Double Cluster, h and $\\chi$ Persei, down to subsolar masses.  By analyzing optical spectra and optical/infrared photometry, we constrain the distance moduli (dM), reddening (E(B-V)), and ages for h Persei, $\\chi$ Persei and the low-density halo population surrounding both cluster cores.  With the exception of mass and spatial distribution, the clusters are nearly identical in every measurable way.   Both clusters have E(B-V) $\\sim$ 0.52--0.55 and dM = 11.8--11.85; the halo population, while more poorly constrained, likely has identical properties. As determined from the main sequence turnoff, the luminosity of M supergiants, and pre-main sequence isochrones, ages for h Persei, $\\chi$ Persei and the halo population all converge on $\\approx$ 14 Myr, thus showing a stunning agreement between estimates based on entirely different physics.   From these data, we establish the first spectroscopic and photometric membership lists of cluster stars down to early/mid M dwarfs.  At minimum, there are $\\sim$ 5,000 members within 10' of the cluster centers, while the entire h and $\\chi$ Persei region has at least $\\sim$ 13,000 and as many as 20,000 members.  The Double Cluster contains $\\approx$ 8,400 M$_{\\odot}$ of stars within 10' of the cluster centers. We estimate a total mass of at least 20,000 M$_{\\odot}$.  We conclude our study by outlining outstanding questions regarding the past and present properties of h and $\\chi$ Persei. From comparing recent work, we compile a list of intrinsic colors and derive a new effective temperature scale for O--M dwarfs, giants, and supergiants. ",
        "introduction": "Known and studied for the past two centuries and perhaps since ancient times, the Double Cluster -- h and $\\chi$ Persei -- presents a rare opportunity to precisely study multiple stages of stellar evolution.  The Double Cluster is located in the Perseus spiral arm and contains an exceptionally high density of evolved stars -- red supergiants and B giants/supergiants -- and early B-type dwarfs, indicating that the cluster is young, $\\lesssim$ 30 Myr old \\citep[cf.][]{Humphreys1978}.  This huge sample of high-mass stars probes post-main sequence stellar evolution, specifically the main sequence turnoff and the evolution of M supergiants.  For any reasonable initial mass function (e.g. Miller-Scalo), h and $\\chi$ Persei must also have thousands of solar/subsolar-mass members, which allow probes of pre-main sequence evolution.  Derived post main sequence and pre-main sequence \\textit{ages} for stars in h and $\\chi$ Persei are an acid test for stellar evolution models and critically affect conclusions about planet formation.  Because all stars in Gyr-old clusters have reached the main sequence, their ages are derived from comparing the main sequence turnoff and the luminosities of giants/supergiants with predictions from post-main sequence isochrones. In young clusters, high-mass stars are typically too few in number to derive post-main sequence ages.  Thus, ages for young, $\\lesssim$ 50 Myr-old clusters are almost always derived from fitting pre-main sequence isochrones to color-magnitude diagrams \\citep[e.g.][]{Mayne07}. However, if post and pre-main sequence age estimates for young, populous clusters significantly disagree, it is not clear that either estimate for any cluster should be considered accurate.  Furthermore, the lack of a reliable absolute age calibration impedes attempts to use circumstellar disk population statistics for different clusters to constrain planet formation empirically and prevents strong comparisons with solar system chronology and models of planet formation \\citep[e.g.][]{He07, Currie2010, jcast07, KenyonBromley2009}.  An extensive study of h and $\\chi$ Persei from early type, high-mass stars to late type, low-mass stars also addresses issues specific to the Double Cluster.  Chief among these issues is membership: formal lists are currently limited to B type dwarfs, giants, and supergiants; A--K supergiants; some stars showing evidence for circumstellar gas accretion or luminous debris disk emission; and x-ray luminous stars \\citep[][]{Uribe2002, Sl02, Bk05, Cu07b, Cu07c, Cu08a, CurrieEvans2009}.  A more uniform membership list would yield better estimates for the cluster mass function, constrain the spatial distribution of stars, and provide a reliable census from which to constrain models of star and planet formation.  In this paper, we describe the first exhaustive photometric and spectroscopic survey of h and $\\chi$ Persei down to subsolar masses.  Comprised of optical photometry for $\\sim$ 47,000 stars and spectra for $\\sim$ 11,000, our survey strongly constrains cluster properties -- reddening, distance, age, and structure -- and produces the first membership list for the Double Cluster that includes main sequence and pre-main sequence stars.  \\S 2 describes our data acquisition, image processing, and basic photometry/spectroscopic analysis.  \\S 3 yields distance and pre/post-main sequence age estimates for each component of the Double Cluster: the h Persei core, the $\\chi$ Persei core, and the low-density halo population surrounding both cores.  In \\S 4, we use the results of previous sections to identify h and $\\chi$ Persei members based on spectroscopy and identify probable members based on photometry. With a membership catalog, we investigate other cluster properties in \\S 5 such as the spatial distribution of members, the cluster mass function, and mass segregation.  In \\S 6, we summarize our results, discuss outstanding questions regarding h and $\\chi$ Per, and suggest fruitful future research programs.  The Appendix lists a new effective temperature scale and intrinsic colors for O--M dwarfs, giants, and supergiants. ",
        "conclusions": "\\subsection{Summary of Analysis and Major Findings} This paper describes the first extensive photometric and spectroscopic survey of main sequence and pre-main sequence stars in h and $\\chi$ Persei. By analyzing optical photometry for 47,000 stars and spectroscopy of 11,000 stars, we derived the median reddening, distance modulus, and age for each component of h and $\\chi$ Per.  We then constructed the first extensive membership list of h and $\\chi$ Per stars, ranging from massive B--M supergiants to pre-main sequence stars whose masses are likely comparable to the hydrogen burning limit.  Our study yields the following major results:  \\begin{itemize} \\item The median reddening values for the h Persei core, $\\chi$ Persei core, and halo region are comparable: E(B-V) $\\sim$ 0.55, E(B-V) $\\sim$ 0.52, and E(B-V) $\\sim$ 0.52, respectively.  \\item The distance moduli for h Per, $\\chi$ Per, and the halo region are also nearly identical: dM$_{h Per}$ = 11.8, dM$_{\\chi Per}$ = 11.85, and dM$_{halo}$ = 11.85.  \\item Ages for all three components of h and $\\chi$ Persei are identical: $\\sim$ 14 Myr. Moreover, post-main sequence ages and pre-main sequence ages for each component are identical. Thus, the properties of h and $\\chi$ Persei are consistent with the coeval cores and the low-density halo population emerging from a single, explosive star-forming event.  \\item Within 10' of the two cluster centers, the Double Cluster contains at least $\\sim$ 5,000 stars; h Persei is about 30\\% more populous than $\\chi$ Persei.  The halo region contains at least $\\sim$ 7,000 stars and as many as 15,000 stars, bringing the total number of stars in h and $\\chi$ Persei to $\\sim$ 20,000.  The estimated masses for h Persei, $\\chi$ Persei, and the halo region are $\\sim$ 4,700 M$_{\\odot}$, $\\sim$ 3,700 M$_{\\odot}$, and $\\sim$ 11,000 M$_{\\odot}$.  \\item Both clusters show clear evidence of mass segregation within 3' of the cluster centers, though it is stronger for h Persei.  The relaxation timescales for both clusters are comparable to or less than their ages within 3'.  Therefore, mass segregation may either be primordial or dynamical. \\end{itemize}  These results support and extend recent studies of h and $\\chi$ Per conducted by \\citet{CurrieEvans2009}, \\citet{Mayne07}, \\citet{Cu07a}, \\citet{Bk05}, \\citet{Sl02}, and \\citet{Ke01}.  Remarkably, four mutually exclusive sets of authors using different techniques converge on essentially the same properties for the clusters' reddening, distance, and age. Combined with previous work, our results clearly refute earlier claims that the Double Cluster and its environs have a substantial age spread \\citep{Wildey1964, Marco2001}, have a substantially different age \\citep{Schild1965, Schild1967, Marco2001}, or are located at substantially different distances \\citep{Schild1967, Kharchenko2005}.   All indications are that the different components of h and $\\chi$ Persei substantially differ only by spatial distribution and mass.  \\subsection{Future Research on the Stellar Population of h and $\\chi$ Persei} In addition to advancing our understanding of h and $\\chi$ Persei properties, this study clearly identifies advances needed to paint a more complete picture of the Double Cluster.  We highlight several promising areas of future research below:  \\begin{itemize} \\item \\textbf{Wide-Field Optical Photometry/Deep Optical Spectroscopy} -- Simple wide-field optical photometric surveys can easily expand our membership list and amplify the scientific output of our work.  The mismatch in survey area between our spectroscopy ( $\\sim$ 1 square degree) and photometry ($\\sim$ 0.6 square degrees) leaves $\\sim$ 3500 stars with spectroscopy but no photometry.  Our results demonstrate the need for high-quality photometry to measure reddening and to establish membership.  Because nearly all of these stars are bright (V $\\lesssim$ 19), they are easily accessible by wide-field cameras on 1--2 meter-class telescopes.  The upper main sequence of halo stars is undersampled largely because of our selection criteria for Hectospec targets (J $\\ge$ 14).  Added to our existing survey, the entire upper main sequence of h and $\\chi$ Persei can probably be probed after a few additional Hectospec or Hydra fiber settings. Because these stars are extremely bright, obtaining high signal-to-noise spectra will be trivial.  Furthermore, our optical photometry identifies many probable cluster members that are beyond the 2MASS detection limit but clearly within range of Hectospec for integration times of $\\sim$ 1--2 hours (V $\\lesssim$ 21--22). A survey establishing the spectral types for many of these stars would provide valuable information on the luminosity and temperatures of pre-main sequence stars at $\\approx$ 14 Myr. Other multi-object spectrographs coming online in the near future (e.g. MODS on the Large Binocular Telescope, Binospec on the MMT) can probe even further down the cluster mass function.  \\item \\textbf{Robust Membership Determinations from \\textit{SIM} and \\textit{GAIA}} --  Proper motion studies are the most definitive method for determining h and $\\chi$ Per membership.  Unfortunately, at $\\sim$ 2.3 kpc distant, h and $\\chi$ Per stars exhibit tiny proper motions, the size of which render ground-based campaigns to verify our membership list hopeless. However, space-based interferometric missions -- specifically \\textit{SIM} and \\textit{GAIA} -- are easily capable of detecting the $\\sim$ $\\mu$-arc second motions of pre-main sequence cluster stars.  \\textit{SIM's} and \\textit{GAIA's} clear ability to yield definitive catalogs of members for h and $\\chi$ Persei and other distant, populous clusters would have an enormous impact on studies of these clusters and stellar evolution in general.  \\item \\textbf{Metallicity} -- Clearly, the greatest source of systematic error in constraining the clusters' properties is metallicity.  While we argue that h and $\\chi$ Per likely have a near-solar metallicity, it is possible that the sources whose properties support our contention are not indicative of the clusters' stars as a whole.  To address metallicity, at least three avenues of research should be explored.  First, sophisticated stellar atmosphere modeling of stars in a variety of evolutionary states should constrain the photospheric chemical abundances \\citep[e.g.][]{Dufton1990, Smartt1997, Venn2002} and help to define the typical chemical composition.  Second, high-resolution echelle spectra of cluster stars (e.g. with Hectochelle), especially G-type pre-main sequence members, should not only constrain the clusters' velocity dispersions (and thus aid membership identification) but will also be capable of yielding metallicity estimates.  Third, building upon work by \\citet{Southworth2004,Southworth2005}, metallicities can be estimated by comparing masses and radii of eclipsing binaries to model predictions. In this regard, results from the MONITOR program \\citep{Ai07} may be crucial for ending the debate on metallicity.   \\item \\textbf{Effective Temperatures} -- As described in Appendix A, the T$_{e}$ scale for early B dwarfs and evolved B stars is more poorly constrained than for O stars and A--M stars. Non-LTE modeling of the many early B stars in h and $\\chi$ Persei may provide better constraints on the T$_{e}$ scale.  More specifically, determining T$_{e}$ for individual B stars strengthens our ability to derive accurate post-main sequence ages.  As was argued by \\citet{Sl02} and can be seen in Figures \\ref{msturnoffchi}--\\ref{msturnoffhalo}, the binning of B stars into discrete T$_{e}$ values for their spectral type potentially introduces a spurious spread in T$_{e}$ at a given V magnitude brightness, which could be misinterpreted as a spread in age.  Individual measurements for T$_{e}$ will tighten the locus of cluster stars at the main sequence turnoff.  Accurate T$_{e}$ estimates may be be especially important for cluster Be stars, which are often difficult to spectral type because of their Balmer line emission.  A large campaign to derive atmospheric properties of Be stars in h and $\\chi$ Persei is underway \\citep{Marsh2010}.  \\item \\textbf{The Halo Population} -- While our results confirm previous claims \\citep[e.g.][]{Cu07a} that h and $\\chi$ Persei contains a large halo population, it is possible that the halo extends beyond  25--30' from either cluster center, or even well beyond the square-degree field covered by our spectroscopic survey. A more spatially extended survey of stars would yield a better estimate for the total mass of the halo population.  A more complete census of young stars within $\\approx$ 5 degrees of the h and $\\chi$ Persei cores better reveals the large star formation history within which h and $\\chi$ Per emerged, including the Double Cluster's relationship to other nearby young clusters such as W3, W4, and W5 located several degrees away.  Specifically, it may be possible to address the relationship between h and $\\chi$ Persei and the Perseus OB1 association, determining if the core and halo populations have properties completely distinct from Per OB1, if h and $\\chi$ Per represents a unique epoch in star formation that propogated through the present day Per OB1 region, or if its properties are characteristic of the Per OB1 association as a whole \\citep[e.g.][]{Sl02, Lee2008}.   \\item \\textbf{The Circumstellar Disk Population: Constraints on Planet Formation} -- Finally, the large sample of stars now confirmed as h and $\\chi$ Persei members can be used to investigate planet formation by using Spitzer and (later) the James Webb Space Telescope. New Spitzer observations reveal an order-of-magnitude increase in the number of stars detected at 3.6--8 $\\mu m$ \\citep{Currie2010}.  Combined with membership information derived here, Spitzer data will yield constraints on the disk population from a sample size easily dwarfing that of the \\textit{FEPS} and \\textit{Cores-to-Disks} Legacy Programs combined and thus providing far superior statistical reliability.  This fundamentally different kind of dataset allows robust probes of planet formation that are impossible with other samples, providing estimates of the frequency of long-lived protoplanetary disks; warm, terrestrial planet-forming debris disks; and extremely luminous cold debris disks. Combined with data from 5--25 Myr-old clusters -- e.g. NGC 2362, Upper Scorpius, Orion OB1, NGC 1960, and NGC 2232 \\citep[][]{CurrieLada2009, Carpenter2006, Hernandez2007,Balog2010, Currie2008b} -- it is possible to reconstruct the time history of terrestrial, gas giant, and icy planet formation to compare with our solar system's chronology and models of planet formation \\citep[e.g.][]{KenyonBromley2009, jcast07}. \\end{itemize}"
    },
    "0911/0911.0936_arXiv.txt": {
        "abstract": "The formation scenarios for {\\it single} low-mass ($M<0.45M_{\\odot}$) white dwarfs include enhanced mass loss from a metal-rich progenitor star or a common envelope phase of a solar-like star with a close-in massive planet or a brown dwarf. Both scenarios suggest that low-mass white dwarfs may have planets. Here, we present a {\\em Spitzer} IRAC search for substellar and planetary mass companions to 14 low-mass white dwarfs. One of our targets, HS 1653+7753, displays near- and mid-infrared flux excess. However, follow-up MMT observations show that this excess is due to a nearby resolved source, which is mostly likely a background object. Another target, PG 2257+162, shows flux excess compatible with a late-type stellar companion. We do not detect substellar companions to any of the remaining targets. In addition, eight of these stars do not show any radial velocity variations, ruling out stellar mass companions including other white dwarfs. We conclude that a significant fraction of the low-mass white dwarfs in our sample do not have stellar or massive brown dwarf companions. ",
        "introduction": "The fate of planetary systems around Sun-like stars can be studied by finding planets around their white dwarf (WD) remnants. WDs are physically small ($R\\sim0.01R_\\odot$), which sharply limits their luminosity in the infrared, where the planet spectrum peaks \\citep{burleigh02,ignace01}. Compared to main-sequence (MS) stars, the significant gain in contrast makes WDs excellent targets for photometric searches for planetary companions. However, to this day, no planet detection around WDs has been reported. {\\em Spitzer} IRAC imaging surveys by \\citet{mullally07}, \\citet{debes07}, \\citet{farihi08}, and \\citet{kilic09b}, and near-infrared imaging surveys by \\citet{friedrich06}, \\citet{burleigh08}, and \\citet{hogan09} were unsuccessful in finding planetary mass companions to WDs with limits typically between 5 and $10 M_J$. Here we focus on using low-mass WDs to provide additional leverage for detecting planetary mass companions.   \\subsection{CURRENT CONSTRAINTS}  Timing measurements of pulsating WDs is an alternative, powerful method to search for planets \\citep{winget03}, but this method can be used only for a limited sample of stars. There is a promising signature of a $M\\sin i= 2.4 M_J$ planet in the pulsation timing measurements of the WD GD 66 \\citep{mullally09}. However, a full orbit has not been observed, and this detection remains provisional (F. Mullally 2009, priv. comm.).  The discovery of circumstellar debris disks around more than a dozen WDs imply that 1\\% to 3\\% of WDs with cooling ages $<$ 500 Myr have disks \\citep[see][and references therein]{farihi09}. These disks most likely formed by tidally disrupted asteroids \\citep{jura03}. Hence, there is indirect evidence that at least a few percent of WDs have remnant planetary systems. Of course, the observed debris disk fraction may represent only the planetary systems that become unstable during post-MS evolution, and the true fraction of WDs with planets may be much higher.  The reason for the lack of known planetary companions to WDs may be a target selection effect \\citep{kilic07b}. \\citet{fischer05} find that less than 3\\% of the stars with $-$0.5 $<$ [Fe/H] $<$ 0.0 have Doppler-detected planets, whereas 25\\% of the stars with [Fe/H] $>$ +0.3 have gas giant planets. Based on the age-metallicity relation derived by \\citet{reid07}, \\citet{kilic07b} find that only 3\\% of the stars in the solar neighborhood have had [Fe/H] $>$ +0.3 in the past 10 Gyr. Therefore, a random selection of WDs would be dominated by WDs that come from low metallicity progenitors. The best way to find planetary companions to WDs may be to search for companions to WDs that form from metal-rich stars.  \\subsection{LOW-MASS WHITE DWARFS}  Low-mass helium-core WDs ($M < 0.45\\ M_\\odot$) can be produced from interacting binary systems, and traditionally all of them have been attributed to this channel. However, a low-mass WD can also result from a single star that experiences severe mass loss on the first ascent giant branch. \\citet{kalirai07} identified several low-mass WDs in the upper part of the WD cooling sequence of the metal-rich cluster NGC 6791, and they argued that there is a mechanism in clusters to produce low-mass WDs without requiring binary star interactions. However, this argument is now challenged by \\citet{bedin08}, who argue that the bright peak observed in the WD luminosity fuction of this cluster is better explained as binary WDs. \\citet{bedin08} also argue that the low-mass WDs found by \\citet{kalirai07} can be explained as the product of binary star evolution.  Unlike the relatively faint low-mass WDs in NGC 6791, it is possible to obtain radial velocity measurements of low-mass WDs in the field and test the binary formation scenario. Around half the low-mass WDs observed by \\citet{maxted00} lack detectable radial velocity companions. The ESO SNIa Progenitor Survey (SPY) searched for radial velocity variations in more than a thousand WDs using the Very Large Telescope \\citep{napiwotzki01}. They found that about 58\\% of the low-mass WDs in their sample do not show radial velocity variations \\citep{napiwotzki07}. Their detailed analysis of the detection probability of MS, WD, and brown dwarf companions show that their observations of low-mass WDs are inconsistent with the idea that all of them reside in close binary systems. The SPY observations are less sensitive to brown dwarf companions, because of their smaller mass, but even in this case 100\\% binary fraction rate can be ruled out (R. Napiwotzki 2007, private communication). \\citet{kilic07b} argue that these single low-mass WDs might come from old metal-rich stars that truncate their evolution prior to the helium flash from severe mass loss. The fraction of single low-mass WDs among the field WD sample is also compatible with the fraction of metal-rich dwarfs among the field dwarf population.  A strong implication of the \\citet{kilic07b} scenario is that single low-mass WDs are good targets for planet searches because they may arise from metal-rich progenitors. This leads to an inversion of prior strategies for seeking planetary companions to WDs, which avoid low-mass WDs on the grounds that they are the products of binary interactions. Instead, some of these objects may be the descendants of metal-rich stars that are known to have a high frequency of planet detection. If the progenitors of these single low-mass WDs have planets, the planets are more likely to survive the post-MS evolution, as these He-core WDs do not go through the asymptotic giant branch phase.  An alternative formation scenario for single low-mass WDs is proposed by \\citet{nelemans98}. In this scenario, low-mass WDs form from Sun-like stars with close planetary companions. A common envelope phase with a massive planet can expel the stellar envelope and form a low-mass WD. The planet may (not) survive the common envelope phase. However, if there are additional planets in wide orbits, they are likely to survive \\citep{livio83,burleigh02,duncan98}.  Both scenarios proposed by \\citet{nelemans98} and \\citet{kilic07b} suggest that low-mass WDs may have planets. Here we present a {\\em Spitzer} IRAC search for substellar companions to 14 low-mass WDs. Our sample selection and observations are discussed in Section 2, whereas the detectability of planets around our targets and the results of our search are discussed in Section 3 and 4. ",
        "conclusions": "Our IRAC survey of 14 low-mass WDs did not reveal any substellar or planetary mass companions. One of our targets, HS 1653+7753, has a nearby source that is likely to be a background object. Another target, PG 2257+162, possibly has a low-mass stellar companion. In addition, the IRAC photometry for one of our targets, WD 0950$-$572 is affected by a nearby star. The remaining 11 low-mass WDs do not show any mid-infrared flux excesses. Adding the three low-mass WDs observed by \\citet{hansen06} and \\citet{mullally07} to our sample brings the sample size to 14 WDs with no detectable mid-infrared flux excess. Eleven of these targets do not show radial velocity variations, ruling out companions more massive than $\\approx0.08M_\\odot$ \\citep{maxted00}. Therefore, these 11 low-mass WDs appear to have no stellar (including another WD) or substellar companions to the limits given in Table 3. The non-detection of stellar, substellar, or planetary companions to our targets is not inconsistent with the idea that single low-mass WDs descend from single metal-rich stars through extreme mass loss.  \\citet{nelemans98} predict that if a substellar object with $M<32 M_J$ goes through a common-envelope phase with a star, it will either evaporate or spiral-in during the process. More massive objects are expected to survive this evolution, as evidenced by WD 0137$-$049B \\citep[$0.053~M_\\odot$,][]{maxted06}, although the theoretical 32 $M_J$ limit needs to be tested by observations. \\citet{nelemans98} use WD 1614+136 to demonstrate that their common envelope phase scenario is feasible. They assume that WD 1614+136 is a remnant of a 1 $M_{\\odot}$ star with a core-mass of 0.33$M_{\\odot}$. They find that the minimum companion mass required to expel the envelope of the progenitor is 21 $M_J$. No companion is detected in IRAC observations of this star, and \\citet{hansen06} rule out $M\\geq20M_J$ companions for ages $\\leq1$ Gyr. This limit goes up to about 40 $M_J$ for 5 Gyr old companions. A well tuned scenario involving a common envelope phase between a $\\sim$5 Gyr old 1 $M_{\\odot}$ star and a 21-40 $M_J$ brown dwarf at a certain orbital separation can explain the single, low-mass WD 1614+136.  Our survey is sensitive to massive brown dwarf companions around the majority of our targets even at 5 Gyr (see Table 3). We rule out $M>40 M_J$ companions around 8 of our targets at 5 Gyr. Like WD 1614+136, only a well tuned common-envelope phase scenario involving 20-40 $M_J$ companions would explain the low mass WDs in our sample. \\citet{grether06} find that the driest part (the minimum number) of the brown dwarf desert occurs at 31 $M_J$. Thus, it seems unlikely that the scenario proposed by \\citet{nelemans98} will explain all single, low-mass WDs. However, given the caveats in understanding the common-envelope phase evolution and our age dependent limits on possible companions, we cannot rule out this scenario.  An alternative scenario for the formation of single low-mass WDs involves binary mergers of even lower mass WDs \\citep{iben97}. \\citet{maxted98} find that such mergers result in rotational velocities on the order of 1000 km s$^{-1}$, unless the WDs are efficient in loosing angular momentum through mass loss during the merging process. \\citet{maxted98} measure rotational velocities of 50 km s$^{-1}$ for two single low-mass WDs and they suggest that efficient angular momentum loss should take place if these WDs formed through binary mergers. \\citet{nelemans98} find this formation scenario unlikely. Currently, there are no known binary low-mass WD progenitor systems with combined mass $\\approx0.4M_{\\odot}$ and with merger time shorter than a Hubble time \\citep[see Table 1 in][]{nelemans05}. Radial velocity studies of extremely low-mass ($0.2M_{\\odot}$) WDs are therefore important to search for such systems \\citep{kilic07a,kilic09a} and to see if this scenario is likely to explain at least a fraction of single, low-mass WDs."
    },
    "0911/0911.3809_arXiv.txt": {
        "abstract": "An \\xmm\\ observation performed in  May 2008 has confirmed that the 13 seconds pulsations in the X-ray binary \\hr\\ are due to a rapidly rotating white dwarf. From the pulse time delays induced by the 1.55 days orbital motion, and the system's inclination, constrained by the duration of the X-ray eclipse discovered in this observation, we could derive a mass of 1.28$\\pm$0.05 $\\msole$ for the white dwarf. The future evolution of this post common envelope binary system will likely involve a new phase of mass accretion through Roche-lobe overflow  that could drive the already massive white dwarf above the Chandrasekhar limit and produce a Type Ia supernova. ",
        "introduction": "White dwarfs have masses in a narrow range centered at about 0.6 $\\Msole$ (see, e.g., \\cite{kep07}), but a few examples of massive white dwarfs ($>$1.2 $\\Msole$) have been reported \\cite{dah04,ven08}. Massive white dwarfs in binary systems, where mass transfer can occur, are particularly interesting since they are good candidates for the formation of Type Ia supernovae.  Using data obtained with the \\xmm\\ satellite, we have recently found compelling evidence \\cite{mer09} for the presence of a  white dwarf with mass $>$1.2 $\\Msole$ in the X--ray binary system \\hr . In the next section we describe the properties of this peculiar binary, composed of a 13 s X--ray pulsar orbiting a hot subdwarf star. We then briefly review the new \\xmm\\ results and discuss some of their implications. ",
        "conclusions": "The above findings indicate that \\rx\\ is one of the most massive white dwarfs currently known  and the one with the shortest spin period. We note that most white dwarf masses are derived with indirect methods, such as surface gravity estimates obtained by spectral modelling, or the measurement of gravitational redshift. These methods, beside depending on models and assumptions, provide combined information on mass and radius, hence the resulting masses cannot be used to determine independently the mass-radius relation. Our mass determination for \\rx\\ is directly based  on a dynamical measurement, thus making this star an ideal target to better constrain the white dwarfs equation of state. The mass lower limit, coupled  with the 13 s spin period,  allows us to set an upper limit of 6000 km on the white dwarf's radius, based only on simple rotational stability considerations \\cite{cha87}.   The high mass and fast spin of \\rx\\ likely result from a previous evolutionary phase in which the accretion of mass and angular momentum took place at a much larger rate than currently observed. The accretion of mass down to the white dwarf surface implies that the magnetic dipole is \\ltsima 10$^{29}$ ($\\Mdot$ / 10$^{-8}$ $\\Msole$ yr$^{-1}$)$^{15/16}$ (1350 km s$^{-1}$ / V$_{WIND}$)$^{15/4}$ Gauss cm$^{3}$,  where we have normalized the wind mass loss $\\Mdot$ and terminal velocity V$_{WIND}$ to the values measured for \\hd\\ \\cite{ham81}. A stronger magnetic field would prevent accretion through the onset of the propeller effect. Thus, \\hr\\ can be seen as  a white dwarf analogous of low mass X-ray binaries in which weakly magnetic neutron stars are spun-up to short spin periods and shine as recycled millisecond radio pulsars when accretion onto them stops \\cite{bha91}.   The future evolution of \\hd\\ will probably lead to a new phase of unstable mass transfer through Roche-lobe overflow \\cite{ibe94}, during which the accretion of helium-rich material  might push the already massive white dwarf above the  Chandrasekhar limit and trigger a Type Ia supernova explosion.      \\begin{theacknowledgments} We acknowledge partial support from contract ASI/INAF I/088/06/0 \\end{theacknowledgments}"
    },
    "0911/0911.3662_arXiv.txt": {
        "abstract": "{ A large fraction of the observed  brown dwarfs may form by gravitational fragmentation of unstable discs. This model reproduces the brown dwarf desert, and provides an explanation the existence of planetary-mass objects and for the binary properties of low-mass objects.  We have performed an ensemble of radiative hydrodynamic simulations and determined the statistical properties of the low-mass objects produced by gravitational fragmentation of discs. We suggest that there is a population of brown dwarfs loosely bound on wide orbits  ($100-5000$ AU) around Sun-like stars that surveys of brown dwarf companions should target. Our simulations also indicate that planetary-mass companions to Sun-like stars are unlikely to form by disc fragmentation.} ",
        "introduction": "\\label{intro} The formation of brown dwarfs (BDs) is not well understood (\\cite{whit2007}). Low-mass objects are difficult to form by gravitational fragmentation of unstable gas, as for masses in the BD regime ($\\stackrel{<}{_\\sim}~80~{\\rm M}_{\\rm J}$, where ${\\rm M}_{\\rm J}$ is the mass of Jupiter) a high density is required for the gas to be Jeans unstable. Thus, if BDs form according to the standard model of star formation, i.e. the collapse of a prestellar core,  then the pre-BD core must have a density of $\\stackrel{>}{_\\sim}10^{-16} {\\rm g\\ cm}^{-3}$.  \\cite{padoan} and \\cite{hennebelle} suggest that these high density cores can be formed by colliding flows in a turbulent magnetic medium. However, this model requires a large amount of turbulence, and has difficulty explaining the binary properties of BDs. Additionally, the large number of brown-dwarf mass cores that the theory predicts  have not been observed.  Another way to reach the high densities required for the formation of BDs is in gravitationally unstable discs (\\cite{whit06}, \\cite{stam07b}). These discs form around newly born stars and grow quickly in mass by accreting material from the infalling envelope. They become unstable if the mass accreted onto them cannot efficiently redistribute its angular momentum outwards in order to accrete onto the central star (\\cite{attwood09}).  This mechanism reproduces many of the observed statistical properties of  low-mass objects that are not satisfactorily explained by other theories, in particular the BD desert, the statistics of low-mass binary systems, and the formation of free-floating planets (\\cite{stam09a},  \\cite{stam09b},  \\cite{stam09c}).  BDs are also thought to form as ejected embryos from star forming regions, i.e. as by-product of the star formation process (\\cite{rei}).  In this model BDs form the same way as low-mass stars, i.e. in collapsing molecular cores,  but shortly after their formation they are ejected from their parental core and stop accreting any further material. Hence, they do not realise their potential to become hydrogen-burning stars.  The above BD formation mechanisms are not exclusive of each other and can work in conjunction in star forming regions. For example, turbulent fragmentation can produce cores which collapse to form discs, the discs may then fragment to produce BDs, and the BDs may avoid accreting additional mass by being ejected.  In this paper we focus on the mechanism of BD formation by fragmentation of gravitationally unstable discs. We discuss the  properties of the BDs produced with this mechanism and  the predictions of the model that can be tested by observations.   \\begin{figure*} \\centerline{ \\includegraphics[angle=-90,width=5.5cm]{stamatellos_fig1.ps} \\includegraphics[angle=-90,width=5.5cm]{stamatellos_fig2.ps} \\includegraphics[angle=-90,width=5.5cm]{stamatellos_fig3.ps} } \\caption{SPH radiative hydrodynamic simulation of the evolution of a $0.7\\,{\\rm M}_\\odot$ disc around a $0.7\\,{\\rm M}_\\odot$ star. The disc is represented using $2\\times 10^6$ SPH particles. 3 snapshots are presented. The disc is gravitationally unstable and can cool efficiently, hence it quickly fragments to form low-mass H-burning stars, BDs, and planetary-mass objects. These objects form at radii from $\\sim 100$ to $\\sim 300\\,{\\rm AU}$ but due to mutual interactions a large fraction of  them escape  from the disc.} \\label{fig:sim} \\end{figure*} ",
        "conclusions": "The mechanism of BD formation by fragmentation of gravitationally unstable discs reproduces the BD desert around Sun-like stars, and can explain the existence of free-floating planetary mass objects and the binary properties of low-mass objects. We suggest that there is a population of BDs on wide orbits around Sun-like stars that could be probed with future observations."
    },
    "0911/0911.3348_arXiv.txt": {
        "abstract": "We present our {\\it Spitzer} IRS  spectroscopic survey from 10$\\mu m$ to 37$\\mu m$ of the Seyfert galaxies of the 12$\\mu m$ Galaxy Sample, collected in high resolution mode (R$\\sim$ 600). The new spectra of 61 galaxies, together with the data we already published, gives us a total of 91 12$\\mu m$ Seyfert galaxies observed, out of 112. We discuss the mid-IR emission lines and features of the Seyfert galaxies, using an improved AGN classification scheme: instead of adopting the usual classes of Seyfert 1's and Seyfert 2's, we use the spectropolarimetric data from the literature to divide the objects into categories  \"AGN 1\" and \"AGN 2\", where AGN 1's include all broad-line objects, including the Seyfert 2's showing hidden broad lines in polarized light. The remaining category, AGN 2's contains only Seyferts with no detectable broad lines in either direct or polarized spectroscopy. We present various mid-IR observables, such as ionization-sensitive and density-sensitive line ratios, the PAH 11.25$\\mu$m feature and the H$_2$ S(1) rotational line equivalent widths, the (60$\\mu m$ - 25$\\mu m$) spectral index and the source extendedness at 19$\\mu m$, to characterize similarities and differences in the AGN populations, in terms of AGN dominance versus star formation dominance.  We find that the mid-IR emission properties characterize all the AGN 1's objects as a single family, with strongly AGN-dominated spectra. In contrast, the AGN 2's can be divided in two groups, the first one with properties similar to the AGN 1's except without detected broad lines, and the second with properties similar to the non-Seyfert galaxies, such as LINERs or starburst galaxies.  We computed a semianalytical model to estimate the AGN and the starburst contributions to the mid-IR galaxy emission at 19$\\mu m$. For 59 galaxies with appropriate data, we can separate the 19$\\mu m$ emission into AGN and starburst components using the measured mid-IR spectral features.  We use these to quantify the brightness thresholds that an AGN must meet to satisfy our classifications: AGN 1 have an AGN contribution  $\\geq$ 73 \\% and AGN 2 $\\geq$ 45 \\% of their total emission at 19$\\mu m$.  The detection of [NeV] lines turns out to be an almost perfect signature of energy production by an AGN.  Only 4 ($\\sim$ 7.5\\% percent) of 55 AGN 1 and 2 (10\\% percent) out of 20 AGN 2 do not have [NeV] 14.3$\\mu$m down to a flux limit of $\\sim 4 \\times 10^{-15} erg s^{-1}cm^{-2}$. We present mean spectra of the various AGN categories. Passing from AGN-dominated to starburst-dominated objects, the continuum steepens, especially at wavelengths shorter than 20$\\mu m$, while the PAH feature increases in its equivalent width and the high ionization lines decrease.  We estimate H$_2$ mass and excitation temperature through the measurement of the S(1) rotational line of this molecule. Finally we derive the first local luminosity functions for the brighest mid-infrared lines and the PAH feature at 11.25$\\mu m$. No statistical difference is apparent in the space densities for Seyfert 1's and 2's of a given line luminosity, nor for the new classes of AGN 1's and 2's. We use the correlation between [Ne V] line and nonstellar infrared continuum luminosity to derive the global output of accretion-powered galactic nuclei in the local universe. ",
        "introduction": "This paper contains the final results of the \\textit{Spitzer} high-resolution IRS spectroscopic survey of the sample of Seyfert galaxies (hereafter 12MSG) included in the {\\it IRAS} 12$\\mu$m galaxy sample  \\citep[hereafter RMS]{rms93}. In \\citet{tom08} (hereafter Paper I) we have presented and analyzed the first 30 high-resolution spectra of 29 Seyfert galaxies of this sample (one IRAS galaxy, Mrk 1034, was coincident with a pair, for which we obtained two spectra). The first spectroscopic observations of active galaxies with the Infrared Spectrometer \\citep[IRS]{hou04} onboard the {\\it Spitzer Space Observatory} \\citep{wer04} have been collected on classical AGNs \\citep{wee05} and ULIRGs \\citep{arm07}. After the work presented in Paper I and the referenced works therein, a few more studies have discussed the Spitzer mid-infrared spectra of Seyfert galaxies, among these \\citet{deo07}, \\citet{mel08a} and \\citet{wu09}. As expected, the mid-infrared spectra of Seyfert galaxies show forbidden lines originating in the Narrow Line Regions, excited by the AGN ionizing flux. This power is thought to be produced by black-hole accretion, i.e., ultimately from the conversion of gravitational into radiative energy. The fine structure lines of [NeV] at 14.32$\\mu$m and 24.31$\\mu$m originate exclusively in the highly ionized gas (with an ionization potential of 97eV) illuminated by the AGN\\footnote{ either currently or some time during the last several hundred years, because, even if the AGN ionizing continuum could be completely switched off, the photoionized NLR could still be detected, due to its large extension and the long recombination time}. As discussed in Paper I, the [OIV] line at 25.88$\\mu$m (ionization potential of 55eV) is most probably excited from the AGN. In fact, \\citet{mel08b} consider this line as an accurate and truly isotropic indicator of AGN activity, even if it could also originate from high excitation starburst emission or in shocks in low-metallicty starbursts \\citep{lut98}.  The [NeIII]15.55$\\mu$m line (Ne+ and Ne++ have ionization potentials of 14eV less than O++ and O+++ respectively) can be excited both from AGN activity \\citep{gor07} and from starbursts \\citep{tho00}. Superimposed on the AGN spectra, the lines of [NeII]12.81$\\mu$m, [SIII]18.71$\\mu$m and 33.48$\\mu$m and [SiII]34.82$\\mu$m originate in gas with moderate ionization and most of their emission is generated by young newly formed stars, even if some contribution from the AGN is also expected \\citep{sm92}. The relationship between the [OIV]25.88$\\mu$m, [NeIII]15.55$\\mu$m and [NeII]12.81$\\mu$m lines in an eterogeneous sample of Seyfert galaxies has been studied by \\citet{mel08a}, who found that Seyfert 1's and Seyfert 2's have different AGN and star formation contributions to the total emission.  The interstellar medium produces the pure rotational lines of molecular hydrogen, as already shown by the early results of \\textit{ISO} spectroscopy in \\citet{ge98} and \\citet{rig02}, respectively. \\citet{wu09} analyzed the IRS low resolution spectra of 103 Seyfert galaxies from the 12MSG and measured the policyclic aromatic hydrocarbons (hereafter PAH) emission features and the Silicate absorption strength. The PAHs have been proposed as star formation tracers by \\citet{pl89}, while the Silicate absorption is sensitive to heavy dust obscuration of the nucleus.  According to the simplest Unified Model for AGN, the Accreting Torus Model (ATM) \\citep{mgt}, Seyfert 1 and 2 galaxies are the same kind of objects, only viewed from different angles.  The strongest demonstration of this is detection via optical spectropolarimetry of the Broad Line Region emission--the defining characteristic of  Seyfert 1'--in a significant minority of Seyfert 2 galaxies \\citep{am85,ant93}. A different scenario postulates an evolutionary difference: that Seyfert 2 are the early stages of the transition of HII/Starburst galaxies into Seyfert 1's. Two suggested evolutionary progressions are HII $\\rightarrow$ Seyfert 2 \\citep{kau03,sb01}, or a fuller scenario of HII $\\rightarrow$ Seyfert 2 $\\rightarrow$ Seyfert 1 \\citep{hm99,kro02,lev01}. Because the radiation due to the star formation processes is roughly isotropic, the ATM predicts no observational difference in the star formation tracers between Seyfert 1's and 2's. If, on the other hand, star formation is stronger in Seyfert 2's, as suggested in \\citet{buc06}, then some evolution from Seyfert 2's to Seyfert 1's could be invoked. If general interstellar extinction towards the center of the galaxy is extremely high, optical data alone might not always provide the correct (intrinsic) classification.  Our sample is described in \\S2, the observations and the data reduction are briefly reported in \\S3, the direct results of our observations, the estimates of the H$_2$ masses and temperatures and the measure of the continuum extendedness are presented in \\S4, the diagnostic diagrams and the semianalytical models to interpret them are presented and discussed in \\S5. In \\S6 the [NeV] is quantified as an unambiguous AGN activity indicator, in \\S7 we show that the mid-infrared diagnostics differentiate AGN 1's from the other populations, and we derive the average spectra for each class of objects that can be used as templates also for predictions and comparisons with high redshift populations. Finally in \\S8 we present the line luminosity functions for our sample, and calculate the total accretion power generated in the local universe. The conclusions are summarized in \\S9. ",
        "conclusions": "Our large sample of Seyfert galaxies, and the accompanying spectropolarimetric classifications have improved and extended the analysis of Paper I. In the IR diagnostic diagrams presented, we find that AGN 1--defined as having broad line emission of some kind--have high values of ionization-sensitive line ratios, relative to the strength of star-formation tracers. In contrast, those galaxies we re-classify as non-Seyfert's have low ionization-sensitive line ratios and high PAH equivalent widths, and lie in the opposite regions of the diagnostic diagrams from AGN. The class of AGN 2's--those without broad lines--spread across both regions of the diagrams.  The simplest and strongest version of the AGN Unified Model requires that any differences between Seyfert 1's and Seyfert 2's should be due to the inclination angle of an obscuring structure (e.g. of a torus) with the line of sight to the observer, the Seyfert 2's being covered by the thick dust in the equator. However, in order to salvage this unification model from the results presented here, one has to explain why AGN 2's show a wider range of mid-IR properties, with similarities both to AGN 1 and the non-Seyfert's. Perhaps this could derive from the particular geometry of the \"torus-BLR-scattering region\": the torus in fact can absorb the radiation not only from the BLR but also from the scattering region where the polarized emission could be produced, depending, for example,  from the precise inclination angle, or the intrinsic clumpiness of the torus. That is the basic suggestion of \\citet{hei97}. However, to suppress high-ionization emission lines we observe in the IR-- which are assumed to be isotropic-- the absorbing structure would also need to block out a significant fraction of the ionizing photons that would otherwise illuminate the NLR, even though this does not seem to happen in Seyferts with any detectable BLR.  But once additional intrinsic differences other than just the viewing angle are admitted, the attractive simplicity of ATM unification is lost.  On the contrary, however, there may be ''genuine'' type 2 Seyfert nuclei.  These would be the half of our AGN2 category which have smaller proportions of AGN emission, and do not have analogs among the Seyfert 1s.  As we discussed, the remaining NLR emission we detect in the optical and infrared spectra of these true Sy2's may be the ''fossil'' remnants of a 10--100 parsec extent of gas that had been ionized by an AGN which was active in previous millenia, but has been effectively ''turned off'' for the last several hundred years.  Given their strong variability over all timescales, AGN which recently ''turned off'' must exist. Depending on the duty cycles of power generation through black hole accretion, they could account for up to half of our AGN2.  This would then imply the existence of a comparable population of ''recently turned on'' AGN.  These could stand out as having unusually strong broad line and near-IR continuum emission, relatively to their NLR strength.  The most likely candidates for these ''young'' or more accurately, ''rejuvenated'' AGN1's are those Seyfert 1's having relatively strong FeII emission and low values of NLR/BLR ratios such as [OIII]5007/H$\\beta$ \\citep{bg92}. Thus it could well turn out that long-term variability is at least as important as viewing angle in unifying the various observed classes of AGN.  \\begin{itemize} \\item[-]The main results of this paper are as follows: We present  the {\\it Spitzer} IRS high-resolution spectra of almost 80\\% of the Seyfert galaxies of the 12$\\mu$m galaxy sample, a total of 91 galaxies; \\item[-] We adopted a spectropolarimetric classification, with the \"AGN 1\" class  broadly defined to include both the Seyfert 1's and the \"hidden broad line Region\" Seyfert 2's, as detected through optical spectropolarimetry. All of our infrared diagnostics are consistent with Sy 1's and HBLR Sy 2's belonging to the same single class. The AGN 2 class contains the remaining Seyfert 2 galaxies without polarized broad lines are likely a mixture of weaker AGN 1, in which the BLR exists but has not yet been detected, and \"true\" AGN 2, galaxies that may not have been producing much hard ionizing radiation for the last several hundred years. Our AGN 1/2 distinction, based solely on reddening-independent IR data, is supported by more of the data than the usual traditional classification scheme that divides the Seyfert 1's and 2's based strictly on the detectability of BLR emission in direct optical spectroscopy. It appears that the mid-IR observed properties characterize the AGN 1 as an homogeneous class of objects. The mid-IR behaviour of AGN 2's, instead, shows objects similar to AGN 1's and others more similar to non-Seyfert galaxies.  \\item[-] Semi-analytic models based on the observed mid-IR spectra are effective in separating the AGN and starburst components in Seyfert galaxies. We find that for 31 AGN 1 the average AGN percentage contribution to the 19$\\mu$m luminosity is 92.2$\\pm$6.6\\%, while for 16 AGN 2 this percentage decreases to 79.3$\\pm$15.7\\% and for 9 non-Seyfert's is 67.6$\\pm$17.2\\%. Although with large scatter, there is a clear trend of decreasing AGN strength from AGN1 to AGN2 to non-Seyfert's. Ionization-sensitive line ratios can discriminate AGN 1's from AGN 2's and non-Seyferts. The diagnostic diagram of [NeV]/[NeII] versus [OIV]/[NeII] provides a measure of the AGN strength in both AGN 1's and AGN 2's. The AGN 2 do not appear to be a homogeneous population with respect to starburst tracers, such as the EW PAH and the EW of [NeII]. Moreover, some AGN 2 nuclei could either contain more ongoing star formation and/or be covered by more dust than any of the AGN 1's.  \\item[-] The 0-0 S(1) H$_2$ rotational transition can be used to estimate the mass of the emitting regions. The average mass value is of the order of 10$^8$ M$_{\\sun}$, which is in agreement with other estimates for Seyfert and starburst galaxies.  \\item[-] The density-sensitive line ratios (the [NeV] and [SIII] doublet ratios), in about 40\\% of the objects imply electron densities below the low density limit. Therefore they can not be used reliably to  estimate the electron density of these regions. The simple interpretation  that the line emission from the NLR and the HII regions can be heavily absorbed by dust even in these mid-IR lines \\citep{dud07} is not confirmed from our data because the galaxies of our sample {\\it do not} show both line ratios affected at the same time in the same objects. In addition, the abnormally low ratios occur with the same frequency in all different types of emission line galaxies.  The other possibility is that some assumed atomic physics parameters need revision for both of these line ratios. \\item[-] We derived for the first time the line luminosity functions for either classical Seyfert 1's and 2's and also for AGN 1 and AGN 2. We do not find significant differences between the various populations, in either the shape or the normalization of their line LFs.  \\item[-] The mid-infrared [NeV] lines are unambiguous tracers of the AGN NLR. We therefore use the [NeV] line luminosity function of all Seyfert galaxies to estimate the accretion power in the local universe within a volume out to z=0.03. We find that the power originating from accretion at 19$\\mu$m is $\\sim$ 2 $\\times$ 10$^{46}$ erg s$^{-1}$ , about 4 times less than the bolometric power. For comparison, the power related to star formation and stellar evolution in the Seyfert galaxies population at 19$\\mu$m is one tenth of that one from accretion.  \\end{itemize}"
    },
    "0911/0911.0382_arXiv.txt": {
        "abstract": "Based on spectroscopic signatures, about one-third of known H$_2$O maser sources in active galactic nuclei (AGN) are believed to arise in highly inclined accretion disks around central engines.  These ``disk maser candidates'' are of interest primarily because angular structure and rotation curves can be resolved with interferometers, enabling dynamical study. We identify five new disk maser candidates in studies with the Green Bank Telescope, bringing the total number published to 30.  We discovered two (NGC\\,1320, NGC\\,17) in a survey of 40 inclined active galaxies ($v_{sys}< 20000$\\,km\\,s$^{-1}$).  The remaining three disk maser candidates were identified in monitoring of known sources: NGC\\,449, NGC\\,2979, NGC\\,3735.  We also confirm a previously marginal case in UGC\\,4203.  For the disk maser candidates reported here, inferred rotation speeds are 130--500\\,\\kms. Monitoring of  three more rapidly rotating candidate disks (CG\\,211, NGC\\,6264, VV\\,340A) has enabled measurement of likely orbital centripetal acceleration, and estimation of central masses (2--7$\\times10^7$\\,M$_\\odot$) and mean disk radii (0.2--0.4\\,pc).  Accelerations may ultimately permit estimation of distances when combined with interferometer data.  This is notable because the three  AGN are relatively distant ($10000<v_{sys}<15000$\\,\\kms), and fractional error in a derived Hubble constant, due to peculiar motion of the galaxies, would be small.  As signposts of highly inclined geometries at galactocentric radii of $\\sim 0.1$--1\\,pc, disk masers also provide robust orientation references that allow analysis of (mis)alignment between AGN and surrounding galactic stellar disks, even without extensive interferometric mapping.  We find no preference among published disk maser candidates to lie in high-inclination galaxies.  This provides independent support for conclusions that in late-type galaxies, central engine accretion disks and galactic plane orientations are not correlated. ",
        "introduction": "Extragalactic H$_2$O maser emission from active galactic nuclei (AGN) in spiral galaxies in some cases traces highly inclined disk structures at radii $\\sim 0.1$--1\\,pc from the central engines.  Cases established by direct interferometric imaging that resolves position-velocity structure include NGC\\,4258 \\citep[e.g.,][]{Miyoshi1995,Argon2007}, NGC\\,1068 \\citep{Greenhill1997}, NGC\\,3079 (Yamauchi et al. 2004; Kondratko, Greenhill, \\& Moran 2005)\\nocite{Yamauchi2004, Kondratko2005},  Circinus \\citep{Greenhill2003}, UGC\\,3789 \\citep{reid09}, and NGC\\,6323 \\citep{Braatz2007}. These ``disk masers''  have common spectroscopic characteristics, highly red- and blueshifted emission lines, called ``high-velocity\" emission lines, that more or less evenly bracket the systemic velocity, and emission close to the systemic velocity, called ``low-velocity\" emission lines. Both are defined in the frame of the rotating disk, projected along the line of sight. We use this terminology of ``high-\" and ``low-\" velocity lines throughout this paper. This template is also observed in other H$_2$O masers  for which imaging data are not available, and it can be a powerful diagnostic.   In total, at least 30 published masers (out of 107  known) share these  characteristics \\citep[e.g.,][and references therein]{Greenhill2008}.  The high-velocity red- and blueshifted emission has been observed to be offset on the order of $10^2$--$10^3$\\,\\kms, among sources. We note that very long baseline interferometric (VLBI) imaging is essential to confirm the identity of disk masers. Hence, we have used the terminology ``disk maser candidates\" for sources with spectroscopic identification only. We also note that VLBI imaging of some disk masers shows signs of additional maser emission arising in outflows (e.g., Circinus) and jets (e.g., NGC\\,1068).  H$_2$O disk masers highlight structures oriented edge-on, or nearly so, and in close proximity to central engines.  This is plausible because large reservoirs of molecular gas along lines of sight are required to generate detectable emission \\citep{eli82}. Among all H$_2$O masers in AGN, $95\\%$ exhibit  X-ray absorption column densities $>10^{23}$\\,cm$^{-2}$ (Greenhill et al.\\ 2008)\\nocite{Greenhill2008}.   Among disk masers, the fraction is nearly as high, and the distribution of obscuring column densities peaks above $10^{24}$\\,cm$^{-2}$.  For thin disks, maser emission is beamed narrowly about the  midplane, into a solid angle $\\ll4\\pi$, as dictated by requisite velocity coherence along lines of sight \\citep[e.g.,][]{Miyoshi1995}.    Geometry limits observable emission to disks that are highly inclined.  Although common association has been well recognized with optical type-2 and partly obscured type-1 objects \\citep[e.g.,][]{Braatz1997b}, disk maser emission is a sharper indicator of (inclination) orientation for nuclear structure than optical classifications are.  In addition, it is a more direct indicator than, e.g., models of Fe K$\\alpha$ emission from hot accretion disks \\citep{nandra97}, the differential brightness of radio jets, and models of kiloparsec-scale ionization structures.  Comparisons of inclination complement analyses of position angles and naturally enable analyses of relative 3-D  orientation \\citep{Clarke1998, Pringle1999, Nagar1999}.  We conducted a survey of AGN to detect new disk masers, selecting galaxies with inclined stellar disks (inclination angle $i>70^\\circ$).  We detected two disk maser candidates and one unclassified maser (Section\\,3.1).  We also conducted deep integrations to identify indicators of disk maser emission among previously  known sources, detecting another three  (Section\\,3.2), and  monitored another three known disk maser candidates to measure centripetal acceleration due to disk rotation (Section\\,3.3).  We discuss estimation of black hole masses and disk radii in Section\\,4 as well as the distribution of inclinations among disk maser candidate hosts and  misalignment between central engines and galactic stellar disks. ",
        "conclusions": "\\label{conclusions}  In a survey with the GBT of 40 inclined AGN, we have detected three new water maser sources (NGC\\,1320, NGC\\,17, and IRAS 16288-3929; Figure~\\ref{new.masers}). We classify two of these (NGC\\,1320 and NGC\\,17) as disk maser candidates based on their spectra. Three more cases of disk maser candidates were discovered through deep integrations toward known masers (NGC\\,449, NGC\\,2979, and NGC\\,3735; Figure~\\ref{other1}).  One previously marginal case was  confirmed (UGC\\,4203; Figure~\\ref{other2}). Inferred rotation speeds are $\\sim 100$--300\\,\\kms. At present, the details of interpretation are sensitive to estimates of systemic velocity at optical wavelengths and  21~cm (the HI line).  These show significant scatter for some galaxies, and in general, comprehensive study is mandatory for each new maser detection to resolve uncertainties in classification. Interferometric follow-up of discoveries reported here will enable confirmation studies of disk geometry (e.g., radii, warping) and estimation of central engine mass, at least  for the sources with large numbers of maser Doppler components.  In general, systematic velocities that result from disk models may provide the best estimates of systemic velocity, and these might usefully be compared to the results of pointed optical and molecular line studies that include dynamical modeling.  For three additional maser sources exhibiting high-velocity maser emission (CG\\,211, NGC\\,6264, VV\\,340A), we report secular drift in multiple low-velocity features---a probable manifestation of centripetal acceleration due to disk rotation. Combined with rotation velocities inferred from spectra, these centripetal accelerations suggest central engine masses of a few times $10^7$\\,M$_\\odot$ and mean disk radii of a few times $10^{-1}$\\,pc. Measurable accelerations of $\\sim$ 1--3\\,\\kms\\ yr$^{-1}$ and large recessional velocities ($>10^4$\\,\\kms) make these three sources particularly attractive for interferometric study in pursuit of robust estimates of geometric distances and H$_0$.  As beacons of highly inclined structures in AGN, at radii of $\\sim 0.1$--1\\,pc,  disk masers also highlight  those galactic nuclei in which we can define a reference plane with high accuracy and study (mis)alignment with surrounding kiloparsec-scale stellar disks. We find no preponderance of edge-on galactic disks among published disk maser candidate hosts. The comparison is in agreement with studies that infer orientations of accretion disks and geometrically thick tori from radio jets and ionization cones, though the maser data refer to a unique range of radii and probe inclination rather than position angle on the sky.  Our studies provide more candidates for the quest to measure H$_0$ with disk masers. The best available estimate of H$_0$ is uncertain by 5\\% and achieved with a composite analysis of Cepheid and supernova data, with a zero point provided by the disk maser distance to NGC\\,4258 \\citep{riess09}. Ultimately, accuracy on the order of  1\\% would match accuracies anticipated for CMB and other data sets in the near future and add a new independent constraint in cosmological parameter estimation. However, achieving this accuracy for H$_0$ will be challenging, probably requiring detailed study of a disk maser sample at least three times larger than what is available today, assuming maser distances with 10\\% accuracy as the norm.  The minimum sample recessional velocity would be $10^4$\\,\\kms, though objects as nearby as 3200\\,\\kms\\ may be useful where peculiar motions are known with uncertainties as small as 160\\,\\kms.  Difficulty in accurately estimating peculiar motions in some volumes may restrict this extended sample, e.g., near the Great Attractor or where uncertainty in Tully-Fisher or Fundamental Plane  distances are large."
    },
    "0911/0911.2664_arXiv.txt": {
        "abstract": "{We investigate the general properties of Unified Dark Matter (UDM)  fluid models where the pressure and the energy density are linked by a barotropic equation of state (EoS)  $p=p(\\rho)$ and the perturbations are adiabatic. The EoS is assumed to admit a future attractor that acts as an effective cosmological constant, while asymptotically in the past the pressure is negligible. UDM models of the dark sector are appealing because they evade the so-called ``coincidence problem'' and ``predict\" what can be interpreted as $w_{\\rm DE} \\approx -1$, but  in general suffer the effects of  a non-negligible Jeans scale that wreak havoc  in the evolution of perturbations, causing a large Integrated Sachs-Wolfe effect and/or changing structure formation at small scales.  Typically,   observational constraints are violated, unless the  parameters of the UDM model are tuned to make it indistinguishable from $\\Lambda$CDM. Here we show how this  problem can be avoided, studying in detail  the functional form  of the Jeans scale in adiabatic UDM perturbations and  introducing a class of models with a fast transition between an early Einstein--de\\! Sitter CDM-like era and a later $\\Lambda$CDM-like phase. If the transition is fast enough, these models  may exhibit  satisfactory  structure formation and CMB fluctuations. To consider a concrete case, we introduce a toy UDM model and show that it can predict CMB and matter power spectra that are in agreement with observations for a wide range of parameter values.} ",
        "introduction": "In the last three decades the flat $\\Lambda$CDM model \\cite{Peebles:1984ge,Efstathiou:1990xe} has emerged as the  standard ``concordance'' \\cite{Spergel:2003cb,Tegmark:2003ud} model of cosmology. It assumes General Relativity (GR) as the correct theory of gravity, with two unknown components  dominating the dynamics of the  late Universe: {\\it i )} a cold collisionless Cold Dark Matter (CDM) describing some weakly interacting particles, responsible for structure formation, {\\it ii )} a cosmological constant $\\Lambda$\\ \\cite{Weinberg:1989, Sahni:1999gb} making up the balance to make the Universe spatially flat and  driving the observed cosmic acceleration \\cite{Riess:1998cb,Perlmutter:1998np,Riess:1998dv,Kowalski:2008ez}. The main alternative to the cosmological constant is a more general dynamic component called Dark Energy (DE) \\cite{Copeland:2006wr,Peebles:2002gy,Padmanabhan:2002ji}. Many independent observations support both the existence of a CDM component and that of a separate DE \\cite{Kowalski:2008ez,Allen:2004cd,Tegmark:2006az,Percival:2006kh,Percival:2007yw,Komatsu:2008hk,Dunkley:2008ie,Hinshaw:2008kr,Percival:2009xn,Reid:2009xm}.  Early proposals \\cite{Peebles:1984ge,Efstathiou:1990xe} of the $\\Lambda$CDM model were adding $\\Lambda$ to CDM in an attempt to conciliate in the simplest possible way the emerging inflationary paradigm, which requires a spatially flat Universe and an almost scale-invariant spectrum of perturbations, with the observed low density of matter. It should however be recognised that, while some form of CDM is independently expected to exist within any modification of the Standard Model of high energy physics, the really compelling reason to postulate the existence of DE has been the  cosmic acceleration measured in the last decade \\cite{Riess:1998cb,Perlmutter:1998np,Riess:1998dv,Kowalski:2008ez,Percival:2007yw,Komatsu:2008hk,Dunkley:2008ie,Hinshaw:2008kr,Percival:2009xn}. It is mainly for this reason that it is worth investigating the hypothesis that CDM and DE are the two faces of a single Unified Dark Matter (UDM) component, thereby also   avoiding the so-called ``coincidence problem'' \\cite{Zlatev:1998tr}.  Other attempts to explain the observed acceleration also exist, most notably by assuming a gravity theory other than GR, or an interaction between DM and DE (see e.g.\\ \\cite{Copeland:2006wr, Durrer:2008, Koyama:2007rx, Capozziello:2007ec, Sotiriou:2008rp} and \\cite{Quercellini:2008vh, Valiviita:2008iv, CalderaCabral:2008bx, Majerotto:2009np, CalderaCabral:2009ja, Valiviita:2009nu, Bassett:2002fe}). In this paper however we focus on UDM models, where this single matter component  provides an  explanation for structure formation and cosmic acceleration.  In general, in the $\\Lambda$CDM model or in most  models with DM and DE, the CDM component is free to form structures at all scales, with  DE  only affecting the general overall expansion \\cite{Copeland:2006wr,Peebles:2002gy,Padmanabhan:2002ji}. Instead, a general feature of UDM models  is the possible appearance of an effective sound speed, which may become  significantly different from zero during the Universe evolution, then corresponding in general to the appearance of a Jeans length (i.e.\\ a sound horizon) below which the dark fluid does not cluster (e.g.\\ see \\cite{Hu:1998kj, Bertacca:2007cv, Pietrobon:2008js}). Moreover, the presence of a non-negligible speed of sound can modify the evolution of the gravitational potential, producing a strong Integrated Sachs Wolfe (ISW) effect \\cite{Bertacca:2007cv}. Therefore, in UDM models it is crucial to study the evolution of the effective speed of sound and that of the Jeans length.  Several adiabatic fluid models and models based on non canonical kinetic Lagrangians have been investigated in the literature. For example, the generalised Chaplygin gas \\cite{Kamenshchik:2001cp, Bilic:2001cg, Bento:2002ps} (see also \\cite{Makler:2002jv, Bento:2002yx, Alcaniz:2002yt, Santos:2006ce, Gorini:2007ta, Sen:2005sk, Makler:2003iw, Carturan:2002si, Amendola:2003bz, Sandvik:2002jz, Giannantonio:2006ij}), the Scherrer \\cite{Scherrer:2004au} and generalised Scherrer solutions \\cite{Bertacca:2007ux}, the single dark perfect fluid with ``affine'' 2-parameter barotropic equation of state (see \\cite{Balbi:2007mz,Pietrobon:2008js} and the corresponding scalar field models \\cite{Quercellini:2007ht}) and the homogeneous scalar field deduced from the galactic halo space-time \\cite{DiezTejedor:2006qh,Bertacca:2007fc}. In general, in order for UDM models to  have a background evolution that fits observations and a very small speed of sound, a severe fine-tuning of their parameters is necessary (see for example \\cite{Pietrobon:2008js, Makler:2003iw, Carturan:2002si, Amendola:2003bz, Sandvik:2002jz, Scherrer:2004au, Giannakis:2005kr, Piattella:2009da}). Finally, one could also easily reinterpret UDM models based on a scalar field Lagrangian in terms of generally non-adiabatic fluids \\cite{DiezTejedor:2005fz, Brown:1992kc} (see also \\cite{Bertacca:2007ux, Bertacca:2008uf}). For these models the effective speed of sound, which remains defined in the context of linear perturbation theory, is not the same as the adiabatic speed of sound (see \\cite{Hu:1998kj}, \\cite{Garriga:1999vw} and \\cite{Mukhanov:2005sc}). In \\cite{Bertacca:2008uf} a reconstruction technique is devised for the Lagrangian, which allows  to find models where the effective speed of sound is small enough, such that the {\\it k}-essence scalar field can cluster (see also \\cite{Camera:2009uz}).  In the present paper we investigate the possibility of constructing adiabatic UDM models where the Jeans length is very small, even when the speed of sound is not negligible. In particular, our study is focused on models that admit an effective cosmological constant and that are characterised by a short period during which the effective speed of sound varies significantly from zero. This  allows a fast transition between an early matter dominated era, which is indistinguishable from an Einstein--de\\! Sitter model, and a more recent epoch whose dynamics, background and perturbative,  are very close to that of a standard $\\Lambda$CDM model.  To consider a concrete example, we introduce a 3-parameter class of  toy UDM adiabatic models with fast transition.  One of the parameters is the effective cosmological constant $\\rho_{\\Lambda}$ or, equivalently, the corresponding density parameter $\\Omega_{\\Lambda}$; the other two are $\\rho_{\\rm s}$ and $\\rho_{\\rm t}$, respectively regulating how fast the transition is and the redshift of the transition. Studying the Jeans scale in these models we find  an approximate analytical relation that sets a constraint on these two  parameters, a sufficient condition that  $\\rho_{\\rm s}$ and $\\rho_{\\rm t}$ have to satisfy in order for the models to be minimally viable.  This relation can be used to fix $\\rho_{\\rm s}$ for any given $\\rho_{\\rm t}$: in this case, with respect to a flat $\\Lambda$CDM, in practice our models have one single extra parameter. With the help of this relation we establish our  main result:  if the fast transition takes place early enough, at a redshift $z\\gtrsim 2$ when the effective cosmological constant is still subdominant, then the predicted background evolution,  Cosmic Microwave Background (CMB) anisotropy and  linear matter power spectrum are in agreement  with observations for a broad range of parameter values.  In practice, in our toy models the predicted CMB  and matter power spectra  do not display significant differences from those computed in the $\\Lambda$CDM model, because the Jeans length remains  small at all times, except for negligibly short periods,  even if during the fast transition the speed of sound can be large. In other words,  this kind of adiabatic UDM models  evade the  ``no-go theorem'' of Sandvik {\\it et al} \\cite{Sandvik:2002jz} who, studying the generalised Chaplygin gas UDM models,  showed that this broad class must have an almost constant negative pressure at all times  in order to satisfy observational constraints,  making these models  in practice indistinguishable from the $\\Lambda$CDM model (see also \\cite{Pietrobon:2008js}).  The paper is organised as follows: in section \\ref{sec:bgperteq} we introduce the basic equations describing the background and the perturbative evolution. In section \\ref{sec:eos} we use the pressure-density plane to analyse  the properties that a general barotropic UDM model has to fulfil in order to be viable. In section \\ref{sec:tghmodel} we introduce our  toy UDM model with fast transition and study its background evolution, comparing it  to a $\\Lambda$CDM. In section \\ref{sec:perts} we analyse the properties of perturbations in this model, focusing on the the evolution of the  effective speed of sound  and that of the Jeans length during the  transition. Then, using the CAMB code \\cite{Lewis:1999bs}, in section \\ref{sec:spectra} we compute the CMB and the matter power spectra. Finally, section \\ref{sec:concl} is devoted to our conclusions. ",
        "conclusions": "\\label{sec:concl}   The last decade of observations of large scale structure \\cite{Allen:2004cd,Tegmark:2006az,Percival:2006kh,Percival:2007yw,Percival:2009xn,Reid:2009xm}, the search for Ia supernovae (SNIa) \\cite{Riess:1998cb,Perlmutter:1998np,Riess:1998dv,Kowalski:2008ez} and the measurements of the CMB anisotropies \\cite{Spergel:2003cb,Komatsu:2008hk,Dunkley:2008ie,Hinshaw:2008kr} are very well explained by assuming  that two dark components govern the dynamics of the Universe. They are DM, thought to be the main responsible for structure formation, and an additional DE component that is supposed to drive the measured cosmic acceleration \\cite{Copeland:2006wr,Peebles:2002gy,Padmanabhan:2002ji}.  At the same time, in the context of General Relativity, it is very interesting to study  possible alternatives. A popular one is that of an interaction between DM and DE,  without violating current observational constraints \\cite{Copeland:2006wr,Quercellini:2008vh,Valiviita:2008iv,CalderaCabral:2008bx,Majerotto:2009np,CalderaCabral:2009ja,Valiviita:2009nu}. This possibility  could alleviate the so called ``coincidence problem\" \\cite{Zlatev:1998tr}, namely, why are the energy densities of the two dark components of the same order of magnitude today. Another attractive, albeit radical, explanation of  the observed cosmic acceleration and structure formation  is to assume the existence of a single dark component: UDM models \\cite{Kamenshchik:2001cp,Bilic:2001cg, Sandvik:2002jz,Carturan:2002si, Amendola:2003bz,Makler:2003iw, Scherrer:2004au, Giannakis:2005kr, Bertacca:2007cv, Bertacca:2007ux, Bertacca:2007fc, Quercellini:2007ht, Bertacca:2008uf, Balbi:2007mz,Pietrobon:2008js, Camera:2009uz} where, by definition, there is no coincidence problem.  In the present paper we have investigated the general properties of UDM fluid models where the pressure and the energy density are linked by a barotropic equation of state (EoS) $p=p(\\rho)$ and the perturbations are adiabatic. Using the pressure-density plane, we have analysed  the properties that a general barotropic UDM model has to fulfil in order to be viable. We have assumed that the EoS of UDM models admits a future attractor which acts as an effective cosmological constant, while asymptotically in the past the pressure is negligible, studying the possibility of constructing adiabatic UDM models where the Jeans length is very small, even when the speed of sound $c_{\\rm s}$ is not negligible. In particular, we have focused on models that admit an effective cosmological constant and that are characterised by a short period during which the effective speed of sound varies significantly from zero. This allows a fast transition between an early epoch that is indistinguishable from a standard matter dominated era, i.e.\\ an Einstein de~Sitter model, and a more recent epoch whose dynamics, background and perturbative, are very close to that of a standard $\\Lambda$CDM model.  In the second part  of the paper, in order to quantitatively investigate observational constraints on UDM models with fast transition, we have introduced and discussed a  toy model based on a hyperbolic tangent EoS [see  Eq.\\ (\\ref{tanh})]. We have shown that if the  transition takes place early enough,  at a redshift $z_{\\rm t}\\ \\gtrsim 2$ when the effective cosmological constant is still subdominant,  being also fast enough, then these models can avoid the oscillating and decaying time evolution of the gravitational potential  that in many UDM models causes  CMB and matter fluctuations incompatible with observations. Consequently, the background evolution, the CMB anisotropy and the linear matter power spectrum predicted by our model do not display significant differences from those computed in a reference   $\\Lambda$CDM  \\cite{Komatsu:2008hk,Hinshaw:2008kr}, because the Jeans length $\\lambda_{\\rm J}=a/k_{\\rm J}$, where $k_{\\rm J}$ is the Jeans wave number [see Eq.\\ (\\ref{kJ2analytic})], remains  small at all times, except for negligibly short periods, even if during the fast transition the speed of sound can be large. In this way, our toy models (and more in general UDM models with a similar fast transition) can evade the ``no-go theorem'' of Sandvik {\\it et al} \\cite{Sandvik:2002jz}, as we discussed in the introduction.  Specifically, we have analysed the properties of perturbations in our toy model, focusing on the the evolution of the effective speed of sound  and that of the Jeans length during the transition. In this way, we have been able to set theoretical constraints on the parameters of the model, predicting sufficient conditions for the model to be viable. Finally, guided by  these predictions  and using  the CAMB code \\cite{Lewis:1999bs}, we have computed the CMB and the matter power spectra  showing that our toy model, for a wide range  of parameters values, fits observation.  The full likelihood analysis for this model and its parameters would be an interesting extension of the study carried out here, which we will address in a future work."
    },
    "0911/0911.2949_arXiv.txt": {
        "abstract": "{ We investigate the spectral evolution of white dwarfs by considering the effects of hydrogen mass in the atmosphere and convective overshooting above the convection zone. Our numerical results show that white dwarfs with $M_{\\rm H}\\sim10^{-16}~M_{\\odot}$ show DA spectral type between $46,000\\lesssim T_{\\rm eff}\\lesssim26,000~{\\rm K}$ and DO or DB spectral type may appears on either side of this temperature range. White dwarfs with $M_{\\rm H}\\sim10^{-15}~M_{\\odot}$ appear as DA stars until they cool to $T_{\\rm eff}\\sim31,000~{\\rm K}$, from then on they will evolve into DB white dwarfs as a result of convective mixing. If $M_{\\rm H}$ in the white dwarfs more than $10^{-14}~M_{\\odot}$, the convective mixing will not occur when $T_{\\rm eff}>20,000~{\\rm K}$, thus these white dwarfs always appear as DA stars. White dwarfs within the temperature range $46,000\\lesssim T_{\\rm eff}\\lesssim31,000~{\\rm K}$ always show DA spectral type, which coincides with the DB gap. We notice the importance of the convective overshooting and suggest that the overshooting length should be proportional to the thickness of the convection zone to better fit the observations. ",
        "introduction": "\\label{sect:intro}  White dwarfs can be classified into several spectral types in terms of the spectral characteristics. The current classification system was introduced by Sion et al. (\\cite{Sc}) and has been modified several times. In this system, the spectral type of a white dwarf is denoted by a letter D plus another letter indicating its spectral characteristics. Sometimes, a suffix is added to indicate some other features (polarization, magnetic field, pulse, etc.). Table \\ref{spec} lists a spectral classification scheme from McCook \\& Sion (\\cite{Ma}).  \\begin{table}[ht] \\begin{center} \\caption{White dwarf spectral types} \\begin{tabular}{p{1pt} l l} \\toprule \\multicolumn{2}{l}{Spectral Type} & \\multicolumn{1}{c}{Characteristics} \\\\ \\midrule DA && Only Balmer lines; no He I or metals present \\\\ DB && He I lines; no H or metals present \\\\ DC && Continuous spectrum, no lines deeper than $5\\%$ in any part of the electromagnetic spectrum \\\\ DO && He II strong; He I or H present \\\\ DZ && Metal lines only; no H or He lines \\\\ DQ && Carbon features, either atomic or molecular in any part of the electromagnetic spectrum \\\\ P & (suffix) & Magnetic white dwarfs with detectable polarization \\\\ H & (suffix) & Magnetic white dwarfs without detectable polarization \\\\ X & (suffix) & Peculiar or unclassifiable spectrum \\\\ E & (suffix) & Emission lines are present \\\\ ? & (suffix) & Uncertain assigned classification; a colon (:) may also be used \\\\ V & (suffix) & Optional symbol to denote variability \\\\ \\bottomrule \\end{tabular} \\label{spec} \\end{center} \\end{table}  Most of white dwarfs are of DA type which have hydrogen-dominated atmospheres. They are found at all effective temperatures from $170,000~{\\rm K}$ down to about $4,500~{\\rm K}$ (Kurtz et al. \\cite{Kb}). DA white dwarfs occupy the vast majority (about $75\\%$) of all known white dwarfs.  About $25\\%$ of the observed white dwarfs have helium-dominated atmospheres which can be divided further into two spectral types. The DO white dwarfs are found between approximately $100,000$ to $45,000~{\\rm K}$, their outer atmospheres being dominated by singly ionized helium (He II). The DB white dwarfs are found between approximately $30,000$ to $12,000~{\\rm K}$, their outer atmospheres being dominated by neutral helium (He I).  There is an interesting fact that few white dwarfs with helium-dominated atmospheres (DO or DB type) are found in the effective temperature range of $45,000\\lesssim T_{\\rm eff} \\lesssim 30,000~{\\rm K}$. This is the so-called DB gap. The reason for this phenomenon is not clear. On one side, there is strong gravitational field in the white dwarf, which may cause a stratified atmosphere. The so-called gravitational settling effect lets the light element (hydrogen) floating to the stellar surface and the heavy element (helium) sinking to the bottom of the stellar envelope. The typical mass fractions of hydrogen and helium in the white dwarfs are $M_{\\rm H}/M_{\\rm tot}\\lesssim 10^{-4}$ and $M_{\\rm He}/M_{\\rm tot}\\lesssim 10^{-2}$, which are the mass threshold for residual nuclear burning (Tremblay \\& Bergeron \\cite{T}). If there is no mixing between hydrogen and helium, we will expect all white dwarfs to be DA stars. On the other side, however, many observational data show that $M_{\\rm H}$ may be significantly lower than the typical value. The existence of a large number of non-DA white dwarfs indicates that some physical mechanisms are competing with the gravitational settling and make the spectral type of some DA white dwarfs changed. The DB gap is suspected to be due to the competition between the gravitational settling and convective mixing.  Fontaine \\& Wesemael (\\cite{Fb}) proposed that when a white dwarf starts cooling from the hot PG $1159$ type star, hydrogen is mixed within the outer helium envelope and the white dwarf shows the DO spectral type. After that, hydrogen gradually floats to the stellar surface in the strong gravitational field. When the DO star cools to $T_{\\rm eff}\\sim45,000~{\\rm K}$, hydrogen has accumulated enough at the surface, and the white dwarf is turned into a DA star. They supposed further that as soon as the DA star cools down to $T_{\\rm eff}\\sim30,000~{\\rm K}$, the He I/II ionization zone in the stellar envelope becomes convective. The convective motion may penetrate into the top hydrogen layer and mix the hydrogen atmosphere into the helium envelope, the less abundant hydrogen being overwhelmed by the more abundant helium and becoming undetectable. So the DA star then appears as a DB white dwarf.  Shibahashi (\\cite{Sa}, \\cite{Sb}) proposed a different scenario for the DB gap. During the early stage of a white dwarf's evolution, the convection in the He II/III ionization zone mixes the hydrogen layer and results in a helium-dominated atmosphere, and the white dwarf appears as a DO star. When the star cools down to around $T_{\\rm eff}\\sim45,000~{\\rm K}$, the He II/III ionization zone becomes deep enough that convection disappears in the DO star's atmosphere, so hydrogen floating to the surface and then the white dwarf being transformed into a DA star. When the white dwarf cools further down to $T_{\\rm eff}\\sim30,000~{\\rm K}$, the He I/II ionization zone generates a convection zone again, as Fontaine \\& Wesemael's proposal, a similar mixing between H and He occurs and the white dwarf appears as a DB star.  The difference between the two scenarios of the DB gap is the time scale of the gravitational settling. In Fontaine \\& Wesemael's assumption, the settling process happens slowly as the cooling of the white dwarf. But in Shibahashi's assumption, it happens quickly as soon as the convection is turned off. However, the recent data of the Sloan Digital Sky Survey (SDSS) suggest that several DB white dwarfs do appear in the DB gap (Eisenstein et al. \\cite{E}). These facts imply that the formation mechanism of the DB gap are not clear and more works should be done.  In the present paper, we calculate a series of white dwarf evolutionary models to investigate the spectral evolution of white dwarfs caused by the convective mixing. The details of model calculations and input physics are presented in Section 2. In Section 3 we discuss the results of our numerical models. Conclusions are summarized in Section 4. ",
        "conclusions": "\\label{sect:conclusion}  From the above investigations we have found that the DB gap could be explained as a consequence of the convective mixing in white dwarfs. DA white dwarfs with $M_{\\rm H}/M_{\\odot}\\sim10^{-16}$ have opportunities to transform into DO white dwarfs at $T_{\\rm eff}\\gtrsim46,000~{\\rm K}$ or DB white dwarfs at $T_{\\rm eff}\\lesssim26,000~{\\rm K}$, respectively. DA white dwarfs with $M_{\\rm H}/M_{\\odot}\\sim10^{-15}$ will can transform into DB stars below $T_{\\rm eff}\\sim31,000~{\\rm K}$. White dwarfs with $M_{\\rm H}$ greater than $10^{-14}~M_{\\odot}$ always appear as DA stars at $T_{\\rm eff}\\gtrsim18,000~{\\rm K}$. It is obvious that in the effective temperature range between $46,000\\lesssim T_{\\rm eff}\\lesssim31,000~{\\rm K}$ almost all of white dwarfs have the DA spectral type, as shown in Figure \\ref{abo}. This scenario substantially coincides with the observation. We can also estimate that the hydrogen mass is $M_{\\rm H}/M_{\\odot}\\sim10^{-16}$ for the DO white dwarfs and is $M_{\\rm H}/M_{\\odot}\\sim10^{-15}$ for the hot DB white dwarfs ($T_{\\rm eff}>20,000~{\\rm K}$).  \\begin{figure}[ht] \\begin{center} \\includegraphics{ms311fig13.eps} \\caption{A schematic representation of the DB gap. $q_{H}=M_{H}/M_{\\odot}$.} \\label{abo} \\end{center} \\end{figure}  Based on our numerical results, the convective overshooting plays a crucial role in the formation of the so-called DB gap, through the convective mixing effect. It allows the convective motion penetrating into the hydrogen layer and makes helium in the deep stellar interior being observable on the stellar photosphere. The overshooting length is an important parameter of the model. According to our results, the overshooting length should be proportional to the thickness of the convection zone, which gives better agreement between the model results and observations.  The hydrogen mass $M_{\\rm H}$ is another important parameter, which is used as a criterion for deciding when helium is dominant in the atmosphere of the white dwarf. It determine decisively the critical effective temperature for the white dwarf changing its spectral type."
    },
    "0911/0911.5044_arXiv.txt": {
        "abstract": "{We present a summary of an asteroseismic signature of helium ionization reported by Houdek \\& Gough (2007, 2008, 2009) for low-degree p modes in solar-type stars, and illustrate its applications for asteroseismic diagnoses.} ",
        "introduction": "Abrupt variation in the stratification of a star (relative to the scale of the inverse radial wavenumber of a seismic mode of oscillation), such as that resulting from the (smooth, albeit acoustically relatively abrupt) depression in the first adiabatic exponent $\\gamma_1=(\\partial {\\ln p}/\\partial{\\ln\\rho})_s$ caused by the ionization of helium, where $p$, $\\rho$ and $s$ are pressure, density and specific entropy, or from the sharp transition from radiative to convective heat transport at the base of the convection zone, induces small-amplitude oscillatory components (with respect to frequency) in the spacing of the cyclic eigenfrequencies $\\nu_i$ of seismic modes, with $i=(n,l)$, where $n$ and $l$ are order and degree respectively. We shall call such abrupt variations in the stellar stratification an acoustic glitch. Our interest is principally in the near-surface glitch caused by helium ionization. ",
        "conclusions": "It was demonstrated by \\citet{Houdek07} that the variation of the sound speed induced by helium ionization enables one to determine from the low-degree eigenfrequencies a measure that is almost proportional to the helium abundance, with little contamination from other properties of the structure of the star. The deviation \\begin{equation} \\delta\\nu_i:=\\nu_i-\\nu_{{\\rm s}i} \\end{equation} of the eigenfrequency from the corresponding frequency $\\nu_{{\\rm s}i}$ of a similar smoothly stratified, i.e. glitch-free, star is the indicator of a measure of the helium abundance $Y$. The individual frequency contributions arising from the acoustic glitches to the total frequency contribution $\\delta\\nu_i$ are: \\begin{equation} \\delta\\nu_i=\\delta_\\gamma\\nu_i+\\delta_{\\rm c}\\nu_i+\\delta_{\\rm u}\\nu_i\\,, \\label{e:all_glitches} \\end{equation} \\begin{figure}[tbp] \\begin{center} \\includegraphics[width=7.8cm]{ghoudek_fig1.eps} \\end{center} \\caption{The symbols denote contributions $\\delta\\nu_i$ to the frequencies $\\nu_i$ produced by the acoustic glitches of the Sun.} \\label{fig:one_column} \\end{figure} where $\\delta_\\gamma\\nu_i$ and $\\delta_{\\rm c}\\nu_i$ arise respectively from the variation in $\\gamma_1$ induced by helium ionization and from the acoustic glitch at the base of the convection zone caused by a near discontinuity in the second derivative of density. Approximate expressions for $\\delta_\\gamma\\nu_i$ and $\\delta_{\\rm c}\\nu_i$ were presented by \\citet{Houdek07, Houdek08}. The additional upper-glitch component $\\delta_{\\rm u}\\nu_i$, which is produced by the ionization of hydrogen and the upper superadiabatic boundary layer of the envelope convection zone, is difficult to model. \\citet{Houdek07, Houdek08} estimated all three glitch components by fitting to second differences $\\Delta_{2i}\\nu:=\\nu_{n-1,l}-2\\nu_{n,l}+\\nu_{n+1,l}$ of the observed frequencies $\\nu_i$ the seismic diagnostic \\begin{equation} \\Delta_{2i}\\nu \\simeq\\Delta_{2i}(\\delta_\\gamma\\nu+\\delta_{\\rm c}\\nu+\\delta_{\\rm u}\\nu)\\,, \\end{equation} and represented the upper-glitch component by a series of inverse powers of $\\nu_i$, i.e. $\\Delta_{2i}\\delta_{\\rm u}\\nu=\\sum_{k=0}^3a_k\\nu_i^{-k}$, from which the four coefficients $a_k$, additional to seven more coefficients in $\\delta_\\gamma\\nu_i$ and $\\delta_{\\rm c}\\nu_i$, were determined. The upper glitch frequency component is represented by \\citep{Houdek09}  \\begin{figure*}[htbp] \\begin{center} \\includegraphics[width=18.0cm]{ghoudek_fig2.eps} \\end{center} \\caption{The symbols in the {\\bf upper panels} denote second differences $\\Delta_{2i}\\nu:=\\nu_{n-1,l}-2\\nu_{n,l}+\\nu_{n+1,l}$ for low-degree modes obtained from the BiSON ({\\bf left panel}; Basu et al. 2007) and from simulated data for a 1\\,M$_\\odot$ model of age 5.54 Gy ({\\bf right panel}); the simulation was based on two periods of 4-months of observation, separated by 1 year, with the SONG network (see Grundahl et al. 2007). The solid curves are fits to $\\Delta_{2i}\\nu$ based on the analysis by Houdek \\& Gough (2007). The dashed curves are the smooth contributions, including a third-order polynomial in $\\nu^{-1}_i$ to represent the upper-glitch contribution from near-surface effects. The {\\bf lower panels} display the remaining individual contributions from the acoustic glitches to $\\Delta_{2i}\\nu$: the dotted and solid curves are the contributions from the first and second stages of helium ionization, and the dot-dashed curve is the contribution from the acoustic glitch at the base of the convective envelope.} \\label{fig:two_columns} \\end{figure*}  \\begin{eqnarray} \\delta_{{\\rm u}}\\nu_i &\\simeq&A+B\\nu_i\\cr &\\!\\!+\\!\\!&{\\left[\\frac{1}{2}a_0\\nu_i^2+a_1\\nu_i(\\ln\\nu_i\\!-\\!1)-a_2\\ln\\nu_i+\\frac{1}{2}a_3\\nu_i\\right]h^{-2}}\\cr &\\equiv&A+B\\nu_i+F_{\\rm u}\\,, \\label{e:delnu_surf} \\end{eqnarray} where $ h=(\\nu_{n+1,l}-\\nu_{n-1,l})/2$. The somewhat arbitrarily chosen power series in $\\nu^{-1}_i$ for the upper-glitch component leads to the two undetermined constants of summation $A$ and $B$, which \\citet{Houdek09} chose by minimizing \\begin{equation} E\\equiv||{\\delta_{\\rm u}\\nu_i}||_2=\\sum_n(A+B\\nu_i+F_{\\rm u})^2\\,. \\end{equation} In Fig.\\,1 is displayed the sum of all acoustic glitches, i.e. Eq.(\\ref{e:all_glitches}), for low-degree solar frequencies observed by BiSON \\citep{Basu07}. Displayed in the top panels of Fig.\\,2 are the second differences of $\\delta\\nu_i$ (solid curve) and of the upper-glitch component $\\delta_{\\rm u}\\nu_i$ (dashed curve) for the Sun and for simulated data of a solar-like star observed with the Danish SONG instrument. The corresponding individual frequency contributions $\\delta_\\gamma\\nu_i$ (dotted and solid curves for the two stages of helium ionization) and $\\delta_{\\rm c}\\nu_i$ (dot-dashed curve) are illustrated in the lower panels of Fig.\\,2.  The amplitudes of the oscillatory contributions from the helium ionization (solid and dotted curves in the lower panels of Fig.\\,2) provide an independent measure of the helium abundance of the star. Moreover, with the help of the smoothed (glitch-free) eigenfrequencies $\\nu_{{\\rm s}i}$ the core-sensitive coefficients of the asymptotic representation of solar-like oscillations \\citep{Tassoul80, Gough86} can be evaluated more accurately. This information stabilizes the age calibration of the Sun and solar-like stars using only low-degree acoustic modes \\citep{Houdek08, Houdek09}."
    },
    "0911/0911.2208_arXiv.txt": {
        "abstract": "We study the dynamics of a non-minimally coupled scalar field cosmology with a potential function. We use the framework of dynamical systems theory to investigate all evolutional paths admissible for all initial conditions. Additionally, we assume the presence of barotropic matter and show that the dynamics can be formulated in terms of an autonomous dynamical system. We have found fixed points corresponding to three main stages of the evolution of the universe, namely, radiation, matter and quintessence domination epochs. Using the linearization of the dynamical systems in the vicinity of the critical points we explicitly obtain formulas determining the effective equation of state parameter for the universe at different epochs. In our approach the form of $w(z)$ parametrisation is derived directly from the dynamical equations rather than postulated {\\it a priori}. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0911/0911.5648_arXiv.txt": {
        "abstract": "{Probability distribution of densities is a fundamental measure of molecular cloud structure, containing information on how the material arranges itself in molecular clouds.} {We derive the probability density functions (PDFs) of column density for a complete sample of prominent molecular cloud complexes closer than $d \\lesssim 200$ pc. For comparison, additional complexes at $d\\approx 250-700$ pc are included in the study.} {We derive near-infrared dust extinction maps for 23 molecular cloud complexes, using the \\textsf{nicest} colour excess mapping technique and data from the 2MASS archive. The extinction maps are then used to examine the column density PDFs in the clouds.} {The column density PDFs of most molecular clouds are well-fitted by log-normal functions at low column densities ($0.5$ mag $< A_V \\lesssim 3-5$ mag, or $-0.5 < \\ln{A_V/\\overline{A}_V} \\lesssim 1$). However, at higher column densities prominent, power-law-like wings are common. In particular, we identify a trend among the PDFs: active star-forming clouds always have prominent non-log-normal wings. In contrast, clouds without active star formation resemble log-normals over the whole observed column density range, or show only low excess of higher column densities. This trend is also reflected in the cumulative forms of the PDFs, showing that the fraction of high column density material is significantly larger in star-forming clouds. These observations are in agreement with an evolutionary trend where turbulent motions are the main cloud-shaping mechanism for quiescent clouds, but the density enhancements induced by them quickly become dominated by gravity (and other mechanisms) which is strongly reflected by the shape of the column density PDFs. The dominant role of the turbulence is restricted to the very early stages of molecular cloud evolution, comparable to the onset of active star formation in the clouds.} {} ",
        "introduction": "Star formation takes place exclusively in molecular clouds, or more precisely, in the most extreme density enhancements of those clouds. In the current view, the structure of molecular clouds, and thereby the occurrence of the density enhancements, is heavily affected by the motions induced by supersonic turbulence \\citep[e.g.][]{sca98}. In parallel, the cloud structure is also crucially affected by the self-gravity of gas and magnetic fields inside the clouds. The relative strengths of these cloud-shaping mechanisms are currently under lively debate and regarded as one of the critical open questions in the physics of the interstellar medium \\citep[for reviews, see][]{mac04, mck07}.   The impact of supersonic turbulence for molecular cloud structure is concretely evidenced by the structural characteristics of molecular clouds that seem to agree with theoretical predictions and numerical simulations of such turbulence \\citep[see e.g. \\S 2.1 in][]{mck07}. % One particularly important statistical property is the probability distribution of densities, which describes the probability of a volume $dV$ to have a density between $[\\rho,\\rho + d\\rho]$. This distribution is expected to take a log-normal shape in isothermal, turbulent media not significantly affected by the self-gravity of gas \\citep[e.g.][]{vaz94, pad97, ost99}. The function plays a fundamental role in current theories of star formation: it is used to explain among others the initial mass function of stars, and the star formation rates and efficiencies of molecular clouds \\citep[e.g.][]{pad02, elm08}.   The log-normality of the probability distributions of density is reflected also in \\emph{column} densities computed from simulations \\citep[e.g.][]{ost01, vaz01, fed09}. Unfortunately, measuring column densities in molecular clouds is a challenge in astrophysics in itself. The commonly used methods for deriving column densities, i.e. measurements of CO line emission, thermal dust emission, and dust extinction, suffer from various model-dependent effects, and often probe only narrow ranges of physical conditions \\citep[e.g.][]{goo09, vas09}.   For the most nearby molecular clouds, dust extinction measurements in the near-infrared provide sensitivity over a relatively wide dynamical range, starting from $N(\\mathrm{H}_2+\\mathrm{H}) \\gtrsim 0.5 \\times 10^{21}$ cm$^{-2}$ \\citep{lom06}. The highest measurable column densities depend on the limiting magnitude of near-infrared data available; using e.g. 2MASS data, column densities of $\\sim 25 \\times 10^{21}$  cm$^{-2}$ are reached \\citep[e.g.][]{kai06, lom06}. This broad range, together with the independency of such data on the dust temperature, makes dust extinction mapping a viable method to study the large-scale, lower-density regions of molecular clouds and thereby to test the predictions from simulations of supersonic turbulence.  In this Letter, we present the first results of a study where we utilise a novel near-infrared dust extinction mapping method to study the structural parameters in a large sample of nearby molecular clouds. In this paper, we focus on the column density PDFs in the clouds, the presentation of the maps and further analysis is left to a forthcoming paper (Kainulainen et al. in prep.). Our cloud sample forms a complete set of prominent cloud complexes within $200$ pc that have an extent of more than $\\sim 4$ pc, or are roughly more massive than $\\sim 10^3$ M$_{\\sun}$. For comparison, the sample includes some additional clouds up to $d\\approx 700$ pc. The method we adopt allows us to determine the column densities over a range that extends to significantly higher column densities than can be probed by CO line emission (due to the freeze-out of molecules), enabling study of the structural parameters in a regime not widely accessed before. ",
        "conclusions": "\\label{sec:discussion}   While supersonic turbulence is expected to develop a density PDF close to a log-normal distribution, prominent deviations from that shape are predicted in strongly self-gravitating systems \\citep[e.g.][]{kle00,fed08}. Recent observational studies have indeed indicated that the column density PDFs of molecular clouds are close to log-normal distributions. For example, \\citet{lom06, lom08} examined the column density PDFs in the Pipe, Rho Oph, and Lupus molecular clouds. They concluded that the PDF of Ophiuchus is satisfactorily fitted by a log-normal function, while the PDF of Lupus is extremely well fitted by it. However, the PDF of Pipe required four log-normal functions, which they suggested originated from physically distinct components along the line of sight.  \\begin{table} \\begin{minipage}[t]{\\columnwidth} \\caption{Molecular cloud properties and the derived parameters} \\centering \\renewcommand{\\footnoterule}{}  % \\begin{tabular}{lcccccc} \\hline \\hline Cloud & $D$\\footnote{References: (1) \\citet{lom08b} (2) \\citet{tor07} (3) \\citet{str96} (4) \\citet{knu98} (5) \\citet{cas98} (6) \\citet{kun98} (7) \\citet{lom06} (8) \\citet{mat79} (9) \\citet{oba98} (10) \\citet{cor97} (11) \\citet{bur93} (12) \\citet{cer93} (13) \\citet{cra70} (14) \\citet{lad09}.} & $M_{H2} \\times 10^4$ & $\\overline{A_V}$ & $\\sigma_{data}$ & $\\sigma_{fit}$ & $N_\\mathrm{YSO}$\\footnote{An order-of-magnitude estimate of the pre-main-sequence star population in the clouds, based on \\citet{rei08}.} \\\\ \\hline \\multicolumn{6}{c}{Star-forming clouds, physical resolution 0.1 pc}\\\\ \\hline \\object{Ophiuchus}        &  119$^1$ & 0.6  & 2   & 4.3 & 0.48\\footnote{No good fit was achieved. A rough estimate is given for reference.} &  300 \\\\ \\object{Taurus}           &  140$^2$ & 0.81 & 1.8 & 1.5 & 0.49 &  300 \\\\ \\object{Serpens}          &  259$^{3}$ & 3.6  & 3 & 2.4 & 0.51 &    300 \\\\ \\object{Cha I}            &  165$^4$ & 0.23 & 1.3 & 1.3 & 0.35 & 200 \\\\ \\object{Cha II}           &  150$^4$ & 0.15 & 1.3 & 1.0 & 0.35 & 50  \\\\ \\object{Lupus III}        &  155$^1$ & 0.11 & 1.3 & 0.9 & 0.35$^c$ & 50 \\\\ \\object{CrA cloud}        &  129$^5$ & 1.2  & 0.8 & 1.0 & 0.44 & tens \\\\ \\object{Lupus I}          &  155$^1$ & 0.29 & 1.0 & 0.7 & 0.43 & tens \\\\ \\object{LDN1228}$^d$      &  200$^{6}$ & 0.23 &  1.1 & 0.7 &  0.32 &  tens  \\\\ \\object{Pipe}             &  130$^7$ & 0.95 & 2.5 & 1.4 & 0.29$^c$ & 15 \\\\ \\object{LDN134}           &  100$^{8}$ & 0.13 & 1.2 & 0.6 & 0.39 &  a few \\\\ \\object{LDN204}\\footnote{The most prominent Lynds dark nebula in the region.} &  119$^1$ & 0.43 & 1.9 & 0.8 & 0.41 &  a few  \\\\ \\object{LDN1333}$^d$      &  180$^9$ & 0.57 & 1.2 & 0.5 & 0.38 &  a few    \\\\ \\hline \\multicolumn{6}{c}{Non-star-forming clouds, physical resolution 0.1 pc}\\\\ \\hline \\object{LDN1719}$^d$      &  120$^4$ & 0.53 & 0.6 & 0.7 & 0.50 &   \\\\ \\object{Musca}            &  150$^4$ & 0.07 & 1.0 & 0.7 & 0.45 &   \\\\ \\object{Cha III}          &  150$^4$ & 0.18 & 1.3 & 0.8 & 0.46 &   \\\\ \\object{Coalsack}         &  150$^{10}$     & 0.5  & 2.7 & 1.4 & 0.28 &   \\\\ \\object{Lupus V}          &  155$^1$ & 0.36 & 1.4 & 0.7 & 0.42 &   \\\\ \\hline \\multicolumn{6}{c}{Star-forming clouds, physical resolution 0.6 pc}\\\\ \\hline \\object{Ori A GMC}        &  450$^{11}$ & 11   & 1.4 & 2.8 & 0.5 &  $>$ 2000 \\\\ \\object{Per cloud}        &  260$^{12}$ & 2.0  & 1.7 & 1.7 & 0.49 &  $>$ 100 \\\\ \\object{Ori B GMC}        &  450$^{11}$ & 9.0  & 1.2 & 1.7 & 0.59 &  $>$ 100 \\\\ \\object{Cepheus A}        &  730$^{13}$ & 3.5  & 1.5 & 1.6 & 0.47 &  $>$ 100 \\\\ \\object{California}       &  450$^{14}$ & 11   & 1.4 & 0.7 & 0.51 &  tens\\\\ \\end{tabular} \\label{table:1} \\end{minipage} \\end{table}   The column density PDFs presented in this Letter show that simple log-normal functions fit the PDFs poorly when considering the whole observed column density range, i.e. $N=0.5- 25 \\times 10^{21}$ cm$^{-2}$. As seen in Figs. \\ref{fig:pdfs} and \\ref{fig:pdfs2}-\\ref{fig:pdfs4}, the PDFs of most clouds deviate from simple log-normality at $s \\gtrsim 0.5-1$, or $A_V \\gtrsim 2-5$ mag. Even though the log-normals fit the peaks of the PDFs well, the excess ``wings'' at higher column densities are persistent features of the molecular cloud structure. Likewise, about half of the clouds show non-log-normal features at low column densities. It is, however, difficult to ascertain whether the low column density features are real. It is quite possible that they are mostly residuals caused by other, physically distinct clouds along the line of sight. Nevertheless, they can also be real signatures of cloud structure; non-log-normal features at low column densities have been predicted, related to intermittent fluctuations in turbulent media \\citep[e.g.][]{fed09}. % We note that the number of pixels in the non-log-normal, higher extinction parts is not small. In fact, the threshold above which the wings become prominent ($A_V \\gtrsim 2-5$ mag) is rather low, suggesting that material related to star formation is dominantly located in the non-log-normal part of the PDF. This is illustrated in the right panel of Fig. \\ref{fig:taurus} which shows the extinction map of Taurus with a contour at $A_V=4$ mag highlighting the regions belonging to the non-log-normal wing of the PDF. The figure also shows the known pre-main-sequence stars that clearly concentrate on the regions of high column density \\citep{reb09}.    Intriguingly, our data appear to show a clear trend. All clouds with active star formation show strong excess of higher column densities (Figs. \\ref{fig:pdfs} and \\ref{fig:pdfs2}). In contrast, almost all quiescent clouds have PDFs that either are well described by log-normal functions over the entire column density range, or show relatively low excess of high column densities (Figs. \\ref{fig:pdfs} and \\ref{fig:pdfs3}). The trend is obvious for the lower-mass clouds in our sample, but an indication of it is seen also among the more massive clouds (Fig. \\ref{fig:pdfs4}): the active star-forming clouds of Orion have prominent wings compared to the California nebula, a massive cloud with significantly lower star-forming activity \\citep{lad09}. In the context of turbulent molecular cloud evolution, these observations are in agreement with a picture in which the structure of a molecular cloud in the early stage of its evolution is decisively shaped by turbulent motions. Hence, its column density PDF is close to log-normal, like it is the case for the non-active clouds in our sample. As the cloud evolves, prominent local density enhancements can become self-gravitating, which also assembles a growing fraction of the gas to higher column density structures. This significantly alters the simple log-normal form of the column density PDF, and is very concretely demonstrated by the cumulative forms of the PDFs (Fig. \\ref{fig:cmf-comb}), showing how dramatically the fraction of material at high column density increases from non-active to active clouds. The cumulative PDFs shown in Fig. \\ref{fig:cmf-comb} generalize this trend, suggested earlier by studies of individual cloud complexes \\citep{cam99, lom06, lom08, lad09}. In the low column density regions of molecular clouds turbulent motions prevail as the dominant structure-shaping mechanism, as indicated by the log-normal-like parts of the PDFs. This appears natural, since those are likely to be the regions where the role of self-gravity remains small.   In this Letter, we have characterised the shape of the column density PDFs in nearby molecular clouds and demonstrated the prevalence of non-log-normalities in them. From the PDFs, we identified a trend that is in agreement with a picture where self-gravity has a significant role in shaping the cloud structure starting from a very early stage, corresponding to the formation of first stars in the cloud. An immediate question following these observations is to what extent similar features are present in the simulations of supersonic turbulence. This can be directly addressed by a comparison of our data to simulations that include self-gravity and follow the evolution of cloud structure as a function of time \\citep[e.g.][]{off08, ban09}. The data presented in this Letter provide a unique set for this purpose, and we are going to address this in a forthcoming paper."
    },
    "0911/0911.3741.txt": {
        "abstract": "The procedure of the construction of a sample of distant ($z>0.3$) radio galaxies using NED, SDSS, and CATS databases for further application in statistical tests is described. The sample is assumed to be cleaned from objects with quasar properties. Primary statistical analysis of the list is performed and the regression dependence of the spectral index on redshift is found. \\mbox{\\hspace{1cm}}\\\\ PACS: 98.52.Eh, 98.54.-h, 98.54.Cr, 98.62.Ve, 98.70.Dk, 98.80.Es %\\pacs{98.52.Eh, 98.54.-h, 98.54.Cr, 98.62.Ve, 98.70.Dk, 98.80.Es} ",
        "introduction": "Radio galaxies, which are among the most powerful cosmic objects, allow us to study the evolution of matter and the dynamics of the expansion of the Universe at different cosmological epochs.  Investigations of these radio sources are linked to several cosmological tests, which permit estimating the parameters and evolutionary characteristics of the Universe. Radio galaxies are identified with giant elliptical galaxies with absolute magnitudes $M\\sim-26$ and have black holes with masses on the of $10^9M_{\\sun}$ as their central engines. These facts allow the radio galaxies to be used as tools for the study of the parameters of the distribution of visible and dark matter, dynamics of the Universe, and the history of the formation of its structure. Let us point out various published methods developed to find cosmological parameters from the observed characteristics of radio galaxies ~\\cite{vo_par1:Verkhodanov2_n,vo_par2:Verkhodanov2_n,vo_par3:Verkhodanov2_n}, such as: \\begin{itemize} \\item ``magnitude--redshift'' Hubble diagrams \\linebreak \\mbox{(``K--$z$'')} ~\\cite{willott_k_z:Verkhodanov2_n}; \\item the ''size--$z$'' diagram ~\\cite{gurvits:Verkhodanov2_n,guerra:Verkhodanov2_n,jack_jan:Verkhodanov2_n}; \\item the ``age--$z$'' diagram ~\\cite{starob:Verkhodanov2_n}; \\item the ``source number--flux' ($\\log N-\\log S$) distribution ~\\cite{condon84:Verkhodanov2_n,condon89:Verkhodanov2_n}; \\item search for clustering and determination of the clustering parameter ~\\cite{blake_wall:Verkhodanov2_n}; \\item search for gravitational lenses ~\\cite{gravlens:Verkhodanov2_n}; \\item investigation of the properties of the radio haloes of clusters of galaxies \\cite{colafrancesko:Verkhodanov2_n}; \\item cosmic background temperature variations caused by the interaction of the large-scale structure (clusters of galaxies) at different redshifts with cosmic microwave background and observed on small angular scales (the Sunyaev--Zel'dovich effect due to inverse \\linebreak Compton scattering) \\cite{sz:Verkhodanov2_n} and on large scales in variable gravitational potential---the Sachs--Wolfe effect \\cite{swe:Verkhodanov2_n}. Radio galaxies can be used as a tool for searching for clusters of galaxies~\\cite{protoclust:Verkhodanov2_n}; \\item and a number of other methods (see, e.g., reviews \\mbox{\\cite{vo_par1:Verkhodanov2_n,vo_par2:Verkhodanov2_n,miley_debreuck:Verkhodanov2_n}}). \\end{itemize}  A striking feature of radio galaxies is that we observe them virtually from their formation epoch, i.e., their radio sources are so powerful that modern radio surveys have catalogued almost all objects of this type~\\cite{vo_par3:Verkhodanov2_n}. Such objects can serve as good probes for the study of the formation of clusters of galaxies. Thus \\mbox{Venemans et al.~\\cite{protoclust:Verkhodanov2_n}} show that 75\\% of the radio galaxies with $z>2$ are associated with protoclusters. Based on this fact, the above authors estimate that there are roughly about 3$\\times$10$^{-8}$ of forming clusters per comoving Mpc$^3$ in the redshift interval \\mbox{$2<z<5.2$}  with an active radio source. They also estimate the density of protoclusters within the same redshift interval to be of the order of 6$\\times$10$^{-6}$ Mpc$^{-3}$.  Identification of radio galaxies with giant elliptical galaxies (gE) born as a result of mergers in the early epoch allows them to be used not only to probe the formation of the large-scale structure and test models of star formation. The (``K--$z$'') diagram confirms the age of the old stellar population in gE-type galaxies~\\cite{willott_k_z:Verkhodanov2_n} and this fact can also be used to compute the ages of these galaxies \\cite{rg_age:Verkhodanov2_n,rg_rc_age:Verkhodanov2_n,age_sys:Verkhodanov2_n} and perform quick photometric redshift estimates for these objects out to  $z\\sim4$ \\cite{verkh_phot_z:Verkhodanov2_n}.  Another important problem associated with the study of radio galaxies is the origin of supermassive black holes. The estimates of energy release in the central engines of radio galaxies at redshifts \\mbox{$4<z<5$}, when the Universe was slightly older than one billion years (in terms of the $\\Lambda$CDM model), yield black-hole masses on the order of  $10^9 M_{\\sun}$, whereas it takes much more than one billion years for such black holes to form \\cite{debreuk_z:Verkhodanov2_n}. This paradox can be resolved in terms of the hierarchical formation model, but the actual number of such objects (black holes with masses on the order of $10^9 M_{\\sun}$ in the redshift interval \\mbox{$4<z<5$)} in the Universe is not known with certainty.  Distant radio galaxies are selected by a criterion based on the spectral index of their continuum. To this end, lists of sources with ultrasteep spectra are compiled, where the parameters needed are the low--frequency data points on the continuum radio spectra. Among such objects a high percentage of distant radio galaxies is found~\\mbox{\\cite{dagkes:Verkhodanov2_n,blum_miley:Verkhodanov2_n,par_big3a:Verkhodanov2_n,par_big3b:Verkhodanov2_n,uss_list:Verkhodanov2_n}}, including the most distant objects with redshifts $z>4.5$: $z=5.199$ \\cite{debreuk_z:Verkhodanov2_n} and with $z=4.514$ \\cite{kopylov_z:Verkhodanov2_n}.  Nowadays databases and identification techniques allow automating search and identification of distant radio galaxies. A number of authors~\\mbox{\\cite{deca_bul:Verkhodanov2_n, deca_aat1:Verkhodanov2_n, deca_aat2:Verkhodanov2_n, deca_azh:Verkhodanov2_n, deca_piter08:Verkhodanov2_n, deca_bul2:Verkhodanov2_n, grg_piter08:Verkhodanov2_n}} applied such approaches to  study radio sources using CATS~\\footnote{\\tt http://cats.sao.ru} \\cite{cats:Verkhodanov2_n,cats2:Verkhodanov2_n} and NED \\footnote{\\tt http://nedwww.ipac.caltech.edu} (NASA/IPAC Extragalactic Database) databases of astronomical catalogs and SkyView virtual telescope~\\footnote{\\tt http://skyview.gsfc.nasa.gov}, which allow performing morphological analysis of objects. \\begin{figure*}[tbp] \\centerline{\\hbox{ \\psfig{figure=fig1_map_sources.ps,width=14cm} %\\includegraphics[width=13.1cm]{Verkhodanov2_fig1.eps} }} \\caption{Sky positions of the selected radio sources in Galactic coordinates. The white circles and gray crosses indicate the SDSS objects and all other sources, respectively.} \\label{f1:Verkhodanov2_n} \\end{figure*}   The catalog of distant radio galaxies can be used  not only to perform just purely cosmological studies, but also carry out statistical analyses of identification lists and of the corresponding populations at different \\mbox{wavelengths \\cite{irtex3:Verkhodanov2_n,irtex4:Verkhodanov2_n,balayan:Verkhodanov2_n,trushkin:Verkhodanov2_n};} to search for and analyze the properties of subsamples of radio galaxies~\\cite{rg_gen:Verkhodanov2_n,rg_zen0:Verkhodanov2_n,rg_zen1:Verkhodanov2_n,rg_zen2:Verkhodanov2_n,grg_piter08:Verkhodanov2_n}, and simulate radio-astronomical surveys at \\mbox{RATAN--600 \\cite{compl:Verkhodanov2_n,rzf_cmb:Verkhodanov2_n,rzf_majorova:Verkhodanov2_n}.} Organizing the databases and identifying the available heterogeneous measurements with particular objects are also tasks of great interest, and, as practical experience shows, such compiled catalogs are popular among the astronomical community \\cite{cats2:Verkhodanov2_n,andernach:Verkhodanov2_n,vollmer:Verkhodanov2_n}.  Moreover, radio galaxies, as extended objects, may distort microwave background and bias the estimates of the angular power spectrum  \\mbox{of CMB~\\cite{huffenberger:Verkhodanov2_n,grg_piter08:Verkhodanov2_n}.}  In this paper we report a catalog of radio galaxies with spectroscopic  \\mbox{redshifts $z\\ge0.3$} based on the data from  the NED, SDSS\\footnote{\\tt http://www.sdss.org}\\cite{sdss:Verkhodanov2_n}, CATS, and other databases and on an analysis of the corresponding bibliography, which allowed us to discard  objects exhibiting quasar properties. There is no yet a consensus  among the research community as to which radio galaxies should be considered  ``distant'', with different authors using this term to denote objects with redshifts above  $0.3$, $0.5$, or $1.0$. Based on observational data, one can set the boundary for distant radio galaxies at  $z\\sim0.6-0.7$ (corresponding the age of the Universe of about 6.5~Gyr), when, according to $\\Lambda$CDM cosmology, the domination of dark energy (DE) began. The lower redshift boundary can then be set at half this $z$ value. Such a choice permits the test to include the population of galaxies after the beginning of the DE domination.  This is the first publication of a series of three papers dedicated to the construction of the sample and statistical analysis of the catalog of radio galaxies, which contains objects with known spectroscopic~$z$ and available radio data and optical photometry. In the future we plan to perform cosmological tests based on the physical parameters of the objects of our catalog. ",
        "conclusions": "We report the first part of the catalog of radio galaxies with $z>0.3$ with complete information on the coordinates of the objects, their spectral indices at radio frequencies, and redshifts. We constructed our catalog based on the lists of objects from the  NED and CATS databases. Our sample includes objects of the SDSS survey and of the ``Big Trio'' program. We perform preliminary statistical analysis and compute the distributions of spectral indices and fluxes and plot diagrams of parameters as functions of  $z$. The catalog forms a basis for further study of the objects and of the properties of distant radio galaxies population. For our objects we construct the spectra and the ``flux--redshift'' and ``spectral index--redshift'' diagrams. On the former diagram we determine the upper boundary of the dependence of flux on  $z$, which can be described by the regression relation \\mbox{$\\log S = 4.16 - 0.45z$.} The latter diagram yields the regression relation \\mbox{$\\alpha(z) = -0.73-0.15z$.}    %\\input{Verkhodan2_tab_copy3.tex} \\noindent {\\small {\\bf Acknowledgments}. This research has made use of the NASA/IPAC Extragalactic Database (NED) which is operated by the Jet Propulsion Laboratory, California Institute of Technology, under contract with the National Aeronautics and Space Administration. We also made use of the CATS database of radio astronomical catalogs \\cite{cats:Verkhodanov2_n,cats2:Verkhodanov2_n} and FADPS\\footnote{\\tt http://sed.sao.ru/$\\sim$vo/fadps\\_e.html} radio-astronomical data reduction system~\\cite{fadps:Verkhodanov2_n,fadps2:Verkhodanov2_n}. This work was supported by the Program of the Support of Leading Scientific Schools in Russia (the School of S.\\,M.\\,Khaikin) and the RFBR grant nos.\\,09-02-00298 and 09-02-92659-Ind. O.V.V. acknowledges partial support from the Russian Foundation for Basic Research (project No\\,07-02-01417-a) and the Foundation for the Support for Russian Science (``Young Doctors of Science of the Russian Academy of Science''). }"
    },
    "0911/0911.1893_arXiv.txt": {
        "abstract": "{ We derive the abundance of 19 elements in a sample of 64 stars with fundamental parameters very similar to solar, which minimizes the impact of systematic errors in our spectroscopic 1D-LTE differential analysis, using high-resolution ($R\\simeq60,000$), high signal-to-noise ratio ($S/N\\simeq200$) spectra. The estimated errors in the elemental abundances relative to solar are as small as $\\simeq0.025$\\,dex. The abundance ratios [X/Fe] as a function of $\\feh$ agree closely with previously established patterns of Galactic thin-disk chemical evolution. Interestingly, the majority of our stars show a significant correlation between [X/Fe] and condensation temperature ($\\tc$). In the sample of 22 stars with parameters closest to solar, we find that, on average, low $\\tc$ elements are depleted with respect to high $\\tc$ elements in the solar twins relative to the Sun by about 0.08\\,dex ($\\simeq20\\,\\%$). An increasing trend is observed for the abundances as a function of $\\tc$ for $900<\\tc<1800$\\,K, while abundances of lower $\\tc$ elements appear to be roughly constant. We speculate that this is a signature of the planet formation that occurred around the Sun but not in the majority of solar twins. If this hypothesis is correct, stars with planetary systems like ours, although rare (frequency of $\\simeq15$\\,\\%), may be identified through a very detailed inspection of the chemical compositions of their host stars. } ",
        "introduction": "Standard spectroscopic abundance analyses suffer from a variety of systematic errors that are difficult to remove. Using the highest quality data, errors in stellar parameters, atomic/molecular data, the use of static/homogeneous model atmospheres, and the assumption of local thermodynamic equilibrium set an optimistic lower limit of about 0.05~dex ($\\sim10$\\,\\%) to the accuracy of abundance determinations \\citep[e.g.,][]{asplund05:review}. A simple way to minimize the impact of these uncertainties is to perform differential analyses of stars that are very similar to each other so that systematic errors are largely cancelled out. Naturally, attempts have been made using stars with parameters that are very similar to solar, the so-called solar twins \\citep[e.g.,][]{melendez06:twins,melendez07:twins}. Recent work on very high-quality spectra of a few solar twins suggests that it is even possible to reach 0.01\\,dex accuracy ($\\sim2$\\,\\%, \\citealt{melendez09:twins}). This level of precision can be useful for revealing the fine details of abundance trends and, perhaps more importantly, to determine whether the solar chemical composition is anomalous.  The use of the Sun as a reference star is understandable. Its basic properties (effective temperature, luminosity, mass, radius, and age) are very well known, and spectra of high quality are available or can be easily acquired. Defining a good sample of solar twins is a more difficult task \\citep{cayrel96}. Stars with fundamental parameters very similar to solar exist \\citep[e.g.,][]{porto97,melendez06:twins,takeda07,melendez07:twins}, yet when more detailed analyses of their chemical compositions or evolutionary states are made, some differences arise, such as the apparently low Li abundance of the Sun and the older age of HIP\\,56948, the best solar twin known to date \\citep{melendez07:twins,takeda09}. However, solar twins with Li abundances and ages similar to those of the Sun most likely exist \\citep{pasquini08}, while the age of HIP\\,56948 may have been overestimated \\citep{donascimento09}. In any case, although the Sun is probably not unique, it does not seem to be a very common object in the solar neighborhood \\citep[e.g.,][]{gustafsson08:sun_unique}.  Determining the detailed chemical composition of the Sun compared to other stars in the solar vicinity will help us determine how unusual our star is, and perhaps even why. Previous studies have been inconclusive because of the systematic errors described above, and because the relevant differences may be small. Here we derive precise abundances for a carefully selected sample of solar twins and analogs, for which systematic errors are minimized using differential analysis, and speculate about the nature of the abundance trends found. ",
        "conclusions": ""
    },
    "0911/0911.3896_arXiv.txt": {
        "abstract": "We present high spectral resolution VLT observations of the BAL quasar SDSS~J0318-0600. This high quality data set allows us to extract accurate ionic column densities and determine an electron number density of $n_e$=10$^{3.3\\pm0.2}$ cm$^{-3}$ for the main outflow absorption component. The heavily reddened spectrum of SDSS~J0318-0600 requires purely silicate dust with a reddening curve characteristic of predominately large grains, from which we estimate the bolometric luminosity. We carry out photoionization modeling to determine the total column density, ionization parameter and distance of the gas and find that the photionization models suggest abundances greater than solar. Due to the uncertainty in the location of the dust extinction, we arrive at two viable distances for the main ouflow component from the central source, 6 and 18 kpc, where we consider the 6 kpc location as somewhat more physically plausable. Assuming the canonical global covering of 20\\% for the outflow and a distance of 6 kpc, our analysis yields a mass flux of 120 M$_{\\sun}$ yr$^{-1}$ and a kinetic luminosity that is $\\sim$0.1\\% of the bolometric luminosity of the object. Should the dust be part of the outflow, then these values are $\\sim$4$\\times$ larger. The large mass flux and kinetic luminosity make this outflow a significant contributor to AGN feedback processes. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0911/0911.4364_arXiv.txt": {
        "abstract": "We present high-resolution spectroscopic measurements of the abundances of the $\\alpha$ element titanium (Ti) and s-process elements yttrium (Y) and lanthanum (La) for 59 candidate M giant members of the Sagittarius (Sgr) dwarf spheroidal (dSph) + tidal tail system pre-selected on the basis of position and radial velocity.  As expected, the majority of these stars show peculiar abundance patterns compared to those of nominal Milky Way stars, but as a group the stars form a coherent picture of chemical enrichment of the Sgr dSph from [Fe/H] = -1.4 to solar abundance. This sample of spectra provides the largest number of Ti, La and Y abundances yet measured for a dSph, and spans metallicities not typically probed by studies of the other, generally more metal-poor Milky Way (MW) satellites. On the other hand, the overall [Ti/Fe], [Y/Fe], [La/Fe] and [La/Y] patterns with [Fe/H] of the Sgr stream plus Sgr core do, for the most part, resemble those seen in the Large Magellanic Cloud (LMC) and other dSphs, only shifted by $\\Delta$[Fe/H]$\\sim$+0.4 from the LMC and by $\\sim$+1 dex from the other dSphs; these relative shifts reflect the faster and/or more efficient chemical evolution of Sgr compared to the other satellites, and show that Sgr has had an enrichment history more like the LMC than the other dSphs. By tracking the evolution of the abundance patterns along the Sgr stream we can follow the time variation of the chemical make-up of dSph stars donated to the Galactic halo by Sgr.  This evolution demonstrates that while the bulk of the stars currently in the Sgr dSph are quite unlike those of the Galactic halo, an increasing number of stars farther along the Sgr stream have abundances like Milky Way halo stars, a trend that shows clearly how the Galactic halo could have been contributed by present day satellite galaxies even if the {\\it present} chemistry of those satellites is now different from typical halo field stars. Finally, we analyze the chemical abundances of a moving group of M giants among the Sgr leading arm stars at the North Galactic Cap, but having radial velocities unlike the infalling Sgr leading arm debris there. Through use of ``chemical fingerprinting'', we conclude that these mostly receding northern hemisphere M giants also are Sgr stars, likely {\\it trailing arm} debris overlapping the Sgr leading arm in the north. ",
        "introduction": "Though it is now clear that accretion of dwarf galaxies likely played a prominent role in creating the Milky Way's (MW) stellar halo (Searle \\& Zinn 1978), with strong observational evidence (e.g., Majewski 1993; Majewski, Munn \\& Hawley 1996; Bell et al. 2008) and a theoretical backing by $\\Lambda$CDM models (e.g., Bullock \\& Johnston 2005; Robertson et al. 2005; Abadi et al.\\ 2006; Font et al.\\ 2006a), it is often highlighted that the chemical abundance patterns of current MW satellites are rather different than those of halo field stars, which typically show significantly higher [$\\alpha$/Fe] than MW dSph stars at the same [Fe/H] (e.g., Fulbright 2002; Shetrone et al.\\ 2003; Tolstoy et al.\\ 2003; Venn et al.\\ 2004; Geisler et al.\\ 2005). The reason for these differences remains a matter of speculation. One interpretation is that the dwarf systems that contributed the bulk of the halo were more massive, Magellanic Cloud-sized systems that were accreted and destroyed very early on, so that the chemistry of the accreted stars was necessarily dominated only by enrichment from Type II supernovae (SN II)(Robertson et al.\\ 2005; Font et al.\\ 2006a). Alternatively (or in addition), if current dwarf satellites have experienced both prolonged chemical evolution and tidal disruption, this will naturally lead to evolution in the types of stars that satellites contribute to the Galactic halo (Majewski et al.\\ 2000, 2002).  In principle, one can test the bridge from dwarf galaxy chemistry to halo star chemistry {\\it directly} if one can identify the stars already contributed by the dwarf galaxy to the halo. Perhaps the easiest way to do this is by exploring elemental abundances along the tidal tails of disrupting dwarf galaxies. Presently, the best known example of a tidally disrupting Milky Way satellite is the Sagittarius (Sgr) dwarf spheroidal (dSph) galaxy, a system that is especially interesting for understanding the issues raised above because the proximity of its core and tails make it particularly accessible to high resolution spectroscopic study, and because the tails produced by its steady assimilation into the MW have been tracked over a substantial length (e.g., Ibata et al.\\ 2001; Majewski et al.\\ 2003, hereafter Paper I; Belokurov et al.\\ 2007) corresponding to several gigayears of tidal disruption (Law et al.\\ 2005, ``Paper IV'' hereafter). Thus the extensive Sgr tidal stream can provide a key link between dwarf galaxies, tidal disruption and capture, and the chemical evolution and origin of the MW halo.  Exploration of the chemistry of any dwarf galaxy can be used to study star formation environments that are quite distinct from those in the MW, but Sgr provides an interesting case study for chemical evolution in its own right. It is the most luminous and massive of the current MW dSph systems, and has a history punctuated by a number of star formation episodes (Sarajedini \\& Layden 1995; Layden \\& Sarajedini 2000; Siegel et al.\\ 2007). This active star formation history quickly elevated Sgr's mean metallicity to almost solar by about 5 Gyr ago (Bellazzini et al.\\ 2006; Siegel et al.\\ 2007), making it the dSph with the most metal-rich stellar populations known associated with the Milky Way. Because Sgr has been tidally disrupting for at least 2.5-3.0 Gyr (Paper IV), in principle this means that Sgr could have contributed stars with a wide metallicity range to the Galactic halo.  This has been confirmed by the recent analysis of Chou et al. (2007, hereafter Paper V), who measured [Fe/H] ranging from -1.4 dex to 0.0 dex for stars identified spatially, dynamically and by spectral type (i.e. the M giants characteristic of the Sgr dSph) with the Sgr stream. Interestingly, the Paper V analysis revealed a strong metallicity gradient along the Sgr stream --- proof of a time dependence in the enrichment level of the stars donated to the halo by Sgr, and evidence for preferential tidal stripping of metal poor stars, which has led to divergent metallicity distribution functions (MDFs) between lost and retained Sgr stars. A similar metallicity gradient in the Sgr leading and trailing streams has also been seen by the high resolution analysis of M giants by Monaco et al. (2007). Because stars are typically stripped from the outer parts of dSphs, the changing MDF along the Sgr stream suggests a strong metallicity gradient within the original Sgr system.   Sgr is also chemically interesting because its abundance patterns reveal that while it is undersolar and depleted with respect to the MW pattern for its $\\alpha$-, odd-$Z$ and iron peak-elements among its [Fe/H] $\\gtrsim$ -1 populations, Sgr's more metal-poor populations seem to have [$\\alpha$/Fe] that {\\it do} resemble those of the Galactic halo (Smecker-Hane \\& McWilliam 2002, Bonifacio et al.\\ 2004, Monaco et al.\\ 2005b, Sbordone et al.\\ 2007). Presently these findings are based on only about 6 Sgr stars having both [Fe/H] $< -1$ and measured [$\\alpha$/Fe] from all of the above studies, as well as a handful of stars from M54 (possibly the nucleus of Sgr --- see discussions by Ibata et al.\\ 1994; Sarajedini et al.\\ 1995; Da Costa \\& Armandroff 1995; Bassino \\& Muzzio 1995; Layden \\& Sarajedini 2000; Paper I, but also cf. Monaco et al.\\ 2005a) in the study by Brown et al.\\ (1999). However, if true, the trend may not be unique to Sgr: Abundance patterns (e.g., [$\\alpha$/Fe]) for some of the very most metal-poor stars in other dSphs also seem to overlap those of halo stars of the same metallicity (Shetrone et al.\\ 2003; Geisler et al.\\ 2005; Tolstoy 2005). Thus, if {\\it these} dSph systems experienced tidal disruption, the few very metal poor stars they now hold with MW-like abundance patterns may only represent the residue of a formerly much larger metal-poor population that may have been predominantly stripped from the satellites over their lifetime (see discussion of this in the case of the Carina dSph in Majewski et al. 2002 and Mu\\~noz et al. 2006). Such a scenario could explain the possible origins of the abundance dichotomy between the present dSphs and the Galactic halo.  In this paper we continue our detailed exploration of Sgr stream stars from Paper V with a focus on the chemical trends of the $\\alpha$-element titanium (Ti) and the s-process elements yttrium (Y) and lanthanum (La) along the stream. Our study is based on the largest sample of high resolution spectra yet obtained of Sgr stars, and includes a significant number for stars with [Fe/H] $< -0.9$ (\\S2-3). Collectively, and when combined with additional data from the literature, these stars present a much more complete picture of the chemical abundance patterns for Sgr stars, including the first view of Sgr s-process abundances for [Fe/H] $<-0.9$. In \\S4 we show that, as seen before with smaller samples, the Sgr [$\\alpha$/Fe] abundances are enhanced and MW-like among the more metal-poor stars, but we also show, for the first time that this trend also arises among the s-process elements we explore. Because of the overall variation of the MDF reported in Paper V, there are increasing numbers of stars with ``MW-like'' abundances with distance along the stream away from the Sgr core. Thus, the Sgr stream provides a direct connection between the unusual, non-MW-like abundance patterns in the Sgr core and stars with halo-like abundance patterns contributed by Sgr several Gyr ago. In this way, Sgr provides a counter-example of the Robertson et al. (2005) and Font et al. (2006a) hypothesis that the $\\alpha$-enhanced stars of the halo were deposited there long ago.  In \\S 5 we also discuss the implications of the Sgr abundance patterns for the chemical evolution of this dSph. By comparison of the Sgr abundance patterns to those of other MW satellites, we show that Sgr more closely resembles the Large Magellanic Cloud in its chemical evolution than the other dSph-type systems.  Finally, throughout this paper we simultaneously analyze the chemical abundances of a moving group of M giants among the Sgr leading arm stars at the North Galactic Cap but having different radial velocities than the infalling Sgr leading arm debris there. In Paper V we showed these stars to have a very similar MDF to leading Sgr arm stars stripped several Gyr ago, a result expected if this moving group were constituted by Sgr stars stripped at about the same time. In a demonstration of the technique of ``chemical fingerprinting'', we conclude from their peculiar Ti, Y and La patterns, which also match those of the extreme Sgr leading arm, that most of these receding northern hemisphere M giants could represent Sgr {\\it trailing arm} stars in the northern hemisphere. ",
        "conclusions": "Chemical abundances in stars are fossil records of the enrichment history of a galaxy and their distinctive patterns provide signatures that, if not uniquely branding its stars, at least allow us a means to test whether particular stars are likely to be associated with that system based on whether their chemistry fits into its overall abundance patterns. We have applied this test here not only to demonstrate the likely high purity of the spatially- and kinematically-selected sample of M giant Sgr tidal tail star candidates used in Paper V to show the existence of a strong metallicity gradient along the Sgr tails, but also to prove the likely Sgr-origin of the somewhat mysterious ``North Galactic Cap moving group'' also discussed in that previous contribution. We conclude that most of the NGC moving group stars are probably from the old trailing debris arm of Sgr (their MDF also supports this conclusion, see Paper V). Thus the present paper demonstrates the applicability of ``chemical fingerprinting'', a technique long discussed as one of the potentially valuable future tools of stellar populations research.  The evolving chemical content of a galaxy depends on many variables --- such as the stellar initial mass function and the SFR --- that lead to the specific distribution of chemical patterns among the stars in the system. Tidal stripping can also shape the {\\it observed} abundance patterns of the stars in a galaxy by preferentially removing certain populations, particularly if the system has spatial variations in metallicity and/or an age-metallicity relation over timescales overlapping the period during which the tidal loss of stars occurs. Both of these situations are at play in the Sgr system. Only by surveying both the stars lost from as well as remaining in a galaxy like Sgr do we have hope of recovering an unbiased view of its chemical history. We have explored the chemical abundance patterns in stars along the Sgr tidal tails, which, when combined with data on stars in the Sgr core, afford us the most complete view of the original chemical abundance distributions of this system to date; but even then, our view here is likely to be rather incomplete and still tells us only part of the story. For example, we still do not have good information on the chemical properties of the oldest, most metal-poor Sgr populations.  In stars associated with Sgr we have found lower values of [Ti/Fe] at [Fe/H] $> -1$ than for MW stars at the same [Fe/H] (Fig.\\ 3). And while the MW exhibits an apparent transition of [Ti/Fe] to solar-like levels at [Fe/H] $\\sim -0.7$, such a transition (and to even lower [Ti/Fe] levels) happens for Sgr stars at [Fe/H] $\\sim -1$. This shows that Sgr had a slower SFR than the MW, like other Galactic satellites. We also find the Sgr stars to have subsolar abundances of the light s-process element yttrium at all [Fe/H] probed, while the heavy s-process element lanthanum is enhanced relative to the MW for [Fe/H] $\\gtrsim -0.5$ (Figs.\\ 4-6). This indicates the importance of low-metallicity AGB nucleosynthesis in the metal-rich Sgr stars (Smecker-Hane \\& McWilliam 2002) and is another signature of a slower SFR than occurred in the MW.  We also find that although the LMC and other dSphs exhibit similar chemical pattern trends as Sgr, these patterns exhibit their significant abundance transitions at lower [Fe/H] than for Sgr. Such differences suggest a faster enrichment and more rapid star formation history in Sgr relative to those in the LMC and other dSphs. After applying a hypothetical shift in [Fe/H] of +1 dex in the abundance patterns of other dSphs and +0.4 dex for the LMC we find that the chemical patterns of Sgr, LMC and other dSphs strongly resemble each other (Fig.\\ 10), which suggests the possibility of a universal chemical enrichment progression among MW satellites that differs only by the SFRs; these SFRs, in turn, are likely correlated to the original masses of these systems. The relative shifts suggest that the relative SFR of these galaxies are, in order from slowest to fastest: the other MW dSphs, the LMC, Sgr and the MW respectively.  Based on their close chemical similarities we suggest that Sgr was probably formerly more similar in nature to the present morphology and structure of the LMC than to those of other MW dSphs. The wide ranges of metallicities in Sgr suggests a large age spread, and implies a long duration of star formation in its central regions (Smecker-Hane \\& McWilliam 2002).  This agrees with the discovery of quite young populations in the Sgr core (Siegel et al.\\ 2007), which shows that Sgr has had even relatively recent star formation activity, like the LMC.  The color-magnitude distribution of Sgr stars demonstrates that its star formation history was highly variable, including some fairly well-defined ``bursts'' (Layden \\& Sarajedini 2000; Bellazzini et al.\\ 1999; Monaco et al.\\ 2002; Siegel et al.\\ 2007). These produced populations with different, but overlapping radial density profiles in the progenitor satellite, and likely a strong internal metallicity gradient. In Paper V we found a strong metallicity gradient along the Sgr tidal arms, with the stars in these arms increasingly more metal-poor with angular separation (i.e. orbital phase difference) from the core. By comparing this gradient to models of the timescale for the disruption that produced these tails (Paper IV) we argued in Paper V that Sgr must have experienced a quite rapid change in its binding energy over the past several Gyr, even if the Sgr progenitor had one of the steepest metallicity gradients among the known MW dSphs. These multiple population bursts also created Sgr's unique chemical patterns, especially at higher metallicities, compared to those of the MW (as well as to those of other dSphs). Nevertheless, while the abundance patterns of stars presently in the Sgr core differ greatly from those of MW stars of the same metallicity, in this paper we have found that the more metal-poor stars in the Sgr tail actually have abundance patterns that more closely resemble those of MW stars at their respective [Fe/H]. The Sgr example vividly demonstrates that while the current populations of stars in dSph satellites are indeed chemically differentiated from the MW field population and that ``one could not build the MW halo from the {\\it present} MW satellites'' (as emphasized by, e.g., Unavane et al. 1996), this point is not very relevant because it is still possible that the MW halo field population could have derived from the {\\it stripped off} populations of these very same satellites. Majewski et al. (2002) and Mu\\~noz et al. (2006) have previously made the same point using the Carina dSph example.   The Sgr system can be seen as an evolutionary bridge from dSphs to the MW in galaxy evolution. On the one hand, recent models (Robertson et al.\\ 2005; Font et al.\\ 2006a,b) predict that the local halo assembled rapidly --- before SN Ia had time to occur --- with the early accretion and dissipation of a few massive satellites to produce the high values of [$\\alpha$/Fe] at low [Fe/H] seen among halo field stars. Font et al.\\ suggest the MW satellites accreted $\\sim$9 Gyr ago have all been disrupted completely, while the surviving satelliites of today were only recently accreted into the MW system within the past few Gyr on relatively circular orbits.  The implication is that these satellite systems have yet to contribute stars to the halo. On the other hand, Sgr clearly provides an example of an ongoing merger event that is not only contributing stars to the halo, but some low metallicity stars of high [$\\alpha$/Fe]; it may therefore be seen as a counterexample to the ''early accretion'' hypothesis. However, Font et al.\\ point out that none of the present MW satellites are presently located in the inner halo except Sgr and because of this Sgr may be an exceptional case. But the recent work on the Ursa Minor (Mu\\~noz et al. 2005), Leo I (Sohn et al. 2007; Mateo et al. 2008) and Carina dSphs (Mu\\~noz et al. 2006, 2008) yielding evidence for tidally stripped stars from these systems suggests that the present satellites may also have already contributed stars to the MW halo. This, coupled with the finding that some of the MW dSph stars at the lowest metallicities are indeed $\\alpha$-enhanced (Shetrone et al. 2003; Geisler et al. 2005; Tolstoy 2005), suggests that Sgr may not be the only current contributor of such stars; at the very least these systems are likely to be {\\it future} contributors of these stars, which shows that they did not all originate in early accretions.  The results of this paper would be considerably strengthened by a careful survey of the chemical abundance patterns of Sgr trailing arm stars. The significant overlap in orbital phase position along the Sgr leading arm (see Fig.\\ 1 in Paper IV) ``fuzzes out'' the resolution of the dynamical stripping time. In contrast to the stronger phase mixing in the leading arm, the dynamics of the longer trailing arm demonstrate much better energy sorting of the debris, so that position along the trailing tail is much better correlated to the amount of time since a star was stripped from the Sgr core. Further scrutiny of cleanly isolated trailing arm stars may reveal even more clear chemical trends with dynamical age. We will investigate this in future work."
    },
    "0911/0911.3772_arXiv.txt": {
        "abstract": "{ We present $^{12}$CO(1--0) and $^{12}$CO(2--1) maps of the interacting Seyfert 2/LINER galaxy \\nnn\\ obtained with the IRAM interferometer at resolutions of 2\\farcs1 $\\times$ 1\\farcs4 and 1\\farcs1 $\\times$ 0\\farcs7, respectively. We also present single-dish IRAM 30\\,m observations of the central region of \\nnn\\ for the $^{12}$CO(1--0), $^{12}$CO(2--1), and HCN(1--0) transitions at resolutions of 22\\arcsec, 12\\arcsec, and $29^{\\prime\\prime}$, respectively. The CO emission is distributed over a disk of diameter $\\sim$16\\arcsec ($\\sim$2.2\\,kpc), within which are several, randomly distributed peaks. The strongest peak does not coincide with the nucleus, but is instead offset from the center, $\\sim$2-3$^{\\prime\\prime}$ ($\\sim$340\\,pc) toward the west/southwest. The kinematics of the molecular component are quite regular, as is typical of a rotating disk. We also compared the $^{12}$CO distribution of \\nnn\\ with observations at other wavelengths in order to study correlations between different tracers of the interstellar medium. The \\hst/F606W WFPC2 images show flocculent spiral structures and an ``S-shape'' feature $\\simgt$60\\,pc in radius, possibly associated with a nuclear bar or with the radio jet. A two-dimensional bulge/disk decomposition of the $H$-band (\\hst/F160W) and 3.6\\,$\\mu$m (\\spitzer/IRAC) images reveals a circumnuclear ``ring'' $\\sim$10-14\\arcsec\\ in diameter, roughly coincident in size with the CO disk and with a star-forming ring previously identified in ionized gas. This ring is not present in the near-infrared (NIR) $J-K$ color image, nor is it present in the ``dust-only'' image constructed from the 8\\,\\micron\\ IRAC map. The implication is that the excess residual ring is stellar, with colors similar to the surrounding disk. We interpret this ring, visible in ionized gas, which appears as stars in the NIR, and with no sign of hot dust, as due to a red super giant population at least 10--15\\,Myr old. However, star formation is still ongoing in the disk and in the ring itself. Using NIR images, we computed the gravity torques exerted by the stellar potential on the gas. The torques are predominantly positive in both $^{12}$CO(1--0) and $^{12}$CO(2--1), suggesting that gas is not flowing into the center, and less than 5\\% of the gas angular momentum is exchanged in each rotation. This comes from the regular and almost axisymmetric total mass and gas distributions in the center of the galaxy. In \\nnn, the AGN is apparently not being actively fueled in the current epoch. ",
        "introduction": "Since molecular gas is the predominant phase of the interstellar medium (ISM) in the inner regions of spiral galaxies, CO lines represent an optimum tracer of nuclear gas dynamics, active galactic nuclei (AGN) fueling mechanisms, and their link with circumnuclear star formation. Although most galaxies host super massive black holes (SMBHs) and the gas accretion phenomenon is usually invoked to explain nuclear activity in galaxies, the nature of this activity is still not well known. The feeding of an AGN through accretion depends on an adequate supply of gas whose angular momentum has been reduced enough to enable inflow on the small spatial scales surrounding the BH. Angular momentum must be removed from the disk gas \\citep[e.g.,][and references therein]{jogee06}, a process that can be accomplished through non-axisymmetric perturbations of internal or external origin. In the first case, they arise from disk instabilities and, in the second, from galaxy collisions, mergers, and mass accretion \\citep{heckman86}. Either way, they usually manifest themselves as density waves, such as large-scale spirals or as bars and their gravity torques \\citep[e.g.,][]{sakamoto99,francoise01}, or as more localized phenomena, including nested nuclear bars \\citep[e.g.,][]{friedli93}, lopsidedness or $m = 1$ perturbations \\citep[e.g.,][]{shu90,santi00}, or warped nuclear disks \\citep[e.g.,][]{schinnerer00a,schinnerer00b}.  To better understand  the mechanisms for gas fueling of AGN, we started a high-resolution and high-sensitivity CO survey of nearby active galaxies at the IRAM Plateau de Bure Interferometer (PdBI), the NUclei of GAlaxies (NUGA) project \\citep[][]{santi03}. The galaxies of the NUGA sample already studied show surprising results. In fact, there is no unique circumnuclear molecular gas feature linked with nuclear activity, but instead a variety of molecular gas morphologies which characterize the inner kpc of active galaxies. These morphologies include one- and two-armed instabilities \\citep{santi03}, well-ordered rings and nuclear spirals \\citep{francoise04,vivi08}, circumnuclear asymmetries \\citep{melanie05} and large-scale bars \\citep{fred07,leslie08}. The analysis of the torques exerted by the stellar gravitational potential on the molecular gas of the NUGA sample has shown that the gas can be driven away from the AGN (e.g. for NGC\\,4321) or toward it (e.g. for NGC\\,2782, NGC\\,3147, and NGC\\,4579). However, the velocities observed for NUGA are too small to correspond to the AGN feedback models where violent molecular outflows and superwinds are expected \\citep[e.g][]{narayanan06,hopkins06}.  The different morphologies we find are probably related to the various timescales \\citep{santi05}. Large-scale bars can transport gas inward efficiently \\citep[e.g.,][]{francoise85,sakamoto99}, and it appears that they can also drive powerful starbursts \\citep[e.g.,][]{knapen02,jogee05}. Nevertheless, a clear correlation between large-scale bars and nuclear activity has not yet been found \\citep[e.g.,][]{mulchaey97}, probably because the timescales for bar-induced gas inflow and AGN duty cycles are very different. Bars drive inflow over timescales \\citep[$\\gtrsim$300 Myr,][]{jogee05} larger than those of AGN accretion-rate duty cycles \\citep[$\\sim$1-10 Myr,][]{heckman04,hopkins06,king07}, and active accretion seems to occur only intermittently over the lifetime of a galaxy \\citep[][]{ferrarese01,marecki03,janiuk04,hopkins06,king07}. This implies that most AGN are in an intermediate phase between active accretion episodes, making the detection of galaxies with nuclear accretion rather difficult.  Viscosity rather than self-gravity can also play a significant role in the fueling process. Viscous torques, generally weak and with timescales quite long at large radii, in combination with gravitational torques can coordinate efforts to produce recurrent episodes of activity during the typical lifetime of any galaxy \\citep{santi05}. Viscous torques can produce gas inflow on scales $\\sim$100-200 pc if they act on a contrasted nuclear ring distribution and in the absence of particularly strong positive gravitational torques. In NGC\\,4579, the efficiency of viscosity may be comparable to the efficiency of gravity torques in the inner $\\sim$50\\,pc \\citep[][]{santi09}.  This paper, dedicated to the galaxy \\nnn, is the latest of the NUGA series where results obtained for the galaxies of the sample are described on a case-by-case basis. \\nnn\\ ($D$ = 28 Mpc for $H_{0}$ = 73 km s$^{-1}$Mpc$^{-1}$) is an interacting galaxy \\citep[e.g.,][]{rampazzo95,vivi04,iono05}, classified as a Seyfert 2 by \\citet{gonzalez96} and as a LINER by \\citet{veilleux95}, and of early and unbarred Hubble type (SAa pec). \\nnn\\ and its late-type companion NGC\\,5954 (LINER/Seyfert 2, SAB(rs)cd pec) constitute a binary system (VV\\,244, Arp\\,91) where the two galaxies are separated by a projected distance of 5.8 kpc ($\\sim$43\\arcsec). They show clear signs of interaction visible in the distorted morphology, the presence of a tidal bridge (or distorted arm) connecting the two galaxies, and of prominent star-forming regions. Both galaxies have circumnuclear starbursts that may have been induced by the interaction \\citep{gonzalez96}.  \\nnn\\ hosts a compact radio core and jet, revealed by high-resolution radio continuum observations with MERLIN \\citep{melanie07a}. The jet is resolved at 18\\,cm, and after beam deconvolution is roughly $\\sim$0\\farcs3 in length (40\\,pc), with a position angle (PA) of $\\sim$10$^\\circ$. The small-scale radio continuum structure at 20\\,cm is similar in orientation, but with lower spatial resolution ($\\sim$1\\farcs5), slightly more toward the east, PA$\\sim$25$^\\circ$ \\citep{jenkins84}. This jet-like elongation as seen at lower resolution was first referred to as a ``jet'' by \\citet{gonzalez96} because it is roughly aligned with the structure in the ``excitation map'', obtained by dividing [O{\\sc iii}] emission by H$\\alpha$.  The NGC\\,5953/54 pair has also been mapped in atomic and molecular gas. Both galaxies are embedded in a common \\hi\\ envelope, with a clear velocity gradient along the \\hi\\ plume extending more than 8\\,kpc to the northwest \\citep{chengalur94,iono05,haan07,haan08}. There is some indication of a faint diffuse optical counterpart to the \\hi\\ plume \\citep{chengalur94,hernandez03}. The \\hi\\ peak is also significantly displaced from the stellar disks \\citep{iono05}. The overall \\hi\\ velocity gradient runs from southeast to northwest, roughly perpendicular to the rotation in the ionized and molecular gas \\citep{hernandez03,iono05}. The most recent \\hi\\ mass determination for \\nnn\\ has been obtained by  \\citet[][]{haan08}, M$_{\\rm H\\,I}=1.1\\times$10$^{9}$\\,M$_{\\odot}$ (value scaled to our adopted distance of $D=28$\\,Mpc), typical of that expected for interacting galaxies of the same morphological type \\citep[][]{vivi04}.  \\begin{table} \\caption[]{Fundamental parameters for \\nnn.} \\begin{center} \\begin{tabular}{lll} \\hline \\hline Parameter  & Value$^{\\mathrm{b}}$ & References$^{\\mathrm{c}}$ \\\\ \\hline $\\alpha_{\\rm J2000}$$^{\\mathrm{a}}$ & 15$^h$34$^m$32.36$^s$ & (1) \\\\ $\\delta_{\\rm J2000}$$^{\\mathrm{a}}$ & 15$^{\\circ}$11$^{\\prime}$37\\farcs70 & (1) \\\\ $\\alpha_{\\rm dyn}$$^{\\mathrm{a}}$   & 15$^h$34$^m$32.38$^s$ & (2) \\\\ $\\delta_{\\rm dyn}$$^{\\mathrm{a}}$   & 15$^{\\circ}$11$^{\\prime}$37\\farcs59 & (2) \\\\ $V_{\\rm sys, hel}$              & 1990 km\\,s$^{-1}$ & (1) \\\\ RC3 Type                        & SAa pec & (3) \\\\ Nuclear Activity                & S2/LINER & (4) $\\&$ (5) \\\\ Inclination                     & 42$^{\\circ}$ & (1) \\\\ Position Angle                  & 45$^{\\circ}$ $\\pm$ 1$^{\\circ}$\\ & (1) \\\\ Distance                        & 28\\,Mpc ($1^{\\prime\\prime} = 136\\,{\\rm pc}$) & (3) \\\\ L$_{B}$                         & $4.9 \\times 10^{9}$\\,L$_{\\odot}$ & (6) \\\\ M$_{\\rm H\\,I}$                  & $1.1 \\times 10^{9}$\\,M$_{\\odot}$ & (7) \\\\ M$_{\\rm H_{2}}$                 & $2.2 \\times 10^{9}$\\,M$_{\\odot}$ & (8) \\\\ M$_{\\rm dust}$(60 and 100\\,$\\mu$m)& $2.7 \\times 10^{6}$\\,M$_{\\odot}$ & (6) \\\\ L$_{\\rm FIR}$                   & $1.4 \\times 10^{10}$\\,L$_{\\odot}$ & (9) \\\\ \\hline \\hline \\end{tabular} \\label{table1} \\end{center} \\begin{list}{}{} \\item[$^{\\mathrm{a}}$] ($\\alpha_{\\rm J2000}$, $\\delta_{\\rm J2000}$) is the phase tracking center of our $^{12}$CO observations, ($\\alpha_{\\rm dyn}$, $\\delta_{\\rm dyn}$) is the dynamical center derived from radio observations for the core of \\nnn\\ by \\citet{melanie07a}. \\item[$^{\\mathrm{b}}$] Luminosity and mass values extracted from the literature have been scaled to the distance of $D=28$\\,Mpc. \\item[$^{\\mathrm{c}}$] (1) This paper; (2) \\citet{melanie07a}; (3) NASA/IPAC Extragalactic Database (NED); (4) \\citet{gonzalez96}; (5) \\citet{veilleux95}; (6) \\citet{vivi04}; (7) \\citet{haan08}; (8) \\citet{iono05}; (9) {\\it IRAS} Catalog. \\end{list} \\end{table}  The molecular gas distribution in \\nnn\\ is symmetric with a short extension pointing toward NGC\\,5954 \\citep[e.g.,][]{yao03,iono05}. The H$_{2}$ mass content estimated by \\citet{iono05} is $2.2\\times10^{9}\\,$M$_{\\odot}$ (scaled to the distance of $D=28$\\,Mpc for \\nnn), higher than the molecular hydrogen mass found for the interacting companion NGC\\,5954 (M$_{\\rm H_{2}}$ = $1.1 \\times 10^{9}\\,$M$_{\\odot}$). Higher-order transitions of the CO molecule have also been detected in \\nnn, including the $^{12}$CO(3--2) line by \\citet{yao03} with the James Clerk Maxwell Telescope (FWHM$\\sim$15\\arcsec), suggesting a high excitation of the carbon monoxide, indicative of dense and hot gas. Table \\ref{table1} summarizes the fundamental characteristics of \\nnn.  The structure of this paper is as follows.  In Sect. \\ref{sec:obs}, we describe our new observations of \\nnn\\ and the literature data with which we compare them. In Sects. \\ref{sec:30m} and \\ref{sec:pdbi}, we present the observational results, both single dish and interferometric, describing morphology, excitation conditions, and kinematics of the molecular gas in the inner kpc of \\nnn. Comparisons between $^{12}$CO observations and those obtained at other wavelengths are given in Sect. \\ref{sec:stars}. In Sect. \\ref{sec:torques}, we describe the computation of the gravity torques derived from the stellar potential in the inner region of \\nnn. Sect. \\ref{sec:conclusions} summarizes our results.  We will assume a distance to \\nnn\\ of $D=28$\\,Mpc \\citep[HyperLeda\\footnote{http://leda.univ-lyon1.fr},][]{paturel03} and a Hubble constant $H_0=73$\\,km\\,s$^{-1}$\\,Mpc$^{-1}$. This distance implies that 1\\arcsec\\ corresponds to 136\\,pc.  \\begin{figure*} \\centering \\includegraphics[width=0.8\\textwidth,angle=-90]{ch-map-co10_merged.eps} \\caption{$^{12}$CO(1--0) velocity channel maps observed with the IRAM PdBI+30\\,m in the nucleus of \\nnn, with a spatial resolution of 2\\farcs1 $\\times$ 1\\farcs4 (HPBW). The maps are centered on the phase tracking center of our observations assumed to be coincident with the dynamical center of the galaxy (see Sect. \\ref{sec:dyncen}). Velocity channels range from $\\Delta V = -130\\,{\\rm km\\,s^{-1}}$ to $+130\\,{\\rm km\\,s^{-1}}$ in steps of $5\\,{\\rm km\\,s^{-1}}$ relative to $V_{\\rm sys, hel}$ = 1990 km\\,s$^{-1}$ (see Sect. \\ref{sec:dyncen}). The contours run from $-0.20\\,{\\rm mJy\\,beam^{-1}}$ to $120\\,{\\rm mJy\\,beam^{-1}}$ with spacings of $20\\,{\\rm mJy\\,beam^{-1}}$.} \\label{channels10} \\end{figure*}   \\begin{figure*} \\centering  \\includegraphics[width=0.8\\textwidth,angle=-90]{ch-map-co21_merged.eps} \\caption{Same as Fig. \\ref{channels10} but for the $^{12}$CO(2--1) line, with a spatial resolution of 1\\farcs1 $\\times$ 0\\farcs7. The contours run from $-60\\,{\\rm mJy\\,beam^{-1}}$ to $150\\,{\\rm mJy\\,beam^{-1}}$ with spacings of $30\\,{\\rm mJy\\,beam^{-1}}$.} \\label{channels21} \\end{figure*} ",
        "conclusions": "} The molecular gas in the Seyfert 2/LINER galaxy \\nnn\\ has been mapped with high resolution (2\\farcs1 $\\times$ 1\\farcs4 for the $^{12}$CO(1--0) line and 1\\farcs1 $\\times$ 0\\farcs7 for the $^{12}$CO(2--1) line) inside a radius of $\\sim$20\\arcsec\\ ($\\sim$2.7\\,kpc). The $^{12}$CO emission is distributed over a disk of diameter of $\\sim$16$^{\\prime\\prime}$ with several peaks randomly distributed. The strongest intensity peak is not located in the center of the galaxy but shifted $\\sim$2-3$^{\\prime\\prime}$ toward the west/southwest. The kinematics of the molecular gas show a general regularity with some local wiggles especially to the west of the nucleus maybe associated with the intensity peak toward the SW.  By comparing the molecular gas distribution with observations at other wavelengths, we studied correlations between different tracers of the ISM. Optical and NIR morphology has been studied by analyzing different images. The F606W/HST  broad-band optical image of NGC\\,5953 shows  a flocculent spiral structure and  an ``S'' like-bar of $\\sim$250\\,pc and $\\simgt$60\\,pc in size, respectively. These structures are absent in the $1.6\\,{\\rm \\mu m}$ (\\hst/NICMOS/F160W) image and in the IRAC $8\\,{\\rm \\mu m}$ ``dust-only'' one; instead in the optical and at $1.6\\,{\\rm \\mu m}$ we identified at $\\sim$3\\arcsec\\ west of center the foreground star studied by \\citet{piraf90}. The stellar structure has been investigated by performing a two-dimensional B/D decomposition on the ground-based $H$-band image and IRAC 3.6\\,$\\mu$m image. Both sets of residuals show a circumnuclear stellar ring with a radius of $\\sim$5-7\\arcsec, approximately coincident in size with the $^{12}$CO disk. The presence of the ring in both sets of residuals suggests that it is not the typical ``doughnut'' associated with the diffraction limit and IRAC PRF incompatibilities, but rather is a real feature. The size of this stellar ring is the same of the circumnuclear star-forming ring seen in H$\\alpha$ with previous observations. In addition, since this ring is visible in the NIR as stars but not as hot dust emission ($J-K$), we interpret it as a RSG population at least 10\\,Myr old.  There is an apparent counter-rotation between gas and stars inside the ring, and stars outside. The ring could be the separation between the two kinematically decoupled components (KDC). The formation of a KDC could be explained as a result of the ongoing interaction, and could be in its early stages. Alternatively, this apparent counter-rotation could be due to a warp of the plane of the disk. The ring would then be the start of the warp.    Using NIR images we found that gravity torques acting on the gas are predominantly positive in the region $\\sim$100--400\\,pc in both  $^{12}$CO(1--0) and $^{12}$CO(2--1), indicating that the gas is not fueling the central $\\sim$100\\,pc region, down to the effective spatial resolution of our observations. In addition, the absolute values of these torques are very small, certainly due to the almost axisymmetric total mass and gas distributions in the center of the galaxy. The AGN in \\nnn\\ is apparently not being actively fueled in the current epoch."
    },
    "0911/0911.1777_arXiv.txt": {
        "abstract": "Computation of the marginal likelihood from a simulated posterior distribution is central to Bayesian model selection but is computationally difficult.  The often-used harmonic mean approximation uses the posterior directly but is unstably sensitive to samples with anomalously small values of the likelihood.  The Laplace approximation is stable but makes strong, and often inappropriate, assumptions about the shape of the posterior distribution.  It is useful, but not general.  We need algorithms that apply to general distributions, like the harmonic mean approximation, but do not suffer from convergence and instability issues.  Here, I argue that the marginal likelihood can be reliably computed from a posterior sample by careful attention to the numerics of the probability integral.  Posing the expression for the marginal likelihood as a Lebesgue integral, we may convert the harmonic mean approximation from a sample statistic to a quadrature rule.  As a quadrature, the harmonic mean approximation suffers from enormous truncation error as consequence .  This error is a direct consequence of poor coverage of the sample space; the posterior sample required for accurate computation of the marginal likelihood is much larger than that required to characterize the posterior distribution when using the harmonic mean approximation. In addition, I demonstrate that the integral expression for the harmonic-mean approximation converges slowly at best for high-dimensional problems with uninformative prior distributions. These observations lead to two computationally-modest families of quadrature algorithms that use the full generality sample posterior but without the instability.  The first algorithm automatically eliminates the part of the sample that contributes large truncation error.  The second algorithm uses the posterior sample to assign probability to a partition of the sample space and performs the marginal likelihood integral directly.  This eliminates convergence issues. The first algorithm is analogous to standard quadrature but can only be applied for convergent problems.  The second is a hybrid of cubature: it uses the posterior to discover and tessellate the subset of that sample space was explored and uses quantiles to compute a representative field value.  Qualitatively, the first algorithm improves the harmonic mean approximation using numerical analysis, and the second algorithm is an adaptive version of the Laplace approximation.  Neither algorithm makes strong assumptions about the shape of the posterior distribution and neither is sensitive to outliers.  Based on numerical tests, we recommend a combined application of both algorithms as consistency check to achieve a reliable estimate of the marginal likelihood from a simulated posterior distribution.  \\par\\bigskip\\noindent {\\bf Keywords:} Bayesian computation, marginal likelihood, algorithm, Bayes factors, model selection ",
        "introduction": "\\label{sec:intro}  A Bayesian data analysis specifies joint probability distributions to describe the relationship between the prior information, the model or hypotheses, and the data.  Using Bayes theorem, the posterior distribution is uniquely determined from the conditional probability distribution of the unknowns given the observed data.  The posterior probability is usually stated as follows: \\begin{equation} P(\\param|\\model,\\data) = \\frac{\\pi(\\param|\\model)L(\\data|\\param,\\model)}{Z} \\label{eq:bayes} \\end{equation} where \\begin{equation} Z \\equiv P(\\data|\\model) = \\int d\\param\\,\\pi(\\param|\\model) L(\\data|\\param,\\model) \\label{eq:evidence} \\end{equation} is the marginal likelihood.  The symbol ${\\cal M}$ denotes the assumption of a particular model and the parameter vector $\\param\\in\\Omega$.  For physical models, the sample space $\\Omega$ is most often a continuous space.  In words, equation (\\ref{eq:bayes}) says: the probability of the model parameters given the data and the model is proportional to the prior probability of the model parameters and the probability of the data given the model parameters.  The posterior may be used, for example, to infer the distribution of model parameters or to discriminate between competing hypotheses or models. The latter is particularly valuable given the wide variety of astronomical problems where diverse hypotheses describing heterogeneous physical systems is the norm \\citep[see][for a thorough discussion of Bayesian data analysis]{Gelman.etal:03}.  For parameter estimation, one often considers $P(\\data|\\model)$ to be an uninteresting normalization constant.  However, equation (\\ref{eq:evidence}) clearly admits a meaningful interpretation: it is the support or \\emph{evidence} for a model given the data.  This see this, assume that the prior probability of some model, $\\model_j$ say, is $P(\\model_j)$.  Then by Bayes theorem, the probability of the model given the data is $P(\\model_j|\\data) = P(\\model_j) P(\\data|\\model_j) / P(\\data)$.  The posterior odds of Model $j=0$ relative to Model $j=1$ is then: \\begin{equation*} \\frac{P(\\model_0|\\data)}{P(\\model_1|\\data)} = \\frac{P(\\model_0)}{P(\\model_1)} \\frac{P(\\data|\\model_0) }{P(\\data|\\model_1)}. \\end{equation*} If we have information about the ratio of prior odds, $P(\\model_0)/P(\\model_1)$, we should use it, but more often than not our lack of knowledge forces a choice of $P(\\model_0)/P(\\model_1)=1$. Then, we estimate the relative probability of the models given $\\data$ over their prior odds by the Bayes factor $P(\\data|\\model_0)/P(\\data|\\model_1)$ \\citep[see][for a discussion of additional concerns]{Lavine.Schervish:99}.  When there is no ambiguity, we will omit the explicit dependence on $\\model$ of the prior distribution, likelihood function, and marginal likelihood for notational convenience.  The Bayes factor has a number of attractive advantages for model selection \\citep{Kass.Raftery:95}: (1) it is a consistent selector; that is, the ratio will increasingly favor the true model in the limit of large data; (2) Bayes factors act as an Occam's razor, preferring the simpler model if the ``fits'' are similar; and (3) Bayes factors do not require the models to be nested in any way; that is, the models and their parameters need not be equivalent in any limit.  There is a catch: direct computation of the marginal likelihood (eq. \\ref{eq:evidence}) is intractable for most problems of interest. However, recent advances in computing technology together with developments in Markov chain Monte Carlo (MCMC) algorithms have the promise to compute the posterior distribution for problems that have been previously infeasible owing to dimensionality or complexity.  The posterior distribution is central to Bayesian inference: it summarizes all of our knowledge about the parameters of our model and is the basis for all subsequent inference and prediction for a given problem. For example, current astronomical datasets are very large, the proposed models may be high-dimensional, and therefore, the posterior sample is expensive to compute.  However, once obtained, the posterior sample may be exploited for a wide variety of tasks.  Although dimension-switching algorithms, such as reversible-jump MCMC \\citep{Green:95} incorporate model selection automatically without need for Bayes factors, these simulations appear slow to converge for some of our real-world applications.  Moreover, the marginal likelihood may be used for an endless variety of tests, ex post facto.  \\citet{Newton.Raftery:94} presented a formula for estimating $Z$ from a posterior distribution of parameters.  They noted that a MCMC simulation of the posterior selects values of $\\param\\in\\Omega$ distributed as \\[ Z\\times P(\\param|\\data) = L(\\data|\\param) \\pi(\\param) \\] and, therefore, \\begin{equation} Z\\times \\int_\\Omega d\\param\\,\\frac{P(\\param|\\data)}{L(\\data|\\param)} = \\int_\\Omega d\\param\\,\\pi(\\param) = 1 \\label{eq:Zdef0} \\end{equation} or \\begin{equation} \\frac1Z = \\int_\\Omega d\\param\\,\\frac{P(\\param|\\data)}{L(\\data|\\param)} = E\\left[\\frac{1}{L(\\param|\\data)}\\right]_{P(\\param|\\data)}, \\label{eq:hmean} \\end{equation} having suppressed the explicit dependence on $\\model$ for notational clarity.  This latter equation says that the marginal likelihood is the harmonic mean of the likelihood with respect to the posterior distribution.  It follows that the harmonic mean computed from a sampled posterior distribution is an estimator for the marginal likelihood, e.g.: \\begin{equation} {\\tilde Z} = \\left[\\frac1N\\sum_{i=1}^N \\frac{1}{L(\\data|\\theta_i)}\\right]^{-1}. \\label{eq:hsamp} \\end{equation} Unfortunately, this estimator is prone to domination by a few outlying terms with abnormally small values of $L_j$ \\citep[e.g. see][and references therein]{Raftery.etal:07}.  \\citet{Wolpert:02} describes convergence criteria for equation (\\ref{eq:hsamp}) and \\citet{Chib.Jeliazkov:01} present augmented approaches with error estimates.  Alternative approaches to computing the marginal likelihood from the posterior distribution have been described at length by \\citet{Kass.Raftery:95}.  Of these, the Laplace approximation, which approximates the posterior distribution by a multidimensional Gaussian distribution and uses this approximation to compute equation (\\ref{eq:evidence}) directly, is the most widely used.  This seems to be favored over equation (\\ref{eq:harmonic}) because of the problem with outliers and hence because of convergence and stability.  In many cases, however, the Laplace approximation is far from adequate in two ways. First, one must identify all the dominant modes, and second, modes may not be well-represented by a multidimensional Gaussian distribution for problems of practical interest, although many promising improvements have been suggested \\citep[e.g.][]{DiCicio.etal:97}.  \\citet{Trotta:07} explored the use of the Savage-Dickey density ratio for cosmological model selection \\citep[see also][for a full review of the model selection problem for cosmology]{Trotta:08}.  Finally, we may consider evaluation of equation (\\ref{eq:evidence}) directly.  The MCMC simulation samples the posterior distribution by design, and therefore, can be used to construct volume elements in $k$-dimensional parameter space, $d\\param$, e.g. when $\\Omega\\subset\\mathbb{R}^k$.  Although the volume will be sparsely sampled in regions of relatively low likelihood, these same volumes will make little contribution to equation (\\ref{eq:evidence}).  The often-used approach from computational geometry, Delaney triangulation, maximizes the minimum angle of the facets and thereby yields the ``roundest'' volumes.  Unfortunately, the standard procedure scales as ${\\cal O}(kN^2)$ for a sample of $N$ points.  This can be reduced to ${\\cal O}(N\\log N + N^{k/2})$ using the flip algorithm with iterative construction \\citep{EdelsbrunnerShah:96} but this scaling is prohibitive for large $N$ and $k$ typical of many problems.  Rather, in this paper, we consider the less optimal but tractable kd-tree for space partitioning.  In part, the difficulty in computing the marginal likelihood from the sampled posterior has recently led \\citet[``nesting sampling'']{Skilling:06} to suggest an algorithm to simulate the marginal likelihood rather than the posterior distribution.  This idea has been adopted and extended by cosmological modelers \\citep{Mukherjee.etal:06,Feroz.Hobson:08}.  The core idea of nesting sampling follows by rewriting equation (\\ref{eq:Zdef0}) as a double integral and swapping the order of integration, e.g. \\begin{equation} Z = \\int_\\Omega d\\theta \\pi(\\theta) \\int_0^{L(\\data|\\theta)} dy = \\int_0^{\\sup\\{L(\\data|\\theta): \\theta\\in\\Omega\\}} dy \\int_{L(\\data|\\theta)>y} d\\theta \\pi(\\theta). \\label{eq:nest} \\end{equation} The nested sampler is a Monte Carlo sampler for the likelihood function $L$ with respect to the prior distrbution $\\pi$ so that $L>y$.  The generalization of the construction in equation (\\ref{eq:nest}) for general distributions and multiple dimensions is the Lebesgue integral (see \\S\\ref{sec:intgr}).  Clearly, this procedure has no problems with outliers with small values of $L(\\data|\\param)$.  Of course, any algorithm implementing nested sampling must still thoroughly sample the multidimensional posterior distribution and so retains all of the intendant difficulties that MCMC has been designed to solve.  In many ways, the derivation of the nested sampler bears a strong resemblance to the derivation of the harmonic mean but without any obvious numerical difficulty.  This led me to a careful study of equations (\\ref{eq:evidence}) and (\\ref{eq:Zdef0}) to see if the divergence for small value of likelihood could be addressed.  Indeed they can, and the following sections describe two algorithms based on each of these equations.  These new algorithms retain the main advantage of the harmonic mean approximation (HMA): direct incorporation of the sampled posterior without any assumption of a functional form.  In this sense, they are fully and automatically adaptive to any degree multimodality given a sufficiently large sample.  We begin in \\S\\ref{sec:intgr} with a background discussion of Lebesgue integration applied to probability distributions and Monte Carlo (MC) estimates.  We apply this in \\S\\ref{sec:evid} to the marginal likelihood computation. This development both illuminates the arbitrariness in the HMA from the numerical standpoint and leads to an improved approach outlined in \\S\\ref{sec:algo}.  In short, the proposed approach is motivated by methods of numerical quadrature rather than sample statistics.  Examples in \\S\\ref{sec:examples} compare the application of the new algorithms to the HMA and the Laplace approximation.  The overall results are discussed and summarized in \\S\\ref{sec:summary}. ",
        "conclusions": "\\label{sec:summary}  In summary, much of the general measure-theoretic underpinning of probability and statistics naturally leads naturally to the evaluation of expectation values.  For example, the harmonic mean approximation \\citep[HMA,][]{Newton.Raftery:94} for the marginal likelihood has large variance and is slow to converge \\citep[e.g.][]{Wolpert:02}.  On the other hand, the use of analytic density functions for the likelihood and prior permits us to take advantage of less general but possibly more powerful computational techniques.  In \\S\\S\\ref{sec:intro}--\\ref{sec:algo} we diagnose the numerical origin of the insufficiencies of the HMA using Lebesgue integrals.  There are two culprits: 1) the integral on the left-hand side of equation (\\ref{eq:Zdef0}) may diverge if the measure function $M(Y=L^{-1})$ from equation (\\ref{eq:measureX}) decreases too slowly; and 2) truncation error may dominate the quadrature of the left-hand side of equation (\\ref{eq:Zdef0}) unless the sample is appropriately truncated.  Using numerical quadrature for the marginal likelihood integral (eqns. \\ref{eq:evidence} and \\ref{eq:measureX}) leads to improved algorithms: the \\emph{Numerical Lebesgue Algorithm} (NLA) and the \\emph{Volume Tessellation Algorithm} (VTA).  Our proposed algorithms are a bit more difficult to implement and have higher computational complexity than the simple HMA, but the overall CPU time is rather modest compared to the computational investment required to produce the MCMC-sampled posterior distribution itself.  For a sample of size $N$, the sorting required by NLA and VTA has computational complexity of ${\\cal O}(N\\log N)$ and ${\\cal O}(N\\log^2 N)$, respectively, rather than ${\\cal O}(N)$ for the harmonic mean. Nonetheless, the computational time is a fraction of second to minutes for typical values of $10^5<N<10^8$ (see \\S\\ref{sec:algo}).  The geometric picture behind NLA is exactly that for Lebesgue integration.  Consider integrating a function over a two-dimensional domain.  In standard Riemann quadrature, one chops the domain into rectangles and adds up their area.  The sum can be refined by subdividing the rectangles; in the limit of infinitesimal area, the resulting sum is the desired integral.  In the Lebesgue approach, one draws horizontal slices through the surface and adds up the area of the horizontal rectangles formed from the width of the slice and the vertical distance between slices.  The sum can be refined by making the slices thinner when needed; in the limit of slices of infinitesimal height, the resulting sum is the desired integral.  In the Riemann case, we multiply the box area in the domain, $dA$, by the function height, $f$.  In the Lebesgue, we multiply the slice height in the range, $df$, by the domain area, $A$ (see Fig. \\ref{fig:geom}). Both algorithms easily generalize to higher dimension.  For the Lebesgue integral, the slices become level sets on the hypersurface implied by the integrand.  Therefore the Lebesgue approach always looks one-dimensional in the level-set value; the dimensionality $k$ is `hidden' in the area of domain (\\emph{hypervolume} $A$ for $k>3$) computed by the measure function $M(Y)$.  The level-set value for the NLA is $Y=1/L(\\data|\\param)$.  Once determined, NLA applies the trapezoidal rule to the sum over slices and compute the upper and lower rectangle sums as bounds.  Clearly, the error control on this algorithm might be improved by using more of the information about the run of $A$ with $f$.  Having realized that the practical failure of the harmonic mean approximation is a consequence of the sparsely sampled parameter-space domain, NLA addresses the problem by determining a well-sampled subset $\\Omega_s\\subset\\Omega$ from the MCMC sample, ex post facto. Restricted to this subset, $\\Omega_s$, the value of the integral $J$ on the right-hand side of equation (\\ref{eq:Zdef0}) is less than unity.  We determine $\\Omega_s$ by a binary space partitioning (BSP) tree and compute $J$ from this partition.  A BSP tree recursively partitions a the k-dimensional parameter space into convex subspaces. The VTA is implemented with a kd-tree \\citep{Cormen.etal:01} for simplicity.  In addition, one may use VTA by itself to compute equation (\\ref{eq:evidence}) directly.  Judged by bias and variance, the test examples do not suggest a strong preference for either the NLA or the VTA.  However, both are clearly better than the HMA or the Laplace approximation.  Conversely, because these algorithms exploit the additional structure implied by smooth, well-behaved likelihood and prior distribution functions, the algorithms developed here will be inaccurately and possibly fail miserably for \\emph{wild} density functions.  The NLA and the VTA are not completely independent since the NLA uses the tessellation from the VTA to estimate the integral $J$.  However, the value of the integral $K$ tends to dominate $Z$, that is $|\\log K|\\gg|\\log J|$, and the contributions are easily checked.  Based on current results, I tentatively recommend relying preferentially on VTA for the following reasons: 1) there is no intrinsic divergence; 2) it appears to do as well as VTA even in a high-dimensional space; and 3) there is no truncation threshold $h_\\ast$.  Figure \\ref{fig:k2d} illustrates the potential for volume artifacts that could lead to both bias and variance.  This error source affects both the VTA and NLA (through the computation of ${\\tilde J}$) but the affect on the NLA may be larger (\\S\\ref{sec:test}).  Additional real-world testing, especially on high-dimensional multimodal posteriors, will provide more insight.  In test problems described in this paper, I explored the effects of varying the threshold $h_\\ast$ and the kd-tree bucket size $\\barm$.  These parameters interact the sample distribution, and therefore, are likely to vary for each problem.  I also recommend implementing both the NLA, VTA, HTM, Laplace approximation and comparing the four for each problem.  We are currently testing these algorithms for astronomical inference problems too complex for a simple example; the results will be reported in future papers.  An implementation of these algorithms will be provided in the next release of the UMass Bayesian Inference Engine \\citep[BIE,][]{Weinberg:10}.  There are several natural algorithmic extensions and improvements not explored here.  \\S\\ref{sec:intgr} describes a smoothed approximation to the computation of $M(Y)$ (eqns. \\ref{eq:nummeas}--\\ref{eq:ItildeN}) rather than the step function used in \\S\\ref{sec:algo}.  The direct integration of equation (\\ref{eq:evidence}) currently ignores the location of field values in each cell volume.  At the expense of CPU time, the accuracy might be improved by fitting the sampled points with low-order multinomials and using the fits to derive a cubature algorithm for each cell.  In addition, a more sophisticated tree structure may decrease the potential for bias by providing a tessellation with ``rounder'' cells.  In conclusion, the marginal likelihood \\[Z=\\int d\\param\\pi(\\param|\\model)L(\\data|\\param,\\model)\\] may be reliably computed from a Monte Carlo posterior sample though careful attention to the numerics.  We have demonstrated that the error in the HMA is due to samples with very low likelihood values but significant prior probability.  It follows that their posterior probability also very low, and these states tend to be outliers.  On the other hand, the converged posterior sample is a good representation of the posterior probability density by construction.  The proposed algorithms define the subdomain $\\Omega_s\\subset\\Omega$ dominated by and well-sampled by the posterior distribution and perform the integrals in equation (\\ref{eq:Zdef0}) over $\\Omega_s$ rather than $\\Omega$.  Although more testing is needed, these new algorithms promise more reliable estimates for $Z$ from an MCMC simulated posterior distribution with more general models than previous algorithms can deliver."
    },
    "0911/0911.1294_arXiv.txt": {
        "abstract": "In reionized regions of the Universe, gas can only collapse to form stars in dark matter (DM) haloes which grow to be sufficiently massive. If star formation is prevented in the minihalo progenitors of such DM haloes at redshifts $z$~$\\ga$~20, then these haloes will not be self-enriched with metals and so may host Population (Pop) III star formation.  We estimate an upper limit for the abundance of Pop III star clusters which thus form in the reionized Universe, as a function of redshift.  Depending on the minimum DM halo mass for star formation, between of the order of one and of the order of a thousand Pop~III star clusters per square degree may be observable at 2 $\\la$ $z$ $\\la$ 7.  Thus, there may be a sufficient number density of Pop~III star clusters for detection in surveys such as the Deep-Wide Survey (DWS) to be conducted by the {\\it James Webb Space Telescope} (JWST). We predict that Pop~III clusters formed after reionization are most likely to be found at $z$ $\\ga$ 3 and within $\\sim$ 40 arcsec ($\\sim$ 1 Mpc comoving) of DM haloes with masses of $\\sim$ 10$^{11}$ M$_{\\odot}$, the descendants of the haloes at $z$ $\\sim$ 20 which host the first galaxies that begin reionization. However, if star formation is inefficient in the haloes hosting Pop~III clusters due to the photoionizing background radiation, these clusters may not be bright enough for detection by the Near-Infrared Camera which will conduct the DWS.  Nonetheless, if the stellar initial mass function (IMF) is top-heavy the clusters may have sufficiently high luminosities in both Ly$\\alpha$ and He~{\\sc ii}~$\\lambda$1640 to be detected and for constraints to be placed on the Pop~III IMF. While a small fraction of DM haloes with masses as high as $\\sim$ 10$^9$ M$_{\\odot}$ at redshifts $z$ $\\la$ 4 are not enriched due to star formation in their progenitors, external metal enrichment due to galactic winds is likely to preclude Pop~III star formation in a large fraction of otherwise unenriched haloes, perhaps even preventing star formation in pristine haloes altogether after reionization is complete at $z$~$\\sim$~6. ",
        "introduction": "How long did the epoch of primordial star formation last? This is an important and timely question, as a primary goal of next generation telescopes, such as the {\\it James Webb Space Telescope} (JWST; e.g. Gardner et al. 2006; Windhorst et al. 2006) and the {\\it European Extremely Large Telescope}\\footnote{http://www.eso.org/sci/facilities/eelt} (E-ELT), is to detect primordial galaxies and star clusters in the early Universe.  While these telescopes may allow to detect individual supernovae (SNe) exploding within the first Population (Pop)~III star clusters and galaxies (see e.g. Wise \\& Abel 2005; Haiman 2008), the emission from the stellar populations within the first dwarf galaxies at $z$ $\\ga$ 10 is likely too dim for even the JWST to detect (e.g. Ricotti et al. 2008; Johnson et al. 2009).  Instead, only very rare, massive primordial galaxies at such high redshifts, or dwarf galaxies and star clusters forming at lower redshift may be bright enough for detection (see e.g. Ciardi \\& Ferrara 2001; Scannapieco et al. 2003). Therefore, whether or not Pop~III star clusters can be found by next generation telescopes critically depends on the amount of primordial star formation that takes place at relatively low redshift.    The hierarchical nature of structure formation dictates that more massive dark matter (DM) haloes are built up from smaller ones, and that more massive haloes generally have progenitors in which star formation may take place at earlier times. In order for a Pop~III star cluster to form in a DM halo, star formation in its progenitor haloes may have to be completely suppressed, lest SNe enrich the primordial gas with heavy elements (but see e.g. Jimenez \\& Haiman 2006; Pan \\& Scalo 2007; Dijkstra \\& Wyithe 2007). Exposure to a persistant photoionizing background, such as takes place in reionized regions of the Universe at high redshifts, can suppress star formation in DM haloes with masses substantially higher than those in which the first stars (e.g. Abel et al. 2002; Bromm et al. 2002; Yoshida et al. 2003) and galaxies (e.g. Greif et al. 2008; Bromm et al. 2009) form, completely preventing the infall of gas into haloes with circular velocities up to $\\sim$ 30 km~s$^{-1}$, corresponding to virial masses of $\\ga$ 3 $\\times$ 10$^8$ M$_{\\rm \\odot}$  at $z$ $\\la$ 10 (e.g. Efstathiou 1992; Thoul \\& Weinberg 1996; see also Kitayama \\& Ikeuchi 2000; Gnedin 2000; Ciardi \\& Ferrara 2005). If star formation in the progenitors of these more massive haloes can be suppressed by the photoionizing background during reionization, then these haloes can remain unenriched until the primordial gas collapses into them to form clusters of Pop~III stars at relatively low redshifts, e.g. at $z$~$\\la$~6.  \\begin{figure*} \\includegraphics[width=4.5in]{Figure1.ps} \\caption{The maximum mass $M_{\\rm max}(z_{\\rm reion})$ of DM haloes which have, on average, no minihalo progenitors massive enough to host star formation before a redshift $z_{\\rm reion}$, as a function of redshift $z$.  Also shown are two possible minimum halo masses, $M_{\\rm min,20}$ and $M_{\\rm min,30}$, required for photoionized gas to collapse and form stars.  For minimum halo masses of $M_{\\rm min,20}$ and $M_{\\rm min,30}$, only haloes in which star formation is suppressed by the reionization of the gas at $z_{\\rm reion}$ $\\ge$ 19 and $\\ge$ 23, respectively, can remain pristine until the reionized gas collapses to form stars. These star-forming haloes may host Pop~III star clusters luminous enough to be detected in deep surveys by the JWST (see Section 4).  } \\end{figure*}  The formation of Pop~III stars at such low redshifts has been a possibility recently studied through the use of large-scale cosmological simulations. In their high-resolution simulations of cosmic reionization, Shin et al. (2008) (also Trac \\& Cen 2007) found that haloes with masses $\\sim$ 10$^8$ M$_{\\odot}$ may host Pop~III star formation at z $\\la$ 6, but their simulation did not resolve the collapse of minihaloes, i.e. haloes with virial temperature $T_{\\rm vir}$ $\\la$ 10$^4$ K, which could have led to the enrichment of the gas at higher redshifts (see also Schneider et al. 2006; Tornatore et al. 2007).  In another recent study, Trenti et al. (2009) model the enrichment of haloes with $T_{\\rm vir}$ $\\ga$ 10$^4$ K by their progenitors, and find that significant Pop~III star formation may take place at $z$ $\\la$ 6, as well.  However, their study did not explicitly account for the suppression of star formation by photoionization, although, especially at redshifts $z$ $\\ga$ 10 (Dijkstra et al. 2004), it is possible that star formation takes place in haloes with $T_{\\rm vir}$ $\\sim$ 10$^4$ K even in the presence of a photoionizing background.  In related analytical work, Wyithe \\& Cen (2007) estimated the contribution of Pop~III stars to the reionization of the Universe, arguing that primordial star formation in neutral minihaloes may be suppressed until these haloes grow to be sufficiently massive, perhaps allowing for the formation of a significant fraction of Pop~III stars up to the end of reionization at $z$ $\\sim$ 6.  In this work, it was assumed that star formation in minihaloes was prevented by the photodissociation of molecules at very high redshifts, e.g. $z$ $\\ga$ 20 (Haiman et al. 1997).  However, more recent work has shown that molecular cooling can operate in minihaloes despite the build-up of the photodissociating background radiation field (e.g. Wise \\& Abel 2007; O'Shea \\& Norman 2008; Johnson et al. 2008; Trenti \\& Stiavelli 2009), providing motivation to relax this assumption and consider alternate scenarios for the formation of Pop~III stars at low redshifts.  In the present work, we account for star formation in minihaloes at z $\\ga$ 20, as well as for the suppression of star formation in DM haloes during reionization, with the aim of evaluating the abundance and detectability of Pop~III star clusters in reionized regions of the Universe. In Section 2, we outline the requirements for Pop~III star cluster formation in the post-reionization Universe.  In Section 3, we calculate an upper limit for the abundance of such clusters and estimate the effect that large-scale metal-enriched galactic winds have in preventing their formation, while in Section 4 we discuss the prospects for their detection in future surveys. We conclude with a discussion of our results in Section 5.  For our calculations, we assume the latest cosmological parameters reported by the {\\it Wilkinson Microwave Anisotropy Probe} (WMAP) collaboration (Komatsu et al. 2009). ",
        "conclusions": "We have shown that the suppression of star formation in DM minihaloes by photoheating during the early stages of reionization can lead to the existence of unenriched DM haloes with masses $\\ga$ 10$^8$ M$_{\\odot}$ at redshifts $z$ $\\la$ 10 which may host Pop~III star clusters.  While the abundance of such haloes is sensitively dependent on the degree to which galactic winds enrich the IGM with metals, as well as on the minimum halo mass for star formation in the reionized regions of the Universe, there may be up to thousands of bright Pop~III star clusters per square degree at redshifts 3 $\\la$ $z$ $\\la$ 6, corresponding to $\\sim$ 100 clusters within the area of the DWS to be carried out by the JWST.  This is, however, likely a strong upper limit to the abundance, as metal-enriched winds expelled by the same galaxies that reionize the primordial gas at early times may easily enrich it by these redshifts, depending on the speed of the winds and on how quickly the metals are mixed into the pristine haloes.  If this enrichment is sufficiently rapid, Pop~III star formation in pristine DM haloes may be completely prevented after reionization is complete at $z$ $\\sim$ 6.  Although the stellar mass, and so the luminosity, of Pop~III clusters formed in reionized regions of the Universe is limited by the fraction of the primordial gas that is first able to collapse into the DM haloes, we estimate that Ly$\\alpha$ emission from these star clusters could be detected by the JWST if the stellar IMF is top-heavy or if the star formation efficiency $f_{\\rm *}$ is high, i.e. of the order of 10 percent.  The He~{\\sc ii}~$\\lambda$1640 emission line, which is an indicator of stellar IMF and metallicity, may be detectable out to at least $z$ $\\sim$ 3 for the case of a top-heavy IMF, perhaps allowing to place constraints on the Pop~III IMF, at least at low redshift. Furthermore, in reionized regions Pop~III clusters may form preferentially within $\\sim$ 40 arcsec of more luminous galaxies at 3 $\\la$ $z$ $\\la$ 6, the descendants of the first galaxies which began reionization at $z$ $\\ga$ 20, perhaps aiding in the identification of these clusters in future surveys.  Current observations already yield some constraints on the abundance of Pop~III star clusters at $z$ $\\sim$ 4 (Dawson et al. 2004; Nagao et al. 2008; Wang et al. 2009; see also e.g. Fosbury et al. 2003; Shapley et al. 2003; Nagao et al. 2005; Jimenez \\& Haiman 2006; Prescott et al. 2009; Bouwens et al. 2009).  The non-detection of strong He~{\\sc ii} $\\lambda$1640 emitters at $z$ $\\sim$ 4 reported by Nagao et al. (2008) yields an upper limit on the number density of Pop~III clusters of $n_{\\rm III}$ $\\sim$ 2500 Gpc$^{-3}$, between the upper limits plotted in Fig. 3.  However, the limiting flux in this survey is only low enough to rule out the presence of the most luminous Pop~III clusters. Nonetheless, this imposes a useful constraint, and suggests that if clusters of massive Pop~III stars do form after reionization, they may in fact form in relatively low-mass haloes in which star formation is inefficient.  We note that the Pop~III IMF in star clusters formed after reionization is likely different than that of the first stars formed in minihaloes at $z$ $\\ga$ 20 (so-called Pop~III.1 stars), due to the onset of turbulence in the more massive haloes in which the clusters form (Wise et al. 2007b; Greif et al. 2008), to the possible enhanced abundances of the coolants H$_{\\rm 2}$ and deuterium hydride (HD) in the reionized primordial gas (e.g. Nagakura \\& Omukai 2005; Johnson \\& Bromm 2006; Yoshida et al. 2007; McGreer \\& Bryan 2008), and to the lower temperature of the cosmic microwave background (e.g. Larson 1998; Tumlinson 2007; Smith et al. 2009). Instead of the characteristic stellar mass of $\\sim$ 100 M$_{\\odot}$ expected for Pop~III.1 stars formed in the earliest minihaloes, Pop~III.2 stars with a characteristic mass of the order of 10 M$_{\\odot}$ may instead be detected in star clusters formed from reionized primordial gas.  Such clusters would have spectral properties intermediate between the two IMFs, Salpeter and top-heavy, discussed in Section 4.   An important next step to better estimate the abundance of Pop~III star clusters will be to carry out cosmological simulations which capture the earliest stages of the inhomogeneous reionization of the Universe, as well as the chemical enrichment from both the first stars formed in minihaloes and large-scale galactic winds.  We note that the cosmological metal-enrichment simulations presented in Tornatore et al. (2007) already provide some indication that Pop~III clusters may form in rare regions after reionization. In addition, cosmological radiative transfer simulations may be necessary to estimate the masses, and so the luminosities, of stellar clusters that form in the DM haloes into which the reionized primordial gas is first able collapse in the course of hierarchical structure formation. Related to this, magnetic fields pervading the reionized gas may have an important dynamical effect on the formation of Pop~III star clusters at low redshifts (see Schleicher et al. 2009a); simulations tracking the growth and evolution of magnetic fields in the IGM would thus also be of interest.  Furthermore, in the present work we have only considered the 'average' merger history of a potential Pop~III star-forming halo, ignoring the range of possible merger histories that such haloes may have. This assumption would clearly be relaxed within the context of a cosmological simulation.  Although we have accounted for the higher average density of regions which are reionized at early times with our choice of $f_{\\rm reion, mass}$ in equation (5), we have otherwise assumed that the haloes in which Pop~III star clusters form are not biased, i.e. that these haloes are otherwise located randomly in space.  In principle this could lead to an overestimate of $n_{\\rm III}$($z$), for example, if haloes with mass $M_{\\rm min}$($z$) typically form in underdense regions, while it is only overdense regions which are reionized at early times.  However, haloes with mass $M_{\\rm min}$($z$) correspond to $\\ga$ 1-$\\sigma$ fluctuations in the cosmological density field at $z$ $\\ga$ 3 (e.g. Barkana \\& Loeb 2001), typically form at 3 $\\la$ $z$ $\\la$ 4 (see e.g. Boylan-Kolchin et al. 2009), and so it does not appear likely that they form preferentially in underdense regions at $z$ $\\ga$ 3, the redshifts at which we find Pop~III star clusters are mostly likely to form in them. Therefore, we expect that accounting for the detailed clustering properties of these haloes around the descendants of the sources of reionization at $z$ $\\sim$ 20 would not change our results substantially, although the effect of bias should be considered in more depth in future studies.  Another implicit assumption throughout this work has been that stellar clusters will form in the pristine DM haloes which are massive enough for the gas to collapse into them after reionization. There is the possibility that this gas could instead collapse to form a $\\ga$ 10$^4$ M$_{\\odot}$ black hole directly, if molecular cooling is suppressed (e.g. Bromm \\& Loeb 2003; Begelman et al. 2006; Lodato \\& Natarajan 2006; Spaans \\& Silk 2006; Regan \\& Haehnelt 2009; see also Schleicher et al. 2009b).  Shang et al. (2009) find that the LW background at $z$ $\\la$ 8 (Ahn et al. 2009) may be high enough to suppress the formation of H$_{\\rm 2}$ molecules and allow for the formation of such a massive single object. However, it is not clear that the relatively small fraction $f_{\\rm coll}$ of the gas which is able to collapse is sufficiently high for this to occur. Also, despite the presence of a persistant LW background, the enhanced free electron fraction in the collapsing reionized gas will boost the formation rate of H$_{\\rm 2}$ to some extent, perhaps aiding to prevent direct collapse to a black hole.  Although we have explored the impact on the abundance of Pop~III clusters of different choices for the speed of metal enrichment by galactic winds, and for the minimum DM halo mass for star formation after reionization, we have considered only one choice for the early reionization history, based on the simulations of Shin et al. (2008). As the abundance of unenriched haloes is dependent on the fraction of the Universe reionized at redshifts $z$ $\\ga$ 20, as well as on the typical size of cosmological H~{\\sc ii} regions at these redshifts, our results are in principle quite sensitive to this choice.  Observational constraints on the early stages of reionization, for example from observations of 21-cm emission by instruments such as the {\\it Square Kilometer Array} (e.g. Carilli et al. 2004), may thus provide indirect constraints on the abundance of unenriched haloes at lower redshifts. In the upcoming decade, next generation telescopes such as the JWST and the E-ELT, however, may first provide direct constraints on the amount of Pop~III star formation taking place in the early Universe, with the exciting possibility of detecting the emission from individual Pop~III star clusters, perhaps even well after the epoch of reionization."
    },
    "0911/0911.4825_arXiv.txt": {
        "abstract": "A recollection of special moments spent with Yakov Borisovich Zeldovich and with the scientists of Soviet Union and abroad. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0911/0911.0017_arXiv.txt": {
        "abstract": "We study structure formation in the presence of primordial non-Gaussianity of the local type with parameters $\\fnl$ and $\\gnl$. We show that the distribution of dark-matter halos is naturally described by a multivariate bias scheme where the halo overdensity depends not only on the underlying matter density fluctuation $\\delta$ but also on the Gaussian part of the primordial gravitational potential $\\varphi$. This corresponds to a non-local bias scheme in terms of $\\delta$ only. We derive the coefficients of the bias expansion as a function of the halo mass by applying the peak-background split to common parameterizations for the halo mass function in the non-Gaussian scenario. We then compute the halo power spectrum and halo-matter cross spectrum in the framework of Eulerian perturbation theory up to third order. Comparing our results against N-body simulations, we find that our model accurately describes the numerical data for wavenumbers $k\\leq 0.1-0.3 \\ h$ Mpc$^{-1}$ depending on redshift and halo mass. In our multivariate approach, perturbations in the halo counts trace $\\varphi$ on large scales and this explains why the halo and matter power spectra show different asymptotic trends for $k\\to 0$. This strongly scale-dependent bias originates from terms at leading order in our expansion. This is different from what happens using the standard univariate local bias where the scale-dependent terms come from badly behaved higher-order corrections. On the other hand, our biasing scheme reduces to the usual local bias on smaller scales where $|\\varphi |$ is typically much smaller than the density perturbations. We finally discuss the halo bispectrum in the context of multivariate biasing and show that, due to its strong scale and shape dependence, it is a powerful tool for the detection of primordial non-Gaussianity from future galaxy surveys. ",
        "introduction": "Measurements of the cosmic microwave background (CMB) anisotropies have confirmed the hypothesis that the present inhomogeneities in the matter density were seeded by small fluctuations at primordial times \\citep{Komatsu:2008hk}. Such perturbations, which are expected to be created as quantum vacuum fluctuations, have been generally modeled with the simple statistical assumption of being a Gaussian random field with nearly scale invariant power spectrum \\cite{Bardeen:1985tr}.  The inflationary mechanism is often invoked to describe the early universe, but the details remain debated \\cite {Lyth:1998xn}. In the simplest single-field, slow-roll model, small curvature (adiabatic) perturbations are generated with a nearly Gaussian distribution \\cite{Maldacena:2002vr, Acquaviva:2002ud}. However in other models such as the curvaton scenario  \\cite{Lyth:2001nq, Linde:1996gt, Lyth:2002my} some additional fields would decay at later times, producing larger non-Gaussianity \\cite{Sasaki:2006kq,  Malik:2006pm, Vernizzi:2006ve}; cyclic or ekpyrotic universes without inflation could also produce large non-Gaussianities during their contracting phase \\cite{Khoury:2001wf,Creminelli:2007aq}. Furthermore, multi-field models can in general produce isocurvature modes of the perturbations \\cite{Polarski:1994rz}. See \\cite{Bartolo:2004if} for a review and \\cite{Komatsu:2009kd} for recent updates and future prospects.  The first observable predictions from inflation --- flatness and the near scale invariance of the power spectrum of the perturbations --- have been under scrutiny for some time from observations of both the large scale structure (LSS) \\cite{Tegmark:2006az} and the CMB \\cite{Komatsu:2008hk}. Fairly strict constraints on the adiabaticity of the primordial perturbations also exist \\cite{Valiviita:2009bp}, while their Gaussian distribution has only recently become testable.  Although other possibilities exist (see \\cite{Fergusson:2008ra} for a review), many models produce primordial non-Gaussianity of the \\emph{local type} where the Bardeen's potential $\\Phi$ can be expressed in terms of an auxiliary Gaussian potential $ \\varphi$ as \\be \\label{eq:basichi} \\Phi (\\mathbf{x}) = \\varphi(\\mathbf{x}) + \\sum_{j=2}^{\\infty} Q_{\\mathrm{NL}j} \\left[ \\varphi^j(\\mathbf{x}) - \\langle \\varphi^j(\\mathbf{x}) \\rangle \\right], \\ee where the series will be in practice truncated at some finite order $N$, and the odd momenta of the Gaussian potential $\\varphi$ vanish by definition. The first parameter $Q_{\\mathrm{NL}2}$, usually dubbed $\\fnl$, quantifies the leading-order departure from purely Gaussian initial conditions through the irreducible three-point function or the bispectrum of the potential. While standard inflation forecasts a slow-roll suppressed, primordial $|\\fnl| \\ll 1$, subsequent evolutionary processes are expected to increase the amount of non-Gaussianity up to $|\\fnl|\\sim 1$ \\cite{Bartolo:2003gh}. On the other hand, more complex models can produce $ |\\fnl| \\gg 1 $, although the actual predicted values vary. The second parameter $Q_{\\mathrm{NL}3}$, generally called $\\gnl$, quantifies the next higher order contribution and is related to the irreducible four-point function or the trispectrum of the potential. Since $\\varphi \\sim 10^{-5} $, this contribution can be important only if $\\gnl$ is big, $ \\gnl \\gtrsim \\fnl^2$. This is plausible in the interactive curvaton model \\cite{Enqvist:2008gk,Assadullahi:2007uw,Huang:2008zj} and in other multi-field scenarios \\cite{Byrnes:2009qy,Huang:2009vk}.  The traditional method to constrain primordial non-Gaussianity has been the three-point statistics of CMB anisotropies. The current limits from WMAP are $-9 < \\fnl < 111 $ at the $95\\%$ c.l. \\cite{Komatsu:2008hk}. Different analyses of the same data found $ -178 < \\fnl < 64 $ using Minkowski functionals \\cite{Komatsu:2008hk}, $-4 < \\fnl < 80$ using an optimized estimator \\cite{Smith:2009jr}, and $ -18  < \\fnl <  80 $ from wavelet decomposition \\cite{Curto:2009pv}, while a detection ($ 27 < \\fnl < 247 $) was claimed by \\cite{Yadav:2007yy}. Constraints on $\\gnl$ are $-5.6 \\cdot 10^5 < \\gnl < 6.4 \\cdot 10^5$ \\cite{Vielva:2009jz}. The Planck satellite is expected to reduce the uncertainty to $ \\sigma (\\fnl) \\sim 5 $ \\cite{Komatsu:2001rj}. This result will be nearly cosmic-variance limited, and further significant improvements from CMB studies will be difficult to achieve. This reason, together with the desire of having independent results, affected by different systematics, provided the motivation to study the detectability of primordial non-Gaussianity from the LSS. In this case, however, the non-linear growth of density perturbations can superimpose a new non-Gaussian signal onto the primordial one \\cite{Scoccimarro:2003wn}, which may be difficult to retrieve. Determining the mass distribution of galaxy clusters at low and high redshift provides a way to circumvent this problem \\cite{Matarrese:2000iz,LoVerde:2007ri}. However, due to the low-number statistics, these methods have been so far less successful than the CMB. Upcoming surveys such as PanSTARRS, DES, LSST, ADEPT, EUCLID, JDEM or eROSITA, WFXT and SPT are expected to substantially improve the situation.  A new technique, based on linear perturbation theory, has been recently introduced by \\cite{Dalal:2007cu} (Dal07 in the following). These authors showed that local non-Gaussianity breaks the independence of small and large scales density fluctuations. As a consequence, the clustering of dark matter halos is altered, becoming enhanced on large scales for a positive $\\fnl$. An analytical derivation of the corresponding scale-dependent bias has been also presented by \\cite{Matarrese:2008nc, Slosar:2008hx,Afshordi:2008ru,Valageas:2009vn}, together with some observational constraints on $\\fnl$ from existing data of the clustering of galaxies and their correlation with the CMB anisotropies. Using luminous red galaxies and quasars from the SDSS, \\citet{Slosar:2008hx} obtained $ -29 < \\fnl < 69 $, competitive with the CMB results. The first constraints on $\\gnl$ from the LSS give $ -3.5 \\cdot 10^5 < \\gnl < 8.2 \\cdot 10^5 $ \\cite{Desjacques:2009jb}, \\emph{assuming $\\fnl = 0$}.  N-body simulations show only approximate agreement with the model by Dal07 \\cite {Desjacques:2008vf, Pillepich:2008ka, Grossi:2009an}. In particular, the power spectrum of dark-matter halos seems to scale with the wavenumber and the $\\fnl$ parameter in a different way than predicted. This discrepancy provides the main motivation for this paper where we study the effect of non-Gaussian initial conditions on the clustering of halos in the weakly non-linear regime of perturbation growth. Applying the peak-background split technique, we show that the halo distribution on large scales is naturally described by a \\emph{bivariate local biasing scheme}, where the halo overdensity is expanded in a Taylor series of both the matter perturbations $\\delta$ and the primordial Gaussian potential $\\varphi$. Since $\\varphi$ and $\\delta$ are related by the Poisson equation, this can be equivalently seen as a non-local description in terms of $\\delta$ only. This reduces to the usual univariate \\emph{local} bias (where halo overdensities are expanded in series of $\\delta$ only) for Gaussian initial conditions and, in general, on small scales. Using standard (Eulerian) perturbation theory (SPT) up to third order to account for the non-linear growth of density fluctuations, we show that our new biasing scheme leads to the presence of several new terms in the halo power spectrum and bispectrum. We show that our results reduce to the usual Gaussian solution in the limit $\\fnl \\rightarrow 0$, and to the results by Dal07 (revised as in \\cite{Slosar:2008hx,Desjacques:2008vf}) if we only consider the leading-order terms. Finally, we compare our theory with the N-body simulations by \\cite{Pillepich:2008ka} (hereafter PPH08) and find that our model can explain the numerical results to a much greater accuracy than both linear and univariate local theories.  Our paper differs from the recent work by \\cite{Taruya:2008pg, Sefusatti:2009qh, Sefusatti:2007ih, Jeong:2009vd} which is based on the assumption that fluctuations in the halo number counts only depend on the local mass density. By ignoring the expansion in the potential $\\varphi$ and only considering the dependence on the matter density perturbations $\\delta$, one would obtain different results which do not reduce to the model by Dal07 to leading order and do not match the simulations as well; in this case higher-order terms in the halo power spectrum grow bigger than the first-order contribution on large scales, thus casting doubts on the validity of the perturbative expansion. Our approach is also different from the work by \\cite{McDonald:2008sc} because we use SPT without applying any renormalization technique. Renormalizing the bias removes any undesired dependence on the cutoff scale introduced to regularize loop corrections in SPT; however, in such a model the bias coefficients cannot be predicted and should be used as fitting parameters to match observations or simulations.  The plan of this paper is as follows. In Section \\ref {sec:tracers} we summarize the main biasing schemes which have been proposed to describe the distribution of different tracers of the LSS. In Section \\ref{sec:massfunc} we introduce some models for the halo mass function arising from non-Gaussian initial conditions and compare them against the N-body results by PPH08. We then describe in Section \\ref {sec:bias} how a multivariate bias scheme naturally emerges by applying the peak-background split technique to compute halo overdensities in the non-Gaussian case. In Section \\ref{sec:statistics} we give a short summary of the statistical properties of non-Gaussian density fields. After computing the halo power spectrum and the halo-matter cross spectrum in our multivariate biasing scheme in Section \\ref{sec:pk}, we test our theoretical models against the N-body simulations by PPH08. We derive the halo bispectrum in Section \\ref{sec:bispectrum}, and conclude in Section \\ref {sec:concl}.  \\section {Tracers of the large-scale structure and biasing } \\label{sec:tracers} The large-scale structure of the Universe can be described in terms of different tracers: mass, luminosity, galaxy counts. In this paper we will consider dark-matter halos and mass but our formalism can be straightforwardly extended to any other tracer. Let us consider the mass overdensity field $ \\delta (\\mb x) $ and the corresponding density contrast of dark-matter halos in a given mass range $\\delta_h(\\mb x)$. After smoothing both fields on a relatively large scale $R$, it is reasonable to expect that $\\delta_h$ is a local function of $\\delta$ that can be expanded in a Taylor series as follows: \\be \\delta_h(\\mb x)=b_0+b_1 \\delta(\\mb x)+b_2\\delta^2(\\mb x)/2!+b_3\\delta^3(\\mb x)/3!+\\dots\\;, \\label{ldb} \\ee where the bias coefficients $b_i$ are in principle scale and mass dependent \\cite{Fry:1992vr}. This approximation neglects stochasticity in the $\\delta_h$ vs. $\\delta$ relation and is thus dubbed local deterministic biasing. Numerical simulations from Gaussian initial conditions show that it is accurate on scales of the order of 10 Mpc and larger \\cite{2001MNRAS.320..289S}. In the following sections we will show that Eq.~(\\ref{ldb}) does not hold in the presence of primordial non-Gaussianity of the local type and we will explain how it should be modified.  Note that, in general, the bias coefficients in Eq.~(\\ref{ldb}) are not independent as the mean halo overdensity must vanish and $\\delta_h$ must assume the value $-1$ when $\\delta=-1$. In order to build a predictive theory, the values for the bias coefficients should be derived from a model. A common approach is to use the peak-background split technique \\cite {Bardeen:1985tr, Cole:1989vx, Mo:1995cs, Catelan:1997qw} where the mass perturbations are divided into fine-grained (peak) and coarse-grained (background) components. The key idea is to ascribe halo formation to the collapse of the high-frequency modes, while the large-scale distribution and motion of these condensations is determined by the low-frequency modes. Starting from a model for the conditional halo mass function (i.e. the mass function in regions where the background density assumes a specific value), the peak-background split gives an expression for the halo distribution in Lagrangian space (i.e. in the linear density field, $\\delta_1$): \\be \\delta_h^L(\\mb q)=b_0^L+b_1^L \\delta_1(\\mb q)+b_2^L\\delta_1^2(\\mb q)/2!+b_3^L\\delta_1^3(\\mb q)/3!+\\dots\\;, \\label{llagbias} \\ee where the bias coefficients $b_i^L$ are obtained from the $i$-th order derivatives of the conditional mass function with respect to the background density contrast (see  Section \\ref{sec:bias} for further details). When the background scale is much larger than the Lagrangian size of the halos, the Lagrangian bias parameters show very little dependence on the background scale and the unconditional mass function can be safely used to derive them \\cite{Sheth:1999mn}. We will follow this approach in this paper.  In the absence of large-scale velocity bias, the halo density in the evolved Eulerian space is given by \\be \\label {eq:eulag4} 1+\\delta_h(\\mb x)=[1+\\delta_h^L(\\mb q)][1+\\delta(\\mb x)] \\ee where $\\mb q$ is the Lagrangian position of the fluid elements that moved to the Eulerian location $\\mb x$ \\cite{Catelan:1997qw}. Note that the conversion between Lagrangian and Eulerian quantities is non-local, non-linear and stochastic as it depends on the displacement $\\mb x-\\mb q$ and on both the initial and the evolved fields $\\delta_1$ and $\\delta$. Therefore, the local Lagrangian biasing scheme given in Eq.~(\\ref{llagbias}) will not generally be compatible with Eq.~(\\ref{ldb}). \\citet{Catelan:2000vn} showed that these two bias models give rise to different shapes of the halo bispectrum that could then be used to distinguish between them using data from observations or simulations. A simplified approach is obtained by assuming that the long-wavelength modes of the density field evolve locally according to the spherical collapse model \\cite{Mo:1995cs, Mo:1996zb}. In this case, Eqs.~(\\ref{llagbias}) and (\\ref{ldb}) are fully compatible and the Eulerian bias parameters can be written in terms of the Lagrangian ones (see Eq.~(\\ref{eq:eubi}) in Section \\ref{sec:bias}). In particular, $b_1=1+b_1^L$. This equation is completely general as it derives from mass conservation \\cite{Mo:1995cs, Catelan:1997qw}. However, the relation between higher-order Eulerian and Lagrangian bias parameters depends on the adopted dynamics for the background density field. A perturbative calculation of the power spectrum for local Lagrangian biasing in the Gaussian scenario has been presented by \\cite{Matsubara:2008wx}. The equivalent result for the local Eulerian bias scheme has been derived by \\cite{Heavens:1998es}. In this paper we generalize this latter result to non-Gaussian initial conditions of the local type and also present a model for the bias coefficients as a function of the halo mass. The derivation of a multivariate bias scheme and the corresponding calculations of the halo power spectrum and bispectrum constitute our main results.  \\section {Halo mass function and primordial non-Gaussianity}  \\label{sec:massfunc}  The number density $\\mathcal {N}$ of halos of mass $M$ at a redshift $z$ is described by the mass function \\be n = \\frac {d \\mathcal {N}} {dM} = f \\left(\\frac {\\delta_c} {\\sigma} \\right) \\frac {\\bar \\rho} {M^2} \\left| \\frac {d \\ln \\sigma^{-1}} {d \\ln (M)} \\right|\\;, \\label{eq:massfun} \\ee where $\\delta_c\\simeq 1.686$ is the threshold for the linear density contrast which corresponds to the collapse of spherical perturbations. In Eq. (\\ref{eq:massfun}), $\\sigma^2$ denotes the variance of the linear density field, calculated as \\be \\sigma^2 (M,z) = \\frac{D^2(z)}{2 \\pi^2}\\int k^2 \\, P_0 (k) \\, W^2_f (k, M) dk, \\ee with $P_0(k)$ the linear matter power spectrum at $z=0$, $D(z)$ the linear growth factor of density fluctuations normalized to unity today, and $W_f(k,M)$ a filter function with mass resolution $M$. We use a top-hat filter in real space with radius $R_f = [3 M / (4 \\pi \\bar \\rho)]^{1/3}$, where $\\bar \\rho$ is the average density of the Universe. The analytical form of the distribution $f(\\delta_c / \\sigma)$ can be derived from a theoretical model or by fitting numerical data. We list below some possible choices for this function.     \\subsection {Gaussian mass functions}  \\paragraph* {\\bf Press-Schechter (PS)} For reference we first consider the Press-Schechter theory \\citep{Press:1973iz}, in which the mass function deriving from Gaussian initial conditions is given by \\beq f_{\\mathrm {PS}} \\left( \\frac {\\delta_c} {\\sigma} \\right) = \\sqrt{\\frac{2}{ \\pi}} \\frac{\\delta_c}{\\sigma} e^{-\\frac{{\\delta_c^2}}{2{\\sigma}^2}}. \\eeq It is well known that this model gives only a rough approximation to numerical data (see Fig.~\\ref{fig:mf}). The Press-Schechter theory can be improved by introducing extra parameters in the mass function and fitting them against numerical simulations. This approach has been followed e.g. by Jenkins et al. \\cite{Jenkins:2000bv}, Warren et al. (W) \\cite{Warren:2005ey}, Tinker et al. \\cite{Tinker:2008ff}.   \\paragraph* {\\bf Sheth-Tormen (ST)}  The Press-Schechter theory is based on the spherical collapse model. This can be improved upon by considering the collapse of ellipsoidal perturbations and fitting some new parameters against numerical simulations. \\citep{Sheth:1999mn, Sheth:2001dp}. The final result, known as the ST mass function, is \\be f_{\\mathrm {ST}} \\left( \\frac {\\delta_c} {\\sigma} \\right) = A \\sqrt {\\frac{2 \\alpha}{\\pi}} \\left[ 1 + \\left( \\alpha \\frac{\\delta_c^2}{\\sigma^2} \\right)^{-p} \\right]  \\frac{\\delta_c}{\\sigma} e^{-\\frac{\\alpha {\\delta_c^2}}{2{\\sigma}^2}}, \\ee where the extra parameters are $\\alpha = 0.707$, $p = 0.3$ and $A$ is obtained by requiring that all the mass is collapsed into halos, which gives $A = 0.322$.   \\subsection {Non-Gaussian mass functions} In the simplest case, local non-Gaussianity is described by truncating Eq.~(\\ref{eq:basichi}) after the second order term, which corresponds to: \\be \\label{eq:basic} \\Phi (\\mathbf{q}) = \\varphi(\\mathbf{q}) + \\fnl \\left[ \\varphi^2(\\mathbf{q}) - \\langle \\varphi^2(\\mathbf{q}) \\rangle \\right]. \\ee It can be shown that the halo mass function is very sensitive to the value of $\\fnl$. For positive (negative) values of $\\fnl$, its high-mass tail becomes more (less) prominent than in the Gaussian case. To first-order in the non-linearity parameter $\\fnl$, it is possible to account for this effect by considering the skewness of the density perturbations, defined as $S_3 (\\sigma) = \\langle \\delta^3 \\rangle/\\sigma^4$.   \\paragraph*{\\bf Matarrese-Verde-Jim\\'enez (MVJ)}  The Press-Schechter theory can be generalized to non-Gaussian initial conditions by using the saddle point approximation to calculate the probability for the linear density field to be above $\\delta_c$ \\citep{Matarrese:2000iz}. In this case, one obtains: \\be f_{\\mathrm {MVJ}} \\left( \\frac {\\delta_c} {\\sigma} \\right) = \\sqrt{ \\frac{2}{\\pi}} e^{-\\delta_{\\star}^2 / (2 \\sigma^2)} \\left|  \\frac {\\delta_c^3}{6 \\, \\sigma \\, \\delta_{\\star} } \\, \\frac{d S_3(\\sigma)}{d \\ln \\sigma} + \\frac {\\delta_{\\star}}{\\sigma} \\right|, \\ee where the new parameter $\\delta_\\star$ is defined as $\\delta_{\\star} \\equiv \\delta_c  \\sqrt{1 - \\delta_c S_3 (\\sigma) / 3} $. Since the ST model outperforms the PS one in the Gaussian case, it is standard practice to define a new mass function as $ f_{\\mathrm {MVJ}} \\to f_{\\mathrm {MVJ}} \\cdot f_{\\mathrm {ST}} / f_{\\mathrm {PS}} $. A further modification which has been suggested by \\cite{Grossi:2009an} to improve the agreement with numerical simulations is to correct the collapse threshold $\\delta_c$ by a factor $\\sqrt a = \\sqrt {0.8}$ in the expression $f_{\\mathrm {MVJ}}/f_{\\mathrm {PS}}$. In what follows we will adopt both corrections. A similar result was derived by \\cite{Valageas:2009vn} using a different approach.  \\paragraph*{\\bf LoVerde et al. (LV)} \\label {sec:LoV} Another way to generalize the PS model is to use the Edgeworth expansion to approximate the probability distribution function for the linear density contrast \\cite {LoVerde:2007ri}. This gives: \\be f_{\\mathrm {LV}} \\left( \\frac {\\delta_c} {\\sigma} \\right) = \\sqrt{ \\frac{2}{\\pi}} e^{-\\delta_c^2 / (2 \\sigma^2)} \\left\\{  \\left[ \\frac{\\delta_c}{\\sigma} + S_3(\\sigma) \\frac{\\sigma}{6} \\left( \\frac{\\delta_c^4}{\\sigma^4} - 2 \\frac{\\delta_c^2}{\\sigma^2} - 1  \\right)   \\right] + \\frac{1}{6} \\frac {d S_3(\\sigma)}{d \\ln \\sigma} \\sigma \\left( \\frac{\\delta_c^2}{\\sigma^2} - 1 \\right) \\right\\} \\;. \\ee As for the MVJ case, we will use an effective form of this mass function expressing the correction to the ST formula:  $ f_{\\mathrm {LV}} \\to f_{\\mathrm {LV}} \\cdot f_{\\mathrm {ST}} / f_{\\mathrm {PS}} $, with the further modification of correcting the collapse threshold $\\delta_c$ by a factor $\\sqrt a =  \\sqrt {0.8}$ in the ratio $f_{\\mathrm {LV}}/f_{\\mathrm {PS}}$.  \\paragraph*{\\bf Maggiore-Riotto (MR)} Maggiore \\& Riotto \\citep {Maggiore:2009rx} computed the halo mass function by solving the excursion set problem for non-Markovian processes with a path-integral approach, and found \\be f_{\\mathrm {MR}} \\left( \\frac {\\delta_c} {\\sigma} \\right) = (1 - \\tilde \\kappa) \\sqrt{ \\frac{2}{\\pi}} \\frac {\\sqrt a \\delta_c}{\\sigma} e^{-a \\delta^2_c / (2 \\sigma^2)} \\left[ 1 + \\frac {\\sigma^2}{6 \\sqrt {a} \\delta_c} h_{NG} (\\sigma)  \\right] + \\frac {\\tilde \\kappa}{\\sqrt {2 \\pi}} \\frac {\\sqrt {a} \\delta_c}{\\sigma} \\Gamma \\left(0, \\frac {a \\delta_c^2}{2 \\sigma^2} \\right). \\ee Here $\\Gamma (0,x)$ is the incomplete Gamma function, the additional parameters are $a \\simeq 0.8$ and $\\tilde \\kappa = a \\kappa$ where $\\kappa \\simeq 0.4562 - 0.0040\\, R_f$ with $R_f$ the smoothing scale. Primordial non-Gaussianity affects the function \\be h_{NG}(\\sigma) = \\frac{a^2 \\delta_c^4}{\\sigma^4} S_3(\\sigma) - \\frac{a \\delta_c^2}{\\sigma^2} \\left[ 2 S_3(\\sigma) + U_3(\\sigma) - \\frac {d S_3}{d \\ln \\sigma} \\right] - \\left[ S_3 (\\sigma) + U_3 (\\sigma) + V_3 (\\sigma) + \\frac {d S_3}{d \\ln \\sigma} + \\frac {d U_3}{d \\ln \\sigma}  \\right] , \\ee where $U_3$ and $V_3$ are given by: \\ba U_3 (\\sigma) &=& \\frac {3}{\\sigma^2} \\left[ \\frac {d}{d(\\sigma_1^2)} \\langle \\delta (\\sigma_1^2) \\delta^2 (\\sigma^2) \\rangle \\right]_{\\sigma_1^2 = \\sigma^2} \\\\ V_3 (\\sigma) &=& \\frac {9}{2} \\left[ \\frac{d^2}{d(\\sigma_1^2)^2} \\langle \\delta (\\sigma_1^2) \\delta^2 (\\sigma^2)  \\rangle \\right]_{\\sigma_1^2=\\sigma^2} + 12 \\left[ \\frac {d}{d(\\sigma_1^2)} \\frac {d}{d(\\sigma_2^2)}\\langle \\delta (\\sigma_1^2) \\delta (\\sigma_2^2) \\delta (\\sigma^2) \\rangle  \\right]_{\\sigma_1^2=\\sigma_2^2=\\sigma^2}\\;. \\ea The MR function does not need any ad-hoc rescaling.  \\begin {figure} \\begin{center} \\includegraphics[width=0.32\\columnwidth,angle=0]{figures/f_mass_fnl_0.eps} \\includegraphics[width=0.32\\columnwidth,angle=0]{figures/f_mass_fnl_80.eps} \\includegraphics[width=0.32\\columnwidth,angle=0]{figures/f_mass_fnl_500.eps} \\end{center} \\caption{Comparison of different models for the halo mass function originating from  Gaussian (dashed) and non-Gaussian (solid) initial conditions with the N-body data by PPH08. From left to right, the different panels refers to $\\fnl = 0, 80, 500$. Note that all non-Gaussian models (with the exception  of MR) have been rescaled by the ratio $f_{\\mathrm {ST}} / f_{\\mathrm {PS}} $, and thus coincide with the ST function in the leftmost panel. } \\label{fig:mf} \\end{figure}  \\paragraph*{\\bf Lam-Sheth (LS)}  An extension of the ST model to primordial non-Gaussianity has been recently proposed by \\cite {2009arXiv0905.1702L}. In this case the mass function is written as: \\be f_{\\mathrm {LS}} \\left( \\frac{\\delta_c}{\\sigma} \\right) = f_{\\mathrm {ST}} \\left( \\frac{\\delta_c}{\\sigma} \\right) \\left\\{ 1 + \\frac{\\sigma S_3}{6} H_3 \\left[ \\frac {b (\\sigma)}{\\sigma}  \\right]  \\right\\}, \\ee where $H_3 \\left(x \\right) = x (x^2 - 3)$, and the mass-dependent collapse barrier is $b(\\sigma) = \\sqrt{a} \\delta_c \\left[ 1 + \\beta \\left( \\sigma / \\sqrt {a} \\delta_c \\right)^{2 \\gamma} \\right]$, with $\\beta = 0.4$, $\\gamma = 0.6$, $a = 0.7$. Note that, in the case of a constant barrier, this mass function reduces to the LV one, if in the latter we neglect the term proportional to the derivative of $S_3$, which is generally small.  \\subsection{Comparison with N-body simulations} In Fig.~\\ref {fig:mf} we compare the theoretical mass functions presented above with the N-body data by PPH08. The halos in the simulations were identified using a friends-of-friends algorithm with linking length $b = 0.2 \\lambda$, where $\\lambda$ is the mean interparticle distance. Consistently with PPH08, for our calculations we assume the WMAP5 $\\Lambda$CDM model, with parameters $h=0.701, \\sigma_8=0.817, n_s = 0.96, \\Omega_m = 0.279, \\Omega_b = 0.0462, \\Omega_{\\Lambda} = 0.721$. We find that all the theoretical mass functions match the numerical output  within $\\sim 10-20\\%$ for $\\fnl<500$. Most of the discrepancy originates from the fact that the ST mass function underestimates the halo counts from the simulations (left panel in Fig.~\\ref {fig:mf}). Indeed, the models are rather accurate in predicting the ratio between the counts in a non-Gaussian model vs. a Gaussian one (see also \\cite {Desjacques:2008vf, Pillepich:2008ka, Grossi:2009an}). We also consider a fitting formula for the mass function that was computed by PPH08 from the very same data plotted in Fig.~\\ref {fig:mf}. Here we rewrite it as \\be f_{\\mathrm {PPH}} \\left( \\frac {\\delta_c} {\\sigma} \\right) = \\left[ D + B \\left( \\frac{\\delta_c}{1.686 \\, \\sigma} \\right)^A  \\right] \\exp \\left( - \\frac{C \\delta_c^2}{1.686^2 \\, \\sigma^2} \\right), \\ee where we have explicitly introduced a dependence on the threshold collapse density $\\delta_c$, and $A, B, C, D$ are fitting parameters which depend on $\\fnl$. We use the values from Table~5 in PPH08.  In the next section, we will use the mass function to compute the halo bias parameters in the non-Gaussian scenario. Given that all the models for $n$ are of the same quality, as a reference, we will only show the results obtained with the LV mass function and the PPH fit.  \\section {Halo bias} \\label {sec:bias}  \\subsection {Peak-background split} \\label{sec:pbsplit} We decompose the Gaussian auxiliary potential $\\varphi$ into the (statistically independent) contributions of long- and short-wavelength modes: \\be \\varphi (\\mathbf{q}) = \\varphi_l(\\mathbf{q}) + \\varphi_s (\\mathbf{q})\\;. \\ee Eq.~(\\ref{eq:basic}) then gives \\begin{eqnarray} \\label{eq:phis} \\Phi_l&=& \\varphi_l + \\fnl  \\varphi_l^2-\\langle \\varphi^2 \\rangle \\nonumber\\\\ \\Phi_{m}&=& 2 \\fnl \\varphi_l \\, \\varphi_s \\\\ \\Phi_s&=& \\varphi_s + \\fnl \\varphi_s^2 \\;, \\nonumber \\end{eqnarray} where the dependence on the spatial position is understood. The mixed term $\\Phi_{m}$  contributes to the short-wavelength part but derives from the coupling of $\\varphi_l$ and $\\varphi_s$. It vanishes for Gaussian initial conditions. When passing from real to Fourier space, the products of two fields become convolutions. Strictly speaking, the terms $\\varphi_l \\varphi_s$ and $\\varphi_s^2$ would also contribute to the long modes $\\Phi_l$, due to the mixing of modes caused by the convolution operation. We have checked however that these additional contributions are completely subdominant.  Using the Poisson equation, $\\nabla^2 \\Phi=A \\, \\delta$ with $A = 3 \\Omega_m H_0^2 / (2 c^2)$,  for the density fluctuations we can then write \\begin{eqnarray} \\label{eq:deltas} \\delta_l&=& \\delta_{Gl} (1+2\\fnl \\varphi_l) +2 A^{-1} \\fnl \\nabla \\varphi_l \\cdot \\nabla \\varphi_l \\nonumber\\\\ \\delta_{m}&=& 2\\fnl (\\delta_{Gs} \\,\\varphi_l+\\delta_{Gl}\\, \\varphi_s)+4 A^{-1} \\fnl \\nabla \\varphi_l \\cdot \\nabla \\varphi_s \\nonumber \\\\ \\delta_s&=& \\delta_{Gs} (1+2\\fnl \\varphi_s)+2 A^{-1} \\fnl \\nabla \\varphi_s \\cdot \\nabla \\varphi_s\\;, \\end{eqnarray} where $\\nabla^2 \\varphi=A \\, \\delta_{G}$. Notice that: \\be \\label{eq:venti} \\delta_m = 2\\fnl\\left[\\frac{\\delta_s-2A^{-1}\\fnl \\nabla \\varphi_s \\cdot \\nabla \\varphi_s}{1+2\\fnl \\varphi_s}\\,\\varphi_l\\,+ \\frac{\\delta_l-2A^{-1}\\fnl \\nabla \\varphi_l \\cdot \\nabla \\varphi_l}{1+2\\fnl \\varphi_l}\\,\\varphi_s \\right] +4 A^{-1} \\fnl \\nabla \\varphi_l \\cdot \\nabla\\varphi_s \\;. \\ee  In the spirit of the peak-background split \\cite {Bardeen:1985tr, Cole:1989vx, Mo:1995cs, Catelan:1997qw}, the short-wavelength modes of the density field collapse to form virialized condensations (dark-matter halos) while the long-wavelength ones modulate the halo counts and are responsible for large-scale motions. In the non-Gaussian case, however, the halo collapse will also be influenced by the long-wavelength modes of $\\varphi$ and $\\nabla \\varphi$ which contribute to $\\delta_m$. In a Press-Schechter approach, modulations in $\\delta_l$ will modify the threshold for halo collapse as in the Gaussian case. However, \\emph {in the presence of non-Gaussian fluctuations, the large-scale modes of the pseudo-potential will also alter the statistical properties of the small-scale modes in the density field}. This provides an additional source of biasing with respect to the Gaussian case. Suppose we want to apply the Press-Schechter algorithm to $\\delta$. The probability that the small-scale fluctuation $\\delta_s+\\delta_m$ is above the collapse threshold $\\delta_c-\\delta_l$ (probability which is obtained by averaging over $\\delta_s$) would then explicitly depend on $\\delta_l$, $\\varphi_l$ and $\\nabla \\varphi_l$. This implies that the resulting large-scale halo overdensity cannot be proportional to $\\delta_l$ as in the Gaussian case (to first order). Rather, in the general case, $\\delta_h\\simeq b_1 \\delta_l+f_1 \\varphi_l +g_1 \\nabla \\varphi_l\\cdot \\nabla \\varphi_l$ plus higher-order terms. The bias coefficients are given by the Taylor expansion of the conditional mass function $n(M|\\delta_l,\\varphi_l, \\nabla \\varphi_l\\cdot \\nabla \\varphi_l)$. Unfortunately, the models for the mass function listed in the previous section have been obtained by averaging over the entire Lagrangian volume and have no memory of the cross-talk between large and small scales. We attempted the calculation of $n(M|\\delta_l,\\varphi_l, \\nabla \\varphi_l\\cdot \\nabla \\varphi_l)$ by adopting a Press-Schechter approach and starting from the Gaussian fields $\\varphi, \\nabla \\varphi$ and $\\nabla^2 \\varphi$ but we could not obtain a closed form due to the complexity of the expressions.  An approximated model can be obtained assuming that halos form from the highest peaks of $\\delta_s$.\\footnote{We only assume that halo formation happens around some of the density peaks. This is different from the approach by \\cite{Matarrese:2008nc,Desjacques:2009jb}, in which all peaks form halos, and a one-to-one correspondence between them is assumed.} We want to implement this requirement in Eq.~(\\ref{eq:deltas}). Let us consider what the labels \"short\" and \"long\" mean in practical terms. The short part of the fields will only include a narrow shell of modes centered around the wavelength corresponding to Lagrangian size of the halos. On the other hand, the long part will be formed with all the Fourier modes with larger wavelengths. In this case, $\\varphi_s$ will be closely tracing $\\delta_{Gs}\\simeq \\delta$. This implies that the high  density peaks will nearly coincide with the maxima of $\\varphi_s$, where $\\nabla \\varphi_s=0$. In this case, \\be \\delta_s+\\delta_m=\\delta_s \\left(1+\\frac{2\\fnl \\varphi_l}{1+2\\fnl \\varphi_s}\\right)+ \\frac{\\delta_l-2A^{-1}\\fnl \\nabla \\varphi_l \\cdot \\nabla \\varphi_l}{1+2\\fnl \\varphi_l}\\,2\\fnl \\varphi_s \\;. \\nonumber \\ee The Lagrangian size of galaxy- and cluster-sized halos ranges between 1 and 10 Mpc. This implies that $\\langle \\delta_s^2 \\varphi_l^2 \\rangle^{1/2} \\gg \\langle \\delta_l^2 \\varphi_s^2 \\rangle^{1/2}$ and  $\\langle \\varphi_l^2 \\rangle \\simeq \\langle \\varphi_s^2\\rangle$, since perturbations in the pseudo-potential are nearly scale invariant. Moreover, $\\fnl \\langle \\varphi_l^2 \\rangle^{1/2}\\ll 1$ for the values of $\\fnl$ of physical interest. We thus obtain \\be \\label{eq:dng} \\delta_s+\\delta_m\\simeq \\delta_s \\left(1+2\\fnl \\varphi_l\\right)\\;, \\ee i.e. the amplitude of small-scale density fluctuations is enhanced in regions where $\\varphi_l$ is large. Therefore, the conditional mass function $n(M|\\varphi_l)$  can be computed from the unconditional one $n(M)$ by simply multiplying the r.m.s. of the density fluctuations by the factor $1+2\\fnl \\varphi_l$.\\footnote{\\citet{Slosar:2008hx} and \\citet{Afshordi:2008ru} derived a similar expression but notice that ours is written in terms of the non-Gaussian density field.} At the same time, we can use the peak-background split to derive $n(M|\\delta_l,\\varphi_l)$ by simply replacing $\\delta_c$ with $\\delta_c-\\delta_l$.  In each point in Lagrangian space, we can thus define a Lagrangian halo density field as \\be \\delta^L_h(\\mathbf{q}) = \\frac {n [M, \\delta_l(\\mathbf{q}), \\varphi_l(\\mathbf{q})]}{\\bar n} - 1, \\ee where the average can simply be taken as $\\bar n = n (M, 0, 0)$. Here it is possible to replace $n$ with $f$ since the proportionality factors cancel out, so that we can write more explicitly \\be \\label{eq:moreexplic} \\delta^L_h(\\mathbf{q}) = \\frac {f \\left( \\frac{\\delta_c - \\delta_l(\\mathbf{q})}{\\left[ 1 + 2 \\fnl \\varphi_l(\\mathbf{q}) \\right] \\sigma}   \\right) }{f \\left( \\frac{\\delta_c}{\\sigma} \\right)} - 1. \\ee We can then expand the perturbations in a Taylor series in terms of \\emph{both} variables $\\delta_l$ and $\\varphi_l$, obtaining \\be \\label{eq:deltah0} \\delta_h^L (\\mathbf{q}) =  \\sum_{j=0}^{\\infty} \\sum_{m=0}^{\\infty} \\frac {b_{jm}^L}{j!m!} \\, \\delta_l^j (\\mathbf{q}) \\, \\varphi_l^m (\\mathbf{q}) \\, . \\ee Up to third order in the perturbations, this gives: \\ba \\label{eq:bfexp0} \\delta^L_{h} (\\mathbf q) &=& b^L_0 + b^L_{10} \\, \\delta + b^L_{01}\\, \\varphi + \\nonumber \\\\ &+& \\frac{1}{2!} \\left( b^L_{20} \\, \\delta^2 + 2\\,b^L_{11}\\, \\delta  \\varphi + b^L_{02} \\, \\varphi^2 \\right)+ \\nonumber \\\\ &+& \\frac {1} {3!} \\left( b^L_{30}\\, \\delta^3 + 3 \\, b^L_{21} \\, \\delta^2 \\varphi + 3 \\, b^L_{12}\\, \\delta \\varphi^2 + b^L_{03}\\, \\varphi^{3} \\right), \\ea where all the density perturbations on the r.h.s. are Lagrangian and non-Gaussian.  Eq.~(\\ref{eq:moreexplic}) implies that not all the coefficients $b_{jm}^L$ are independent. In particular, all the $b_{jm}^L$ with $m \\ne 0$ can be written in terms of the $b_{j0}^L$. Up to third order we have: \\ba \\label{eq:bsaintindep} b_{01}^L &=& 2 \\, \\fnl \\, \\delta_c \\, b_{10}^L \\nonumber \\\\ b_{11}^L &=& 2 \\, \\fnl \\, ( - b_{10}^L + \\delta_c \\, b_{20}^L) \\nonumber \\\\ b_{02}^L &=& 4 \\, \\fnl^2 \\, (- 2 \\delta_c \\, b_{10}^L + \\delta_c^2 \\, b_{20}^L )  \\nonumber \\\\ b_{21}^L &=& 2 \\, \\fnl \\, (- 2 b_{20}^L + \\delta_c \\, b_{30}^L ) \\nonumber \\\\ b_{12}^L &=& 4 \\, \\fnl^2 \\, ( 2 b_{10}^L - 4 \\delta_c \\, b_{20}^L + \\delta_c^2 \\, b_{30}^L ) \\nonumber \\\\ b_{03}^L &=& 8 \\, \\fnl^3 \\, ( 6 \\delta_c \\, b_{10}^L - 6 \\delta_c^2 \\, b_{20}^L   + \\delta_c^3 \\, b_{30}^L ). \\ea It is important to remember that the functional form of the halo mass function accounts for the effect of non-Gaussianity on the short wavelength modes $\\delta_s, \\delta_m$. The bias coefficients $b_{j0}^L$ will then depend implicitly on $\\fnl$ through the shape of the mass function.  \\subsection {Extension to higher-order non-Gaussianity} \\label{sec:gnl}  If the model for non-Gaussianity is extended to higher order, then Eq.~(\\ref{eq:basic}) is replaced by Eq.~(\\ref{eq:basichi}). If we consider cubic corrections, we have the following additional contributions  to Eqs.~(\\ref{eq:phis}): \\begin{eqnarray} \\Delta \\Phi_l&=& \\gnl \\, \\varphi_l^3 \\nonumber\\\\ \\Delta\\Phi_{m}&=& 3 \\gnl  \\left( \\varphi_l^2 \\, \\varphi_s + \\varphi_l  \\, \\varphi_s^2 \\right) \\\\ \\Delta \\Phi_s&=& \\gnl \\, \\varphi_s^3 , \\nonumber \\end{eqnarray} which correspond to the following additions in Eqs.~(\\ref{eq:deltas}): \\begin{eqnarray} \\Delta \\delta_l&=& 3 \\gnl \\, \\varphi_l^2 \\, \\delta_{Gl} + 6 A^{-1} \\gnl \\, \\varphi_l \\, \\nabla \\varphi_l \\cdot \\nabla \\varphi_l \\nonumber\\\\ \\Delta \\delta_{m}&=& 3 \\gnl \\, \\delta_{Gl} \\,  \\varphi_s (2  \\varphi_l +  \\varphi_s) + 6 \\gnl A^{-1} \\, \\varphi_s \\, \\nabla \\varphi_l \\cdot \\nabla \\varphi_l  \\\\ \\Delta \\delta_s&=& 0 \\nonumber, \\label{eq:deltashi} \\end{eqnarray} where we have already imposed the peak condition. In analogy with the previous section we thus identify the leading term as: \\be \\label{eq:gnl} \\delta_s+\\delta_m \\simeq \\delta_s \\left( 1 + 2 \\fnl \\varphi_l + 3 \\gnl \\varphi_l^2  \\right). \\ee It follows that the r.m.s. of the small-scale density fluctuations will now be altered by a factor $\\left( 1 + 2 \\fnl \\varphi_l + 3 \\gnl \\varphi_l^2  \\right)$ with respect to the Gaussian case. Therefore, considering $\\gnl \\ne 0$ introduces additional terms in the bias coefficients in Eq.~(\\ref{eq:bsaintindep}), given by \\ba \\label{eq:dbgnl} \\Delta b_{02}^L &=& 6 \\, \\gnl \\, \\delta_c \\, b_{10}^L   \\nonumber \\\\ \\Delta b_{12}^L &=& 6 \\, \\gnl \\, ( - b_{10}^L + \\delta_c \\, b_{20}^L), \\ea while all the other bias coefficients remain unchanged (apart from the modifications due to the implicit dependence of the mass function on $\\gnl$, which we do not calculate here).  Eq.~(\\ref{eq:gnl}) can be finally generalized to an arbitrary order $N$ as \\be \\delta_s+\\delta_m \\simeq \\delta_s \\sum_{j=2}^N j Q_{\\mathrm{NL}j} \\varphi_l^{j-1}. \\ee This equation shows that the leading contribution of each successive order will depend on a higher power of the potential, and its effects will therefore be smaller in amplitude. Note that $h_{\\mathrm{NL}} \\equiv Q_{\\mathrm{NL}4}$ is the highest-order term that can explictly modify the $b_{ij}$ parameters (up to third order in the bias expansion), although all the $ Q_{\\mathrm{NL}j}$ will introduce implicit dependences by modifying the halo mass function.    \\subsection {Bias from a mass function} We want now to explicitly calculate the halo bias corresponding to a given mass function. As discussed above, in the non-Gaussian case the mass function will also be explicitly dependent on the potential $\\varphi$, and the halo overdensities can now be derived from  Eq.~(\\ref{eq:deltah0}), as a bivariate series expansion in terms of $\\delta_l$ and $\\varphi_l$. Since the effect of the short-wavelength modes is taken into account by the functional form of the mass function, we will henceforth drop the $l$ indices and use the symbols $\\delta$ and $\\varphi$ to denote the long-wavelength parts of the perturbations.  \\paragraph*{\\bf PS model} For reference, we derive the bias coefficients corresponding to the simple PS mass function (see also \\cite{Mo:1995cs}) \\ba \\label {eq:b1} b^L_{10} &=& -\\frac{1}{\\delta_c} + \\frac {\\delta_c}{\\sigma^2},  \\nonumber \\\\ b^L_{20} &=& \\frac {\\delta_c^2}{\\sigma^4} - \\frac{3}{\\sigma^2} \\nonumber \\\\ b^L_{30} &=& \\frac{\\delta_c^3}{\\sigma^6} - \\frac{6 \\delta_c}{\\sigma^4} + \\frac{3}{\\delta_c \\sigma^2} \\ea and show their mass dependence in the right panel of Fig.~\\ref{fig:b1b2b3}. The remaining coefficients can be obtained using Eq.~(\\ref{eq:bsaintindep}): \\ba b^L_{01} &=& \\left( \\frac {2 \\delta_c^2 }{\\sigma^2} - 2 \\right) \\fnl  \\nonumber \\\\ b^L_{11} &=&   \\left( \\frac{2}{\\delta_c} +  \\frac {2 \\delta_c^3}{\\sigma^4} -  \\frac {8 \\delta_c}{\\sigma^2} \\right) \\fnl  \\nonumber \\\\ b^L_{02} &=&  2 \\left( \\frac{ 2 \\delta_c^4}{\\sigma^4} - 10 \\frac {\\delta_c^2} {\\sigma^2} + 4  \\right) \\fnl^2  \\nonumber \\\\ b^L_{21} &=& 2 \\left(\\frac {\\delta_c^4}{ \\sigma^6} - \\frac {8 \\delta_c^2}  {\\sigma^4} + \\frac {9} {\\sigma^2} \\right) \\fnl  \\nonumber \\\\ b^L_{12} &=& 2 \\left( \\frac {2 \\delta_c^5}{ \\sigma^6} - \\frac {20 \\delta_c^3 }{ \\sigma^4} + \\frac {34 \\delta_c }{ \\sigma^2} - \\frac {4}{ \\delta_c}  \\right) \\fnl^2    \\nonumber \\\\ b^L_{03} &=& 6 \\left( \\frac {4 \\delta_c^6}{3  \\sigma^6} - \\frac{16 \\delta_c^4 }{  \\sigma^4} + \\frac{36 \\delta_c^2 }{  \\sigma^2} - 8  \\right)  \\fnl^3. \\ea Notice that the ``usual'' bias coefficients ($b^L_{10}, b^L_{20}, b^L_{30}$) are in this case independent from $\\fnl$. This does not hold in general, since an implicit dependence on $\\fnl$ will be introduced by any non-Gaussian mass function.  \\paragraph*{\\bf General case} To derive the bias coefficients up to the third order, we can repeat the same procedure for any other mass function. We have calculated these coefficients for all the mass functions listed in Section \\ref{sec:massfunc} finding an overall agreement in the trends with mass and with the non-linearity parameter $\\fnl$. The analytic form of the bias parameters is much more complex than for the PS mass function and we will not write it explicitly. As an example, in Fig.~\\ref{fig:b1b2b3} we show how the bias coefficients $b_{j0}^L$ depend on $\\fnl$ and halo mass for the LV mass function and the PPH fit.  \\subsection {Lagrangian and Eulerian bias} \\begin {figure} \\begin{center} \\includegraphics[width=0.45\\columnwidth,angle=0]{figures/fig_b1b2b3.eps} \\includegraphics[width=0.45\\columnwidth,angle=0]{figures/fig_b1b2b3_M.eps} \\end{center} \\caption{Lagrangian bias factors at $z=0$ as a function of $\\fnl$ for a halo mass of $2 \\cdot 10^{14} M_{\\odot}/h$ (left), and as a function of the halo mass, for $\\fnl = 500 $ (right). The results obtained from the PPH and LV mass functions are shown in both cases. For reference, in the right panel we also show the results obtained from the PS mass function, which does not depend on $\\fnl$.} \\label{fig:b1b2b3} \\end{figure}  A model for the formation of the LSS provides a relationship between the density perturbations in Lagrangian and Eulerian space. This relation is generally non-local \\cite{Catelan:1997qw} but a local approximation (where $\\mb x \\equiv \\mb q$) suffices to approximately describe the evolution of large-scale perturbations. In this case one writes \\cite{Mo:1995cs,Mo:1996zb} \\be \\label {eq:series} \\delta^L = \\sum_{j=1}^{\\infty} a_j \\left( \\delta^E \\right)^j, \\ee where the $a_j$'s parameterize the evolution of mass-density fluctuations. For the simple case of spherical collapse, we have \\cite{Bernardeau:1994zd} \\be a_1 = 1 \\, ; \\:\\: a_2 = -17 / 21 \\, ; \\:\\: a_3 = 341 / 567. \\ee  Starting from the Lagrangian halo density perturbations $\\delta_{h}^L$  given in Eq.~(\\ref{eq:bfexp0}) we want to use Eqs.~(\\ref{eq:eulag4}) and (\\ref{eq:series}) to write the Eulerian halo overdensity in terms of Eulerian density perturbations. This gives: \\ba \\label{eq:bfexpEuler} \\delta_{h} (\\mathbf x) &=& b_0 + b_{10} \\, \\delta + b_{01}\\, \\varphi+ \\nonumber \\\\ &+& \\frac{1}{2!} \\left(b_{20} \\, \\delta^2 + 2 \\, b_{11} \\, \\delta \\varphi + b_{02} \\, \\varphi^2  \\right) + \\nonumber \\\\ &+& \\frac{1}{3!} \\left( b_{30} \\, \\delta^3 + 3 \\, b_{21} \\, \\delta^2 \\varphi + 3 \\, b_{12} \\, \\delta \\varphi^2 + b_{03} \\, \\varphi^{3} \\right), \\ea where all the density perturbations on the r.h.s. are Eulerian and non-Gaussian (from now on we drop the superscript $E$ and all densities will be Eulerian unless explicitely stated otherwise) and the bias coefficients are given by the following expressions: \\ba \\label{eq:eubi} b_{10} &=&  1 + a_1 \\, b^L_{10} \\nonumber   \\\\ b_{20} &=&  2 (a_1 + a_2) \\, b^L_{10} + a_1^2 \\, b^L_{20} \\nonumber \\\\ b_{30} &=&  6 (a_2 + a_3) \\, b^L_{10} +  3 \\left(  a_1^2 + 2 a_1 a_2 \\right) \\, b^L_{20} + a_1^3 \\, b^L_{30} \\\\ b_{01} &=&  b^L_{01}  \\nonumber \\\\ b_{11} &=&  b^L_{01} + a_1 \\,  b^L_{11} / 2  \\nonumber \\\\ b_{02} &=&  b^L_{02} \\nonumber \\\\ b_{21} &=&  (a_1 + a_2) \\, b^L_{11}  + a_1^2 \\, b^L_{21} / 3 \\nonumber \\\\ b_{12} &=&  b^L_{02} + a_1 \\, b^L_{12} / 3 \\nonumber \\\\ b_{03} &=&  b^L_{03}.  \\nonumber \\ea Note that Eq.~(\\ref{eq:bfexpEuler}) differs from Eq.~(\\ref{ldb}) due to the presence of extra terms which are proportional to different powers of the Gaussian auxiliary potential $\\varphi$. Since $\\varphi$ is nearly scale invariant while $\\delta$ has a larger variance on smaller scales, the additional terms will only affect the statistics of the halo distribution on the largest scales. Also, $\\varphi$ does not evolve with time while $\\delta$ does (according to $\\delta(z) = D(z) \\delta(z=0)$ at linear order) thus implying that, for a given set of bias coefficients, the new terms will become less and less important over time. A third peculiarity of  Eq.~(\\ref{eq:bfexpEuler}) is that the halo overdensity can differ from zero also in regions with mean mass density.  \\subsection {Perturbative expansion} In order to account for the non-linear evolution of mass-density fluctuations in Eq.~(\\ref{eq:bfexpEuler}) we use standard Eulerian perturbation theory (see \\cite{Bernardeau:1994zd} for a review). We therefore expand the density fields to third order as $\\delta = \\delta_{1} + \\delta_{2} + \\delta_{3} + {\\cal O}(\\delta_4) $ where $\\delta_n$ is ${\\cal O}(\\delta_1^n)$. Each term can be expressed as \\be \\tilde \\delta_n(\\mb k)=\\int \\frac{d^3 \\mb q_1}{(2 \\pi)^3} \\dots \\frac{d^3 \\mb q_n}{(2 \\pi)^3} \\,\\delta_D\\left(\\mb k-\\sum_{i=1}^n \\mb q_i\\right)\\, J_n(\\mb q_1,\\dots,\\mb q_n)\\, \\tilde \\delta_1(\\mb q_1) \\dots \\tilde \\delta_1(\\mb q_n)\\;, \\label{eq:perturbn} \\ee where $\\delta_D$ is the Dirac delta distribution, the tilde denotes Fourier transformation, and the $J_n$ are specific kernel functions. On the other hand, since the Gaussian potential $\\varphi$ is the primordial one, there is no need to expand it, and it fully coincides with its first-order part $\\varphi \\equiv \\varphi_{1} $.  We can now explicitly rewrite Eq.~(\\ref{eq:bfexpEuler}) up to the third perturbative order as \\ba \\label{eq:deltaE} \\delta_{h} (\\mathbf x) &=& b_0 + q_{10} \\delta_{1}  + q_{11} \\varphi_{1}  + \\nonumber  \\\\ &+& q_{20} \\delta_{2} + q_{21} \\delta_{1}^2   + q_{22} \\delta_{1} \\varphi_{1} + q_{23} \\varphi_{1}^2  + \\nonumber  \\\\ &+& q_{30} \\delta_{3} + q_{31} \\delta_{1} \\delta_{2} + q_{32} \\delta_{2} \\varphi_{1} + q_{33} \\delta_{1}^2 \\varphi_{1} + q_{34} \\delta_{1} \\varphi_{1}^2 + q_{35} \\delta_{1}^3  + q_{36} \\varphi_{1}^3 , \\ea where, to simplify the notation and facilitate bookkeeping of the terms which will appear in the perturbative expression for the power spectrum, we have replaced the biases $b_{ij}$ with new coefficients $q_{ij}$. The explicit form of these  is given in Table~\\ref{tab:coef}.  Eq.~(\\ref{eq:deltaE}) fully describes the Eulerian halo bias at the third perturbative order in the non-Gaussian case. The leading order term,  $\\delta_h\\simeq [1+a_1b_{10}^L (\\fnl) ] \\, \\delta+ 2 \\fnl b_{10}^L (\\fnl) \\, \\varphi$, was already recognized by Dal07, \\cite{Afshordi:2008ru} and \\cite{Slosar:2008hx}.  \\begin{table} \\begin{tabular}{|c| c|  c|  c| c|  c|  c|} \\hline $q_{10} = b_{10}$  &  $q_{11} = b_{01}$  &   & &   &    &   \\\\ $q_{20} = b_{10}$  &  $q_{21} = b_{20} / 2 $  & $q_{22} = b_{11} / 2 $   & $q_{23} = b_{02} / 2 $  &   &    &   \\\\ $q_{30} = b_{10}$  &  $q_{31} = b_{20}$  &  $q_{32} = b_{11} / 2$ &  $q_{33} = b_{21} / 6$ &  $q_{34} = b_{12} / 6$  & $q_{35} = b_{30} / 6$   & $q_{36} = b_{03} / 6$  \\\\ \\hline \\end{tabular} \\caption{Mapping of the bias coefficients in the full perturbative expansion.} \\label{tab:coef} \\end{table}       \\section {Clustering statistics and non-Gaussianity} \\label{sec:statistics}  Statistical analysis of random fields, such as the  Bardeen potential $\\Phi (\\mb x)$, can be performed by studying the irreducible $N$-point correlation functions $\\langle \\Phi (\\mb x_1) \\Phi (\\mb x_2) ... \\Phi (\\mb x_N) \\rangle$, or alternatively the $N$-spectra \\be (2 \\pi)^3 {\\cal S}_N (\\mathbf k_1,\\mathbf k_2, ..., \\mathbf k_N) \\,  \\delta_D (\\mathbf k_1 + \\mathbf k_2 + ... + \\mathbf k_N) = \\langle \\tilde \\Phi (\\mathbf k_1) \\tilde \\Phi (\\mathbf k_2) ... \\tilde \\Phi (\\mathbf k_N)  \\rangle\\;. \\ee For Gaussian fields all odd-order spectra vanish. On the other hand, thanks to Wick's theorem, the reducible even-order correlators can be decomposed as products of the power spectrum, \\be (2 \\pi)^3 P_{\\Phi}(k) \\, \\delta_D (\\mathbf k_1 + \\mathbf k_2) = \\langle \\tilde  \\Phi (\\mathbf k_1) \\tilde  \\Phi (\\mathbf k_2) \\rangle, \\ee which, in this case, encodes all the information. If some non-Gaussianity is instead introduced, then higher-order statistics become important, such as the bispectrum \\be (2 \\pi)^3 B_{\\Phi}(\\mathbf k_1, \\mathbf k_2, \\mathbf k_3) \\, \\delta_D (\\mathbf k_1 + \\mathbf k_2 + \\mathbf k_3) = \\langle \\tilde  \\Phi (\\mathbf k_1) \\tilde  \\Phi (\\mathbf k_2) \\tilde  \\Phi (\\mathbf k_3) \\rangle, \\ee and the irreducible trispectrum \\be (2 \\pi)^3 T_{\\Phi}(\\mathbf k_1, \\mathbf k_2,\\mathbf k_3,\\mathbf k_4) \\, \\delta_D (\\mathbf k_1 + \\mathbf k_2 + \\mathbf k_3 + \\mathbf k_4) = \\langle \\tilde  \\Phi (\\mathbf k_1) \\tilde \\Phi (\\mathbf k_2) \\tilde \\Phi (\\mathbf k_3) \\tilde \\Phi (\\mathbf k_4) \\rangle. \\ee Using our simplest model of non-Gaussianity given in Eq.~(\\ref{eq:basic}) to define the non-Gaussian potential $\\Phi$ in terms of the Gaussian one $\\varphi$, we obtain: \\be \\label{eq:simplifme} P_{\\Phi}(k) = P_{\\varphi}(k) + 2 \\fnl^2 \\int \\frac{d^3 \\mb{q}}{(2 \\pi)^3} P_{\\varphi}(q) P_{\\varphi}(|\\mb{k} - \\mb{q}|)\\simeq P_{\\varphi}(k) \\ee \\be B_\\Phi(\\mathbf k_1,\\mathbf k_2,\\mathbf k_3) \\simeq 2 \\fnl \\left[P_\\varphi(k_1) P_\\varphi(k_2) + \\mathrm {(2\\ cyclic)}\\right] \\label{eq:bispPhi} \\ee \\be T_{\\Phi}(\\mathbf k_1, \\mathbf k_2, \\mathbf k_3, \\mathbf k_4) \\simeq 4 \\fnl^2 \\left\\{P_\\varphi(k_1) P_\\varphi(k_2) \\left[P_\\varphi(|\\mb k_1+\\mb k_3| )+ P_\\varphi(|\\mb k_1+ \\mb k_4|)\\right]+\\mathrm{ (5\\ cyclic)}\\right\\} \\label{eq:trispPhi} \\ee where we dropped a sub-leading term proportional to $\\fnl^N$ for each of the $N$-spectra (this is why we used the symbol of approximate equality).\\footnote{In multi-field inflationary models the non-Gaussian contribution to the trispectrum may scale independently from the bispectrum. For this reason, the factor $\\fnl^2$ in Eq.~(\\ref{eq:trispPhi}) is sometimes re-labeled $\\tau_{\\mathrm{NL}}$, and treated as an independent parameter. Observational constraints on the trispectrum of the Bardeen's potential should then discriminate between such models and the simplest inflationary scenarios.} We have checked that the discarded terms are indeed negligibly small. For instance, the sub-leading contribution to $P_\\Phi$ contributes less than $1\\%$ of the total for $|\\fnl| \\sim 1000$ in the $k$-range of interest.  Considering also the third-order term in Eq.~(\\ref{eq:basichi}) introduces additional contributions to the power spectrum of the Bardeen's potential. The leading-order term can be written as: \\be \\Delta P_{\\Phi}(k) =  6 \\gnl P_{\\varphi}(k) \\int \\frac{d^3 \\mb{q}}{(2 \\pi)^3} P_{\\varphi}(q)\\;. \\ee For  $P_{\\varphi}(k)\\propto k^{n_s-4}$, the integral above presents an ultraviolet $(k \\to \\infty)$ divergence if $n_s\\geq 1$ and an infrared $(k \\to 0)$ divergence if $n_s\\leq 1$. In general, this is not a problem as the physical process creating the fluctuations will automatically introduce cutoffs in $P_\\varphi$ at small and large wavelengths. For example, cosmic inflation will generate perturbations with characteristic sizes comprised between the reheating scale and the present-day horizon \\cite{Matarrese:2000iz}. However, if $\\Delta P_\\Phi$ is non-negligible with respect to the leading-order contribution $P_\\varphi$ (i.e. if $6|\\gnl| \\langle \\varphi^2 \\rangle$ is not much less than unity), the results of the perturbative expansion are of limited use unless artificial cutoffs are introduced and the parameters of the theory are renormalized. The condition above reduces to $|\\gnl| \\ll 10^7$ if the currently favored values for the amplitude and the spectral index of primordial perturbations are plugged in. Present-day observational limits on $\\gnl$ \\cite{Desjacques:2009jb,Vielva:2009jz} therefore suggests that $\\Delta P_\\Phi$ should contribute at the percent level or less to the power spectrum of the potential. Note that in numerical simulations \\cite{Desjacques:2009jb}, non-physical infrared and ultraviolet cutoffs are introduced by the use of a finite volume with periodic boundary conditions. Considering a non-vanishing $\\gnl$ also adds another leading-order correction to the trispectrum of the Bardeen potential: \\be \\label{eq:deltaT} \\Delta T_{\\Phi}(\\mathbf k_1, \\mathbf k_2, \\mathbf k_3, \\mathbf k_4)\\simeq 6 \\gnl P_\\varphi(k_1) P_\\varphi(k_2) P_\\varphi(k_3) + \\mathrm{(3\\ cyclic)}\\;, \\ee while it does not modify the bispectrum of $\\Phi$ at leading order.  Linear perturbations in the density at redshift $z$ are related to those in the primordial potential (formally at $z\\to \\infty$) by the Poisson equation \\be \\label{eq:poissalpha} \\tilde \\delta_1(k) = \\alpha(k) \\tilde  \\Phi(k) \\ee with \\be \\alpha(k) = \\frac {2 c^2 k^2 T(k) D(z)} {3 \\Omega_m H_0^2} \\frac{g(0)}{g(\\infty)}, \\ee where the matter growth factor $D(z)$ and the transfer function $T(k)$ have been introduced to account for the linear evolution of $\\delta_1$. The function $g(z)\\equiv (1+z)D(z)$ is the linear growth factor for the potential, and $g(\\infty)/g(0)\\simeq 1.3$ in the currently favored cosmology. Therefore, we can relate the power spectrum of linear density fluctuations to the power spectrum of the primordial potential by writing \\be P_0 (k) \\equiv P_{\\delta_1}(k) = \\alpha^2(k) \\, P_{\\Phi} (k)  \\simeq  \\alpha^2(k) \\, P_{\\varphi} (k), \\label{eq:convspectra} \\ee where the last approximation follows from Eq.~(\\ref{eq:simplifme}). Similar equations can be written for the three- and four-point correlators of the linear density perturbations, which we will label $B_0$ and $ T_0$ respectively, by combining Eq.~(\\ref{eq:poissalpha}) with Eq. (\\ref{eq:bispPhi}) and (\\ref{eq:trispPhi}).  The fact that the Bardeen's potential decays with time proportionally to $g(z)$ implies that the actual values of the coefficients $ Q_{\\mathrm{NL}j}$ depend on the cosmic epoch at which  Eq.~(\\ref{eq:basichi}) is applied (see Section 2.2 in \\cite{Pillepich:2008ka}). Here we apply it at early times (which is sometimes called the ``CMB convention'') while other authors use the fields linearly extrapolated at $z=0$ (the ``LSS convention''). In general, \\be Q_{\\mathrm{NL}j}^{\\mathrm{LSS}}=Q_{\\mathrm{NL}j}^{\\mathrm{CMB}} \\left[\\frac{g(\\infty)}{g(0)} \\right]^j \\ee so that $\\fnl^{\\mathrm{LSS}}\\simeq 1.3\\, \\fnl^{\\mathrm{CMB}}$ and $\\gnl^{\\mathrm{LSS}}\\simeq 1.7\\, \\gnl^{\\mathrm{CMB}}$. This conversion factors should be taken into account when comparing papers using different conventions.  \\section {Power spectra} \\label{sec:pk} We are now ready to calculate the two-point statistics of the LSS arising from non-Gaussian initial conditions. We compute the halo-halo power spectrum and the halo-matter cross spectrum as follows. First, we take the Fourier transform of Eq.~(\\ref{eq:deltaE}) and build the corresponding two-point correlators $\\langle \\tilde\\delta_h(\\mathbf k_1) \\tilde \\delta_h(\\mathbf k_2)\\rangle$ and $\\langle \\tilde\\delta_h(\\mathbf k_1) \\tilde \\delta(\\mathbf k_2)\\rangle$. These are composed of many pieces and we only consider terms up to the fourth perturbative order. For instance the leading contribution to halo-halo spectrum (second order in terms of the perturbations) is composed of 3 terms, the third order correction is made of 8 pieces (of which one identically vanishes because $\\varphi$ is a Gaussian field), and the fourth-order one contains 30 terms (14 of which are obtained multiplying a linear perturbation by a third-order one -- indicated by the subscript $_{(13)}$ hereafter -- and 16 are originated by the product of two second-order terms -- subscript $_{(22)}$ hereafter). To proceed we then: (a) use Eq.~(\\ref{eq:perturbn}) and write the density perturbations of order $n>1$ as convolutions of $n$ linear perturbations and a kernel $J_n$; (b) express the linear density perturbations in terms of the potential $\\Phi$ using the Poisson Eq.~(\\ref{eq:poissalpha}); (c) take the ensemble averages by using the expressions for the power spectrum, bispectrum and trispectrum given in Eqs.~(\\ref{eq:bispPhi}), (\\ref{eq:trispPhi}), (\\ref{eq:deltaT}), and (\\ref{eq:convspectra}). While $\\langle \\tilde  \\varphi (\\mathbf k_1) \\tilde  \\varphi (\\mathbf k_2) \\tilde  \\varphi (\\mathbf k_3) \\rangle=0$, attention must be payed to the mixed terms in $\\Phi$ and $\\varphi$ as: \\be \\langle \\tilde  \\Phi (\\mathbf k_1) \\tilde  \\Phi (\\mathbf k_2) \\tilde  \\Phi (\\mathbf k_3) \\rangle\\simeq \\langle \\tilde  \\varphi (\\mathbf k_1) \\tilde  \\Phi (\\mathbf k_2) \\tilde \\Phi (\\mathbf k_3) \\rangle+ \\mathrm{(2 \\ cyc.)} \\simeq \\langle \\tilde  \\varphi (\\mathbf k_1) \\tilde  \\varphi (\\mathbf k_2) \\tilde  \\Phi (\\mathbf k_3) \\rangle +\\mathrm{(2 \\ cyc.)} \\propto \\fnl \\ee where equalities only hold at leading order in $\\varphi^n$ (i.e. $\\varphi^4$) as the central correlator also contains a sub-leading term proportional to $\\fnl \\gnl$ (which scales as $\\varphi^6$) and the leftmost one some terms proportional to $\\fnl^3$, $\\fnl \\gnl$ (both scaling as $\\varphi^6$), and $\\fnl \\gnl^2$ ($\\propto \\varphi^8$). Similarly, to leading order in $\\varphi^n$ (i.e. $\\varphi^6$), \\ba \\langle \\tilde  \\Phi (\\mathbf k_1) \\tilde \\Phi (\\mathbf k_2) \\tilde \\Phi (\\mathbf k_3) \\tilde \\Phi (\\mathbf k_4) \\rangle&\\simeq& \\fnl^2 T_A+ \\gnl T_B\\\\ \\langle \\tilde  \\varphi (\\mathbf k_1) \\tilde \\Phi (\\mathbf k_2) \\tilde \\Phi (\\mathbf k_3) \\tilde \\Phi (\\mathbf k_4) \\rangle + \\mathrm{(3 \\ cyc.)} &\\simeq& 2 \\,\\fnl^2 T_A+ 3 \\,\\gnl T_B\\\\ \\langle \\tilde  \\varphi (\\mathbf k_1) \\tilde \\varphi (\\mathbf k_2) \\tilde \\Phi (\\mathbf k_3) \\tilde \\Phi (\\mathbf k_4) \\rangle+ \\mathrm{(5 \\ cyc.)} &\\simeq& \\fnl^2 T_A+ 3\\, \\gnl T_B\\\\ \\langle \\tilde  \\varphi (\\mathbf k_1) \\tilde \\varphi (\\mathbf k_2) \\tilde \\varphi (\\mathbf k_3) \\tilde \\Phi (\\mathbf k_4) \\rangle + \\mathrm{(3 \\ cyc.)} &\\simeq& \\gnl T_B \\ea where $T_A(\\mb k_1,\\mb k_2,\\mb k_3,\\mb k_4)$ and $ T_B(\\mb k_1,\\mb k_2,\\mb k_3, \\mb k_4)$ are defined in Eqs. (\\ref{eq:trispPhi}) and (\\ref{eq:deltaT}).   \\subsection {Results and comparison with N-body simulations}  \\paragraph* {\\bf Matter} If we set $b_{10}=1$ and all the other bias coefficients to zero, we obtain an expression for the power spectrum of mass-density perturbations which coincides with the result by \\cite {Taruya:2008pg}: $P^{mm}(k, z) = D^2(z) P_{11}(k) + D^3(z) P^{mm}_{12}(k) + D^4(z) \\left[ P^{mm}_{22}(k) + P^{mm}_{13}(k) \\right] $, where \\ba P^{mm}_{11} (k) &=& P_0 (k) \\nn \\\\ P^{mm}_{12} (k) &=& 2 \\int \\frac{d^3 \\mathbf q}{(2 \\pi)^3} J^{(s)}_2 (\\mathbf q, \\mathbf k - \\mathbf q) B_0 (- \\mathbf k, \\mathbf q, \\mathbf k - \\mathbf q) \\nn \\\\ P^{mm}_{22} (k) &=& 2 \\int \\frac{d^3 \\mathbf q}{(2 \\pi)^3} \\left[ J^{(s)}_2 (\\mathbf q, \\mathbf k - \\mathbf q) \\right]^2  P_0(q) P_0(|\\mb k - \\mb q|) + \\nn \\\\ &~& + \\int \\frac{d^3 \\mathbf p \\, d^3 \\mathbf q}{(2 \\pi)^6} \\, J^{(s)}_2  (\\mathbf p, \\mathbf k - \\mathbf p) \\, J^{(s)}_2  (\\mathbf q, -\\mathbf k - \\mathbf q) \\,  T_0(\\mathbf p, \\mathbf k - \\mathbf p, \\mathbf q, -\\mathbf k - \\mathbf q) \\nn \\ea \\ba P^{mm}_{13} (k) &=& 6 \\int \\frac{d^3 \\mathbf q}{(2 \\pi)^3} J^{(s)}_3 (\\mb k, \\mb q, -\\mb q) P_0(q) P_0(k) +  \\nn  \\\\ &~& + 2 \\int \\frac{d^3 \\mathbf p \\, d^3 \\mathbf q}{(2 \\pi)^6} \\, J^{(s)}_3 (\\mb p, \\mb q, \\mb k - \\mb p -\\mb q) \\,  T_0(-\\mb k, \\mb p, \\mb q, \\mb k - \\mb p -\\mb q), \\ea and $J_n^{(s)}$ indicates a kernel which has been symmetrized with respect to its arguments. We have checked that the two-loop contributions proportional to $T_0$ are, in general, negligibly small and will not be considered hereafter if $\\gnl=0$. In the left panel of  Fig. \\ref {fig:Pmmk} we plot the ratio between the matter power spectra originating from a non-Gaussian model (with $\\fnl \\neq 0$ and $\\gnl=0$) and from Gaussian initial conditions. We consider several values of $\\fnl$ and we compare our analytical results with data from the N-body simulations by PPH08 at both redshift 0 and 1. Primordial non-Gaussianity alters the matter power spectrum at the few percent level for $k<0.2 \\ h$ Mpc$^{-1}$ and these deviations are remarkably well reproduced by the one-loop corrections.  \\begin {figure} \\begin{center} \\includegraphics[width=0.45\\columnwidth,angle=0]{figures/fig_Pm_fNL_wdata_z0z1.eps} \\includegraphics[width=0.45\\columnwidth,angle=0]{figures/fig_Pk_z0z1_Dalal_M5M3.eps} \\end{center} \\caption{\\textbf{Left:} Deviation of the matter power spectrum in models with different $\\fnl$ (and $\\gnl=0$) from the Gaussian case at $z=0$ (top) and $z=1$ (bottom). The lines indicate our one-loop calculation for different values of  $\\fnl$ while points with error bars correspond to the N-body simulations by PPH08. \\textbf{Right:} Halo-matter cross spectrum at $z=0$ for a narrow bin of halo masses centered around $M = 2 \\cdot 10^{14} M_{\\odot}/h$ (top) and at $z=1$ for $M = 5 \\cdot 10^{13} M_{\\odot}/h$ (bottom). The solid and dashed lines have been obtained using our model with the bias parameters from the LV and PPH mass functions, respectively. The dotted lines indicate the model by Dal07, while points with error bars correspond to the simulations by PPH08. } \\label{fig:Pmmk}  \\label{fig:Phmk} \\end{figure}  \\paragraph* {\\bf Halos} The halo-halo power spectrum (and similarly the halo-matter cross spectrum) deriving from our multivariate biasing scheme can be written as \\be P^{hj}(k,z) =  D^2(z) \\, P^{hj}_{11}(k) + D^3(z) \\, P^{hj}_{12}(k) + D^4(z) \\, \\left[ P^{hj}_{22}(k) + P^{hj}_{13}(k) \\right]\\;, \\ee where the superscript $j$ indicates either matter ($m$) or halo ($h$) fluctuations. The full expressions of the different terms are lengthy and we report them only in the Appendix. We highlight that our calculation reduces to: (a) the linear result by Dal07 if we only consider the leading-order terms (and further assume that the mass function does not depend on $\\fnl$); (b) the usual one-loop Gaussian expression derived e.g. by \\cite {Heavens:1998es} if we set $\\fnl=0$, and (c) the non-Gaussian result by \\cite{Taruya:2008pg} if we ignore the terms which are proportional to the potential perturbations $\\varphi$ in our multivariate biasing scheme. Notice that for $k \\rightarrow 0$ we obtain $P^{hh}(k) \\propto P_0(k)/\\alpha^2(k)$ and $P^{hm}(k) \\propto P_0(k)/\\alpha(k)$. In the right panel of Fig.~\\ref{fig:Phmk}, we test our theoretical predictions for the halo-matter cross spectrum as a function of $\\fnl$ and for $\\gnl=0$ (solid lines) against the N-body data by PPH08. The bias factors in the models have been calculated from the LV and PPH mass functions. We consider halos with mass $M \\simeq 2 \\cdot 10^{14} M_{\\odot}/h$ at $z=0$ and $M \\simeq 5 \\cdot 10^{13} M_{\\odot}/h$ at $z=1$. We have chosen two different mass bins to keep the number of halos at each redshift large enough to avoid substantial shot noise contamination in the simulations. Our analytical results are in very good agreement with the simulation data for the whole range of $\\fnl$ and up to scales $k \\lesssim 0.2 \\, h$~Mpc$^{-1}$. Note that our spectra differ from the linear result by Dal07 (over-plotted with dotted lines) both on large and small scales. The small-scale departure is due to the non-linear growth of perturbations which we take into account up to the third perturbative order. The large-scale discrepancy is discussed in detail in the next subsection. In Fig.~\\ref{fig:PhmM} we show how the cross spectrum depends on the halo mass at three different wavenumbers and for two redshifts. Independently of halo mass, at $z=1$ our model accurately matches the outcome of the simulations for $k<0.2 \\, h$~Mpc$^{-1}$. Likewise, at $z=0$, the theory agrees well with the numerical data on the largest scales  while it tends to over-predict the cross power for $k>0.1 \\, h$~Mpc$^{-1}$ and $M<10^{14} M_\\odot/h$.   We have checked that, for $\\gnl = 0$, the terms proportional to the trispectra of the potentials $\\Phi$ and $\\varphi$ are generally subdominant even when they generate contributions to the halo power spectrum which diverge as $k\\to 0$. For instance, the trispectrum contribution to the term $\\langle \\widetilde{\\delta_1^2} \\widetilde{\\delta_1^2} \\rangle$ in the halo-halo power spectrum scales as $P_0(k)/\\alpha^2(k)$ as like as the leading term. However, for $|\\fnl|<500$, this correction contributes at most at percent level and only on very large scales. Also the trispectrum contribution in $ \\langle \\widetilde{\\delta_1^3} \\tilde\\delta_1 \\rangle$ which scales as $P_0(k)/\\alpha(k)$ is subdominant ($\\ll 1\\%$) for both the halo-halo and the halo-matter cases. Similar conclusions can be drawn for the terms including $\\varphi$: we have checked that the contributions arising from averages where one or more $\\delta_1$ are replaced by $\\varphi$ are subdominant.  \\paragraph* {\\bf The effect of $\\mathbf {\\gnl}$} The situation becomes more complicated if we consider also the third-order term in Eq.~(\\ref{eq:basichi}) with realistic values of $\\gnl$. In this case, there will be new contributions to the halo and halo-matter power spectra coming from two sources: the trispectrum gets the additional term of Eq.~(\\ref{eq:deltaT}), and the bias factors become altered as described in Eq.~(\\ref{eq:dbgnl}); each of these modifications is linear in $\\gnl$.  Considering first the effect of $\\Delta T_{\\Phi}$, we have found that this adds a (negligible) constant contribution to $\\langle \\widetilde{\\delta_1^2} \\widetilde{\\delta_1^2} \\rangle$ but generates another term which scales as  $P_0(k)/\\alpha(k)$ in $\\langle \\tilde\\delta_1  \\widetilde{\\delta_1^3} \\rangle$. The latter can become the dominant contribution on large scales and for high values of $\\gnl \\gtrsim 10^5$. In the limit $k \\to 0$, this two-loop term (whose full expression is given in Eqs.~(\\ref{eq:PII13},\\ref{eq:PII13b}) in the Appendix for the halo and halo-matter cases) reduces to \\ba \\label{eq:triII} P_{13}^{hh, II} (k) &\\to& \\gnl \\, b_{10} b_{30} \\, \\sigma^4 (R) \\, \\Sigma_3(R) \\, \\frac{P_0(k)}{\\alpha(k)} \\propto \\gnl \\, k^{n_s-2} \\, \\nn \\\\ P_{13}^{hm, II} (k) &\\to&  \\frac 1 2 \\,\\gnl \\, b_{30} \\, \\sigma^4 (R) \\, \\Sigma_3(R) \\, \\frac{P_0(k)}{\\alpha(k)} \\propto \\gnl \\, k^{n_s-2} \\, , \\ea where we define \\be \\Sigma_3(R) \\equiv \\int \\frac{d^3 \\mb p \\, d^3 \\mb q}{(2 \\pi)^6} \\, \\frac{P_0(q)}{\\alpha(q)}\\, \\frac{P_0(p)}{\\alpha(p)} \\, \\alpha (|\\mb p + \\mb q|) \\, W(q R) \\, W(p R) \\, , \\ee such that $S_3 = \\fnl \\, \\Sigma_3$. In our formalism, $P_{13}^{hh, II} (k)$ corresponds to the leading contribution found by \\cite {Desjacques:2009jb}, who used it to derive observational constraints on $\\gnl$ (assuming $\\fnl=0$). In the limit of high peaks, the scaling $b_{10} b_{30} \\to (\\delta_c / \\sigma)^4$ is recovered. Note, however, that the amplitude of this term depends on the adopted smoothing scale $R$ first introduced in Section~\\ref{sec:tracers}, and further discussed in Section~\\ref{sec:smoo}.  Second, we look at the effect of $\\Delta b$. As shown in Eq.~(\\ref{eq:dbgnl}), only the coefficients $b_{02}$ and $ b_{12}$ are altered by $\\gnl$. In the halo-matter case these new terms are subdominant with respect to the trispectrum contribution of Eq.~(\\ref{eq:triII}). However, in the halo-halo spectrum, the leading term for $k \\to 0$ is $P^{hh}_{(23)(23)} \\propto \\langle \\widetilde {\\varphi^2} \\widetilde {\\varphi^2} \\rangle \\propto b_{02}^2$, which scales as \\be P^{hh}_{(23)(23)} \\to \\frac 1 2 \\, b_{02}^2 \\, \\sigma^2_{\\varphi} (R) \\, \\frac{P_0(k)}{\\alpha^2(k)} \\propto  \\gnl^2 \\, k^{n_s-4} \\, , \\ee where $\\sigma^2_{\\varphi} = \\int d q \\, q^2 \\, [P_0 (q) / \\alpha^2 (q)] \\, W^2 (q R) / (2 \\pi^2)$. Because of the quadratic dependence on $\\gnl$, this term dominates on very large scales for high values of $\\gnl$ and small $\\fnl$. Its dependence on the smoothing radius $R$ is very weak because the potential $\\varphi$ is nearly scale invariant. Note that the terms in $\\fnl \\gnl$ are generally subdominant.  \\begin {figure} \\begin{center} \\includegraphics[width=0.45\\columnwidth,angle=0]{figures/fig_Dalal_Phm_M_z0_500.eps} \\includegraphics[width=0.45\\columnwidth,angle=0]{figures/fig_Dalal_Phm_M_z1_500.eps} \\end{center} \\caption{The halo-matter cross spectrum as a function of halo mass at $z=0$ (left) and $z=1$ (right) is plotted for three different values of the comoving wavenumber $k$ (in $h$ Mpc$^{-1}$). Colored solid and dashed lines indicate the results of our perturbative calculation using the LV and PPH mass functions, respectively. Data points with error bars correspond to the N-body simulations by PPH08. The thin lines show the linear theory by Dal07 (black), also corrected with the factor $\\beta$ (magenta) introduced by PPH08 (see Section \\ref{sec:asymp} for further details). At $z=0$  structure has evolved further into the non-linear regime, so that the range of validity of both the linear and one loop theories is reduced. } \\label{fig:PhmM} \\end{figure}   \\subsection{Bias and asymptotic behavior on large scales} \\label{sec:asymp} Let us define the effective bias function \\be b_{\\mathrm{eff}} (k, \\fnl) \\equiv \\frac {P^{hm} (k, \\fnl)}{P^{mm} (k, \\fnl)}\\;, \\ee and compare it with the standard Gaussian bias by introducing the bias deviation \\be \\Delta b (k, \\fnl) = b_{\\mathrm{eff}} (k, \\fnl) - b_{\\mathrm{eff}} (k, 0). \\ee In this section we will only consider the case $\\gnl =0$. In the limit $k \\to 0$ and for large $R$ ($\\sigma_R^2 \\ll 1$), the dominant contribution to the halo power spectrum is given by the tree-level term, and we thus obtain \\be \\label{eq:dbaaa} \\Delta b_{\\mathrm{linear}} (k) = b_{10} (\\fnl) - b_{10} (\\fnl = 0) + 2\\fnl \\delta_c \\, [b_{10}(\\fnl)-1]/\\alpha(k)\\;. \\ee The scale-dependent non-Gaussian correction is proportional to the factor $b_{10}-1$ as originally shown by Dal07, although there is also an additional scale-independent correction, due to the fact that in our model $b_{10}$ is a function of $\\fnl$ (similar conclusions have been reached by \\cite{Slosar:2008hx, Desjacques:2008vf, Pillepich:2008ka, Desjacques:2009jb, Valageas:2009vn} following different approaches). In the simplest model by Dal07 the scale-independent term is missing: \\be \\Delta b_{\\mathrm{Dal07}} (k) =  2\\fnl \\delta_c \\, [b_{10}(\\fnl)-1]/\\alpha(k)\\;. \\label{eq:dalal} \\ee In the left panel of Fig.~\\ref{fig:db}, we plot $\\Delta b (k)$ for $\\fnl=500$ and test the different models against the N-body simulations by PPH08. We can see that considering only the scale-dependent term as in Dal07 does not match the simulations very well, since, contrary to the N-body data, $\\Delta b$ cannot change sign with increasing $k$ (see also  \\cite{Pillepich:2008ka}). The agreement vastly improves if we use Eq.~(\\ref{eq:dbaaa}) with the bias parameters  computed from a non-Gaussian mass function (LV) as this adds a constant negative shift (see the left panel of Fig.~\\ref{fig:b1b2b3}) to the bias deviation. Considering the full calculation to third perturbative order further improves the agreement with the simulations for $k>0.1\\ h$ Mpc$^{-1}$ up to a maximum value of the wavenumber which depends on the adopted smoothing scale for the perturbative calculations (see Section~\\ref{sec:smoo} for further details). \\begin {figure} \\begin{center} \\includegraphics[width=0.45\\columnwidth,angle=0]{figures/fig_deltab_alllin3.eps} \\includegraphics[width=0.45\\columnwidth,angle=0]{figures/fig_deltab_ratio.eps} \\end{center} \\caption{ \\textbf{Left:} Change in the effective bias $\\Delta b$  for fixed $\\fnl=500$, $M = 2 \\cdot 10^{14} M_{\\odot}/h$, at $z=0$.  Different models are compared with N-body simulations: the simple $\\Delta b_{\\mathrm{Dal07}}$ (dotted), $\\Delta b_{\\mathrm{linear}}$ (dashed) using the LV mass function, and the full one-loop theory (solid), which yields the best match. \\textbf{Right:} Fractional difference between the one-loop prediction for $\\Delta b$ (with LV and PPH mass functions) and the Dal07 linear theory, compared with the simulations, for different values of $\\fnl$, at the same mass and redshift. } \\label{fig:db} \\label{fig:dbratio} \\end{figure} \\begin {figure} \\begin{center} \\includegraphics[width=0.45\\columnwidth,angle=0]{figures/fig_deltab_fnl.eps} \\includegraphics[width=0.45\\columnwidth,angle=0]{figures/fig_deltab_Dalal_b1_z0.eps} \\end{center} \\caption{\\textbf{Left:} Bias deviation at $z=0$ as a function of $\\fnl $ at three different wavenumbers, and for $ M = 2 \\cdot 10^{14} M_{\\odot}/h$. The results for LV (solid) and PPH (dashed, within its range of validity) mass functions are shown.  We also plot the prediction for the linear (Dal07) theory. \\textbf{Right:} As in the left panel but as a function of $b_1^L $ and for $\\fnl = 500$. The normalization factor $\\alpha/\\Gamma \\equiv \\alpha / (2 A)$ is chosen to reproduce figure 11 in PPH08. } \\label{fig:dbfnl} \\label{fig:dbb} \\end{figure} To highlight the importance of the non-linear and scale-independent corrections, in the right panel of Fig.~\\ref{fig:dbratio} we plot the ratio \\be \\frac {\\Delta b (k, \\fnl)} {\\Delta b_{\\mathrm{Dal07}} (k, \\fnl)} \\ee for several values of $\\fnl$ and using both the LV and PPH mass functions. We consider halos with mass $M = 2 \\cdot 10^{14} h^{-1} M_{\\odot}$ at $z=0$. The simulation data by PPH08 are in good agreement with our third-order calculation, while the linear (Dal07) model cannot reproduce them on scales $k > 0.02 \\, h$~Mpc$^{-1}$. The impact of primordial non-Gaussianity on the halo bias is further explored in the left panel of Fig.~\\ref{fig:dbfnl}, where we show the bias deviation as a function of $\\fnl$ at three selected scales, again for halos with $M = 2 \\cdot 10^{14} h^{-1} M_{\\odot}$ at $z=0$. Note that, contrary to what predicted by the linear (Dal07) model, the relationship between $\\Delta b$ and $\\fnl$ is non-linear, as first observed by PPH08 in their N-body simulations, and this is now fully explained by our perturbative calculation to third order. We finally study the mass dependence of the bias deviation by showing, in the right panel of Fig.~\\ref{fig:dbb}, how $\\Delta b$ changes as a function of the first Lagrangian bias coefficient $b_{10}^L$. We consider three values of the wavenumber in the quasi-linear and mildly non-linear regime, $\\fnl =500 $, and $z=0$. In the Dal07 model, $\\Delta b$ depends linearly on $b_1^L$ and the relation between these two quantities is independent of $k$. This is indicated by the solid black line which departs more and more from the simulation data with increasing $k$. In order to better describe the N-body results, PPH08 presented a fitting function, $\\beta(k,\\fnl)$, in the form of a multiplicative (scale-dependent) correction to the Dal07 model (magenta lines). Note that our third-order calculations match well the numerical outcome. Some discrepancy is noticeable for $k=0.2\\ h$ Mpc$^{-1}$ and $b_{10}^L<1$, where non-linear effects become more important (due to the small first bias coefficient) and perturbation theory becomes less accurate. The difference between the bias deviations obtained with the LV and PPH mass functions at large halo masses emphasize the need for accurate parameterizations of the halo counts for the rarest objects.   Our result in Eq.~(\\ref{eq:dbaaa}) differs from the perturbative calculations based on the univariate local bias by \\citep{Taruya:2008pg, Sefusatti:2009qh} where the scale-dependent part of the bias deviation was found to scale as $b_{20}$ times the variance of the mass density field. Strictly speaking, for $\\fnl \\neq 0$ and $k\\to 0$ our result gives $P^{hm}(k,\\fnl) \\to (b_{01}+2 b_{20} \\fnl \\sigma^2_R) P_0(k)/\\alpha(k)$ but the term proportional to $b_{20}$ is suppressed by smoothing on the scale $R$ (which is necessary to truncate the bias expansion at third order in a meaningful way). Note that, using the PS expression for the bias parameters in the limits of high peaks, $\\delta_c/\\sigma \\gg 1$, this reduces to $P^{hm}(k,\\fnl) \\to 2 \\fnl (\\delta_c^2/\\sigma^2) [1+ (\\sigma^2_R/\\sigma^2)]P_0(k)/\\alpha(k)$. In this case, the two contributions are identical if $R$ is chosen to be the Lagrangian radius of the halos (i.e. $R=R_f$) and the same window function is used to compute $\\sigma$ and in the calculation of the perturbative power spectra. In general, however, using the Lagrangian radius of the halos gives $\\sigma$ of order unity and this is too large to allow the truncation of the bias expansion at a finite order. For this reason in our calculations we use $\\sigma_R<\\sigma$ and the smoothing-dependent contribution proportional to $b_{20}$ is subdominant.\\footnote{ It is interesting to see what alternative approaches find regarding this discrepancy. For instance, Matarrese \\& Verde \\cite{Matarrese:2008nc} computed the two-point correlation function of regions where the density exceeds a high threshold $\\delta_c\\gg \\sigma$ and found that, for large separations, the non-Gaussian correction scales as $(\\delta_c^3/\\sigma^6)\\, \\xi_3(\\mb x_1,\\mb x_1, \\mb x_2)$ with $\\xi_3$ the three-point correlation function of the mass density. In Fourier space this coincides with the high-peak limit of our result but, provided that $\\sigma^2=\\sigma^2_R$, it also matches the result by \\citep{Taruya:2008pg}. This happens because for high peaks both $\\delta_c (b_{10}^L)^2$ and $b_{10}^L b_{20}^L$ are proportional to $\\delta_c^3$. The same ambiguity applies to the higher-order calculations in Desjacques \\& Seljak \\cite{Desjacques:2009jb}.}  As highlighted in Eq.~(\\ref{eq:bfexpEuler}) the leading order for $\\delta_h$ in our multivariate biasing scheme includes a term proportional to $\\varphi$ and this generates the scale-dependent correction in $\\Delta b$. The proportionality with $b_{20}$ found by other authors derives from the assumption that a local deterministic bias scheme holds true also in the presence of non-Gaussian perturbations. In this case, for $k\\to 0$, second-order terms dominate over the tree-level contribution to the power spectrum which casts some doubts on the validity of the perturbative expansion. In the left panel of Fig.~\\ref{fig:important} we show the difference between our multivariate approach and the standard local bias. The asymptotic scale dependence $ P^{hm}(k) \\propto \\alpha^{-1}(k) \\, P_0(k) \\propto k^{n_s-2} $ for $k \\rightarrow 0$ is recovered in both cases, but the amplitude of the diverging term in the standard local bias model depends on the smoothing length that has to be introduced to cure the ultraviolet divergence of the mass variance. In  Fig.~\\ref{fig:important} we use a smoothing scale of 10 $h^{-1}$ Mpc for both models and the asymptotic term deriving from the standard local bias is strongly subdominant with respect to the correction given in Eq.~(\\ref{eq:dbaaa}). As discussed by \\cite{Taruya:2008pg}, the result of the univariate local model depends strongly on the smoothing scale, as it is proportional to $\\sigma^2(R)$. For instance, using $R = 2 h^{-1}$Mpc boosts the amplitude of the scale-dependent bias (see  Fig.~\\ref{fig:important} (left)). This is why only the multivariate model can reproduce the results from N-body simulations by PPH08 without tuning additional parameters. The main practical advantage in this case is that the bias parameters can be predicted from a model for the mass function, while the results from the standard local bias can only be used after ``renormalizing'' the bias coefficients \\cite{McDonald:2006mx} and using them as fitting functions. However, even though one can play with the parameters of the theory to fit some data, one should not forget that the physical origin of the scale-dependent bias is that large-scale fluctuations in $\\delta_h$ trace perturbations in $\\varphi$ and this is not accounted for by the standard local bias model. The differences between the models are further highlighted in the right panel of Fig.~\\ref{fig:importantrelat}, where we plot the ratio of the halo-matter cross spectra obtained with different approximations with respect to our full non-Gaussian one-loop calculation, at $\\fnl=500$ and at $z=0$ and $1$. This figure summarizes all the conclusion we have reached in this section: (1) the univariate local biasing assumption yields the correct $k$-dependence but the wrong amplitude of the spectrum on large scales; (2) the linear approximation in Eq.~(\\ref{eq:dbaaa}) lacks small-scale power; (3) the simpler model by Dal07 also features a scale-independent offset in the effective bias. Similar results can be obtained for the halo-halo power spectrum, for which the asymptotic scale dependence is $ P^{hh}(k) \\propto \\alpha^{-2}(k) \\, P_0(k) \\propto k^{n_s-4} $ for $k \\rightarrow 0$ .    \\begin {figure} \\begin{center} \\includegraphics[width=0.45\\columnwidth,angle=0]{figures/fig_Pk_locnonloc.eps} \\includegraphics[width=0.45\\columnwidth,angle=0]{figures/fig_Pk_relativecorr2.eps} \\end{center} \\caption{ \\textbf{Left:} The standard univariate local bias prescription (local) is compared with our multivariate scheme for the halo bias (LV full). We consider the same halo masses as in Fig.~\\ref{fig:Phmk} (right) and we plot the halo-matter cross spectra deriving from the two bias models at $z=0,1$ and with a smoothing scale $R=10 h^{-1}$ Mpc. In the Gaussian case ($\\fnl=0$) the cross spectra coincide, but they are very different for $\\fnl=500$. While the asymptotic scale dependence for $k \\rightarrow 0$ is recovered in both cases, only the multivariate model can reproduce the results from the N-body simulations. We also show  the dependence on $R$ in the univariate model, by overplotting the result with  $R=2 h^{-1}$ Mpc. \\textbf{Right:} Ratio between the halo-matter cross spectra obtained with different approximations and our full perturbative calculation, using $\\fnl=500$ and for the same redshifts and halo masses as in Fig~\\ref{fig:Phmk}. Assuming univariate local biasing (red, solid) severely under-predicts the large-scale power. On the other hand, the linear approximations in Eqs.~(\\ref{eq:dbaaa}) (green, dashed) and (\\ref{eq:dalal}) (blue, dotted) lack small-scale power. Notice that the Dal07 model also features a constant offset on large scales due to the missing scale-independent correction discussed in the main text. } \\label{fig:important} \\label{fig:importantrelat} \\end{figure}   \\subsection {On the smoothing and other perturbative approaches} \\label{sec:smoo}  The series expansion in Eq.~(\\ref{ldb}) only applies when $\\delta_h$ and $\\delta$ have been smoothed on a scale $R$. The reason is twofold. First, the locality of the bias is expected to degrade progressively when smaller and smaller scales are considered. Second, we truncate both the bias and the perturbative expansions to finite order which is a good approximation only if the neglected terms give a small contribution. This requires that typically $\\delta \\ll 1$, i.e. $R\\gg 5 h^{-1}$ Mpc.  As explicitely indicated in the Appendix, in this paper we have used a Gaussian kernel $W(kR)=\\exp[-(kR)^2/2]$ with $R=10\\ h^{-1}$ Mpc to smooth the evolved fields $\\delta_1$, $\\delta_2$ and $\\delta_3$ before applying Eq.~(\\ref{ldb}). An obvious consequence of this procedure is that also the resulting halo power spectrum is suppressed on scales $ k \\gtrsim R^{-1}$. A common prescription to lessen this effect and extend the theory to (slightly) larger wavenumbers is to divide out $W^2(kR)$ from the perturbative result for $P^{hh}(k)$ and $P^{hm}(k)$, as introduced by \\cite {Smith:2006ne}. Here we have followed this approach to consider wavenumbers up to $k\\sim 0.3 \\ h$ Mpc$^{-1}$.  Some of the one-loop corrections to the halo power spectrum present ultraviolet divergences that are automatically cured by using a finite value of $R$. However, some of these integrals give rise to scale-independent contributions for $k\\to 0$ whose amplitude depends on $R$, as shown already by \\cite {Heavens:1998es}. This is somewhat unsatisfactory, since it makes the results dependent upon a non-fundamental quantity, and it is amongst the reasons which have led to the application of renormalization techniques to the theory, often borrowed from other areas of physics. The existing approaches, as recently reviewed by \\cite {Carlson:2009it}, include the renormalized perturbation theory \\cite{Crocce:2005xy,Crocce:2005xz}, the closure theory \\cite{Taruya:2007xy}, the time renormalization group flow model \\cite{Pietroni:2008jx}, and the renormalization group perturbation theory \\cite{McDonald:2006hf,McDonald:2006mx}. Most of these approaches do not include a bias model and just apply to the matter density field. The renormalization of the bias parameters, included in some of the models, makes the theory free from any undesired dependence on the smoothing scale. This can be achieved by grouping different perturbative terms together and relabeling some parameters to include the smoothing-dependent factors, a procedure which is not uniquely defined. An altogether different approach which does not need such an operation is the Lagrangian resummation theory by \\cite{Matsubara:2007wj,Matsubara:2008wx}.  On the other hand, SPT has the advantage of remaining a fully predictive theory, where the bias coefficients can be calculated as a function of halo mass. Furthermore, the choice of the smoothing scale $R$ is not completely arbitrary, but confined to a rather narrow range around $R \\simeq 10\\ h^{-1}$~Mpc. Indeed, the smoothing needs to be $R \\gtrsim 8 \\ h^{-1}$~Mpc in order not to break the validity of the perturbative expansion in a significant fraction of the volume ($\\sigma \\ll 1$ for matter and $\\sigma \\ll b_{20}/b_{10}$ for halos) and, on the other hand, $R$ needs to be as close as possible to this limit if we want to prevent the smoothing from wiping out the non-linear corrections at the wavenumbers of interest.  We have checked numerically that different choices of the smoothing scale $R$ within a reasonable range larger than the Lagrangian size of the halos do not affect our results significantly. Notice that, for non-Gaussian perturbations, the $k$-independent -- but $R$-dependent -- terms arising in SPT on large scales  are less important, since the halo power spectrum grows with decreasing $k$.   \\section {Bispectra} \\label{sec:bispectrum}  The leading contributions to the halo bispectrum from non-Gaussian initial conditions of the local type have been recently computed in the framework of the local bias model given in Eq.~(\\ref{ldb}) \\cite{Sefusatti:2007ih,Sefusatti:2009qh,Jeong:2009vd}: \\be B_h(\\mb k_1, \\mb k_2, \\mb k_3)=b_1^3 B_\\delta(\\mb k_1,\\mb k_2,\\mb k_3)+b_1^2b_2\\left[P_\\delta(k_1)P_\\delta(k_2)+ \\mathrm{(2 \\ cyc.)}+\\frac{1}{2}\\int\\frac{d^3q}{(2\\pi)^3} T_\\delta(\\mb q,\\mb k_1-\\mb q,\\mb k_2,\\mb k_3) + \\mathrm{(2\\ cyc.)}\\right]. \\label{eq:bisploc} \\ee Using Eulerian perturbation theory to follow the growth of density perturbations gives up to fourth order \\be B_\\delta(\\mb k_1,\\mb k_2,\\mb k_3)\\simeq B_0(\\mb k_1,\\mb k_2,\\mb k_3)+ 2 F_2(\\mb k_1, \\mb k_2) P_0(k_1) P_0(k_2) + \\mathrm{ (2\\ cyc.)}\\;, \\ee where the second term is generated by non-linear gravity while \\be B_0(\\mb k_1,\\mb k_2,\\mb k_3)\\,\\simeq\\, 2 \\fnl \\,    \\left[  \\alpha (k_3) \\, \\frac {P_0(k_1)\\,P_0(k_2)}{\\alpha (k_1) \\, \\alpha (k_2)} +  \\mathrm{2 \\ cyc.} \\right] \\, \\ee is the linear matter bispectrum due to primordial non-Gaussianity. Similarly, the term between square brackets in Eq.~(\\ref{eq:bisploc}) reduces to \\be P_0(k_1) P_0(k_2) + \\mathrm{(2\\ cyc.)}+\\frac{1}{2}\\int \\frac{d^3q}{(2\\pi)^3} T_0(\\mb q,\\mb k_1-\\mb q,\\mb k_2,\\mb k_3) + \\mathrm{(2\\ cyc.)}\\;. \\ee  In full analogy with the power spectrum calculation discussed above, it is straightforward to show that our new expansion of $\\delta_h$ in terms of both $\\delta$ and $\\varphi$ gives rise to many additional contributions. The most compact form is obtained by writing the expansion of the product of three $\\delta_h$ evaluated in real space at three different locations. For the term which generates the contribution to the halo bispectrum which is proportional to the linear matter bispectrum, we have: \\be b_1 \\delta_1 \\delta_1 \\delta_1 \\to b_{10} \\delta_1 \\delta_1 \\delta_1 + b_{10}^2 b_{01}\\varphi \\delta_1 \\delta_1 + \\mathrm{(2\\ cyc.)} + b_{10} b_{01}^2 \\varphi \\varphi \\delta + \\mathrm{(2\\ cyc.)} +b_{01}^3 \\varphi \\varphi \\varphi\\;. \\label{eq:bisp3ord} \\ee On the other hand, for the term accounting for the non-linear evolution of the density, one finds: \\be b_1^3 \\delta_1 \\delta_1  \\delta_2 + \\mathrm{(2\\ cyc.)} \\to b_{10}^3 \\delta_1 \\delta_1 \\delta_2 + \\mathrm{(2\\ cyc.)} + b_{10}^2 b_{01} \\varphi\\delta_1 \\delta_2+ \\mathrm{(5\\ cyc.)}+ b_{10} b_{01}^2 \\varphi \\varphi \\delta_2 + \\mathrm{(2\\ cyc.)}\\;. \\label{eq:bisp4orda} \\ee Finally, for the source of the term between square brackets in in Eq.~(\\ref{eq:bisploc}), we get: \\ba \\frac{1}{2} b_1^2 b_2 \\delta_1 \\delta_1 \\delta_1^2 + \\mathrm{(2\\ cyc.)} &\\to& \\frac{1}{2}  b_{10}^2 b_{20} \\delta_1 \\delta_1\\delta_1^2+ \\mathrm{(2\\ cyc.)} + b_{10}^2 b_{11} \\delta_1 \\delta_1 (\\varphi \\delta_1)+ \\mathrm{(2\\ cyc.)} + \\frac{1}{2} b_{10}b_{20}b_{01} \\delta_1\\varphi \\delta_1^2+ \\mathrm{(5\\ cyc.)} + \\nonumber \\\\ &&\\frac{1}{2}  b_{10}^2 b_{02}\\delta_1\\delta_1 \\varphi^2 + \\mathrm{(2\\ cyc.)} +b_{10}b_{01}b_{11}\\delta_1\\varphi (\\varphi \\delta_1) + \\mathrm{(5\\ cyc.)} +\\frac{1}{2} b_{01}^2 b_{20} \\varphi \\varphi \\delta_1^2+ \\mathrm{(2\\ cyc.)} +\\nonumber \\\\ &&\\frac{1}{2} b_{10}b_{01} b_{02} \\delta_1\\varphi \\varphi^2 + \\mathrm{(5\\ cyc.)}+ b_{01}^2b_{11} \\varphi\\varphi(\\varphi\\delta_1)+ \\mathrm{(2\\ cyc.)}+ \\frac{1}{2} b_{01}^2 b_{02} \\varphi \\varphi\\varphi^2 + \\mathrm{(2\\ cyc.)}\\;. \\label{eq:bisp4ordb} \\ea  The halo bispectrum can be computed by Fourier transforming the expressions above. For instance, from Eq.~(\\ref{eq:bisp3ord}) we obtain: \\be \\label{eq:bisp} B_{h}(\\mb k_1,\\mb k_2,\\mb k_3) \\simeq \\left[ b_{10} + \\frac {b_{01}} {\\alpha(k_1)}  \\right] \\left[ b_{10} + \\frac {b_{01}} {\\alpha(k_2)}  \\right] \\left[ b_{10} + \\frac {b_{01}} {\\alpha(k_3)}  \\right]  B_0(\\mb k_1, \\mb k_2, \\mb k_3) \\;, \\ee while assuming a local-bias scheme would have given $B_h(\\mb k_1,\\mb k_2,\\mb k_3) \\simeq b_{10}^3 \\, B_0(\\mb k_1, \\mb k_2, \\mb k_3)$. The mixed matter-halo bispectra can be obtained in a similar way, and their expressions differ from Eq.~(\\ref{eq:bisp}) only for the presence of a reduced number of scale-dependent bias factors. We plot the equilateral configuration of the bispectra in Fig.~\\ref{fig:Bk}, for two values of $\\fnl$. Since we are only considering the tree-level contribution, the Gaussian bispectrum is vanishing. Note that our leading-order result of Eq.~(\\ref{eq:bisp}) can be reproduced by taking  the analogous formula obtained with the univariate local bias and simply replacing $b_{10}$ with the scale-dependent bias $b_{10}+b_{01}/\\alpha(k)$. More complex equations relate the scale-dependent terms obtained by taking the Fourier transform of Eqs. (\\ref{eq:bisp4orda}) and (\\ref{eq:bisp4ordb}). We will not discuss them in detail here.  It is important to notice that the scale-dependent bias changes the shape dependence of the halo bispectrum. The tree-level term $B_0$ is dominated by the squeezed configurations (where one of the wavevectors is small) and the $\\fnl$-dependent term $b_{01}/\\alpha$ makes it even more so. Comparing our results with the figures in \\cite{Jeong:2009vd} suggests that, in strict analogy with the result for the power spectrum, the terms proportional to $b_{01}/\\alpha$ give by far the dominant contribution to the bispectrum on large scales. The shapes of the bispectrum parts coming from the non-linear growth of perturbations and second-order biasing are very different \\cite{Sefusatti:2007ih, Jeong:2009vd} and this makes the bispectrum a promising tool to measure $\\fnl$.  We can see that the expanded form of the bispectrum deriving from Eqs. (\\ref{eq:bisp4orda}) and (\\ref{eq:bisp4ordb}) depends on all the non-Gaussianity parameters in a non-trivial way: all terms involving the matter bispectrum $B_0$, i.e. all terms involving an average over three $\\delta$'s or $\\varphi$'s bring in a linear dependence on $\\fnl$. Then, all terms involving the matter trispectrum $T_0$, i.e. averages over four $\\delta$'s or $\\varphi$'s, have two contributions, depending  on $\\gnl$ and $\\fnl^2$ (or in general $\\tnl$) respectively. In addition to this, we have additional dependences on $\\fnl$ and $\\gnl$ implicit in the bias factors, as described in Section \\ref{sec:bias}. This complex shape and scale dependence of the bispectrum offers an unique opportunity to simultaneously constrain $\\fnl$, $\\gnl$ and $\\tnl$ and thus distinguish between inflationary models. We will investigate this in the near future.  \\begin {figure} \\begin{center} \\includegraphics[width=0.45\\columnwidth,angle=0]{figures/fig_Bk.eps} \\end{center} \\caption{Theoretical bispectra at tree level in the equilateral configuration: $B^{mmm}(k,k,k), B^{hmm}(k,k,k), B^{hhm}(k,k,k)$, $ B^{hhh}(k,k,k)$ for two values of $\\fnl$.  In the Gaussian case ($\\fnl = 0$) the tree-level bispectrum vanishes.}  \\label{fig:Bk} \\end{figure} ",
        "conclusions": "\\label {sec:concl} We have studied the growth of structure from non-Gaussian initial conditions of the local type. In particular, we have shown that the spatial distribution of dark-matter halos is naturally described by a multivariate local bias scheme where the halo number density depends on the underlying values of the density field $\\delta$, the auxiliary Gaussian potential $\\varphi$, and (possibly) also on its gradient. This bivariate local approach can be equally interpreted as a non-local description in terms of $\\delta$ only, since $\\varphi$ and $\\delta$ are related by the Poisson equation. Adopting the peak-background split, some common parameterizations of the halo mass function, and a local model for the evolution of large-scale density perturbations, we have derived the coefficients of this multivariate expansion as a function of the halo mass and of the parameters quantifying the level of primordial non-Gaussianity.   Using SPT to approximate the non-linear growth of density perturbations, we have computed the halo power spectrum and the halo-matter cross spectrum up to the third non-vanishing perturbative order. For unbiased tracers our result coincides with the matter power spectrum presented by \\cite{Taruya:2008pg}. However, in the most general (biased) case, it differs from what is obtained adopting the standard local bias expansion in terms of the density field \\cite{Taruya:2008pg, Sefusatti:2009qh, Sefusatti:2007ih, Jeong:2009vd}. The most remarkable feature is that the scale-dependent bias first discussed in Dal07 appears at leading order in our model for the power spectrum. This is because in our multivariate biasing scheme halo fluctuations on large scales trace the Gaussian potential $\\varphi$ rather than $\\delta$. However, our model reduces to the usual univariate case on larger scales, where the variance of density fluctuations is much larger than that of the potential. Note that both the multivariate and the univariate models predict that the dominant contribution to the halo power spectrum scales as $\\fnl\\, P_0(k)/\\alpha(k)$ for $k\\to 0$. However, for the standard univariate biasing, the amplitude of this term is given by a badly behaved integral which strongly depends on the assumed smoothing scale. Renormalization of the second bias coefficient (which should then be treated as a fitting parameter when comparing the theory to observation or simulations) is unavoidable in this case, while it is not needed in our model.  We have then tested our results against the N-body simulations by PPH08, finding excellent agreement for both the matter and the halo two-point functions. Focusing on the scale-dependent bias generated by primordial non-Gaussianity, we have shown that our model accounts for the discrepancies previously found between the predictions by Dal07 and the outcome of numerical simulations \\cite {Desjacques:2008vf, Pillepich:2008ka, Grossi:2009an}. Corrections to the simpler model by Dal07 arise for two main reasons: (a) the bias coefficient $b_{10}$ depends on $\\fnl$ due to the fact that the shape of the mass function is altered by primordial non-Gaussianity (see also \\cite{Slosar:2008hx, Desjacques:2008vf, Desjacques:2009jb}), and this adds a scale-independent offset to the bias deviation $\\Delta b$; (b) considering perturbation theory up to third order generates numerous additional corrective factors that become important on intermediate and small scales. With our one-loop calculation of the power spectrum, the range of validity of the theory extends up to scales $k \\sim 0.1-0.3\\ h$ Mpc$^{-1}$  depending on halo mass and redshift.  We have also shown how our calculations can be extended to include higher-order terms of primordial non-Gaussianity, for instance by considering a non-vanishing primordial trispectrum proportional to the parameter $\\gnl$. In this case, the halo power spectrum includes an additional contribution proportional to $\\gnl^2 P_0(k)/\\alpha^2(k)$ which, depending on the values of $\\fnl$ and $\\gnl$, may be dominant on the largest scales. This term only appears in our multivariate expansion and originates from the bias parameter $b_{02}$ which includes a correction proportional to $g_{\\rm NL}$. On the other hand, in agreement with \\cite{Desjacques:2009jb}, we have found that both the halo-halo and halo-matter spectra acquire a dependence on $\\gnl$ from the trispectrum of the linear density field. This term scales as $\\gnl P_0(k)/\\alpha(k)$ but its normalization depends on the assumed smoothing scale and cannot be robustly predicted by the theory.  Finally, we have calculated the halo bispectrum deriving from our multivariate biasing scheme. At tree level, our result corresponds to the usual bispectrum deriving from Gaussian initial conditions but where the linear bias $b_{10}$ is replaced by  $b_{10} + b_{01}/\\alpha(k)$. This is different from what has been found assuming univariate local biasing \\cite{Sefusatti:2007ih, Jeong:2009vd}. Therefore the analysis of three-point statistics represents a promising tool to test the different biasing schemes against observations. Also, the complex shape and scale dependence of the halo bispectrum offers an unique opportunity to simultaneously constrain $\\fnl$, $\\gnl$ and $\\tau_{\\mathrm NL}$ and thus put entire classes of inflationary models under scrutiny. We will explore this in more detail in a forthcoming paper."
    },
    "0911/0911.0684_arXiv.txt": {
        "abstract": "The Smith Cloud is a massive system of metal-poor neutral and ionized gas (M$_{\\rm gas}$ $\\ga 2\\times 10^6$~M$_\\odot$) that is presently moving at high velocity (V$_{\\rm GSR}$ $\\approx$~300~km~s$^{-1}$) with respect to the Galaxy at a distance of 12~kpc from the Sun. The kinematics of the cloud's cometary tail indicates that the gas is in the process of accretion onto the Galaxy, as first discussed by \\citet{Lockman08}. Here, we re-investigate the cloud's orbit by considering the possibility that the cloud is confined by a dark matter halo. This is required for the cloud to survive its passage through the Galactic corona. We consider three possible models for the dark matter halo (NFW, Einasto, Burkert) including the effects of tidal disruption and ram-pressure stripping during the cloud's infall onto and passage through the Galactic disk. For the NFW and Einasto \\footnote{The Einasto profile is sometimes referred to as the Sersic profile \\citep[e.g.][]{Zait08}. These have the same functional form but since the Sersic profile arises from a projected distribution, we prefer to use a radial function. We note, however, that the projection of the Einasto function is nontrivial analytically \\citep{Mazure02}.} dark-matter models, we are able to determine reasonable initial conditions for the Smith Cloud, although this is only marginally possible with the Burkert model. For all three models, the progenitor had an initial (gas+dark matter) mass that was an order of magnitude higher than inferred today. In agreement with \\citeauthor{Lockman08}, the cloud appears to have punched through the disk $\\approx$~70~Myr ago. For our most successful models, the baryon to dark matter ratio is fairly constant during an orbital period but drops by a factor of $2-5$ after transiting the disk. The cloud appears to have only marginally survived its transit, and is unlikely to retain its integrity during the next transit $\\approx$~30~Myr from now. ",
        "introduction": "The Smith Cloud is perhaps the best known and most studied member of the high-velocity clouds (HVCs). While covering over a third of the sky, little is known about the nature of these clouds, with many suggestions as to their origin. While we now know the distances to a number of HVCs, their orbits are more difficult to determine. Since its discovery in \\citeyear{Smith63}, the Smith Cloud was originally believed to be an extension of the Galactic disk \\citep{Smith63} but it was later shown to be separated kinematically from the Galactic disk \\citep{Lockman84} leading to its subsequent classification as an HVC \\citep{Wakker97}.  While the origin of the cloud is still unknown, with suggestions of gas belonging to an infalling dwarf galaxy \\citep{Bland-Hawthorn98} or gas lifted from the Galactic disk \\citep{Sofue04}, the distance, mass and velocity of the cloud are now well established. Distance indicators  based on stellar absorption-line bracketing, H$\\alpha$ flux, and the kinematic distance of gas interacting with the cloud, provide a consistent estimate of $12.4\\pm1.3$~kpc \\citep{Lockman08, Putman03, Wakker08}. The H{\\small{}I} mass of the cloud is found to be $1\\times 10^6$~M$_\\sol$ \\citep{Lockman08}, with WHAM data indicating a comparable or greater amount in ionized hydrogen \\citep{Hill09}. The velocity components have also been determined \\citep{Lockman08} allowing the orbit of the cloud to be calculated for the first time. The Smith Cloud's importance is that it provides insight into gas at a more advanced stage of accretion than the Magellanic Stream.  Earlier studies have shown how easily gas disrupts as it passes through a hot corona \\citep{Moore94,Quilis01,Bland-Hawthorn07,Heitsch09}. It is therefore remarkable that a massive, high-velocity gaseous system has survived so close to the disk. In light of recent developments, we consider the prospect of a confining dark halo surrounding at least some HVCs to be plausible, as first proposed by \\citet{Blitz01}, and therefore worthy of further investigation. A number of ultra-faint dwarf galaxies have now been discovered in the Galactic halo \\citep{Simon07,Kirby08,Geha09}. What is particularly striking is that the inferred dark halo mass exceeds $\\sim 10^7$~M$_\\odot$ even for systems with very few stars \\citep{Strigari08}. These observations hint at the prospect of dark matter dominated galaxies containing \\HI\\ that have not experienced star formation. Complex H is one such `dark galaxy' candidate, with a mass dominated by dark matter if it is self-gravitating \\citep{Simon06}.  Here we suggest that the Smith Cloud is another candidate for a dark galaxy, illustrated in Fig. \\ref{fig:SC}, with the gas partly stabilised by a dark matter component. We calculate the subhalo required for the survival of a gas cloud with properties similar to those observed in the Smith Cloud today. In \\S\\ref{sec:DM} we discuss the properties of three possible dark matter profiles that may encompass the Smith Cloud. In \\S\\ref{sec:MS} we discuss the orbit of the Smith Cloud within a realistic Galactic potential. In \\S\\ref{sec:SAC} we present the equations that are used to calculate the ram pressure stripping of the subhalos along their orbit. In \\S\\ref{sec:MR} we present the results of the semi-analytic models and discuss the implications of these results for the Smith Cloud. In \\S\\ref{sec:conc} we summarize these results and discuss the limitations of our model.  \\begin{figure} \\center \\includegraphics[width=0.45\\textwidth]{f1.eps} \\caption{H{\\scriptsize I} image of the Smith Cloud (Fig. 1. in \\citet{Lockman08}) with superimposed core and halo for three different dark matter models (NFW, Einasto, Burkert) discussed further in \\S\\ref{sec:MR}. The tidally stripped mass today is $1-2\\times10^8$~M$_\\sol$ in dark matter. The core size assumes an initial gas mass at apogalacticon of $1.1\\times10^7$~M$_\\sol$, although most of this gas is distributed along the orbit of the Smith Cloud today. We note that the Burkert profile has no surviving core, and the NFW core is displayed in white for visibility.} \\label{fig:SC} \\end{figure} ",
        "conclusions": "\\label{sec:conc} We have modelled the Smith Cloud by assuming that it is supported against total disruption by a dark-matter subhalo. By considering the evolution of this system over the past $150$~Myr, we are able to compare our predicted properties to the known properties of the Smith Cloud today. Under the assumption of dark-matter confinement, the subhalo encompassing the Smith Cloud has a tidal mass of $\\sim 3\\times10^8$~M$_\\sol$ at the present time. Even before this subhalo arrived at its present orbit, it will have lost $50-90$\\% of its dark matter and a comparable amount of progenitor gas due to tidal and ram pressure stripping by the Galaxy. In this respect, the likelihood of an infalling (dark-matter confined) ``Smith Cloud'' event is comparable to the rate of dwarf galaxy infall, i.e. roughly one event per Gyr.  If the NFW and Einasto profiles are appropriate descriptions of dark matter halos, we find that the Smith Cloud can only exist in a narrow strip of the available parameter space defined by the initial dark matter subhalo mass and the initial gas mass. This remains true with and without dark matter stripping. This narrow strip also appears, albeit shifted with respect to the present axes, if the Galactic halo is considered to be $10\\times$ over or underdense from our model setup, or if the dark matter containing subhalo was virialized at an earlier redshift and correspondingly is denser than a subhalo virialized today.  Consistent with our picture, the observed wake in Fig. 1 presumably resulted from ablation processes due to the impulsive shock of the disk transit. We note with interest that the wake flares (i.e. expands transverse to the direction of the wake) beyond the putative dark matter core (Fig. 1). In our picture, this can be understood in the context of a transsonic flow. The Smith Cloud is moving at up to 300\\kms\\ with respect to a hot halo with adiabatic sound speed $\\sim$200\\kms\\ making the cloud-halo interface mildly supersonic. Once the stripped gas is left behind, it must undergo mixing with the hot halo through Kelvin-Helmholtz instabilities, causing the stream to expand perpendicular to the flow. The temperature of the \\HI\\ gas is predicted to rise along the wake. Unless the confining core-dominated, dark matter halo is very massive, we do not expect the declining potential seen by the wake material to influence the degree of flaring. These issues are being addressed in new hydrodynamical simulations of the Smith Cloud that we will present elsewhere.  With respect to further constraints on shock interactions, we expect that shock signatures at UV to x-ray wavelengths will have largely faded away, and the \\HI\\ ``hole'' at the crossing point will have been substantially stretched by differential shear. For cloudlets smaller than 100 pc, thermal conduction due to the halo corona \\citep{McKee77} and the Galactic radiation field convert the ablated gas to a clumpy plasma. We therefore strongly encourage more extensive mapping of the \\HI\\ and the ionized gas.  As we discuss elsewhere \\citep{Bland-Hawthorn08}, it is very difficult to see how the Smith Cloud, like several other large HVCs, could have come in from, say, the distance of the Magellanic Stream. This was our motivation for considering a cloud stabilised against disruption by a confining dark matter halo. In essence, we have considered the prospect that the Smith Cloud constitutes a `dark galaxy' where star formation never took place. Any evidence of a co-moving stellar population would have profound consequences. An essential requirement for star formation to occur appears to be the presence of molecular gas \\citep[e.g.][]{Calzetti09}. We envisage much deeper searches for the presence of molecular gas in compact HVCs, in particular, with the newly commissioned {\\em Cosmic Origins Spectrograph} on the {\\em Hubble Space Telescope}.  Our model is not unique. The cloud may have been dislodged from the outer disk or confining potential of an infalling dwarf galaxy. A cloud metallicity of [Fe/H]$\\approx$-1 may be appropriate in either scenario but there are no obvious candidates at the present time. The interloper must be on a prograde orbit which rules out some infalling dwarfs \\citep[e.g. $\\omega$Cen;][]{Bekki03} but conceivably implicates disrupting dwarfs like Canis Major, assuming these were still losing gas in the recent past."
    },
    "0911/0911.0198_arXiv.txt": {
        "abstract": "The active RS CVn star II Peg has been spectroscopically monitored for almost 18 years with the SOFIN spectrograph at NOT, La Palma, Spain. In this paper we present five new surface temperature maps of the object for the years 1999 (two maps), 2001 (one map) and 2002 (two maps). ",
        "introduction": "II Peg is one of the very active RS CVn stars and it is the brightest X-ray star with 50pc of the Sun. RS CVn stars are closely detached binaries where the more massive component is a G-K giant or subgiant and the secondary a subgiant or dwarf of spectral class G to M. Because of the low luminosity of the secondary many RS CVn systems appear as single-line binaries which are suitable for spectral analysis (Berdyugina et al. 1998a). In close binaries the rapid rotation is maintained by tidal forces due to the close companion.  Large amplitude brightness variation of RS CVn stars imply the presence of enormous star spots on their surfaces covering up to 50\\% of the visible disk. Also coronal X-ray and microwave emissions, strong flares in the optical, UV, radio and X-ray are seen. Cool spots on the stellar surface will locally alter the photospheric absorption lines and continuum intensities.  Previous investigations on the temperature distribution over the surface of II Peg include the study of Berdyugina et al. (1998b), who presented surface temperature maps for 1992-1996, and Bergyugina et al. (1999c) with surface maps for 1996-1999, both studies were based on observations with the SOFIN-spectrograph at NOT. Gu et al. (2003) presented surface images of II Peg for 1999-2001 based on observations with the Coude echelle spectrograph at the Xinglong station of the National Astronomical Observatories, China, but the spectral lines used for inversions were different to that of SOFIN observations. Photometric light curve variations of the object were analysed by Berdyugina \\& Tuominen (1998c), and by Rodon\\`{o} et al. (2000).  The results of Berdyugina et al. (1998a,b, 1999a,b,c) and Berdyugina \\& Tuominen (1998c) consistently show that there are two active longitudes separated approximately by 180$^{\\circ}$, migrating in the orbital reference frame, and that a switch of activity level occurs periodically with a period of 4.65 years. In the surface images of Gu et al. (2003) the general spot pattern was quite similar, but the drift with respect of the orbital frame was less obvious, and the switch of the activity level appeared to occur earlier than predicted by Berdyugina et al (1999c). Rodon\\`{o} et al. (2000) found a much more complicated spot pattern from their analysis of photometry: they report on the existence of a longitudinally uniformly distributed component together with three active longitudes, with complicated cyclic behavior. Carroll et al. (2009) have also applied a Zeeman Doppler imaging technique to derive the magnetic field configuration on the surface of II Peg during 2007 using spectropolarimetric (Stokes I and V) observations with SOFIN. Their maps show a very similar spot pattern as found by Berdyugina et al. (1998b, 1999c); moreover, the radial field direction is opposite on different active longitudes. ",
        "conclusions": "The following spectral lines were used in the surface temperature inversions: Fe I 6173 \\AA, Ni I 6175 \\AA, Ni I 6177 \\AA, and Fe I 6180 \\AA. Spectral line parameters were obtained from the Vienna Atomic Line Database (Kupka et al. 1999); the $log(gf)$ values of the two included Ni I lines were modified from the standard value of $-0.53$ to $-0.2$ as the standard values were found to produce much weaker absorption lines than the actually observed ones. This is equivalent to an increase in the Ni abundancy, which is probably the real reason for the observed strong Ni lines.  Figures 1-5 show the results of the inversions. During July-August 1999 (Fig. 1) only one spot is seen approximately at latitude 40$^{\\circ}$. The minimum temperature inside the spot was measured to occur at the approximate longitude of 170$^{\\circ}$ or phase 0.47. Our image is quite different from the one obtained by Gu et al. (2003) for almost a simultaneous observing period, but a different spectral region. Their inversions gave much larger, longitudinally extended, spot structures around 170-290$^{\\circ}$ at roughly 60$^{\\circ}$ of latitude.  In September 1999 our inversions also show only one strong spot that has barely moved in the orbital reference frame (latitude 44$^{\\circ}$ and longitude 160$^{\\circ}$ or phase 0.44). There is a weaker, longitudinally extended feature visible between 220-270$^{\\circ}$ or phase 0.61-0.75.  Almost two years later, in August 2001 (Fig. 3), the star exhibits much more surface structures: three spots are visible in our image (longitudes 70$^{\\circ}$, 140$^{\\circ}$, and 200$^{\\circ}$ or phases 0.19, 0.39 and 0.56) at an approximate latitude of 40$^{\\circ}$. The inversions of Gu et al. (2003) for an observing run 5 months later show much less structure, and the maximal activity seems to have moved roughly to the other side of the star than what was seen in their images during 1999 and 2000.  One year later, in August 2002 (Fig. 4) our maps show one strong spot at the latitude 40$^{\\circ}$ and longitude 80$^{\\circ}$ or phase 0.22, and a weaker one at 300$^{\\circ}$ or phase 0.83. In November 2002 (Fig. 5) only the stronger spot is seen approximately at the same location. Comparing the surface temperature distribution during the observing seasons 1999 and 2002, the maximal spot activity seems to have moved by roughly 100$^{\\circ}$ in the orbital reference frame, while very little drift of the spots can be seen during the consequtive observing runs during 1999 and 2002. In between these two 'states' of only one strong spot at different location on the surface, a much more complex distribution could be seen during August 2001. Our images give some support to the conclusion of Gu et al. (2003) of a significant change of the longitudinal spot distribution having occured sometime during 2001 (Fig. 6).  We plan to continue to study the spot evolution on II Peg by analysing photometric and spectropolarimetric observations of the object, both to invert the surface temperature distribution, but also the magnetic field configuration of the object."
    },
    "0911/0911.5068_arXiv.txt": {
        "abstract": "{ Realistic 3D radiative MHD simulations reveal the magneto-convective processes underlying the formation of the photospheric fine structure of sunspots, including penumbral filaments and umbral dots. Here we provide results from a statistical analysis of simulated umbral dots and compare them with reports from high-resolution observations. A multi-level segmentation and tracking algorithm has been used to isolate the bright structures in synthetic bolometric and continuum brightness images. Areas, brightness, and lifetimes of the resulting set of umbral dots are found to be correlated: larger umbral dots tend to be brighter and live longer. The magnetic field strength and velocity structure of umbral dots on surfaces of constant optical depth in the continuum at 630~nm indicate that the strong field reduction and high velocities in the upper parts of the upflow plumes underlying umbral dots are largely hidden from spectro-polarimetric observations. The properties of the simulated umbral dots are generally consistent with the results of recent high-resolution observations. However, the observed population of small, short-lived umbral dots is not reproduced by the simulations, possibly owing to insufficient spatial resolution. ",
        "introduction": "Recent realistic MHD simulations have revealed that the complex brightness and flow structure of sunspots results from convective energy transport dominated by a strong vertical (in the umbra) or inclined (in the penumbra) magnetic field \\citep{Heinemann:etal:2007, Scharmer:etal:2008, Rempel:etal:2009a, Rempel:etal:2009b}. In the umbra, such magneto-convection occurs in the form of narrow upflow plumes with adjacent downflows \\citep[][henceforth referred to as SV06]{Schuessler:Voegler:2006}.  In computed intensity images, these plumes appear as bright features that share various properties with observed umbral dots.  Here we provide a systematic study of the umbral simulation results with specific emphasis on the properties of the bright structures and their comparison with recent results from high-resolution observations. ",
        "conclusions": ""
    },
    "0911/0911.3825_arXiv.txt": {
        "abstract": "In order to know the formation epoch of the oldest elliptical galaxies as a function of mass and observed redshift, a statistical analysis for 333 extremely red objects (EROs) classified as old galaxies (OGs) at $0.8\\le z\\le 2.3$ is carried out. Once we get $M_V$ and $(B-V)$ at rest for each galaxy, we calculate the average variation of this intrinsic color with redshift and derive the average age through a synthesis model (the code for the calculation of the age has been made publicly available). The average gradient of the $(B-V)$ color at rest of EROs/OGs is 0.07--0.10 Gyr$^{-1}$ for a fixed luminosity. The stars in these extremely red elliptical galaxies were formed when the Universe was $\\sim 2$ Gyr old on average. We have not found a significant enough dependence on the observed redshift and stellar mass: $\\left(\\frac{dt_{\\rm formation}}{dt_{\\rm observed}}= -0.46\\pm 0.32\\right)$, $\\left( \\frac{dt_{\\rm formation}}{d\\log _{10}M_*}=-0.81\\pm 0.98\\right)$ Gyr. This fits a scenario in which the stellar formation of the objects that we denominate as EROs-OGs is more intense at higher redshifts, at which the stellar populations of the most massive galaxies form earlier than or at the same time as less massive galaxies. ",
        "introduction": "It is usually sustained that galaxies at high redshift are intrinsically bluer than at low redshift (e.g., Dickinson et al. 2003; Rudnick et al. 2006; Labb\\'e et al. 2007), as would expected if their populations were younger and with lower mass/luminosity ratios. However, analysis of this intrinsic color evolution is not free from caveats, due mainly to the difficulty in disentangling selection effects. For instance, Dickinson et al. (2003, Fig. 2) showed a clear lack of red objects between $2<z<3$ with $(m_{1700\\AA }-m_B)_{AB,rest}>3$, while there were many of these red objects at $z<2$; so they consequently claimed that there is an evolution in color. However some red galaxies were probably missed in that analysis. As a matter of fact, if we take the subsample of old elliptical galaxies without extinction ($f_{\\rm old}\\ge 0.95$) with $2<z<3$ from Miyazaki et al. (2003) ($N=12$ galaxies) we find that all of them have $2.5<(m_{1700\\AA }-m_B)_{rest}<4.5$ with an average of 3.6. Miyazaki et al.'s galaxies are much redder than Dickinson et al.'s, which shows that the Dickinson et al. (2003) sample does not contain the reddest galaxies.   On the question of the ages of galaxies and when were they formed, there also some uncertainties. There are  many observations at different redshifts of galaxies nearly as old as the Universe, although the density of those galaxies at high redshift is significantly lower (Renzini 2006; Abraham et al. 2007). Early-type massive galaxies with early formation were found at $z\\sim 1$ (di Serego Alighieri et al. 2006; 2007), $z=1-2$ (Spinrad et al. 1997; Daddi et al. 2005; Longhetti et al. 2005; Toft et al. 2005; Trujillo et al. 2006), $z=2-3$ (Daddi et al. 2005; Labb\\'e et al. 2005; Cassata et al. 2008, Kriek et al. 2008, 2009), $z=3-4$ (Toft et al. 2005; Chen \\& Marzke 2004), $z=4-5$ (Chen \\& Marzke 2004; Rodighiero et al. 2007) and $z>5$ (Wiklind et al. 2008). Most of this information is obtained from photometry, but there are also some spectra  massive elliptical galaxies at $z=1.4-2.2$ (Cimatti et al. 2008, Kriek et al. 2009) revealing them to be old galaxies. Semi-analytical $\\Lambda $CDM models (De Luc\\'\\i a et al. 2006) claim that the formation of low mass galaxies is first to give way to mergers of very massive galaxies, which is apparently at odds with observation of these massive galaxies at high redshift, even if they were formed through dry mergers in a downsizing scenario. Ferreras et al. (2009) also found that very massive galaxies do not have significant evolution at $z<1.2$. However, Schiavon et al. (2006), who have taken spectra of red galaxies ($U-B>0.25$) at $z\\sim 0.9$, derived their age  to be on average 1.2 Gyr, which means that some galaxies were formed at lower redshifts. Arnouts et al. (2007) also show evidence for a major build up of the red sequence between $z=2$ and $z=1$.  In this paper, we pay further attention to the determination of intrinsic color and age variation for different redshifts in a statistical way for two different samples of very red elliptical galaxies in order to constraint their formation epoch. ",
        "conclusions": "\\label{.form}  The average epoch of star formation (the epoch of formation of the first stars, might be lower) is $t_{\\rm form.}=t_{\\rm obs.}-t_{\\rm gal.}$, shown in Fig. \\ref{Fig:ages_form}. Separating the evolution from the mass dependence, \\begin{equation} t_{\\rm form.}=d_1+d_2(t_{\\rm obs.}-5)+ d_3\\log_{10}(M_*) ,\\end{equation} with $d_1=1.94\\pm 0.51$, $d_2=-0.46\\pm 0.32$, $d_3=-0.81\\pm 0.98$ on average for the ECDFS+Miyazaki samples. The observed EROs/OGs are all formed within a narrow range of epochs, when the Universe was less than 4 Gyr old ($z>1.7$). The present fit is compatible with all EROs/OGs formed at the same time at, on average, $t_{\\rm form.}=2.0\\pm 0.3$ Gyr ($z\\sim 3-4$). Again, we remind the reader that the criterion to select EROs/OGs is more restrictive at lower redshift, picking out only very old galaxies, while at higher redshift the range of allowed ages is wider. We might therefore expect that the oldest EROs at high $z$ (low $t_{\\rm obs.}$) will have an even lower formation age. The age $t_{\\rm form.}=2.0$ Gyr is a conservative lower limit representing the average sample; there must be some EROs/OGs formed beforehand. Given that $\\frac{dt_{\\rm form.}}{dt_{\\rm obs.}}=-0.46\\pm 0.32<<+1$ for a given stellar mass, this means that the galaxies of the present sample are not formed continuously with the same rate, but more intensely at higher redshifts.  The stellar populations of most massive galaxies are not formed much later than the less massive ones ($d_3\\le 0$). This agrees the results mentioned in the introduction that very massive evolved galaxies detected at redshifts 1.5--6 were formed in the very early Universe (Daddi et al. 2005: Chen \\& Marzke 2004; Rodighiero et al. 2007; Wiklind et al. 2008). This might appear in contradiction with the result of \\S \\ref{.colors} that galaxies are redder for lower luminosities, but it does not. As said, the mass--luminosity ratio does not remain constant giving higher luminosity for younger objects so it is not contradictory that older/redder objects correspond to lower luminosities and higher masses. It is in fact observed if we compare Figs. 4 and 5 of Taylor et al. (2009b) for $z>1.25$: clearly, a strong dependence on stellar mass does not mean a strong dependence on luminosity.   \\begin{figure} \\begin{center} \\vspace{1cm} \\mbox{\\epsfig{file=ages.eps,height=6cm}} \\end{center} \\caption{Average age of the stellar populations of EROs which are passively evolving early-type galaxies. The dashed line stands for the limiting maximum $t_{gal}=t_{\\rm obs.}$ within the standard cosmology.} \\label{Fig:ages} \\end{figure}  \\begin{figure} \\begin{center} \\vspace{1cm} \\mbox{\\epsfig{file=ages_form.eps,height=6cm}} \\end{center} \\caption{Average age of the Universe at which the stellar populations of the galaxies observed at age $t_{\\rm obs.}$ have been formed.} \\label{Fig:ages_form} \\end{figure}   \\  {\\bf Acknowledgments:} I thank the anonymous referee for helpful comments in correcting and improving this paper, to A. Vazdekis (IAC) for helpful suggestions on the use of his V96 model and comments on a draft of this paper, and to T. J. Mahoney (IAC) for proof-reading this paper. The author was supported by the {\\it Ram\\'on y Cajal} Programme of the Spanish Ministry of Science."
    },
    "0911/0911.3222_arXiv.txt": {
        "abstract": "We assess the importance of the magneto-rotational instability in core-collapse supernovae by an analysis of the growth rates of unstable modes in typical post-collapse systems and by numerical simulations of simplified models.  The interplay of differential rotation and thermal stratification defines different instability regimes which we confirm in our simulations.  We investigate the termination of the growth of the MRI by parasitic instabilities, establish scaling laws characterising the termination amplitude, and study the long-term evolution of the saturated turbulent state. ",
        "introduction": "\\label{Sek:Intro}  At the end of its hydrostatic evolution, the inner core of a star of more than about eight solar masses, consisting of about 1.5 solar masses of iron or oxygen, neon, and magnesium loses support against its self-gravity and collapses to supranuclear density $\\rho_{\\mathrm{nuc}} \\geq 2 \\times 10^{14} ~ \\mathrm{g\\,cm^{-3}}$, releasing a gravitational binding energy $E_{\\mathrm{grv}} \\sim 10^{53} ~ \\mathrm{erg}$.  When this density is reached, a shock wave is launched and starts to propagate outward.  Still within the core, the shock wave stalls and turns into a stationary accretion shock as energy is lost when nuclei are dissociated in the post-shock region: the prompt explosion fails.  Most of the gravitational binding energy liberated during collapse leaves the star in form of neutrinos, but a small ($\\sim 1\\%$) fraction is used to unbind the envelope and power a supernova (SN) explosion.  How a fraction of $E_{\\mathrm{grv}}$ is transferred to the shock wave and the surrounding envelope is still not fully understood, and forms the central issue of SN theory.  This is not a problem of energy budget but of energy transfer.  The gravitational binding energy of the envelope ($\\sim 10^{50} ~ \\mathrm{erg}$) as well as the kinetic and electromagnetic energies of a SN ($\\sim 10^{49}$ and $\\sim 10^{51} ~ \\mathrm{erg}$, respectively) are much smaller than $E_{\\mathrm{grv}}$.  Consequently, the transfer of a small fraction of $E_{\\mathrm{grv}}$ suffices to revive the shock wave and to trigger the explosion.  However, it is extremely difficult to tap efficiently the energy carried by weakly interacting neutrinos. The mechanisms proposed to account for shock revival and explosion can be grouped into the following categories: \\begin{description} \\item[The spherical neutrino mechanism] Heating of the post-shock matter by interactions with the neutrinos diffusing out of the core revives the shock \\citep[e.g.,][]{Bethe__1990__RvMP__SN_mech}.  This mechanism works for stars around $8 ~ \\mathrm{M}_{\\odot}$ due to a steep density profile facilitating shock propagation \\citep{Kitaura_Janka_Hillebrandt__2006__aap__Explosions_of_O-Ne-Mg_cores_theCrab_supernova_and_subluminous_typeII-P_supernovae}. \\item[Hydrodynamic instabilities] Without the restriction to spherical symmetry, hydrodynamic instabilities can develop: convection in the proto-neutron star (PNS) and behind the shock and the standing accretion shock instability (SASI) can enhance the efficiency of neutrino heating and the transfer of binding energy of the accreting gas to the surrounding matter \\citep[][]{Blondin_Mezzacappa_DeMarino__2003__apj__Stability_of_Standing_Accretion_Shocks_with_an_Eye_toward_Core-Collapse_Supernovae,Foglizzo_et_al__2007__apj__Instability_of_a_Stalled_Accretion_Shock_AAC}. They may also account for the pronounced asphericities of the explosion and the observed pulsar kicks \\citep[e.g.,][]{Scheck_et_al__2006__AA__2d_Aniso}.  While this mechanism has proven successful in detailed 2d simulations of stars up to 15 solar masses employing state-of-the-art microphysics \\citep{Marek_etal__2009__AA__EOS-depnc_nu_GW}, its robustness and applicability to stars of higher mass is still uncertain. \\item[Energy transfer by waves] Apart from neutrinos, waves may carry energy from the central region to the surroundings.  Unstable g-modes in the PNS excited by accreting matter can lead to the emission of acoustic waves travelling upwards and depositing energy near the shock \\citep[][]{Burrows_et_al__2006__apj__A_New_Mechanism_for_Core-Collapse_Supernova_Explosions}. A similar effect can be caused by Alfv\\'en waves excited in the convective region of a magnetised PNS \\citep[][]{Suzuki_Sumiyoshi_Yamada__2008__ApJ__Alfven_driven_SN}. The importance of acoustic waves has been suggested in simulations, whereas the heating by Alfv\\'en waves has not been explored thoroughly. \\item[Rotational mechanisms] Any rotation of the progenitor will be amplified during the collapse; additionally, differential rotation will be created even for an initially rigidly rotating progenitor. If sufficiently rapid, rotation may provide sufficient additional energy to unbind the post-shock matter.  In contrast to the energy of neutrinos, rotational energy can be tapped fairly easily, e.g., by viscosity or magnetic fields \\citep[][]{Thompson_Quataert_Burrows__2004__ApJ__Vis_Rot_SN}. Additionally, a differentially rotating magnetised fluid can be unstable against the magneto-rotational instability (MRI) \\citep[][]{Balbus_Hawley__1991__ApJ__MRI,Balbus_Hawley__1998__RMP__MRI}, leading to field amplification, turbulence and enhanced transport of angular momentum. \\end{description}  None of these mechanisms has been able to provide a robust explanation for SNe across the entire mass range.  Because of the diversity and complexity of the underlying physics, the problem is studied using complementary approaches: very detailed simulations including the best treatment of the microphysics that can be afforded and, on the other hand, simplified models neglecting most of these complications for the benefit of covering a larger parameter space of initial conditions. The study of (magneto-)rotational mechanism relies currently mostly on simplified models.  \\cite{Akiyama_etal__2003__ApJ__MRI_SN} have pointed out that the most basic instability condition of the MRI, i.e., a negative radial gradient of the rotational profile, $\\Omega (r)$, is typically fulfilled in rotating post-collapse cores.  They have, based on simplified spherical modelling, estimated growth times of a few milliseconds and saturation fields of the order of $10^{15} ~ \\mathrm{G}$.  Simplified global simulations \\citep[e.g.,][]{Obergaulinger_Aloy_Mueller__2006__AA__MR_collapse} have shown that fields of this range (reached only for already strongly magnetised progenitors with fields of the order $10^{12} ~ \\mathrm{G}$) can have significant effects on the dynamics of the explosion, e.g., by extracting rotational energy from the inner core or launching jet-like outflows.  Additionally, signs for the growth of the MRI were found in a few of those simulations.  Based only on such simplified studies, a thorough assessment of the influence of the MRI is still not possible. ",
        "conclusions": ""
    },
    "0911/0911.4599_arXiv.txt": {
        "abstract": "{ The current deepest radio surveys detect hundreds of sources per square degree below $0.1\\,$mJy. There is a growing consensus that a large fraction of these sources are dominated by star formation although the exact proportion has been  debated in the literature. However, the low luminosity of these galaxies at most other wavelengths makes determining the nature of individual sources difficult. If future, deeper surveys performed with the next generation of radio instrumentation are to reap high scientific reward we need to develop reliable methods of distinguishing between radio emission powered by active galactic nuclei (AGN) and that powered by star formation. In particular, we believe that such discriminations should be based on purely radio, or relative to radio, diagnostics. These diagnostics include radio morphology, radio spectral index, polarisation, variability, radio luminosity and flux density ratios with non-radio wavelengths e.g. with different parts of the infrared (IR) regime. We discuss the advantages and limitations of these various diagnostics methods with current and future surveys. However, weeding AGN out of deep radio surveys can already provide several insights into the star formation at high redshift. As well as reproducing the well known rise with redshift in the comoving star formation rate density, we also see evidence for the continued dominance of LIRGs and ULIRGs to the total star forming budget across redshifts $1-3$. Additionally, while we see that the IR-radio relation for star forming galaxies does hold to high redshifts ($z>1$) there is a mild deviation depending on the IR waveband used and the range of IR SEDs found. We will discuss the possible reasons behind this change in properties.}  \\FullConference{Panoramic Radio Astronomy: Wide-field 1-2 GHz research on galaxy evolution\\\\ June 2-5 2009\\\\ Groningen, the Netherlands}  \\begin{document} ",
        "introduction": "Traditionally radio surveys have been dominated by powerful radio-loud active galactic nuclei (AGN) and radio galaxies were the first probe of the distant Universe [1], and references therein. However, surveys by the next generation of radio telescopes we will obtain vast samples of star forming galaxies (SFGs). In fact, as studies of star formation, these surveys will have several advantages over current surveys at other wavelengths. Firstly, radio luminosity is a direct tracer of star formation rate (SFR) once AGN activity has been accounted for and secondly, the combination of area and the sensitivity will provide a fast survey speed allowing for an unprecedented census of star formation in the distant Universe.  \\begin{figure} \\includegraphics[width=0.35\\textwidth, angle=-90]{fig1b.eps} \\includegraphics[width=0.35\\textwidth, angle=-90]{fig1a.eps} \\caption{{\\bf Left} The $1.4\\,$GHz Euclidean normalised source counts from the $13^{\\rm H}$ field divided into SFGs and AGN [2]. Models of the contribution from different populations are overlaid: S04 [3], J08 [4]. {\\bf Right} Compilation from the literature of the fraction of SFGs as a function of $1.4\\,$GHz flux density [2,5,6,7,8,9].} \\label{fig1} \\end{figure}  The sub-mJy regime of the radio source counts has been studied for several decades now and the up-turn below $\\sim1\\,$mJy has been well characterised albeit with some scatter. Currently the debate is over the normalisation of the source counts and type of sources that are responsible for the up-turn. However, the most important question is how do we distinguish between AGN and SFGs in the deepest surveys?  We believe that as some AGN have very low emission in the radio, compared to even low Milky Way rates of star formation, the simple presence of an AGN does not necessarily infer that the radio emission of a galaxy is from an accretion driven source. Hence, methods to discriminate between whether AGN or star forming activity contribute to the observed radio flux density should be related to the radio emission, and we list some of these methods here:  \\begin{itemize} \\item Radio morphology \\item Radio spectral index/radio Spectral Energy Distribution (SED) \\item Radio variability \\item Radio polarisation \\item Flux density ratios/full SED modeling \\end{itemize}  This list is not exhaustive and several of these methods have already been used in various combinations (e.g. [2,5,6,7,8,9]). We have found that a combination of the radio morphology and spectral index is the most effective method if you have data of sufficient quality, but in practice for the deepest surveys the flux ratio method, based on SEDs of SFGs and AGN, is the most productive. So employing these techniques, and on the assumption that one process dominates we can separate the faint radio source counts by galaxy type (e.g. Fig.~\\ref{fig1} left).  We have demonstrated that while AGN dominate the 1.4\\,GHz source counts around 0.4\\,mJy, SFGs make up $70-80\\%$ below 0.1\\,mJy [2] with the transition between the two populations occuring between $0.1$ and $0.2\\,$mJy. So while SFGs dominate at the faintest flux densities there remains a non-negligible contribution from radio quiet AGN around $50\\,\\mu$Jy. This result is consistent with determinations and constraints in the literature (Fig.~\\ref{fig1} right) which are reaching a consensus about the nature of the faint radio source counts. The observed differences in this plot are easily explained by the different methods used in discriminating between AGN and SFGs as well as sample variance between the different deep narrow survey fields used in each study.  Our knowledge of the faint radio source population will be greatly improved with surveys from up-coming radio facilities (see other contributions to these proceedings). At fainter flux densities the modeled evolution of the luminosity function of both populations suggests that the fraction of SFGs will asymptote towards $100\\%$ [4]. The eVLA and eMERLIN will push to such flux densities where star formation completely dominates and provide samples of radio-selecting SFGs out to the highest redshifts. Other projects, like the Evolutionary Map of the Universe (EMU, see [10] in these proceedings) on the Australian Square Kilometer Array Pathfinder, will probe the star forming regime well below $0.1\\,$mJy over a large fraction of the whole sky. However, in preparation for such surveys we must obtain a better understanding of radio emission from star formation in the local Universe and a better understanding of the radio luminosity to star formation rate (SFR) conversion.  \\begin{figure} \\includegraphics[width=0.35\\textwidth, angle=-90]{fig3a.eps} \\includegraphics[width=0.35\\textwidth, angle=-90]{fig3b.eps} \\caption{{\\bf Left} Star formation rate plotted against redshift for a sample of radio selected SFGs from the $13^{\\rm H}$ field (coloured symbols indicating host galaxies of different stellar masses). Local radio-selected star forming galaxies are illustrated by black dots. Equivalent IR luminosity is plotted on the right axis. At high redshift we see massive galaxies with star formation rates in excess of that seen in the local Universe. {\\bf Right} Derived comoving star formation rate density (SFRD) as a function of redshift compared with the mean SFRD presented in [11]. The results in the radio agree well with those at other wavelengths which bodes well for future radio surveys making detailed studies of star formation at high redshift.} \\label{fig2} \\end{figure}  \\newpage ",
        "conclusions": "Radio observations of the distant Universe have traditionally been used to study AGN, but we are now begining to obtain a census of star formation from deep, wide radio surveys. There are three crucial issues in exploiting such data: (i) distinguishing between AGN and SFGs, (ii) calibrating the radio luminosity/SFR relation across all redshifts, radio luminosities and type of galaxy, and (iii) obtaining redshifts from ancillary data. The radio/IR relation for SFGs appears to depend on IR SED and hence waveband used in the IR. We must better understand this relationship locally before applying it to galaxies at high redshift."
    },
    "0911/0911.3146_arXiv.txt": {
        "abstract": "We present a deep \\textit{Chandra X-ray Observatory} study of the peculiar binary radio millisecond pulsar PSR J1740--5340 and candidate millisecond pulsars (MSPs) in the globular cluster NGC 6397.  The X-rays from PSR J1740--5340 appear to be non-thermal and exhibit variability at the binary period.  These properties suggest the presence of a relativistic intrabinary shock formed due to interaction of a relativistic rotation-powered pulsar wind and outflow from the unusual ``red-straggler/sub-subgiant'' companion. We find the X-ray source U18 to show similar X-ray and optical properties to those of PSR J1740--5340, making it a strong MSP candidate.  It exhibits variability on timescales from hours to years, also consistent with an intrabinary shock origin of its X-ray emission. The unprecedented depth of the X-ray data allows us to conduct a complete census of MSPs in NGC 6397. Based on the properties of the present sample of X-ray--detected MSPs in the Galaxy we find that NGC 6397 probably hosts no more than 6 MSPs. ",
        "introduction": "Globular clusters are well known for their abundance of rotation-powered millisecond pulsars \\citep{Camilo00, Freire03, Ran05}\\footnote{For an up-to-date list of all known globular cluster MSPs see http://www.naic.edu/~pfreire/GCpsr.html.}.  These objects are believed to have been produced by the evolution of low-mass X-ray binaries (LMXBs) that were themselves produced through dynamical interactions of binaries and neutron stars.  The superb sub-arcsecond X-ray imaging capability of the \\textit{Chandra X-Ray Observatory} has allowed X-ray studies of MSPs in the dense cores of globular clusters for the first time.  In addition to establishing the X-ray properties of the Galactic population of MSPs, X-ray studies of cluster MSPs offer constraints on stellar and binary evolution and the internal dynamical evolution of globular clusters.  To date, X-ray counterparts of these systems have been detected with \\textit{Chandra} in 47 Tuc \\citep{Grind02,Bog06}, NGC 6397 \\citep[][and this work]{Grind02}, M28 \\citep[][S.~Bogdanov et al. in preparation]{Beck03}, M4 \\citep{Bass04a}, NGC 6752 \\citep{DAmico02}, M71 \\citep{Els08}, and Ter 5 \\citep{Hein06}.  NGC 6397 is the nearest apparently core-collapsed and the second closest globular cluster, with $D\\approx2.4$ kpc \\citep[][and references therein]{Han07,Strick09}.  It hosts one known MSP, PSR J1740--5340 \\citep{DAmico01}, which is bound to an unusual ``red-straggler'' or ``sub-subgiant''\\footnote{The terms red straggler and sub-subgiant are used for cluster stars that lie in a location below the base of the red giant branch that is difficult to understand in the context of standard evolution scenarios for single stars or binaries in a dynamically inactive environment \\citep{Oro03}.} companion in a 32.5-hour binary orbit \\citep{Ferr03}. This is at odds with the standard MSP formation theory that predicts either a white dwarf or a very-low-mass ($\\sim$0.03 M$_{\\odot}$) degenerate companion \\citep[see][for a review]{Bhatt91}, implying that either this binary has just emerged from the ``recycling'' phase in a LMXB \\citep{Bur02} or has been involved in a close dynamical binary-binary encounter in which the original companion to the pulsar was exchanged for the current one \\citep[][and references therein]{Cam05}.   Observations of MSPs, both in the field of the Galaxy and in globular clusters, have revealed that their X-ray emission can be of thermal and/or non-thermal character \\citep[see][for a review]{Zavlin07,Bog08a}.  In many MSPs, the X-rays appear to be predominantly due to surface emission from the magnetic polar caps of the neutron star \\citep{Bog06,Zavlin06,Bog09}. Non-thermal pulsed X-ray emission, on the other hand, can arise from particle acceleration processes in the pulsar magnetosphere as seen from the most energetic MSPs \\citep[see, e.g.][]{Rut04}. Alternatively, non-thermal X-rays can be produced via interaction of the rotation-powered pulsar wind with the ambient interstellar medium, as seen for PSR B1957+20 \\citep{Stap03} and J2124--3358 \\citep{Hui06}, or material from a close binary companion as in PSR J0024--7204W in 47 Tuc \\citep{Bog05}.  \\begin{figure*}[!t] \\begin{center} \\includegraphics[width=0.85\\textwidth]{f1.ps} \\end{center} \\caption{Coadded \\textit{Chandra X-ray Observatory} 295.1-ks ACIS-S3 image of the core of the globular cluster NGC 6397 in the 0.3--6 keV band. The $1.5\\arcsec$ circles are centered on the 79 X-ray sources detected within the $2.3\\arcmin$ half-mass radius of the cluster (outer dashed circle), with the positions of PSR J1740--5340 (U12) and CXOGlb J174041.6--534027 (U18) marked. The inner dashed circle shows the $5.5\\arcsec$ core radius of the cluster. The grayscale corresponds to number of counts increasing logarithmically from 0 (\\textit{white}) to 3515 (\\textit{black}). North is up and east is to the left.} \\end{figure*}  In this paper, we present \\textit{Chandra} deep imaging spectroscopic observations of NGC 6397, with a particular focus on PSR J1740--5340 and CXOGlb J174041.6--534027 (U18), a plausible MSP candidate. We also investigate the X-ray source population of NGC 6397 in an attempt to identify plausible MSP candidates and constrain the MSP content of this cluster.  The work is organized as follows. In \\S 2 we describe the data reduction and analysis procedures. In \\S 3 we focus on the properties of PSR J1740--5340, while in \\S 4 we investigate the MSP candidate U18. In \\S5 we attempt to place interesting constraints on the total number of MSPs in the cluster based on the available X-ray and optical data. We present a discussion in \\S6 and conclusions in \\S7. ",
        "conclusions": "We have presented \\textit{Chandra X-ray Observatory} ACIS-S observations of the nearby core-collapsed globular cluster NGC 6397. The depth of the available data has allowed interesting insight into the MSP population of this cluster.  For instance, it has provided a more complete multi-wavelength picture of the PSR J1740--5340 binary. This peculiar system appears to exhibit predominantly non-thermal emission modulated at the orbital phase. Given the pulsar energetics, the variable non-thermal X-ray emission is indicative of interaction between the relativistic pulsar wind and material from the companion, with evidence for an occultation of the resulting shock by the secondary star. Ascertaining the detailed geometry of the intrabinary shock would require much more sensitive orbital-phase resolved spectroscopic X-ray and optical observations of this system than presently available.  Based on its X-ray and optical properties, U18 is a strong candidate for a MSP, though one that may be perpetually eclipsed at radio frequencies. As noted by \\citet{Tav91} such MSPs may be relatively common both in globular clusters and in the field of Galaxy. Moreover, these sources may be mis-classified as LMXBs or CVs based on their X-ray to optical flux ratio alone \\citep[see in particular][]{Homer06,Arch09}.  Further detailed multi-wavelength observations of PSR J1740--5340, U18, and similar systems may reveal more information concerning the physics of the shock, which, may, in principle offer insight into the properties of MSP winds, collisionless relativistic shocks, and particle acceleration mechanisms.  Moreover, this and similar MSP binaries provide important clues about the behavior of accreting neutron stars transitioning to rotation power, which are believed to also contain an active radio pulsar and an overflowing companion \\citep{Camp04,Hein07,Hein09}.  The deep X-ray observations of NGC 6397 supplemented by optical observations, also provide insight into the MSP population of this cluster. Specifically, if the X-ray luminosity function of the NGC 6397 MSP population resembles that of the Galactic sample of MSPs, then this cluster contains no more than $\\sim$6 MSPs. This range of values is generally consistent with those inferred based on scaling of stellar encounter rates in globular clusters."
    },
    "0911/0911.5283_arXiv.txt": {
        "abstract": "{} {We approach the study of Saturn and its environment in a novel way using X-ray data, by making a systematic and uniform spectral analysis of all the X-ray observations of the planet to date. } {We present the results of the two most recent (2005) XMM-Newton observations of Saturn together with the re-analysis of an earlier (2002) observation from the XMM-Newton archive and of three Chandra observations in 2003 and 2004. While the XMM-Newton telescope resolution does not enable us to resolve spatially the contributions of the planet's disk and rings to the X-ray flux, we can estimate their strengths and their evolution over the years from spectral analysis, and compare them with those observed with Chandra.} {The spectrum of the X-ray emission is well fitted by an optically thin coronal model with an average temperature of 0.5 keV. The addition of a fluorescent oxygen emission line at $\\sim$0.53 keV improves the fits significantly. In accordance with earlier reports, we interpret the coronal component as emission from the planetary disk, produced by the scattering of solar X-rays in Saturn's upper atmosphere, and the line as originating from the Saturnian rings. The strength of the disk X-ray emission is seen to decrease over the period 2002 -- 2005, following the decay of solar activity towards the current minimum in the solar cycle. By comparing the relative fluxes of the disk X-ray emission and the oxygen line, we suggest that the line strength does not vary over the years in the same fashion as the disk flux. We consider possible alternatives for the origin of the line. The connection between solar activity and the strength of Saturn's disk X-ray emission is investigated and compared with that of Jupiter. We also discuss the apparent lack of X-ray aurorae on Saturn; by comparing the planet's parameters relevant to aurora production with those of Jupiter we conclude that Saturnian X-ray aurorae are likely to have gone undetected because they are below the sensitivity threshold of current Earth-bound observatories. A similar comparison for Uranus and Neptune leads to the same disappointing conclusion, which is likely to hold true also with the planned next generation International X-ray Observatory. The next step in advancing this research can only be realised with in-situ X-ray observations at the planets.} {} ",
        "introduction": "In recent years X-ray observations have provided a novel way of furthering the study of planets in our solar system and our understanding of the physical processes taking place on them and in their environments (for a review see Bhardwaj et al. \\cite{Bhardwaj_07} and references therein). For planets with a significant magnetic field, namely Jupiter and the Earth, the X-ray emission has been found to separate essentially into two components: a high-latitude auroral one, and a low-latitude disk component, which has different spectral properties, relating it to solar X-rays (Maurellis et al. \\cite{Mau_00}, Gladstone et al. \\cite{Glad_02}, Branduardi-Raymont et al. \\cite{BR_04}, Bhardwaj et al. \\cite{Bhardwaj_solcon}, Elsner et al. \\cite{El_05}, Branduardi-Raymont et al. 2007a,b, Bhardwaj et al. \\cite{Bhardwaj_07}). Naively we would expect to observe a similar dichotomy on Saturn, given the strength of its magnetic field and its fast rotation (as for Jupiter) and its powerful UV aurorae (G\\'{e}rard et al. \\cite{Gerard_04}, G\\'{e}rard et al. \\cite{Gerard_05}), but this is not the case. While the characteristics of Jupiter's UV aurorae are dictated by the rotation of plasma internal to its magnetosphere (Bhardwaj and Gladstone \\cite{BG_00}, Clarke et al. \\cite{Clarke_04}), those of Earth are solar wind driven (Clarke et al. \\cite{Clarke_05}); Saturn's UV aurorae appear to be at variance with both these scenarios, rather than being intermediate between the two, as originally expected (Bhardwaj and Gladstone \\cite{BG_00}). In Saturn's case the UV aurora responds strongly to the solar wind dynamic pressure (Crary et al. \\cite{Crary_05}), whose compression of the magnetosphere may induce reconnection leading to auroral brightenings (Cowley et al. \\cite{Cowley_05}); compared to Earth, though, where aurora brightenings are very rapid ($\\sim$ tens of minutes), those on Saturn can last for days (Clarke et al. \\cite{Clarke_05}). Moreover, UV auroral activity has also been observed on Saturn during a period (Oct. - Nov. 2005) of quiet solar wind conditions, suggesting an intrinsically dynamic magnetosphere with injections of plasma occasionally taking place in the night or dawn sectors (G\\'{e}rard et al. \\cite{Gerard_06}). Recent work by Clarke et al. (\\cite{Clarke_09}) has added to this picture: from an extensive campaign of observations by the Hubble Space Telescope (HST), coupled with planetary spacecraft and solar wind data, they find evidence for a direct relationship between Saturn's auroral activity and solar wind conditions, while the correlation is not so strong at Jupiter.  \\begin{table*} \\caption{{\\it XMM-Newton} and {\\it Chandra} observations of Saturn} \\label{table:1} \\centering \\begin{tabular}{c c c c c c l }   % \\hline\\hline & Observation & Observation & Apparent & Heliocentric & Geocentric & \\\\ Satellite & mid-epoch & duration & diameter & distance & distance & References \\\\ &  &  (ks)  &  &  (AU)  &  (AU) &  \\\\ \\hline \\\\ {\\it XMM-Newton} (1) & 2002 Oct. 01, 14:00 UT & 21 & 18.8'' & 9.0 & 8.8 & Ness et al. \\cite{Ness_x}, this paper \\\\ {\\it Chandra} (1) & 2003 Apr. 14, 18:00 UT & 66 & 17.5'' & 9.0 & 9.5 & Ness et al. \\cite{Ness_c} \\\\ {\\it Chandra} (2) & 2004 Jan. 20, 05:30 UT & 37 & 20.5'' & 9.0 & 8.1 & Bhardwaj et al. \\cite{Bhardwaj_05a}, \\cite{Bhardwaj_05b}, \\\\ &  &    &  &    &   &  this paper \\\\ {\\it Chandra} (3) & 2004 Jan. 26, 20:00 UT & 36 & 20.4'' & 9.0 & 8.2 & Bhardwaj et al. \\cite{Bhardwaj_05a}, \\cite{Bhardwaj_05b} \\\\ {\\it XMM-Newton} (2) & 2005 Apr. 22, 10:00 UT & 84 & 18.1'' & 9.1 & 9.2 & This paper  \\\\ {\\it XMM-Newton} (3) & 2005 Oct. 28, 21:00 UT & 74 & 18.2'' & 9.1 & 9.1 & This paper  \\\\ \\\\ \\hline \\end{tabular} \\end{table*}  Early observations of Saturn with {\\it XMM-Newton} and {\\it Chandra} (Ness et al. \\cite{Ness_x}, Ness et al. \\cite{Ness_c}) unambiguously detected X-ray emission, found it to be concentrated in the equatorial regions of the planet, and indicated that any X-ray auroral component, if present at all, is not easily detectable. More recent {\\it Chandra} observations have provided us with a remarkable view of the X-ray morphology of the Saturnian system, and of its temporal behaviour in the X-rays: both the planet's disk and the rings have been separately detected, and strong variability, closely following the behaviour of the solar X-ray flux, has been discovered in the disk X-ray emission (Bhardwaj et al. \\cite{Bhardwaj_05a}, Bhardwaj et al. \\cite{Bhardwaj_05b}). This suggests that its origin may lie in the scattering of solar X-rays in the upper layers of Saturn's atmosphere, as is the case for Jupiter's disk emission (Bhardwaj et al. \\cite{Bhardwaj_solcon}, Bhardwaj et al. \\cite{Betal_06}, Cravens et al. \\cite{Cravens_06}). Unlike Jupiter, though, no difference is found in the spectral characteristics of Saturn's disk and polar emissions, thus no evidence for X-rays of auroral origin is provided by the {\\it Chandra} data. In fact, Bhardwaj et al. (\\cite{Bhardwaj_05a}) remark that the X-ray emission from the polar regions appears to be anticorrelated with the strength of the FUV aurora, measured by the HST Imaging Spectrograph simultaneously with the {\\it Chandra} observations. Moreover, only relatively little detail has been published so far about Saturn's X-ray spectral characteristics: an optically thin, coronal model is a good match for the disk emission (Ness et al. \\cite{Ness_x}, Ness et al. \\cite{Ness_c}, Bhardwaj \\cite{Bhardwaj_06}), while the rings spectrum is practically made up of a single O K$\\alpha$ emission line, which has been attributed largely to scattering of solar X-rays on the icy water ring material (Bhardwaj et al. \\cite{Bhardwaj_05b}). However, solar fluorescence alone could not explain the ring X-ray brightness completely. To follow up on these issues, we made more extensive {\\it XMM-Newton} observations of Saturn in 2005: this was a particularly appropriate time to take a deeper look at the planet, while a campaign of HST observations was taking place, and the {\\it Cassini} spacecraft (which entered orbit around Saturn in July 2004) was making coordinated in-situ measurements of a variety of physical parameters in its environment.  In this paper we report the results from our 2005 observations, together with a re-analysis of the earlier (2002) data extracted from the {\\it XMM-Newton} archive (sec.s 2 and 3). We also model the {\\it Chandra} X-ray spectra from the 2003 and 2004 observations (which have already been described by Ness et al. \\cite{Ness_c}, Bhardwaj et al. \\cite{Bhardwaj_05a} and Bhardwaj et al. \\cite{Bhardwaj_05b} $-$ see sec. 3). We discuss our findings in sec. 4, also making comparisons with Jupiter and Earth, as well as Uranus and Neptune, and draw our conclusions in sec. 5. ",
        "conclusions": "\\begin{enumerate} \\item Statement 1. We report the results of all the X-ray observations of Saturn carried out to date. We present a detailed spectral analysis of the available data which reinforces the earlier conclusion (Bhardwaj et al. \\cite{Bhardwaj_05a}) that the planet's disk acts as a reflector of solar X-rays, much in the same way as observed for Jupiter, Earth, Venus and Mars (cf. Bhardwaj et al. \\cite{Bhardwaj_07}). \\item Statement 2. Our results also suggest that the strength of the X-ray emission from the Saturnian rings, in the form of a fluorescence oxygen line, does not systematically decrease, like that from the disk, while the solar activity cycle evolves towards its minimum. As a possible alternative to scattering of solar X-rays in the rings icy material, we suggest that the origin of the line may be related to the formation of the spokes, which appear on the East ansa of the rings as do the bright X-ray spots seen by {\\it Chandra}, by lightning-induced electron beams. \\item Statement 3. Unlike Jupiter and Earth, we do not find evidence for X-ray aurorae on Saturn. This can be explained by the limited sensitivity of the observations to date, if the auroral X-ray and FUV emissions on Saturn have relative strengths similar to those on Jupiter. Other factors, such as magnetic field strength, magnetospheric particle densities and internal plasma sources may also contribute. \\item Statement 4. After comparison with Jupiter and Earth, we conclude that Saturn is unique in its X-ray characteristics as it is in the FUV (Clarke et al. 2005). We still have much to learn about the precise mechanisms operating in its aurorae, and in its rings, leading (or not, as it may be) to X-ray production. The effectiveness of further observations with {\\it Chandra} and {\\it XMM-Newton} will increase as we move out of the current (2009) minimum in the solar activity cycle, providing invaluable additional information on the evolution of both, disk and rings emissions, and of their response to the solar X-ray flux and the solar wind. However, a much more efficient way to pursue and enhance these studies, and relate the electromagnetic output of the planet to particle and magnetic field measurements, would involve flying an X-ray spectrometer to operate in-situ, onboard a spacecraft mission to the Saturnian system. We look forward to the possible realisation of this project in association with the missions to the outer planets currently being studied by both ESA and NASA. Such a mission to Saturn is most likely to involve a close encounter with Titan, and an X-ray spectrometer would return unique data on this enigmatic satellite, and its interactions with Saturn and the solar wind as well, especially at times when it is outside Saturn's magnetosphere. In-situ X-ray observations are also the only feasible way to search for auroral X-ray emission from Uranus and Neptune.  \\end{enumerate}"
    },
    "0911/0911.0690_arXiv.txt": {
        "abstract": "We use multi-wavelength, matched aperture, integrated photometry from {\\it GALEX}, SDSS and the RC3 to estimate the physical properties of 166 nearby galaxies hosting 168 well-observed Type Ia supernovae (SNe~Ia).  The ultra-violet (UV) imaging of local SN~Ia hosts from {\\it GALEX} allows a direct comparison with higher redshift hosts measured at optical wavelengths that correspond to the rest-frame UV.  Our data corroborate well-known features that have been seen in other SN~Ia samples. Specifically, hosts with active star formation produce brighter and slower SNe~Ia on average, and hosts with luminosity-weighted ages older than 1 Gyr produce on average more faint, fast and fewer bright, slow SNe~Ia than younger hosts.  New results include that in our sample, the faintest and fastest SNe~Ia occur only in galaxies exceeding a stellar mass threshhold of $\\sim10^{10}$ M$_{\\odot}$, leading us to conclude that their progenitors must arise in populations that are older and/or more metal rich than the general SN~Ia population.  A low host extinction sub-sample hints at a residual trend in peak luminosity with host age, after correcting for light-curve shape, giving the appearance that older hosts produce less-extincted SNe~Ia on average.  This has implications for cosmological fitting of SNe~Ia and suggests that host age could be useful as a parameter in the fitting.  Converting host mass to metallicity and computing $^{56}$Ni mass from the supernova light curves, we find that our local sample is consistent with a model that predicts a shallow trend between stellar metallicity and the $^{56}$Ni mass that powers the explosion, but we cannot rule out the absence of a trend.  We measure a correlation between $^{56}$Ni mass and host age in the local universe that is shallower and not as significant as that seen at higher redshifts.  The details of the age -- $^{56}$Ni mass correlations at low and higher redshift imply a luminosity-weighted age threshhold of $\\sim3$ Gyr for SN~Ia hosts, above which they are less likely to produce SNe~Ia with $^{56}$Ni masses above $\\sim0.5$ M$_{\\odot}$. ",
        "introduction": "}  A fundamental motivation for studying Type Ia supernovae (SNe~Ia) is to arrive at a physical explanation for the observed fact that these objects have well-behaved and calibratable explosions \\citep{Phillips:93:L105}.  This property gives cosmographers a tool for measuring universal expansion and provided the first direct observational evidence for cosmic acceleration \\citep{Riess:98:1009, Perlmutter:99:565}.  Knowing the physical process that generates SNe~Ia provides strong constraints on the nature of the SN~Ia progenitor.  This, in turn, allows us to predict the evolution of the SN~Ia population \\citep[e.g.,][]{Howell:07:L37}.  This results in tighter control of the systematic errors arising from population evolution with redshift, an important uncertainty in measuring cosmological parameters \\citep[e.g.,][]{Astier:06:31,Conley:09}.  A well-constrained physical model for SNe~Ia would have further utility in tracing high-redshift star formation and for predicting the effects of SNe~Ia on the chemical enrichment of their host galaxies \\citep[e.g.,][]{Kobayashi:07:215}.  In order to calibrate SNe~Ia and make them useful for cosmology, their intrinsic variation must be measured and accounted for.  The peak absolute magnitude of a given SN~Ia is a strong function of its initial decline rate \\citep{Phillips:93:L105}, and a weaker function of its peak color \\citep{Riess:96:88, Tripp:98:815, Tripp:99:209, Parodi:00:634}.  These empirical correlations reduce the intrinsic variation of $\\sim$1 magnitude in the B-band to $\\sim0.1$ magnitudes and thus provide the accurate luminosity distance estimates required for measuring cosmological parameters \\citep[e.g.,][]{Riess:96:88, Tonry:03:1, Guy:05:781, Prieto:06:501, Jha:07:122, Conley:08:482}.  The goal of SN~Ia progenitor studies is to use these measures of SN~Ia intrinsic variation to explore trends with host properties that could shed light on the underlying population that produces SNe~Ia \\citep{Phillips:93:L105, Hamuy:95:1, Hamuy:00:1479, Howell:01:L193, Gallagher:05:210, Sullivan:06:868, Gallagher:08:752}.  Since SN~Ia explosions are powered by the radioactive decay of $^{56}$Ni, it is widely accepted that the intrinsic variation in SN~Ia brightness and decline rate is primarily driven by the amount of $^{56}$Ni present in the SN explosion, with more luminous and slower declining explosions being powered by more $^{56}$Ni \\citep{Truran:67:2315, Colgate:69:623}.  One avenue of exploration would be to connect this theoretical idea to the progenitor population of SNe~Ia by comparing host galaxy properties with the measures of SN~Ia light curve variation.  This requires a mechanism that varies the amount of $^{56}$Ni as a function of some property of the host galaxy.  \\citet{Timmes:03:L83} presented a model relating progenitor metallicity to produced SN~Ia $^{56}$Ni mass for a constant (Solar) O/Fe ratio.  This model was expanded and tested by \\citet[hereafter H09]{Howell:09:661} using an intermediate redshift ($0.2 < z < 0.75$) sample of SNe~Ia.  They used spectral energy distribution (SED) fitting of optical integrated host photometry to infer the host stellar masses for their sample.  These masses were used as a proxy for host metallicities through the \\citet{Tremonti:04:898} mass-metallicity relation.  They estimated the peak bolometric luminosity and rise time from the SN~Ia photometry to calculate the required mass of $^{56}$Ni to power the explosion and compared the calculated host metallicity versus Ni mass trend to the \\citet{Timmes:03:L83} model.  Their data are consistent with the model, although with a higher scatter than predicted by the theory.  They also found that varying the O/Fe ratio according to thin or thick-disk models produced no substantial change in their results.  If the dependance of $^{56}$Ni mass on host metallicity is real, then fainter and faster declining SNe~Ia should be associated with higher metallicity, and thus more massive, host galaxies.  If a decline in the rate of fast fading, fainter SNe~Ia with redshift could be detected, it would be evidence for the relation between Ni mass and metallicity because we expect the more distant universe to consist of lower-mass and lower-metallicity objects on average when compared with the local universe. At this time, however, Malmquist biases complicate an accurate determination of the rate evolution for low-luminosity SNe~Ia \\citep{Foley:08,Gaitan:09}.  We therefore look to host properties of low-luminosity SNe~Ia in the local universe, where they are easier to detect and apparently more abundant.  We use this larger sample to see if the fainter SNe do appear in higher mass, higher metallicity galaxies on average.  To compare host properties with SN~Ia variation, we construct a low-redshift sample of SNe~Ia with well observed light curves that provide a light curve shape parameter, called stretch, the observed maximum light color, and the peak luminosity of the SN in the rest-frame B-band \\citep{Conley:08:482}.  We characterize the hosts of these SNe using newly generated and catalog galaxy integrated UV and optical magnitudes measured within matched apertures as inputs to a SED fitting program that estimates host stellar mass, luminosity-weighted age and star formation rate (SFR). We first use these results to explore the relation between the observed light curve parameters and the derived host properties and compare our results to studies at higher redshift using the same methods \\citep[][H09]{Sullivan:06:868}.  We isolate a subset of hosts with low internal extinction in order to explore the relationship between SN peak brightness and host properties where the source of line-of-sight color for the SNe is minimized.  We then extend the work of H09 by applying their techniques to our local sample to derive SN~Ia $^{56}$Ni-mass, from the light-curve photometry, and compare these with host metallicities, from our derived host stellar masses, and host luminosity-weighted ages.  Throughout this paper we express host masses in M$_{\\odot}$ and SFR in M$_{\\odot}$ yr$^{-1}$ and assume a Hubble constant of $H_0 = 73$ km s$^{-1}$ Mpc$^{-1}$. ",
        "conclusions": "}  The implications of these results must be tempered by the limitations of the techniques.  In particular, the SED-fitting method employed was optimized for higher redshift galaxies and we do not employ infra-red data to constrain the SEDs of our low-redshift hosts.  Our aim is to compare hosts at high and low redshifts using the same methods.  We focus here on trends and this comparison and will pursue more precise determinations of host parameters with infra-red data and more detailed methods in a future paper.  We reiterate that the ages estimated with the PEGASE.2 fits are luminosity weighted and the masses are current stellar masses.  The star formation rates only measure current activity based on UV and optical data and have not been modified by FIR data to account for dust.  The global host extinctions are derived by extincting the entire SED and do not account for variable thermal components in the IR.  \\subsection{Host Properties at High and Low Redshift}  Our results show features similar to those presented in higher redshift host studies using the same methods \\citep[][H09]{Sullivan:06:868, Sullivan:09:539}.  Low stretch SNe~Ia appear to be rare in high sSFR hosts and the average stretch of SNe~Ia in the lowest sSFR hosts is smaller than in hosts with higher sSFR (see Figure~\\ref{fig_str_ssfr}).  The average stretch of SNe in hosts younger than $\\sim$2 Gyr is close to 1 and for older hosts the average stretch appears to drop below 1 (see Figure~\\ref{fig_str_age}).  H09 were the first to plot host stellar mass against stretch, but their sample does not include SNe~Ia with $s < 0.75$ (see their Figure 4).  The interesting feature in our plot (Figure~\\ref{fig_str_mass}) is the limited mass range for lower-stretch SNe.  This feature is present in the H09 figure, but it is less pronounced due to having fewer low-stretch SNe in their sample.  The apparent stellar mass threshhold of $10^{10}$ M$_{\\odot}$ for the hosts of low-luminosity SNe~Ia is fairly robust against selection effects.  Our sample includes many lower-mass hosts within which it is easier to detect low-luminosity SNe.  That these SNe prefer higher-mass hosts supports the idea that they arise from an older progenitor population. Figures~\\ref{fig_str_ssfr}, \\ref{fig_str_age} and \\ref{fig_str_mass} show that the one high sSFR host (E445-G66) of a SNe~Ia with $s < 0.80$ (SN1993H) is also massive, implying that the recently formed populations are mixed with the older populations that could be responsible for the low-stretch SN~Ia.  The correlation between mass and metalicity implies that metallicity could play a role as well (see \\S\\ref{sec_metal_ni}).  We see no low-mass, high sSFR hosts of low-stretch SNe~Ia and we also see few older, higher-mass hosts of the highest stretch SNe~Ia.  A method that can more accurately asses the star formation history of these hosts and the relative contributions from older and younger populations will produce a more definitive constraint on the progenitors of the low-luminosity SNe~Ia \\citep{Schawinski:09}.  Our plot of SN peak color versus host age illustrates the difficulty of disentangling the source of the color (see Figure~\\ref{fig_clr_age}) since it shows no correlation.  We are also challenged in our comparison with higher redshift samples by the bias against SNe with very red peak colors intrinsic to cosmology surveys (see \\S\\ref{sec_sample}).  Even in the local universe, our sample of six SNe~Ia with $\\mathcal{C}>0.7$ is too small for definitive interpretation.  We point out one common feature of our figure and Figure~3 in \\citet{Sullivan:09b}: very red SNe appear to be rare in hosts less than 0.5 Gyr in luminosity-weighted age.  Based on Figure~\\ref{fig_ni_metal}, we conclude that the model of \\citet{Timmes:03:L83}, relating metallicity to produced $^{56}$Ni mass, appears consistent with the derived host metallicities and SN~Ia $^{56}$Ni masses for SNe in the local universe as well as at higher redshifts (H09), although the local data are also consistent with no trend in these properties.  The degeneracy between age and metallicity makes it difficult to decide which is the more important factor in determining the $^{56}$Ni mass of hosted SNe~Ia.  The apparent transition at a luminosity-weighted host age of 3 Gyr (see Figure~\\ref{fig_ni_age} and H09, Figure 8) could place interesting constraints on progenitor scenarios given that this implies a fairly long delay time for low $^{56}$Ni mass, hence low-luminosity and low-stretch SNe~Ia.  \\subsection{Host Properties and Cosmological Fitting}  It is encouraging that the features in plots of light curve properties against host properties in local hosts are similar when compared with the SNLS sample out to $z=0.75$ \\citep{Astier:06:31, Howell:07:L37, Sullivan:09b}, although a detailed comparison using Hubble diagram residuals will be needed to determine if any significant evolution of the population can be detected \\citep{Conley:09, Sullivan:09b}.  These trends do signal potential systematics for SN cosmology: as we extend the samples beyond redshift $z=0.75$, we know the mix of host populations will move toward lower-mass, higher sSFR galaxies.  Our examination of low-extinction hosts suggests a trend in brightness with host age such that SNe~Ia in older hosts appear brighter after correcting for stretch and color (see Figure~\\ref{fig_age_mb}).  While the correlation we see is hardly significant (2.1$\\sigma$), it does suggest that luminosity-weighted host age might be an interesting parameter to explore in cosmological fits of SNe~Ia.  A larger sample of SNe~Ia from low-extinction hosts would allow a more definitive test of this correlation.  While there is no apparent trend in a similar plot in \\citet[see their Figure~9]{Gallagher:08:752}, we see a trend that is marginally significant. There are several possible explanations for this difference.  Since the hosts in the \\citet{Gallagher:08:752} sample were selected morphologically, their sample may not be as free from dust extinction as originally thought. In addition, their corrections for light curve shape are derived with a method that assumes a Milky Way-like dust is responsible for residual color after a quadratic color-decline relation is subtracted \\citep[see comparison of fitting methods in][]{Conley:08:482}, which may not be appropriate for these hosts.  The trend we see in Figure~\\ref{fig_age_mb} is in the right sense to be caused by residual line-of-sight extinction since we would expect the SNe hosted by the oldest galaxies to have the least non-local extinction.  These data suggest that either the SNe~Ia or the dust produced in galaxies of different host properties are different or both.  \\citet{Hicken:09a} showed a weak trend in the Hubble residuals when using samples divided into three host morphology bins, after removing the reddest SNe.  An exploration based on more physical host properties could provide a more accurate assessment of this effect \\citep{Sullivan:09b}.  Further progress may be made by extending host photometry farther into the infra-red to help constrain the host dust content and by more detailed fitting of host star-formation history to constrain the most recent episode of star formation \\citep[e.g.,][]{Schawinski:09}."
    },
    "0911/0911.0003_arXiv.txt": {
        "abstract": "We study the properties of a sample of 211 heavily-obscured Active Galactic Nucleus (AGN) candidates in the Extended Chandra Deep Field-South selecting objects with $f_{24\\mu m}$/$f_R$$>$1000 and $R$-$K$$>$4.5. Of these, 18 were detected in X-rays and found to be obscured AGN with neutral hydrogen column densities of $\\sim$10$^{23}$~cm$^{-2}$. In the X-ray undetected sample, the following evidence suggests a large fraction of heavily-obscured (Compton Thick) AGN: (i) The stacked X-ray signal of the sample is strong, with an observed ratio of soft to hard X-ray counts consistent with a population of $\\sim$90\\% heavily obscured AGN combined with 10\\% star-forming galaxies. (ii) The X-ray to mid-IR ratios for these sources are significantly larger than that of star-forming galaxies and $\\sim$2 orders of magnitude smaller than for the general AGN population, suggesting column densities of $N_H$$\\gtrsim$5$\\times$10$^{24}$~cm$^{-2}$. (iii) The Spitzer near- and mid-IR colors of these sources are consistent with those of the X-ray-detected sample if the effects of dust self-absorption are considered. Spectral fitting to the rest-frame UV/optical light (dominated by the host galaxy) returns stellar masses of $\\sim$10$^{11}$M$_\\odot$ and $<$E(B-V)$>$=0.5, and reveals evidence for a significant young stellar population, indicating that these sources are experiencing considerable star-formation. This sample of heavily-obscured AGN candidates implies a space density at $z$$\\sim$2 of $\\sim$10$^{-5}$Mpc$^{-3}$, finding a strong evolution in the number of  $L_X$$>$10$^{44}$~erg/s sources from $z$=1.5 to 2.5, possibly consistent with a short-lived heavily-obscured phase before an unobscured quasar is visible. ",
        "introduction": "Understanding the processes of galaxy formation and evolution is one of the outstanding problems of modern astronomy. It is now clear that the growth and feeding of the central black hole of an active galactic nucleus (AGN) and the formation of its host galaxy must be intimately connected, as evidenced by the striking correlations between the properties of supermassive black holes and the stellar systems in which they reside \\citep{magorrian98, gebhardt00, ferrarese01}. Furthermore, it is now strongly suspected that AGN activity is critical in the regulation of star formation in the host galaxy either by removing all the gas (e.g., \\citealp{hopkins06,menci06}), by heating it (e.g. \\citealp{croton06}) or both \\citep{schawinski06}. However, an observational link between AGN activity and quenching of star formation has not yet been established (e.g., \\citealp{schawinski09}). To understand how the AGN and galaxy formation processes are connected requires a detailed census of the AGN population as well as detailed study of individual objects.  The visibility and detectability of an AGN is determined by the interplay among black hole mass, accretion rate, gas distribution, dust geometry and the amount of star formation in the host galaxy. Thus, no single selection method provides a complete view of how galaxies and AGN co-evolve. Many surveys (e.g., \\citealp{barger01,cowie02}; see \\citealp{brandt05} for a review) have focused on X-ray selection because X-ray emission is a universal feature of AGN accretion. However, X-ray selection has limited sensitivity to heavily obscured and/or low accretion rate sources, even in the deepest surveys \\citep{treister04}. In particular, examples of the most heavily-obscured AGN, the so-called Compton-Thick (CT) AGN, which have $\\tau$$\\sim$1 for Compton scattering or $N_H$$\\sim$10$^{24}$~cm$^{-2}$, are mostly found in the local Universe.  Only a few of these sources have been identified at higher redshifts, although this is simply due to selection effects \\citep{dellaceca08b}. In contrast, early AGN population synthesis models that can fit the X-ray background (XRB) spectral shape and intensity \\citep{treister05b,gilli07} require $\\sim$20-30\\% CT AGN in order to explain the observed peak at $\\sim$30~keV in the XRB radiation. Recently, \\citet{treister09b} combined observations of local ($z$$<$0.05) AGN at very hard X-rays (E$>$20~keV), using the International Gamma-Ray Astrophysics Laboratory (INTEGRAL) and Swift satellites, with XRB models and concluded that the CT AGN fraction is only $\\sim$10\\% at these redshifts. They conclude that CT AGN represent a similarly small fraction of the total XRB radiation and total black hole growth. However, because only $\\sim$2\\% of the XRB comes from CT AGN at $z$$>$2 and only a few sources are known at such high redshifts, this population still remains basically unconstrained, with up to an order of magnitude uncertainty in number density \\citep{treister09b}.  Because the energy absorbed at optical to X-ray wavelengths is later re-emitted in the mid-IR, it is expected that AGN, in particular the most obscured ones, should be very bright mid-IR sources (\\citealp{treister06a} and references therein), easily detectable in observations with the Spitzer observatory. Sources having mid-IR excesses, relative to their rest-frame optical and UV emission, have been identified as potential CT AGN candidates at $z$$\\sim$2 (\\citealp{daddi07,fiore08,georgantopoulos08} and others). However, because of the strong connection between vigorous star formation and AGN activity in the most luminous infrared sources \\citep{sanders88}, the relative contribution of these processes is still uncertain and remains controversial (e.g., \\citealp{donley08,pope08}).  In order to detect CT AGN up to high  redshifts, deep multiwavelength coverage is critical. The Extended Chandra Deep Field South (ECDF-S) is one of the best fields to carry out this study. The rich multiwavelength data available in the ECDF-S includes four $\\sim$250~ksec Chandra pointings covering an area of $\\sim$0.3~deg$^2$, and deep Spitzer data in the IRAC bands from the Spitzer IRAC/MUSYC Public Legacy in ECDF-S (SIMPLE) survey and at 24~$\\mu$m from the Far-Infrared Deep Extragalactic Legacy Survey (FIDEL). The 24~$\\mu$m depth in this field is $\\sim$35~$\\mu$Jy, comparable to the deep Great Observatories Origins Deep Survey (GOODS) observations.  In this paper, we present a study of the properties of the CT AGN candidates in the ECDF-S selected from their 24~$\\mu$m to optical flux ratio and rest-frame optical to UV colors. In Section 2 we describe the multiwavelength data used in this work, while our selection scheme is presented in Section 3. The properties of the sources individually detected in X-rays are discussed in Section 4, while the remaining sources are studied in Section 5. The near and mid IR colors, X-ray to mid-IR flux ratios and space density of our sources are presented in Section 6, and our conclusions are reported in Section 7. We assume a $\\Lambda$CDM cosmology with $h_0$=0.7, $\\Omega_m$=0.3 and $\\Omega_\\Lambda$=0.7, in agreement with the most recent cosmological observations \\citep{spergel07}. ",
        "conclusions": "We presented a study of heavily-obscured (CT) AGN candidates  selected from their optical, near and mid-IR colors. These sources were selected by requiring $f_{24\\mu m}$/$f_R$$>$1000 and $R$-$K$$>$4.5. Using this selection method, we found a total of 211 infrared-red excess sources in the ECDF-S. Of these, 18 were individually-detected in the Chandra 250-ksec observations of this field.  These sources are moderately-obscured, with $N_H$=10$^{22}$-10$^{23}$~cm$^{-2}$, according to their X-ray spectra. Two of these sources have $N_H$$>$10$^{-24}$~cm$^{-2}$, and therefore are classified as CT AGN, including one at $z$=4.65. The average redshift for the X-ray-detected sample is $<$$z$$>$$\\sim$2.37 (median=1.85), in agreement with previous studies. By performing spectral fitting to the observed optical to near-IR fluxes of the X-ray-detected sources we found that the host galaxies have typical stellar masses of 10$^{11.7}$M$_\\odot$ and are subject to moderate extinction, with $<$E(B-V)$>$=0.6.  For the X-ray undetected sources, we found a slightly lower average redshift of $<$$z$$>$$\\sim$1.9, based on photometric measurements only. We found that in general the X-ray-detected sources are brighter in both the 24~$\\mu$m and optical bands compared to the undetected ones. Taking advantage of the low Chandra background, we performed X-ray stacking, finding a significant ($\\gtrsim$3$\\sigma$) detection in both the soft and hard bands. The average X-ray flux of these sources is $\\sim$8$\\times$10$^{-17}$~erg~cm$^{-2}$s$^{-1}$, indicating that in order to detected these sources individually, exposure times of $\\sim$10 Msec with Chandra would be required. Taken at face value, the observed average HR of 0.13 corresponds to a $N_H$ of $\\sim$2$\\times$10$^{23}$~cm$^{-2}$. However, this can also be interpreted as the effects of combining a population of 90\\% CT AGN and 10\\% star-forming galaxies, which in general have softer X-ray spectra. From this analysis we found marginal evidence for a small dependence of the fraction of AGN relative to star-forming galaxies on 24~$\\mu$m flux, going from $\\sim$95\\% for the brightest sources to $\\sim$80\\% at the lowest flux bin. The spectral fitting analysis performed to these sources indicate that in general there is evidence for substantial young stellar populations, younger than 100 Myrs. This suggests that these sources are simultaneously experimenting significant star-formation and heavily-obscured AGN activity. We did not found significant differences in the stellar masses and extinction values for the host galaxies of the X-ray detected and undetected IR-red excess sources.  The X-ray undetected IR-red excess sources have redder $f_{24}$/$f_8$ flux ratios than those detected in X-rays. The average $f_{24}$/$f_8$$\\simeq$25 is hard to explain assuming the observed spectra of ultra-luminous infrared galaxies, however it can be well-explained by dust re-radiation models in which the effects of self-absorption are important. Similarly, while the IRAC colors of most of these sources are outside the typical AGN region, this can be explained by the effects of absorption and/or a dominance of the near-IR emission by the host galaxy. In contrast, the X-ray detected sources have IRAC colors consistent with those of less-obscured AGN. The X-ray to mid-IR ratio for the X-ray-detected sources is similar to the average value for the overall X-ray population. For the X-ray-undetected sources this ratio is $\\sim$2 orders of magnitude lower than for the sample individually detected in X-rays. This ratio is consistent with obscuration levels of $\\sim$5$\\times$10$^{24}$ to 10$^{25}$~cm$^{-2}$, while it is significantly larger than the relative X-ray emission expected from star-formation activity alone. Using a constant threshold in X-ray to mid-IR flux ratio, we found that the fraction of AGN in our sample should be  $\\sim$80\\%, consistent with the value found from the analysis of the stacked X-ray signal. We also found a similar dependence of the AGN fraction on 24~$\\mu$m luminosity ranging from $\\sim$90\\% at the bright end to only $\\sim$10\\% for the faintest sources.  Finally, we studied the space density of sources implied by our sample of 24~$\\mu$m-selected heavily-obscured AGN candidates. While at lower redshifts and luminosities we found a good agreement with population synthesis models and extrapolations of existing X-ray luminosity functions to higher obscuration, we found significant excesses for high-luminosity sources at high redshifts. Furthermore, we found a strong evolution in the number of sources with $L_X$$>$10$^{44}$~erg~s$^{-1}$ from $z$=1.5 to 2.5. This can be interpreted as evidence for a relatively short-lived heavily-obscured phase before the strong radiation pressure removes the surrounding dust and turns the sources into an unobscured quasar."
    },
    "0911/0911.2006_arXiv.txt": {
        "abstract": "We present new photometric and spectroscopic observations for 2M\\,1533+3759 (=~NSVS\\,07826147), the seventh eclipsing subdwarf B star + M dwarf (sdB+dM) binary ever found.  It has an orbital period of 0.16177042 day, or $\\sim$3.88~h, significantly longer than the 2.3--3.0 hour periods of the other known eclipsing sdB+dM systems. Spectroscopic analysis of the hot primary yields $T_{\\rm eff} = 29230 \\pm 125$~K, $\\log g = 5.58 \\pm 0.03$ and log $N$(He)/$N$(H) = $-2.37 \\pm 0.05$.  The sdB velocity amplitude is $K_{\\rm 1} = 71.1 \\pm 1.0$~km~s$^{-1}$.  The only detectable light contribution from the secondary is due to the surprisingly strong reflection effect, whose peak-to-peak $BVRI$ amplitudes are 0.10, 0.13, 0.15, and 0.19 magnitudes, respectively.  Light curve modeling produced several solutions corresponding to different values of the system mass ratio, $q$ ($M_{\\rm 2}$/$M_{\\rm 1}$), but only one is consistent with a core helium burning star, $q$ = 0.301.  The orbital inclination is $86.6\\degr$.  The sdB primary mass is $M_{1} = 0.376 \\pm 0.055~M_{\\sun}$ and its radius is $R_{1} = 0.166 \\pm 0.007~R_{\\sun}$. 2M\\,1533+3759 joins PG\\,0911+456 (and possibly also HS\\,2333+3927) in having an unusually low mass for an sdB star.  SdB stars with masses significantly lower than the canonical value of $0.48~M_{\\sun}$, down to as low as $0.30~M_{\\sun}$, were theoretically predicted by Han et al.\\ (2002, 2003), but observational evidence has only recently begun to confirm the existence of such stars.  The existence of core helium burning stars with masses lower than 0.40--0.43~$M_{\\sun}$ implies that at least some sdB progenitors have initial main sequence masses of 1.8--2.0$~M_{\\sun}$ or more, {\\it i.e.}\\ they are at least main sequence A stars. The orbital separation in 2M\\,1533+3759 is $a = 0.98 \\pm 0.04~R_{\\sun}$.  The secondary has $M_{2} = 0.113 \\pm 0.017~M_{\\sun}$, $R_{2} = 0.152 \\pm 0.005~R_{\\sun}$ and $T_{\\rm eff_{2}} = 3100 \\pm 600$~K, consistent with a main sequence M5 star.  If 2M\\,1533+3759 becomes a cataclysmic variable (CV), its orbital period will be 1.6~h, below the CV period gap. ",
        "introduction": "Subdwarf B (sdB) stars are evolved, hot, compact stars ($23,000$ K $< T_{\\rm eff} < 37,000$ K; $5.2 < \\log g < 6.0$), commonly found in the disk and halo of our Galaxy \\citep{Saffer94}.  They are believed to ascend the first red giant branch (RGB) following the exhaustion of central hydrogen, somehow experiencing sufficient mass loss prior to the RGB tip to remove nearly all of their envelopes. They subsequently evolve blueward from the RGB before igniting helium in their cores. From an evolutionary point of view, sdB stars are also known as extreme horizontal branch (EHB) stars \\citep{Heber86}. Their helium burning cores, generally expected to be just under $0.5~M_{\\sun}$, are essentially identical to those of normal horizontal branch (HB) stars. However, their hydrogen envelopes are too thin and inert ($< 0.01~M_{\\sun}$) \\citep{Saffer94,Heber86} to support double shell burning, so they never make it to the asymptotic giant branch. Following core helium exhaustion, they evolve directly into sdO stars before proceeding down the white dwarf cooling track \\citep{Dorman93}.  In the context of understanding Galaxy evolution and cosmology, sdB stars play an important role because their large UV flux appears to be the dominant source of the ``UV upturn'' phenomenon observed in elliptical galaxies and the centers of spiral bulges \\citep{DB82,Greg99,Brown97}.  The UV excess in old stellar populations has been used as an age indicator in evolutionary population synthesis \\citep{Yi97,Yi99}, although more recent work has begun to consider alternative binary scenarios that would have quite different effects \\citep{P08}.  Various evolutionary scenarios have been proposed for sdB stars, but the details of the formation mechanisms are not yet well determined. Possible formation channels can be divided into single star evolution with enhanced mass loss at the tip of RGB \\citep{Castellani93,DCruz96} and close binary evolution, first suggested by \\citet{Mengel76}. Recently, Han et al.\\ (2002, 2003) conducted an in-depth theoretical investigation through binary population synthesis. They found that common-envelope evolution, resulting from dynamically unstable mass transfer near the tip of the first RGB, should produce short-period binaries ($P \\approx 0.1-10$~day) with either a main sequence (MS) or white dwarf (WD) companion.  If a red giant star loses nearly all of its envelope prior to the red giant tip via stable mass transfer, a long-period sdB binary with a MS companion can be produced ($P \\approx 10-500$~day).  A most interesting feature of Han et al.'s models is that they predict a much larger range of sdB progenitor masses than had previously been considered, including stars sufficiently massive to avoid a helium flash and instead undergo quiescent helium ignition in non-degenerate cores (see also \\citealp{Hu07}; \\citealp{Politano08}).  Binary formation scenarios appear likely to be responsible for the majority of observed field sdB stars, as a large fraction are observed to occur in binaries (e.g., \\citealp{Lisker05}; \\citealp{MR03}; \\citealp{Maxted01}; \\citealp{Saffer01}; \\citealp{Green97}; \\citealp{Allard94}). Nevertheless, the same studies show that there are a sizable fraction of sdB stars, 30\\% or more, that do not now appear to be in binaries: there is no sign of a companion in high S/N optical spectra or infrared colors, and their radial velocities are constant to within the observational errors (a few km s$^{-1}$) over many months. \\citet{MB08} also found a significant fraction, 96\\%, of sdB stars in globular clusters to be single stars, in contrast to observed field sdB stars. Han et al.\\ (2002, 2003) investigated the possibility of forming single sdB stars by merging two helium white dwarfs, which would allow the formation of more massive sdB stars ($0.4-0.65~M_{\\sun}$), and \\citet{Politano08} considered the possibility that some sdB stars might form from mergers during common envelope evolution, followed by rotationally-induced mass loss. Still, unusually high mass loss in single red giant stars cannot yet be ruled out.  The distribution of sdB masses is clearly one of the most important constraints on the several possible formation channels.  Different observational techniques provide different windows of opportunity for investigating these masses.  More sdB masses have been derived by asteroseismology than by any other method to date.  Asteroseismology provides an extremely high level of precision (and is the only way to determine envelope masses, in addition to total masses), but it has so far been successfully applied only to the relatively rare short period sdB pulsators.  Two different types of multimode sdB pulsators have been discovered: short period V361\\,Hya pulsators (originally, EC\\,14026 stars, \\citealp{Kilkenny97}) which comprise a rather small percentage of the hotter sdB stars, and longer period V1093\\,Her pulsators (PG\\,1716 stars, \\citealp{Green03}), which seem to be fairly common among cooler sdB stars. The rapid oscillations of V361\\,Hya stars are interpreted as low-order pressure modes (p-modes) that are driven by a $\\kappa$-mechanism associated with the radiative levitation of iron in the thin diffusion-dominated envelopes \\citep{Charpinet96, Charpinet97}.  The same mechanism has also been shown to explain the excitation of high-order gravity modes (g-modes) in the V1093\\,Her stars \\citep{Fontaine03}.  Asteroseismological modeling has so far been extremely successful with p-mode pulsations in the envelopes of sdB stars, and the resulting stellar parameters are generally in very good agreement with theoretical expectations (e.g., \\citealp{Fontaine08}; \\citealp{Charpinet07}, and references therein). On the other hand, g-mode pulsations, which extend much more deeply into the stellar cores, will require more sophisticated interior models before they can be satisfactorily analyzed by asteroseismology \\citep{Randall07}.  The list of p-mode pulsators whose parameters have been derived by asteroseismology is presented in Table 1.  Most of the derived masses are within a few hundredths of a solar mass of the canonical sdB mass of $0.48~M_{\\sun}$, except for PG\\,0911+456 \\citep{Randall07}, which will be discussed further in \\S 7. Interestingly, the only post-common envelope binaries in this list are Feige\\,48 \\citep{VanGrootel08a} and PG\\,1336$-$018 \\citep{Charpinet08}.  Indeed, the large majority of V1093\\,Her stars exhibit low or negligible radial velocity variations, of the order of a few km s$^{-1}$ or less, and thus must be single stars, or have extremely low mass companions, or else occur in long period binaries with a main sequence F, G, or K star primary.  This is not surprising, since sdB stars whose radial velocity variations are clearly indicative of post-common envelope binaries are preferentially found at temperatures cooler than most V1093\\,Her stars \\citep{Green08}.  Traditional methods of deriving masses by exploiting binary properties are therefore quite important.  For one thing, binaries provide a vital test of asteroseismology in the rare cases where the sdB primary is a pulsator.  More importantly, until improved asteroseismic models and extensive satellite observations make it possible to successfully model g-mode sdB pulsators, the only way to derive masses for a larger sample of post-common envelope sdB stars is to analyze their binary properties.  Finally, there are simply a large number of binaries that contain non-pulsating sdB stars.  The difficulty with most sdB stars in post-common envelope systems is that they are single-lined spectroscopic binaries with essentially invisible compact secondaries.  In principle, precise measurements of the sdB surface gravity and rotational velocity in a tidally locked system will yield the orbital inclination, allowing the individual component masses to be determined from the mass function (e.g., \\citealp{Geier08}), but the accuracy of this approach has not yet been proven. There are, however, a small number of rare post-common envelope sdB+dM binaries \\citep{Maxted04}, which have been identified by their reflection effects -- e.g., the sinusoidal variation observed in the light curve due to reradiated light from the heated side of the tidally locked M dwarf -- that are more promising.  The known sdB+dM systems are summarized in Table 2.  If narrow lines originating from the cool secondary could be detected, then masses of both components could be derived from the double-lined spectroscopic solution.  Again, this should be possible in principle, especially in binaries with the shortest orbital periods, where the heated face of the secondary is brighter than it otherwise would be, but results so far have been ambiguous.  \\citet{Vuckovic08} detected emission lines from the secondary in PG\\,1336-018, by subtracting the spectrum of the hot primary from spectra taken at other phases, but the S/N of the spectra were only sufficient to claim general consistency with the orbital solution described in \\citet{Vuckovic07}. Using much higher S/N spectra of a similar sdO+dM binary, AA~Dor, \\citet{Vuckovic08} were able to determine a velocity amplitude for the secondary, but their derived primary mass has now been vigorously disputed by \\citet{Rucinski09}.  \\citet{WS99} presented a good argument for the detection of $H\\alpha$ absorption lines from the secondary in HW~Vir, again by subtracting the spectrum near minimum light from spectra near maximum light, and obtained reasonable velocities, but it is perplexing that absorption lines and no emission lines should have been seen.  An apparently more successful method is to model the light variations in sdB+dM binaries exhibiting reflection effects, especially the eclipsing systems, in order to determine the system parameters.  This is a very complex endeavor.  The models have many free parameters, and there are large uncertainties that typically require additional information to constrain the solution.  Often, the light curves provide more than one high quality solution.  For example, \\citet{Drechsel01} had to make use of a mass--radius relation for the secondary star to decide between two solutions that implied quite different sdB masses for HS\\,0705+6700 (0.483 and $< 0.3~M_{\\sun}$). \\citet{Heber04} needed to use their spectroscopic $\\log g$ and mass--radius relations to discriminate between two solutions with different secondary albedos and inclinations in HS\\,2333+3927. \\citet{Vuckovic07} found three possible solutions modeling the light curves PG\\,1336$-$018, and it was not possible to choose between two of them until \\citet{Charpinet08} derived a consistent primary mass by asteroseismological modeling.  Furthermore, even when a single family of solutions can be identified, there still remain unavoidable ambiguities in choosing one ``best'' model \\citep{Drechsel01}.  Even in the most favorable cases of eclipsing sdB+dM binaries, the eclipses are not flat-bottomed, leading to a small range of nearly equivalent solutions in the vicinity of the deepest minimum.  The resulting small variations in the mass mass ratio, $q$, lead to a significant range in the derived sdB mass.  The uncertainties are obviously larger when there is no eclipse.  Still, light curve modeling provides valuable information, and when the derived sdB mass can be verified -- rarely by asteroseismology, more often from consistency with the spectroscopic surface gravity or projected rotational velocity -- our confidence in the results is greatly increased.  It is clearly important to investigate as many sdB+dM binaries as possible, especially the eclipsing systems, in order to build up a more comprehensive picture of sdB masses produced by post-common envelope evolution and to compare with the distribution of masses from other formation channels.  In this paper, we report on the system parameters of 2M\\,1533+3759 ($15^{h}33^{m}49.44^{s}$, $+37\\degr59\\arcmin28.2\\arcsec$, J2000), a new eclipsing sdB+dM binary with a longer orbital period than any eclipsing sdB+dM discovered so far.  This star was first recognized as an sdB in 2005 (although it remained unpublished) during a continuing spectroscopic survey \\citep{Green08} of bright blue stellar candidates selected from a variety of sources, including the 2MASS survey \\citep{Skrutskie06}.  The current investigation was motivated by \\citet{KS07}, who discovered that 2M\\,1533+3759 is an eclipsing binary, NSVS\\,07826147, through their work with the Northern Sky Variability Survey (NSVS; \\citealp{Wozniak04}).  \\citet{KS07} identified a group of nine eclipsing binaries with short periods and relatively narrow eclipse widths, indicating very small radii for the components.  Since their list includes the well-known HW~Vir (\\citealp{Lee09} and references therein), as well as 2M\\,1533+3759, which we confirmed to be a spectroscopic near-twin of HW~Vir, \\citet{KS07} proposed that the other objects in their Table 3 might also be sdB+dM binaries.  \\S 2 presents the results from our follow-up spectra for these stars.  In \\S 3, we describe new spectroscopy and photometry for 2M\\,1533+3759.  The data analyses are given in \\S 4 and \\S 5, and the system parameters are derived in \\S 6.  We discuss possible selection effects and consider the unusually low derived mass for the sdB mass in \\S 7.  \\S8 looks at the evolution of 2M\\,1533+3759, and \\S 9 contains our conclusions. ",
        "conclusions": "The sdB star 2M\\,1533+3759 is the seventh eclipsing sdB+dM binary discovered to date.  Its orbital period of 0.16177042 $\\pm$ 0.00000001~d is 29\\% longer than the 0.12505 day period of the next longest eclipsing sdB+dM, BUL$-$SC16~335.  The amplitude of the reflection effect in 2M\\,1533+3759 is surprisingly strong, only about 0.05 mag weaker than the amplitude observed in HW~Vir, in spite of the longer orbital period and the fact that the temperatures of the primary stars are similar.  2M\\,1533+3759 is the only new sdB binary among the eclipsing systems that were proposed to be sdB+dM by \\citet{KS07} on the basis of their narrow eclipse widths.  This result is consistent with the 2MASS colors of other known reflection-effect sdB+dM systems, all of which have $J-H < 0$. 2M\\,1533+3759 and the archetypal HW~Vir \\citep{MM86} are the only two binaries in Kelley \\& Shaw's (2007) Table~3 that have similarly blue IR colors, and the only two that contain sdB stars.  Spectroscopic parameters 2M\\,1533+3759 were derived by fitting Balmer and helium line profiles in high S/N spectra to a grid of zero-metallicity NLTE model atmospheres. The effective temperatures derived from low (9\\AA) and medium (1.9\\AA) resolution spectra exhibit clear variations with orbital phase.  Phase variations are much less significant for the surface gravities, and completely negligible for the helium abundance fraction.  Our adopted parameters for the sdB star, $T_{\\rm eff} = 29230 \\pm 125$\\,K, $\\log g = 5.58 \\pm 0.03$, log $N$(He)/$N$(H) =$ -2.37 \\pm 0.05$, were determined from medium resolution spectra taken when the reflection effect was near minimum.  The inferred rotational velocity has a negligible affect on the derivation of these parameters.  Light curve modeling with the MORO code produced only one well-fitting solution consistent with a core helium burning primary.  The system mass ratio, $q$ ($M_{\\rm 2}$/$M_{\\rm 1}$), is $0.301 \\pm 0.014$ and the inclination angle, $i$, is $86.6\\degr \\pm 0.2\\degr$.  The robustness and precision of these numbers are due to the high precision of the light curves and the fact that the system is eclipsing.  Radial velocities for the sdB component were used to derive the velocity amplitude, K$_{\\rm 1}$ = $71.1 \\pm 1.0$~km~s$^{-1}$, leading to component masses of $M_{\\rm 1}$ = $0.376 \\pm 0.055~M_{\\sun}$ and $M_{\\rm 2}$ = $0.113 \\pm 0.017~M_{\\sun}$.  The errors in the masses are dominated by the uncertainty in $q$.  Since the mass ratio and inclination are even more uncertain in non-eclipsing systems, our inability to more tightly constrain the primary mass provides a strong illustration for why useful sdB masses from light curve modeling can usually be obtained only from eclipsing binaries.  The orbital separation derived from the masses and the period is $a = 0.98 \\pm 0.04~R_{\\sun}$.  The individual radii, $R_{\\rm 1}$ = $0.166 \\pm 0.007~R_{\\sun}$, $R_{\\rm 2}$ = $0.152 \\pm 0.005~R_{\\sun}$ were then calculated from the relative radii, $R_{\\rm 1}/a$ and $R_{\\rm 2}/a$, determined by the light curve solution.  Both radii are consistent with theoretical expectations, and the resulting sdB surface gravity, $\\log g = 5.57 \\pm 0.07$, is completely consistent with the adopted spectroscopic value above.  We constructed a synthetic line-blanketed spectrum to investigate potential systematic effects caused by our use of zero metallicity NLTE atmospheres to derive the spectroscopic parameters.  If 2M\\,1533+3759 had solar abundances of C, N, O, S, and Fe in its atmosphere, our assumption of zero metals would have overestimated the effective temperature by about 2000\\,K, and the surface gravity by almost 0.2 dex.  Thus, the true $T_{\\rm eff}$ and $\\log g$ abundances would have been about 27300\\,K and 5.40, respectively.  The modeled light curve solution at this lower temperature is only negligibly different from our original solution, and thus the resulting system parameters remain essentially unchanged.  However, in this case, the calculated sdB surface gravity, $\\log g = 5.57$, would be much less consistent with the expected value of 5.40.  This suggests that the full correction to solar metallicites assumed above is not appropriate for 2M\\,1533+3759.  An important conclusion is that there is now significant observational evidence, from two completely independent techniques, asteroseismology (PG\\,0911+456) and modeling of eclipsing/reflection effect light curves (2M\\,1533+3759, and perhaps HS\\,2333+3927), for the existence of sdB stars with masses significantly lower than the canonical $0.48 \\pm 0.02~M_{\\sun}$.  2M\\,1533+3759 must have formed via the first common-envelope channel, since the companion is an M dwarf.  With a probable sdB mass in the range $0.32-0.43~M_{\\sun}$, this star is expected to have evolved from a main sequence A star with an initial mass $> 1.8-2.0~M_{\\sun}$.  The existence of such a binary might support recent theoretical predictions that sdB stars can be produced by such massive progenitors, including the assumption that the ionization energy of the red giant envelope contributes to the ejection of the common envelope \\citep{Han02,Han03}.  If the primary mass of 2M\\,1533+3759 could be measured more precisely, or if the separation between the two components could be measured independently, this system ought to provide a very useful observational constraint for the upper limit to the main sequence mass of an sdB progenitor.  If 2M\\,1533+3759 becomes a cataclysmic variable (CV) after orbital shrinkage due to gravitational radiation brings the Roche lobe into contact with the M dwarf secondary, its orbital period of the CV at the onset of mass transfer will be 1.6 hours, below the CV period gap."
    },
    "0911/0911.5626_arXiv.txt": {
        "abstract": "{ We investigate the dust size and dust shell structure of the bipolar proto-planetary nebula M 1--92 by means of radiative transfer modeling. Our models consists of a disk and bipolar lobes that are surrounded by an AGB shell, each component having different dust characteristics. The upper limit of the grain size $a_\\mathrm{max}$ in the lobes is estimated to be $0.5~\\mu$m from the polarization value in the bipolar lobe. The $a_\\mathrm{max}$ value of the disk is constrained with the disk mass (0.2~$M_{\\sun}$), which was estimated from a previous CO emission line observation. We find a good model with $a_\\mathrm{max}=1000.0~\\mu$m, which provides an approximated disk mass of 0.15~$M_{\\sun}$.  Even taking into account uncertainties such as the gas-to-dust mass ratio, a significantly larger dust of $a_\\mathrm{max}>100.0~\\mu$m, comparing to the dust in the lobe, is expected. We also estimated the disk inner radius, the disk outer radius, and the envelope mass to be 30~$R_\\star$(=9~AU), 4500~AU, and 4~$M_{\\sun}$, respectively, where $v_\\mathrm{exp}$ is the expansion velocity. If the dust existing in the lobes in large separations from the central star undergoes little dust processing, the dust sizes preserves the ones in the dust formation. Submicron-sized grains are found in many objects besides M 1--92, suggesting that the size does not depend much on the object properties, such as initial mass of the central star and chemical composition of the stellar system. On the other hand, the grain sizes in the disk do. Evidence of large grains has been reported in many bipolar PPNs, including M 1--92. This result suggests that disks play an important role in grain growth. } ",
        "introduction": "Planetary Nebula (PN) morphology represents a history of physical processes in the stellar$/$circumstellar environments. A large sample ($\\sim$500) of galactic PNs has shown that 20--25~\\% of them have spherical structures, but the others have asymmetric (elliptic, bipolar, and point-symmetric structures) \\citep[e.g.][]{za86,stanghellini93,cs95,manchado96,soker97,st98}. Several theoretical studies have been carried out to explain the PN morphologies \\citep[see review by][]{bf02}. Widely accepted facts and physical processes are cool spots on the stellar surface \\citep{frank95,sc99}, the stellar rotation \\citep{bc93}, magnetic field \\citep{pascoli92,mm98,mm99,matt00}, and binary interaction \\citep{lss79,morris81,morris87,bl90,sl94}. While details are still unknown, these mechanisms one way or another cause the equatorially enhanced mass loss in the AGB phase \\citep[e.g.][]{meixner99,ueta00}. In the subsequent post-AGB phase, the central star blows away the low density fast wind, which inflates the lobe in the polar direction more than in the equatorial direction \\citep{kwok78,kwok82}. This interacting stellar wind (ISW) model explains various morphologies of elliptic and bipolar PNs \\citep{balick87}. In the above scenario, the binary interaction is thought to be one of the most promising mechanism to shape PNs with a narrow waist between the bipolar lobes because this can amplify the mass loss in the equatorial plane and make a disk-like structure \\citep[e.g.][]{sr00}. In addition, the binary interaction is also thought to form asymmetric structures such as jets, ansae \\citep[e.g.][] {soker90,soker92}, and point-symmetric shapes such as a spiral and quadruple by precessing motion \\citep{manchado96,mh06}.  It is difficult to directly detect the effect of binary interaction on PN shaping using the current observing techniques. However, it would be possible to infer the presence of the effects due to the aforementioned scenario if the inner part of the circumstellar environment is investigated closely enough. Near-infrared imaging polarimetry is a powerful method to probe dust shells and to provide important information that can be derived in more details than non-polarimetric imaging. Previous observations and dust scattering modeling have revealed disk and bipolar lobe structures of evolved stars \\citep[e.g.][hereafter UMM07]{ss95,su03,gledhill05,ueta05,murakawa05,murakawa07,ueta07}. However, most previous experiments treated only a single grain model in the entire dust shell or SEDs were not considered simultaneously. These shortcomings have prohibited us from investigating into the grain and disk properties in detail. To resolve the situation, we considered both polarimetric images and SEDs in our recent radiative transfer modeling of the bipolar proto-PN (PPN) \\object{Frosty Leo}, assuming two different dust models, one in the equatorial region (disk) and the other in the bipolar lobes \\citep{murakawa08a}. From this modeling, we were able to derive evidence for grain growth in the disk region. In this paper, we use this technique for the bipolar PPN: \\object{M 1--92} (=IRAS 19343+2926), whose characteristics were discussed previously by us (UMM07). In Sect.~\\ref{rtmodel}, details of our radiative transfer modeling are described. We will discuss the implication of our model result in Sect.~\\ref{discuss}. ",
        "conclusions": "We performed two-dimensional radiative transfer modeling of the dust shells of the bipolar PPN M 1--92. Our modeling applied geometries with a disk and bipolar envelope surrounded by an AGB wind shell, each of which has different dust characteristics. The model parameters were constrained by comparing them with the previously observed SEDs, the intensity and polarization images from the HST NICMOS 2 archived data (UMM07), and a previous radio observation in the CO emission line \\citep{bujarrabal98b,alcolea07}. With a waterlily-shaped hollow envelope, the bright bipolar lobes of M 1--92 were reproduced. For the dust sizes, we found submicron-sized grains ($a_\\mathrm{max}=0.5~\\mu$m) in the bipolar lobes.  The dust size in the disk was constrained with the disk mass, which was estimated from the CO emission line data. We obtained $a_\\mathrm{max}=1000.0~\\mu$m. Although this estimate includes some uncertainties in the gas-to-dust mass ratio, the grain size distribution, and the geometry of the region considered to be the disk, an $a_\\mathrm{max}\\la100.0~\\mu$m is hard to explain the disk mass and a significantly large size of $>100.0~\\mu$m is expected.  We conclude that grain growth is likely to occur in the M 1--92 disk. Further works including submillimeter and millimeter wavelength imaging and photometry would provide better interpretations on grain growth.  The small grains in the bipolar lobes are consistent with results of many other AGB and PPNs, suggesting that dust formed by AGB mass loss has a similar size, which does not depend on the object much. The formation of such submicron-sized dust is explained with dust formation and stellar wind theories.  The presence of large grains in the disk should be explained with another mechanism, which is probably the long-lived disk hypothesis. The grain growth possibly depends on the disk geometry and the stellar temperature. Better understanding will be provided if a detailed systematic analysis is made for many other samples."
    },
    "0911/0911.2230_arXiv.txt": {
        "abstract": "We present the baryon fractions of 2MASS groups and clusters as a function of cluster richness using total and gas masses measured from stacked ROSAT X-ray data and stellar masses estimated from the infrared galaxy catalogs. \u00a0We detect X-ray emission even in the outskirts of clusters, \u00a0beyond $r_{200}$ for richness classes with X-ray temperatures above 1 keV. This enables us to more accurately determine the total gas mass in these groups and clusters. We find that the optically selected groups and clusters have flatter temperature profiles and higher stellar-to-gas mass ratios than the individually studied, X-ray bright clusters. We also find that the stellar mass in poor groups with temperatures below 1~keV is comparable to the gas mass in these systems.  Combining these results with individual measurements for clusters, groups, and galaxies from the literature, we find a break in the baryon fraction at $\\sim 1$~keV. Above this temperature, the baryon fraction scales with temperature as $f_{b} \\propto T^{0.20\\pm0.03}$. \u00a0\u00a0We see significantly smaller baryon fractions below this temperature, and the baryon fraction of poor groups joins smoothly onto that of systems with still shallower potential wells such as normal and dwarf galaxies where the baryon fraction scales with the inferred velocity dispersion as $f_{b} \\propto \\sigma^{1.6}$. The small scatter in the baryon fraction at any given potential well depth favors a universal baryon loss mechanism and a preheating model for the baryon loss. The scatter is, however, larger for less massive systems. Finally, we note that although the broken power-law relation can be inferred from data points in the literature alone, the consistency between the baryon fractions for poor groups and massive galaxies inspires us to fit the two categories of objects (galaxies and clusters) with one relation. ",
        "introduction": "There is a missing baryon problem in the nearby universe (see the review by Bregman 2007). We can see this by comparing the baryon fraction from direct measurements of galaxies and galaxy clusters to the value determined from the WMAP 3-year data (Spergel et al.\\ 2007), where the baryon fraction, the ratio of the baryonic mass to the total mass, is $f_b = \\Omega_b/\\Omega_m = 0.175\\pm0.012$\\footnote{This is consistent with the value, $f_b = 0.171\\pm0.009$, determined from the WMAP 5-year data (Dunkley et al.\\ 2009).} independent of the Hubble constant. \u00a0The baryon fraction of galaxies can also be measured by comparing the gravitating mass measured through gravitational lensing or the dynamics of satellites to the stellar and gas mass (e.g. Hoekstra et al.\\ 2005; Heymans et al.\\ 2006; Mandelbaum et al.\\ 2006; Gavazzi et al.\\ 2007; Jiang \\& Kochanek 2007).\u00a0 These studies find similar results. For example, Hoekstra et al.\\ (2005) used isolated nearby galaxies to find $f_b = 0.056$ for spirals and $f_b = 0.023$ for ellipticals -- on average, spirals have lost 2/3 of their initial baryons and ellipticals have lost 6/7 of their initial baryons. \u00a0For galaxy clusters, where the gas mass dominates the baryon content, the missing baryon problem is less severe, particularly in the most massive clusters. Recent studies of relaxed massive clusters (e.g., Vikhlinin et al.\\ 2006; Allen et al.\\ 2008) show that after adding the stellar baryon component, the baryon fraction in most massive, relaxed clusters is close to the cosmological value. \u00a0In the less massive galaxy groups, only a few are suitable for measuring the baryon fraction out to the outskirts of the groups, and they generally show significantly lower baryon fractions (e.g., Sun et al.\\ 2009).  These observations indicate that the extent of baryon loss depends on the depth of the gravitational potential well: rich clusters retain their cosmological allotment of baryons, while galaxies are baryon-poor. \u00a0The mass scale of the transition in the baryon fraction is not well-constrained because there are so few measurements on the mass scales of groups. \u00a0In groups, the baryon content should still be dominated by the gas mass, but the weakness of the X-ray emission from the outskirts of groups means that few groups have sufficiently deep X-ray observations to measure the baryon fraction out to a large enough radius to represent a complete inventory. Few are measured to \\rfive\\ (Sun et al.\\ 2009), where the over density inside the radius is 500 times the critical density of the universe $\\rho_c = 3H(z)^2/8\\pi G$, and even fewer are observed to \\rtwo, close to the virial radius. \u00a0The alternative to individual measurements of the baryon fractions in clusters and groups, is to stack the \\rosat\\ All-Sky Survey (RASS; Voges et al.\\ 1999) X-ray data to measure the average properties of groups and clusters selected by other means. In particular, Dai et al.~(2007) were able to detect gas emission well beyond $r_{200}$ for galaxy groups with temperatures $T>1.5$~keV, significantly further out than the measurements of individual groups, using groups and clusters selected from the 2MASS (Skrutskie et al.\\ 2006) catalogs using a matched filter algorithm (Kochanek et al.\\ 2001). For other richness classes, the gas emission is detected beyond $r_{500}$, whereas X-ray observations of these individual poorest groups only have detections of their central regions. \u00a0The stacking method has also been applied to the other large optically-selected cluster catalogs (Rykoff et al.\\ 2008; Shen et al.\\ 2008).  In this paper, we combine the results from the Dai et al.~(2007) stacking analysis with results for individual galaxies, groups and clusters to summarize the dependence of the baryon fraction on potential well depth. \u00a0This provides a more complete picture of baryon losses across a broad range of systems, which should enable us to better understand the origin of the losses. We calibrate all the measurement to the three year WMAP cosmology of $H_0 = 73~\\rm{km~s^{-1}~Mpc^{-1}}$, $\\Omega_{\\rm m} = 0.26$, and $\\Omega_{\\Lambda}= 0.74$ (Spergel et al.\\ 2007).  \\begin{deluxetable}{cccccccccc} \\tabletypesize{\\scriptsize} \\tablecolumns{10} \\tablewidth{0pt} \\tablecaption{Average Properties of the RASS-Stacked 2MASS Groups and Clusters I\\label{tab:bfi}} \\tablehead{ \\colhead{Richness} & \\colhead{$N_{*666}$} & \\colhead{$\\sigma$} & \\colhead{$\\langle z \\rangle$} & \\colhead{$T$} & \\colhead{$R_c$} & \\colhead{$\\beta$} & \\colhead{$r_{200}$} & \\colhead{$r_{500}$} \\\\ \\colhead{Class} & \\colhead{} & \\colhead{(\\kms)} & \\colhead{} & \\colhead{(keV)} & \\colhead{(Mpc)} & \\colhead{} & \\colhead{(Mpc)} & \\colhead{(Mpc)} } \\startdata 0 & 16.6 & 840 & 0.082 & $4.7^{+1.4}_{-0.7}$ \u00a0\u00a0\u00a0& $0.42\\pm0.03$ & $0.75\\pm0.05$ & 1.57 & 0.94 \\\\ 1 & 5.27 & 440 & 0.068 & $1.7^{+0.5}_{-0.3}$ \u00a0\u00a0\u00a0& $0.29\\pm0.04$ & $0.59\\pm0.06$ & 0.84 & 0.49 \\\\ 2 & 1.80 & 330 & 0.054 & $1.09^{+0.09}_{-0.05}$ & $0.19\\pm0.06$ & $0.59\\pm0.10$ & 0.70 & 0.41 \\\\ 3 & 0.60 & 290 & 0.039 & $0.91^{+0.10}_{-0.05}$ & $0.12\\pm0.08$ & $0.55\\pm0.10$ & 0.65 & 0.40 \\\\ 4 & 0.20 & 220 & 0.027 & $0.60^{+0.11}_{-0.20}$ & $0.13\\pm0.10$ & $0.68\\pm0.14$ & 0.58 & 0.35 \\\\ \\enddata \\tablecomments{The $\\beta$ and $R_c$ values are slightly different from the results from Dai et al.\\ (2007), since we re-fit the surface brightness profiles without the residual background component.} \\end{deluxetable}  \\begin{deluxetable}{cccccccccc} \\tabletypesize{\\scriptsize} \\tablecolumns{10} \\tablewidth{0pt} \\tablecaption{Average Properties of the RASS-Stacked 2MASS Groups and Clusters II\\label{tab:bfii}} \\tablehead{ \\colhead{$M_{g, 500}$} & \\colhead{$M_{500}$} & \\colhead{$f_{g, 500}$} & \\colhead{$M_{g, 200}$} & \\colhead{$M_{s, 200}$} & \\colhead{$M_{200}$} & \\colhead{$r_{sg, 200}$} & \\colhead{$f_{g, 200}$} & \\colhead{$f_{b, 200}$} \\\\ \\colhead{($10^{12} M_{\\odot}$)} & \\colhead{($10^{12} M_{\\odot}$)} & \\colhead{} & \\colhead{($10^{12} M_{\\odot}$)} & \\colhead{($10^{12} M_{\\odot}$)} & \\colhead{($10^{12} M_{\\odot}$)} & \\colhead{} & \\colhead{} & \\colhead{} } \\startdata $24.0\\pm1.0$ \u00a0\u00a0& $320\\pm70$ & $0.08\\pm0.02$ & 50.2$\\pm$2.1 \u00a0& 6.2$\\pm$0.4 \u00a0\u00a0& 600$\\pm$140 & 0.12$\\pm$0.01 & 0.08$\\pm$0.02 \u00a0\u00a0& 0.093$\\pm$0.021 \u00a0\u00a0\\\\ $2.879\\pm0.27$ & $43\\pm10$ \u00a0& $0.07\\pm0.02$ & 8.04$\\pm$0.75 & 2.41$\\pm$0.06 & 91$\\pm$22 \u00a0\u00a0& 0.30$\\pm$0.03 & 0.09$\\pm$0.02 \u00a0\u00a0& 0.115$\\pm$0.029 \u00a0\u00a0\\\\ $0.935\\pm0.08$ & $26\\pm5$ \u00a0\u00a0& $0.037\\pm0.007$ & 2.30$\\pm$0.19 & 1.25$\\pm$0.02 & 49$\\pm$9 \u00a0\u00a0\u00a0& 0.54$\\pm$0.05 & 0.047$\\pm$0.009 & 0.073$\\pm$0.015 \u00a0\u00a0\\\\ $0.367\\pm0.08$ & $21\\pm4$ \u00a0\u00a0& $0.017\\pm0.005$ & 0.81$\\pm$0.17 & 0.65$\\pm$0.01 & 37$\\pm$7 \u00a0\u00a0& 0.80$\\pm$0.17 & 0.022$\\pm$0.006 & 0.040$\\pm$0.012 \u00a0\u00a0\\\\ $0.159\\pm0.02$ & $14\\pm5$ \u00a0\u00a0& $0.012\\pm0.004$ & 0.33$\\pm$0.04 & 0.63$\\pm$0.02 & 26$\\pm$9 \u00a0\u00a0& 1.87$\\pm$0.25 & 0.013$\\pm$0.005 & 0.037$\\pm$0.013 \\\\ \\enddata \\tablecomments{$r_{sg}$, $f_g$, and $f_b$ are the ratio of stellar mass to gas mass, the gas fraction, and the baryon fraction, respectively. The derived quantities scale with the Hubble constant as $r\\propto h^{-1}$, $M_g \\propto h^{-5/2}$, $M_s \\propto h^{-2}$, $M \\propto h^{-1}$, $f_{sg} \\propto h^{1/2}$, and $f_g \\propto h^{-3/2}$. The baryon fraction $f_b$ does not have a simple dependence on $h$, and is calculated assuming $h=0.73$.} \\end{deluxetable}  \\begin{figure} \\epsscale{1} \\plotone{f1.eps} \\caption{Normalized surface brightness profiles of the stacked RASS images of the 2MASS clusters multiplied by $R^2$ and in arbitrary units, where we have shifted the five profiles for clarity. \u00a0In particular, we detect extended cluster/group emission beyond \\rtwo\\ for richness classes 1 and 2. The dashed lines are the best fit $\\beta$ models. \\label{fig:sb}} \\end{figure}  \\begin{figure} \\epsscale{1} \\plotone{f2.eps} \\caption{Normalized temperature profiles for the stacked RASS data for the 2MASS clusters in relative units. \u00a0We also show the theoretical predictions from numerical simulations for a range of cluster mean temperatures with dotted lines (Stanek et al.\\ 2009). \u00a0The solid line is our best fit model to the data points within \\rtwo, and the dashed line is the universal temperature model from Vikhlinin et al.\\ (2005) for relaxed X-ray clusters. \\label{fig:t}} \\end{figure} ",
        "conclusions": "Baryon loss, compared to the global average, as a function of potential well depth over the range from rich clusters to dwarf galaxies should provide a good basis for understanding the causes of the losses. \u00a0For deep potential wells, rich clusters with $T \\gs 5$~keV, baryon loss is not significant. \u00a0The baryon fractions of these clusters, including both the gas and stellar components and corrected to \\rtwo, are close to the cosmological value measured by WMAP. This is consistent with recent studies (e.g, Vikhlinin et al.\\ 2006; Giodini et al.\\ 2009; Stanek et al.\\ 2009) of individual clusters. \u00a0For clusters between 5 and 1 keV, baryon loss becomes increasingly important, a trend that can be modeled as a power-law. \u00a0Near $\\approx 1$~keV, there is a sharp decline in the baryon fraction of groups, and this decline smoothly joins onto the observed baryon fractions in galaxies. Groups and galaxies of the same potential well depth appear to have the same baryon fractions. This suggests that baryons missing from galaxies do not reside in the potential wells of their parent groups.  The small scatter in the baryon fractions of galaxies, galaxy groups, and galaxy clusters with $V_c \\gs 30\\kms$ about this general trend indicates that baryon loss mechanism of different systems must be fairly universal and primarily controlled by the depth of the system's potential well. \u00a0Such a mechanism could be the pre-heating of baryons before they collapse. In this picture, the baryons never fell into galaxies and groups, but remain well beyond $r_{200}$. \u00a0Alternatively, the gas fell into the potential wells and was subsequently removed by feedback provided from SNe and AGNs. \u00a0It is hard to reconcile this scenario with similar baryon fractions in systems with similar potential wells but very different stellar-to-gas mass ratios. \u00a0For example, if the feedback is dominated by SNe, whose rate is proportional to the stellar content, the gas rich systems should have lower baryon losses than systems dominated by stellar mass at the same potential well depth. \u00a0In some galaxies, Anderson and Bregman (2010) show that the energy from feedback is inadequate to expel the gas. \u00a0It is possible that feedback processes are responsible for increasing the scatter of the mean relation, since the total feedback energy depends on a variety of aspects such as the AGN fraction in clusters or groups, the radio AGN fraction, the stellar-to-gas ratio, and the stellar metallicity. \u00a0 For massive clusters, to interpret the small scatter of baryon fractions, feedback must be dominated by one process or the feedback energy cannot be significant compared to the depth gravitational potential well so that it will have little effect on the baryon fraction of clusters. \u00a0 For less massive systems, such as galaxies, it is easier for the baryons to escape their potential wells, and we do observe a larger scatter around the mean relation. For dwarf galaxies, one would expect the effect to be more pronounced, and we observe a very large scatter.  It has been well established that the scaling relations in clusters and groups of galaxies deviate from the ``self-similar'' model (e.g., Kaiser 1986; Navarro et al.\\ 1995). This includes the observed steeper slopes of $L$--$T$ (e.g., White et al.\\ 1997; Wu et al.\\ 1999; Rosati et al.\\ 2002) and $M$--$T$ (e.g., Xu et al.\\ 2001; Sanderson et al.\\ 2003; Arnaud et al.\\ 2005) relations, the break in the $L$--$T$ relation (e.g., Helsdon \\& Ponman 2000; Xue \\& Wu 2000) at the group regime, and an entropy floor in the $S$--$T$ relation (e.g., Helsdon \\& Ponman 2000). In particular, the breaks in these observed scaling relations occur at similar places as our break for baryon fractions ($\\sim$ 1~keV).  This leads us to contemplate that the causes for these breaks are from the same origin. There are a number of models proposed to explain the deviations from the ``self-similar'' model (see Voit 2005 for a review), which involve either pre-heating (e.g., Evrard \\& Henry 1991; Kaiser 1991; Bialek et al.\\ 2001; Kay et al.\\ 2007; Gottlober \\& Yepes 2007; Stanek et al.\\ 2009) or combinations of cooling (e.g., Muanwong et al.\\ 2001; Borgani et al.\\ 2002; Dave et al.\\ 2002; Kay et al.\\ 2003; Valdarnini et al.\\ 2003) and heating from AGNs (McCarthy et al.\\ 2008) or galactic winds (e.g., Dave et al.\\ 2008). Some of these models only attempt to solve the problems in a particular mass range, such as in groups and clusters. We note that the baryon fraction measurement can extend all the way to low mass systems to dwarf galaxies, which provides a much longer baseline to constrain models. The small scatter about the broken power-law relation for baryon fractions from dwarf galaxies to clusters suggests a universal baryon loss mechanism. Therefore, it is unlikely that models with ingredients only involving physical processes in a limited mass range will explain such a relation.  The hierarchical structure growth of the CDM theory has been successful to explain many observations; however it also encounter difficulties such as over-predicting low mass halos compared to observation (e.g., Kauffman et al.\\ 1993; Klypin et al.\\ 1999; Moore et al.\\ 1999). The discrepancy is reduced to be within a factor of $\\sim4$ when counting the ultra-faint dwarf galaxies recently discovered from SDSS (e.g., Simon \\& Geha 2007). One hypothesis to solve this discrepancy is that re-ionization at high redshift ($z\\sim10$) could suppress formation of dwarf galaxies (e.g., Bullock et al.\\ 2000; Somerville 2002; Benson et al.\\ 2002; Ricotti \\& Gnedin 2005; Moore et al.\\ 2006; Simon \\& Geha 2007; Koposov et al.\\ 2009). The entropy imposed by re-ionization make it more difficult for low mass halos to constrain gas and form stars. Therefore, a pre-heating model is capable of significantly affecting both the very low and high mass end of the systems.  However, detailed modeling is needed to test whether such models can reproduce the baryon fractions in a wide range of systems.  In addition, our results also present another constraint to the hierarchical growth theory.  We find that the baryon fractions are higher for more massive systems.  If massive clusters are formed through merging from smaller structures, then the theory need to explain why the baryon fraction increases after the mergers. Presumably, a large amount of cold or warm gas, which is outside the potential wells of smaller structures, is accreted during the merging process, and falls in the potential well of the larger structure afterwards. Alternatively, it is possible that the smaller structures that we observe today are quite different from those that merge into larger systems.  The baryon fraction measurements of our stacked groups ($T<1$~keV) have filled in the gap between measurements of individual galaxies and cluster/groups above 1~keV. \u00a0The consistency between the baryon fraction measured in these poor groups and massive galaxies suggests a universal baryon loss mechanism that primarily depends on the depth of the system's potential well. \u00a0More individual measurements of the baryon fractions for groups in this regime are needed to more accurately measure the average and begin to characterize its scatter. \u00a0Optimally, the X-ray emission should be measured at quite large radius, beyond \\rfive, and the groups should be selected using a range of different methods, including both the X-ray and optically selected groups. \u00a0For example, our stacked data for optically selected clusters show flatter surface brightness and temperature profiles at large radii than the X-ray bright clusters. The associated optical imaging data is also important to constrain the stellar content of poor groups, since our stacking analysis show that in this regime the stellar mass is comparable to the gas mass, as also suggested by other studies (e.g., David 1997; Giodini et al.\\ 2009)."
    },
    "0911/0911.1625_arXiv.txt": {
        "abstract": "Loop Quantum Gravity (L.Q.G.) is one of the two most promising tentative theory for a quantum description of gravity. When applied to the entire universe, the so-called Loop Quantum Cosmology (L.Q.C.) framework offers microscopical models of the very early stages of the cosmological history, potentially solving the initial singularity problem via bouncing solutions or setting the universe in the appropriate initial conditions for inflation to start, via a phase of super-inflation. More interestingly, L.Q.C. could leave a footprint on cosmological observables such as the Cosmic Microwave Background (CMB) anisotropies. Focusing on the modified dispersion relation when holonomy and inverse-volume corrections arising from the L.Q.C. framework are considered, it is shown that primordial gravity waves generated during inflation are affected by quantum corrections. Depending on the type of corrections, the primordial tensor power spectrum is either suppressed or boosted at large length scales, and strongly departs from the power-law behavior expected in the standard scenario. ",
        "introduction": "Building a complete quantum theory of gravitation is one of the greatest challenges in theoretical physics since the advent of general relativity and quantum mechanics. Since many decades, different proposals emerge and among the theories willing to reconcile Einstein gravity and quantum mechanics, Loop Quantum Gravity (L.Q.G.) is appealing as it is based  on a non-perturbative quantization of the 3-space geometry \\cite{rovelli1}. Loop Quantum Cosmology (L.Q.C.) is a finite, symmetry reduced model of L.Q.G.  suitable for the study of the whole Universe as a physical system (see, {\\it e.g.}, \\cite{bojo0}).   Such a challenge is currently a pure theoretical open question as all the observed phenomena do not {\\it a priori} need such a quantum theory of gravity to be explained (at least alternatives to a quantum-gravity-based explaination still exist). However, it is well admitted that even in the absence of observational facts to guide theoretical investigations, very high energetic phenomena can be used not only as {\\it gedanken experiments} to test theoretical proposals, but also as future probes of quantum gravitational physics. With black holes evaporation, primordial cosmology is probably the most powerful place in the cosmos where quantum gravity play an important role.  In its very early stages, our universe went through a phase of cosmic inflation \\cite{linde}. Although still debated, this high energetic phenomena (the energy scale of inflation could be as high as $10^{16}$~GeV and most of the inflationary models are inspired by high energy/particle physics or by quantum gravity) has received strong support from many experiments, including from the WMAP 5-Years results \\cite{wmap}, and by solving many theoretical issues in cosmology. This exponential growth of the scale factor allows cosmologists of the eighties to solve most of the problems of the non-inflationary Big-Bang models as such an additionnal phase explains {\\it e.g.} the curvature of the universe or its homogeneity. Moreover, it provides a mechanism to generate cosmological perturbations in the primordial universe. In most of the inflationary models, only scalar and tensor perturbations are produced. On the one hand, scalar perturbations are the primordial seeds for galaxies and large scale structures formation by gravitational collapse and they leave their footprint on the temperature and E-mode polarization anisotropies of the Cosmic Microwave Background (CMB). On the other hand, tensor perturbations correspond to primordial gravity waves and could be observed via their unique footprint on the undetected yet B-mode anisotropies of the CMB. Different inflationary scenarios, some of them inspired or emerging from quantum gravity settings, {\\it a priori} lead to different statistical properties for this cosmological perturbations, which inevitably lead to different statistical properties of the CMB anisotropies and large scale structures. Inversely, those cosmological observables can therefore be used as probes of inflation making primordial cosmology a laboratory to test high energetic and/or quantum gravity proposals.  Current observations of the CMB anisotropies or of the large scale structures are not precise enough to distinguish between different inflationary scenarios. But many new experiments\\cite{cmb-exp} such as the PLANCK satellite\\cite{planck} or the EBEX balloon-borne experiment\\cite{ebex}, are planned, or already flying to make precise measurements of the statistical properties of both temperature and polarization anisotropies of the CMB. In the perspective of this new class of experiments and taking the fact that quantum gravity is now in the playground of cosmology, it becomes promising to study the impact of quantum gravity proposals on the generation of cosmological perturbations during inflation.  Following this guideline, we consider the influence of L.Q.C. corrections to general relativity on the production of  gravitational waves during inflation. Quite a lot of work has already been devoted to gravitational waves in L.Q.C. \\cite{L.Q.C.gen} and in this paper, we follow the path adopted in \\cite{grainL.Q.G.2,grainL.Q.G.3,grainL.Q.G.4}. Our approach assumes the background to be described by the standard slow-roll inflationary scenario whereas L.Q.C. corrections are taken into account to compute the propagation of tensor modes within such an inflationary background. This approach is heuristically justified (to decouple the physical effects) and intrinsically plausible (as, on the one hand, the L.Q.C.-driven superinflation can only be used to set the proper initial conditions to a standard inflationary stage if the horizon {\\it and} flatness problems are both to be solved \\cite{tsuji} and as, on the other hand, it seems that the quantum bounce can trigger on a standard inflationary phase \\cite{jakub}). ",
        "conclusions": "Testing quantum theories of gravity is probably one of the most important challenge of current fundamental physics. Loop Quantum Gravity has not yet been experimentally probed but it seems that cosmological observations could allow for a clear signature of L.Q.G. effects. In this communication, we have shown how the tensor power spectrum is affected by L.Q.G. corrections via a modified dispersion relation for the propagation of primordial gravity waves in a inflationary background. This constitutes a first step towards a full derivation of the statistical properties of the cosmological perturbations produced in the early universe where L.Q.G. corrections on the background dynamics are also included. Although B-mode detection could be achieved by the PLANCK satellite and is the main goal of several dedicated experiments for the forthcoming decade, the signal remains very difficult to extract from the lensing background and some further refinements are required to quantify the amplitude of the expected L.Q.G. effects. Finally, the formalism established in this article should be used for L.Q.G. corrections to scalar perturbations that are just being investigated \\cite{bojo3} and could be even more promising from the observational viewpoint."
    },
    "0911/0911.1280_arXiv.txt": {
        "abstract": "Many  analyses  and  parameter  estimations  undertaken  in  astronomy require  a  large set  (\\largesample)  of non-analytical,  theoretical spectra, each  of these defined  by multiple parameters.   We describe the  construction of  an \\ND-dimensional  grid which  is  suitable for generating  such spectra.   The  theoretical spectra  are designed  to correspond  to  a targeted  parameter  grid  but  otherwise to  random positions in the parameter space, and they are interpolated on-the-fly through  a  pre-calculated  grid  of  spectra.  The  initial  grid  is designed to  be relatively  low in parameter  resolution and  small in occupied hard disk space and therefore can be updated efficiently when a  new model  is  desired.  In  a  pilot study  of stellar  population synthesis  of  galaxies,  the  mean  square errors  on  the  estimated parameters are  found to decrease  with the targeted  grid resolution. This  scheme of generating  a large  model grid  is general  for other areas of studies, particularly  if they are based on multi-dimensional parameter space and are focused on contrasting model differences. ",
        "introduction": "Various analyses and parameter estimations performed in growing number of applications  in astronomy require  a large set  of non-analytical, theoretical spectra,  where each spectrum in-principle  can be defined by multiple  free parameters.   In stellar population  synthesis using the                          Bayesian                         approach \\citep[e.g.,][]{2003MNRAS.341...33K,2005MNRAS.362...41G,2007ApJS..173..267S}, as many as $10^{4} -  10^{5}$ theoretical spectra and the derived line indices are  used in  order to  cover a wide  range of  star formation histories from  early- to  late-type galaxies.  With  about 5  or more free parameters,  each theoretical spectrum is a  single stellar burst defined by metallicity and age,  superimposed at a given mass fraction on a  spectral component with exponentially  decreasing star formation rate characterized by  the age of the oldest  stars, e-folding time of star formation, and both spectral components could be dust-attenuated. Similarly,  $10^{5} -  10^{6}$ theoretical  spectra are  considered in parameter estimation on stellar spectra \\citep[e.g.,][who considered a model   defined  by   14   parameters]{2007ApJS..169..328R},  and   on \\ion{H}{2}  regions \\citep[e.g.,][who  studied a  model defined  by 15 parameters]{2009MmSAI..80..397M}.  One approach  to prepare and manage a non-analytical model is to  store all of the pre-calculated spectra, fixed at both  the parameter choice and the  parameter resolution.  As models are  becoming more sophisticated and are  growing in varieties, however, updating and storing $10^{5}$  or more spectra may not be the most flexible approach  for many astronomers. To handle  large sets of theoretical spectra at  multi-dimensional parameter space is therefore a question that cannot be neglected.  We  consider here current  integrated stellar  population models  as a case  study.    These  models,  together   with  parameter  estimation techniques  or  other analyses,  were  shown  by  many authors  to  be invaluable for deriving composition and star formation rate/history of different types of galaxies through their observed spectra, regardless whether     or    not     a     large    model     grid    is     used \\citep[e.g.,][]{2000AJ....119.1645T,2001MNRAS.327..849R,2003MNRAS.341...33K,2003MNRAS.343.1145P,2003ApJ...587...55G,2004MNRAS.355..273C,2004ApJ...613..898T,2004MNRAS.351.1151B,2005MNRAS.358..363C,2005MNRAS.362...41G,2006MNRAS.365..385M,2006MNRAS.365...74O,2007MNRAS.378.1550P,2007MNRAS.381..263A,2007MNRAS.381.1252T,2009MNRAS.393..406C,2009MNRAS.tmp.1327R}. Firstly,  there  is  a  variety of  spectral  synthesis  computational programs which use different flavors in initial stellar mass function, stellar     types,      and/or     stellar     evolutionary     tracks \\citep[e.g.,][]{1997A&A...326..950F,1999ApJS..123....3L,2003MNRAS.344.1000B,2005MNRAS.362..799M,2007MNRAS.382..498C}. The variations in input  ingredients exist not only among  the models, but also  within a  single  model in  the  form of  additional freedom  in parameter  choice.  Secondly, increasingly  more parameters  are being included in  the analyses, meaning  that the parameter  space defining the  theoretical   spectra  is  getting  larger.   So   much  so  that \\citet{2009ApJ...694..902L}   have  recently  constructed   a  stellar population model  to the level of individual  element abundances, with the goal  to understand the  effect of each  and every element  to the integrated spectrum of a  stellar population.  However, there are very few studies in the astronomy literature which address the case where a multi-dimensional  parameter  space  defines  a  model.   Two  of  the fundamental questions  are: (1)  how do we  obtain, or  interpolate, a theoretical  spectrum   in-between  the  parameter   values,  that  is originally  unavailable?   (2)  what  is  the best  way  for  such  an interpolation, taking into  account of the noise in  the data, and the possible non-linear dependence \\citep{2009AJ....138.1365V} between the theoretical spectrum  and its underlying physical  parameters?  One of the goals of this work is to address the first question.  Here we describe  a novel approach which increases  the flexibility of handling  a large set  (\\largesample) of  theoretical spectra,  and is general  for  \\ND-dimension (\\ND{D})  parameter  space.  The  approach adopts a  hypercube, an \\ND{D}  analog of a  cube of length  unity for each side \\citep[e.g.,][]{Anthony87}, to represent the parameter space underlying each  and every theoretical spectrum.  The  spectrum at any intermediate parameter point is generated on-the-fly\\footnote{We refer the on-the-fly ability to be  the generation of theoretical spectra at run time  as required  by the actual  computation, e.g.,  in parameter estimation.} through multi-linear interpolation, upon an initial model grid  that  is  designed  to  be lower  in  parameter  resolution  and therefore  can be  updated efficiently.   We examine  a pilot  case in stellar population synthesis of galaxies, to show that the approach is applicable to studies  in astronomy, and well manageable  by a typical personal computer -- a computer with Intel(R) Pentium(R)D 3.20~GHz CPU and 3.19~GHz RAM is used in this work.  All of  the spectra  considered in this  work are expressed  in vacuum wavelength, and are re-sampled  to within the optical wavelength range \\restwaverange \\ at  a resolution of 1~\\AA \\  per wavelength bin.  The Cartesian  coordinate   system  is   used  throughout  this   work  in representing the parameter space. ",
        "conclusions": "\\label{S:summary}  In   applications   where   a   large  (sample   size   \\largesample), non-analytical model  is needed,  the hypercube representation  of the parameter space, together with multi-linear interpolation for deriving the theoretical value at an intermediate grid point, are shown here to be feasible to improve the  flexibility in handling the models.  Using stellar population  synthesis of galaxies  as a pilot study,  we found that the  increase in targeted parameter grid  resolution improves the parameter estimates in terms of the mean square error. This example is given by incorporating  the Bayesian parameter estimation.  Naturally, our  model generating approach  can be  combined with  other parameter estimation techniques, such as the Markov Chain Monte Carlo.  Models of large sample size can be applied to studies in various areas and disciplines.  In the context of spectral analyses, a large grid of theoretical  spectra is  expected  to be  applicable  to mock  catalog construction, studying the  relationship between spectral features and model parameters,  spectral fitting, and  stellar population synthesis of galaxies.   Our approach is particularly suitable  for studies that are  based on multi-dimensional  parameter space,  and are  focused on investigating  differences among  results  obtained through  different models."
    },
    "0911/0911.1249_arXiv.txt": {
        "abstract": "{We constrain the parameters describing the kinematical state of the universe using a cosmographic approach, which is fundamental in that it requires a very minimal set of assumptions (namely to specify a metric) and does not rely on the dynamical equations for gravity.  On the data side, we consider the most recent compilations of Supernovae and Gamma Ray Bursts catalogues. This allows to further extend the cosmographic fit up to $z=6.6$, i.e. up to redshift for which one could start to resolve the low $z$ degeneracy among competing cosmological models.  In order to reliably control the cosmographic approach at high redshifts, we adopt the expansion in the improved parameter $y=z/(1+z)$. This series has the great advantage to hold also for $z>1$ and hence it is the appropriate tool for handling data including non-nearby distance indicators.  We find that Gamma Ray Bursts, probing higher redshifts than Supernovae, have constraining power and do require (and statistically allow) a cosmographic expansion at higher order than Supernovae alone. Exploiting the set of data from Union and GRBs catalogues, we show (for the first time in a purely cosmographic approach parametrized by deceleration $q_0$, jerk $j_0$, snap $s_0$) a definitively negative deceleration parameter $q_0$ up to the $3 \\sigma$ confidence level. We present also forecasts for realistic data sets that are likely to be obtained in the next few years.}  \\begin{document} ",
        "introduction": "A huge amount of cosmological data sets with increasing reliability has been collected during the last decade, providing new insights on the dynamical state of our universe. The interpretation of the Hubble diagram for Type Ia Supernovae (SNeIa) and the observations of polarization and anisotropies in the power spectrum of the Cosmic Microwave Background (CMB) showed that the universe is undergoing an accelerated phase of expansion. Furthermore, the availability of new cosmological probes, such as high redshift SNe and Gamma-Ray Bursts (GRBs), enables to improve the knowledge of the cosmological parameters and allows to distinguish among different models: any new test should be a challenge for all the attempts providing an explanation of the current cosmic acceleration \\cite{exp}. On the other hand, the recent development of a model independent approach (the cosmographic one \\cite{chiba, vc_taylor, vc, Rapetti:2006fv}) gained increasing interest for catching as much information as possible directly from observations, retaining the minimal priors of isotropy and homogeneity and leaving aside any other assumption. Let us stress that the idea supporting cosmography is rather clean: here we are not providing any new model for Dark Energy, but just a deeper, even if simpler, analysis of the cosmological data sets with the aim to give a fiducial frame against which any theoretical proposal should be tested. The only ingredient taken into account \\textit{a priori} is the FLRW line element obtained from kinematical requirements \\be ds^2=-c^2dt^2+a^2(t)\\l[\\frac{dr^2}{1-kr^2}+r^2d\\Omega^2\\r]\\,; \\ee exploiting this metric, it is possible to express the luminosity distance $d_L$ as a power series in the redshift parameter $z$, the coefficients of the expansion being functions of the scale factor $a(t)$ and its higher order derivatives.  In this paper we will discuss the use of luminosity distance indicators in the high redshift universe in order to constrain the values of fundamental cosmological parameters that describe the model we are interested in.  Even though the prominent role of SNe (in the high-$z$ version, too) in doing the job is well-known, the potentiality of GRBs as cosmological standard candles has been recently explored as a possible proposal to increase the number of high redshift distance ladders. Data coming from observations of both SNe and GRBs are used to fit directly the expression for the luminosity distance $d_L$.  Following \\cite{wei}, $d_L$ can be defined from the relation between the apparent luminosity $l$ of an object and its absolute luminosity $L$ \\bea l=\\frac{L}{4 \\pi r^2_1 a^2(t_0)(1+z)^2}= \\frac{L}{4\\pi d_L^2}\\,, \\eea where $r_1$ is the comoving radius of the light source emitting at time $t_1$, $t_0$ is the later time an observer in $r=0$ is catching the photons, and redshift $z$ is, as usual, defined as $1+z=a(t_0)/a(t_1)$. The radial coordinate $r_1$ in a FLRW universe can be written for small distances as \\cite{v1} \\be r_1=\\int_{t_1}^{t_0}\\frac{c}{a(t)}dt-\\frac{k}{3!}\\l[\\int_{t_1}^{t_0}\\frac{c}{a(t)}dt\\r]^3+\\mathcal{O}(5)\\,, \\ee with $k=-1, 0, +1$ respectively for hyperspherical, Euclidean or spherical universe. In such a way, it is possible to recover the expansion of $d_L$ for small $z$ \\bea d_L(z)=&&cH_0^{-1}\\l\\{z+\\frac{1}{2}(1-q_0)z^2-\\frac{1}{6}\\l(1-q_0-3q_0^2+j_0+ \\frac{kc^2}{H_0^2a^2(t_0)}\\r)z^3+\\r.\\nonumber \\\\ && \\nonumber \\\\ &&\\hspace{-1cm}\\l.+\\frac{1}{24}\\l[2-2q_0-15q_0^2-15q_0^3+5j_0+10q_0j_0+ s_0+\\frac{2kc^2(1+3q_0)}{H_0^2a^2(t_0)}\\r]z^4+...\\r\\}\\,, \\eea having defined the cosmographic parameters as: \\bea H_0&\\equiv&\\l.\\frac{1}{a(t)}\\frac{da(t)}{dt}\\right|_{t=t_0}\\equiv\\l.\\frac{\\dot{a}(t)}{a(t)}\\r|_{t=t_0}~,\\label{eq:H0}\\\\ q_0&\\equiv&-\\l.\\frac{1}{H^2}\\frac{1}{a(t)}\\frac{d^2a(t)}{dt^2}\\right|_{t=t_0}\\equiv\\l.-\\frac{1}{H^2}\\frac{\\ddot{a}(t)}{a(t)}\\r|_{t=t_0}~,\\label{eq:q0}\\\\ j_0&\\equiv&\\l.\\frac{1}{H^3}\\frac{1}{a(t)}\\frac{d^3a(t)}{dt^3}\\right|_{t=t_0}\\equiv\\l.\\frac{1}{H^3}\\frac{a^{(3)}(t)}{a(t)}\\r|_{t=t_0}~,\\label{eq:j0}\\\\ s_0&\\equiv&\\l.\\frac{1}{H^4}\\frac{1}{a(t)}\\frac{d^4a(t)}{dt^4}\\right|_{t=t_0}\\equiv\\l.\\frac{1}{H^4}\\frac{a^{(4)}(t)}{a(t)}\\r|_{t=t_0}~.\\label{eq:s0} \\eea  It is important to stress the {existence} of other  definitions for distance indicators \\cite{vc} rather than the more often used $d_L$: depending on which physical quantity one is measuring, it could be more convenient to extract from data one of these instead of the luminosity distance. These different quantities admit different expressions of the Taylor expansion in the redshift, such that it could be more natural to estimate cosmographic parameters, whose expression instead does not depend on the analytic expansion, in one of these particular frameworks. {This ambiguity could lead to a misunderstanding even about the proper definition of distance one should retain}. From now on we will refer only to luminosity distance since it is the most direct choice in the case of measures of distance for SNe and GRBs.  In recent works \\cite{lucariello,wang-dai}, a first attempt to fit luminosity distance by data using more distant objects as ``standard candles'' has been performed. However, the ill-behaviour at high $z$ of the redshift expansion used there is known to strongly affect the results. We are no longer going to use the standard relation linking the luminosity distance to the ordinary-defined redshift. As already pointed out in \\cite{vc_taylor}, the lack of validity of the Taylor-expanded expression for $d_L$ could be settled down approximately at $z\\sim1$. In order to avoid problems with the convergence of the series for the highest redshift objects as well as to control properly the approximation induced by truncations of the expansions, it is useful to recast $d_L$ as a function of an improved parameter $y=z/(1+z)$~\\cite{vc_taylor,CPL}. In such a way, being $z\\in(0,\\infty)$ mapped into $y\\in(0,1)$, we will be in principle able to retrieve the right behaviour for series convergence at any distance. The introduction of this new redshift variable will not affect the definition of cosmographic parameters, while the luminosity distance at fourth order in the $y$-parameter becomes: \\bea d_L(y)=&&\\frac{c}{H_0}\\left\\{y-\\frac{1}{2}(q_0-3)y^2+\\frac{1}{6}\\left[12-5q_0+3q^2_0-(j_0+\\Omega_0)\\right]y^3+\\frac{1}{24}\\left[60-7j_0-\\right.\\right.\\nonumber \\\\ &&\\left.\\left.-10\\Omega_0-32q_0+10q_0j_0+6q_0\\Omega_0+21q^2_0-15q^3_0+s_0\\right]y^4+\\mathcal{O}(y^5)\\right\\}~, \\eea where $\\Omega_0=1+kc^2/H_0^2a^2(t_0)$ is the total energy density. For the flat universe, $\\Omega_0=1$. On the other hand, luminosity distance can also be expressed as ``logarithmic Hubble relations\": \\bea \\ln{\\left[\\frac{d_L(y)}{y~\\rm Mpc}\\right]}=&&\\frac{\\ln{10}}{5}\\left[\\mu(y)-25\\right]-\\ln{y}=\\ln{\\left[\\frac{c}{H_0}\\right]} -\\frac{1}{2}(q_0-3)y+\\frac{1}{24}\\left[21-4(j_0+\\Omega_0)+\\right.\\nonumber \\\\ &&\\left.+q_0(9q_0-2)\\right]y^2+\\frac{1}{24}\\left[15+4\\Omega_0(q_0-1)+j_0(8q_0-1)-\\right.\\nonumber\\\\ &&\\left.-5q_0+2q^2_0-10q^3_0+s_0\\right]y^3+\\mathcal{O}(y^4)~;\\label{eq:y-exp} \\eea and then, the distance modulus (in which are expressed current data) is given by the expression: \\bea \\mu(y)=&&25+\\frac{5}{\\ln{10}}\\left\\{\\ln{\\left[\\frac{c}{H_0}\\right]}+\\ln{y} -\\frac{1}{2}(q_0-3)y+\\frac{1}{24}\\left[21-4(j_0+\\Omega_0)+q_0(9q_0-2)\\right]y^2+\\right.\\nonumber \\\\ &&\\hspace{-.5cm}+\\left.\\frac{1}{24}\\left[15+4\\Omega_0(q_0-1)+j_0(8q_0-1)-5q_0+2q^2_0-10q^3_0+s_0\\right]y^3+\\mathcal{O}(y^4)\\right\\}~. \\label{dm} \\eea  {It is clear that the higher the order reached in the redshift expansion of $d_L$, the better the data fitting will be (more free parameters).  However, for a given data set there will be an upper bound on the order of the expansion which is statistically significant in fitting the data.  In the following sections we shall fit the data with two different truncations of expansion (\\ref{eq:y-exp}) without and with the third order term ($y^3$). For the sake of conciseness we shall respectively label them as Cosmography I and Cosmography II. While we shall see that the higher redshift data make the latter truncation the most statistically significant among the two, we shall show in the end that no improvement is at the moment achievable by extending the truncation to fourth order in $y$.} ",
        "conclusions": "Even though the intriguing era of ``precision cosmology\" has now started, the huge flow of data with increasing resolution is not yet able to solve definitely the contest among different theoretical proposals for the evolution of the universe at late times. The suggestion carried on in this work is a return to a more conservative approach to cosmology, in the meaning of a proper link between observations and cosmographic parameters. In this sense cosmography allows to regain a comprehensive bird's-eye view on the problem, since it gives an unbiased interpretation to the collected data sets.  We have here included high redshift data ($z\\gtrsim1$) to provide new constraint on the main cosmographic parameters. The lack of validity of the Taylor expansion of the luminosity distance in the usual redshift definition has been circumvented, following the insight of \\cite{vc}, by the translation of the distance definition in a new redshift variable $y$, spanning the range $(0, 1)$ whereas $z$ runs into the interval $(0, +\\infty)$; the improvement obtained with this procedure is that now we are able to fit the data of any possible observable candidate to be a cosmological ``standard candle\", no matter what its distance is.  Working with high redshift objects means taking into account both distant SNe (up to $z\\sim1.55$) and GRBs (up to $z\\sim6.6$); in particular, observations of GRBs are quickly approaching even larger redshift (redshifts up to $z\\sim10$ are expected in next few years). This might make of GRBs the crucial players in the near future, as such high redshifts will allow to determine the higher order parameters of the cosmographic series with unprecedented accuracy. Of course, we do not yet understand GRBs so well that we can use them as proper standard candles. However, it has been proposed to use correlations between GRB observables to standardize their energetics \\cite{grb1}. Furthermore, in this work we have also use the insight of \\cite{xiagrb} in order to solve the ``circular problem\" in the determination of the GRB redshifts.  We performed the calculation for two successive truncations of the Taylor expansion of the luminosity distance (up to the third and fourth order in $y$); further terms are, at the moment, not statistically relevant, as shown by running an F-test.  The interest about the possibility to reach the highest possible order is strictly related to the main aim of cosmography: the ability to test and discriminate meaningfully among competing cosmological models. Given that most of the models on the market are build in order to recover (or simulate) Dark Energy at low redshift, their expansion histories are basically degenerate at late times. Breaking such a degeneracy would imply knowledge of the early universe expansion curve and hence would require, within a cosmographic approach, an accurate determination of the higher order parameters jerk $j_0$ and snap $s_0$. This can be achieved only via significant data at high redshift.  Our results can be summarized in few points: we have found that the inclusion of high $z$ GRB data might lead to misleading results for early truncations of the cosmographic series as in Cosmography I.  We showed that the cosmographic series truncated at the fourth order in $y$, Cosmography II, is the most accurate and statistically significant one given current data. The latter allows not only a more accurate determination of the lowest order parameters $q_0$ and $j_0$ but also a first meaningful estimate of the snap parameter $s_0$. All the results are compatible with $\\Lambda$CDM within $2\\sigma$. Most importantly, the complete set of Union and GRB data allows us to introduce for the first time a $q_0$ definitively negative in the limit of the $3 \\sigma$ confidence level.  We have also discussed the future perspective to further ameliorate such results with planned or proposed experiments. In this case we showed that a further reduction of the error bars (by a factor two) will be possible."
    },
    "0911/0911.1555_arXiv.txt": {
        "abstract": "The aim of this talk is to present the most recent advances in establishing plausible planetary system architectures determined by the gravitational tidal interactions between the planets and the disc in which they are embedded during the early epoch of planetary system formation. We concentrate on a very well defined and intensively studied process of the disc-planet interaction leading to the planet migration. We focus on the dynamics of the systems in which low-mass planets are present. Particular attention is devoted to investigation of the role of resonant configurations. Our studies, apart from being complementary to the fast progress occurring just now in observing the whole variety of planetary systems and uncovering their structure and origin, can also constitute a valuable contribution in support of the missions planned to enhance the number of detected multiple systems. ",
        "introduction": "Multiple-planet systems are of particular interest to test theoretical models of planet formation and the evolution of planetary systems. Just a few days before this conference, Jason Wright (Wright {\\em et al.\\/} \\cite{wright}) sent a paper in which he and his colleagues generated a catalogue of the 27 published multiple-planet systems around stars within 200 pc from the Sun to the other participants. The authors made several interesting observations about the systems' statistical properties. At least five of these systems harbour planets which are in or near mean-motion resonance. This fact makes the importance of orbital migration even stronger than immediately after the discovery of giant planets very close to their host stars (e.g. Kley \\cite{kley}; Papaloizou {\\em et al.\\/} \\cite{fivebig}). These arguments, together with the increasing number of known planets with masses in the range of a few Earth masses, have stimulated our study of migration-induced resonances (e.g. Papaloizou \\& Szuszkiewicz \\cite{ps}, Podlewska \\& Szuszkiewicz \\cite{edytas}) in systems containing low-mass planets.  Our present goal is to investigate migration-induced architectures of planetary systems, focusing mainly on resonant configurations. In order to achieve this goal we have performed a series of numerical simulations of planets migrating in resonance. The main tools used in our studies are two-dimensional hydrodynamic simulations, simple analytic modelling and N-body investigations. By constructing a simple analytic model, we were able to verify the reliability of our numerical calculations. By combining the hydrodynamic simulations with the N-body technique we were able to follow the dynamical evolution of the planets for a substantial amount of time, comparable with the estimated lifetimes of gaseous discs. ",
        "conclusions": "Particular attention has been devoted here to the investigation of the occurrence of resonant configurations in systems containing low-mass planets. Studies of commensurabilities have already been applied successfully to analyse the motion of pulsar planets and they proved to be a powerful tool in that context. They have the  potential to be similarly useful in our search for Earth-like planets in other systems. Our prediction of the occurrence of the commensurabilities in such systems may turn out to be particularly valuable in the case of the TTV (Transit Timing Variation) observations, because the differences in the time intervals between successive transits, caused by planet-planet interactions, are largest near mean-motion resonances (e.g. Agol {\\em et al.\\/} \\cite{agol}). Finally, detection of such resonances can also yield useful information about orbital migration as a process operating during planet formation."
    },
    "0911/0911.3550_arXiv.txt": {
        "abstract": "In hybrid inflationary models, inflation ends by a sudden instability associated with a steep ridge in the potential. Here we argue that this feature can generate a large contribution to the curvature perturbation on observable scales. This contribution is almost scale-invariant but highly non-Gaussian.\t The degree of non-Gaussianity can exceed current observational bounds, unless the inflationary scale is extremely low or the hybrid potential contains very large coupling constants. Non-linear effects on small scales may quench the non-Gaussian signal, and while we find no compelling evidence that this occurs, full lattice simulations are required to definitively address this issue.  \\vspace{3mm} \\begin{flushleft} \\textbf{Note added}: We now believe that nonlinear effects will invalidate the original computation in this paper essentially instantaneously after the short-wavelength modes reach the minimum of their potential. This means that the mechanism described in this paper will not lead to appreciatable curvature perturbations on long wavelengths, and no useful constraints on hybrid inflation will result.  We have inserted a brief calculation on p2 of this manuscript to explain this fact, but have otherwise left the manuscript unchanged. \\end{flushleft}   \\vspace{3mm} \\begin{flushleft} \\textbf{Keywords}: Inflation, cosmological perturbation theory, physics of the early universe, quantum field theory in curved spacetime. \\end{flushleft} ",
        "introduction": "\\label{sec:intro}  Recently, a rapid accumulation of data has made clear that the orthodox scenario of structure formation---in which small fluctuations oscillate under the influence of gravity within a almost-uniform primeval plasma---% is in excellent agreement with observation. The initial conditions for these fluctuations are determined by the primordial curvature perturbation, $\\zeta$, usually supposed to have been synthesized during an earlier cosmic epoch, perhaps during inflation or ekpyrosis. Many versions of inflation are characterized by slow, smooth evolution and yield very Gaussian fluctuations. In other versions evolution is a dramatic process, allowing larger non-Gaussianities.  In this paper, we show that significant power and non-Gaussianity can be generated when a set of inflationary trajectories fall from a ridge in field space. Such a process is a crucial element of the hybrid inflationary scenario, introduced by Linde, in which inflation ends through a sudden instability \\cite{Linde:1991km,Linde:1993cn,Copeland:1994vg}. Quantum fluctuations lead to slight differences in field values at the top of the ridge, causing different regions of the universe to follow one of a narrow bundle of trajectories. Trajectories near the edge of the bundle are ejected from the ridge early in their evolution, whereas those in the core remain on the ar\\^{e}te longer. The dispersion generated in this way induces a variation in expansion history (``$\\delta N$'') from trajectory to trajectory within the bundle, resulting in a nearly scale-invariant but highly non-Gaussian spectrum of density perturbations. A careful analysis of this process reveals the surprising fact that the final dispersion in the bundle depends only weakly on the initial dispersion. Hence falling from the ridge can have a large effect on the final curvature perturbation, even if the bundle of trajectories is initially very tightly focused.  Ultimately these contributions to the power spectrum can be traced to the conversion of isocurvature modes into adiabatic ones during descent from the ridge. In the early days of the inflationary paradigm, it was assumed that the curvature perturbation was primarily determined by\t quantum fluctuations in the inflaton field ``freezing out'' as successive physical scales left the cosmological horizon. Later, when inflationary models containing two or more fields were constructed, it was understood that the curvature perturbation could grow or decay as neighbouring horizon volumes diverged along adjacent but inequivalent field-space trajectories. The evolution is caused by light isocurvature degrees of freedom which play no dynamical role while cosmic microwave background (CMB) scales are leaving the horizon, but later come to influence the expansion history of the universe. Although models with isocurvature effects are more complicated, they admit a richer phenomenology. In models with canonically normalized fields the reprocessing of isocurvature modes can lead to a significant non-Gaussian component of $\\zeta$ \\cite{Maldacena:2002vr,Lyth:2005fi,Seery:2005gb}. In the present case, the isocurvature modes are massive while CMB scales leave the horizon and are suppressed by cosmological expansion. Nevertheless, falling from the ridge enormously amplifies these perturbations and can produce a significant effect.  Our analysis applies to the original hybrid scenario and many related models. The qualitative conclusions may apply during a phase of two-field ekpyrosis \\cite{Koyama:2007mg,Koyama:2007ag,Buchbinder:2007at, Koyama:2007if,Lehners:2007wc,Lehners:2008my}, which also depends on dispersion near a ridge. The inflationary production of non-Gaussianity in dispersing models has been source of interest in recent years. Alabidi determined the non-Gaussian yield in a hybrid model, beginning from a range of locations near but not coincident with the ridge \\cite{Alabidi:2006hg}. Byrnes {\\etal} performed a similar analysis, showing that the initial conditions could be tuned to produce a large effect \\cite{Byrnes:2008wi,Byrnes:2008zy}. Related studies have appeared in the literature \\cite{Cogollo:2008bi,Rodriguez:2008hy}.  The generation and evolution of fluctuations in hybrid inflation was studied by Randall, Solja\\v{c}i\\'{c} \\& Guth \\cite{Randall:1995dj}, and later Garc\\'{\\i}a-Bellido, Linde \\& Wands \\cite{GarciaBellido:1996ke,GarciaBellido:1996qt}. More detailed calculations were undertaken by Copeland, Pascoli \\& Rajantie \\cite{Copeland:2002ku}. The field whose instability ultimately ends inflation is termed the \\emph{waterfall field} and at early times its potential is designed so that it is very heavy, causing its perturbations to decay. Only fluctuations orthogonal to the waterfall make an appreciable contribution to the number of e-folds, $N$, required to reach a subsequent surface of uniform energy density. The result is a spectrum of density perturbations which can be evaluated by ignoring the waterfall field. In the simplest scenario, where only one canonically normalized field $\\phi$ is relevant, the single-field formula $\\delta N \\simeq H \\delta \\phi / |\\dot{\\phi}|$ applies. There is an added complication, because the waterfall field rolls very slowly at the epoch of ejection. Therefore fluctuations in $N$ may be large on scales leaving the horizon at that time, and one must verify that there is not an unacceptable synthesis of topological defects or primordial black holes \\cite{GarciaBellido:1996qt}. The original hybrid model produced a blue spectrum of perturbations, but a larger class of so-called ``hill-top'' models \\cite{Ovrut:1984qp,Boubekeur:2005zm} yield red spectra compatible with CMB constraints \\cite{Komatsu:2008hk}.  In this paper, we point out that the traditional calculation of curvature perturbations in this model should be augmented with a new contribution sourced by the waterfall transition. While the trajectories remain on the ridge, fluctuations in $\\delta N$ are well-described by conventional perturbation theory applied to the transverse directions, as explained above. At the epoch of ejection, however, fluctuations in the waterfall field become important, \\emph{no matter how small they are}. To see that this can be so, note that $\\delta N$ measured between the central trajectory (which does not leave the ridge) and any other trajectory (which falls into a non-inflating minimum) tends to infinity. Thus, even exponentially suppressed fluctuations can be amplified. It follows that a tiny fluctuation between neighbouring horizon volumes may ultimately give rise to a large density fluctuation.  The basic mechanism by which large non-Gaussianity is generated can be illustrated using an inverted simple harmonic oscillator. Suppose we consider a particle of unit mass and position $x$, for which the Lagrangian \\begin{equation} L = \\frac{1}{2}\t \\dot x^2 + \\frac{1}{2} \\omega^2 x^2 \\end{equation} describes an inverted harmonic oscillator of natural frequency $\\omega$. We wish to compute the time $t$ required for the particle to roll down to some specified final position, $x_F$, as a function of its initial position $x_0$. We assume that $\\omega x_F \\gg 1$, $x_0 \\ll x_F$, and that the particle has zero initial velocity. At late times $x(t) \\sim (x_0/2) e^{\\omega t}$, so \\begin{equation} t(x_0) = \\frac{1}{\\omega} \\ln \\frac{2 x_F}{|x_0|} \\end{equation} Next we study the statistics of $t$ for an ensemble of particles in which $x_0$ is Gaussian distributed with variance $\\sigma^2$, assuming $\\omega \\sigma \\ll 1$. The mean rolling time is \\be \\langle t \\rangle = \\int_{-\\infty}^{\\infty} \\frac{\\e{-x^2_0 / 2\\sigma^2}}{\\sqrt{2\\pi} \\sigma} \\frac{1}{\\omega} \\ln\t \\frac{2 x_F}{|x_0|} \\; \\d x_0 = \\frac{1}{\\omega} \\left[ \\frac{\\gamma+ 3 \\ln 2}{2} + \\ln \\left( \\frac{x_F}{\\sigma} \\right) \\right] \\ee where $\\gamma$ is the Euler--Mascheroni constant.\tThe average time $\\langle t \\rangle$ depends only weakly on the initial distribution width, $\\sigma$, with narrower distributions taking longer to roll off the ar\\^{e}te.  More surprising results obtain if we compute higher statistics of $t$, such as its variance \\be \\sigma_{t}^2 = \\langle (t\t- \\langle t \\rangle )^2 \\rangle = \\frac{\\pi^2}{8 \\omega^2} \\ee third moment \\be \\mu_3 = \\langle (t - \\langle t \\rangle )^3 \\rangle = \\frac{7\\zeta (3)}{4 \\omega^3} \\ee and fourth moment \\be \\mu_4 = \\langle (t - \\langle t \\rangle )^4 \\rangle =\t \\frac{7 \\pi^4}{64 \\omega^4} \\ee where $\\zeta(z)$ is Euler's zeta function and $\\zeta(3) \\sim 1.202$ is Ap\\'{e}ry's constant. Unlike $\\langle t \\rangle$, these moments are completely independent of the initial variance $\\sigma^2$, and depend only on the curvature of the potential. Therefore, no matter how small the variance associated with $x_0$, the variance in arrival times is the same. Moreover, one typically measures departures from Gaussianity by computing quantities such as the skew $\\mu_3/\\sigma_{t}^3$, and excess kurtosis $\\mu_4/\\sigma_{t}^4 - 3$. In our case these are pure numbers, of order unity, which are completely independent of $\\omega$. Therefore, our purely Gaussian initial distribution always becomes highly non-Gaussian through the process of rolling off the hill.  This toy example can be translated into the inflationary context using a simple dictionary.\tRoughly speaking, the coordinate $x$ represents the waterfall field which mediates the end of hybrid inflation,\t the end point $x_F$ represents the location of the reheating minimum, and the time $t$ represents the number of e-foldings required to reach reheating, at which point inflation ends. The initial Gaussian distribution of release points models the quantum flucutations in the waterfall field just before the hybrid transition. Fluctuations in the e-folding history are related to the curvature perturbation, $\\zeta$.\tTherefore, the moments of $t$ represent the moments of $\\zeta$.\tThe fact that they are large suggests that large non-Gaussianity can be generated by the hybrid waterfall.\tOf course, the full inflationary story is much more complicated, and we will support our conclusions with semi-analytic arguments and numerical calculations using realistic forms of the inflationary potential. However, the basic mechanism is sufficiently general that its essence is captured by the inverted harmonic oscillator.  This mechanism generates power on all scales. In the traditional picture of hybrid inflation, effects occuring near the waterfall transition could only generate significant power on short scales. For example, Randall {\\etal} and Garc\\'{\\i}a-Bellido {\\etal} identified large fluctuations which collapse to black holes or topological defects. These are generated on the horizon scale at the epoch of ejection, yielding a spike in the power spectrum at small wavelengths. In contrast, a careful analysis of the process described above indicates that its power spectrum is nearly scale-invariant. This is possible because the final variance between expansion histories depends very weakly on the initial dispersion of release points, as illustrated by the toy model above. During inflation, long-wavelength fluctuations in the waterfall field are inevitably produced, although these are strongly suppressed by cosmic expansion.\tIn the picture we employ, these tiny fluctuations are encoded in the width of the bundle of field-space trajectories followed by cosmological-scale patches of the universe. After the hybrid transition, the tilt of the power spectrum can be determined by comparing the dispersion of the bundle of trajectories followed by large-scale patches with that of horizon-scale patches. Since the final dispersion of trajectories is very weakly dependent on the initial dispersion, the final dispersion in each bundle is roughly the same. This is the hallmark of a scale-invariant process.  Could this new component in the curvature perturbation become observationally detectable? If the growth of perturbations during the waterfall is not curtailed it will be amplified indefinitely, yielding a bispectrum in conflict with observation. However, any hybrid model automatically contains a cut-off on the amplification which can be achieved. While traversing the ridge, small fluctuations push horizon-scale patches to one side or the other, which roll down and reheat into radiation. This so-called tachyonic preheating \\cite{Felder:2000hj,Felder:2001kt,GarciaBellido:2002aj} is non-perturbative, and does not admit a description in terms of a homogeneous coarse-grained background field with small fluctuations. Beyond some limiting time, of order that required for the field to reach the minimum, the evolution of large-scale modes becomes dominated by radiation produced during the reheating phase rather than the inflationary potential. Because the ridge in field space no longer drives trajectories to diverge, the growth of isocurvature modes from the waterfall ceases.  Ultimately, numerical simulations may be required to determine whether a significant effect can be generated before perturbation theory breaks down. One might have expected that a very long time would be required before isocurvature fluctuations from the waterfall could compete with the adiabatic mode. In fact, our analysis suggests that the field smoothed on large scales typically reaches the reheating minimum only a fraction of an e-fold after the field smoothed on the horizon scale. Whether a given model can generate acceptably Gaussian fluctuations depends on a subtle competition between the perturbative growth of isocurvature modes and a non-perturbative cut-off from tachyonic preheating. In our opinion, it is not clear that the outcome can be decided on the basis of perturbative calculations. Our aim is to show that there is serious interest in obtaining the answer. We hope that the quantitative intricacies of the competition can eventually be settled using lattice simulations.  The structure of this paper is as follows. In \\S\\ref{sec:background}, we fix our notation by reviewing familiar material concerning inflationary perturbations and the $\\delta N$ formalism. In \\S\\ref{sec:evolution} we give a phenomenological description of evolving field fluctuations near a sharp ridge, and use this to derive some approximate expressions for the spectrum and bispectrum sourced in the curvature perturbation. In \\S\\ref{sec:validity} we discuss the validity of our calculation. We work in natural units where $c = \\hbar = 1$, and the Planck mass associated with Newton's gravitational constant is set to unity, $\\Mp = (8 \\pi G)^{-1/2} = 1$, although $\\Mp$ will occasionally be restored to clarify the relative magnitude of terms. ",
        "conclusions": "\\label{sec:conclusions}  We have argued that the transition ending hybrid inflation can generate a large contribution to the curvature perturbation. This contribution arises from fluctuations in the ``waterfall'' field, $\\chi$, which are created as modes exit the inflationary horizon. These perturbations are subsequently suppressed by the cosmological expansion. Although these fluctuations are very small when inflation ends, they are amplified by the tachyonic instability in $\\chi$ and can provide a significant contribution to the curvature perturbation. To calculate the enhancement we use a gradient expansion in which we treat the moments of the initial distribution non-perturbatively. This method is similar to the method of ``moment transport'' which has recently been used to calculate the curvature perturbation in inflationary models \\cite{Mulryne:2009kh}. Here, the procedure is simplified owing to our ability to calculate an approximate analytic solution. This technique predicts that the fluctuations generated by the waterfall are highly non-Gaussian, and can easily exceed observational constraints. If present in a given model, these effects lead to new constraints on viable hybrid models together with a unique observational signature.  There are two phenomena which enable the tiny fluctuations in $\\chi$ to generate significant effects. First, the variance in $\\delta N$ following completion of the waterfall depends extremely weakly on the variance in $\\chi$ before the waterfall begins.  While surprising, this phenomenon is sufficiently general that it can be illustrated using an inverted simple harmonic oscillator. Similar behaviour is seen in the semi-analytic and numerical examples given above. Therefore, even very tiny fluctuations in $\\chi$ can lead to large fluctuations in $\\delta N$.  This means that the exponential suppression of $\\chi$-fluctuations during inflation does not guarantee that these fluctuations are harmless when hybrid inflation completes. In a single-field model, this kind of amplification would be forbidden because the curvature perturbation $\\zeta$ would be conserved on super-horizon scales. Since hybrid inflation necessarily involves multiple fields, the presence of isocurvature modes means that $\\zeta$ is no longer conserved. As we have argued above, the hybrid waterfall can cause it to grow significantly.  Second, the spectrum of fluctuations in $\\delta N$, and hence the spectrum of the curvature perturbation $\\zeta$, is nearly scale-invariant.\tNaively, one might have expected that any process operating at the end of inflation could only affect modes of wavelength smaller than the horizon. This intuition is borne out for some processes: for example, the production of black holes or defects at the end of hybrid inflation results in a very blue power spectrum, which would be unimportant on cosmological scales.  In the present case, it is true that $\\chi$ has a strongly blue power spectrum immediately prior to the hybrid transition. This is because fluctuations in $\\chi$ are suppressed during inflation. However, the waterfall process ensures that the final fluctution in $\\delta N$ is nearly independent of that in $\\chi$.  Therefore, the \\emph{final} curvature fluctuation is nearly the same, independently of scale, even though it is seeded by fluctuations which, on long wavelengths, are much smaller. Thus, the hybrid transition can reprocess the blue spectrum of $\\chi$ into a nearly scale-invariant one. The key difference in behaviour between large and small scales is the transit time from the onset of the waterfall to kinetic domination.  The analysis we have presented here depends crucially on the validity of the separate universe picture while CMB scales are being ejected from the ar\\^{e}te. On small scales, it is known that the hybrid transition creates large (nonlinear) gradients in the scalar fields. This rapid growth in gradients is responsible for the ``tachyonic preheating\" effect studied by Felder {\\etal} \\cite{Felder:2000hj,Felder:2001kt}. It is possible that short wavelength (horizon-size) perturbations reach this nonlinear regime very rapidly, before significant curvature perturbations are generated on longer wavelength (cosmological-size) scales. If so, this would prevent the dispersion of trajectories and lead to effects which are much weaker than those estimated here. An estimate of the onset of this gradient instability can be obtained by studying the time taken for different scales to reach the kinetic-dominated regime, at which time the curvature pertubation $\\zeta$ ceases to evolve. Our estimates indicate that long and short wavelength perturbations reach the kinetic-dominated regime at nearly the same time, within a tiny fraction of an e-folding from the onset of the hybrid transition. The hierarchy between the transit time for long and short wavelengths is typically a number of order unity. In this regime a subtle interplay exists between effects on short and long length scales, and the condition for validity of the separate universe picture is not known. We expect that a definitive resolution of these questions requires numerical simulations of the hybrid transition. We conclude that perturbation theory provides no compelling reason to believe that waterfall contributions to the curvature perturbation should be strongly suppressed.  Even if tachyonic preheating prevents the generation of a significant curvature perturbation on CMB scales, it seems impossible to prevent the appearance of strong non-linearities on the shortest scales --- those associated with the horizon scale at the transition. If these fluctuations are large, they will collapse to form black holes or topological defects. The analysis we have presented indicates that the formation of these non-perturbative objects should be taken to proceed from very non-Gaussian initial conditions. It is known that the final mass fraction contained in such collapsed objects can depend on the detailed profile of initial fluctuations \\cite{Bullock:1996at,Ivanov:1997ia,Hidalgo:2007vk}, which potentially provides another means to constrain models containing a waterfall transition.  \\ack DS would like to thank the Center for Particle Cosmology at the University of Pennsylvania for their hospitality during the completion of this work.\tWe would also like to thank Christian Byrnes, Justin Khoury, Louis LeBlond, Jean-Luc Lehners, Paul Steinhardt, and Mark Trodden for useful comments and discussions. DS is supported by STFC. DM acknowledges support from the Centre for Theoretical Cosmology.  \\appendix"
    },
    "0911/0911.1963.txt": {
        "abstract": "%context { Long-term precise Doppler measurements with the CORALIE spectrograph revealed the presence of two massive companions to the solar-type star HD202206. Although the three-body fit of the system is unstable, it was shown that a 5:1 mean motion resonance exists close to the best fit, where the system is stable. It was also hinted that stable solutions with a wide range of mutual inclinations and low O-C were possible. } %aims { We present here an extensive dynamical study of the HD202206 system aiming at constraining the inclinations of the two known companions, from which we derive possible ranges of value for the companion masses. } %methods { We consider each inclination and one of the longitude of ascending node as free parameters. For any chosen triplet of these parameters, we compute a new fit. Then we study the long term stability in a small (in terms of O-C) neighborhood using Laskar's frequency map analysis. We also introduce a numerical method based on frequency analysis to determine the center of\tlibration mode inside a mean motion resonance. } %results { We find that acceptable coplanar configurations (with low $\\chi^2$ stable orbits) are limited to inclinations to the line of sight between $30\\degr$ and $90\\degr$. This limits the masses of both companions to roughly twice the minimum: $m_b \\in [16.6\\,M_{Jup};\\,33.5\\,M_{Jup}]$ and $m_c \\in [2.2\\,M_{Jup};\\,4.4\\,M_{Jup}]$. Non coplanar configurations are possible for a wide range of mutual inclinations from $0\\degr$ to $90\\degr$, although $\\Delta\\Omega = 0 [\\pi]$ configurations seem to be favored. We also confirm the 5:1 mean motion resonance to be most likely. In the coplanar edge-on case, we provide a very good stable solution in the resonance, whose $\\chi^2$ does not differ significantly from the best fit. Using our method to determine the center of libration, we further refine this solution to obtain an orbit with a very low amplitude of libration, as we expect dissipative effects to have dampened the libration. } %conclusion {} ",
        "introduction": "The CORALIE planet-search program in the southern hemisphere has found two companions around the HD202206 star. The first one is a very massive body with $17.5\\mjup$ minimum mass \\citep{Udry_etal_2002}, while the second companion is a $2.4\\mjup$ minimum mass planet \\citep{Correia_etal_2005}. The parent star has a mass of $1.044$ solar masses, and is located $46.3$ pc from the Solar System. The HD202206 planetary system is an interesting case to investigate the brown dwarf desert since the more massive companion  can be either a huge planet (formed in the circumstellar disk) or a low-mass brown dwarf candidate. %    A more precise mass determination than the minimum given by radial velocity %    technique is however required to delve further into this issue , but %    unfortunately the distance of HD202206 (46.3 pc) prevents the HIPPARCOS %    astrometric data from constraining its visual orbit.  \\citet{Correia_etal_2005} found that the orbital parameters obtained with best fit for the two planets leads to a catastrophic events in a short time (two keplerians fit and full three-body fit alike). This was not completely unexpected given the very large eccentricities ($0.435$ and $0.267$) and masses of the two planets. Using frequency analysis \\citep{Laskar_1990,Laskar_1993,Laskar_1999} they performed a study on the global dynamics around the best fit, and found that the strong gravitational interactions with the first companion made the second planet evolution very chaotic, except for initial conditions in the 5:1 mean motion resonance. Since the associated resonant island actually lies close to the $\\chi^2$ minimum value of the best fit, they concluded that the system should be locked in this 5:1 resonance. Later on, Gozdziewski and co-workers also looked for stable solutions in this system using their GAMP algorithm \\citep{Godziewski_etal_2006}. They provided two possible solutions (among many others), one coplanar, and one with a very high mutual inclination.  Since then, new data have been acquired using the CORALIE spectrograph. A new reduction of the data changed some parameters, including the mass of the HD202206 star. A new fit assuming a coplanar edge-on configuration was derived from the new set of radial velocity data. This new solution still appears to be in the 5:1 mean motion resonance, but is also still unstable. The most striking difference from \\citet{Correia_etal_2005} is the smaller eccentricity of HD202206\\,c.  In the present work, we will continue in more detail the dynamical study started in \\citet{Correia_etal_2005}, using the new fit as a starting point. We also aim to find constraints on the orbital parameters of the two known bodies of this system, in particular the inclinations (and thus the real masses of the planets). \\citet{Godziewski_etal_2006} already showed that stable fits could be obtained with different inclinations, using a particular fitting genetic algorithm that adds stability computation to select its populations (GAMP). Although very effective in finding a stable fit, this algorithm cannot find all possible solutions. We prefer here an approach which separates the fitting procedure from dynamical considerations, as it allows for a better assessment of the goodness of the fit, and whether the model is a good description of the available data. The trade-off is a more difficult handling of the high number of parameters.  %%%%%%%%%%%%%%% % PLAN %%%%%%%%%%%%%%%  We briefly present the new set of data in Sect.~\\ref{sec:data}, the numerical methodology in Sect.~\\ref{sec:method}. We review in details the dynamics in the coplanar edge-on case in Sect.~\\ref{sec:cop_edge}. We then release the constraint on the inclination of the system from the line of sight in Sect.~\\ref{sec:cop}, and finally briefly investigate the mutually inclined configurations in Sect.~\\ref{sec:mutual}. ",
        "conclusions": "%  \\subsection{Mass limits} %TODO %  \\subsection{Mass limits}   Assuming that the system is coplanar, we performed a systematic study of the dynamics of the system for different inclinations to the line of sight. We are able to find constraints for the inclination to the line of sight: $30\\degr \\le i \\le 90\\degr$. This means that the companions' masses are most likely not greater than twice their minimum values: \\begin{itemize} \\item $1 \\le m_b / m_b^{(0)} \\le 2$ \\item $1 \\le m_c /m_c^{(0)} \\le 2$ \\end{itemize} We also studied the influence of mutual inclination for two different inclinations of the planet $b$ ($i_b = 90\\degr$ and $i_b = 30\\degr$), but did not find any clear correlation other than that retrograde potential stable solution consistent with the radial velocity data seem to be limited to mutual inclinations close to $180\\degr$ (i.e. nearly coplanar orbits). As \\citet{Godziewski_etal_2006}, we find possible stable solutions with low $\\chi^2$ for a wide range of mutual inclinations between $0\\degr$ and $90\\degr$. The current data cannot yield more precise constraints. Also the masses determination is dependent on the stellar mass which is not well established. %    But ratio should hold \\textbf{develop a bit}.  %  \\subsection{Mean motion resonance} Although all published dynamical studies of HD202206 suggest that $b$ and $c$ are in a 5:1 mean motion resonance, it is still a debated question. For instance, \\citet{Libert_Henrard_2007} assume that it is just close to the commensurability. Libration of $\\Delta\\varpi$ occurs for particular initial values of this angle, providing a stabilizing mechanism outside the mean-motion resonance, not far from the best fits. In most cases other than retrograde coplanar configurations, those orbits in near-commensurability are worse solutions than the ones in resonance, but they could be more probable if the eccentricities are overestimated (especially for $b$). We find that all significant stable zones with the best O-C are in the 5:1 mean motion resonance. In fact the minimum $\\chi^2$ is almost always in the resonance or very close to it, and stable orbits in the resonance can be found with $\\chi^2$ not significantly higher than the best fit. In addition the O-C level curves tend to follow roughly the resonant island, even though the agreement is not as perfect as for the HD45364 system \\citep{Correia_etal_2009}. This is an improvement from \\citet{Correia_etal_2005} where the best fit lay outside the resonant island, and the $\\chi^2$ had to be degraded to find a stable solution. We thus believe that the resonant configuration is the most probable. We provide a stable solution (\\textbf{S2}, Table \\ref{tab:90_stable}) in the coplanar edge-on case. This solution shows a high amplitude resonant mode in the libration of the critical angle. We believe that this resonant mode is probably dampened by dissipative processes. We use frequency analysis to find a tore on which such orbits exist. Although the specific orbit we give in Table \\ref{tab:s3} does not have a very low $\\rsc$ at 1.55, we expect that the true orbit will be close to it with a low libration amplitude.   Note that for retrograde configurations, the picture is quite different. The best fit lies in a very stable region just outside the mean motion resonance. While these orbits are valid candidates from the dynamical and the observational points of view, we do not favor them as the formation of these systems is hard to explain.   We investigated the possibility of undetected companions. We found that planets with masses smaller than approximately one Neptune mass can exist for semi-major axis lower than 0.12 AU. $0.5\\,M_{Jup}$ planets are also possible beyond 6.5 AU. No planets are possible between 0.12 AU and 6.5 AU as they would be unstable. The two planets model may prove to be wrong in the future, but these hypothetical new companions should not have a big impact on the already detected ones.  %  \\subsection{Discussion} %TODO %\\textsl{ %    \\begin{itemize} %     \\item discussion sur le fait que notre procedure de fit reste local autour du %    best fit du cas coplanaire edge-on. on loupe peut etre d autre config. poss. %     \\item peut-on expliquer e1 par une autre planete en res 1:2 ? (Anglada-Escude, Lopez-Morales, Chambers 2008) %     \\item Formation: Nelson: protoplanet autour d une binaire eccentrique ($e>0.2$), %    la migration s arrete a la resonance 4:1 (le disque deviens eccentrique) %    $q>0.02$ (20Mjup) suffit pour exciter le disque (ici on a $q>0.015$) (Papaloizu, Nelson, Masset 2001). Mais le systeme final a $e_{bin}<0.2$ (ici 0.4) %    \\end{itemize} %}"
    },
    "0911/0911.0309_arXiv.txt": {
        "abstract": "{Magneto-rotational instability (MRI) has been suggested to lead to a rapid growth of the magnetic field in core collapse supernovae and produce departures from spherical symmetry that are important in determining the explosion mechanism.} {We address the problem of stability in differentially rotating magnetized proto-neutron stars at the beginning of their evolution.} {To do this we consider a linear stability taking into account non-linear effects of the magnetic field and strong gravity.} {Criteria for MRI are derived without simplifying assumptions about a weak magnetic field. In proto-neutron stars, these criteria differ qualitatively from the standard condition $d \\Omega/d s <0$ where $\\Omega$ is the angular velocity and $s$ the cylindrical radius. If the magnetic field is strong, the MRI can occur only in the neighbourhood of the regions where the spherical radial component of the magnetic field vanishes. The growth rate of the MRI is relatively low except for perturbations with very small scales which usually are not detected in numerical simulations. We find that MRI in proto-neutron stars grows more slowly than the double diffusive instability analogous the Goldreich-Schubert-Fricke instability in ordinary stars.} {} ",
        "introduction": "There is a growing amount of evidence that core-collapse supernovae are asymmetric and that the core-collapse mechanism itself is responsible for the asymmetry (see Buras et al. 2003, Akiyama et al. 2003 for more details). Several possibilities are explored to account for this observed asymmetry. One is associated with the influence of rotation on convection, which seems to be inevitable during the early evolution of proto-neutron stars (PNS). Convective motions in PNSs are very fast ($\\sim 10^{8}-10^{9}$ cm/sec) and, therefore, the convective turnover time is short, $\\sim 1-10$ ms (see, e.g., Burrows \\& Lattimer 1986). Nevertheless, if angular momentum is conserved, the collapsing core can spin up to very short periods $\\sim 5-10$ ms and generate strong differential rotation. Such fast rotation modifies convection and makes convective motions anisotropic and constrained to the polar regions (Fryer \\& Heger 2000, Miralles et al. 2004). This mechanism is a natural way to create anisotropic energy and momentum transport by convective motions, that only requires that the angular velocity be of the order of the Brunt-V\\\"ais\\\"al\\\"a frequency.  The other possibility to create asymmetry is the effect of jets (see, e.g., Khokhlov et al. 1999, Wheeler et al. 2002). Even though the mechanism of jet formation is still unclear, it seems that MHD jets are common in systems where a central body accretes matter with angular momentum and magnetic field (see, e.g., Meier at al. 2001), and a core-collapse supernova is such a system. Calculations have established that nonrelativistic axial jets originating within the collapsed core can initiate a bipolar asymmetric supernova explosion that is consistent with observations (Hwang et al. 2000).  Another way to generate asymmetry is associated with the magnetic field that can be an important ingredient of the explosion mechanism (Bisnovatyi-Kogan 1971, Kundt 1976). The toroidal magnetic field can be amplified by differential rotation to such high values that it becomes dynamically important (Ardelyan et al. 2005). The effect of the magnetic field on asymmetry of supernovae was considered by Wheeler et al.(2000, 2002) who found that it is possible to produce both a strong toroidal field and an axial jet. Two-dimensional MHD simulations of core collapse indicate that the shape of shock waves and the neutrinosphere can be modified by the effect of the magnetic field (Kotake et al. 2004, Takiwaki et al. 2004).  The possible presence of a magnetic field and differential rotation in a core-collapse supernova favours magnetorotational instability (MRI), which can enhance turbulent transport and amplify the magnetic field. This instability was considered by Akiyama et al.(2003) in the context of core collapse. The authors argued that instability must occur in core collapse and that it has the capacity to produce fields that are sufficiently strong to affect, if not cause, the explosion. Thompson et al. (2005) constructed one-dimensional models, including rotation and magnetic fields, to study the mechanism of energy deposition. They explored several mechanism for viscosity and argue that turbulent viscosity caused by the MRI can be most effective. Numerical simulations provide contradictory conclusions regarding the importance of MRI in core collapse. Moiseenko et al. (2006) claim that MRI has been found in their 2D simulations, and that it is responsible for a strong amplification of the poloidal magnetic flux. However, what these authors call MRI is different to standard MRI considered by Velikhov (1959) (see also Balbus \\& Hawley (1991)). For instance, the instability found by Moiseenko et al (2006) starts to develop only when the ratio of the toroidal and poloidal fields reaches a value of $\\sim$ a few tens. On the other hand, the onset of standard MRI in 2D does not depend on the toroidal field at all. The dependence on the ratio of the toroidal and poloidal field is more typical for Tayler instability (Tayler 1973), which is more relevant to the topology of the magnetic field than differential rotation. Therefore, it is possible that Moiseenko et al. (2006) incorrectly identify instability, and MRI does not occur in their simulations. Apart from that, Moiseenko et al. (2006) attributed a rapid growth of the toroidal and poloidal fields to a dynamo driven by the magnetorotational instability. This also rise some doubts because of Cowlings's anti-dynamo theorem which states that an axisymmetric dynamo cannot exist (see, e.g., Shercliff 1965). Two-dimensional simulations of core collapse with a strong magnetic field have been performed by Sawai et al. (2005, 2008). They found that the magnetic field can play an important role in the dynamics of the core only if the poloidal field of the progenitor is strong enough ($\\sim 10^{12}-10^{13}$ G), but MRI was not seen in the considered models. The magnetic field is amplified mainly by field compression and field wrapping in these simulations. On the contrary, Shibata et al. (2006) claim that they found MRI in their simulations of magnetorotational core collapse in general relativity. These authors paid attention to resolution in order to resolve unstable MRI modes and they claim that amplification of the magnetic field in the considered models is caused by MRI. Note, however, that the poloidal field obtained in their models is of the order of that estimated from conservation of magnetic flux in core collapse, and a more refined analysis is required to determine the mechanism of amplification. Fryer \\& Warren (2004) argued that it is difficult to produce magnetic fields in excess of $10^{14}$ G even if MRI occurs in core collapse because rotation is not sufficiently fast. A detailed study of the magnetorotational core collapse has been performed by Obergaulinger et al. (2006a, 2006b, 2009). The initial magnetic field was purely poloidal in their models with a strength ranging from $10^{10}$ to $10^{13}$ G. Such fields are much higher than those estimated to exist in realistic stellar cores, but the authors wanted to investigate the principal effects of a magnetic field. The initial magnetic field is amplified by differential rotation in these simulations, giving rise to a strong toroidal component. The poloidal component grows mainly by compression during collapse and does not change significantly after core bounce. The authors also found that extended regions exist where the criterion of MRI is satisfied at various epochs. However, the growth rate of this instability is typically too small, except for a few models with a strong initial magnetic field $B=10^{12}$ G.  In this paper, we study the effect of MRI on core-collapse supernova. Since the magnetic field can be sufficiently strong, the criterion of instability is derived taking into account terms depending on the Alfven frequency. We derive the criterion that applies to any rotation profile but a special consideration is made for the case of shellular rotation that often is used to mimic rotation of proto-neutron stars. We address only the axisymmetric instability because numerical simulations of a magnetic core-collapse are usually done in 2D. The main goal of this study is to show that the effect of MRI on proto-neutron stars often is overestimated. ",
        "conclusions": "Interest in MRI  in core collapse supernova is due to the fact that it can generate a strong magnetic field and produce significant departures from spherical symmetry. These departures are crucial for the explosion mechanism. In the present paper, we have derived criteria of MRI in proto-neutron stars taking into account the effect of a non-axial magnetic field. Criteria have been obtained in a form analogous to the Solberg--H{\\o}iland criteria but including terms containing the magnetic field. It turns out that the criterion of MRI in proto-neutron stars can differ from the standard condition $\\partial \\Omega / \\partial s < 0$ even in a weak magnetic field due to strong gravity and the gradient of the lepton fraction. If the Alfven frequency $\\omega_{A0}$ is small compared to the angular velocity then the instability occurs in the region where the component of $\\nabla \\Omega^2$ perpendicular to $\\vec{C}= -(\\beta/T) \\Delta \\nabla T + \\delta \\nabla Y$ has a negative projection on $\\vec{e}_s$ whereas the component $\\nabla \\Omega^2$ parallel to $\\vec{C}$ is unimportant for instability. In the case of slow rotation, when departures from sphericity are small ($g \\gg s \\Omega^2$), the criterion reduces to a simple inequality $\\partial \\Omega/ \\partial \\theta < 0$ (see, e.g., Urpin 1996). For instance, shellular rotation with $\\Omega= \\Omega(r)$ which is often used in modelling proto-neutron stars does not satisfy this condition. Therefore, MRI does not occur if the proto-neutron star rotates shellularly and the wavelength of perturbations is such that the Alfven frequency is smaller than $\\Omega$. Only detailed calculations of rotational core collapse can give the answer to whether the condition of MRI is fulfilled in proto-neutron stars and, generally, this answer should depend on rotation of the progenitor. Note that the velocity of unstable perturbations is approximatelly perpendicular to the radius because gravity strongly suppresses motion in the radial direction and, likely, the radial turbulent transport should be suppressed when the instability saturates.  However, MRI can arise in proto-neutron stars even if the neccesary condition that determines the onset of instability in a weak field is not fullfied. This occurs if the magnetic field is very strong or the wavelength of perturbations is small such that $\\omega_{A0} > \\Omega$. The latter inquality is equivalent to \\begin{equation} \\lambda < \\lambda_c =2 \\pi B/ \\Omega \\sqrt{4 \\pi \\rho} \\approx 1.8 \\times 10^3 B_{13} \\Omega_3^{-1} \\rho_{14}^{-1/2} \\;\\; {\\rm cm}, \\end{equation} where $\\lambda =2\\pi/k$ is the wavelength, $B_{13}=B/10^{13}$G, $\\Omega_3 = \\Omega/10^3$ s$^{-1}$, $\\rho_{14}= \\rho/10^{14}$ g/cm$^3$. For perturbations, satisfying condition (35), the instability can occur if $\\Omega_s < 0$ but only in the neighbourhood of the cone (or cones depending on the field topology) in which the radial component of $\\vec{B}$ is vanishing. The opening angle of the unstable region around this cone is of the order of $\\Omega/\\omega_0$ (see Eq.(34)). Note that unstable motion in this case are perpendicular to the radius as well since gravity suppresses radial motions. The reason why the instability is not entirely suppressed even by a very strong magnetic field is qualitatively very simple. Perturbations with $\\vec{k}=(k_r, 0, 0)$ cannot be suppressed by gravity, and such perturbations do not feel the stabilizing influence of the magnetic field in the region where $B_r \\approx 0$ as well, because this influence  is proportional to the Alfven frequency ($\\propto \\vec{k} \\cdot \\vec{B}$).  It appears that the importance of MRI in core collapse can be overappreciated since its growth rate is relatively low and reaches the value $\\sim \\Omega$ only for perturbations with a wavelength $\\sim \\lambda_c$ (see Eq.(35)). Apart from this, gravity strongly suppresses development of any perturbations with non-radial wavevectors. Only perturbations with $k \\approx k_r$ can be unstable but hydrodynamic motion for such perturbations has a small radial component and turbulent transport should be inefficient radially. Also, it is possible that  MRI occurs only in not very extended regions of the proto-neutron star that may diminish substantially its effect on core collapse. Note that since the wavelength $\\lambda_c$ is small, one needs a very high resolution in numerical simulations to see the most rapidly growing modes (Obergaulinger et al. 2009). Perturbations with longer $\\lambda$ grow substantially slower.  Our conclusion is in contrast to the results obtained by Masada et al. (2007) who considered axisymmetric and nonaxisymmetric magnetorotational instability of PNSs taking into account dissipative processes. These authors also used a local approximation in the stability analysis. In the local analysis, however, the shape of perturbations is assumed to be unchanged and, therefore, this approach applies only if the growth rate of instability is greater than the rate with which perturbations change their shape. In the case of MRI, the growth rate is smaller than (or, at maximum, comparable to) $s |\\nabla \\Omega|$ whereas perturbations change their shape because of differential rotation with a rate $\\sim s |\\nabla \\Omega|$ (see, e.g., Balbus \\& Hawley 1992 for details). Therefore, the results obtained for a nonaxisymmetric MRI in the local approximation raise some doubts. As far as axisymmetric instability is concerned, Masada et al. (2007) were confused when identifying different modes in the dispersion equation. In the dissipative case, they obtained a dispersion equation of the seventh order which describes seven different modes. Certainly, different modes can be unstable in different conditions, but the authors considered only one criterion (Eq.(32) of their paper). Unfortunately, this criterion does not correspond to MRI but describes a secular instability that is a magnetic analogy of the well known Goldreich-Schubert-Fricke (GSF) instability. The magnetic analogy of this instability was first considered by Urpin (2006) for the case of ordinary stars, and Masada et al. (2007) obtained the criterion of the same instability modified for the conditions of PNSs. In contrast to MRI, this instability is dissipative and disappears if diffusive coefficients go to zero. Stratification has a weak impact on this dissipative instability, and it can arise in PNSs. Note, however, that diffisive processes can reduce the stabilizing influence of stratification on the MRI as well.   As was noted by Masada et al. (2007), diffusion of heat and leptons can influence buoyancy if the characteristic diffusion timescale is shorter than the buoyant frequency $\\omega_0$. Since the diffusion timescale is determined mainly by lepton diffusion which is slower than thermal diffusion, this condition is equivalent to $\\xi k^2 \\geq \\omega_0$ where $\\xi$ is the coefficient of lepton diffusivity. In deep layers with a density $\\sim 10^{14}$ g/cm$^3$, the value of $\\xi$ is $\\sim 3 \\times 10^7$ cm$^2$s$^{-1}$ (see Eq.(45) by Masada et al. (2007)). Hence, diffusion begins to suppress the stabilizing effect of stratification when the wavelength of perturbations is shorter than $\\sim 2 \\pi \\sqrt{\\xi / \\omega_0} \\sim 10^3$ cm (we suppose $\\omega_0 \\sim 10^3$ s$^{-1}$). The effects considered in this paper are important for longer wavelengths for the standard pulsar magnetic field.   As was mentioned above, MRI is not the only instability caused by differential rotation in proto-neutron stars. The analogy of the GSF dissipative instability can occur in both magnetic and non-magnetic PNSs. In non-magnetic stars, this instability arises if the angular velocity depends on the vertical coordinate $z$ (Urpin 2007), and the shellular rotation is a particular case of such rotation. Since neutrino transport in PNSs is much more efficient than radiative transport in ordinary stars, the dissipative instability can be rather fast. If $\\Omega_z= \\partial \\Omega / \\partial z \\sim \\Omega/ s$, the condition of the instability reads \\begin{equation} \\frac{|s \\Omega_z|}{\\sqrt{\\Omega_e^2}} \\sim \\Omega \\gg 2 \\omega_0 \\sqrt{\\frac{\\nu}{\\kappa_T}}, \\end{equation} where $\\nu$ and $\\kappa_T$ are the coefficients of viscosity and thermal diffusivity. The difference to the standard criterion of GSF instability ($\\partial \\Omega/ \\partial z \\neq 0$) is caused by the more complex character of neutrino trancport that involves diffusion of both heat and the lepton fraction. This instability can arise soon after collapse even in dense central regions of the star where convection does not occur. It turns out that in all cases, when dissipative instability arises, its growth rate is higher than that of MRI. Indeed, the maximum growth rate of MRI is $\\sim \\Omega$ and can be reached for perturbations with  wavelength $\\sim \\lambda_c$. The growth rate of the dissipative instability is approximately $\\kappa_T k^2 (\\Omega / \\omega_0)^2$. The maximum wavevector for which the instability occurs is $\\sim \\sqrt{\\Omega/\\nu}$. For perturbations with such wavevector, the growth rate is of the order of $\\Omega (\\kappa_T /\\nu)(\\Omega/\\omega_0)^2$ and is always higher than the growth rate of  MRI if condition (36) is satisfied. Therefore, dissipative instability seems to be  more efficient in proto-neutron stars than MRI."
    },
    "0911/0911.1738_arXiv.txt": {
        "abstract": "A number of substellar companions to evolved cool stars have now been reported.  Cool giants are distinct from their progenitor Main Sequence (MS) low-mass stars in a number of ways.  First, the mass loss rates of cool giant stars are orders of magnitude greater than for the late-type MS stars.  Second, on the cool side of the Linsky-Haisch ``dividing line'', K and M giant stars are not X-ray sources, although they do show evidence for chromospheres.  As a result, cool star winds are largely neutral for those spectral types, suggesting that planetary or brown dwarf magnetospheres will not be effective in standing off the stellar wind.  In this case one expects the formation of a bow shock morphology at the companion, deep inside its magnetosphere.  We explore radio emissions from substellar companions to giant stars using (a) the radiometric Bode's law and (b) a model for a bow shock morphology.  Stars that are X-ray emitters likely have fully ionised winds, and the radio emission can be at the milli-Jansky level in favorable conditions. Non-coronal giant stars produce only micro-Jansky level emissions when adjusted for low-level ionisations.  If the largely neutral flow penetrates the magnetosphere, a bow shock results that can be strong enough to ionise hydrogen.  The incoherent cyclotron emission is sub-microJansky.  However the long wavelength radio emission of solar system objects is dominated by the cyclotron maser instability (CMI) mechanism.  Our study leads to the following two observational prospects.  First, for coronal giant stars that have ionised winds, application of the radiometic Bode's law indicates that long wavelength emission from substellar companions to giant stars may be detectable or nearly detectable with existing facilities.  Second, for the non-coronal giant stars that have neutral winds, the resultant bow shock may act as a ``feeder'' of electrons that is well-embedded in the companion's magnetosphere.  Incoherent cyclotron emissions are far too faint to be detectable, even with next generation facilities; however, much brighter flux densities may be achievable when CMI is considered. ",
        "introduction": "The discovery of extrasolar planets around solar-type stars has now become a fairly regular occurrence (e.g., see the Extrasolar Planets Encyclopaedia maintained by J.~Schneider at www.obspm.fr/encycl/encycl.html).  In addition to these, there have now been numerous detections of planetary companions to cool {\\em giant} stars, (early detections include Frink \\etal\\ 2002, Setiawan \\etal\\ 2003, and Sato \\etal\\ 2003).  With the ansatz that the formation of planetary systems around solar type stars is a relatively common occurrence, plus the recognition that such stars will evolve to become giants that experience significant mass loss on the way to becoming white dwarf stars, it becomes interesting to consider how giant star winds might affect substellar companions.  This paper is not the first to entertain questions about the eventual evolution of planetary systems during late stellar phases.  Several authors have considered the angular momentum transfer between the orbit of a planet or Brown Dwarf with a red giant and/or its wind (Livio \\& Soker 1984; Livio, Soker, \\& Harpaz 1984; Livio \\& Soker 2002).  Such effects could lead to ``sculpting'' of Planetary Nebulae (e.g., Soker 2001), which could explain the non-spherical but axisymmetric shapes that are so commonly observed (e.g., Balick 1987).  Rasio \\etal\\ (1996) have discussed the tidal decay of planetary orbits during the red giant phase.  Indeed, even as far back as 1924, Jeans examined the evolution of binary orbits due to stellar mass-loss (only in that paper, the mass loss being considered was in the form of the conversion of matter to energy via nuclear fusion).  Moreover, several authors have considered the possibility of detecting planetary companions to white dwarf stars, with the goal of constraining the evolution of planetary systems through empirical means (Zuckerman \\& Becklin 1987; Livio, Pringle, \\& Saffer 1992; Li, Ferrario, \\& Wickramasinghe 1998; Chu \\etal\\ 2001; Ignace 2001; Burleigh, Clarke, \\& Hodgkin 2002).  Recently, a substellar companion to a subdwarf in a post-red giant phase of evolution was reported (Geier \\etal\\ 2009).  This is particularly interesting as an example of a low mass companion that avoided being engulfed by the bloated red giant star to survive in a small period orbit around the remnant.  Such occurrences bolster the need for the consideration of detecting substellar companions to red giants to understand better the connections between stellar evolution and planetary evolution.  In this contribution one single fact is emphasized:  when stars like the Sun evolve to become red giants, the mass loss increases substantially by several orders of magnitude which may have observational consequences for detecting substellar companions and detailing their physical and orbital properties.  Already the influence of a strong wind on a substellar companion has been considered by Struck, Cohanim, \\& Willson (2003), who discuss the possibility of wind accretion by brown dwarf companions during the asymptotic giant branch (AGB) phase.  In terms of detecting extrasolar planets, a number of researchers have explored radio emissions from substellar companions to stars by extrapolating the radio properties of solar system planets to other star systems (e.g., see Zarka 2007 and references therein). The radio emission of the Earth and other solar system planets with significant magnetic fields are dominated by the cyclotron maser instability (CMI) process in which the low frequency spectrum below about 100~MHz is dominated by coherent cyclotron emission from mildly relativistic electrons (Gurnett 1974).  A key point is that electrons are fed to the magnetic poles where the CMI is strong, and that the coherent cyclotron emission dominates the incoherent emission by several orders of magnitude.  A greater understanding of the detailed energetics and emissive mechanism has come about fairly recently in relation to the terrestrial auroral kilometric radiation (AKR) by Mutel \\etal\\ (2007) and Mutel, Christopher, \\& Pickett (2008).  In applications to extrasolar planets, Jupiter has been prominent in these considerations, since it shows bursting behavior in the radio band that can be quite bright, partly in relation to interactions occuring between Jupiter and Io (e.g., Zarka 1998, 2004).  Griessmeier, Zarka, \\& Spreeuw (2007) have summarized various mechanisms that may contribute to radio emissions in the case of exoplants around main sequence stars. For example, coronal mass ejections and in some cases magnetic field interactions between the star and a short-period planet can produce strong radio signals.  Our contribution is to extend the use of the radiometric ``Bode's Law'' for late type main sequence stellar winds to those of giant stars for which several new considerations must be taken into account. This law is an empirically determined, and theoretically motivated, power-law relation between planetary radio emissions and the solar wind kinetic energy flux.  The radio emissions occur at low frequencies of about 0.3--40~MHz, of which the upper end is observable with LOFAR (Farrell \\etal\\ 2004). A lunar radio observatory such as the proposed Dark Ages Lunar Interferometer (DALI) would be sensitive to almost that entire range (see Lazio \\etal\\ 2007 for a summary of the capabilities of lunar based radio observatories).  Giant star winds differ from those of low-mass main sequence stars (i.e., solar-type winds) in important ways.  Giants that are warmer than early K tend to be X-ray emitters, whereas those that are cooler are X-ray faint.  The distinction in the X-ray properties is known as the Linsky-Haisch ``dividing line'' (Linsky \\& Haisch 1979).  Although these X-ray faint giants possess chromospheres, their winds are largely neutral (e.g., Reid \\& Menten 1996), whereas the earlier giants and the winds of low-mass main sequence stars are ionised plasma flows.  Consequently, in the case of the neutral giant star winds, the flow can {\\em penetrate} the magnetosphere of a substellar companion, and the supersonic wind flow will set up a bow shock in the vicinity of that companion. Cassinelli \\etal\\ (2008) have considered a closely related scenario. Interested in X-ray emissions, those authors used hydrodynamic simulations to model the bow shock for an early-type stellar wind interacting with a dense spherical gaseous clump.  With a wind speed of order $10^3$ km s$^{-1}$, and assuming adiabatic cooling, they were able to derive a power-law dependence for the differential emission measure with temperature of the form $dEM/dT \\propto T^{-7/3}$, with peak hot gas temperatures of order $10^7$~K.  That simulation remains relevant for the case of a substellar companion to a cool giant star.  Although cool star winds are far slower than hot star winds, with values of 15 km~s$^{-1}$ for AGB stars up to perhaps 100 km~s$^{-1}$ for some red giants, the winds are themselves much cooler, with sound speeds of a few km~s$^{-1}$.  As a result, the wind is still quite supersonic.  However, this does not mean that bow shocks to substellar companions will produce X-ray emissions for the cooler giant star systems. These winds are indeed much slower than those of early-type stars, and so the peak post-shock temperature will be considerably smaller.  For a strong shock, as considered by Cassinelli \\etal, the post-shock temperature at the bow shock head (or ``apex'') $T_A$ follows the well-known Rankine-Hugoniot relation, with  \\begin{equation} T_A = \\frac{3}{16}\\,\\frac{\\mu m_H}{k}\\,v_{\\rm rel}^2 = 14\\,{\\rm MK}\\, \\left(\\frac{v_{\\rm rel}}{1000~{\\rm km\\; s^{-1}}} \\right)^2. \\label{eq:Ta} \\end{equation}  \\noindent where $\\mu$ is the mean molecular weight for the gas, $m_H$ is the mass of hydrogen, and $v_{\\rm rel}$ is the relative speed between the wind and the companion.  The orbital motion can be non-trivial compared to the wind, but considering only the wind speed for the moment, one expects apex temperatures on the order of $T_A \\approx 3,000 - 140,000$ K for speeds of 15--100 km~s$^{-1}$. At the low end, orbital motion will likely dominate the incident shock speed in the rest frame of the companion (e.g., the Earth orbits at about 30 km~s$^{-1}$).  As a result, post-shock temperatures of $\\sim 10^4-10^5$~K are expected, sufficient to ionise hydrogen, thus leading to a substantial reservoir of electrons that can interact with a magnetosphere to produce radio emissions.  The scenario for the giant stars is thus significantly different than for main sequence stars.  First, known low-mass companions to giants have no examples of extremely short period orbits (a week or less) as in the main sequence case.  Second, giants with the most massive winds will be neutral, and thus have no counterpart in the solar system or to current applications among late-type main sequence stars.  Third, for the X-ray emitters (or ``coronal'' giants), there has been no evidence for coronal mass ejections as in the solar case, so that the application of the radiometric Bode's law is likely the most important emissive mechanism in the long wavelength radio band.  It is worth noting that Griessmeier \\etal\\ (2007) have evaluated the expected radio emissions for all known exoplanets at that time.  Their list includes companions to giant stars that had been reported at that time, and these do not produce detectable radio emissions.  However, those authors appear to have applied a main sequence solar wind model in relation mass-loss rate and wind speed which is inapplicable to giant star winds. Here we use empirical relations that are more appropriate for evolved cool star winds.  To explore the radio emissions from companions to giant stars, the next section summarizes the properties of the host stars and the orbital properties of their companions.  In Section~3, we apply the radiometric Bode's law to giant star systems, comparing emissions between stars on either side of the Linsky-Haisch dividing line. In Section~4, we consider the bow shock scenario for the case of neutral winds that penetrate deep into a magnetosphere.  We use the results of Cassinelli \\etal\\ (2008) to estimate the incoherent cyclotron emission. Results of our study are summarized in Section~5, where we conclude that new simulations are required to consider properly the operation of the CMI in the bow shock scenario.  An Appendix details our derivation for the cyclotron emission from a bow shock.  \\begin{table*} \\centering \\begin{minipage}{140mm} \\caption{Nominal Stellar and Planetary Data on Evolved Stars with Planetary Companions.} \\begin{tabular}{@{}lcccccccccr@{}} \\hline & \\multicolumn{4}{c}{\\it Host Star Data}  & &\\multicolumn{4}{c}{\\it Planet Data} \\\\ Name & Spectral & Distance & Mass & Radius & Luminosity & $\\dot M$ & $M\\sin i$& $a$ & $e$ & Ref.$^a$ \\\\ & Type  & (pc)  & $(M_\\odot)$  & $(R_{\\odot})$  & $(L_{\\odot})$ & $(10^{-9} \\,M_\\odot$ yr$^{-1})$ & $(M_J)$  & (AU)  &       \\\\ &  &  &  &  &  &       &     &       \\\\ \\hline HD185269       &   G0 IV       &  47     &  1.28   &  1.88  &   4.0$^b$  &  0.002 &  0.94       &   0.077   &  0.30     &  J06     \\\\ 81 Cet        &   G5 III:     &   97    &   2.4   &   11   &   60    &  0.11 &   5.3    &   2.5    &   0.206   &   S08b \\\\ HD 11977       &   G5 III      &  67     &  1.91   &  10    &   55$^b$  &  0.12 &  6.54       &   1.93    &  0.4      &  Se05    \\\\ 18 Del        &   G6 III      &   73    &   2.3   &   8.5  &   40    &  0.059 &   10.3   &   2.6    &   0.08    &   S08a \\\\ 11 Com        &   G8 III      &   112   &   2.7   &   19   &   170   &  0.48 &   19.4   &   1.29   &   0.231   &   L08 \\\\ HD175541       &   G8 IV       &  128   &  1.65    &   3.80  &   8.6     &  0.008 &    0.61    &    1.03    &   0.33     &  J07     \\\\ HD192699       &   G8 IV       &  67     &   1.68  &   3.9   &   10.2     &  0.009 &    2.5    &    1.16    &   0.149    &  J07     \\\\ HD 104985      &   G9 III      &  102    &  1.6    &  11    &   59    &  0.16 &   63     &  0.78    &    0.03   &   Sa03    \\\\ $\\xi$ Aql       &   K0 III      &   63    &   2.2   &   12   &   69    &  0.15 &   2.8    &   0.68   &   0$^d$      &   S08a \\\\ $\\epsilon$ Tau  &  K0 III      &  47.5   &   2.7   &   14   &   97    &  0.20 &   7.6    &   1.93   &   0.151   &   S07 \\\\ HD102272       &   K0 III      &  360    &   1.9   &  10.1  &   53$^b$  &  0.11 &  5.9        &   0.61       &  0.05   &  N09b    \\\\ 14 And        &   K0 III      &   76    &   2.2   &   11   &   58    &  0.12 &   4.8    &   0.83   &   0$^d$      &   S08b \\\\ HD17092        &   K0 III      &  110    &   2.3   &  10.1  &   43$^b$     &  0.076 &  4.6        &   1.3       &  0.17    &  N07     \\\\ $\\beta$ Gem    &   K0 III      &  10.3   &   1.86  &   9    &   34$^b$   &  0.066 &    2.9    &   1.69   &   0.06      &  H06, R06     \\\\ HD81688        &   K0 III-IV   &   88    &   2.1   &   13   &   72    &  0.18 &   2.7    &   0.81   &   0$^d$      &   S08a \\\\ $\\kappa$ Cr B  &   K0 IV       &  31.1   &   1.8   &   4.71  &   12.3      &  0.013 &   1.8    &    2.7     &   0.146    &  J08    \\\\ 6 Lyn         &   K0 IV       &   57    &   1.7   &  5.2   &   15    &  0.018 &   2.4    &   2.2    &   0.134   &   S08b \\\\ HD32518        &   K1 III      &   120   &   1.13  &  10.2  &   41$^b$  &  0.15 &  3.04  &  0.59     &   0.01      &   D09b\\\\ 4 U Ma         &   K1 III      &  62     &  1.23   &  18    &   110$^b$   &  0.64 &  7.1        &   0.87    &  0.43     &  D07     \\\\ HD167042       &   K1 III      &  50.    &   1.64  &   4.30  &   10.5     &  0.011 &    1.7    &    1.3     &   0        &  J08     \\\\ HD 47536       &   K1 III      &  121    &   1-1.5  &  23   &   4380   &  40 &   5-6   &   2     &    0.20      &   Se03       \\\\ HD167042       &   K1 IV       &   50    &   1.5   &   4.5  &   10    &  0.012 &   1.6    &   1.3    &   0.101   &   S08b \\\\ $\\gamma$ Ceph  &   K1 IV       &  13.8   &   1.6   &   4.66  &   11$^b$  &  0.013 &   1.7     &   2.13   &   0.12     &  H03     \\\\ HD210202       &   K1 IV       &  56    &  1.85    &   4.45  &   11.3     &  0.011 &    2.0    &    1.17    &   0.152    &  J07     \\\\ 42 Dra         &   K1.5 III    &   97    &   0.98  &   22   &   130$^b$  &  1.2 &  3.88  &  1.19    &  0.38       &  D09a \\\\ BD +20 2457    &   K2 II       &   200+$^c$  &  2.8  &  49  &   610$^b$     &  4.3 &  21.42  &  1.45  &  0.15   &  N09a \\\\ BD +20 2457    &   K2 II       &   200+$^c$  &  2.8  &  49  &   610$^b$     &  4.3 &  12.47  &  2.01  &  0.18   &  N09a \\\\ HD 13189       &   K2 II?      &  --- &  2-7    &   --- &   4000   &  --- &   8-20   &  1.5-2.2   &   0.27    &   H05     \\\\ $\\iota$ Dra    &   K2 III      &  31     &  1.05   &  12.9  &   70    &  0.34 &   8.9    &   1.3    &    0.70   &   F02    \\\\ HD24210        &   K3 III      &   140   &   1.25  &  56    &   950$^b$  &  17 &  6.90  &  1.33  &  0.15    &  N09 \\\\ HD139357       &   K4 III      &   118   &   1.35  &  11.5  &   58$^b$   &  0.20 &  9.76  &  2.36    &  0.10       &  D09a \\\\ 11 U Mi       &   K4 III      &   120   &   1.8   &   24   &   180$^b$  &  0.96 &  11.2  &  1.54     &   0.08    &   D09b\\\\ \\hline \\end{tabular} $^a$References:  D09 = Dollinger et al. (2009b); D09a = Dollinger et al. (2009a); N09a = Niedzielski et al. (2009a); L08 = Liu et al. (2008); S08a = Sato et al. (2008a); S08b = Sato et al (2008b); S07 = Sato et al (2007) N09b = Niedzielski et al. (2009b); N07 = Niedzielski et al. (2007); D07 = Dollinger et al. (2007); J06 = Johnson, et al. (2006); Se05 = Setiawan et al. (2005); H05 = Hatzes et al. (2005); Sa03 = Sato et al. (2003); F02 = Frink et al. (2002); Se03 = Setiawan et al. (2003); H06 = Hatzes et al. (2006); R06 = Reffert et al. (2006); H03 = Hatzes et al. (2003); J07 = Johnson et al. (2007); J08 = Johnson et al. (2008). \\\\ $^b$ Luminosity computed from values of $T_{\\rm eff}$ and $\\log g$ given by respective reference. \\\\ $^c$ Luminosity estimate based on spectral type. \\\\ $^d$ Eccentricity fixed to zero for orbital solution. \\\\  \\end{minipage} \\end{table*} ",
        "conclusions": "}  On the whole our analysis for the detection of radio emissions from the interaction between the wind of a red giant star and its substellar companion is largely negative.  Cyclotron emissions appear to be at the micro-Jansky level, orders of magnitude below current or near-future detection thresholds. However, we have taken the most conservative and pessimistic approach in our estimates.  We have considered red giant winds that are largely neutral (i.e., those on the cool side of the ``dividing line''), yet the winds of earlier type red giants may be highly ionised with large gain factors in the resultant radio emissions.  There have been radio-wavelength surveys (typically at 2 and 6 cm) of red giant stars including stars on the asymptotic giant branch of the H-R diagram (see Drake \\& Linsky 1983, 1986; Drake, Linsky, \\& Elitzur 1987; Drake \\etal\\ 1991; Luttermoser \\& Brown 1992).  Typically evolved stars on the cool side of the ``dividing line'' are either not detected at these wavelengths or are very weak sources --- giant stars of M6 or warmer have been detected (Drake \\etal\\ 1991).  The weak radio emission of these non-Mira giants are thought to be due to partially ionised winds. There have been some detections of giants on the cool side of M6~III. For example $o$~Ceti (Mira) appears variable in its radio emission. Mira was detected at 6~cm by Spergel, Giuliana, \\& Knapp (1983) as a $3\\sigma$ event; however, Drake \\etal\\ (1987) failed to detect this pulsating long-period variable (LPV) at either 2~cm or 6~cm.  Since Mira has a hot companion, it is not clear if the emission arose from the LPV or from the companion. In a survey of seven N-type carbon stars at 3.6~cm, Luttermoser \\& Brown (1992) only detected V~Hya, which is a peculiar carbon star that shows evidence of bipolar outflow (Tsuji \\etal\\ 1988).  In the case of a neutral wind, we have argued that a bow shock will result in the vicinity of the companion object.  If the relative flow speed is enough, one can expect post-shock ionisation of the wind material.  For incoherent cyclotron emission, the radio flux of the ionised component will be quite weak.  However, it may be possible that some fraction of these newly created electrons are fed to the polar regions of the companion where the CMI mechanism can operate leading to much stronger radio emission.  In addition, analogous to the Jupiter-Io interaction, moons of substellar companion may also produce bow shocks, and these might provide a source of electrons that could be accelerated in the magnetosphere and feed the CMI. Our model for the bow shock is not adequate to address these possibilities.  New self-consistent MHD calculations will be needed to model the interaction of the bow shock with the companion's magnetosphere and to assess the channeling of electrons for use with the CMI.  Of course, all of the effects described so far would be affected by orbital eccentricity of the companion.  For orbits of a fixed semi-major axis, the radio flux density will be significantly higher near periapse as compared to apapse for increasingly eccentric orbits.  But then of course, gains in signal strength near periapse will occur over a diminishing fraction of the orbital period as eccentricity increases.  Although these effects should be explored, they are of secondary importance to substantial task of modeling the flow dynamics more accurately.  Finally, it is possible that the AGB stars might be targeted for long wavelength radio studies.  The AGB winds are certainly neutral so that the bow shock scenario would be of interest.  The mass-loss rates can be extremely large at around $10^{-5}\\,M_\\odot$ yr$^{-1}$.  Since the radio emission scales with $\\dot{M}^2$, there would be a gain factor of 6 orders of magnitude above our estimates for the radio emission, and that is just for incoherent cyclotron emission.  Of course, AGB stars are relatively rare and so tend to be more distant.  Plus the companion will probably need to be in an orbit that is greater than 1~AU, otherwise it could be engulfed by the star.  With a wind speed of just 15 km s$^{-1}$, the astute reader may recognize that the apex temperature will hardly be large enough to ionise hydrogen.  However, there are actually two effects to mitigate the low wind speed.  First is that the orbital speed could be larger by perhaps a factor of 2 or more. Second, the hydrodynamical simulations of Cassinelli \\etal\\ (2008) are for clumps treated as blunt obstacles without self-gravity.  Jovian and brown dwarf mass objects have surface escape speeds of 30 km s$^{-1}$ and higher.  Given that the bow shock apex is at around 1.5 times the radius of the blunt object in the Cassinelli \\etal\\ simulations, a slow AGB wind can double and triple in speed as it falls into the gravitational potential of the substellar companion (i.e., similar to accretion described by Bondi \\& Hoyle 1944).  Of course, inclusion of gravity changes the details of the bow shock, but the main point is that significant ionisation would still be viable.  Collecting the various factors, $S_\\nu$ of around 0.1~mJy might be achievable.  It is clear that long wavelength observations will provide new opportunities for detecting substellar companions and for studying their magnetospheres.  Companions to red giants should not be neglected in these efforts.  An interesting result of our study has been the relatively novel consideration of neutral winds interactions with substellar companions, completely unlike the solar system case.  The prospects for future observations are certainly exciting, but there is a clear need for more detailed modeling of MHD bow shocks in relation to the CMI to explore the viability of radio detection."
    },
    "0911/0911.1481_arXiv.txt": {
        "abstract": "We present analysis of both the short-term optical and long-term multiwavelength variability of CTA 102. In 2004, this object was observed in an intense optical flaring state. Extensive R-band microvariability observations were carried out during this high state. In 2005, we obtained several weeks of contemporaneous radio, optical, and X-ray observations of CTA 102. These observations recorded distinct flaring activity in all three wavebands. Subsequent analysis revealed that this object may appear redder when in a brighter optical state, and that the X-ray, optical and radio activity do not appear to be correlated. The shape of the observed spectral energy distributions suggests that both synchrotron-related and external inverse Compton processes may contribute to the X-ray emission. Our results are also compared to other results on this object and archival microvariability observations. It appears that more rapid, dramatic microvariability events occur when CTA 102 is in an elevated optical flux state. ",
        "introduction": "Blazars are a radio-loud subclass of active galactic nuclei (AGNs), typically featuring extreme core-dominated radio morphologies. Their optical continua are markedly steep, they exhibit a high degree of optical polarization (up to about 20\\%), and their fluxes vary on timescales of an hour to several years at all observed wavelengths. The most extreme and unique property of blazars is their highly beamed continuum, most likely produced by a jet of relativistic material oriented close to the observer's line of sight \\citep[][ and references therein]{urr95}.  Blazars can be subdivided into two categories: BL Lac objects and Flat Spectrum Radio Quasars (FSRQs). BL Lac objects have particularly weak spectral lines, so their redshifts, and hence luminosities, are difficult to determine. FSRQs are distinguished by the presence of broad optical emission lines. As the continuum flux varies, the distinction between the two classes may become blurred. Spectral lines indicating an FSRQ may be present during a low-continuum state, but such lines may become weak or disappear during a high-continuum state, thus resembling a BL Lac object \\citep[][ and references therein]{ulr97}.  Blazars exhibit the most extreme variability observed for any class of AGN. Optical microvariability has been detected on timescales as short as minutes to hours, with similar variations observed at all wavelengths except radio \\citep{ mil96}. The earliest microvariability results were from BL Lac, one of the most frequently studied blazars \\citep{rac70,mil89}. Observations of close correlations between flares observed in different wavebands strongly indicated that these variations are intrinsic to the blazars themselves \\citep{wag95}. Possible extrinsic explanations, such as interstellar scintillation and gravitational microlensing, do not match the observations.  FSRQs and low-frequency peaked BL Lac objects (LBLs) typically exhibit strong optical microvariability on timescales as short as several minutes. Among recent observations of BL Lac, \\citet{bot04} observed $\\Delta$R of $\\sim$ 0.35 magnitudes (or, in terms of flux F, $\\Delta$F/F of $\\sim$ 0.4) over 1.5 hours, \\citet{nik05} observed a change in R of 0.10 $\\pm$ 0.01 magnitudes within 1 hour, and \\citet{ben06} observed $\\Delta$R of 0.147 $\\pm$ 0.020 magnitudes over 2.5 hours. Observations of 3C 279 from 1989 to 2002 found several instances of R-band microvariability covering many different overall flux states \\citep{bal02a,bal02b}. \\citet{how04} observed R-band microvariability in 3C 345. They also note that there was no apparent correlation between the occurrence of microvariability and the overall flux state of the object.  The spectral energy distribution (SED) of blazars is distinguished by two peaks: one in the radio/UV regime and the other in the X-ray/$\\gamma$-ray regime. The spectrum in the radio to UV range is generally agreed to arise from synchrotron emission from relativistic electrons spiraling around the jet's magnetic field lines \\citep{ulr97}. The X-ray/$\\gamma$-ray spectrum is most likely due to inverse Compton (IC) emission in FSRQs and LBLs, while most of the X-rays are the high-energy tail of the synchrotron emission in high-frequency peaked BL Lac objects (HBLs) \\citep{ulr97}. A major unsolved puzzle of blazars is what supplies the ``seed photons\" which are upscattered to produce the IC emission. In order to fully understand blazars, we must understand the physics of the region near the central engine, since this is where jet particles are collimated and accelerated to relativistic speeds.  CTA 102, also known as PKS J2232+1143, has a redshift of 1.037 \\citep{sch65}, and was first identified as a radio source by the Caltech Radio Survey \\citep{har60}. \\citet{san65} confirmed that it was a quasar, and observations of a blue stellar optical counterpart were published the same year \\citep{wyn65}. CTA 102's optical polarization, observed at nearly 11\\%, put it in the class of High Polarization Quasars (HPQs), objects which exhibit optical polarization of at least 3\\% \\citep{moo81}. Its spectral properties identify it as an optically violent variable (OVV) quasar \\citep{mar86}. Variability in CTA 102 was observed in the radio \\citep{hun72} and optical bands \\citep[e.g.,][]{pic88} as early as the late 1960s \\citep{bar78}. \\citet{tor99} observed a simultaneous radio and optical flare in CTA 102 in 1997.  In this work, we present the results of an intensive optical microvariability campaign performed on CTA 102 during a very high flux state. The results of this campaign are compared with microvariability observed in this object at other times and different flux levels. We also present the results of the first simultaneous radio, optical, and X-ray campaign performed for this object. We discuss the correlations, or absence thereof, found between the various wavebands and the variation in optical color present during an optical flare. We present and discuss the broadband SEDs and compare them with previous results on CTA 102 and similar objects. ",
        "conclusions": "Our analysis of the optical microvariability of CTA 102 suggests that the character of the microvariability changes in response to the overall optical state. In higher flux states, microvariability events feature more rapid changes in flux compared to events observed at lower flux states (see Figure 6). Our results appear to contradict those of \\citet{sta04}. The results of \\citet{sta04} suggest that the mechanisms behind optical microvariability and long-term variability are different, and thus should not appear to exhibit correlated behavior. Investigating correlations between flux state and the character of microvariability is an important step toward understanding the underlying physics behind blazar microvariability. \\citet{how04} suggest that microvariability can be produced by a variety of mechanisms, including instabilities in the particle acceleration mechanism, variations in the supply of electrons, or small inhomogeneities in the object's magnetic fields or the ambient medium. Others suggest that accretion disk instabilities and fluctuations can cause microvariations \\citep{man93,sta04,gup05}. However, \\citet{how04} argue that accretion disk phenomena are not responsible for microvariations because such phenomena would be washed out by the highly beamed jet emission. Additional campaigns are needed on CTA 102 and other blazars to more comprehensively examine possible relationships between the optical flux state and optical microvariability activity in these objects.  Significant optical microvariability has been observed in this object, but not in conjunction with remarkable radio or X-ray activity. In 2004, the most dramatic microvariability observed occurred simultaneously with no flaring activity in the radio regime. In 2005, significant X-ray, optical, and radio flaring was observed, but with no significant optical microvariability events. This could indicate that the processes causing optical microvariability are not related to larger scale jet events affecting the synchrotron emission in this object, which would likely cause correlated optical and radio events. The cause of the lack of microvariability with X-ray activity needs further investigation in future campaigns. One possible explanation for these missing correlations would be that the optical microvariability is at times not related to jet processes, thus arguing in favor of accretion disk activity causing optical microvariability. However, the 2005 campaign was the only campaign where X-ray, optical, and radio data were all available. More broadband campaigns with more even sampling across regimes are needed to properly investigate the microvariability-jet and microvariability-accretion disk connections.  Investigations of the optical emission show that CTA 102 appears redder when in a brighter flux state. \\citet{gu06} suggest that this behavior could be a result of the ``blue bump\" observed in many FSRQs, thought to be thermal accretion disk radiation. Depending on how large this thermal contribution is, it could be significant in lower optical flux states and overwhelmed by the jet emission in higher states, thus making the FSRQ appear redder when brighter. Recent campaigns on similar objects have found significant blue bump contributions to the observed spectra, e.g., \\citet{kat08,rai08}. Future campaigns on FSRQs should include UV observations to better show how strong the blue bump component is compared to the continuum.  Broadband spectral analysis of CTA 102 suggests that the X-ray emission could be due to either the synchrotron or the IC process. Multiyear \\emph{Swift}/BAT transient monitor results provided by the \\emph{Swift}/BAT team\\footnote{http://swift.gsfc.nasa.gov/docs/swift/results/transients/index.html, CTA 102 is listed as 4C 11.69.} indicate that CTA 102 was in a relatively low X-ray flux state at the time of our 2005 campaign. Additional campaigns at varying levels of X-ray flux are needed to better determine the source of the X-ray emission in this object. Such campaigns would greatly benefit from longer temporal coverage in order to find clearer correlations between the synchrotron and X-ray behavior. It is also important to include UV and $\\gamma$-ray observations in future campaigns in order to better investigate the high-energy synchrotron tail and IC peak of CTA 102 and similar objects."
    },
    "0911/0911.4571_arXiv.txt": {
        "abstract": "We present a radio survey carried out with the Australia Telescope Compact Array.  A motivation for the survey was to make a complete inventory of the diffuse emission components as a step towards a study of the cosmic evolution in radio source structure and the contribution from radio-mode feedback on galaxy evolution.  The Australia Telescope low-brightness survey (ATLBS) at 1388~MHz covers 8.42~deg$^{2}$ of the sky in an observing mode designed to yield images with exceptional surface brightness sensitivity and low confusion. The survey was carried out in two adjacent regions on the sky centred at RA: $00^{\\rm h}~35^{\\rm  m}~00^{\\rm s}$, DEC: $-67\\degr~00\\arcmin~00\\arcsec$ and RA: $00^{\\rm h}~59^{\\rm m}~17^{\\rm s}$, DEC: $-67\\degr~00\\arcmin~00\\arcsec$ (J2000.0). The ATLBS radio images, made with 0.08~mJy~beam$^{-1}$ rms noise and $50\\arcsec$ beam, detect a total of 1094 sources with peak flux exceeding 0.4~mJy~beam$^{-1}$. The ATLBS source counts were corrected for blending, noise bias, resolution, and primary beam attenuation; the normalized differential source counts are consistent with no upturn down to 0.6~mJy. The percentage integrated polarization $\\Pi_{0}$ was computed after corrections for the polarization bias in integrated polarized intensity; $\\Pi_{0}$ shows an increasing trend with decreasing flux density. Simultaneous visibility measurements made with longer baselines yielded images, with $5\\arcsec$ beam, of compact components in sources detected in the survey. The observations provide a measurement of the complexity and diffuse emission associated with mJy and sub-mJy radio sources. 10\\% of the ATLBS sources have more than half of their flux density in extended emission and the fractional flux in diffuse components does not appear to vary with flux density, although the percentage of sources that have complex structure increases with flux density. The observations are consistent with a transition in the nature of extended radio sources from FR-{\\sc ii} radio source morphology, which dominates the mJy population, to FR-{\\sc i} structure at sub-mJy flux density. ",
        "introduction": "Our understanding of galaxy evolution across cosmic time depends on multi-wavelength data, which are products of ultra-deep surveys across the electromagnetic spectrum that have been made using the most sensitive observatories in operation. Instrument design constraints and resource limitations usually lead to survey strategies that range from all-sky surveys with low angular resolution and sensitivity to ultra-deep surveys of small sky regions that are made with high angular resolution; different survey strategies address different components of source populations and different aspects of galaxy evolution.  The radio component of these multi-wavelength campaigns that target small sky regions has most often been done using interferometer telescopes configured to give images with sub-arcsec resolution that are comparable to or better than the corresponding optical surveys. These radio surveys usually have extremely good flux sensitivity and are capable of detecting $\\mu$Jy emission from distant galaxies. However, Fourier synthesis imaging that is done with widely spaced interferometer elements---in order to image with high angular resolution---tend to lack surface brightness sensitivity, which is the ability to detect faint extended emission components. This is partly due to missing short spacings, which implies missing information on extended emission, and partly because of incomplete visibility coverage, which results in increased confusion. Furthermore, imaging with an interferometer array that has complete visibility coverage will still have less (redundant) short spacings than a filled aperture and hence less sensitivity to extended structure. Consequently, the ultra-deep radio images might fail to reproduce extended emission components associated with galaxies, although extremely faint compact components are represented in the images.  The normalized differential radio source counts are observed to show an upturn below flux density of about 1~mJy \\citep{wi85} indicating a rapid increase in the number of faint sources at these flux density levels, which might constitute a new population.  The bulk of these faint radio sources in the 0.1--1~mJy range in flux density are identified with early type galaxies \\citep{ma08}, with radio structure believed to be of the FR-{\\sc i} \\citep{fr74} type \\citep{pa07}, which often have associated extended emission components. Therefore a complete census of the radio emission associated with faint galaxies, at these flux density levels and at intermediate redshifts of $z=1$--2, requires imaging with good surface brightness sensitivity. Additionally, the radio morphology of the relatively lower surface brightness extended emission is a clue to the nature of the radio source and is a means of distinguishing between active galactic nuclei (AGNs) of different classes.     The Square Kilometre Array (SKA\\footnote{http://www.skatelescope.org}) has the potential to detect AGNs, including the radio quiet population, out to redshift $z \\approx 6$. As compared to X-ray and optical surveys, radio imaging with the next generation telescopes are likely to become the instrument of choice for identifying high redshift AGNs, in particular the obscured AGNs \\citep{ja04}. In this context, it is important to quantify the expectations for radio source confusion from both compact and extended radio source components at faint flux density levels as this is a potential limitation to imaging with sub-$\\mu$Jy sensitivity. A characterization of the radio sky at faint flux density---in total intensity and polarization---is a useful input to simulations of the radio sky and optimization of SKA array configurations as well as observing strategies.  The survey presented here has been made with the specific goal of providing a view of the low surface brightness radio emission associated with mJy and sub-mJy radio sources, to provide a database for their characterization, assessment of the cosmic evolution of extended radio components and their influence on galaxy evolution. The surface brightness sensitivity of this survey, which we refer to as the Australia Telescope Low-Brightness Survey (ATLBS), is about a factor of five better than any previous survey with comparable resolution \\citep{sub07}. In this first paper we present the ATLBS survey together with the source counts and a population study of the radio structures. Forthcoming papers will present detailed radio structures, optical identifications, polarization analysis, and explore in depth the open problems like, for example, cosmic evolution of low power radio galaxies, evolution of the radio source structure with cosmic epoch and the role of kinetic-mode feedback from AGNs on galaxy environment and galaxy evolution, where progress depends on our understanding of the low surface brightness radio sky. ",
        "conclusions": "We have used the Australia Telescope Compact Array to survey 8.42~deg$^{2}$ sky area at a radio frequency of 1388~MHz.  The interferometer observations were made in a mode designed to mosaic image the wide field with complete visibility coverage, and hence low confusion, and with exceptional surface brightness sensitivity.  The data were used to reconstruct (a) a low resolution image, with beam FWHM $53\\arcsec \\times 47\\arcsec$ and rms noise 0.08~mJy~beam$^{-1}$, and (b) a high resolution image with beam FWHM of $4\\farcs6$. Whereas the low resolution image reproduces the extended and diffuse radio emission associated with sources in the fields, the high resoluion image resolves out structure on scales exceeding the beam size. Together, the images provide an estimate of the structural properties of the mJy and sub-mJy radio sources. We refer to our radio survey as the Australia Telescope low brightness survey, and use the acronym ATLBS.  A total of 1094 radio sources with peak flux density exceeding 0.4~mJy~beam$^{-1}$ in the low resolution image were cataloged and their source properties estimated. The source detections correspond to a density of 130 sources per square degree. The studies presented herein have considered only sources with peak flux density exceeding about 5 times the rms noise in the image.  The normalized differential source counts derived from the ATLBS shows no evidence for an upturn down to about 0.6~mJy; the ATLBS counts are consistent with the ATESP source counts, but relatively low compared to many other surveys, including the PDS. This result suggests that there is no substantial population of low surface brightness sources or source components at these flux densities that have been missed by previous surveys.  The ATLBS counts---as also ATESP counts---are relatively low compared to many other counts perhaps because our counts are based on a source catalog, rather than a component catalog. As far as we know, blending has been considered and corrected for the first time in the work presented herein; blending is of concern in surveys such as the ATLBS that aim to go deep in surface brightness sensitivity. The derived ATLBS source counts have been corrected for blending, noise bias, resolution and primary beam attenuation over the survey area.  The derived source counts are consistent with that of the ATESP survey suggesting that the relatively large sky coverage of these surveys is key to robust measurement of source counts. The considerable scatter in the derived differential source counts at about 1~mJy flux density between the various surveys that were made with relatively smaller sky coverage is perhaps owing to genuine field-to-field variations in counts.  To summarize, the work presented here emphasises the importance of constucting source catalogs, accounting for blending and noise bias, and making large area surveys in order to refine our understanding of differential source counts at sub-mJy and mJy flux densities.  The main results presented in this paper concern the statistical properties of the radio structure and polarization in sub-mJy sources compared to the mJy radio source population. We have defined a complexity parameter $\\chi$ as the ratio of the total flux density of the source to the peak flux density in the high resolution image, this parameter is a measure of the source complexity and the departure of the source structure from an unresolved single compact component. Additionally, we have defined a diffuseness parameter $\\delta$ as the ratio of total flux density to the sum of the flux density in compact components, this parameter is a measure of the flux density in diffuse emission components.  These measures of morphology have been computed for all ATLBS sources. The points arising from an examination of the sources above and below 1~mJy flux density are listed below.  \\begin{enumerate}  \\item In the low-resolution images made with $50\\arcsec$ beam, the fraction of extended sources rises from 15\\% for the sub-mJy population to 28\\% for 1--10~mJy sources and to 70\\% for 10--100~mJy sources. At this resolution, only 2\\% of sub-mJy sources are observed to have composite structure where as this fraction rises to 15\\% for 1--10~mJy sources and 50\\% for 10--100~mJy sources.  \\item Less than 1\\% of the sub-mJy ATLBS sources have been observed to have multiple compact components.  However, about 10\\% of 1--10~mJy sources have multiple compact components and this fraction rises to 36\\% in sources in the 10--100~mJy range.  \\item The median complexity $\\chi$ for ATLBS sources is 1.28 and the median diffuseness $\\delta$ is 1.09.  20\\% of ATLBS sources have $\\chi$ exceeding 2.0, implying that a fifth of the sources are doubles or triples or have more than half their flux density in extended emission. 10\\% of the ATLBS sources have $\\delta$ exceeding 2.0, implying that a tenth of the sources have more than half their flux density in diffuse emission.  \\item We observe no significant trend in median $\\delta$ with flux density. However, $\\chi$ rises significantly with increasing flux density.  Whereas only 8\\% sub-mJy ATLBS sources have $\\chi$ exceeding 2.0, 28\\% of 1--10~mJy sources have $\\chi$ exceeding 2.0 and this fraction rises to 55\\% for the 10--100~mJy sources.  \\end{enumerate}  The sub-mJy ATLBS sources, with 10\\% sources having more than half the flux density in extended emission, almost always have a single compact component, if present. On the other hand, sources with higher flux density tend to have a greater fraction of multiple compact components, although the fractional flux density in the diffuse emission might be the same.  This is consistent with population synthesis models for the differential source counts wherein the mJy radio source population is dominated by the relatively powerful radio sources, which are often of the hot-spot type with FR-{\\sc ii} structure, and the sub-mJy sources are dominated by the relatively lower power radio sources, which often manifest the FR-{\\sc i} structure with a single compact component.  We have computed the percentage integrated polarized intensity for the ATLBS sources and examined their variation with total flux density.  As far as we know, we have formulated herein a polarization bias correction for integrated polarized emission for the first time. We observe an increase in the percentage polarization in the ATLBS sources with decreasing flux density.  The median percentage polarization is above 10\\% in the sub-mJy sources and declines to a few percent in sources with 100~mJy flux density. We have been unable to find any correlation between the percentage polarization and the source properties like size, complexity or the fractional flux in the diffuse emission.  Since we do observe a decrease in source complexity in fainter sources, it may be that the increased percentage polarization in the fainter sources is owing to a transition from FR-{\\sc ii} dominated population at higher flux densities to FR-{\\sc i} dominated population. Radio source populations dominated by edge-darkened FR-{\\sc i} jets or relaxed doubles and relict sources with relatively homogeneous magnetic field orientation may suffer less depolarization, due to averaging over the spatial extent of the source, as compared to the FR-{\\sc ii} sources that have more complex spatial structure in their field distributions. Alternately, the increased fractional polarization in the fainter sources may be due to lower internal Faraday depolarization if the fainter sources are at relatively higher redshift and the emission frequency in the rest frame of the source is higher in the case of the fainter sources. Higher resolution radio imaging and redshift measurements of the optical identifications of the ATLBS sources---both of which are currently underway---might shed light on this issue.  The Very Large Array (VLA) survey of the Chandra Deep Field South \\citep{ke08}, together with optical identifications of the radio sources \\citep{ma08}, suggest that the the sub-mJy radio sources above 0.08~mJy are dominated by non-thermal emission associated with early type galaxies hosting AGNs.  \\citet{pa07} suggest that the dominant population is low-luminosity AGNs of the FR-{\\sc i} type.  The ATLBS study of radio source morphology are consistent with this view.  The ATLBS survey has indicated that the sub-mJy radio source population does indeed have a non-negligible fraction of their integrated flux density in diffuse emission. Therefore, it is vital that deep radio surveys be made with adequate surface brightness sensitivity to measure the total radio luminosity associated with the AGNs, and accurately quantify any associated mechanical feedback.  The ATLBS survey regions are being re-observed using extended array configurations of the ATCA to reconstruct source structures with higher angular resolution, to confirm the findings presented in this paper and explore further the radio structural properties of the sub-mJy radio source population.  Additionally, the wide fields are being mosaic observed in the optical and IR bands using the CTIO Blanco Telescope and the Anglo-Australian Telescope in a study of the host galaxies and their environments. These observations and the consequent refinements in our understanding of the evolution in extended radio sources will be the subject of forthcoming papers."
    },
    "0911/0911.4747_arXiv.txt": {
        "abstract": "We have constructed a {\\em detailed} spectral atlas covering the wavelength region 930\\,\\AA~ to 1225\\,\\AA~ for 10 sharp-lined B0-B9 stars near the main sequence.  Most of the spectra we assembled are from the archives of the {\\it FUSE} satellite, but for nine stars wavelength coverage above 1188\\,\\AA~ was taken from high-resolution {\\it IUE} or echelle {\\it HST/STIS} spectra. To represent the tenth star at  type B0.2\\,V we used the {\\it Copernicus} atlas of $\\tau$\\,Sco. We made extensive line identifications in the region 949\\,\\AA~ to 1225\\,\\AA~ of all atomic features having published oscillator strengths at types B0, B2, and B8. These are provided as a supplementary data product, -  hence the term {\\em detailed} atlas. Our list of found features totals 2288, 1612, and 2469 lines, respectively. We were able to identify 92\\%, 98\\%, and 98\\% of these features with known atomic transitions with varying degrees of certainty in these spectra. The remaining lines do not have published oscillator strengths. Photospheric lines account for 94\\%, 87\\%, and 91\\%, respectively, of all our identifications, with the remainder being due to interstellar (usually molecular H$_2$) lines. We also discuss the numbers of lines with respect to the distributions of various ions for these three most studied spectral subtypes. A table is also given of 167 least blended lines that can be used as possible diagnostics of physical conditions in B star atmospheres. ",
        "introduction": "\\label{int}   To understand the physical conditions of O and B stars and their immediate environments there can be no better mode of study than high-dispersion ultraviolet spectroscopy. The spectral energy distributions of hot stars peak in this regime, and it is also here that absorption lines arising from a wide range of excited atomic states appear in high concentration.  These lines provide a rich and rather unexplored area for the study of markers of atmospheric temperature, pressure, kinematics, and chemical composition. In addition, atoms undergoing far-UV transitions ``see\" down to deep layers of the atmosphere because of the local relative transparency. The resulting intense UV radiation field is responsible for creating departures from LTE in these transitions and influencing the calibration of spectral diagnostics. Once atmospheres of single hot stars are understood, this knowledge can be extended to the interpretation of spectra of distant ensembles of hot stars and star formation regions.  At about the turn of this century the era dawned in which specific far-UV spectral lines observed in OB clusters embedded in highly redshifted galaxies can be studied from the ground at optical wavelengths (e.g., de Mello 2000 et al., Pettini et al.  2001, Dawson et al. 2002, Pettini et al. 2002, Mehlert et al. 2002, Croft et al. 2002, Robert et al. 2003, Shapley et al. 2003, de Mello et al. 2004). In the two decades leading up to and including this milestone, enormous progress was made in understanding the atmospheres and origin of OB stars. First came high-resolution spectroscopic observations of the middle-UV during the extended period (1978-1996) of the {\\it International Ultraviolet Explorer (IUE)} satellite. Next, the 1999 launch of the {\\it Far Ultraviolet Spectroscopic Explorer (FUSE),} brought access to the far-UV for hundreds of hot stars in the Milky Way and Magellanic Clouds.  This era of grand exploration and surveys closed with the decommissioning of {\\it FUSE} in 2007. The final reprocessing of the {\\it FUSE} spectra (covering the wavelength region 920\\,\\AA~ to 1188\\,\\AA) now permits the examination of a wealth of homogeneous high-quality material.  For O and B stars one consolidation of these gains was the addition of a fine set of middle-UV and far-UV spectral atlases. Examples of the former are the {\\it Copernicus} atlas of $\\tau$\\,Sco (Rogerson et al. 1977) and {\\it IUE} pictorial atlases of O and B stars (Walborn, Nichols-Bohlin, \\& Panek 1985, Rountree \\& Sonneborn 1993, Walborn, Parker, \\& Nichols 1995). Pellerin et al. (2002) inaugurated a second group of digitized atlases with their atlases of O and early-B type stars of all luminosity classes. This work had the goal of exhibiting the general behavior of the strong photospheric lines and the so called UV wind lines with effective temperature T$_{eff}$ and luminosity class. Valuable supplemental information about the molecular H$_2$ lines formed in the interstellar medium (ISM) was also provided. From the point of view of stellar researchers, these features contaminate many of the far-UV photospheric lines of most Galactic stars. This atlas was closely followed by a second spectral atlas on OB stars in the two Magellanic Clouds by the same group (Walborn et al. 2002). This work concentrated on the behavior of wind lines with respect to metallicity as well as temperature and luminosity.  In addition to these atlases, Blair et al. (2009) have published a compendium of {\\it FUSE} spectra of hot stars in the Magellanic Clouds. This atlas focused on the identification of far-UV resonance lines arising from the ISM that are sometimes found at multiple velocities. Specialized atlases are now underway for the purposes of displaying the behavior of specific elements in particular types of stars. For example, an atlas depicting lines of heavy metal elements in {\\it IUE} spectra of late-B and early B stars has recently been published by Adelman et al. (2004).  The first far-UV high-dispersion spectral coverage published for a B star was the {\\it Copernicus} atlas of $\\tau$\\,Scorpii (B0.2\\,V; Walborn 1971) by Rogerson \\& Upson (1977; ``RU77\"). Rogerson \\& Ewell (1985; ``RE85\") published a detailed tabulation of photospheric and ISM lines identified in this atlas. The RE85 work was undertaken at a time when spectral synthesis tools were not commonly available, and when reliable identifications lay in the hands of a comparative handful of experienced spectroscopists who were also specialists in atomic physics. Even so, in the absence of commonly available spectral syntheses programs at that time, it was difficult to make wholesale line identifications without significant numbers of errors. This situation has changed dramatically in the intervening years with the development of spectral line synthesis tools that make use of extensive atomic line libraries (as well as the organization of the line libraries themselves).  Inspired by the {\\it Copernicus} atlas, we decided to address the need for the identification of lines in high-dispersion far-UV spectra along the B main sequence. Thus, the first goal of our work is to select and provide far-UV spectral data for main sequence B0 to B9 stars. The atlas spectra should be sharp-lined enough to permit the resolution of a majority of dominant individual atmospheric lines and to provide spectral data products in a common, continuous-spectrum format ranging in wavelength from near the short-wavelength limit of the {\\it FUSE} regime to the red wing (1225\\,\\AA) of Lyman\\,$\\alpha$.  Such an assemblage requires not only the cojoining of all {\\it FUSE} spectral detector segments but also (for B0) of second-order scans of {\\it Copernicus's} ``U1\" multiplier, and echelle spectra from {\\it IUE} or the {\\it Hubble Space Telescope Imaging Spectrograph (STIS)}. This is necessary to extend our coverage from the {\\it FUSE} long-wavelength limit of 1188\\,\\AA~ to 1225\\,\\AA, as justified below.  Our second goal is to choose three spectra among our atlas sample that capture the low, high, and midpoints of the excitation range of ions in B star atmospheres and to identify all possible absorption lines that make a noticeable contribution to these spectra according to published atomic line libraries. This addition changes the emphasis of the atlas from the classical high-level montage presentation to a ``detailed\" attribution of the spectral features. Our third goal is to discuss those metallic lines identified from the second goal in the context of diagnostics of effective temperature, luminosity, and heavy element abundances.  This paper is organized as follows. The selection of spectra and the methodology for data conditioning and line identification are set forth in $\\S$2. Section $\\S$3 displays portions of the atlas and a detailed list of all (several thousand) of our line identifications as well as a list of ``clean\" lines across much of the B-star domain. We also discuss an example of the important C\\,III 1176\\,\\AA~ region used as a temperature diagnostic for this spectral type. In $\\S$4 we give relevant statistics from our identifications for the B0, B2, and B8 exemplars we call spectral templates. We provide a comparison of our results with an earlier critique of the Rogerson \\& Ewell (1985) line attributions for the $\\tau$\\,Sco atlas. We finally give commentary on a number of possible spectral markers for physical conditions in these stars' atmospheres. ",
        "conclusions": "\\subsection{Line statistics}  The number of lines found in our spectral templates for B0, B2, and B8p is 2288, 1612, and 2469, respectively. Of these, 7.9\\%,\\footnote{This rate may be compared to Rogerson \\& Ewell's (1985) stated unidentified rate of 39\\%, a rate that is surely low because many of their identified lines still do have not published log\\,$gf$'s.}, 2.1\\%, and 2.2\\% could not be identified from our line library. The percentages of lines that are overpredicted are 12.5\\%, 5.5\\%, and 5.6\\%, respectively.  Conversely, similar percentages, namely 13\\%, 7\\%, and 6\\% of our found features, have ``uncertain\" (low log\\,$gf$) criterion. We expect that the great majority of these are actually valid identifications and their log\\,$gf$s are too low. For example, random spot checks of the log\\,$gf$'s required to match the observed strengths with our spectral syntheses suggest that the log\\,$gf$ error distribution is likewise consistent with a broad Gaussian with two approximately equal ``weak\" and ``strong\" tails.  Photospheric lines account for 93.7\\%, 87.0\\%, and 91.3\\% of the identifications in the respective templates. The few percent complement of these are due to ISM features, which become increasing numerous with decreasing wavelengths. For example, shortward of 1148\\,\\AA, we identified a total of 200 H$_2$ lines that are either dominant or secondary contributors to absorption features. Of these only two identifications are clearly H$_2$ primary lines in the $\\tau$\\,Sco spectrum, a star located along the ``Scorpius ISM tunnel.\" More typically, H$_2$ lines comprise the overwhelming contribution of the ISM component in our B2 and B8p template spectra and account for the lower two percentages (87.0\\%, 91.3\\%) for these cases compared to the overall identification rates.    As to photospheric features, we made a total of 5800 photospheric atomic line identifications in the three template stars, of which 4216 are unique. We expect that synthetic photospheric and ISM spectra could be computed solely from this list of unique lines. This also means that altogether our list includes 13.7\\% of our adopted line library tabulations over its range of 949 to 1225\\,\\AA. We believe this high rate (particularly considering the inclusion of excited ions in our line list) speaks to the general success of our identification efforts.  We estimate that this percentage would be as high as 15-16\\% if we excluded wavelength regions blanketed by saturated resonance, hydrogen, or H$_2$ ISM lines.  From our line statistics we find a few trends with spectral type. The first of these is the a comparatively high rate of uncertain identifications and nonidentifications for B0 spectral features, suggesting an incompleteness in log\\,$gf$'s  for transitions arising from thrice ionized Fe-group elements. However, at least part of this is likely a residual of the problem that interfered with the identification efforts by RE85. Many of RE85's still unconfirmed identifications may well prove to be correct at some future point when quantitative measurements of these lines are made.  A second trend with spectral type is in the number of far-UV photospheric lines identifed across the range of B spectra. We found a local minimum at type $\\sim$B2 in the number of detected lines.  Some 20-25\\% of the comparative dearth of identifications at type B2 arises from obscuration of photospheric lines by the larger wavelength regions dominated by strong ISM H$_2$ and photospheric resonance lines arising from ions like Si\\,III. The rest of this dearth is a consequence of the ionizations changing from thrice and single ionization states to second ionizations as one moves toward B2 from  B0 and B8. Thus, the addition of lines of twice-ionized states does not compensate for the loss of lines from the other two ion stages. An additional detail is that over 91\\% of the identified photospheric lines in the B2 spectrum are also visible (as primary or secondary lines) in either the B0 or B8 spectrum. Conversely, some 67\\% and 53\\% of the lines in the B0 and B8 spectrum are {\\em not} visible in the B8 and B0 spectrum, respectively. Some 608 primary or secondary lines remain visible across the entire B main sequence domain.  A third trend is that the percentage of isolated primary lines in resolution wavelength bin groups declines dramatically across the early B spectral range (from 81\\% at B0, 66\\% at B2, to 59\\% at B8p) and also with decreasing wavelength. Some of these ``isolated\" primary lines are barely resolved, and these can be utilized only in a sharp-lined spectrum. For this reason the list of least contaminated lines (Table\\,4) runs to only 5\\% of the total line list.  The general expectation that the Fe- group lines dominate the far-UV spectra of hot stars is borne out by our identifications. Indeed iron itself dominates this population. Numbers of iron lines identified in the associated ion stages are exhibited in Table\\,5. Lines of chromium, followed closely by manganese, are a distant second to the incidence of iron lines.     \\subsection{Previous critique of Rogerson-Ewell identifications in $\\tau$ Sco atlas}  Cowley \\& Merritt (1987; hereafter ``CM87\") employed Monte Carlo techniques to test for the presence of several ions identified in the $\\tau$\\,Sco atlas by Rogerson \\& Ewell (1985). CM87 have critiqued RE85's claims that lines of certain ions in particular are present in this spectrum. Here we offer comments on these differences based on our own identifications specifically in the region 949\\,\\AA~ to 1225\\,\\AA:  \\no {\\it N II:~ } CM87 questioned whether excited N\\,II features are present. However, since they did not list the earlier claimed identifications by RE85, we cannot address their objections directly. Overall, we found 6 excited lines from this ion in our photospheric syntheses (including 3 resonance lines apparently not in dispute). Since these excited N\\,II lines arise from levels at 13.5\\,eV, they can hardly be ISM lines. Thus, the presence of photospheric N\\,II lines in this star's far-UV spectrum is likely.  \\no {\\it N I:~ }  Our {\\sc SYNSPEC} syntheses predicted no detectable N\\,I photospheric resonance lines. This implies that features found at these wavelengths are formed in the ISM. However, the photospheric syntheses also predict 8 or 9 N\\,I lines (one is designated uncertain). All of these arise from levels at 3.6\\,eV. Therefore at least several photospheric lines from this ion are likely to be present in the far-UV $\\tau$\\,Sco spectrum as well as later types. Lines with this excitation potential are generally not primarily formed by the ISM, but we cannot rule out a secondary contribution.  \\no {\\it Ne II:~ } RE85 noted that most of their identifications of Ne\\,II lines are blended features. One of these is the line at 1181.3\\,\\AA.~ We confirm this, the only possible detection of this ion.  Nonetheless, we have designated 1181.3\\,\\AA~ as an uncertain identification.  \\no {\\it P IV:~ } CM87 reported that identifications of lines of this ion were insecure in this spectrum. However, we found 17 lines as P\\,IV in our photospheric syntheses, all but one of which we consider are likely identifications. In addition, we confirm the detection of at least two P\\,IV lines in the {\\it IUE} wavelength range that CM87 considered ``marginally significant.\" We were able to do this on the basis of many more {\\it IUE} observations made subsequent to the CM87 study.  The presence of this ion is secure.  \\no {\\it P III:~} CM87 questioned RE85's claim of 91 P\\,III lines in this spectra. In the far-UV, we found only four lines of this ion.  Of these only 1184.2 and 1194.7\\,\\AA~ are excited lines. Because these features are among the strongest predicted subordinate lines for this ion in our spectral syntheses, we are able to confirm the identification in strong P\\,III lines in the photospheric spectrum.  \\no {\\it S II:~ }  CM87 agreed with RE85 that lines of this ion are present, but not on the basis of as many lines they claimed. We identified only 4 lines for this ion, of which 2 are marked ``uncertain.\"  The presence of this ion is probably secure, but (as also for P\\,III) there is at best a marginal hope of using the few available S\\,II lines for diagnostic purposes.  \\no {\\it Mn III:~ }  CM87 stated that they found ``marginal support... for the presence a few of the strongest Mn\\,III lines.\" They suggested that a future study might find the abundance of this element to be subnormal.  However, we have identified 135 lines of this ion in the same spectrum, of which no more than ${\\frac 13}$ are uncertain identifications. The presence of this ion in this spectrum is certainly secure.  \\no {\\it Zn III:~} Like CM87 we are unable to confirm identifications of this ion in the far-UV.     \\subsection{Temperature and chemical Indicators} \\label{chem}  In this section we indicate lines that may be used in far-UV spectra of {\\it sharp-lined} B stars near the main sequence to refine diagnostics of effective temperature, chemical composition, and occasionally log\\,$g$. For the B V stars most of the interesting abundances are likely to relate to products of the CNO-cycle and to evidence of Bp compositional anomalies, usually Si, Cr, and/or Mn. However, we include lines of other elements that may serve as metallicity indicators.  In our descriptions below we give weight to the ``clean\" spectral lines listed in Table\\,4. However, we also include some lines that have minor blend contributions if they are present over a range of spectral types. We made our judgments by explicitly surveying the B0, B1, B2, B5, B8, and B9 spectra from this atlas. In general, those lines of singly-ionized species exhibit the greatest variations in strength along the B sequence. Lines arising from thrice-ionized atoms, except for abundant elements like silicon, generally exist only in at types B0 and B1. Those lines exhibiting the smallest changes in equivalent width have moderate excitations (4-8 eV) of doubly ionized atoms - a delicate balance between excitation and ionization effects. Being least sensitive to changes in photospheric temperature, and not being sensitive to wind conditions, these lines are the best indicators of abundance. We prefix below the ion name of Si\\,IV with an asterisk because the members of the UV resonance doublet (1394\\,\\AA, 1403\\,\\AA) exhibit an obvious wind component in the great majority of early B-type spectra.  Because wind lines can be strongly contaminated by both emission and absorptions deeper in the star's atmosphere, lines listed for Si\\,IV are especially valuable diagnostics of temperature and/or abundance.  Reasonable ionization criteria can often be framed from the following ion ratios in our B star atlas spectra: carbon (IV/III/II/I), nitrogen (IV/III/II/I), oxygen (III/II/I), silicon (IV/III/II), phosphorus (V/IV/III/II), sulfur (IV/III/II), chlorine (III/II), chromium (IV/III/II), manganese (III/II), iron (IV/III/II), and cobalt (II/II). In those cases where lines of three ion stages can be detectable, all three are generally visible only down to types B1 or B2. Exceptions to this statement are silicon and sulfur. For these elements lines of three ionization states may be found in spectra for type B5 and even later.  We list lines in order of ion atomic sequence that are usually visible in two of our three major template spectra (B0, B2, and B8) and which usually can be found in our Table\\,3. Lines that appear in only our B0 or B0-B1 spectra include He\\,I 958.6\\,\\AA, C\\,IV 1107.5 and 1168.9\\,\\AA, P\\,V 1117.9\\,\\AA, 1128.0\\,\\AA, and Ca\\,III 1116.0\\,\\AA. Pellerin et al. (2002) have discussed the special case of He\\,II 1084.8\\,\\AA~ line, whose diagnostic value is compromised by blends of N\\,II lines. We will not repeat these in the following ion listing.   \\no {\\bf C III:~ } The feature at 977.0\\,\\AA~ is a well known pressure-sensitive resonance line in B spectra, although in practice its profile in Galactic B star spectra is marred by the appearance of molecular H$_2$ lines (Pellerin et al. 2002, Walborn et al. 2002). Nonetheless, researchers can make use of this line along with the C\\,III 1176\\,\\AA~ complex to determine both an early B star's spectral type and luminosity class, particularly in Magellanic Cloud stars well out of the Galactic plane (e.g., Walborn et al. 2002). As noted in $\\S$\\ref{node1176}, the weak C\\,III line complex between 1165.6\\,\\AA~ and 1165.9\\,\\AA~ extends from O8 to B1-B2 on the main sequence.  \\no {\\bf C II:~ } Strengths of the subordinate (9.6\\,eV) 961.2, 972.3, and 996.3\\,\\AA~ lines of this low-ionization species increase slowly and persist through B-types. Strengths of the high excitation (13.7\\,eV) 1138.9 and 1139.3\\,\\AA~ lines persist to B5 before weakening to invisibility. Lines at 1009.9 and 1010.3\\,\\AA~ attain maximum strengths at B2-B3, but they then decrease slowly and are still visible at type B8.  \\no {\\bf C I:~} A resonance feature, 1135.8\\,\\AA~ is the only line clearly visible across all the B subtypes.  A second resonance line at 1138.5\\,\\AA~ appears at B2 and remains visible for A types and later.  \\no {\\bf N IV:~ } The high excitation 1131.4\\,\\AA~ feature is visible from at least O9 through type B1. This appears to be a good indicator of T$_{eff}$ above 25,000\\,K. However, note that this line is contaminated by O\\,III 1131.0\\,\\AA, which also fades at about the same rate as N\\,IV. It is absent by type B2.  \\no {\\bf N III:~ }  The N\\,III complex near 979.8\\,\\AA~ fades very fast with spectral type. As noted by Pellerin et al. (2002), the lines at 987.7 and 991.5\\,\\AA~ fade into nearby atmospheric or H$_2$ blends at type B2, and 1184.5\\,\\AA~ does so at $\\sim$B5.  \\no {\\bf N II:~ }  The features at 1083.9\\,\\AA~ and in the range 1085.5\\,\\AA~ to 1085.7\\,\\AA~ are resonance lines that undergo a maximum strength at about B2 and remain visible at B8.  \\no {\\bf N I:~ } The photospheric N\\,I lines at 1099.0, 1172.0, 1177.6, and 1200.2\\,\\AA~ have excitations in the range 0-3.6 eV and increase rapidly with spectral type.  The 1184.2\\,\\AA~ line first makes its appearance at type B5 and strengthens at least to early A star spectra.  Of all the N\\,I lines 1200\\,\\AA~ is the cleanest in our atlas spectra.  \\no {\\bf O III:~ }  All five lines from this ion are highly excited and decrease to only marginal detectability at B2. These are at 1016.7, 1149.6, 1157.0, 1196.7, and 1199.9\\,\\AA. Of these 1199\\,\\AA~ is the least blended and most suitable for quantitative measurement in early B spectra. Pellerin et al. (2002) have noted the utility of 1149\\,\\AA~ to hotter stars and particularly in oxygen-rich WC stars.  \\no {\\bf O II:~ } Only one line, the high excitation line at 1130.1\\,\\AA, remains visible through the whole spectral range, attaining a maximum at about type B8.  \\no {\\bf O I:~ } The line O\\,I 990.2\\,\\AA~ is visible in all our templates. For the $\\tau$\\,Sco spectrum the feature is probably a photospheric/ISM blend Lines at 1041.6 and 1127.4\\,\\AA~ appear at B2 and steadily increase in strength.  \\no {\\bf Al II:~ } The moderate excitation line at 1048.5\\,\\AA~ appears in B2 spectra and increases in strength with advancing type.  \\no {\\bf $^*$Si IV:~} The most visible far-UV photospheric line of this important ion is 1066.6\\,\\AA. This feature decreases in strength and fades at type B6. The line at 1122.5\\,\\AA, as also noted by Pellerin et al. 2002, can be a good temperature diagnostic through B5, but it becomes contaminated by a strong C\\,I aggregate.  \\no {\\bf Si III:~ } This ion is represented by nicely contrasting excitation potentials (15-16 eV and 6.5 eV). The former are at 967.9, 1032.8, and 1207.5\\,\\AA, which decrease in strength and generally disappear beyond type B5. Two saturated lines at 1206.5\\,\\AA, already strong at B0, become major features in late B spectra. The high excitations of the lines at 993.5, 1109.9, and 1113.2\\,\\AA~ balance changes in silicon ionization (also noted by Pellerin et al. 2002), causing their strengths to undergo a maximum at B0-B2.  \\no {\\bf Si II:~ }  The moderate excitation (5 eV) lines 1224.2 and 1224.9\\,\\AA~ steadily increase in strength with type. The only readily visible Si\\,II line in the far-UV wavelength range is 1193.2\\,\\AA~ and 1194.4\\,\\AA, which remain visible at B2 and later types.  \\no {\\bf P II:~ } Two resonance lines at 1149.9 and 1152.8\\,\\AA~ steadily increase through the B types. This makes them prominent diagnostics in late B spectra.  \\no {\\bf S IV:~ } The resonance line at 1073.5\\,\\AA~ strengthens throughout the B range (see also Pellerin et al. 2002).  \\no {\\bf S III:~ } The relatively low ionization potential of this ion allows the resonance lines at 1201.7 and 1202.2\\,\\AA~ to increase only slowly through the B types, making them useful abundance indicators. The only S\\,III line in the {\\it FUSE} range, 1143.8\\,\\AA, is marginally resolved. It also increases slowly through the B range.  \\no {\\bf S II:~ } Several 3\\,eV lines at 1095.0, 1116.1, 1124.4, 1124.9 (strong), and 1131.6\\,\\AA~ make their appearances at B1-B2 and strengthen with advancing type. The less excited lines at 1006.2 and 1045.7\\,\\AA~ increase more slowly, making them possible abundance indicators (along with the just noted S\\,III lines).  \\no {\\bf Cl III:~ } The resonance line at 1015.0\\,\\AA~ undergoes a maximum at type B2, beyond which it becomes contaminated by an Fe II line. The resonance line at 1009.7\\,\\AA~ is visible only near type B0.  \\no {\\bf Cl II:~ } The line at 1087.3\\,\\AA, arising from a level at 3.5\\,eV, is the only Cl\\,II feature available in the far-UV, but is blended with an Fe\\,II and Ni\\,II line at B2. The far-UV Cl lines are not suitable for constraining thermodynamic information about B-star atmospheres. Abundances of chlorine from far-UV lines of Cl\\,II or Cl\\,III may be estimated only with great care.  \\no {\\bf Ti III:~ } A weak line at 1080.7\\,\\AA~ increases strength from B0 to B2 but beyond B5 becomes at best semiresolved between nearby Fe\\,II and Mn\\,III lines.  \\no {\\bf V III:~ } A single resonance line at 1148.4\\,\\AA~ is visible throughout the B types.  \\no {\\bf Cr IV:~ } The high excitation (31 eV) 1103.3\\,\\AA~ line is visible in late O to B2 spectra. However, the line can be resolved in sharp-lined spectra.  \\no {\\bf Cr III:~ } The number density of Cr$^{2+}$ peaks at about B2, and the strengths of several available low-excitation lines increase noticeably. These lines include 967.5, 1040.0, 1041.3, 1047.0, 1055.8, 1065.3, 1073.7, 1119.9, 1132.7, and 1146.3\\,\\AA. Arising at 4 eV, the 1181.7 and 1187.5\\,\\AA~ lines slowly {\\it decrease}, suggesting that they would make the best chromium abundance diagnostics. This is a fortunate circumstance, e.g. for the study of abundances in late-type chromium-rich B peculiar spectra, as these lines are just within the wavelength limits of both {\\it IUE} and {\\it FUSE} high dispersion spectra.  \\no {\\bf Cr II:~ } There are no reliable lines in the {\\it FUSE} regime for this ion that are visible across the entire B domain. The single 1219.5\\,\\AA~ line becomes visible at about O9.5 and becomes contaminated by blends beyond type B2.  \\no {\\bf Mn III:~ } The only Mn\\,III line (5 eV) visible at all B types is 1052.7\\,\\AA. It increases with type at a moderate rate.  \\no {\\bf Mn II:~ } Just one semi-resolved line, 1161.2\\,\\AA, is visible for this ion.  To the extent it can be resolved from nearby lines its strength increases slowly with type. Therefore it can potentially serve as an abundance diagnostic, particularly in late B spectra where it is likely to be used to distinguish between normal B and Mn-rich Bp populations.  \\no {\\bf Fe IV:~ } A few lines with excitations of 19-20 eV are visible in the B0 to B2 range, and they decrease in strength with type: 1022.6, 1047.2, 1135.2, 1156.2, and 1157.4\\,\\AA. Of these lines 1157\\,\\AA~ is among the most useful iron ionization diagnostics in early B spectra, first because of its nearly constant strength and second because of its proximity to the strengthening Fe\\,III 1157.5\\,\\AA~ line, which is a serviceable line for middle B spectra.  \\no {\\bf Fe III:~ } For the most part the numerous lines of Fe\\,III decrease in strength within the B range, e.g. 961.7, 997.5, 1005.1, 1030.9, 1045.9, 1047.9, 1057.5, 1058.5, 1076.5, and 1117.3\\,\\AA. However, the features at 959.3 and 979.0\\,\\AA~ increase in strength slowly, attaining a maximum at B2, and then decrease very slowly by 5-10\\% at B8.  Contrasted with other Fe\\,III lines, these can be used for abundances. Arising at 6.1 eV, the 1063.4\\,\\AA~ line increases in strength to B2 and remains roughly constant at later types. Given this dependence, the line might be used as a secondary abundance indicator in late B stars. A few other lines exhibit more complicated dependences. For example 1164.7\\,\\AA~ (9.9 eV) decreases in strength from B0 to B2 and disappears at later types. In addition to the Fe\\,IV lines noted above, this line is a useful temperature diagnostic for the B0-B2 stars.  Because the 1221.0\\,\\AA~ line undergoes a maximum at B2, it may serve as a secondary abundance indicator in this narrow domain and a secondary temperature indicator for early and late B spectra.  \\no {\\bf Fe II:~ } All photospheric Fe II lines increase in strength through the B spectral domain (but note the presence of ISM lines of both Fe\\,II and Fe\\,III in some B and O-type stars (Walborn et al. 2002). Only isolated Fe\\,II lines at 994.5, 1075.6, 1115.8, 1150.6, 1160.9, and 1190.8\\,\\AA~ are visible for all B types. Several are visible only near type B2: 1091.5, 1112.0, 1116.9, 1121.9, 1124.1, 1125.4, and 1195.4\\,\\AA.  \\no {\\bf Co IV:~ }  Seven high-excitation (16-17 eV) lines are probably visible in late O stars and at type B0, of which three are uncertain identifications. The likely identifications are at 1115.1, 1116.0, 1118.2, and (to type B1) 1188.0\\,\\AA.  \\no {\\bf Co III:~ } Lines at 1049.7 and 1116.8\\,\\AA~ arise from low (1.9 eV) and high (9 eV) excitation states, respectively. The contrast this pair of lines offers to one another suggests that they could be used to form a temperature diagnostic. Both lines increase in strength with advancing type.  \\no {\\bf Ni III:~ } Because 973.7\\,\\AA~ increases strength through B2 to become blended at B5, it does not offer a reliable ionization diagnostic. The 979.2\\,\\AA~ line shows only a slight increase in strength through the B types. This suggests it might be used as an abundance indicator in B stars.  \\no {\\bf Ni II:~ }  Lines of this ion first appear at B2 and then strengthen through at least B8. The primary lines are at 1076.0, 1134.5, 1173.4, and 1181.6\\,\\AA.    It can be added that a Pt\\,III line at 1080.0\\,\\AA~ in the HD\\,182308 spectrum (B8\\,Vp) is the only platinum line we have found. As this is the among the strongest unblended lines of this element in the far-UV spectrum of a B8-B9 star, it is not surprising to find it in a Hg-Mn B peculiar-type spectrum, It will certainly be present in others.  Finally, the Ar\\,I 1048.2\\,\\AA~ line is strongly contaminated by the ISM in all or nearly all of the atlas spectra, including $\\tau$\\,Sco. In general, a number of atomic ISM lines appear in all three spectra, including singly ionized atoms of C, N, O, and Cl. In spectra later than B2 it is often difficult to assess the relative strengths of the photospheric and ISM contributions, particularly since either contribution alone can saturate the core.   We wish to express our appreciation to Dr. John B. Rogerson for permitting us to include the Rogerson-Upson ({\\it Copernicus}) atlas of $\\tau$\\,Sco in this study. We also thank Drs. S. Adelman and C. Proffitt for providing the author with an unpublished, preliminary {\\it IUE/SWP} atlas of the B2\\,IV star $\\gamma$\\,Peg. The quality of this paper was substantially improved thanks to a number of comments by an anonymous referee. This work was supported by a NASA Astrophysics Data Analysis grant NNX07AH57G and a grant from the {\\it FUSE} project NNX08AG99G to the Catholic University of America.  \\clearpage"
    },
    "0911/0911.0678_arXiv.txt": {
        "abstract": "We present a survey, using the \\chandra\\ X-ray observatory, of the central gravitating mass profiles in a sample of 10 galaxies, groups and clusters, spanning $\\sim$2 orders of magnitude in virial mass. We find the total mass distributions from $\\sim$0.2--10\\reff, where \\reff\\ is the optical effective radius of the central galaxy, are remarkably similar to powerlaw density profiles. The negative logarithmic slope of the mass density profiles, $\\alpha$, systematically varies with \\reff, from $\\alpha \\simeq$2, for systems with \\reff$\\sim$4~kpc to $\\alpha \\simeq$1.2 for systems with \\reff\\gtsim 30~kpc. Departures from hydrostatic equilibrium are likely to be small and cannot easily explain this trend. We show that the conspiracy between the baryonic (Sersic) and dark matter (NFW/ Einasto) components required to maintain a powerlaw {\\em total} mass distribution naturally predicts an anti-correlation between $\\alpha$ and \\reff\\ that is very close to what is observed. {The systematic variation of $\\alpha$ with \\reff\\ } implies a dark matter fraction within \\reff\\ that varies systematically with the properties of the galaxy in such a manner as to reproduce, without fine tuning, the observed tilt of the fundamental plane. We speculate that establishing a nearly powerlaw total mass distribution is therefore a fundamental feature of galaxy formation and the primary factor which determines the tilt of the fundamental plane. ",
        "introduction": "\\label{sect_introduction} The global optical properties of giant elliptical galaxies are often parametrized by three key quantities, the effective (half-light) radius (\\reff) of the stellar light, the mean surface brightness (or luminosity) and the central line-of-sight velocity dispersion (\\sigmac). In this three-dimensional parameter space, they occupy a narrow ``fundamental'' plane \\citep{dressler87a,djorgovski87a}, the orientation of which differs significantly from naive expectations from the virial theorem (assuming homology and a constant mass-to-light ratio, $\\gamma$). The relative importance which deviations from homology, variations in the stellar population properties and changing dark matter fractions play in producing this ``tilt'' have been hotly debated \\citep[\\eg][and references therein]{renzini93b,hjorth95a,ciotti96a,prugniel97a,graham97a,gerhard01a,padmanabhan04a,trujillo04a,cappellari06a,bolton07a,tortora09a}, with recent work tending to emphasize the importance of dark matter \\citep[\\eg][]{cappellari06a,bolton07a}. To maintain the thinness of the fundamental plane, however, the dark matter fraction within $\\sim$\\reff\\ must be tightly correlated with the optical properties of the galaxy \\citep[\\eg][]{ciotti96a}.   Strong observational evidence for the existence of massive dark matter halos around early-type galaxies has been provided by independent observational constraints from X-ray studies, lensing and stellar dynamics (\\eg\\ \\citealt{humphrey06a}, hereafter \\citetalias{humphrey06a}; \\citealt{mathews03a} and references therein; \\citealt{kochanek95a,griffiths96a,gavazzi07a,gerhard01a,thomas07a}). Although the total mass distribution within \\reff\\ is dominated by the stars, the dark matter fraction is non-negligible, and stellar dynamics and lensing studies (most of which are restricted to within $\\sim$\\reff) suggest that it establishes a conspiracy with the luminous matter to produce total mass density ($\\rho_m$) profiles close to  $\\rho_m \\propto R^{-2}$ (where R is the radius), \\ie\\ the ``singular isothermal sphere'' \\citep[\\eg][]{kochanek95a,treu04a,koopmans06a,koopmans09a,kronawitter00a}. This ``bulge-halo conspiracy'' is similar to that which establishes flat rotation curves in disk galaxies.  Total mass distributions approximately consistent with $\\rho_m \\propto R^{-2}$ have long been reported from hydrostatic X-ray analysis of nearby elliptical galaxies at scales much larger than \\reff\\ \\citep[\\eg][]{trinchieri86a,thomas86a,serlemitsos93a,nulsen95a,kim95a,rangarajan95a,matsushita98a} and, in their \\rosat\\ imaging analysis of two galaxies, \\citet{buote94,buote98d} actually found an isothermal sphere potential to be preferred over some other mass distributions. With the improved capabilities of \\chandra\\ and \\xmm, much tighter constraints on radial mass profiles have now been reported for a wider array of galaxies \\citep[\\eg][]{humphrey06a,humphrey08a,humphrey09d,fukazawa06a,osullivan04b,osullivan07b,buote02a}, which similarly resemble $\\rho_m \\propto R^{-2}$ \\citep[as pointed out by][]{gavazzi07a}, although not exactly so over all radial scales \\citep{romanowsky09a}. \\citet{fukazawa06a} fitted a model of the form $\\rho_m \\propto R^{-\\alpha}$ to the combined mass profiles of their galaxy sample outside 10~kpc, finding $\\alpha=1.67\\pm0.33$.  At the scale of massive galaxy clusters, hydrostatic X-ray analysis has revealed much less cuspy mass distributions \\citep[$\\alpha \\sim$1--1.4; \\eg][]{buote04b,voigt06a,arabadjis02a,david01a}. In particular, \\citet{lewis03a} and \\citet[][]{zappacosta06a} studied the very relaxed clusters A\\thin 2029 and A\\thin 2589, finding that the {\\em total} mass profiles were in good agreement with the NFW shape expected for the dark matter {\\em only}, once again suggesting some kind of conspiracy between the luminous and dark matter to produce an approximately powerlaw {\\em total} mass distribution in the core of the system. Stellar dynamics and lensing studies of clusters have, similarly, found less cuspy total mass profiles than would be expected for an unmodified NFW component plus the stellar mass \\citep[\\eg][]{sand04a,kelson02a}.  In order to investigate this disparate behaviour at different mass scales, in this paper we carry out a uniform X-ray analysis of a sample of relaxed galaxies, groups and clusters to investigate the shape of the innermost mass distributions. X-ray analysis is ideally suited for this study, since it allows straightforward mass measurements in individual systems over a wide radial range, from the baryon dominated regime (R\\ltsim \\reff) to regions where the dark matter dominates the gravitating mass (R$\\gg$\\reff), thus providing adequate leverage to elucidate any conspiracy between these different components. For relaxed systems, hydrostatic equilibrium is believed to be an excellent approximation, with nonthermal effects contributing no more than $\\sim$20\\%\\ of the total pressure \\citep[\\eg][hereafter \\citetalias{humphrey09d}]{churazov08a,nagai07a,piffaretti08a,humphrey09d}. In our previous papers (\\citetalias{humphrey06a,humphrey09d}; \\citealt{gastaldello07a,zappacosta06a}) we have carried out detailed decompositions of the mass distribution of these systems into baryonic and non-baryonic components. We have not, however, investigated in detail the relationship between these two components. In light of the apparent conspiracies between them at both cluster and galaxy scales, in this present paper, we adopt the more pragmatic approach of examining whether the {\\em inner} part of the mass profile can be modelled as a powerlaw, and investigating whether its slope varies systematically.  Throughout this work, we assume a cosmology of \\hnought=${\\rm 70\\ km\\ s^{-1}\\ Mpc^{-1}}$, \\omegam=$0.3$ and \\omegalambda=0.7. All error-bars, unless stated otherwise, correspond to 1-$\\sigma$ confidence regions. ",
        "conclusions": "\\subsection{A luminous-dark matter conspiracy} \\begin{figure} \\includegraphics[width=2.3in,angle=270]{f3.eps} \\caption{{\\em Left:} Comparison of the mass profiles of NGC\\thin 4261 (circles), MKW4 (triangles) and A\\thin 2589 (stars), normalized to have the same mass at 100~kpc. Note that the shapes of the mass profiles are different even over overlapping radial ranges. {\\em Right:} The same, but scaling the radius with respect to the virial radius ($R_{vir}$) and renormalizing each profile to have the same mass at 0.001$R_{vir}$.\\label{fig_mass_profiles2}} \\end{figure} In general, the stars are believed to dominate the gravitational potential within $\\sim$\\reff, with the dark matter dominating outside this scale \\citep[\\eg][]{brighenti97a,gerhard01a}; in particular, we have shown this for the systems studied here \\citep[\\citetalias{humphrey06a,humphrey09d}][]{gastaldello07a,lewis03a,buote04b,zappacosta06a}. Therefore,  the approximately powerlaw mass distributions shown in Fig~\\ref{fig_mass_profiles} indicate an apparent conspiracy between the luminous and dark matter to produce a scale-free total mass distribution, at least within $\\sim$10\\reff. At much larger scales, there is evidence that the profiles deviate from this simple shape \\citep[\\eg][]{lewis03a,vikhlinin06b,humphrey06a,gastaldello07a,romanowsky09a}. Nevertheless, in the inner regions of the systems, this effect is similar to the ``bulge-halo'' conspiracy found at much smaller scales (\\ltsim\\reff) in lensing and stellar dynamics studies \\citep[\\eg][]{treu04a,koopmans06a,koopmans09a}. {It is important to appreciate that the observed powerlaw mass profiles cannot simply arise from the arbitrary superposition of any reasonable stellar and dark mass components. To illustrate this quantitatively, we simulated a powerlaw-like mass profile and fitted it with a model comprising Sersic and dark matter components (as described in detail in \\S~\\ref{sect_toy_model}). For any given galaxy luminosity, we found that both \\reff\\ and the normalization of the dark matter component cannot vary by more than $\\sim$25\\%\\ from their ``best'' values while maintaining mean absolute fractional deviations between the model and data of less than 10\\%\\ (as indicated by our fits in \\S~\\ref{sect_mass_slope}).}  The origin of this conspiracy is unclear, but it is likely tied to the complex interaction of baryons and dark matter in the centres of galaxies (for example, \\citealt{robertson06b} point out the importance of gas physics in maintaining the tilt of the fundamental plane, which may be related to this conspiracy; \\S~\\ref{sect_fp}). Unfortunately these processes are very poorly understood; the predicted central dark matter density cusps due to ``adiabatic contraction'' \\citep{blumenthal86a,gnedin04a} have not been observationally confirmed \\citep[\\eg][]{humphrey06a,zappacosta06a,gnedin07a,dutton07a}, suggesting that other effects, such as dynamical friction, may be important \\citep[\\eg][]{elzant04a,dutton07a,abadi09a}. It is unclear, therefore, that any model is so far robustly predicting a powerlaw {\\em total} mass distribution arising from these effects.  Since the data were not all fitted over the same physical scales, it is important to assess whether the systematic variation in $\\alpha$ we observe is simply an artefact of these different fit ranges. In Fig~\\ref{fig_mass_profiles2} we show the mass profiles of three representative systems with different \\reff, arbitrarily scaled for clarity, and shown as a function of physical radius (left) and fraction of the virial radius \\citep[right; deduced from the virial masses reported in][]{buote07a}. In conjunction with Fig~\\ref{fig_mass_profiles} (which shows the same profiles as a function of fraction of \\reff), it is immediately clear that {\\em throughout overlapping radial scales} the slopes of the three mass distributions are all very different, so that the fit range is unlikely to be the cause of the trends observed in Fig~\\ref{fig_mass_slope}.  Although all of the systems we considered are morphologically relaxed, which should eliminate objects with the largest departures from hydrostatic equilibrium, it is still important to consider whether deviations from hydrostatic equilibrium could be responsible for the observed trends. In order to affect the {\\em shape} of the mass profile, any source of nonthermal pressure must vary radially in a finely balanced manner, given that the gas pressure profile is consistent with a powerlaw mass density profile (Fig~\\ref{fig_mass_profiles}) that varies systematically with the host properties (Fig~\\ref{fig_mass_slope}). For the cluster A\\thin 2589, \\citet{zappacosta06a} concluded that, if the true mass distribution is much cuspier than inferred from the X-rays (for example, an unmodified NFW plus a stellar mass component), no plausible source of turbulent or magnetic pressure could make the X-ray measurements imply such a flat total density profile. \\citetalias{humphrey09d} carried out full mass decompositions for three of our galaxies, finding stellar mass to light ratios in excellent agreement with the predictions of stellar population synthesis models, suggesting that nonthermal pressure is very small (\\ltsim 10--20\\%) within \\reff. This was further supported by the good agreement between the X-ray constraints on the central black hole masses, and those obtained from optical methods for these objects. Furthermore, nine of our objects were included in the sample of \\citet{buote07a}, who showed that their dark matter halos inferred from X-ray mass analysis approximately exhibit the relationship between halo ``concentration'' (\\cvir) and total mass predicted from numerical structure formation simulations. This supports the idea that the shapes of the mass profiles for these systems are not significantly  in error due to non-hydrostatic effects.  In their survey of 15 lensing systems, \\citet{koopmans06a} found that $\\alpha$ ranged from $\\sim$1.8--2.4, while the \\reff\\ spanned $\\sim$2--17~kpc. In the left hand panel of Fig~\\ref{fig_mass_slope}, we overlay their measured $\\alpha$ {\\em versus} \\reff. We find that they are roughly consistent with the points from our X-ray sample, but at given \\reff, $\\alpha$ is generally larger. For strict consistency with our adopted photometric bands, their data may need to be adjusted, however. Since \\reff\\ in the K-band is typically smaller than in bluer bands (as an illustration of this effect, for the galaxies in the catalogue of \\citealt{labarbera08a}, we find that that \\reff\\ is on average $\\sim$50\\%\\ larger in the R-band than the K-band), this would likely involve shifting the \\citeauthor{koopmans06a} to the left, bringing them into better agreement with our results. By inspection, there is  a slight hint in their data  of the anti-correlation we observe between $\\alpha$ and \\reff, although it is not statistically significant. However, since their fits were all restricted to regions \\ltsim\\ 0.9\\reff, where the stellar mass dominates the potential, they are more likely to have been subject to local curvature in the mass distribution than the X-ray data were, given the better radial coverage of the latter.   \\subsection{Decomposing the mass profiles: a toy model} \\label{sect_toy_model} \\begin{figure} \\includegraphics[width=2.5in,angle=270]{f4.eps} \\caption{Sample simulated mass profile and associated mass decomposition. The data-points correspond to a powerlaw mass distribution with $\\alpha=1.8$, while the solid line indicates the best-fitting two-component (dark plus luminous matter) model. Error-bars, corresponding to a 10\\%\\ fractional uncertainty, are shown to guide the eye. Also shown are the corresponding best-fitting dark matter (dotted line) and stellar mass (dashed line) models. Note how the sum of two scale-dependent models produces an approximately powerlaw total mass distribution over the fitted range.\\label{fig_mass_decomposition}} \\end{figure} To investigate quantitatively the implications of the apparent conspiracy between dark and luminous mass components implied by Fig~\\ref{fig_mass_profiles}, we constructed a simple toy model by decomposing a pure powerlaw mass distribution into an NFW \\citep{navarro97} dark matter profile plus a de Vaucouleurs stellar mass component \\citep[using the deprojected Sersic approximation of][and fixing the Sersic index, $n=4$]{prugniel97a}. {In our previous work (\\citealt{buote07a}; \\citetalias{humphrey06a,humphrey09d}; \\citealt{gastaldello07a,zappacosta06a}), we have shown that decompositions of this type for the systems studied in the present work give parameters for the NFW model which are consistent with theory. The measured NFW concentration parameters are slightly higher than predicted median from N-body simulations, but this can be reasonably attributed to relaxed X-ray selected systems having formed early}. To perform the decomposition, we first constructed a set of mass ``data-points'' from the powerlaw model in 10 approximately logarithmically-spaced bins between $\\sim$0.4--200~kpc, which we then fitted with the two-component model by minimizing the rms of the fractional residuals. The normalization and scale radius of the dark matter component were allowed to vary freely, as were \\reff\\ and the normalization of the stellar model. We allowed the slope of the input mass distribution to vary systematically from $\\alpha=2.2$ to $\\alpha=1.35$. Below this limit, we found that the derived relationships (discussed below) became non-monotonic, which may represent the limitations of our simple parametrization rather than any physical insight, and so we prefer not to discuss this regime. While one might {\\em a priori} expect significant degeneracies in this decomposition given the number of free parameters, in fact the shapes of the Sersic and NFW mass profiles are sufficiently different that we found a unique decomposition for any input powerlaw, at least for the range of profiles and models we investigated. We show an example simulated profile, and the best-fitting dark and stellar mass components, in Fig~\\ref{fig_mass_decomposition}, clearly illustrating how they can conspire to produce an approximately scale-free model.  We overlay in Fig~\\ref{fig_mass_slope} (left panel) the best-fitting relation between \\reff\\ and $\\alpha$ obtained from our simple mass decompositions (solid line), which can be approximated to better than 1\\%\\ by $\\alpha = 2.31-0.54\\log R_e$. Without any fine tuning, it clearly captures the overall behaviour of the observed data reasonably well. Since the stellar light profiles of real galaxies may not be exactly de Vaucouleurs in form, we have also experimented with varying the Sersic index, n, from 3--5. Similarly, the NFW model used to parameterize the dark matter halo over the fitting range and may not be strictly correct if the baryonic and non-baryonic components interact gravitationally. Therefore, we have tried replacing it with the revised mass model of  \\citet{navarro04a}, allowing the mass slope index to fit freely and thus provide more freedom in parameterizing the radial distribution of dark matter. The dotted lines in Fig~\\ref{fig_mass_slope} bound the range of recovered relations, given these changes, which clearly bracket the observed data. In practice, the impact of changing the stellar light profile dominated this estimate of our model uncertainty.   Since our simple toy model can be freely renormalized, we cannot use it alone to make any explicit predictions for the relationship between $\\alpha$ and $M_{75}$. One way to make further progress, however, is to exploit the ``Kormendy relation'' between a galaxy's luminosity and \\reff\\ \\citep{kormendy77a}. For a given stellar mass-to-light ratio ($\\gamma_*$), $M_{75}$ must be adjusted  to keep the stellar component consistent with the Kormendy relation while maintaining the powerlaw total mass distribution for a given $\\alpha$ (and hence \\reff). We adopted a form for this relation which is appropriate for our data, estimated from the K-band galaxy luminosities and \\reff\\ tabulated by \\citet{gastaldello07a} and \\citetalias{humphrey06a}. We found a mean relation $\\log L/L_\\odot = 10.95 + 0.79 \\log R_e$, with an intrinsic scatter of $\\sim$0.11~dex, where $L$ is the total luminosity of the galaxy in the K-band and $L_\\odot$ is the luminosity of the Sun, assuming a K-band absolute Solar magnitude of 3.41. We assumed $\\gamma_*=1$, appropriate for an old stellar population \\citep[\\eg][]{maraston05a}. In Fig~\\ref{fig_mass_slope} (right panel) we overlay the predictions of this simple model (solid line), and an estimate of the range of uncertainty, which factors in the effects of varying the Sersic index of the stellar light by $\\pm 1$, the intrinsic scatter in the Kormendy relation and a $\\pm$50\\%\\ variation in $\\gamma_*$ (which dominates the error budget). Clearly the model captures the overall behaviour of the data fairly well.   {To verify that our model is self-consistent, it is important to determine if the parameters of the fitted dark matter model it predicts are reasonable. Since the toy model cannot be used alone to set the normalization of the mass profile, to do this we considered the simple model, folding in the Kormendy relation, described above. Starting with the $\\alpha$-$M_{75}$ relation predicted by this model, we simulated corresponding powerlaw mass profiles, and decomposed them into Sersic and dark matter components. We found monotonic relations between\\ \\mvir\\ and $\\alpha$ and between NFW \\cvir\\ and $\\alpha$, allowing us to recast this as a \\cvir-\\mvir\\ relation, which we can approximate by \\cvir$=c_{14} \\left( M_{vir}/10^{14} M_\\odot\\right)^{-\\beta}$, where $c_{14} = 10$ and $\\beta=0.17$. This is very close to the best-fitting \\cvir-\\mvir\\ relation found by \\citet{buote07a}, for which $c_{14}=9.0\\pm0.4$ and $\\beta=0.172\\pm0.026$. Exploring the estimated  range of uncertainty on the model linking $\\alpha$ and $M_{75}$, we found $c_{14}$ to lie in the range 6--16 and $\\beta$ from 0.13--0.18.  The total range of \\mvir\\ spanned by our predicted \\cvir-\\mvir\\ relation ($\\sim 12.5<log_{10} M_{vir} < 15.0$) is similar to the range spanned by the systems in our sample. It appears, therefore, that the parameters of the NFW component implied by our simple model are appropriate for the kinds of system studied in this work. }  \\subsection{Implications for the fundamental plane} \\label{sect_fp} \\begin{figure} \\includegraphics[width=2.5in,angle=270]{f5.eps} \\caption{The predicted ratio between the total mass and stellar mass enclosed within \\reff\\ from our toy model  ($\\gamma_m$), shown as a function of \\reff\\ (solid line). The dot-dash lines indicate our estimate of the uncertainty on the model, and the dashed line indicates the best-fitting linear approximation in the range $\\log R_e=$0.2--1.0. The data-points are the observed K-band M/L ratios from the data of \\citet{labarbera08a} (see text). Note the good agreement between the data and the toy model. \\label{fig_mass_to_light}} \\end{figure} The existence of a conspiracy between the dark and luminous matter to produce a powerlaw  {\\em total} mass profile should have implications for the tilt of the fundamental plane. The fundamental plane may be written in the form \\begin{equation} \\log R_e = A \\log \\sigma_0 + B \\mu + const.  \\label{eqn_fp} \\end{equation} where  $\\sigma_0$ is the central velocity dispersion, $\\mu$ is the mean surface brightness within \\reff, and A and B are proportionality constants. One way to gain insight into this relation is to interpret it in terms of the scalar virial theorem \\citep[\\eg][]{binney08a}. We can write this in the form: \\begin{eqnarray} \\log R_e & = & 2\\log \\sigma_0 - \\log\\left(\\frac{L_{1/2}}{\\pi R_e^2}\\right) - \\log \\gamma + 5.47 \\nonumber \\end{eqnarray} Here \\reff\\ is in kpc, $\\sigma_0$ is in $km\\ s^{-1}$, $L_{1/2}$ is (approximately) the light within \\reff\\ (in Solar units) and  $\\gamma = \\gamma_* \\gamma_m$  is the mass to light (M/L) ratio in Solar units, where $\\gamma_*$ is the stellar M/L ratio and $\\gamma_m$ the ratio of total mass to stellar mass within \\reff. {The value of the additive constant on the right hand side is useful only in setting the absolute normalization of $\\gamma_*$ and does not affect our conclusions. For convenience we chose to determine it from the theoretical relation of \\citet{wolf09a}\\footnote{We replaced the total velocity dispersion in their relation with the central velocity dispersion, which is only strictly valid for a flat velocity dispersion profile.}, since this implies a mean $\\gamma_*$ of $\\sim$1 for the \\citet{labarbera08a} galaxy sample (see below), which is consistent with the expectations for an old stellar population. Alternative, empirical relations have been found by \\citet[][assuming mass follows light]{cappellari06a} and \\citet[][assuming a singular isothermal sphere mass distribution]{bolton07a}, which would imply lower $\\gamma_*$, by as much as $\\sim$40\\%\\ (for the \\citeauthor{cappellari06a} relation). Nevertheless, for our work the absolute value of $\\gamma_*$ is unimportant, and so this choice does not affect our conclusions.}     In the K-band, assuming the absolute Solar magnitude is 3.41 and neglecting cosmological effects, we can write $\\log (L_{1/2}/\\pi R_e^2) \\simeq 16.00 -0.4\\mu$, and hence \\begin{equation} \\log R_e  = 2\\log \\sigma_0 +0.4 \\mu -\\log \\gamma - 10.53 \\label{eqn_virial} \\end{equation} If structural non-homology does not significantly affect this equation, the tilt of the fundamental plane arises through the dependence of $\\gamma$ on the other properties of the galaxy. Eliminating $\\sigma_0$ between Eqns~\\ref{eqn_fp} and \\ref{eqn_virial} and rearranging, we obtain: \\begin{eqnarray} \\log \\gamma & = & \\left( \\frac{2}{A}-1 \\right) \\log R_e - \\left( \\frac{2B}{A} - 0.4 \\right) \\mu + const. \\nonumber \\\\ & = & (0.33\\pm0.04) \\log R_e + (0.007\\pm0.011) \\mu + const. \\nonumber\\\\ & \\simeq &  (0.33\\pm0.04) \\log R_e + const. \\label{eqn_gamma} \\end{eqnarray} where we have used the ``$\\log \\sigma_0$ fit'' values for A and B from the recent K-band analysis of \\citet{labarbera08a}. The $\\mu$ term is consistent with zero, within the error, suggesting that the tilt of the fundamental plane can arise if $\\gamma$ depends only on \\reff, which  is exactly the prediction of the toy model, as we discuss below. Nevertheless, since the $\\mu$ term may be non-negligible, to obtain a more accurate assessment of how $\\gamma$ varies with \\reff, we computed $\\log \\gamma$ directly from the data of \\citeauthor{labarbera08a} by using Eqn~\\ref{eqn_virial}. We grouped these data into a series of narrow \\reff\\ bins\\footnote{for consistency with our X-ray analysis, we ignored those systems with $\\log$\\reff \\ltsim 0.2}, within each of which we calculated the mean $\\gamma$ and its error (by applying the central limit theorem to the scatter of the data), as shown in Fig~\\ref{fig_mass_to_light}.   We can use our toy model not only to predict the relationship between \\reff\\ and $\\alpha$, but also the dark matter fraction within \\reff\\ (and hence $\\gamma_m$), as is immediately clear from inspection of Fig~\\ref{fig_mass_decomposition}. Since the dark matter {\\em fraction} does not depend on the overall normalization of the mass profile, we do {\\em not} need to fold in any other relations (such as the Kormendy relation). Furthermore, since we find that \\reff\\ varies monotonically with $\\alpha$ (Fig~\\ref{fig_mass_slope}), it is possible to recast this as a relation between $\\gamma_m$ and \\reff. We show the resulting relation (roughly $\\log \\gamma_m \\simeq -0.03\\log R_e +0.32 \\left( \\log R_e \\right)^2$) in Fig~\\ref{fig_mass_to_light}, along with an estimate of the uncertainty that arises from from varying the Sersic index by $\\pm$1, or adopting the \\citet{navarro04a} model for the dark matter profile. It is immediately clear that there is good agreement between the model and the \\citet{labarbera08a} data, {implying the mean $\\gamma_* \\simeq 1$, as noted above}. Since the stellar populations of early-type galaxies (and hence $\\gamma_*$) do not strongly vary with \\reff\\ \\citep[at least for \\reff \\gtsim 1~kpc:][]{graves09b}, this {suggests that assuming the ubiquity of approximately powerlaw density profiles (supported by Fig~\\ref{fig_mass_profiles}) is sufficient (albeit not necessary) to explain the tilt of the fundamental plane. While it is possible that not all giant early-type galaxies have density profiles which so closely resemble a powerlaw (although we are not aware of any convincing such counter-examples found by X-ray methods), to lie on the fundamental plane, Fig~\\ref{fig_mass_to_light} indicates that their $\\gamma_m$ must still vary with \\reff\\ in a very similar way to the predictions of the toy model. Nevertheless, since it is actually the flattening of the density profile with increasing \\reff\\ that is responsible for the $\\gamma_m$-\\reff\\ relation predicted by the toy model, it is plausible that similar flattening of the mass distributions in a system without a powerlaw-like mass distribution might produce approximately the same trend.}  One intriguing prediction of our toy model is that the relationship between $\\log \\gamma_m$ and $\\log R_e$ has some curvature, which would also imply curvature in the fundamental plane. There is, in fact, some evidence that brightest cluster galaxies (BCGs) may not lie on the same fundamental plane as smaller systems \\citep[\\eg][]{oegerle91a,vonderlinden07a,bernardi07a}, but it is not clear that the observed behaviour could be explained in terms of the predicted curvature. Indeed, there is a hint that the observed $\\gamma$ {\\em versus} \\reff\\ relation begins to deviate from our simple predictions at \\reff\\gtsim 20~kpc (Fig~\\ref{fig_mass_to_light}), which may arise due to deviations from structural homology among BCGs, or it may represent the limitations of our simple model.  To make a more quantitative comparison between the traditional fundamental plane and our toy model, we  adopted a linear approximation for the toy model relationship between $\\gamma_m$ and \\reff\\ ($\\log \\gamma_m \\simeq -0.11 +0.37 \\log R_e$; this is accurate to about $\\sim$0.03~dex over the range $0.2<\\log R_e<1.0$, as shown in Fig~\\ref{fig_mass_to_light}, assuming $\\gamma_*=1$). Substituting this into Eqn~\\ref{eqn_virial} (assuming $\\gamma_*$ is constant) and rearranging into the form of  of Eqn~\\ref{eqn_fp}, we found  A=1.46 and B=0.29. These are in good agreement with recent K-band measurements of the fundamental plane \\citep[A$\\sim$1.4--1.5 and B$\\sim$0.30--0.33:][]{labarbera08a,jun08a,mobasher99a,pahre95a}. Repeating the analysis, but allowing the Sersic index to vary by $\\pm$1, or using the alternative dark matter model of \\citet{navarro04a}, we obtained A in the range $\\sim$1.2--1.6 and B$\\simeq$0.23--0.31.  Thus, assuming only a powerlaw total mass distribution, we can clearly reproduce the shape of the K-band fundamental plane, without any fine-tuning of the model parameters.  We therefore speculate that the establishment of a powerlaw total mass distribution is a fundamental feature of galaxy formation and the principal factor in determining the tilt of the fundamental plane. Since the relative importance of gas-rich and ``dry'' merging for the assembly of present-day giant ellipticals is still under debate \\citep[\\eg][]{hopkins08a,naab09a,nipoti09a}, the requirement of producing a nearly powerlaw total mass distribution in the core of low-redshift systems will therefore be a key test of these models."
    },
    "0911/0911.2357_arXiv.txt": {
        "abstract": "It is known since the seminal study of \\citet{Las1989} that the inner planetary system is chaotic with respect to its orbits and even escapes are not impossible, although in time scales of billions of years. The aim of this investigation is to locate the orbits of Venus and Earth in phase space, respectively to see how close their orbits are to chaotic motion which would lead to unstable orbits for the inner planets on much shorter time scales. Therefore we did numerical experiments in different dynamical models with different initial conditions -- on one hand the couple Venus-Earth was set close to different mean motion resonances (MMR), and on the other hand Venus' orbital eccentricity (or inclination) was set to values as large as $e=0.36$ ($i=40^{\\circ}$). The couple Venus-Earth is almost exactly in the 13:8 mean motion resonance. The stronger acting 8:5 MMR inside, and the 5:3 MMR outside the 13:8 resonance are within a small shift in the Earth's semimajor axis (only 1.5 percent). Especially Mercury is strongly affected by relatively small changes in eccentricity and/or inclination of Venus in these resonances. Even escapes for the innermost planet are possible which may happen quite rapidly. ",
        "introduction": "Mean Motion Resonances are essential in Solar System Dynamics not only for the planets but also for the motion of the asteroids. The appearance of chaotic motion in the asteroid belt detected by \\citet{Wis1981} -- the 3:1 MMR of an asteroid with Jupiter -- was the first one discovered after the seminal discovery by \\citet{Hen1964} for galactic dynamics. But the structure of the asteroid belt is also created by the secular resonances (SR), where the motion of the perihelia and the nodes can be in resonance for different planets in combination with the asteroid's perihelion respectively node. For the planets it was found by Laskar in different papers (\\citet{Las1988,Las1990,Las1996}) that especially the Inner Solar System (ISS) is in a chaotic state. The appearance of SR may even shift the orbit of Mars out from a stable orbit into an unstable one with high probability when one does not take into account relativity. The large planets seem to be in quite a 'safe' region, although many MMR are acting: e.g. the 5:2 MMR between Jupiter and Saturn. Recent investigations by \\citet{Mic2001,Gal2006} show how close their motion is to a chaotic state with severe implications for the other planets.  Because the orbits of Venus and Earth are quite close to the 13:8 MMR\\footnote{the periods of Venus and Earth are almost in the ratio 13:8} and they are strongly coupled with respect to the inclinations and eccentricities (Fig. \\ref{Fig1}) -- comparable to Jupiter and Saturn -- we investigated the phase space close to this high order resonance. The order of a resonance is defined as the difference $q$ when we characterize a MMR by $p:(p+q)$. The value $q$ is connected to the exponent in front of the fourier term in the expansion of the disturbing function; the higher the exponent in a small quantity in front of the Fourier term is, the less strong is the influence on the dynamics, unless there is a resonance acting (e.g. Murray and Dermott). Thus for the 13:8 MMR we have to deal with the order 5. It is interesting to note that a relatively small shift of the Earth to a smaller semimajor axis would bring the couple Venus--Earth into the 8:5 MMR (order 3); a shift to a larger semimajor axis would bring both planets into the 5:3 MMR (order 2) (Fig. \\ref{Fig2}). In this plot the results of the computations in the simple three body problem Sun-Venus-Earth are presented to show the location of the resonances and also their shape.       \\begin{figure} \\begin{center} \\includegraphics[width=3.0in,angle=270]{fig1.ps} \\caption{The coupling in the eccentricities of Venus and Earth; from -100 to 95 Myrs (upper graph), from -2.5 to 2.5 Myrs (middle graph) and from 95 to 100 Myrs (lower graph), after \\citet{Dvo2003}.} \\label{Fig1} \\end{center} \\end{figure}    \\begin{figure} \\begin{center} \\includegraphics[width=4.5in,angle=0]{fig2.eps} \\caption{The MMRs between Venus and the Earth; along the x-axis the initial semimajor axis of the Earth is changed, along the y-axis the initial values for the eccentricity of Venus are varied. The greytone indicates the maximum eccentricity, the vertical white lines show the position of the MMR up to high orders.} \\label{Fig2} \\end{center} \\end{figure}      Besides extensive numerical integration of the motions of the planets in different dynamical models, we also used chaos indicators and analysed the frequencies involved. Furthermore simplified mapping models were constructed to understand the structure of phase space where Venus and Earth are embedded.  The content of the paper is follows: after a careful test of the dynamical model which we used (section 2) we roughly describe the three MMR mentioned above and show the structure determined with the aid of a chaos indicator (section 3). We then report on the results of different numerical experiments where we show the dependency of the stability of the inner planetary system on the eccentricities and the inclinations of Venus (section 4). In the conclusions we discuss the possible consequences for the stability of the inner planetary system. In an appendix first results of an applied mapping for motions in MMRs are shown. ",
        "conclusions": "In this investigation we explored the environment of the 13:8 MMR, which is very close to the actual position of Venus and Earth. After a careful test of a simplified dynamical model of our planetary system appropriate for our task we discussed the main structures of the three MMR in the vicinity of the two planets using the results of a study of the {\\sc FLI}.  In another step, using extensive numerical integrations we changed the ratio of the semimajor axes Venus--Earth\\footnote{Venus' semimajor axis was fixed.}; the orbit of the Earth was given different values of $a$ to cover the proximate neighborhood of the MMR where the twin-planets are in. Then, in two different runs the eccentricity of Venus was set to eccentricities up to $e = 0.36$ and in the dynamical model, consisting of the inner Solar System with Jupiter and Saturn, the equations of motion were integrated for $10^7$ years. The same was done in a second run where we changed the inclination of Venus up to $i=40^{\\circ}$. These two experiments were undertaken for each of the three MMR 8:5, 13:8 and 5:3 and the results analysed with respect to the maximum eccentricties of the inner planets. It turned out that only Mercury is the planet which would suffer from such perturbations, that it would achieve very large values of its eccentricity and thus would be thrown out from its current orbit. According to our results Mars is not as strongly affected by a larger inclination or eccentricity of Venus as Mercury; this can be understood by the presence of Jupiter and Saturn keeping this planet almost in its original plane of motion. Venus and also the Earth are only affected for larger initial inclinations of Venus, but the planets stay more or less in their orbits although sometimes achieving eccentricities as large as $e=0.3$.  In an attempt to understand the dynamical structure of the resonances a mapping approach (appendix) was constructed in the simple elliptic restricted three body problem Sun--Venus and massless Earth for the 3:1 MMR. The same approach was then used to construct the surface of section of the 5:3 MMR. Until now we failed to get a similar graph for the other two resonances which ask for a development of the Hamiltionian up to order 3 and order 5, but we work on it.   As a final statement we may say that according to our results although the couple of Venus and Earth is close to a MMR we do not expect big changes of the orbits of the inner planets unless Venus will come into an orbit of an inclination of about 7 degrees or will achieve an eccentricity of 0.2. Then, in fact, Mercury would be on an unstable orbit which could finally throw it either into the Sun or far out into the main belt of asteroids!"
    },
    "0911/0911.2950_arXiv.txt": {
        "abstract": "We developed a method to separate a long-term trend from observed temporal variations of polarization in blazars using a Bayesian approach.  The temporal variation of the polarization vector is apparently erratic in most blazars, while several objects occasionally exhibited systematic variations, for example, an increase of the polarization degree associated with a flare of the total flux.  We assume that the observed polarization vector is a superposition of distinct two components, a long-term trend and a short-term variation component responsible for short flares.  Our Bayesian model estimates the long-term trend which satisfies the condition that the total flux correlates with the polarized flux of the short-term component.  We demonstrate that assumed long-term polarization components are successfully separated by the Bayesian model for artificial data.  We applied this method to photopolarimetric data of OJ~287, S5~0716$+$714, and S2~0109$+$224. Simple and systematic long-term trends were obtained in OJ~287 and S2~0109$+$224, while no such a trend was identified in S5~0716$+$714. We propose that the apparently erratic variations of polarization in OJ~287 and S2~0109$+$224 are due to the presence of the long-term polarization component. The behavior of polarization in S5~0716$+$714 during our observation period implies the presence of a number of polarization components having a quite short time-scale of variations. ",
        "introduction": "Blazars are believed to be active galactic nuclei (AGN) with relativistic jets pointing toward us (e.g. \\cite{bla78blazar}). Doppler-boosted non-thermal emission from the jet dominates from radio to $\\gamma$-rays in blazars.  The radio---optical emission is dominated by synchrotron emission, and hence, highly polarized. Polarimetric observations have been extensively performed since they provide a clue for the structure of magnetic field in the jet (e.g. \\cite{ang80blazar_pol}; \\cite{mea90blazar_pol}).  Another important feature of blazars is a violent variability in a wide range of time-scales from a few minutes to years (e.g. \\cite{huf92variation}).  Since the polarization degree (PD) and angle (PA) are also variable with time, dense and long-term polarimetric observations are essential to understand the variation mechanism of blazars and the magnetic field structure in the jet.  \\citet{moo82bllac} conducted an intensive observation campaign for BL~Lac with multi-longitude optical observatories over a week period, and found that the polarization changes trace out a random walk in the Stokes $QU$ plane.  In general, erratic variations have been observed in polarization of most blazars (e.g. \\cite{ang80blazar_pol}).  On the other hand, systematic variations in PD and PA were also occasionally detected in blazars, for example, flares associated with the increase of PD (\\cite{smi863c354}; \\cite{tos98mrk421}; \\cite{efi98on231}) and the color variation correlating with the variation in PD (\\cite{cel07ao0235}).  \\citet{mar08bllac} have recently reported a smooth rotation of the polarization vector in BL~Lac, which indicates a propagation of the emitting region through a helical magnetic field in the jet.  Thus, both erratic and systematic variations of polarization have been observed in blazars.  The erratic variation of the polarization vector can be explained by a scenario that a number of independent sources with randomly oriented and strong polarizations blink (\\cite{moo82bllac}; \\cite{imp893c273}; \\cite{jon85rotation}).  The systematic variations are, however, difficult to be explained only by the random variation. Alternatively, a number of systematic variations can be apparently overlooked in PD and PA if short-term polarization variations are superimposed on a long-term polarization trend.  A schematic example is shown in figure~\\ref{fig:example}.  The upper and lower panels show the light curve of the total and polarized fluxes and the temporal variation of the polarization vector in the Stokes $QU$ plane, respectively.  In this example, the polarization vector is assumed to be a superposition of two components; one shows a short-term variation associated with a flare of the total flux, and the other is a long-term trend.  In the case that there is no long-term trend, PF correlates with the total flux, independent of the PA of the short-term flare.  In the case that the long-term trend has a significant PF as shown in the lower panel of figure~\\ref{fig:example}, however, the PF can show no positive correlation with the light curve.  If the PA of the short-term flare is random and observations are sparse, the variation of polarization can be apparently erratic.  \\begin{figure} \\begin{center} \\FigureFile(80mm,80mm){quexam.eps} \\end{center} \\caption{Example of the short-term flare in the artificial data for our Bayesian model.  The upper panel shows the light curve of the total and polarized fluxes indicated by the open and filled circles, respectively.  The lower panel shows the temporal variation of polarization in the Stokes $QU$ plane.  A given error for the points is shown in the  right-lower corner of the panel.}\\label{fig:example} \\end{figure}  The presence of the long-term component has been suggested by several observations.  The short-term variability of blazars is characterized by a combination of shots (\\cite{huf92variation}).  A typical time-scale of the shot was estimated to be $\\sim 1$~d (\\cite{tak00mrk421}; \\cite{kat01variation}).  In addition to these short-term variations, it is well known that blazars show variations with a time-scale of months--years.  It is possible that the long- and short-term variation components have different polarization features. \\citet{mar02vlba} reported VLBI radio observations of blazars, which show two or more components in polarization maps.  If the polarization components in the radio range can be observed also in the optical range, an observed optical polarization vector should be a composition of those distinct radio polarization components since the angular resolution of optical observations is much lower than that of VLBI radio images.  \\citet{hag02bllac} reported a long polarimetric monitoring of BL~Lac in 1969--1991, and found the existence of the preferred direction of the polarization vector.  They propose a model that the polarization of BL~Lac has two kinds of sources, a stationary component and a large number of randomly polarized components (also see, \\cite{bli09bllac}).  This stationary component could be considered as a long-term trend if it varies gradually.  The two-component model shown in figure~\\ref{fig:example} is one of the simplest one among models with multiple polarization components. It deserves further investigation if such a simple model can extract a systematic trend from erratic variations of polarization in blazars. Here, we present an exploratory Bayesian approach to separate a long-term trend from observed temporal variations of polarization in blazars.  A description of the model and method is shown in the next section.  We apply the model to observed photopolarimetric data of blazars and discuss the implications of the results in section~3.  In section~4, we discuss the validity of the model and results.  Finally, we summarize our findings in section~5. ",
        "conclusions": "\\subsection{Estimation of the Hyperparameter, $w$, by Maximizing Marginal Likelihood}  \\begin{table} \\caption{Marginal likelihood, $M(w)$, calculated with different $w/\\sigma$.}\\label{tab:mar} \\begin{center} \\begin{tabular}{cccc} \\hline $w/\\sigma$ & OJ~287 & S5~0716$+$714 & S2~0109$+$224\\\\ & \\multicolumn{3}{c}{$\\times 10^3$} \\\\ \\hline 0.10 & 2.06 & 2.86 & 0.97\\\\ 0.50 & 2.96 & 3.96 & 1.92\\\\ 1.00 & 3.20 & 4.44 & 2.16\\\\ 2.00 & 3.21 & 4.61 & 2.21\\\\ 3.00 & 2.99 & 4.53 & 2.22\\\\ 5.00 & 2.60 & 4.15 & 2.17\\\\ \\hline \\end{tabular} \\end{center} \\end{table}  As mentioned in subsection~2.3, we made the criterion to determine an appropriate value of the hyperparameter, $w$, in our Bayesian model. In general, a hyperparameter of an empirical Bayesian model can be estimated by maximizing a marginal likelihood.  The marginal likelihood, $M$, is equivalent with the constant, $C$, in equation~(2), defined as: \\begin{eqnarray} M(w) = \\frac{L(\\mathbf{y}|\\mathbf{x})\\pi(\\mathbf{x}|w)} {P(\\mathbf{y}|\\mathbf{x})}. \\end{eqnarray} We calculated $M(w)$ based on the method described in \\citet{chi01ml}. The model parameter, $\\mathbf{y}$, can be arbitrary since $M(w)$ is independent of $\\mathbf{y}$.  The numerator of equation~(14) was calculated with the estimated best parameters of $\\mathbf{y}$. The denominator can be estimated by the Monte Carlo integration drawing a sample from $P(\\mathbf{y}|\\mathbf{x})$ which were obtained by our Bayesian analysis.  The results are summarized in table~\\ref{tab:mar}.  $M(w)$ takes the maximum with $w/\\sigma =2.00$--$3.00$ for all cases of the three objects.  A long-term component estimated with such a large $w/\\sigma$ definitely draws a quite complicated path in the $QU$ plane, as can be seen from figure~\\ref{fig:samplequ}.  As a result, it is not acceptable in our criterion about the distance ratio, $R$.  The ratio of two $M(w)$ is called the ``Bayes factor'' (BF), which used to evaluate the models.  The BF are calculated to be $\\sim 1.6$, $\\sim 1.6$, and $\\sim 2.3$ with $M(0.10)$ and the maximum of $M$ for OJ~287, S5~0716$+$714, and S2~0109$+$224, respectively.  These BF are too small to conclude that the model with $w/\\sigma =2.00$--$3.00$ is significantly better than that with $w/\\sigma =0.10$ (\\cite{kas95bf}).  Thus, the maximization of $M(w)$ is not suitable to determine $w$ in our model.  This is mainly because the prior distribution is not the real distribution for $Q_{\\rm L}$ and $U_{\\rm L}$.  Another form of the prior distribution may be more suitable for the case of polarization variations in blazars.  The practical criterion defined in subsection~2.3 is enough for our exploratory model in this paper.  \\subsection{Application in the ($Q/I,U/I$) Plane}  \\begin{figure*} \\begin{center} \\FigureFile(160mm,160mm){samplequd.eps} \\end{center} \\caption{Same as figure~\\ref{fig:samplequ}, but for $Q/I$ and $U/I$.  The errors for each points of the estimated long-term trends are typically less than 1.0, 1.3, and 0.9~\\% for OJ~287, S5~0716$+$714, and S2~0109$+$224, respectively.}\\label{fig:samplequd} \\end{figure*}  Our Bayesian model can find a long-term trend in the $QU$ plane.  It can be also applied in the $(q,u)=(Q/I,U/I)$ plane, while strong limitations are imposed for the application.  In our two-component model, the observed $(q,u)$ is described as: \\begin{eqnarray} q = \\frac{Q_{\\rm L} + Q_{\\rm S}}{I_{\\rm total}}, u = \\frac{U_{\\rm L} + U_{\\rm S}}{I_{\\rm total}}\\\\ I_{\\rm total} = I_{\\rm L} + I_{\\rm S}. \\end{eqnarray} The two polarization components, namely, $(q_{\\rm L},u_{\\rm L})=(Q_{\\rm L}/I_{\\rm L},U_{\\rm L}/I_{\\rm L})$ and $(q_{\\rm S},u_{\\rm S})=(Q_{\\rm S}/I_{\\rm S},U_{\\rm S}/I_{\\rm S})$, can be separated in the $qu$ plane when $q_{\\rm L}$ and $u_{\\rm L}$ are much smaller than $q_{\\rm S}$ and $u_{\\rm S}$ (e.g. \\cite{mes97polari}).  If the PD of the long-term trend is not negligible, the polarization components cannot be separated without the estimation of $I_{\\rm L}$ and $I_{\\rm S}$.  In the case that the PD of the long-term trend is large, the application of the Bayesian model in the $qu$ plane means the estimate of $(Q_{\\rm L}/I_{\\rm total}, U_{\\rm L}/I_{\\rm total})$ that lead to a positive correlation between the light curve and the modified PD calculated from $(Q_{\\rm S}/I_{\\rm total}, U_{\\rm S}/I_{\\rm total})$. Since $I_{\\rm L}$ is assumed to be time-independent in our model, the temporal behavior of the estimated $(Q_{\\rm S}/I_{\\rm total}, U_{\\rm S}/I_{\\rm total})$ is probably analogous to that of $(Q_{\\rm S}/I_{\\rm S}, U_{\\rm S}/I_{\\rm S})$.  In this sense, the Bayesian analysis in the $qu$ plane is not completely nonsense, and can still provide meaningful results.  A result in the $qu$ plane would be noteworthy, particularly if the long-term trend in the $qu$ plane draws an analogous path to that in the $QU$ plane.  Figure~\\ref{fig:samplequd} is the same as figure~\\ref{fig:samplequ}, but for the $qu$ plane. We found that the long-term trends estimated with $w/\\sigma =0.20$ are quite analogous to those in figure~\\ref{fig:samplequ} in the case of S2~0109$+$224.  The long-term trend in OJ~287 also exhibits a common feature with that in the $QU$ plane in terms of the oscillation of the polarization vector within a narrow range of PA.  The similarity between the long-term trends in the $QU$ and $qu$ planes implies that the short-term flares were associated with the increases of PD in OJ~287 and S2~0109$+$224. The short-term flares may originate from a region where the direction of the magnetic field is more aligned than the whole emitting region.  \\subsection{Future Studies}  We showed that the polarization vector in OJ~287 and S2~0109$+$224 can be separated into two components, the long-term trend and the short-term variation component.  The result, however, provides no clear evidence for the existence of the two components in these objects.  Our Bayesian model just gives us one of possible views for temporal variations of polarization in blazars. This view deserves to be discussed since simple and systematic variations can be extracted from apparently erratic variations. On the other hand, further investigations are needed to confirm the two-component view.  One of the most direct ways to obtain evidence for the two-component scenario would be VLBI radio observations.  \\citet{mar02vlba} reported the presence of multiple polarization components even within radio cores in blazars.  VSOP-2 is a space VLBI project, which will resolve sub-pc regions around the central blackhole of AGN (\\cite{VSOP2}).  VSOP-2 will reveal the detailed structure of the polarization components in the radio core.  A part of the radio polarization components might exhibit temporal variations synchronized with the optical polarization components resolved by our Bayesian model.  If this is the case, our Bayesian model, for the first time, provides a tool to investigate the temporal evolution of each polarization component observed both in the optical and radio regimes.  Our model assumes that the PF of the short-term component correlates with the light curve of the total flux (condition ``ii)'' in section~2.1).  In other words, we assumes the total flux exhibits a flare whenever a short-term polarization flare occurs.  This assumption gives the condition that the application of our Bayesian model is validated.  Our model definitely yields a false result if it applies to, for example, short-term polarization flares without significant variations of the total flux.  If such variations are typical in a blazar, our model is not validated for it.  Our assumption should, hence, be evaluated case by case for objects.  In our picture, a short-term flare with a larger amplitude is expected to show a better correlation with polarization because the specific polarization-feature of a smaller flare would be readily diluted by another small flares.  Hence, the polarization behavior of short and large flares would provide a key for the validation of our assumption.  \\citet{mar08bllac} reported a rotation of the optical polarization vector in BL~Lac and interpreted it as a shift of the emitting region through a helical magnetic field in the jet.  Such a rotation event can have been missed if the observed polarization has a long-term trend as considered in our model.  Our Bayesian model would provide a tool to search the rotation events of the polarization vector both in the observed, long- and short-term components.  No systematic rotation event can be seen in the long- and short-term components in the three objects, as shown in subsection~3.2."
    },
    "0911/0911.0955_arXiv.txt": {
        "abstract": "We present new integrated light  spectroscopy of globular clusters in NGC 5128, a nearby giant elliptical galaxy  less than 4 Mpc away, in order to  measure  radial velocities  and  derive  ages, metallicities,  and alpha-element  abundance  ratios.   Using  the  Gemini  South  8-meter telescope with  the instrument GMOS,  we obtained spectroscopy  in the range  of $\\sim 3400-5700$  $\\rm{\\AA}$ for  72 globular  clusters with signal-to-noise  greater  than 30  $\\rm{\\AA}^{-1}$  and  we have  also discovered 35  new globular clusters  within NGC 5128 from  our radial velocity measurements.  We measured and compared the Lick indices from H$\\delta_A$ through  Fe5406 with the single  stellar population models of \\cite{tmb03}  and \\cite{tmk04} in order to  derive age, metallicity and [$\\alpha$/Fe] values.   We also measure Lick indices  for 41 Milky Way globular  clusters from \\cite{puzia02}  and \\cite{schiavon05} with the  same methodology for  direct comparison.   Our results  show that 68$\\%$ of  the NGC  5128 GCs have  old ages  ($> 8$ Gyr),  14$\\%$ have intermediate ages ($5-8$ Gyr), and 18$\\%$ have young ages ($< 5$ Gyr). However, when we look at the metallicity of the globular clusters as a function  of age,  we  find 92$\\%$  of  metal-poor GCs  and 56$\\%$  of metal-rich GCs  in NGC 5128 have  ages $> 8$ Gyr,  indicating that the majority  of  both  metallicity  subpopulations of  globular  clusters formed early,  with a significant  population of young  and metal-rich globular  clusters forming later. Our   metallicity  distribution   function  generated   directly  from spectroscopic  Lick  indices  is  clearly  bimodal, as  is  the  color distribution of the same set  of globular clusters. Thus the  metallicity bimodality  is real and  not an artifact  of the color to metallicity conversion. However, the metallicity distribution function obtained  from comparison with the  single stellar population models is  consistent with a unimodal, bimodal, or multimodal shape.  The [$\\alpha$/Fe]   values   are  supersolar   with   a   mean  value   of $0.14\\pm0.04$, indicating a fast  formation timescale.  However, the GCs in NGC 5128 are not as [$\\alpha$/Fe] enhanced  as the Milky Way GCs also examined in this  study.  Our  measured indices  also indicate  that  the globular clusters in NGC 5128 may  have a slight overabundance in nitrogen and a wider range  of calcium strength compared to  the Milky Way globular clusters.  Our results support a rapid, early formation of the globular cluster system in NGC 5128, with subsequent major accretion and/or GC and star forming events in more recent times. ",
        "introduction": "\\label{sec:intro}  The  star formation history  of early-type  galaxies is  quite complex with evidence supporting episodes of star formation at both high and low redshifts.  At high redshifts (z $> 2$),  the formation of the bulk of the  galactic stellar component is supported  by, for example, the fundamental  plane relation \\citep{djorgovski87,dressler87,treu99} in          the          monolithic         collapse          scenario \\citep{eggen62,tinsley72,larson74,larson75,silk77,arimoto87}.       The hierarchical                      merging                     scenario \\citep{toomre77,white78,peebles80,whitefrenk91,kauffmann93,baugh98,cole00,somerville01} on the  other hand, predicts a  range of formation  timescales for the stellar component down to much  lower redshift values.  In the extreme case,  gas-rich  merging galaxies  are  seen  today  where young  star clusters,  for  example in  NGC  4038/4039 \\citep{whitmore95},  and/or young  and  intermediate  aged   globular  clusters  (GCs)  have  been detected, indicating  that GCs, as well  as a fraction  of the stellar component of the galaxy can form at z $\\leq 1$.  The  key to  understanding  the  formation epochs  of  these types  of galaxies  is through  their star  formation history.   Most early-type galaxies are  too distant to resolve  stars from which  we can measure their ages  directly. Unresolved light studies are  difficult as well, as they  are hindered  by the age-metallicity  degeneracy \\citep[][and references  therein]{mould80} as  well as  internal extinction  in the galaxies.  GCs, on  the other hand,  are essentially coeval structures  that form with nearly single  age and single metallicity.  They  have been used, therefore,  as powerful  tools  to trace  the  complex star  formation history of  early-type galaxies.   Their internal properties,  such as ages  and metallicities  are used  as representations  of  the overall stellar component of the galaxy.  We  also benefit from the use of GCs as   they   do   not   suffer   from   problems   involving   internal reddening.   GCs can be used as good tracers of the time period  of early-type galaxy formation because we can determine their ages, metallicities, and alpha-enhancements through their spectra.  By examining the ratio of GCs that are old to GCs with younger ages, as well as comparing the age difference between metal-rich and metal-poor GCs, we can determine which formation scenario is dominant.  Previous studies using photometry  to determine ages and metallicities of  GCs are  plagued with  age-metallicity degeneracies, especially in the optical regime \\citep{wortheyg94}. One way  to avoid the  age-metallicity degeneracy using  photometry is through the use of (near-)infrared  data, which is mainly sensitive to the metallicity of the red  giant branch, combined with optical data, sensitive to both age  and metallicity.  This technique has been shown to be effective at  distinguishing between  GC populations that   appear  to   be  single-aged   and  old  such as NGC  7192 \\citep{hempel03},   NGC    1399   \\citep{kundu05},   and    NGC   3115 \\citep{puzia02}, versus those with a wide spread of ages  such as NGC    4365 \\citep{puzia02,kundu05} and  NGC 5846 \\citep{hempel03, hempel07}. However, a sample of GCs in NGC 4365 with intermediate ages were found to have old ages using spectroscopic techniques \\citep{brodie05}. These color indices are also not precise enough to be useful for individual GCs or to derive genuine age distribution functions.  Spectroscopy allows us  to determine the internal properties of  individual GCs  as  long  as we  have  high signal-to-noise  data. The measurement  of  the  strength  of  absorption line  features  in  the spectra enables us to obtain ages, metallicities, and [$\\alpha$/Fe] through comparison with SSP  models. This technique enables the  collection of a significant fraction of both metal-rich  and metal-poor GC data in one environment that can show the formation epoch(s) and chemical composition at the time of formation of the GCs within one early-type galaxy.  The proximity of NGC 5128 makes this early-type galaxy extremely attractive to conduct a large, homogeneous age and  metallicity study \\citep[$3.8\\pm0.1$  Mpc][]{harris09}. Features within this galaxy (the  dust lane and faint shells \\citep{malin78}, visible  star formation  from the interaction  of the radio jet  with the shells of  HI \\citep{graham98}, and  a young tidal stream in the  halo \\citep{peng02}) are recent phenomena which will have little bearing on the much older GC population.  Previous studies have been able to obtain spectroscopy for GCs in many different galaxies  to determine  global properties of  the metal-poor versus             metal-rich clusters       \\citep[][among others]{puzia05,strader05,sharina06}.  Within  NGC 5128, we  should be able to determine whether there  is a significant fraction of young or intermediate-aged GCs,  or if the majority are  old, similar to their Milky Way  GC counterparts.  Our primary  goal is to  use a large and high signal-to-noise (S/N) sample of  GC spectra within NGC 5128  to obtain the  relative age difference between  the metal-rich  and  the metal-poor  clusters.  Currently, we have the advantage of knowing over 415 GCs in NGC 5128 which have been confirmed with  either radial  velocity measurement and/or  HST images cataloged in  \\cite{woodley07}.  This  catalog has provided  the database from which we selected our target objects.  We present our results in the sections below:  section \\ref{sec:data} describes our data set and its reduction; section \\ref{sec:rv} describes  the radial velocity measurements  obtained  and   the  confirmation  of  newly  discovered GCs; section   \\ref{sec:results}   describes   the calibration of  our data and our resulting ages, metallicities,  and [$\\alpha$/Fe]; section \\ref{sec:dis} discusses our findings; and we conclude in Section \\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  Using the Gemini  South 8-meter telescope with GMOS, we obtained spectra of GCs in NGC 5128 to obtain radial velocities, ages, metallicities, and alpha-to-iron abundance ratios.  The main results from this study are:  \\begin{itemize} \\item We discovered 35 new  GCs from radial  velocity measurements  as  well as remeasuring radial velocities for  101 previously confirmed GCs.  \\item Our quantitative age estimates  suggest $68\\%$ of our GCs have  old ages ($>8$ Gyr), $14\\%$ have  intermediate ages ($5-8$  Gyr), and $18\\%$  have young ages ($< 5$ Gyr), suggesting several, or at least extended, formation epochs.  \\item Separating the GCs  into metal-rich and metal-poor  subpopulations by  the model-derived  [Z/H]  quantity, we find that 92$\\%$ of  metal-poor ([Z/H]$<-1$) and 56$\\%$ of metal-rich ([Z/H]$\\geq -1$)  GCs   in  NGC  5128   have  ages  $>  8$   Gyr.  \\item We  also find both  the Washington  color and  spectroscopic synthetic metallicity  index  [MgFe]$^\\prime$   show  a bimodal distribution for  the GCs, indicating  that there are  two  distinct  metallicity   subpopulations.  \\item The [$\\alpha$/Fe] for  the GCs  are super solar, but not as large as in the Milky Way.  We suggest that the GCs in NGC 5128 formed on slower timescales than GCs in the Milky Way and other massive elliptical galaxies \\citep{puzia06}. \\end{itemize}  Our results support a scenario with a rapid, early formation of the bulk of the GCS and significant subsequent major accretion and/or GC- and star-forming events in more recent times. We note that  our relatively high fraction of  metal-rich and  young GCs is  likely biased  because we probed the inner regions of  NGC  5128 and observed mainly  the brightest GCs.  We  consider our results, therefore, to  be a fractional upper  limit on the number  of metal-rich and  young GCs in NGC 5128."
    },
    "0911/0911.4810_arXiv.txt": {
        "abstract": "We consider an extension of the holographic Ricci dark energy model by introducing an interaction between dark energy and matter. In this model, the dark energy density is given by $\\rho_{\\Lambda}=-\\frac{1}{2}\\alpha M_{p}^{2}{\\cal R}$, where ${\\cal R}$ is the Ricci scalar curvature, $M_{p}$ is the reduced Planck mass, and $\\alpha$ is a dimensionless parameter. The interaction rate is given by $Q=\\gamma H \\rho_{\\Lambda}$, where $H$ is the Hubble expansion rate, and $\\gamma$ is a dimensionless parameter. We investigate current observational constraints on this model by applying the type Ia supernovae, the baryon acoustic oscillation and the CMB anisotropy data. It is shown that a nonvanishing interaction rate is favored by the observations. The best fit values are $\\alpha=0.45\\pm 0.03$ and $\\gamma=0.15\\pm 0.03$ for the present dark energy density parameter $\\Omega_{\\Lambda 0}=0.73\\pm0.03$. ",
        "introduction": "Recent observations of type Ia supernovae (SNIa) have revealed that the present Universe is undergoing accelerated expansion~\\cite{riess1}. This indicates that the energy density of the Universe at present epoch is dominated by dark energy with equation of state $w_{\\Lambda} = \\frac{p_{\\Lambda}}{\\rho_{\\Lambda}}<-\\frac{1}{3}$, where $\\rho_{\\Lambda}$ is the energy density of dark energy, and $p_{\\Lambda}$ is its pressure. Combining with observations of CMB anisotropy~\\cite{WMAP} and the baryon acoustic oscillation (BAO)~\\cite{SDSS}, the density parameter is determined as $\\Omega_{\\Lambda 0} = \\rho_{\\Lambda 0}/\\rho_{c 0} \\sim 0.7$, where $\\rho_{c 0} = 8.0992h^{2}\\times 10^{-47}{\\rm GeV}^{4}$ is the critical density today and $H_{0}=100h~{\\rm km/s/ Mpc}$ is the present Hubble parameter.  The simplest way to explain dark energy is to introduce a cosmological constant~\\cite{einstein}. However, this model suffers from two problems~\\cite{weinberg}. The first one is the fine-tuning problem: the observed cosmological constant is extremely small compared to the fundamental Planck scale $\\rho_{\\Lambda}$ $\\sim$ $10^{-120} M_{p}^{4}$, requiring an incredible fine-tuning. The second one is the cosmic coincidence problem: why the cosmological constant and matter have comparable energy density today even though their time evolution is so different. A variety of models have been proposed to solve these problems, such as quintessence~\\cite{quin}, phantom~\\cite{phantom}, quintom~\\cite{quintom} and holographic dark energy~\\cite{li, gao, cosmicconstraintsRDE, RDE}.  In particular, the holographic dark energy (HDE) models have been discussed extensively in recent years. These models are motivated by the holographic principle of quantum gravity~\\cite{quantumgravity}. From the condition that a system with size $L$ would not form a black hole, it is required that the total vacuum energy should not exceed the mass of the black hole of the same size. Therefore, the dark energy density $\\rho_{\\Lambda}$ must satisfy $L^{3} \\rho_{\\Lambda}\\lesssim L M_{p}^{2}$~\\cite{cohen}, where $M_{p} = 1/\\sqrt{8\\pi G}$ is the reduced Planck mass. Saturating this inequality, $\\rho_{\\Lambda}$ is defined by~\\cite{li} \\begin{eqnarray} \\rho_{\\Lambda}=3c^{2}M_{p}^{2}L^{-2}, \\end{eqnarray} where $c$ is a dimensionless parameter. In Ref.~\\cite{cohen}, the size $L$, which is regarded as an IR cut-off, was chosen to be the inverse Hubble expansion rate $L\\sim H^{-1}$ to naturally explain the observed vacuum energy density $\\rho_{\\Lambda 0} \\sim 3 M_{p}^{2}H_{0}^{2}$. However, it was shown that this choice can not explain the accelerated expansion of the Universe at present~\\cite{hsu}. This problem was solved in Ref.~\\cite{li} by choosing $L$ to be the future event horizon $R_{h}$. Extensive studies on this model have been done, and the HDE model with the future event horizon as the IR cut-off is found to be consistent with current observational data~\\cite{consistentobs}. Various aspects on the HDE model have been discussed in Ref.~\\cite{aspects}. There have also been further developments on the HDE model by introducing an interaction between dark energy and matter~\\cite{IHDE, newinteraction}. However, it was pointed out that this model with $L=R_{h}$ has a conceptual problem that the future event horizon, which determines $\\rho_{\\Lambda}$, depends on the future evolution of the Universe, hence violates causality~\\cite{causality}.    Subsequently, inspired by the HDE model, the holographic Ricci dark energy (RDE) model~\\cite{gao} was proposed in which $\\rho_{\\Lambda}$ is proportional to the Ricci scalar curvature ${\\cal R}$. It was shown that this model does not only avoid the causality problem and is phenomenologically viable, but also naturally solves the coincidence problem. Also, it was found that the causal connection scale $R_{\\rm CC}$ consistent with cosmological observations is given by $R_{\\rm CC}^{-2}=\\dot{H}+2 H^{2}$ which is proportional to ${\\cal R}$ in a flat universe~\\cite{Rcc}. This may provide us with a physical motivation for the RDE model. Cosmological constraints on this model were studied in Ref.~\\cite{cosmicconstraintsRDE}. Similar models have also been studied in Ref.~\\cite{RDE}.   In this paper, we consider the RDE model with an interaction between dark energy and matter, and call it the interacting Ricci dark energy (IRDE) model. We investigate the observational constraints on this model obtained from SNIa, CMB and BAO data. We organize this paper as follows. In section 2, we describe the IRDE model, and obtain analytic expressions for cosmic time evolution. In section 3, we discuss the observational constraints on this model. We summarize our results in section 4. ",
        "conclusions": "We have considered the IRDE model where the interaction rate is given by eq. (\\ref{intrate}). We have derived the analytic expressions for the Hubble parameter (\\ref{hubbleparameter}) and the energy density of dark energy (\\ref{rdes}) and matter (\\ref{rms}). Both $\\rho_{\\Lambda}$ and $\\rho_{m}$ include contributions with non-standard time evolution. We have also investigated current observational constraints on this model from SNIa, CMB and BAO observations. In particular, the CMB constraint is strongly affected by the interaction. We have shown that a nonvanishing interaction rate is favored by the observations, giving a reduction of $\\chi_{\\rm min}^{2}$. The best fit values are $\\Omega_{\\Lambda 0}=0.73\\pm0.03$, $\\alpha=0.45\\pm0.03$ and $\\gamma=0.15\\pm0.03$."
    },
    "0911/0911.4217_arXiv.txt": {
        "abstract": "{The observations of gamma-ray bursts (GRBs) such as 980425, 031203 and 060218, with luminosities much lower than those of other classic bursts, lead to the definition of a new class of GRBs -- low-luminosity GRBs. The nature of the outflow responsible for them is not clear yet. Two scenarios have been suggested: one is the conventional relativistic outflow with initial Lorentz factor of order of $\\Gamma_0\\ga 10$ and the other is a trans-relativistic outflow with $\\Gamma_0\\simeq 1-2$. Here we compare the high energy gamma-ray afterglow emission from these two different models, taking into account both synchrotron self inverse-Compton scattering (SSC) and the external inverse-Compton scattering due to photons from the cooling supernova or hypernova envelope (SNIC). We find that the conventional relativistic outflow model predicts a relatively high gamma-ray flux from SSC at early times ($<10^4 {\\rm s}$ for typical parameters) with a rapidly decaying light curve, while in the trans-relativistic outflow model, one would expect a much flatter light curve of high-energy gamma-ray emission at early times, which could be dominated by both the SSC emission and SNIC emission, depending on the properties of the underlying supernova and the shock parameter $\\epsilon_e$ and $\\epsilon_B$. The Fermi Gamma-ray Space Telescope should be able to distinguish between the two models in the future.} ",
        "introduction": "Long duration gamma-ray bursts are generally believed to result from the death of massive stars, and their association  with core-collapse supernovae (SNe of Type Ib/c) has been observed over the last decade. The first hint for such a connection came with the discovery of a nearby SN 1998bw in the error circle of GRB 980425 (Galama et al. 1998; Iwamoto et al. 1998) at distance of only about $40 {\\rm Mpc}$. The isotropic gamma-ray energy release is  of the order of only $10^{48} {\\rm erg}$ (Galama et al. 1998) and the radio afterglow modelling suggests an energy of $10^{49}-10^{50} {\\rm erg}$ in a mildly relativistic ejecta (Kulkarni et al. 1998; Chevalier \\& Li 1999). Recently {\\it Swift} discovered GRB 060218, which is the second nearest GRB identified so far (Campana et al. 2006; Cusumano et al. 2006; Mirabal \\& Halpern 2006; Sakamoto 2006). It is also an under-energetic burst with energy in prompt $\\gamma$/X-rays about $5\\times10^{49} {\\rm erg}$ and is associated with SN 2006aj. { Another burst, GRB 031203, is a third example of this group (Sazonov et al. 2004)}.  {If one assumes that GRB 060218-like bursts follow the logN-logP relationship of high-luminosity GRBs then, as argued by Guetta et al. (2004), no such burst with redshifts $z < 0.17$ should be observed by a HETE-like instrument within the next 20 years. Therefore, the unexpected discovery of GRB 060218 may suggest that these objects form} a different new class of GRBs from the conventional high-luminosity GRBs (Soderberg et al. 2006; Pian et al. 2006; Cobb et al. 2006; Liang et al. 2006; Guetta \\& Della Valle 2007; Dai 2008), although what distinguishes such low-luminosity GRBs (LL-GRBs) from the conventional high-luminosity GRBs (HL-GRBs) remains unknown. Their rate of occurrence may be one order of magnitude higher than that of the typical ones (e.g. Soderberg et al. 2008). They are rarely recorded because such { intrinsically dim GRBs can only be detected from relatively short distances} with present gamma-ray instruments.  The \"compactness problem\" of GRBs requires that the outflow of normal GRBs should be highly relativistic with a bulk Lorentz factor of $\\Gamma_0\\ga 100$ (Baring \\& Harding 1997; Lithwick \\& Sari 2001). The low-luminosity and softer spectra of low-luminosity GRBs relax this constraint on the bulk Lorentz factor. It was suggested that the softer spectrum and low energetics of GRB060218 (or classified as X-ray flash 060218) may indicate a somewhat lower Lorentz factor of the order of $\\sim 10$ (e.g. Mazzali 2006; Fan et al. 2006; Toma et al. 2006).    An alternative possibility is that low-luminosity GRBs are driven by a trans-relativistic outflow with $\\Gamma_0\\simeq2$ (Waxman 2004; Waxman, M\\'esz\\'aros \\& Campana, 2007; Wang, Li, Waxman \\& M\\'esz\\'aros  2007; Ando \\& M\\'esz\\'aros 2008). The flat light curve of the X-ray afterglow of the nearest GRB, GRB 980425/SN 1998bw, up to 100 days after the burst, has been argued to result from the coasting phase of a mildly relativistic shell with an energy of a few times $10^{49}$ erg (Waxman 2004). From the thermal energy density in the prompt emission of GRB060218/SN2006aj, Campana et al. (2006) inferred that the shell driving the radiation-dominated shock in GRB 060218/SN 2006aj must be mildly relativistic. This trans-relativistic shock could be driven by the outermost parts of the envelope that get accelerated to a mildly relativistic velocity when the supernova shock accelerates in the density gradient of the envelope of the supernova progenitor (Colgate 1974; Matzner \\& McKee 1999; Tan et al. 2001), or it could be due to a choked relativistic jet propagating through the progenitor (Wang et al. 2007).   In this paper, we investigate the high energy afterglow emission from low luminosity GRBs for both {relativistic and trans-relativistic} models and explore whether the Fermi Gamma-ray Space Telescope can distinguish between these two models with future observations of low-luminosity GRBs. Since photons from the underlying supernova are an important seed photon source for inverse-Compton (IC) scattering, we consider both synchrotron self-inverse Compton (SSC)  and external IC scattering due to supernova photons (denoted by SNIC hereafter). At early times (within a few days after the burst), a UV-optical SN component was recently detected from SN2006aj and SN2008D (Campana et al. 2006; Soderberg et al. 2008), which has been interpreted as the cooling SN envelope emission after being heated by the radiation-dominated shock\\footnote{Although there is a disagreement on the origin of the early UV-optical emission from SN2006aj, an agreement has been reached for that of SN2008D (Waxman et al. 2007; Soderberg et al. 2008; Chevalier \\& Fransson 2008).} (Waxman et al. 2007; Soderberg et al. 2008; Chevalier \\& Fransson 2008). In our calculation, we take into account this seed photon source in addition to the late-time supernova emission, which peaks after ten days.   Recently, Ando \\& Meszaros (2008) discussed the broadband emission from SSC and SNIC for a trans-relativistic ejecta in a low-luminosity GRB at a particular time-- the ejecta deceleration time. Here we study the time evolution of the high-energy gamma-ray emission resulted from such IC processes and consider both the trans-relativistic ejecta and the highly relativistic ejecta scenarios for low-luminosity GRBs.  The paper is organized as follows. First, we describe the dynamics of shock evolution in the two models in $\\S$ \\ref{models}.  In $\\S$3, we present the formula for the calculation of the inverse-Compton emissivity. For the SNIC emission, we take into account the anisotropic scattering effect.  Then we present the results of the spectra and light curves of SSC and SNIC emission for the two different models and explore the detectability of these components by Fermi Large Area Telescope (LAT) in $\\S$ 4. Finally, we give the conclusions and discussions. ",
        "conclusions": "The external stellar wind provides a source of Thomson opacity to scatter the supernova emission, thus a quasi-isotropic, back-scattered SN radiation field is present.  Let's study whether this component is important. For a GRB shock locating at radius $R$ and moving with a Lorentz factor $\\gamma$, the Thomson scattering optical depth of the wind is $\\tau_{\\rm w}=\\sigma_TRn=\\sigma_TKR^{-1}=2.0\\times10^{-4}\\dot{m}R_{15}^{-1}$. The scattered SN energy density by the wind in the comoving frame of the blast wave is $u'^{\\rm w}_{\\rm SN}=(L_{SN}/{4\\pi c R^2}) \\tau_{\\rm w} \\gamma^2$ due to the relativistic boosting effect. On the other hand, the energy density of the supernova photons impinging the shock from behind is $u'_{\\rm SN}=\\gamma^{-2}(L_{SN}/{4\\pi c R^2})$ (see Eq.21). The ratio between these two energy densities is $u'^{\\rm w}_{\\rm SN}/u'_{\\rm SN}=2.0\\times10^{-4}\\dot{m}R_{15}^{-1} \\gamma^4$. For the trans-relativistic ejecta model with $\\Gamma_0=2$, the radius of the GRB ejecta is $R=1.9\\times10^{14}$ $\\rm{cm}$ at $t=10^3 \\rm{s}$, so $u'^{\\rm w}_{\\rm SN}\\ll u'_{\\rm SN}$ for typical wind parameters. For the conventional relativistic ejecta model, the Lorentz factor and radius of the GRB ejecta are respectively $\\gamma\\sim4$ and $R=1.0\\times10^{15}$ $\\rm{cm}$ at $t=10^3 \\rm{s}$, so we also have $u'^{\\rm w}_{\\rm SN}\\ll u'_{\\rm SN}$. This means that the wind-scattered supernova radiation field is subdominant compared to the direct impinging supernova photon field and hence we neglect its contribution to the high-energy gamma-ray emission.  Our estimate of the pair-production opacity for high-energy photons in $\\S$ \\ref{tao} is based on the common assumption that the colliding photons are isotropic in the rest frame. However, as we have shown above, the low-energy photons from the supernova essentially move radially outward before colliding with high-energy photons (e.g. Wang, Li \\& \\Meszaros 2006). Therefore, the collision process between the high-energy photons and soft supernova photons is anisotropic, which would decrease the pair-production opacity. This will be the subject of a more detailed future calculation and would be useful to explore whether $\\rm{TeV}$ photons can escape from the source, which is important for checking the detectability by ground-based Cherenkov detectors such as Magic, VERITAS, Milago, HESS, ARGO etc. Additionally, besides the high-energy gamma-ray emission discussed in this work, the high-energy neutrino emission arising from $p\\gamma$ interactions between shock-accelerated protons and photons from the supernova may also provide a constraint on the model for low-luminosity GRBs (e.g. Yu et al. 2008).  Trans-relativistic ejecta may also exist in the usual high luminosity long GRBs, besides in low luminosity ones, since {accelerating shocks are} expected to accompany the supernova. Berger et al. (2003) and Sheth et al. (2003) found that the radio and optical afterglow indicates a low velocity component more than $1.5 \\rm{days}$ after the explosion in GRB030329/SN2003dh. However, since the  highly relativistic ejecta is much more energetic than the trans-relativistic component, high-energy gamma-ray emission from the latter component {could easily remain hidden}.  In summary, we have calculated the spectra and light curves of the high-energy gamma-ray afterglow emission from low luminosity GRBs for the two main models in the literature, i.e. the trans-relativistic ejecta model ($\\Gamma_0\\simeq2$) and the conventional highly relativistic ejecta model ($\\Gamma_0\\ga 10$), considering both synchrotron self inverse-Compton (SSC) and the external inverse-Compton due to photons from the underlying supernova/hypernova. Our analysis takes into account a full Klein-Nishina cross section for inverse Compton scatterings, the anisotropic scattering of supernova photons and the opacity for high energy photons due to annihilation with low-energy photons (mainly from the supernova/hypernova).   We find that for the supernova case the conventional relativistic outflow model predicts a relatively high gamma-ray flux from SSC at early times ($<10^5-10^6 {\\rm s}$ for typical parameters) with a rapidly decreasing flux, while in the trans-relativistic outflow model, a much flatter light curve of high-energy gamma-ray emission dominated by the SSC emission, is expected at early times. For the hypernova case, the SSC emission also dominates and decays sharply at early time in the conventional relativistic ejecta model, while for the trans-relativistic ejecta model, both the SSC emission and the SNIC emission could be dominant at early time, depending on the shock parameters $\\epsilon_e$ and $\\epsilon_B$. The main difference between these two models arises from their different initial Lorentz factors, which induces a different dynamical evolution of the shock at early times. As a result, different observational features arise, such as different light curve shapes, different flux levels and different dominant components (as detailed in $\\S$ \\ref{spectra and light curves}). As shown in $\\S$ \\ref{detectability of Fermi LAT}, high-energy gamma-ray emission can be detected in both models as long as $\\epsilon_e$ is large enough, although detection from the conventional relativistic ejecta  is much easier. Thus, with future high energy gamma-ray observations by Fermi LAT, one can expect to be able to distinguish between the two models based on the above differences in observational features.    {"
    },
    "0911/0911.4484_arXiv.txt": {
        "abstract": "We present an analysis of the distribution of structural properties for Milky Way-mass halos in the Millennium-II Simulation (MS-II).  This simulation of structure formation within the standard \\lcdm\\ cosmology contains thousands of Milky Way-mass halos and has sufficient resolution to properly resolve many subhalos per host.  It thus provides a major improvement in the statistical power available to explore the distribution of internal structure for halos of this mass.  In addition, the MS-II contains lower resolution versions of the Aquarius Project halos, allowing us to compare our results to simulations of six halos at a much higher resolution.  We study the distributions of mass assembly histories, of subhalo mass functions and accretion times, and of merger and stripping histories for subhalos capable of impacting disks at the centers of halos.  We show that subhalo abundances are {\\it not} well-described by Poisson statistics at low mass, but rather are dominated by intrinsic scatter.  Using the masses of subhalos at infall and the abundance-matching assumption, there is less than a 10\\% chance that a Milky Way halo with $\\mvir =10^{12}\\,\\msun$ will host two galaxies as bright as the Magellanic Clouds.  This probability rises to $\\sim 25\\%$ for a halo with $\\mvir=2.5 \\times 10^{12} \\,\\msun$.  The statistics relevant for disk heating are very sensitive to the mass range that is considered relevant.  Mergers with infall mass : redshift zero virial mass greater than 1:30 could well impact a central galactic disk and are a near inevitability since $z=2$, whereas only half of all halos have had a merger with infall mass : redshift zero virial mass greater than 1:10 over this same period. ",
        "introduction": "Our understanding of galaxy formation is strongly influenced by our own Galaxy. Several of the apparent points of tension between observations and the predictions of the now-standard $\\Lambda$ Cold Dark Matter (\\lcdm) structure formation model -- for example, the abundance of dark matter subhalos in comparison to the relative paucity of observed Local Group satellites and the existence of the thin Galactic disk in light of the ubiquity of dark matter halo mergers -- are most clearly demonstrated from Milky Way data.  While it is clearly an excellent testing ground for a wide array of phenomena related to galaxy formation, the Milky Way (MW) is nonetheless a single object and may not be representative in its properties.  This is particularly true with respect to infrequent, stochastic phenomena such as major mergers: while \\lcdm\\ robustly predicts the mean number of mergers onto halos of a given mass, the halo-to-halo scatter is large.  The situation is similar for cosmological re-simulations of individual MW-mass dark matter halos.  In the quest to understand the Milky Way's formation, numericists have been simulating individual MW-mass dark matter halos at increasingly high resolution.  By investing substantial computational resources, the most recent generation has reached one billion particles per halo, with particle masses $\\sim\\! 1000 \\,\\msun$ \\citep{diemand2008, springel2008, stadel2009}.  Only a small number of such simulations can be performed and there is no guarantee that the simulated halos will be ``typical.''  In order to explore whether various characteristics of the Milky Way (e.g., its satellite galaxy population, thin disk, or merger history) are consistent with \\lcdm\\ expectations, we must understand the statistical properties predicted for MW-mass halos.  The advent of modern galaxy surveys has enabled such an undertaking observationally: the Sloan Digital Sky Survey (SDSS; \\citealt{york2000}) has provided a huge sample of galaxies with stellar masses similar to that of the MW, although only mean halo properties averaged over large numbers of similar galaxies can readily be measured \\citep{mandelbaum2006b, more2009}.  Predicting the formation histories and redshift zero properties of the halos of a representative population of MW-mass galaxies requires numerical simulations that combine large volume with high mass resolution.  Resolving the internal structure of $\\sim$ ten subhalos per host halo requires approximately one hundred thousand particles per host, which is not trivial when combined with the goal of obtaining a statistical sample of MW-mass halos.  This would, for example, require $10^{12}$ particles in the volume probed by the Millennium Simulation (MS; \\citealt{springel2005b}), a 100-fold increase over the particle number actually used.  The only computationally feasible approach at present is to use a smaller simulation volume combined with higher mass resolution.  In this paper, we use the Millennium-II Simulation (MS-II; \\citealt{boylan-kolchin2009}), which was {\\it designed} to meet these requirements.  It allows us to investigate statistical properties of MW-mass halos, including their assembly histories and their subhalo content.  Throughout this work, we also compare results for our representative sample to the six halos from the Aquarius simulation series of \\citet{springel2008}.  These cosmological resimulations of Milky Way-mass dark matter halos have approximately one thousand times better mass resolution than the \\millen\\ (see Section~\\ref{subsec:aquarius}).  The basic properties of the \\millen\\ and of the sample of MW-mass dark matter halos we investigate are given in Section~\\ref{sec:simulations}, together with a convergence study for subhalos.  Section~\\ref{sec:mainHalos} discusses the assembly history of MW-mass halos in terms both of their virial mass and of their central gravitational potential, and relates these to their $z=0$ concentrations and angular momenta.  Section~\\ref{sec:subhalos} investigates the subhalo population of MW-mass halos at $z=0$, including the subhalo mass function and whether its scatter is Poissonian, the distribution of masses for the most massive subhalo in each host, and typical accretion times for subhalos as a function of their current mass.  In Section~\\ref{sec:mergers}, we explore the merger histories of MW-mass halos, focusing on mergers that are likely to result in galaxy-galaxy mergers.  Section~\\ref{sec:discussion} contains a discussion of our results and their implications, while Section~\\ref{sec:conclusions} gives our conclusions.  Some properties of satellite subhalos are quantified in the Appendix.  Throughout this paper, all logarithms without a specified base are natural logarithms. ",
        "conclusions": "\\label{sec:conclusions} We have presented a statistical study of Milky Way-mass halos from the Millennium-II Simulation, which contains over 7000 halos in the mass range $10^{11.5} \\le \\mvir \\le 10^{12.5}\\,\\hmsun$.  Our principal results can be summarized as follows: \\begin{itemize} \\item As several previous studies have shown, halo growth proceeds in an ``inside-out'' fashion: the central gravitational potential of a typical MW-mass halo reached half of its present-day value by $z=4$, on average, whereas the virial mass reached half of its present-day value at $z=1.2$. \\item The ratio of $z=0$ mass to the host halo mass for the most massive subhalo of a MW-mass halo has a broad distribution that peaks at 1\\%.  The corresponding quantity computed using infall mass $\\macc$ instead of $\\msubo$ peaks at 3.5\\% \\item The differential and cumulative subhalo mass functions, computed in terms of $\\mu \\equiv \\msub/\\mvir$, are both well-fitted by a power law with an exponential cut-off at large $\\mu$.  We find both the abundance per log decade in $\\mu$ and the number of subhalos with mass greater than $\\mu$ to be proportional to $\\mu^{-0.935}$ for small $\\mu$. \\item The scatter in the cumulative mass function at fixed $\\mu$ is only well-approximated by a Poisson distribution at large $\\mu$ (corresponding to mean occupation number $\\langle N \\rangle \\la 4-5$). Intrinsic scatter of approximately 18\\%, independent of the host halo mass, becomes increasingly important at lower $\\mu$ (larger $\\langle N \\rangle$) and is dominant over Poisson scatter at $\\mu \\la 3\\times 10^{-4}$ ($\\langle N \\rangle \\ga 20$). \\item The Negative Binomial distribution with variance in $\\cmf$ that is equal to the sum of a Poisson term and a fractional intrinsic scatter of 18\\% matches the subhalo occupation data well at all $\\mu$ for host halo masses between $10^{12}$ and $10^{15} \\,\\hmsun$. \\item The statistic $\\alpha_2(\\mu)$ [equation~(\\ref{eq:alpha})], which is frequently used to characterize deviations of the HOD from Poisson, is not discriminating: distributions that differ strongly from Poisson can lead to few percent differences of $\\alpha_2$ from unity. \\item Accretion redshifts of $z=0$ subhalos vary systematically with host halo mass and with the ratio of subhalo to host masses.  Massive subhalos at $z=0$ were typically accreted more recently.  This is a dynamical effect: massive subhalos accreted at early times either merge via dynamical friction or are heavily tidally truncated.  At fixed $\\mu$, subhalos of more massive host halos were accreted more recently, reflecting the hierarchical build-up of dark matter halos. Subhalos with $\\mu=0.001$ have a median accretion redshift of $\\zacc=1.1$ for hosts with $\\mvir \\approx 10^{12}\\,\\hmsun$ and $\\zacc=0.6$ for halos with $\\mvir > 10^{13.5}\\,\\msun$. \\item The frequency of central mergers -- that is, the merger of a satellite with the central region of a MW-mass halo -- is a strong function of the satellite's $\\macc$.  Since $z=2$ 90\\% of MW-mass halos have experienced a central merger with a satellite having $\\macc/\\mviro > 0.03$ while 50\\% have had such a merger with a $\\macc/\\mviro > 0.1$ satellite, and only 30\\% with a $\\macc/\\mviro > 0.15$ satellite.  The compatibility of thin stellar disks with the frequent mergers expected in \\lcdm\\ thus depends strongly on exactly which mergers destroy disks.  If $\\macc/\\mviro > 0.03$ events are sufficient, the Milky Way must have an accretion history that lies among the quietest 10\\% for halos of similar mass, while if $\\macc/\\mviro > 0.1$ is required, then the Milky Way may have had a fairly typical merger history. \\end{itemize}"
    },
    "0911/0911.1574_arXiv.txt": {
        "abstract": "Microwave zebra pattern structure is an intriguing fine structure on the dynamic spectra of solar type IV radio burst. Up to now, there isn't a perfect physical model for the origin of the solar microwave zebra pattern. Recently, Ledenev, Yan and Fu (2006) put forward an interference mechanism to explain the features of microwave zebra patterns in solar continuum events. This model needs a structure with a multitude of discrete narrow-band sources of small size. Based on the model of current-carrying plasma loop and the theory of tearing mode instability, we proposed that the above structure does exist and may provide the main conditions for the interference mechanism. With this model, we may explain the frequency upper limit, the formation of the parallel and equidistant stripes, the superfine structure and intermediate frequency drift rate of the zebra stripes. If this explanation is valid, the zebra pattern structures can reveal some information of the motion and the inner structures of the coronal plasma loops. ",
        "introduction": "During much of solar flares, it is very frequently found that a kind of intriguing fine structure pattern superposed on the solar radio broadband spectrum of type IV bursts, and behaved as a series of almost parallel and equidistant stripes in the dynamic spectrum. Such structure is called zebra pattern. Most often, zebra patterns are observed in meter and decimeter frequency range, and with up to 10 and more stripes (Slottje, 1972). It is seldom to observe zebra patterns in microwave frequency range in the early observations, and even if we found them, there always only 3 or 4 stripes in a zebra pattern structure (Ning et al, 2000; Ledenev, Yan and Fu, 2001). However, in recent microwave observations, some remarkable zebra patterns are also found with up to 30 stripes in the frequency range of 2.60 -- 3.80 GHz (Chernov et al, 2005). Fig. 1 is an example of zebra pattern occurred in the frequency of 2.90 -- 3.80 GHz observed at Chinese Solar Broadband Radiospectrometer (SBRS/Huairou) in 02:42:55 -- 02:43:20 UT, 13 Dec. 2006 in the famous flare event, with strong right polarization, 5 stripes, and the duration is about 10 seconds, cognizably. Up to now, the upper limit frequency of Zebra pattern structure is below 6 GHz, and the corresponding wavelength is about 5 cm (Altyntsev et al, 2005). The duration of zebra pattern event is from several to a few decades of seconds. It is uncommon to last for more than 30 seconds.  \\begin{figure} \\begin{center} \\includegraphics[width=8cm]{fig1.ps} \\caption{An example of Zebra pattern structure occurred in the frequency of 2.90 -- 3.80 GHz observed at Chinese Solar Broadband Radiospectrometer (SBRS/Huairou) in 02:42:55 -- 02:43:20 UT, 13 Dec. 2006. The upper and lower panels are left and right polarization components, respectively.} \\end{center} \\end{figure}  If we fix the time and plot the profile of the emission flux with respect to frequency in the Zebra pattern structure, we may find that the flux profile behaves periodic feature, and the period is the frequency gap between two stripes. Fig. 2 gives an example profile of the emission flux at 02:43:05 UT, 13 Dec. 2006 of the Zebra pattern showing in Fig.1. There are 5 peaks in this profile, and each represents one stripe.  \\begin{figure} \\begin{center} \\includegraphics[width=8cm]{fig2.eps} \\caption{A profile of the zebra pattern emission flux with respect to the frequency at a fixed time of 02:43:05 UT, 13 Dec. 2006. The solid and the dash-ploted curves are indicated the right and left polarization components, respectively.} \\end{center} \\end{figure}  The another main feature of Zebra pattern structure is the intermediate frequency drift against the continuum emission of type IV radio outbursts (Slottje, 1972; Mollwo, 1983; etc). Fig. 1 shows that frequency drift rate is about 55 MHz/s during 02:42:59 -- 02:43:03 UT, and -45 MHz/s during 02:43:03 -- 02:43:10 UT around the frequency of 3.50 GHz, negatively or positively.  Chernov et al (2003, 2008) have found that the Zebra pattern stripes in the microwave range often have some superfine structures, consisting of separate spike-like pulses with millisecond duration. Chen and Yan (2007) also found that Zebra stripes consist of periodically narrow band pulsating superfine structures, and the period is about 30 milliseconds. Because of the saturation around the center of the Zebra stripes associated with the flare event occurred in 13 Dec. 2006 (see in Fig. 1), we could not distinguish the obvious superfine structures in this event. However, from the limb parts of the Zebra stripes, we may also find some evidences of the quasi-periodic narrow band pulsating superfine structures. Fig. 3 shows the quasi-periodic narrow band pulsating superfine structures of the Zebra stripe in a segment of 02:43:02.8 -- 02:43:03.8 UT at frequency of 3.65 GHz, near the limb of the fourth stripe (numbered from low frequency to high frequency). And the period of the pulsating superfine structures is about 30 -- 35 milliseconds.  \\begin{figure} \\begin{center} \\includegraphics[width=8.cm]{fig3.eps} \\caption{The quasi-periodic narrow band pulsating superfine structures of the zebra stripe in a segment of 02:43:02.8 -- 02:43:03.8 UT at frequency of 3.65 GHz, near the limb of one stripe.} \\end{center} \\end{figure}  Generally, people think that the microwave Zebra pattern may provide some useful information about the kernel of the flaring regions. It is necessary for any model of zebra pattern to explain the following features: (1) the upper limit of frequency of Zebra pattern, (2) almost parallel and equidistant stripes, (3) superfine structures, (4) intermediate frequency drift rate. In order to interpret the formation of Zebra pattern, a great number of theoretical models were proposed. These models can be classified simply into two groups:  (1) Isogenous models, which proposed that all the stripes in a Zebra pattern are generated from a single emission source, and the emission mechanism is a kind of nonlinear coupling between two Bernstein modes, or Bernstein mode and some electrostatic upper hybrid mode waves. The frequency gap between two stripes is very close to the frequency of electron cyclotron emission (Rosenberg, 1972; Chiuderi et al, 1973; Zaitsev and Stepanov, 1983).  (2) Heterogenous models, which proposed that all the stripes in a Zebra pattern are generated from different emission source regions, and different source region are located at different positions in the magnetic flux tube. In each source region, some resonant conditions are satisfied, and the emission mechanism is probably the coupling between whistler wave and electrostatic upper hybrid mode wave which triggered by some plasma micro-instabilities. The frequency gap between two stripes is greatly different to the frequency of electron cyclotron emission (Kuijpers, 1975; Fomichev and Fainshtein, 1981; Mollwo, 1983; Ledenev, Yan, and Fu, 2001; Chernov et al, 2005; Altyntsev et al, 2005).  However, up to now, there is no perfect theoretical model which can explain Zebra pattern satisfactorily. Recently, Ledenev, Yan, and Fu (2006) proposed that Zebra pattern is possibly formed from some interference mechanism in the propagating processes. They assumed that there are some inhomogeneous layers with small size in solar coronal plasma, and such structure will change the radio waves into direct and reflected rays. When the direct and reflected rays meet at the position of observer, interference will take place and form Zebra pattern structure.  When the dimensional size of an emission source region is smaller than the minimum wavelength of emission spectrum, the source region can be treated as a point source. At the same time, when the emission spectrum is continuum, then it is possible to form an interference pattern. The dimensional size of the source region is much smaller than the characteristic size of plasma density gradient, and can be regarded as a point source. However, the dimensional size of the source region is always much larger than the emission wavelength. The interference condition requires that the source region has a narrow, zonal inhomogeneous interior structure to keep some definite phasic difference between the direct and the reflected rays. At the same time, in order to generate definite interference strength, the number of the point sources should be abounded. Then, is there plenty of such abounded point sources in the solar flaring region? What mechanism can generate such structures?  In fact, the abounded point sources are necessary not only in the interference model of Ledenev, Yan, and Fu (2006), but also in other zebra pattern models. For example, the model of LaBelle et al (2003), in which assumed that the emission is generated from the double plasma resonant layer in coronal loop. The presence of localized density irregularities within the type IV source region leads to trapping of the upper hybrid Z-mode waves in density enhancements, transforms into electromagnetic waves by the electron-cyclotron maser mechanism, and forms the zebra pattern. In such case, the localized density irregularity is a necessary condition, and a number of point sources can meet this condition naturally.  Briefly, the structure with abounded point sources is an important condition for the formation of zebra pattern. Based on the analysis of current-carrying plasma loop model and the related resistive tearing-mode instability, this work proposed that the tearing mode magnetic islands in the current-carrying plasma loop can form a reasonable structure to generate the interference mechanism and produce the Zebra patterns. We introduce the formation of the tearing-mode magnetic islands in current-carrying plasma loop in section 2. Then in section 3, we present a detailed explanation of the interferential rays and the formationtions of Zebra patterns. At the final, some summaries and discussions are given in section 4. ",
        "conclusions": "the interference model can give a reasonable explanation of the radio emission with Zebra pattern structures, and the current-carrying plasma loop model can provide all the necessary conditions for the interference model, and can be applied to explain the other features, such as the the superfine structure, the upper frequency limit, the intermediate frequency drifting rate of the zebra stripes, and the durations. If this model is really valid, the zebra pattern can provide an useful tool for studying the current-carrying plasma systems in the solar flaring region. From the scrutinizing of the microwave Zebra pattern structures, we may get much of information about the motion and the inner structure feature of the plasma loop, and can reconstruct the space configuration of the emission source region, because it is believed that they are closely related to the flare primary-energy releasing processes. However, there are much of works need to do theoretically and observationally.  On the other hand, the large numbers of magnetic islands inside the plasma loops can provide much more opportunities to accelerate the large numbers of energetic particles in flare events. According to the previous works (Miller et al, 1997; Vlahos, Isliker \\& Leperti, 2004; etc), the number of accelerated electrons with energy above 20 keV is roughly $10^{34} - 10^{37}$ electrons $s^{-1}$ in a flare event, and in some X-class flares the number can be reached to $10^{37}$ electrons $s^{-1}$. The current-carrying plasma loop model indicates that the particle accelerations may occur not only in the cusp configuration above the coronal loop, but also can take place inside the whole loop. With this model, because the total volume of magnetic islands is much lager than that of the current sheet near the cusp configuration above the loop, and the plasma is denser in the loops, there are much more electrons which can be accelerated by the induced electric field generated from the tearing magnetic reconnections. As for the number of magnetic islands associated with the accelerated electrons, we may give a roughly estimation: the width of the magnetic island is about 2 km, thickness about 5 cm, length can be assumed as 1\\% of the loop's length (assumed as about $10^{5}$ km). Then the volume of one island is in the order of $10^{14}$ cm$^{3}$. The plasma density could be assumed $10^{10}-10^{11}$ cm$^{-3}$. Then the number of accelerated electrons around one island is about $10^{24}-10^{25}$. In order to meet the number of $10^{37}$ electrons s$^{-1}$ accelerated in big flares, the number of islands should be $10^{12}-10^{13}$. If we suppose the number of flux tubes in the flaring region is about 10 -- 100, then there are about $10^{10}-10^{11}$ magnetic islands in one plasma loop. We may find that the volume of all islands is only a small fraction ($\\sim$ 0.1\\%) in the flaring region. In fact, it was a big problem to explain the number of $10^{37}$ electrons s$^{-1}$ accelerated in some X-class flares perfectly. However, in this work, it is not our main task to investigate the particle acceleration, the above discussion is only a roughly estimation.  Excited by these energetic particles accelerated around the tearing-mode magnetic islands in the current-carrying plasma loops, a series of fine structures will be formed in the microwave spectrograms. In fact, by using this model, we explained the fast quasi-periodic pulsations occurred in the famous flare event of 13 Dec. 2006 (Tan et al, 2007). This work is an another attempt. However, it needs to study in more detailed."
    },
    "0911/0911.4956_arXiv.txt": {
        "abstract": "{We introduce and publicly release a new standalone code, {\\sc adaptsmooth}, which serves to smooth astronomical images in an adaptive fashion, in order to enhance the signal-to-noise ratio (S/N). The adaptive smoothing scheme allows to take full advantage of the spatially resolved photometric information contained in an image in that at any location the minimal smoothing is applied to reach the requested S/N. Support is given to match more images on the same smoothing length, such that proper estimates of local colours can be done, with a big potential impact on multi-wavelength studies of extended sources (galaxies, nebulae). Different modes to estimate local S/N are provided. In addition to classical arithmetic-mean averaging mode, the code can operate in median averaging mode, resulting in a significant enhancement of the final image quality and very accurate flux conservation. To this goal also other code options are implemented and discussed in this paper. Finally, we analyze in great detail the effect of the adaptive smoothing on galaxy photometry, in particular to check for conservation of surface brightness (SB) and aperture-integrated fluxes: deviations in SB with respect to the original image can be limited to $<0.01$~mag, with flux difference in apertures of less than 0.001 mag.  With respect to already existing adaptive smoothing codes, mainly developed by the X-ray community, {\\sc adaptsmooth} uniquely provides the median averaging mode and a greater flexibility to adapt to different noise regimes and degrees of image contrast, which can occur especially in optical and near-IR imaging.}   Keywords: {techniques: image processing, photometric.}    \\bibliographystyle{aa} ",
        "introduction": "Extended astronomical sources, like galaxies and nebulae, are characterized by a huge dynamical range of surface brightness (SB) throughout their extent. This range spans several orders of magnitude and can be captured by current astronomical imaging devices at optical and near-IR wavelengths. However, as a consequence the signal-to-noise ratio (S/N) of the information contained in the image varies wildly depending on the (surface) brightness at any given position. Therefore the problem arises how to extract the information by optimally balancing the S/N of the flux measurement and the effective spatial resolution.  Imaging devices are typically operated in background noise limited conditions, i.e. the main source of noise comes from the background rather than from the Poisson photon noise of the source itself, which becomes dominant only at very high S/N levels. While the background noise puts a well defined lower limit to the brightness of very localized/point-like sources that can be detected at a given S/N, when imaging diffuse objects the detection limit and the S/N at which a given surface brightness is measured can be improved by combining the flux measured in increasingly more pixels. Two approaches are available to do this: binning and filtering (or smoothing). In the first case, a number $N$ of pixels that share a similar intensity level and/or match given spatial constraints are averaged, resulting in an improvement in S/N of $\\sqrt N$, in the ideal case of uniform uncorrelated noise. Examples of the binning approach are azimuthally averaged profiles of the surface brightness of galaxies along elliptical isophotes \\citep[e.g. the algorithm of][as implemented in the {\\sc iraf} task {\\it ellipse}]{busko_96,jedre87} and the tessellation algorithms that are extensively utilized in the analysis of integral field spectroscopic observations \\citep[for an implementation see][]{cappellari_copin_2003}. As a result of binning, the number of resolution elements is decreased (hence the spatial resolution is worsened or even the dimensionality is lowered) and all the available information in the image is compressed at a higher S/N level. The filtering (or smoothing) approach consists in cutting the highest spatial frequencies that contribute most of the noise. The result of filtering is a smoothed image, where the intensity of each pixel is replaced by a weighted average of the neighbouring pixels, over an extent and with a weighting scheme that depend on the particular kernel adopted. Popular choices of kernels are 2D gaussians and top-hat functions (resulting in the so-called `boxcar' smoothing). As with the binning, also in the filtering approach the S/N is increased by a factor of the order of $\\sqrt N$ (where $N$ is the number of pixels in the kernel). Although the number of spatial elements (pixels) is preserved, the information in the smoothed image is correlated and spatial resolution is lost, proportionally to the kernel size. In fact, by cutting the highest spatial frequencies not only the noise but also the sharpest features of the image are thrown away.  Both approaches have of course their own pros and cons. Filtering/smoothing, in general, offers the best 2D rendition of the information, although it is often non-optimal in terms of compromise between spatial resolution and S/N enhancement, which, reversely, is a point of strength of the tessellation algorithms. In the regions of high surface brightness (SB) the S/N is high, such that no or minimal smoothing is requested, while smoothing with increasingly larger kernels is required to enhance the S/N of lower and lower SB regions to acceptable levels. The optimal solution is to apply to the image a smoothing with variable kernel size, which matches the local S/N: this is what is called {\\it adaptive smoothing}.  The X-ray community, pushed by the necessity to deal with diffuse sources in very low S/N regimes, has produced a number of codes to implement this basic concept. A (possibly incomplete) list includes: {\\sc fadapt} \\citep[distributed as part of the package {\\sc ftools},][]{blackburn_95}; \\cite{huang_sarazin_96}; {\\sc asmooth} \\citep[included in the Chandra data reduction package CIAO, also known as {\\sc csmooth} in the IDL implementation,][]{ebeling+06}; \\cite{sanders_06}; {\\sc asmooth} in the XMM-Newton Science Analysis System.  The new, standalone, adaptive smoothing code {\\sc adaptsmooth} presented in this paper has been developed with in mind mainly typical optical and near-IR imaging (although its application at other wavelengths is well supported too) and allows to easily handle different S/N regimes, even in cases where very limited information about the noise properties of the images are available, as detailed in Sec. \\ref{subsec:noise}. We introduce {\\it median} adaptive smoothing in addition to the {\\it mean} adaptive smoothing, which is the only method implemented in the other codes mentioned above. The improvement deriving from using the median as average estimate is especially notable in presence of strong intensity gradients (e.g. bright point-like sources overlaid on an area low surface brightness) and is thoroughly illustrated in Sec. \\ref{sec:testgal}. We also provide different methods to prevent cross-talking between pixels at different S/N levels, which can be switched on or off by the user (see Sec. \\ref{subsec:xtalk}) as opposed to the rigid implementation \\citep[e.g. {\\sc a/csmooth},][]{ebeling+06} or lack of them in other codes. Similarly to {\\sc a/csmooth} in CIAO and {\\sc asmooth} in the XMM package, {\\sc adaptsmooth} outputs maps of the smoothing length at each location (``smoothing masks'') and allows to use them as input for smoothing other images: this feature allows matching smoothing scales across different images and is therefore essential to produce, e.g., color maps (see Sec. \\ref{subsec:inmask}).  {\\sc adaptsmooth} is made publicly available via this URL \\texttt{http://www.wiki-site.com/index.php/ADAPTSMOOTH} or on request to the author. In the following sections we describe the concept of the algorithm and the actual code implementation with the available options. We show a couple of examples on astronomical images. Finally we analyze in detail the reliability of the code output in terms of conservation of flux and surface brightness in order to perform accurate surface photometry (specifically of galaxies) in different conditions. ",
        "conclusions": "We have introduced {\\sc adaptsmooth}, a new standalone code to perform adaptive smoothing of astronomical images, which is made publicly available. {\\sc adaptsmooth} uses either median or arithmetic-mean average within apertures whose size is variable and determined in order to provide pixel surface brightness measurements above the minimum requested S/N. In this way the best compromise between photometric accuracy and spatial resolution is obtained at any location of an image. The use of smoothing masks allows to work with multi-band imaging and hence to spatially resolve the spectral energy distribution of extended objects. Adaptive smoothing extends to much dimmer surface brightness the ability of measuring local fluxes and colours and hence can have a very strong impact on the study of extended objects, such as galaxies \\citep[as shown, e.g., in][]{zibetti_charlot_rix_09a,zibetti_charlot_rix_10} and nebulae (see Fig. \\ref{fig:SFregion}).  We provide and discuss different ways to estimate local S/N in an image (based on Poisson+background noise, on background noise only, or on local signal fluctuations) and options to reduce artifacts (smoothing level cuts). Finally, we analyze in detail the qualitative and quantitative effects of adaptive smoothing on galaxy images. In the best operation mode for a typical optical image (i.e. median smoothing and smoothing level cut 1) the azimuthally averaged SB profile of the smoothed images deviates by less than 0.01 mag from that of the original image, while the flux growth curve deviates by less than 0.001 mag.  Compared to other existing codes for adaptive smoothing, {\\sc adaptsmooth} provides the broadest variety of options to operate in different regimes. In particular, {\\sc adaptsmooth} uniquely implements a median-averaging mode, which turns out to provide best results in terms of flux conservation and minimization of artifacts originating from bright sources when the photon statistics are well populated. The use of smoothing level cuts is much more flexible than in other codes (e.g. in {\\sc csmooth}). Finally, the three possible choices for S/N estimate cover essentially all regimes for typical optical and near-IR images (as well as for shorter wavelengths) and guarentee good results even with minimal knowledge of the data error model. Such support is only partly provided by other adaptive smoothing codes (e.g. {\\sc csmooth} supports Poisson and background dominated noise, but no direct r.m.s. estimate on the image, while XMM {\\sc asmooth} only supports Poisson errors or user provided variance maps)."
    },
    "0911/0911.0602_arXiv.txt": {
        "abstract": "A new scheme for numerical integration of the 1D2V relativistic Vlasov-Maxwell system is proposed. Assuming that all particles in a cell of the phase space move with the same velocity as that of the particle located at the center of the cell at the beginning of each time step, we successfully integrate the system with no artificial loss of particles. Furthermore, splitting the equations into advection and interaction parts, the method conserves the sum of the kinetic energy of particles and the electromagnetic energy. Three test problems, the gyration of particles, the Weibel instability, and the wakefield acceleration, are solved by using our scheme. We confirm that our scheme can reproduce analytical results of the problems. Though we deal with the 1D2V relativistic Vlasov-Maxwell system, our method can be applied to the 2D3V and 3D3V cases. ",
        "introduction": "In a wide range of plasma processes operating in laboratories or astrophysical phenomena, interactions between relativistic particles and electromagnetic fields play vital roles. For instance, recent laser experiments revealed that a high intensity laser can accelerate particles to ultra-relativistic speed (see, e.g., \\citet{e96}). Non-thermal components found in spectra of active astrophysical objects, e.g., supernova remnants, gamma-ray bursts, and jets from active galactic nuclei, are interpreted as synchrotron radiation emitted by charged, relativistic particles gyrating about magnetic field lines.  There are two approaches in order to model such plasmas. One is the fluid approach based on the relativistic magnetohydrodynamics(RMHD) and the other is the kinetic approach based on the Boltzmann equation coupled with the Maxwell equations. Because the fluid approach implicitly assumes that the distribution of particles in the momentum space is the Maxwell-J\\\"uttner distribution, which is the relativistic extension of the classical Maxwell-Boltzmann distribution, the kinetic approach is indispensable for dealing with the momentum distribution deviating from the thermal equilibrium. Especially, dilute plasmas in which collisions between particles composing the plasmas are absent, often called collisionless plasmas, are known to be modeled by the so-called Vlasov-Maxwell system \\cite{s94}.  At present, the most reliable and reasonable method to simulate dynamical behaviors of collisionless plasmas is the particle-in-cell (PIC) simulation (see, e.g., \\citet{BL91}), which calculates the orbits of charged particles by solving the equation of motion and the configuration of electromagnetic fields by solving the Maxwell equations. In this method, the momentum distribution of plasmas is approximated by an ensemble of the momentum of each particle placed in the physical space. Although the number of particles in virtual plasmas produced by a PIC simulation is much smaller than that in real plasmas, it is known that behaviors of plasmas are well reproduced by the method. Nevertheless, we cannot avoid significant numerical noises due to the shortage of particles when we focus on the high momentum tail of the distribution function of plasma particles.  On the other hand, the direct numerical integration of the Vlasov-Maxwell system (referred to as \"Vlasov simulation\"), which discretizes the momentum space as well as the physical space, does not suffer from such noises. Therefore, some methods to perform the Vlasov simulation have been developed \\cite{ck76,f99,ny99,mcct02}. While Vlasov simulations require higher computational performance than PIC simulations do, recent developments of computational technology allow us to study plasma processes by using Vlasov simulations. For example, \\citet{mcct02} presented a scheme to numerically integrate the 2D3V Vlasov-Maxwell system (the term \"2D3V\" means the two dimensional space associated with the three dimensional velocity space) and demonstrated that the scheme could simulate the Weibel instability with high accuracy. \\citet{vvm05} provided a scheme for the integration of the electrostatic 1D2V Vlasov-Poisson system in a uniform magnetic field. They adopted the polar coordinates in the velocity space, which allows them to perform simulations with a good energy conservation. The scheme presented by \\citet{vtchm07} integrates the Vlasov-Maxwell system in the hybrid approximation, i.e., they solve the 2D3V electromagnetic Vlasov equation for ions and fluid equations for electrons, based on the current advance method introduced by \\citet{matthews}. \\citet{SS09} investigated non-linear behavior of the Weibel instability in detail by using a scheme similar to that of \\citet{mcct02}. \\citet{SG06} performed a series of simulations for the magnetic reconnection and confirmed that their results are consistent with those of some PIC simulations. However, the attempts stated above treated only non-relativistic plasmas. Investigations into the numerical integration of the relativistic Vlasov-Maxwell system are still rare. Although \\citet{blgsb08} presented a scheme for the 1D2V relativistic Vlasov-Maxwell system, they assumed that particles have no dispersion in the transverse momentum space. Furthermore, there exists another problem that the mass and energy conservations are not always guaranteed unlike PIC simulations.  In this paper, we propose a new conservative scheme for the numerical integration of the 1D2V relativistic Vlasov-Maxwell system that allows particles to have dispersions in the momentum space. The scheme is based on the semi-Lagrangian approach, which is extensively used to solve the Vlasov-Maxwell system \\cite{blgsb08,s96,bs03}. In Sec. \\ref{5:formulation}, we introduce the 1D2V relativistic Vlasov-Maxwell system and some characteristic scales, and then transform the equations for convenience of the subsequent sections. Sec. \\ref{5:scheme} describes the method for the numerical integration of the system. In Sec. \\ref{5:test}, we calculate three test problems using the scheme proposed in Sec. \\ref{5:scheme}, the gyration of particles, the Weibel instability, and the wakefield acceleration. We conclude this paper in Sec. \\ref{5:conclusion}. ",
        "conclusions": "} In this paper, we have proposed a new conservative scheme for numerical integration of the relativistic Vlasov-Maxwell system and performed three test problems, the gyration of particles, the Weibel instability, and the wakefield acceleration. Adopting semi-Lagrange method, we succeed in developing a scheme that conserves the number of particles and the sum of the energies of particles and electromagnetic fields. Since the previous scheme \\cite{blgsb08} solving the relativistic Vlasov-Maxwell system do not treat the dispersion of the lateral momentum of particles, our scheme is the first one that can treat the dispersion correctly. Results of the simulations clearly indicate that our method succeeds in reproducing detailed behaviors of the distribution functions in the phase space. Especially, the tail of the distributions where only a tiny fraction of particles reside seems to be solved with considerably high accuracy, while PIC simulations would suffer from large statistical error there.  As we noted above, Vlasov simulations generally require more computational resources than PIC simulations do. Furthermore, as previous works \\cite{mcct02,cmvv07} have investigated, Vlasov simulations suffer from so-called \"filamentation problem\". Ref.\\cite{cmvv07} studied wave-particle interactions of a plasma approaching to an equilibrium state using PIC simulation and showed that the equilibrium is realized through a phase mixing accompanied by formation of filamentary structures in the phase space. In Vlasov description, particles composing a plasma are treated as a continuous medium, which means that Vlasov equation can not take account of essential discreteness of plasma. As a result, artificial entropy may arise when a structure with the characteristic scale smaller than the mesh size is generated in the phase space. Ref.\\cite{cmvv07} argued that this artificial entropy prevents the plasma treated by the Vlasov simulation from following the correct path toward the statistical equilibrium. Therefore, we need careful studies of long-term evolutions of plasmas by using Vlasov simulations.  Nevertheless, Vlasov simulations provide us detailed dynamics of plasmas in the phase space. Though we deal with the 1D2V relativistic Vlasov-Maxwell system, our method can be applied to the 2D3V and 3D3V cases Although our scheme proposed in this paper suffer from numerical diffusion, there is a plenty room for improvement. For example, in order to integrate advection part of the Vlasov equation, we can use the orbit of the particle located at each vertex of a cell. In other words, taking account for the deformation of the cell at each time step improves the accuracy of the scheme. In theoretical investigations into complex behaviors of collisionless plasmas, Vlasov simulations must be a attractive tool to compensate defects of PIC simulations."
    },
    "0911/0911.2431_arXiv.txt": {
        "abstract": "{ The analysis of flare start-times confirms the periods found years ago (Aschenbach et al., 2004) in the near-infrared and X-ray light-curves related to the Sgr A* black hole. The assignment of the frequencies found to radial and vertical epicyclic frequencies $\\nu\\sb{\\rm r}$, $\\nu\\sb{\\rm v}$, respectively, as well as to the Kepler orbital frequency $\\nu\\sb{\\rm K}$ reveals resonances of $\\nu\\sb{\\rm v}$/$\\nu\\sb{\\rm r}$=7:2, and $\\nu\\sb{\\rm K}$/$\\nu\\sb{\\rm v}$=3:1. The highest observed frequency of 10~mHz  is identified as the Kepler frequency corrected by the rotational frame-dragging frequency, as expected from the  Lense-Thirring effect. These frequency assignments conclude a black hole mass of M~=~(4.10--4.34)$\\times$10$\\sp{6} M_\\odot$ and a spin of a~=~(0.99473--0.99561). ",
        "introduction": "The center of our Galaxy is suspected to contain a supermassive black hole. Observationally this conjecture is based on the analysis of motions and orbit sizes of stars moving around some central mass (Sch\\\"odel et al. 2002; Ghez et al. 2005; Gillessen et al. 2009). These observations suggest the presence of a fairly small volume containing a mass of several million solar masses. Mass and spin of the black hole have also been estimated from quasi-periodic oscillations (Abramowicz et al. 2004; Aschenbach et al. 2004; Aschenbach 2004) which have been suggested to be present in the near-infrared (Genzel et al. 2003) and X-ray (Aschenbach et al. 2004) light-curves of Sgr A*.  These candidate periods are suspect because of their enhanced power spectral density in at least three Sgr A* observations, which include the October 26, 2000 ({\\it{Chandra}}, Baganoff et al. 2001), the October 3, 2002 ({\\it{XMM-Newton}}, Porquet et al. 2004) and the near-infrared (NIR) observations of June 16, 2003 (Genzel et al. 2003). These observations were particularly interesting because they showed fairly large outbursts of Sgr A*. I have summarized the results, and I have proposed candidate period-like structures, centered around 110~s, 219~s and 1173~s (Aschenbach et al. 2004). Later measurements, both in the NIR and the X-ray band, were not conclusive on these periods. In most of the cases there was no indication of a period at all, but when a time structured signal was suggested by the data, it happened to be close to the candidate period of 1100~s. This includes a  1330~s period (B\\'elanger et al. 2006), and a $\\sim$1500~s period (Meyer et al. 2008). However, each one of these observations, taken as a single event, was not and cannot be considered to be of that significance which would justify a statistically unambiguous claim of the detection of a period. But the fact that these coincidences among the light-curve observations exist as far as the putative periods are concerned, is encouraging further study.  The indication of more than one period in the light-curve data suggests the possibility that the light-curve is not dominated by just one period but that the signal is actually modulated by one ore even more frequencies. So far, the analyses had to deal with light-curves which are usually subject to a large amount of noise (be it white, or red or even pink). Therefore I took the opportunity to analyse the starting times of a sequence of X-ray flares which were observed with {\\it{XMM-Newton}} between 31 March and 5 April 2007, which in principle provide an independent access to periodic patterns. ",
        "conclusions": "The periods indicated in earlier observations of the near-infrared and X-ray light-curves from the accretion disk around the Sgr A* black hole are confirmed by the analysis of the flare start-times. The frequencies found imply resonances of 3:1 and 7:2 suggesting that radial and vertical wave oscillations as well as the Kepler frequency are involved. Taking into account the rotational frame-dragging, which should anyhow be included in such an analysis (the Lense-Thirring effect, the Cunningham-Bardeen approach), the data are consistent with the presence of a supermassive, M~=~(4.10--4.34)$\\times$10$\\sp{6} M_\\odot$, and very rapidly spinning, a~=~(0.99473--0.99561), black hole. The inclusion of rotational frame-dragging, which is one of the key features of General Relativity, is mandatory in this analysis and, of course, interpretation.  An open issue is the ad-hoc assumption that hot-spots are created from pre-existing fluctuations. The analysis presented shows that the spin of the Sgr A* black hole is very close to ${\\it a\\sb{\\rm{c}}}$=0.0.9953, above which the anomalous orbit velocity effect or 'Aschenbach' effect (Stuchl\\'ik et al. 2005) occurs (Aschenbach 2004). At ${\\it a\\sb{\\rm{c}}}$=0.9953, the radius derivative of the orbital velocity $\\partial{v\\sp{(\\Phi)}}/\\partial r = 0$. This condition might cause an instability in the disk, like a shear-instability, which is being amplified by the radial and vertical waves."
    },
    "0911/0911.0420_arXiv.txt": {
        "abstract": "{% Globular clusters (GCs) are tracers of the gravitational potential of their host galaxies. Moreover, their kinematic properties may provide clues for understanding the formation of GC systems and their host galaxies.  We use the largest set of GC velocities obtained so far of any elliptical galaxy to revise and extend the previous investigations (Richtler et al.~2004) of the dynamics of NGC\\,1399, the central dominant galaxy of the nearby Fornax cluster of galaxies. The GC velocities are used to study the kinematics, their relation with population properties, and the dark matter halo of NGC 1399. We have obtained 477 new medium--resolution spectra (of these, 292 are spectra from 265 individual GCs, 241 of which are not in the previous data set). with the VLT FORS\\,2 and Gemini South GMOS multi--object spectrographs. We revise velocities for the old spectra and measure velocities for the new spectra, using the same templates to obtain an homogeneously treated data set. Our entire sample now comprises velocities for almost 700 GCs with projected galactocentric radii between 6 and 100\\,kpc. In addition, we use velocities of GCs at larger distances published by Bergond et al.~(2007).  Combining the kinematic data with wide--field photometric Washington data, we study the kinematics of the metal--poor and metal--rich subpopulations.  We discuss in detail the velocity dispersions of subsamples and perform spherical Jeans modelling.  The most important results are: The red GCs resemble the stellar field population of NGC 1399 in the region of overlap. The blue GCs behave kinematically more erratic.  Both subpopulations are kinematically distinct and do not show a smooth transition. It is not possible to find a common dark halo which reproduces simultaneously the properties of both red and blue GCs.  Some velocities of blue GCs are only to be explained by orbits with very large apogalactic distances, thus indicating a contamination with GCs which belong to the entire Fornax cluster rather than to NGC\\,1399.  Also, stripped GCs from nearby elliptical galaxies, particularly NGC\\,1404, may contaminate the blue sample.  We argue in favour of a scenario in which the majority of the blue cluster population has been accreted during the assembly of the Fornax cluster. The red cluster population shares the dynamical history of the galaxy itself. Therefore we recommend to use a dark halo based on the red GCs alone.  The dark halo which fits best is marginally less massive than the halo quoted by Richtler et al.~(2004). The comparison with X-ray analyses is satisfactory in the inner regions, but without showing evidence for a transition from a galaxy to a cluster halo, as suggested by X-ray work. } ",
        "introduction": "\\subsection{The globular cluster systems of central elliptical galaxies }  Shortly after M87 revealed its rich globular cluster system (GCS) \\citep{baum1955,racineGCLF} it became obvious that bright ellipticals in general host globular clusters in much larger numbers than  spiral galaxies \\citep{harrisracine79}. Moreover, the richest GCSs are found for elliptical galaxies in the centres of galaxy clusters, for which M87 in the Virgo cluster and NGC\\,1399 in the Fornax cluster are the nearest examples. For recent reviews of the field, see \\cite{bs06}, and \\cite{richtler06}.  The galaxy cluster environment may act in different ways to produce these very populous GCSs. Firstly, there is the paradigm of giant elliptical galaxy formation by the merging of disk galaxies (e.g. \\citealt{toomre77}, \\citealt{renzini06}). That early--type galaxies can form through mergers is evident by the identification of merger remnants and many kinematical irregularities in elliptical galaxies (counter-rotating cores, accreted dust and molecular rings).  In fact, starting from the bimodal colour distribution of GCs in some giant ellipticals, \\cite{az92} predicted the efficient formation of GCs in {{spiral--spiral}} mergers, before this was confirmed observationally \\citep{schweizer93}.  In their scenario, the blue clusters are the metal--poor GCSs of the pre--merger components while the red (metal--rich) GCs are formed in the material which has been enriched in the starbursts accompanying the early merger {{ (gaseous merger model).  }}\\par {{However, \\cite{forbes97} pointed out that the large number of metal--poor GCs found around giant ellipticals cannot be explained by the gaseous merger model. These authors proposed a `multi-phase collapse model' in which the blue GCs are created in a pre--galactic phase along with a relatively low number of metal--poor field stars. The majority of field stars, i.e.~the galaxy itself and the red GCs, are then formed from the enriched gas in a secondary star formation epoch.  }} {{ This scenario is also supported by the findings of \\cite{spitler08} who studied an updated sample of 25 galaxies spanning a large range of masses, morphological types and environments. They confirmed that the spirals generally show a lower fraction of GCs normalised to host galaxy stellar mass than massive ellipticals, thus ruling out the possibility that the GCSs of massive ellipticals are formed through major wet mergers.}}  {{Further, Spitler et al.~suggest that the number of GCs per unit halo mass is constant -- thus extending the work by \\cite{blakeslee97} who found that in the case of central cluster galaxies the number of GCs scales with the cluster mass -- pointing towards an early formation of the GCs.}}   \\par {{ In some scenarios (e.g.~\\citealt{cote98,hilker99III,cote02,beasley02}), the accretion of mostly metal--poor GCs is responsible for the richness of the GCSs of central giant ellipticals.}}  \\par But only during the last years it has become evident that accretion may be important even for GCSs of galaxies in relatively low density regions like the Milky Way (e.g.~\\citealt{helmi08}), the most convincing case being the GCs associated with the Sagittarius stream {\\citep{ibata01,bellazzini03}}. Therefore, GC accretion should plausibly be an efficient process for the assembly of a GCS in the central regions of galaxy clusters.   Given this scenario one expects huge dark matter halos around central giant ellipticals, perhaps even the sum of a galaxy-size dark halo and a cluster dark halo (e.g.~\\citealt{ikebe96}).  However, dark matter studies in elliptical galaxies using tracers other than X-rays were long hampered by the lack of suitable dynamical tracers. Due to the rapidly declining surface brightness profiles, measurements of stellar kinematics are confined to the inner regions, just marginally probing the radial distances at which dark matter becomes dominant. One notable exception is the case of NGC\\,6166, the cD galaxy in Abell\\,2199, for which \\cite{kelson02} measured the velocity dispersion profile out to a radius of 60\\,kpc. Another one is the case of NGC\\,2974 where an H{\\textsc{i}} disk traces the mass out to 20\\,kpc \\citep{weijmans08}. \\par Only with the advent of 8m--class telescopes and multi--object spectrographs it has become feasible to study the dynamics of globular cluster systems (GCSs) of galaxies as distant as 20\\,Mpc. Early attempts (\\citealt{hb87,grillmair94,cohen97,minniti98,kisslerpatig98}) were restricted to the very brightest GCs, and even today there are only a handful of galaxies with more than 200 GC velocities measured.  Large ($N_{\\rm{GC}}>200$) samples of GC radial velocities have been published for M87 \\citep{cote01} and NGC\\,4472 \\citep{cote03} in Virgo and \\object{NGC\\,1399} in Fornax (\\citealt{richtler04}, hereafter Paper\\,I).  {{Also the GCS dynamics of the nearby ($\\sim4\\,\\rm{Mpc}$) disturbed galaxy Cen\\,A (NGC\\,5128) has been studied extensively \\citep{peng04gcs,woodley07}, with 340 GC velocities  available to date.}}    Investigating the kinematics and dynamics of GCSs of elliptical galaxies  has a twofold objective.  Firstly, kinematical information together with the population properties of GCs promise to lead to a deeper insight into the formation history of GCSs with their different GC subpopulations. Secondly, GCs can be used as dynamical tracers for the total mass of a galaxy and thus  allow the determination of the dark matter profile, out to large galactocentric distances which are normally inaccessible to studies using the integrated light. These results then can be compared to X--ray studies.  A large number of probes is a prerequisite for the analysis of these dynamically hot systems. Therefore, giant ellipticals, known to possess extremely populous and extended GCSs with thousands of clusters, have been the preferred targets of these studies.    Regarding the formation history of GCSs, a clear picture has not yet emerged.  Adopting the usual bimodal description of a GCS by the distinction between metal--poor and metal--rich GCs, the kinematical properties seem to differ from galaxy to galaxy. For example, in M87 the blue and red GCs do not exhibit a significant difference in their velocity dispersion (\\citealt{cote01}, with a sample size of 280 GCs).  Together with their different surface density profiles, C{\\^o}t{\\'e} et al. concluded that their orbital properties should be different: the metal--poor GCs have preferentially tangential orbits while the metal--rich GCs prefer more radial orbits.  In NGC\\,4472, on the other hand, the metal--rich GCs have a significantly lower velocity dispersion than their blue counterparts (found for a sample 250 GCs, \\citealt{cote03}), and these authors conclude that the cluster system as a whole has an isotropic orbital distribution.  NGC\\,1399, the object of our present study, is a galaxy, which, as a central cluster galaxy, is similar to M87 in many respects (see \\citealt{dirsch03}). Here, the metal--poor GCs show a distinctly higher velocity dispersion than the metal--rich GCs, but more or less in agreement with their different density profiles (with a sample size of $\\sim\\!470$ GCs). A similar behaviour has been found for NGC\\,4636 \\citep{schuberth06}.  In most other studies, the sample sizes are still too small to permit stringent conclusions or the separate treatment of red and blue GCs, but dark matter halos have been found in almost all cases.    \\subsection{The case of NGC 1399}  NGC\\,1399 has long been known to host a very populous globular cluster system (e.g.~\\citealt{dirsch03} and  references therein). With the photometric study by \\cite{bassino06} it became clear that the GCS of NGC\\,1399 extends to about 250\\,kpc, which is comparable to the core radius of the  cluster \\citep{ferguson89}. Accordingly, it has always been an attractive target for studying the dynamics of its GCS. One finds there the largest sample of GC velocities (469) available so far (\\citealt{richtler04} (Paper\\,I), \\citealt{dirsch04}).  It was shown in Paper\\,I  that blue and red GCs are kinematically different, as was expected from their different number density profiles: the red GCs exhibit a smaller velocity dispersion than the blue GCs in accordance with their respective density profiles. Evidence for strong anisotropies has not been found. The radial velocity dispersion profile was found to be constant for red and blue GCs.  However, as we think now, this could have been a consequence of a velocity cut introduced to avoid outliers.  A dark halo of the NFW type under isotropy reproduced the observations satisfactorily.  No rotation was detected apart from a slight signal for the outer blue GCs.  It was shown that some of the extreme radial velocities in conjunction with the derived dark halo were only understandable if they were being caused by orbits with very large apogalactic distances.  In this paper, we extend our investigation of the NGC\\,1399 GCS to larger radii (80 kpc). We simultaneously revise the old velocities/spectra in order to have an homogeneously treated sample. The case of Modified Newtonian Dynamics (MOND) has already been  discussed in \\cite{richtler08}, where it has been shown that MOND still needs additional dark matter of the order of the stellar mass. We do not come back to this issue in the present contribution.  \\par  The GCS of NGC 1399 is very extended. One can trace the blue GCs out to about 250\\,kpc, the red GCs only to 140\\,kpc \\citep{bassino06}. Regarding total numbers, there are only half as many red GCs as there are blue, suggesting that the formation of GCs in mergers is not the dominant mechanism producing a high specific frequency, even if in the central regions of a proto--cluster the merger rate is supposed to be particularly high.   \\par Following Paper\\,I, we assume a distance modulus of $31.40$.  At the distance of 19\\,Mpc, $1\\arcsec{}$ corresponds to 92\\,pc, and $1\\arcmin{}$ corresponds to 5.5\\,kpc. \\par   This paper is organised as follows: In Sect.~\\ref{sect:obs}, we describe the observations and the data reduction.  The velocity data base is presented in Sect.~\\ref{sect:database}. In Sect.~\\ref{sect:sample}, we present the photometric properties and the spatial distribution of our GC velocity sample. The contamination by interlopers is discussed in Sect.~\\ref{sect:interlopers}. The properties of the line--of--sight velocity distribution are studied in Sect.~\\ref{sect:losvd}. In Sect.~\\ref{sect:rot} we test our GC sample for rotation. The line--of--sight velocity dispersion and the higher order moments of the velocity distributions are calculated in Sect.~\\ref{sect:dispersion}. The Jeans modelling and the derived mass models are described in Sects.~\\,\\ref{sect:jeans} and \\ref{sect:massmodels}. The results are discussed and summarised in Sections\\,\\ref{sect:discussion} and \\ref{sect:conclusions}. ",
        "conclusions": "\\label{sect:conclusions} Using the largest data set of globular cluster (GC) radial velocities available to date, we revisit and extend the investigation of the kinematical and dynamical properties of the NGC\\,1399 GC system with respect to Richtler et al.~(\\citeyear{richtler04}, Paper\\,I). We measure about 700 GC radial velocities out to approximately $14.5\\arcmin\\simeq 80\\,\\rm{kpc}$ (in Paper\\,I we reached 40 kpc). To this sample we add 56 GC velocities from Bergond et al. (\\citeyear{bergond07}, B+07) which go as far out as 200 kpc.\\\\ Our main findings are the following: There is no significant rotation signal among the red (metal--rich) subpopulation.  We find rotation around the minor axis for the blue (metal--poor) clusters in a radial interval of $4\\arcmin{} < R < 8\\arcmin{}$ (i.e.~$22\\,\\rm{kpc}\\lesssim R\\lesssim44\\,\\rm{kpc}$), however weak. The blue and red clusters form two {{\\emph{kinematically distinct}} subpopulations rather than showing a continuum of kinematical properties.  The red clusters correspond under various kinematical aspects to the stellar field populations.  Their velocity dispersion declines outwards and their velocity distribution suggests {orbital isotropy}.    The jump to the higher dispersion of the blue clusters occurs at a colour ($C\\!-\\!R = 1.55$), which is also suggested by the morphology of the colour distribution.  The blue GCs, however, show a more complex behaviour. Their velocity dispersion profile is not as smooth as that of the red GCs. There exist very high/low individual velocities, which suggest very large apogalactic distances, as already found in Paper\\,I. These objects seem to belong to an intergalactic cluster population, which {may be made up of GCs stemming from disrupted dwarfs and GCs stripped off neighbouring galaxies during close encounters}.  We performed a Jeans analysis in order to constrain the mass profile (stellar plus dark).  We found it difficult to find a dark halo which simultaneously accounts for the red and the blue clusters, which again argues for a different dynamical history.  The dark NFW--halo which was found to represent the kinematics of the red GCs ($\\beta_{\\mathrm{GCs}}=0$) and the stellar velocity dispersion profile presented by \\cite{saglia00} ($\\beta_{\\mathrm{stars}}=+0.5$), our Model\\,{\\it{a10}}, has the parameters $\\varrho_{\\mathrm{dark}}=0.0088\\,M_\\odot\\,\\rm{pc}^{-3}$ and $r_{\\mathrm{dark}}=34\\,\\rm{kpc}$. The virial mass is $M_{\\textrm{vir}}=8.0\\times10^{12} M_\\odot$. \\\\ Including the velocities from B+07, (Model\\,{\\it{b10}}) yields a slightly more massive halo ($\\varrho_{\\mathrm{dark}}=0.0077\\,M_\\odot\\,\\rm{pc}^{-3}$ and $r_{\\mathrm{dark}}=38\\,\\rm{kpc}$, and a virial mass $M_{\\textrm{vir}}=9.5\\times10^{12} M_\\odot$).   These NFW halos, which were found to represent the observations, are marginally less massive than the one from Paper\\,I.  The total mass profile fits agree reasonably well with the X--ray based mass profiles out to 80\\,kpc. However, our model halos when extrapolated stay significantly below the X--ray masses. Moreover, we do not see the transition from the galaxy halo to the cluster halo claimed to be present in the X--ray data.  We argue that these findings are consistent with a scenario where the red GCs are formed together with the bulk of the field population, most probably in {early} multiple mergers with many progenitors involved. {This is consistent with the NGC\\,1399 GCs being predominantly old \\citep{kissler98,forbes01,kundu05,hempel07}, with ages very similar to the stellar age of NGC\\,1399 ($11.5\\pm2.4$\\,Gyr, \\citealt{trager00}) } \\par A large part of the blue GC population has been acquired by accretion processes, most plausibly through dwarf galaxies during the assembly of the Fornax cluster. Moreover, there should be a significant population of intra-cluster GCs."
    },
    "0911/0911.5205_arXiv.txt": {
        "abstract": "The nearby ($z=0.03015$) cluster of galaxies Abell~2199 was observed by Suzaku in X-rays, with five pointings for $\\sim 20$ ks each. From the XIS data, the temperature and metal abundance profiles were derived out to $\\sim 700$ kpc (0.4 times virial radius). Both these quantities decrease gradually from the center to peripheries by a factor of $\\sim 2$, while the oxygen abundance tends to be flat. The temperature within $12'$ ($\\sim 430$ kpc) is $\\sim 4$ keV, and the 0.5--10 keV X-ray luminosity integrated up to $30'$ is $(2.9 \\pm 0.1) \\times 10^{44}$ erg s$^{-1}$, in agreement with previous XMM-Newton measurements. Above this thermal emission, no significant excess was found either in the XIS range below $\\sim 1$ keV, or in the HXD-PIN range above $\\sim 15$ keV. The 90\\%-confidence upper limit on the emission measure of an assumed 0.2 keV warm gas is (3.7--7.5) $\\times 10^{62}$ cm$^{-3}$ arcmin$^{-2}$,  which is 3.7--7.6 times tighter than the detection reported with XMM-Newton. The 90\\%-confidence upper limit on the 20--80 keV luminosity of any power law component is $1.8 \\times 10^{43}$ erg s$^{-1}$, assuming a photon index of 2.0. Although this upper limit does not reject the possible 2.1$\\sigma$ detection by the BeppoSAX PDS, it is a factor of 2.1 tighter than that of the PDS if both are considered upper limits. The non-detection of the hard excess can be reconciled with the upper limit on diffuse radio emission, without invoking the very low magnetic fields ($< 0.073 \\mu$G) which were suggested previously. ",
        "introduction": "The intra-cluster medium (ICM) in clusters of galaxies, i.e, hot thermal plasmas in collisional ionization equlibria confined within their gravitational potential, constitutes the most dominant form of cosmic baryons that has ever been detected. Although the ICM radiates predominantly optically-thin thermal X-ray emission, some clusters have been reported to exhibit excess signals above the thermal emission at the lowest or highest energy ends of their X-ray spectra. One interpretation of the soft excess is emission from Warm Hot Intergalactic Medium (WHIM), which has been predicted from cosmological simulations (e.g. \\cite{cen-1999}) and is expected to solve the so-called missing baryon problem (e.g \\cite{fukugita-1998}). Another interpretation of the spectral soft and/or hard excess is non-thermal emission from accelerated particles in galaxy clusters (e.g. \\cite{lieu-1999}; \\cite{fusco-1999}): a power-low shaped spectrum can exceed the dominant thermal emission at both sufficiently high and low energies. This provides a possibility that galaxy clusters are giant accelerators in the universe. Thus, the search for soft and hard excess signals in clusters forms a very important research subject. However, the instrumental sensitivity in these photon energies has been insufficient, and hence the existence of these excess components has remained controversial (e.g. \\cite{bregman-2006}; \\cite{nevalainen-2007}; \\cite{rossetti-2004}; \\cite{fusco-2007}).   Among those galaxy clusters from which the excess emission has been reported, Abell~2199 is particularly interesting, because it has been suspected to exhibit both the soft and hard excess components. In fact, Kaastra et al. (1999) reported the detection of both, based on the BeppoSAX, Extreme Ultraviolet Explorer, and ROSAT observations. The two excess components were simultaneously expressed by a single power law with a photon index $\\sim 1.8$ which is superposed on the thermal emission spectrum, and interpreted as inverse Compton (IC) photons produced when the cosmic microwave background (CMB) photons are scattered up by relativistic electrons in the cluster. However, as argued by \\citet{kempner-2000}, the magnetic field in Abell~2199 would have to be unusually weak ($< 0.073\\; \\mu$G) in order for the relativistic electrons not to produce synchrotron emission beyond measured upper limits of diffuse radio emission. As an alternative interpretation, \\citet{kempner-2000} proposed non-thermal bremsstrahlung by supra-thermal electrons which are being accelerated in the cluster.  In the era of XMM-Newton, the possibility of thermal emission from a warm gas also became another popular idea to explain the soft excess in several clusters. Kaastra et al. (2004) reported a redshifted O~VII emission line from a 0.2 keV warm gas in Abell~2199, along with Coma, Abell~1795, Sersic~159-03, MKW~3s, Abell~2052, and Abell~3112, and hence ascribed their soft excess to emission from WHIM. At the same time, Chandra has made progress on the ICM physics of Abell~2199 especially in the central region, such as an asymmetric temperature distribution in the direction perpendicular to the jets within $30''$ of the center (Kawano et al. 2003), and a weak isothermal shock associated with the central active galactic nucleus (Sanders \\& Fabian 2006).   In order to examine the soft and hard excess phenomena of Abell~2199 with a better sensitivity, we observed the object with Suzaku (\\cite{mitsuda-2007}). The excellent low-energy capability of XIS-BI (back illuminated CCD, \\cite{koyama-2007}) and the low background of the silicon PIN detectors in the HXD (\\cite{takahashi-2007}; \\cite{kokubun-2007}) have provided the best data in the soft and hard X-ray bands, respectively. As a result, we have obtained tight upper limits on both excess components, and strengthened the dominance of the thermal emission from this cluster.   In the present paper, quantities quoted from the literature are rescaled to the Hubble constant of $H_0 = 71\\; \\rm{km\\; s^{-1}\\; Mpc^{-1}}$, and errors are given at 90 \\% confidence range unless otherwise stated. ",
        "conclusions": "}  Using the five pointing observations with Suzaku, we measured the temperature and metal abundances of Abell~2199 out to $\\sim 700 $ kpc (0.4 times virial radius), and searched the XIS and HXD-PIN data for soft and hard excess emissions, respectively, both above the thermal component. We however found no significant excess emissions in either energy bands, and derived their upper limits. Over a broad energy band of $\\sim 0.4$ keV to $\\sim 30$ keV, the emission is dominated by the thermal components with temperatures of a few keV.  The upper limits on soft excess, expressed as the emission measure per unit solid angle from a 0.2 keV warm gas (table~\\ref{table:warm-upper}), are more stringent, typically by more than a factor of 3, than the reported positive detection by XMM-Newton (\\cite{kaastra-2004}), $(28 \\pm 13) \\times 10^{62}$ cm$^{-3}$ arcmin$^{-2}$. One likely cause of this inconsistency between Suzaku and XMM-Newton is confusion with solar wind charge exchange (SWCX) emission, which has been observed with Chandra (\\cite{wargelin-2004}), XMM-Newton (\\cite{snowden-2004}), and Suzaku (\\cite{fujimoto-2007}). In fact, proton flux near the Earth, observed with WIND-SWE\\footnote{http://web.mit.edu/space/www/wind.html}, enhanced by a factor of 3--4 for a part of the XMM-Newton observatons compared with that of Suzaku. The relevant proton light curves are shown in Appendix~\\ref{appen:proton-flux}. Similarly, in Suzaku observations of outskirt regions of the Coma cluster, Takei et al. (2008) found no significant excess O~VII or O~VIII emissions which were previously reported with XMM-Newton (\\cite{finoguenov-2003}), and concluded that the excess emission is likely due to an enhancement in the SWCX emission during the XMM-Newton observations.   The Suzaku data have constrained the 20--80 keV luminosity of any excess hard emission to be $< 1.8 \\times 10^{43}$ erg s$^{-1}$, assuming that it has a power-law spectrum of photon index 2.0 and employing the most conservative estimates of the CXB. Using the BeppoSAX PDS, \\citet{nevalainen-2004} detected non-thermal emission  from Abell~2199 at 2.1$\\sigma$ level, with 20--80 keV luminosity of $(1.7 \\pm 1.3) \\times 10^{43}$ erg s$^{-1}$ (90\\%-confidence errors) assuming a power-law spectrum of photon index 2.0. This detection by the PDS is not rejected by our result. However, the BeppoSAX results, with a significance of only 2.1$\\sigma$, would be considered as an upper limit, like the present result. In this case, our upper limit is a factor of 1.7 more stringent than that of the PDS. From a calibration analysis using the Crab Nebula, the 20--80 keV flux derived from the PDS is known to be systematically lower by 21\\% than that derived from HXD-PIN (\\cite{nakazawa-2009}). If we take into account this systematic effect, the difference in the upper limits between HXD-PIN and PDS increases to a factor of 2.1.  Considering the difference between the FOVs of HXD-PIN ($34'$ in FWHM) and PDS ($1.3^{\\circ}$ in FWHM), the claimed detection by the PDS becomes consistent with our upper limit if the excess hard X-ray emission is much more extended than the HXD-PIN FOV. However, the total FOV of HXD-PIN in our pointing observations ($\\sim 43'= 1.5$ Mpc radius in FWHM) mostly cover the whole cluster region, $\\sim 0.9$ times the Virial radius of Abell~2199  ($\\sim 1.7$ Mpc, see \\S~\\ref{subsection:bgd:sys}). Therefore, to reconcile the two results, we would have to assume that the emission is spatially distributed much beyond the virial radius.   According to Kempner \\& Sarazin (2000), a radio synchrotron flux density at 327 MHz, $S_{\\nu}$, expected from Abell~2199, is related to the 10--100 keV IC X-ray flux, $S_{\\rm HXR}$, as \\begin{equation} S_{\\nu} = 234 \\;\\left( \\frac{S_{\\rm HXR}}{10^{-11}\\; {\\rm erg}\\; {\\rm cm}^{-2}\\; {\\rm s}^{-1}} \\right) \\left( \\frac{B}{1\\;\\mu {\\rm G}} \\right)^{1.81} \\left( \\frac{\\nu}{327\\; {\\rm MHz}} \\right)^{-0.81} {\\rm Jy}, \\label{eq:} \\end{equation} assuming the photon index of the hard X-ray emission to be 1.81. Here, $B$ is the cluster magnetic field, and $\\nu$ is the observed frequency. The upper limit they quoted on the diffuse radio flux, $S_{\\nu} < 3.25$ Jy at 327 MHz, and the hard X-ray emission detected by BeppoSAX, required a very weak magnetic field ($< 0.073\\; \\mu$G). However, given our non-detection of such an excess hard X-ray flux, this is no longer the case. When the HXD-PIN data are used to constrain a power-law component with a photon index of 1.81 (instead of 2.0), its 10--100 keV luminosity becomes $< 1.5 \\times 10^{-11}$ erg cm$^{-2}$ s$^{-1}$. As a result, effectively any strength of magnetic field  is allowed as shown in figure~\\ref{fig:mag-limit}.     Regardless of the magnetic field strength, the present result constrain the number of relativistic electrons in the system, so that their IC emission should not exceed the present hard X-ray upper limit. If we assume a power-law spectrum of synchrotron emission with a photon index 2.0, the corresponding electron number spectrum has a power-law index 3.0, and is written as \\citep{petrosian-2006} \\begin{equation} N(\\gamma) = 2 N_{\\rm tot} \\gamma^{2}_{\\rm min} \\gamma^{-3} \\; \\; \\;  (\\gamma > \\gamma_{\\rm min}), \\label{eq:} \\end{equation} where $\\gamma$ is the Lorentz factor, $N_{\\rm tot}$ is the total number of electrons, and $\\gamma_{\\rm min}$ is a lower-cutoff in $\\gamma$ which is conservatively estimated to be $\\sim 10^3$ (\\cite{petrosian-2006}). These relativistic electrons produce IC emission with a 20--80 keV luminosity of \\begin{equation} L_{\\rm HXR} = 1.0 \\times 10^{45} \\left( \\frac{N_{\\rm tot}}{10^{65}}\\right) (1+z)^4 \\; {\\rm erg}\\;  {\\rm s}^{-1}, \\label{eqn:hxr-lumi} \\end{equation} where $z$ is the redshift. From equation~\\ref{eqn:hxr-lumi} and our upper limit, we obtain $N_{\\rm tot} <  1.6 \\times 10^{63}$. Then, the integrated electron kinetic energy $K_{\\rm e}$ for $\\gamma > \\gamma_{\\rm min}$ in Abell~2199 becomes \\begin{equation} K_{\\rm e} = N_{\\rm tot} m_{\\rm e} c^2 (2 \\gamma_{\\rm min} - 1) < 2.6 \\times 10^{60} \\left( \\frac{\\gamma_{\\rm min}}{10^{3}}\\right) \\; {\\rm erg}, \\end{equation} where $m_{\\rm e}$ is the electron mass. Since electrons with $10^3 \\lesssim  \\gamma \\lesssim 10^4$, which contribute to the $20$--$80$ keV IC emission, have a liftime of $\\sim 10^{8-9} $ yr in typical ICM conditions (\\cite{petrosian-2001}), the acceleration luminosity in Abell~2199 can be constrained as $< 8.4 \\times 10^{44} \\; (\\gamma_{\\rm min}/10^3)  $ erg s$^{-1}$. This upper limit is rather loose, and is three times higher than the 0.5--10 keV luminosity of the thermal emission. If we assume that the magnetic field of Abell~2199 is $\\sim 1\\; \\mu$G as suggested by Faraday rotation measurements (\\cite{ge-1994}), the radio upper limit gives the dominant constraint on the relativistic electrons, with the predicted  $S_{\\rm HXR}$ two orders of magnitude below the present upper limit (figure~\\ref{fig:mag-limit}). In this case, the constraint of acceleration luminosity becomes $\\lesssim 10^{43} \\; (\\gamma_{\\rm min}/10^3)  $ erg s$^{-1}$ which is about an order of magnitude lower than the 0.5--10 keV luminosity.   The total mass of Abell~2199 is $1.3 \\times 10^{14} M_{\\odot}$ when integrated to 0.8 Mpc (\\cite{fukazawa-2004}). If Abell~2199 was formed via a major merger between two clusters with similar masses $M$ ($\\sim 7 \\times 10^{13} M_{\\odot}$ each) at a relative speed of $v \\sim 3000$ km s$^{-1}$, the kinetic energy deposited onto the merged system becomes \\begin{equation} E_{\\rm merger} \\sim \\frac{1}{2}Mv^2 =  6 \\times 10^{63} \\left( \\frac{ M }{ 7 \\times 10^{13} \\; M_{\\odot}} \\right) \\left( \\frac{ v }{ 3000\\; {\\rm km}\\; {\\rm s}^{-1}} \\right)^2\\; {\\rm erg}. \\end{equation} Since shocks produced in cluster mergers are thought to last typically $\\tau \\sim 10^9$ yr (\\cite{takizawa-2000}), the energy input rate to non-thermal electrons becomes $\\dot{E}_{\\rm in} \\sim 1 \\times 10^{46} \\; (f/0.05)$ erg s$^{-1}$, where $f$ is the fraction of energies given to them. Values of $\\dot{E}_{\\rm in} = 10^{46-47}$ erg s$^{-1}$ are also expected from theoretical models of merging clusters such as Coma (e.g. \\cite{takizawa-2004}). Then, our upper limit on the acceleration luminosity is an order of magnitude lower than  $\\dot{E}_{\\rm in}$,  assuming $f \\sim 0.05$. We hence infer that Abell~2199 have been free from major merger events for more than $\\sim 10^9$ yr, which is a typical time scale on which the IC emission decays (\\cite{takizawa-2004}).    MK acknowledges support from a grant based on the Special Postdoctoral Researchers Program of RIKEN. The present work is supported in part by the Grant-in-Aid for Scientific Research (S), No. 18104004.   \\appendix"
    },
    "0911/0911.2049_arXiv.txt": {
        "abstract": "\\fontsize{10}{10.6}\\selectfont  The HH 1--2 region in the L1641 molecular cloud was observed in the near-IR $J$, $H$, and $K_s$ bands, and imaging polarimetry was performed. Seventy six point-like sources were detected in all three bands. The near-IR polarizations of these sources seem to be caused mostly by the dichroic extinction. Using a color-color diagram, reddened sources with little infrared excess were selected to trace the magnetic field structure of the molecular cloud. The mean polarization position angle of these sources is about 111\\arcdeg, which is interpreted as the projected direction of the magnetic field in the observed region of the cloud. The distribution of the polarization angle has a dispersion of about 11\\arcdeg, which is smaller than what was measured in previous studies. This small dispersion gives a rough estimate of the strength of the magnetic field to be about 130 $\\mu$G and suggests that the global magnetic field in this region is quite regular and straight. In contrast, the outflows driven by young stellar objects in this region seem to have no preferred orientation. This discrepancy suggests that the magnetic field in the L1641 molecular cloud does not dictate the orientation of the protostars forming inside. ",
        "introduction": "Magnetic fields play a crucial role in various astrophysical processes, including the evolution of interstellar molecular clouds and star formation (Shu et al. 1987; Bergin \\& Tafalla 2007; McKee \\& Ostriker 2007). One of the problems related to star formation concerns the competition between magnetic and turbulent forces (Mac Low \\& Klessen 2004). The magnetic field direction can be measured by observing the dichroic polarization of background stars in the optical and near-IR bands and/or the linearly polarized emission from the dust grains in the mid-IR and far-IR bands (Davis \\& Greenstein 1951; Matthews \\& Wilson 2000). The large-scale alignment of dust grains with the magnetic field is known to be the cause of the dichroic extinction and the interstellar polarization seen in the direction of background sources. Because of the low extinction, near-IR imaging polarimetry is particularly useful in tracing the dichroic polarization of background stars and embedded sources seen through dense clouds (Vrba et al. 1976; Wilking et al. 1979; Tamura et al. 1987; Kandori et al. 2007). Since both dichroic extinction and scattering processes can contribute to the polarization of embedded sources, multiwavelength polarimetry can be useful in discriminating between the two mechanisms (Casali 1995).  The L1641 cloud is one of the nearest giant molecular clouds and is a site of active star formation (Kutner et al. 1977; Maddalena et al. 1986; Strom et al. 1989; Morgan \\& Bally 1991; Sakamoto et al. 1997; Zavagno et al. 1997). The role of magnetic field in the star formation activity of L1641 is complicated. With visual polarimetry of background stars, Vrba et al. (1988) found that the dispersion in position angles is large (33\\arcdeg) and suggested that the role of magnetic field in the global scale is only incidental. However, they also found that the outflows in L1641 tend to be parallel to the field direction and suggested that the role of magnetic field is important in the local scale. In contrast, Casali (1995) performed near-IR polarimetry of young stellar objects (YSOs) in L1641 and found that the alignment of polarization vectors is poor, which suggests that the magnetic field was not dominant in the collapse dynamics. To understand the situation better, it was suggested that more extensive polarimetry in the vicinity of each outflow and YSO is necessary (Vrba et al. 1988).  One of the well studied parts of the L1641 cloud is the region around the reflection nebula NGC 1999 and the Herbig-Haro objects HH 1--2. This region contains several YSOs and outflows (Herbig 1951; Haro 1952; Warren-Smith et al. 1980; Strom et al. 1989; Corcoran \\& Ray 1995; Choi \\& Zhou 1997; Rodr{\\'\\i}guez et al. 2000). The magnetic field structure in the HH 1--2 region has been studied based on optical polarizations of point sources (Strom et al. 1985; Warren-Smith \\& Scarrott 1999, hereafter WS). They found that the local magnetic field is directed roughly along the axis of the HH 1--2 outflow, but the number of detectable stars was too small because of the large obscuration in this region.  In this paper, we present a wide-field near-IR polarimetry of the HH 1--2 region. In Section 2 we describe the observations and data reduction. In Section 3 we present the results of the polarimetry of point-like sources. In Section 4 we discuss the magnetic field structure and the star-forming activity in the HH 1--2 region. A summary is given in Section 5. ",
        "conclusions": ""
    },
    "0911/0911.2755_arXiv.txt": {
        "abstract": "We present  deep Hubble Space Telescope (HST) NICMOS 2 \\filter{F160W} band observations  of the central 56\\arcsec$\\times$57\\arcsec\\ (14pc$\\times$14.25pc) region around R136  in  the  starburst  cluster 30 Dor (NGC 2070) located in the Large Magellanic Cloud. Our aim is to derive the stellar Initial Mass Function (IMF)  down to $\\sim$1 M$_\\odot$ in order to test whether the IMF in a massive metal--poor cluster is similar to that observed in nearby young clusters and the field  in our Galaxy. We estimate the mean age of the cluster  to be 3 Myr by combining our \\filter{F160W} photometry with previously obtained HST WFPC2 optical \\filter{F555W} and \\filter{F814W}  band photometry  and  comparing the stellar locus in the color--magnitude diagram with  main sequence and pre--main sequence isochrones. The color--magnitude diagrams show the presence of differential extinction and  possibly  an  age spread of a few Myr. We convert the magnitudes into masses adopting both a single mean age of 3 Myr isochrone and  a constant star formation history from 2 to 4 Myr. We derive the  IMF after correcting for incompleteness due to crowding. The faintest stars detected have a mass of 0.5 M$_\\odot$ and the data are more than  50\\%\\ complete outside a radius of  5 pc down to a mass limit of 1.1 M$_\\odot$ for 3 Myr old objects. We find an IMF  of $\\frac{dN}{d\\log M }\\propto M^{-1.20\\pm 0.2}$ over the mass range 1.1--20 M$_\\odot$ only slightly shallower than a Salpeter IMF\\@. In particular, we find no strong evidence for a flattening of the IMF down to 1.1 M$_\\odot$ at a distance of 5 pc from the center, in contrast to a flattening at 2 M$_\\odot$ at a radius of 2 pc, reported in a previous optical HST study. flattening at 2 M$_\\odot$ at a radius of  2 pc previously found. We examine several possible reasons for the different results including the possible presence of mass segregation and the effects of differential extinction, particularly for the pre--main sequence sources. If the IMF determined here applies to the whole cluster, the cluster would be massive enough to remain bound and evolve into a relatively low--mass globular cluster. ",
        "introduction": "The shape of the stellar Initial Mass Function (IMF) and whether it is universal or not are key issues  in astrophysics. For clusters within 2 kpc, there is no  compelling evidence  for variations in the stellar IMF \\citep[e.g.][]{meyer00,kro02,chabrier_conf} or the brown dwarf IMF \\citep[e.g.][]{andersen08}. However, these clusters only span a limited range in total cluster mass ($10^2-10^3$ M$_\\odot$) and all  have a metallicity similar to the solar value. Thus, we are forced to observe more extreme regions of star formation in search of  variations in the IMF as a function of environment. It has been suggested that the shape of the IMF and in particular the characteristic mass where the IMF flattens from a Salpeter power--law could depend on the  metallicity  in the  molecular cloud out of which the stars are formed. \\citet{low}, \\citet{larson}, and \\citet{omukai} suggest that a lower metallicity results in higher temperatures in the molecular cloud which would increase the Jeans mass. This would in turn result in a top heavy IMF relative to the solar metallicity IMF\\@.   The closest place with massive metal--poor young star clusters is the Large Magellanic Cloud (LMC). The metallicity is only $\\frac{1}{3}-\\frac{1}{2}$  the solar value \\citep{smith}  and   star clusters can be studied in some detail despite a distance of $\\sim$50 kpc \\citep{westerlund}. Of particular interest is the 30 Dor cluster which is powering the most luminous HII region in the Local Group \\citep{kennicutt}. The cluster has a mass of at least 2.2$\\times 10^4$ M$_\\odot$ within a  radius of 4.7 pc \\citep{hunter95} and is a relatively low-mass analog to the more distant starburst clusters. R136 lies at the center of the 30 Dor cluster and has long commanded significant attention: Once thought to be a single $\\sim$1000 M$_\\odot$ star \\citep{cassinelli}, the region is now known to host  numerous O  stars \\citep{melnick85,weigelt,pehlemann,campbell}.   The  whole   30 Dor region, with a size of 200 pc, appears to have an age spread of $\\sim$20 Myr \\citep{mcgregor,selman2} with stars still forming \\citep{rubio,maercker}. R136 appears to have a much smaller age spread of at most a   few Myr \\citep{melnick85,brandl96,masseyhunter}. An  age of  2 Myr or less  is inferred from   spectroscopy of the O stars in the very cluster center \\citep{masseyhunter}, whereas the  intermediate mass population is thought  to be  $\\sim$3--4 Myr old \\citep{hunter95}.  \\citet{masseyhunter}  obtained   HST spectroscopy  of the 65 bluest and most luminous sources within 17\\arcsec\\ of the cluster center. They derived the IMF over the mass range 15--120 M$_\\odot$ and found it  to be well approximated by a power--law $\\frac{dN}{dlog M}\\propto M^{\\Gamma}$ with a slope of $\\Gamma=-1.3\\pm0.1$, consistent with a Salpeter slope IMF\\@ \\citep{salpeter}. \\citet{hunter95,hunter96} obtained  \\filter{F555W} (\\filter{V}) and \\filter{F814W} (\\filter{i}) band optical photometry utilizing  HST/WFPC2 in order to resolve the cluster's intermediate mass stellar population. The IMF  derived for  different annuli out to a radius of 4.7 pc was found to be in the range $-1.46 < \\Gamma < -1.17$ for the mass range 2.8--15 M$_\\odot$, again consistent with a Salpeter slope IMF\\@. \\citet{masseyhunter} combined their results for the high--mass IMF with the results from \\citet{hunter95,hunter96} in order to constrain the IMF from 2.8 M$_\\odot$ up to 120 M$_\\odot$. Comparing the number of high--mass stars predicted by the intermediate--mass IMF from \\citet{hunter96}, they found the number of massive stars was consistent with a single power--law IMF with a Salpeter slope, i.e. $\\Gamma=-1.35$.    Combining the two data sets used in \\citet{hunter95,hunter96}, \\citet{sirianni} derived the IMF between 1.35 M$_\\odot$ and 6.5 M$_\\odot$, extending the IMF determination  into the mass range where the stars are still in their pre--main sequence phase. The IMF was derived in a box with the dimensions $\\sim$30\\farcs4$\\times$26\\farcs8\\arcsec\\ (7.6pc$\\times$6.7pc), but excluding the inner most 13\\farcs6$\\times$8.6\\arcsec\\ (3.5pc$\\times$2.2pc). Again, a Salpeter slope was found down to  2 M$_\\odot$, but the IMF was found to be flatter than Salpeter, $\\Gamma=-0.27\\pm0.08$, between 1.35 M$_\\odot$ and 2 M$_\\odot$, suggesting the characteristic mass is higher in this massive, metal--poor cluster than $\\sim$ 0.5 M$_\\odot$ as found  in the Galactic field \\citep[][]{kro02}.   The foreground (A$_\\mathrm{V}=0.7$ mag) and  differential extinction (A$_\\mathrm{V}\\sim0-2$ mag) within the cluster \\citep{brandl96} makes it desirable to observe the cluster in the infrared, for example the \\filter{H} band where the extinction is less than 20\\% that of the \\filter{V} band. % In addition, pre--main sequence stars are often associated with circumstellar disks and outflows which will introduce additional extinction for the clusters low--mass content.  We have observed R136 with HST/NICMOS Camera 2 through the \\filter{F160W} band, which is  similar to a ground--based \\filter{H} filter. The observations were aimed at being sensitive to  objects  below 1 M$_\\odot$ for a stellar population with an age of 3 Myr. Preliminary results have previously been presented in \\citet{HZ1,HZ2}, and \\citet{MA}.  The paper is structured as follows. The data and their reduction is described in Section 2. Section 3 shows the results for the \\filter{F160W} band imaging. The IMF is derived in Section 4 and compared  with the IMF  derived by \\citet{sirianni}. We point out several plausible reasons for the different results in the optical and near--infrared, including mass segregation, and differential extinction. Finally, our conclusions  are presented in Section 5. ",
        "conclusions": " \\begin{itemize} \\item{From the color--magnitude diagram obtained by combining our photometry with previously published HST/WFPC 2 \\filter{F555W} data  we constrain the age of the lower--mass stellar content in the cluster to be 2--4 Myr, consistent with previous estimates. We derive individual masses for the objects detected adopting a 3 Myr isochrone.} \\item{We have detected stars in  the cluster down  to  0.5 M$_\\odot$  at $r > 5$ pc, assuming an age of 3 Myr.} \\item{The derived IMF is consistent with a Salpeter slope IMF with no evidence for a flattening at low masses down to the 50\\%\\ completeness limit corresponding to a mass  of 1.1 M$_\\odot$  outside a radius of 5 pc for a 3 Myr population and 1.4 M$_\\odot$ if the oldest stars are 4 Myr. } \\item{The result is in disagreement with the flattening of the IMF below 2 M$_\\odot$ observed by \\citet{sirianni} using optical data covering a region closer to the cluster center. We suggest two possible reasons for the discrepancy: differential extinction and mass segregation.} \\item{We find no evidence for mass segregation outside 3 pc, but with the current data, we cannot rule out that closer to the center the low--mass stars are segregated.} \\item{From the radial surface brightness  profile we have derived a  core radius for the cluster of 0.025 pc (0\\farcs1),  consistent with previous estimates by \\citet{hunter95}.} \\item{The mass of the cluster within 7 pc between 25 M$_\\odot$ and down to 2.1 M$_\\odot$ is estimated to be 5$\\cdot10^4$ M$_\\odot$. If the IMF continues with  a Salpeter slope down to 0.5 M$_\\odot$ the total  mass estimate will double.} \\item{The total mass of the cluster combined with the  large number of low--mass stars suggests that  the 30 Dor cluster may survive to become   a proto--globular cluster depending on the cluster velocity dispersion.} \\end{itemize}"
    },
    "0911/0911.1756_arXiv.txt": {
        "abstract": "Our universe contains a great number of extremely compact and massive objects which are generally accepted to be black holes.  Precise observations of orbital motion near candidate black holes have the potential to determine if they have the spacetime structure that general relativity demands.  As a means of formulating measurements to test the black hole nature of these objects, Collins and Hughes introduced ``bumpy black holes'': objects that are almost, but not quite, general relativity's black holes.  The spacetimes of these objects have multipoles that deviate slightly from the black hole solution, reducing to black holes when the deviation is zero.  In this paper, we extend this work in two ways.  First, we show how to introduce bumps which are smoother and lead to better behaved orbits than those in the original presentation.  Second, we show how to make bumpy Kerr black holes --- objects which reduce to the Kerr solution when the deviation goes to zero.  This greatly extends the astrophysical applicability of bumpy black holes.  Using Hamilton-Jacobi techniques, we show how a spacetime's bumps are imprinted on orbital frequencies, and thus can be determined by measurements which coherently track the orbital phase of a small orbiting body.  We find that in the weak field, orbits of bumpy black holes are modified exactly as expected from a Newtonian analysis of a body with a prescribed multipolar structure, reproducing well-known results from the celestial mechanics literature.  The impact of bumps on strong-field orbits is many times greater than would be predicted from a Newtonian analysis, suggesting that this framework will allow observations to set robust limits on the extent to which a spacetime's multipoles deviate from the black hole expectation. ",
        "introduction": "\\label{sec:intro}  \\subsection{Motivation: Precision tests of the black hole hypothesis}  Though observations of gravitational systems agree well with the predictions of general relativity (GR), the most detailed and quantitative tests have so far been done in the weak field.  (``Weak field'' means that the dimensionless Newtonian potential $\\phi \\equiv GM/rc^2 \\ll 1$, where $M$ is a characteristic mass scale and $r$ a characteristic distance.)  This is largely because many of the most precise tests are done in our solar system (e.g., {\\cite{bertotti}}). Even the celebrated tests which use binary neutron stars (e.g., {\\cite{pulsars}}) are essentially weak-field: for those systems, $M \\sim\\,{\\rm several}\\,M_\\odot$, $r = {\\rm orbital\\ separation} \\sim\\,{\\rm several}\\,R_\\odot$, so $\\phi \\sim\\,{\\rm a\\ few}\\times GM_\\odot/R_\\odot c^2 \\sim \\, {\\rm a\\ few}\\times 10^{-6}$.  This situation is on the verge of changing.  Observational technology is taking us to a regime where we either are or soon will be probing motion in strong gravity, with $\\phi \\gtrsim 0.1$.  Examples of measurements being made now include radio studies of accretion flows near the putative black hole in our galactic center {\\cite{shep}} and x-ray studies that allow precise measurements of accretion disk geometries {\\cite{psaltis_livrev,nms08}}.  Future measurements include the possible discovery of a black hole-pulsar system, perhaps with the Square Kilometer Array {\\cite{smits}}, and gravitational-wave (GW) observations of small bodies spiraling into massive black holes due to the backreaction of GW emission {\\cite{emri}}.  For weak-field studies, a well-developed paradigm for testing gravity has been developed.  The {\\it parameterized post-Newtonian} (PPN) expansion {\\cite{mtw, will93}} quantifies various measurable aspects of relativistic gravity.  For example, the PPN parameter $\\gamma$ (whose value is 1 in GR) quantifies the amount of spatial curvature produced by a unit of rest mass.  Other PPN parameters quantify a theory's nonlinearity, the degree to which it incorporates preferred frames, and the possible violation of conservation laws.  See Ref.\\ {\\cite{will93}}, Chap.\\ 4 for a detailed discussion.  Unfortunately, no similar framework exists for strong-field studies.  If we hope to use observations as tools for testing the nature of strong gravity objects and strong-field gravity, we need to rectify this.  Black holes are of particular interest for studying strong-field gravity.  Aside from having the strongest accessible gravitational fields of any object in the universe\\footnote{More accurately, they have the largest potential $\\phi \\sim GM/Rc^2$.  One may also categorize weak or strong gravity using spacetime curvature or tides; arguably this is a more fundamental measure for assessing whether GR is likely to be accurate or not.  From this perspective, black holes are actually {\\it not} such ``strong gravity'' objects; indeed, the tidal field just outside a $10^8\\,M_\\odot$ black hole's event horizon is not much different from the tidal field at the surface of the Earth.  See Ref.\\ {\\cite{psaltis_livrev}} for further discussion of this point.}, within GR they have an amazingly simple spacetime structure: the ``no-hair'' theorems {\\cite{nohair1,nohair2,nohair3,nohair4,nohair5}} guarantee that the exterior spacetime of any black hole is completely described by only two numbers, its mass $M$ and spin parameter $a$.  {\\it Any} deviation from that simplicity points to a failure either in our understanding of gravity or in the nature of ultracompact objects.  Recent work by Brink {\\cite{brink1}} has reviewed in detail the challenges involved in testing the spacetime structure of massive compact bodies when one relaxes the assumption that the spacetime is Kerr.  As in Brink's discussion, we will focus on spacetimes that are stationary, axisymmetric, and vacuum in some exterior region.  An especially useful parameterization is provided by the work of Geroch {\\cite{geroch70}} and Hansen {\\cite{hansen74}}.  They demonstrate that such a spacetime is completely specified by a set of mass multipole moments $M_l$ and current multipole moments $S_l$.  For a fluid body, $M_l$ describes the angular distribution of the body's mass, and $S_l$ describes the angular distribution of mass flow.  In Newtonian theory, $M_l$ labels a piece of the gravitational potential that varies as $Y_{l0}(\\theta,\\phi)/r^{l+1}$.  There is no Newtonian analog to $S_l$. In the weak-field limit of GR, they label a magnetic-like contribution to gravity and are reminiscent of magnetic moments.  In general, the moments $\\{M_l, S_l\\}$ can be arbitrary.  For black holes, they take special values: a Kerr spacetime has {\\cite{hansen74}} \\begin{equation} M_l + i S_l = M(ia)^l\\;, \\label{eq:Kerr_moments} \\end{equation} (in units with $G = 1 = c$; we neglect the astrophysically uninteresting possibility of the hole having macroscopic charge).  In other words, only two moments out of the set $\\{M_l, S_l\\}$ are independent: $M_0 = M$, and $S_1 = aM$.  Once those two moments have been specified, all other moments are fixed by Eq.\\ (\\ref{eq:Kerr_moments}) {\\it if} the spacetime is Kerr.  Testing the hypothesis that an object is a Kerr black hole can thus be framed as a null experiment.  First, measure the putative black hole spacetime's multipoles.  Using the moments $M_0$ and $S_1$, determine the parameters $M$ and $a$.  {\\it If the spacetime is Kerr, all moments for $l \\ge 2$ must be given by Eq.\\ (\\ref{eq:Kerr_moments}).}  The null hypothesis is that any deviation from those Kerr moments is zero. Failure of the null hypothesis means the black hole candidate is not a Kerr black hole, and may indicate a failure of strong-field GR.  \\subsection{Bumpy black holes: Past work}  Performing this null experiment requires strong-field spacetimes with multipoles differing from those of black holes.  Given the tremendous success of general relativity and the black hole hypothesis at explaining a wide span of data, a reasonable starting point is to consider spacetimes whose multipoles deviate only slightly.  Imagine starting with a spacetime whose Geroch-Hansen moments satisfy \\begin{equation} M_l + i S_l = M(ia)^l + \\delta M_l + i\\delta S_l\\;. \\label{eq:Bumpy_Kerr_moments} \\end{equation} The null hypothesis --- black hole candidates are described by GR's Kerr metric --- means $\\delta M_l = 0$, $\\delta S_l = 0$ for all $l$.  To this end, Collins and Hughes {\\cite{ch04}} (hereafter CH04) introduced the {\\it bumpy black hole}: a spacetime that deviates in a small, controllable manner from the exact black holes of GR.  By construction, the bumpy black hole includes ``normal'' black holes as a limit.  This is central to testing the black hole hypothesis by a null experiment.  Spacetimes of other proposed massive, compact objects (for example, boson stars {\\cite{csw86, ryan97}}) typically do not include black holes as a limiting case.  This limits their utility if black hole candidates are, in fact, GR's black holes.  Measurements of black holes using observables formulated in a bumpy black hole spacetime should simply measure the spacetime's ``bumpiness'' (defined more precisely below) to be zero.  Though a useful starting point, the bumpy black holes developed in CH04 had two major shortcomings.  First, the changes to the spacetime that were introduced to modify its multipole moments were not smooth. The worked example presented in CH04 is interpreted as a Schwarzschild black hole perturbed by an infinitesimally thin ring of positive mass around its equator and by a pair of negative mass infinitesimal points near its poles.  Though this changes the spacetime's quadrupole moment (the desired outcome of this construction), it gives the spacetime a pathological strong-field structure.  This is reflected in the fact that non-equatorial strong-field orbits are ill-behaved in the CH04 construction {\\cite{dimitrios_private}}.  Second, CH04 only examined bumpy Schwarzschild black holes. Although their approach differed in many details from that used in CH04, Glampedakis and Babak {\\cite{gb06}} (hereafter GB06) rectified this deficiency with their introduction of a ``quasi-Kerr'' spacetime. Their construction uses the exterior Hartle-Thorne metric {\\cite{h67,ht68}} describing the exterior of any slowly rotating, axisymmetric, stationary body.  It includes the Kerr metric to ${\\cal O}(a^2)$ as a special case.  Identifying the influence of the quadrupole moment in the Hartle-Thorne form of the Kerr metric, they then introduce a modification to the full Kerr spacetime that changes the black hole's quadrupole moment from its canonical value.  \\subsection{This analysis}  The goal of this work is to rectify the deficiencies of CH04 and to extend the bumpy black hole concept to spinning black holes.  Dealing with the non-smooth nature of the bumps presented in CH04 is, as we show in Secs.\\ {\\ref{sec:spacetimes}}, {\\ref{sec:results_schw}}, and {\\ref{sec:results_kerr}}, quite straightforward.  It simply requires introducing perturbations to the black hole background that are smooth rather than discontinuous.  In essence, rather than having bumps that correspond to infinitesimal points and rings, the bumps we use here are smeared into pure multipoles.  Extending CH04 to spinning black holes is more of a challenge.  Given that GB06 already introduced a bumpy-black-hole-like spacetime that encompasses spacetimes with angular momentum, one might wonder why another construction is needed.  A key motivation is that we would like to be able to make an {\\it arbitrary} modification to a black hole's moments.  Although showing that a black hole candidate has a non-Kerr value for the quadrupole moment would be sufficient to falsify its black hole nature (at least within the framework of GR), one can imagine scenarios in which the first $L$ moments of a black hole candidate agree with the Kerr value, but things differ for $l > L$.  For example, Yunes et al.\\ have shown that in Chern-Simons modifications to GR, slowly rotating black hole solutions have the multipolar structure of Kerr for $l < 4$, but differ for $l \\ge 4$ {\\cite{csbh}}.  There are many ways in which black hole candidates might differ from the black holes of general relativity; we need to develop a toolkit sufficiently robust that it can encompass these many potential points of departure.  Our technique for making bumpy Kerr holes is based on the Newman-Janis algorithm {\\cite{nj}}.  This algorithm transforms a Schwarzschild spacetime into Kerr by ``rotating'' the spacetime in a complex configuration space\\footnote{This complex ``rotation'' is made most clear by framing the discussion using the Ernst potential {\\cite{ernst}}.  Transforming from Schwarzschild to Kerr corresponds to adding an imaginary part to a particular potential.  Further discussion is given in a companion paper {\\cite{vigeland}}.}.  In this paper, we construct bumpy Kerr black holes by applying the Newman-Janis algorithm to bumpy Schwarzschild black holes.  The outcome of this procedure is a spacetime whose mass moments are deformed relative to Kerr.  When the bumpiness is set equal to zero, we recover the Kerr metric.  A companion paper {\\cite{vigeland}} examines the multipolar structure of bumpy Kerr black holes in more detail, demonstrating that one can construct a spacetime in which the mass moments deviate arbitrarily from Kerr (provided that they are small).  Once one has constructed a bumpy black hole spacetime, one then needs to show how its bumps are encoded in observables.  The most detailed quantitative tests will come from orbits near black hole candidates. As such, it is critical to know how orbital frequencies change as a function of a spacetime's bumpiness.  More generally, we are faced with the problem of understanding motion in general stationary axisymmetric vacuum spacetimes.  Brink {\\cite{brink2}} has recently published a very detailed analysis of this problem, with a focus on understanding whether and for which situations the spacetimes admit integrable motion.  She has found evidence that geodesic motion in such spacetimes may, in many cases, be integrable.  If so, the problem of mapping general spacetimes (not just ``nearly black hole'' spacetimes) may be tractable.  Gair, Li, and Mandel {\\cite{glm08}} have similarly examined orbital characteristics in the Manko-Novikov spacetime {\\cite{mn92}}, which has a particular tunable non-Kerr structure.  They show how orbits change in such spacetimes, and how its bumpiness colors observable characteristics.  For this analysis, we confine ourselves to the simpler problem of motion in bumpy black hole spacetimes, addressing this challenge in Secs.\\ {\\ref{sec:results_schw}} and {\\ref{sec:results_kerr}} using canonical perturbation theory.  As is now well known (and was shown rather spectacularly by Schmidt {\\cite{schmidt}}), Hamilton-Jacobi methods let us write down closed-form expressions for the three orbital frequencies ($\\Omega^r$, $\\Omega^\\theta$, $\\Omega^\\phi$) which completely characterize the behavior of bound Kerr black hole orbits. Since bumpy black hole spacetimes differ only perturbatively from black hole spacetimes, canonical perturbation theory lets us characterize how a spacetime's bumps shift those frequencies, and thus are encoded in observables.  Similar techniques were used in GB06 to see how frequencies are shifted in a quasi-Kerr metric, and were also used by Hinderer and Flanagan {\\cite{hf08}} in a two-timescale analysis of inspiral into Kerr black holes.  \\subsection{Organization and overview}  We begin in Sec.\\ {\\ref{sec:spacetimes}} with an overview of the spacetimes that we study.  We start with the axially symmetric and stationary Weyl line element.  We first review the Einstein field equations in this representation, introduce the Schwarzschild limit, and then describe first order perturbations.  The spacetime's bumpiness is set by choosing a function $\\psi_1$ which controls how the spacetime deviates from the black hole limit.  We initially leave $\\psi_1$ arbitrary, except for the requirement that it be small enough that terms of order $(\\psi_1)^2$ can be neglected.  Later in the paper, we will take $\\psi_1$ to be a pure multipole in the Weyl representation.  Following our discussion of bumpy Schwarzschild spacetimes, we show how to use the Newman-Janis algorithm to build bumpy Kerr black holes.  We study geodesic motion in these spacetimes in Sec.\\ {\\ref{sec:motion}}.  We begin by reviewing the most important properties of normal black hole orbits, reviewing Kerr geodesics in Sec.\\ {\\ref{sec:kerr_geodesics}}, and then describing how to compute orbital frequencies using Hamilton-Jacobi methods in Sec.\\ {\\ref{sec:normal_bh_freqs}}.  The discussion of frequencies is largely a synopsis of Schmidt's pioneering study, Ref.\\ {\\cite{schmidt}}.  We then show how these techniques can be adapted, using canonical perturbation theory, to bumpy black holes (Sec.\\ {\\ref{sec:bumpy_bh_freqs}}).  Canonical perturbation theory requires averaging a bump's shift to an orbit's Hamiltonian.  This averaging was developed in Ref.\\ {\\cite{dh04}}, and is summarized in Appendix {\\ref{app:averaging}}.  We present our results for bumpy Schwarzschild and bumpy Kerr orbits in Secs.\\ {\\ref{sec:results_schw}} and {\\ref{sec:results_kerr}}, respectively.  We take $\\psi_1$ to be a pure multipole in the Weyl sector, construct the spacetime, and numerically compute the shifts in $\\Omega^r$, $\\Omega^\\theta$, and $\\Omega^\\phi$.  Detailed results are given for $l = 2$, $l = 3$, and $l = 4$.  There is no reason in principle to stop there, though the results quickly become repetitive. Section {\\ref{sec:results_schw}} gives our results for bumpy Schwarzschild black holes, and Sec.\\ {\\ref{sec:results_kerr}} gives results for bumpy Kerr.  A few results are worth highlighting.  First, we find that the exact numerical results for frequency shifts correctly reproduce the Newtonian limit as our orbits are taken into the weak field.  (We develop this limit in some detail in Appendix {\\ref{app:newton}} to facilitate the comparison.)  The frequency shifts are substantially enhanced in the strong field.  The shift to the radial frequency is particularly interesting: it tends to oscillate, shifting between an enhancement and a decrement as orbits move into the strong field. This behavior appears to be a robust signature of non-Kerr multipole structure in black hole strong fields.  Interestingly, it turns out that black hole spin does not have a very strong impact on the bumpiness-induced shifts to orbital frequencies. Spin's main effect is to change the location of the last stable orbit. For large spin, orbits reach deeper into the strong field, amplifying the bumps' impact on orbital frequencies.  Aside from the change to the last stable orbits, the impact of a particular multipolar bump looks largely the same across all spin values.  Some examples of the frequency shifts we find are shown in Figs.\\ {\\ref{fig:dOm_l2}} ($l = 2$, Schwarzschild), {\\ref{fig:dOm_l4}} ($l = 4$, Schwarzschild) and {\\ref{fig:dOm_l2_Kerr}} ($l = 2$, Kerr).  (We have no figures for $l = 3$ since the secular shifts to orbital frequencies are zero in that case, as they are for all odd values of $l$.  We also do not show results for $l = 4$ bumps of Kerr black holes since, as discussed above, they are very similar to the $l = 4$ Schwarzschild results.)  Finally, we summarize our analysis and suggest some directions for future work in Sec.\\ {\\ref{sec:summary}}.  Among the points we note are that, in this analysis, we only consider changes to the mass moments of the black hole spacetime.  Adding ``current-type'' bumps to a spacetime is discussed in a companion paper {\\cite{vigeland}}.  We also do not discuss in this analysis the issue of measurability. Turning these foundations for mapping the multipole moment structure of black holes into a practical measurement program (for instance, via gravitational-wave measurements, timing of a black hole-pulsar binary, or precision mapping in radio or x-rays of accretion flows) will take a substantial effort.  Throughout this paper, we work in geometrized units with $G = c = 1$; a useful conversion factor in these units is $1\\,M_\\odot = 4.92\\times 10^{-6}$ seconds.  When we discuss bumpy black holes, we will always use a ``hat'' accent to denote quantities which are calculated in the pure black hole background spacetimes.  For example, an orbital frequency is written $\\Omega = \\hat\\Omega + \\delta\\Omega$; $\\hat\\Omega$ is the frequency of an orbit in the black hole background, and $\\delta\\Omega$ denotes the shift due to the black hole's bumps. ",
        "conclusions": "\\label{sec:summary}  This analysis significantly improves on the earlier presentation of bumpy black holes given in CH04, producing bumps with well-behaved strong-field structure, and extending the concept to Kerr black holes. These extensions greatly expand the astrophysical relevance of these spacetimes.  We have also demonstrated how Hamilton-Jacobi theory can be applied to orbits in bumpy spacetimes to categorize the anomalous precessions arising from their bumps in a reasonably straightforward manner.  Bumpy black holes can now, at least in principle, be used as the foundation for strong-gravity tests with astrophysical data.  It's worth re-emphasizing why we propose to use bumpy black holes, rather than using exact solutions which include black holes as a limit (for example, the Novikov-Manko spacetime {\\cite{mn92}} used in Ref.\\ {\\cite{glm08}}).  In large part, our choice is a matter of taste.  Our goal is to tweak a black hole's moments in an arbitrary manner, so that the non-Kerr nature is entirely under our control.  From the standpoint of formulating a null experiment, it arguably makes no difference whether the Kerr deviation takes one particular form or another.  {\\it Any} falsifiable non-Kerr form is good enough to formulate the test.  To our minds, the nice feature of this approach is that if, for example, a theory of gravity specifies that a black hole should have the same moment structure as general relativity up to some $l = L$, it is simple to design a spacetime tailored to testing that theory.  The case of Chern-Simons gravity discussed in Ref.\\ {\\cite{csbh}} shows that this motivation is not merely academic, but is motivated by plausible alternatives to general relativity.  This analysis provides a complete description of ``mass-type'' bumps, i.e., perturbations to the mass moments $M_l$.  We have not examined ``current-type bumps,'' shifts to the spin moments $S_l$.  It is also worth bearing in mind that a pure multipole in the Weyl sector, $\\psi^l_1 \\propto Y_{l0}(\\theta)$ does not correspond to a shift in a pure Geroch-Hansen multipole.  A companion paper by one of us (SJV) addresses these issues {\\cite{vigeland}}.  That analysis shows how spin moments can be adjusted from their Kerr values by perturbing the Ernst potential describing the spacetime {\\cite{ernst}}.  It also shows that a pure Weyl-sector multipole $\\psi^l_1$ changes no Geroch-Hansen moments {\\it lower} than $l$.  As such, a pure Geroch-Hansen moment perturbation can be assembled by combining multiple Weyl-sector multipole perturbations.  Aside from these more formal issues, the next major step in this program will be to use these foundations to formulate actual strong-field gravity tests that can be applied to astrophysical data. We imagine several directions that would be interesting to follow:  \\begin{itemize}  \\item {\\it Extreme mass ratio inspiral (EMRI)}: The capture and inspiral of stellar mass compact into massive black holes at galaxy centers is one of the original motivations of this work.  Much of the recent literature on testing and mapping black hole spacetimes has centered on understanding the character of orbits in non-Kerr black hole candidate spacetimes {\\cite{brink1,brink2,haris}}, with an eye on application to gravitational-wave measurement of EMRIs.  The full analysis of EMRIs in non-Kerr spacetimes is, in principle, quite complicated since their non-Kerr-ness breaks the Petrov Type-D character of Kerr black holes.  As such, it may be quite difficult to accurately compute their radiation emission.  It may not be quite so difficult in bumpy black hole spacetimes.  Thanks to the smallness of their non-Kerr character, it may be fruitful to use a ``hybrid'' approach in which the short timescale motion is computed in the bumpy spacetime, but the radiation generation and backreaction is computed in the Kerr spacetime; a similar idea was suggested in the context of modeling EMRI events in Chern-Simons gravity {\\cite{ys09}}.  Given that our goal is to formulate a null experiment, this hybrid may be good enough for a useful test, in lieu of solving the entire radiation reaction problem in non-Kerr spacetimes.  \\item {\\it Black hole-pulsar systems}: One of the goals of the planned Square Kilometer Array (SKA) {\\cite{smits}} is the discovery of a black hole-pulsar binary system.  If such a system is discovered, detailed observation over many years should be able to tease the multipole structure of the black hole from the data.  Similar observations of neutron star-pulsar binary systems have already allowed us to make exquisite measurements of neutron star properties and gravitational-wave emission {\\cite{stairs_lr}}.  The tools developed here may already be adequate for doing this analysis since such binaries will have a relatively slow inspiral time.  \\item {\\it Accretion flows on black hole candidate}: Programs to observe the (presumed) black hole at the center of our galaxy are maturing very quickly; programs to study accretion flows onto stellar mass black holes in x-rays are already quite mature.  In our galactic center, the most precise measurements come from millimeter wavelength radio emission from gas accreting onto this central object.  The precision of these measurements is increasing to the point where we will soon be able to use them to map the detailed strong-field spacetime structure of the spacetime near Sagittarius A*.  It would be a worthwhile exercise to repeat analyses of the appearance of these flows in Kerr spacetimes {\\cite{avery}} for bumpy spacetimes to see how accurately such measurements may be able to probe the central object's multipoles.  Such an analysis will require developing imaging maps of bumpy black holes, going beyond the orbit frequency analysis we have given here.  In the domain of x-rays, it would be very interesting to extend work on quiescent accretion flows (as in, for example, Ref.\\ {\\cite{nms08}}) to include accretion in bumpy black hole spacetimes.  It may be possible to extend those analyses to fit for multipoles beyond the black hole's mass and spin.  \\end{itemize}"
    },
    "0911/0911.3279_arXiv.txt": {
        "abstract": "{}{We present an elegant method of determining the eigensolutions of the induction and the dynamo equation in a fluid embedded in a vacuum.}{The magnetic field is expanded in a complete set of functions. The new method is based on the biorthogonality of the adjoint electric current and the vector potential with an inner product defined by a volume integral over the fluid domain. The advantage of this method is that the velocity and the dynamo coefficients of the induction and the dynamo equation do not have to be differentiated and thus even numerically determined tabulated values of the coefficients produce reasonable results.}{We provide test calculations and compare with published results obtained by the classical treatment based on the biorthogonality of the magnetic field and its adjoint. We especially consider dynamos with mean-field coefficients determined from direct numerical simulations of the geodynamo and compare with initial value calculations and the full MHD simulations.}{} ",
        "introduction": "The generation and evolution of magnetic fields in cosmic bodies like the planets and stars is generally thought to be governed by induction processes due to motions in their electrically conducting fluid interior. The magnetic field $\\bB$ is described by the induction equation \\be \\frac{\\p\\bB}{\\p t} = \\bn \\t D\\bB \\label{ie} \\ee where \\be D\\bB = \\bu\\t\\bB - \\eta\\bn\\t\\bB \\;. \\label{io} \\ee Here $\\bu$ represents the velocity and $\\eta$ the magnetic diffusivity.  In the framework of mean-field theory \\citep[e.g.][]{m78,kr80}, $\\bu$ and $\\bB$ are considered as mean, e.g. ensemble averaged, quantities, whereas the action of the small-scale turbulent flow on the mean magnetic field is parametrised by the so-called dynamo coefficients, $\\balpha$ and $\\bbeta$. They are, in general, tensors of second and third rank, respectively. We use the following compact notation of the mean field coefficients, which include the so-called $\\gamma$, $\\delta$, and $\\kappa$-effects, see e.g. \\citet{r80}. Then, the operator $D$ reads \\be D\\bB = \\bu\\t\\bB + \\balpha\\cd\\bB - \\bbeta:(\\bn\\bB) - \\eta\\bn\\t\\bB \\label{do} \\ee instead of (\\ref{io}), and acts on the mean magnetic field. Except for the additional $\\balpha$ and $\\bbeta$ terms in the $D$ operator, the induction and the dynamo equation are formally equivalent. Thus, the new method presented here equally applies to both.  The dynamo region is located in a flow domain $V$ with exterior vacuum $E$. In this work we assume $V$ to be either a sphere or a spherical shell. The magnetic field $\\bB$ is continuous through the boundary $\\p V$ and potential in $E$.  In kinematic dynamo theory, all coefficients ($\\bu$, $\\balpha$, $\\bbeta$, and $\\eta$) are assumed given and independent of the magnetic field. Thus the dynamo equation is linear in the magnetic field and can be solved by considering an eigenvalue problem \\be \\lam\\bB = \\bn \\t D\\bB \\label{dep} \\ee with eigenvalues $\\lam$ describing the time evolution proportional to $\\exp(\\lam t)$ of the magnetic field $\\bB$.  Many studies have been made of the eigenvalues of the dynamo operator for various celestial bodies and with many forms of the dynamo coefficients \\citep[e.g.][]{bg54,r60,sk69,ds71,r72,rs72,g73,kr75,ss89,dj89,dgs93,g00,sz01, lj04,lj05,jw06,jw07}. Often the coefficients are approximated by simple analytical functions of position, and their tensorial character is disregarded. Recently, the test-field method, developed by \\citet{srsrc05,srsrc07} \\citep[see also][]{osb01,osbr02}, allows one to determine all tensorial components of $\\balpha$ and $\\bbeta$ directly from self-consistent numerical simulations \\citep{b08,k09}. These coefficients are sometimes strongly varying functions of position. This may introduce large errors because the dynamo operator $\\bn\\t D$ involves differentiation of the dynamo coefficients, and these are only available as numerically determined tabulated values.  In this paper we present a new method that does not require differentiation, so it is also applicable to numerically determined dynamo coefficients. The method is based on the biorthogonality of the electric current and the vector potential with an inner product defined by a volume integral over the fluid. This property has already been noted by \\citet{rb87}, \\citet{h88}, \\citet{frs93}, \\citet{hs95}, and \\citet{rraf02}.  The method is described in detail in Sect.~\\ref{sec_ep}. Extensive test calculations have been performed and compared with published results by other eigenvalue methods. Some of these tests are presented in Sect.~\\ref{sec_tr}. In Sect.~\\ref{sec_gd} we apply the new method to eigenmodes of the dynamo operator with coefficients obtained from geodynamo models. Our conclusion are drawn in Sect.~\\ref{sec_concl}. ",
        "conclusions": "\\label{sec_concl}  We presented a new method for computing the eigenvalues and eigenfunctions of the induction and the dynamo equation. The method is based on the biorthogonality of the adjoint electric current and the vector potential with an inner product defined by a volume integral over the fluid domain. The advantage of the method is that the velocity and dynamo coefficients do not have to be differentiated. The method is therefore well-suited for spatially strongly variable dynamo coefficients. %  We tested the new method against the classical treatment and proved that it works correctly and efficiently. We applied it to two cases with dynamo coefficients derived from direct numerical simulations of the geodynamo. The obtained dynamo eigenmodes are promising candidates for decomposing the magnetic field of the numerical simulations and for analysing the statistical properties of the mode coefficients as proposed by \\citet{h09}."
    },
    "0911/0911.4345_arXiv.txt": {
        "abstract": "The delayed detonation model describes the observational properties of the majority of type Ia supernovae very well. Using numerical data from a three-dimensional deflagration model for type Ia supernovae, the intermittency of the turbulent velocity field and its implications on the probability of a deflagration-to-detonation (DDT) transition are investigated. From structure functions of the turbulent velocity fluctuations, we determine intermittency parameters based on the log-normal and the log-Poisson models. The bulk of turbulence in the ash regions appears to be less intermittent than predicted by the standard log-normal model and the She-L\\'{e}v\\^{e}que model. On the other hand, the analysis of the turbulent velocity fluctuations in the vicinity of the flame front by R\\\"opke suggests a much higher probability of large velocity fluctuations on the grid scale in comparison to the log-normal intermittency model. Following Pan et al., we computed probability density functions for a DDT for the different distributions. The determination of the total number of regions at the flame surface, in which DDTs can be triggered, enables us to estimate the total number of events. Assuming that a DDT can occur in the stirred flame regime, as proposed by Woosley et al., the log-normal model would imply a delayed detonation between $0.7$ and $0.8$ seconds after the beginning of the deflagration phase for the multi-spot ignition scenario used in the simulation. However, the probability drops to virtually zero if a DDT is further constrained by the requirement that the turbulent velocity fluctuations reach about $500\\,\\mathrm{km\\,s^{-1}}$. Under this condition, delayed detonations are only possible if the distribution of the velocity fluctuations is not log-normal. From our calculations follows that the distribution obtained by R\\\"opke allow for multiple DDTs around $0.8$ seconds after ignition at a transition density close to $1\\times 10^{7}\\,\\mathrm{g\\,cm^{-3}}$. ",
        "introduction": "\\label{Intro}  Although deflagration models of Type Ia supernovae \\citep[e.g.][]{Rein02, Gam03, Roep05} reproduce bulk properties of fainter Type Ia supernovae, delayed detonations \\citep[see][for a review of explosion scenarios]{HilleNie00} seem to be the the most promising way of modeling the majority of the observed events \\citep{RoepNie07, Mazzali07}.  However, a theoretical explanation for the deflagration-to-detonation (DDT) transition that gives rise to this type of thermonuclear explosion has been missing so far. Originally postulated by \\citet{Khokhlov3} and \\citet{WoosWeav94}, the possibility of a DDT in type Ia supernovae was questioned by \\citet{Niemeyer}, because the preconditioning required to trigger a self-sustaining detonation front would be extremely unlikely to occur in the deflagration phase. This conclusion was based on ensemble-average scaling arguments. \\citet{Pan}, on the other hand, pointed out that highly intermittent turbulent fluctuations act in favour of a DDT, because intermittency allows for velocities much higher than the ensemble average.  As proposed by \\citet{Niemeyer2}, \\citet{Niemeyer3} and \\citet{Khokhlov4}, a necessary condition for triggering a DDT is that the distributed burning regime has to be entered. Based on log-normal \\citep{Kolmogorov} and log-Poisson \\citep{She} intermittency models with typical parameters expected for turbulence in thermonuclear supernovae, Pan et al. calculated probabilities for this DDT condition. Their calculations implied transition densities in the range $3.8\\times 10^{7}\\ldots 2.7\\times 10^{8}\\,\\mathrm{g\\,cm^{-3}}$. \\citet{LisHille00}, \\citet{Woosley07}, and \\citet{WoosKer09}, on the other hand, provided numerical evidence that there are stronger constraints on a DDT than just entering the regime of distributed burning. As a consequence, lower transition densities would be favored. Analyzing the small-scale turbulent velocity distribution predicted by the subgrid-scale model in several thermonuclear supernova simulations, \\citet{Roep07} found that intermittent velocity fluctuations up to $1000\\,\\mathrm{km\\,s^{-1}}$ on a length scale of $10\\,\\mathrm{km}$ are possible. Fluctuations of this magnitude could trigger a detonation in accordance with \\citet{LisHille00}.  The ambiguity of conditions for DDTs in type Ia supernovae makes predictions difficult. Here, we follow a pragmatic approach and investigate different DDT criteria, assuming simple approximations to the microphysical conditions. Following the approach of Pan et al., we calculate the probability of a DDT using data sets from a highly resolved three-dimensional deflagration model \\citep{Roepke}. \\citet{CirSchm08}, in the following referred to as paper I, determined characteristic scales and scaling exponents from the structure functions of the turbulent velocity fluctuations up to order six. We fitted log-normal and log-Poisson intermittency models to these scaling exponents. In addition, we determined probability density functions (pdfs) of the mass density constrained to the flame surface in order to calculate the probability of a DDT at different stages of the deflagration. There are no further adjustable parameters apart from the initial conditions of the deflagration model. Assuming a certain size for DDT regions, the total number of DDTs can be obtained from the cumulative probability and the number of regions in the volume occupied by the flame.\\footnote{ Down to the Gibson scale, the flame front can be considered as a fractal surface, which is numcerically represented by a zero level set. Thus, the volume occupied by the flame is defined by all grid cells within a certain distance to the zero level set.}  Basically, \\citet{Pan} assumed that a DDT is triggered once the Karlovitz number $\\mathrm{Ka}$ (the square root of the ratio of the laminar flame width to the Gibson scale as defined in Section~\\ref{sc:ddt_prob}) exceeds a certain value \\citep{Niemeyer3,WoosKer09}, i.~e., burning has advanced sufficiently far into the regime of distributed burning.  Because the critical size of a preconditioned region is extremely sensitive to the mass density, it is difficult to estimate the total number of DDTs. Nevertheless, our results suggest that a delayed detonation would result at an early stage of the explosion if the conditions for a DDT were constrained by the Karlovitz number only. A further implication would be a transition density well above $1\\times 10^{7}\\,\\mathrm{g\\,cm^{-3}}$, which was found in the study by \\citet{Pan}.  Whereas the criterion for the onset of distributed burning can be expressed in terms of laminar flame properties, \\citet{WoosKer09} argued that the conditions for a DDT cannot be parameterized on the basis of the laminar flame speed and thickness. Using numerical data for the nuclear time scale in the so-called stirred flame regime \\citep{Kerst01} and following the same procedure as outlined above, we calculated the DDT probability. Although there are substantial uncertainties in this approach, a DDT at later time and lower density compared to the estimates following from laminar flame properties is implied. Most important, we find that a DDT can be excluded if the magnitude of the turbulent velocity fluctuations is required to exceed about $500\\,\\mathrm{km\\,s^{-1}}$, which is suggested by microphysical studies \\citep{Woosley07}. Delayed detonations are possible, however, if the distribution of the velocity fluctations significantly deviates from a log-normal shape. According to \\citet{Roep07}, this appears to be the case for regions near the flame front.  The paper is organized as follows. In the next Section, we briefly review the log-normal and log-Poisson intermittency models. Moreover, the results on the intermittency of turbulence in the simulation are discussed. The procedure of calculating the DDT probability is described in Section~\\ref{sc:ddt_prob}. In Section~\\ref{sc:number_ddt}, the estimation of the total number of DDTs is explained. We present our numerical results in this Section~\\ref{sc:results}, followed by a discussion in the last Section. ",
        "conclusions": "We investigated the intermittency properties of turbulence in the numerical simulation of a thermonuclear supernova by \\citet{Roepke}. \\citet{Pan} proposed that the probability of entering the distributed burning regime can be computed from the log-normal model for intermittent turbulence, and from this probability the incidence of a DDT can be inferred. Evaluating the characteristic scales of turbulence and the probability density functions of the mass density in the vicinity of the flame front at different instants, we calculated the probability of a DDT for various Karlovitz numbers. We also investigated the influence of more restrictive criteria following from the numerical studies by \\cite{WoosKer09}.  Assuming that a detonation can be triggered after the onset of distributed burning, our calculations indicate that a delayed detonation would commence at an early stage of the explosion. In contrast to \\citet{Pan}, where a spherical flame is assumed, the highly wrinkled and folded flame front in the numerical simulation greatly increases the number of regions of critical size. Although we cannot precisely determine the total number of regions of size $\\ell_{\\mathrm{c}}$, because the critical size greatly varies with the mass density \\citep{Niemeyer2}, and $\\ell_{\\mathrm{c}}\\ll\\Delta$, where $\\Delta$ is the grid resolution, it appears that $N_{\\mathrm{DDT}}$ greatly exceeds unity already $0.6$ seconds after ignition. Then the transition density would be signifcantly higher than $10^{7}\\,\\mathrm{g\\,cm^{-3}}$. This possibility was pointed out by \\citet{Pan}.  If DDT conditions are constrained by an interval of Damk\\\"{o}hler numbers in the stirred flame regime \\citep{WoosKer09}, a DDT earlier than $0.7$ seconds after ignition can be excluded. Assuming a typical value of the Damk\\\"{o}hler number, for which a DDT can be triggered, a delayed detonation could occur around $t=0.8$ seconds, and the transition density would be close to $10^{7}\\,\\mathrm{g\\,cm^{-3}}$. However, assuming a log-normal distribution, the typcial velocity fluctuations would be implausibly low. If turbulent velocity fluctuations greater than $500\\,\\mathrm{km s^{-1}}$ are required \\citep{WoosPriv}, then non-log-normal distributions are vital for delayed detonations.  One of the problems of determining log-normal distributions from scaling exponents of structure functions is that it is difficult to calculate higher-order two-point statistics. Consequently, the computed scaling exponents are sufficient to constrain the central part of the distribution only, whereas the tails are based on an extrapolation. A further caveat is that we are not able to separate ash from fuel regions, because two-point statistics of turbulence can only be computed in convex regions. For DDTs, however, turbulence in fuel close to the flame front is signicant. Indeed, \\citet{Roep07} showed that the tail of the pdf of turbulent velocity fluctuations on a length scale of about $10\\,\\mathrm{km}$ in the vincity of the flame front is much flatter than what is expected on the basis of log-normal models. With this distribution, DDTs between $0.7$ and $0.8$ seconds after ignition at a transition density close to $10^{7}\\,\\mathrm{g\\,cm^{-3}}$ are definitely possible. Of course, the detonation time also depends on the ignition scenario and the subsequent evolution of the deflagration \\citep[see, for example,][]{SchmNie06}. In any case, it is crucial to settle the question of the relevant distribution of turbulent velocity fluctuations in the future. Given a certain distribution, the minimal velocity fluctuation that is necessary to trigger a DDT is an extremely important parameter. Depending on this threshold, delayed detonations might be theoretically confirmed or excluded. Consequently, further studies of the microphysics, particularly in the range of densities below $10^{7}\\,\\mathrm{g\\,cm^{-3}}$, will be essential for a more accurate evaluation of the DDT probability in type Ia supernovae.  Based on the insights provided by such calculations, it may be possible to devise an algorithm that determines the local probability of a DDT in large-scale simulations of thermonunclear supernovae. Since the processes causing a DDT occur on scales that are difficult to resolve in such simulations, a subgrid scale model is likely to play some role \\citep{SchmNie06b}. Apart from this, the propagation of the burning zone has to be treated beyond the onset of distributed burning. \\citet{Schm07} proposed that the level set technique can be extended at least into the broken-reaction-zones regime \\citep{KimMen00}, provided that the burning time scale does not exceed some fraction of the eddy turn-over time corresponding to the numerically unresolved velocity fluctuations (i.~e., the ratio of the grid cell size to the square root of the specific subgrid scale turbulence energy). This approach also calls for further microphysical studies. An entirely different approach might be the use of adaptive mesh refinement and the \\emph{in situ} calculation of the processes triggering a DDT, for instance, by means of LEM \\citep{Kerst91}. Once all numerical challenges are met, quantitative theoretical arguments in favor or against delayed detonations as an explanation for type Ia supernovae will be within reach."
    },
    "0911/0911.4942_arXiv.txt": {
        "abstract": "We compare the results of the numerical simulation of the viscous-like interaction of the solar wind with the plasma tail of a comet, with velocities of H$_2$O+ ions in the tail of comet Swift-Tuttle determined by means of spectroscopic, ground based observations. Our aim is to constrain the value of the basic parameters in the viscous-like interaction model: the effective Reynolds number of the flow and the interspecies coupling timescale. We find that in our simulations the flow rapidly evolves from an arbitrary initial condition to a quasi-steady state for which there is a good agreement between the simulated tailward velocity of H$_2$O+ ions and the kinematics derived from the observations. The fiducial case of our model, characterized by a low effective Reynolds number ($Re \\approx 20$ and selected on the basis of a comparison to {\\em in situ} measurements of the plasma flow at comet Halley, yields an excellent fit to the observed kinematics. Given the agreement between model and observations, with no {\\it ad hoc} assumptions, we believe that this result suggests that viscous-like momentum transport may play an important role in the interaction of the solar wind and the cometary plasma environment. ",
        "introduction": "The interaction of the solar wind with the plasma environment of a comet has been the subject of numerous studies for the past 50 years. The basic model to describe how this interaction takes place, stems from the work of Biermann (1951), Alfven (1957), Biermann et al. (1967) and Wallis (1973), and involves primarily the effect of mass loading of the solar wind by the picked-up, newly born cometary ions from the expanding neutral envelope of the comet, and the draping of the interplanetary magnetic field into its magnetic tail. Cravens and Gombosi (2004) and Ip (2004) review the current understanding of the interaction of the solar wind with the plasma environment of comets on the basis of these processes.  In addition to these effects, P\\'erez-de-Tejada et al. (1980) and P\\'erez-de-Tejada (1989) proposed that, as it happens in other ionospheric, unmagnetized obstacles to the solar wind, such as Venus and Mars, several features of the large scale flow dynamics in the plasma environment of comets, can be attributed to the action of viscous-like forces as the solar wind interacts with cometary plasma. As suggested by Reyes-Ruiz et al. (2009, hereafter referred to as RR09), momentum and heat transport coefficients in cometosheath plasmas can be greatly increased with respect to their normal, ``molecular'' values, owing to the turbulent character of the flow (Baker et al. 1986, Scarf et al. 1986, Klimov et al. 1986, Tsurutani \\& Smith 1986) and/or to the development of plasma instabilities leading to effective wave-particle interactions, as discussed by Shapiro et al. (1995) and  Dobe et al. (1999 and references therein) in the ionosheath of Venus.  {\\em In situ} measurements by the Giotto spacecraft indicate that, as in Venus and Mars, the solar wind flow in the ionosheath of comet Halley exhibits an intermediate transition, also called the ``mystery transition'', approximately half-way between the bow shock and the cometopause (Johnstone et al. 1987, Goldstein et al. 1986, Reme 1991). Below this transition, as we approach the cometopause, the antisunward velocity of the shocked solar wind decreases in a manner consistent with a viscous boundary layer (P\\'erez-de-Tejada 1989). Also  indicative of the presence of viscous-like dissipation processes, the temperature of the gas increased and the density decreased, as the spacecraft moved from the intermediate transition to the cometopause.  In a recent study, following Perez-de-Tejada (1999), RR09 have argued that the superalfvenic character of the flow in the cometosheath, downstream from the nucleus and in the tail region, as found from {\\em in situ} measurements at comet Giacobinni-Zinner, suggests that viscous-like effects are, at least, as important as ${\\bf J} \\times {\\bf B}$ forces throughout the region. The fact that $M_A^2 >> 1$ in the cometosheath means that the magnetic energy density is much smaller than the kinetic energy associated with the inertia of the plasma. This implies that ${\\bf J} \\times {\\bf B}$ forces are not the dominant dynamical factor responsible for the large scale properties of the flow in the region. By comparing the magnitude of the terms corresponding to momentum transport due to viscous-like forces and ${\\bf J} \\times {\\bf B}$ forces in the momentum conservation equation, Perez-de-Tejada (1999) argued that downstream from the terminator in the ionosheath of Venus, a scenario analogous to the one considered in this paper, the fact that the flow is superalfvenic, as found from the {\\em in situ} measurements of the Mariner 5 and Venera 10 spacecraft, indicates that viscous-like forces may dominate over ${\\bf J} \\times {\\bf B}$ forces in the flow dynamics in the boundary layer formed in the interaction of solar wind and ionospheric plasma. If the flow is characterized by a low effective Reynolds number, $Re$, this layer extends over a significant portion of the ionosheath of the solar wind obstacle. RR09 have constrained the value of the effective Reynolds number of the flow, $Re$, by comparing the results of numerical simulations of the interaction between the solar wind and a cometary plasma tail, taking into account the effect of viscous-like forces, with {\\em in situ} measurements in the ionosheath of comets Halley and Giacobinni-Zinner. They conclude that a value of $R  \\lesssim 50$ is necessary to reproduce the relative location of the cometopause, intermediate transition and bow shock in these comets.  In this paper we continue our study of the hypothesis that viscous-like forces are important in the solar wind-comet plasma interaction, by comparing the flow properties along the plasma tail of a comet, as derived from the results of our numerical simulations, with the kinematics of H$_2$O+ ions determined from spectroscopic, ground-based observations of comet Swift-Tuttle during late 1992 (Spinrad et al. 1994). Since the precise origin of the viscous-like momentum transfer processes invoked in our simulations is not yet clear, our study will contribute to the understanding of this process by placing constraints on the parameters, such as the effective Reynolds number, that control the flow dynamics. This being the first attempt to include viscous-like effects in models of the solar wind-comet interaction, several, potentially important factors have been neglected, such as the effects of magnetic fields, possibly important in the midplane of the plasma tail (see above), and ion-neutral collisions, important closer to the comet nucleus. These issues are further discussed in RR09 and shall be considered in future studies.  The paper is organized as follows: in section II we review our numerical simulations of the viscous interaction between the solar wind and the plasma tail of a comet. In section III we present the observations of Spinrad et al. (1994), and the inferred velocity profiles that will be used in the comparison. The results of the comparison between model and observations are presented in section IV. In section V we discuss several issues related to our results, such as time dependence of the flow on the timescale of the different observations. Finally, a summary of our results and our main conclusions are presented in section VI. ",
        "conclusions": "In RR09 we have presented the first numerical simulations of the interaction of the solar wind with the plasma tail of a comet including the effects of viscous-like forces  as originally proposed by Perez-de-Tejada et al. (1980). In this paper we have carried out a comparison of the tailward velocity profiles of the cometary ions derived in our numerical simulations, with the kinematics of H$_2$O+ ions determined from spectroscopic, ground based observations along the tail of comet Swift-Tuttle (Spinrad et al. 1994).  Our main result is that the case we have chosen as fiducial in a previous study, on the basis of a comparison of our model with {\\it in situ} measurements at comets Halley and Giacobinni-Zinner (RR09), characterized by a low Reynolds number ($Re^{a,b} = 30$) and moderate interspecies coupling parameter ($\\nu_o = 0.1$), gives one of the best fits to the flow kinematics observed by Spinrad et al. (1994).  Cases with a much larger Reynolds number (smaller viscosity) are ruled out by the observations. Hence, on the basis of these results, we conclude that an efficient viscous-like momentum transfer between the solar wind and cometary material in the plasma tail, is necessary to explain the observed kinematics in comet Swift-Tuttle. The fact that a similar conclusion is reached from the analysis at comets Halley and Giacobinni-Zinner (RR09) suggests that viscous-like momentum transfer may be a dynamically important process in the interaction of the solar wind and the plasma environment of comets in general.  We finalize by pointing out several issues still to be addressed, that may have important consequences on the results of our simulations and on the conclusions of this study. Our simplified treatment does not consider the ongoing creation of H$_2$O+ ions in the tail region due to the ionization of cometary neutrals. At this stage in our ongoing modeling effort, we have also neglected the effect of the draped IMF on the flow dynamics (although as discussed in RR09 we do not expect these to be dominant), 3D effects, and time dependence of the incoming flow. The effect of the precise form of the viscous-like transport and effective interspecies coupling terms in the equations of motion, is another pending issue. We believe that the further assessment of the relevance of these factors is beyond the present study. They are the subject of work currently in progress and will be reported in future contributions."
    },
    "0911/0911.2496_arXiv.txt": {
        "abstract": "In this work we present a new and efficient Bayesian method for nonlinear three dimensional large scale structure inference. We employ a Hamiltonian Monte Carlo (HMC) sampler to obtain samples from a multivariate highly non-Gaussian lognormal Poissonian density posterior given a set of observations. The HMC allows us to take into account the nonlinear relations between the observations and the underlying density field which we seek to recover. As the HMC provides a sampled representation of the density posterior any desired statistical summary, such as the mean, mode or variance, can be calculated from the set of samples. Further, it permits us to seamlessly propagate non-Gaussian uncertainty information to any final quantity inferred from the set of samples. The developed method is extensively tested in a variety of test scenarios, taking into account a highly structured survey geometry and selection effects. Tests with a mock galaxy catalog based on the millennium run show that the method is able to recover the filamentary structure of the nonlinear density field. The results further demonstrate the feasibility of non-Gaussian sampling in high dimensional spaces, as required for precision nonlinear large scale structure inference. The HMC is a flexible and efficient method, which permits for simple extension and incorporation of additional observational constraints. Thus, the method presented here provides an efficient and flexible basis for future high precision large scale structure inference. ",
        "introduction": "Modern large galaxy surveys allow us to probe cosmic large scale structure to very high accuracy if the enormous amount of data can be processed and analyzed efficiently. Especially, precision reconstruction of the three dimensional density field from observations poses complex numerical challenges. For this reason, several reconstruction methods and attempts to recover the underlying density field from galaxy observations have been presented in literature \\citep[see e.g.][]{1989ApJ...336L...5B,1991ASPC...15...67B,1994ApJ...423L..93L,HOFFMAN1994,1995MNRAS.272..885F,BISTOLAS1998,WEBSTER1997,1999AJ....118.1146S,2002MNRAS.331..901Z,2004MNRAS.352..939E,KITAURA2009,JASCHE2009}. Recently, \\citet{KITAURA2009} presented a high resolution Wiener reconstruction of the Sloan Digital Sky Survey (SDSS) matter density field, and demonstrated the feasibility of precision large scale structure analysis. The Wiener filtering approach is based on a linear data model, which takes into account several observational effects, such as survey geometry, selection effects and noise \\citep{kitaura,KITAURA2009,JASCHE2009}. Although, the Wiener filter has proven to be extremely efficient for three dimensional matter field reconstruction, it still relies on a Gaussian approximation of the density posterior. While this is an adequate approximation for the largest scales, precision recovery of nonlinear density structures may require non-Gaussian posteriors. Especially, the detailed treatment of the non-Gaussian behavior and structure of the Poissonian shot noise contribution may allow for more precise recovery of poorly sampled objects. In addition, for a long time it has been suggested that the fully evolved nonlinear matter field can be well described by lognormal statistics \\citep[see e.g.][]{HUBBLE1934,PEEBLES1980,COLES1991,GAZTANAGA1993,KAYO2001}. These discussions seem to advocate the use of a lognormal Poissonian posterior for large scale structure inference. Several methods have been proposed to take into account non-Gaussian density posteriors \\citep[see e.g.][]{SAUNDERS2000,kitaura,ENSSLIN2008,KJ2009}.  However, if the recovered nonlinear density field is to be used for scientific purposes, the method not only has to provide a single estimate, such as a mean or maximum a postiori reconstruction, but it should also provide uncertainty information, and the means to nonlinearly propagate this uncertainty to any final quantity inferred from the recovered density field.  For this reason, here we propose a new Bayesian method for nonlinear large scale structure inference. The developed computer program HADES (HAmiltonian Density Estimation and Sampling) explores the posterior distribution via an Hamiltonian Monte Carlo (HMC) sampling scheme. Unlike conventional Metropolis Hastings algorithms, which move through the parameter space by a random walk, and therefore require prohibitive amounts of steps to explore high dimensional spaces, the HMC sampler suppresses random walk behavior by introducing a persistent motion of the Markov chain through the parameter space \\citep[][]{DUANE1987,NEAL1993,NEAL1996}. In this fashion, the HMC sampler maintains a reasonable efficiency even for high dimensional problems \\citep{HANSON2001}. The HMC sampler has been widely used in Bayesian computation \\citep[see e.g.][]{NEAL1993}. In cosmology it has been employed for cosmological parameter estimation and CMB data analysis \\citep{HAJIAN2007,TAYLOR2008}.  In this work we demonstrate that the HMC can efficiently be used to sample the lognormal Poissonian posterior even in high dimensional spaces. In this fashion, the method is able to take into account the nonlinear relationship between the observation and the underlying density which we seek to recover. The scientific output of the HMC is a sampled representation of the density posterior. For this reason, any desired statistical summary such as mean, mode and variance can easily be calculated from the HMC samples. Further, the full non-Gaussian uncertainty information can seamlessly be propagated to any finally estimated quantity by simply applying the according estimation procedure to all samples. This allows us to estimate the accuracy of conclusions drawn from the analyzed data.  In this work, we begin, in section \\ref{PDF_LOGNORMAL}, by presenting a short justification for the use of the lognormal distribution as a prior for nonlinear density inference, followed by a discussion of the lognormal Poissonian posterior in section \\ref{PDF_LOGPOIS}. Section \\ref{HAMILTONIAN_SAMPLING} outlines the HMC method. In section \\ref{equations_of_motion} the Hamiltonian equations of motion for the lognormal Poissonian posterior are presented. Details of the numerical implementation are described in section \\ref{Numerical_implementation}. The method is extensively tested in section \\ref{HADES_testing} by applying HADES to generated mock observations, taking into account a highly structured survey geometry and selection effects. In section \\ref{Discussion} we summarize and conclude. ",
        "conclusions": "\\label{Discussion} In this work we introduced the Hamiltonian Monte Carlo sampler for nonlinear large scale structure inference and demonstrated its performance in a variety of tests. As already described above, according to observational evidence and theoretical considerations, the posterior for nonlinear density field inference is adequately represented by a lognormal Poissonian distribution, up to overdensities of \\(\\delta \\sim 100\\). Hence, any method aiming at precision estimation of the fully evolved large scale structure in the Universe needs to handle the nonlinear relation between observations and the signal we seek to recover. The Hamiltonian Monte Carlo sampler, presented in this work, is a fully Bayesian method, and as such tries to evaluate the lognormal Poissonian posterior, given in equation \\ref{eq:LOGNORMALPOISSONIAN_POSTERIOR}, via sampling. In this fashion, the scientific output of the method is not a single estimate, but a sampled representation of the multidimensional posterior distribution. Given this representation of the posterior any desired statistical summary, such as mean, mode or variances can easily be calculated. Further, any uncertainty can seamlessly be propagated to the finally estimated quantities, by simply applying the according estimation procedure to all Hamiltonian samples.  Unlike conventional Metropolis Hastings algorithms, which move through the parameter space by random walk, the Hamiltonian Monte Carlo sampler suppresses random walk behavior by following a persistent motion. The HMC exploits techniques developed to follow classical dynamical particle motion in potentials, which, in the absence of numerical errors, yield an acceptance probability of unity. Although, in this work we focused on the use of the lognormal Poissonian posterior, the method is more general. The discussion of the Hamiltonian sampler in section \\ref{HAMILTONIAN_SAMPLING}, demonstrates that the method can in principle take into account a broad class of posterior distributions.  In section \\ref{HADES_testing}, we demonstrated applications of the method to mock test cases, taking into account observational uncertainties such as selection effects, survey geometries and noise. These tests were designed to study the performance of the method in real world applications.  In particular, it was of interest to establish intuition for the behavior of the Hamiltonian sampler during the initial burn-in phase. Especially, the required amount of samples before the sampler starts drawing samples from the correct posterior distribution was of practical relevance. The tests demonstrated, that for a realistic setup, the initial burn-in period is on the order of \\(\\sim 100\\) samples.  Further, the tests demonstrated that the Hamiltonian sampler produces unbiased samples, in the sense that each sample possesses correct power. Unlike a filter, which suppresses the signal in low signal to noise regions, the Hamiltonian sampler nonlinearly augments the poorly or not observed regions with correct statistical information. In this fashion, each sample represents a complete matter field realization consistent with the observations.  The convergence of the Markov chain was tested via a Gelman\\&Rubin diagnostic. We compared the intra and inter chain variances of two Markov chains each of length \\(20000\\) samples. The estimated PSRF indicated good convergence of the chain. This result demonstrates, that it is possible to efficiently sample from non-Gaussian distributions in very high dimensional spaces.  In a final test the method was applied to a realistic galaxy mock observation based on the millennium run \\citep{CROTON2006}. Here we introduced again survey geometry and selection effects and generated \\(20000\\) samples of the lognormal Poissonian posterior. The results nicely demonstrate that the Hamiltonian sampler recovers the filamentary structure of the underlying matter field realization. For this test we also calculated the ensemble mean and the corresponding ensemble variance of the Hamiltonian samples, demonstrating that the Hamiltonian sampler also provides error information for a final estimate.  To conclude, in this paper we present a new and numerically efficient Bayesian method for large scale structure inference, and its numerical implementation HADES. HADES provides a sampled representation of the very high dimensional non-Gaussian large scale structure posterior, conditional on galaxy observations. This permits us to easily calculate any desired statistical summary, such as mean, mode and variance. In this fashion HADES is able to provide uncertainty information to any final quantity estimated from the Hamiltonian samples. The method, as presented here, is very flexible and can easily be extended to take into account additional nonlinear observational constraints and joint uncertainties.  In summary, HADES, in its present form, provides the basis for future nonlinear high precision large scale structure analysis."
    },
    "0911/0911.4753_arXiv.txt": {
        "abstract": "We have used over 10,000 early-type galaxies from the 6dF Galaxy Survey (6dFGS)  to construct the Fundamental Plane across the optical and near-infrared passbands. We demonstrate that a maximum likelihood fit to a multivariate Gaussian model for the distribution of galaxies in size, surface brightness and velocity dispersion can properly account for selection effects, censoring and observational errors, leading to precise and unbiased parameters for the Fundamental Plane and its intrinsic scatter. This method allows an accurate and robust determination of the dependencies of the Fundamental Plane on variations in the stellar populations and environment of early-type galaxies. ",
        "introduction": "One of the most important empirical scaling relations of early-type galaxies is the Fundamental Plane - a  tight correlation in the log space of effective radius $R_{e}$, central velocity dispersion $\\sigma$,  and mean effective surface brightness $\\langle \\mu_{e} \\rangle$  (\\cite[Dressler et al. 1987]{Dressler_etal87}).  Fundamental Plane studies aim to understand what physical processes constrain early-type galaxies to this plane rather than being uniformly scattered.  However such studies raise as many questions as they answer, not least being how the Fundamental Plane varies with stellar population and environment, and the origin of the observed \\emph{tilt} of the plane with respect to virial expectations.  These questions will be addressed in our analysis of early-type galaxies from the 6dF Galaxy Survey.  As well as determining the Fundamental Plane across a range of masses and environments, we will in future determine precise Fundamental Plane distances and peculiar velocities with which to map the local peculiar velocity field out to $\\sim$16,500 km\\,s$^{-1}$.  Recent studies have found various correlations of the Fundamental Plane with age and stellar population (\\cite[Graves et al. 1987]{Graves_etal09}, \\cite[Gargiulo et al. 2009]{Gargiulo_etal09}).  Therefore, as part of the 6dFGS Fundamental Plane study we will measure Lick indices to give ages, metallicities and $\\alpha$-enhancements and quantify the effect of stellar populations on the Fundamental Plane.  We will investigate age-metallically effects and observe any trends in the mass-to-light ratio along the Fundamental Plane. \\clearpage ",
        "conclusions": ""
    },
    "0911/0911.1947_arXiv.txt": {
        "abstract": "{}{To understand stellar magnetism and to test the validity of the Babcock-Leighton flux transport mean field dynamo models with stellar activity observations} {2-D mean field dynamo models at various rotation rates are computed with the STELEM code to study the sensitivity of the activity cycle period and butterfly diagram to parameter changes and are compared to observational data. The novelty is that these 2-D mean field dynamo models incorporate scaling laws deduced from 3-D hydrodynamical simulations for the influence of rotation rate on the amplitude and profile of the meridional circulation. These models make also use of observational scaling laws for the variation of differential rotation with rotation rate.} {We find that Babcock-Leighton flux transport dynamo models are able to reproduce the change in topology of the magnetic field (i.e. toward being more toroidal with increasing rotation rate) but seem to have difficulty reproducing the cycle period vs activity period correlation observed in solar-like stars if a monolithic single cell meridional flow is assumed. It may however be possible to recover the $P_{cyc}$ vs $P_{rot}$ relation with more complex meridional flows, if the profile changes in a particular assumed manner with rotation rate.} { The Babcock-Leighton flux transport dynamo model based on single cell meridional circulation does not reproduce the $P_{cyc}$ vs $P_{rot}$ relation unless the amplitude of the meridional circulation is assumed to increase with rotation rate which seems to be in contradiction with recent results obtained with 3-D global simulations.} ",
        "introduction": "The Sun is a magnetic star that possesses a well defined cyclic activity period of 22 years. It is believed that the origin of this magnetism and regular activity is linked to a global scale dynamo operating in and at the base of the solar convection zone (Parker 1993). In order to model and understand this global scale dynamo, a useful approach has been to make use of the mean field dynamo theory (\\cite{Mof78}; Krause \\& Raedler 1980). This method has the advantage that it only deals with the large scale magnetic field, assuming some parameterization of the underlying small scale turbulence and magnetism. Computing those physical processes self-consistently would otherwise require costly full 3-D magnetohydrodynamical (MHD) simulations (Cattaneo 1999; Brun, Miesch \\& Toomre 2004). Among the various mean field dynamo models, the Babcock-Leighton flux transport dynamo models have recently been the favored ones and have demonstrated some success at reproducing solar observations assuming either an advection or a diffusion dominated regime for the transport of field from the surface down to the tachocline (Dikpati \\& Charbonneau 1999; Nandy \\& Choudhuri 2002; Dikpati et al. 2004; Chatterjee, P. et al. 2004; Charbonneau 2005; Bonanno et al. 2005; Jouve \\& Brun 2007; Yeates et al. 2008).  It is legitimate to ask if such models can be applied to other stars and lead to the same agreement with the observations of for instance solar type stars that may possess different rotation rates and activity levels.  \\subsection{Observational Constraints}  Observations of the magnetism of 'solar type' stars, i.e. stars possessing a deep convective envelope and a radiative interior (late F, G, K and early M spectral type) are becoming more and more available (Giampapa 2004) and may provide additional constraints on our understanding of global scale stellar dynamos. One difficulty of such observational programs is that they require long term observations since stellar cycle periods are likely to be commensurate to the solar 11-yr sunspots cycle period (or 22-yr for a full cycle including two polarity reversals of the global poloidal field). Thanks to the data collected at Mount Wilson Observatory since the late 60's, such data is available (Wilson 1978; Baliunas et al. 1995). Among the sample of 111 stars (including the Sun as a star) originally observed between the F2 to M2 spectral type, it is found that about 50\\% of the stars possess a cyclic activity, with cycle (starspot) periods varying roughly between 5 to 25 yrs, i.e. between half to twice the sunspot cycle period. They further indicate that among the inactive stars of the sample some are likely to be in a quiet phase (as was the Sun during the Maunder minimum). Overall activity cycles seems to be more frequent for less massive stars K than for F stars. More recent observational programs have been pursued that now even provide information on the field topology as a function of the rotation rate, such as the one using the Espadons and Narval instruments and the Zeeman Doppler Imaging technique (Donati et al. 2006). Applying this observational technique over a sample of four solar analogues with rotation rate varying from one to three times solar, Petit et al. (2008) have shown that the field amplitude increases as a function of the star's rotation rate and, more importantly, becomes more and more dominated by its toroidal component (modulo possible bias in the observational technique used). If such a trend is confirmed, i.e. that the field topology is becoming more toroidal with increasing rotation rate, it is a very important and instructive result and puts strong constraints on the dynamo models.  The systematic analysis of stellar magnetism data revealed that for solar type stars there is a good correlation between the cycle and rotation periods of the stars and that correlation is even stronger when using the Rossby number ($Ro = P_{rot}/\\tau$) that takes into account the convection turnover time $\\tau$ at the base of the stellar convective envelope (Noyes et al. 1984; Soon et al. 1993; Hempelmann et al. 1995; Baliunas et al. 1996). As the star rotates faster, its cycle period is found to be shorter.  Typically Noyes et al. (1984) found that $P_{cyc} \\propto P_{rot}^n$, with $n = 1.25 \\pm 0.5$. Based on an extended stellar sample Saar \\& Brandenburg (1999) and Saar (2002) have argued that there is actually two branches when plotting the cycle period vs the rotation period of the stars. They make the distinction between the primary (starspot) cycle and Gleissberg or grand minima type modulation of the stellar activity. For the active branch they found an exponent $n \\sim 0.8$ and for the inactive stars $n\\sim 1.15$ (Charbonneau \\& Saar 2001). It is also found that this correlation breaks at high rotation rate with the possible appearance of a super active branch. Further at very high rotation rate saturation of the X-ray luminosity seems to limit the validity of the scaling found at more moderate rotation rates (Pizzolato et al. 2003). For G type stars this saturation is found for rotation rate above 35 ${\\bf \\rm km s^{-1}}$, for K type stars at about 10 ${\\bf \\rm km s^{-1}}$ and for M dwarfs around 3-4 ${\\bf \\rm km s^{-1}}$ (Browning 2008; Reiners, Basri \\& Browning 2009). How stellar magnetic flux scales with rotation rate is thus also important to understand since it is telling us how the magnetic field generated by dynamo action inside the stars emerges and imprints the stellar surface (Rempel 2008) and if it actually saturates.  \\subsection{Possible theoretical interpretation}  Theoretical considerations to interpret stellar magnetism based on classical mean field $\\alpha$-$\\omega$ dynamo models  (Durney \\& Latour 1976, Knobloch et al. 1981; Baliunas et al. 1996) naturally find correlation between rotation rate and stellar activity. In particular it is found that both magnetic field generation and the dynamo number $D$ (i.e. a Reynolds number characterizing the mean field $\\alpha$ and $\\omega$ dynamo effects used in the models) vary with the rotation period of the star. This is due to the fact that in these models both effects are sensitive to the rotation rate of the star. The $\\omega$-effect is a direct measure of the differential rotation $\\Delta \\Omega$ established in the star. It is well known both theoretically and observationally that the differential rotation in the convective envelope of solar-type stars is directly connected to the star's rotation rate $\\Omega_0$ (Donahue et al. 1996; Barnes et al. 2005; Ballot et al. 2007; Brown et al. 2008; K\\\"uker \\& R\\\"udiger 2008). However, the exact scaling $n_r$ (i.e. $\\Delta \\Omega \\propto \\Omega_0^{n_r}$) is still a matter of debate among the observers and theoreticians, being sensitive to the observational techniques used and to the modelling approach.  The $\\alpha$-effect used in mean field theory is related to helical turbulence and represents a measure or parameterization of the mean electromotive force (emf) (Moffatt 1978; Pouquet et al. 1976). It is thus also linked to the rotation rate of the star and the amount of kinetic helicity present in its convective envelope. However in Babcock-Leighton flux transport dynamo models the standard $\\alpha$-effect is replaced by a surface term linked to the tilt of the active regions with respect to the east-west direction (i.e. Joy's law, Kosovichev \\& Stenflo 2008). This tilt is thought to be due to the action of the Coriolis force during the rise of the toroidal structures that emerge as active regions (D'Silva \\& Choudhuri 1993). Recent 3-D simulations in spherical shells (Fan 2008; Jouve \\& Brun 2007b, 2009) seem to indicate that this is not the only effect responsible for the observed tilt and that the twist and arching of the toroidal structures as well as the continuous action of the surface convection during the emergence have some influence on the resulting tilt. Thus the dependence with rotation rate of the Babcock-Leighton surface source term used in this class of mean field dynamo models is not straightforward to assess and needs to be studied in detail.  Another important ingredient in flux transport models is the large scale meridional circulation (MC) used to connect the surface source term generating the poloidal field to the region of strong shear at the base of the convection zone (i.e. the tachocline) where it will be subsequently sheared by the $\\omega$-effect in order to close the global dynamo loop (i.e. $B_{pol} \\rightarrow B_{tor} \\rightarrow B_{pol}$) . The meridional flow (or \"conveyor belt\") thus plays an important role in setting the cycle period of the global dynamo in this class of flux transport model. As a direct consequence it is natural to ask how the meridional circulation amplitude and profile change with the rotation rate and how these may influence the magnetic cycle period.  Early work by Dikpati et al. (2001) have assumed that the amplitude of the meridional flow is proportional to the rotation rate of the star. With such scaling they seem to be able to reproduce some of the rotation vs cycle period correlation observed in solar type stars. Charbonneau \\& Saar (2001) have also studied the sensitivity of various mean field dynamo models in reproducing stellar activity observations assuming for the Babcock-Leighton type that the meridional flow amplitude increases either as $\\Omega_0$ or as $\\log(\\Omega_0)$. Similarly Nandy (2004) and Nandy \\& Martens (2007) have explored the solar-stellar connection with B-L mean field dynamo models and have",
        "conclusions": "The idea of this work was to better understand stellar magnetism, using both 2-D and 3-D calculations. To do so, we have introduced into Babcock-Leighton (B-L) flux transport dynamo models, results from 3-D global hydrodynamical simulations  by Brown et al. (2008). To be more specific we have used these simulations to determine the variation of the amplitude of the meridional circulation with rotation rate and complemented our study with observations of the variation of the angular velocity contrast $\\Delta \\Omega$ with rotation rate. Using these scaling laws  we have built a series of 2-D mean field stellar dynamo models in which we have modified in turn the rotation rate, the amplitude of the meridional circulation and of the angular velocity gradient. This has allowed us to test the validity of this type of models on solar-like stars with faster rotations and intense magnetism.  Our study shows that an increase in the rotation rate leads to a larger $\\omega$-effect which in turn modifies the ratio between poloidal and toroidal fields with a trend similar to that found observationally in a small sample of rapidly rotating suns (Petit et al. 2008). Other authors have also studied such trends using solar mean field B-L flux transport dynamo models (see for instance Charbonneau \\& Saar 2001; Dikpati et al. 2001; Nandy \\& Martens 2007). They also find that the toroidal field grows in amplitude with the star's rotation rate. The difference between our work and previous studies comes from the choice made in varying the meridional circulation amplitude with rotation rate. In 3-D global calculations an increase of the rotation rate leads to a decrease in the meridional flow amplitude whereas in previous studies the opposite was assumed since such knowledge was not yet available. Given that the cycle period in advection-dominated B-L dynamo models is almost inversely proportional to the intensity of the meridional circulation, we find that our models predict slower magnetic cycles for more rapidly rotating stars. This is not consistent with the observational data collected at Mount Wilson since the 1960's (e.g. Noyes et al. 1984, Baliunas et al. 1995). It thus seems difficult to reconcile advection-dominated B-L models widely used in the solar dynamo community with stellar activity observations if we use scaling laws relating rotation rates and meridional flow amplitudes given by consistent 3-D models. Of course the systematic observation of the meridional circulation realized in solar-like stars rotating at various rotation rate would be of great use and could solve the apparent discrepancy if we indeed found that such flows are not so sensitive to the rotation rate. Unfortunately such observations are not yet available. Hopefully with the Corot and Kepler satellites now in orbit this situation may change in the near future.  We have thus tried to investigate ways to improve the agreement between 2-D models and observations in refining our calculations. A least-square fit on the various parameters of the standard model shows that we cannot rely on a modifications of $\\Omega_0$, $s_0$ and $\\eta$ to act on the cycle period, as the latter is almost completely determined by the amplitude of the meridional flow (see Equation 12). We thus computed more sophisticated models where the differential rotation and meridional circulation profiles were modified.  First, we computed a model at 5 times the solar rotation rate where the modified contrast in the angular velocity was taken into account. Some uncertainties both from an observational or a theoretical standpoint still exist concerning the scaling between $\\Delta\\Omega$  and $\\Omega_0$ but we find that the effect on the cycle period is negligible, even when we consider the most extreme observation of the variation of the differential rotation contrast with rotation rate (Donahue et al. (1996)). Modifying the tachocline thickness as a function of rotation rate according to the scaling law of Brun et al. (1999) does not help to improve the relationship between the rotation rate of the star and the cycle period either. We show that the main impact of a modification of the profile of differential rotation is to decrease the percentage of the poloidal component in the total magnetic field again because of the increased efficiency of the source for toroidal field: the $\\omega$-effect. The modification of the source term for poloidal field (decay of active regions in Babcock-Leighton models and the kinetic helicity of the flow in $\\alpha$-$\\omega$ models) due to an increase of the rotation rate still needs to be addressed to study the evolution of the poloidal to toroidal field ratio in those refined cases.  Given the strong dependency of the magnetic cycle period on the meridional circulation and the fact that 3-D models exhibit multicellular patterns of the flow in latitude, it is reasonable to think that modifying the profile of the MC according to 3-D models may help to get closer to the observational data. Indeed, we show that adding a counter cell in latitude which extends until about $20\\degr$ in both hemispheres reduces drastically the cycle period. We find a quasi-linear relationship between the extension of the counter cell and the cycle period, which is very promising to reconcile Babcock-Leighton models with observations of stellar activity. Another solution could be to change the way poloidal magnetic fields are transported from the surface down to the tachocline. In Yeates et al. (2008) the transport is dominated  by diffusion to link the distant magnetic field generation regions. In their case the cycle period seems less sensitive to the meridional flow amplitude and profile assumed and depends more on the diffusion profile. Guerrero \\& Gouveia Dal Pino (2008) have introduced in Babcock-Leighton dynamo equations an extra term representing the pumping of magnetic fields due to turbulent convection (inspired by the work of K\\\"{a}pyl\\\"{a} et al. 2006). This new process leads to vertical and horizontal transport of the fields. Their model seems to show a significant decrease of the cycle period dependency on the meridional flow amplitude (from $v_0^{-0.9}$ to $v_0^{-0.12}$, although only a shallow meridional flow is considered, which may influence the scaling law) and that the turbulent pumping is now responsible for regulating the cyclic activity. Reintroducing the meridional flow amplitude scalings deduced from 3-D calculations in this type of models would thus have much less impact on the cycle period. The changes in the pumping coefficients with respect to the rotation rate would have to be investigated to see if some of the observations of stellar magnetic activity could be recovered.  In conclusion, advection-dominated Babcock-Leighton flux-transport models may offer insights into the dynamo action occurring within the convection zones of solar-type stars.  However, it is necessary to make serious modifications to these models if we wish to reconcile their predictions with observations of real stellar magnetism."
    },
    "0911/0911.4086_arXiv.txt": {
        "abstract": "Directional detection of Galactic Dark Matter is a promising  search strategy for discriminating WIMP events from  background. Technical progress on gaseous detectors and read-outs has permitted the design and construction of competitive experiments. However,  to take full advantage of this powerful detection method, one need to be able to extract information from an observed recoil map to identify a WIMP signal. We present a comprehensive formalism, using a map-based likelihood method allowing to recover the main incoming direction of the signal and its significance, thus  proving  its Galactic origin. This is a blind analysis intended to be used on any directional data. Constraints are deduced in the ($\\sigma_n, m_\\chi$) plane and systematic  studies are presented in order to show that, using this analysis tool, unambiguous Dark Matter detection can be achieved on a large range of exposures and background levels. ",
        "introduction": "\\label{sec:intro}  A substantial body of astrophysical evidence supports the existence of non-baryonic Dark Matter :  locally, from the rotation curves of spiral galaxies \\cite{rubin}, on cluster scale from the observation of galaxy cluster collisions \\cite{clowe} and on the largest scale from  cosmological observations \\cite{wmap}. Candidates for this class of   to-be-discovered particles (WIMP: Weakly Interacting Massive Particles) naturally arise from extensions of the standard model of particle physics, such as supersymmetry~\\cite{susydm} or extra-dimensions~\\cite{kkdm}.\\\\ Like most spiral galaxies, the Milky Way is supposed to be immersed in a halo of WIMPs which outweighs the luminous component by at least an order of magnitude. Tremendous experimental efforts on a host of techniques have been made in the field of direct search of WIMP Galactic Dark Matter. With really small expected event rates (${\\cal O} \\rm (10^{-5}-1) \\ day^{-1}kg^{-1}$), the challenge is to  distinguish  a genuine WIMP signal from  background events. Two main Dark Matter search strategies may be identified : the detector may be designed to reach an extremely  high level of background rejection \\cite{cdms,revue} and/or to provide an unambiguous positive WIMP signal. The second strategy can be applied by searching for a correlation of the WIMP signal either with the motion of the Earth around the Sun, observed as an annual modulation of the number of events \\cite{freese}, or with the direction of the Solar motion around the Galactic center \\cite{spergel}, which happens to be roughly in the direction of the Cygnus constellation. The latter is generally referred to as directional detection of Dark Matter and several project of detectors are being developed for this goal \\cite{Drift,newage,MIMAC,mit,white}.\\\\ Directional detection of Dark Matter has first been proposed in \\cite{spergel}, highlighting the fact that a strong forward/backward asymmetry is expected in the case of an isothermal spherical Galactic halo, which could, in principle, be observed   even with a poor angular resolution. Since then, several phenomenological studies have been performed \\cite{copi1,copi2,copi3,morgan1,morgan2,green1,green2, host,vergados,gondolo}. Using an unbinned likelihood  method \\cite{copi1,copi2,copi3} or non-parametric statistical tests on unbinned data \\cite{morgan1,morgan2,green1,green2}, it has been shown that a few number of events  ${\\cal O} (10)$ is required to reject the isotropy, which is characteristic of the background event distribution. However, these methods do not allow to identify and characterize a WIMP signal and hence do not take full advantage of data from upcoming directional detectors \\cite{white}.    In this Letter,  we use a map-based likelihood analysis in order to extract from an observed recoil map the main incoming direction of the events and its significance. In this way, the Galactic origin of the signal can be proved by showing its correlation with the direction of the Solar motion with respect to the Dark Matter halo.  Hence, the goal of this new method is not to reject the background hypothesis, but rather to identify a genuine WIMP signal. This method is tested on realistic simulated cases defined by a low number of WIMP events ${\\cal O} (10)$ hidden  by a non-negligible background   and measured with a rather low angular resolution $\\sim 15^\\circ$ (FWHM) \\cite{santos}. This blind analysis is intended to be applied to directional data of any detector.\\\\ The outline of this Letter is the following : after introducing the directional detection framework and phenomenology, the  map-based likelihood analysis is presented and applied to a realistic simulated recoil map. Main results are presented and  constraints are deduced in the  $(\\sigma_n, m_\\chi)$  plane. Eventually, systematic  studies are presented in order to show that this analysis tool gives satisfactory results on a large range of exposures and background levels. ",
        "conclusions": "\\label{sec:conclusion} We have presented a powerful statistical analysis tool to extract information from  a  data sample of a directional detector in order to identify a Galactic WIMP signal. As a proof of principle, it has been tested within the framework of an isothermal spherical halo model. We have shown the feasibility to extract from an observed map the main incoming direction of the signal and its significance, thus  proving  its Galactic origin. Evidence in favor of Galactic Dark Matter may be within reach of upcoming directional detectors even at low exposure. In a second step, with increasing exposure and by using an extended method currently under development, Dark Matter properties (triaxility, halo rotation, WIMP mass) may also be constrained.\\\\   For this first study of directional detection as a strategy to discover Galactic Dark Matter, the standard halo model has been considered. It may be noticed that   recent results from high resolution numerical simulations seem to be in favour of non-smooth WIMP velocity distributions \\cite{vogelsberger,kuhlen}. However, as their resolution is many orders of magnitude larger than the scale of the ultra-local Dark Matter distribution probed by current and future detectors, this result is not fully relevant for direct detection. Even if applied at ultra-local scale, the result from \\cite{kuhlen} is that high velocity WIMPs ($v > 500$~km.s$^{-1}$) could deviate from the Solar motion direction ($\\ell_\\odot ,  b_\\odot $) within about 10$^{\\circ}$ due to the presence of subhalos, tidal streams or an anisotropic velocity distribution. But, as the minimal speed is 130~km.s$^{-1}$ in this study and as the Maxwellian distribution of WIMP velocity is peaked around 300~km.s$^{-1}$ the contribution of high speed WIMPs to the WIMP flux is negligible. For instance, this high speed feature would be present in data from detector with a higher energy threshold. In such case, this will mildly affect the result presented in this Letter by shifting the main recoil direction recovered from the likelihood analysis by a few degrees. However, the discovery proof will  still be  reached as long as the WIMP signal is still directional, thus  remaining different from the background.  Effect of non-standard halo models (triaxial, with stochastic features or streams) will be addressed in a dedicated forthcoming paper."
    },
    "0911/0911.3942_arXiv.txt": {
        "abstract": "As is common for antenna arrays in radio astronomy, the output of the MWA's correlator is the intensity measured in visibility space.  In addition, the final power spectrum will be created in visibility space.  As such, correcting for the ionosphere in visibility space instead of real space saves the computation required to inverse Fourier transform to real space and then Fourier transform back (a significant decrease in computation for systems operating in real time such as the MWA.)  In this paper, we explore this problem of correcting for ionospheric distortions in the $uv$-plane.  The mathematical formula for obtaining the unperturbed data from that reflected by the ionosphere is non-local, which in any practical application creates edge effects because of the finite nature of the $uv$-plane (section \\ref{ic}).  In addition, obtaining an analytic solution for the unperturbed intensity is quite difficult, and can only be done using very specific expansions of ionospheric perturbations.  We choose one of these models (with perturbations as sinusoidal modes, section \\ref{sinmodes}) and run numerical codes to further study the correction.  Numerically implementing this correction to too few orders distorts the data in such a way as to be worse than not correcting at all (section \\ref{nmax}).  It is therefore critical to correct to the correct number of orders, and we present an analytic estimate for the optimal order (section \\ref{theorynmax}).  This analytic estimate shows that the optimal number of orders varies with $\\vec{u}$, and in particular increases as $\\vec{u}$ increases along the direction of an ionospheric distorting mode.  Based on this observation, we then investigate a couple of methods which save computation (section \\ref{feasible}).  These methods are (a) eliminating the intensity at values of $\\vec{u}$ which require too many orders, and (b) correcting to different orders at different $\\vec{u}$.  Both methods prove successful, although the first creates a loss of some precision in the real space sky.  We conclude by considering an alternate form with which to model ionospheric perturbations (section \\ref{choice2}).  This alternate form was once again chosen because it lends itself to an analytic solution, but contains as many (if not more) downfalls than the original choice. ",
        "introduction": "\\label{back}  \\hspace{4 mm} We begin with the basic background on the ionospheric operator, followed by that of the ionospheric correction operator.  These sections expand upon the mathematical framework created and briefly outlined in Morales, et. al., \\cite{MM}.  \\subsection{Ionospheric Operator}  \\hspace{4 mm} The ionospheric operator {\\bf A}($\\vec{\\theta'}, \\vec{\\theta}$) is the operator which takes an unperturbed map of the sky $I(\\vec{\\theta})$ and maps it to a perturbed map $\\tilde{I}(\\vec{\\theta'})$ which has been distorted by the ionosphere, \\begin{equation} \\label{Adef} \\tilde{I}(\\vec{\\theta'}) = {\\bf A}(\\vec{\\theta'},\\vec{\\theta})I(\\vec{\\theta}) . \\end{equation} Put simply, {\\bf A}($\\vec{\\theta'},\\vec{\\theta}$) is a generalized coordinate change from $\\vec{\\theta}$ to $\\vec{\\theta'}$.  (The order of the arguments of {\\bf A}($\\vec{\\theta'},\\vec{\\theta}$) here, and with other operators later, is indicative of the direction of change.)  In the regime of MWA, it is a very good approximation that the mapping of angles is approximately linear with only a small deviation $\\vec{\\delta \\theta}(\\vec{\\theta}, t)$, \\begin{equation} \\label{delthetadef} \\vec{\\theta'} = \\vec{\\theta} + \\vec{\\delta \\theta}(\\vec{\\theta}, t) . \\end{equation} The perturbed intensity $\\tilde{I}(\\vec{\\theta'})$ at $\\vec{\\theta'}$ is the summed contribution from the intensities $I(\\vec{\\theta_i})$ at all $\\vec{\\theta_i}$ where this relation holds; ie, \\begin{equation} \\label{discrete} \\tilde{I}(\\vec{\\theta'}) = \\sum_i   I(\\vec{\\theta_i}) \\delta(\\vec{\\theta'} - \\vec{\\theta_i} - \\vec{\\delta \\theta}(\\vec{\\theta_i}, t)) \\end{equation} where $\\delta$ represents the Dirac delta function.  For example, if (only) $\\vec{\\theta_1}$ and $\\vec{\\theta_2}$ are mapped to $\\vec{\\theta_1'}$ according to equation \\ref{delthetadef}, then $\\tilde{I}(\\vec{\\theta'_1}) = I(\\vec{\\theta_1}) + I(\\vec{\\theta_2})$.  In the limit that $\\vec{\\theta}$ is a continuous variable, this discrete sum for $\\tilde{I}(\\vec{\\theta'})$ is altered to an integral, \\begin{equation} \\label{A1} \\tilde{I}(\\vec{\\theta'}) = \\int{ d^2\\theta A(\\vec{\\theta'},\\vec{\\theta}) I(\\vec{\\theta}) } , \\end{equation} where \\begin{equation} \\label{A2} A(\\vec{\\theta'},\\vec{\\theta}) = \\delta(\\vec{\\theta'}-\\vec{\\theta}-\\vec{\\delta\\theta}(\\vec{\\theta},t)) . \\end{equation} (Notice that we've used that the magnitude of the delta-function's argument's gradient is approximately 1 here).  From this equation it is evident that in the limit of continuous $\\vec{\\theta}$, the ionospheric operator equation becomes \\begin{equation} \\label{Aop} \\tilde{I}(\\vec{\\theta'}) = {\\bf A}(\\vec{\\theta'},\\vec{\\theta})I(\\vec{\\theta})  = \\int{ d^2\\theta A(\\vec{\\theta'},\\vec{\\theta}) I(\\vec{\\theta}) } . \\end{equation}  So far our calculations have been confined to real space, using the variables \\{ $\\vec{\\theta'}, \\vec{\\theta}$ \\}.  However, as is common for antenna arrays in radio astronomy, the output of the correlator for MWA will actually be the Fourier transform of the real space sky intensity, $I(\\vec{u})$.  In addition, the final power spectrum of the sky will also be measured in this Fourier transfrom space (also known as \\emph{visibility space}, or \\emph{$uv$-space}).  Therefore, correcting for the ionosphere in real space requires inverse Fourier transforming to real space, making the correction, and then Fourier transforming back to visibility space.  The problem with this is that Fourier transforming is computationally expensive, especially for a system operating in real time, such as the MWA's \\emph{Real Time System} \\cite{DM}.  As such, correcting for the ionosphere in the $uv$-plane would greatly reduce computation.  We now study the nature of this $uv$-plane correction.  Let $I(\\vec{u})$ represent the Fourier transform of $I(\\vec{\\theta})$.  Define this Fourier transform by \\begin{equation} \\label{Fourierdef} I(\\vec{u}) \\equiv {\\bf F}(\\vec{u}, \\vec{\\theta})I(\\vec{\\theta}) \\equiv \\int{ d^2 \\theta e^{- i \\vec{u} \\cdot \\vec{\\theta} } I(\\vec{\\theta}) }, \\end{equation} and the corresponding inverse Fourier transform by \\begin{equation} \\label{InverseFourierdef} I(\\vec{\\theta}) \\equiv {\\bf F}^{-1}(\\vec{\\theta},\\vec{u}) I(\\vec{u}) \\equiv \\int{ \\frac{d^2u}{(2\\pi)^2} e^{i\\vec{u} \\cdot \\vec{\\theta}} I(\\vec{u}) } . \\end{equation} Now define {\\bf A}$(\\vec{u'},\\vec{u})$ to be the ionospheric operator in the $uv-$plane; that is, the operator that maps the unperturbed map $I(\\vec{u})$ to the perturbed map $\\tilde{I}(\\vec{u'})$, \\begin{equation} \\label{Audef} \\tilde{I}(\\vec{u'}) = {\\bf A}(\\vec{u'},\\vec{u}) I(\\vec{u}) . \\end{equation}   One way to obtain $\\tilde{I}(\\vec{u'})$ from $I(\\vec{u})$ is to inverse Fourier transform $I(\\vec{u})$ to $I(\\vec{\\theta})$ using {\\bf F}$^{-1}(\\vec{\\theta}, \\vec{u})$, apply the ionospheric operator {\\bf A}$(\\vec{\\theta'}, \\vec{\\theta})$ in real space to obtain $\\tilde{I}(\\vec{\\theta'})$, and then Fourier transform $\\tilde{I}(\\vec{\\theta'})$ to $\\tilde{I}(\\vec{u'})$ using {\\bf F}$(\\vec{u'}, \\vec{\\theta'})$.  In all, \\begin{equation} \\label{Audef2} \\tilde{I}(\\vec{u'}) = {\\bf F}(\\vec{u'}, \\vec{\\theta'}) {\\bf A}(\\vec{\\theta'}, \\vec{\\theta}) {\\bf F}^{-1}(\\vec{\\theta}, \\vec{u}) I(\\vec{u}) . \\end{equation} Comparison of this to the definition of {\\bf A}$(\\vec{u'},\\vec{u})$ shows that \\begin{equation} \\label{Auopdef} {\\bf A}(\\vec{u'},\\vec{u}) = {\\bf F}(\\vec{u'}, \\vec{\\theta'}) {\\bf A}(\\vec{\\theta'}, \\vec{\\theta}) {\\bf F}^{-1}(\\vec{\\theta}, \\vec{u}) . \\end{equation} This expression may be thought of as simply a basis change of {\\bf A} from \\{$\\vec{\\theta}, \\vec{\\theta'}$\\} to \\{$\\vec{u}, \\vec{u'}$\\}.   The three operators on the right here have all been previously given.  Plugging in these predetermined expressions (equations \\ref{Aop}, \\ref{Fourierdef}, and \\ref{InverseFourierdef}) and simplifying as much as possible, we obtain \\begin{eqnarray} \\label{ } \\tilde{I}(\\vec{u'}) & = & {\\bf A}(\\vec{u'},\\vec{u}) I(\\vec{u}) \\\\ & = &  \\int{ \\frac{d^2u}{(2\\pi)^2} \\left( \\int{ d^2\\theta e^{i(\\vec{u}-\\vec{u'})\\cdot \\vec{\\theta}-i\\vec{u'} \\cdot \\vec{\\delta\\theta}} } \\right) I(\\vec{u}) } . \\end{eqnarray} Notice that the integral over $\\vec{\\theta'}$ has been evaluated by using the delta function from the expression for ${\\bf A}(\\vec{\\theta'},\\vec{\\theta})$ (see equations \\ref{A2} and \\ref{Aop}; here again we use that the magnitude of the delta-function's argument's gradient is approximately 1).  One interesting characteristic of this expression is the non-local nature, by which finding the value of the perturbed intensity $\\tilde{I}(\\vec{u'})$ at one particular $\\vec{u'}$ requires knowing the value of the pure intensity $I(\\vec{u})$ at other $\\vec{u} \\neq \\vec{u'}$.  This property will also appear in the ionospheric correction operator, found below.   \\subsection{Ionospheric Correction Operator} \\label{ic}  \\hspace{4 mm} The previous section corresponds to the operator which distorts the pure data into the perturbed data, but the reverse process is what actually interests us -- we want to correct for the effect of the ionosphere to obtain the pure data from the perturbed data.  Define {\\bf A$^{T}$}$(\\vec{\\theta}, \\vec{\\theta'})$ to be this ionospheric correction operator which corrects for the influence of the ionosphere by mapping the perturbed map of the sky $\\tilde{I}(\\vec{\\theta'})$ back to the unperturbed map $I(\\vec{\\theta})$, \\begin{equation} \\label{ } I(\\vec{\\theta}) = {\\bf A^{T}}(\\vec{\\theta},\\vec{\\theta'}) \\tilde{I}(\\vec{\\theta'}) . \\end{equation} The MWA will not run during periods of scintillation (at which times multiple values of $\\vec{\\theta}$ are perturbed to the same $\\vec{\\theta'}$), but will instead run during times when it is a very good approximation that the mapping from $\\vec{\\theta'}$ to $\\vec{\\theta}$ is one-to-one and approximately linear with only a small correction, \\begin{equation} \\label{ } \\vec{\\theta} = \\vec{\\theta'} + \\vec{\\delta \\theta'} (\\vec{\\theta'}, t) . \\end{equation} The derivation of an expression for the ionospheric correction operator in the $uv$-plane ${ \\bf A^{T} }(\\vec{u},\\vec{u'})$ follows analogously to the derivation of ${\\bf A}(\\vec{u'},\\vec{u})$, so I'll merely quote the result: {\\small \\begin{eqnarray} I(\\vec{u}) & = & {\\bf A^{T}}(\\vec{u},\\vec{u'}) \\tilde{I}(\\vec{u'}) \\\\ & = &  \\int{ \\frac{d^2u'}{(2\\pi)^2} \\left( \\int{ d^2\\theta' e^{i(\\vec{u'}-\\vec{u}) \\cdot \\vec{\\theta'}-i\\vec{u} \\cdot \\vec{\\delta\\theta'}} } \\right) \\tilde{I}(\\vec{u'}) } . \\label{IUpure} \\end{eqnarray} } % This is the most general expression for obtaining $I(\\vec{u})$ from $\\tilde{I}(\\vec{u'})$; proceeding further requires knowledge of the ionospheric perturbation $\\vec{\\delta\\theta'}$.  Numerically solving for $I(\\vec{u})$ using this equation is an incredibly daunting task for an arbitrary choice of $\\vec{\\delta\\theta'}$, as it involves a double integral over all space for every value of $\\vec{u}$.  Therefore, unless we find a choice of $\\vec{\\delta\\theta'}$ which offers an analytic solution for $I(\\vec{u})$, the correction for the ionosphere in the $uv$-plane will actually be more computationally expensive than the two Fourier transforms necessary to correct for the ionosphere in real space.  Unfortunately, choices of $\\vec{\\delta\\theta'}$ which lend themselves to analytic solutions are hard to come by.  There are a couple, however, and they will be discussed in the following sections. ",
        "conclusions": "\\hspace{4 mm}  The $uv$-plane correction only makes computational sense if the model for the ionospheric perturbation allows for an analytic solution to $I(\\vec{u})$ (equation \\ref{IUpure}).  One such model is a sum over sinusoidal modes (equation \\ref{dtsine}).  By running numerical codes with this choice, the most important result discovered was that under correcting in the $uv$-plane is worse than not correcting at all (section \\ref{nmax}).  But in addition to this, correcting to too many orders or requiring too many modes to model the effect of the ionosphere may lead to a computationally unreasonable problem (section \\ref{numterms}).  To help avoid this issue, a theoretical estimate of the number of orders of correction necessary (which agrees well with the sample sky provided in this paper) may be used (section \\ref{theorynmax}).  This estimate reveals that the number of orders of correction necessary varies in the $uv$-plane.  This, however, suggests two methods for alleviating the problem:  eliminating those points in the $uv$-plane which are particularly troublesome at the cost of precision for the real space sky (section \\ref{meth1}) and correcting to different orders at different points in the $uv$-plane (section \\ref{meth2}).  Both techniques prove successful and make the problem of correcting in the $uv$-plane more feasible.  In addition, depending on how often the ionosphere's effect is updated compared to the timescales of change in the ionosphere, it may be compuationally favorable to update the previously determined effect of the ionosphere rather than rederive its full effect from scratch each time."
    },
    "0911/0911.3804_arXiv.txt": {
        "abstract": "We analyze $V$-band photometry of the aperiodic variability in T~CrB. By applying a simple idea of angular momentum transport in the accretion disc, we have developed a method to simulate the statistical distribution of flare durations with the assumption that the aperiodic variability is produced by turbulent elements in the disc. Both cumulative histograms with Kolmogorov-Smirnov tests, and power density spectra are used to compare the observed data and simulations. The input parameters of the model $R_{\\rm in}$ and $\\alpha$ are correlated on a certain interval and the most probable values are an inner disc radius of  $R_{\\rm in} \\simeq 4 \\times 10^9$\\,cm and a viscosity of $\\alpha \\simeq 0.9$. The disc is then weakly truncated. We find that the majority of turbulent events producing flickering activity are concentrated in the inner parts of the accretion disc. ",
        "introduction": "Cataclysmic variables (CVs) are interacting binaries with a white dwarf as the primary component and a late type main sequence star as the secondary star which fills its Roche-lobe (see Warner 1995 for review). Consequently mass is transferred to the white dwarf through the inner Lagrangian point (L$_{\\rm 1}$) and an accretion disc is formed (not in the case of magnetic polars). The accretion disc is the site where several known physical mechanisms manifest themselves in the form of short-term stochastic variations in brightness, generally known as flickering. These oscillations are one of the most characteristic observational properties of accreting systems with typical amplitudes ranging from a few dozen of milli-magnitudes up to more than one magnitude on timescales from seconds to minutes. Symbiotic stars are related systems to CVs where the secondary star is an evolved red giant where the  mass loss can be due to either the wind from the giant star or the  Roche-lobe secondary star, as in \"classical\" CVs. Thus accretion discs can be formed with the associated  stochastic variations.  The first systematic study of mechanisms generating the flickering was performed  by Bruch (1992), where he compared theoretically estimated energies and timescales for individual mechanisms, with observationally derived values for a selection of CVs and related symbiotic stars. The author proposed four possible mechanisms responsible for the observed variations in brightness: (1) unstable mass transfer from the L$_{\\rm 1}$ point and interaction with the disc edge, (2) dissipation of magnetic loops, (3) turbulences in the accretion disc and (4)  unstable mass accretion onto the white dwarf. Apparently it was not enough to restrict the interpretation of the flickering in the whole family of CVs to only these four mechanisms. For example, in the case of a truncated accretion disc (intermediate polars) or in the case of the total absence of a disc (polars) the source of such variations is known to be the accretion stream threaded by the magnetic field and its subsequent interaction with the surface of the white dwarf -- accretion shock.  As mentioned earlier, Bruch (1992) performed a systematic study of flickering  activity in a sample of CVs and symbiotic systems. The promising models which do not contradict the empirical results, are unstable mass accretion onto the central accretor or turbulence in the inner accretion disc. Bruch (1996) studied the variations of photometric data as a function of orbital phase in Z\\,Cam. The source of the flickering was localized in the vicinity of the central accretor, but the outer disc edge and its interaction with the gas stream from the L$_{\\rm 1}$ point cannot be neglected (hot-spot). Bruch (2000) studied 4 eclipsing systems: HT\\,Cas, V2051\\,Oph, UX\\,UMa and IP\\,Peg.  In these systems, similar to Z\\,Cam, the flickering source  is located in the central part of the disc or the hot-spot. For V2051\\,Oph the region of the flickering was confirmed by Baptista and Bortoletto (2004) using eclipse mapping. The hot-spot is the source of low frequency flickering and the central disc of the high frequency flickering. Finally, T\\,CrB case was studied by Zamanov and Bruch (1998). The $U$ photometric data showed fast variability located near the central white dwarf. Following these observational results it is clear that the dominant source of flickering variability in CVs and T\\,CrB is the central part of the disc or the vicinity of the white dwarf. A promising scenario to explain this phenomenon is the existence of magneto-hydrodynamic turbulence transporting the angular momentum outwards (Balbus and Hawley 1998). However why only the central part of the disc is the source of the observed turbulence?  In this paper we focus on the analysis of turbulent transport of angular momentum in the disc. We have decided to study T\\,CrB due to its rich phenomenology of rapid variability (i.e. flickering activity) and because of it's disc properties are suitable for our model.  In Sec.~\\ref{observations} we present our observations. Our treatment of turbulence in the disc as the source of flickering is explained in Section~\\ref{source} and the simulations are described in Section~\\ref{simulations}. Finally the results are analyzed in Sec.~\\ref{analysis} and discussed in Sec.~\\ref{discussion}. ",
        "conclusions": "\\label{discussion}  The 90\\% confidence level contours of the PDS fits demonstrates that the inner disc radius and the viscosity parameter are correlated in the interval $R_{\\rm in} \\simeq 1.1 - 4.3 \\times 10^9$\\,cm for all aceptable $\\alpha$ values for a disc in hot branch, i.e. 0.1 -- 1.0. Our best solution $R_{\\rm in} \\simeq 3.9 \\times 10^9$\\,cm and $\\alpha \\simeq 0.9$ is in agreement with the inner disc radius obtained by Skopal (2005) through modeling of UV/optical/IR spectra of T\\,CrB; i.e. $0.05 - 0.10$\\,R$_{\\rm \\odot} \\simeq 3.5 - 7.0 \\times 10^9$~cm. This is larger than the estimated white dwarf radius and hence the disc seems to be truncated. Furthermore, the best model viscosity parameter is in agreement with expectations for a disc in the permanent hot branch, i.e. $0.1 < \\alpha < 1.0$.  Our best values are derived from the minimal value of $\\chi^2$. The 90\\% contours levels suggest a large uncertainty in parameters which are correlated. The determination of one of the two parameters requires an independent measurement of the other. Taking the inner disc radius from Skopal (2005) to lie in the interval of $3.5 < R_{\\rm in} < 4.3 \\times 10^9$\\,cm then gives a viscosity interval of $0.7 < \\alpha < 1.0$. King et al. (2007) report, for fully ionized geometrically thin accretion discs, that the observational evidence suggests $\\alpha \\sim 0.1 - 0.4$. In particular, the simulations for SS\\,Cyg (Schreiber et al. 2003) yield a viscosity value of about $0.1 - 0.2$, whereas Schreiber and G\\\"ansicke (2002) obtain $\\alpha$ = 0.5 for SS\\,Cyg in the ionized hot state. Our derived viscosity value of $0.7 - 1.0$ is considerably higher than previous values. However, Schreiber et al. (2003) used a maximum mass transfer rate of $\\sim 1 \\times 10^{17}$~g~s$^{-1}$ in his simulations. This is one order of magnitude lower than mass accretion rate in T\\,CrB. The high mass transfer rate through the disc together with the irradiation by the central source may explain the higher temperature and viscosity than in standard dwarf nova system. Furthermore the lower viscosity values were determined from the long-term light curve behaviour where the considered timescales are typical for the largest disc radii, whereas our constraints apply mainly at the inner disc radius. This gives an idea that the viscosity parameter $\\alpha$ can also be a function of the distance $r$. In the standard disc instability model (see e.g. Lasota 2001 for review) the main difference between the viscous parameter in the cold and hot branch is due to the ionisation of hydrogen. Once the material is in the hot state, the ionization state of hydrogen does not change with increasing temperature and the viscosity parameter is supposed to be constant in the whole disc. Intensive studies of the disc behaviour with constant  $\\alpha$ were performed in the case of dwarf novae and low mass X-ray binaries.  But the typical disc radii of these binaries are of the order of $10^{10} - 10^{11}$\\,cm. The case of symbiotic systems, such as T\\,CrB is different  because of the  large binary separation and the possible outer disc radii up to $10^{12}$\\,cm. But we showed that our simulations are not sensible to such large disc radii in the acceptable $\\alpha$ regime. The outer disc radius used for T\\,CrB is $10^{11}$\\,cm, and hence is in agreement with the radius of low mass X-ray binaries modelled by Lasota (2001) with constant $\\alpha$ values. Therefore we conclude that radial variations of the viscous parameter is not required. An alternative explanation for the different $\\alpha$ results could be a different model other than the Shakura and Shunyaev (1973) model for the large outer discs in symbiotic systems. An application of our simulations to dwarf novae with disc radii of $10^{10}$\\,cm can test whether the $\\alpha$ differences remain or disappear. This case is problematic because dwarf novae discs are not in a steady-state. Another possibility is to study dwarf novae in outburst when the disc is fully ionised and is in a quasi steady-state (see e.g. Lasota 2001) and in a permanent hot state except for in the lightcurve minima e.g. in VY\\,Scl systems (see e.g. Warner 1995).  Another interesting question is the location of the flickering events in the inner disc. This can be understood from the behavior of the scaling parameter $k(r)$ (Fig.~\\ref{disc_prof}). For very small $r$ values, in the vicinity of the central star, $k(r)$ rises very steeply because of the small disc scale height (Fig.~\\ref{disc_prof}) and hence small turbulent eddies. Furthermore, $k(r)$ decreases with increasing $r$ and the disc scale height, and thus less events are needed to transport the angular momentum. This strongly suggests that the inner parts of the disc are the source of the fast flickering, which is in agreement with observations (see e.g. Bruch 1992, Zamanov \\& Bruch 1998, Bruch 1996, Bruch 2000, Baptista \\& Bortoletto 2004, Zamanov 2004). A similar behavior is deduced from the sensitivity of our modeling to the input parameters $R_{\\rm in}$ and $R_{\\rm out}$, which is very high for $R_{\\rm in} \\sim 10^9 - 10^{10}$\\,cm but not for $R_{\\rm out} \\sim 10^{11} - 10^{12}$\\,cm. This suggests that the majority of turbulent events are located in the central part of the disc up to radial distance $\\sim 10^{10}$\\,cm.  \\begin{figure} \\includegraphics[width=80mm]{disc_profile.eps} \\caption{A simulation of the scaling parameter $k$, disc scale height $H$ and radial velocity $v_{\\rm r}$ for $M_1$ = 1.37~M$_{\\rm \\odot}$, $\\dot{M}_{\\rm acc} = 1.9 \\times 10^{18}$~g~s$^{-1}$ and $\\alpha = 0.5$.} \\label{disc_prof} \\end{figure}  The inner disc radius and the viscosity is correlated only for low values of $r$ comparing to the disc radius. This can be explained because the viscosity must increase with higher $R_{\\rm in}$ in order to get the same $\\chi^2$ residual. This means that by truncating the central part of the disc, we are removing many fast events which should be compensated by the increase of the radial velocity of slower events at larger radii. The radial velocity $v_{\\rm r} \\sim \\alpha^{4/5}$ strongly decreases at very small $r$ (see Fig.~\\ref{disc_prof}). An alternative explanation relies on the total number of events in the central disc. The scaling parameter $k$ gives an estimate of the number of events and this is very large at small $r$. Hence by truncating the central part of the disc we are removing many events. Whereas at larger $r$ the decrease of $k$ and $v_{\\rm r}$ is flatter and the correlation between $R_{\\rm in}$ and $\\alpha$ is not as strong as for small values of $r$.  Fig.~\\ref{synt_runs} shows three examples of simulated lightcurves computed for different viscosities. It is obvious that high frequency fluctuations increase with $\\alpha$, whereas low-frequency fluctuations decrease with $\\alpha$. This disappearance of large events is visible in our PDS (Fig.~\\ref{pds_examples}). The low viscosity PDS has a strong peak near frequency $10^{-3}$~Hz which is caused by the quasi-periodic oscillations present in the simulated data. Such variations has a timescale of $\\sim 0.01$ day clearly visible on Fig.~\\ref{synt_runs} for $\\alpha = 0.01$ case. The timescale of the data is too short to distinguish whether the fluctuations are stochastic or quasi-periodic. For higher viscosities, large events disappear and so does the evidence for quasi-periodic signals in the simulated data. Hence the peak at $10^{-3}$~Hz dissapear.  In addition, the observed PDS in Fig.~\\ref{pds_examples} are scattered which produce a high $\\chi^2$. This behavior tells us something about the observational demands. Our observed PDS is computed from 9 runs, whereas the simulated PDS is computed from about 100 synthetic runs. Smoother observed PDS (made from a higher number of observational data sets) are required for the analysis to get lower $\\chi^2$ and then make the method more sensitive to the input parameters. The mean duration of our data is $\\sim 2.6$ hours. Longer observations could remove the quasi-periodic signals at low viscosities and reduce the scatter in the observed and simulated PDS at low frequencies ($\\sim 10^{-4}$~Hz). Using of a single filter can improve the data resolution to $\\sim 10$ seconds and hence the quality of the PDS at higher frequencies ($\\sim 10^{-2}$~Hz). Also the observational gaps needed to observe a comparison star (clearly visible on Figs.~\\ref{runs} and \\ref{synt_runs}) affect the PDS. Therefore CCD observations with fast readout speeds are preferred.  \\begin{figure} \\includegraphics[width=80mm]{synt_runs.eps} \\caption{ Three examples of the synthetic data compared to an observed lightcurve. The simulated runs are computed for $M_1$ = 1.37~M$_{\\rm \\odot}$, $\\dot{M}_{\\rm acc} = 1.9 \\times 10^{18}$~g~s$^{-1}$ and $R_{\\rm in} = 3.5 \\times 10^9$~cm and three different viscosities $\\alpha = 0.01$, 0.50 and 10.0.} \\label{synt_runs} \\end{figure}  Our best model yields a parameters in reasonable agreement with other works. We therefore conclude that our simple idea of angular momentum transport through turbulent eddies is sufficient to explain the observed flickering statistics. No other mechanisms are then required to generate the observed flickering activity."
    },
    "0911/0911.1801_arXiv.txt": {
        "abstract": "We used an N-body smoothed particle hydrodynamics algorithm, with a detailed treatment of star formation, supernovae feedback, and chemical enrichment, to perform eight simulations of mergers between gas-rich disc galaxies. We vary the mass ratio of the progenitors, their rotation axes, and their orbital parameters and  analyze the kinematic, structural, and chemical properties of the remnants. Six of these simulations result in the formation of a merger remnant with a disc morphology as a result of the large gas-fraction of the remnants. We show that stars formed during the merger (a sudden starburst occur in our simulation and last for  $0.2-0.3\\,{\\rm Gyr}$) and those formed after the merger have different kinematical and chemical properties. The first ones are  located in thick disc or the halo. They are partially supported by velocity dispersion and have high \\al ratios even at metallicities as high as $\\rm[Fe/H]=-0.5$. The former ones -- the young component -- are located in a thin disc rotationally supported and have lower \\al ratios. The difference in the rotational support of both components results in the rotation of the thick disc lagging that of the thin disc by as much as a factor of two, as recently observed. We find  that, while  the kinematic and structural properties of the merger remnant depends strongly upon the orbital parameters of the mergers, there is  a remarkable uniformity in the chemical properties of the mergers. This suggests that general conclusions about the chemical signature of gas-rich mergers can be drawn. ",
        "introduction": "\\subsection{Disc-Disc Mergers}  In a $\\Lambda$ Cold Dark Matter Universe ($\\Lambda$CDM), the formation of structures generally involves the merging of smaller structures. Mass is built up predominantly by the merger of objects of masses between $3-30\\%$ of the galaxy's total mass at redshift $z=0$ \\citep{stewartetal08}. Since $z=2$, approximately 60\\% of galaxies have experienced a major merger, which we will take in this paper as those involving a ratio of masses 3:1 or higher. Also, an important fraction of the baryonic mass (30 - 50\\%) of $z=0$ galaxies is provided by major mergers \\citep{stewartetal09}. Recent observations show that nearly 20\\% of the massive disc galaxies have undergone a major merger since $z\\sim1.5$ \\citep{lopezetal09} This process is virtually scale-free, meaning that low-mass dark matter halos, in which disc galaxies presumably reside, have undergone similar merging histories to high-mass dark matter halos in which, predominantly, ellipticals reside. This at first is puzzling, as mergers have long been associated with forming spheroidal structures, i.e. early-type galaxies \\citep{toomre77}. Following this line of reasoning, every galaxy in the nearby universe should be an elliptical or an irregular, which is clearly not the case. Furthermore, observed elliptical galaxies can be described by simple relations relating colors, luminosities, masses, etc. These simple relations are difficult to reproduce with a scenario involving collisions \\citep{bs98,bs99}.   The resolution of this, of course, lies with the baryonic physics. Studies using two stable discs as initial conditions (\\citealt{sh05,robertsonetal06}; \\citealt[][hereafter B07]{brooketal07}; \\citealt{rb08}) as well as fully cosmological simulations \\citep{gv08} have shown that binary gas-rich mergers can result in disc morphologies. Hence, major mergers between gas-rich disc galaxies can possibly play an important role in the formation of a large number of the spiral galaxies we observe today. These facts bring an addition to the original scenario for the formation of disc galaxies, which involved the assembly of a halo of dark matter and ionized gas by hierarchical clustering and mergers of sub-halos, followed by the collapse and virialization of the dark matter halo, and finally the dissipative collapse of the baryonic component, resulting in a rotationally-supported disc, that will later fragment to form stars \\citep{wr78,fe80,blumenthaletal84}.  One must also consider a formation scenario involving major collisions between spiral galaxies. The first numerical simulations of major collisions between disc galaxies were performed under the assumption that dissipation by gasdynamical effects was negligible. The results showed that the remaining galaxy had several properties attributed to elliptical galaxies, including luminosity profiles \\citep{nb03,barnes02,bh92,hernquist92,hernquist93,bs98,bs99,nbh99}. Even with the quite recent inclusion of feedback processes, the results remained mostly unchanged \\citep{bh96,mh96,springel00,naabetal06}. Minor collisions have also been studied in the past and results showed that this kind of collision tends to form galaxies whose disc is destroyed or very perturbed \\citep{hm95,bekki98}. All the aforementioned studies assumed a ratio gas/stars more representative of present low redshift mergers. Of course, stars have not always been present inside galaxies. Going back in time, we can expect to find fewer and fewer stars, and therefore a more important gaseous component \\citep{stewartetal09}.  \\subsection{Thick Discs}  The Milky Way's thick disc is made predominantly of old stars, with ages in the range of $8-12\\,{\\rm Gyr}$, significantly older than the typical star found in the thin disc \\citep{gwj95}. The thick and thin disc differ in their structural, chemical, and kinematical properties: (1) The thick disc has a scale height of about 900 pc, compared to 300 pc for the thin disc \\citep{juric08}. (2) The stars in the thick disc have a low metallicity, $\\rm[Fe/H]\\sim-0.6$ on average, and a ratio \\al larger than the stars in the thin disc. (3) Stars in the thick disc tend to rotate slower than stars in the thin disc, by about $40\\,{\\rm km\\,s^{-1}}$ \\citep{gwk89} and have an average vertical velocity dispersion of $45\\,{\\rm km\\,s^{-1}}$ compared to $\\sim10-20\\,{\\rm km\\,s^{-1}}$ for the stars in the thin disc \\citep{delhaye65,cb00}.  The formation of the Milky Way's thick disc remains an open question, with popular theories including the heating of a thin disc by minor mergers \\citep{quinnetal93,kazantzidisetal08,kzb08}, or the direct accretion of stars, which are tidally torqued into the plane of the existing thin disc \\citep{abadietal03}, or dynamical heating by destruction of very massive star clusters \\citep{kroupa02,ee06}. We favor a scenario in which the Milky Way's thick disc is formed during the high-redshift epoch in which mergers are most common in $\\Lambda$CDM, and accretion rates are high, with the disc which formed during this turbulent period being born thick; the thin disc subsequently grows during the relatively quiescent period which follows, which in the case of the Milky Way is approximately the last 10 Gyrs \\citep{brooketal04b,brooketal05,brooketal06}.  Recent observations suggest that all spiral galaxies are surrounded by a red, flattened envelope of stars, and these structures have been interpreted as showing that thick discs are ubiquitous in spiral galaxies \\citep{db02}. Hence, understanding the formation of the thick disc will provide major insights into the formation processes of spiral galaxies. If we make the hypothesis that thick disc are formed during gas-rich major mergers, it would mean that a significant number of spiral galaxies have experienced major mergers. Indeed, simulations have shown that gas-rich mergers tend to form two disc components (\\citealt{robertsonetal06}; B07). These simulations included the effect of star formation and supernovae feedback, which allowed to study the kinematical and structural properties of galaxy merger remnants in detail. However, they did not include a detailed treatment of chemical enrichment and, therefore,  could not predict the abundances of metals in stars and the interstellar medium \\citep[but see also][who studied the formation of ellipticals galaxies due to a more violent major merger]{bs98,bs99} . It is well-established that the thick disc of the Milky Way possesses a unique chemical signature, which is believed to be directly related to the formation process and it is important to analyze if the merger scenario for the formation of the thick disc component is compatible with  these chemical signatures. This is the basis of our study here: we believe that the remnants of major mergers may be associated with a significant fraction of the ``red envelopes'' observed in \\citet{db02}, i.e. with extragalactic thick discs. Since we expect the conditions for merging galaxies during hierarchical clustering at a high redshift to vary, we need to consider a wide range of parameters, such as mass ratio, orbit, and gas fraction of the merger progenitors. In this paper, we present a series of 8 simulations of mergers of gas-rich disc galaxies, with different mass ratios and orbital parameters. We analyze the kinematic and structural properties of the merger remnants. The results for one of the simulations have already been presented in a previous paper (B07). We now complete this work by presenting the entire suite of simulations.  For kinematic properties, we determine the relative importance of rotation and velocity dispersion in providing support. We also investigate if a counter-rotating component arises naturally in a gas-rich merger, as such components are observed in nature \\citep{yd06}. For structural properties, we determine if the luminosity profiles of these thick discs formed in major mergers follow an exponential law, and if so, what their respective scale lengths are. We determine if the thick and thin discs in merger remnants are coplanar, or if there is a significant angle between them, something that was not included in our previous work.   However, the main goal of the present paper is to study the viability of the merger mechanism for the formation of the thick disc from the chemical evolution's point of view, but without neglecting the kinematical and dynamical properties. With this in mind, we choose initial conditions likely to produce a discy remnant.  In this paper, we use the numerical algorithm GCD+, which includes a detailed treatment of chemical enrichment, to investigate whether key abundance ratios such as \\al and \\fe in simulated remnants of gas-rich disc galaxy mergers have values and distributions similar to the ones found in the observed disc galaxies, including the Milky Way.   The reminder of this paper is organized as follows: in \\S2, we briefly describe the numerical simulations employed, including the N-body/SPH code adopted (GCD+), the initial conditions software (GalactICS), and the basic parameters of our 8 realizations. The kinematical and structural properties of the merger remnants are presented in \\S3 and discussed in \\S4. Conclusions are presented in \\S5. ",
        "conclusions": "Using GCD+, we have performed 8 simulations of major mergers between gas-rich spiral galaxies. We have analyzed the kinematic, structural, and chemical properties of stars formed before and during the collision (the old population) and stars formed after the collision (the young population). We used the star formation rate to define these two populations. A fraction of the old stars end up in the halo of the merger remnant, while the remaining stars form a thick disc which is partly supported by velocity dispersion, partly supported by rotation, and sometimes includes a significant counter-rotating component. The young stars form a thin disc that is supported by rotation, and, in many cases, a ring that might or might not be coplanar with the thin disc. The disc themselves tend to be coplanar, the angle between them being $\\sim6^\\circ$ or less. With rare exceptions, both discs are well-fitted by an exponential profile, and the scale-length of the thick disc exceeds the one of the thin disc a few percents up to a factor of 1.60.  The starburst occurring during the collision rapidly enriches the gas in various metals. Explosions of Type~II SNe follow rapidly the start of the collision, owing to the short lifetime of their progenitors, and enrich the intergalactic medium in $\\alpha$-elements. Enrichment by Type~Ia SNe is spread over a large period of time, enabling the enrichment in iron of both stellar populations, which results in an old population having a \\al ratio higher than the young population, even at relatively large metallicities ($\\hbox{\\fe}=-0.5$). This could explain the high \\al ratio observed in stars in the thick disc and the halo of the Milky Way. This result do not depend strongly upon the initial conditions, since it was found in all simulations.  \\begin{figure*} \\begin{center} \\includegraphics[width=4.8in]{figure19.eps} \\caption{Star formation epoch versus final radius for old stars (solid line) and young stars (dot-dashed line), for the 8 simulations. The horizontal dashed lines indicate the beginning and the end of the starburst.} \\label{ageradM12} \\end{center} \\end{figure*}  \\begin{figure*} \\begin{center} \\includegraphics[width=4.8in]{figure20.eps} \\caption{Toomre diagram for stars located near the disc plane ($Z<1\\,{\\rm kpc}$), for simulation M12, divided according to metallicity. The metallicity range is indicated on top of each panel.} \\label{toomreM12} \\end{center} \\end{figure*}  Our main conclusion is that the morphological, kinematical and chemical properties of the thick and thin disc can be reproduced in an scenario where thick disc formed in a gas rich merger between two disc galaxies. Furthermore, while the structural and kinematical properties of the merger remnants are strongly dependent on the initial conditions --features such as the ratio of the disc scale-lengths, the extent of the discs, the presence of rings or of a counter-rotating component, vary from simulation to simulation-- the chemical abundances show a remarkable consistency among the various simulations. The key result of this study is that the ratio \\al remains constant with increasing \\fe for old stars, up to $\\rm[Fe/H]=-0.5$ while it decreases with increasing metallicity for young stars. This is true in each of the 8 simulations we considered. The observed  chemical signatures in our merger remnants require that mergers happen before the onset of the majority of Type~Ia SNe, regardless of the orbits or mass ratios of the progenitors. Furthermore, the mergers need to be gas rich to lead to a remnant with a rotationally supported disc. These characteristics are consistent with the mergers we expect to find at high redshift and, therefore, our results support this scenario for the formation of the different disc components in the Milky Way."
    },
    "0911/0911.3035_arXiv.txt": {
        "abstract": "{ The aim of this project is to determine the consistency of an assumed cosmological model by means of a detailed analysis of the brightness profiles of distant galaxies. Starting from the theory developed by Ellis and Perry (1979) connecting the angular diameter distance obtained from a relativistic cosmological model and the detailed photometry of galaxies, we assume the presently most accepted cosmological model with non-zero cosmological constant and attempt to predict the brightness profiles of galaxies of a given redshift. Then this theoretical profile can be compared to observational data already available for distant, that is, high redshift, galaxies. By comparing these two curves we may reach conclusions about the observational feasibility of the underlying cosmological model. ",
        "introduction": "The most basic goal of cosmology is to determine the spacetime geometry and matter distribution of the Universe by means of astronomical observations. Accomplishing this goal is not an easy or simple task, and due to that, since the early days of modern cosmology several methods have been advanced such that theory and observations are used to check one another. Detailed analysis of the cosmic microwave background radiation, galaxy number counts and supernova cosmology are just a few of the methods employed nowadays in cosmology, deriving results that complement one another. In this work we aim at discussing one of these methods, namely the connection between galaxy brightness profiles and cosmological models.  Thirty years ago, Ellis and Perry (1979) advanced a very detailed discussion where such a connection is explored. Their aim was to determine the spacetime geometry of the universe by connecting the angular diameter distance, also known as area distance, obtained from a relativistic cosmological model, and the detailed photometry of galaxies. They then discussed how the galaxy brightness profiles of high redshift galaxies could be used to falsify cosmological models as the angular diameter distance could be determined directly from observations.  Nevertheless, to carry out this program to its full extent, one would need detailed information about galaxy evolution. Without a consistent theory on how galaxies evolve, it is presently impossible to analyze cosmological observations without assuming a cosmological model. In addition, brightness profiles are subject to large observational errors, making it difficult to achieve Ellis and Perry's aim of possibly using the angular diameter distance determination to distinguish cosmological models.  This work is based on Ellis and Perry (1979) theory, although our aim is more limited in the sense that we do not seek to determine the underlying cosmological model by directly measuring the angular diameter distance, but to assume the presently most favored cosmology, deriving cosmological distances from it and seeking to discuss the consistency between its predictions and detailed observations of surface brightness of distant galaxies. Our goal is to obtain a theoretical brightness profile by means of the assumed cosmological model and compare it with its observational counterpart at various redshift ranges, and for different galaxy morphologies.  The outline of the paper is as follow. In \\S2 we introduce the cosmological distances and their connections to astrophysical observables. In \\S3 we describe the parameters that determine the surface brightness structure and in \\S4 we discuss the criteria for selecting galaxies, in view of the importance of evolutionary effects in galactic surface brightness. ",
        "conclusions": ""
    },
    "0911/0911.0053_arXiv.txt": {
        "abstract": "{ We present a comprehensive {\\ro analysis of weak gravitational lensing by large-scale structure} in the \\textit{Hubble Space Telescope} Cosmic Evolution Survey (COSMOS), in which we combine space-based galaxy shape measurements with ground-based photometric redshifts to study the redshift dependence of the lensing signal and constrain cosmological parameters. After applying our weak lensing-optimized data reduction, principal component interpolation for the spatially and temporally  varying ACS point-spread function, and improved modelling of charge-transfer inefficiency, we measure a lensing signal which is consistent with pure gravitational modes and no {\\ro significant shape} systematics. We carefully estimate the statistical uncertainty from simulated COSMOS-like fields obtained from ray-tracing through the Millennium Simulation, including the  full non-Gaussian sampling variance. We test our lensing pipeline on simulated space-based data, recalibrate  non-linear power spectrum corrections using the ray-tracing analysis, employ photometric redshift information to reduce potential contamination by intrinsic galaxy alignments, and  marginalize over systematic uncertainties. We find that the weak lensing signal scales with redshift as expected from General Relativity for a concordance $\\Lambda$CDM cosmology, including the full cross-correlations between different redshift bins. Assuming a flat $\\Lambda$CDM cosmology, we measure \\mbox{$\\sigma_8\\left(\\Omega_\\mathrm{m}/0.3\\right)^{0.51}=0.75\\pm0.08$} from lensing, in perfect agreement with WMAP-5, yielding joint constraints \\mbox{$\\Omega_\\mathrm{m}= 0.266^{+0.025}_{-0.023}  $}, \\mbox{$\\sigma_8=0.802^{+0.028}_{-0.029}$} (all 68.3\\% conf.). Dropping the assumption of flatness and using priors from the HST Key Project and Big-Bang nucleosynthesis only, we find a negative deceleration parameter $q_0$ at 94.3\\% confidence from the tomographic lensing analysis, providing independent evidence for the accelerated expansion of the Universe. For a flat $w$CDM cosmology and prior \\mbox{$w\\in[-2,0]$}, we obtain \\mbox{$w<-0.41$} (90\\% conf.). Our dark energy constraints are still relatively weak solely due to the limited area of COSMOS. {\\ro However, they provide an important demonstration for the usefulness of tomographic weak lensing measurements from space.} } ",
        "introduction": "During the last decade strong evidence for an accelerated expansion of the Universe has been found with several independent cosmological probes including type Ia supernovae  \\citep[][]{rfc98,pag99,rsc07,kra08,hwb09}, cosmic microwave background \\citep{bab00,svp03,kdn09}, galaxy clusters \\citep{ars08,mae08,mar09,vkb09}, baryon acoustic oscillations \\citep{ezh05,pce07,pre09}, integrated Sachs-Wolfe effect \\citep{gsc08,gns08,hhp08}, and strong gravitational lensing \\citep{sma09}. Within the standard cosmological framework this can be described with the ubiquitous presence of a new constituent named dark energy, which counteracts the attractive force of gravity on the largest scales and contributes  \\mbox{$\\sim 70\\%$} to the total energy budget today. There are various attempts to explain dark energy, ranging from Einstein's cosmological constant, via a dynamic fluid named quintessence, to a possible breakdown of General Relativity \\citep[e.g.][]{hul07,aab09}, all of which lead to profound implications for fundamental physics. In order to make substantial progress and to be able to distinguish between the different scenarios, several large dedicated surveys are currently being designed.  One  technique holding particularly high promise to constrain dark energy \\citep{abc06,pse06,aab09} is weak gravitational lensing, which utilizes the subtle image distortions imposed onto the observed shapes of distant galaxies while their light bundles pass through the gravitational potential of foreground structures \\citep[e.g.][]{bas01}. The strength of the lensing effect depends on the total foreground mass distribution, independent of the relative contributions of luminous and dark matter. Hence, it provides a unique tool to study the statistical properties of large-scale structure directly \\citep[for reviews see ][]{sch06,hoj08,mvw08}.  Since its first detections by \\citet{bre00,kwl00,wme00} and \\citet{wtk00}, substantial progress has been made with the measurement of this cosmological weak lensing effect, which is also called cosmic shear. Larger surveys have significantly reduced statistical uncertainties \\citep[e.g.][]{hyg02,btb03,jbf03,mrb05,wmh05,hmw06,smw06,hss07,fsh08}, while tests on simulated data have led to a better understanding of PSF systematics \\citep[][and references therein]{hwb06,mhb07,bbb09}. Finally, being a geometric effect, gravitational lensing depends on the source redshift distribution, where most earlier measurements had to rely on external redshift calibrations from the small \\textit{Hubble} Deep Fields. Here, the impact of sampling variance was demonstrated by \\citet{bhs07}, who recalibrated earlier measurements using photometric redshifts from the much larger CFHTLS-Deep, significantly improving derived cosmological constraints.   Dark energy affects the distance-redshift relation and suppresses the time-dependent growth of structures. Being sensitive to both effects,  weak lensing is a powerful probe of dark energy properties, also providing important tests for theories of modified gravity \\citep[e.g.][]{beb01,bew04,stu07,dmm07,jaz08,sch08}. Yet, in order to significantly constrain these redshift-dependent effects, the shear signal must be measured as a function of source redshift, an analysis often called weak lensing tomography or 3D weak lensing \\citep[e.g.][]{hu99,hu02,hut02,jat03,hea03,huj04,bej04,sks04,taj04,hkt06,tkb07}. Redshift information is additionally required to eliminate potential contamination of the lensing signal from intrinsic galaxy alignments \\citep[e.g.][]{kis02,his04,hwh06,jos08}. In general, weak lensing studies have to rely on photometric redshifts \\citep[e.g.][]{ben00,iam06,hwb08} given that most of the studied galaxies are too faint for spectroscopic measurements.   {\\mt So far, tomographic cosmological weak lensing techniques were applied to real data by \\citet{btb05,smw06,kht07,mrl07}}. Dark energy constraints from previous weak lensing surveys were limited by the lack of the required individual photometric redshifts \\citep{jjb06,hmw06,smw06,kbg09} or small survey area \\citep{kht07}. The currently best data-set for 3D weak lensing is given by the COSMOS Survey \\citep{saa07}, which is the largest continuous area ever imaged with the \\textit{Hubble Space Telescope} (HST), comprising 1.64 deg$^2$ of deep imaging with the Advanced Camera for Surveys (ACS). Compared to ground-based measurements, the HST point-spread function (PSF) yields substantially increased number densities of sufficiently resolved galaxies and better control for systematics due to smaller PSF corrections. Although HST has  been used for earlier cosmological weak lensing analyses \\citep[e.g.][]{rrg02,rrc04,meh05,hbb05,ses07}, these studies lacked the area and deep photometric redshifts which are available for COSMOS \\citep{ics09}. This combination of superb space-based imaging and ground-based photometric redshifts makes COSMOS the perfect test case for 3D weak lensing studies. \\citet{mrl07} conducted an earlier 3D weak lensing analysis of COSMOS, in which they {\\mt correlated the shear signal between three redshift bins } and constrained the matter density $\\Omega_\\mathrm{m}$ and power spectrum normalization $\\sigma_8$. In this paper we present a new analysis of the data, with several differences compared to the earlier study: we employ a new, exposure-based model for the spatially and temporally varying ACS PSF, which has been derived from dense stellar fields using a principal component analysis (PCA). Our new parametric correction for the impact of charge transfer inefficiency (CTI) on stellar images eliminates earlier PSF modelling uncertainties caused by confusion of CTI- and PSF-induced stellar ellipticity. Using the latest photometric redshift catalogue of the field \\citep{ics09}, we split our galaxy sample into five individual redshift bins and additionally estimate the redshift distribution for very faint galaxies forming a sixth bin without individual photometric redshifts, doubling the number of galaxies used in our cosmological analysis. We study the redshift scaling of the shear signal between these six  bins in detail, employ an accurate covariance matrix obtained from ray-tracing through the Millennium Simulation, which we also use to recalibrate non-linear power spectrum corrections, and marginalize over parameter uncertainties. In addition to $\\Omega_\\mathrm{m}$ and  $\\sigma_8$, we also constrain the dark energy equation of state parameter $w$ for a flat $w$CDM cosmology, and the vacuum energy density $\\Omega_\\Lambda$ for a general (non-flat) $\\Lambda$CDM cosmology, yielding constraints for the deceleration parameter $q_0$.   This paper is organized as follows. We summarize the most important information on the data and photometric redshift catalogue in Sect.\\thinspace\\ref{se:data}, while further details on the ACS data reduction are given in App.\\thinspace\\ref{app:data}. Section \\ref{se:wl} summarizes the weak lensing measurements including our new correction schemes for PSF and CTI, for which we provide details in App.\\thinspace\\ref{app:psf_cti}. We conduct various tests for shear-related systematics in Sect.\\thinspace\\ref{se:wl:tests}. We then present the weak lensing tomography analysis in Sect.\\thinspace\\ref{se:tomo}, and cosmological parameter estimation in Sect.\\thinspace\\ref{se:cosmo}. We discuss our findings and conclude in Sect.\\thinspace\\ref{se:dis}.   Throughout this paper all magnitudes are given in the AB system, where  $i_{814}$ denotes the \\texttt{SExtractor} \\citep{bea96} \\mbox{$\\mathtt{MAG\\_AUTO}$} magnitude measured from the ACS data (Sect.\\thinspace\\ref{se:data:acs}), while $i^+$ is the {\\ro \\texttt{MAG\\_AUTO}} magnitude determined by \\citet{ics09} from the  Subaru data (Sect.\\thinspace\\ref{se:tomo:zcatsbright}). In several tests we employ a reference {\\ro WMAP-5-like \\citep{dkn09}} flat $\\Lambda$CDM cosmology characterized by \\mbox{$\\Omega_\\mathrm{m}=0.25$}, \\mbox{$\\sigma_8=0.8$}, \\mbox{$h=0.72$},  \\mbox{$\\Omega_\\mathrm{b}=0.044$}, \\mbox{$n_\\mathrm{s}=0.96$}, where we use the transfer function by  \\citet{eih98} and non-linear power spectrum corrections according to \\citet{spj03}. ",
        "conclusions": "\\label{se:dis}  We have measured weak lensing galaxy shear estimates from the HST/COSMOS data by applying a new model  for the spatially and temporally varying ACS PSF, which is  based on a principal component analysis of PSF variations in dense stellar fields. We find that most of the PSF changes can be described with a single parameter related to the HST focus position. Yet, we also correct for additional PSF variations, which are coherent for neighbouring COSMOS tiles taken closely in time. We employ updated parametric corrections for charge-transfer inefficiency, for both galaxies and stars, removing earlier modelling uncertainties due to confused PSF- and CTI-induced stellar ellipticity. Finally, we employ a simple correction for signal-to-noise dependent shear calibration bias, which we derive from the STEP2 simulations of ground-based weak lensing data. Tests on simulated space-based data confirm a relative shear calibration uncertainty \\mbox{$|m|\\le 2\\%$} over the entire used magnitude range if this correction is applied. We  decompose the measured shear signal into curl-free E-modes and curl-component B-modes. As expected from pure lensing, the B-mode signal is consistent with zero for all second-order shear statistics, providing an important confirmation for the success of our  correction schemes for instrumental systematics.  We combine our shear catalogue with excellent ground-based photometric redshifts from \\citet{ics09} and carefully estimate the redshift distribution for faint ACS galaxies without individual photo-$z$s. This allows us to study weak lensing cross-correlations in detail between six redshift bins, demonstrating that the signal indeed scales as expected from General Relativity for a concordance $\\Lambda$CDM cosmology.   We employ a robust covariance matrix from 288 simulated COSMOS-like fields obtained from ray-tracing through the Millennium Simulation \\citep{hhw09}. Using our  3D weak lensing analysis of COSMOS, we derive constraints \\mbox{$\\sigma_8(\\Omega_m/0.3)^{0.51}=0.79\\pm 0.09$} for a flat $\\Lambda$CDM cosmology, using non-linear power spectrum corrections from \\citet{spj03}. A recalibration of these predictions based on the ray-tracing analysis changes our constraints to \\mbox{$\\sigma_8\\left(\\Omega_\\mathrm{m}/0.3\\right)^{0.51}=0.75\\pm0.08$} (all 68.3\\% conf.). Our results are perfectly   consistent with WMAP-5, yielding joint constraints \\mbox{$\\Omega_\\mathrm{m}= 0.266^{+0.025+0.057}_{-0.023-0.042}  $}, \\mbox{$\\sigma_8=0.802^{+0.028+0.055}_{-0.029-0.060}$} (68.3\\% and 95.4\\% confidence). { \\mt They also agree with weak lensing results from the CFHTLS-Wide \\citep{fsh08} and recent galaxy cluster constraints from \\citet[][]{mar09} within \\mbox{$ 1\\sigma$}. } Our errors include the full statistical uncertainty including the non-Gaussian sampling variance, Gaussian photo-$z$ scatter, and marginalization over remaining parameter uncertainties, including the redshift calibration for the faint \\mbox{$i^+>25$} galaxies.  Our results are consistent with the 3D lensing constraints \\mbox{$\\sigma_8(\\Omega_m/0.3)^{0.44}=0.866\\pm 0.033\\, (\\mathrm{stat.})^{+0.052}_{-0.035}\\,(\\mathrm{sys.})$} from \\citet{mrl07} assuming non-linear power spectrum corrections according to \\citet{spj03}, at the \\mbox{$\\sim 1 \\sigma$} level. The analyses differ systematically in the treatment of PSF- and CTI-effects, where the success of our methods is confirmed by the vanishing B-mode. Furthermore,  \\citet{mrl07} employ earlier photo-$z$s based on fewer bands \\citep{mcs07}. Note that the analysis of \\citet{mrl07} yields tighter statistical errors, which may be a result of their covariance estimate from the variation between the four COSMOS quadrants. This potentially introduces a bias in the covariance inversion due to too few independent realisations \\citep{hss07b}. While the absolute calibration accuracy  of the shear measurement method was estimated to be the dominant source of uncertainty in their error budget, we were able to reduce it well below the statistical error level. As a further difference, our analysis employs photometric redshift information to reduce potential contamination by intrinsic galaxy alignments, where we exclude the shear-shear auto-correlations for the relatively narrow redshift bins 1 to 5 to minimize the impact of physically associated galaxies. In addition, we exclude  luminous red galaxies, which were found to carry the strongest intrinsic alignment with the density field of their large-scale structure environment causing the shear \\citep[][]{hmi07}. Finally, we  do not include  angular scales \\mbox{$\\theta<0\\farcm5$} due to {\\ro increased } modelling uncertainties for the non-linear power spectrum.  {\\mt Similarly to \\citet{mrl07}, we obtain a lower estimate \\mbox{$\\sigma_8(\\Omega_m/0.3)^{0.62}=0.68\\pm0.11$} for a non-tomographic (2D) analysis, assuming \\citet{spj03} power spectrum corrections. The lower signal compared to the 3D lensing analysis is expected, given that the most massive structures in COSMOS are located at \\mbox{$0.7\\lesssim z \\lesssim 0.9$} \\citep{sab07}, creating a strong shear signal for high redshift sources only, which is detected by the 3D analysis. In contrast, the bulk of the galaxies in the 2D lensing analysis are located at too low redshifts to be substantially lensed by these structures, yielding a relatively low estimate for $\\sigma_8$. Nonetheless, as sampling variance  is properly accounted for in our error analysis, the constraints are still consistent.   }   For a general (non-flat) $\\Lambda$CDM cosmology, we find a negative deceleration parameter \\mbox{$q_0<0$} at 96.0\\% confidence using our default priors, and at 94.3\\% confidence if only priors from the HST Key Project and BBN are applied. Hence, our tomographic weak lensing measurement provides independent  evidence for the accelerated expansion of the Universe. For a flat $w$CDM cosmology we constrain the (constant) dark energy equation of state parameter to \\mbox{$w<-0.41 \\,(90\\%\\,\\mathrm{conf.})$} for a prior \\mbox{$w\\in[-2,0]$}, fully consistent with $\\Lambda$CDM. Our dark energy constraints are still weak compared to recent results from independent probes \\citep[e.g.][]{kra08,hwb09,ars08,mae08,mar09,vkb09,kdn09}. This is solely due to the limited area of COSMOS, leading to a dominant contribution to the error budget from sampling variance.  {\\ro While the area covered by COSMOS is still small \\mbox{(1.64\\,$\\mathrm{deg}^2$)}, the high resolution and depth of the HST data allowed us to obtain cosmological constraints which are comparable to results from substantially larger ground-based surveys. However, note that HST was by no means designed for cosmic shear measurements. In contrast, future space-based lensing mission such as Euclid\\footnote{\\url{http://sci.esa.int/euclid}} or JDEM\\footnote{\\url{http://jdem.gsfc.nasa.gov/}} will be highly optimised for weak lensing measurements. High PSF stability, a much larger field-of-view providing thousands of stars for PSF measurements, carefully designed CCDs which minimize charge-transfer inefficiency, and improved algorithms will remove the need for some of the empirical calibrations employed in this paper. }  {\\mt In order to fully exploit the information encoded in the weak lensing shear field, second-order shear statistics, as used here, can be complemented with higher-order shear statistics to probe the non-Gaussianity of the matter distribution \\citep[e.g.][]{bar09,vlw10}. Based on  our COSMOS shear catalogue, \\citet{ssw10} present such a cosmological analysis using combined second and third-order shear statistics. }   Finally, we  stress that weak lensing can only provide precision constraints on cosmological parameters if sufficiently accurate models exist to compare the measurements to. Our analysis of the relatively small COSMOS Survey is still limited by the statistical measurement uncertainty, for which our approximate model recalibration using the Millennium  Simulation is sufficient. {\\ro Most of the cosmological sensitivity in COSMOS comes from  quasi-linear and non-linear scales. We cut our analysis only at highly non-linear scales \\mbox{$\\theta<0\\farcm5$}, corresponding to a comoving separation of \\mbox{$\\sim 360$ kpc} at \\mbox{$z=0.7$} (roughly the redshift of the most massive structures in COSMOS). At these scales non-linear power spectrum corrections have substantial  uncertainties, in particular  due to the influence of baryons \\citep[e.g.][]{rzk08}. Given that  our results are basically unchanged if even smaller scales are included (insignificant increase in $\\sigma_8$ by \\mbox{$<1\\%$}), we expect that the model uncertainty for the larger scales should still be sub-dominant compared to our statistical errors. } However, analyses of large future surveys {\\ro will urgently require improved model predictions including corrections for baryonic effects, also for dark energy cosmologies with \\mbox{$w\\ne -1$},} and optionally also for theories of modified gravity. {\\mt Once these are available, careful analyses of large current and future weak lensing surveys will deliver precision constraints on cosmological parameters and dark energy properties. }"
    },
    "0911/0911.2260_arXiv.txt": {
        "abstract": "We study the profiles of 75~086 elliptical galaxies from the Sloan Digital Sky Survey (SDSS) at both large ($70-700$ $h_{70}^{-1}$kpc) and small ($\\sim 4$ $h_{70}^{-1}$kpc) scales.  Weak lensing observations in the outskirts of the halo are combined with measurements of the stellar velocity dispersion in the interior regions of the galaxy for stacked galaxy samples. The weak lensing measurements are well characterized by a Navarro, Frenk and White (NFW) profile. The dynamical mass measurements exceed the extrapolated NFW profile even after the estimated stellar masses are subtracted, providing evidence for the modification of the dark matter profile by the baryons. This excess mass is quantitatively consistent with the predictions of the adiabatic contraction (AC) hypothesis. Our finding suggests that the effects of AC during galaxy formation are stable to subsequent bombardment from major and minor mergers. We explore several theoretical and observational systematics and conclude that they cannot account for the inferred mass excess. The most significant source of systematic error is in the IMF, which would have to increase the stellar mass estimates by a factor of two relative to the Kroupa IMF to fully explain the mass excess without AC, causing tension with results from SAURON \\citep{2006MNRAS.366.1126C}. We demonstrate a connection between the level of contraction of the dark matter halo profile and scatter in the size luminosity relation, which is a projection of the fundamental plane. Whether or not AC is the mechanism supplying the excess mass, models of galaxy formation and evolution must reconcile the observed halo masses from weak lensing with the comparatively large dynamical masses at the half light radii of the galaxies. ",
        "introduction": "\\label{sec:intro} In hierarchical models of structure formation, matter organizes itself into gravitationally bound halos that are formed through mergers of smaller halos, and by accretion along filamentary large scale structure.  The properties of these halos have been studied extensively in N-body Cold Dark Matter (CDM) simulations.  Dark matter halos have been shown to exhibit a universal density profile, of which the most common parametrization is referred to as the NFW profile \\citep{1997ApJ...490..493N}. More recent studies have determined that dark matter profiles deviate on average slightly from this two parameter description, and are better described by the Einasto profile, which introduces one extra parameter that mitigates the tendency for the measured concentration to spuriously depend on the innermost radius in the fit to the density profile (\\citealt{2004MNRAS.349.1039N,2005ApJ...624L..85M,2008MNRAS.387..536G}). The regularity in dark matter halos is one of the most robust predictions of hierarchical structure formation.  Galaxies are believed to form inside dark matter halos via the cooling and condensation of baryons to their centers.  Hydrodynamic simulations of hot gas in halos indicate that as this process occurs, baryonic physics affects the shape of the dark matter profile.  As the baryons cool to the center, they begin to dominate the mass in the halo interior.  The dark matter feels the change in the gravitational potential and is drawn into the center, steepening the central profile.   The response of the dark matter to the cooling baryons has been modeled as the  adiabatic contraction (AC) of a series of spherical shells under the assumptions of spherical symmetry, circular particle orbits, and conservation of angular momentum \\citep{1986ApJ...301...27B}.  This treatment has been modified to account for the orbital eccentricity of particles in \\cite{2004ApJ...616...16G}, and has been shown to reproduce the profiles of halos in hydrodynamical simulations at the 10--20 per cent level.  Realistically, the AC process is synchronous with other more violent processes of merging and accretion.  Older galaxies, in which the gas has been converted to stars long ago, continue to suffer dissipationless merging that has been argued to remove the influence of AC \\citep{2004ApJ...614...17G,2009ApJ...697L..38J}.  If a galaxy has had some portion of its baryonic content shock heated and removed to the warm-hot intergalactic medium where it is no longer seen, then potentially the contraction of the dark matter halo will be smaller, while the halo mass will remain the same.  The hydrodynamic simulations of \\cite{2008ApJ...672...19R} suggest that the AC of the halo persists at a level that affects the matter power spectrum at halo-sized scales.  On the other hand, \\cite{2007ApJ...658..710N} suggest that the galaxy formation history, and specifically whether recent growth is due to star formation versus due to mergers and/or accretion, influences whether adiabatic contraction occurs or whether the dark matter halo in fact becomes less concentrated.  Due to the many theoretical uncertainties, an observational handle on the mass distribution in both the inner and outer regions of the halo is crucial.  The most reliable way to probe the dark matter is through its gravitational effects. Weak gravitational lensing (for a review, see \\citealt{2001PhR...340..291B}) of background galaxies by a foreground dark matter halo, or galaxy-galaxy lensing, is one robust way to constrain the halo density profile. This method has been used to constrain halo density profiles outside of the regions where the influence of baryons is significant, for stacked galaxy samples (due to the lower signal) and for stacked and individual clusters (e.g. \\citealt{2003ApJ...588L..73D,2006MNRAS.372..758M,2008JCAP...08..006M,2009arXiv0903.1103O,2009ApJ...693.1570R}).  However, the inner regions of the halo where AC is important are typically inaccessible to weak lensing except perhaps for very deep space-based observations (e.g., \\citealt{2007ApJ...667..176G} measure the weak lensing signal down to 10 \\hkpc).  In this regime, the number of available background galaxies is small, confusion of light from the lens galaxy can make the measurement of background galaxy shapes unreliable, and the weak lensing approximation may fail entirely on small scales.  Progress can be made by combining weak lensing with another probe such as strong lensing or kinematics \\citep[e.g.,][]{2003ApJ...598..804K,2007ApJ...667..176G,2007ApJ...668..643L, 2008ApJ...681..187B,2009MNRAS.393.1235C,2009ApJ...699.1038O, 2009ApJ...694.1643U}.  In this paper, we combine weak lensing observations in the halo, outside of the central regions ($>70$ $h_{70}^{-1}$kpc), with dynamical measurements that probe the central regions ($<10$ $h_{70}^{-1}$kpc) of a stacked sample of elliptical galaxies from the Sloan Digital Sky Survey (SDSS). This approach allows us to constrain the density profile over a large dynamic range of scales. An alternate approach is to combine strong and weak lensing measurements, as in \\cite{2007ApJ...667..176G}; however, our approach allows the use of a statistical sample of elliptical galaxies, rather than being constrained only to those that are strong lenses. Finally, integral-field spectroscopy \\citep[e.g., ][]{2002MNRAS.329..513D,2006MNRAS.366.1126C} can be used to probe the galaxy density velocity distribution and therefore infer the matter profile in detail within the effective radius, which in comparison with stellar population synthesis models applied to the light distribution, can be used to infer the amount of dark matter within the effective radius.  However, the connection to the profile on much larger scales is unclear, so it is impossible on the basis of IFU observations alone to distinguish between an adiabatically contracted halo versus a halo that is very massive overall.  In section \\ref{sec:obs}, we describe the sample of galaxies used in the analysis. In section \\ref{sec:results}, we begin with the dynamical mass measurement and then combine it with weak lensing observations.  We further divide the sample along the mean size-luminosity relation and show that the scatter from this relation correlates with properties of the underlying dark matter profile.  We conclude in section \\ref{sec:conclusions}.  Throughout, we use comoving distances unless specifically noted, with lowercase $r$ representing the three dimensional separation and uppercase $R$ representing projected quantities.  We have assumed a flat LCDM cosmology with $\\Omega_m=0.3$ and $\\Omega_\\Lambda=0.7$. Since we must compare stellar masses, which scale like $h^{-2}$ (for Hubble parameter $H_0=100 h$km/s/Mpc), with dynamical and halo masses, which scale like $h^{-1}$, our results are by necessity dependent on the choice of $h$.  We use $h=0.7$ for all quantities, and explicitly indicate their $h$-dependence using $h_{70}$; e.g., halo masses are indicated with $h_{70}^{-1}M_{\\odot}$, stellar masses with $h_{70}^{-2}M_{\\odot}$, distances with $h_{70}^{-1}$kpc, ratios of stellar to dynamical mass with $h_{70}^{-1}$, and similarly for other quantities.  When relating results from other works that use $h=1$, we simply use \\hkpc\\ or $h^{-1}M_{\\odot}$. ",
        "conclusions": "\\label{sec:conclusions} We have combined weak lensing and velocity dispersion observations to study the dark matter profiles of elliptical galaxies from the SDSS.  The radial profile of the weak lensing shear is consistent with that expected from an NFW halo profile.  We have fit two NFW parameters to the weak lensing data (mass and concentration), and extrapolated the profile inward to smaller radii not accessible to weak gravitational lensing.  We have deduced the dynamical mass at these smaller radii by measuring the velocity dispersion of the stars within the half light radius of the galaxy.   We compare the measured dynamical mass to the extrapolation of the NFW profile, and find that there is a significant excess of mass in the interior.   Using estimates of the stellar mass of the galaxy, and assuming that all non-stellar mass is dark matter, we find that the dark matter contribution to the dynamical mass is still far in excess of the NFW prediction.  This result is in support of the model of adiabatic contraction (AC), which predicts that the cooling and condensation of baryonic material will deepen the gravitational potential well of the galaxy, and pull the dark matter towards the halo center.   These results suggest that the effects of AC are stable to the subsequent bombardment of major and minor mergers suffered by these objects since the time their gas cooled.  We compare the observation to a theoretical model of adiabatic contraction \\citep{2004ApJ...616...16G} and find good agreement between the two.  However, we acknowledge that systematic uncertainties in the determination of the stellar mass make it difficult to prove that the excess mass at the half light radius is indeed dark matter.  An alternative interpretation may be that the excess mass is comprised of some other dark baryonic form, for example it may indicate an IMF with more of the mass concentrated in dim, low mass stars relative to the Kroupa IMF used in our stellar mass estimates.  In order to fully explain the observation without adiabatic contraction, the stellar masses would need to increase by a factor of two.  A change of this magnitude in the stellar masses would be difficult to reconcile with recent results from the SAURON collaboration. They conclude that a 30 per cent increase in stellar mass (from masses obtained with the Kroupa IMF) is inconsistent with their dynamical data derived from integral-field spectroscopy.  Finally, we divide the sample of galaxies along the mean size-luminosity relation and show that the scatter from the mean relation is related to the underlying properties of the dark matter profile.  We find that due to a higher baryon fraction with more of the baryons concentrated toward the center, the smaller/brighter subsample experiences more lasting effects of adiabatic contraction than the larger/dimmer subsample.  If adiabatic contraction is not the explanation, the most likely alternative is different IMFs for the two subsamples.   Also, we find that the two subsamples have comparable stellar masses but different halo masses, suggesting that stellar mass does not trace halo mass, possibly due to different formation histories that result in differing dark and stellar components (with larger halo mass being associated with a dimmer, less concentrated stellar component).  Ultimately, regardless of whether the excess mass is dark matter or baryonic, this set of observations indicates that the total mass at the half light radius far exceeds what would naively be expected from an NFW or Einasto profile. Contemporary models of galaxy formation and evolution must accommodate this relatively high ratio of dynamical mass in the interior to total halo mass in early type elliptical galaxies."
    },
    "0911/0911.0265_arXiv.txt": {
        "abstract": "We present the results from the spectral analysis of more than 7,500 {\\it RXTE} spectra of 10 AGN. Our main goal was to study their long term X--ray spectral variability. The sources in the sample are nearby, X-ray bright, and they have been observed by {\\it RXTE} regularly over a long period of time ($\\sim$7--11 years). High frequency breaks have been detected  in their power spectra, and these  characteristic frequencies imply time scales of the order of a few days or weeks. Consequently, the {\\it RXTE} observations we have used most probably sample most of the flux and spectral variations that these objects exhibit. Thus, the {\\it RXTE} data are ideal for our purpose. Fits to the individual spectra were performed in the 3--20 keV energy band. We modeled the data in a uniform way using simple phenomenological models (a power-law with the addition of Gaussian line and/or edge to model the iron K$\\alpha$ emission/absorption features, if needed) to consistently parametrize the shape of the observed X--ray {\\it continuum} of the sources in the sample. We found that the average spectral slope does not correlate with source luminosity or black hole mass, while it correlates positively with the average accretion rate. We have also determined the (positive) ``spectral slope -- flux\" relation for each object, over a flux range larger than before. We found that this correlation is  {\\it similar} in all objects, except for NGC~5548  which displays limited spectral variations for its flux variability. We discuss this global  ``spectral slope -- flux'' trend in the light of current models for spectral variability. We consider (i) intrinsic variability, expected e.g. from Comptonization processes, (ii) variability caused by absorption of X-rays by a single absorber whose ionization parameter varies proportionally to the continuum flux variations, (iii) variability resulting from the superposition of  a constant reflection component and  an intrinsic power-law which is variable in flux but constant in shape,  and, (iv) variability resulting from the superposition of a constant reflection component and an intrinsic power-law which is  variable both in flux and shape. Our final conclusion is that scenario (iv) provides the best fit to the data of all objects, except for NGC~5548. ",
        "introduction": "The current paradigm for the central source in AGN postulates a black hole (BH) with a mass of $10^6-10^9$ M$_{\\odot}$, and a geometrically thin, optically thick accretion disc that presumably extends to the innermost stable circular orbit around the black hole. This disc  is supposed to be  responsible for the broad, quasi-thermal emission component in the optical-UV spectrum of AGN (the so-called  Big Blue Bump).  At energies above $\\sim$2 keV, a power-law like component is observed. This is attributed to emission by a hot corona ($T \\sim 10^8 - 10^9$ K) overlying the thin disc. The corona up-Comptonizes the disc soft photons to produce the hard ($E \\sim 2 - 100$ keV) X-ray emission. An   Fe K$\\alpha$ line is also detected, produced either by fluorescence on cold matter at $E = 6.4$ keV or by photoionization of H-like or He-like Fe at 6.9 keV and 6.7 keV, respectively. This feature is believed to be emitted on the surface of the accretion disc, from the reprocessing of the X-rays. In many cases the line is resolved, and its width can, in theory, constrain the position of the inner radius of the disc and the spin of the central BH.  The AGN X--ray emission is strongly variable, even on time scales as short as a few hundred seconds, both in flux and in the shape of the observed energy spectrum. These variations can provide information on the  physical conditions, the size, and the geometry of the X--ray emitting region. Regarding the flux variability, although the broad band power spectra of AGN are not well understood, the identification of characteristic time scales in them, combined with the large range in the masses of the BHs that have been studied so far, can constrain (to some degree) the size and the geometry of the X--ray source \\citep[e.g.][]{mchardyea:2006}. On the other hand, recent detailed studies of the spectral variations observed in a few objects provide valuable details of the physical conditions of the X--ray reprocessing \\citep[see e.g.]{miniuttiea:2007} and/or on the physical properties of absorbing material which may exist close to the central source \\citep*[see e.g.][]{millerea:2008}.  Progress in settling alternative interpretations of the observed spectral variability can be made by correlating the spectral and timing properties in AGN. Along this line, in a recent paper \\citep[][]{papadakisea:2009} we studied 11 nearby, X--ray bright Seyferts, which have been intensively observed over the last two decades by the {\\it Rossi X-ray Timing Explorer} ({\\it RXTE}). We detected a positive correlation between their mean spectral slope and the characteristic break frequency, $\\nu_{\\rm br}$, at which the power spectrum changes its slope from -1 to -2. To the extent that this frequency corresponds to a (still unknown) characteristic  time scale of the X--ray source, this correlation suggests that the mean X--ray spectral slope in AGN is determined by a fundamental parameter of the system, most probably its accretion rate. If true, this result has important implications on current Comptonization models.  In this work we use the results from the model fitting of thousands of {\\it RXTE} spectra of 10 AGN that were included in the sample of \\citet[][]{papadakisea:2009}. Our prime goal is to study in detail the X--ray spectral variability of each object. Although the spectral resolution of {\\it RXTE} is not as good as that of {\\it Chandra}, {\\it XMM-Newton} or {\\it Suzaku}, the {\\it RXTE} data sets that we used are ideal for our purpose. First, the sources in the sample are nearby and X-ray bright. As a result, their X--ray spectral shape can be accurately determined even from relatively short ($\\sim$1--2 ksec) {\\it RXTE} observations. Second, each source in the sample was observed regularly (more than 200--300 times, and in a few cases more than a thousand times) over a long  period of time (of the order of 7--11 years). This is contrary to typical {\\it Chandra} and {\\it XMM-Newton} observations of individual AGN, which usually span a period of a few days (maximum). Third,  the characteristic frequencies of the objects imply time scales of the order of a few days or weeks. Consequently, the {\\it RXTE} observations we have used most probably sample most of the flux and spectral variations that they exhibit.  A subsample of the observations we study have already  been studied by \\citet[][]{papadakisea:2002}. These authors used hardness ratios in various energy bands to study the spectral variability of four objects, namely NGC~5548, NGC~5506, NGC~4051 and MCG~-6-30-15. In the present work we studied many more observations of these objects and we have also analysed data for six more Seyferts.  Furthermore, instead of computing hardness ratios, we fitted each spectrum with the same phenomenological model of a  ``power law plus a Gaussian line\"  to derive the observed photon index, $\\Gamma_{\\rm obs}$. In this way, we parametrized all available 3--20 keV {\\it RXTE} spectra of the sources in a uniform way. We then used the resulting best model fitting values to: i) determine the mean spectral slope for each object and investigate its correlation with the main source parameters such as BH mass, accretion rate and luminosity, and ii)  construct their long term ``spectral slope vs. flux\" relations. For the reasons explained above, we believe that we  determined the flux related spectral variability in these sources more accurately than before.  Several possible reasons for the observed X--ray spectral variability in AGN have been proposed in the past. The simplest possibility is that the $\\Gamma_{\\rm obs}$ variations correspond to intrinsic variations in continuum slope   \\citep*[e.g.][]{haardtea:1997,coppi:1999,beloborodov:1999}. For example, \\citet{nandrapapadakis:2001} have showed that  the X-ray spectral changes in NGC~7469 are intrinsic and  respond to the UV-soft photon input variations, exactly as one would expect in the case of thermal Comptonization models. However, there has also been  certain observational evidence that the $\\Gamma_{\\rm obs}$ variations originate as a result of superposition of different spectral components with different variability properties. One such possibility is a spectrum composed of a constant reflection component and a variable, in normalization only, power-law like continuum of constant slope \\citep*[e.g.][and references therein]{taylorea:2003,pontiea:2006,miniuttiea:2007}. Another  possibility is a constant slope power law which varies in flux and complex variations in an absorber, e.g. its ionization state and/or the covering factor of the source \\citep[see recent review by][]{turnermiller:2009}.  One of our main result is that the X--ray spectral slope steepens with increasing flux in a {\\it similar} way for all the objects in the sample. This similarity implies that, although all the above mentioned mechanisms may operate in AGN, it is just one of them which is mainly responsible for the observed spectral variability in AGN.  The paper is organized as follows. In Section 2, we discuss the data selection and reduction, and we also describe the  phenomenological models that we used to parametrize the shape of AGN continuum. In Section 3, we report the spectral fitting results, and the results from the spectral variability study.  In Section 4, we discuss our findings and compare them with predictions of several spectral variability models. We give our conclusions in Section 5. ",
        "conclusions": "We analyzed more than 7,500 \\rxte\\ PCA data of 10 nearby AGN with the goal of studying their spectral variability properties. The objects were observed regularly with \\rxte\\ over many years since 1996. Given the large number of observations, spread over many years, we were able  to determine accurately their average observed spectral slope, luminosity and accretion rate.  Using the best model-fit results from the hundreds of spectra of each source, we also determined their long-term ``observed spectral slope -- X--ray flux'' relation. This kind of long term variability analysis can not be accomplished with the rather short AGN observations currently performed by {\\it Chandra}, {\\it XMM-Newton} or {\\it Suzaku}, despite their higher effective area and superior spectral resolution. Our main results are as follows:  \\begin{list}{\\labelitemi}{\\leftmargin=0em\\itemsep=0.5em}  \\item[1.] The average spectral slope is {\\it not} the same in all objects. It does not correlate either with BH mass or luminosity. However, $\\overline{\\Gamma}$ correlates positively with the mass accretion rate: objects with higher accretion rate have steeper spectra, and $\\Gamma \\simeq 2.7\\dot{m}^{0.08}$. This result can be explained if the Compton amplification factor, $A$, decreases proportionally with the accretion rate in AGN.  \\item[2.] We detected the iron K$\\alpha$ line in many spectra of seven sources in the sample. The average EW of the line anti-correlates with the average 2--10 keV luminosity (i.e. the so-called Iwasawa-Taniguchi effect). Our results are in agreement with recent results which are based on recent studies of large samples of AGN \\citep[][]{bianchiea:2007}.  \\item[3.] The spectra of each source  become steeper with increasing flux. This is a well known result. What we have shown in this work that this trend holds over many years, and over almost the full range of flux variations that AGN exhibit. We also found that the $\\Gamma_{\\rm obs}-F_{\\rm 2-10}$ relations are {\\it similar} in all objects, except for NGC~5548. There are many reasons which can cause apparent and/or intrinsic  variations in the X--ray spectrum of AGN. It is possible that they all operate, to a different extent in various objects. However, the common spectral variability pattern we detected implies that just a single mechanism is responsible for the bulk of the observed spectral variations.  \\item[4.] NGC~5548 displays limited spectral variations for its flux variability. Although uncommon, this behaviour has also been observed in other Type I Seyfert, PG 0804+761 \\citep*[][]{papadakisea:2003}. This behaviour, different to what is observed in most (Type I) AGN, raises the issue of different spectral states in AGN, just like in GBHs. Investigation of this issue is beyond the scope of the present work.  \\item[5.] The scenario in which the spectral variability is caused by absorption of X-rays by a {\\it single} medium whose ionization parameter varies proportionally to the continuum flux variations fails to account for the observed spectral variations.  \\item[6.] A ``power-law, with constant $\\Gamma_{\\rm intr}$ and variable flux, plus a constant reflection component\" can fit the observed ``$\\Gamma -F_{\\rm 2-10}$'' relations of most (but not all) objects in the sample. The flux of the reflection component necessary to explain the observed spectral variability is rather large, implying average reflection amplitudes of the order of $R\\sim 5$. If true, this result implies that the primary source of X-rays is located close to a maximally rotating Kerr black hole, where relativistic effects (such as light bending) are expected to be strong.  \\item[7.] The observed ``$\\Gamma -F_{\\rm 2-10}$'' relations (except for NGC~5548) can be explained if we assume that the power-law continuum varies both in flux and shape as $\\Gamma_{\\rm intr}\\propto F_{\\rm 2-10}^{0.08}$ (as implied by the {\\it average} ``spectral slope - accretion rate\" relation that we observed) {\\it and} the reflecting material responds to the average continuum spectrum, with $R=1$.  \\end{list}"
    },
    "0911/0911.0924_arXiv.txt": {
        "abstract": "The first  study \\langeditorchanges{of} migration-induced resonances in a pair of Earth-like planets has been performed by Papaloizou and Szuszkiewicz (\\cite{ps}). \\langeditorchanges{They} concluded that in \\langeditorchanges{the} case of disparate masses embedded in a disc with the surface density expected for a  minimum mass solar nebula at 5.2 au, the most likely resonances are \\langeditorchanges{ratios of large integers, such as} 8:7. For equal masses, planets tend to enter into the 2:1 or 3:2 resonance. In Papaloizou and Szuszkiewicz (\\cite{ps}) the two low-mass planets have \\langeditorchanges{masses} equal to 4 Earth masses, chosen to mimic the very well known example of two  pulsar planets which are close to the 3:2 resonance. That study has stimulated quite a few interesting questions. One of them is considered here, namely how the behaviour of the planets close to the \\langeditorchanges{mean-motion} resonance depends on the actual \\langeditorchanges{values} of the \\langeditorchanges{masses} of the planets. We have chosen a 3:2 commensurability and investigated the outcome of an orbital migration in the vicinity of this resonance in the case of \\langeditorchanges{a pair} of equal mass \\langeditorchanges{super-Earths,} whose mass is \\langeditorchanges{either 5 or 8} Earth \\langeditorchanges{masses.} ",
        "introduction": "In the first study \\langeditorchanges{of} migration-induced resonances in a pair of low-mass planets Papaloizou and Szuszkiewicz (\\cite{ps}) \\langeditorchanges{performed} simulations of two interacting planets embedded in a disc with which they also interact (see also Szuszkiewicz, this volume). The simulations \\langeditorchanges{were} carried out for a variety of planet masses, initial orbital separations and initial surface densities in order to explore the possible outcomes. In the case of planets with equal masses when the  relative migration is slow, Papaloizou and Szuszkiewicz (\\cite{ps}) have found that the planets become trapped in the nearest available resonance. In their \\langeditorchanges{study,} the two low-mass planets \\langeditorchanges{each have a mass of} 4 Earth masses, chosen to mimic the very well known example of two  pulsar planets which are close to the 3:2 resonance. In this work we have modelled \\langeditorchanges{the} evolution of \\langeditorchanges{a pair} of planets of equal mass in a gaseous protoplanetary \\langeditorchanges{disc, again} in the vicinity of a 3:2 \\langeditorchanges{mean-motion} resonance, but \\langeditorchanges{for two different, larger, values of the mass, namely,} 5 and 8 Earth \\langeditorchanges{masses}.  %  We have used the hydrodynamical code \\langeditorchanges{NIRVANA,} originally written by Ziegler (\\cite{Ziegler}) and then adapted to disc-planet simulations. More details about the code can be found in the work by Nelson {\\em et al.\\/} (\\cite{nelson}). The kinematic viscosity of the disk is zero and no magnetic field is present. The accretion of matter onto  planets is not included. The units used in the code are the following. The mass of the central object is taken to be the mass unit. The initial distance $r_{p2}$ of the inner planet from the central body  is \\langeditorchanges{the} length unit. The time is measured in $\\left(GM_{\\odot}/{{r_{p2}}^3}\\right)^{-1/2}$. The profile of the initial surface density in the disc is essentially flat and it is described in \\langeditorchanges{detail} in Papaloizou and Szuszkiewicz (\\cite{ps}). The initial value of the surface density \\langeditorchanges{$\\Sigma_0$ at the locations of the planets} is $\\Sigma_0=6.0\\times 10^{-4}$ \\langeditorchanges{in our units}. The computational domain in polar coordinates \\langeditorchanges{extends from 0.33 to 4} length units in the radial direction and \\langeditorchanges{covers} all azimuthal angles from 0 to \\langeditorchanges{2$\\pi$, and} was divided into $384$ radial slices and $512$ angular sectors. The disc model is locally isothermal with aspect ratio $H/r=0.05$ ",
        "conclusions": "The planets with $5$ and $8\\mathrm{M}_\\oplus$ entered the resonance and continued evolution in this state till the end of our experiment. The mean values of the orbital eccentricities are excited to the value of $\\sim 0.008$ in both cases and keep oscillating around this value. \\langeditorchanges{The} eccentricities of the internal and external planets are roughly the same. The ratio of the orbital semi-major axes \\langeditorchanges{slowly approaches} the exact \\langeditorchanges{resonance}. \\langeditorchanges{A} full discussion of these results will appear in \\langeditorchanges{a} forthcoming paper."
    },
    "0911/0911.3320_arXiv.txt": {
        "abstract": "{We present initial result of a large spectroscopic survey aimed at measuring the timescale of mass accretion in young, pre-main-sequence stars in the spectral type range K0 -- M5. Using multi-object spectroscopy with VIMOS at the VLT we identified the fraction of accreting stars in a number of young stellar clusters and associations of ages between 1 -- 50 Myr. The fraction of accreting stars decreases from $\\sim$ 60~\\% at 1.5 -- 2 Myr to $\\sim$ 2~\\% at 10 Myr. No accreting stars are found after 10 Myr at a sensitivity limit of $10^{-11}$ \\myr. We compared the fraction of stars showing ongoing accretion (\\facc) to the fraction of stars with near-to-mid infrared excess (\\fIRAC). In most cases we find \\facc ~$<$ \\fIRAC ~, \\ie, mass accretion appears to cease (or drop below detectable level) earlier than the dust is dissipated in the inner disk. At 5 Myr, 95~\\% of the stellar population has stopped accreting material at a rate of $\\gtrsim 10^{-11}$ \\myr, while $\\sim$ 20~\\% of the stars show near-infrared excess emission. Assuming an exponential decay, we measure a mass accretion timescale (\\tacc) of 2.3 Myr, compared to a near-to-mid infrared excess timescale (\\tIRAC) of 2.9 Myr. Planet formation, and/or migration, in the inner disk might be a viable mechanism to halt further accretion onto the central star on such a short timescale.} ",
        "introduction": "Circumstellar disks around young, pre-main-sequence stars are the natural loci of planet formation \\citep[\\eg][]{henning08}. Such a {\\em protoplanetary} disk is formed during the collapse of the molecular cloud to conserve the initial angular momentum. It is made of interstellar gas and dust. The conventional planet formation model is the so-called {\\it core-accretion} model \\citep[\\eg][]{pollack96,mordasini08}. In the simulations of \\citet{pollack96}, the formation of solar-system like planets takes up to 16 Myr. Giant planet formation is much faster ($\\sim$ 3 Myr) if planet migration is included \\citep{alibert04,alibert05}. Gravitational instability in the disk was also proposed as a viable scenario to form planets, especially in the outer disk \\citep{boss97, boss98, durisen07}. If gravitational instabilities occur, the disk may fragment and form dense self-gravitating clumps which are the precursor of gas giant planets. In this model planet formation is very fast ($\\lesssim$ 1 Myr).  Planet formation models face the fast evolution of protoplanetary disks. Star forming regions show a steady decline with time in the fraction of stars having infrared excess emission (e.g. \\citealp{haisch01, haisch05, bouwman06, hillenbrand08}). This is thought to be caused by gradual clearing of dust in the inner disk. The growing consensus is that the warm small dust grains disappear within the first 5$\\div$9 Myr (\\tdust). The quantity \\tdust ~is sometimes adopted as the disk lifetime. Although dust is essential to form a planet, in terms of disk mass, and hence dynamics and evolution, it accounts only for a small percentage. The bulk of the disk mass is thought to be gaseous. At present, it is unclear whether the gas dissipation timescale (\\tgas) is similar to \\tdust.  In this paper we report on a study aimed at determining the timescale for mass accretion (\\tacc) in protoplanetary disks. The timescale \\tacc, i.e. the time at which the disk accretion phase ceases, provides a strong constraint on \\tgas. Gas may still be present after \\tacc, however, the amount of remaining gas, and hence of disk mass, may be too low to be able to form giant planet(s). Our observational strategy, observations and data reduction are presented in section 2. Analysis and results are presented in section 3 and 4 respectively. In section 5 we compare our result with literature data. Conclusions are drawn in section 6.   \\begin{table*} \\caption{Regions observed with VIMOS. The surveyed area for each cluster is 16\\arcmin$\\times$18\\arcmin  ~centered at the coordinates listed below. The limiting magnitude is V = 21 mag. MOS observations were carried out with the ``HR-Orange'' grism (maximum spectral coverage: 5200 -- 7600 \\AA; $\\lambda / \\Delta\\lambda$ = 2150) and a slit width of 1$\\arcsec$.} \\label{tab:obs} \\centering \\begin{tabular}{lllllll} \\hline\\hline Cluster          & Observing period & RA(J2000) &DEC(J2000)& Distance & A$_V$ & Sources \\\\%& min. M \\\\ & & [hh:mm:ss]&[dd:mm:ss]& [pc]     & [mag] & [N]     \\\\%& [\\msun]\\\\ \\hline $\\sigma$ Orionis & Oct. 2006 -- Mar. 2007 & 05:38:00 & -02:37:00 & 360     & 0.15  & 216  \\\\%   & 0.03 \\\\ NGC 6231         & Apr. 2008 -- Sep. 2008 & 16:54:10 & -41:49:30 & 1240    & 1.36  & 573  \\\\%   & 0.25 \\\\ NGC 6531         & Apr. 2008 -- Sep. 2008 & 18:04:13 & -22:29:24 & 1210    & 0.87  & 627  \\\\%   & 0.3  \\\\ ASCC 58          & Oct. 2006 -- Mar. 2007 & 10:15:07 & -54:58:12 & 600     & 0.28  & 370  \\\\%   & 0.1  \\\\ NGC 2353         & Oct. 2006 -- Mar. 2007 & 07:14:30 & -10:16:00 & 1120    & 0.22  & 316  \\\\%   & 0.3  \\\\ Collinder 65     & Oct. 2006 -- Mar. 2007 & 05:26:05 & +15:41:59 & 310     & 0.40  & 198  \\\\%   & 0.15 \\\\ NGC 6664         & Apr. 2008 -- Sep. 2008 & 18:36:42 & -08:13:00 & 1440    & 2.20  & 508  \\\\%   & 0.5  \\\\ \\hline\\hline \\end{tabular} \\end{table*} ",
        "conclusions": "In this paper we presented an analysis of the evolution of mass accretion in PMSs in the spectral type range K0 -- M5. These are the main results: (1) the fraction of stars with ongoing mass accretion decreases fast with time, going from $\\sim$ 60\\% at 1.5 -- 2 Myr down to $\\sim$ 2\\% at 10 Myr; (2) this fraction is systematically lower than the fraction of stars showing near-to-mid infrared excess; (3) mass accretion and dust dissipation in the inner disk appear to have different characteristic timescales, 2.3 and 2.9 Myr respectively; (4) within 5 Myr the mass accretion rate of 95\\% of the stellar population drops below our detection limit of $10^{-11}$ \\myr. While viscous evolution and photoevaporation might be unable to slow down accretion (and leave a substantial dust mass) on such a short timescale, planet formation, and/or migration, in the inner disk (few AU) might be a viable mechanism to halt further accretion onto the central star."
    },
    "0911/0911.1113_arXiv.txt": {
        "abstract": "It has been claimed recently that Swift has captured for the first time a late-time afterglow re-brightening of clear nonflaring origin after the steep decay of the prompt emission in a long gamma-ray burst (GRB), which may have been produced by a narrow jet viewed far off-axis. However, this interpretation of the observed re-brightening of the X-ray afterglow (AG) of GRB081028 is unlikely in view of its large equivalent isotropic gamma-ray energy. Moreover, we show that the late-time re-brightening of the AG of GRB081028 is well explained by the cannonball (CB) model of GRBs as a synchrotron flare emitted when the jet that produced the GRB in a star formation region (SFR) in the host galaxy crossed the SFR boundary into the interstellar medium or the halo of the host galaxy. We also show that all the other observed properties of GRB081028 and its afterglow are well reproduced by the CB model. On the other hand, we demonstrate that far-off axis GRBs, which in the CB model are `soft' GRBs and XRFs and consequently have much smaller isotropic equivalent gamma-ray energies, have slowly rising afterglows with a late-time power-law decay identical to that of ordinary GRBs, in good agreement with the CB model predictions. ",
        "introduction": "A smoothly rising X-ray afterglow (AG) following the steep decay of the prompt emission in a long gamma ray burst (GRB) was first observed in GRB970508 with the Narrow Field Instrument (NFI) aboard the BeppoSAX satellite (Piro et al.~1998). Contemporaneous optical observations with ground based telescopes detected a corresponding re-brightening of its optical afterglow beginning around $6\\times 10^4$ sec after burst (e.g., Galama et al.~1998 and references therein). This re-brightening was interpreted by Piro et al.~(1998) as due to a delayed activity of the central GRB engine.  Chiang \\& Dermer~(1997) and Dar \\& De R\\'ujula~(2000) showed that a rise in the afterglow of a GRB produced by a jet of highly relativistic plasmoids (Shaviv \\& Dar 1995) can result also from its deceleration if its viewing angle was initially far outside its relativistic beaming cone (far off-axis viewing angle). Both interpretations, however, failed to reproduce well enough the lightcurve of the AG of GRB970508. Consequently, Dado, Dar \\& De R\\'ujula (hereafter DDD) 2002 have investigated alternative explanations and showed that the late-time re-brightening of the AG of GRB970508 could be explained by the encounter of the jet with a density bump such as that expected at the boundary of the star formation region (SFR) where the GRB took place. A similar re-brightening of the X-ray afterglow of GRBs was later observed with the Swift X-ray telescope (XRT) in a few GRBs, e.g., in GRB060614 (Mangano et al.~2007) and in several other GRBs where gaps in the data sampling and low statistics prevented a firm conclusion (e.g., GRBs 051016B and 070306, Swift lightcurve repository, Evans et al.~2007,2009).  Overlooking the above, Margutti et al.~(2009) have recently claimed that Swift has captured for the first time a smoothly rising X-ray re-brightening of clear non-flaring origin after the steep decay of the prompt emission in a long gamma-ray burst GRB081028, which may have been produced by a narrow jet viewed far off-axis. But, the `far off-axis viewing' of GRB081028 is not supported by its large equivalent isotropic gamma-ray energy: The radiation emitted by a highly relativistic jet with a bulk motion Lorentz factor $\\gamma\\!\\gg\\!1$, which is observed from a small angle ($\\theta\\!\\ll\\!1$) relative to its motion, is amplified by a factor $\\delta^3$ resulting from relativistic beaming and Doppler boosting, where the Doppler factor $\\delta$ is well approximated by $\\delta(t)\\!=\\![\\gamma(t)\\,(1-\\beta(t)\\, cos\\theta]^{-1}\\!\\approx 2\\!\\gamma/(1\\!+\\!\\gamma^2\\,\\theta^2)$ (e.g. Shaviv \\& Dar 1995; Dar 1998). Hence, for the typical viewing angle of GRBs, $\\theta\\!\\sim\\!1/\\gamma$, $\\delta\\!\\sim\\! \\gamma$, and the observed GRB fluence is amplified by a factor $\\gamma^3$ relative to that in the jet rest frame. For `far off-axis viewing' angles where $\\gamma^2\\,\\theta^2\\!\\gg\\!1$, this amplification of $E_{iso}$ is reduced by a factor $[(1\\!+\\!\\gamma^2\\,\\theta^2)/2]^{-3}$. For example, for $\\theta\\!=\\! 3/\\gamma$, $E_{iso}$ is reduced by a factor $8\\times 10^{-3}$ and GRBs viewed at $\\theta\\geq 3/\\gamma$ appear as X-ray flashes (XRFs) rather than ordinary GRBs (see, e.g., Dar \\& De R\\'ujula 2000; DDD2004). From the observations of GRB081028, with the Burst Alert Telescope (BAT) aboard Swift, Margutti et al.~(2009) inferred that $E_{iso}\\!=\\!(1.1\\!\\pm\\!0.1) \\times 10^{53}$ erg, which is typical of ordinary GRBs and is inconsistent with the far off-axis interpretation of the re-brightening of its AG.  Despite the many serious discrepancies between the observed properties of GRB081028 and the predictions of the fireball (FB) model, which were pointed out by Margutti et al.~(2009) and are quite common in many other GRBs, alternative models of GRBs were ignored in their paper. One such model is the cannonball (CB) model, which has been shown repeatedly to reproduce well the main observed properties of GRBs and their AGs (DDD2009a and references therein). Thus, in this paper, we demonstrate that the main observed properties of GRB081028 and its afterglow are well reproduced by the CB model of GRBs. In particular, we show that the observed late-time re-brightening of the AG of GRB081028 is well explained as being a late-time synchrotron radiation (SR) flare. Such flares are expected to be emitted by the jet which produced the GRB, when it crosses the boundary of the star formation region, where presumably the GRB took place, into the ISM or the halo of the host galaxy. Moreover, we show that GRBs which have much larger viewing angles than those of ordinary GRBs, and thus appear as X-ray flashes (XRFs) or `soft' GRBs, have rather slowly rising late-time afterglows which are well described by the CB model. ",
        "conclusions": "In the CB model, most XRF and soft GRBs are ordinary GRBs observed from a far off-axis viewing angle, i,e, viewing angles much larger than those of typical GRBs. This implies smaller value of $(\\!1+\\!z)\\,E_p$, much smaller value of $E_{iso}$, prompt pulses, flares and afterglows much more stretched in time, and a slowly rising late-time afterglow, which turns into an asymptotic power-law decay $\\!\\sim\\! t^{\\!-\\!\\Gamma\\!+\\!1/2}$, as demonstrated here for XRF080707 and the soft GRBs 091029 and 091127.  GRB081028 was observed from space with the BAT and XRT aboard Swift and from the ground with optical telescopes. The observations provided detailed information on its properties, which differ from those expected from a far off-axis GRB. The equivalent isotropic energy $E_{iso}$ and peak energy $(1\\!+\\!z)\\,E_p$ of its prompt emission are typical of ordinary GRBs and satisfy the [(1+z)Ep,Eiso] correlation for ordinary GRBs (Fig.~\\ref{f0}):  The measured value of $E_{iso}$ yields $(1+z)\\,E_p\\!\\sim\\! 280$ keV, in good agreement with its measured value $(1\\!+\\!z)\\,E_p\\!=\\!222^{+81}_{-36}$ keV, which is much larger than the typical values of $E_p$ in XRFs. Its lightcurves are well described by the CB model assuming an ordinary GRB:  The BAT 15-150 keV lightcurve of the prompt emission is well described by a sum of two pulses/flares produced by two CBs by ICS of glory photons (Fig.~\\ref{f1}). Its 0.3-10 keV X-ray lightcurve, which begins during the second peak, shows a steep decay of the prompt emission accompanied by a rapid spectral softening with flares superimposed. It is well described by a sum of three X-ray flares produced by ICS of glory photons by the electrons in CBs late ejected  by a weakening central engine (Fig.~\\ref{f2}). The smoothly rising X-ray AG, which follows the steep decay, peaks around 22 ks, after which it turns into a power-law decay. This behaviour of the AG is well described by a late-time SRF (Fig.~\\ref{f2}). In particular, the photon spectral index during the late time flare remains constant, $\\Gamma_X\\!=\\!2.04\\!\\pm\\!0.06$, as expected in the CB model for SRFs. Moreover, the power-law decay of the SRF has an index $\\alpha$=2.09 which satisfies within errors the CB model prediction (DDD2009a) $\\alpha_X\\!=\\!\\beta_X\\!+\\!1=\\!\\Gamma\\!=\\!2.091\\!\\pm\\!0.063)$, Had it been an ordinary synchrotron AG of a far-off axis GRB, i.e.,  like that of XRFs, its temporal index would have been $\\alpha_X\\!=\\!\\beta_X\\!+\\!1/2\\!=\\!\\Gamma\\!-\\!1/2 \\!\\approx\\!1.6 $ (DDD2009a), independent of viewing angle, in contradiction with the observed value, $\\alpha\\!\\sim\\! 2.1$. These temporal and spectral properties of the late-time flare in the AG of GRB081028 and the normal values of $E_{iso}$ and $E_p$ of its prompt emission agree well with the CB model interpretation of the re-brightening of its afterglow - i.e., a late-time SR flare produced by the same jet that produced GRB081028, presumably during its passage through the boundary of the star formation region where GRB081028 took place, into the ISM or the halo of the host galaxy.   \\noindent {\\bf Acknowledgment} The authors are grateful to Raffaella Margutti for kindly providing the tabulated data used in Figs.~1,3,4."
    },
    "0911/0911.3785_arXiv.txt": {
        "abstract": "{X-ray data analysis have found that fairly complex structures at cluster centres are more common than expected. Many of these structures have similar morphologies, which exhibit spiral-like substructure.} {It is not yet well known how these structures formed or are maintained. Understanding the origin of these spiral-like features at the centre of some clusters is the major motivation behind this work.} {We analyse deep \\textit{Chandra} observations of 15 nearby galaxy clusters (0.01 $ < z < $ 0.06), and use X-ray temperature and substructure maps to detect small features at the cores of the clusters.} {We detect spiral-like features at the centre of 7 clusters: A85, A426, A496, Hydra A cluster, Centaurus, Ophiuchus, and A4059. These patterns are similar to those found in numerical hydrodynamic simulations of cluster mergers with non-zero impact parameter. In some clusters of our sample, a strong radio source also occupies the inner region of the cluster, which indicates a possible connection between the two. Our investigation implies that these spiral-like structures may be caused by off-axis minor mergers. Since these features occur in regions of high density, they may confine radio emission from the central galaxy producing, in some cases, unusual radio morphology.} {} ",
        "introduction": "Previous studies have shown that many clusters continue to be in the process of forming, groups or individual galaxies being accreated from the outer large scale structure filaments \\citep[e.g.,][]{Durret03}. Temperature maps of the X-ray emitting gas have shown that even clusters with apparently relaxed structures in X-ray can have perturbed temperature maps, which are indicative of recent ongoing mergers \\citep[e.g.,][]{Finoguenov05,Durret08}. The detection of these small structures is one of the main achievements of the present era of X-ray telescopes such as {\\it Chandra} and XMM-{\\it {\\textit Newton}}. However, the presence of cooling cores suggest that recent merger(s) have not had enough time to destroy the cooling core.  More than 70\\% of cooling-flow clusters contain a central cD galaxy that is a radio source \\citep[e.g.,][]{Burns90,Eilek04}.\\textit{Chandra} observations of these systems found that the core of many clusters exhibit morphological complexities that are probably produced by the interaction between the intracluster medium (ICM) and the central radio galaxy.  The most important connection between the X-ray ICM and the non-thermal relativistic electrons that produce radio emission is dynamical. In the case of very energetic jets, the ICM is heated and compressed because of supersonic shocks \\citep{Heinz98,Reynolds01}. However, there are weaker jets that do not provide an efficient shock heating, but instead a bubble of low-density gas, which can be identified as a depression in the X-ray surface brightness, that rises because of buoyancy \\citep[e.g.,][]{Churazov00,Bruggen02,McNamara05}.  On the other hand, the ICM can also affect the radio emission from radio galaxies by placing upper limits on the radio emission. The unusual radio morphology of central cluster galaxies that fill X-ray cavities of low density has been reported previously \\citep{Bohringer93,Taylor94,Fabian00,Fabian02,Taylor02}. The most extreme case of radio distortion is the amorphous radio galaxy PKS 0745-191 at the core of the Centaurus cluster \\citep{Taylor94}.  Another interesting feature is the spiral-like feature at the core of Perseus \\citep{Churazov03,Fabian00,Taylor02}, A2204 \\citep{Sanders08}, and A2029 \\citep{Clarke04}. This spiral pattern was also found in hydrodynamical simulations of merging clusters with non-zero impact parameters \\citep{Gomez02,Ascasibar06}. In these simulations, the infalling sub-cluster causes a disturbance in the mass peak as it passes through the cluster centre. Thus, the central cool gas of the main cluster acquires angular momentum, producing its spiral pattern.  To understand more clearly the ICM-radio interaction and to characterize the spiral feature, we present a sample of 15 nearby clusters, among which 7 exhibit this feature in temperature and substructure maps, and a large cooling-flow and an embedded radio galaxy at the centre. The paper is organized as follows: in Sect.~\\ref{sample}, we describe the data sample and the {\\it Chandra} data analysis; in Sect.~\\ref{imaging}, we explain the image procedures used to detect the spiral-like pattern; in Sect.~\\ref{res}, we present our results and notes for individual clusters; in Sect.~\\ref{disc}, we discuss a possible scenario to explain the spiral-like feature formation and the ICM-radio connection; in Sect.~\\ref{SyntImg}, we present a synthetic image produced to reproduce the surface brightness of the Perseus cluster, and in Sect.~\\ref{conc}, we summarize our findings.  All distance-dependent quantities are derived by assuming the Hubble constant $H_{0}=70$~km~s$^{-1}$~Mpc$^{-1}$, $\\Omega_{M}=0.3$ and, $\\Omega_{\\Lambda}$= 0.7. ",
        "conclusions": "\\label{conc}  In the most widely accepted scenario of structure formation, clusters of galaxies form hierarchically by means of mergers. Besides the modification of galaxy distribution and triggering of star formation, these merger events produce variations in the physical properties, such as density and temperature variations that can be observed in X-rays. In this sense, the substructure maps have provided a powerful means of identifying hidden structures that are the last record of the merger history.  We have presented 7 (A85, A426, A496, Hydra A cluster, Centaurus, Ophiuchus and A4059) out of 15 clusters that exhibit for spiral-like substructure correlated spatially with the boundaries of cold-fronts and cavities. The cavities inside these patterns are, in most of cases, associated with radio emission. The aim of this study has been to determine the frequency of this morphological pattern at the centres of nearby clusters. Our analysis has shown how observational studies can improve the constraints on numerical simulations of structure formation and improve our understanding of ICM physical properties.  Comparing our results to numerical simulations \\citep{Ascasibar06}, we have discussed the suitability of a model in which this spiral-like pattern is produced by an off-axis minor-merger that displaces the central cool gas. Since the minor merger does not disrupt the cool core, the central cool gas is still driven towards the central object that may produce the radio emission. This radio emission is bounded by the ICM inhomogeneities that will probably create an overall radio appearance that is quite different from most radio galaxies. Thus, we suggest that these distorted radio-emission morphologies may be caused by the confinement of dense X-ray gas in the spiral pattern. However, our interpretation of the spiral-like pattern and its relation to X-ray cavities and radio emission should be tested by additional more sophisticated simulations that ascertain possible merger signatures."
    },
    "0911/0911.3050_arXiv.txt": {
        "abstract": "A gravitational lens system can be perturbed by ``rogue systems\" in angular proximities but at different distances. A point mass perturbed by another point mass can be considered as a large separation approximation of the double scattering two point mass (DSTP) lens. The resulting effective lens depends on whether the perturber is closer to or farther from the observer than the main lens system. The caustic is smaller than that of the large separation binary lens when the perturber is the first scatterer; the caustic is similar in size with the large separation binary lens when the perturber is the last scatterer. Modelling of a gravitational lensing by a galaxy requires extra terms other than constanst shear for the perturbers at different redshifts. Double scattering two distributed mass (DSTD) lens is considered. The perturbing galaxy behaves as a monopole -- or a point mass -- because the dipole moment of the elliptic mass distribution is zero. ",
        "introduction": "The Galactic bulge is being surveyed for gravitational microlensing in search of microlensing planets. The lensing systems are the standard single scattering $n$-point mass lenses, where the bound system may consist of a single star, a multiple star, or a planetary system with one or two host stars. The Galactic bulge microlensing probability is $\\sim 10^{-6}$ and the probability for an unbound system to be aligned with the main lensing system can be ignored because it is $\\sim 10^{-12}$. The number of stars being surveyed is less than $10^9$. The lensing probability is proportional to the Einstein ring radius square, and the Einstein ring radius of the lensing toward the Galactic Bulge is characteristically $\\sim 1$ mas. Thus, if we consider the probability of a ``rogue\" system to be within $1$ as from the main lensing system, the probability is $\\sim 1$. In fact, ground-based microlensing events are known to be ``plagued\" with blending of light, and some of them other than the main lens itself can also be gravitationally relevant for the photon path. The perturbers can be dark as well.  The probability for the ``rogue\" system to be at the same distance as the main lens system is small where the ``same distance\" should be understood in the context that the two lens systems are within the coherence scale in which the lensing can be considered a single scattering lensing. Within the coherence scale, the probability for the photon path to weave through the two lenses can be ignored \\citep{RCB} (RB10 from here on).  Thus it is most reasonable to assume in general that the two gravitationally unbound microlens systems are at different distances, and they would be best considered as a double scattering lensing system. Here it is assumed that the two lens elements are widely separated in the sky based on the argument in the previous paragraph and calculate the effects of the ``rogue\" system  on the main lens. The double scattering lensing is a time-sequential process, and it matters whether the ``rogue\" system is farther away than the main lens from the observer or closer. The two cases are schematically shown in figure \\ref{fig_twocases}. By wide separation it is implied that the separation $\\ell$ is much larger than the Einstein ring radius of the main lens.  We assume that the ``rogue\" perturbing system is a single point mass. The main lens of interest will be a single star, a multiple star, a planet system with one or two host stars, or a wide binary stars one of which hosts planets. Here we consider the most common and simplest case of a signle star and study the wide separation approximation of the double scattering two point mass (DSTP) lens.  Then the most important effect of the perturbation is to break the degeneracy of the point caustic of the single lens to an extended caustic curve, and the size of the caustic will be the indicator of the influence of the perturber. It will be shown that the caustic size depends on the distances of the lenses and whether the perturber is in the front or in the back. When the perturber is the first scatterer, the caustic is smaller than that of the binary lens, and it is similar in size when the main lens is perturbed by a ``rogue\" system in front. A binary lens forms when the two lenses have the same distance -- or within the coherence length.   The binary lens at large separation is made of a point mass and a constant shear (and plus the source shift), and it is briefly discussed in the appendix.  A multiple-point mass lens perturbing a single point mass main lens is approximated by the same form of the approximate DSTP lens equation (with effective coefficients) and can be concluded to behave in the similar manner to the single point mass perturber.   In lensing by a galaxy, modelling is done customarily by assuming an elliptic mass (sometimes replaced by an elliptic potential) and a constant shear. The galaxy lensing is of order $1$ arcsecond, and there are often other galaxies in the angular vicinity of the main lens system.  The perturbers can be group memebers of the main lens or galaxies at different distances. If there is a perturbing galaxy at the same distance, its monopole  will add an external shear as is the case with the binary lens. However, if the perturbing masses are at different distances, the perturbation should include deflection terms other than the constant external shear. We consider the wide separation approximation of the double scattering two distributed mass (DSTD) lens. The other terms depend on the double scattering parameter and vanish when the perturbers are at the same distance as the main lens because the double scattering parameter vanishes. The perturbation by an elliptic mass galaxy is approximated by the perturbation by a point mass (monopole) because the dipole moment of the elliptic mass distribution is zero. The main lens galaxy, assumed to have an elliptic mass distribution, has a finite size caustic, and the effect of the perturber is to change its shape, size, and position. Even in the simpler case of the main galaxy as the monopole-quadrupole lens requires numerical calculations. We leave the perturbations of the finite size caustics for future work.  It should be necessary to point out that \\citet{keeton} uses Taylor expansion and concludes that the effects of a perturbing mass on the galaxy lens is to add constant convergence and constant shear irrelevantly of whether the perturber is the first scatterer or the last. There may be a problem in the expansion cutoff. In the region of interest around the critical curve of the galaxy lens, the first term in the Taylor expansion is small because the Jacobian determinant is zero or small and is likely to be smaller than the second order term. The second order term is well known for the square root behavior of the lensing near the critical curve or caustic crossing. It is not clear whether the Taylor expansion can be used at all. We use power expansion around the critical curve of the main lens which is the region of interest.   The DSTD lens equation is an obvious extension of the DSTP lens equation in which the delta function integral for the 2-d gravitational field of a point mass is generalized to the density function integral for the 2-d gravitational field of the distributed mass.   The DSTP lens equation is known since 1986 (Blandford and Narayan) and have been studied \\citep{KA86, ES93, werner}. Here the derivation of the DSTP lens equation studied in RB10 is briefed for clarity and convenience. Instead of using the formula for the time delay and Fermat principle, the well-known derivation of the single lens equation from an exact solution of the general relativity, the Schwarzschild metric, with the assumption of the linear gravity and small angle approximation is used. The Schwarzschild metric is asymptotically flat and the scattering planes can be joined easily in the asymptotic regions. The DSTP lens equation is obtained by joining two scattering planes with the freedom to rotate. ",
        "conclusions": ""
    },
    "0911/0911.1864_arXiv.txt": {
        "abstract": "We apply the empirical method built for redshift $z=0$ in the previous work of Wang et al. to a higher redshift, to link galaxy stellar mass directly with its hosting dark matter halo mass at redshift of around $0.8$. The $M_{stars}$-$M_{infall}$ relation of the galaxy stellar mass $M_{stars}$ and the host halo mass $M_{infall}$ is constrained by fitting both the stellar mass function and the correlation functions at different stellar mass intervals of the VVDS observation, where $M_{infall}$ is the mass of the hosting halo at the time when the galaxy was last the central galaxy. We find that for low mass haloes, their residing central galaxies are less massive at high redshift than those at low redshift. For high mass haloes, central galaxies in these haloes at high redshift are a bit more massive than the galaxies at low redshift. Satellite galaxies are less massive at earlier times, for any given mass of hosting haloes. Fitting both the SDSS and VVDS observations simultaneously, we also propose a unified model of the $M_{stars}$-$M_{infall}$ relation, which describes the evolution of central galaxy mass as a function of time. The stellar mass of a satellite galaxy is determined by the same $M_{stars}$-$M_{infall}$ relation of central galaxies at the time when the galaxy is accreted and becomes a sub-component of a larger group. With these models, we study the amount of galaxy stellar mass increased from z$\\sim 0.8$ to the present day through galaxy mergers and star formation. Low mass galaxies ($<3\\times 10^{10}h^{-1}M_{\\odot}$) gain their stellar masses from z$\\sim 0.8$ to $z=0$ mainly through star formation. For galaxies of higher mass, we find that the increase of stellar mass solely through mergers from $z=0.8$ can make the massive galaxies a factor $\\sim 2$ larger than observed at $z=0$, unless the satellite stellar mass is scattered to intra-cluster stars by gravitational tidal stripping or to the extended halo around the central galaxy that is not counted in the local observation. We can also predict stellar mass functions of redshifts up to $z\\sim 3$, and the results are consistent with the latest observations. Future more precise observational data will allow us to better constrain our model. ",
        "introduction": "\\label{sec:intro}  To study how galaxies form and evolve in their hosting dark matter haloes, a lot of efforts have been made to link galaxy properties with the properties of dark matter haloes which they reside in.  The usual methods used include galaxy kinematics\\citep{erickson1987} and galaxy lensing\\citep{mandelbaum2005,mandelbaum2006}, which measure the mass of hosting dark matter haloes directly. Semi-analytic models trace the gas cooling, star formation, and feedback processes `ab initio' to get the properties of galaxies of present day\\citep{lucia2006,bower2006}. Halo occupation distribution models study the galaxy-halo connection empirically, to model galaxy properties using certain assumed formula to describe the galaxy-halo relation\\citep{jing1998,berlind2002,yang2003,vale2004,conroy2006,wang2006}.  For the models that describe the formation and evolution history of galaxies, the statistics that are commonly used to constrain these models include: number statistics such as number density, luminosity function and stellar mass function\\citep{bullock2002,zehavi2005,moster2009}, spatial clustering properties described by correlation functions\\citep{jing1998, yang2003, zehavi2005}, void probability distribution\\citep{vale2004}, pairwise velocity dispersion\\citep{jing1998}, the Tully-Fisher relation\\citep{yang2003}, the colour distribution of galaxies and the clustering dependence on galaxy colour\\citep{bosch2003, kravtsov2004, zheng2004,wang2007}.   The current models of galaxy-halo connection mainly focus on low redshift study, particularly of the present day, simply because we can handle well the observational statistics only at low redshifts.  With the development of large scale galaxy redshift surveys, observations are obtained not only for the local Universe, but also toward higher redshifts. Surveys aiming at studying the properties of high redshift galaxies include DEEP2 survey\\citep{davis2003}, the COMBO-17 survey\\citep{wolf2004}, the VIMOS-VLT Deep Survey(VVDS)\\citep{lefevre2005}, the Cosmic Evolution Survey (COSMOS)\\citep{scoville2007} and zCOSMOS survey\\citep{lilly2009}. With the help of these surveys, luminosity functions of different types of galaxies are obtained up to redshift $z>7$\\citep{reddy2008,bouwens2008}. Correlation functions in luminosity bins reaches redshift of $z\\sim 1$\\citep{coil2006,pollo2006}.  Stellar mass functions have been detected for galaxies up to redshift of $z\\sim 5$\\citep{drory2005, fontana2006, elsner2008}. Correlation functions in bins of stellar mass have been studied for galaxies of redshift $z\\sim 1$\\citep{meneux2008,meneux2009} for VVDS and zCOSMOS observations.   Based on the observational data obtained at high redshifts, several works have been done to model the properties of galaxies at early epoch.  Some works use the HOD models to study high z galaxy properties\\citep{cooray2005}, as well as galaxy clustering properties\\citep{bullock2002,yan2003,cooray2006,conroy2006,white2007}, which focus mainly on the clustering dependence of galaxy luminosity. Most recently, \\citet{zheng2007} uses the HOD method to model the clustering of DEEP2 galaxies as a function of luminosity, which reaches redshift of $z\\sim 1$.  While luminosity is the most studied and easily got property of a galaxy, stellar mass is nevertheless a more fundamental property, since luminosity may be affected a lot by dust attenuation. \\citet{moster2009} uses a statistical approach to determine the relation between galaxy stellar mass and its hosting dark matter halo, and constrains the evolution of galaxy stellar mass by fitting the stellar mass functions at different redshifts taken from \\citet{drory2004}.  In the previous work of \\citet{wang2006}, an empirical method has been used to link galaxy stellar mass directly with its hosting dark matter halo mass. The method falls in between the semi-analytic approach and the halo occupation distribution approach. Positions of galaxies are predicted by following the merging histories of halos and the trajectories of subhaloes in the Millennium Simulation\\citep{springel2005}. The stellar mass of galaxies at redshift $0$ is related to the quantity $M_{infall}$ by a double power law function. $M_{infall}$ is defined as the mass of the halo at the time when the galaxy was last the central dominant object. Parameters describing the function are constrained by fitting both the stellar mass function and the correlation functions at different stellar mass intervals from SDSS observation\\citep{li2006a}. The derived $M_{stars}$-$M_{infall}$ relation is in excellent agreement with the determination from galaxy-galaxy weak lensing measurement of \\citet{mandelbaum2006}. In this study, we will apply this method to an earlier epoch. By fitting the statistic results of VVDS observation, we can study the connection between galaxy mass and its hosting halo mass at redshift of around $0.8$. Based on the relations obtained both of today and of higher redshift, we can study the evolution of galaxy stellar mass from $z\\sim 0.8$ to present day.  After a galaxy falls into a larger group and becomes a satellite, its surrounding gas is shock-heated and star formation ceases in a short timescale. The stellar mass of satellite galaxies should remain about the same as the time when they are accreted. In this case, the stellar mass of a satellite galaxy is determined by the $M_{stars}$-$M_{infall}$ relation of central galaxies at the time of infall. This inspires us to describe the $M_{stars}$-$M_{infall}$ relation for all galaxies at any redshift in a uniform way, by modelling the evolution of $M_{stars}$-$M_{infall}$ relation of central galaxies. Assuming that satellite stellar mass does not change after infall, the relation for satellite galaxies follows the relation of central galaxies at an earlier epoch when it is accreted. We will explore this unified model in \\S 4. With this model, we can also test if a significant amount of satellite disruption by tidal forces is required by current observations.  This paper is organised as follows: in Sec.~\\ref{sec:model}, we present the model for fitting VVDS observations to get the relation between galaxy stellar mass and the hosting halo mass at $z\\sim 0.8$. Based on the models describing galaxy stellar masses both at low and high redshifts, we analyse in Sec.~\\ref{sec:massincrease} how galaxies gain their masses from redshift of $0.8$ to today. In sec.~\\ref{sec:uniform} we build a unified model to fit observations of both low and high redshifts simultaneously, assuming that satellite stellar mass is determined by the $M_{stars}$-$M_{infall}$ relation of central galaxies at the time of its accretion, and study the mass growth of galaxies from $z=0.8$ based on this model. In sec.~\\ref{sec:SMF} we predict the stellar mass functions of higher redshifts based on our two best-fit models, and compare our results with recent observations.  This work is based on the Millennium Simulation\\citep{springel2005}. ",
        "conclusions": "\\label{sec:conclusion}   We apply the empirical model of \\citet{wang2006} to higher redshift, to link galaxy stellar mass with its hosting dark matter halo mass at redshift of around $0.8$. The $M_{stars}$-$M_{infall}$ relation is constrained by fitting both the stellar mass function and the correlation functions at different stellar mass intervals from VVDS observation. We find that the difference of $M_{stars}$-$M_{infall}$ relation between central and satellite galaxies at $z\\sim 0.8$ are much larger than the galaxies at present day. At all mass scales of haloes, satellite galaxies at $z= 0.8$ have much lower stellar masses than satellites within the same mass haloes at $z=0$, while the mass of central galaxies changes less from high redshift to today. Central galaxies are less massive in low mass haloes, and more massive in high mass haloes at $z=0.8$ than the galaxies at $z=0$.  Under the assumption that the mass of satellite galaxy remains about the same as it is a central before it infalls to a larger group, we build a unified model for the $M_{stars}$-$M_{infall}$ relation, which describes the evolution of galaxy mass as a function of halo mass at any given redshift. Satellite stellar mass is determined by the $M_{stars}$-$M_{infall}$ relation of central galaxy at the time of its infall. The best-fit model of both SDSS and VVDS stellar mass functions and clustering functions gives an obvious evolution of $M_{stars}$-$M_{infall}$ relation from $z=0.8$ to $z= 0$, with the mass of galaxies at higher redshift being lower than the galaxy mass at the present day, for all masses of hosting haloes. The $M_{stars}$-$M_{infall}$ relation provided by the unified model is different as the relations in the models when SDSS and VVDS data are fitted separately, especially for galaxies with massive haloes and at high redshift. The central galaxies are less massive and satellite galaxies are more massive in the unified model than the separate model. Different sets of parameters can both give reasonable fits to the observed data, because the statistics of high mass galaxies are not tightly constrained in observation at high redshift.  We study the amount of galaxy stellar mass growth from $z\\sim0.8$ to today, in either way of galaxy merger or star formation. For the models when SDSS and VVDS data are fitted separately, we find that for galaxies that reside in haloes of mass less than $10^{12}h^{-1}M_{\\odot}$, the masses of their most massive progenitors at $z=0.83$ vary from about 20 percent to 60 percent of the present day mass. Meanwhile, galaxy mergers contribute only a small fraction to the galaxy mass of today. This indicates that a large fraction of $z=0$ stellar masses comes from star formation during the period between $z=0.83$ and $z=0$. For galaxies within massive haloes, the total amount of stellar mass from the main progenitor at $z=0.83$ and from that increased from mergers with other galaxies actually exceeds the present-day galaxy mass. Although there could be a difference of the definitions for stellar mass of central galaxies in observation and in the models based on simulation (for example, the stars in the envelope of the central galaxies may not be counted in observation), this indicates that there may be little star formation ongoing in these massive galaxies. The unified model basically tells the same story, except that the contribution of the most massive progenitors at $z=0.83$ to the mass of the present day is no more than 75 percent even for the most massive galaxies.  Based on the models we build, we give predictions of higher redshift stellar mass functions, with redshift up to $z\\sim 3$. It is encouraging that our predictions are consistent with recent observations of stellar mass functions presented by \\citet{marchesini2008} and \\citet{kajisawa2009}. However, the amplitude of the stellar mass functions predicted by the unified model is always higher at low mass end, and lower at high mass end, compared with the prediction of the model that fits SDSS and VVDS data separately. Besides, in both of our models, we assume that for the $M_{stars}$-$M_{infall}$ relation, its slopes at high and low mass end and scatter around the median value do not change with time, which may not actually be true. The current models can be tested by future observational data, and a more accurate model can be constructed with better data.   Besides the observational data of VVDS survey used to constrain our model in this work, galaxy stellar mass function and correlation functions in different stellar mass bins are also obtained for zCOSMOS survey by \\citet{pozzetti2009} and \\citet{meneux2009}. As pointed out in \\citet{meneux2009}, the zCOSMOS field is centred on an extreme overdensity region. We therefore choose to use the VVDS results in this modelling work, although it is possible that the VVDS field is a bit under-dense on the other hand. Hopefully this cosmic variance effect can be significantly reduced with upcoming larger samples of galaxies, based on which the models can be constrained much more tightly.  Recent study of \\citet{moster2009} determined the galaxy stellar mass - halo mass connection at high redshift constrained with stellar mass functions only. They claim that stellar mass function alone is enough to constrain the relation between galaxy stellar mass and their hosting halo mass. We have shown at the end of Sec.~\\ref{sec:model} that this may not be true for high redshift situation, although this was proved viable when considering models at the local Universe. However, their result shows the same trend as our $M_{stars}$-$M_{infall}$ relation in the separate model, although the quantitatively value of the relation is somewhat different. For low mass haloes, galaxies at higher redshifts have less stellar mass than galaxies at a lower redshift. For high mass haloes, galaxies at high redshift is a bit more massive than galaxies at low redshift. In our unified model, on the other hand, galaxy mass is always smaller at high redshift than that at low redshift, but the difference for galaxies within massive haloes is quite small. The precision of all these models can be tested only by improved observational results both on stellar mass functions and on the clustering properties at high redshifts."
    },
    "0911/0911.1345.txt": {
        "abstract": "{cosmic strings, superstrings, early universe, branes, braneworlds} Cosmic strings are predicted by many field-theory models, and may have been formed at a symmetry-breaking transition early in the history of the universe, such as that associated with grand unification.  They could have important cosmological effects.  Scenarios suggested by fundamental string theory or M-theory, in particular the popular idea of brane inflation, also strongly suggest the appearance of similar structures.  Here we review the reasons for postulating the existence of cosmic strings or superstrings, the various possible ways in which they might be detected observationally, and the special features that might discriminate between ordinary cosmic strings and superstrings. ",
        "introduction": "\\label{intro}  There are many examples of linear topological defects in low-temperature con\\-densed-matter systems, including vortices in superfluid helium and in atomic Bose-Einstein condensates, and magnetic flux tubes in superconductors.  These structures typically appear when a system goes through a phase transition into a low-temperature ordered phase in which the underlying symmetry is spontaneously broken.  They arise because in certain circumstances the ordering may be frustrated.  Cosmic strings are analogous objects that may have been formed in the early universe.  They could be produced during one of the early symmetry-breaking phase transitions predicted by many particle-physics models, for example one associated with the breaking of a `grand unification' symmetry.  Similar point-like or planar defects --- monopoles or domain walls --- may also be formed.  Superstrings are the supposed basic constituents of matter in fundamental string theory or M-theory, which is at present the leading contender as a unified theory of all particle interactions including gravity.  Both cosmic strings and superstrings are still purely hypothetical objects.  There is no direct empirical evidence for their existence, though there have been some intriguing observations that were initially thought to provide such evidence, but are now generally believed to have been false alarms.  Nevertheless, there are good theoretical reasons for believing that these exotic objects do exist, and reasonable prospects of detecting their existence within the next few years.  Initially, cosmic strings and superstrings were regarded as two completely separate classes.  Cosmic strings, if they exist, stretch across cosmological distances and though exceedingly thin are sufficiently massive to have noticeable gravitational effects.  On the other hand superstrings were seen as minute objects, even on the scale of particle physics, far too small to have any directly observable effects.  But in the last few years it has emerged that under certain circumstances they too can grow to macroscopic size and play the same role as cosmic strings.  Indeed, it is now very possible that observations of a cosmic superstring might provide the first real direct evidence for M-theory.  (For reviews see  Polchinski 2005, Davis \\& Kibble 2005, Sakellariadou 2009, Myers \\& Wyman 2009.)  We begin, in Section \\ref{cosmicstr}, by describing the simplest examples of cosmic strings, which arise in the breaking of an abelian U(1) symmetry, and discuss the reasons for thinking that cosmic strings might appear during phase transitions in the very early history of the universe, in a process closely analogous to what happens when a superfluid or superconductor is cooled rapidly through its transition temperature.  In Section \\ref{dynamics} we review the dynamics of cosmic strings, and in Section \\ref{evolution}, discuss how a network of cosmic strings, once formed, would evolve as the universe expands.  In Section \\ref{sustr}, we review the ideas behind fundamental superstring theory and the characteristics of the five known consistent superstring theories.  In all five cases, consistency requires a spacetime of ten rather than four dimensions.  A mechanism is therefore required to reduce the number of dimensions of our spacetime to the familiar four.  One such is the Kaluza--Klein mechanism of \\emph{compactification}: the unobserved dimensions are in fact curled up so small as to be unobservable on ordinary scales.  Another idea, the \\emph{brane-world} mechanism, has emerged more recently.  A key feature of superstring theories is the appearance of \\emph{branes}, surfaces of various dimensions embedded in the larger spacetime.  A $p$-brane is a surface of spatial dimension $p$, or spacetime dimension $p+1$.  The basic idea of the brane world is that all the constituents of our universe are bound to a 3-brane that can move within the larger nine-dimensional space; they in fact consist of short segments of fundamental string with their ends firmly rooted on the brane.  In Sections \\ref{tension}-\\ref{stability}, we explain why the initial rejection of the idea of cosmic superstrings has been overturned by showing how they can have sufficiently small tensions, survive a period of inflation and be cosmologically stable, and use a series of models to demonstrate the point, using them to also demonstrate the similarities and differences  between these cosmic superstrings and ordinary cosmic strings (often called \\emph{solitonic} strings).  One particularly important difference, discussed in Section \\ref{reconnect}, lies in the probability of `intercommuting' or exchange of partners when strings cross, and also in the possibility of forming three-string junctions. The observational signatures of the different types of strings are the subject of Section \\ref{observation}, where we also discuss the current observational bounds on their existence.  Our conclusions are summarized in Section \\ref{conc}. ",
        "conclusions": "\\label{conc} Many early-universe scenarios predict the appearance of cosmic strings or superstrings.  The most promising perhaps is brane-world inflation.  Models in which inflation is driven by a collision between  D3- and $\\OL{\\mbox{D3}}$-branes generically predict the formation of strings of various types, F- and D-strings and composite $(p.q)$-strings.  The details of which strings are stable or metastable are complicated but the conditions are not overly restrictive. Even in models where inflation does not involve branes, cosmic superstrings may be formed by other mechanisms.  Moreover, ordinary cosmic strings may appear later in the evolutionary process.  There are many possible ways by which cosmic strings or superstrings may be detected.  Their direct lensing effect may be observable if $G\\mu$ is reasonably close to the present observational upper bound.  Loops offer the possibility of being seen if their tension is low enough, because they last longer, hence their number density will increase. This is particularly the case if they were to cluster near our galaxy (Chernoff\\ \\& Tye 2007). In fact these authors claim that the planned astronomical survey mission GAIA could observe strings with tensions as low as $G\\mu \\geq 10^{-10}$ because the loops generate microlensing events that will be within GAIA's reach. Small-scale observations of the CMB using the Atacama Cosmology Telescope (ACT) and the South Pole Telescope (SPT) may reveal the characteristic discontinuities created by cosmic strings for tensions satisfying $G\\mu \\geq 10^{-7}$ (Fraisse \\ea\\ 2008). The fact that strings can directly source vector mode perturbations at last scattering implies that even light strings can generate strong B-mode polarisation in the CMB, providing a signal which is distinguishable from the type of spectrum expected from inflation (Pogosian \\& Wyman 2008, Bevis \\ea\\ 2008). Current experiments searching for  characteristic patterns in the B-mode polarisation, include BICEP, PolarBear, and Spider and these have just been joined by Planck in the search. Gravitational radiation emitted from cusps, as well as from kinks or junctions, may be detectable by future experiments, such as LISA and remain probably the most sensitive test for cosmic strings and cosmic superstrings, as they could detect loops with tensions as low as $G\\mu \\geq 10^{-13}$ (Damour\\ \\& Vilenkin 2005).  Meanwhile, the introduction of Pulsar Timing Arrays will provide ever more constraining limits on the allowed string tensions as will direct gravity wave searches like LIGO2 or LISA.  Success in any of these observations would trigger huge interest, and allow us to plan further observations to probe the detailed nature of the strings, in particular to distinguish ordinary cosmic strings from superstrings.  A positive result would yield a lot of information about the nature of the underlying fundamental theory, information that would be very hard to gather in any other way."
    },
    "0911/0911.5056_arXiv.txt": {
        "abstract": "The \\emph{Planck} On-Fligh Forecaster (POFF) is a tool to predict when a position in the sky will be within a selected angular distance from any receiver direction of the \\emph{Planck} satellite according to its pre-programmed observational strategy. This tool has been developed in the framework of the \\emph{Planck} LFI Core Team activities, but it is now used by the whole collaboration.  In this paper we will describe the tool and its applications to plan observations with other instruments of point sources which are expected to enhance the possibilities of scientific exploitation of the \\emph{Planck} satellite data, once they will be publicly available. Collecting simultaneous multi-frequency data, like those that can be planned with the POFF, will help, on one hand, to investigate variability of point sources and, on the other, to reconstruct point source spectral energy distributions on wide frequency ranges minimizing the effects due to source variability.  POFF is a combination of IDL routines which combine the publicly available information about the \\emph{Planck} scanning strategy and focal plane shape in order to identify if a given (list of) position(s) can be observable by the satellite at a given frequency and/or by selected receivers in a given time range. The output can be displayed with the desired time resolution and selecting among various sorting options.  The code is not a Planck product, but it has been validated within the \\emph{Planck} LFI pipeline, looking for sources in the first satellite datasets. It will be implemented among the general tools of the LFI Data Processing Center. The code format and the large number of options make it flexible and suitable for many applications, allowing to get results quickly.  POFF is currently successfully used to plan activities within the \\emph{Planck} collaboration, including observations with several ground-based facilities, and it is distributed outside it. ",
        "introduction": "The ESA's \\emph{Planck}\\footnote{Planck \\emph{(http://www.esa.int/Planck)} is a project of the European Space Agency - ESA - with instruments provided by two scientific Consortia funded by ESA member states (in particular the lead countries: France and Italy) with contributions from NASA (USA), and telescope reflectors provided in a collaboration between ESA and a scientific Consortium led and funded by Denmark.} satellite, launched on May the 14th, in addition to improving anisotropy measurements of the Cosmic Microwave Background (CMB), is surveying the sky in nine frequency bands (33, 40, 70 GHz observed by the array of radiometers of the Low Frequency Instrument (LFI, Mandolesi et al. 1998 and 2009), and 100,143, 217, 353, 545, 857 GHz for the bolometric array of the High Frequency Instrument, (HFI, Puget et al. 1998, Lamarre et al. 2009)  with FWHM ranging from 5 to 33 arcmin). It will provide the first all-sky survey above 100 GHz (The Planck Collaboration 2006). It is expected to detect thousands of Galactic and extragalactic sources. Most of the extragalactic sources will be dusty galaxies, detected at high frequencies, but, given its sensitivity, it is expected to detect up to several hundred extragalactic radio sources (L\\'{o}pez-Caniego et al. 2006, Leach et al. 2008, Massardi et al. 2009) in one or more of its lower frequency bands (30, 44, 71, 100, 143 GHz) and to measure polarisation (L\\'{o}pez-Caniego et al. 2009) up to 353 GHz: this means that it will observe from 2 to 3 times more radio sources than detected in WMAP\\footnote{http://lambda.gsfc.nasa.gov/product/map/current/} maps at similar frequencies (23-94 GHz, Wright et al.\\ 2008, Massardi et al. 2009). A large fraction of these sources may also be detected up to 353 GHz.  Some of the radio sources will be very bright objects (prominent quasars and other flat spectrum Galactic or extragalactic sources) and will have been frequently observed in the past and well characterized. Others will be rare inverted spectrum sources with $\\alpha > 0$ ($S\\propto \\nu^\\alpha$), or other unusual sources not systematically studied or even present in low frequency catalogs. Still, others may be flaring sources with temporarily strong high-frequency emission (see Burigana 2000, Terenzi et al. 2002, 2004 for preliminary studies of \\emph{Planck} capabilities to address source variability issues).  To date, the AT20G (Murphy et al. 2009, Massardi et al. 2008, 2009) provides the best ground-based sample of the high frequency sky. Given its high resolution, it mainly collects extragalactic objects. \\emph{Planck} will improve on this, by providing a blind survey at even higher frequencies and at several epochs. We expect that \\emph{Planck}'s all sky, high frequency and multi-frequency survey will provide a less biased sample of blazars and other extreme extragalactic sources than past ground-based works, and will consequently allow us to check features of the jet-shock paradigm introduced by most models for AGN emission. Studying the spectral energy distribution (SED) and the variability of a large sample of AGNs, including various subclasses, is a way to probe the physics of the innermost regions of the sources. With a substantial, unbiased sample, unification models of AGN (Urry \\& Padovani 1995) can be tested. \\emph{Planck} data together with lower frequency information may help to differentiate between unifying schemes based on Doppler and obscuration mechanism, exploiting the \\emph{Planck} High Frequency Instrument to characterize the IR emission from galaxies and quasars. Poisson fluctuations given by undetected extragalactic radio sources (i.e. sources at fluxes smaller or of the order of $200-300$ mJy) should dominate the CMB angular power spectrum at $\\ell$ greater than or of the order of 1500, at \\emph{Planck} LFI frequencies (Toffolatti et al. 1998, De Zotti et al. 1999, Toffolatti et al. 2005). As a consequence, it will be crucial to reconstruct the emission properties of bright extragalactic radio sources.  During outbursts, warious classes of variable Galactic radio sources show flux density levels high enough to be observed with a good S/N by \\emph{Planck}. As a remarkable example, the massive X-ray binary system Cyg X-3 is one of most active and variable Galactic radio source with quiescent periods and strong outbursts interpreted as synchrotron emission in a jet-like structure, reaching up to 20 Jy at 8.4~GHz. More than 200 Galactic X-ray binaries are known and $10 \\%$  of these are radio loud (Mirabel and Rodrigues 1999). Furthermore, Luminous Blue Variable stars (van Genderen 2001), like Eta Carinae, are characterized by variability and by sudden outbursts, during which a large amount of mass is ejected from the star and the flux density can reach several Jy at centimeter wavelengths and tens of Jy in the millimeter. Finally, active binary stars, like RS CVn and Algol, have alternate periods of quiescence, with flux densities of few tens of mJy, and active periods characterized by flares, lasting several weeks, every 2-3 months, reaching several hundreds of mJy up to Jy at cm wavelengths (Umana et al. 1995).  Radio spectra show the maximum of emission at frequencies higher than 10~GHz and, during the impulsive phases, up to about 100 GHz.  In order to fully exploit and enhance the scientific impact of \\emph{Planck} data, a set of supporting possible coeval observations of the same sources, with particular attention to those that show unusual spectra or flaring events, across as much of the electromagnetic spectrum as possible are being carried out with many facilities all around the world in the framework of the \\emph{Planck} non-CMB science-driven activities.  \\emph{Planck} observations may help identifying the properties of the SED in the high frequency radio and FIR bands, and will benefit from the ground-based observations at longer wavelength for the reconstruction of the properties of different emission components.  After the satellite launch, ESA has publicly released the information about the Planned Pointing Lists (PPL) which enclose the expected sequence of pointing positions of the satellite spin axis and the Spacecraft/Instrument Alignment Matrix (SIAM) which describes the angles of each receiver pointing direction with respect to the spin axis.  By combining these pieces of information, the POFF (\\emph{Planck} On-Flight Forecaster\\footnote{This tool is not an official Planck product, but it is produced on a best-effort basis by individual members of the Planck collaboration.}) code allows us to predict when a given (set of) position(s) in the sky will be within a given angular distance from any satellite receiver pointing direction.  Hence, the code can select, among a list of given positions, those that are expected to be observed within a given period by the satellite. These capabilities make it extremely useful for several purposes, and is already frequently used within the \\emph{Planck} community to plan ground-based activities (and most of all observations) almost simultaneously with the satellite.  Its main use is to plan 'coeval observations' for selected samples of objects. With the term 'coeval' we mean that the positions scanned through by the satellite are observed with other instruments within a small enough period so that, if the position is associated to a source, the observation is not affected by the intrinsic variability of the source. In fact, the allowed time delay between two observations to be still considered \"simultaneous\" depends on the physical properties of the observed source. It is also a powerful support to the activities of LFI Data Processing Center aimed at identifying point sources in the actual satellite time ordered data, providing constrains on the observing time for given positions.  In the same spirit that led ESA to make available the PPL and SIAM information, POFF is now publicly available to plan almost simultaneous observational activities to be compared with the results of the satellite once they will be publicly released. A preliminary all-sky catalogue of compact and point sources, extracted from \\emph{Planck}'s data, will be released to the scientific community approximately 10 months after the first full sky survey is finished (i.e. February 2010), in time for the first post-launch call for Herschel observing proposals.  A tool analogous to the POFF has been developed in the framework of the \\emph{Planck} collaboration to track the transits of Solar System moving objects on \\emph{Planck} receivers ad to plan coeval observations of those bodies detectable by the satellite (Maris and Burigana 2009).  In this paper we will describe the fundamental aspects of the \\emph{Planck} scanning strategy (\\S \\ref{sec:Planck_ss}) and focal plane (\\S \\ref{sec:Planck_fp}), the condition to define when a position will be observed by a receiver (\\S \\ref{sec:Planck_obscond}), the structure of the POFF code (\\S \\ref{sec:code}) and its performances (\\S \\ref{sec:performances}), and some representative applications (\\S \\ref{sec:application}). Finally, in \\S \\ref{sec:conclusions} we summarize our results and illustrates the main perspective of POFF applications. In the appendices we collected a brief user manual for each input keyword and few examples of application. ",
        "conclusions": "\\label{sec:conclusions} The \\emph{Planck} On-Fligh Forecaster (POFF) is a tool to predict when a position in the sky will be within one (by default) beam unit from any selected receiver of the \\emph{Planck} satellite according to the preprogrammed \\emph{Planck} scanning strategy. This tool has been developed in the framework of the \\emph{Planck} LFI Core Team activities and is going to be implemented among the LFI Data Processing Center general tools. The code provides the dates in which the satellite will observe each position given in an input list for each frequency channel and/or satellite receiver, so that it is possible to plan a priori the observations of known objects.  The POFF code has demonstrated to be a precious tool in planning ground-based observations simultaneous with the satellite. Many activities are on-going exploiting several facilities of the two hemispheres, in order to get the advantage of the pointing knowledge on one hand to minimize the effects of source variability, and to investigate the source variability on the other, once the ground-based observational results will be compared with \\emph{Planck} data.  The code is available in IDL pseudo-executable binary format on the dedicated webpage http://web.oapd.inaf.it/rstools/POFF\\footnote{Under construction}."
    },
    "0911/0911.2489_arXiv.txt": {
        "abstract": "Recent detection of blazar 3C279 by MAGIC has confirmed previous indications by H.E.S.S. that the Universe is more transparent to very-high-energy gamma rays than currently thought. This circumstance can be reconciled with observations of nearby blazars provided that photon oscillations into a very light Axion-Like Particle occur in extragalactic magnetic fields. The emerging ``DARMA scenario'' can be tested in the near future by the satellite-borne {\\it Fermi} LAT detector as well as by the ground-based Imaging Atmospheric Cherenkov Telescopes H.E.S.S., MAGIC, CANGAROO III, VERITAS and by the Extensive Air Shower arrays ARGO-YBJ and MILAGRO. \\vspace{1pc} ",
        "introduction": "Imaging Atmospheric Cherenkov Telescopes (IACTs) are providing us with an impressive amount of information about the Universe in the energy interval $100 \\, {\\rm GeV} - 100 \\,{\\rm TeV}$. Observations carried out by these IACTs concern gamma-ray sources over an extremely wide interval of distances, ranging from the parsec scale for Galactic objects up to the Gigaparsec scale for the fartest detected blazar 3C279. This circumstance allows not only to infer the intrinsec properties of the sources, but also to probe the nature of photon propagation throughout cosmic distances.  The latter fact is of paramount importance for very-high-energy (VHE) gamma-ray astrophysics, since the horizon of the observable Universe rapidly shrinks above $100 \\, {\\rm GeV}$ as the energy further increases. This is due to the fact that photons from distant sources scatter off background photons permeating the Universe, thereby disappearing into electron-positron pairs~\\cite{stecker1971}. It turns out that the corresponding cross section $\\sigma (\\gamma \\gamma \\to e^+ e^-)$ peaks where the VHE photon energy $E$ and the background photon energy $\\epsilon$ are related by $\\epsilon \\simeq (500 \\, {\\rm GeV}/E) \\, {\\rm eV}$. As far as observations performed by IACTs  are concerned, the cosmic opacity is dominated by the interaction with ultraviolet/optical/infrared diffuse background photons~\\footnote{Frequency band $1.2 \\cdot 10^{3} \\, {\\rm GHz} - 1.2 \\cdot 10^{6} \\, {\\rm GHz}$, corresponding to the wavelength range $0.25 \\, \\mu {\\rm m} - 250 \\, \\mu {\\rm m}$.}, usually called Extragalactic Background Light (EBL), which is produced by galaxies during the whole history of the Universe. Owing to such an absorption process, photon propagation is controlled by the optical depth ${\\tau}(E,D)$, with $D$ denoting the source distance. Hence, the observed photon flux $\\Phi_{\\rm obs}(E,D)$ is related to the emitted one $\\Phi_{\\rm em}(E)$ by \\begin{equation} \\label{a0} \\Phi_{\\rm obs}(E,D) = e^{- \\tau(E,D)} \\, \\Phi_{\\rm em}(E)~. \\end{equation} Neglecting evolutionary effects on the EBL spectral energy distribution for simplicity, the optical depth reads ${\\tau}(E,D) \\simeq D/{\\lambda}_{\\gamma}(E)$, where ${\\lambda}_{\\gamma}(E)$ is the photon mean free path for $\\gamma \\gamma \\to e^+ e^-$ referring to the present cosmic epoch. As a consequence, Eq.~(\\ref{a0}) simplifies as \\begin{equation} \\label{a1} \\Phi_{\\rm obs}(E,D) \\simeq e^{- D/{\\lambda}_{\\gamma}(E)} \\ \\Phi_{\\rm em}(E)~. \\end{equation} The function ${\\lambda}_{\\gamma}(E)$ decreases like a power law from the Hubble radius $4.3 \\, {\\rm Gpc}$ around $100 \\, {\\rm GeV}$ to $1 \\, {\\rm Mpc}$ around $100 \\, {\\rm TeV}$~\\cite{CoppiAharonian}. Now, Eq.~(\\ref{a1}) entails that the observed flux is {\\it exponentially} suppressed both at high energy and at large distances, so that sufficiently far-away sources become hardly visible in the VHE range and their observed spectrum should anyway be {\\it much steeper} than the emitted one.  Yet, observations carried out by IACTs have failed to detect such a behaviour. A first indication in this respect was reported by the H.E.S.S. collaboration in connection with the discovery of the two blazars H2356-309 ($z = 0.165$) and 1ES1101-232 ($z = 0.186$) at $E \\sim 1 \\, {\\rm TeV}$~\\cite{aharonian:nature06}. Stronger evidence comes from the observation of the blazar 3C279 ($z = 0.538$) at $E \\sim 0.5 \\, {\\rm TeV}$ by the MAGIC collaboration~\\cite{3c}. In particular, the signal from 3C279 collected by MAGIC in the region $E<220$ GeV has more or less the same statistical significance as the one in the range 220 GeV $< E <$ 600 GeV ($6.1 \\sigma$ in the former case, $5.1 \\sigma$ in the latter).  Turning the argument around and assuming {\\it standard} photon propagation as described above, the observed spectrum $\\Phi_{\\rm obs}(E,D)$ can only be reproduced by an emission spectrum $\\Phi_{\\rm em}(E)$ {\\it much harder} than for any other blazar observed so far.  A way out of this difficulty relies upon a modification of the emission spectrum $\\Phi_{\\rm em}(E)$. A possibility involves the presence of strong relativistic shocks, which can substantially harden $\\Phi_{\\rm em}(E)$~\\cite{Stecker2007+2008}. A different option invokes photon absorption inside the blazar, which has been shown to produce again an emission spectrum $\\Phi_{\\rm em}(E)$ considerably harder than previously thought~\\cite{Aharonian2008}. While successful at increasing the fraction of VHE emitted photons, these attempts fail to explain why {\\it only} for the most distant blazars do these mechanisms become important.  A very different solution was recently proposed by the present authors and is usually referred to as the ``DARMA scenario''~\\cite{drm}. Its characteristic feature is the presence of Axion-Like Particles (ALPs) (more about this, later) and rests upon the mechanism of photon-ALP oscillation in cosmic magnetic fields, whose existence has definitely  been proved by AUGER observations~\\cite{auger}. Once ALPs are produced close enough to the source, they travel {\\it unimpeded} throughout the Universe -- since they do not undergo EBL absorption -- and can convert back to photons before reaching the Earth. As a consequence, the {\\it effective} photon mean free path ${\\lambda}_{\\gamma , {\\rm eff}} (E,D)$ gets {\\it increased} so that the observed photons cross a distance in excess of ${\\lambda}_{\\gamma}(E)$. Moreover, it has been shown that the DARMA scenario works for an ALP lighter than about $10^{- 10} \\, {\\rm eV}$~\\footnote{Somewhat similar ideas are discussed in~\\cite{sim}.}.  A deeper insight into the DARMA mechanism can be achieved by introducing the probability $P_{\\gamma \\to \\gamma}(E,D)$ that a photon remains a photon after propagation over a distance $D$, so that we have \\begin{equation} \\label{a1bisa} \\Phi_{\\rm obs}(E,D) = P_{\\gamma \\to \\gamma}(E,D) \\ \\Phi_{\\rm em}(E)~. \\end{equation} When only photon absorption is operative, Eq.~(\\ref{a1}) can similarly be rewritten as \\begin{equation} \\label{a1bisaQ} \\Phi_{\\rm obs}(E,D) = P_{\\gamma \\to \\gamma}^{(0)}(E,D) \\ \\Phi_{\\rm em}(E)~, \\end{equation} with \\begin{equation} \\label{a1trisaW} P_{\\gamma \\to \\gamma}^{(0)}(E,D) \\simeq e^{- D/{\\lambda}_{\\gamma}(E)}~. \\end{equation} In the presence of photon-ALP oscillations, Eq.~(\\ref{a1trisaW}) gets replaced by \\begin{equation} \\label{a1trisb} P_{\\gamma \\to \\gamma}(E,D) \\simeq e^{- D/{\\lambda}_{\\gamma}(E)} \\, X(E,D) \\end{equation} and the above discussion entails $X(E,D) > 1$. Moreover, Eq.~(\\ref{a1}) presently becomes \\begin{equation} \\label{a1bis} \\Phi_{\\rm obs}(E,D) \\simeq e^{- D/{\\lambda}_{\\gamma , {\\rm eff}}(E,D)} \\ \\Phi_{\\rm em}(E)~, \\end{equation} with \\begin{equation} \\label{a1tris} {\\lambda}_{\\gamma , {\\rm eff}}(E,D) = - \\frac{D}{{\\rm ln} \\, P_{\\gamma \\to \\gamma}(E,D)}~, \\end{equation} so as to guarantee consistency with Eq.~(\\ref{a1bisa}). Next, by inserting Eq.~(\\ref{a1trisb}) into Eq.~(\\ref{a1tris}) we get \\begin{equation} \\label{a1tric} \\frac{{\\lambda}_{\\gamma , {\\rm eff}}(E,D)}{{\\lambda}_{\\gamma}(E)} \\simeq \\frac{D}{D - {\\lambda}_{\\gamma}(E) \\,  {\\rm ln} \\, X(E,D)} \\end{equation} and since $X(E,D) > 1$ we find ${\\lambda}_{\\gamma , {\\rm eff}}(E,D) > {\\lambda}_{\\gamma}(E)$, which is just a formal restatement of our previous conclusion. Still, Eq.~(\\ref{a1bis}) possesses the advantage to explicitly show that even a {\\it small} increase of ${\\lambda}_{\\gamma , {\\rm eff}}(E,D)$ gives rise to a {\\it large} enhancement of the observed flux $\\Phi_{\\rm obs} (E,D)$. As we shall see, the DARMA mechanism makes ${\\lambda}_{\\gamma , {\\rm eff}}(E,D)$ shallower than ${\\lambda}_{\\gamma}(E)$, although it remains a decreasing function of $E$. So, the resulting observed spectrum is {\\it much harder} than the one predicted by Eq.~(\\ref{a1}), thereby ensuring agreement with observations even by adopting for far-away sources the {\\it same} emission spectrum characteristic of nearby ones.  Our aim is to review the main features of the DARMA scenario as well as its application to blazar 3C279. ",
        "conclusions": ""
    },
    "0911/0911.0390_arXiv.txt": {
        "abstract": "We derive the axis ratio distribution of X-ray clusters using the XMM-Newton catalogue \\citep{2008A&A...478..615S}. By fitting the contour lines of the X-ray image by ellipses, we confirm the X-ray distribution is well approximated by the elliptic distribution with a constant axis ratio and direction. We construct a simple model describing the axis ratio of the X-ray gas assuming the hydrostatic equilibrium embedded in the triaxial dark matter halo model proposed by \\citet{2002ApJ...574..538J}  and the hydrostatic equilibrium. We find that the observed probability density function of the axis ratio is consistent with this model prediction. ",
        "introduction": "Galaxy clusters have been playing an important role in the determination of cosmological parameters such as the matter density, the amplitude of the initial fluctuation, the Hubble constant. Upcoming surveys of galaxy clusters via optical, X-ray, lensing and the Sunyaev-Zel'dovich effect will enable one to probe new aspects of cosmology such as the nature of the dark energy \\citep[e.g.][]{2006astro.ph..9591A} and the initial non-Gaussianity. Therefore, the physical modeling of galaxy clusters is important to properly interpret the cluster data. In particular, non-sphericity adds significant uncertainty to various cosmological applications of the clusters \\citep[e.g.][]{2003ApJ...599....7O,2004ApJ...609...50D}.  The statistical nature of the non-sphericity has been investigated using N-body simulations by many authors \\citep[e.g.][]{2002ApJ...574..538J,2005ApJ...629..781K,2006MNRAS.366.1503P,2006MNRAS.367.1781A,2007MNRAS.376..215B}.  Among them, \\citet[][hereafter JS02]{2002ApJ...574..538J} proposed a triaxial model of the dark matter halo based on the detailed analysis of N-body simulations. They also provided the fitting formula of the probability distribution of the axis ratio.  Their phenomenological model has been applied to many studies that investigate the systematic errors of the cosmological applications with clusters. The observational test for the axis ratio distribution, however, is quite limited. \\citet{2009ApJ...695.1446E} analyzed 4281 clusters in the Sloan Digital Sky Survey using the gravitational lensing. By stacking all the clusters, they were merely able to conclude that the axis ratio of the projected dark matter halo $f = 0.48^{+0.14}_{-0.09}$. Due to large uncertainty, however, the measurement of the distribution of the axis ratio was difficult. Future surveys will make a significant improvement of the direct measurement of dark matter haloes with new techniques \\citep[e.g.][]{2009MNRAS.400.1132H}.  On the other hand, X-ray clusters also represent the non-sphericity, which makes the large statistical error to measurement of the Hubble constant \\citep[e.g][]{2008ApJ...674...11K}. The theoretical predictions of the probability density function (PDF) of the axis ratio using X-ray clusters have been considered by several authors \\citep[e.g.][]{2004ApJ...617..847W, 2003ApJ...585..151L}. These are based on the assumption that the non-sphericity of X-ray clusters originated from the non-spherical gravitational potential due to the dark matter halo. It is crucial to see if the non-sphericity of X-ray clusters can be explained by the non-sphericity of the underlying dark matter halo.  There are a lot of papers which measured the axis ratio of the X-ray clusters observed by Einstein  \\citep[e.g.][]{1989ApJS...70..723M, 1992ApJ...400..385B, 1995ApJ...447....8M}, ROSAT \\citep[e.g.][]{1996ApJ...457..565B,1997MNRAS.292..920W,2001MNRAS.320...49K} and Chandra \\citep[e.g.][]{2007A&A...467..485H,2009ApJ...691.1648F}. \\cite{2007MNRAS.377..883F} have constructed the theoretical prediction of the gas distribution and justified their model by comparing with 8 simulated clusters. Applying this model to 46 simulated dark matter haloes, they predicted the average and variance of the axis ratio of X-ray clusters. Comparing with ROSAT samples they have found good agreement between the predicted and observed average and variance.   Because the XMM-Newton has the largest effective area among present X-ray satellites and its field of view is also large, one can investigate the details of the X-ray shape or isocontour lines of clusters including the radial dependence of the axis ratio and the direction of semi major axis for each cluster.  While the axis ratio constraints are not expected to be much better than have been obtained previously, it is useful to have an independent computation to serve as a check on systematic errors between different data sets and different computational methods. In this paper, we perform detailed comparisons of the PDF derived by the X-ray analysis using high quality data set of the XMM-Newton cluster catalogue \\citep{2008A&A...478..615S} with the simple theoretical prediction of the PDF based on the phenomenological model of shape of underlying dark matter haloes provided by JS02.  The rest of the paper is organized as follows. In \\S 2, we describe the cluster data we use and the data processing such as fitting clusters and estimating the statistical errors. In \\S 3, in order to model the non-sphericity, we investigate the radial dependence of the axis ratio and the axis direction. After describing a simple model of the gas non-sphericity, we compare the non-spherical properties of X-ray clusters with the theoretical prediction. Finally, we summarize our results in \\S 4. ",
        "conclusions": "In this paper, we have examined the axis ratio distribution of X-ray halos using the XMM-Newton cluster catalogue. By fitting the surface brightness contours by ellipses, we confirmed that the radial dependence of the axis ratio and the axis direction is relatively small, that is, the X-ray halo is well approximated by the triaxial ellipsoid. Constructing the simple model based on the hydrostatic equilibrium of the underlying triaxial dark matter halo proposed by  \\citet{2002ApJ...574..538J}, we found that the observed PDF of the axis ratio of X-ray halo agrees well with the theoretical prediction.  We have shown that the axis ratio distribution of a sample of X-ray clusters predicted by the $\\Lambda$ CDM model agrees with that measured by XMM-Newton, consistent with the findings of previous studies using different data sets, which is encouraging given the simplifying assumptions adopted in all such model comparisons, including our own."
    },
    "0911/0911.3500_arXiv.txt": {
        "abstract": "CAKE (Cosmic Abundances below Knee Energies) was a prototype balloon experiment for the determination of the charge spectra and of abundances of the primary cosmic rays with nuclear charge Z$>$10. It was a passive instrument made of layers of CR39$^{\\textregistered}$ and Lexan$^{\\textregistered}$/Makrofol$^{\\textregistered}$ nuclear track detectors; it had a geometric acceptance of $\\sim$0.7 m$^2$sr for Fe nuclei. Here, the scanning and analysis strategies, the algorithms used for the off-line filtering and for the tracking in automated mode of the primary cosmic rays are presented, together with the resulting CR charge distribution and their abundances. ",
        "introduction": "\\label{sec:intro} The determination of the elemental abundances of cosmic rays (CRs) observed in the solar system provides valuable information on the nature of their sources, acceleration mechanisms and propagation in the Interstellar Medium (ISM). While there is a general consensus for supernova shocks being the main acceleration sites of CR particles with energy per nucleon up to $10^{15}$ eV, still unresolved is the question of which are their sources.  From measurements of the abundances made near the Earth it is possible, by modelling the propagation processes in the ISM, to infer the elemental abundances at CR sources as well as to place constraints on the nature of their original environments. The present data cannot select one model between coronal material of later-type stars (which are rich of ions in a gas-phase) and Interstellar Medium in the proximity of Supernovae Remnants. Two proposed CR selection mechanisms based on the elemental First Ionization Potential (FIP) and on the volatility/refractory state, \\emph{i.e.} on the condensation temperature $T_c$, could justify  the first scenario with respect to the second. The measured CR abundances show a strong prevalence of the refractory elements (which are generally characterized by low FIP and large $T_c$, e.g. Mg ($Z=12$), Ca ($Z=20$), Fe ($Z=26$) \\cite{waddington}) with respect to the volatile ones (characterized by high FIP and small $T_c$, e.g. H ($Z=1$), He ($Z=2$), O ($Z=8$)). The apparent anti-correlation between FIP and $T_c$ prevents to choose between the two mechanisms. Such degeneracy could be solved by measuring the abundances of the few volatile low-FIP elements and the rare refractory with high-FIP ones \\cite{meyer}; e.g. Ga ($Z=31$), Ge ($Z=32$), Rb ($Z=37$), Pb ($Z=82$), Bi ($Z=83$) \\cite{waddington}.  A number of signatures exists in the abundance spectrum of heavy nuclei ($30<Z<74$) \\cite{waddington}; in order to reach confident conclusions it is necessary to achieve high-resolution charge measurements and high statistics. Elements with $Z>28$ are rare and large exposure factors (collecting area $\\times$ exposure time $\\times$ angular acceptance) are required. For example the abundances of $Z>30$ nuclei is 1000$\\div$10000 times smaller than Fe nuclei, whose flux is $\\sim$1 cm$^{-2}$sr$^{-1}$day$^{-1}$.  Since many years high altitude balloon flights provided an alternative to more expensive (even in terms of realization time) satellite based experiments. Flight durations have been growing from few days to several weeks on recent Long Duration Balloons (LDB) flown from Antartic regions  e.g. \\cite{tiger}, and plans are underway to provide missions up to months (Ultra LDB).  Plastic passive detectors are promising since they can provide good charge resolution, and are certainly suitable for balloon flights, as they can be used in large collecting areas and limited weight. The main limitation is that the detector has to be recovered after the flight for the analysis with microscopes. Such type of technique was already used succesfully in several experiments on balloons [4-6] % and in space [7-10] %  Nuclear Track Detectors (NTDs) like CR39$^{\\textregistered}$ can detect charged particles with $Z/\\beta\\geq5$ with very good charge resolution \\cite{bologna}, whereas Lexan$^{\\textregistered}$/Makrofol$^{\\textregistered}$ polycarbonates have higher thresholds, $Z/\\beta\\geq50$ \\cite{makrofol0}. Due to the large flux of low charged particles the analysis of CR39$^{\\textregistered}$ detectors and the determination of the relative abundances can be time consuming. Thus the necessity to develop an automatic system for reading and measuring the track characteristics.  The purpose of the CAKE balloon experiment from the ASI (\\emph{Agenzia Spaziale Italiana}) was to test the implementation of new scanning and analysis techniques for future LDB flights of large area detectors (tens of m$^2$sr) \\cite{cake1}. We report here about the scanning procedures and algorithms developed to analyse the CR39$^{\\textregistered}$ sheets exposed during a trans-mediterranean flight.  The obtained primary charge spectrum at the top of the atmosphere is presented. \\vspace{-0.3cm} ",
        "conclusions": "\\label{sec:conclusions} This paper was focused on the development of automatic analysis procedures for determining the chemical composition of cosmic ray nuclei using nuclear track detectors. The NTD technique offers good charge resolution and is attractive for the exploration/determination of the charge spectrum of cosmic particles at reasonable cost. The explotation of this technique in large area experiments relies on automatic image analyzers for scanning and measuring the detector foils as well as on the availability of LDB and ULDB flights.  By making use of a short stratospheric balloon flight we had the opportunity to study and develop several tools to obtain a charge spectrum of cosmic ray particles with $Z/\\beta\\geq10$ at the top of the atmosphere. Our measurement is in fair agreement with other observations. \\vspace{-0.3cm}"
    },
    "0911/0911.4927_arXiv.txt": {
        "abstract": "We search for stars with proper motions in a set of deep Subaru images, covering about 0.48 square degrees to a depth of $i' \\simeq 26$, taken over a span of five and a half years. We follow the methods described in \\citet{Richmond2009} to reduce and analyze this dataset. We present a sample of 69 stars with motions of high significance, and discuss briefly the populations from which they are likely drawn. Based on photometry and motions alone, we expect that 14 of the candidates may be white dwarfs. Our candidate with the largest proper motion is surprisingly faint and likely to prove interesting: its colors and motions suggest that it might be an M dwarf moving at over 500 km/sec or an L dwarf in the halo. ",
        "introduction": "This paper continues our effort to use the Subaru telescope to search for proper motions among very faint stars. Our first paper (\\cite{Richmond2009}) examined a series of images of the Subaru Deep Field (SDF) (\\cite{Kashikawa2004}), located close to the Northern Galactic Pole ($l = 37{\\rlap.}^{\\circ}6$, $b = +82{\\rlap.}^{\\circ}6$). We report here on a similar analysis of images of the Subaru/XMM-Newton Deep Survey (SXDS) field (\\cite{Sekiguchi2005}, \\cite{Furusawa2008}, \\cite{Morokuma2008}), which lies at high galactic latitude in the southern galactic hemisphere ($l \\simeq 170^{\\circ}$, $b \\simeq -60^{\\circ}$). As in our earlier work, we are taking advantage of a dataset compiled for the study of very distant extragalactic objects; the long time coverage necessary to detect supernovae and AGN provides us with the baseline needed to measure proper motions for a significant number of stars in our own Milky Way. We pay special attention to members of the faintest stellar populations, white dwarfs.  Since we will follow very closely the methods used in our analysis of the SDF, we urge the reader to consult \\citet{Richmond2009} for a detailed description of some procedures we may mention only briefly here. However, the SXDS dataset differs from the SDF dataset in one very important way: it consists of a mixture of images taken through two passbands, the Suprime-Cam $R_c$ and $i'$ (\\cite{Miyazaki2002}), while the SDF data was all taken through $i'$. This inhomogeneity complicates efforts to determine the selection effects which define our sample of moving objects; therefore, we discuss in depth our tests of the magnitudes and motions to which our analysis of the SXDS is reasonably complete. One of our goals in this project is to compare observed sets of moving stars to those predicted by various models of stellar populations in the Milky Way. We will need a good understanding of the selection effects on the observations in order to compare them fairly to models in a future paper.  Section 2 describes the observations and the steps we took to convert the raw images into clean, seamless mosaic images. We list in section 3 our procedures for finding and measuring the properties of stars in the images, yielding a list of formal motions for tens of thousands of stars. Most of these motions, of course, were not significantly different from zero, and so we discuss in section 4 our techniques for selecting a small subset of stars with significant proper motions. We used artificial stars inserted into our images to estimate the completeness of our sample as a function of magnitude and motion, and, to a limited degree, color. In section 5, we briefly compare our sample of moving stars with those found in the SDF and predicted in the the Besan\\c{c}on model (\\cite{Robin2003}). We also highlight one very interesting star, which combines a large proper motion with a very faint apparent magnitude. ",
        "conclusions": "Following \\citet{Richmond2009}, we computed the reduced proper motion for the candidates in our sample in order to separate stars belonging to different populations. Figure \\ref{fig:reducedpm} compares stars in our sample to stars in a simulation of the region based on the Besan\\c{c}on model. The regions drawn in the diagram are exactly the same as those shown in Figures 7 and 8 of \\citet{Richmond2009}. We adopt their relationship \\begin{equation} (V - I) = 0.391 + 1.1145(V_s - i') \\end{equation} to compare our results to the models.   \\begin{figure*} \\begin{center} \\FigureFile(150mm,240mm){Figure7.eps} \\end{center} \\caption{Reduced proper motion diagram for observed objects in the SXDS and simulated objects based on the Besan\\c{c}on model.  Real objects have errorbars in both directions, though some are too small to see. The object at lower right is placed according to lower limits to its $V$-band magnitude; see text for details. \\label{fig:reducedpm} } \\end{figure*}  Of the 68 objects in our sample, 13 lie within the ``white dwarf'' (WD) region of the diagram, 32 within the ``halo'' region, 12 within the ``disk'' region, 10 within the ``halo sd?'' region, and 1 falls far from all the rest (we discuss that outlier below). The fraction of stars in each region agree reasonably well with the fractions found in our survey of the SDF. Although the fraction of objects in the WD region (19\\%) is somewhat higher than in the SDF (9\\%), we do not consider the difference significant, given the small number of objects in each set, and the number of objects lying very close to the somewhat arbitrary boundaries of the WD region in each sample.  We performed multiple simulations using the Besan\\c{c}on model to generate stars in the SXDS over an effective area of two square degrees. We applied the selection criteria derived in Section 3 to these simulated catalogs, then scaled the results to the actual area we observed (0.48 square degrees). We found a total of 12.5 WD in the simulation, which agrees well with the 13 candidates in the WD region of our reduced proper motion diagram. Over half (56\\%) of the WDs in the simulation were members of the halo, with most of the remainder (34\\%) drawn from the thick disk.    One of the candidates, SXDSPM J021840.0-053623, deserves individual attention. It has the highest proper motion in our sample, $0{\\rlap.}^{''}186$ per year (see Figures \\ref{fig:candmotiondiag} and \\ref{fig:candpanels}), and yet is very faint: $R_c = 26.09$, $i' = 24.38$, $z = 22.62$. The object is very red: our attempt to measure a magnitude at the object's position in the $V_s$-band image yielded a formal value of $V_s \\sim 28.3$, but we choose instead to quote a lower limit of $V_s \\geq 27.3$ based on secure measurements of faint stars nearby. Its colors $(R_c - i') = 1.69$ and $(i' - z) = 1.76$ place it far from the locus of cool main sequence stars. If we use the more reliable $(i' - z)$ color alone, the star might be an M9 dwarf (\\cite{Hawley2002},\\cite{West2005}) with an absolute magnitude $M_i \\simeq 15.5 \\pm 0.5$. However, in that case, the distance to the object would be $\\sim 600 \\pm 120$ pc, and the tangential velocity $v_t \\sim 520 \\pm 180$ km/s. The $(i' - z)$ color of an L2 dwarf is not very different from our value, and the absolute magnitude of such a star, $M_i \\simeq 17.0 \\pm 0.5$, would yield a lower velocity, $v_t \\sim 260 \\pm 90$ km/s, consistent with the halo population. The residuals from a simple linear fit to the object's motion show no sign of parallax, which indicates that it must be at least $\\sim 20$ pc away from the Sun.   \\begin{figure} \\begin{center} \\FigureFile(80mm,80mm){Figure8.eps} \\end{center} \\caption{Motion of SXDSPM J021840.0-053623 in Right Ascension over the period Sep 29, 2002 to Sep 28, 2005. We could not determine a precise position of the object in the $R_c$-band image of Jan 9, 2008. \\label{fig:candmotiondiag} } \\end{figure}   \\begin{figure*} \\begin{center} \\FigureFile(150mm,240mm){Figure9.eps} \\end{center} \\caption{Images of SXDSPM J021840.0-053623 (at center) in $i'$-band over a three-year period. North is up and East to the left. Vertical lines have been drawn through the centers of objects in the 2002 image to show relative east-west motion clearly. \\label{fig:candpanels} } \\end{figure*}    The proper motions computed in this paper complement those we described in \\citet{Richmond2009}, since the SXDS lies in the southern galactic hemisphere and the SDF in the northern. Together, they comprise a set of 167 stars with reliable motions and multicolor photometry. We have characterized the selection effects of both samples as a function of magnitude, proper motion, and (for the SXDS) color. Our next step will be to compare rigorously the properties of the stars in our samples with the predictions made by current models of galactic populations, paying special attention to the objects likely to be white dwarfs.   \\medskip We thank the staff at the Subaru Telescope for their assistance with the observations used in this project. MWR gratefully acknowledges grant S-03031 from the JSPS Invitation Fellowship for Research in Japan. TM is financially supported by the Japan Society for the Promotion of Science (JSPS) through the JSPS Research Fellowship. Data analysis was in part carried out on the common use data analysis computer system at the Astronomy Data Center, ADC, of the National Astronomical Observatory of Japan."
    },
    "0911/0911.4861_arXiv.txt": {
        "abstract": " ",
        "introduction": "Understanding how and when the first stars and galaxies formed is one of the most fundamental problems in astronomy. Furthermore, whilst many sophisticated models of early galaxy formation and evolution have been constructed, it is clear that observations of the most distant galaxies are mandatory to really test, refine, or refute such models. Indeed, considerable manpower and telescope time have been invested in such observations, with the detection of a Gamma Ray Burst (GRB) at $z\\approx 8.2$ \\citep[][]{Tanvir} being one of the most recent highlights of this extraordinary endeavor. However, despite the recent success in using GRBs to find the most distant sources, the current samples of high redshift galaxies have been mostly assembled using two methods: the broad-band drop-out technique and narrow-band imaging surveys.  The widely-used drop-out technique \\citep[pioneered at $z\\sim3$ by][]{steidel96} requires very deep broad-band imaging, and can identify $z>7$ galaxies as z-band drop-outs \\citep[e.g.][]{bouwens08,richard08}. Furthermore, the use of this technique, combined with the recent installation of the Wide Field Camera 3 (WFC3) on the Hubble Space Telescope (HST), has led to the identification of roughly 20 $z\\approx7-8$ candidates \\citep[e.g.][]{2009arXiv0909.1803B,Oesch,McLure09}. While this is an efficient method for identifying candidates, it still requires detailed spectroscopic follow-up to confirm them, especially to rule out contributions from other populations with large z$-J$ breaks, such as dusty or evolved $z\\sim2$ galaxies and ultra-cool galactic stars \\citep[e.g.][]{mclure06}. Confirming the candidates is actually quite a significant challenge, since the typical $z>7$ candidates found so far are just too faint for spectroscopic follow-up.  The narrow-band imaging technique has the advantage of probing very large volumes looking for Ly$\\alpha$ in emission. Whilst it can only detect sources with strong emission lines, and still depends on the Lyman-break technique to isolate very high-redshift emitters, it can yield suitable targets for follow-up spectroscopy with the current instrumentation. Narrow-band Ly$\\alpha$ searches at $3<z<7$ have been extremely successful in detecting and confirming emitters \\citep[e.g.][]{hu99} and, so far, this technique has resulted in the spectroscopic confirmation of the highest redshift galaxy \\citep[$z=6.96$:][]{Iye}. Even more recently, \\cite{Hibon} identified 7 candidate Ly$\\alpha$ emitters at $z=7.7$. There have been attempts to detect Ly$\\alpha$ emitters at $z\\sim9$ \\citep[e.g.][]{willis05,cuby07,willis08}, but all such studies have been unsuccessful to date, having surveyed very small areas (a few tens of square arcmins at most).  With the advent of wide-field near-IR detectors it is now possible to increase the sky area studied by over 2 to 3 orders of magnitude and reach the regime where one can realistically expect to detect $z \\sim 9$ objects. This is a key aim of, for example, the narrow-band component of the UltraVISTA Survey \\citep[c.f.][]{Nilsson07}. It is also an aim of HiZELS, the Hi-Z Emission Line Survey \\citep[c.f.][]{Geach,Sobral,Sobral2,Garn}, that we are carrying out using the WFCAM instrument on the 3.8-m UK Infrared Telescope (UKIRT). HiZELS is using a set of existing and custom-made narrow-band filters in the $J$, $H$ and $K$ bands to detect emission lines from galaxies at different redshifts over $\\sim$ 10 square degrees of extragalactic sky. In particular, the narrow-band $J$ filter (hereafter NB$_{\\rm J}$) is sensitive to Ly$\\alpha$ at $z=8.96$. ",
        "conclusions": "\\begin{itemize}  \\item Deep narrow-band imaging in the $J$ band ($\\lambda=1.211\\pm0.015\\mu$m) has been used to search for bright Ly$\\alpha$ emitters at $z=8.96$ over an area of 1.4 deg$^2$. No Ly$\\alpha$ emitter was found brighter than ${\\rm L}\\approx7.6\\times10^{43}$\\,erg\\,s$^{-1}$. \\item The Ly$\\alpha$ luminosity function constraints at $z\\sim9$ have been significantly improved for $L>10^{43.8}$\\,erg\\,s$^{-1}$ and combined with constraints from deeper but much smaller previous surveys. The results rule out significant positive evolution of the Ly$\\alpha$ LF beyond $z\\sim6$; they are in line with recent semi-analytic \\& phenomenological model predictions, rejecting some extreme models. \\item It has been shown that for narrow-band searches, T-dwarfs can mimic Ly$\\alpha$ emitters at $7.7<z<8.0$, $9.1<z<9.5$ and $11.7 < z < 12.2$; they will not contaminate the future UltraVISTA narrow-band survey (and can even be identified via a narrow-band {\\it deficit}), but they may contaminate narrow-band Ly$\\alpha$ searches within the $z=7.7$ and $z=9.4$ atmospheric windows. \\end{itemize}  These results show that bright ${\\rm L}>10^{43.8}$\\,erg\\,s$^{-1}$ Ly$\\alpha$ emitters are extremelly rare. Although the area coverage is absolutely important, a depth+area combination is likely to be the best approach for gathering the first sample of these very high-redshift galaxies. In fact, that is the strategy of the narrow-band component of the UltraVISTA survey \\cite[c.f.][]{Nilsson07}, using the VISTA telescope, which will map 0.9 deg$^2$ of the COSMOS field to a planned 5$\\sigma$ luminosity limit of $L=10^{42.53}$\\,erg\\,s$^{-1}$ and a surveyed volume of 5.41$\\times10^5$ Mpc$^3$ (see right panel of Figure \\ref{fig2}) at $z=8.8$. This combination lies below all current predictions for the $z\\sim9$ Ly$\\alpha$ LF and the survey is expected to detect 2-20 Ly$\\alpha$ emitters at $z=8.8\\pm0.1$. Furthermore, the continuation of HiZELS on UKIRT and the extension of the narrow-band $J$ survey to a wider area offers a complementary approach which might be able to detect one of the brightest Ly$\\alpha$ emitters at $z\\sim9$, perfectly suited for spectroscopic follow-up and potentially enabling the detailed studies which simply won't be possible for much fainter emitters, even if they are detected.   \\bigskip  \\small"
    },
    "0911/0911.3994_arXiv.txt": {
        "abstract": "{The possibility that the non-minimal coupling inflation could be eternal is investigated. We calculate the quantum fluctuation of the inflaton in a Hubble time and find that it has the same value as in the minimal case in the slow-roll limit. Armed with this result, we have studied some concrete non-minimal inflationary models including the chaotic inflation and the natural inflation while the inflaton is non-minimally coupled to the gravity and we find that these non-minimal inflations could be eternal in some parameter regions. } ",
        "introduction": "Inflation has been remarkably successful in explaining the properties of the universe and the origin of the primordial perturbation \\cite{inflation}, which is thought of as the seed of the large scale structures. In the context of new inflation, the early inflating universe is  driven by a scalar field called inflaton, which is at first living at an unstable state like the false vacuum state, then it slowly rolls down to a stable state like the true vacuum state. An interesting phenomenon in the inflationary scenario is the eternal inflation, which means the inflation never ends.  The eternal new inflation was first discovered by Steinhardt \\cite{steinhardt-nuffield}, and Vilenkin \\cite{Vilenkin:1983xq} showed that new inflationary models are generically eternal. As we known, the decay of the false vacuum  to a true vacuum is an exponential process, however, it also exponentially expands when it decays during the inflation time. In fact, the rate of exponential expansion is always much faster than the rate of exponential decay in any successful inflationary models. Therefore, the false vacuum never disappears and the total volume of the false vacuum grow exponentially with time when inflation starts, see ref \\cite{Guth:2007ng}.  Actually, some inflationary models like chaotic inflation does not have a false vacuum state, but it can also be eternal, which is called slow-roll eternal inflation \\cite{Linde:1986fd}. In these models, the inflaton is classically rolling down the hill, and the change in the field during some time interval is influenced by the quantum fluctuations, which can drive the field upward or downward relative to the classical trajectory. There is always some probability that the classical evolution is smaller than the quantum fluctuations, then the inflaton will fluctuate up and not down. Therefore, this process will continues forever and inflation will never ends. The recent progress on eternal inflation, see ref. \\cite{recent pro eternal} and for recent reviews see ref. \\cite{Guth:2007ng, Linde:2007fr}.  There is a class of inflationary models called non-minimal inflations, in which the inflaton ($\\varphi$) is non-minimally coupled to the gravity and terms like $f(R,\\varphi)$ are introduced in the effective action, where $R$ is the Ricci scalar. The cosmological effects of the non-minimal inflationary models have been well studied. It shows that the power spectrum is generally blue, and the problem of getting a running spectral index is eased in the non-minimal inflation \\cite{miaoLi}. For other recent studies on the non-minimal inflation, see ref. \\cite{recent nonminimal} . However, the eternal property of them has not studied in these literatures. In this paper, we have investigated that whether the non-minimal coupling inflation could be eternal or not, and we focus on the simplest model with  $ f = \\zeta R\\varphi^2$ .  The paper is organized as follows. In Section \\ref{s bd}, we briefly review the background dynamics for the inflaton and the metric, impose the slow-roll condition and define three slow-roll parameters including the ordinary two and a new one. In Section \\ref{flu inf}, we calculate the quantum fluctuation and the classical motion  of the inflaton during a Hubble time. Then, in Section \\ref{example}, we focus on some concrete non-minimal inflation models and find that they could be eternal in some parameter regions. In the last section, we give some discussion the conclusions. ",
        "conclusions": "Many inflationary models have eternal properties like the new inflation and the chaotic inflation. In this paper, we have studied the  eternal properties of  non-minimal inflationary models, in which there is a non-minimal coupling between the inflaton and the gravity. We have calculated the quantum fluctuation of the inflaton at the first order, and find that it has the same value as  in the minimal case, namely, the quantum fluctuation of inflaton in a Hubble time is roughly proportional to the Hubble parameter during the inflation. If the quantum fluctuation overcome the classical motion of the inflaton, the inflation could  never end.  We have studied some concrete non-minimal inflationary models including the chaotic inflation with power law potentials and the natural inflation, derived the conditions under which the inflation could be eternal and found that in some parameter spaces, it could be eternal inflation. Our results could be simply generalized to the case of $f \\sim R\\varphi^n$, where $n=2$ is studied in this paper."
    },
    "0911/0911.2454_arXiv.txt": {
        "abstract": "{} {Cosmic shear is a powerful method to constrain cosmology, provided that any systematic effects are under control. The intrinsic alignment of galaxies is expected to severely bias parameter estimates if not taken into account. We explore the potential of a joint analysis of tomographic galaxy ellipticity, galaxy number density, and ellipticity-number density cross-correlations to simultaneously constrain cosmology and self-calibrate % unknown intrinsic alignment and galaxy bias contributions.} {We treat intrinsic alignments and galaxy biasing as free functions of scale and redshift and marginalise over the resulting parameter sets. % Constraints on cosmology are calculated by combining the likelihoods from % all two-point correlations between galaxy ellipticity and galaxy number density. The information required for these calculations is already available in a standard cosmic shear data set. % We include contributions to these functions from cosmic shear, intrinsic alignments, galaxy clustering and magnification effects.} {In a Fisher matrix analysis we compare our constraints with those from cosmic shear alone in the absence of intrinsic alignments. For a potential future large area survey, such as Euclid, % the extra information from the additional correlation functions can make up for the additional free parameters in the intrinsic alignment and galaxy bias terms, depending on the flexibility in the models. For example, the Dark Energy Task Force figure of merit is recovered even when more than 100 free parameters are marginalised over. We find that the redshift quality requirements are similar to those calculated in the absence of intrinsic alignments.} {} ",
        "introduction": "\\label{sec:intro}  On its way to Earth the light from % distant galaxies is continuously deflected by the matter density inhomogeneities which it passes by. This induces distortions of the shapes of the projected galaxy images on the sky, causing modifications of these shapes of order $1\\,\\%$, known as weak gravitational lensing, or cosmic shear. % Hence, to measure this distortion effect on distant galaxies, one requires statistical methods; see e.g. \\citet{bartelmann01} and \\citet{schneider06} for detailed reviews.  Detected in 2000 \\citep{bacon00,kaiser00,vwaer00,wittman00}, cosmic shear has since rapidly evolved into a mature technique that produces increasingly stringent constraints on cosmological parameters \\citep[see e.g.][]{jarvis06,hoekstra06,semboloni06,hetterscheidt07,benjamin07,massey07,schrabback07,fu07,schrabback09}. Probing both the evolution of structure and the geometry of the Universe at low redshifts, it is complementary to observations of the cosmic microwave background \\citep[e.g.][]{tereno05,das08} % and considered as potentially % the most powerful method to pin down the properties of dark energy \\citep{hu02b,albrecht06,peacock06}. Upcoming and future surveys like Pan-STARRS\\footnote{Panoramic Survey Telescope \\& Rapid Response System, \\texttt{http://pan-starrs.ifa.hawaii.edu}}, DES\\footnote{Dark Energy Survey, \\texttt{https://www.darkenergysurvey.org}}, LSST\\footnote{Large Synoptic Survey Telescope, \\texttt{http://www.lsst.org}}, JDEM\\footnote{Joint Dark Energy Mission, \\texttt{http://jdem.gdfc.nasa.gov}}, and Euclid\\footnote{\\texttt{http://sci.esa.int/science-e/www/area/\\\\index.cfm?fareaid=102}} will produce deep imaging over large fractions of the sky and thereby yield unprecedented insight into the properties of dark matter, dark energy and gravitation \\citep[e.g.][]{takada04,refregier04,refregier06,kitching08b,thomas08}. %  The great % statistical power of cosmic shear demands a careful % assessment of possible systematic errors that might bias the results. A serious limitation may arise from a physical systematic caused by the intrinsic alignment of galaxies. The matter structure around galaxies can modify their intrinsic shape and their orientation. Firstly, this can result in correlations between the intrinsic shapes of galaxies which are close both on the sky and in redshift (intrinsic ellipticity correlations, or II correlations). Moreover, a dark matter halo can intrinsically align a physically close galaxy in the foreground and at the same time contribute to the lensing signal of a background object, which induces gravitational shear-intrinsic ellipticity correlation (GI, \\citealp{hirata04}).  Detailed investigations have been performed on the alignment between haloes \\citep{croft00,heavens00,lee00,catelan01,crittenden01,jing02,mackey02,hirata04,bridle07a,schneiderm09}, as well as the alignment of the spin or the shape of a galaxy with its own dark matter % halo (e.g. \\citealp{pen00,bosch02,okumura08,okumura09,brainerd09}; see also \\citealp{schaefer08}). Intrinsic alignments have also been investigated observationally, where non-vanishing II and GI signals have been detected in several surveys \\citep{brown02,heymans04,mandelbaum06,hirata07,brainerd09}.  The results of both theoretical studies and observations show large variations, but most are consistent with a contamination of the order $10\\,\\%$ by both II and GI correlations for future surveys that further divide the galaxy sample into redshift slices (cosmic shear tomography). Hence, the control of intrinsic alignments in cosmic shear studies is crucial to obtain unbiased results on cosmological parameters. Accurate models would solve the problem, % but progress is hampered due to the dependence of intrinsic alignments on the intricacies of galaxy formation and evolution within their dark matter environment. Currently, the level of models is crude, and partly only motivated phenomenologically (see e.g. \\citealp{schneiderm09} for recent progress).  The II contamination can be controlled relatively easily by excluding close pairs of galaxies from the analysis \\citep{king02,king03,heymans03,takada04b}. \\citet{joachimi08b,joachimi09} introduced a nulling technique which transforms the cosmic shear data vector and discards all entries of the transformed data set that are potentially contaminated by the GI signal. While this approach only relies on the well-known redshift dependence of gravitational lensing, \\citet{king05} projects out the GI term by making use of template functions. Furthermore the work by \\citet{mandelbaum06} and \\citet{hirata07} suggests that the intrinsic alignment may be dominated by luminous red galaxies which could be eliminated from the cosmic shear catalogues. All these removal techniques require excellent redshift information, and still they can cause a significant % reduction in the constraints on cosmology.  Deep imaging surveys not only provide information about the shape of galaxies, but allow in addition for a measurement of galaxy number densities, as well as cross-correlations between shape and number density information. This substantial extension of the set of observables increases the cosmological information to be extracted and, more importantly, enables one to internally calibrate systematic effects \\citep{hu03,bernstein08}. By adding galaxy number density information one adds signals that are capable of pinning down the functional form of intrinsic alignments, but one also introduces as another systematic, the galaxy bias, which quantifies the lack of knowledge about how galaxies, i.e. the visible baryonic matter, follow the underlying dark matter distribution.  It is the scope of this work to elucidate the performance of a joint analysis of galaxy shape and number density information as regards the ability to constrain cosmological parameter in presence of general and flexible models of intrinsic alignments and galaxy bias. In doing so we incorporate several cosmological signals which have been considered before as promising probes of cosmology themselves, including galaxy clustering from photometric redshift surveys % \\citep[][]{blake05,dolney06,zhan06,blake07,padmanabhan07} % galaxy-galaxy lensing \\citep[e.g.][]{schneider97,guzik01,guzik02,seljak02,seljak05,yoo06,johnston07,cacciato08} % and lensing magnification \\citep[][]{broadhurst95,zhang05,zhang06,vwaer09}. % We follow the ansatz outlined in \\citet{bernstein08} and extend the investigation by \\citet{bridle07} who considered the residual information content in galaxy shape correlations after marginalising over the parameters of two log-linear grid models representing the II and GI terms. We quantify the cross-calibration properties of the joint set of observables and determine the requirements on cosmological surveys to efficiently apply this joint approach.  This paper is organised as follows: In Sect.$\\,$\\ref{sec:method} we give an overview on the two-point correlations that form part of the galaxy shape and number density observables, and we derive their explicit form. Two appendices provide further details. Section \\ref{sec:modelling} demonstrates how we model the different signals and their dependence on cosmology. We introduce a general grid parametrisation for the intrinsic alignments and the galaxy bias. Furthermore we summarise our Fisher matrix formalism and the figures of merit we employ. In Sect.$\\,$\\ref{sec:results} we present our results on the dependence of the parameter constraints on the freedom in the model of intrinsic alignments and galaxy bias, the characteristics of the redshift distributions, and the priors on the different sets of nuisance parameters. Finally, in Sect.$\\,$\\ref{sec:conclusions} we summarise our findings and conclude. ",
        "conclusions": "\\label{sec:conclusions}  In this work we studied the joint analysis of galaxy number density and shape correlations to constrain cosmological parameters in presence of contaminations by the intrinsic alignment of galaxies and the galaxy bias. We considered the signals due to gravitational shear, intrinsic shear, intrinsic galaxy clustering, and lensing magnification, explicitly computing all possible two-point correlations thereof. We introduced a two-dimensional grid parametrisation to account for the unknown scale % and redshift dependence of both intrinsic alignments and galaxy bias. Further nuisance parameters were used to describe the uncertainty in the width of the photo-z bin-wise redshift distributions and in the slope of the galaxy luminosity functions within each bin.  Our Fisher matrix analysis demonstrates that the simultaneous use of ellipticity correlations, number-density correlations, and in particular ellipticity-number-density cross-correlations allows for a substantial amount of internal calibration of the bias terms. With flexible models that contain in total more than 200 nuisance parameters we can recover the volume of the error ellipsoid in parameter space when compared to assuming a pure gravitational lensing signal and using just the ellipticity correlation information. % The dark energy parameters $w_0$ and $w_a$ suffer more than other parameters on % marginalisation over nuisance parameters, so that $56\\,\\%$ of the ${\\rm FoM}_{\\rm DETF}$ are retained for a Euclid-like survey in this most flexible setup considered. % The ${\\rm FoM}_{\\rm DETF}$ for the combined set of shape and number density correlations is close to the pure lensing ${\\rm FoM}_{\\rm DETF}$ if we choose a model which uses about 100 nuisance parameters to describe intrinsic alignments and galaxy bias. Our approach also proves beneficial for upcoming ground-based surveys with DES-like survey characteristics. % The benefit is greatest for the more ambitious survey. % In addition, we assumed the slopes of the galaxy luminosity functions in each photo-z bin to be unknown nuisance parameters and found that they are well calibrated internally from the data.  The information which we added on top of the standard cosmic shear analysis comes without any extra cost since galaxy number density measurements are directly available from imaging data. Given our encouraging findings, we hence suggest that the joint consideration of galaxy shape and number density information could become the standard technique whenever intrinsic alignments are suspected to make a significant contribution to the cosmic shear signal.  When interpreting our results, one has to keep in mind that our grid parametrisation of the bias terms has limited flexibility. On increasing the number of nuisance parameters in the grids, the curves for the figures of merit flatten off. However we are at present not able to say whether they approach a lower limit or continue to decrease for a large number of grid nodes. This issue, which is of considerable theoretical and practical interest, is currently under investigation, as well as a comparison with the performance of removal techniques such as nulling \\citep{joachimi08b,joachimi09}. %  Due to the finite number of grid points the resulting angular power spectra will not span perfectly the range of possible physical models. % Moreover, we found that of order 10 nuisance parameters in $k$ per grid node in redshift do not suffice to represent all relevant functional forms of the galaxy bias. Therefore we removed all observables from the parameter estimation which have a significant contribution from the galaxy clustering signal at large $k$ % where the signal is strongest and the bias least known. However, since in the near future we expect to have at least coarse, but reliable knowledge about the functional forms of the galaxy bias on linear scales and the intrinsic alignment signals, the models used in our approach should still yield realistic results.  While the log-linear grid parametrisation is fairly general and intuitive, it may not be the most efficient way to represent freedom in intrinsic alignments and the galaxy bias. As all bias terms originate from physical processes, they are smooth and do not feature strong oscillations or isolated peaks. Thus, we presume that the bias terms can efficiently be parametrised in terms of complete sets of smooth functions such as the Fourier, Chebyshev, or Legendre series. Truncating the higher orders of these series will only limit the model to represent highly oscillatory or small-scale features. These parametrisations might therefore be more efficient in comprising the full set of realistic bias terms \\citep[see also][]{kitching08} for a given number of nuisance parameters, this number in turn being dictated by the available computational power. The latter will play a critical role when the approach suggested here is performed in a full likelihood analysis with several hundreds of nuisance parameters. % We add that it might be necessary to furthermore consider the bias terms individually for different galaxy types and luminosities % because it is known % that both the intrinsic alignments and the galaxy bias vary considerably with galaxy type and luminosity % \\citep[for recent examples see][]{mandelbaum06,hirata07,mccracken07, swanson08, simon08, cresswell09, wang07}. %  Redshift information is crucial for discriminating % the different signals that contribute to the observables. We investigated the dependence of the parameter constraints from % the joint set of correlations on characteristics of the redshift distributions. We confirm the observation by foregoing works that the number of photo-z bins needed to retrieve the bulk of information about cosmology increases by at least a factor of two when using only % ellipticity correlations and marginalising % over the intrinsic alignment signals. Using the complete set of correlations, the figures of merit do not level off so quickly, % but continue to increase as the number of tomographic bins is increased. % This can be explained by the narrower kernel in the redshift integrations that link the projected and the three-dimensional power spectra involving galaxy number-density signals.  Moreover the figures of merit decrease substantially more slowly % as a function of the photometric redshift dispersion $\\sigma_{\\rm ph}$ when adding galaxy number density information to the data. However, in our marginalisation we only used a single parameter accounting for the uncertainty in the spread of all bin-wise redshift distributions. This way the redshift cross-terms of the number-density correlations efficiently calibrate the shape of the redshift distributions. Therefore this result alone cannot % be interpreted as a potential relaxation of the requirements for photometric redshift accuracy in cosmological surveys featuring cosmic shear. Conversely, the continuing increase in the figure of merit as a function of $N_{\\rm zbin}$ makes an even larger number of well-separated and compact redshift distributions desirable in the case of the joint data set.  The recent work by~\\citet{zhang09} uses density-ellipticity correlations to self-calibrate photometric redshift parameters but a full approach may need to simultaneously deal with intrinsic alignments and photometric redshift properties if sufficient spectra cannot be obtained to calibrate the photometric redshifts independently. % In a forthcoming paper we will investigate into a more realistic approach including uncertainty and outliers in the redshift distributions of each individual photo-z bin, as well as the benefits of spectroscopic redshift information for a subsample of the galaxy catalogue. %  We emphasise that in the approach suggested here all considered correlations help in % constraining cosmological parameters. Hence, none of the contributions is regarded as a systematic signal, and the different levels of uncertainty concerning the exact form of the signals are represented by nuisance parameters. By means of the joint analysis we increase the statistical power via the cross-calibration abilities of the signals, and reduce the risk of undetected systematic effects one faces when considering signals such as cosmic shear, galaxy clustering, or lensing magnification individually.  This integrative ansatz is not limited to two-point correlations, but can be generalised to the three-point level in a straightforward way. Future surveys will provide excellent data for studying three-point correlations whose exploitation can break parameter degeneracies and improve constraints considerably \\citep[e.g.][]{takada04}. Our knowledge about intrinsic alignments and galaxy bias at the three-point level is currently even more limited than in the two-point case (however see \\citealp{semboloni08}), % so that a joint investigation of shape and number density observables, including a general parametrisation of the various bias terms, may be an appropriate way forward. However, limitations due to computational power will most likely play a dominant role in this case.  One can also think of incorporating further sets of observables into the analysis. For instance, planned surveys like Euclid will also include a spectroscopic survey of a subset of the galaxies in order to determine the underlying redshift distributions of the photo-z bins. The spectroscopic data can be used to measure baryonic acoustic oscillations (note that the acoustic peaks are not included in our galaxy number-density correlations) and galaxy peculiar velocities. The latter allow for the measurement of redshift space distortions. Similar to the case considered here, the joint analysis of galaxy number density, shape, and velocity information \\citep{guzik09} will efficiently cross-calibrate nuisance terms such as the galaxy bias and tighten constraints on the cosmological model."
    },
    "0911/0911.2986_arXiv.txt": {
        "abstract": "We study magnetothermal instability in the ionized plasmas including the effects of Ohmic, ambipolar and Hall diffusion. Magnetic field in the single fluid approximation does not allow transverse thermal condensations, however, non-ideal effects highly diminish the stabilizing role of the magnetic field in thermally unstable plasmas. Therefore, enhanced growth rate of thermal condensation modes in the presence of the diffusion mechanisms speed up the rate of structure formation. ",
        "introduction": "\\label{sec:1} Role of the magnetic field in the dynamics of the gaseous astrophysical systems is generally studied within the framework of {\\it ideal} MHD equations. In this simplified approach, which is good under certain conditions and circumstances, it is assumed that the coupling between the charged and neutral species of the system is perfect. But the perfect coupling assumption can be violated, in particular, when the density of the charged particles can be much lower than that of the neutral species. For example, one should note to this fact in dense molecular clouds \\citep[e.g.,][]{ciolek2002}. Some authors have also criticized application of ideal MHD equation in modeling accretion discs, particular for the discs around young stellar objects \\citep[e.g.,][]{wardle2004}.   The very existence of various species  with different masses and electrical charges and their collisions and the possible momentum transfer between the particles should be considered in any theory of structure formation in interstellar medium (ISM). However, interest towards modeling astrophysical plasmas within a multifluid approach has been raised over recent years \\citep[e.g.,][]{falle03, turlough07, galli08, pandey2008, inutsuka2008}. In particular, simulating multifluid systems is a challenging area. For example, \\cite{kim09}  incorporating ambipolar diffusion in the strong coupling approximation into a multidimensional magnetohydrodynamic (MHD) code based on the total variation diminishing scheme. More generalized Multifluid numerical schemes are also studied during recent years by some authors \\citep[e.g.,][]{falle03, turlough07}.  On the theoretical side, efforts to understand physics of multifluid plasmas are in progress \\citep[e.g.,][]{tassis, kunz09}. Recently,  \\cite*{pandey2008} (hereafter PW) clarified the relationship between the fully ionized and weakly ionized limits by developing a unified single-fluid framework for the dynamics of the plasmas of arbitrary ionization. In another related study, general expressions for the resistivities, the diffusion timescales and the heating rates in a three-fluid medium are obtained by \\cite*{galli08}.  They showed that the value of the Ohmic resistivity is increased in a collapsing cloud and the ambipolar diffusion occurs on a time scale comparable to the dynamical time scale.  Among various physical mechanisms responsible for the density inhomogeneities in the ISM, it has been realized for a long time that thermal instability can be an efficient processes \\citep[e.g.,][]{field65,burkert00,heneb07,inutsuka2008,gazol09}. Thermal instability in a cooling medium can also be affected when dynamics of different charged species (e.g., dust particles) are included \\citep{ibanez02}. It is shown that negatively charged particles stimulate the thermal instability in the sense that the conditions for the instability to hold are wider than similar conditions in a single-fluid description \\citep{Kopp,shukla03}.  Thermal instability within ambipolar regime also studied by \\cite*{asghar03} and \\cite*{asghar07}, where in the former the frictional heating by the ion-neutral is included. In another similar study, \\cite*{fabian06} revisited the problem of clump formation due to thermal instabilities in a weakly ionized plasma. However, dynamics of ions and their contribution to the net cooling function are not included in these studies of thermal instability in ambipolar regime. Therefore, \\cite*{fukue07} extended the linear classical thermal instability to a case, in which dynamics of neutrals and ions and their interactions are considered. They showed that ion-neutral friction with the magnetic field affects the morphology and evolution of the interstellar matter. Thermal instability for a system obeying generalized Ohm's law has also been studied by \\cite*{bora}. They found that the instability criterion involves the field strength, resistivity and electron inertia terms for transverse perturbations, but for the parallel to the ambient magnetic field the instability criterion is independent of all these non-ideal effects. However, they did not apply their analysis to the structure formation due to the thermal instability in the interstellar medium.  In this paper, our goal is to include non-ideal effects in a magnetized thermally unstable plasma. We follow a recent approach by PW, in which ions, electrons and neutrals are included. Then, non-ideal Ohmic, Hall and ambipolar terms appear in a generalized form of the induction equation. We study role of these terms in magnetothermal  instability in the linear regime. The basic assumptions and the equations are presented in the next section.  Linearized equations and the dispersion relation are derived in section \\ref{sec:Linear}. Analyzing effects of non-ideal mechanisms on the thermally unstable modes are done in section \\ref{sec:Analysis}. We conclude by a summary of the results and possible implications in the final section. ",
        "conclusions": "We studied magnetothermal instability with the effect of non-ideal Ohmic, ambipolar and Hall diffusion. Our linear analysis shows that the criteria of instability does not change comparing to the ideal case as long as charged species do not contribute to the net cooling function. Also, the system becomes more unstable in the presence of non-ideal effects and it is more probable to have larger clouds comparing to the ideal case.    Although the vital role of magnetic field dissipation in very dense interstellar clouds is a key process in standard theories of star formation \\citep[e.g.][]{nakano86}, our results show that such non-ideal mechanisms may operate in thermally unstable systems such as HI regions or warm ISM.  \\cite*{inutsuka2008} studied formation of structures  in a weakly ionized and magnetized interstellar medium using two-fluid magnetohydrodynamic simulations. When orientation of magnetic field is perpendicular to the flow, the rate of formation of clouds slows down significantly according to \\cite*{inutsuka2008}. Actually, linear ideal magnetothermal instability shows that magnetic field prevents structure formation transverse to the field lines \\citep{field65}. However, our results show that the stabilizing effect of magnetic field drastically diminishes because of the dissipative processes like Ohmic or ambipolar. Therefore, we think, the opportunity of fast molecular cloud formation directly from the warm neutral medium would highly increase due to the non-ideal effects, at least in the linear regime.  One should note that enhanced growth rate due to the non-ideal terms is independent of the true mechanisms of the operating processes. But we think this independency is valid in the linear regime and the nonlinear evolution of the system will depend on the type of dominant diffusion process. In particular, Hall diffusion significantly differs from Ohmic and ambipolar diffusions regarding to the energy considerations. In fact, it is known that Ohmic and ambipolar diffusion are dissipative processes and reduce magnetic energy of the system. But Hall diffusion does not contribute to the dissipation of the magnetic energy and so, it is not a dissipative mechanism. Its main role goes back to redistributing the current within the system and it may lead to enhanced dissipation of the magnetic energy because of Ohmic or ambipolar diffusions. Numerical simulations show that the dissipation rate of the MHD turbulence is strongly affected by the strength of ambipolar diffusion \\citep{kim09}. But it is an open question to explore possible role of Hall diffusion in nonlinear evolution of thermally unstable system \\citep[see also,][]{wardle2004}."
    },
    "0911/0911.1088_arXiv.txt": {
        "abstract": "{The blazar PG\\,1553+113 is a well known TeV $\\gamma$-ray emitter. In this paper we determine its spectral energy distribution through simultaneous multi-frequency data to study its emission processes. An extensive campaign was carried out between March and April 2008, where optical, X-ray, high-energy (HE) $\\gamma$-ray, and very-high-energy (VHE) $\\gamma$-ray data were obtained with the KVA, Abastumani, REM, \\textit{Rossi}XTE/ASM, AGILE and MAGIC telescopes, respectively. We combine the data to derive the source's spectral energy distribution and interpret its double-peaked shape within the framework of a synchrotron self-Compton model.} ",
        "introduction": "The transformation  of gravitational energy in an accretion disk around a supermassive black hole into radiation is the widely believed underlying cause of emission in active galactic nuclei (AGN). Furthermore, the emission is beamed from the jet perpendicular to the disk by a mechanism that although not fully understood yet, most likely relates to the focusing properties of the rotating, fully ionized accretion disk (e.g. Blandford \\& Znajek 1977). The viewing angle of the observer determines the observed phenomenology of AGN (Urry \\& Padovani 1995). The AGN whose relativistic plasma jets point towards the observer are called blazars. The blazar class includes flat spectrum radio quasars (FSRQs) and BL Lac objects, the main difference between the two classes is in their emission lines, which are strong and quasar-like for FSRQs and weak or absent in BL Lacs.   The overall (radio-to-$\\gamma$-ray) spectral energy distribution (SED) of blazars shows two broad non-thermal continuum peaks. For high-peaked BL Lac objects (HBLs), the first peak of the SED is in the UV/X-ray bands [as opposed to IR/optical  for low peaked BL Lac objects, LBLs], whereas the second peak is in the multi-GeV band (multi MeV for LBLs). The low-energy peak is thought to arise from electron synchrotron emission. The leptonic model sugests that the second peak forms due to inverse Compton emission. This can be due to scattering of the synchrotron photons by the non-thermal population of electrons responsible for the synchrotron radiation (synchrotron self-Compton, SSC; eg: Maraschi 1992) or of photons that originate outside the relativistic plasma blob (external Compton, EC; an external source of these \\textquoteleft seed\\textquoteright photons could be the accretion disk (eg: Dermer 1993) and/or the broadline region (eg: Sikora et al. 1994)). Blazars often show violent flux variability, which may or may not be correlated between the different energy bands. Strictly simultaneous observations are crucial to investigate these correlations and understand the underlying physics of blazars.  The HBL source PG\\,1553+113 was firmly detected at very high-energy $\\gamma$-rays (VHE; photon energy E$>$100 GeV) by the MAGIC telescope at a significance level of 8.8\\,$\\sigma$ above 200 GeV, based on data from April - May 2005 and January - April 2006 (Albert et al. 2007). Observations with the H.E.S.S. telescope array in 2005 yielded a tentative detection in the VHE band, at the level of 4\\,$\\sigma$ (5.3\\,$\\sigma$ using a low energy threshold analysis; Aharonian et al. 2006), which was confirmed later with the combination of the 2005 and 2006 datasets (Aharonian et al. 2008). After the first detection of PG\\,1553+113 with MAGIC, a  multi-frequency campaign on this source was conducted in July 2006 (Albert et al. 2009). The main difference between our present and the previous campaign is the use of X-ray and the high-energy (HE; photon energy E$>$100 MeV) flux.     The lack of detection of spectral lines (neither in emission nor in absorption) in the optical spectrum of PG\\,1553+113 makes it impossible to directly measure its redshift (Falomo \\& Treves 1990). However, an ESO-VLT spectroscopic survey of unknown-redshift BL\\,Lac objects suggests $z>0.09$ (Sbarufatti et al. 2006), while the absence of host galaxy detection in HST images raises this lower limit to $z>0.25$ (Treves et al. 2007). On the other hand, the absence of a break in the VHE spectrum can be interpreted as suggesting $z<0.42$ (Mazin \\& Goebel 2007). The absence of a Ly-$\\alpha$ forest (Falomo \\& Treves) in the the spectrum also constrains a lower redshift. The data obtained in the far-UV by the Cosmic Origins Spectrograph installed in the Hubble Space Telescope is of sufficient quality to select $\\sim 40$ Ly-$\\alpha$ absorbers at low redshift including a strong line at z=0.395, which constrains the resdshift of the source to be $z>0.395$ (Danforth, private comm.). ",
        "conclusions": "Analyzing the MAGIC data, an excess of 415 $\\gamma$-like  events, over 1835 normalized background events was found, yielding  a significance of 8.0\\,$\\sigma$. The resulting differential VHE spectrum of PG\\,1553+113 averaged over all observing intervals is plotted in Fig.1 (filled circles). It can be described by a power law ${dN \\over dE} = F_{0} \\Big( {E \\over 200GeV} \\Big) ^{\\Gamma}$ m$^{-2}$\\,s$^{-1}$\\,TeV$^{-1}$, where $F_{0}$ is the normalization flux at 200 GeV and $\\Gamma$ is the photon index during our observation, which are both given in Table 1. The test on a possible spectral cut-off was also performed. However, fewer points of the spectrum do not favour a cut-off power law over a simple power law. The lowest point of the spectrum is at 82 GeV, mainly because of the losses of low-energy events from the cleaning and $\\gamma$-selection cuts.  The values obtained during our previous observations (Albert et al. 2007) are also given in Table 1.  \\begin{table} \\label{table:spec} \\begin{center} \\begin{tabular}{|p{1.in}|c|c|} \\hline Observation period &  $F_{0}$ [ph\\,TeV$^{-1}$s$^{-1}$m$^{-2}$] &  $\\Gamma$   \\\\ \\hline  March-April\\,2008 & $2.0\\pm0.3 \\times  10^{-6}$ & $-3.4\\pm0.1$ \\\\ \\hline March\\,2008  & $1.9\\pm0.4 \\times 10^{-6}$& $-3.5\\pm0.2$ \\\\ \\hline April\\,2008 & $2.1\\pm0.4 \\times  10^{-6} $ & $-3.3\\pm0.2$ \\\\ \\hline April-May\\,2005 +  January-April\\,2006 & $1.8\\pm0.3 \\times  10^{-6}$  & $-4.2\\pm0.3$ \\\\ \\hline  \\end{tabular} \\end{center} \\caption{$F_{0}$ and $\\Gamma$ during the MAGIC current observations and the previous observation. The errors are statistical only. The systematic uncertainty is estimated to be 35\\% in the flux level and 0.2 in the photon index (Albert et al. 2008a).} \\end{table}  The interaction of VHE $\\gamma$-rays with the extragalactic background light (EBL; a recent review can be found in Mazin \\& Raue 2007) leads to an attenuation of the VHE $\\gamma$-ray flux via $e^+/e^-$ pair production. We computed the de-absorbed (i.e., intrinsic) fluxes with a specific \\textquoteleft low star formation model\\textquoteright of the EBL (Kneiske et al. 2004), assuming a source redshift of $z=0.3$. The resulting de-absorbed points are represented as empty squares in Fig.1.  \\begin{figure}[h] \\label{fig:spec} \\centering \\epsfig{file=spectrum_final.eps,height=6cm,width=9cm,clip=} \\caption{MAGIC measured spectrum of PG\\,1553+113 (filled circles). The statistical significance of the points from left to right are 2.7, 4.2, 3.4, 4.3, and 3.0 $\\sigma$ respectively. The EBL-corrected points are shown as empty squares.  The spectrum obtained during our first observation is shown as a dashed line.} \\end{figure}  The HE data reduction results from AGILE are summarized in Table 2. The 2\\,$\\sigma$ upper limits obtained by AGILE are consistent with the average flux point observed by the $Fermi$-LAT for this source during 2008 August-Octorber (Abdo et al. 2009). The upper limit obtained in the third time interval was used for the modeling of the SED. The fluxes and corresponding effective photon frequencies of the other telescopes which contribute to this multi-frequency campaign are reported in Table 3.   \\begin{table}[htdp] \\label{table:spec} \\centering \\begin{tabular}{|c|c|c|} \\hline Time interval &  Energy &  U.L.Flux [ph\\,m$^{-2}$s$^{-1}$] \\\\ \\hline   \\multirow{2}{*}{ March 16-21} &  $>$ 100\\,MeV & $5.6 \\times 10^{-3}$ \\\\ & $>$ 200\\,MeV &  $3.6 \\times 10^{-3}$  \\\\ \\hline \\multirow{2}{*}{March 25-30} &  $>$ 100\\,MeV & $5.5 \\times 10^{-3}$ \\\\ & $>$ 200\\,MeV &  $2.8 \\times 10^{-3}$  \\\\ \\hline \\multirow{2}{*}{April 10-30} &   $>$ 100\\,MeV & $3.4 \\times 10^{-3}$ \\\\ & $>$ 200\\,MeV &  $2.1 \\times 10^{-3}$  \\\\ \\hline \\end{tabular} \\caption{2\\,$\\sigma$ Upper limit calculated from the AGILE data in three different time intervals.} \\end{table}        The SED of PG\\,1553+113 is shown in Fig.2. The VHE and HE $\\gamma$-ray flux points are from MAGIC and AGILE respectively. The X-ray point, provided by {\\it R}XTE/ASM, represents the average flux between March 1 and May 31, 2008. The optical R-band point, provided by the KVA telescope, is the average flux obtained on 2008 March 18 and 19. The flux provided by Abastumani is the average flux of the 2008 April 1 - May 17 observations. In addition to these  data we also used the NIR flux from REM. To assess the soundness of this addition, we checked the optical variability of the source during this period using Abastumani data, and found that the source was essentially stable (minimum and maximum values of $log(\\nu F_{\\nu})$ are $-10.14$ and $-10.02$ respectively). For a comparison of the HE flux, we included the flux points from the \\textit{ Fermi} $\\gamma$-ray Space Telescope (Flux, $F(E>100 MeV)= 8\\pm1 \\times 10^{-4} ph\\,m^{-2}\\,s^{-1}$ and photon index, $\\Gamma=1.7\\pm0.6$; Abdo et al. 2009). The average flux (15-30\\,keV) obtained from the X-ray satellite $Swift$/BAT during 39 months (2004 December - 2008 February) of observation (Cusumano et al. 2010) is also included.  \\begin{figure}[h] \\label{fig:ssc} \\centering \\vspace{-1cm} \\epsfig{file=1553_abastumani.ps,height=11cm,width=11cm,clip=} \\vspace{-1cm} \\caption{Average SED of PG\\,1553+113 measured in 2008 March-April. The empty triangles denote the REM data, the open square represents the KVA data, the open circle denotes the Abastumani data, and the open square denotes $R$XTE/ASM data.  The arrow at HE denotes the  AGILE upper limit. The empty squares in the VHE range are the de-absorbed MAGIC data. We also show the non-simultaneous flux points from {\\it Fermi} (bowtie) and $Swift$/BAT (small filled circle).} \\end{figure}  We fitted the resulting simultaneous SED  with a homogeneous one-zone SSC model (Tavecchio et al. 2001). The model assumes that  the source is a spherical plasmon of a radius R, moving with a Doppler factor $\\delta$ towards the observer at an angle $\\theta$  with respect to the line of sight threaded with a uniforming distributed tangled magnetic field of the strength $B$. The injected relativistic particle population is described as a broken power-law spectrum with the normalization $K$, extending from $\\gamma_{\\rm min}$ to $\\gamma_{\\rm max}$ with the indices  $n_{\\rm 1}$ and $n_{\\rm 2}$ below and above the break Lorentz factor $\\gamma_{\\rm br}$. By fitting the observed flux with the model, we obtained the following parameters: $\\gamma_{\\rm min}$ = 1, $\\gamma_{\\rm br} = 3 \\times 10^4$, $\\gamma_{\\rm max} = 2 \\times 10^5$, $K=0.5 \\times 10^4$ cm$^-3$, $n_{\\rm 1}=2$, $n_{\\rm 2}=4.7$, $B=0.7\\,$G, $R=1.3 \\times 10^{16}$ cm, and $\\delta=23$. The optical and X-ray flux constrain on the slope of the electron energy distribution (EED), while the X-ray and VHE spectrum fix the Lorentz factors.  The difference between the current SED and the previous one published in Albert et al. (2007) is due to flux variation in the X-ray and a small variation of the slope of VHE spectrum. We fitted the previous result with the Tavecchio et al. (2001) SSC model to compare the physical parameters of the SED. The difference arises from the EED, but the slopes and $\\gamma_{\\rm max}$ remain constant. The $\\gamma_{\\rm min}$ and $\\gamma_{\\rm br}$ of the previous observation are found to be $3\\times10^3$ and  $2.7 \\times 10^4$ respectively.   \\begin{table}[htdp] \\label{table:spec} \\centering \\begin{tabular}{|c|c|c|} \\hline Instrument & log($\\nu$ [Hz]) &  log($\\nu$F($\\nu$)) [erg cm$^{-2}$ s$^{-1}$] \\\\ \\hline KVA     &    14.63   &  -10.17  \\\\ \\hline Abastumani     &    14.63   &  -10.08  \\\\ \\hline \\multirow{3}{*}{REM} &   14.38 & -10.33 \\\\ &14.27 & -10.34 \\\\ &14.13 & -10.38 \\\\ \\hline XTE & 18.03 & -10.3 \\\\ \\hline \\end{tabular} \\caption{Effective frequencies and corresponding fluxes from PG\\,1553+113 from KVA, Abastumani, REM and RXTE instruments obtained during this campaign.} \\end{table}  During this campaign, no significant variability of the VHE flux is found. The integral flux (E$>$200\\,GeV) during these observations is $1.3 \\pm 0.3 \\times 10 ^{-7}$cm$^{-2}$s$^{-1}$  while during the first observations it was  $1.0 \\pm 0.4 \\times 10^{-7}$cm$^{-2}$s$^{-1}$. The X-ray flux\\footnote{Note that the X-ray data used in Albert et al. (2007) was not taken simultaneously with VHE and optical data.} increases by about a factor of two, while the averaged X-ray flux during 39 months of $Swift$/BAT observations agrees with our SED. The optical flux during our first observation and the current observation does not show any significant variability. The {\\it Fermi} bowtie and lowest-energy MAGIC data points together with the model fit indicate a variability at HE or VHE $\\gamma$-rays.   Our results suggest that the variability of PG\\,1553+113 at different frequencies is time-dependent: hence only a simultaneous multi-frequency monitoring campaign over a large time span will give more information on the source. Relative to this it is worth mentioning that the AGILE and MAGIC data presented here constitute the first simultaneous broad-band $\\gamma$-ray observation (and ensuing SED) of any blazar, though the first simultaneous detection accomplished during the multi-frequency campaign of Mkn\\,421 (Donnarumma et al. 2009), and the first broad-band $\\gamma$-ray spectrum was obtained from PKS\\,2155-304 (Aharonian et al. 2009) by H.E.S.S. and $Fermi$."
    },
    "0911/0911.3044_arXiv.txt": {
        "abstract": "{Open clusters are excellent tracers of the structure, kinematics, and chemical evolution of the disc and a wealth of information can be derived from the spectra of their constituent stars. } {We investigate the nature of the chemical composition of the outer disc open cluster Tombaugh 2. This has been suggested to be a member of the GASS/Mon substructure, and a recent study by Frinchaboy et al. (2008) suggested that this was a unique open cluster in possessing an intrinsic metal abundance dispersion. We aim to investigate such claims. } {High resolution VLT+GIRAFFE spectra in the optical are obtained and analyzed for a number of stars in the Tombaugh 2 field, together with independent UBVI$_{\\rm C}$ photometry. Radial velocities and position in the color-magnitude diagram are used to assess cluster membership.  The spectra, together with input atmospheric parameters and model atmospheres, are used to determine detailed chemical abundances for a variety of elements in 13 members having good spectra. } {We find the mean metallicity to be [Fe/H]$=-0.31 \\pm 0.02$ with no evidence for an intrinsic abundance dispersion, in contrary to the recent results of Frinchaboy et al. (2008). We find Ca and Ba to be slightly enhanced while Ni and Sc are solar. The r-process element Eu was found to be enhanced, giving an average [Eu/Ba]=+0.17. The Li abundance decreases with T$_{\\rm eff}$ on the upper giant branch and maintains a low level for red clump stars. The mean metallicity we derive is in good agreement with that expected from the radial abundance gradient in the disc for a cluster at its Galactocentric distance. } {Tombaugh 2 is found to have abundances  as expected from its Galactocentric distance and no evidence for any intrinsic metallicity dispersion. The surprising result found by Frinchaboy et al. (2008), that is the presence of 2 distinct abundance groups within the cluster, implying either a completely unique open cluster with an intrinsic metallicity spread, or a very unlikely superposition of a cold stellar stream and a very distant open cluster, is not supported by our new result. } ",
        "introduction": "\\begin{figure*} \\centering \\includegraphics[width=\\columnwidth]{f1a.eps} \\includegraphics[width=\\columnwidth]{f1b.eps} \\caption{Left panel: The CMD of Tombaugh 2 in the optical and infrared. Members stars studied in this paper are indicated as filled circles. Right panel: Distribution on the sky of our targets (filled points) and F08 objects (empty squares).} \\label{f1} \\end{figure*}  One of the many regions in our Galaxy for which we lack detailed knowledge is the outer disk. How and when was it formed? Via an outside-in or inside-out process \\citep{chi01}? What is the nature of the metallicity gradient? Does it show a constant slope with Galactocentric distance or is there a leveling out in the outer disk \\citep{twa97,car04,car07,magr09} What if any is the role of mergers? \\\\ A particularly intriguing development in recent years has been the suggestion that there is a merger remnant lying near the Galactic plane. The structure was first identified as the Monoceros stream \\citep[Mon;][]{new02,iba03,yan03}, also known as the Galactic anticenter stellar structure  \\citep[GASS;][]{roc03,cra03} GASS/Mon was discovered as an overdensity of stars near the Galactic plane that seems to wrap around the outer parts of the Galactic disc and is frequently explained as tidal debris from the disruption of a dwarf galaxy on a low inclination orbit \\citep[e.g.,][]{cra03,yan03,pen05}.\\\\ Further arguments have been made that the reputed dwarf galaxy in Canis Major \\citep[CMa;][]{bel04} is the progenitor of the Mon/GASS structure \\citep{martin04}. However, the nature and/or reality of the proposed Mon ``tidal debris stream'' and the CMa overdensity have been called into question and are currently a matter of great debate. Momany \\citeyearpar{mom04,mom06} have argued that much of the observed stellar overdensity associated with Mon --- and particularly all of that associated with CMa (at $l \\sim 240^{\\circ}$) --- is due to the warping and flaring of the Galactic disc, and that no ``extra-Galactic'' component is needed to account for the apparent overdensities in the third quadrant. The presence of `blue plume'' stars in this part of the sky has been used to argue further for the presence of a dwarf galaxy nucleus in CMa \\citep{bel04,martinez05,din05,but07}. However, these young stars have also been more prosaically attributed to the presence of previously unknown features of spiral arm structure \\citep{car05,moi06}.  Open clusters have traditionally played a critical role in the study of the structure, kinematics, formation and chemical evolution of the disk. For this reason alone, the old outer Galactic disk open cluster Tombaugh 2 (To2; at Galactic coordinates [$l,b$] = [232.8,$-$6.9]$^{\\circ}$), located at a Galactocentric distance of $\\sim$15 kpc and 700 pc below the Galactic plane, is worthy of study and was indeed the subject of several previous metallicity investigations. \\citet{bro96} obtained high resolution spectra of 3 stars, finding [Fe/H]$=-0.4\\pm 0.25$, while \\citet{frie02} derived [Fe/H]$-0.44 \\pm 0.09$ from much lower resolution spectra of 12 members. Both of these studies were conducted with the CTIO 4m telescope. Subsequently, \\citet{frin04} suggested that To2 may possibly be associated with Mon/GASS, as well as with CMa \\citep{bel04}. Note that its Galactic longitude is very close to that of the CMa overdensity. This provided a  strong additional incentive  for a followup study of its chemistry and kinematics in greater detail with a larger telescope, as even the giants are faint due to its distance of over 13kpc. \\citet[][ - hereafter F08]{frin08} used this motivation to obtain VLT/FLAMES spectra (both with UVES at high resolution and with GIRAFFE at somewhat lower resolution). They were able to derive velocities and detailed abundances of a number of elements for 18 To2 cluster members. Their most surprising result was an apparent large spread in metallicity : $\\Delta$[Fe/H] $> 0.2$. They were unable to account for this spread given their observational errors and presented a number of possible scenarios, including the likelihood that To2 possessed an intrinsic metallicity spread. They argued for the possible presence of 2 populations in the cluster, distinguished by their different mean chemical characteristics --- with a metal-rich, Ti-normal group and a metal-poor, Ti-enhanced group, namely ($\\langle$[Fe/H]$\\rangle$, $\\langle$[Ti/Fe]$\\rangle$) $\\sim$ ($-$0.06, $+$0.02) and ($-$0.28, $+$0.36). The more metal-poor group appeared more centrally concentrated, and they suggested that this group represented the true To2 clusters stars and the metal-rich population was   an overlapping, and kinematically associated, but ``cold''  stellar stream.  If true, this would be the first such metallicity spread uncovered in an open cluster, and/or the first evidence of multiple populations in a cluster of this age (about 2 Gyr) in the Galaxy, making To2 an extremely intriguing object. Although multiple populations within globular clusters are now known to exist, these clusters are all much more massive than an open cluster such as To2, as expected if this phenomenon is due to a cluster being massive enough to retain ejecta from a first generation of stars in order to make a second, chemically enriched generation.      To date, the F08 study is the only one which suggests multiple populations in an open cluster. However, as F08 pointed out, further observations are desperately required to corroborate their surprising results and make a more definitive investigation.  Given its above history, a new metallicity study of To2 is imperative. These points motivated the present study, wherein we wished to obtain independent high-resolution spectra with a large telescope of a number of To2 stars in order to investigate the reality of the putative metallicity dispersion and the nature of its chemical composition. In 2, we describe our observations and reduction procedures. In 3, we discuss the determination of radial velocities and membership for our observed sample. In 4, we present the details of our abundance analysis, including the derivation of the input atmospheric parameters and of the internal errors. In 5, our basic abundance results are presented, and in 6 we compare these to previous investigations, especially that of F08 in light of the above. Finally, we summarize and emphasize the importance of our main results in 7.  \\begin{table*} \\caption{Target stars and parameters} \\label{t1} \\centering \\begin{tabular}{llccccccccccccc} \\hline \\hline ID1 & ID2 & $\\alpha$($^0$) & $\\delta$($^0$) & V & U-B & B-V & V-I$_{\\rm C}$ & J$_{\\rm 2MASS}$ &\\ H$_{\\rm 2MASS}$ & K$_{\\rm 2MASS}$ & T$_{\\rm eff}$ & log g & v$_{\\rm t}$ & RV$_{\\rm H}$\\\\ \\hline 1512 & 2138 & \\tiny{105.744536} & \\tiny{-20.848913} & 16.857 & 1.122 & 1.206 &  1.294 & 14.870 & 14.278 & 14.164 & 5210 & 3.29 & 1.10 & 117.6\\\\ 1672 & 1100 & \\tiny{105.724763} & \\tiny{-20.836042} & 15.726 & 0.818 & 1.231 &  1.319 & 13.511 & 12.850 & 12.666 & 4820 & 2.63 & 1.26 & 119.1\\\\ 1827 &  837 & \\tiny{105.748951} & \\tiny{-20.825960} & 15.773 & 0.877 & 1.233 &  1.347 & 13.521 & 12.865 & 12.708 & 4900 & 2.70 & 1.24 & 120.9\\\\ 1886 & 1532 & \\tiny{105.730038} & \\tiny{-20.822460} & 15.711 & 0.638 & 1.116 &  1.251 & 13.675 & 13.024 & 12.876 & 5010 & 2.72 & 1.24 & 120.0\\\\ 2184 &  485 & \\tiny{105.793143} & \\tiny{-20.799812} & 15.857 & 0.702 & 1.118 &  1.262 & 13.755 & 13.214 & 13.055 & 4980 & 2.77 & 1.23 & 119.8\\\\ 236 & 2895 & \\tiny{105.864625} & \\tiny{-20.854195} & 15.919 & 0.669 & 1.141 &  1.281 & 13.804 & 13.234 & 13.031 & 5050 & 2.83 & 1.21 & 121.0\\\\ 238 & 1802 & \\tiny{105.731226} & \\tiny{-20.854284} & 15.699 & 0.889 & 1.258 &  1.341 & 13.431 & 12.771 & 12.601 & 4780 & 2.60 & 1.27 & 120.5\\\\ 2846 & 2430 & \\tiny{105.742321} & \\tiny{-20.734119} & 16.017 & 0.547 & 1.166 &  1.187 & 14.072 & 13.526 & 13.368 & 5160 & 2.92 & 1.19 & 122.9\\\\ 299 & 2902 & \\tiny{105.775540} & \\tiny{-20.841908} & 15.523 & 1.002 & 1.327 &  1.426 & 13.107 & 12.397 & 12.277 & 4780 & 2.53 & 1.29 & 121.8\\\\ 3574 & 2074 & \\tiny{105.766791} & \\tiny{-20.818871} & 16.138 & 0.678 & 1.161 &  1.321 & 13.919 & 13.311 & 13.200 & 5150 & 2.96 & 1.18 & 121.7\\\\ 3763 & 2894 & \\tiny{105.761083} & \\tiny{-20.806566} & 16.051 & 0.783 & 1.041 &  1.288 & 13.883 & 13.367 & 13.175 & 5000 & 2.86 & 1.20 & 120.1\\\\ 3836 &  494 & \\tiny{105.769301} & \\tiny{-20.803593} & 16.501 & 0.704 & 1.133 &  1.382 & 14.357 & 13.750 & 13.617 & 5110 & 3.09 & 1.15 & 120.9\\\\ 591 & 2442 & \\tiny{105.835422} & \\tiny{-20.779896} & 15.944 & 0.706 & 1.293 &  1.355 & 13.690 & 13.018 & 12.909 & 5020 & 2.82 & 1.21 & 125.1\\\\ \\hline \\end{tabular} \\end{table*} ",
        "conclusions": "The existence of multiple populations in Galactic globular clusters and some Large Magellanic Cloud clusters is now well established \\citep{pio08,marino08,marino09,mil09}. Indeed, $\\omega$ Centauri is now known to possess at least 5 populations showing different ages and abundances \\citep{vil07}. However, certainly not all such clusters show this phenomenon, at least not in a clear way. Indeed, the dominant characteristic that seems to be in common among such clusters is their mass - only the most massive clusters are involved. Given the current paradigm concerning the origin of these multiple populations - retention of gas left over and polluted by a first generation of star formation to form a second, chemically distinct population - obviously requires sufficient cluster mass to retain the required SN and AGB or massive star wind ejecta. The required mass is estimated to be some 10$^5$M$_\\odot$ \\citep{mie08}. Thus, only the most massive Galactic globulars fulfill this requirement (although note that at least one less massive cluster, M4 at 10$^4$ M$_\\odot$, also shows multiple populations, \\citealt{marino08}). The case of the LMC clusters is more uncertain, but it is clear that the phenomenon is correlated with mass. On the other hand, Galactic open clusters generally do not exceed masses of $10^3 - 10^4M_\\odot$ so that the possibility of retaining any first generation ejecta is very unlikely - any gas that remains after the first star formation episode is subsequently quickly removed by the ejecta itself. Thus, one does not expect to find multiple populations in open clusters. To2 is a typical open cluster in this respect and certainly does not exceed a few $\\times 10^3M_\\odot$.  The finding by F08 that To2 possessed an intrinsic metallicity spread (or was a superposition of an outer disc cluster well out of the plane and a cold stellar stream with exactly the same velocity but distinct chemistry) was then understandably met with some incredulity and at the very least required independent corroboration. Hence the present study.  We obtained independent data but using the same telescope and spectrograph as they did. However, we deliberately achieved superior data quality - in both S/N and by observing in a longer wavelength regime where line crowding was significantly reduced. Our reduction and analysis procedures were virtually identical to theirs.  To summarize, we obtained high resolution VLT+GIRAFFE spectra in the optical for a number of stars in the Tombaugh 2 field. Radial velocities and position in the color-magnitude diagram are used to assess cluster membership. The spectra, together with input atmospheric parameters and model atmospheres, are used to determine detailed chemical abundances for a variety of elements in 13 stars confirmed as members.  We derive a mean metallicity  [Fe/H]$=-0.31 \\pm 0.02$, with no evidence for an intrinsic abundance dispersion. We find Ca, Ba, and Eu to be enhanced while Ni and Sc are solar. Li abundances decrease with T$_{\\rm eff}$ on the upper giant branch and maintain a low level for red clump stars. The mean metallicity we derive is in good agreement with that expected from the radial abundance gradient in the disc for a cluster at its Galactocentric distance. The surprising result found by F08, viz. for 2 distinct abundance groups within the cluster, implying either a completely unique open cluster with an intrinsic metallicity spread, or a very unlikely superposition of a cold stellar stream and a very distant open cluster, is not supported by our new data. We suspect that the F08 data was subject to substantially larger errors than they estimated, especially given the low S/N of their UVES spectra and their much bluer wavelength range, which was plagued by line crowding. To2, instead of being a unique cluster, is found to be a normal representative of its class."
    },
    "0911/0911.1794_arXiv.txt": {
        "abstract": "While it has long been known that a large number of short-lived transient spirals can cause stellar migration, here we report that another mechanism is also effective at mixing disks of barred galaxies. The resonance overlap of the bar and spiral structure induces a nonlinear response leading to a strong redistribution of angular momentum in the disk. We find that, depending on the amplitudes of the perturbers, the changes in angular momentum, $\\Delta L$, could occur up to an order of magnitude faster than in the case of recurrent spirals. The signature of this mechanism is a bimodality in $\\Delta L$ with maxima near the bar's corotation and its outer Lindblad resonance; this is independent of the properties of the spiral structure. For parameters consistent with the Milky Way the disk mixes in about 3~Gyr and the stellar velocity dispersion increases with time in a manner roughly consistent with observations. This new mechanism could account for both the observed age-velocity relation and the absence of age-metallicity relation in the solar neighborhood. Spiral-bar interaction could also explain observations showing that strongly barred galaxies have weaker metallicity gradients than weakly barred or non-barred galaxies. ",
        "introduction": "The metallicity of the interstellar medium (ISM) in a galaxy is expected to increase with time due to the progressive enrichment through stellar feedback. Consequently, one would expect that younger stars formed at the same galactic radius would have higher metallicities. In addition, the ISM metallicity has been well established to decrease with increasing galactic radius, $r$, due to a more efficient star formation and enrichment of the ISM in the central regions of galaxies \\citep{wilson94,daflon04}. This in turn would result in coeval stars being progressively more metal-poor as $r$ increases.  However, the metallicities of stars in the solar neighborhood (SN) have been observed to lack the expected correlation with age \\citep{edvardson93,haywood08}. Re-examining the age-metallicity distribution of the Geneva-Copenhagen survey \\citep{nordstrom04}, \\cite{haywood08} found that the data are only consistent with the age-metallicity relation (AMR) for stars younger than 3~Gyr. For other samples the width of the metallicity distribution was found to be larger than what could be accounted for by measurement errors and/or inhomogeneous ISM. Because of the age-velocity relation (AVR), where older stellar samples are observed to have higher velocity dispersion than younger ones (e.g. \\citealt{holmberg09}), stars from the inner and outer Milky Way (MW) disk can enter the SN and thus blur the expected AMR. However, this can only account for about 50 \\% of the observed scatter \\citep{nordstrom04}.  A possible explanation for this discrepancy is the radial migration of stars from their birth guiding radii \\citep{haywood08,schonrich09}. Indeed, there is now considerable evidence that the non-axisymmetry of the Galactic potential can cause significant perturbations in the motion of both stars and gas: observational evidence for this comes from (i) the observed non-circular motions of gas flows in the inner MW \\citep{bissantz03}, (ii) the large non-axisymmetric motions of star-forming regions \\citep{xu06}, (iii)  the moving groups in the SN \\citep{dehnen98} containing stars of very different ages \\citep{famaey05,antoja08}, suggesting that their clumping in velocity space is most likely not due to irregular star formation, but rather to dynamical perturbations from the bar \\citep{dehnen00,minchev10} and/or the spirals (e.g., \\citealt{qm05}, \\citealt{antoja09}), or an orbiting satellite galaxy \\citep{quillen09,minchev09}.  A numerically well-studied effect of this non-axisymmetry of the Galactic potential is the radial migration of stars due to resonant scattering by transient spiral structure (SS) (\\citealt{sellwood02}, hereafter SB02; \\citealt{roskar08}), which could explain the aforementioned absence of AMR in the SN. N-body simulations based on cosmological initial conditions (e.g., \\citealt{heller07}) would a priori be the most natural framework in which to study radial mixing in galactic disks. However, these simulations cannot at present produce large enough thin galactic disks due to the transfer of angular momentum from the disk to the live dark matter halo. The resulting galaxies are typically thick and present large stellar velocity dispersions, so that they are closer to S0 galaxies than to MW-like ones \\citep{sanchez09}. On the other hand, \\cite{roskar08} used an N-body hydrodynamical model in which the initial conditions were designed to mimic the stage following a last major merger when the thin disk formation is supposed to start. Both SB02 and \\cite{roskar08} compare the effects of radial migration in the MW disk, for which the existence of a bar is firmly established, to a simulated galaxy lacking a central bar. Unless the effect of the bar is unimportant to the mixing process, their results would only poorly approximate the radial migration in our Galaxy. In this paper we check whether the effect of galactic bars can indeed be neglected in the context of radial migration in galactic disks, as assumed in a number of previous studies. For this purpose we study the combined effect of SS and a central bar (and more generally of multiple patterns) on the radial mixing of stars in a stellar disk.  There is observational evidence for multiple patterns in external galaxies, such as asymmetries in the spiral structure (e.g. \\citealt{henry03,meidt08}). By expanding galaxy images in Fourier components, \\cite{elmegreen92} noted that many galaxies exhibit hidden three-armed components and suggested that multiple spiral density waves can propagate simultaneously in galaxy discs. In addition, by means of the recently developed \\citep{merrifield06} Radial Tremaine-Weinberg (TWR) method, \\cite{meidt09} analyzed the high-quality HI and CO data cubes available for four spiral galaxies. The authors found direct evidence for the presence of resonant coupling of multiple patterns, including spiral-spiral and spiral-bar components.  The results of N-body simulations provide additional motivation for studying the dynamical effects of multiple patterns. \\cite{sellwood85} and  \\cite{sellwood88} presented N-body experiments in which a bar coexists with a spiral pattern moving at a much lower angular velocity. \\cite{tagger87} and \\cite{sygnet88} explained this as a non-linear mode coupling between a bar and SS. Sellwood's simulations, and the theoretical explanation of \\cite{tagger87} and \\cite{sygnet88}, were later confirmed by \\cite{masset97} and \\cite{rautiainen99}.   How do multiple patterns affect the dynamics of galactic disks? \\cite{quillen03} considered the dynamics of stars that are affected by perturbations from both SS and a central bar by constructing a one-dimensional Hamiltonian model for the strongest resonances in the epicyclic action-angle variables. Quillen pointed out that when two perturbers with different pattern speeds are present in the disc, the stellar dynamics can be stochastic, particularly near resonances associated with one of the patterns. Similar findings were presented more recently by \\cite{jalali08}. All these results are not surprising since it has already been shown by \\cite{chirikov79} that in the case of resonance overlap the last KAM surface between the two resonances is destroyed, resulting in stochastic behavior. We therefore expect that resonance overlap could give rise to both velocity dispersion increase and radial migration.   In this paper we study the combined effect of a central bar and SS on the dynamics of a galactic disk. We concentrate on a steady-state SS in order to assess the pure effect of non-linear coupling between the two perturbers, noting that while there is strong circumstantial evidence that SS is mostly not-stationary in unbarred galaxies, the situation is less clear in barred ones \\citep{bt08}. In any case, the following model can also serve as a proxy for studying the effect of transients, and is complementary to the transient spiral mechanism developed by SB02.  \\begin{figure}[t!] \\epsscale{1.1} \\plotone{fig1.eps} \\figcaption{ Resonances in a galactic disk for nearly circular orbits and a flat rotation curve. Corotation occurs along the dotted black curve and is given by $\\Omega_{b,s}=\\Omega(r)$, where $\\Omega_{b,s}$ is the bar or spiral pattern speed, $\\Omega(r)=v_0/r$ is the local circular frequency, and $v_0$ is the constant circular velocity. The 2:1 outer and inner Lindblad resonances (OLR and ILR) occur along the solid black curves, computed as $\\Omega_{b,s}=\\Omega(r)\\pm\\kappa/2$, where $\\kappa$ is the local radial epicyclic frequency. The outer and inner 4:1~LRs occur along the dashed black curves, given by $\\Omega_{b,s}=\\Omega(r)\\pm\\kappa/4$. The red horizontal line indicates the bar pattern speed used in this paper, $\\Omega_b=55.5$~km/s/kpc, and the blue horizontal line shows a spiral pattern speed of $\\Omega_s=24$~km/s/kpc, such as the one used in the simulations shown in Figs.~\\ref{fig:tevol} and \\ref{fig:eps}. The vertical red and blue lines give the radial positions of each resonance for the bar and spiral structure, respectively (solid lines: 2:1; dashed lines: 4:1, dotted lines: CR). \\label{fig:om_r} } \\end{figure} ",
        "conclusions": "In this work we have examined the combined effect of a central bar and spiral structure on the radial migration in galactic disks. We have found that:  \\noindent (i) This new mechanism for galactic disk mixing could be up to an order of magnitude more effective than the transient spirals one (see Sec.~\\ref{sec:strength}). We note that the effect is non-linear, strongly dependent on the strengths of the perturbers. The non-linearity of the mechanism is confirmed by the fact that an increase in the amplitudes yields a much larger increase in the angular momentum changes in a simulation including {\\it both} bar and spiral structure, compared to the case of a single perturber. In other words, the individual effects do not add up linearly. Another indication of non-linearity is that the effect increases with time only in the case of the two perturbers propagating simultaneously.  \\noindent (ii) The signature of this mechanism is a bimodality in the changes of angular momentum in the disk with maxima near the bar's corotation and its outer Lindblad resonance (see Figs.~\\ref{fig:om} and \\ref{fig:rhist}). This is true regardless of the spiral pattern speed and can be used in identifying the spiral-bar resonance overlap migration in N-body simulations, where the disk dynamics is much more complicated.  \\noindent (iii) For bar and spiral parameters consistent with Milky Way observations we find that it takes $\\sim$~3~Gyr to achieve the mixing for which transients require 9~Gyr. In addition to radial mixing, spiral-bar coupling can account for the age-velocity relation (AVR) observed in the solar neighborhood (Fig.~\\ref{fig:sig}, Sec.~\\ref{sec:sig}). Note however that the estimates for the Milky Way spiral and bar strengths derived from observations {\\it now} may simply be irrelevant for our understanding of the current state of mixing in the Galactic disk (see Sec.~\\ref{sec:mw}).   \\noindent (iv) This mechanism could explain observations showing that strongly barred galaxies have weaker metallicity gradients than weakly barred or non-barred galaxies (e.g., \\citealt{zaritsky94,martin94,perez09}).    The effect of resonance overlap we have described in this paper appears to be quite strong. In order for it to work, a galactic disk must contain a central bar\\footnote{Note that resonance overlap, and thus migration, will also result from multiple spiral structure \\citep{mig2}.}. Unlike in our test-particle simulations, where the disk responds solely to the bar potential, as seen in N-body simulations and deduced from observations, a bar is expected to always drive spiral structure (e.g., \\citealt{salo10}). Thus, any barred galaxy would be affected strongly by this mechanism. As the disk heats and SS diminishes in strength, the effect on $\\Delta L$ will decrease. It is therefore likely that resonance overlap migration will be most vigorous at certain periods during a galaxy lifetime.  In this work we have only considered steady-state spirals in order to assess the pure effect of the coupling between bar and spiral structure. Considering transient spirals would increase the efficiency, since the SB02 mechanism would come into play as well. Therefore, a combination of a central bar and short-lived transients is expected to be extremely efficient at mixing barred spiral galaxies.  As discussed in Sec.~\\ref{sec:mw}, the most promising technique to put constraints on this mechanism in the MW is \"chemical tagging\" \\citep{freeman02, joss10} which will become possible with the forthcoming spectroscopic survey HERMES \\citep{freeman08,freeman10}. More ongoing and planned large Galactic surveys, such as RAVE, SEGUE, SIM Lite, GYES, LAMOST and GAIA, can search for signatures of the mechanism (see also \\citealt{mq08}).  In a follow-up paper \\citep{mig2} we study the effect of resonance overlap in fully self-consistent, TreeSPH/N-body simulations including a gaseous component and star formation. The migration mechanism described in this paper can provide an explanation for the observed flatness and spread in the age-metallicity relation. Our results, however, need to be incorporated into a chemo-dynamical galactic model to properly assess the effect on the AMR, and possibly also on the formation of the thick disk, similarly to the work by \\cite{schonrich09}."
    },
    "0911/0911.3272_arXiv.txt": {
        "abstract": "Giant pulses observed in a number of pulsars show a record brightness temperature, which corresponds to the high energy density of $U \\approx 10^{15}$erg/cm$^{3}$. Comparable densities of energy in the radio-frequency region are attainable in a cavity-resonator being the pulsar internal vacuum gap. Emission of energy through the breaks accidentally appearing in the magnetosphere of open field lines corresponds to the giant pulses. The emitted energy is determined by the break area, which causes a power dependence of break occurrence probability. The observed localization of the giant pulses relative to the average pulse is explained by the radiation through a waveguide near the magnetic axis or through a slot on the border of the open field lines. Separate discharges may be superimposed on the radiation through the breaks, thus forming the fine structure of giant pulses with duration up to some nanoseconds. Coulomb repulsion of particles in the puncture spark in the gap leads to spark rotation around its axis in the crossed fields, which provokes the appearance of observed circular polarization of the giant pulses. The correlation between the giant pulse phase and the phase of the hard pulsar emission (X-ray and gamma) is naturally explained as well.  Thus, a wide range of events observed at the giant pulses can be explained from the viewpoint that the internal vacuum gap is a cavity-resonator stimulated by discharges and radiating through the breaks in the magnetosphere. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0911/0911.0311.txt": {
        "abstract": "By constructing a global model based on 3D local magnetohydrodynamical (MHD) simulations, we show that the disk wind driven by magnetorotational instability (MRI) plays a significant role in the dispersal of the gas component of proto-planetary disks. Because the mass loss time scale by the MRI-driven disk winds is proportional to the local Keplerian rotation period, a gas disk dynamically evaporates from the inner region with possibly creating a gradually expanding inner hole, while a sizable amount of the gas remains in the outer region. The disk wind is highly time-dependent with quasi-periodicity of several times Keplerian rotation period at each radius, which will be observed as time-variability of protostar-protoplanetary disk systems. %should be observable as a variability due to the motion of the effective %surface of the disk. These features persistently hold even if a dead zone exists because the disk winds are driven from the surface regions where ionizing cosmic rays and high energy photons can penetrate. %.ionizing cosmic rays and high energy photons can penetrate. Moreover, the predicted inside-out clearing significantly suppresses the infall of boulders to a central star and the Type I migration of proto-planets %the solid component, %inward drift of boulders and migration of planets are which are favorable for the formation and survival of planets. ",
        "introduction": "%\\section*{} %{\\color{red} Planets are believed to form in protoplanetary disks consisting of gas and dust components around newly born stars. %} The evolution of the gas component is crucial in determining the final state of the planetary system, such as the number and locations of terrestrial (rocky) and Jovian (gas-giant) planets \\citep{il04}. %(e.g., Ida \\& Lin 2004). %If the gas dissipate quickly, Jovian (gas-giant) %planets cannot form, while if it remains long, newly formed planets %infall to a central star by the interaction with the gas. %Recent observations show that the typical lifetime of protoplanetary disks %is $10^6 - 10^7$ yr. The amount of the gas that is captured by Jovian planets %in various planetary systems is generally much smaller than the total gas of the disk. Therefore, the gas component should dissipate via other mechanisms %besides the accretion to planets with the observationally inferred timescale of $10^6 - 10^7$ yr \\citep[e.g.,][] {hll01,her08}. %, which is inferred from the observations of various star clusters %(Haisch et al.2001). The currently favored scenario is that the gas dissipates by the combination of photoevaporation by Ultraviolet (UV) flux from a central star and viscous accretion \\citep[e.g.][] {shu93,mat03,tak05,alx06,cm07}. % (e.g. Shu et al. 1993; Matsuyama et al.2003; Alexander et al.2006). %to a central star %Recently, the photoevaporation by UV radiation from the chromosphere of a %central star and the accretion shock is widely discussed as a possible %mechanism of the dispersal of the outer gas disk ($\\gtrsim 10$ AU for the %protosolar disk), while the inner gas disk dissipates by the accretion to %a central star (Shu et al. 1993; Matsuyama et al.2003; Alexander et al.2006). However, the time-evolution of the luminosity and spectrum of the UV radiation is quite uncertain, and moreover some observed transitional disks with inner holes are inconsistent with the photoevaporation mechanism \\citep{cal05,esp08,hug09}. %(Calvet et al.2005; Espaillat et al.2008; Hughes et al.2009).   By performing local 3D ideal MHD simulations in the shearing box approximation \\citep{hgb95,mt95}, \\citet[][SI09 hereafter]{si09} have shown that MHD turbulence in protoplanetary disks effectively drives disk winds. %{\\color{red} %SI09 have further shown that The local timescale of the dynamical evaporation by magnetorotational instability (MRI) driven disk winds is defined as \\begin{equation} \\tau_{\\rm ev} =\\Sigma/(\\rho v_z)_{\\rm w}, \\label{eq:evtm} \\end{equation} where $(\\rho v_z)_{\\rm w}$ is the sum of the mass fluxes from the upper and lower disk surfaces and $\\Sigma$ is the surface density. We can estimate $\\tau_{\\rm ev} \\sim$ several thousand years at 1 AU of the typical protosolar disk \\citep[e.g., minimum mass solar nebula, or MMSN in short hereafter;][]{hay81};  here we use the typical values, $\\Sigma = 1700$ g cm$^{-2}$ and $(\\rho v_z)_{\\rm w}\\sim 10^{-8}$ g cm$^{-2}$s$^{-1}$ (see \\S\\ref{sec:ndzls} for the scaling of the disk wind flux). While the actual dispersal time is much longer as shown in this paper because the radial mass accretion is not taken into account here, this estimate shows that the MRI-driven disk wind should play a significant role in the dispersal of protoplanetary disks. In this paper, we investigate the evolution of gas density with disk winds and radial mass accretion in a global model. %}  In SI09, we did not take into account the effects of a so-called dead zone which is inactive with respect to MRI because of the insufficient ionization for the coupling between gas and magnetic fields. The dead zone is believed to form around the midplane in the inner parts of protoplanetary disks because ionizing cosmic rays and X-rays cannot penetrate from the disk surfaces (Sano et al.2000). %\\footnote{{\\color{red}There is still a debate on the %formation of dead zones mainly because it is not straightforward to accurately %determine the ionization degree. Besides the ionization by cosmic rays and %X-rays discussed here, \\citet{is05} introduced a self-sustained mechanism %that keeps the enough ionization degree in MRI turbulence by current carrying %electrons.}} In this paper, we improve the models of SI09 by investigating the effects of dead zones by performing resistive MHD simulations. %we also investigate the effects of dead zones by resistive MHD simulations. % in addition to the model calculations for no dead zones. %We will show that the MRI-driven disk winds are blown out intermittently from %a protoplanetary disk with a dead zone and they are effective in the %dispersal of the gas disk because the winds are driven from the regions near %the surfaces with sufficient ionization degree.  \\begin{figure*} \\figurenum{1} \\epsscale{0.9} \\begin{center} \\plotone{ttmsev5.eps} %\\plotone{f1.eps} \\end{center} \\caption{%{\\color{red} Dynamical evaporation of a protoplanetary gas disk by local 3D ideal MHD simulation without mass supply by radial accretion. %} We impose weak vertical magnetic field, $B_z$, with the plasma $\\beta$ value, $\\beta_{z,{\\rm mid}}=8\\pi \\rho_{\\rm mid} c_s^2/B_z^2 = 10^6$, at the midplane. The lower right panel shows the total mass normalized by the initial mass in the local simulation box as a function of rotation time. %Since we do not take into account mass accretion, the decrease of the mass %is slower in the realistic situation. The four color panels show the snapshots of the local protoplanetary disk simulation at $t=0$ (initial condition), 200 rotations, 2200 rotations, and 4130 rotations. The $x$, $y$, and $z$ components respectively correspond to radial, azimuthal, and vertical components. The unit of each component is scale height, $H\\equiv \\sqrt{2}c_s/\\Omega$. The box size is $(x,y,z)=(\\pm 0.5 H, \\pm 2H, \\pm 4H)$, which is resolved by (32,64,256) grid points. The colors indicate density %in logarithmic scale normalized by the initial value at the midplane, the white solid lines illustrate magnetic field lines, and the arrows show velocity field. A small number of magnetic field lines (vertical lines) in the panel of $t=0$ reflect that the initially imposed magnetic field is weak (the number of field lines is scaled by magnetic field strength). We should emphasize that the actual dispersal is much slower because of the mass supply by accretion (see text).} \\label{fig:ttmsev} \\end{figure*}   The construction of the paper is as follows : We firstly describe the MHD simulations in local shearing boxes with and without dead zones in \\S \\ref{sec:lsb}. In \\S \\ref{sec:mdl} we describe our global model. In \\S \\ref{sec:res}, we show how the MRI disk winds disperse the gas component of protoplanetary disks from the inner part. Also, we show its effects on the planet formation. In \\S \\ref{sec:dis}, we discuss the properties of disk evolution mainly focusing on the escape of the disk winds from the gravity of a central star.  % ",
        "conclusions": "\\label{sec:dis}  %\\subsection{Comparison with Photoevaporation Mechanism}  %\\subsection{Escape of Disk Winds} %When doing so we should take a special care first because %the local simulations cannot directly capture the radial mass inflow and %second because the velocities of the disk winds at the upper and lower %simulation boundaries are slower than the escape speed. % %Therefore, it is not trivial whether the disk winds %further flow out and escape from the disk in the actual situations. %In Paper II we extensively discuss the escape of %the disk winds firstly from the energetics of the liberated gravitational %energy by accretion %%an energetics point of view %and secondly by using the local simulations with larger vertical simulation %boxes. %The energetics consideration shows that the gravitational energy released %by the accretion is sufficient for the escape of the disk winds except %in the very inner region, $< 1$ AU. In the simulations with larger %vertical box sizes, the disk winds are continuously accelerated %with height. The wind mass flux becomes moderately smaller but is bounded %by the floor value, $C_{\\rm w}$. We can conclude that the wind mass flux, %$C_{\\rm w}=2\\times 10^{-5}$, adopted in this paper gives a reasonable %estimate. %%Summarizing the discussions in Paper II, we can conclude that at least %%the wind materials that corresponds to the adopted mass flux parameter, %%$C_{\\rm w} =2\\times 10^{-5}$, can escape from the disk.  %\\subsection{Escape of Disk Winds} %\\label{sec:esc} %The method adopted in the paper is to consider the evolution of global %protoplanetary disks by using the results of the local 3D MHD simulations. %The local simulation box extends up to $\\pm 4 H$ in the vertical direction, %and the velocity of the disk winds is an order of the sound speed %\\citep{si0}. %Considering a typical protoplanetary disk such as MMSN, %this velocity is much smaller than the escape velocity. %For example, at 1 AU of the MMSN the sound velocity is $\\approx 1$ km s$^{-1}$, %while the escape velocity, $v_{\\rm esc}=\\sqrt{2}r\\Omega$, is 42 km s$^{-1}$. %It is not a trivial matter whether the disk winds are further accelerated and %escape from the disk. %In this section, we discuss the escape of the disk winds first from %the energetics of accretion disks and second by using the local simulations %with larger vertical boxes.  \\subsection{Energetics} \\label{sec:eng} In this paper we have applied results of the local simulations to the global models. We use the mass flux of the disk winds at the upper and lower boundaries, $z=\\pm 4H$, of the simulation box. Since the wind velocities at the boundaries are still smaller than the escape speed from a central star, we should carefully examine whether the disk winds really escape from the disks. In this subsection, we examine the energetics of the disk winds to see whether the accretion energy can potentially drive the disk winds that can escape from the gravity of a central star. %investigate the escape of the disk winds from an energetics point of view. %The energy source of the disk winds is the gravitational energy liberated %by mass accretion.  The starting point here is an equation for the conservation of total energy : %with the only $r$ derivative is $$ \\hspace{-2cm} \\frac{\\partial}{\\partial t}\\left[\\rho\\left(\\frac{1}{2}v^2 + \\frac{1}{\\gamma-1}\\frac{p}{\\rho} - \\frac{G M_{\\star}}{r}\\right) + \\frac{B^2}{8\\pi}\\right] $$ $$ \\hspace{-0.5cm} %\\eqnum{S19} + %\\frac{1}{r}\\frac{\\partial}{\\partial r}\\left[r \\mbf{\\nabla\\cdot} \\left[\\left\\{\\rho \\mbf{v}\\left(\\frac{1}{2}v^2 + \\frac{\\gamma}{\\gamma-1}\\frac{p}{\\rho} -\\frac{G M_{\\star}}{r}\\right) \\right.\\right. %$$ %\\begin{equation} $$ \\begin{equation} \\left.\\left. + \\mbf{v} \\frac{B^2}{4\\pi} - \\frac{\\mbf{B}}{4\\pi}(\\mbf{B}\\cdot\\mbf{v})\\right\\}\\right] = - q_{\\rm loss}, \\label{eq:eng3d} \\end{equation} where $q_{\\rm loss}$ is energy loss which is modeled below and here we consider the only $r$ derivative. % among the spatial derivatives of the three components and We integrate Equation (\\ref{eq:eng3d}) with $\\int dz$ %On the integration, we by neglecting the small terms except the anisotropic stress, $\\delta v_{\\phi} v_r - B_{\\phi}B_r/4\\pi\\rho = \\alpha c_s^2$, whereas we separate $v_{\\phi}$ into Keplerian rotation plus small perturbation, $v_{\\phi}= r\\Omega + \\delta v_{\\phi}$. We also assume that accretion velocity and sound speed ($c_s=\\sqrt{\\gamma p/\\rho}$) are smaller than rotation velocity, %the gas pressure is much smaller than the gravity of a central star, $v_r, c_s\\ll r\\Omega$. Then, the total energy of a ring at $r$ changes according to \\begin{equation} %\\eqnum{S20} \\frac{\\partial}{\\partial t}\\left[-\\Sigma\\frac{r^2\\Omega^2}{2} \\right] + \\frac{1}{r}\\frac{\\partial}{\\partial r}\\left[r\\Omega \\frac{\\partial}{\\partial r}(r^2 \\Sigma \\alpha c_s^2) +r^2\\Sigma \\Omega \\alpha c_s^2\\right] = - Q_{\\rm loss} , \\label{eq:tteng} \\end{equation} where $Q_{\\rm loss} = \\int q_{\\rm loss} dz$ and we have used $\\frac {G M_{\\star}}{r^2} = r \\Omega^2$ and $r\\Sigma v_r = \\frac{2}{r\\Omega} \\frac{\\partial}{\\partial r}(r^2 \\Sigma \\alpha c_{s}^2)$. % and Equation (\\ref{eq:macev}). % and $r\\Sigma v_r=-\\frac{2}{r\\Omega} \\frac{\\partial}{\\partial r} %(r^2\\Sigma \\alpha c_s^2)$; %we wrote $v_r$ for simplicity instead of the mass-weighted average, %$\\int v_r \\rho dz / \\Sigma$. The spatial derivative term on the left-hand side represents the energy liberated by mass accretion and viscous heating.    We investigate whether this liberated gravitational energy is sufficient to drive disk winds. For simplicity, we only consider the kinetic energy of disk winds for the loss term and neglect additional effects such as heating by UV/X-ray radiation (energy input), acceleration by stellar winds (momentum input), %from a central star  and and radiative cooling (energy loss). %The contribution from disk winds can be included in $Q_{\\rm loss}$. Since we consider the disk winds from a Keplerian rotating disk, we can write $Q_{\\rm loss}=\\frac{1}{2}\\rho v_{z}(v_z^2 + r^2 \\Omega^2)$. %We consider the disk wind with vertical velocity, $v_z$, from a Keplerian %rotating disk. Then, the kinetic energy flux of the disk wind is %$\\frac{1}{2}\\rho v_{z}(v_z^2 + r^2 \\Omega^2)$, which is substituted into %$-Q_{\\rm loss}$ in Equation (\\ref{eq:tteng}). Disk winds can escape to infinity if $v_z$ exceeds the escape speed, $v_{\\rm esc} = \\sqrt{2}r\\Omega$; namely if the condition, $$ %\\begin{equation} %\\eqnum{S21} \\frac{\\partial}{\\partial t}\\left[\\Sigma\\frac{r^2\\Omega^2}{2} \\right] - \\frac{1}{r}\\frac{\\partial}{\\partial r}\\left[r\\Omega \\frac{\\partial}{\\partial r}(r^2 \\Sigma \\alpha c_s^2) +r^2\\Sigma \\Omega \\alpha c_s^2\\right] $$ \\begin{equation} - \\frac{3}{2}(\\rho v_z)_{\\rm w} r^2 \\Omega^2 \\ge 0 \\label{eq:engcd} \\end{equation} is satisfied, disk winds can be driven solely by the liberated gravitational energy of accretion, because the left-hand side of Equation (\\ref{eq:engcd}) is the energy flux of disk winds at infinity. %Unless the condition is fulfilled, disk winds cannot be accelerated %to infinity without additional acceleration mechanism. % (Figure\\ref{fig:ctn}). %It is more difficult for disk winds from inner regions %to satisfy the condition because of larger gravity. The time derivative term in Equation (\\ref{eq:engcd}) can be eliminated by using Equation (\\ref{eq:sgmev}). Then, using the relation of Keplerian rotation, $\\Omega \\propto r^{-3/2}$, Equation (\\ref{eq:engcd}) is rewritten as %The left panel of Figure \\ref{fig:engacc} shows the left hand side of %Equation (\\ref{eq:engcd}) in unit of $r_0^2 \\Omega_0^3 \\Sigma_{\\rm 0,init}$ %at $10^6$ yr of Model II, %our standard disk wind case %(the constant $\\alpha = \\alpha_{\\rm fl} = 8\\times 10^{-3}$ and %$C_{\\rm w} = C_{\\rm w,fl} = 2\\times 10^{-5}$ with, $\\delta_{\\rm sw}=0$; %Model II and Figure \\ref{fig:sgev1} in the main paper), %where the subscript, `0', denotes the location of 1 AU and $\\Sigma_{\\rm 0,init}$ %is the initial surface density at 1 AU. %The sign changes at $r=1.5$ AU $\\equiv r_{\\rm dw}$, which indicates that %in the outer region of $r> r_{\\rm dw}$ the accretion energy can potentially %accelerate the disk winds to the escape velocity but in the inner region %the accretion energy is not sufficient. \\begin{equation} \\frac{3}{2}\\Omega \\Sigma \\alpha c_s^2 \\ge 2r^2 \\Omega^2 (\\rho v_z)_{\\rm w}, \\end{equation} or more specificly, \\begin{equation} r\\ge \\frac{8}{9\\pi}\\frac{C_{\\rm w}^2 (r_0 \\Omega_0)^4}{\\alpha^2 c_{s,0}^4}r_0 \\equiv r_{\\rm dw}, \\end{equation} where we are using Equation (\\ref{eq:ndms}) and the relation of the MMSN ($c_s\\propto r^{-1/4}$) with Keplerian rotation. Substituting the standard values, $\\alpha=8\\times 10^{-3}$ and $C_{\\rm w} =2\\times 10^{-5}$, in the global model we have $r_{\\rm dw} = 1.4$ AU.  The disk winds in the inner region, $r<r_{\\rm dw}$, do not have sufficient energy to escape from a disk by accretion. The fates of these wind materials %with the insufficient energy in the inner region are (i) accreting directly to a central star if the angular momentum is removed, (ii) blown away by the stellar winds (see \\S \\ref{sec:stw}), or (iii) returning back to the disk (see Figure \\ref{fig:ctn} for the schematic picture). If the most of the wind gas follows the processes (i) or (ii), our calculations in \\S \\ref{sec:res} give the correct surface gas densities. On the other hand, if the process (iii) is dominant, we overestimate the escaping mass flux of the disk winds. %However, we think the fraction that returns back to the disk is not large %as shown in the next section. We further discuss the escape of the disk winds in \\S \\ref{sec:esc}.   \\subsection{Global Modeling with Energetics} \\begin{figure} \\figurenum{13} \\epsscale{1.} \\begin{center} \\plotone{glbdsk_fig6.ps} %\\plotone{f12.eps} \\end{center} \\caption{Time evolution of disk surface density. The black thin lines are the results without disk winds (Model I); the blue thin lines are the results with the disk winds of the constant mass flux, $C_{\\rm w}=C_{\\rm w,fl}$ (Model II); the red thick lines the results with the disk winds which take into account the energetics limiter for $C_{\\rm  w}$ (Equation \\ref{eq:clmt}). % ($\\delta_{\\rm sw}$). The dotted lines are the initial values. The dashed, solid, and dot-dashed lines are the results at $10^5, 10^6$, and $10^7$yrs. } \\label{fig:sgev2} \\end{figure} \\vspace{1cm} We can take into account the energetics argument (\\S \\ref{sec:eng}) in the global model. %in the disk wind mass flux in the calculation of the surface density %evolution, Equation (\\ref{eq:sgmev}). % in the main paper. We adopt the following limiter for disk wind mass flux, $C_{\\rm w}$ : \\begin{equation} %\\eqnum{S24} C_{\\rm w} = \\min(C_{\\rm w,fl}, C_{\\rm w,eng}), \\label{eq:clmt} \\end{equation} where $C_{\\rm w,eng}$ corresponds to the mass flux that gives the left-hand side of Equation (\\ref{eq:engcd}) equal to zero, namely %\\begin{equation} $$ %\\eqnum{S25} \\hspace{-4cm}C_{\\rm w,eng} = \\left\\{\\frac{\\partial}{\\partial t} \\left[\\Sigma\\frac{r^2\\Omega^2}{2}\\right] \\right . $$ \\begin{equation} \\left. - \\frac{1}{r}\\frac{\\partial}{\\partial r}\\left[r\\Omega \\frac{\\partial}{\\partial r}(r^2 \\Sigma \\alpha c_s^2) +r^2\\Sigma \\Omega \\alpha c_s^2\\right] \\right\\} \\left[\\frac{3}{2} r^2 \\Omega^2 \\rho_{\\rm mid} c_s\\right]^{-1}. \\end{equation} Equation (\\ref{eq:clmt}) reduces the mass flux in $r<r_{\\rm dw}$ to the value available from the liberated gravitational energy by accretion. In $r\\ge r_{\\rm dw}$, Equation (\\ref{eq:clmt}) gives $C_{\\rm w} = C_{\\rm w,fl}$.  Figure \\ref{fig:sgev2} displays the results which take into account the limiter for $C_{\\rm w}$ (thick red lines) in Model II, in comparison with Model I (no disk wind; black thin lines) and Model II (constant $C_{\\rm w}$; blue thin lines). %and that without disk winds (black thin lines). As expected, the difference between the red and blue lines appears only in $r<r_{\\rm dw}$($=1.4$ AU), the surface densities in the outer region are the same. Although the surface density in the inner region becomes larger in the case with the $C_{\\rm w}$ limiter, it is still much lower than the surface density of the no wind case. Then, the disk winds still play an essential role in the dispersal of protoplanetary disks even though taking into account the energetics of the disk winds. %Accordingly, the timescale of type I migration, $\\tau_{\\rm mig,I}$, %(Equation \\ref{eq:tpImg}) is also much longer than that of the no wind case %even in the inner region. %For example, for the initial surface density with $f_{\\rm g}=1$, the migration %time of an Earth-mass planet is $\\tau_{\\rm mig,I} > 3\\times 10^6$ yr in the %entire region at time $=10^6$ yr. %%We can conclude that %The planet migration is unimportant after $10^6$ yr even in this case. %%though the mass flux %%of the disk winds in the inner region is reduced.  \\subsection{Mass Loss Rate} \\begin{figure} \\figurenum{14} \\epsscale{1.} \\begin{center} \\plotone{dotm_fig3.ps} %\\plotone{f13.eps} \\end{center} \\caption{Structure of mass loss/accretion rates of Model II in the main paper at $t=10^6$ yr. %{\\it Left} : Energetics of the disk winds. The solid line plots the %left hand side of Equation (\\ref{eq:engcd}) in unit of $r_0^2 \\Omega_0^3 %\\Sigma_{\\rm 0,init}$ and the dotted line illustrates %the zero level. The negative sign indicates that the disk wind can be %potentially accelerated to the escape speed. {\\it right} : %Comparison of the mass loss and accretion rates. The solid line is the mass loss by the disk winds (Equation \\ref{eq:mdwd}) and the dashed line is the mass accretion rate (Equation \\ref{eq:macc}) of the model taking into account the disk winds. The dotted line is the mass accretion rate of the model without disk winds for comparison. The arrow indicates the location of $r_{\\rm dw}$. } \\label{fig:engacc} \\end{figure}  %In Figure \\ref{fig:dmnd}, we present the fraction of mass loss by disk winds %which can escape from a disk ($\\dot{M}_{z,{\\rm esc}}$) or trapped by the %gravity of a central star ($\\dot{M}_{z,{\\rm tr}}$), in comparison with %mass accretion rate, $\\dot{M}_r$, at the inner boundary. Here $\\dot{M}_z$ %and $\\dot{M}_r$ are respectively written as Figure \\ref{fig:engacc} shows the mass accretion rate, $\\dot{M}_r$, and the mass loss rate of disk winds, $\\dot{M}_z$, which are respectively defined as \\begin{equation} %\\eqnum{S22} \\dot{M}_r = - 2\\pi r\\Sigma v_r , \\label{eq:macc} \\end{equation} and \\begin{equation} %\\eqnum{S23} \\dot{M}_z(r) = 2\\pi \\int_{r}^{r_{\\rm out}} r' dr' (\\rho v_z)_{\\rm w}. \\label{eq:mdwd} \\end{equation} %where $(\\rho v_z)_{\\rm w}$ is the sum of mass flux from the upper and lower %disk surfaces and $\\dot{M}_z$ is the integration of $(\\rho v_z)_{\\rm w}$ %from $r$ to the outer boundary $r_{\\rm out}$($=10000$ AU). In the case with disk winds, the mass accretion rate decreases for decreasing $r$ because the mass is lost by the disk winds. The total mass loss rate by the disk winds is $\\dot{M}_z(r_{\\rm in}) =2.5\\times 10^{-9}$ $M_{\\odot}$ yr$^{-1}$ and the mass loss rate from $r>r_{\\rm dw}$ that can be accelerated to infinity by accretion energy is $\\dot{M}_z(r_{\\rm dw}) = 1.4\\times 10^{-9}$ $M_{\\odot}$ yr$^{-1}$; more than the half of the total wind mass loss can escape from the disk by the liberated gravitational energy. %Since this is more than the half of the total %$M_{\\rm z}$, the disk winds play a significant role even if the disk winds %from $r<r_{\\rm dw}$ ($1.1\\times 10^{-9} M_{\\odot}/{\\rm yr}$) all returns back %to the disk. The mass loss by the disk winds is larger than the accretion rate. %Although the inner boundary of the simulation is $r_{\\rm in} =$ (0.01 AU), %the disk truncates %at several stellar radii by stellar magnetic field in realistic %situations. % (see footnote of the main paper). Assuming the truncation of the disk at 5-10 stellar radii ($\\sim 0.1$ AU), %($\\approx 10R_{\\odot} \\approx 5$ stellar radii), $\\dot{M}_r \\approx 2\\times 10^{-10}$ $M_{\\odot}$ yr$^{-1}$, which can be regarded as the mass accretion rate through magnetic flux tubes directly connecting to the central star \\citep{ken96}. %(Kenyon et al. 1996). %The total accretion can be larger than this value by the contribution %from the direct accretion of the inner disk wind material. %On the other hand,  \\begin{figure} \\figurenum{15} \\epsscale{1.1} \\begin{center} \\plotone{glbdsk_tev_fig1.ps} %\\plotone{f14.eps} \\end{center} \\caption{Time evolution of the mass loss rate of the disk wind (solid) and accretion rate (dashed) of Model II. Here, the wind mass loss rate is the integration of $r>r_{\\rm dw}$, $\\dot{M}_z(r_{\\rm dw})$, that can potentially escape from a central star gravity by accretion energy. The shown accretion rate is the maximum value at each time (e.g. for $t=10^6$ yrs in Figure \\ref{fig:engacc}, the maximum accretion rate is obtained at $r\\approx 15$ AU). } \\label{fig:tmevmd} \\end{figure}   %The mass loss rate by the disk winds is $2.5\\times 10^{-9}f_{\\rm g} %M_{\\odot}$ yr$^{-1}$ at $10^6$yr (\\S\\ref{sec:eng}). % (Figure \\ref{fig:engacc} in {\\it SOM}). %For the standard MMSN ($f_{\\rm g}=1$), this is The obtained mass loss rate by the MRI-driven disk winds is larger than the mass loss rate, $\\lesssim 10^{-10} - 10^{-9}$ $M_{\\odot}$ yr$^{-1}$, predicted by the UV photoevaporation (e.g. Matsuyama et al.2003). Recently, photoevaporation by X-rays from a central star is also proposed. It is reported that the mass loss rate due to X-rays may be larger %, up to $\\sim 10^{-8}$ $M_{\\odot}$ yr$^{-1}$, than that by the UV photoevaparation, where the calculated mass loss rates are rather uncertain %if sufficient X-ray flux is supplied \\citep{ecd09,gdh09,owe10}. Significant difference of the disk wind mass loss from the photoevaporation processes by UV or X-rays is time evolution. % (e.g. Matsuyama et al.2003). The mass loss rate by the disk winds is correlated with the accretion rate (Figure \\ref{fig:tmevmd}), because the energy source of the disk winds is the gravitational energy liberated by accretion. %(\\S\\ref{sec:esc} of {\\it SOM}). Therefore, the wind mass loss rate is larger at earlier times when the accretion rate is larger; the disk winds significantly contribute to the dispersal of the gas component of a disk from the beginning. As a result, an inner hole forms from the early epoch and its size is gradually expanding. %This is a clear contrast to On the other hand, the UV photoevaporation mechanism is significant at the later times after the sufficient mass dissipates by accretion. Therefore, an inner hole is anticipated at later time when the evaporating mass flux becomes comparable to the mass accretion rate (Alexander et al. 2006). The X-ray photoevaporation, which may be effective from slightly earlier time, is also expected to give the similar trend to the UV photoevaporation \\citep{gdh09}. %\\citep{alx06}.%(Alexander et al.2006). %The MRI-driven disk winds play an essential role in the dispersal %of the gas component of protoplanetary disks, which is true even if %the dead zone exists initially (Paper II). %(\\S \\ref{sec:dz} of {\\it SOM}). %The disk winds also naturally predict %inside-out evaporation. A sharp inner hole may form if the net vertical %magnetic flux is rather strong, %The disk winds are highly structured and become quasi-periodically %strong because of the breakup of %large channel flows (Suzuki \\& Inutsuka 2009). These features are the %predictions of the MRI-driven disk winds, which are observable by ALMA in %the near future.    \\subsection{Escape of Disk Winds} \\label{sec:esc} %Following the energetics arguments in \\S \\ref{sec:eng}, We further continue the discussions on the escape of disk winds.  \\subsubsection{Local Simulations with Larger Vertical Boxes} \\label{sec:lslvb} %In the local simulations we have shown that the average velocities of the %disk winds become an order of the sound speed at the simulation boundaries, %$z=\\pm 4H$. %Considering a typical protoplanetary disks such as MMSN, this velocity is %smaller than the escape velocity. For example, at 1 AU of the MMSN the sound %velocity is $\\approx 1$ km s$^{-1}$, while the escape velocity is %42 km s$^{-1}$. %In the previous subsections, we discuss the escape of the disk winds from %the energetics of accretion disks. In order to study the acceleration of the disk winds at higher altitudes, %we need to calculate the global dynamics. Although the high resolution %global 3D simulation that can %resolve small scale turbulence is the most direct way, performing such %simulations until the saturation of MHD turbulence is still computationally %very expensive.   Here, we perform the local 3D MHD simulations with larger vertical boxes. %than that used in the main paper (Figure \\ref{fig:btdep}) %to quantitatively examine the effect of the mass returning to the disk from %higher altitudes. %In the simulation, We here use the realistic vertical gravity, \\begin{equation} %\\eqnum{S26} g_z = \\frac{G M_{\\star} z}{(r^2+z^2)^{3/2}} = \\Omega^2 z \\frac{r^3}{(r^2+z^2)^{3/2}}, \\label{eq:zgr} \\end{equation} where $r$ is radial location from a central star. In the usual local simulations, %used in the main paper as well as most of previous works, the only leading term, $g_z\\simeq \\Omega^2 z$, is considered, %the tidal force, $3\\Omega^2 \\mbf{x}$, in $x$ direction and because $r$ does not appear explicitly and it can be treated more easily % in this treatment (Hawley et al.1995). On the other hand, the escape velocity is not defined at the expense of the simplification, and the vertical gravity is overestimated by a factor of $\\frac{r^3}{(r^2+z^2)^{3/2}}$, which makes density at a high altitude unrealistically lower. % in Equation (\\ref{eq:zgr}). %by which it is more difficult to simulate disk winds at higher altitudes %because the densities become lower than in the real situations. %Equation (\\ref{eq:zgr}) shows that one overestimates $g_z$ by a factor of %$\\frac{r^3}{(r^2+z^2)^{3/2}}$, if one takes the leading order, $\\Omega^2 z$. To avoid these shortcomings %run the simulations in larger vertical boxes we use Equation (\\ref{eq:zgr}) by explicitly setting $r$ % and box sizes which are tabulated in (Table \\ref{tab:lzb}). Note that $r=20H$ corresponds to the location of $r\\approx 1$ AU for the MMSN. In these runs, we set the net vertical fields with $\\beta_{{z,\\rm mid}}=10^6$ at the midplanes. %We set $r=20H$, which is for the value at 1 AU of the MMSN, and the vertical %box of $-8 H < z < 8H$.   \\begin{figure} \\figurenum{16} \\epsscale{1.2} \\begin{center} \\plotone{var_t_av_fig36_cmpstr12.ps} %\\plotone{f15.eps} \\end{center} \\caption{Comparison of the disk wind structures with the different box sizes. On the left, we display the time-averages of vertical velocity and density. On the right we compare the time-averages of magnetic energies and $\\alpha$ values only in the $-6H<z<6H$ region. The solid lines are the results with $r=20H$ and box sizes, $-6H < z < 6H$ (Model VI) and $-8H < z < 8H$ (Model VII). The dashed lines are the results with $r=10H$ and box size, $-12 H < z < 12 H$ (Model VIII). The dotted lines are the results of the reference case (Model II; $r\\rightarrow \\infty$ and $-4H < z < 4H$). On the right we use the thick lines for Model VII.} \\label{fig:lzb} \\end{figure}    \\begin{table}[h] %\\begin{deluxetable}{|c|c|c|} \\tablenum{2} \\begin{tabular}{|c|c|c|} \\hline Model & $r$ & Box Size \\\\ %& $C_{\\rm w}$\\\\ \\hline \\hline II(Reference) & -- & $-4H < z < 4H$ \\\\ %& $7.5\\times 10^{-5}$ \\\\ \\hline VI & 20 & $-6H < z < 6H$ \\\\ %& $3.5\\times 10^{-5}$ \\\\ \\hline VII & 20 & $-8H < z < 8H$ \\\\ %& $3.5\\times 10^{-5}$ \\\\ \\hline VIII & 10 & $-12H < z < 12H$ \\\\ %& $6\\times 10^{-5}$ \\\\ \\hline \\end{tabular} \\caption{Radial positions and vertical box sizes of the local simulations. } \\label{tab:lzb} \\end{table} %\\end{deluxetable}  %In Models VI and VII, we adopt $r=20H$, which corresponds to the location %of 1 AU for the MMSN. %We perform the simulation runs in the two different vertical %box sizes, $- 6H < z < 6H$ and $- 8H < z < 8H$, to see the dependence %on the box size. %For the larger box size case, $\\pm 12H$, %we reduce the gravity by adopting $r=10 H$. %, which corresponds to the location of 16 AU for the MMSN. Figure \\ref{fig:lzb} compares the time averaged vertical structures of these cases (Table \\ref{tab:lzb}). %In all the cases, the disk winds are accelerated as elevating to larger %$|z|$, and the velocities at the boundaries are higher for larger box cases, %although the velocities are still slower than the escape speeds, %$\\sqrt{2}r\\Omega$. The top left panel shows that the onsets of the disk winds take place at higher altitudes in larger box cases; the results depend on the simulation box size. However, the mass flux of the disk winds ($\\rho v_z$) do not show a monotonic behavior as presented in the bottom panel of Figure \\ref{fig:lzmf}. By increasing the box size from $\\pm 4H$ to $\\pm 6H$, the mass flux decreases at first \\footnote{Although we do not take into account $r$ in the case with the vertical box of $-4H < z < 4H$, this effect is negligible for sufficiently small $z$.}. However, increasing the box size from $\\pm 6H$ to $\\pm 8H$, the mass flux increases slightly. The largest box size case with the smaller gravity, $r=10H$, shows larger mass flux than these smaller box cases. %From this tendency, we expect that The mass flux of the disk winds is bound by a lower limit and dose not becomes further smaller even though we use a larger vertical box.  \\begin{figure} \\figurenum{17} \\epsscale{0.9} \\begin{center} \\plotone{beta_dep_fig11.ps} %\\plotone{f16.eps} \\end{center} \\caption{The $\\bar{\\alpha}$ values (top) and the mass flux of the disk winds (bottom) for the different sizes of the vertical simulation boxes. The triangles are the results of Model II ($r\\rightarrow \\infty$), the squares are the results with $r=20 H$ (Models VI \\& VII), and the circles are the results with $r=10H$ (Model VIII). The dotted line is the level of $C_{\\rm w,fl}$. } \\label{fig:lzmf} \\end{figure}  %The main reason why the mass flux of the disk winds is not reduced so much %in the sufficiently large vertical boxes is that the magnetic fields are %more amplified in these cases (the top-right panel of Figure \\ref{fig:lzb}). The dependence of the mass flux on the simulation box size can be explained by the saturation of the amplified magnetic fields (the top right panel of Figure \\ref{fig:lzb}). When we increase the box size from $\\pm 4H$ to $\\pm 6H$, the escaping mass from the disk surfaces becomes smaller at first. %(e.g. increasing the box size from $\\pm 4H$ to $\\pm 6H$). Accordingly, the escaping magnetic flux of the toroidal ($y$) and radial ($x$) components decrease, because the magnetic fields are frozen in the gas. The amplification of magnetic fields %by MRI and winding is balanced with the escape with the disk winds in addition to magnetic reconnections. In the larger box cases the escaping flux becomes smaller and the saturated level of magnetic fields increases. %($\\approx$ toroidal field strength) %is stronger. As a result of the larger magnetic pressure, the mass is lift up to higher locations %and the density becomes higher there (the bottom left panel of Figure \\ref{fig:lzb}). The increase of the density inhibits further decrease of the mass flux of the disk winds, $\\rho v_z$, in a self-regulated manner. %This is a self-regulation mechanism; the decrease of the disk wind flux itself %inhibits the further decrease of the mass flux through the amplification of %magnetic fields.  In spite of the increase of the magnetic field strength for larger boxes, $\\alpha_z$ does not increase so much (the lower right panel of Figure \\ref{fig:lzb}). This is because the magnetic field of the large box cases around $z\\approx 1.5H$ (the locations of the peaks) is dominated by the coherent toroidal component which does not contribute to the anisotropic Maxwell stress. Therefore, the integrated $\\bar{\\alpha} (= \\int dz \\rho \\alpha_z/\\int dz\\rho )$, which directly determines the global mass accretion rate, does not depend on the vertical box size (the top panel of Figure \\ref{fig:lzmf}).   %Figure \\ref{fig:lzb} shows that The top left panel of Figure \\ref{fig:lzb} shows that the average velocities of the disk winds do not still reach the escape speeds ($=\\sqrt{2}r\\Omega = \\sqrt{2}\\left(\\frac{r}{H}\\right)H\\Omega = 2 \\left(\\frac{r}{H}\\right)c_s$). However, our conservative choice of the floor values, $C_{\\rm w,fl}$, probably gives a reasonable estimate %of the mass flux for the global disk calculations (the dotted line of Figure \\ref{fig:lzmf}), first because the mass flux seems bounded by the lower limit as explained above, and second because there are a couple of mechanisms that favor the escape of the disk winds but are not included in this paper (see below).  \\subsubsection{Magnetocentrifugal Winds} %we cannot tell whether the disk winds %really escape from the disk surfaces. However, we would like to note that Our local simulations do not take into account the acceleration (momentum input) of the disk winds by centrifugal force with global magnetic fields. If poloidal ($r-z$ components) magnetic field lines sufficiently incline with respect to an accretion disk, the gas can flow out along with the field lines by the centrifugal force \\citep{bp82,ks98}. %When one considers, as a %simplest case, the force balance between the gravity by a central star and %centrifugal force along with the field line which rotates with Kepler %frequency, the gas inevitably streams upward if the inclination angle of %the field line is smaller than % $60^{\\circ}$ \\citep[$|B_r|/|B_z|>1/\\sqrt{3}$;][] {bp82}. %Blandford \\& Payne 1982). Such global magnetic fields will be common in protoplanetary disks as a result of the contraction of cold molecular clouds with %more or less coherent interstellar magnetic fields. Therefore, our local simulations probably underestimate the momentum input to the disk winds.  %Summing up the above discussions, %although the local simulations presented in the main paper %(Figure \\ref{fig:btdep}) may overestimate the mass flux of the disk winds %by up to a factor of 3,    \\subsubsection{Stellar Winds} \\label{sec:stw}  \\begin{figure} \\figurenum{18} \\epsscale{1.} \\begin{center} \\plotone{dskwdct3.eps} %\\plotone{f17.eps} \\end{center} \\caption{Schematic picture of dispersal of a protoplanetary disk by disk winds. The wind material in the outer region, $r>r_{\\rm dw}$, can stream out solely by the gravitational energy by accretion. %The disk wind process works more effectively in an inner region, so %an inner hole gradually grows. Because On the other hand, the wind material in the inner location cannot escape by itself from the gravity of a central star. A fraction of the winds directly accrete to the central star after lift up. A fraction of the disk winds may be accelerated by the dynamical pressure of the stellar winds.  If the stellar wind flux is not strong enough, the disk wind material returns back to the disk after transported outward. %The shape of an accretion disk is supposed to be structured especially around %the surfaces because the disk winds are driven with a quasi-periodic cycle %of several rotation periods by the breakups of channel flows (SI09). } %(Suzuki \\& Inutsuka 2009).  } \\label{fig:ctn} \\end{figure}  %So far, we have focused on the dispersal by the MRI-driven disk winds %by using the results of the local 3D MHD simulations. %In this treatment, we do not take into account the magnetocentrifugal force %by global magnetic fields (Blandford \\& Payne 1982; Kudoh \\& Shibata 1997), %which probably contributes to the acceleration of the winds at high %altitudes.  So far, we have focused on the disk winds driven by MHD turbulence. Central T Tauri stars are also expected to drive stellar winds by the mass accretion from disks \\citep{hir97,mp05} and %(Hirose et al.1997; Matt \\& Pudritz 2005) and the surface convections \\citep{cra08}. %(Cranmer 2008). Observations of T Tauri stars show that the outflow rates range from $10^{-11}$ to $10^{-8}$ $M_{\\odot}$ yr$^{-1}$, which is typically 0.01 - 1 times of the mass accretion rates \\citep{wh04}. %(White \\& Hillenbrand 2004). Although it is difficult to determine the exact launching locations of the observed outflows, some fractions are expected to come from central stars \\citep{ed06}. %(Edwards et al. 2006), %and the mass loss rates are considerably larger than the mass loss rate of the %present solar wind ($\\approx 2\\times 10^{-14}$ $M_{\\odot}$ yr$^{-1}$). %The mass loss rate of such stellar winds is %typically up to $\\sim 0.1$ of the accretion rate (Matt \\& Pudritz 2005), which %is $10^{3-6}$ larger than the mass loss rate of the present %solar winds. %Then, the ram pressure of such massive stellar winds is also expected to %play a role in blowing off the lift up material by the disk winds in the %inner region, while  At present the stellar wind is regarded to give a minor effect to the dispersal of the gas component of protoplanetary disks \\citep[e.g.][]{shu93} %(e.g. Shu et al 1993) except in the very surface regions because the stellar winds almost slide along the disk surfaces \\citep{mat09}. %(Matsuyama et al.2009). However, if disk materials are lifted up as we have shown so far, the ram pressure of stellar winds can directly push away the lifted up materials because of the large elevation angle.  We discuss the role of stellar winds from a simple momentum balance. %The comparison of the mass of the lift-up gas and the momentum of the %stellar winds %(i.e. mass loss rate and velocity) %determines whether the lift-up material is totally blown away. %As an example, we take our reference case with the disk winds at $t=10^6$ yr %(Figure \\ref{fig:engacc}). %In the inner region, $r < r_{\\rm dw}$, $\\dot{M}_{z,{\\rm in}}=1.1\\times %10^{-9}$ $M_{\\odot}$ yr$^{-1}$ is lift up by the disk winds but the energy %is not sufficient to escape from the disk. %Typical velocity of T Tauri star winds is $v_{\\star} = 200 - 400$ km s$^{-1}$, %which is an order of the escape speed at the stellar surface. Let us consider a situation, in which the lifted up gas by disk winds floats above a disk and the radial stellar winds hit the floating gas. If we assume that both lifted up gas and stellar wind gas move together radially outward after the hitting, we can estimate the radial velocity of the moving gas, $v_{\\rm cmb}$, from the momentum balance as \\begin{equation} %\\eqnum{S27} (\\dot{M}_{z,{\\rm in}} + W \\dot{M_\\star})v_{\\rm cmb} = W \\dot{M_\\star} v_{\\star}, \\label{eq:stw} \\end{equation} where $\\dot{M}_{z,{\\rm in}}$ is the rate of the mass that is supplied from disk winds but float above a disk without sufficient energy, $W$ is the fraction of the solid angle obscured by the lifted up gas by the disk winds, and $\\dot{M}_{\\star}$ and $v_{\\star} (\\approx 200-400\\; {\\rm km\\;s^{-1}})$ are the mass loss rate and velocity of the stellar winds. As a typical example, we consider the result of Model II at $t=10^6$ yr (Figure \\ref{fig:engacc}). Then, $\\dot{M}_{z,{\\rm in}}=1.1\\times 10^{-9}$ $M_{\\odot}$ yr$^{-1}$; we assume the lifted up gas fill up to the height of $r/2$, which gives $W = \\int_{\\pi/2}^{\\cot^{-1}(1/2)} d\\cos\\theta \\approx 0.45$. %We evaluate $v_{\\rm cmb}$ in comparison with the escape velocity %at 1AU of the MMSN ($\\approx 42$ km s$^{-1}$). From Equation (\\ref{eq:stw}), The gas that is lifted up by the disk winds but does not have the sufficient energy will distribute in $r<r_{\\rm dw}(=1.4\\;{\\rm AU})$. Here, we compare $v_{\\rm cmb}$ with the escape velocity at 1 AU, $v_{\\rm esc, 0}\\approx 42$ km s$^{-1}$ as a typical condition. Substituting these values into Equation (\\ref{eq:stw}), if $\\dot{M}_{\\star}>4\\times 10^{-10}$ $M_{\\odot}$ yr$^{-1}$, $v_{\\rm cmb} > v_{\\rm esc,0}$ is satisfied and the lifted up material can be blown away by the stellar winds.   %, which is rather %small in comparison with the observationally inferred value, %$10^{-11}$ to $10^{-8}$ $M_{\\odot}$ yr$^{-1}$ (White \\& Hillenbrand 2004).  %If the stellar winds are mainly powered by the accretion, the mass loss rate %is generally smaller than the accretion rate. %As discussed in the subsection \\ref{sec:eng}, %the accretion rate at the disk truncation radius is $2 \\times 10^{-10}$ %$M_{\\odot}$ yr$^{-1}$, whereas it could be larger if the inner disk winds %directly accrete to the star. Then, it is inferred that the stellar wind %mass loss rate is $\\dot{M}_\\star\\lesssim 10^{-10}$ $M_{\\odot}$ yr$^{-1}$. %In this case, If $\\dot{M}_{\\star}$ is smaller than this value, the stellar winds can blow away the disk wind material at sufficiently high altitudes where the density is low. The disk wind material at lower heights will move outward after hit by the stellar winds but again return back to the disk without sufficient energy. When the returning location becomes $r > r_{\\rm dw}$, the gas finally flow out by the disk winds. (see Figure \\ref{fig:ctn} for the schematic picture).   \\vspace{1cm}"
    },
    "0911/0911.0792_arXiv.txt": {
        "abstract": "{{\\bf stars: abundances; stars: evolution; Galaxy: evolution; Galaxy: formation; globular clusters: general; open clusters and associations: general}} I discuss the chemical evolution of star clusters, with emphasis on old globular clusters, in relation to their formation histories. Globular clusters clearly formed in a complex fashion, under markedly different conditions from any younger clusters presently known. Those special conditions must be linked to the early formation epoch of the Galaxy and must not have occurred since.  While a link to the formation of globular clusters in dwarf galaxies has been suggested, present-day dwarf galaxies are not representative of the gravitational potential wells within which the globular clusters formed. Instead, a formation deep within the proto-Galaxy or within dark-matter minihaloes might be favoured. Not all globular clusters may have formed and evolved similarly. In particular, we may need to distinguish Galactic halo from Galactic bulge clusters. ",
        "introduction": "Clusters of stars are ideal witnesses of past star formation and chemical evolution. This is because the stars in a cluster outline a sequence in photometric diagrams which can be compared with theoretical models for stellar evolution to infer an age and metallicity for the cluster. This way, young ($<100$ Myr) clusters can be used to investigate the mechanisms of triggering and propagation of star formation in detail and the dependence of star formation on local environmental conditions. Intermediate-age ($\\approx0.1$--10 Gyr) clusters trace much of the history of galaxies, while old ($>10$ Gyr) globular clusters (GCs) probe the formation epoch of galaxies.  The `chemical' (really: `elemental') composition of a cluster is a powerful legacy of the entire history of star formation, stellar evolution, dynamical and interstellar medium (ISM) processes pre-dating the formation of the cluster. The cluster age tags this imprint to a specific time in the history of a galactic system.  Complications arise from the fact that the position and kinematics of a cluster may no longer be directly connected to its place of origin and birth kinematics. In the Milky Way, dynamical heating leads to outward migration away from the disc midplane and the Galactic Centre. Clusters dissolve over time, and those that survive for many billions of years generally have masses around $10^4-10^6$ M$_\\odot$. Their formation requires conditions which are uncommon in the present, nearby Universe. In addition, attempts to quantify the star-formation rate are hampered by the uncertain relationship between the conditions under which stars were formed and the number and size of clusters which were produced. Nevertheless, their chemical composition can constrain all of these open issues.  I concentrate on new lines of evidence, emerging in a fast-paced field of research, complementing and complemented by the excellent reviews of Gratton {\\it et al.}\\ (2004) and Friel (1995), Freeman \\& Bland--Hawthorn (2002) and Brodie \\& Strader (2006). ",
        "conclusions": ""
    },
    "0911/0911.4779_arXiv.txt": {
        "abstract": "Recent measurements of cosmic ray leptons by PAMELA, ATIC, HESS and Fermi revealed interesting excesses. Many authors suggested particle Dark Matter (DM) annihilations could be at the origin of these effects. In this paper, we critically assess this interpretation by reviewing some results questioning the naturalness and robustness of such an interpretation. Natural values for the DM particle parameters lead to a poor leptons production so that models often require signal enhancement effects that we constrain here. Considering DM annihilations are likely to produce antiprotons as well, we use the PAMELA antiproton to proton ratio measurements to constrain a possible exotic contribution. We also consider the possibility of an enhancement due to a nearby clump of DM. This scenario appears unlikely when compared to the state-of-the-art cosmological N-body simulations. We conclude that the bulk of the observed signals most likely has no link with DM and is rather a new, yet unconsidered source of background for searches in these channels. ",
        "introduction": "Recent measurements of cosmic ray positrons and electrons in the GeV-TeV region have rised a lot of interest as they exhibit unexpected features. The PAMELA satellite found evidence for a rise of the positron fraction ($e^+/(e^-+e^+)$) above 10 GeV~\\cite{pamela}, as some previous experiments found some hints for~\\cite{ams,heat}. The data points are shown on the left panel of Fig.~\\ref{fig1}. In an independent way, the ATIC, HESS and Fermi collaborations reported measurements of the $e^-+e^+$ flux, between dozens of GeV and a few TeV~\\cite{atic,hess,fermi}. These seem to indicate a deviation from a pure power law above 100 GeV, as shown on the right panel of Fig.~\\ref{fig1}. Note that sizable systematic errors affect these data so that it is still unclear whether the feature is a sharp peak or a smoother bump.  In the standard picture, primary cosmic rays (mostly protons) are produced and accelerated in supernovae and their remnants, and secondaries result from the spallation of primaries onto the interstellar medium. As conventional sources carry the same baryon number as the whole Universe, electrons are mostly primaries and positrons are usually though to be secondaries. It is extremely difficult to reproduce the PAMELA rise or the leptonic sum feature in conventional diffusion models. As an illustration, the left panel of Fig.~\\ref{fig1} displays a collection of predictions from a diffusion model. The different curves correspond to distinct sets of parameters allowed by cosmic ray measurements. None of propagation setups can reproduce the data (for more details on diffusion models, see~\\cite{delahaye} and references therein). It is yet unclear if the $e^-+e^+$ feature and the positron rise have a common origin, but they both seem to indicate the presence of a nearby primary source of cosmic ray electrons and positrons (leptons), which is not accounted for in conventional cosmic ray models. Due to the efficient energy losses of the electrons and positrons as they propagate in the Galaxy, this primary production must take place within a short $\\sim$1 kpc radius around the Earth.  Dealing with nearby primary sources of positrons, one can generally invoke two classes of models: pure astrophysics and particle physics inspired models. In the first case, it is assumed that some astrophysical source exists in our Galactic neighbourhood which produces an abnormal amount of positrons. This could be for instance a supernova remnant with high level of secondary particles produced at the source, or a pulsar. It is actually quite easy to reproduce the leptonic signals with conventional sources but there is still no obvious candidate. Particle physics inspired interpretations of the data relying on new physics beyond the Standard Model, the second possibility is more speculative. Most studied possibilities consider DM annihilation or decays as exotic source of cosmic rays. For a review of the measurements and possible interpretations, from both astrophysics and particle physics, see~\\cite{tango}. In this paper, we focus on the annihilating DM case, we question in particular the naturalness of such an interpretation. As some conventional interpretations for the cosmic ray leptons excesses exist, it can seem odd to invoke more exotic scenarios. This is however an extensively studied possibility and it is now important to study how natural and robust the DM hypothesis can be.  \\begin{figure} \\centering \\epsfig{figure=Figure1.eps,height=2.5in} \\epsfig{figure=Figure2.eps,height=2.5in} \\caption{Cosmic ray positron fraction (left) and electrons+positrons fluxes (right). \\label{fig1} } \\end{figure} ",
        "conclusions": "In this paper, it is shown that most likely the lepton excesses observed by PAMELA, ATIC, Fermi and HESS are not caused by conventional secondary cosmic rays, meaning that a nearby source significantly contributes to the local flux. The main classes of candidate are astrophysical ($e.g.$ a nearby pulsar) or more exotic (DM annihilations). The first case is clearly favored as one has to invoke non conventional scenarios for DM in order to avoid contraints from different observables. In conclusion, though very exciting, the DM interpretation of the cosmic ray excesses lacks naturalness, and a more conventional source might be at the origin of this phenomenon.   It is definitely possible to reproduce the observed cosmic ray data with the help of dark matter signals. Within this hypothesis however, there will always be some tension between the different channels and observables, or quite a high level of fine-tuning. It could be that we are encircling the dark matter particle properties but most likely, the bulk of the observed leptons come from a nearby astrophysical source producing a large fraction of electron/positron pairs. In that case, it would constitute an additional background for indirect searches for dark matter through lepton channels that had never been accounted for before. A big step forward will be the measurement of the small anisotropy in the arrival directions of the cosmic ray leptons, if any. If it is observed and points towards a known pulsar, the conclusion will be quite clear, whereas if it points at some Fermi unidentified source outside the Galactic disk, dark matter might be back in the game. It also urges to separate electrons from positrons at higher energies and to increase statistics in all channels. Future results from PAMELA and especially AMS-02 on leptons but also all nuclei fluxes will be of great help in feeding the cosmic ray propagation models. Then, the indirect searches for dark matter through charged channels can go on, in particular by looking for fine structures in the spectra. It will then be very interesting (and challenging!) to try to interpret future data and weight them against LHC and direct detection results. Whatever the nature of the source, we might be witnessing the first direct observation of a nearby source of cosmic rays with energies in the range of GeV to TeV. These are exciting times and one might have to wait a little more for the final solution to this cosmic puzzle. The answer will certainly come from a convergence of different messengers information. Thanks to its large field of view, the Fermi gamma-ray telescope will hopefully say something about a nearby source, should it be a pulsar or a more exotic one. Eventually, future large observatories in neutrino and gamma rays (km3net, CTA) will certainly offer the great opportunity to have a deep look into this brainteaser."
    },
    "0911/0911.3968_arXiv.txt": {
        "abstract": " ",
        "introduction": "So far, there has been extensive works devoted to studying the effects of quark matter on astrophysical phenomena; the gravitational wave radiations~\\cite{lin06, yasutake07, abdikamalov08},  cooling processes~\\cite{page00,blaschke00,blaschke01,grigorian05}, neutrino emissions~\\cite{nakazato08, sagert08}, rotational frequencies~\\cite{burgio03}, the maximum energy release by conversions from neutron stars to quark/hybrid stars~\\cite{yasutake05, zdunik07}, etc.. However, uncertainties of equation of state~(EOS) have been still left.  For such studies on compact stars, general relativistic effects are fundamentally important, since baryon density and strong magnetic field are comparable to pressure, $\\rho_0 c^2 \\sim B^2/8\\pi \\sim P$. Here we report the effects of quark-hadron phase transition on the structures of general relativistic stars with purely toroidal magnetic field. For the mixed phase, we take into account of the finite-size effects, which lead to  non-uniform ``Pasta\" structures. The star with pure toroidal magnetic field is unstable, but it becomes another stable star in which toroidal magnetic field is dominant in the dynamical simulation~\\cite{kiuchi08b}. Moreover, Heger et al. have suggest that the toroidal magnetic field may dominate $10^5$ times lager than the poloidal magnetic field at the last stage of the main sequence~\\cite{heger05}.  The organization of this report is as follows. Adopted EOS is briefly discussed in Sec.~2 with  equilibrium models of magnetized rotating compact stars. In Sec.~3, we show our numerical results. In Sec.~4, we discuss the consequence of our calculations. ",
        "conclusions": "\\label{sec:summary} In this report, we have investigated the effects of the quark-hadron mixed phase on the magnetized rotating stars with the general relativistic equilibrium configuration. As a result, we find that the distribution of the magnetic field for hybrid stars has a discontinuity in the quark-hadron mixed phase.  It was pointed out that the strong magnetic field may change EOSs~\\cite{broderick00}. In particular for quark matter, it has been found that the energy gaps in the magnetic Color-Flavor-Locked phase are the oscillating functions with respect to the magnetic field~\\cite{noronha07, fukushima08}. All of their effects are important, and may change our results.  \\bigskip  We thank K.~Kashiwa, M.~Hashimoto, S.~Yamada, S.~Yoshida, and Y.~Eriguchi for informative discussions. This study was supported in part by the Grants-in-Aid for the Scientific Research from the Ministry of Education, Science and Culture of Japan (No. 20540267, 21105512, 19540313)."
    },
    "0911/0911.3535_arXiv.txt": {
        "abstract": "The early part of the gravitational wave signal of binary neutron star inspirals can potentially yield robust information on the nuclear equation of state. The influence of a star's internal structure on the waveform is characterized by a single parameter: the tidal deformability $\\lambda$, which measures the star's quadrupole deformation in response to the companion's perturbing tidal field.  We calculate $\\lambda$ for a wide range of equations of state and find that the value of $\\lambda$ spans an order of magnitude for the range of equation of state models considered.  An analysis of the feasibility of discriminating between neutron star equations of state with gravitational wave observations of the early part of the inspiral reveals that the measurement error in $\\lambda$ increases steeply with the total mass of the binary. Comparing the errors with the expected range of $\\lambda$, we find that Advanced LIGO observations of binaries at a distance of 100~Mpc will probe only unusually stiff equations of state, while the proposed Einstein Telescope is likely to see a clean tidal signature. ",
        "introduction": "The observation of inspiraling binary neutron stars (NSs) with ground-based gravitational-wave detectors such as LIGO and Virgo may provide significantly more information about neutron-star structure, and the highly uncertain equation of state (EOS) of neutron-star matter, than is currently available. The available electromagnetic observations of neutron stars provide weak constraints from properties such as the star's mass, spin, and gravitational redshift (see for example~\\cite{LattimerPrakash2007, ReadLackeyOwenFriedman2009}).  Simultaneous measurements of both the mass and radius of a neutron star~\\cite{Ozel2006, LeahyMorsinkCadeau2008, OzelGuverPsaltis2008, GuverOzeletal2008,Leahyet2009}, on the other hand, have the potential to make significantly stronger constraints.  These measurements, however, depend on detailed modeling of the radiation mechanisms at the neutron-star surface and absorption in the interstellar medium, and are subject to systematic uncertainties.  Another possibility for obtaining information about the neutron star EOS is from the inspiral of binary neutron stars due to gravitational radiation. The tidal distortion of neutron stars in a binary system links the EOS describing neutron-star matter to the gravitational-wave emission during the inspiral. Initial estimates showed that for LIGO, tidal effects change the phase evolution only at the end of inspiral, and that point particle post-Newtonian waveforms can be used for template-based detection~\\cite{Kochanek1992, BildstenCutler1992, LaiWiseman}. With the projected sensitivities of later-generation detectors, however, effects which can be neglected for the purpose of detection may become measurable in the strongest observed signals.  While EOS effects are largest during the late inspiral and merger of two neutron stars where numerical simulations must be used to predict the signal, Flanagan and Hinderer showed that a small but clean tidal signature arises in the inspiral below 400~Hz~\\cite{FlanaganHinderer2008}. This signature amounts to a phase correction which can be described in terms of a single EOS-dependent tidal deformability parameter $\\lambda$, namely the ratio of each star's induced quadrupole to the tidal field of its companion. The parameter $\\lambda$ depends on the EOS via both the NS radius $R$ and a dimensionless quantity $k_2$, called the Love number~\\cite{BrookerOlle1955, MoraWill2004, BertiIyerWill2008}: $\\lambda= 2/(3G) k_2 R^5$.  The relativistic Love numbers of polytropic\\footnote{Polytropes are often written in two forms. The first form is expressed as $p=K\\epsilon^{1+1/n}$, where $p$ is the pressure, $\\epsilon$ is the energy density, $K$ is a pressure constant, and $n$ is the polytropic index.  The second form, is given by $p=K\\rho^{1+1/n}$, where $\\rho$ is the rest-mass density, defined as the baryon number density times the baryon rest mass.  The first form was mainly used in the recent papers~\\cite{Hinderer2008, DamourNagar2009, BinningtonPoisson2009}. However, the second form is more commonly used in the neutron-star literature and is more closely tied to the thermodynamics of a Fermi gas. We will use both forms as was done in Ref.~\\cite{DamourNagar2009}.} EOS were examined first by Flanagan and Hinderer~\\cite{FlanaganHinderer2008, Hinderer2008} and later by others in more detail~\\cite{DamourNagar2009, BinningtonPoisson2009}. Flanagan and Hinderer also examined the measurability of the tidal deformability of polytropes and suggested that Advanced LIGO could start to place interesting constraints on $\\lambda$ for nearby events. However, they used incorrect values for $k_2$, which overestimated $\\lambda$ by a factor of $\\sim 2-3$ and were therefore overly optimistic about the potential measurability. In addition, polytropes are known to be a poor approximation to the neutron star equation of state, and there may be significant differences in the tidal deformability between polytropes and ``realistic'' EOS. In this paper, we calculate the deformability for realistic EOS, and show that a tidal signature is actually only marginally detectable with Advanced LIGO.  In Sec.~\\ref{sec:calculation} we describe how the Love number and tidal deformability can be calculated for tabulated EOS.  We use the equations for $k_2$ developed in \\cite{Hinderer2008}, which arise from a linear perturbation of the Oppenheimer-Volkoff (OV) equations of hydrostatic equilibrium.  In Sec.~\\ref{sec:candidateeos} we then calculate $k_2$ and $\\lambda$ as a function of mass for several EOS commonly found in the literature. We find that, in contrast to the Love number, the tidal deformability has a wide range of values, spanning roughly an order of magnitude over the observed mass range of neutron stars in binary systems.  As discussed above, the direct practical importance of the stars' tidal deformability for gravitational wave observations of NS binary inspirals is that it encodes the EOS influence on the waveform's phase evolution during the early portion of the signal, where it is accurately modeled by post-Newtonian (PN) methods.  In this regime, the influence of tidal effects is only a small correction to the point-mass dynamics.  However, when the signal is integrated against a theoretical waveform template over many cycles, even a small contribution to the phase evolution can be detected and could give information about the NS structure.  Following~\\cite{FlanaganHinderer2008}, we calculate in Sec.~\\ref{sec:radiation} the measurability of the tidal deformability for a wide range of equal- and unequal- mass binaries, covering the entire expected range of NS masses and EOS, and with proposed detector sensitivity curves for second- and third- generation detectors. We show that the measurability of $\\lambda$ is quite sensitive to the total mass of the system, with very low-mass neutron stars contributing significant phase corrections that are optimistically detectable in Advanced LIGO, while larger-mass neutron stars are more difficult to distinguish from the $k_2=0$ case of black holes~\\cite{DamourNagar2009, BinningtonPoisson2009}. In third-generation detectors, however, the tenfold increase in sensitivity allows a finer discrimination between equations of state leading to potential measurability of a large portion of proposed EOSs over most of the expected neutron star mass range.  We conclude by briefly considering how the errors could be improved with a more careful analysis of the detectors and extension of the understanding of EOS effects to higher frequencies.  Finally, in the Appendix we compute the leading order EOS-dependent corrections to our model of the tidal effect and derive explicit expressions for the resulting corrections to the waveform's phase evolution, extending the analysis of Ref.~\\cite{FlanaganHinderer2008}. Estimates of the size of the phase corrections show that the main source of error are post-1 Newtonian corrections to the Newtonian tidal effect itself, which are approximately twice as large as other, EOS-dependent corrections at a frequency of 450 Hz. Since these point-particle corrections do not depend on unknown NS physics, they can easily be incorporated into the analysis. A derivation of the explicit post-Newtonian correction terms is the subject of Ref.~\\cite{Vinesflanagan}.  \\textit{Conventions}: We set $G=c=1$. ",
        "conclusions": "We have calculated the relativistic $l=2$ Love number $k_2$ and resulting tidal deformability $\\lambda$ for a wide range of realistic EOS in addition to polytropes.  These EOS have tidal deformabilities that differ by up to an order of magnitude in the mass range relevant for binary neutron stars. However, the estimated uncertainty $\\Delta\\tilde\\lambda$ for a binary neutron star inspiral at 100~Mpc using the Advanced LIGO sensitivity below 450~Hz is greater than the largest values of $\\tilde\\lambda$ except for very low-mass binaries.  The uncertainty for the Einstein Telescope, on the other hand, is approximately an order of magnitude smaller than for Advanced LIGO, and a measurement of $\\tilde\\lambda$ will rule out a significant fraction of the EOS.  Advanced LIGO can place a constraint on the space of possible EOS by obtaining a $95 \\%$ confidence upper limit of $\\tilde\\lambda(\\mathcal{M}, \\eta) \\lesssim 2 \\Delta\\tilde\\lambda(\\mathcal{M}, \\eta)$.  The tables in Sec.~\\ref{sec:radiation} can also be scaled as follows: For a network of $N$ detectors the uncertainty scales roughly as $\\Delta\\tilde\\lambda/\\sqrt{N}$, and for a closer signal we have $\\Delta\\tilde\\lambda (D/100\\textrm{ Mpc})$."
    },
    "0911/0911.4912_arXiv.txt": {
        "abstract": "{Outflows of young stars drive shocks into dusty, molecular regions. Most models of such shocks are restricted by the assumptions that they are steady and propagating in directions perpendicular to the magnetic fields. However, the media through which shocks propagate are inhomogeneous and shocks are not steady. Furthermore, only a small fraction of shocks are nearly perpendicular.} {We identify features that develop when a shock encounters a density inhomogeneity and ascertain if any part of the precursor region of a non-steady multifluid shock ever behaves in a quasi-steady fashion. If it does, some time-dependent shocks may be modelled approximately without solving the time-dependent hydromagnetic equations.} {We use the code employed previously to produce the first time-dependent simulations of fast-mode oblique C-type shocks including a self-consistent calculation of the thermal and ionisation balances and a fluid treatment of grains.} {Simulations were made for initially steady oblique C-type shocks, each of which encounters one of three types of density inhomogeneities. For a semi-finite inhomogeneity with a density larger than the surrounding medium's, a transmitted shock evolves from being of J-type to a steady C-type shock on a timescale comparable to the ion-flow time through it. A sufficiently upstream part of the precursor of an evolving J-type shock is quasi-steady. The ion-flow timescale is also relevant for the evolution of a shock moving into a region of decreasing density. The models for shocks propagating into regions in which the density increases and then decreases to its initial value cannot be entirely described in terms of the results obtained for monotonically increasing and decreasing densities.} {We present the first time-dependent simulations of dusty C-type shocks interacting with density perturbations. We studied the transient evolution of the shock structure and find that the initial interaction always produces a transition to a J-type shock. Furthermore, the long-term evolution back to a C-type shock cannot always be approximated by quasi-steady models.} ",
        "introduction": "Jets and winds associated with low and high mass proto-stellar objects interact with the molecular clouds surrounding them \\citep[e.g.][]{Aal07}. Shocks are driven into the clouds sweeping up cloud material and producing large scale (0.1 - 1~pc) molecular outflows with velocities of 10-100 km s$^{-1}$ \\citep[e.g.][]{BT99,RB01,SG09}. These outflows are often bipolar and are extremely common in young, low-mass stars \\citep{BL83}, as well as high-mass stars \\citep{C05,S05,LS09}.   As the fractional ionisation in molecular clouds is low \\citep[$\\chi < 10^{-6}$;][and reference therein]{D06}, the neutral gas and  magnetic field are weakly coupled. This coupling is mediated via ion-neutral and grain-neutral collisions, as the magnetic field forces the charged particles to move through the bulk of neutral particles. When a shock forms, the ion-neutral collisions accelerate, compress and heat the neutral gas upstream of the shock, forming a magnetic precursor. If the shock velocity is low and the postshock cooling efficient, the shock becomes continuous in all fluids and is then referred to as a C-type shock \\citep[][]{M71,D80}. Observational studies of the velocity widths of SiO emission suggest that most shocks near low mass proto-stellar objects are C-type \\citep[e.g.][]{MPal92}.  Previous studies of C-type shocks in dusty plasmas have focused on the shock structure under steady state conditions \\citep[e.g.][]{DRD83, PHH90, PH94, W98, G08}. However, SiO observations suggest that proto-stellar jets and winds interact with clumpy structures along their propagation axis \\citep{Mal92, Lal98, Jetal04}. Therefore, the steady state assumption must be relaxed. \\cite{CR02} were the first\\footnote{Time-dependent multifluid shock studies prior to \\cite{CR02},  e.g. those of \\cite{T94, SM97, S97} and \\cite{CPDFF98}, did not include grain dynamics.} to use a time-dependent multifluid magnetohydrodynamics (MHD) code, including dust grains dynamics, to study fast-mode C-type shocks. However, their study is limited to perpendicular shocks. The multifluid approach developed by \\cite{F03} overcomes this restriction.  Using the Falle method, \\citet[][hereafter Paper~I]{vLal09} produced the first time-dependent simulations of fast-mode, oblique C-type shocks that evolve to steady state  and include a selfconsistent calculation of the thermal and ionisation balances and a fluid treatment of grain dynamics.   While the time-dependent simulations of Paper I still focus on steady state shock structures, we now extend this work by studying the temporal evolution of fast-mode C-type shocks interacting with regions of inhomogeneous density. In Sect.~\\ref{sect:model} we describe the numerical model with the initial conditions. We then apply the code to an oblique shock propagating into a region of varying density (Sect.~\\ref{sec:results}) and, in Sect.~\\ref{sect:conc}, we discuss these results and give our conclusions. ",
        "conclusions": "\\label{sect:conc} In this paper we present the first time-dependent simulations of oblique C-type shocks in dusty plasmas interacting with density perturbations. We studied the transient evolution of the shock structure for each of several different types of density perturbations, i.e. increasing, decreasing and clump-like, and examined the dependence of this interaction on shock speed and density contrast.  When a steady C-type shock encounters a density perturbation, the shock reacts to the changing upstream condition by breaking up into multiple waves. While a J-type shock moves into or out of a dense region, a shock or rarefaction wave, respectively, is reflected and propagates in the reverse direction. This reflected shock or wave is present because the post-shock flow behind the transmitted shock is different from the post-shock flow of the initial shock. The relative strength of the transmitted and reflected components depends on the initial shock speed and the density contrast of the perturbation, as these control  the post-shock properties of the transmitted shock, e.g. a lower density contrasts leads to a smaller difference between the far downstream flow and the transmitted post-shock flow. Hence, a weaker rarefaction wave or shock is needed to adjust this change. For the interaction with a clump, where the density decreases again, the situation gets more complicated as the transmitted J-type shock forms a second subshock in the neutral flow of the precursor. It is thus important to realise that we expect a range of different J- and C-type shocks and multifluid rarefaction waves in inhomogeneous clouds.  In our models, we have focused on the evolution of the transmitted shocks. Such a shock is initially J-type and contains a subshock in the neutral flow. It again becomes a steady C-type shock but with a different shock width and velocity. The timescale for it to become steady is of the same order as the ion-flow time through the final shock structure, i.e. about 0.5-1.5 $\\times 10^5$ yr. An important consequence of these timescales is that, if a shock interacts with density inhomogeneities on timescales shorter than this, no steady C-type shock exists. All shocks are thus J-type, although some shocks will have weak subshocks and are then only marginally distinguishable from C-type shocks.  Our simulations show that the steady dusty C-type shock develops from a shock in which the sufficiently far upstream parts of the precursor are steady. Such an evolution was also observed for multifluid shocks in dustless, weakly-ionised gases \\citep{FP99}. Following the \\cite{FP99} result, \\cite{Lal04} developed a quasi-steady method to follow the temporal evolution of multifluid shocks. This method involves the treatment of an evolving non-stationary J-type shock as a sequence of truncated C-type shocks each of which has a neutral subshock at the point where the ion-flow time corresponds to the age of the shock. While this approach was only validated for shocks in dustless, weakly ionised gases \\citep{Lal04}, our result confirms quasi-steady models can also be used for dusty shocks and justifies the use of such methods for the temporal evolution of dusty shocks by \\cite{Getal08}. However, we also find some limitations of the quasi-steady approach when applied to the interaction of shocks with density perturbations. Firstly, this method can only be used after an early adaptation phase in which the shock adjusts to the changing upstream conditions. Secondly, this approach is only valid for a C-type shock interacting with a density perturbation. For a J-type shock, a second neutral subshock forms in the precursor. Behind this secondary shock, the velocities of the neutral and charged fluids are not the same. Such a situation, which is likely to occur in a clumpy medium, cannot be modelled with a quasi-steady approach.  In quiescent cold clouds, silicon is nearly depleted from the gas phase and stored in the core of dust grains. However, observations of SiO emission near protostellar objects show that some silicon is returned to the gas-phase. It is thought that this is due to sputtering of the dust grains and grain-grain collisions in shocks \\citep{MPal92}. Although these processes are not included in the current code, the change of the grain-neutral relative velocity serves as an indication. Grain mantle sputtering requires a grain-neutral relative velocity of $\\approx$ 6 km s$^{-1}$ whereas core sputtering requires a much higher relative velocity of $\\approx$ 19 km s$^{-1}$ \\citep{Cal97}. For our model parameters, with a shock speed of 25~km\\ s$^{-1}$, the threshold for core sputtering is only just attainable. Therefore, we do not expect much SiO emission from these models. However, grain mantle sputtering occurs in all the initial C-type shocks. \\cite{Jal08} have suggested that if the grain mantle consists of a small fraction of silicon, sputtering by heavy molecules such as CO could erode the mantles even in low velocity shocks and account for the narrow SiO line emission observed in some young outflows \\citep[e.g. L1448-mm][]{Jal04}. Silicon erosion from the mantle would be saturated for shock velocities between 10 and 20 km s$^{-1}$. This implies that the SiO enhancement from mantle erosion is saturated in all our models. Variations of the SiO emission then arise as the spatial area over which sputtering occurs changes (i.e. the shock width increases/decreases). Furthermore, we expect a small contribution to mantle sputtering from reflected waves. In subsequent papers we will include grain mantle and core sputtering as well as grain-grain collision terms to describe the release of SiO from the grains to the gas phase and use this to calculate the SiO emission \\citep{Cal97,Jal08}."
    },
    "0911/0911.0536_arXiv.txt": {
        "abstract": "{In a previous paper we presented a homogeneous and 92 \\% complete optical spectral dataset of the 3CR radio sources with redshift $<$ 0.3. Here we use the emission line measurements to explore the spectroscopic properties of the sample. The 3CR sources show a bimodal distribution of Excitation Index, a new spectroscopic indicator that measures the relative intensity of low and high excitation lines. This unveils the presence of two main sub-populations of radio-loud AGN to which we refer to, following previous studies, as High and Low Excitation Galaxies (HEG and LEG, respectively).  In addition to the two main classes, we find one source with a spectrum typical of star forming galaxies, and 3 objects of extremely low level of excitation.  All broad-line objects are HEG from the point of view of their narrow emission line ratios and all HEG are FR~II radio-galaxies with log $L_{178} $ [erg s$^{-1}$] $ \\gtrsim 32.8$.  Conversely LEG cover the whole range of radio power encompassed by this 3CR subsample (30.7 $\\lesssim$ log $L_{178} \\lesssim$ 35.4) and they are of both FR~I and FR~II type. The brightest LEG are all FR~II. HEG and LEG obey to two (quasi) linear correlations between the optical line and extended radio luminosities, with HEG being brighter than LEG in the [O~III] line by a factor of $\\sim$ 10. HEG and LEG are offset also in a plane that compares the black hole mass and the ionizing nuclear luminosity.  However, although HEG are associated with higher nuclear luminosities, we find LEG among the brightest radio sources of the sample and with a clear FR~II morphology, indistinguishable from those seen in HEG.  This suggests that LEG are not simply objects with a lower level of accretion.  We speculate that the differences between LEG and HEG are related to a different mode of accretion: LEG are powered by hot gas, while HEG require the presence of cold accreting material. The high temperature of the accreting gas in LEG accounts for the lack of ``cold'' structures (i.e. molecular torus and Broad Line Region), for the reduced radiative output of the accretion disk, and for the lower gas excitation. ",
        "introduction": "\\label{intro}  Radio galaxies (RG) are an important class of extragalactic objects for many reasons. Studies of radio-loud AGN are the key to understand the processes leading to the ejection of material in relativistic jets and its connection with gas accretion onto the central black holes, the way in which different levels of accretion are related to the process of jets launching, the origin of the AGN onset, and its lifetime. But the intense nuclear activity can also influence the star formation history and the properties of the ISM and ICM, thus representing a fundamental ingredient for the evolution of their hosts and their large scale environment.  RG have been historically classified according to their radio morphology, following the \\citet{fanaroff74} criteria: a FR~I source has bright jets rising from the nucleus, while a FR~II has two bright hot spots far from it. They also noted that FR~II are mostly found at high radio luminosities, while FR~I are associated to weaker radio sources.  The two FR classes also differ, at least statistically, from several other points of view, such as the environment \\citep{zirbel97} and host luminosities \\citep{govoni00}. However, it soon became apparent that the transition between the two classes is continuous and objects of intermediate radio structure do exist \\citep[e.g. ][]{capetti95}. A class of hybrid double sources, with a FR~I jet on one side and a FR~II lobe on the other, was also unveiled by \\citet{gopal00}. This supports explanations for the FR dichotomy based upon jet interaction with the external medium, arguing against interpretations based on intrinsic differences in the central engine. Furthermore, the multi-wavelength behavior of the nuclear emission in RG does not appear to be directly related to the differences between the two FR classes, but it is more closely linked to their spectroscopic nuclear properties \\citep[e.g.][]{chiaberge00,chiaberge:fr2}.  Optical spectroscopic information can clearly play a major role in gaining a better understanding of the properties of the central engines of RG. \\citet{heckman80} and \\citet{baldwin81} proposed to use optical line ratios as diagnostic tools to classify emission-line objects in general and AGN in particular.  They introduced diagnostic diagrams comparing selected emission line ratios, able to distinguish H~II regions ionized by young stars from gas clouds ionized by nuclear activity. Moreover AGN were separated into Seyferts and Low Ionization Nuclear Emission-line Regions \\citep[LINERs,][]{heckman80} based on the relative ratios of the optical oxygen lines ([O I]$\\lambda$6364, [O II]$\\lambda$3727, and [O III]$\\lambda$5007). Subsequently \\citet{veilleux87} revised the definition of the diagnostic diagrams, using only ratios of lines with small separation in wavelength, thus reducing the problems related to reddening as well as to uncertainties on the flux calibration of the spectra.  They used the following line combinations: [O~III]/H$\\beta$ as a function of [N~II]$\\lambda$6583/\\Ha, [S~II]$\\lambda\\lambda$6716,6731/\\Ha, and [O~I]/\\Ha.  The separation between AGN and HII regions, initially introduced empirically, was calibrated theoretically by \\citet{kewley01}. More recently, \\citet[ hereafter K06]{kewley06} selected a sample of $\\sim 85000$ emission line galaxies from the SDSS, finding that Seyferts and LINERs form separated branches on the diagnostic diagrams. They suggested that the observed dichotomy corresponds to the presence of two sub-populations of AGN associated with different accretion states.  An attempt to adopt a similar scheme for the optical classification focusing on radio-loud galaxies was made by \\citet{laing94} on a sub-sample of 3CR radio-galaxies, selected imposing z$<$0.88, V$<$20, and 0$^h <$ RA $< 13^h$. They put on firmer ground the original suggestion by \\citet{hine79} that FR~II sources can be distinguished into subclasses.  They proposed a separation into high excitation galaxies (HEG, defined as galaxies with [O~III]$/$\\Ha$>0.2$ and equivalent width (EW) of [O III] $>$ 3 \\AA) and low excitation galaxies (LEG).  \\citet{tadhunter98} found a similar result from an optical spectroscopic study of the 2Jy sample, in which a sub-class of Weak-Line Radio Galaxies (the sources with EW of [O III])$<10$\\AA) stands out due to a low ratio between emission line and radio luminosities as well of [O III]/[O II] line ratio.  Emission line luminosities show a broad connection with radio power, as verified by many studies \\citep[e.g.][]{baum89a,baum89b,rawlings89,rawlings91}.  This relation holds also for compact steep spectrum (CSS) \\citep{morganti97} and GHz peaked spectrum (GPS) sources \\citep{labiano08}.  \\citet{willott99} demonstrated that the dominant effect is a strong positive correlation between $L_{\\rm radio}$ vs $L_{\\rm line}$ and not to a common dependence of these quantities on redshift. This relationship points to a common energy source for both the optical line and the radio emission, and suggests the radio and line luminosities of RG are determined, to first order, by the properties of their central engines.  \\citet{baum89b} noticed that the large scatter (roughly one order of magnitude) in $L_{\\rm radio}$ for a given $L_{\\rm line}$ suggests that other factors do play a secondary role, such as the environment.  Another possibility is that the line and radio luminosities may be independently correlated with a third parameter, e.g. the amount of cold gas present on the kiloparsec scale.  A correlation between emission lines and core radio powers, although weaker than that observed for $L_{\\rm radio}$ and $L_{\\rm line}$, is also found \\citep[e.g. ][]{baum89b,rawlings89,rawlings91} suggesting that the total radio luminosity is separately correlated with the [O~III] and the core radio power.  \\citet{morganti92} extended the line-radio connection towards lower radio luminosities considering sources from the B2 sample, predominantly FR~I, noting a flattening in the correlation. This result is supported also by the analysis by \\citet{zirbel95}, who also report that FR~II sources produce about 5-30 times more emission line luminosity than FR~I for the same total radio power \\citep[see also][]{tadhunter98,wills04}.  In addition to the diagnostics derived from the narrow lines, optical spectra also reveal the presence of prominent permitted broad lines in a significant fraction of RG.  Unifying models for active galaxies ascribe the lack of broad lines to selective obscuration and differences of orientation between narrow and broad lined AGN \\citep[see ][]{antonucci93,urry95}.  However, despite the massive amount of spectroscopic data collected over the last decades for RG, there are still several key questions waiting for clear answers: are there indeed two (or more) distinct populations of RG? On which basis they can be separated? Which are the physical parameters setting the division? Which is the relationship between the optical and radio properties? When are broad emission lines observed?  Which processes set the presence and detectability of broad lines and how do they fit in the unified models for active nuclei? How do the spectroscopic properties of radio-loud and radio-quiet AGN compare?  A necessary ingredient required to enable a detailed exploration of spectroscopic properties of radio-loud AGN is a dataset as homogeneous and complete as possible.  This is one of the aims of the optical spectroscopic survey that we performed with the Telescopio Nazionale Galileo, presented in \\citet[ hereafter Paper I]{buttiglione09}.  We considered the 113 3CR sources at redshift $<$ 0.3, with an effective coverage of the sub-sample of 92 \\% (i.e. 104 sources).  For 18 sources the spectra are available from the Sloan Digital Sky Survey (SDSS) database \\citep{york00, st02, yip04}, Data Release 4,5,6. For these sources the observations provide uniform and uninterrupted coverage of the key spectroscopic optical diagnostics ratios. The data quality is such that the \\Ha\\ line is detected in all but 3 galaxies and in the majority ($\\sim 75$\\%) of the objects all key emission lines (i.e.  H$\\beta$, [O~III]$\\lambda\\lambda$ 4959,5007 \\AA, [O~I]$\\lambda\\lambda$ 6300,64 \\AA, H$\\alpha$, [N~II]$\\lambda\\lambda$ 6548,84 \\AA, [S~II]$\\lambda\\lambda$ 6716,31 \\AA) required to construct diagnostic diagrams are detected.  The attractiveness of the 3CR catalog of radio sources as a basis for such a study is obvious, being one of the best studied sample of RG. Its selection criteria are unbiased with respect to optical properties and orientation, and it spans a relatively wide range in redshift and radio power. A vast suite of observations is available for this sample, from multi-band HST imaging to observations with Chandra, Spitzer and the VLA, that can be used to address the issues listed above in a multiwavelength approach.  The paper is organized as follows: in Sect. \\ref{optclas} we derive and explore the spectroscopic diagnostic diagrams for the 3CR sample of sources with $z \\leq 0.3$. In Sect. \\ref{or} we compare optical (line and host) luminosity and radio emission. Results are discussed in Sect. \\ref{discussion} and summarized in Sect. \\ref{fine}, where we also present our conclusions.  Throughout, we adopted $H_o = 71$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{\\Lambda} = 0.73$ and $\\Omega_m = 0.27$. ",
        "conclusions": "\\label{fine}  We used measurements of emission lines for a 92 \\% complete survey of the 3CR sample of radio-galaxies with z $<$ 0.3 to explore their spectroscopic properties. The different steps used to classify the 3CR sources in the various sub-populations can be summarized as follows:  $\\bullet$ {\\it Excitation Index diagram:} we first considered a new spectroscopic indicator, the Excitation Index, defined as E.I. = log [O~III]$/$\\Hb - 1/3 (log [N~II]$/$\\Ha + log [S~II]$/$\\Ha + log [O~I]$/$\\Ha) and representing the relative intensity of high and low excitation lines. The presence of a bimodal distribution of E.I. allows us to define two main spectroscopic classes, LEG and HEG, with LEG being sources with E.I. $\\lesssim$ 0.95.  The drawback of this definition is that it requires the measurements of all 6 diagnostic lines.  $\\bullet$ {\\it Diagnostic diagrams:} the close correspondence between the excitation index and the location in the 3 `standard' diagnostic diagrams enabled us to define a class membership in a less ideal situation, i.e. when in addition to \\Ha, \\Hb\\ and [O III], only one or two lines among [S II], [O I] or [N II] can be measured, defining as LEG all sources for which:  log [O~III]/\\Hb\\ - log  [N~II]/\\Ha\\ $\\lesssim$ 0.7,  log [O~III]/\\Hb\\ - log  [S~II]/\\Ha\\  $\\lesssim$ 0.9, or  log [O~III]/\\Hb\\ - log  [O~I]/\\Ha\\ $\\lesssim$ 1.4.  $\\bullet$ {\\it Optical line-radio correlation.} We also used the $L_{\\rm [OIII]}$ vs $L_{178 MHz}$ correlation to extend the classification for objects with the only requirement of a [O~III] line measurement.  This is due to the rather well defined separation in line emission, at fixed radio power, between HEG and LEG. In particular we found  $-1 \\lesssim {\\rm log} ({L_{\\rm [O III]}}/\\nu L_{178}) \\lesssim  \\,\\, 0.5$ for HEG  $-2 \\lesssim {\\rm log} ({L_{\\rm [O III]}}/\\nu L_{178}) \\lesssim -0.5$ for LEG.  Using this strategy we were able to associate a spectroscopic class to 87 sources, representing 84 \\% of the sample. The breakdown in the various sub-classes is of 46 HEG (16 of which are BLO) and 37 LEG. In addition to the two main sub-classes, we found one object with a spectrum typical of star-forming galaxies, and 3 sources of extremely low excitation (ELEG). 17 galaxies remain spectroscopically unclassified.  None of the unclassified sources is consistent with being a HEG, as they all have very low values of $L_{\\rm [OIII]}/\\nu L_{178}$.  However, we cannot conclude in general that they are LEG, since they could belong to the ELEG class.  The ELEG are particularly intriguing, being well offset from the rest of 3CR sample in all diagnostic diagrams. They are also characterized by very low values of the ratio between [O~III] and radio luminosity as well as of core dominance. We defer to a forthcoming paper a more detailed analysis of these sources.  The dual population of HEG and LEG is reminiscent of the Seyfert and LINERs subclasses found in the SDSS emission line galaxies. Indeed, at zero-th order, the spectroscopic properties of radio-loud AGN of the 3CR and the (mostly) radio-quiet AGN of the SDSS are rather similar. However, looking at finer details, the separation between LEG and HEG shows an upward scatter by $\\sim$0.2 dex in the [O III]/\\Hb\\ ratio with respect to the transition from LINER to Seyfert. Furthermore, the 3CR sources are concentrated along the edges of the SDSS density distribution in the spectroscopic diagnostic diagrams. Both results could be related to the substantial mismatch in luminosity between the two samples, the 3CR being brighter by an average factor of 30 in emission line than the SDSS sources. However, we cannot exclude that the location of SDSS and 3CR source is an indication of a genuine difference in the spectroscopic behavior of radio-quiet and radio-loud AGN. The analysis of radio-loud AGN of lower luminosity can be used to test these alternative hypothesis.  Let us now summarize the similarities and differences between the two main spectroscopic classes of RG.  All broad-line objects are HEG from the point of view of their narrow emission line ratios, but no LEG shows broad line in its spectrum.  While all HEG are associated with FR~II radio-galaxies, LEG are of both FR~I and FR~II type.  HEG are only found in relatively bright radio sources, log $L_{178} \\gtrsim 32.8$. Instead LEG cover the whole range of radio power (30.7 $\\lesssim$ log $L_{178} \\lesssim$ 35.4), the brightest of them being FR~II.  HEG and LEG obey to two linear correlations between line and radio emission, with a slope consistent with unity and a rms $\\lesssim 0.5$ dex. HEG are brighter than LEG by a factor of $\\sim 10$ in [O~III] line.  Conversely, the substantial superposition of the host galaxy luminosity of the two classes suggests that the distribution of their black hole masses is not strongly different, with the vast majority of the objects confined in the range 8.5 $\\lesssim$ log $M_{\\rm BH}$ / $M_{\\odot} \\lesssim 9.5$. Considering only the LEG with a FR~II morphology, a further similarity with HEG emerges from the analysis of their radio structure. LEG/FR~II extend over the same range of radio power of HEG and have the same level of radio core dominance. The fact that the radio properties of LEG/FR~II and HEG are essentially indistinguishable is an indication that the two classes share the same range of jet power.  While the differences in terms of line intensities and line ratios (as well as in optical and IR nuclear luminosities) can be interpreted in a framework in which LEG are associated with a lower accretion rate than HEG, the similarity of the radio properties of the two sub-populations (having established that they do not differ in terms of black hole mass) would require that the process of jet launching is essentially decoupled from the level of accretion. This is the case, for example, of the Blandford-Znajek process that envisages jets powered by the extraction of rotational energy from a spinning black hole.  We instead speculated that the separation between LEG and HEG is due to a different mode of accretion: while HEG are powered by cold gas, LEG accrete hot gas, a mechanism that has been already demonstrated to be able to account for the activity of FR~I radio-galaxies.  The high initial temperature of the in-flowing gas in LEG prevents the formation of the various ``cold'' AGN structures, such as molecular tori and Broad Line Regions, seen in HEG but not in LEG.  The radiative output of a ``hot'' accretion disk is reduced with respect to a standard geometrically thin disk, at a given accretion rate, due to the lower cooling at these temperatures and it will be emitted mostly at high energies. The harder spectrum and the lower number of ionizing photons gives rise to emission lines of lower luminosity and excitation, producing the spectral separation in lines properties between LEG and HEG. Since the mechanism of jet launching operates in the innermost regions of the accretion flow, it is likely that the memory of the initial gas conditions are effectively canceled and that the jet power is sensitive only to the final accretion rate.  The attractiveness of this scheme is that it can reproduce simultaneously, by only varying the initial temperature of the accreting gas, the various differences between LEG and HEG. Clearly this hypothesis must be analyzed in greater detail, exploring the properties of the proposed hot accretion flows combined to a high accretion rate. In particular a crucial requirement is that the coronal accreting gas must be able to remain hot during the inflow. The proposed scenario can also be tested and constrained using the limits on the nuclear LEG emission in the various observing bands. It is also of great interest to establish whether LEG can be found in objects of at most moderate redshift (and, consequently, power) such as those considered here, or, conversely, can be associated even with radio galaxies of the highest luminosity.  Summarizing, we found that the radio-galaxies in the 3CR sample are composed by two main populations, separated on the basis of their emission line properties. However, this does not correspond simply to a difference in their radio properties. We propose a scenario in which the two classes are associated with a different mode of accretion, set by the initial temperature of the inflowing gas.  {\\sl Acknowledgments.} SB and ACe acknowledge the Italian MIUR for financial support. This publication makes use of data products from the Two Micron All Sky Survey, which is a joint project of the University of Massachusetts and the Infrared Processing and Analysis Center/California Institute of Technology, funded by the National Aeronautics and Space Administration and the National Science Foundation."
    },
    "0911/0911.0700_arXiv.txt": {
        "abstract": "We report on recent XMM-Newton observations, archival radio continuum and CO data, and SED modeling of the unidentified Galactic plane source HESS\\,J1708-410. No significant extended X-ray emission is observed, and we place an upper limit of $3.2\\times10^{-13} \\rm \\, erg \\, cm^{-2} \\, s^{-1}$ in the 2-4 keV range for the region of TeV emission.  Molonglo Galactic Plane Survey data is used to place an upper limit of 0.27 Jy at 843 MHz for the source, with a 2.4 GHz limit of 0.4 Jy from the Parkes survey of the southern Galactic plane. $^{12}$CO (J 1$\\rightarrow$0) data of this region indicates a plausible distance of 3 kpc for HESS\\,J1708-410. SED modeling of both the H.E.S.S.\\ detection and flux upper limits offer useful constraints on the emission mechanisms, magnetic field, injection spectrum, and ambient medium surrounding this source. ",
        "introduction": "Since beginning operation in 2004 the H.E.S.S. Cherenkov telescope array has proven a valuable tool in the investigation of extremely energetic phenomena, yielding $\\sim 50$ very-high-energy (VHE) Galactic gamma-ray sources. Roughly half of these sources lack definitive counterparts at longer wavelengths, and a few lack \\emph{any} plausible counterpart. The firmly identified Galactic TeV sources consist of Supernova remnants (SNRs), Pulsar Wind Nebulae (PWNe) and high-mass X-ray binaries (HMXBs). There are also plausible associations with young and massive stellar clusters \\citep{aet07}, and dense molecular clouds \\citep{aet06a}. Those TeV sources with no low energy counterpart may be underluminous (at longer wavelengths) members of these source classes, or may represent an entirely new class of TeV source. Understanding these so called `dark accelerators' with no identified counterpart at other wavebands is one of the more intriguing challenges of high energy astrophysics.  Gamma-rays in the VHE regime invariably arise from nonthermal particle populations, though for a given source determining if this population is dominantly leptonic or hadronic often proves quite challenging. For leptons in low density environments, inverse-Compton (IC) upscattering of background photons off highly relativistic electrons dominates, while bremsstrahlung becomes important only amongst very dense surroundings. For any scenario in which the gamma-ray emission arises from relativistic electrons, significant X-ray synchrotron emission can be expected, with a peak synchrotron flux comparable to the peak IC flux for typical galactic magnetic field strengths and radiation densities. This synchrotron emission usually extends down to radio energies with emission from lower energy, $\\cal{O}$(1 GeV), electrons. Yet if the electron spectrum within a source is truncated at low energies a bright TeV source may be unaccompanied by a radio counterpart. Such a situation is plausible for PWNe, given that the lowest particle energy is determined primarily by the four velocity of the pulsar wind upstream of the termination shock \\citep{kc84b} which is in turn linked to observable shock properties such as synchrotron luminosity and shock radius.  A hard electron spectrum with index greater than the canonical value of $-2$ might also result in an underluminous radio counterpart given the relative lack of low energy electrons. Hadronic acceleration results in the production of neutral pions via proton-proton interactions, which decay to produce VHE gamma rays.  Proton interactions also produce charged pions and hence secondary electrons that in turn produce synchrotron radiation in the radio to X-ray range, albeit usually at a much lower level than in leptonic scenarios. Nevertheless, in both scenarios, at least a weak lower energy counterpart is expected. Probing this lower energy regime provides insights into the particle population, ambient medium properties, and magnetic field strength.  HESS\\,J1708$-$410 (hereafter J1708) is one of the relatively few VHE gamma-ray sources for which \\emph{no} plausible explanation has yet been put forward. It was discovered during the H.E.S.S.\\ survey of the inner Galaxy, as described by \\citep{aet06c}. Subsequent observations improved the statistical significance to 11$\\sigma$ over a total livetime of 39 hours \\citep{aet08}.  J1708 is only slightly extended with respect to the H.E.S.S.\\ point spread function, consistent with a marginally elongated Gaussian of approximately $0.06^{\\circ} \\times 0.08^{\\circ}$.  \\citet{aet08} report no significant emission beyond 0.3$^{\\circ}$, ruling out an association with the nearby SNR G\\,345.7$-$0.2. The nearest known pulsar (PSR J1707--4053) lies exterior to our XMM field of view $\\sim 15 \\arcmin$ northwest of the H.E.S.S.\\ centroid, and is several orders of magnitude too weak ($\\dot E \\sim 4\\times 10^{32} \\, \\rm erg \\, s^{-1}$) and too old ($\\tau_c \\sim 5 \\times 10^{6}$ years) to conceivably power the TeV emission. \\citet{aet08} fit the H.E.S.S.\\ source spectrum over the energy range of 0.50--60 TeV and an aperture of 0.3$^{\\circ}$ to fully enclose the source.  This yields a power-law index of $-2.5$ and a flux of $\\approx$ 8$\\times 10^{-12} \\rm \\,erg \\, cm^{-2} \\, s^{-1}$.  An archival ASCA pointing of this region shows only a single point source located over a degree from the H.E.S.S.\\ source.  In addition, a previous XMM-Newton exposure of  G\\,345.7$-$0.2 revealed no X-ray emission at the H.E.S.S.\\ position, though this pointing placed J1708 close to the edge of the field of view.  In this paper we report on new XMM-Newton observations of this enigmatic VHE source. ",
        "conclusions": "Independent of any model, the VHE luminosity of $L \\approx  9 \\times 10^{33} \\, d_{3}^{2} \\rm \\, erg \\, s^{-1}$ indicates that J1708 is a prodigious gamma-ray emitter. For comparison, the Vela X nebula surrounding the Vela pulsar is observed to have gamma-ray luminosity of only $L_{0.6-65 \\, \\rm TeV} \\approx 10^{33} \\, \\rm erg \\, s^{-1}$ \\citep{aet06b}. This would imply a high gamma-ray luminosity if J1708 originates from a PWN. We compare to the Vela PWN rather than the standard candle of VHE astronomy, the Crab, since the lack of X-rays and hence low magnetic field implies an older PWN than 1 kyr. The leptonic scenario requires fine-tuning of low and high energy cutoffs and magnetic field to match the observed H.E.S.S.\\  flux and lower energy upper limits. While hadronic injection requires less manipulation of injection spectrum and source properties, the required energy in protons of $E_p = 8\\times10^{50} \\, (n/1 \\, \\rm cm^{-3}) \\,$$d_3^2$ erg is quite high for low density environs. Yet if the H.E.S.S.\\ detection stems from protons interacting with a molecular clouds of density $\\sim 100 \\rm \\, cm^{-3}$, this energy requirement is significantly relaxed.  Regardless of injection species the diffusion and cooling timescales for VHE particles hint at an age of 10-50 kyr for magnetic fields in the range of $3-5 \\rm \\, \\mu G$.  The currently available data do not allow us to conclusively pin down the nature of the emitting particles in J1708. A Fermi LAT detection would provide a significant boost to understanding this source, and implies a hadron accelerator as is apparent in Figures 4 and 5. Deeper radio images would also help us to elucidate the processes at work in J1708.  Further lowering the radio upper limit would constrain the spectral index of the source, and/or the low energy cutoff for leptons; for hadrons the product of age and $B^{3/2}$ would be constrained. J1708 is particularly interesting among unidentified TeV sources as it has no lower energy candidate counterparts whatsoever and is rather compact. These features, combined with its high gamma-ray luminosity, could point to a very different type of VHE system than those we currently know.  We thank NASA for supporting this research with grant NNX06AH66G."
    },
    "0911/0911.3779_arXiv.txt": {
        "abstract": "We report on two quantitative, morphological estimators of the filamentary structure of the Cosmic Web, the so-called \\emph{global} and \\emph{local} skeletons. The first, based on a global study of the matter density gradient flow, allows us to study the connectivity between a density peak and its surroundings, with direct relevance to the anisotropic accretion via cold flows on galactic halos. From the second, based on a local constraint equation involving the derivatives of the field, we can derive predictions for powerful statistics, such as the differential length and the  relative saddle to extrema counts of the Cosmic web as a function of density threshold (with application to percolation of structures and connectivity), as well as a theoretical framework to study their cosmic evolution through the onset of gravity-induced non-linearities. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0911/0911.4645_arXiv.txt": {
        "abstract": "We use \\textit{Hubble Space Telescope} Fine Guidance Sensor astrometry and high-cadence radial velocities for HD\\,136118 from the HET with archival data from Lick to determine the complete set of orbital parameters for HD\\,136118\\,b.  We find an orbital inclination for the candidate exoplanet of $i_{b} = 163.1\\degr\\pm3.0\\degr$. This establishes the actual mass of the object, $M_{b} = 42^{+11}_{-18} M_{J}$, in contrast to the minimum mass determined from the radial velocity data only, $M_{b}\\sin{i} \\sim12\\,M_{J}$. Therefore, the low-mass companion to HD 136118 is now identified as a likely brown dwarf residing in the ``brown dwarf desert''. ",
        "introduction": "\\par Among the hundreds of exoplanets detected so far, only a few of them have their actual mass known. The widely used Doppler spectroscopy technique yields the radial component of the stellar perturbation velocity only. Consequently, the inclination of the orbital plane is unknown and only the minimum mass of a companion may be determined. To obtain a companion's true mass it is necessary to make use of additional techniques.  \\par The first precise determination of an exoplanet mass was made for a transiting system \\citep{henry2000}. However, transits are observed to occur only for systems that are oriented edge-on, and they have only a reasonable probability of occurrence for close-in planets (semimajor axes less than about 0.1\\,AU). Another way to determine the orbital inclination of an unseen companion is by measuring the stellar reflex motion astrometrically. The first astrometrically determined mass of an exoplanet by \\cite{benedict2002b} was possible thanks to the high precision of the Fine Guidance Sensor (FGS) instrument on the \\textit{Hubble Space Telescope} (\\textit{HST}). The FGS provides per observation precisions of better than $1$\\,mas for small angle relative astrometry. This unique capability enables the detection of stellar perturbations due to planetary mass companions in wide orbits.  \\par We were granted observing time with the \\textit{HST} to measure the perturbation and determine the true mass of HD\\,136118\\,b, which is an exoplanet candidate found by radial velocity measurements \\citep{fischer2002}.  To supplement previously published data and provide better constrain on the companion's spectroscopic orbital parameters we also obtained high-cadence radial velocity measurements with the Hobby-Eberly Telescope (HET). In this paper we present the results of our analysis and it is arranged as follows.  in Section \\ref{sec:1} we review stellar properties for HD\\,136118. In Section \\ref{sec:2} we discuss the instrumental set up and data reduction for both spectroscopy and astrometry. In Section \\ref{sec:3} we describe the orbital model used to analyze this data. In Section \\ref{sec:4} we discuss the strategy employed to obtain the system parameters and present our results. In Section \\ref{sec:5} we summarize and discuss the consequences of our results. ",
        "conclusions": "\\label{sec:5}  HD\\,136118 is a solar type star with a brown-dwarf companion. Table \\ref{tab:allpars} shows a summary of all observed orbital elements of the system.  We found that HD\\,136118\\,b has an orbital inclination of $i = 163.1\\degr\\pm3.0\\degr$, nearly perpendicular to the inferred inclination of the stellar spin axis. This misalignment is an intriguing result since conservation of angular momentum would favor alignment between stellar spin and companion orbital axes, assuming both were formed in the same primordial cloud.  HD\\,136118\\,b is likely a brown dwarf companion orbiting at $2.36\\,$AU that falls in the driest region of the so called `brown dwarf desert' \\citep{grether2006}.  They showed that the frequency of companions in the stellar mass range follows a slope with gradient $-9.1\\pm2.9$, while in the planetary mass region, the gradient is $24.1\\pm4.7$. These two separate linear fits intersect below the abscissa at $M=43^{+14}_{-23}\\,M_{J}$. Surprisingly, HD\\,136118\\,b mass is  $M_{b}=43^{+11}_{-18}\\,M_{J}$. \\cite{reffert2006} measured  astrometric masses for two exoplanet candidates HD\\,38529\\,c ($M = 37^{+36}_{-19}\\,M_{Jup}$) and HD\\,168443\\,c ($M = 34\\pm12\\,M_{Jup}$). Both are likely brown dwarf companions around solar type stars like HD\\,136118\\,b. These objects are important cases for studying the mass function at the brown dwarf mass range.  According to the evolutionary dusty model for brown dwarfs of \\cite{baraffe2001} and \\cite{chabrier2000}, assuming $M_{b} = 0.041^{+0.010}_{-0.017}\\,M_{\\odot}$  and the age of the brown dwarf as $5$\\,Gyr, HD\\,136118\\,b has a temperature of about $T_{b} = 900$\\,K and a radius of $R_{b} = 0.086\\,R_{\\odot}$. If one considers the uncertainty in the age of the system (see Table \\ref{tab:starpar}), this brown dwarf may be much younger, therefore considering the age as $1$\\,Gyr, HD\\,13118\\,b has a temperature of about $T_{b} = 1200$\\,K and  a radius of $R_{b} = 0.1\\,R_{\\odot}$. These characteristics classifies HD\\,136118\\,b as a T dwarf. Using these values we calculate the emission and reflection spectra, and the flux ratio between the brown dwarf and the parent star as show in Fig \\ref{fig:flux-ratio}. The flux ratio increases toward the far infrared (L, M and N bands), where it can get as high as $10^{-4}$.  The astrometric determination of the mass of a low mass companion can decisively characterize it as a planet.  A good illustration of this fact can be seen from the results of our group for three objects that were previously listed as exoplanet candidates: Gliese\\,876\\,b, HD\\,136118\\,b and HD\\,33636\\,b.  Surprisingly each object has been found to belong to a different class: a giant planet, a brown dwarf and a M dwarf star, respectively, \\cite{benedict2002b}, this paper, and \\cite{bean2007}. These results demonstrate the importance of the application of complementary techniques in observing extrasolar planetary systems."
    },
    "0911/0911.1126_arXiv.txt": {
        "abstract": "We use $\\sim$8,600 $>5\\times10^{10}~M_\\odot$ COSMOS galaxies to study how the morphological mix of massive ellipticals, bulge-dominated disks, intermediate-bulge disks, bulge-less disks and irregular galaxies evolves from $z=0.2$ to $z=1$. The morphological evolution depends strongly on mass. At $M>3\\times10^{11}~M_\\odot$, no evolution is detected in the morphological mix: ellipticals dominate since $z=1$, and the Hubble sequence has quantitatively settled down by this epoch. At the $10^{11} M_\\odot$ mass scale, little evolution is detected, which can be entirely explained with major mergers. Most of the morphological evolution from $z=1$ to $z=0.2$ takes place at masses $5\\times10^{10} - 10^{11}~M_\\odot$, where: $(i)$ The fraction of spirals substantially drops and the contribution of early-types increases. This increase is mostly produced by the growth of bulge-dominated disks, which vary their contribution from $\\sim10\\%$ at $z=1$ to $>30\\%$ at $z=0.2$ (cf.\\ the elliptical fraction grows from $\\sim15\\%$ to $\\sim20\\%$). Thus, at these masses, transformations from late- to early-types result in disk-less elliptical morphologies with a statistical frequency of only $30\\% - 40\\%$. Otherwise, the processes which are responsible for the transformations either retain or produce a non-negligible disk component. $(ii)$ The bulge-less disk galaxies, which contribute $\\sim15\\%$ to the intermediate-mass galaxy population at $z=1$, virtually disappear by $z=0.2$. The merger rate since $z=1$ is too low to account for the disappearance of these massive bulge-less disks, which most likely grow a bulge via secular evolution. ",
        "introduction": "Over the last decade, it has been realized that halo and galaxy mass play a crucial role in galaxy evolution. This has been found both observationally \\citep[see e.g.][]{kauffmann03,thomas05,baldry06,bamford09, tasca09,cucciati09,bolzonella09,iovino09,kovac09} and also theoretically \\citep[][]{hahn07a,hahn07b,hahn09,skibba09,crain09}. In particular, the mass assembly rate of galaxies peaks at steadily lower masses \\citep[e.g.][]{cimatti06,scar07a,scar07b,pozzetti09}. Early-type galaxies  at mass scales $\\sim5\\times10^{10}$ \\msol\\ progressively increase their contribution to the global galaxy mass function (MF) from $z=1$ down to today, while the most massive early-type galaxies, with masses at and above $\\sim10^{11}$ \\msol\\ are mostly already in place by $z\\sim1$ \\citep[e.g.][]{zucca06,bundy06,francescini06,caputi06,scar07a,scar07b,ilbert09,pozzetti09,aller09}.  In this paper\\footnote{Throughout the paper we adopt a standard $\\Lambda$CDM cosmology with $\\Omega_M=0.25, ~\\Omega_\\Lambda=0.75$, and $h=0.7$. Magnitudes are given in the AB system \\citep{okeg83}.} we push the investigation of the mass growth  of massive galaxies one step forward, and specifically ask {\\it (i)} what is the detailed morphological mix,  since $z\\sim1$, at all mass scales above $\\sim5\\times10^{10}$ \\msol\\ (the threshold mass at which  $z=0$ galaxies transition between predominantly late- and early-type properties, e.g., Kauffmann et al.\\ 2003); and, {\\it (ii)} at what epoch is the massive galaxy population already fully assembled into the final morphological mix that we observe in the $z=0$ Hubble sequence.  Specifically, exploiting the large area and extensive ancillary datasets of the Cosmic Evolution Survey \\citep[COSMOS;][]{scoville07}, and the detailed morphological binning provided by the {\\it Zurich Estimator of Structural Types (ZEST)} classification (Scarlata et al.\\ 2007), we  split the $\\sim 8600$ $I_{AB}\\le24$ COSMOS galaxies above a completeness mass cut of $5\\times10^{10} M_\\odot$ into five separate morphological bins, and we  study the evolution of the galaxy  morphological mix across the redshift range $z\\sim1-0.2$ as a function of galaxy mass {\\it and} detailed galactic structure. ",
        "conclusions": "Our analysis shows that the evolution of the morphological distribution in galaxies since $z\\sim1$ depends strongly on stellar mass, with little or no evolution at mass scales above $M\\sim10^{11}$ \\msol\\ and detectable evolutionary trends at lower masses ($M=5\\times10^{10}-10^{11}$ \\msol).   At $z\\sim1$ between $5\\times10^{10}$ and $10^{11}~M_\\odot$, there exists a non-negligible fraction (15\\%) of bulge-less disk galaxies. Their relative abundance decreases rapidly towards lower redshift, essentially disappearing by $z\\sim0.4$. In contrast with all other main types, these bulge-less systems grow considerable amounts of stellar mass since $z\\sim1$ through internal star-formation.  At all masses above our completeness limit,  the early-type population grows in importance towards lower redshifts. Interestingly, we find that this growth is dominated by galaxies containing a non-negligible disk component.  Although traditionally grouped together into the single class of {\\it early-type} galaxies, the disk-less spheroids (i.e., elliptical galaxies, $E$) and massive bulges embedded in a non-dominant disk component (the $B$ galaxies in our paper) most likely have a different origin. In the simple-minded framework in which late-type galaxies are directly transformed into `early-type' galaxies, about $60-70$\\% of such transformations result in bulge-dominated $B$-type disk galaxies. The physical mechanisms responsible for these late-to-$B$ transformations must be capable of building or retaining a dynamically-cold component and different mechanisms must be at play in the remaining $\\sim30\\%$ of cases in order to produce the disk-less $E$ galaxies.  At mass scales above  $\\sim10^{11}$\\msol, a simple but robust merger-tree model can reproduce the redshift evolution of the mass fraction of galaxies of all morphological types simultaneously: Mergers with stellar mass ratios down to 1/4 can, without any need of invocation of more complex scenarios, reproduce the morphological mix at $z\\sim0.2$ at the high end of the galactic mass scales, where the bulgeless disk galaxies are only a negligible contribution.  In contrast, at mass scales below $10^{11}$\\msol, neither the evolution of the fractions of early type populations  ($E$ and $B$), nor that of  the bulge-less disks, can be reproduced within a $\\Lambda$-CDM hierarchical picture including only mergers as the responsible mechanism for morphological transformations. Most prominently, the disappearance of the bulge-disk galaxies is not explained by mergers of any mass ratio, down to small accretion events. Global disk dynamical instabilities are thus likely to be active in building bulge components through inward transfer of mass \\citep[see e.g.][]{kk04}, thereby adding substantial stellar mass while also inducing the galaxy to migrate to an earlier-type morphological class. Secular evolution thus appears to be key towards achieving a complete recipe for the formation of the $z=0$ Hubble sequence at mass scales below $10^{11}$ \\msol."
    },
    "0911/0911.3409_arXiv.txt": {
        "abstract": "{} {We investigate the structure of the interstellar medium (ISM) and identify the location of possible embedded excitation sources from far-infrared (FIR) line and mid-infrared continuum emission maps.} {We carried out imaging spectroscopic observations of four giant Galactic star-forming regions with the Fourier Transform Spectrometer (FTS) onboard AKARI.   We obtained \\oiii\\ 88\\micron\\ and \\cii\\ 158\\micron\\ line intensity maps of all the regions: G3.270-0.101, G333.6-0.2, NGC~3603, and M17.} {For G3.270-0.101, we obtained high-spatial-resolution \\oiii\\ 88\\micron\\ line-emission maps and a FIR continuum map for the first time, which imply that \\oiii\\ 88\\micron\\ emission identifies the excitation sources more clearly than the radio continuum emission.  In G333.6-0.2, we found a local \\oiii\\ 88\\micron\\ emission peak, which is indicative of an excitation source.  This is supported by the 18\\micron\\ continuum emission, which is considered to trace the hot dust distribution.  For all regions, the \\cii\\ 158\\micron\\ emission is distributed widely as suggested by previous observations of star-forming regions.} {We conclude that \\oiii\\ 88\\micron\\ emission traces the excitation sources more accurately than the radio continuum emission, especially where there is a high density and/or column density gradient.  The FIR spectroscopy provides a promising means of understanding the nature of star-forming regions.} ",
        "introduction": "Far-infrared (FIR) spectroscopy of star-forming regions provides us with a great deal of information about the interstellar medium (ISM) without being affected by the uncertainty of the extinction correction.  FIR forbidden lines are less sensitive to the electron temperature than optical forbidden lines in \\hii\\ regions, and observations with balloon-borne and airborne telescopes have been used to determine physical properties such as the electron density and elemental abundance \\citep{Rubin88,Simpson95,Simpson04}.  The Long-Wavelength Spectrometer \\citep[LWS;][]{Clegg96} onboard the Infrared Space Observatory \\citep[ISO;][]{Kessler96} provided continuous spectra of the wavelength range 43--197\\micron, enabling us to investigate the physical conditions of the gas in the ionized and/or photodissociation regions (PDRs) by measuring several emission lines and the continuum emission (\\citealt{MartinHernandez02}, see \\citealt{Abergel05} and \\citealt{Peeters05} for reviews).  Since the LWS was a single-beam instrument, raster-scan observations were needed to investigate the spatial distribution of emission spectra and physical properties \\citep[e.g.][]{Rodriguez01,Mizutani02,Okada03,Okada06}.  In the present paper, we present FIR imaging spectroscopic observations with the Fourier Transform Spectrometer (FTS) onboard AKARI \\citep{Murakami07,Kawada08}, the Japanese infrared astronomical satellite launched in 2006.  The FTS is the spectroscopic part of the Far-Infrared Surveyor \\citep[FIS;][]{Kawada07}.  Since the spectroscopic capability of the Multiband Imaging Photometer for {\\it Spitzer} (MIPS) onboard {\\it Spitzer} \\citep{Werner04} is restricted to wavelengths of $<97$\\micron\\ at low spectral resolutions of $R=15$--$25$, the FIS-FTS is a unique FIR spectroscopic instrument in space for studying FIR forbidden-line emission after ISO/LWS and before the Herschel Space Observatory \\citep{Pilbratt08}.  With two dimensional arrays, the FIS-FTS has the capability of obtaining spectral maps over a wide area in a few pointing observations. ",
        "conclusions": "\\centering \\begin{tabular}{ccccccc} \\hline Target & Number of & Observation ID & Observation Date & AOT & Reset interval & Resolution\\\\ & pointings & & (UT) && [s] & \\\\ \\hline G3.270-0.101 & 1 & 5110044 & Mar. 20 2007 & FIS03 & 0.25 & Full \\\\ G333.6-0.2 & 2 & 5110052/53 & Mar. 3 \\& 4 2007 & FIS03 & 0.1 & Full\\\\ NGC3603 & 3 & 1403056/57/59 & Jul. 21,22,23 2007 & FIS03 & 0.1 & Full\\\\ M17 & 1 & 5110024 & Sep. 27 2006 & FIS03 & 0.1 & Full\\\\ \\hline \\end{tabular} \\end{table*}  The observations are summarized in Table~\\ref{obs_summary}.  All observations were carried out with the Astronomical Observation Template (AOT) of FIS03, which is a spectroscopic mode of the AKARI/FIS.  The FIS-FTS adopts a Martin-Puplett-type interferometer \\citep[see][for details of the instrument]{Kawada08}.  To investigate emission lines, we use the full-resolution mode with the spectral resolution of 0.19\\pcm\\ without apodization and 0.38\\pcm\\ with `Hanning' apodization.  We obtain 7 sets of forward and backward scans of the moving mirror with a total $\\sim 11$ minute integration time for each pointing observation.  The FIS has a short wavelength array, SW, and long wavelength array, LW, each of which has narrow and wide bands.  In the spectroscopic mode, only a wide band of each array is used.  They are $3\\times 20$ and $3\\times 15$ pixel arrays that cover 85--130\\pcm\\ and 60--88\\pcm, on pixel scales of 26.8\\arcsec and 44.2\\arcsec, respectively \\citep[][Murakami et al. in prep.]{Kawada07,Kawada08}.  Some of the present observations were carried out as Director's Time observations, and the others as the Mission Program of ISMGN (Interstellar dust and gas in various environments of out Galaxy and nearby galaxies: P.I. H. Kaneda).  The field-of-view (FOV) of SW and LW is $1.5^\\prime \\times 9.8^\\prime$ and $2.4^\\prime \\times 12.2^\\prime$, respectively.  Because of the misalignment of the arrays, the FOVs of SW and LW are not completely matched with each other and cover slightly different areas of the sky \\citep{Kawada07}. The point spread function (PSF) of the FIS-FTS can be described well by a combination of two Gaussian profiles, which are located along the major axis of the detectors with the separation of $14^{\\prime\\prime}$ in each detector array \\citep{Kawada08}.  The FWHM of each Gaussian is $39^{\\prime\\prime}$ and $44^{\\prime\\prime}$ along the minor and major axis, respectively, for SW, and $53^{\\prime\\prime}$ and $57^{\\prime\\prime}$ along the minor and major axis, respectively, for LW.  The position accuracy of the FTS mode is less than a half of the pixel size (Murakami et al. in prep.).  \\subsection{Targets}  To take advantage of the imaging spectroscopy of the FIS-FTS, we selected four active star-forming regions that are bright and extended enough to detect the interferogram with a high signal-to-noise ratio (S/N) over an area of $\\sim 10$\\arcmin. %  G3.270-0.101 is a giant \\hii\\ region located 14.3~kpc from the Sun.  It was observed in the radio and in the mid-infrared (MIR) with the Mid-course Space Experiment (MSX) \\citep{Conti04}.  The number of the Lyman continuum photons $N(\\mathrm{Lyc})$ is estimated to be $3\\times 10^{50}$\\,s$^{-1}$.  \\citet{Okada08} detected several MIR ionic fine-structure lines with the Infrared Spectrometer (IRS) onboard {\\it Spitzer} in this region.  We completed one pointing observation, covering the positions observed by the IRS.  G333.6-0.2 is a bright \\hii\\ region at radio and infrared wavelengths \\citep{Goss70,Becklin73,Hyland80,Fujiyoshi98,Fujiyoshi01,Fujiyoshi05,Fujiyoshi06} that is totally obscured in the visible.  It is located at 3.1~kpc \\citep{Conti04} and likely to be excited by a cluster of O and B stars \\citep{Fujiyoshi98}, behind which dust grains with a large variety of temperatures emit strong infrared emission \\citep{Hyland80}. \\citet{Fujiyoshi06} discovered the complex structure of the central 10 arcsec region using high spatial resolution radio observations.  \\citet{Colgan93} and \\citet{Simpson04} detected several MIR and FIR forbidden lines and examined the electron density and the abundance ratio at the center and several northern positions.  \\citet{Okada08} performed MIR spectroscopy with the {\\it Spitzer}/IRS as for G3.270-0.101, and detected 14 emission lines from metal ions and molecular hydrogen, indicating the coexistence of gas in various excitation states.  NGC~3603 is one of the most massive, optically visible \\hii\\ regions in our Galaxy.  It has an integrated H$\\alpha$ flux of $L$(H$\\alpha$) $\\sim 1.5\\times 10^{39}$\\ erg\\ s$^{-1}$ , equivalent to 20 O5 V stars \\citep{Kennicutt84}.  It is located in the Carina spiral arm at a distance of about 7~kpc \\citep[][and references within]{Nurnberger02}.  The central cluster, HD~97950, contains 3 Wolf-Rayet (WR) stars \\citep{Moffat94} and it is understood it experienced a single burst of star formation 1 Myr ago \\citep{Stolte04}.  It shows a remarkable similarity to the dense core R136 in 30 Dor in the LMC \\citep{Moffat94}.  \\citet{Nurnberger02} determined the properties and structure of the molecular gas using CS observations.  The structure of the ISM and infrared sources were examined with MIR observations \\citep{Nurnberger03,Nurnberger_Stanke03}.  MIR spectroscopic observations were performed using the {\\it Spitzer}/IRS and spatial variations in both polycyclic aromatic hydrocarbons (PAHs) and elemental abundances were examined \\citep{Lebouteiller07,Lebouteiller08}.  M17 is also one of the brightest giant \\hii\\ regions in our Galaxy and has been extensively studied at various wavelengths. It contains a dozen of O- or early B-type stars \\citep{Hanson97} and the age of the central cluster is $<3$~Myr \\citep{Hanson97,Jiang02}.  \\citet{Felli84} presented radio maps of M17 and discussed the structures and nature of the ISM.  The northern bar is clearly seen in optical emission, which is indicative of a small amount of foreground extinction, whereas the southern bar has no clear optical counterpart and is highly obscured.  Emission at \\oiii\\ 88\\micron\\ was detected by \\citet{Ward75}, and the electron density and nature of the ionizing sources were investigated based on observations of the FIR \\oiii\\ and \\niii\\ emission lines \\citep{Emery83}.  The \\cii\\ 158\\micron\\ emission is widely distributed in this region \\citep{Stutzki88,Matsuhara89}.  \\citet{Stutzki88} investigated the clumpiness of the edge-on \\hii\\ region/PDR interface using the spatial distribution of \\cii\\ 158\\micron\\ and CO rotational line emission.  The PDR structures and morphology were also investigated by \\citet{Ando02} and \\citet{Povich07}.  \\subsection{Data reduction}  \\begin{figure*} \\centering \\resizebox{0.45\\linewidth}{!}{\\includegraphics{f1a.eps}} \\resizebox{0.45\\linewidth}{!}{\\includegraphics{f1b.eps}} \\caption{(a) An example of the obtained full-resolution spectra in G333.6-0.2 region.  Three emission lines are labeled.  For the spectrum of 60--89\\pcm, obtained by LW detector, we apply the defringing process with fringe parameters of around \\nii\\ 122\\micron\\ line and do not apply the apodization.  For the spectrum of 88-135\\pcm, obtained by SW detector, we apply the defringing process with fringe parameters around \\oiii\\ 88\\micron\\ line and apply the `Hanning' apodization (see Appendix \\ref{app:linefit}). (b) The lower resolution version of (a) obtained by the Fourier transformation of the shorter range of the interferogram (see text).}\\label{ex_spectrum} \\end{figure*}  We performed the data reduction of the time series data of the FIS-FTS using the official pipeline.  Details of the processes and calibration are described in Murakami et al. (in prep.).  The time series data are provided as integrated ramp signals.  The interferogram is derived from the differentiation of the time series data after correcting for the non-linearity of the ramp curve.  We remove glitches in the interferogram and apply the discrete Fourier transform with the self-determined zero optical path difference to obtain the spectra.  After the flux calibration, we have a spectrum of each forward and backward scan of the driving mirror.  We then take an average of 7 scans of the individual scanning directions.  \\oiii\\ 88\\micron\\ is detected by the SW detector and \\nii\\ 122\\micron\\ and \\cii\\ 158\\micron\\ are detected by the LW detector.  The spectra obtained by both detectors are affected by fringe patterns \\citep[][Murakami et al. in prep.]{Kawada08}.  The defringing process and the fitting function are optimized for each line.  Details of the line-fitting process are described in Appendix \\ref{app:linefit}.  Figure~\\ref{ex_spectrum}a shows an example of a spectrum (an average of the forward and backward scans).  Since the parameters of the defringing process are optimized for fringes around emission lines, the residual of the fringe pattern becomes large at wavenumbers far from the emission lines.  After obtaining the line intensities for each pixel within the detectors, we combine the results from 2 or 3 pointing observations of the same region (table~\\ref{obs_summary}) and compile maps with a regular grid of half the size of the original pixel scale.  Details of constructing line maps and estimating the uncertainties are described in Appendix \\ref{app:const_map}.  The final line maps are shown in Figs.~\\ref{linemap_G3}, \\ref{linemap_G333}, \\ref{linemap_NGC3603}, and \\ref{linemap_M17} for individual regions, which are discussed in the next section.  The maps of the continuum emission at specific wavenumbers are created by a similar procedure to that of the line maps.  Instead of the defringed full-resolution spectra, we use low-resolution spectra, which are obtained by the Fourier transformation of a narrower range of the optical path ($\\pm 0.35$~cm for SW and $\\pm 0.4$~cm for LW) to avoid the effect of fringes (Fig.~\\ref{ex_spectrum}b).  The obtained continuum maps are also shown in Figs.~\\ref{linemap_G3}, \\ref{linemap_G333}, \\ref{linemap_NGC3603}, and \\ref{linemap_M17}.   We have presented the results of imaging spectroscopic observations of 4 giant Galactic star-forming regions with FIS-FTS onboard AKARI.  We obtained \\oiii\\ 88\\micron\\ and \\cii\\ 158\\micron\\ line intensity maps of all the regions.  In G3.270-0.101, we found that the \\oiii\\ 88\\micron\\ emission peaks at the excitation source, whereas the radio continuum emission shows a peak at a region 0.8\\arcmin (3~pc) away.  In G333.6-0.2, we found a local \\oiii\\ 88\\micron\\ emission peak 3\\arcmin (3~pc) away from the central cluster, which is likely to be an embedded O5 star.  In NGC~3603, the distribution of the MIR, FIR, and radio continuum emission as well as the \\oiii\\ 88\\micron\\ emission are similar to each other.  For all regions, \\cii\\ 158\\micron\\ emission is widely distributed, as suggested by previous observations of star-forming regions.  We have pointed out that the \\oiii\\ 88\\micron\\ emission shows a stronger correlation with the 18\\micron\\ continuum emission than 9\\micron\\ emission, which indicates that the 18\\micron\\ emission originates in hot dust grains close to early-type stars, in contrast to the 9\\micron\\ emission, which traces PAH emission, and PAHs are destroyed in the relatively hard UV environments in which O$^{2+}$ forms.  This work indicates that the \\oiii\\ 88\\micron\\ emission traces the excitation sources more closely than the radio continuum emission, especially when there is a large density and/or column density gradient.  The FIR spectroscopy provides a powerful tool for understanding the nature of star-forming regions including the position of the excitation sources."
    },
    "0911/0911.5196_arXiv.txt": {
        "abstract": "We demonstrate the capability of {\\it AKARI} for mapping diffuse far-infrared emission and achieved reliability of all-sky diffuse map. We have conducted an all-sky survey for more than 94 \\% of the whole sky during cold phase of {\\it AKARI} observation in 2006 Feb. -- 2007 Aug. The survey in far-infrared waveband covers 50 $\\mu$m -- 180 $\\mu$m with four bands centered at 65 $\\mu$m, 90 $\\mu$m, 140 $\\mu$m, and 160 $\\mu$m and spatial resolution of 30\\arcsec\\ -- 40\\arcsec (FWHM).% This survey has allowed us to make a revolutionary improvement compared to the IRAS survey that was conducted in 1983 in both spatial resolution and sensitivity after more than a quarter of a century. Additionally, it will provide us the first all-sky survey data with high-spatial resolution beyond 100 $\\mu$m. Considering its extreme importance of the {\\it AKARI} far-infrared diffuse emission map, we are now investigating carefully the quality of the data for possible release of the archival data. Critical subjects in making image of diffuse emission from detected signal are the transient response and long-term stability of the far-infrared detectors. Quantitative evaluation of these characteristics is the key to achieve sensitivity comparable to or better than that for point sources ($<20$ -- 95 [MJy/sr]). We describe current activities and progress that are focused on making high quality all-sky survey images of the diffuse far-infrared emission. ",
        "introduction": "satellite and the {\\it AKARI} all-sky survey} \\label{doiy3_sec:intro}  {\\it AKARI} is the first Japanese satellite dedicated for infrared (IR) astronomy \\citep{doiy3:murakami07}. The satellite was launched in February 2006. Its telescope with a primary mirror of a diameter of $\\phi = 68.5$ cm as well as the focal plane instruments were cooled down with liquid Helium cryogen and mechanical JT coolers as a support \\citep{doiy3:nakagawa07}. During its cold operational phase with liquid Helium (cold phase; March 2006 -- August 2007), infrared observations in 2 -- 180 $\\mu$m wavelength were carried out.  One of the key observations of the {\\it AKARI} mission was to conduct all-sky survey in mid- and far-IR (MIR and FIR) wavelengths centered at 9 $\\mu$m, 18 $\\mu$m, 65 $\\mu$m, 90 $\\mu$m, 140 $\\mu$m, and 160 $\\mu$m. In this contribution, we will discuss FIR part of the survey observation.  Specifications of the FIR wavebands are shown in Table~\\ref{doiy3_tab:spec}. \\begin{table*}[bht] \\caption{Specifications of {\\it AKARI} FIR detectors (\\citeauthor{doiy3:kawada07} \\citeyear{doiy3:kawada07}, \\citeauthor{doiy3:shirahata09} \\citeyear{doiy3:shirahata09}).} \\label{doiy3_tab:spec} \\begin{center} \\leavevmode \\footnotesize \\begin{tabular}[h]{lccccl} \\hline \\\\[-5pt] Band name & N60 & WIDE-S & WIDE-L & N160\\\\[+5pt] \\hline \\\\[-5pt] Center wavelength & 65 & 90 & 140 & 160 & [$\\mu$m]\\\\ Wavelength range & 50 -- 80 & 60 -- 110 & 110 -- 180 & 140 -- 180 & [$\\mu$m]\\\\ Array format & $20\\times 2$ & $20\\times 3$ & $15\\times 3$ & $15\\times 2$ & [pixels]\\\\ Pixel scale\\footnotemark[1] & \\multicolumn{2}{c}{$26\\arcsec \\negthinspace\\negthinspace .8 \\times 26\\arcsec \\negthinspace\\negthinspace .8$} & \\multicolumn{2}{c}{$44\\arcsec \\negthinspace\\negthinspace .2 \\times 44\\arcsec \\negthinspace\\negthinspace .2$} & \\\\ Pixel pitch\\footnotemark[1] & \\multicolumn{2}{c}{$29\\arcsec \\negthinspace\\negthinspace .5 \\times 29\\arcsec \\negthinspace\\negthinspace .5$} & \\multicolumn{2}{c}{$49\\arcsec \\negthinspace\\negthinspace .1 \\times 49\\arcsec \\negthinspace\\negthinspace .1$} & \\\\ FWHM & $32\\arcsec \\negthinspace\\negthinspace .05 \\pm 0\\arcsec \\negthinspace\\negthinspace .10$ & $30\\arcsec \\negthinspace\\negthinspace .17 \\pm 0\\arcsec \\negthinspace\\negthinspace .08$ & $40\\arcsec \\negthinspace\\negthinspace .85 \\pm 0\\arcsec \\negthinspace\\negthinspace .10$ & $38\\arcsec \\negthinspace\\negthinspace .23 \\pm 0\\arcsec \\negthinspace\\negthinspace .15$ & \\\\ Detector device & \\multicolumn{2}{c}{Monolithic Ge:Ga array} & \\multicolumn{2}{c}{Stressed Ge:Ga array} & \\\\ \\hline \\\\[-8pt] \\multicolumn{6}{l}{\\footnotesize{\\footnotemark[1]At the array center.}}\\\\ \\end{tabular} \\end{center} \\end{table*} Four wavebands continuously cover 50 -- 180 $\\mu$m as shown in Figure~\\ref{doiy3_fig:spectra} (also see \\citeauthor{doiy3:kawada07} \\citeyear{doiy3:kawada07}). \\begin{figure}[bt] \\begin{center} \\includegraphics[width=7 cm]{doiy3_spect} \\end{center} \\caption{Spectral responses of the {\\it AKARI} FIR four bands centered at 65 $\\mu$m (N60), 90 $\\mu$m (WIDE-S), 140 $\\mu$m (WIDE-L), and 160 $\\mu$m (N160). Note that {\\it AKARI} has continuous wavebands covering 50--180 $\\mu$m so that we can make precise evaluation of total FIR intensity from inband flux of {\\it AKARI} observations (see $\\S$\\ref{doiy3_sec:discussion}).} \\label{doiy3_fig:spectra} \\end{figure} During the survey observation, the telescope direction is being kept opposite to the earth centre. Since the {\\it AKARI} satellite was launched into a sun-synchronous polar orbit and is revolving around the earth in every 100 minutes, strips of the sky along a great circle with 8\\arcmin -- 12\\arcmin\\ width, which is the width of detector arrays, are scanned with a constant scan speed of 3\\arcmin\\negthinspace.6/sec.  Due to the yearly revolution of the earth, the scan direction is shifted 4\\arcmin~per satellite revolution in longitudinal direction or in the cross-scan direction. This shift angle corresponds to about half or one-third of the detector widths so that we can observe a region in the sky twice or more with a continuous survey observation.  The whole sky is covered with half-a-year continuous observation except for the regions that cannot be surveyed due to moon-shine, heavy bombardment of high-energy particles due to anomalous geomagnetic field above South America (South-Atlantic Anomaly: SAA), and so on.  Pointed observations of specific sources and/or regions also hinder the survey observation by blocking 30-minute observational time per each observation. In these observations, FIR 4 detectors map the sky with a scan speed of 7\\arcsec\\negthinspace.5/sec, 15\\arcsec/sec, or 30\\arcsec/sec. These speeds are much slower (slow-scan observations) than that of the survey observation, so that the data have better quality than the survey data with much more spatial samplings and smaller artifacts due to non-linear behaviour of the detectors (slow-response; see also $\\S$\\ref{doiy3_sec:analysis}). We make comparisons between survey data and slow-scan data taken at the same region of the sky to evaluate the reliability of the data in $\\S$\\ref{doiy3_sec:cal}.  Observational gaps in the first six-month operation were compensated in later periods within the limit of satellite attitude control; cross-scan offset $< \\pm 1\\degr$. After about 17 months of survey period (cold phase), more than 94 \\% of the sky have been observed twice or more so that we have conducted a FIR all-sky survey that virtually covers the whole sky as shown in Fugure~\\ref{doiy3_fig:cov}. \\begin{figure}[ht] \\begin{center} \\includegraphics[width=7 cm]{doiy3_cov} \\end{center} \\caption{Sky coverage of the {\\it AKARI} all-sky survey in ecliptic coordinates (top panel). Cumulative fraction of the sky coverage as a function of the scan number is shown in the bottom panel.} \\label{doiy3_fig:cov} \\end{figure} ",
        "conclusions": "\\label{doiy3_sec:discussion}  One of the key characteristics of the {\\it AKARI} FIR all-sky survey is its good spatial resolution of 30\\arcsec\\ -- 40\\arcsec\\ (Table~\\ref{doiy3_tab:spec}) for covering the whole sky. This advantage is illustrated in Figure \\ref{doiy3_fig:M51} as the spatial resolution of {\\it AKARI} is nearly comparable to that of Spitzer.% \\begin{figure}[ht] \\begin{center} \\includegraphics[width=7 cm]{doiy3_M51} \\end{center} \\caption{M51 images observed with {\\it AKARI}, Spitzer, and Herschel; top-left: {\\it AKARI} 90 $\\mu$m survey image, top-right: {\\it AKARI} 140 $\\mu$m survey image; bottom images are taken from ESA Herschel press release on 19 June 2009; bottom-left: Spitzer/MIPS 160 $\\mu$m image (NASA/JPL-Caltech/SINGS), bottom-right: Herschel/PACS 160 $\\mu$m image (ESA \\& the PACS consortium).} \\label{doiy3_fig:M51} \\end{figure}  Wide and continuous spectral coverage of 50 -- 180 $\\mu$m (Table~\\ref{doiy3_tab:spec}, Figure~\\ref{doiy3_fig:spectra}) is another important advantage of the {\\it AKARI} FIR all-sky survey. Observed spectra, which are averaged in individual intensity bins, are shown in Figure~\\ref{doiy3_fig:LMC}. \\begin{figure}[ht] \\begin{center} \\includegraphics[width=7 cm]{doiy3_LMC} \\end{center} \\caption{FIR spectra observed in LMC. Average spectra in intensity slices are shown. Gray lines are gray-body single temperature fittings with a spectral index $\\beta = 1$.} \\label{doiy3_fig:LMC} \\end{figure} Single-temperature gray-body spectra are also shown. Four spectral bands cover both sides of the spectral peak in shorter wavelengths (65 $\\mu$m and 90 $\\mu$m) and longer wavelengths (140 $\\mu$m and 160 $\\mu$m) enabling us to make precise evaluation of colour temperature, as shown in the figure.  It is important to notice that excess emission is prominent in the 65 $\\mu$m band for all the intensity bins, showing significant contribution of emission from very small grains (VSGs; \\citeauthor{doiy3:dbp90} \\citeyear{doiy3:dbp90}). Thus 90 $\\mu$m / 140 $\\mu$m intensity ratio is a good indicator of the colour temperature of big grains (BGs), which are in a thermal equilibrium with interstellar radiation field (\\citeauthor{doiy3:dbp90} \\citeyear{doiy3:dbp90}), although possible excess emission even in 90 $\\mu$m band is found at the lowest intensity bin as the estimated colour temperature shows a small rise comparing to higher intensity bins.  Since its continuous coverage of the {\\it AKARI} bands in 50 -- 180 $\\mu$m with 90 $\\mu$m and 140 $\\mu$m bands (WideS band and WideL band; Figure~\\ref{doiy3_fig:spectra}), total intensity in 50 -- 180 $\\mu$m ($I_{50-180}$) can be estimated simply by summing up 90 $\\mu$m intensity $I_\\nu ({\\rm WideS})$ and 140 $\\mu$m intensity $I_\\nu ({\\rm WideL})$ as follows: \\[ I_{50-180} = \\Delta \\nu I_\\nu ({\\rm WideS}) + \\Delta \\nu I_\\nu ({\\rm WideL}). \\]  \\citet{doiy3:doi09} make numerical evaluation of the fraction between $I_{50-180}$ and $I_{30-800}$. The fraction is estimated to be 0.4 -- 0.7 in variety of conditions of interstellar space, leading them to conclude that the {\\it AKARI} total FIR intensity, $I_{50-180}$, is a good indicator of the total infrared intensity (TIR).  Takeuchi et al. (\\citeyear{doiy3:takeuchi09}) confirmed this indication by finding good luminosity correlation between $L_{50-180}$ and $L_{\\rm TIR}$ for nearby galaxies, expressed by the following equation: \\[ {\\rm log} L_{\\rm TIR} = 0.964\\ {\\rm log} L^{\\rm 2band}_{\\rm AKARI} + 0.814, \\]\\[ r = 0.989. \\]  It is widely accepted that $I_{\\rm TIR}$ is a good indicator of star-formation activities (e.g. \\citeauthor{doiy3:kennicutt98} \\citeyear{doiy3:kennicutt98}) so that the all-sky data of the {\\it AKARI} FIR survey can be a powerful tool to investigate the star-formation activities in the various sources including near-by star-formation regions as well as distant galaxies.  As a result, the {\\it AKARI} FIR all-sky survey can be a new basic database that is improved from the IRAS survey with its higher spatial resolution and wider spectrum coverage. It can be utilized as a fundamental database for variety of astronomical studies as well as a pilot survey for the current and the future infrared astronomical missions including Herschel, Planck and SPICA.  The data will be released among the {\\it AKARI} project members after the completion of the calibration ($\\S$\\ref{doiy3_sec:cal}) and the public will have access to the data through collaborations with project members. We are planning to release the data to the public after some proprietary time as the {\\it AKARI} FIR point source catalogue \\citep{doiy3:yamamura09}."
    },
    "0911/0911.0585_arXiv.txt": {
        "abstract": "The measurements of the linear polarisation of visible light from quasars give strong evidence for large-scale coherent orientations of their polarisation vectors in some regions of the sky. We show that these observations can be explained by the mixing of the photons with very light pseudoscalar (axion-like) particles in extragalactic magnetic fields during their propagation. We present a new treatment in terms of wave packets and discuss the circular polarisation. ",
        "introduction": "In this work~\\cite{paper}, we are interested in the effect that axion-photon mixing can have on the polarisation of light coming from distant astronomical sources. In particular, the observations of redshift-dependent large-scale coherent orientations of AGN polarisation vectors can, at least qualitatively, even in very simple models, be reproduced as a result of such a mixing of incoming photons with extremely light axion-like particles in external magnetic fields. These observations, presented in the second edition of this conference, were based on good quality measurements of the linear polarisation for a sample of 355 measured quasars in visible light~\\cite{hutsemekers}.  This has been discussed in terms of axion-photon mixing by several authors, in the case of plane waves~\\cite{planewaves} and a prediction from this mixing is an observable circular polarisation comparable to the linear one. Here, we present the case in which light is described by wave packets and show that the circular polarisation can be suppressed with respect to the plane wave case. ",
        "conclusions": "We have briefly presented axion-photon mixing with the use of wave packets. The main consequence of this treatment is the net decrease of circular polarisation with respect to what is predicted using plane waves. Hence, the lack of circular polarisation in the light from AGN does not rule out the ALP-photon mixing."
    },
    "0911/0911.0244_arXiv.txt": {
        "abstract": "{The implication of primordial magnetic-field-induced  structure formation for the HI signal from the epoch of reionization is studied. Using semi-analytic models, we compute both the  density and ionization inhomogeneities in this scenario. We show that: (a) The global HI signal can only be seen in emission, unlike in the standard $\\Lambda$CDM models, (b) the density perturbations induced by primordial fields, leave distinctive signatures of the magnetic field Jeans' length on the HI two-point correlation function, (c) the length scale of ionization inhomogeneities is $\\la 1 \\, \\rm Mpc$. We find that the peak expected signal (two-point correlation function) is $\\simeq 10^{-4} \\, \\rm K^2$ in the range of scales $0.5\\hbox{--}3 \\, \\rm Mpc$ for  magnetic field strength in the range $5 \\times 10^{-10} \\hbox{--}3 \\times 10^{-9} \\, \\rm G$. We also discuss the detectability of the HI signal. The   angular resolution of the on-going and planned radio interferometers allows one to probe only the largest magnetic field strengths that we consider. They   have the sensitivity to detect the  magnetic field-induced features. We show that the future SKA has both the angular resolution and the sensitivity to detect the magnetic field-induced signal in the entire range of magnetic field values we consider, in an integration time  of one week.}   \\begin{document} ",
        "introduction": "Understanding the epoch of reionization is one of the outstanding challenges of modern cosmology. In recent years, a partial understanding  of this important  stage in the history of the universe has emerged. Gunn-Peterson tests have revealed the existence of neutral hydrogen (HI)  at redshifts $z \\ga 5.7$ (\\cite{fan,white,djorgo,fan02,becker}). The detection of temperature-polarization cross-correlation and the polarization-polarization correlation at large angular scales ($\\ell \\la 10$) in the WMAP data have given firm evidence that the universe reionized around $z \\simeq 10$ (e.g. \\cite{wmap}).  The most direct approach to studying this epoch is to attempt to detect the redshifted 21-cm line of the neutral hydrogen during the period when   the universe makes a transition from being fully neutral to fully ionized. There is substantial on-going effort in that direction and this also remains one of the primary goals of upcoming and future radio interferometers (e.g.  MWA, for a detail description see \\cite{bmh06}, LOFAR \\footnote{\\tt www.lofar.org}, and  SKA \\footnote{\\tt www.skatelescope.org/pages/page\\_sciencegen.htm}). Primordial tangled magnetic fields can substantially alter the scenario of structure formation of the universe (\\cite{ss05}, hereafter Paper I and references therein). It has been shown that magnetic field-induced structure formation can cause early reionization (Paper I, \\cite{ts06a}) and might leave detectable signatures in the HI signal from the epoch of reionization (\\cite{ts06b,sbk}).  In this paper, we study the implications of the existence of the primordial magnetic fields for possible studies of the reionization epoch. We compute both the global HI signal and its fluctuating component. In particular, we study in detail the nature of ionization inhomogeneities in the presence of magnetic fields, as their precise understanding is needed to correctly set the scale of the effect.   The detailed quantification  of this aspect constitutes the main difference between our present work and earlier studies. We also discuss in detail the detectability of the magnetic field-induced HI signal.   In \\S 2, we briefly review the salient features of the structure formation process in the presence of primordial tangled magnetic fields. In \\S 3, we give results for the reionization process. In \\S 4, we compute and discuss the global HI signal. In \\S 5 and 6, the fluctuating component of the HI signal is derived and discussed. The detectability of the resultant signal is also discussed. Section 7 is reserved for summary and conclusions. In this paper we use  the parameters of the   FRW  model as favoured by the recent WMAP results \\cite{wmap}: spatially flat FRW model with $\\Omega_m = 0.3$ and $\\Omega_\\Lambda = 0.7$ with  $\\Omega_b h^2 = 0.022$  and $h = 0.7$ \\cite{freedman}. ",
        "conclusions": "Tangled primordial magnetic field can have interesting consequences for cosmology. The presence of these fields leave detectable signatures in the CMBR temperature and polarization anisotropies (\\cite{subramanian3,subramanian4}, \\cite{durrer}, \\cite{seshadri},\\cite{mack}, \\cite{gs05},\\cite{lewis},\\cite{kr05},\\cite{gio},\\cite{gk08},\\cite{yam},\\cite{fin08,pao}). The statistics of this signature will also have distinctive non-Gaussian features \\cite{brown,seshadri2,caprini1}. From these considerations, one can obtain an upper limit on the magnetic fields strength $B_0 \\la 4 \\times 10^{-9} \\, \\rm G$ \\cite{gs05}, for models with  magnetic field power spectrum index $n \\simeq -3$, which are the only models that are consistent with current state of observations \\cite{caprini,ss05}.  Sethi \\& Subramanian (Paper I) considered many aspects of these fields in the post-recombination universe. In particular, they showed that magnetic field-induced structure formation might lead to early reionization and the dissipation of these fields can substantially alter the thermal and ionization history of the universe.  Magnetic field-induced reionization and its possible impact on the detectable HI signal has been considered by several authors \\cite{ts06b,sbk}. In this paper, we study  these aspects further within the framework of semi-analytic models which allow us to compute and add in the impact of reionization inhomogeneities on the HI signal.  Our main findings can be summarized as: \\begin{itemize} \\item[1.] Owing to magnetic field dissipation in the post-recombination, the HI signal is not expected to show any absorption features unlike the standard case (Figure~2) (see also \\cite{sbk}). \\item[2.] The matter power spectrum owing to magnetic-field induced structure formation is peaked around the magnetic Jeans' scale.  This scale is imprinted on the HI signal as typical oscillations in the HI angular correlation function (Figure~5) \\item[3.] Figure~5  also captures the impact of  ionization inhomogeneities. The scale of the ionization inhomogeneities is $\\la 1 \\, \\rm Mpc$, for a wide range of magnetic field strengths (Figure~4). \\item[4.] We show that the presently-operational and upcoming radio interferometers are sensitive to the magnetic field strength $B_0 \\simeq 3 \\times 10^{-9} \\, \\rm G$. However, the future SKA will be able to detect the HI signal for  the entire range of magnetic field strength: $5 \\times 10^{-10} \\la B_0 \\la 3 \\times 10^{-9} \\, \\rm G$. A  detection of a characteristic scale and nature of  oscillations of the HI correlation function  could help measure the magnetic field at these scales directly. \\end{itemize}  The upcoming SKA will also have the capability of direct detection of primordial tangled magnetic field by measuring the Faraday rotation of $\\simeq 10^7$ radio sources.  We note here that the main advantage of a potential detection of the HI signal is that it is sensitive to far smaller values of magnetic field strength as compared to the other probes."
    },
    "0911/0911.5461_arXiv.txt": {
        "abstract": "We numerically investigate orbital evolution of star clusters (SCs) under the  influence of dynamical friction by field stars of their host disk galaxies embedded in dark matter halos. We find that SCs with masses larger than $\\sim 2 \\times 10^5 {\\rm M}_{\\odot}$ can show significant orbital decay within less than 1 Gyr due to dynamical friction by disk field stars in galaxies with disk masses ($M_{\\rm d}$) less than $10^9 {\\rm M}_{\\odot}$. We also find that orbital decay of SCs due to dynamical friction is more remarkable in disk galaxies  with smaller $M_{\\rm d}$ and higher mass-ratios of disks to dark matter halos. The half-number radii ($R_{\\rm h, sc}$) and mean masses within $R_{\\rm h, sc}$ of the SC systems (SCSs)  in low-mass disk galaxies with $M_{\\rm d} \\le  10^9 {\\rm M}_{\\odot}$ are found to evolve significantly with time owing to dynamical friction of SCs. More massive SCs that can spiral-in to the central regions of disks can form multiple SC systems with smaller velocity dispersions so that they can merge with one another to form single stellar nuclei with their masses comparable to $\\sim 0.4$\\% of their host disk masses. Based on these results, we suggest that luminosity functions (LFs)  for more massive globular clusters (GCs) with masses larger than $2\\times 10^5 {\\rm M}_{\\odot}$   can steepen owing to transformation of the more massive  GCs into single stellar nuclei through GC merging in less luminous  galaxies. We also suggest that the half-number radii of GC systems can evolve owing to dynamical friction only for galaxies with their total masses smaller than $\\sim 10^{10} {\\rm M}_{\\odot}$. ",
        "introduction": "Dynamical friction of SCs in galaxies have been discussed in many different contexts, such as formation of stellar galactic nuclei via merging of old GCs (e.g., Tremaine et al. 1975), transformation from non-nucleated dwarfs into nucleated ones (e.g., Oh \\& Lin 2000), physical meanings for the presence of the GC system (GCS) in the Fornax dwarf galaxy (e.g., Oh et al. 2001), and dark matter distributions of dwarf galaxies (e.g., Hernandez \\& Gilmore 1998; Goerdt et al. 2006; Inoue 2009). Orbital decay of GCs due to dynamical friction can significantly change spatial distribution of GCs and thus increase possibilities of them to be destroyed by strong galactic tidal fields (e.g., Vesperini 2000). Dynamical friction of SCs (including GCs) in galaxies is thus considered to be important for better understanding the evolution of GC luminosity functions (GCLFs) and that of half-number radii of the GCSs  in galaxies (e.g., Vesperini 2000). Recently, van de Ven \\& Chang (2009) have suggested the importance of dynamical friction in the dynamical evolution of SCs around nuclear rings in galaxies.  \\begin{table*} \\centering \\begin{minipage}{185mm} \\caption{The ranges of model parameters.} \\begin{tabular}{cccccccc} Parameters & {$M_{\\rm gal}$ ($\\times 10^{10} M_{\\odot}$) \\footnote{The total mass of a galaxy. }} & {$M_{\\rm d}$ ($\\times 10^{9} M_{\\odot}$) \\footnote{The total mass of a disk. }} & {$f_{\\rm d}$ \\footnote{The mass fraction of stellar disk ($=M_{\\rm d}/M_{\\rm gal}$) in a galaxy.}}& {$f_{\\rm b}$ \\footnote{The mass fraction of a bulge ($=M_{\\rm b}/M_{\\rm d}$) in a galaxy .}}& {Galaxy type \\footnote{HSB and LSB represent low-surface brightness and high-surface brightness stellar disks, respectively. }}& {$M_{\\rm V, low}$ \\footnote{The lower luminosity cut-off  in the SC luminosity function. This parameter is fixed at  $-6$ mag in all models (see the main text for main reasons for this).}}& {$M_{\\rm V, high}$ \\footnote{The higher luminosity cut-off in the SC luminosity function.}} \\\\ Value ranges & 0.1--10.0 & 0.1--10.0 & 0.01--0.1  & 0.0--1.0  & LSB or HSB & $-6$ (mag) & $-12 \\sim -8$ (mag) \\\\ \\end{tabular} \\end{minipage} \\end{table*}   Dynamical friction processes of SCs against  {\\it disk field stars in disk galaxies} are important for the following three reasons. Firstly, recent cosmological hydrodynamical simulations on the possible formation sites of first GCs (Kravtsov \\& Gnedin 2005) have shown that the present-day GCs can be formed within disk galaxies at high redshifts ($z\\sim 3$). Secondly,  less luminous disk galaxies like the Large Magellanic Cloud (LMC) have disky GC systems in which GCs have disky spatial distributions and rotational kinematics (e.g., Freeman et al. 1983; Olsen et al. 2004). Thirdly, previous theoretical studies suggested that the present-day dwarf ellipticals (dEs) were previously less luminous disk galaxies with no/little bulges (Mastropietro et al. 2004). Furthermore, previous observational results on physical properties of dwarf galaxies (e.g., Stiavelli et al. 2001) suggest similarity in nuclear density profiles of stars  between dEs and spirals with exponential bulges: these appear to imply that understanding nucleus formation processes due to SC migration caused by dynamical friction in spirals can further lead to better understanding those in dwarfs. Although dynamical friction processes of SCs {\\it against dark matter halos}  have been well investigated by previous works (e.g., Goerdt et al. 2006), those by {\\it field stars in disk galaxies} have not been done by self-consistent numerical simulations  so far.    The purpose of this paper is thus to investigate dynamical friction of SCs against  disk field stars based on fully self-consistent numerical simulations on dynamical evolution of  disk galaxies with SCs. We particularly investigate how dynamical friction of SCs against disk field stars depend on physical properties of their host galaxies, such as disk masses,  bulge-to-disk-ratios,  and  mass-ratios of disk to dark matter halo. Based on the present numerical results, we discuss the observed luminosity functions of GCs dependent on luminosities of their host galaxies (e.g., Jord\\'an et al. 2006),  formation processes of stellar galactic nuclei, and evolution of physical properties of GCSs in galaxies.    The plan of the paper is as follows: In the next section, we describe our  numerical models  for orbital evolution of SCs in disk galaxies. In \\S 3, we present the numerical results mainly on (i)  orbital evolution of SCs and (ii) the physical properties of SCSs for variously different models. In \\S 4, we discuss wide implications of the present results such as formation of stellar galactic nuclei by SC merging and evolution of mass and luminosity functions of SCs due to orbital decay of SCs caused by dynamical friction. We summarize our  conclusions in \\S 5. ",
        "conclusions": "We have investigated orbital  evolution of SCs being influenced by dynamical friction of disk field stars based on galaxy-scale numerical simulations on dynamical evolution of disk galaxies. We have investigated models with variously different model parameters in order to understand dependences of the effectiveness of dynamical friction of SCs on physical properties of disk galaxies. We summarize our principle results as follows.  (1) Dynamical friction of SCs against  disk field stars is much more effective in orbital decay of SCs in comparison with that against  galactic halos in disk galaxies. For example, the half-number radius of the SCS  ($R_{\\rm h, sc}$) in a galaxy with $M_{\\rm d}=10^9 {\\rm M}_{\\odot}$ can decrease by $\\sim 30$\\% owing to orbital decay of SCs caused by dynamical friction  within well less than $10^9$ yr for a reasonable MF of SCs. However orbital decay of SCs due to dynamical friction can be more clearly seen only for more massive SCs (e.g., those with $m_{\\rm sc} \\ge 2 \\times 10^5 {\\rm M}_{\\odot}$).   (2) Dynamical friction processes of SCs against disk field stars depend on physical properties of their host galaxies.  For example, the dynamical friction is much more effective in disks with lower $M_{\\rm d}$ owing to smaller stellar velocity dispersions. As a result of this, $R_{\\rm h, sc}$ can more rapidly evolve with time in disks with lower  $M_{\\rm d}$. Evolution of SCSs in more massive disks (with $M_{\\rm d} \\sim 10^{10} {\\rm M}_{\\odot}$) can not be clearly seen in the present study, which implies dynamical friction is important only for evolution of  SCSs (e.g., $R_{\\rm h, sc}$) in less massive disk galaxies.   (3) Dynamical friction of  SCs by disk field stars is more effective in disks with higher degrees of self-gravitating of the disks (e.g., higher $f_{\\rm d}$), which implies  that higher stellar densities and non-axisymmetric structures such as bars and spiral arms can  promote dynamical friction. Orbital decay of SCs due to dynamical friction in disks does not depend so strongly on bulge masses of the disks for $0 \\le f_{\\rm b} \\le 1$ in the present study. Evolution of SCSs due to orbital decay of SCs caused by dynamical friction is more important  in HSBs than in LSBs.    (4) More massive SCs can spiral-in to the nuclear regions ($R<0.1R_{\\rm d}$) of disk galaxies owing to dynamical friction so that they can form multiple SC systems there in less massive disk galaxies with $M_{\\rm d} \\le  10^{9} {\\rm M}_{\\odot}$. These systems  show velocity dispersions comparable to or less than internal stellar velocity dispersions of SCs so that they are likely to merge quickly with one another to form single massive SCs there. Thus formation of stellar galactic nuclei via merging of more massive SCs is possible in less massive disk galaxies.   (5) The observed $M_{\\rm nuc}-M_{\\rm g}$ relation (i.e., $M_{\\rm nuc} \\sim  0.003 M_{\\rm g}$) for less luminous galaxies can be closely associated with formation of stellar galactic nuclei by SC merging.  The LF/MFs  for more massive  SCs in SCSs of disk galaxies can steepen owing to the formation of stellar galactic nuclei through merging of more massive SCs inwardly transferred to the nuclear region of galaxies. This steeping of SCLF/MFs due to dynamical friction can be more significant in less massive  galaxies."
    },
    "0911/0911.5711_arXiv.txt": {
        "abstract": "Although recent work in numerical relativity has made tremendous strides in quantifying the gravitational wave luminosity of black hole mergers, very little is known about the electromagnetic luminosity that might occur in immediate conjunction with these events.  We show that whenever the heat deposited in the gas near a pair of merging black holes is proportional to its total mass, and the surface density of the gas in the immediate vicinity is greater than the (quite small) amount necessary to make it optically thick, the characteristic scale of the luminosity emitted in direct association with the merger is the Eddington luminosity {\\it independent} of the gas mass.  The duration of the photon signal is proportional to the gas mass, and is generally rather longer than the merger event.  At somewhat larger distances, dissipation associated with realigning the gas orbits to the new spin orientation of the black hole can supplement dissipation of the energy gained from orbital adjustment to the mass lost in gravitational radiation; these two heat sources can combine to augment the electromagnetic radiation over longer timescales. ",
        "introduction": "The merger of two supermassive black holes has been a topic of lively astrophysical speculation for many years \\citep{BBR80}. Recent developments in galaxy formation theory have made the prospect more plausible and suggest an environment for such events: the centers of galaxies that underwent major mergers a few hundred million years in the past \\citep{HaehKauff02,Volont03}.  Mergers may be particularly likely when the galaxy contains a relatively rich supply of interstellar gas, which may help binary black holes overcome the ``last parsec problem\" and approach each other close enough for gravitational wave emission to compress the orbit to merger within a Hubble time \\citep{GR00,AN02,Kazantz04,Escala05,Cuadra09,Dotti09} (but for a contrary view see \\cite{Lodato09}).  The presence of sizable quantities of interstellar gas in the parsec-scale environment then raises the question of how much gas might find itself even closer to the pair at their moment of merger.  This is a subject of great uncertainty.  It has been argued, for example, that there should be little gas closer to the merging pair than $\\sim 100$--$1000r_g$ ($r_g \\equiv GM/c^2$, where $M$ is the total mass of the system) because eventually the timescale for shrinkage of the binary orbit by gravitational wave radiation becomes shorter than the timescale for mass inflow due to internally generated fluid stresses \\citep{MP05}.  On the other hand, the mass of such a circumbinary disk might be as large as $\\sim 100 M_{\\odot}$ or more \\citep{MP05,AN02,Rossi09,Corrales09}; if even $1 M_{\\odot}$ were close enough to the merging black holes to be given heat equal to $1\\%$ of its rest mass, the total energy---$\\sim 10^{50}$~erg---might well be large enough to produce observable radiation.  It is therefore a worthwhile exercise to estimate what sort of light might be generated if even a small fraction of the surrounding gas were able to make its way in close to the merging black holes.  Because the amount of mass near the merging black holes is so difficult to estimate at present, the plan of this paper is to explore prompt electromagnetic radiation in a way that is scaled to whatever gas mass is there.  Thus, we will first estimate the heat per unit mass that might be deposited in this gas, then, in order to find the luminosity, estimate the timescale on which the energy is radiated.  Next, the more model-dependent subject of the spectrum will be broached.  Lastly, having seen how the light emitted depends on gas mass, we will discuss the issues related to whether an ``interesting\" amount of mass may be present.  In order to avoid additional complications, we will ignore any luminosity due to accretion through the circumbinary disk. ",
        "conclusions": "To summarize, when a pair of supermassive black holes merge, provided only that the gas very close to the merging pair has at least a small electron scattering optical depth, we expect the prompt EM signal to likely have a luminosity comparable to the Eddington luminosity of the merged system, $\\sim 10^{45}\\epsilon^{1/2} M_7$~erg~s$^{-1}$.  The duration of this bright phase is proportional to the mass of gas at $\\sim 10r_g$ from the merged black hole, a quantity that is at present extremely difficult to estimate.  If there is only the very small amount of gas necessary to be optically thick, the duration would be only $\\sim 10M_7$~min; on the other hand, quantities several orders of magnitude larger are well within the range of plausibility.  Gas at somewhat greater distance ($10 < r/r_g < 10^3$) should continue to radiate at a level possibly as high as $\\sim 0.1 L_E$, but for a longer time.  Because gravitational wave observatories like the {\\it Laser Interferometric Space Antenna} (LISA) are expected to give approximate source localization days or even weeks in advance of merger \\citep{LH09}, one could hope for synergistic observing campaigns that might catch the entire EM signal. Alternatively, it may be feasible to use EM surveys with large solid-angle coverage to search for source flaring having the characteristics predicted here in order to identify candidate black hole merger systems before any gravitational wave detectors are ready.  The spectrum of this light is, at present, much more difficult than its luminosity to predict with any degree of confidence.  It may peak in the ultraviolet, but if it does, it is likely to be rather bluer than the familiar UV-peaking spectra of AGN (if the spectrum does peak in the UV, extinction in the host galaxy may obscure some number events).  While the inner region still shines, its luminosity will likely be greater than that from greater radii, making the spectrum (to the degree it is thermalized) closer to that of a single-temperature system.  If it is only incompletely thermalized, a still harder spectrum might be expected.  At later times, when the outer disk dominates, the temperature corresponding to a given radius is $\\propto \\eta^{1/8} (r/r_g)^{-1/2} M_7^{1/4}$. Compared to the temperature profile of an accretion disk in inflow equilibrium ($T \\propto (r/r_g)^{-3/4}$), this is a slower decline outward, suggesting a spectrum that might be softer than during the initial post-merger phase, but still harder than that of typical AGN.  In addition, over time its high frequency cut-off will move to lower frequencies.  Because our understanding of what determines the spectra of accretion disks around black holes is far less solid than our understanding of their gross energetics (consider, for example, the still-mysterious fact that ordinary AGN radiate a significant fraction of their bolometric luminosities in hard X-rays), these spectral predictions are far shakier than the predicted bolometric luminosity scale.  After the merger heat has been radiated, the disk should revert to more normal accretion behavior and, beginning at the smallest radii, should display standard AGN properties \\citep{MP05}. Unfortunately, the pace of fading as a function of radius depends on the run of surface density with radius, which at the moment is highly uncertain.  It is therefore far beyond the scope of even this speculative paper to guess how rapidly that may occur."
    },
    "0911/0911.3906_arXiv.txt": {
        "abstract": "The GammeV experiment has constrained the couplings of chameleon scalar fields to matter and photons.  Here we present a detailed calculation of the chameleon afterglow rate underlying these constraints.  The dependence of GammeV constraints on various assumptions in the calculation is studied.  We discuss GammeV--CHASE, a second-generation GammeV experiment, which will improve upon GammeV in several major ways.  Using our calculation of the chameleon afterglow rate, we forecast model-independent constraints achievable by GammeV--CHASE.  We then apply these constraints to a variety of chameleon models, including quartic chameleons and chameleon dark energy models.  The new experiment will be able to probe a large region of parameter space that is beyond the reach of current tests, such as fifth force searches, constraints on the dimming of distant astrophysical objects, and bounds on the variation of the fine structure constant. ",
        "introduction": "\\label{sec:intro}  Increasingly strong evidence has emerged over the past decade that the expansion of  the universe is accelerating, a phenomenon which can be explained by a scalar field ``dark energy'' with negative pressure~\\cite{Komatsu_etal_2009,Dunkley_etal_2009,Kowalski_etal_2008,Percival_etal_2009,Ratra_Peebles_1988,Peebles_Ratra_1988,Caldwell_Dave_Steinhardt_1998}.  Couplings between such a scalar and Standard Model particles could lead to fifth forces observable at laboratory or solar system scales, that must be hidden in order for the dark energy to satisfy constraints on the non-observation of such forces.  There are three known ways to hide dark energy-mediated fifth forces: weak couplings between dark energy and matter~\\cite{Wilczek_1982,Frieman_Hill_Stebbins_Waga_1995,Kim_Nilles_2009}; effectively weak couplings locally, as in the Dvali-Gabadadze-Porrati (DGP) brane world model and its generalizations~\\cite{Dvali_Gabadadze_Porrati_2000,deRham_etal_2008,deRham_Hofmann_Khoury_Tolley_2008}; and an effectively large mass locally, as in chameleon theories~\\cite{Khoury_Weltman_2004a,Khoury_Weltman_2004b,Brax_etal_2004}.  The chameleon mechanism makes the effective mass of a scalar field grow substantially as the density of the surrounding matter is increased.  Chameleon theories are a particularly well-studied class of dark energy models~\\cite{Carlson_Garretson_1994,Burrage_Davis_Shaw_2009,Brax_etal_2009,Brax_etal_2009b,Brax_etal_2007,Brax_etal_2007b,Chou_etal_2009,Ahlers_etal_2008,Gies_Mota_Shaw_2008,Mirizzi_Redondo_Sigl_2009}.  Since fifth forces at least as strong as gravity are already ruled out from solar system scales to submillimeter scales, the chameleon mechanism must raise the effective field mass above about $(0.1\\textrm{ mm})^{-1} \\sim 10^{-3}$~eV locally.   In addition to chameleon dark energies, the ``scalaron'' field, resulting from the transformation to Einstein frame of an $f(R)$ gravity theory, is able to satisfy local and stellar tests of gravity only through the chameleon mechanism~\\cite{Hu_Sawicki_2007,Babichev_Langlois_2009,Upadhye_Hu_2009,Babichev_Langlois_2009b}.  It was shown in \\cite{Mota_Shaw_2007,Mota_Shaw_2006} that chameleons which couple strongly to matter have correspondingly large chameleon effects, allowing them to evade laboratory fifth force constraints such as \\cite{Adelberger2007} on matter couplings $\\bmat \\sim 1$.  If these chameleons also couple strongly to photons, then they are ideally suited to probes of the afterglow phenomenon, the first of which was the GammeV experiment at Fermilab~\\cite{Chou_etal_2009}.  Consider a closed, evacuated cylindrical chamber, with glass windows at the ends and a magnetic field in the interior.  Photons streamed through the windows will occasionally oscillate into chameleon particles in the background magnetic field.  If the mass of one of these chameleons in the walls of the chamber is greater than its total energy inside the chamber, then it will reflect from the wall; such a chameleon will be trapped in the chamber.  After the photon source has been turned off, any remaining chameleons will oscillate back into photons in the magnetic field, producing an observable ``afterglow'' of photons.  Here, we study the oscillation between photons and chameleon particles in the cylindrical vacuum chambers used in afterglow experiments such as GammeV.  We compute the afterglow and decay rates per chameleon particle as a function of the dimensionless chameleon-photon coupling constant $\\bgam$.  Our calculation allows for absorptivity in the chamber walls, chameleon-photon phase shifts due to a nonzero chameleon mass $\\meff(\\textrm{chamber})$, and additional phase shifts $\\xiref$ due to reflections from chamber walls, as found in \\cite{Brax_etal_2007}.  For a particle in a superposition of chameleon and photon states, the glass windows at the ends of the chamber act as quantum measurement devices; photons are transmitted, while chameleons are reflected.  The chameleon decay and afterglow rates can be computed by keeping track of the decline in chameleon amplitude between successive quantum measurements by the windows.  Armed with these rates, we proceed to compute the expected afterglow flux as a function of time in GammeV, as well as the resulting constraints on the chameleon-photon coupling as a function of the chameleon mass, a result that was discussed in ref.~\\cite{Chou_etal_2009}.  We then forecast the constraints that will be attained by an upcoming experiment, GammeV--CHASE, expected to take data in the winter of 2009-2010.  GammeV--CHASE was designed to improve in several ways upon the first GammeV chameleon experiment, hereafter referred to as GammeV.  We show that GammeV--CHASE will improve constraints on large photon couplings $\\bgam$ by several orders of magnitude, bridging the gap between GammeV constraints and the collider constraints of ref.~\\cite{Brax_etal_2009}.  Furthermore, it will reach chameleon masses several times greater than those excluded by GammeV, allowing us to probe masses at the dark energy scale $\\rho_\\mathrm{de}^{1/4} = 2.4\\times 10^{-3}$~eV.  Finally, we use our afterglow computation to forecast GammeV--CHASE constraints on specific chameleon models.  We begin with power law potentials, devoting most of our attention to $\\phi^4$ potentials, which are well-understood and ubiquitous in particle physics.  GammeV--CHASE constraints on $\\phi^4$ chameleons will complement the fifth force experiments of E\\\"ot-Wash \\cite{Adelberger2007} by probing strong matter couplings.  Next, we study two different models of chameleon dark energy, with inverse power law potentials and exponential potentials.  GammeV--CHASE will probe large areas of the $\\bmat$,~$\\bgam$ parameter space of each of these dark energy models. These parameter regions are largely unexplored by the current data, including constraints on the dark energy equation of state, variations in the electromagnetic fine structure constant, and the dimming of distant objects through photon-scalar oscillation in the Galactic magnetic field.  Thus, the GammeV--CHASE laboratory search for afterglow will complement current cosmological probes of dark energy, vastly extending our constraints on dark energy couplings to matter and photons.  The paper is organized as follows.  After a brief discussion of chameleon physics in Sec.~\\ref{sec:chameleons}, we study chameleon-photon conversion in afterglow experiments in Sec.~\\ref{sec:afterglow}, and compute the decay and afterglow rates for GammeV and GammeV--CHASE in Sec.~\\ref{sec:gammev}.  We constrain chameleons using GammeV and forecast constraints for GammeV--CHASE in Sec.~\\ref{sec:constraints}, before concluding in Sec.~\\ref{sec:conclusion}. ",
        "conclusions": "\\label{sec:conclusion}  GammeV and GammeV--CHASE are chameleon afterglow experiments which probe the coupling between photons and chameleon particles that are trapped inside a vacuum chamber.  Using the physics of chameleon-photon oscillation, we have computed the rate $\\Gdec$ at which a chameleon converts to a photon inside the vacuum chamber of such an experiment, and the rate $\\Gaft$ of production of photons which escape the chamber to reach an external detector.  We have computed these rates as functions of the chameleon-photon coupling $\\bgam$, the effective chameleon mass $\\meff(\\textrm{chamber})$ inside the vacuum chamber, and the phase $\\xiref$ that is introduced between chameleons and photons as they reflect from the walls of the vacuum chamber.  The decay and afterglow rates for the GammeV geometry are plotted as functions of $\\meff$ in Fig.~\\ref{f:vary_mass} and $\\xiref$ in Fig.~\\ref{f:vary_xiref}.  With these rates, we have computed the number of chameleons in the chamber and the expected afterglow flux in our detector, as a function of $\\bgam$, $\\meff(\\textrm{chamber})$, and $\\xiref$.  In a previous result, GammeV \\cite{Chou_etal_2009} used the decay rate computed above, as well as the afterglow rate due to non-bouncing trajectories, to exclude a region of the chameleon parameter space.  Here, we have included bouncing trajectories in our afterglow calculation, resulting in constraints that are less conservative and that depend on the phase shift $\\xiref$.  These model-independent constraints are shown in Fig.~\\ref{f:gammev1}.  For $\\xiref=0$ and small masses, the new constraints extend from half of the minimum photon coupling excluded by \\cite{Chou_etal_2009} to twice the maximum coupling excluded by that reference.  For nonzero phases, additional suppression of the chameleon decay rate to photons allows us to push to even higher $\\bgam$, excluding photon couplings ten times as high as those ruled out by \\cite{Chou_etal_2009}.  However, since GammeV was only able to probe chameleon models whose masses scaled rapidly with density, even this new analysis does not allow us to constrain the most common types of chameleon potentials.  GammeV--CHASE is an improved version of GammeV that is expected to take data in the winter of 2009-2010.  In Fig.~\\ref{f:gammev-chase} we forecast the model-independent chameleon constraints that will be achieved by GammeV--CHASE.  These constraints span many orders of magnitude in $\\bgam$, bridging the gap between the constraints from GammeV and from particle accelerators.  We show that the improved control of PMT systematics will allow GammeV--CHASE to probe smaller couplings than GammeV.  Rapid PMT switching and multiple runs at lower magnetic fields will allow us to constrain $\\bgam$ as high as $10^{16}$.  Our forecasts show that GammeV--CHASE will be able to probe a large range of couplings even at masses as high as the dark energy mass scale, $2.4\\times 10^{-3}$~eV.  Furthermore, improvements to the pumping system allow GammeV--CHASE to probe many commonly used and well understood chameleon potentials, including quartic potentials, as well as inverse power law and exponential potentials that could explain the observed cosmic acceleration.  Quartic chameleons have been discussed extensively in the literature~\\cite{Gubser_Khoury_2004,Upadhye_Gubser_Khoury_2006,Mota_Shaw_2006,Mota_Shaw_2007}. Torsion pendulum experiments such as \\cite{Adelberger2007} have ruled out matter couplings up to $\\bmat = 1$, but are insensitive to chameleons with stronger matter couplings.  It is precisely these chameleons which will be trapped in the GammeV--CHASE chamber.  We have shown that GammeV--CHASE will complement laboratory searches for quartic chameleons by probing strongly coupled chameleons, as shown in Fig.~\\ref{f:phi4}.  For self-couplings a few orders of magnitude smaller than unity, constraints on quartic chameleons will span seven orders of magnitude in $\\bmat$ and five orders of magnitude in $\\bgam$, extending from the upper bound $\\bgam \\sim 10^{16}$ provided by particle accelerators down to $\\bgam \\sim 10^{11}$.  GammeV--CHASE is also complementary to searches for a varying fine structure constant.  If we extend a simple model of $\\alpha_\\mathrm{EM}$ variation such as \\cite{Bekenstein1982} to a chameleon theory, by adding a $\\phi^4$ potential and a Yukawa matter coupling with coupling constant proportional to mass, then constraints from GammeV--CHASE will be more powerful than those from laboratory or cosmological tests in a large portion of the parameter space.  Chameleon dark energy, with runaway potentials of the form $V(\\phi) = M_\\Lambda^4 f(\\phi/M_\\Lambda)$ and $M_\\Lambda = 2.4\\times 10^{-3}$~eV, have also been studied in the literature \\cite{Brax_etal_2004}.  In such potentials, $f(\\phi/M_\\Lambda) \\rightarrow 1$ as $\\phi$ runs off to large values, corresponding to low matter densities such as the current cosmological background density.  We have applied our forecast GammeV--CHASE constraints to chameleon dark energy models with inverse power law potentials $V(\\phi) = M_\\Lambda^4 [ 1 + (M_\\Lambda/\\phi)^n]$ as well as exponential potentials $V(\\phi) = M_\\Lambda^4 [ 1 + \\exp(-\\phi/M_\\Lambda)]$.  Figure~\\ref{f:de_pwr} shows the constraints that GammeV--CHASE will be able to place on inverse power law chameleon dark energy.  For the simplest model, $n=1$, we will be able to probe matter couplings ranging from $\\bmat \\lesssim 10^{5}$ to $\\bmat \\lesssim 10^{16}$ and from $\\bgam \\gtrsim 10^{11}$ all the way up to the accelerator constraints, $\\bgam \\sim 10^{16}$.  At the largest matter ouplings, $\\bmat > 10^{12}$, we will probe parameters inaccessible to astrophysical dimming constraints and bounds on variations in the fine structure constant.  GammeV--CHASE constraints on dark energy with an exponential potential, shown in Fig.~\\ref{f:de_exp}, will be similarly powerful, covering matter couplings in the range $10^6 < \\bmat < 10^{16}$.  Furthermore, this range of matter couplings is completely inaccessible to astrophysical dimming constraints, which apply to chameleons with low masses in the galaxy, and bounds from $\\alpha_\\mathrm{EM}$ variation, applicable to models with large variations in $\\phi$.  Thus we have shown that GammeV--CHASE will probe large ranges of previously unexplored parameter space for the simplest models of chameleon dark energy.   \\subsection*"
    },
    "0911/0911.3889_arXiv.txt": {
        "abstract": "In addition to producing loud gravitational waves (GW), the dynamics of a binary black hole system could induce emission of electromagnetic (EM) radiation by affecting the behavior of plasmas and electromagnetic fields in their vicinity. We here study how the electromagnetic fields are affected by a pair of orbiting black holes through the merger. In particular, we show how the binary's dynamics induce a  variability in possible electromagnetically induced emissions as well as an enhancement of electromagnetic fields during the late-merge and merger epochs. These time dependent features will likely leave their imprint in processes generating detectable emissions and can be exploited in the detection of electromagnetic counterparts of gravitational waves. ",
        "introduction": "The promise of detecting and analysing compact systems with both gravitational and electromagnetic waves stands out as one of the most exciting prospects in the coming decades. As already pointed out in a number of works (e.g.~\\cite{Sylvestre:2003vc,Stubbs:2007mk}), most astrophysical systems which produce strong gravitational waves likely emit copiously in the electromagnetic band. Indeed the strong, and possibly highly dynamical, gravitational fields around compact objects affect the dynamics of plasmas and matter which in turn induce different emission mechanisms.  One example of such  a system is a black hole  surrounded by an accretion disk. Strong emission from these systems is understood as the result of radiative processes within jets powered by the extraction of rotational and binding energy. While the latter is qualitatively understood in terms of Newtonian-rooted arguments based on the potential of the central object, the former relies on extracting energy from a rotating black hole in the most efficient energy convertion process we know of.\\\\ \\indent The pioneering works of Penrose~\\cite{Penrose:1969pc} and Blandford and Znajek~\\cite{Blandford:1977ds}, together with  a large body of subsequent work, has provided a basic understanding of possible mechanisms to explain highly energetic emissions from single black hole systems interacting with surrounding plasmas (e.g.~\\cite{2000PhR...325...83L}). The interaction of electromagnetic field lines with the strong gravitational field of a rotating black hole is the fundamental component of these mechanisms to explain the  acceleration of particles that traverse the black hole's ergosphere. This scenario of a pseudo-stationary, single black hole  interacting with an accretion disk is reasonably well understood, and it is employed to explain energetic phenomena such as gamma ray bursts, AGNs, quasars, blazars, etc. However, a highly dynamical stage may occur prior to such a  pseudo-stationary regime which could give rise to strong emissions. In the context of galaxy mergers, such a stage would naturally occur as individual black holes in each galaxy eventually collide in the galaxy resulting from the merger~\\cite{1980Natur.287..307B}.\\\\ \\indent Gravitational waves from such collision would be detectable by the Laser Interferometric Space Antenna (LISA) and, as pointed out in e.g.~\\cite{Milosavljevic:2004cg,Haiman:2009te}, the late orbitting and merger phases would take place within a circumbinary disk and possibly even interacting with some residual plasma inside the orbiting black holes \\cite{Chang:2009rx}.  Possible emissions from these systems are only understood at late times, when the the pseudo-stationary picture mentioned above is applicable. However the intermediate regime has only recently been approached by a few works~\\cite{Palenzuela:2009yr,Chang:2009rx,vanMeter:2009gu} and in all cases with important simplifications introduced to track the system. In this work, as a follow up of~\\cite{Palenzuela:2009yr}, we concentrate on understanding the dynamics of possible electromagnetic fields anchored in the circumbinary disk in the presence of the merging black holes. In particular we examine the field configuration, possible energy enhancement and time variability of these fields as the merger take place and point out possible process that could give rise to a signal around the merger time. To this end, we consider the Einstein-Maxwell system in a setup that describes a pair of black holes close to the merger epoch, study the electromagnetic field behavior and compare to the single black hole case. While we do not consider a plasma in our current study, our analysis helps to understand the possible behavior in its presence and lay the foundations to future work in this direction.\\\\ This work is organized as follows, in Section~\\secapproach, we briefly review our formulation and numerical implementation of the problem. Section III describes the physical set up to study single and binary black hole configurations. Section IV discusses the results obtained for both scenarios considered and highlight main features which could induce emission with particular patterns. We conclude with section V which offers some final considerations and discussions. ",
        "conclusions": "We have analyzed the behavior of electromagnetic fields influenced by the dynamics of a binary black hole system. Our study illustrates several interesting aspects of such systems that emit not only gravitational waves, but can also radiate electromagnetically. Gravitational waves would be emitted through the different dynamical stages of the system --inspiral, merger and ringdown-- through a rather smooth manner with the peak occurring at the merger stage. Electromagnetic waves on the other hand, will be emitted through diverse processes driven by the interaction of the EM fields with surrounding plasma, gas or matter. In the current work we have studied the behavior of the EM fields and illustrated their ``radiative'' behavior as energy that can propagate outwards from the system as the black holes influence them. This energy will likely be absorbed and re-emitted by the surrounding plasma and would, in turn, be possible to observe. The orbiting behavior however, would leave its mark in a time variability naturally induced in the emission process.  Indeed, as we have illustrated here, the EM fields have a clearly discernible pattern tied to the dynamics of the system, making them possible  {\\em tracers of the spacetime} --in the electromagnetic sector-- as these features would imprint particular characteristics in processes producing observable EM signals. In particular, in the pre-merger stages, the black hole dynamics induce EM flux oscillations with a period  half that of the dominant GW signal produced by the system, i.e., a fourth of the orbital period, and a gradual enhancement of the energy in the electromagnetic field. This enhancement, together with a flux of electromagnetic energy would impact surrounding plasma in a stronger way than would be the case for a single stationary black hole as the latter would neither exhibit an enhancement, nor would it give rise to an outward flux of EM energy. Emissions in this later case require mechanisms like accretion or Blandford-Znajek to take place. Certainly, the  same requirements will apply to binary black hole systems at late times as they give rise (generically) to a single spinning black hole. Furthermore, as the merger takes place fields are significantly twisted and stirred opening up the possibility of interesting emissions through magnetic reconnection. A study of such scenario requires a resistive treatment of the problem which is beyond the scope of our current work.\\\\   Perhaps even more exciting is the possibility of inducing a Blandford-Znajek analog for binaries as the merger proceeds. As we have seen, the system's dynamics induces a configuration consistent with the basic picture of this process and, by extracting rotational energy from the system, more powerful emissions could be expected. For this to take place however, an important conservative requirement should hold. Namely, an ergosphere must be present so that rotational energy can be extracted from the merging black holes. The best case scenario is for the ergosphere to form before the plunging phase begins so that the black holes can orbit and sufficient time is left for the extraction to occur.  An estimate for when this takes places can be drawn from an ``effective one body'' approach, as described in the appendix, which indicates highly spinning configurations are required, at least in the particle limit, for an orbiting behavior to exist within the ergosphere. \\\\  Finally, and at a rather academically interesting level, it is interesting to ask what conditions would be required to extract so much energy that the end state is essentially a non-spinning black hole after the merger. For this to take place, the timescale of the BZ process ($\\tau_{BZ}$) should be comparable to that of the merger ($\\tau_M$). The latter typically lasts $\\tau_M \\simeq 50 M$, the former can be estimated as in \\cite{2000PhR...325...83L} giving rise to \\begin{eqnarray} \\tau_{BZ} &\\simeq& 5 \\times 10^8 (10^{15} G/B)^2 (M_{\\odot}/M)\\,M_{\\odot} , \\nonumber \\\\ & \\simeq& 10^7 (10^{15} G/B)^2 (M_{\\odot}/M)^2 \\tau_M \\, ; \\nonumber \\end{eqnarray} thus, for $B > 10^{10} G$ both times result comparable for BH masses $\\geq 10^{8} M_{\\odot}$. Alternatively, one can estimate the amount of energy that could be extracted as a function of field strength within  $\\tau_M=50 M$. If $M_p$ is the final irreducible mass of the black hole after rotational energy has been extracted via the BZ mechanism, $M_p$ obeys \\begin{equation} \\frac{M_p}{M-M_p} \\simeq 10^8 (10^{15} G/B)^2 (M_{\\odot}/M)^2 \\, . \\end{equation} Thus, for $M=10^9 M_{\\odot}$ and $B=\\{10^7 / 10^8\\} G$ about $\\{10^{-4}/ 10^{-2}\\} M_{\\odot} c^2$ (i.e. $\\simeq \\{10^{48} / 10^{50}\\}$ ergs) is released in roughly a day.\\\\  While fields of these magnitudes might be unlikely, it is interesting that the strengths are not completely out of nature's ability to manifest. Beyond this possibility the impact of the dynamics on the electromagnetic fields is to induce a distinct variability as fields are dragged by the black holes. This suggests tantalizing prospects to detect pre-merger electromagnetic signals from systems detectable in the gravitational wave band. However, a complete description of the problem requires the incorporation of gas and radiation effects. Notwithstanding these missing --and important-- ingredients, the main qualitative features---driven by the orbiting behavior of the black holes, whose inertia is many orders of magnitude above all else---would intuitively remain unaltered.   At a more speculative level, these combined signals could be exploited to shed light on possible observations to analyse alternative theories of gravity where photons and gravitons might propagate at  different speeds or gravitational energy could propagate out of our possible 4-dimensional brane (for a recent discussion of some possibilities see~\\cite{Haiman:2008zy,Bloom:2009vx}).  As a final comment we stress that while our work is a step towards understanding possible emissions induced by binary black hole merger processes, we have only scratched the surface of possible phenomenology. For instance, within the current approach, scenarios with unequal masses and/or spins must be investigated and work is in progress to address them~\\cite{MPLRPY09}.  Still, a complete understanding of associated phenomena will require investigating, in particular, the interaction with a surrounding plasma and associated possible emission mechanisms. A preliminary related step in this direction has been considered in~\\cite{Chang:2009rx,2008ApJ...672...83M} where the possibility of emission by a fossil gas in between the black holes has been considered.  \\\\  \\noindent{\\bf{\\em Acknowledgments:}} We would like to thank M. Anderson, P. Chang, J. Frank, L. Rezzolla, S. Liebling, K. Menou, P. Moesta and D. Neilsen for stimulating  discussions as well as to P. Grandclement for assistance with Lorene. This work was supported by the NSF grants PHY-0803629 and PHY-0653375 and also NSERC through a Discovery Grant. Computations were done at TeraGrid. LL acknowledges the Aspen Center for Physics for hospitality where this work was started. Research at Perimeter Institute is supported through Industry Canada and by the Province of Ontario through the Ministry of Research \\& Innovation."
    },
    "0911/0911.3592_arXiv.txt": {
        "abstract": "Mirror matter is a stable self-collisional dark matter candidate. If exact mirror parity is a conserved symmetry of nature, there could exist a parallel hidden (mirror) sector of the Universe which has the same kind of particles and the same physical laws of our (visible) sector. The two sectors interact each other predominantly via gravity, therefore mirror matter is naturally ``dark''. Here I briefly review the cosmological signatures of mirror dark matter, as Big Bang nucleosynthesis, primordial structure formation and evolution, cosmic microwave background and large scale structure power spectra, together with its compatibility with the interpretation of the DAMA annual modulation signal in terms of photon--mirror-photon kinetic mixing. Summarizing the present status of research and comparing theoretical results with observations/experiments, it emerges that mirror matter is not just a viable, but a promising dark matter candidate. ",
        "introduction": "Mirror matter is a stable self-interacting dark matter candidate that emerges if one, instead of (or in addition to) assuming a symmetry between bosons and fermions (supersymmetry), assumes that nature is parity symmetric. It is a matter of fact that the weak nuclear force is not parity symmetric. The main theoretical motivation for the mirror matter hypothesis is that it constitutes the simplest way to restore parity symmetry in the physical laws of nature. When Lee and Yang proposed the non-parity of weak interactions in 1956, they mentioned also the possibility to restore parity by doubling the number of particles in the Standard Model \\cite{Lee:1956qn}. Thereby the Universe is divided into two sectors with the same particles and interactions, but opposite handedness, that interact mainly by gravity. In the minimal parity-symmetric extension of the Standard Model \\cite{Foot:1991bp,Pavsic:1974rq}, the group structure is $G\\otimes G'$, where $G=SU(3) \\otimes SU(2) \\otimes U(1)$ and the prime $(')$ denotes, as usual, the mirror sector. In this model the two sectors are described by the same lagrangians, but where ordinary particles have left-handed interactions, mirror particles have right-handed interactions. Besides gravity, mirror matter could interact with ordinary matter via the so-called kinetic mixing of gauge bosons, or via unknown fields that carry both ordinary and mirror charges. If such interactions exist they must be weak and are, therefore, negligible for many cosmological processes, but could be determinant for the detection of mirror dark matter in some contexts. Since photons do not interact with mirror baryons, or interact only via the weak kinetic mixing, mirror matter constitutes a natural candidate for the dark matter in the Universe.  Many consequences of mirror matter for particle physics and astrophysics have been studied during the last decades. The reader can refer to \\cite{Okun:2006eb} for a review of the history of mirror matter and a list of most relevant papers published before 2006.  Like their ordinary counterparts, mirror baryons can form atoms, molecules and astrophysical objects such as planets, stars and globular clusters. However, even though the microphysics is the same in both sectors, the initial conditions, and then the macrophysics, of the mirror sector should be different, because the cosmology must be different. In particular, Big Bang nucleosynthesis (BBN) requires that the mirror sector has a lower temperature than the ordinary one \\cite{Berezhiani:1995am,Berezhiani:2000gw}. This has implications for the thermodynamics of the early Universe \\cite{Berezhiani:2000gw,Ciarcelluti:2008vs} and for the key cosmological epochs. Analytical results and numerical calculations of BBN \\cite{Berezhiani:2000gw,Ciarcelluti:2008vm} show that the mirror sector should be helium dominated, and the abundance of heavy elements is expected to be higher than in the ordinary sector. Invisible stars made of mirror baryons are candidates for Massive Astrophysical Compact Halo Objects (MACHOs), which have been observed via microlensing events \\cite{Blinnikov:1996fm,Foot:1999hm,Mohapatra:1999ih}. Mirror stars contain more helium and less hydrogen, and therefore have different properties than ordinary stars \\cite{Berezhiani:2005vv}. The accretion of mirror matter onto celestial objects \\cite{Blinnikov:1983gh,Khlopov:1989fj,Blinnikov:1982eh} and the presence of mirror matter inside compact stars, as for example neutron stars, could have interesting observable effects \\cite{Sandin:2008db,Blinnikov:2009nn}. The consequences of mirror matter on primordial structure formation, cosmic microwave background (CMB) and large scale structure (LSS) distribution of matter have been studied \\cite{Ignatiev:2003js,Berezhiani:2003wj,Ciarcelluti:2003wm,Ciarcelluti:2004ij, Ciarcelluti:2004ik,Ciarcelluti:2004ip}. These studies provide stringent bounds on the mirror sector and prove that it is a viable candidate for dark matter. In addition, mirror matter provides one of the few potential explanations for the recent DAMA/LIBRA annual modulation signal \\cite{Bernabei:2008yi,Foot:2008nw,Ciarcelluti:2008qk,Foot:2009mw}, and possibly other experiments, as suggested for low energy electron recoil data from the CDMS collaboration \\cite{Foot:2009gk}.  If the mirror (M) sector exists, then the Universe along with the ordinary (O) particles should contain their mirror partners, but their densities are not the same in both sectors. In fact, the BBN bound on the effective number of extra light neutrinos implies that the M sector has a temperature lower than the O one, that can be naturally achieved in certain inflationary models \\cite{Berezhiani:1995am}. Then, two sectors have different initial conditions, they do not come into thermal equilibrium at later epoch and they evolve independently, separately conserving their entropies, and maintaining approximately constant the ratio among their temperatures.  All the differences with respect to the ordinary world can be described in terms of only two free parameters: \\begin{eqnarray}\\label{mir-param} x \\equiv \\left( s' \\over s \\right)^{1/3} \\approx {T' \\over T} ~~~~~~ {\\rm and} ~~~~~~ \\beta \\equiv {\\Omega'_{\\rm b} \\over \\Omega_{\\rm b}} ~~, \\end{eqnarray} where $T$ ($T'$), $\\Omega_{b}$ ($\\Omega'_{b}$), and $s$ ($s'$) are respectively the ordinary (mirror) photon temperature, cosmological baryon density, and entropy density. The lightest bounds on the mirror parameters are $ x < 0.7 $ and $ \\beta > 1 $, the first one coming from the BBN limit and the second one from the hypothesis that a relevant fraction of dark matter is made of mirror baryons.  As far as the mirror world is cooler than the ordinary one, $x < 1$, in the mirror world all key epochs (as are baryogenesis, nucleosynthesis, recombination, etc.) proceed in somewhat different conditions than in ordinary world. Namely, in the mirror world the relevant processes go out of equilibrium earlier than in the ordinary world, which has many far going implications.  In this paper I consider the cosmological effects of the gravitational interaction between O and M matter, and I add the effects of photon--mirror-photon kinetic mixing only when showing the compatibility of the mirror dark matter interpretation of DAMA/LIBRA annual modulation signal with the present cosmological scenario. ",
        "conclusions": "The current situation of the astrophysical research in presence of mirror dark matter is shown in Figure \\ref{resume}, with emphasis on the connections between the different theoretical studies and the predicted observable signatures. Solid lines mark what is already done (and is compatible with current observational and experimental constraints), while dashed ones mark what is still to do. For the last point a special importance is covered by star formation and N-body simulation studies, towards which people interested in mirror dark matter should address their efforts in the near future.  At present, we have already investigated the early Universe (thermodynamics and Big Bang nucleosynthesis), and the process of structure formation in linear regime, that permit to obtain predictions, respectively, on the primordial elements abundances, and on the observed cosmic microwave background and large scale structure power spectra. In addition, we have studied the evolution of mirror dark stars, which, together with the mirror star formation, are necessary ingredients for the study and the numerical simulations of non linear structure formation, and of the formation and evolution of galaxies. Furthermore, in future studies they can provide predictions on the observed abundances of MACHOs and on the gravitational waves background. Using the results of mirror BBN and stellar evolution, we have verified that the interpretation, in terms of photon--mirror-photon kinetic mixing, of the DAMA annual modulation signal is compatible with the inferred consequences on the physics of the early Universe. Ultimately, we will be able to obtain theoretical estimates of gravitational lensing, galactic dark matter distribution, and strange astrophysical events still unexplained (as for example dark galaxies, bullet galaxy, ...), that can be compared with observations.  Concluding, the astrophysical and experimental tests so far used show that mirror matter is not just a viable, but a promising candidate for dark matter. However, we still need to obtain a complete picture of the Mirror Universe.  \\begin{figure}\\label{resume} \\includegraphics[height=.56\\textheight]{resume.eps} \\caption{Current status of the astrophysical research with mirror dark matter: solid lines mark what is already done, while dashed ones mark what is still to do. For a more extensive explanation see the text.} \\end{figure}   \\begin{theacknowledgments} This work was supported by the Belgian Science Policy Office Inter University Attraction Pole VI/11 ``Fundamental Interactions''. \\end{theacknowledgments}"
    },
    "0911/0911.0728.txt": {
        "abstract": "% \\noindent {\\bf Abstract}  We use a 1D thermochemical and photochemical kinetics model to predict % the disequilibrium stratospheric chemistries of warm and hot Jupiters ($800\\!<\\!T\\!<\\!1200$ K). Thermal chemistry and vertical mixing are generally more important than photochemistry. At 1200 K, methane is oxidized to CO and CO$_2$ by OH radicals from thermal decomposition of water. At $T<1000$ K, methane is reactive but stable enough to reach the stratosphere, while water is stable enough that OH levels are suppressed by reaction with H$_2$. These trends raise the effective C/O ratio in the reacting gases above unity. Reduced products such as ethylene, acetylene, and hydrogen cyanide become abundant; further polymerization should lead to formation of PAHs (Poly-Aromatic Hydrocarbons) and soots. Parallel shifts are seen in the sulfur chemistry, with CS and CS$_2$ displacing S$_2$ and HS as the interesting disequilibrium products. Although lower temperature is a leading factor favoring hydrocarbons, higher rates of vertical mixing, lower metallicities, and lower incident UV radiation also favor organic synthesis. Acetylene (the first step toward PAHs) formation is especially favored by high eddy diffusion coefficients $K_{zz}>10^{10}$ cm$^2$/s. In most cases planetary compositions inferred from transit observations will differ markedly from those inferred from reflected or emitted light from the same planet. The peculiar properties of HD 189733b compared to other hot Jupiters --- a broadband blue haze, little sign of Na or K, and hints of low metallicity --- can be explained by an organic haze if the planet is cool enough. Whether this interpretation applies to HD 189733b itself, many organic-rich warm Jupiters are sure to be discovered in the near future. ",
        "introduction": "The most easily studied transiting hot Jupiters are HD 209458b and HD189733b (Fortney et al 2009). HD 209458b is typical of known hot Jupiters (Burrows et al.\\ 2008). Transit observations show a visible spectrum dominated by the 590 nm sodium resonance lines (Charbonneau et al 2002), the visibility of which implies a clean, haze-free atmosphere (Sing et al 2008). %Dayside observations show a pronounced six micron excess that probably comes from hot high altitude H$_2$O (Machalek et al 2008). % HD 189733b is atypical, although its details are controversial (Fortney et al 2009). In visible light HD 189733b presents a featureless blue haze with little sign of the expected Na and K lines (Pont et al 2008). Lecavalier des Etangs et al (2008) suggest that these lines are obscured by enstatite grains mixed up to microbar levels. Swain et al (2008, 2009) fit transit data between 1.5 and 2.5 $\\mu$m with CO, CO$_2$, H$_2$O, and CH$_4$, all with distinctly subsolar abundances corresponding to a metallicity of $[{\\rm M}/{\\rm H}]\\approx -0.5$. They suggested that adding C$_2$H$_2$, C$_2$H$_4$, or possibly NH$_3$, would improve their model's fit. On the other hand Fortney et al (2009) fit D{\\'e}sert et al's (2009) dayside spectrum at 4.5 $\\mu$m with supersolar CO and CO$_2$ corresponding to $0.5<[{\\rm M}/{\\rm H}]<1.5$.  An organic haze offers another way to explain the appearance of HD 189733b. Here we use a photochemical model % that we used to address S chemistry (Zahnle et al 2009) to address the threshold for organic haze formation in hot or warm Jupiters. In Zahnle et al (2009) we noted that temperatures below 1200 K change the kinds of chemistry that take place. In hot atmospheres methane is efficiently oxidized to CO. Cooler atmospheres are functionally more reduced because H$_2$O is more stable: organic molecules are favored over CO, PAHs (polycyclic aromatic hydrocarbons) and soots (disorganized agglomerations of bigger PAHs) might form and precipitate.  %, CO becomes less abundant yet the CO$_2$/CO ratio increases. The PAHs and soots would be the source of the haze. ",
        "conclusions": "We find that vertical mixing in hot Jupiters with temperatures below 1000 K can provide environments suitable for turning methane into ethylene, acetylene, and hydrogen cyanide. We speculate that such atmospheres would be conducive to generating PAHs and perhaps soots. We also reiterate the point that planetary compositions inferred from transit observations, which probe high altitudes (low pressures), can differ markedly from those inferred from reflected or emitted light from the same planet. In general, thermal chemistry and vertical mixing are more important to disequilibrium chemistry in warm or hot Jupiters than is photochemistry. Photochemistry plays a bigger role at low temperatures.  The presence of abundant PAHs or soots at high altitudes could account for the up-to-now unique properties of HD 189733b, if HD 189733b proves to be cool enough, or its metallicity low enough, or its C/O ratio high enough, or its vertical mixing as vigorous as Showman et al (2009) suggest. This would make HD 189733b the first of a new class of warm Jupiters. As noted, Swain et al (2009) reported that HD 189733b's metallicity might be rather low. %Another factor that might be germane is that HD 189733 itself (the star), is very active and therefore relatively young (Knutson 2010). %Hence the planet is also relatively young, and if still cooling it is likely to be more vigorously convective within. On the other hand HD 189733 is a relatively strong source of UV (Knutson 2010), which makes photolysis rates high, a factor that we have shown will suppress high altitude soot formation. Whether our thermometric interpretation of HD 189733b actually applies to HD 189733b (recent models suggest that the planet is too hot if it has solar C/O) is not the main point:  many organic-rich hazy warm Jupiters are sure to be discovered in the near future."
    },
    "0911/0911.5307_arXiv.txt": {
        "abstract": "{The recent observations of solar-like oscillations in many red giant stars with the CoRoT satellite stimulated the theoretical study of the effects of various physical processes on the modelling of these stars.} {The influence of rotation on the properties of red giants is studied in the context of the asteroseismic modelling of these stars.} {The effects of rotation on the global and asteroseismic properties of red giant stars with a mass larger than the mass limit for degenerate He burning are investigated by comparing rotating models computed with a comprehensive treatment of shellular rotation to non-rotating ones.} {While red giants exhibit low surface rotational velocities, we find that the rotational history of the star has a large impact on its properties during the red giant phase. In particular, for stars massive enough to ignite He burning in non-degenerate conditions, rotational mixing induces a significant increase of the stellar luminosity and shifts the location of the core helium burning phase to a higher luminosity in the HR diagram. This of course results in a change of the seismic properties of red giants at the same evolutionary state. As a consequence the inclusion of rotation significantly changes the fundamental parameters of a red giant star as determined by performing an asteroseismic calibration. In particular rotation decreases the derived stellar mass and increases the age. Depending on the rotation law assumed in the convective envelope and on the initial velocity of the star, non-negligible values of rotational splitting can be reached, which may complicate the observation and identification of non-radial oscillation modes for red giants exhibiting moderate surface rotational velocities. By comparing the effects of rotation and overshooting, we find that the main-sequence widening and the increase of the H-burning lifetime induced by rotation ($V_{\\rm ini}=150$\\,km\\,s$^{-1}$) are well reproduced by non-rotating models with an overshooting parameter of 0.1, while the increase of luminosity during the post-main sequence evolution is better reproduced by non-rotating models with overshooting parameters twice as large. This illustrates the fact that rotation not only increases the size of the convective core but also changes the chemical composition of the radiative zone.} {} ",
        "introduction": "The solar five-minute oscillations have provided a wealth of information on the internal structure of the Sun. These results stimulated various attempts to detect a similar signal on other stars. In past years, solar-like oscillations have been observed for a handful of stars either from the ground by using stabilized spectrographs developed for extra-solar planet searches or from photometric data obtained from space \\citep[see e.g.][]{bed07}. While these asteroseismic observations mainly focused on solar-type stars, solar-like oscillations have also been detected for a few red giants \\citep{fra02, bar04, bar07, der06}. These first detections together with the recent clear identification of non-radial oscillations for many red giant stars (G--K giants) with CoRoT \\citep{der09} stimulated the theoretical study of the asteroseismic properties of red giant stars and of the effects of various physical processes on the modelling of these stars. Rotation is one of the key processes that changes all outputs of stellar models with a peculiarly strong impact on the physics and evolution of massive stars \\citep[see e.g.][]{mae00b}. In this paper, we focus on the effects of rotation on the asteroseismic modelling of red giants massive enough to ignite He burning in non-degenerate conditions by comparing stellar models including shellular rotation (i.e. a strongly anisotropic turbulence leads to an essentially constant angular velocity on the isobars, see Sect.~\\ref{phys_rot} for more details) to non-rotating models. The influence of rotation on the asteroseimic features of red giants and on the determination of the global properties of a red giant star obtained by performing an asteroseimic modelling is studied.  The physical description of rotation is first briefly presented in Sect.~2. The results of the comparisons between rotating and non-rotating models of red giants are discussed in Sect.~3, while the conclusion is given in Sect.~4. ",
        "conclusions": "While red giants generally exhibit low surface rotational velocities, they may have been rotating much more rapidly during the main sequence. Since rotation significantly changes the stellar properties during the main sequence, the further evolution in the red giant phase is found to depend on the rotational history of the star. In particular, the comparison of rotating and non-rotating models computed with exactly the same initial parameters except for the inclusion of shellular rotation shows that rotation induces a significant increase of the stellar luminosity. The change of the evolutionary tracks in the HR diagram due to rotation is mainly due to rotational mixing with only a very limited contribution from hydrostatic corrections induced by rotation. In the case of red giant stars massive enough to ignite He burning in non-degenerate conditions, rotation shifts the location of the core helium burning phase to higher luminosity in the HR diagram. This of course results in a change of the global asteroseismic properties of red giants situated at the same evolutionary stage. Lower values of the large separation of radial modes as well as of the small separation between $\\ell=0$ and $\\ell=2$ modes trapped in the acoustic cavity of the star are then found for rotating models than for the non-rotating ones.  The effects of rotation and overshooting have also been compared. We find that both processes increase the stellar luminosity and the determined age and shift also the location of the end of the main sequence to lower effective temperatures. However, the changes of the stellar global properties induced by rotation cannot be reproduced by a single non-rotating model with a given value of the overshooting parameter. This is due to the fact that, contrary to overshooting, the changes of the global stellar properties induced by rotation are not only due to the increase of the size of the convective core, but also to the transport of helium and H-burning products by shear mixing and meridional circulation in the radiative layers.  Finally, the effects of rotation on the determination of the fundamental stellar parameters and asteroseismic properties for red giants sharing the same location in the HR diagram have been studied. Rotation is then found to decrease the determined value of the mass of a red giant located at a given luminosity in the HR diagram (by about 10\\% for a typical value of the rotation velocity) and to increase the value of the age (by about 40\\% for a typical rotation rate). We thus conclude that although red giants are slow rotators, the inclusion of rotation significantly changes the fundamental parameters determined for a star from an asteroseismic calibration. This is especially interesting in the context of the recent observations of solar-like oscillations in red giants obtained with CoRoT \\citep{der09} and of new asteroseismic data expected from ground based observations and upcoming space missions."
    },
    "0911/0911.2315.txt": {
        "abstract": "{The origin of large scale magnetic fields in astrophysical rotators, and the conversion of gravitational energy into radiation near stars and compact objects via accretion have been  subjects of active research for a half century. Magnetohydrodynamic turbulence  makes both problems highly nonlinear, so both subjects have  benefitted from numerical simulations.However, understanding the key principles and  practical modeling of observations warrants testable  semi-analytic mean field theories that distill the essential physics.  Mean field dynamo (MFD) theory  and alpha-viscosity accretion disc theory exemplify this pursuit.  That the latter is a mean field theory is not always made explicit but the combination of  turbulence and global symmetry imply such.    The more commonly explicit presentation of assumptions in  20th century textbook MFDT  has exposed it to arguably more widespread criticism than incurred by 20th century alpha-accretion theory despite complementary weaknesses.   In the  21st century however,  MFDT has experienced a breakthrough with a dynamical saturation theory  that consistently agrees with simulations. Such has not yet occurred in accretion disc theory, though progress is emerging. Ironically however, for accretion engines,  MFDT and accretion theory are presently two artificially uncoupled pieces of what should be a single coupled theory. Large scale fields and accretion flows are dynamically intertwined because large scale fields likely play a key role in angular momentum transport. I discuss and synthesize aspects of recent  progress in MFDT  and  accretion disc theory to  suggest why the two likely conspire in a unified theory.} ",
        "introduction": "Large scale dynamo (LSD) theory  in astrophysics is aimed at quantitatively understanding the in situ physics of magnetic field  growth, saturation, and  sustenance  on time or spatial scales large compared to the turbulent scales of the host. The presence of magnetohydrodynamic  turbulence makes the problem highly nonlinear, exacerbating  the importance of numerical simulations.  However, large scale spatial symmetry and slow evolution of  large scale fields  compared to turbulent fluctuations  time scales motivates  a mean field approach in which statistical, spatial, or temporal averages are taken  and evolution of the mean field studied (Moffatt 1978). An important goal of 21st century semi-analytic mean field dynamo theory  (MFDT) is to capture the essential  nonlinear physics in a minimalist way that is simple enough for phenomenological models, but also accounts for  the nonlinear saturation found in simulations  (for technical reviews see: Brandenburg \\& Subramanian 2005; Blackman 2007).  The past decade has revealed significant progress in this endeavor as the  growth and saturation of LSDs seen in simulations seems to be reasonably matched by mean field dynamo (MFD) theories that follow the time dependent dynamical evolution of magnetic helicity (e.g. Blackman \\& Field 2002). %(Blackman \\& Field 2002).   Stars show that in situ LSDs must operate in nature.  The sun exhibits sign reversals  of the large scale field over each half  solar cycle (e.g. Wang \\& Sheely 1993;  Schrijver  \\&  Zwaan 2000).   If the field were simply that frozen into to the ISM during initial formation, there would  be no reversals. LSDs are also likely  operating in galaxies, despite whatever initial fields may have been seeded cosmologically because  supernova driven turbulence requires sustenance of the large field against turbulent diffusion.  Turbulent diffusion in 3-D is not likely suppressed as it may be in in 2-D  simulations  for weak magnetic fields (Cattaneo 1994) and there is little evidence that turbulent diffusion in interstellar media is quenched.  Detailed modeling of the geometry of large scale fields can be matched by mean field theories (\\cite{beck96,vallee04,shukurov02}), though incorporation of the nonlinear LSD principles of the past decade for idealized dynamos into more geometrically realistic dynamos of astrophysical rotators with realistic boundary conditions setting is an ongoing avenue of research (e.g. Sur et al. 2007,2009).  Given the role of LSDs in the turbulent rotators of stars and galaxies, their presence  in turbulent accretion discs is also likely. The large scale fields required by prevailing models of astrophysical jets may be produced by an in situ dynamo, although understanding the relative importance of in situ production vs. flux accretion  in that context has been a long-standing topic of research (e.g. Lovelace et al. 2009). As in  stellar and galactic environments of LSDs, the likely  presence of  magnetohydrodynamic turbulence in accretion discs makes the problem  fundamentally nonlinear and a focus of numerical simulations. %It is part for this reason that the focus in non-linearly MHD accretion disc theory %has been on local transport, rather than on angular momentum transport %from large scale fields. Computational resources have demanded local shearing boxes for thin disc studies but, as discussed in section 4, even localized stratified shearing box simulations show evidence for LSDs and the influence of large scale fields (Brandenburg et al. 1995; Lesur \\& Ogilive 2008; Davis et al. 2009).   Non-thermal coronal luminosity from accretion engines also hints at the role of large scale field  because larger scale structures more easily survive the buoyant rise before dissipating in  coronae (Blackman \\& Pessah 2009).  While LSD theory has always been presented with  explicit mean field theories that expose the approximations made, axisymmetric accretion disc models such as the Shakura Sunyaev  (\\cite{ss73}, hereafter SS73) viscous $\\alpha_{ss}$ model are not always recognized as mean field theories by those using them to model spectra.  In this sense, 21st century LSD theory is more progressed than axisymmetric mean field accretion disc theory: The $\\alpha_{ss}$ disc prescription  does not yet offer predictive power for numerical simulations, whereas 21st century LSD theory  predicts the saturation of LSDs seen in a range of  simulations.  Ironically, because large scale fields likely play a fundamental role in accretion discs, both MFD and accretion  theory are really two faces of a single coupled theory.  Herein, I  first explore some conceptual progress and underlying principles of 21st century LSD/MFD theory in broad brush strokes. I then discuss parallels between large scale dynamo theory and accretion disc theory as separate endeavors. Finally,  I bring the two together and identify evolving connections and the importance of integrating the two theories into one. ",
        "conclusions": "For any turbulent MHD  system in which there is both turbulence on short time scales but ordered large scale patterns,  there has to  be a mean field theory that dynamically captures essential physics and key principles in a useful framework. Note that MHD itself is already a mean field theory which averages out the noise of individual particle motions.  The question for turbulent rotators involving  accretion and large scale dynamos is not \"is mean field theory is correct?\"  but rather \"do we have the correct mean field theory?\" Progress in large scale dynamo theory is farther along that accretion theory in this regard.  21st century mean field dynamo theory has incurred a paradigm shift compared to that of the 20th century  by incorporating the dynamical evolution of magnetic helicity. This has led to  mean field models that successfully predict the non-linear saturation of the large scale field growth seen in numerical simulations.  In that way, 21st century dynamo theory has overcome an long standing challenge of 20th century dynamo theory. More work is still needed in applying the principles learned to realistic systems.  The analogous paradigm shift has not yet occurred in accretion theory. The single $\\alpha_{ss}$ viscosity parameter of SS73 does not capture the quantitative properties of angular momentum transport in numerical simulations or provide testable dynamical predictions   for either  local turbulent stresses produced by the MRI, or large scale magnetic stresses.  Some principles behind a generalized semi-analytic accretion model must include are emerging, and more effective closures than the $\\alpha_{ss}$ prescription are showing some promise.  Ultimately, the need to incorporate both the large scale magnetic field growth and its stresses along with those provided by small scale field, implies that semi-analytic mean field dynamo theory and semi-analytic accretion disc theory are two components of what is a unified mean field theory of accretion. Developing this unified theory is a challenging and meaningful goal for the present century."
    },
    "0911/0911.4482_arXiv.txt": {
        "abstract": "We study cosmic-ray acceleration in a supernova remnant (SNR) and the escape from it. We model nonthermal particle and photon spectra for the hidden SNR in the open cluster Westerlund~2, and the old-age mixed-morphology SNR W~28. We assume that the SNR shock propagates in a low-density cavity, which is created and heated through the activities of the progenitor stars and/or previous supernova explosions. We indicate that the diffusion coefficient for cosmic-rays around the SNRs is less than $\\sim 1$\\% of that away from them. We compare our predictions with the gamma-ray spectra of molecular clouds illuminated by the cosmic-rays (Fermi and H.E.S.S.). We found that the spectral indices of the particles are $\\sim 2.3$. This may be because the particles were accelerated at the end of the Sedov phase, and because energy dependent escape and propagation of particles did not much affect the spectrum. ",
        "introduction": "Supernova remnants (SNRs) are canonically considered the main sources of cosmic-rays in the Galaxy. The detection of non-thermal X-ray emission from SNRs clearly indicates that electrons are actually accelerated around the SNR shocks \\citep{koy95}, and the observations can constrain the electron spectra. On the other hand, observational confirmation of accelerated protons is not as easy as that of electrons. One way to study the acceleration and spectrum of protons is to study gamma-ray emission through $pp$-interactions and the decay of neutral pions \\citep*[e.g.][]{nai94,dru94,stu97,gai98,bar99,ber00}. In particular, molecular clouds are efficient targets of cosmic-ray protons because of their high density. Thus, clouds illuminated by the protons accelerated in a nearby SNR could be bright gamma-ray sources \\citep*[e.g.][]{aha96,fat05,fat06,gab09}.  Theoretical studies have suggested that old SNRs could be appropriate objects to investigate gamma-ray emission through $pp$-interactions, because the radiation from the accelerated electrons (primary electrons) disappears as the SNR evolves, owing to their short cooling time \\citep{yam06,fan08}. In other words, we could ignore the gamma-rays from primary electrons via inverse-Compton (IC) scattering of ambient soft photon fields and/or non-thermal bremsstrahlung from the interaction of electrons with dense ambient matter.  In this letter, we consider the evolution of an SNR surrounded by molecular clouds. We calculate the spectrum of cosmic-rays accelerated in the SNR and the photon spectrum of the molecular clouds illuminated by the cosmic-rays. We assume that a supernova explodes in a low-density cavity, because the progenitor star expels ambient dense gas via strong UV-radiation and stellar winds \\citep{che99}. The cavity may also have been created through previous supernova explosions. What differentiates this study is that we consider whether high-energy cosmic-rays illuminating molecular clouds were accelerated even after the SNR became old or they were accelerated only when the SNR was young. We also discuss the influence of particle diffusion on the cosmic-ray spectrum. We construct specific models for the open cluster Westerlund~2 and the SNR W~28, and compere the results with latest observations.  Westerlund~2 is one of the young open clusters from which TeV gamma-ray emission has been detected with H.E.S.S. \\citep{aha07}. It is surrounded by molecular clouds \\citep{fur09}. The gamma-ray emission is extended ($\\sim 0.2^\\circ$) and covers the molecular clouds \\citep{fuk09}. Noticeable objects such as pulsars that can be the source of the gamma-ray emission have not been observed in this region. \\citet{fuj09a} proposed that the gamma-ray emission comes from an old SNR, although there is no clear signature of SNRs in the cluster.  W~28 is a mixed-morphology SNR interacting with a molecular cloud \\citep{woo81}. It is an old SNR and TeV gamma-rays have been detected from molecular clouds around the SNR \\citep{aha08b}. ",
        "conclusions": "In this letter, we constructed a model of gamma-ray emission from molecular clouds surrounding an SNR. We considered old SNRs because gamma-rays mainly come through $pp$-interactions and pion decay, and those from primary electrons can be ignored. In fact, the synchrotron cooling time of electrons with $E\\sim 1$~TeV is only $\\sim 4\\times 10^3$~yr for $B\\sim 60\\:\\rm \\mu G$ \\citep{gab09}, and thus they cool quickly after their acceleration finishes. We assumed that the SNR shock travels in a low-density cavity, and we indicated that the diffusion coefficient of cosmic-ray particles around the SNR needs to be much smaller than that away from it (less than $\\sim 1$\\%), The predicted broad-band spectra for Westerlund~2 and W~28 are consistent with observations (Figs.~\\ref{fig:wd2} and~\\ref{fig:w28}). The proton spectra can be fitted with a power law with indices $\\sim 2.3$, which is the value at cosmic-ray sources expected in cosmic-ray propagation models for the Galaxy \\citep{str07}. The indices suggest that particles were accelerated near the molecular clouds even after the SNR became old and large for these objects. The soft spectra may be because the particles were accelerated at the end of the Sedov phase. While the low Mach number of the shock makes the spectra soft in our model, neutral hydrogens in the cavity may also do that \\citep*{ohi09}. On the other hand, for another SNR IC~443, the broad-band photon spectrum indicates that the index is 2 \\citep{zha08}. For this object, protons accelerated at an earlier stage of the Sedov phase (when the Mach number of the shock is large) may be illuminating molecular clouds. Moreover, if molecular clouds are distributed with various sizes and distances from the SNR, the gamma-ray spectrum may be the superposition of the spectrum of each cloud and may reflect the energy-dependence of particle diffusion (equation~[\\ref{eq:DE}]; see \\citealt{aha96,tor08})."
    },
    "0911/0911.2601_arXiv.txt": {
        "abstract": "I consider some of the issues we face in trying to understand dark energy.  Huge fluctuations in the unknown dark energy equation of state can be hidden in distance data, so I argue that model-independent tests which signal if the cosmological constant is wrong are valuable. These can be constructed to remove degeneracies with the cosmological parameters. Gravitational effects can play an important role. Even small inhomogeneity clouds our ability to say something definite about dark energy. I discuss how the averaging problem confuses our potential understanding of dark energy by considering the backreaction from density perturbations to second-order in the concordance model: this effect leads to at least a 10\\% increase in the dynamical value of the deceleration parameter, and could be significantly higher. Large Hubble-scale inhomogeneity has not been investigated in detail, and could conceivably be the cause of apparent cosmic acceleration. I discuss void models which defy the Copernican principle in our Hubble patch, and describe how we can potentially rule out these models.  This article is a summary of two talks given at the Invisible Universe Conference, Paris, 2009. ",
        "introduction": "The standard LCDM model of cosmology has reached a maturity where we may think of it as a paradigm: more than just a model, it is rather a worldview underlying the  theories and methodology of our particular scientific subject. In many respects it is fantastic, with a handful of constant parameters describing the main features of the model, which lies neatly within the 1- or 2-$\\sigma$ errors bars of most data sets. It is tempting to think of it as a science which is nearing completion, with remaining work simply reducing errors on the parameters.  However, it requires (at least) 3 pieces of physics we don't yet understand: inflation of some sort for the initial conditions; dark matter, and dark energy.  Depending on who you speak to, the dark energy problem ranges from being not a problem at all, other than the old cosmological constant problem (landscape lovers), to being the greatest mystery/calamity in all of science~-- ever! (if you're speaking to someone writing a grant proposal). Whether it turns out to be a storm in a teacup, or a revolution in our understanding of the cosmos, it's going to be difficult getting a handle on it, and ruling out many of the possibilities may be very hard or even impossible for the foreseeable future. In this article I discuss three different aspects of the dark energy problem, and I describe some subtle observable properties of the FLRW models which can help us see if we're on the wrong track. I consider first dark energy in the exact background FLRW models, then in perturbed FLRW models, and finally discuss non-FLRW models with no dark energy in them at all. \\begin{itemize} \\item Attempting to reconstruct the dark energy dynamics poses many difficulties which are well known, not least of which are degeneracies with the other parameters of the model. The physical motivation for many models is dubious at best, yet it is clear that we have to keep an open mind about what the dynamics could be. However, there are consistency relations we may use to test specifically for deviations from $\\Lambda$, which may be used without specifying a model at all.  \\item How does structure in the universe get in the way of our interpretation of a `background' model at all? That is, how do we smooth the matter in the universe to give an FLRW model? This interferes which what we think our background model should be, and hence contaminates our understanding of the dark energy. This effect is surprisingly large~-- at least 10\\% in the deceleration parameter.  \\item It's conceivable that dark energy has nothing to do with new physics at all, and that we're using the wrong background solutions in the first place because the universe has a large Hubble-scale inhomogeneity. These models aren't fully developed, but deserve further investigation, even though they nominally break with the Copernican principle. Again, there are observational consistency relations which let us test for and potentially falsify these sorts of deviations from the standard paradigm.  \\end{itemize} ",
        "conclusions": "If the cosmological constant, $\\Lambda$, is the underlying cause of dark energy then the viewpoint that cosmology is nearly complete is probably true. If $\\Lambda$ is wrong, however, then all bets are off, as `dark energy' could then be all sorts of things, from $x$-essence to modified gravity to large-scale inhomogeneity.   Attempts to justify the value of $\\Lambda$ using landscape arguments along with the multiverse, necessarily combined with the anthropic principle, open two tenable doors: one is that the universe as a whole is colossal or even infinite, and the other is that our Hubble volume is both tiny and incredibly special. If this is indeed the case, then the multiverse breaks with the Copernican principle in a spectacular way: we exist in a highly exceptional corner of the universe, and are not typical observers except in our little patch where the fundamental constants are just so (which might be rather large in terms of Hubble volumes of course, but small in terms of all that there is).  If we're happy to break with a `global Copernican principle' to explain the value of $\\Lambda$, it is not philosophically unreasonable to break with a `local Copernican principle' instead, in order to~-- perhaps~-- preserve a global one. The void models offer the possibility of describing dark energy as radial inhomogeneity in a dynamically predictable model, rather than as an unknown dynamical degree of freedom in a postulated homogeneous model. One can even argue that inflation predicts such a scenario~\\cite{linde}! That they break with the Copernican principle on our Hubble scale is a cause for concern, and even a sophisticated inhomogeneous model which drops the spherical symmetry may well suffer a similar problem. On the scale of the multiverse anything goes, so it's easy to imagine a universe with inhomogeneous fluctuations on Hubble and super-Hubble scales~\\cite{uzan}. Of course almost all speculation beyond a few times our Hubble sphere is really `fiction science' at the moment, requiring wild assumptions and extrapolations in all sorts of ways. Nevertheless, such speculation is important, and tests, such as those presented here, which can decide the scale of homogeneity out to the Hubble scale without assuming \\emph{a priori} the local Copernican principle, may play an useful role in discussing such questions.     All the different possibilities for dark energy typically involve \\emph{functional} degrees of freedom. From an observational point of view, we are really in the position where we have to try to rule out the very simplest models, as we can't investigate function spaces observationally in a meaningful way. It is useful in my view to try to construct model-independent tests of different models where we can. Here I've presented a few which utilise background observations. These can probe different things: $\\mathcal{O}_m(z)$ is non-zero if flat LCDM is wrong; $\\mathcal{O}^{(2)}_{m,k}(z),\\mathcal{L}_\\text{gen}(z)$ are non-constant if $\\Lambda$ is wrong, independently of all parameters of the model; and $\\mathcal{O}_k(z)$ is non-constant and $\\mathcal{C}(z)$ is non-zero if the FLRW models themselves are incorrect, independently of the theory of gravity used, or fluid used to model the dark energy. These are useful precisely because they can be implemented without specifying a background model at all, if the observables $D(z)$ and $H(z)$ are constructed directly from the data.  A important issue lies in the backreaction of perturbations which is tied up with the averaging problem. How do we smooth structure to connect to the background model at all? We have seen that backreaction in the concordance model renormalises our background, which gives a change in the deceleration parameter of at least 10\\%, apparently far higher than would be guessed by hand-waving arguments. Thus, even if dark energy is the cosmological constant, it might be difficult to see it as such until this problem is further quantified and understood.     \\paragraph{Acknowledgements} I would like to thank Kishore Ananda, Bruce Bassett, Tim Clifton, Marina C\\^ortes, George Ellis, Sean February, Ren\\'ee Hlozek, Julien Larena, Teresa Lu, Mat Smith, Jean-Philippe Uzan and Caroline Zunckel  for collaboration and discussions which went into the work presented here. I would like to thank Sean February for the plots in Fig.~\\ref{fig:void}, and Roy Maartens for comments. This work is supported by the NRF (South Africa)."
    },
    "0911/0911.0929_arXiv.txt": {
        "abstract": "Most experiments that search for direct interactions of WIMP dark matter with a target can distinguish the dominant electron recoil background from the nuclear recoil signal, based on some discrimination parameter. An acceptance region is defined in the parameter space spanned by the recoil energy and this discrimination parameter. In the absence of a clear signal in this region, a limit is calculated on the dark matter scattering cross section. Here, an algorithm is presented that allows to define the acceptance region a priori such that the experiment has the best sensitivity. This is achieved through optimized acceptance regions for each WIMP model and WIMP mass that is to be probed. Using recent data from the CRESST-II experiment as an example, it is shown that resulting limits can be substantially stronger than those from a conventional acceptance region. In an experiment with a segmented target, the algorithm developed here can yield different acceptance regions for the individual subdetectors. Hence, it is shown how to combine the data consistently within the usual Maximum Gap or Optimum Interval framework. ",
        "introduction": "Although the existence of dark matter is now well established (see e.g.~\\cite{amsler2008b}), we have not yet succeeded in determining its nature. A highly motivated class of models predicts dark matter to be in the form of Weakly Interacting Massive Particles (WIMPs) (see e.g.~\\cite{jungman1996}). Direct searches for scatterings of such WIMP dark matter particles off nuclei (see e.g.~\\cite{gaitskell2004}) probe and constrain models in highly relevant regions of parameter space. Limits on the cross section of the scattering process are calculated based on the exposure of the experiment as well as events that are observed in an acceptance region.  The question arises how this acceptance region should be defined. Here, an algorithm is presented that allows to define an optimum acceptance region such that the experiment has the best sensitivity for a given WIMP signal. Hence, one can expect to obtain the most stringent limit on the dark matter scattering cross section. The algorithm is illustrated using recent data from CRESST-II~\\cite{angloher2009}, but can be employed in any experiment. It is also shown how to consistently combine data from individual detectors or experiments with differing spectral acceptance in the framework of Yellin's Maximum Gap or Optimum Interval methods~\\cite{yellin2002}. ",
        "conclusions": "Previously, direkt dark matter searches would constrain various WIMP models using one common acceptance region. Here it was shown that by optimizing the acceptance region for each WIMP model, one can improve the sensitivity of an experiment by orders of magnitude. This has been demonstrated with recent data from CRESST-II as an example, where a drastically improved limit resulted in particular for low mass WIMPs. At the same time, the algorithm introduced here removes the ambiguity from defining the acceptance region in a rather ad hoc way. It was shown how to make full use of this optimization within the Maximum Gap or Optimum Interval frameworks to achieve a combined limit for individual subdetectors of a segmented-target experiment, or different experiments altogether."
    },
    "0911/0911.5398_arXiv.txt": {
        "abstract": "The formation of a solar system such as ours is believed to have followed a multi-stage process around a protostar and its associated accretion disk.  Whipple first noted that planetesimal growth by particle agglomeration is strongly influenced by gas drag, and Cuzzi and colleagues have shown that when midplane particle mass densities approach or exceed those of the gas, solid-solid interactions dominate the drag effect.  The size dependence of the drag creates a ``bottleneck'' at  the meter scale with such bodies rapidly spiraling into the central star, whereas much smaller or larger particles do not.    Independent of whether the origin of the drag is angular momentum exchange with gas or solids in the disk, successful planetary accretion requires rapid planetesimal growth to km scales.  A commonly accepted picture is that for collisional velocities $V_c$ above a certain threshold value, ${V_{th}} \\sim$ 0.1-10 cm s$^{-1}$, particle agglomeration is not possible; elastic rebound overcomes attractive surface and intermolecular forces. However, if perfect sticking is {\\em assumed}  for all ranges of interparticle collision speeds the bottleneck can be overcome by rapid planetesimal growth. While previous work has dealt with the influences of collisional pressures and the possibility of particle fracture or penetration, the basic role of the phase behavior of matter--phase diagrams, amorphs and polymorphs--has been neglected. Here it is demonstrated for compact bodies that novel aspects of surface phase transitions provide a physical basis for efficient sticking through collisional melting/amphorphization/polymorphization and subsequent fusion/annealing to extend the collisional velocity range of primary accretion to $\\Delta V_c \\sim$ 1-100 m s$^{-1} \\gg {V_{th}}$, which encompasses both typical turbulent RMS speeds and the velocity differences between boulder sized and small grains $\\sim$ 1-50 m s$^{-1}$.  Therefore, as inspiraling meter sized bodies collide with smaller particles in this high velocity collisional fusion regime they grow sufficiently rapidly to $\\sim$ 0.1 - 1 km scale and settle into stable Keplerian orbits in $\\sim$ 10$^5$ years before photoevaporative wind clears the disk of source material. The basic theory applies to low and high melting temperature materials and thus to the inner and outer regions of a nebula. ",
        "introduction": "\\label{sec:intro}    \\subsection{Cosmogonical Context}  The origin of the solar system and the formation of solar type stars are wed through contemporary studies of primitive solar nebula from a wide range of perspectives including observational and theoretical astrophysics, nucleosynthesis, planetary dynamics, accretion physics and meteoritics \\citep{Whipple:64, Adachi:76, Weidenschilling:77, LissauerAR:93, Cuzzi:06, Quitte:2007, Scott:07, Dullemond:07, Blum:08, Armitage:09}.  The general scenario \\citep{Armitage:09} includes a sequence of stages: (a) the initial collapse of interstellar gas nucleating the central protostar ($\\sim$ 0.1 Myr) (b) slow mass accretion onto the star and primary planetesimal formation around the evolving accretion disk ($\\sim$ Myr) followed by (c) a phase ($\\sim$ Myr) during which the accretion rate drops significantly allowing the photoevaporative wind to ``sever'' the disk into two portions at a radius that depends on the ratio of the stellar accretion rate to the mass loss rate due to photoevaporation, and finally (d)  a rapid {\\em clearing phase} ($\\sim$ 0.1 Myr) during which the inner disk is accreted onto the central star and the lightest elements of outer disk are removed by direct exposure to photoevaporative UV flux. Radiometric dating reveals that Ca-Al-rich inclusions (CAIs) within carbonaceous chondrite meteorites are about 4.57 Gy old, thereby constituting the first planetary materials \\citep[see e.g.,][]{Scott:07}. Several My later, well after protoplanetesimals had already been processed, chondrules within chondrites were still forming \\citep{Scott:07}.  Whereas very small particles move ostensibly as does the slightly sub-Keplerian nebular gas,  planet sized objects couple very weakly to the  gas.  Therefore, the build up of protoplanets reaches an impasse at the meter scale due to the exchange of angular momentum with the gas; this relative motion is responsible for a drag that causes them to drift inward and they are lost to the central star on time scales much less than the CAI to chondrule interval \\citep{Blum:08, Weidenschilling:08, Armitage:09}. The competition between gravitational settling and turbulence can lead to a disk midplane enhanced by solid matter and thus the same basic impasse is operative due to solid-solid angular momentum exchange which dominates the net drag effect  \\cite[e.g.,][and refs. therein]{Cuzzi:1993}.  Therefore, an outstanding problem in cosmogony is to understand how, when objects can agglomerate to the critical $\\sim$ meter scale  by known {\\em low} collisional velocity accretion processes \\citep{Dominik:07, Blum:08, Guttler:10}, they are  not then rapidly lost into the central star?  Hence, it is sought to reveal the basic mechanisms through which matter can accrete sufficiently rapidly to slow their radial motion and thwart their demise.  \\subsection{Theories of Solid Aggregation in Disks}  \\subsubsection{Gravitational Collapse}  Two general approaches are advanced; one focuses on building planetesimals from grain-grain accretion and the other proposes gravitational collapse of the disk---treated as either a one or two-component (solid and gas) ``fluid''---on length scales much larger than grain scale processes.  The latter, Safronov-Goldreich-Ward \\citep{Goldreich:04} and related \\citep{Youdin:02} theories, overcome the bottleneck without the need to treat the microphysics of particle agglomeration, by studying the gravitational collapse of the one (solid or gas) or two-component (solid and gas) disk.  The idea is that when the solids dominate, as matter settles to the midplane the density exceeds a critical value and self gravity induces an instability with a characteristic length scale that  produces planetesimals either sufficiently large to settle into stable Keplerian orbits or of sufficient abundance to drive further gravitational accretion eventually leading to the same large planetesimal state.  When gases dominate, collapse is predicted to form giant gaseous protoplanetesimal cores that drive further accretion of gas \\citep{Boss:97}.  Gravitational collapse by the settling of solids to the midplane is criticized \\citep{Weidenschilling:95, Cuzzi:06, Dominik:07} on the basis that it ignores the effects of gas pressure that drive turbulent shear and maintain the disk density below the critical value.  However, the nature of disk turbulence is invoked both to assist and suppress various types of aggregation.  For example \\cite{Johansen:07} argue that the turbulence driven ``streaming instability'', associated with the relative motion between solids and gas, can localize solids in the midplane and facilitate agglomeration whereas \\cite{Wilkinson:08} argue that in ``Stokes trapping'' turbulence suppresses the aggregation of solids.  \\subsubsection{Sequential Aggregation of Particles}  Agglomeration through surface effects, studied in many experiments and simulations \\citep{Chokshi:93, Dominik:07, Blum:08, Guttler:10}, posits that above a certain threshold collisional velocity, ${V_{th}} \\sim$ 0.1-10 cm s$^{-1}$, particle agglomeration is not possible; elastic rebound overcomes attractive intermolecular forces.  Additionally, the experimental agglomerates are generally of high porosity and their fragility is often cited as a central difficulty with such an approach to planetary accretion \\citep{Youdin:02, Wilkinson:08}.   While, for example, van der Waals interactions can accrete highly porous bodies \\citep{Dominik:07, Blum:08}, both the long nebular time scales and temperature variations will facilitate densification, and thus increased strength, through the {\\em sintering} and {\\em annealing} of polycrystalline matter---well known processes in condensed matter physics that has power law scaling in time \\citep{Dash:06}.  Annealing and sintering are apparently not commonly discussed in regards to the issue of porosity in laboratory experiments of accretion.  Therefore, a basic understanding of surface and interfacial physics and the processes of cohesion and inelasticity is required to assess accretion scenarios and the range of validity of perfect sticking \\citep{Cuzzi:06, Cuzzi:07}.  \\eject  \\subsection{\\label{sec:Nebular}Nebular Thermodynamic Regimes}  Outside the framework of the Minimum Mass Solar Nebula (MMN), it is a challenging exercise to calculate the relevant phase fronts in an evolving accretion disk \\citep{Stevenson:88, Lecar:06, CieslaCuzzi:06, Garaud:07, Armitage:09}.  For example, depending on, {\\em inter alia},  turbulent mixing, the gas density, the nebular optical depth and the epoch of evolution, the snowline can vary many AU with a temperature $\\sim$ 170 $\\pm$ 20 $K$, moving inward as the disk evolves.   Thus, separate from the feasibility of computing it, knowledge of the entire thermodynamic history of a particle undergoing a two-body collision is highly model dependent.  However, it is important to put the situation discussed here in thermodynamic context for particle collisions.  For a given material the basic kinetics of growth at a fixed temperature in a nebula can strongly influence its bulk and surface properties and hence collisional physics.  Based solely on equilibrium considerations, ice alone exercises an enormous area of its phase diagram in a disk.  Regardless of the material, it is well known that a host of basic equilibrium and disequilibrium phenomena will influence the bulk and interfacial properties \\citep{Dash:06}.  Therefore,  depending on the detailed  history of a particle it may, over the long times associated with disk evolution, have annealed to obtain normal compact laboratory properties \\citep{Dash:06}, or it may have a highly porous structure \\citep{Blum:08, Guttler:10}.  These issues are discussed in section \\ref{sec:Liquidity}, but we note that, in the absence of fracture, the mechanism described here provides a wide range of high sticking probability collisions for a very broad range of particle densities and thermodynamic conditions.  Hence, this new mechanism acts in concert with other mechanisms of particle destruction and sticking \\citep{Guttler:10}.  \\subsection{\\label{sec:Overview}Overview of Results}  In this paper a new thermodynamic mechanism of mass agglomeration, termed {\\em collisional fusion} is described.  The mechanism involves the confluence of collisional energetics, intermolecular interactions and structural surface and bulk phase transitions. While the stresses in solid/solid collisions have been embodied in a range of planetesimal studies through examination of, among other things, the restitution coefficient \\citep{Chokshi:93, Bridges:96, Higa:98, Guttler:10}, their influence on the bulk and {\\em interfacial phase behavior} has heretofore received little attention.  Due to a combination of accessibility and the dual role of relevance to both the formation of the gas giants and the dynamics of Saturn's rings, there have been a number of experiments on the collisions of ice particles under various conditions  \\citep[e.g.,][]{Bridges:84, Bridges:96, Higa:98, Dominik:07, Blum:08}, but none at the very high pressures and low temperatures where this new mechanism is operative.   Depending on the temperature and pressure, many materials can take on multiple equilibrium crystal structures ({\\em polymorphs}) and moreover, under other conditions the crystalline order of one of the polymorphs can be lost and the system becomes {\\em amorphous} such as is the case with glass.  Water substance exhibits a rich polymorphism and amorphism in its phase behavior. Here, classical collision physics is modified by an extension of the idea of {\\em damage assisted interfacial melting} \\citep{Dash:06} to the very high pressures that drive the fusion of matter by momentary liquifaction/amorph-polymorphization and freezing/annealing. The theory is quantitatively borne out in solid ice (Ih and Ic) down to $T \\sim$ 150 $K$, low density amorphous ice (LDA) and high density amorphous ice (HDA) at lower temperatures.   Moreover, the process occurs in silicon, which melts at $\\sim$ 1690 $K$, and where the range of collisional fusion speeds is $\\Delta V_c \\sim$ 100-1000 m s$^{-1}$, as seen in the recent molecular dynamics simulations of  \\cite{Suri:08}.  Finally, bodies of inhomogenous composition can fuse in this manner.  For example, a body with an icy mantle of sufficient thickness can fuse with another or other multicomponent materials, such as olivine, with similar melting temperatures can act in the same manner so long as one has access to the relevant phase diagrams.  Hence the approach provides a framework for collisions from the inner to the outer nebula.  For clarity and brevity the focus here is on the latter, where it is found that perfect sticking is extended to a range of collisional speeds $\\Delta V_c \\sim$ 1-100 m s$^{-1}$ which encompass both typical turbulent RMS speeds and the velocity differences between boulder sized and small grains $\\sim$ 1-50 m s$^{-1}$ \\citep{Johansen:07}.  The development of the paper is as follows.  In section \\ref{sec:Summ} the basic physical processes at play and an outline of the logic structure of the mechanism of collisional fusion is summarized. The detailed treatment of the thermodynamic state theory of collisional fusion is described in  \\ref{sec:Thermo} after which, in \\ref{sec:Drift} the ideas are used to address the bottleneck problem by calculating the fate of drifters beginning from several places in nebulae experiencing a wide range of turbulent enhancement/suppression of midplane solids.  A guide to the variables and symbols used is provided in table \\ref{table:Variables}.  Conclusions are drawn in \\ref{sec:Conclude}. ",
        "conclusions": "\\label{sec:Conclude}  Collisional fusion is a new mechanism for high velocity particle agglomeration.  It provides a physical basis for the high sticking probability necessary, but heretofore assumed, to resolve the ``bottleneck'' problem in primary planetary accretion.   Thus, the radial drift of pre-planetesimals can indeed overcome the meter scale bottleneck based on a fundamental description of {\\em how},  {\\em why} and under what {\\em dynamic} and {\\em thermodynamic conditions} collisions with midplane particles can lead to perfect fusion. Operational regimes range from outer nebular regions, where collisions of ice, or ice covered, particles dominate and inner regions where higher melting temperature solids persist.  The frequency of central collisions in the latter will be reduced due to Stokes drag and would have to be incorporated in a detailed numerical model.  The process is qualitatively the same for cases in which the liquid phase is unavailable but a structurally favorable high pressure amorph/polymorph can be accessed during collisions; for water substance at low temperatures \\citep{Mishima:96, Straessle:07}, or silicon at high temperatures \\citep{Suri:08}.  In the case of liquifaction while some melt may be lost to the surroundings and is thus unavailable for fusion, under low vapor pressures it is known that splashing is suppressed \\citep{Nagel:05}, whereas during polyamorphism this is obviously not an issue. While the range  $\\Delta V_c \\sim$ 1-100 m s$^{-1} \\gg {V_{th}}$ calculated here captures astrophysical values \\citep{Johansen:07}, we note that (i) the fracture strength of ice increases as the small particle radius decreases \\citep{Higa:98}, so the relative size of the colliding particles is important, and (ii) even if perfect fusion does not occur the reduction in relative speed may be so dramatic that particles drop below the lower bound ${V_{th}} \\sim$ 0.1-10 cm s$^{-1}$ of \\cite{Chokshi:93}.  Hence, this new mechanism can act in concert with other mechanisms of particle destruction and sticking \\citep{Guttler:10}.  For example, in the inner nebula, where Stokes drag is prevalent, the mechanism proposed here can coexist with that of \\cite{Wurm:01} in which coupling to the gas phase is sufficiently strong that it can lead to reconnection of collisional fragments to a large object.  Therefore, the combined action of collisional fusion and shattering may underlie the perfect sticking required for rapid planetesimal accretion from the inner to the outer nebula."
    },
    "0911/0911.5351_arXiv.txt": {
        "abstract": "Radiative transfer and radiation hydrodynamics use the relativistic Boltzmann equation to describe the kinetics of photons.  It is difficult to solve the six-dimensional time-dependent transfer equation unless the problem is highly symmetric or in equilibrium. When the radiation field is smooth, it is natural to take angular moments of the transfer equation to reduce the degrees of freedom. However, low order moment equations contain terms that depend on higher order moments.  To close the system of moment equations, approximations are made to truncate this hierarchy.  Popular closures used in astrophysics include flux limited diffusion and the $M_1$ closure, which are rather \\textit{ad hoc} and do not necessarily capture the correct physics.  In this paper, we propose a new class of closures for radiative transfer and radiation hydrodynamics.  We start from a different perspective and highlight the consistency of a fully relativistic formalism.  We present a generic framework to approximate radiative transfer based on relativistic \\citeauthor{Grad1949}'s moment method.  We then derive a 14-field method that minimizes unphysical photon self-interaction. ",
        "introduction": "Radiative transfer and radiation hydrodynamics use the ultra-relativistic Boltzmann transport equation to describe the kinetics of photon.  The radiative intensity, which is proportional to the photon distribution function, is a seven-dimensional hypersurface embedded in eight-dimensional phase space \\citep{1973grav.book.....M}. Unless the radiative intensity is in equilibrium, or the problem is highly symmetric, it is difficult to solve the radiative transfer equation either analytically or numerically \\citep[see standard textbook such as][]{1960ratr.book.....C, 1984oup..book.....M, 1986rpa..book.....R, 1991par..book.....S, 2001irtm.book.....P, 2004rahy.book.....C, 2006cmt..conf.....G}.  On the one hand, there is no successful theory to reduce the complexity of radiative transfer.  On the other hand, although numerical algorithms such as direct discretization of the radiative transfer equation and Monte Carlo methods exist, they are computationally too expensive.  In astrophysics, very often we are interested in the radiative energy and flux instead of the intensity.  If the radiation field is smooth (in terms of directions), we simply take angular moments of the radiative transfer equation and solve only for the frequency-dependent moment equations.  The $\\mbox{zeroth-}$, $\\mbox{first-}$, and second order angular moments of the intensity carry clear physical meanings. They are the radiative energy density, radiative flux, and radiative stress tensor, respectively.  For many classic problems in astrophysics like stellar atmospheres, the global symmetry of the system is used to further reduce the degrees of freedom.  Analytical or numerical solutions are then obtained by solving the reduced frequency-dependent moment equations.  This approach is naturally extended to radiation hydrodynamics, which not only describes how the (moving) media radiates, but also how radiation feeds back to matter \\citep{1984oup..book.....M, 2001JQSRT..71...61M, 2004rahy.book.....C, 2007ApJ...667..626K}.  The dynamic equations of the radiative stress tensor contain terms that are related to the third order moment of the intensity, while the equations of the third order moments depend on the fourth order moments and so on.  In order to close the system, we need to make approximations and truncate the moment hierarchy.  This is known as the \\emph{closure problem} in radiative transfer.  Popular closure schemes in astrophysics include variant forms of flux limited diffusion \\citep[FLD,][]{1970JCoPh...6....1C, 1981ApJ...248..321L, 1983JQSRT..29..223P, 1984JQSRT..31..149L}, $P_N$ approximations \\citep[sometimes with diffusion corrections, see][and reference therein]{2000JQSRT..64..619O, 2009CMT....21..511S, 2008JCoPh.227.2864M, 2010JQSRT.111.2052R}, the $M_1$ closure \\citep[equivalent to maximal entropy closure, see][]{1992A&A...265..345J, 2000A&A...356..559S, 2008PhRvD..78b4023F}, and variable Eddington factors \\citep[VEF,][]{1964SAOSR.167..108F, 1969JQSRT...9..407P, 1970MNRAS.149...65A, 2008PASJ...60..377F, 2008PASJ...60.1209F, 2009PASJ...61..367F}\\footnote{The term VEF has recently been abused in the literature.  The original \\citet{1970MNRAS.149...65A} algorithm uses an iterative solver for the Eddington factors, which is technically not a closure relation.  The closures such as $M_1$ should really be called moment methods with \\emph{prescribed Eddington factors} \\citep{1987ApJ...323..227F, 1990ApJ...354...83P, 1992A&A...266..613K}.}.  They have been successfully applied to many astrophysics problems, yet their underlying assumptions (what is the optimal form of a flux limiter for a particular problem?  why is the entropy of photons maximized?  how many moments do we need?) are rather \\textit{ad hoc} and do not guarantee to capture the correct physics.  In hydrodynamics, the Chapman-Enskog theory expands the particle distribution in power series of the Knudsen number, which is the ratio between the particle mean free path and the typical length scale in the problem.  The zeroth order expansion gives rise to ideal hydrodynamics, while the first order terms allow us to calculate transport coefficients such as kinematic viscosity and thermal conductivity \\citep{1965ApJ...141..201H, 1970mtnu.book.....C, 1979ittk.book.....L, 2002rbet.book.....C}.  We would like to apply a similar technique to radiative transfer in order to solve the closure problem.  Unfortunately, photons do not self-interact.  The photon mean free path is determined by the extinction coefficient, which is a material property.  In the cases when the photon density is high but the media is optically thin, radiation dominates the energy budget but cannot be thermalized.  The Chapman-Enskog series for photon distribution function fails to converge \\citep{2001JQSRT..71...61M}.  In addition to the closure problem, as studies of astrophysical objects become more detailed, complicated sub-structures such as convection in stars and turbulence in accretion disks are always found.  The radiating medium moves non-uniformly, which makes the computation of spectra more complicated due to the spatial-dependent Doppler effect.  The standard approach is to Taylor expand the extinction and emission coefficients and to derive moment equations to at least $\\mathcal{O}(v/c)$, where $v$ and $c$ are the speed of the media and the speed of light.  The resulting equations have terms that are physically important even in the limit $v/c \\rightarrow 0$ \\citep{1984oup..book.....M}.  Keeping track of these higher order terms in different physical regimes is a challenging task \\citep{2007ApJ...667..626K}.  Moreover, there are astrophysical systems such as the inner regions of accretion disks and ultra-relativistic jets that move with relativistic speed so the described approach converges too slowly.  If radiative feedback is weak, we can solve the hydrodynamic equations and compute the radiative spectra using post-processing \\citep[see][]{2007CQGra..24..259N, 2009ApJ...701..521C}.  However, there is a large set of problems for which the feedback is important. We have to solve the frequency-dependent three-dimensional moment equations at every single time step.  The full system of radiation hydrodynamics is highly non-linear.  Studies of turbulence show that reducing the number of spatial dimensions and enforcing symmetry in these systems can produce fundamentally wrong results \\citep{1967PhFl...10.1417K, 1971JFM....47..525K}.  To summarize, we have listed three major difficulties\\footnote{Another fundamental difficulty of radiation hydrodynamics comes from the large separation between the radiative time scale and the (hydro-) dynamic time scale.  Nevertheless, this difficulty raises not because of the moment method.  It is a generic property of non-relativistic astrophysical systems.  We will therefore leave it out from this paper.} in using moment methods for radiative transfer and radiation hydrodynamics: \\begin{enumerate} \\item Because of the collisionless nature of photons, standard methods in kinetic theory such as the Chapman-Enskog theory fails to provide a physical closure. \\item When the non-uniform motion of media is considered, it is difficult to take properly the spatial-dependent Doppler effect into account. \\item If the radiative feedback is important, the naive approach is computationally too expensive to couple radiation back to hydrodynamics. \\end{enumerate}  Hydrodynamics depends only on the frequency-integrated radiative force.  Many existing codes, therefore, reduce the degrees of freedom and resolve problem~(iii) by using frequency-integrated equations \\citep{1992ApJS...80..819S, 2001ApJS..135...95T, 2003ApJS..147..197H, 2007A&A...464..429G, 2007ApJ...667..626K, 2008MNRAS.387..295A, 2008PhRvD..78b4023F}.  If the radiative spectrum is needed, it is always possible to post-process the numerical solution.  Once we integrate over frequency, the moment equations are fully relativistic.  The radiative energy density, momentum, and stress tensors together form a covariant stress-energy tensor.  We can then ignore the complication of mixed frame formalism.  This relativistic approach is originally developed to solve radiative transfer in general relativity \\citep{1972ApJ...171..127A, 1978Ap&SS..56..191S, 1981MNRAS.194..439T, 1982MNRAS.199.1137U}.  It is well adopted in numerical general relativistic hydrodynamics \\citep{2005gr.qc.....3085A, 2006MNRAS.367.1739P, 2007MNRAS.382.1041T, 2008PhRvD..78b4023F}.  However, we believe that its true advantage is in the covariant equations, which trivially addresses problem~(ii).  \\citet{Grad1949} proposed a moment method which expands the ratio between the particle distribution and the equilibrium distribution in a multi-dimensional polynomials of the momentum.  Instead, it treats the moments as fundamental fields and keeps track of their evolutions. This method is independent of the mean free path.  Once the reference frame and temperature are chosen, \\citeauthor{Grad1949}'s coefficients are simply linear transforms of the momentum moments, which converge exponentially fast when the distribution function is well behaved\\footnote{It is a simple application of Darboux's principle. See, for example, \\citet{Boyd1999}, on convergent properties of series expansions.}.  Because photon transport is linear, \\citeauthor{Grad1949}'s moments method provides the most natural way to resolve problem~(i).  The idea of applying \\citeauthor{Grad1949}'s moment method to study radiation is not new.  A classic paper by \\citet{1981MNRAS.194..439T} presented very detailed moment formalisms for relativistic radiative transfer by using projected, symmetric, trace-free tensors \\citep{1980RvMP...52..299T}.  Shortly after that, \\citet{1982MNRAS.199.1137U} derived a 14-field approximation for general relativistic radiative transfer based on \\citeauthor{Grad1949}'s method.  However, as far as we know, these attempts were only used in one-dimensional problems \\citep[e.g.][]{1996MNRAS.281.1183Z}.  On the one hand, this is possibly due to the general concept that Newtonian formalisms are always simpler than relativistic formalisms.  On the other hand, \\citeauthor{1981MNRAS.194..439T} and \\citeauthor{1982MNRAS.199.1137U}'s 14-field methods are indeed complicated compared to the standard $\\mathcal{O}(v/c)$ equations in \\citet{1984oup..book.....M}.  Their advantages only appear when we consider more physical regimes such as the ones described in \\citet{2007ApJ...667..626K}.  Based on the above points and some additional physical understanding of moment methods, we propose a new class of closure schemes for radiative transfer and radiation hydrodynamics in this paper.  The paper is organized as the following.  In the next section, we introduce our notations and review the standard radiation hydrodynamics equations.  In section~\\ref{sec:Grad}, we describe \\citeauthor{Grad1949}'s moment method and its ultra-relativistic generalization.  We also summarize the linear transformations that we derive in the appendixes.  In section~\\ref{sec:physics}, we study carefully ideal hydrodynamics and some existing closures for radiative transfer.  They provide us important physical insight, which leads to a new class of closure schemes.  Although the scheme is generic, we specifically look at the 14-field approximations and derive the closure equations in section~\\ref{sec:14}.  Finally, we conclude this paper in section~\\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  In this paper, we propose a class of physically motivated closures for radiative transfer and radiation hydrodynamics.  We start by showing the advantages of frequency-integrated schemes, which are fully relativistic.  The transport terms and extinction terms can be easily evaluated in different reference frames.  We then apply \\citeauthor{Grad1949}'s moment method to compute the flux terms and the extinction terms from the fundamental fields.  The truncated \\citeauthor{Grad1949}'s series is energy scale and fiducial frame dependent.  For the extinction terms, the fluid comoving frame and temperature are the natural choices.  For the flux terms, however, we propose using high order decomposition to reduce unphysical photon self-interaction as well as using photon number density to obtain the energy scale.  We believe that this paper clarifies some implicit assumptions in standard closures and points out the importance of fiducial frame. Although we only present the 14-field method, it is straightforward to derive arbitrarily high order closures from our formulas (see the Appendixes).  With minimal modifications, our closures are also applicable to neutrino transport and relativistic rarefied gas dynamics.  We expect the methods derived from our new framework to outperform existing ones.  Of course, there are many open questions associated with the proposed schemes such as limiting behavior and linear stability of the theory. We will address these issues in subsequent papers.  We are also implementing the 14-field method and integrating it with some hydrodynamic solvers.  We will perform verifications and validations before applying the algorithms to study astrophysical systems.  As a final remark, moment methods are not optimal for solving all problems \\citep{2006cmt..conf.....G}.  For example, for situations with strong beaming, we are better off if we use other techniques such as ray tracing \\citep[see][and reference therein]{2002MNRAS.330L..53A, 2007ApJ...671....1T, 2009MNRAS.393.1090F}."
    },
    "0911/0911.5484_arXiv.txt": {
        "abstract": "The spectral signatures of asymmetry in Type Ia Supernova (SN Ia) explosions are investigated, using a sample of late-time nebular spectra. First, a kinematical model is constructed for SN Ia 2003hv, which can account for the main features in its optical, Near-Infrared (NIR), and Mid-Infrared (Mid-IR) late-time spectra. It is found that an asymmetric off-center model can explain the observed characteristics of SN 2003hv. This model includes a relatively high density, Fe-rich region which displays a large velocity off-set, and a relatively low density, extended $^{56}$Ni-rich region which is more spherically distributed. The high density region consists of the inner stable Fe-Ni region and outer $^{56}$Ni-rich region. Such a distribution may be the result of a delayed-detonation explosion, in which the first deflagration produces the global asymmetry in the innermost ejecta, while the subsequent detonation can lead to the bulk spherical symmetry. This configuration, if viewed from the direction of the off-set, can consistently explain the blueshift in some of the emission lines and virtually no observed shift in other lines in SN 2003hv. For this model, we then explore the effects of different viewing angles and the implications for SNe Ia in general. The model predicts that a variation of the central wavelength, depending on the viewing angle, should be seen in some lines (e.g., [Ni~II]~$\\lambda$7378), while the strongest lines (e.g., [Fe~III] blend at $\\sim 4700$\\AA) will not show this effect. By examining optical nebular spectra of 12 SNe~Ia, we have found that such a variation indeed exists. We suggest that the global asymmetry in the innermost ejecta, as likely imprint of the deflagration flame propagation, is a generic feature of SNe~Ia. It is also shown that various forbidden lines in the NIR and Mid-IR regimes provide strong diagnostics to further constrain the explosion geometry and thus the explosion mechanism. ",
        "introduction": "Type Ia supernovae (SNe~Ia) provide a powerful tool to investigate the cosmological parameters and the nature of the dark energy (Riess et al. 1998; Perlmutter et al. 1999). Thanks to the uniformity of their optical peak luminosity, once a phenomenological relation between the light curve shape and the peak luminosity is applied, SNe~Ia can be used as reliable cosmological distance indicators (Phillips 1993; Phillips et al. 1999).  There is a general consensus that SNe~Ia are thermonuclear explosions of a carbon-oxygen white dwarf (WD) (e.g., Nomoto et al. 1994, 1997; Wheeler et al. 1995; Branch 1998; Hillebrandt \\& Niemeyer 2000). A Chandrasekhar-mass WD has been favored as a progenitor for the majority of SNe~Ia (e.g., H\\\"oflich \\& Khokhlov 1996; Nugent et al. 1997; Mazzali et al. 2007).  The thermonuclear runaway starts with the ignition of deflagration bubbles (e.g., Nomoto et al. 1976, 1984). It has been suggested that the deflagration flame may turn into a detonation wave (delayed detonation model, or deflagration-detonation-transition model; Khokhlov 1991; Yamaoka et al. 1992; Woosley \\& Weaver 1994; Iwamoto et al. 1999; R\\\"opke \\& Niemeyer 2007b), although the details of the transition have not been clarified yet.  \\begin{deluxetable*}{llll} \\tabletypesize{\\scriptsize} \\tablecaption{SN~Ia Optical Nebular Spectra Sample\\tablenotemark{a} \\label{tab:tab1}} \\tablewidth{0pt} \\tablehead{ \\colhead{SN} & \\colhead{Phase (Days)} & \\colhead{[Ni~II]\\tablenotemark{b}} & \\colhead{References} } \\startdata 1986G & +103, +257 & Yes & Cristiani et al. (1992) \\\\ 1990N & +186, +227, +255, +280, +333 & Yes & G\\'omez et al. (1996)  \\\\ 1991bg & +117, +199 & No & Turatto et al. (1996) \\\\ 1991T & +258, +284, +316 & No & G\\'omez et al. (1996) \\\\ 1992G & +106, +128 & No & G\\'omez et al. (1996) \\\\ 1994D & +106 & Yes & G\\'omez et al. (1996) \\\\ 1996X & +246 & No & Salvo et al. (2001) \\\\ 1998aq & +211, +231, +241 & Yes & Branch et al. (2003) \\\\ 1998bu & +249, +329 & Yes & Cappellaro et al. (2001) \\\\ 2000cx & +125, +147, +360 & Yes & Li et al. (2001); Sollerman et al. (2004) \\\\ 2001V & +106 & Yes & Matheson et al. (2008) \\\\ 2001el & +398 & Yes & Mattila et al. (2005) \\\\ 2002dj & +275 & Yes & Pignata et al. (2008) \\\\ 2002er & +216 & No & Kotak et al. (2005) \\\\ 2003cg & +385 & No & Elias-Rosa et al. (2006) \\\\ 2003du & +138, +141, +209, +221, +272, +377 & Yes & Anupama et al. (2005); Stanishev et al. (2007) \\\\ 2003hv & +110, +143, +320, +358, +398\\tablenotemark{c} & Yes & Leloudas et al. (2009); Gerardy et al. (2007); Motohara et al. (2006) \\\\ 2004dt & +152 & Yes &  Altavilla et al. (2007) \\\\ 2004eo & +228 & No & Pastorello et al. (2007) \\\\ 2005cf & +267 & No & Leonard (2007) \\\\ \\enddata \\tablenotetext{a}{The data are compiled from the {\\it SUSPECT} data base, except for SNe 2000cx, 2001el and 2003hv. For SN 2003hv, NIR data are also given in this Table.} \\tablenotetext{b}{Whether the emission line feature at the red end of [Fe~II]~$\\lambda$7155, which we identify as [Ni~II]~$\\lambda$7378 (\\S 4 \\& 5), is detected or not at least at one epoch.} \\tablenotetext{c}{The spectra at $358$ days and $398$ days are the Mid-IR and NIR ones, respectively. } \\end{deluxetable*}   The first deflagration phase may well proceed in a very asymmetric way (Niemeyer et al. 1996; Garc\\'ia-Senz \\& Bravo 2005; Livne et al. 2005; Jordan et al. 2008). The deflagration wave propagates under the work of the buoyancy force, and thus a small perturbation in the progenitor structure could result in a global asymmetry. Rotation and convection in the progenitor WD could provide the seed for this asymmetric deflagration propagation. For example, there is a possibility that the convection in the progenitor WD is dominated by a dipole mode (Woosley et al. 2004), which likely results in highly off-axis ignition and propagation of the deflagration flame (R\\\"opke et al. 2007c; Kasen et al. 2009).   It is, however, observationally challenging to put constraints on the geometry of the explosion. Measurements of polarization suggest that a large global asymmetry does not exist in SNe~Ia (Wang et al. 1996), with only a few exceptions (Howell et al. 2001). However, the polarization probes mainly the outer regions of the expanding SN ejecta, at least with the existing telescopes and instruments. The signature of the possible asymmetry in the deflagration phase, however, can only be probed by looking deeper into the innermost regions. In this respect, late-phase ($\\sim 1$ year past the explosion) spectroscopy can provide an important diagnostics. Following the homologous expansion, the SN ejecta become transparent to optical and longer wavelengths, and thus emission line profiles can be used to probe the distribution of elements that emit the light or that deposit the energy. This strategy has been applied to investigate the aspherical nature of core-collapse SNe from massive stars (e.g., Maeda et al. 2002; Chugai et al. 2005; Mazzali et al. 2005; Maeda et al. 2008; Modjaz et al. 2008; Taubenberger et al. 2009; but see also Milisavljevic et al. 2009).  Probing the explosion geometry of SNe Ia using late-time spectroscopy is still a fairly young field. H\\\"oflich et al. (2004) presented Near-Infrared (NIR) spectra for SN Ia 2003du taken with the {\\it Subaru} telescope. They discussed that a flat-topped profile of  the [Fe~II]~1.644~$\\micron$ emission feature is likely a result of a hole in the distribution of $^{56}$Ni. Motohara et al. (2006) added two other Near-Infrared (NIR) spectra for SNe~Ia taken with the {\\it Subaru} telescope (SNe 2003hv, and 2005W). The [Fe~II]~1.257~$\\micron$ and 1.644~$\\micron$ emission lines displayed different profiles, and especially, those in SN~2003hv were blueshifted by $\\sim 2,000 - 3,000$~km~s$^{-1}$. Motohara et al. (2006) suggested that this can be interpreted as evidence of a global asymmetry in SN~2003hv. Gerardy et~al.~(2007) also found a similar blueshift, in the [Co~III]~$11.88~\\micron$ line detected in their Mid-Infrared (Mid-IR) spectrum of SN 2003hv, taken with the {\\it Spitzer Space Telescope}. On the other hand, the late-phase optical spectra, available for more than 20 normal SNe~Ia, all look very similar to each other, and seem to show little evidence for any asymmetry.  Recently, Leloudas et al. (2009) reported on observations of SN~2003hv, including late-time optical spectra. Thus, it is now for the first time possible to examine a late-time spectrum of a SN~Ia all the way from the optical to the Mid-IR. We want to investigate if the kinematical interpretation for the blueshifts in the NIR and Mid-IR forbidden lines of SN~2003hv is consistent with the optical spectrum.  In this paper, we examine late-time spectra of SN 2003hv, using multi-dimensional radiation transfer calculations and a simple kinematic model, and arrive at a plausible explosion geometry of SN 2003hv. Using the structure that was successfully applied to SN~2003hv, we then investigate signatures of the asymmetry and the expected diversity in the late-phase spectra, resulting from various viewing orientations. We also compile published late-time optical spectra of SNe~Ia, and investigate to what extent the expected diversity exists in the data. In doing this, we identify a possible signature of the explosion asymmetry in late-time optical spectra of SNe~Ia.  This paper is organized as follows. In \\S 2, the observational data of SN~2003hv and other SNe~Ia are summarized. In \\S 3,  our model and the method of multi-dimensional spectrum synthesis are presented. The results are shown in \\S 4 \\& 5; \\S 4 focuses on the geometry of SN 2003hv, while the implications for SNe Ia in general are presented in \\S 5. The paper is closed with a discussion and our conclusions in \\S 6. Future observing strategies to further constrain the explosion mechanism are also presented in \\S 6. ",
        "conclusions": "\\subsection{Summary}  In this paper, we have investigated the geometry of the innermost region of  SNe Ia, which should provide information on the explosion mechanism. We have shown that some enigmatic properties observed in the late-time optical, NIR, and Mid-IR spectra of SN Ia 2003hv -- that some emission lines show blueshift while others do not -- can in fact be naturally explained by a simple kinematic model, if the viewing angle is close to the direction of the off-set. In this model, the high density regions (the ECAP-zone dominated by stable Fe-peak elements, and the HD-zone dominated by $^{56}$Ni) are off-set with respect to the center of the expansion by $\\sim 3,500$ km s$^{-1}$, and are surrounded by a less asymmetric,  low density $^{56}$Ni-rich region (the LD-zone).  Table 4 summarizes the expected qualitative behavior of strong emission lines in late-time nebular SN Ia spectra.  Forbidden lines from ions with high ionization stage and/or high excitation temperature are mostly emitted from the outer, \"spherical\" LD-zone.  Such lines show little variation in the central wavelength, irrespective of the viewing angle.  On the other hand, lines from ions with low ionization stage and/or  low excitation temperature are generally emitted from the inner, \"off-set\" ECAP- or HD-zones. These lines show a strong variation in the central wavelength as a function of the viewing angle.  \\begin{figure} \\epsscale{1.0} \\plotone{f10.eps} \\caption {The number distribution of SNe Ia as a function of the central wavelength of the [Ni~II]~$\\lambda$7378 feature. The observed number distribution, using 12 SNe~Ia, is shown by shaded bars. The expected distribution from the present model (as a function of the viewing angle) is shown by the line. \\label{fig:fig10}} \\end{figure}    Encouraged by this finding, we have investigated available optical late-time spectra of SNe Ia compiled mostly from the {\\it SUSPECT} archive. Figure 9 summarizes our results from the model calculations, and the measurements in the 12 observed SNe~Ia sample. From the model, we see that the variation in the central wavelength of [Ni~II]~$\\lambda$7378 basically follows the simple kinematic expectation. Note, however, for a precise prediction of the behaviors of emission lines, relative contributions from the different zones have to be considered, which further depend on the ionization structure, the electron temperature and the density. For [Ni~II]~$\\lambda$7378, the effects of these details turn out to be relatively small, but some lines are indeed affected by these complications (\\S 5). We have provided an explanation on the observational behavior that [Ni~II]~$\\lambda$7378 shows the diversity in the central wavelength while [Fe~III] blend at 4,700\\AA\\ does not: they trace different zones, and the inner ECAP and HD zones have the large off-set while the outer LD zone does not.  We suggest that [Ni~II]~$\\lambda$7378 to be one of the few lines at optical wavelengths to provide strong diagnostics of the innermost region of SNe Ia. Among the 20 SNe Ia, about half of them (12 SNe) show a probable detection of this feature. Similar diagnostics can be obtained by [Fe~II]~$\\lambda7155$ as well. Thus, the spectra in the wavelength range of $\\sim 7,000 - 7,500$\\AA\\ should provide useful diagnostics of the explosion asymmetry. With present instrumentation, these features can be more easily observed than other strong indicators in the NIR and Mid-IR regime. A late-time spectrum of an overluminous SN Ia 2006gz (Hicken et al. 2007) also shows a hint of the possible blueshift in these features (Maeda et al. 2009b). However, the spectrum is very noisy and the identification is not secure. Also, the spectrum was peculiar in a sense that it did not show the strong features in the blue (i.e., at $\\sim 4,700$\\AA). For these reasons, we have not included SN 2006gz in our analysis, despite the possible importance of asphericity in this object (Hillebrandt et al. 2007; Sim et al. 2007; Maeda \\& Iwamoto 2009a).    The geometry we derived could naturally arise in SN Ia explosions. The initial deflagration may proceed in a very asymmetric way, while the subsequent detonation is expected to leave $^{56}$Ni more spherically distributed. An important question is whether the geometry suggested for SN 2003hv represents a special case or a general structure of SN Ia explosions. Figure~10 shows the distribution of 12 SNe Ia, as a function of the velocity shift in [Ni~II]~$\\lambda$7378. Although the sample is small, the general agreement between the observations and the expectation assuming the geometry of SN 2003hv as a generic feature of SN Ia explosions, is very encouraging.  From this analysis, it seems that SN 2003hv may not be very different from other SNe~Ia, and its largest velocity shifts can be attributed to the viewing angle effect.  \\subsection{Temporal Evolution}  \\begin{figure} \\epsscale{1.0} \\plotone{f11.eps} \\caption {Examples of the temporal evolution in (a) the region around the [Fe~III] blend at 4700\\AA\\ and (b) that around [Ni~II]~$\\lambda$7378. The data are taken from Cristiani et al. (1992) for SN~1986G, Leloudas et al. (2009) for SN~2003hv, and Li et al. (2001) and Sollerman et al. (2004) for SN~2000cx. \\label{fig:fig11}} \\end{figure}    We have found that the line center of [Ni~II]~$\\lambda$7378 does not correlate with the epoch at which the spectrum was taken. Therefore, the line velocity shift should reflect the intrinsic geometry and variation in the viewing angle. On the contrary, the central wavelength of the [Fe~III] blend at 4700\\AA\\  does show a clear correlation with the epoch; The line always shows a blueshift before $\\sim 200$ days, while there is no significant wavelength shift in spectra taken after $\\sim 200$ days (Fig. 9). This is  exemplified by the temporal evolution of the nebular spectra, available for a few SNe Ia (Fig. 11): Thus, we suggest that the [Fe~III] traces the intrinsic geometry only after $\\sim 200$ days, and that the outer, relatively low density region is less aspherical than the innermost region.  Figure~11 shows that there is also diversity in the temporal evolution of [Ni~II]~$\\lambda$7378. In SNe~1986G and 2003hv, [Ni~II]~$\\lambda$7378 persisted over a long period, while it only shows a transient feature in SN~2000cx. This may indeed be consistent with the model calculation; in our calculation, the strength of [Ni~II]$\\lambda$7378 is largely reduced at $\\sim 350$ days (Figs. 2 \\& 4), following the temperature drop and the so-called infrared catastrophe in the ECAP zone (see also \\S 6.3).  The central wavelength of the [Ni~II] line does not evolve in SN~1986G, while it shows a somewhat larger blueshift in the later phase in  SN~2003hv. The difference is likely related to the distribution of different species, and the resulting  $\\gamma$-ray deposition and potentially the positron escape. An increasing degree of the blueshift in SN~2003hv as a function of time might indicate that $^{56}$Ni is distributed inhomogeneously within the ECAP zone, being relatively concentrated at the high velocity \"edge\" of this zone (e.g., $\\sim 2,000 - 3,500$ km s$^{-1}$).  In such a scenario, $\\gamma$-rays irradiate the whole ECAP zone early on, while later the positrons heat only the relatively \"high-velocity\" region.  The number of SNe~Ia with a good temporal coverage at late phases is still small, and it is strongly encouraged to perform multi-epoch spectroscopy. Such intensive observations are useful to understand (1) the detailed distribution of $^{56}$Ni, and (2) detailed thermal processes (e.g., $\\gamma$-ray deposition, positron escape, IR catastrophe) in the SN Ia nebulae. It is also important to investigate (3) a possible evolution in the line profiles if we are to use them to investigate the details of the explosion physics (see \\S 6.3 for more details).  \\subsection{Future Perspectives}  The conclusion in the present paper that the late-time nebular spectra can be used to investigate the geometry of ejecta, thus the explosion mechanism of SNe Ia, provides a new strategy to clarify the nature of SN Ia explosions. Presently, the sample is still limited: The late-time nebular spectra are available for $\\sim 20$ SNe Ia in the optical wavelength. There are only a few examples in the NIR and Mid-IR. For most of them, a temporal sequence of late-time spectra are not available. In this section, we highlight the importance of expanding the sample of late-time nebular spectra of SNe Ia.  In doing this, we summarize predictions from recent theoretical models. Two scenarios have been proposed to create the ECAP zone in the context of the asymmetric ignition of the initial deflagration bubbles. In the first scenario, many ignition points are clustered near the center of the white-dwarf with possible bulk off-set (\\S 1), producing the ECAP zone by the deflagration flame. In the second scenario, one or a few ignition points are placed far away from the center of the white dwarf. The deflagration is weak and does not produce much neutron-rich materials. The deflagration flame rises to the surface without initiating the detonation.  The detonation may then be triggered near the white dwarf surface, at the opposite side of the initial deflagration ignition (so-called Gravitationally Confined Detonation model; Jordan et al. 2008). In this case, the initial deflagration is so weak and the white dwarf suffers from virtually no expansion before the detonation passes through the central region. Thus, the electron capture reactions in the detonation are not negligible, producing the ECAP zone (Meakin et al. 2009). This model predicts the off-set of the ECAP zone toward the opposite direction of the initial deflagration ignition. This second scenario could actually be regarded as an extreme case of the first scenario in terms of the distribution of the initial deflagration bubbles.  Although these two scenarios are qualitatively different to one another for the origin of the ECAP zone, the resulting configuration (e.g., the off-center ECAP zone) may be similar to what we derived in this paper. In the second scenario, however, the detonation should produce a large amount of $^{56}$Ni in order to produce the ECAP zone more massive than $\\sim 0.1 M_{\\odot}$ (Meakin et al. 2009), at least their model sequence presented to date (e.g., $\\sim 1.1 M_{\\odot}$ of $^{56}$Ni in their 2D models). This scenario therefore may not be favored as the origin of the ECAP zone for the particular case of SN 2003hv, which is relatively faint as a normal SN Ia (Leloudas et al. 2009). This second scenario is still a possibility to explain the wavelength shift in the nebular emission lines seen in bright SNe Ia, which also show the signature of the asymmetry in the ECAP zone (\\S 5 \\& \\S 6.1). Although we expect that the Gravitationally Confined Detonation (GCD) model results in bright SNe Ia if the detonation is to produce the ECAP zone, the detailed relation between the brightness and the mass of the ECAP zone should be dependent on details of the initial conditions such as the ignition density (which depends on the accretion rate of a white-dwarf before the explosion; e.g., Nomoto 1982; Nomoto et al. 1984). Further study based on the GCD scenario is therefore important.  We suggest that future observations of late-time SN Ia emission can provide important information to understand the details of the explosion mechanism(s). Hereafter, we discuss importance of the following (future) observations, especially in view of the available theoretical ideas as mentioned above. \\begin{itemize} \\item Expanding the sample of the line velocity shift measurement, especially in the optical wavelengths. \\item Studying the detailed line profiles, especially in the NIR. \\item Investigating the mass of the ECAP zone. \\end{itemize}   \\subsubsection{Expanding the sample for the line velocity shift measurement}  A series of hydrodynamic simulations based on the first scenario provide interesting predictions (Kasen et al. 2009). More asymmetric distribution of the initial deflagration results in a larger amount of $^{56}$Ni produced by the detonation, yielding brighter SNe. For a larger degree of the off-set in the initial deflagration bubbles, we should also expect a larger degree of the off-set in the ECAP zone produced by the deflagration (Maeda et al. 2009c). We therefore expect that brighter SNe Ia show a larger off-set velocity in the ECAP zone in this scenario. Since the wavelength shift seen in nebular spectra of individual SNe Ia is a result from a combination of the degree of the asymmetry and the viewing angle, this relation should be observationally investigated in a statistic way with a large number of SNe Ia. We suggest to investigate the distribution of the line-velocity (Figure 10) as a function of the luminosity of SNe Ia (or so-called $\\Delta m_{15}$ in the light curve shape which correlates with the luminosity). The scenario predicts that brighter SNe Ia should show a wider distribution in the velocity space, while fainter SNe Ia should show more narrowly peaked distribution at zero velocity.  The GCD model is an extreme end of the model sequence, predicting bright SNe Ia. The mechanism to create the ECAP zone is different from the first scenario. Thus, it may well give the distribution of the velocity shift as a function of the luminosity different from the first scenario.  These issues cannot be addressed with the current sample of SNe Ia for which nebular spectra are available -- the current data are consistent with the idea that SNe Ia have the generic off-set as derived for SN 2003hv (Fig. 10), but it does not reject a possibility that the off-set velocity is dependent on the luminosity because of the small sample. Increasing the number of the sample will tell us (1) whether brighter SNe Ia have larger off-set in the ECAP zone (a test for the first explosion scenario), and (2) whether the distribution of off-set of the ECAP zone is explained by a single scenario or hybrid scenarios, especially in bright SNe Ia (a test for the GCD model).  \\subsubsection{Detailed Line Profiles especially in the NIR}  In this paper, we have focused on the line wavelength shift, but have not tried to fit the details of the line profiles. H\\\"oflich et al. (2004) and Motohara et al. (2006) suggested to use the detailed line profile of [Fe~II]~$1.644~\\micron$ to probe the distribution of electron capture products (i.e., the ECAP-zone in our model). H\\\"oflich et al. (2004) suggested that the apparently flat-topped profile of this line in SN 2003du is a signature of a \"hole'' in the region emitting the [Fe~II]. They attributed this hole to the existence of an inner $^{56}$Ni-free, electron capture products-rich region, since such a region lacks the heating source when the ejecta become transparent to $\\gamma$-rays. The same suggestion was also made for SN 2003hv based on better data (Motohara et al. 2006). There are also other possibilities for the flat-topped profile, e.g., the infrared catastrophe (Leloudas et al. 2009).  Understanding the origin of the details of the line profiles requires further studies (e.g., H\\\"oflich et al. 2006). For the \"hole\" interpretation of the flat-topped profiles, it should be examined whether such a configuration can result from hydrodynamic explosion models.  An intensive model survey based on the first scenario has been presented by Kasen et al. (2009) in 2D. According to their model sequence, we would expect that more asymmetric distribution of the initial deflagration bubbles could result in a larger amount of $^{56}$Ni produced by the detonation near the zero-velocity. As such, we expect that bright SNe Ia tend to show the peaked line profiles. Faint SNe Ia may tend to show the flat-topped profile (see below). The most solid case showing the flat-top profile in the NIR [Fe~II] so far is SN 2003hv. This is a relatively faint SN Ia, thus being consistent with this theoretical expectation. Also, the bright SN 1991T showed the peaked profile (Bowers et al. 1997) and it is also consistent with the expectation.  Apart from the above prediction, an issue remains whether the first scenario can produce the flat-topped profiles. The current state-of-art 3D deflagration model according to this scenario indicates some amount of unburned C+O is mixed down to the center (e.g., R\\\"opke et al. 2007a). These materials may be later detonated to produce $^{56}$Ni (R\\\"opke \\& Niemeyer 2007b; Kasen et al. 2009). Thus, it is expected that the ECAP zone and the HD zone are mixed, rather than form the distinctly layered structure. An interesting question is whether future simulations can lead to the \"hole\" configuration with some initial conditions for the thermonuclear bubbles, and whether different initial conditions can result in a different degree of the mixing (to explain the diversity in the detailed profile).  In the second, GCD scenario, the ECAP zone would not suffer from the mixing process since it is created by the detonation. It is expected to be always a \"hole\", but the size of the zone may well depend on the strength of the initial deflagration. For the weaker deflagration, the density of the white-dwarf at the detonation is higher, and thus the size and mass of the ECAP zone are larger (Meakin et al. 2009). This raises an interesting possibility that the variation in the size of the ECAP zone may explain the variation in the detailed line profile; brighter SNe Ia are expected to more likely produce a flat-top profile. This tendency is different from the expectation from the first scenario.    In our kinematic model, we have the \"hole\" in the form of the ECAP-zone, which is assumed to be macroscopically separated from the outer $^{56}$Ni-rich HD/LD-zones. In this configuration,  our model makes the following predictions; (1) Emission lines formed in the ECAP-zone have narrow widths ($\\sim 3,000$~km~s$^{-1}$),  and have a peaked profile (with respect to the \"off-set\" wavelength). This is exemplified by [Ni~II]~$\\lambda$7378, [Ni~II]~$6.634~\\micron$, [Ni~III]~$7.350~\\micron$. (2) Emission lines originating from the HD-zone have flat-topped profiles. Typical examples are [Fe~II]~$\\lambda$7151, [Fe~II]~$1.257~\\micron$, [Fe~II]~$1.644~\\micron$, and [Co~III]~$11.88~\\micron$. (3) Emission lines dominated by the LD-zone have broad, peaked profiles, since the size of the \"hole\" is much smaller than the emitting radius.  Typical examples are [Fe~III]~$2.218~\\micron$, and [Fe~III]~$2.904~\\micron$. The [Fe~III] blend at 4700\\AA\\ also has the same tendency, although the blended nature of this feature makes the situation complicated. Some lines show a combination of these characteristics; for example, [Fe~II]~$17.93~\\micron$ is predicted to show a peaked, relatively broad profile with a  wavelength shift depending on the viewing angle (\\S 5.4). Thus, looking at various emission lines, we will be able to obtain different, independent information on the  explosion physics.   The best lines for investigating the details of the distribution of $^{56}$Ni and the hole (i.e., the ECAP zone) are [Fe~II]~1.257~$\\micron$ and [Fe~II]~1.644~$\\micron$, thanks to their low excitation temperature and the relatively isolated nature of these lines. Adding to this, [Ni~II]~$\\lambda$7378 and some Mid-IR lines (Table 4) can give the direct view of the ECAP zone distribution. Thus, expanding the sample of late-time NIR spectra is highly encouraged, preferentially with the optical and (if possible) mid-IR data taken almost simultaneously. This will tell us (1) whether brighter SNe Ia have a larger amount of the low-velocity $^{56}$Ni (a test for the first explosion scenario), (2) whether the fainter SNe tend to show the flat-topped profile (a constraint on the distribution of the deflagration bubbles), and (3) whether the brightest SNe Ia show the hole of the $^{56}$Ni distribution (a test for the GCD model as the origin of the brightest SNe Ia).    \\subsubsection{Mass of Neutron-rich Fe-peak Elements}  In \\S 6.2, we have mentioned an importance of multi-epoch spectroscopy to clarify details of thermal processes within SN Ia nebulae. This is also important to clarify the explosion physics.  Different scenario should result in the different mass of the ECAP zone. The first scenario predicts a smaller mass of the ECAP zone for brighter SNe Ia (since the weaker detonation is followed by the stronger detonation; Kasen et al. 2009). The trend should be opposite in the GCD model sequence (since the neutron-rich materials and $^{56}$Ni are both created by the detonation). Deriving the mass of neutron-rich materials (i.e., ECAP zone) as a function of the SN Ia peak luminosity will therefore give us a strong test for these scenarios. No significant variation of the mass of neutron-rich materials as a function of the luminosity (or $\\Delta m_{15}$) has been found so far (Mazzali et al. 2007). Understanding the thermal properties will provide the better accuracy in the mass estimate so that the investigation of the details of the explosion physics becomes possible."
    },
    "0911/0911.1955_arXiv.txt": {
        "abstract": "The FIRAS data are independently recalibrated using the WMAP data to obtain a CMB temperature of $2.7260\\pm 0.0013$. Measurements of the temperature of the cosmic microwave background are reviewed. The determination from the measurements from the literature is cosmic microwave background temperature of \\tf$\\pm 0.00057$~K. ",
        "introduction": "The Wilkinson Microwave Anisotropy Probe (WMAP) data presents an opportunity to recalibrate the Far InfraRed Absolute Spectrophotometer (FIRAS) experiment and produce an independent check of the other measurements of the cosmic microwave background (CMB) temperature. In sections 2~\\&~3 the WMAP data will be presented and combined with the FIRAS data to make an independent estimate of the CMB temperature. In sections 4~\\&~5 this new estimate will be combined with others from the literature to generate an improved estimate for the CMB temperature.  The WMAP data only measures the difference in intensity between different points on the sky. However, the precision is sufficient such that the velocity of the WMAP spacecraft can be used to calibrate the velocity to various points on the surface of last scattering of the CMB. This velocity in turn is used to form a differential spectrum of the CMB. The differential spectrum is then fit with a single parameter which is the CMB temperature. ",
        "conclusions": "The calibration methods for the Far Infrared Absolute Spectrophotometer (FIRAS) have been described and the accuracy estimated. All of the recent precision estimates of the CMB temperature agree within 2.5 times their uncertainties. These estimates were made with a variety of methods from different platforms and different frequencies. Combining all of the estimates results in a very modestly elevated $\\chi^2$ and an improved absolute temperature estimation of $\\tf\\pm 0.00057$~K."
    },
    "0911/0911.3161_arXiv.txt": {
        "abstract": "We have developed a semi-analytic method of parameterizing $N$-body simulations of self-gravity wakes in Saturn's rings, describing their photometric properties by means of only six numbers:  three optical depths and three weighting factors.  These numbers are obtained by fitting a sum of three gaussians to the results of a density-estimation procedure that finds the frequencies of various values of local density within a simulated ring patch.  Application of our parameterization to a suite of $N$-body simulations implies that rings dominated by self-gravity wakes appear to be mostly empty space, with more than half of their surface area taken up by local optical depths around~0.01.  Such regions will be photometrically inactive for all viewing geometries.  While this result might be affected by our use of identically-sized particles, we believe the general result that the distribution of local optical depths is trimodal, rather than bimodal as previous authors have assumed, is robust.  The implications of this result for the analysis of occultation data are more conceptual than practical, as we find that occultations can only distinguish between bimodal and trimodal models at very low opening angles.  Thus, the only adjustment needed in existing analyses of occultation data is that the model parameter $\\tau_{gap}$ should be interpreted as representing the area-weighted average optical depth within the gaps (or inter-wake regions), keeping in mind the possibility that the optical depth within those inter-wake regions may vary significantly.  The most significant consequence of our results applies to the question of why ``propeller'' structures observed in the mid-A~ring are seen as relative-bright features, even though the most prominent features of simulated propellers are regions of relatively low density.  Our parameterization of self-gravity wakes lends preliminary quantitative support to the hypothesis that propellers would be bright if they involve a local and temporary disruption of self-gravity wakes.  Even though the overall local density is lower within the propeller-shaped structure surrounding an embedded central moonlet, disruption of the wakes would flood these same regions with more ``photometrically active'' material (i.e., material that can contribute to the rings' local optical depth), raising their apparent brightnesses in agreement with observations.  We find for a wide range of input parameters that this mechanism indeed can plausibly make propellers brighter than the wake-dominated background, though it is also possible for propellers to blend in with the background or even to remain dark.  We suggest that this mechanism be tested by future detailed numerical models. ",
        "introduction": "} Saturn's dense A~and B~rings are pervaded by a microstructure, dubbed ``self-gravity wakes,'' of alternating dense and rarified regions that arise due to a rough balance between the clumping together of particles under their mutual self-gravity and their shearing apart again due to tidal forces \\citep{JT66,Salo92,Richardson94}.  The characteristic alignment of elongated self-gravity wake structures causes variation with viewing geometry in the brightness of the rings as seen in ground-based optical \\citep{Camichel58,Salo04,French07} and microwave observations \\citep{Dunn02,Dunn04} as well as spacecraft images \\citep{Franklin87,Porco08}, and in the rings' opacity as mesaured by stellar occultations \\citep{Colwell06,Colwell07,Hedman07}, which in the latter case has led to empirical determination of some wake properties.  Existing analyses of stellar occultation data from \\Cassit{}~UVIS \\citep{Colwell06,Colwell07} and \\Cassit{}~VIMS \\citep{Hedman07,NH09} explain the observations through the use of simple models that assume an optical-depth dichotomy, with nearly-opaque wakes (with optical depth $\\tau_{wake}$) and a low but relatively constant optical depth in the spaces between the wakes ($\\tau_{gap}$). However, while some preliminary work has shown that such a model can produce brightness profiles that are consistent with data \\citep{Hedman07,SaloLondon08}, it has not been verified that such a bi-modal model describes the actual nature of simulated wakes, not to mention real ones. And if it does, how do $\\tau_{wake}$ and $\\tau_{gap}$ relate to environmental parameters such as overall occultation optical depth and coefficient of restitution?  What do the values of $\\tau_{gap}$ inferred from observations tell us about the conditions under which wakes occur?  \\citet{Salo04} and \\citet{Porco08}, on the other hand, produce plots of ring brightness at various geometries for various simulation input parameters, inferring the best input parameters by finding simulated brightnesses that best compare to observations.  But where is the brightness coming from?  For images in reflected light, is the brightness varied primarily by changes in the opacity of dense wake structures, or by the fractional area covered by them?  For images in transmitted light, do more photons get scattered into the camera from the rarefied regions between the wakes (the wakes themselves being largely too opaque to allow much light to pass through them), or are more photons being scattered by the edges of the wake structures?  Better understanding of these questions would help to guide future simulations, and might allow one to infer more from the available observational data.  Yet another current question concerns the nature of ``propeller'' structures observed in images of the mid-A ring \\citep{Propellers06,Propellers08,Sremcevic07}, each of which is inferred to be a disturbance surrounding an unseen moonlet embedded within the ring.  Observed propellers generally show the disturbed region as brighter than the background ring brightness, though theoretical models \\citep{SS00,SSD02,Seiss05,LS09} indicate that the most prominent feature of the disturbance is a decrease in local density.  Among other hypotheses for this curious behavior, \\citet{Propellers06,Propellers08} have proposed that self-gravity wakes tend to lock up ring material into a photometrically inactive state, and that propellers can release this material by locally and temporarily disrupting the wake structure, thus causing the disturbed area to contain more photometrically active material even if it contains less material overall.  Our parameterization of self-gravity wakes allows us to evaluate this hypothesis in a more quantitative fashion than has previously been possible.  A major reason why these questions remain largely unanswered lies in the difficulty of \\textit{local density estimation}, the quantification of an underlying density function based on the locations of a finite number of particles, for a strongly heterogeneous density distribution.  The simplest method of density estimation is to divide a simulation space into bins of equal size and to count the number of particles in each bin, but this fails to give meaningful results for current simulations of self-gravity wakes.  This is because the rarefied regions require large bins in order to accumulate enough particles in each bin to give good statistics, but such large bins will badly smear the sharp boundaries between the dense wakes and the rarefied regions.  Contrariwise, if the bins are small enough to resolve the sharp boundaries, nearly all bins in the rarefied regions will contain zero particles (with a few containing one particle), giving no sense of the average density in the rarefied regions.  We have constructed a density estimation method, to be applied to self-gravity wake simulations, based on circular bins that expand in rarefied regions and contract in dense regions.  A given circular bin is only used if the particles within it satisfy certain criteria (described in detail below) to ensure that they are sufficiently evenly distributed; if those criteria are met, then the bin grows until they are no longer met.  The overlapping circular bins are then projected onto a finely-meshed rectangular grid; nearly all bins in the latter grid contain multiple overlapping circular bins, and the value of the density at each location is taken from the average of the overlapping circular bins.  Applying our density estimation method to a set of numerical simulations of a ring patch characterized by self-gravity wakes, we parameterize the general properties of each simulation as a weighted combination of three optical depth values.  This semi-analytic treatment allows us to comment on trends that are applicable to a wide range of input conditions.  We specifically discuss the implications of our results for previous interpretations of occultation data, as well as for the question of why ``propellers'' appear as relative-bright features.  Section~\\ref{WakeSims} describes our suite of numerical simulations, the results of which constitute the input data for our semi-analytic method.  Section~\\ref{DensEst} describes our density estimation method.  Section~\\ref{Results} describes our results, and Section~\\ref{Discussion} provides further discussion.  Section~\\ref{Conclusions} presents a summary and conclusions. ",
        "conclusions": "} We have developed a semi-analytic method of parameterizing simulations of self-gravity wakes in Saturn's rings, describing their photometric properties by means of only six numbers:  three optical depths and three weighting factors.  These numbers are obtained by fitting a sum of three gaussians to the results of a density-estimation procedure that finds the frequencies of various values of local density within a simulated ring patch.  In order to account for the conversion from dynamical optical depth to photometric optical depth, we use the expression $\\tau_{phot}/\\tau_{dyn} \\simeq 1 + kD$ \\citep{SK03}, where $D$ is the volume filling factor (which we estimate); we find the best results by setting the scalar $k=1$.  Our first surprising result is that rings dominated by self-gravity wakes appear to be mostly empty space, with more than half of their area taken up by local optical depths around~0.01.  Such regions will be photometrically inactive for all viewing geometries.  While this result might be affected by our use of a monodisperse size distribution, we suggest that it might be robust due to the lower coefficient of restitution that we used in our $N$-body simulations, following \\citet{Porco08}.  If our models are in fact robust, then the bimodal density distribution assumed in the interpretation of occultation data by previous investigators should be replaced by a \\textit{tri}modal distribution.  The practical results of this turn out to be minimal, as occultations can only distinguish between bimodal and trimodal models at very low opening angle.  Existing analysis of occultation data should be interpreted with $\\tau_{gap}$ representing the area-weighted average optical depth within the gaps (or inter-wake regions), keeping in mind the possibility that there may be strong variations in optical depth within those inter-wake regions.  Applying our parameterization of self-gravity wakes lends preliminary quantitative support to the hypothesis of \\citet{Propellers08} that ``propellers'' observed in the mid-A ring are bright because of a local and temporary disruption of self-gravity wakes.  Even though the overall local density is lower within the propeller-shaped structure surrounding an embedded central moonlet, disruption of the wakes would flood these same regions with more photometrically active material, raising their apparent brightnesses in agreement with observations.  We find that this mechanism can plausibly work for a wide range of input parameters.  We hope that this hypothesis will eventually be directly tested by detailed numerical simulations.  However, this is presently very difficult due to the high expense in terms of computational resources necessary to account for the small lengthscales appropriate for self-gravity wakes (a few meters) while simultaneously accounting for the very large lengthscales appropriate for propellers (several km).  The former essentially requires a large number of particles per unit area, while the latter requires a large patch size.  As computational resources increase to the point that such simulations become feasible, it will become possible to determine whether propellers really are characterized by a local and temporary disruption of self-gravity wakes, as well as whether the photometric properties of propellers can be explained using the hypothesis we have outlined.  In the meantime, the results of our semi-analytic method give us reason to hope that the solution to the problem indeed lies in this direction."
    },
    "0911/0911.1214_arXiv.txt": {
        "abstract": "Long gamma-ray bursts (GRBs) are associated with the explosion of massive stars in star forming regions. A large fraction of GRBs show intrinsic absorption as detected in optical spectra but absorption signatures are also detectable in afterglow X--ray spectra. We present here a comprehensive analysis the full sample of 93 GRBs with known redshift promptly observed by Swift XRT up to June 2009. The distribution of X--ray column densities clearly shows that GRBs are heavily absorbed indicating that they indeed occur in dense environments. Furthermore, there is a  lack of heavily absorbed GRBs at low redshift ($z\\lsim 1-2$) that might therefore be candidates for the missing `dark' GRB population. However, there is no statistically significant correlation between the amount of X--ray absorption and the `darkness' of a GRB. Finally, we compare the hydrogen column densities derived in the optical with those derived from X--ray absorption. The two distributions are different, with the optical column densities being lower than the X--ray ones, which is even more apparent when correcting for metallicity effects. The most likely explanation is photoionization of hydrogen in the circumburst material caused by the radiation field of the burst. ",
        "introduction": "(Long) gamma--Ray Bursts (GRBs) are recognized as being powerful cosmological tools, being detected from the local Universe to the highest redshifts reached today ($z\\sim 8.2$) (Salvaterra et al. 2009; Tanvir et al. 2009). GRBs are so bright that they can illuminate, for a short time, regions of distant galaxies otherwise inaccessible and probe the cosmic chemical evolution of the Universe (e.g. Fiore et al. 2005; Prochaska, Chen \\& Bloom 2006; Fox et al. 2008). The GRB light shines through a medium that is very different (i.e. denser) from the typical interstellar medium (ISM) observed with traditional tools in distant galaxies High-resolution optical spectroscopic observations allow us to probe the relative abundances of different elements highlighting the physical state of the gas, its chemical enrichment and dust content, and the nucleosynthesis processes in the GRB star-forming environment (Savaglio 2006; Prochaska et al. 2007). Despite the poor number statistics, a comparison between the metallicities obtained by GRBs and those from studies of absorbing systems in the sightlines of quasar show that the former tend to be more metal-rich (on average by a factor of 5 at high-$z$) and to have a flatter or no redshift evolution (see, e.g., Fynbo et al. 2008, 2009; Savaglio, Glazebrook \\& LeBorgne 2009; Chen et al. 2009). Variations in the intensity of optical absorption lines were detected in a few cases of very optically bright GRBs. These variations help to constrain the distance of the absorbing gas to the GRB. This has been done in only two cases where distances larger than $\\sim 2$ kpc have been derived. Matter within this radius has been photoionized by the GRB radiation flux, which means that optical observations can only probe the outermost part of the region surrounding the GRB (D'Elia et al. 2009; Vreeswijk et al. 2007).  Matter along the line of sight also attenuates X--ray photons. X--rays are photoelectrically absorbed by metals, producing a well known and well identifiable pattern in the observed X--ray spectra (Morrison \\& McCammon 1983). In contrast to absorption in the optical range, the absorption of X--rays is not sensitive to the state of the element (either gas or solid), therefore these measurements provide an unbiased view of the total absorbing column density of material. As demonstrated by several studies the X--ray absorbing column towards GRBs is high (Stratta et al. 2004; Campana et al. 2006; Watson et al. 2007), which implies in principle, i.e. for any known dust to gas ratio, a high dust extinction. In contrast, GRB optical afterglows, when present, often show a lack of reddening. Several suggestions have been advanced to explain this dichotomy but a clear understanding has not been reached yet (e.g. Galama \\& Wijers 2001, see also Nardini et al. 2009).  Stratta et al. (2004) presented a systematic analysis of a sample of 13 bright afterglows observed with the BeppoSAX narrow field instruments. In only two cases they found a significant detection of additional intervening material in excess of the one in our Galaxy. However, due to the limited photon statistics, they could not exclude that intrinsic X--ray absorption was also present in the other bursts. Campana et al. (2006) presented a systematic study of the intrinsic column density absorption in 6 GRBs with known redshift observed by Swift. Adding these GRBs  to others found in the literature with known redshift (mainly Chandra and XMM-Newton observations), they concluded that there is strong evidence that GRBs are intrinsically absorbed (Campana et al. 2006). The distribution of absorbing column densities is consistent with GRBs occurring within molecular clouds, which supports the link between (long) GRBs and star formation. Watson et al. (2007) found a large discrepancy between optical and X--ray derived GRB column densities and suggested photoionization of the gas cloud as the most likely cause.  Here we consider the full sample of 93 GRBs detected by Swift (Gehrels et al. 2004) XRT (Burrows et al. 2005) up to June 2009 which were observed within 1,000 s (to ensure relatively high signal) from the burst onset and have a known redshift. In section 2 we briefly describe the data analysis and the way we evaluated the column density. In section 3 we discuss our results and in section 4 we draw our conclusions. ",
        "conclusions": "We considered a complete sample of 93 Swift long GRBs with prompt XRT pointing ($\\lsim 1,000$ s) and with a spectroscopic redshift. The distribution of the intrinsic X--ray absorption column density is consistent with a lognormal distribution with mean $\\log N_H(z)=21.9\\pm0.1$ ($90\\%$ confidence level) and standard deviation $\\sigma_{\\log N_H(z)}=0.5\\pm0.1$. This confirms that long GRBs occur in a dense environment, which is consistent with the expected distribution of absorptions for GRBs occurring randomly in molecular clouds (and not with long GRBs occurring randomly in any place of a galaxy similar to own Galaxy).  Looking at the distribution of X--ray column densities versus redshift we find a lack of non-absorbed GRBs at high redshift and a lack of heavily absorbed GRBs at low redshift. The former might be explained in terms of more compact and dense star formation regions in the young Universe (or to a sizable contribution from intervening systems). The latter might be interpreted as due to a change in the dust properties with redshift, with GRBs at redshift $z\\lsim 2-3$ having a higher dust to gas ratio for the same X--ray column density. This will naturally provide a lack of heavily (X--ray) absorbed GRBs at small redshifts, i.e. dark GRBs. We also investigated in more detail the relation between X--ray absorption and darkness of a GRB. We considered two different samples of long GRBs selected at different times (1,000 s and 11 hr). For none of the two we found a statistically significant relation between X--ray absorption and darkness, nor in a two-dimensional test including X--ray absorption and redshift. We note however that the sample of dark GRBs with known redshift is small and the weakness of our correlation might also depend on this.  Finally, we compare the hydrogen column densities derived in the optical and those from X--ray absorption. The two distributions are markedly different, with the optical column densities being lower than X--ray ones (and even more if correcting for metallicity effects). The most likely explanation is in terms of photoionization (depletion) of hydrogen in the circumburst material caused by the burst radiation field (Lazzati et al. 2001;  Watson et al. 2007)."
    },
    "0911/0911.0738_arXiv.txt": {
        "abstract": "We present prospects of two experiments using the Tokyo Axion Helioscope. One is a search for solar axions. In the past measurements, axion mass from 0 to $0.27 \\ \\mathrm{eV}$ and from 0.84 to $1.00 \\ \\mathrm{eV}$ have been scanned and no positive evidence was seen. We are now actively preparing a new phase of the experiment aiming at axion mass over $1 \\ \\mathrm{eV}$. The other is a search for hidden sector photons from the Sun. We have been designing and testing some additional equipments, which have to be installed on the helioscope to search for hidden photons with mass of over $10^{-3} \\ \\mathrm{eV}.$ ",
        "introduction": "The Sun could copiously emit weakly interacting particles, that could eventually be detected inside a sensitive detector at the Earth.  \\begin{wrapfigure}{r}{6cm} \\includegraphics[scale=0.09]{ohta_ryosuke.fig1.eps} \\caption{Detection schematics of solar axions and hidden photons from the Sun.} \\label{Fig:incident} \\end{wrapfigure}  Axion is one of such particles. The existence of axion is implied to solve the strong CP problem ~\\cite{Peccei:1977hh}. Axions are expected to be produced in the solar core through their coupling to photons. This process is called Primakoff process. The outgoing axion has average energy of about 4 keV~\\cite{Raffelt:2005mt}. Sikivie proposed an ingenious experiment to detect such axions~\\cite{Sikivie:1983ip}. A detection schematics of solar axion is shown in Fig.~\\ref{Fig:incident}. The detection device called axion helioscope is a system of a strong magnet and an X-ray detector, where the solar axions are transformed into X-ray photons through the inverse Primakoff process 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 axion becomes feasible.  Hidden sector photon is another kind of weakly interacting particles. The existence of hidden photons is predicted by several extensions of Standard Model. If light hidden photons exist, they could be produced through kinetic mixing with solar photons~\\cite{Redondo:2008aa, Gninenko:2008pz}. Therefore it is natural to consider the Sun as a source of low energy hidden photons. A detection schematics of hidden photon from the Sun is also shown in Fig.~\\ref{Fig:incident}. Unlike the case of the axion, no magnetic field is required to transform photons into hidden photons.  In this paper we report the current status of two experiments. One is the search for solar axions and the other is the search for hidden photons from the Sun. ",
        "conclusions": ""
    },
    "0911/0911.5437_arXiv.txt": {
        "abstract": "FRIDA (inFRared Imager and Dissector for the Adaptative optics system of the GTC) will be a NIR (1-2.5$\\mu$m) imager and Integral Field Unit spectrograph to operate with the Adaptative Optics system of the 10.4 m GTC telescope. FRIDA will offer broad and narrow band diffraction-limited imaging and integral field spectroscopy at low, intermediate and high spectral resolution. The Extragalactic Astrophysics and Astronomical Instrumentation group of the Universidad Complutense de Madrid (GUAIX) is developing the Data Reduction Pipeline for FRIDA. Specific tools for converting output, reduced datacubes to the standard Euro3D FITS format will be developed, in order to allow users to exploit existing VO applications for analysis. FRIDA is to be commissioned on the telescope in 2011. ",
        "introduction": "FRIDA\\footnote{More information on FRIDA at:\\\\ FRIDA Project URL:http://www.astroscu.unam.mx/ia\\_cu/proyectos/frida/ \\\\ FRIDA Instrument URL: http://www.iac.es/project/frida/\\\\ FRIDA GUAIX URL: http://guaix.fis.ucm.es/instrumentation} will be a second generation instrument for GTC, and the first proposed for the Adaptative Optics (AO) system of the telescope \\citep{Lopez07}. FRIDA is being designed and constructed by an international consortium, lead by the Instituto de Astronom\\'{\\i}a of the UNAM. It is an integral field spectrograph with imaging capabilities, optimized for the 1.1-2.4 $\\mu$m range, and it will provide broad and narrow band imaging and Integral Field Spectroscopy (IFS) capabilities with low, intermediate and high spectral resolutions, and with different spatial resolutions in the NIR (see Table \\ref{tab:table}). Its combination of high spectral and spatial resolutions will provide FRIDA with unique capabilities among other existing Integral Field Units (IFUs). FRIDA will be able to tackle a large number of astrophysical problems, from Solar system bodies to cosmological surveys. Its characteristics and its operating wavelength range will expose dusty environments, rich in stellar absorption and emission lines from many different species and gas phases, and sample kinematics and stellar or gaseous content of high redshift galaxies.  FRIDA will perform imaging at the diffraction limit of the telescope, thanks to the correction performed to the wavefront by the GTC AO module. In the IFS mode, a monolithic image slicer composed by a set of mirrors will break down a small portion of the AO-corrected field into 30 slices, that will play the role of 1D spectroscopy slits. The light of each slice is dispersed by the spectrograph onto the detector. With these data, a 3D datacube can be built up containing a pseudo-monocromatic image of the IFU field-of-view (FoV) in each plane (see Fig.~\\ref{fig:fig1} for a explanatory sketch of the datacube obtention). The FRIDA IFU design is based on FISICA IFU for FLAMINGOS-II \\citep{Eikenberry04}.  \\begin{figure} \\centering \\includegraphics[width=0.95\\linewidth]{eliche1.eps} \\caption{FRIDA IFS mode diagram. Credits: A.~L\\'{o}pez, S.~Cuevas.\\label{fig:fig1}} \\end{figure}   Figures \\ref{fig:fig2} and \\ref{fig:fig3} show the light path through the different FRIDA optical components for the imaging and IFS modes, respectively \\citep{Cuevas08}. In both cases, the light path starts at the cryostat window (on the left), coming directly from the GTC-AO module. It crosses the mask and filter wheels. The camera wheel selects the plate scale. In the IFS mode, the IFU dissects a small portion of the FoV into the 30 slices, and sends them to the grating carrousel for wavelength dispersion (Fig.~\\ref{fig:fig3}). The dispersed light of each slice is finally focussed on the detector (\"slitlet\"). In the imaging mode, the beam avoids entering the IFU and a flat mirror replaces the grating in the spectrograph (Fig.~\\ref{fig:fig2}).  \\begin{figure} \\centering \\includegraphics[width=\\linewidth]{eliche2.eps} \\caption{FRIDA optical layout in imaging mode. Credits: A.~L\\'{o}pez, S.~Cuevas.\\label{fig:fig2}} \\end{figure}  \\begin{figure} \\centering \\includegraphics[width=\\linewidth]{eliche3.eps} \\caption{FRIDA optical layout in IFS mode. Credits: A.~L\\'{o}pez, S.~Cuevas. \\label{fig:fig3}} \\end{figure}  \\vspace{-0.3cm} ",
        "conclusions": ""
    },
    "0911/0911.0176_arXiv.txt": {
        "abstract": "Previous submillimetre (submm) observations detected $\\rm 0.7\\,M_{\\odot}$ of cool dust emission around the Luminous Blue Variable (LBV) star $\\eta$ Carinae. These observations were hindered by the low declination of $\\eta$ Carinae and contamination from free-free emission orginating from the stellar wind.  Here, we present deep submm observations with \\laboca~at 870\\,$\\mu$m, taken shortly after a maximum in the 5.5-yr radio cycle.  We find a significant difference in the submm flux measured here compared with the previous measurement: the first indication of variability at submm wavelengths.  A comparison of the submm structures with ionised emission features suggests the 870\\,$\\mu$m is dominated by emission from the ionised wind and not thermal emission from dust. We estimate $0.4\\pm 0.1\\,\\rm M_{\\odot}$ of dust surrounding $\\eta$ Carinae.  The spatial distribution of the submm emission limits the mass loss to within the last thousand years, and is associated with mass ejected during the great eruptions and the pre-outburst LBV wind phase; we estimate that $\\eta$ Carinae has ejected $>\\rm 40\\,M_{\\odot}$ of gas within this timescale. ",
        "introduction": "$\\eta$ Carinae ($\\eta$ Car), one of the most luminous infra-red (IR) objects in our Galaxy ($\\rm > $$10^6L_{\\odot}$), is well known for dramatic outbursts over the last thousands of years, in which vast amounts of material is ejected outwards from the star.  The great eruption phase during the last 200 years are responsible for creating the dusty bipolar nebula known as the Homunculus, as seen in the {\\it Hubble Space Telescope (HST)} images.  The recent discovery of a fast blast wave from the 1843 eruption (Smith 2008) indicates that the mass loss in $\\eta$ Car is not simply due to a stellar wind, but could be more akin to a low-energy supernova remnant. Extensive IR imaging has revealed important clues about the mass loss history; the Infrared Space Observatory (ISO) discovered a massive cool dust component (Morris et al.\\ 1999; Smith et al.\\ 2003) with a total of $\\sim$$\\rm 0.1\\,M_{\\odot}$ of dust proposed to be present in the nebula, suggesting a mass loss rate of $0.5\\,\\rm M_{\\odot}$\\,yr$^{-1}$ during the last 200 years.  Submillimetre (submm) observations (Gomez et al.\\ 2006, hereafter G06) of $\\eta$ Car at 450 and 850\\,$\\mu$m with the Submillimetre Common User Bolometer Array (SCUBA) detected the presence of a massive component of dust, with $\\sim$$\\rm 0.7\\,M_{\\odot}$ needed to reproduce the IR-submm Spectral Energy Distribution (SED).  Their work suggested that up to four times more mass had been ejected from $\\eta$ Car during its recent, violent history than previously proposed.  The SCUBA data also indicated that the cool dust component was extended along the mid-plane, far beyond the inner optical and IR region; this was interpreted as mass which had been ejected on a much longer timescale than the well-known, smaller features such as the torus and Homunculus.  The previous analysis of the SCUBA data was hindered by two unknowns: whether the submm emission originated from free-free and not thermal emission from dust and whether the extended structure was reliable.  Firstly, the submm fluxes are likely to be contaminated by strong free-free emission from the extended stellar wind which is thought to vary with frequency as $\\nu^{0.6-1.3}$ (e.g. Cox et al. 1995; Brooks et al.\\ 2005, hereafter B05). This is further compounded since the radio and millimetre (mm) fluxes are highly variable, with $\\eta$ Car at 3\\,cm changing appearance from a point source in 1992 to an extended region with five times more flux in 1996 (Duncan \\& White 2003). The X-ray, radio and mm variability of $\\eta$ Car has been well documented (Pittard et al.\\ 1998; Abraham et al.\\ 2005; Damineli et al.\\ 2008), with the star undergoing periodic variability over a 5.5-yr period.  This variability has been associated with shocks from an extended disc surrounding the $\\eta$ Car binary system (e.g. Duncan \\& White 2003) and/or due to episodic mass loss which, in turn, decreases the number of ionising photons. Since there were no available mm observations to match the epoch of the original 1998 SCUBA data (taken during the minimum phase of the radio cycle), G06 were not able to estimate the contribution to the submm due to free-free emission.  Second, the SCUBA data were taken at very large airmass ($A\\sim 5.5$) due to the declination of $\\eta$ Car.  This meant that the beam shape at 450\\,$\\mu$m was distorted due to the deformation of the dish at such low elevations and the morphology of the source was difficult to ascertain from this dataset.  Accurately determining the distribution and mass of dust in the stellar wind has enormous consequences not only on determining the mass loss history of $\\eta$ Car, but also understanding the progenitors of `exotic' core-collapse supernovae (SNe; Pastorello et al.\\ 2007; Smith et al.\\ 2009).  In this Letter, we present deep submm \\laboca~observations of $\\eta$ Car, taken shortly after the maximum phase in the cycle.  In \\S\\ref{sec:obs} we discuss the data reduction and compare the 870\\,$\\mu$m emission to well-known features seen in X-ray, optical, H$\\alpha$, 8\\,$\\mu$m and 1.2\\,mm in \\S\\ref{sec:res}.  In \\S\\ref{sec:disc} we estimate the dust and gas mass for the stellar wind and our conclusions are presented in \\S\\ref{sec:conc}. ",
        "conclusions": "\\label{sec:conc}  We found that the submm extended emission surrounding $\\eta$ Car, originally identified in G06, is unresolved by the \\laboca~beam at 870\\,$\\mu$m, and is highly variable, following the well-documented cycle seen at mm and radio wavelengths.  Accounting for the flux difference from G06 and associating the 800-1000\\,$\\mu$m with free-free emission rather than thermal emission from grains, we revise the dust mass in the stellar wind to $0.4\\pm 0.1\\,\\rm M_{\\odot}$. This translates to a mass loss greater than $\\rm 40\\,M_{\\odot}$ over the past 1000\\, years.  Future observations with Herschel and ALMA are needed to constrain the submm and mm variability and multi-epoch data will allow us to separate out the dust and free-free components throughout the cycle. Higher angular resolution data are needed to separate out the multiple phases of mass loss, this will provide crucial information about the evolutionary path to massive-star SN explosions."
    },
    "0911/0911.4613_arXiv.txt": {
        "abstract": "Measurement of the 21cm hyperfine transition of neutral hydrogen provides a unique probe of the epoch of reionization and the Dark Ages. Three major mechanisms are believed to dominate the radiation process: emission from neutral hydrogen surrounding the ionized bubbles of first galaxies and/or quasars, emission from neutral hydrogen inside minihalos, and absorption of diffuse neutral hydrogen against the cosmic microwave background. In the present work, by simply combining the existing analytic models for the three mechanisms, we investigate the contribution of cross-correlation between these three components to the total 21cm angular power spectrum, in the sense that neutral hydrogen associated with different radiation processes traces the large-scale structures of underlying density perturbations. While the overall 21cm power spectrum remains almost unchanged with the inclusion of the cross-correlations, the cross-correlation may play a key role in the determination of the 21cm power spectrum during the transition of 21cm radiation from emission-dominated phase to absorption-dominated phase at redshift $z\\approx 20$. A significant suppression in the 21cm angular power spectrum during this transition is anticipated as the result of negative contribution of the cross-correlation between the absorption of diffuse neutral hydrogen and the emission components. Therefore, an accurate prediction of the cosmic 21cm power spectrum  should take the cross-correlation into account especially at the transition phase. ",
        "introduction": "The 21cm hyperfine transition of neutral hydrogen in the form of either emission or absorption against the cosmic microwave background (CMB) provides a unique probe of the history of our universe between the surface of last scattering ($z\\approx1000$) and the end of reionization ($z\\approx6$). While the expected brightness temperature of these 21cm signals is two orders of magnitude smaller than the CMB temperature and can be easily swamped by extremely strong foreground such as the Milky Way, many ambitious radio facilities (e.g. 21CMA\\footnote{See http://cosmo.bao.ac.cn}, LOFAR\\footnote{See http://www.lofar.org}, MWA\\footnote{See http://www.haystack.mit.edu/arrays/MWA}, PAPER\\footnote{See http://astro.berkeley.edu/$\\sim$dbacker/eor}, SKA\\footnote{See http://www.skatelescope.org}, etc) have been constructed and are being planned to search for the weak signals primarily from the epoch of reionization (EOR) at $6<z<20$. Indeed, the Dark Ages and the EOR may constitute the last frontier of observational cosmology.  There have been many efforts - both theoretical and numerical - over the past few years aimed at understanding various physical processes in the intergalactic medium (IGM) during the EOR [see \\citealt{furlanetto06} and \\citealt{morales} for a review]. In order to make the 21cm signals visible, the spin temperature of neutral hydrogen must differ from the temperature of the CMB, $T_{\\rm CMB}$, and an effective mechanism exists to couple the spin temperature to the gas temperature. Apart from some exotic scenarios such as dark matter particle decay, three mechanisms are believed to play dominant roles in the generation of observable 21cm signals from the Dark Ages to the EOR: scattering of UV photons of first stars primarily through the so-called Wouthuysen-Field effect (\\citealt{wouthuysen};  \\citealt{field}), heating of IGM by collision mainly within minihalos (\\citealt{madau}; \\citealt{iliev02}), and absorption of uncollapsed IGM against the CMB radiation prior to the formation of first cosmic structures (\\citealt{scott}; \\citealt{loeb}; Madau et al. 1997; \\citealt{lewis}). Among these three mechanisms, the foreground absorption of neutral hydrogen against the CMB dominates the 21cm background through the epoch of the Dark Ages down to redshift $z\\approx20$, while for $6<z<20$ the energy release of first luminous objects (stars and/or quasars) plays a leading role in generating the 21cm emission in excess of the CMB (e.g. \\citealt{scott}; \\citealt{barkana}). During the transition phase or the 'Grey Ages' between the end of the Dark Ages and the dawn of reionization at $z\\approx20$, a combined mechanism may come into effect in the redshifted 21cm background.  If we decompose the surface brightness temperature of 21cm background from the Dark Ages through the EOR into three components: emission from neutral hydrogen surrounding the ionized bubbles of first galaxies and/or quasars $T_{\\rm em}$, emission inside minihalos  $T_{\\rm mh}$, and absorption of IGM against the CMB  $T_{\\rm ab}$,  the 21cm power spectrum can be formally determined by three auto-correlation terms $\\langle T_{\\rm em}T_{\\rm em}\\rangle+ \\langle T_{\\rm mh}T_{\\rm mh}\\rangle+\\langle T_{\\rm ab}T_{\\rm ab}\\rangle$, and three cross-correlation terms $2\\langle T_{\\rm em}T_{\\rm mh}\\rangle+ 2\\langle T_{\\rm em}T_{\\rm ab}\\rangle+2\\langle T_{\\rm mh}T_{\\rm ab}\\rangle$. In particular, two of the cross-correlations relevant to the absorption have negative amplitudes as a result of $T_{\\rm S}<T_{\\rm CMB}$, where $T_{\\rm S}$ is the spin temperature of neutral hydrogen. Although quite a lot of work has been done in recent literature on the theoretical estimate of the three auto-correlation terms (e.g. \\citealt{iliev02}; \\citealt{zaldarriaga}; \\citealt{loeb}; \\citealt{furlanetto04a, furlanetto04b}; \\citealt{furlanetto06}; \\citealt{santos}; \\citealt{lewis}; \\citealt{mao}), it has remained unclear so far to what extent the three cross-correlations contribute to the total 21cm power spectrum, especially in the Grey Ages. The essentials of three cross-correlations arise from the fact that both collapsed and uncollapsed neutral hydrogen may trace underlying gravitational potentials of large-scale density perturbations, through which their 21cm emission/absorption signals are correlated at large-scales. \\cite{furlanetto06} have actually included the cross-correlation between $T_{\\rm em}$ and $T_{\\rm mh}$ in the study of the power spectrum of 21cm fluctuations, although their work focused  mainly on the observational distinction between signatures of minihalos and those of IGM, in which massive halos (galaxies) are ignored. In the present study we will concentrate on the evaluation of each of the cross-correlations in terms of the 21cm angular power spectrum for different frequencies or redshifts.   Note, however, that two of the cross-correlations, $\\langle T_{\\rm em}T_{\\rm ab}\\rangle$ and $\\langle T_{\\rm em}T_{\\rm mh}\\rangle$, should be safely treated as a mathematical elegance rather than a physically realistic situation through most of the period of the Dark Ages to the EOR. The latter requires the co-existence of 21cm emission from 'hot' IGM around first galaxies and 21cm absorption from 'cold' IGM, represented by $\\langle T_{\\rm em}T_{\\rm ab}\\rangle$, and the co-existence of 21cm emission from 'hot' IGM around first galaxies and 21cm emission from 'hot' IGM inside minihalos, represented by  $\\langle T_{\\rm em}T_{\\rm mh}\\rangle$, within the same regions. Unless an inhomogeneous heating scenario is considered (\\citealt{barkana}; \\citealt{pritchard07}; \\citealt{santos}; \\citealt{mesinger}), the two correlation terms $\\langle T_{\\rm em}T_{\\rm ab}\\rangle$ and $\\langle T_{\\rm em}T_{\\rm mh}\\rangle$ may become physically meaningless except for the very rapid transition phase or the Grey Ages. In this regard, the last cross-correlation between the 21cm emission from minihalos and the 21cm absorption from the diffuse IGM,  $\\langle T_{\\rm mh}T_{\\rm ab}\\rangle$, is probably the only physically plausible term that is worth investigating during and even before the Grey Ages.   The main difficulty of evaluating accurately the cross-correlations at present within the framework of analytic description of the reionization process is perhaps the absence of a unification scheme in which the three mechanisms of generating cosmic 21cm signals can be incorporated and treated at any redshifts through the Dark Ages to the EOR. The emission scenario associated with first galaxies which has been extensively discussed in recent literature applies only to lower redshifts at $6<z<20$, while the absorption model of uncollapsed IGM is restricted to higher redshifts beyond $z\\approx20$. During the Grey Ages when the effect of the cross-correlations between different mechanisms is expected to be visible in the cosmic 21cm power spectrum, our theoretical prediction relies on the extrapolation of the current emission and absorption models unless we make an attempt at exploring the unification scheme in which the distribution and evolution of both neutral hydrogen and its reionization can be correctly described through the entire period of the Dark Ages and the EOR including the transition phase.  It is likely that numerical simulation (e.g. \\citealt{zahn}; \\citealt{trac}; \\citealt{iliev08}; \\citealt{santos}) should be eventually used so that various physical processes can be incorporated in the estimate of the cross-correlations in the cosmic 21cm power spectrum.  Having realized these difficulties and limitations,  we still intend to take the three cross-correlations into account in the calculation of the cosmic 21cm power spectrum. By simple combining the existing analytic scenarios of generating the cosmic 21cm signals at the Dark Ages and the EOR, we address the question of what and how large observational effect one may expect to see if the three correlation terms are included. Our goal in the present paper is to call the attention of the 21cm cosmology community to the possible role of the cross-correlations in the 21cm power spectrum especially at the Grey Ages rather than the detailed modelling of the evolution of cosmic neutral hydrogen and reionization process. It is hoped that future numerical simulation can provide a more sophisticated treatment of the problem. ",
        "conclusions": "Using a simple analytic prescription of both 21cm emission and absorption of neutral hydrogen in the epoch of the Grey Ages, we have calculated the cross-correlation between three major mechanisms of cosmic 21 cm radiation, namely, the emission from neutral hydrogen surrounding the ionized bubbles of first galaxies, the emission of neutral hydrogen inside minihalos, and the absorption of diffuse neutral hydrogen against the CMB radiation. It has been shown that inclusion of the three cross correlations does not alter the overall angular power spectrum of cosmic 21cm background. However, during the transition of 21cm radiation field from emission-dominated phase to absorption-dominated phase, the cross-correlation may play a key role in the determination of the shape and magnitude of the 21cm power spectrum. This arises primarily from the negative contribution of cross-correlation between the absorption of diffuse IGM and the other two emission mechanisms, which leads to a significant power suppression on the fluctuation of 21cm radiation field during the transition phase. In our naive models for the above three mechanisms, especially the oversimplified scenario for the 21cm emission from the IGM associated with first galaxies, the transition occurs around redshift $z\\approx17-18$, which varies slightly with multipole $\\ell$.  The effect of cross-correlations on the 21cm angular power spectrum could be detected as a remarkable drop of power at a certain frequency range corresponding to the transition phase. The precise value of the transition frequency depends on the combined effects of various mechanisms of 21cm radiation in the epoch of the Grey Ages. An extension of the present work to the adoption of more sophisticated models for the 21cm radiation field at the Grey Ages may provide a more accurate prediction of the 21cm angular power spectrum.  Finally, we must caution that the effect of the cross-correlations on the total cosmic 21cm power spectrum could be considerably exaggerated as a result of our 'toy model', and some of our numerical results may even be artificial because of the 'artificial' split between the three major mechanisms of 21cm radiation. Numerical simulation should be applied to clarify the issue."
    },
    "0911/0911.4339_arXiv.txt": {
        "abstract": "{The evolutionary stage of the $\\delta$ Scuti star 44 Tau has been unclear. Recent asteroseismic studies have claimed models on the main sequence, as well as in the expansion phase of the post-main sequence evolution. However, these models could not reproduce all of the observed frequencies, the mode instability range, and the fundamental stellar parameters simultaneously. A recent photometric study has increased the number of detected independent modes in 44~Tau to 15, and a newly found gravity mode at 5.30~cd$^{-1}$ extends the observed frequency range.} { One of the possible evolutionary stages of 44~Tau has not yet been considered: the overall contraction phase after the main sequence. We computed asteroseismic models to examine whether models in this evolutionary stage provide a better fit of the observed frequency spectrum.} {We used Dziembowski's pulsation code to compute nonadiabatic frequencies of radial and nonradial modes. Observation of two radial modes and an avoided crossing of dipole modes put strong constraints on the models. A two-parametric overshooting routine is utilized to determine the efficiency of element mixing in the overshoot layer above the convective core.} {We find that pulsation models in the post-MS contraction phase successfully reproduce the observed frequency range, as well as the frequency values of all individual radial and nonradial modes. The theoretical frequencies of the mixed modes at 7.79~cd$^{-1}$ and 9.58~cd$^{-1}$ are in better agreement with the observations if efficient element mixing in a small overshoot layer is assumed.} {} ",
        "introduction": "\\begin{table*}[t!] \\caption{Observed frequencies and mode identifications.} \\label{table1} \\footnotesize \\begin{center} \\begin{tabular}{lcccccc} \\noalign{\\smallskip} \\hline\\hline \\noalign{\\smallskip} \\multicolumn{2}{c}{Frequency} & A$_y$  & A$_{<v^1>}$  & $(\\ell, m)$ & A$_v$/A$_y$  &  $\\phi_v$-$\\phi_y$   \\\\ & [cd$^{-1}$] & [millimag]  & [km $s^{-1}$] & & &  [$^\\circ$]    \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} $f_{1}$ &\t6.8980  &  27.32 & 2.22\t& (0,0) &  1.442 $\\pm$ 0.003 &  +2.89 $\\pm$ 0.13\t\\\\ $f_{2}$ &\t7.0060\t&   9.38 & 0.46 & (1,1)\t&  1.465 $\\pm$ 0.011 &  -1.57 $\\pm$ 0.43 \\\\ $f_{3}$ &\t9.1174\t&   8.39 & 0.44 & (1,1)\t&  1.463 $\\pm$ 0.011 &  -1.19 $\\pm$ 0.45 \\\\ $f_{4}$ &\t11.5196\t&   9.93 & 0.73 & (1,0)\t&  1.421 $\\pm$ 0.009 &  -2.23 $\\pm$ 0.36  \\\\ $f_{5}$ &\t8.9606\t&   9.31 & 1.03 & (0,0)\t&  1.463 $\\pm$ 0.012 &  +1.92 $\\pm$ 0.45 \\\\ $f_{6}$ &\t9.5611\t&   8.50 & 0.30 & (1,-)\t&  1.451 $\\pm$ 0.012 &  -0.94 $\\pm$ 0.51 \\\\ $f_{7}$ &\t7.3031\t&   4.74 & 0.70 & (2,0)\t&  1.449 $\\pm$ 0.025 &  -6.91 $\\pm$ 0.90 \\\\ $f_{8}$ &\t6.7955\t&   2.58 & 0.28 & (2,0)\t&  1.342 $\\pm$ 0.041 &  -7.26 $\\pm$ 2.01 \\\\ $f_{9}$ &\t9.5828\t&   1.84 & 0.12 & (2,-)\t&  1.517 $\\pm$ 0.054 &  -7.82 $\\pm$ 2.02 \\\\ $f_{10}$ &\t6.3390\t&   1.69 & 0.21 & (2,-)\t&  1.482 $\\pm$ 0.064 &  -6.55 $\\pm$ 2.24 \\\\ $f_{11}$ &\t8.6391\t&   1.78 & 0.32 & (-,0)\t&  1.383 $\\pm$ 0.054 &  -4.87 $\\pm$ 2.19 \\\\ $f_{12}$ &\t11.2947\t&   1.03 & 0.11 & \t--\t&       --                    &   --               \\\\ $f_{13}$ &\t12.6915\t&   0.28 & --   &   --  &       --                    &  --                \\\\ $f_{14}$ &\t5.3047\t&   0.70 & --   &\t--  &      --                     &  --                \\\\ $f_{15}$ &\t7.7897\t&   0.28 & --   &\t--  &      --                     &  --                \\\\ \\noalign{\\smallskip} \\hline \\end{tabular} \\end{center} \\end{table*}   The $\\delta$ Scuti stars are a class of pulsators exhibiting low-order acoustic and gravity modes driven by the opacity mechanism acting in the HeII ionization zone. In evolved $\\delta$ Scuti stars mixed modes, i.e., modes with acoustic behaviour in the envelope and gravity behaviour in the interior, are common. Since such modes probe different layers inside a star than purely acoustic modes, the $\\delta$ Scuti stars are excellent targets for asteroseismology.  44~Tau is a $\\delta$~Scuti pulsator of spectral type F2 IV. Its exceptionally slow intrinsic rotation of $V_{\\rm rot}$ = 3 $\\pm$ 2 km s$^{-1}$ (Zima et al. 2007) is much lower than the average value for $\\delta$ Scuti stars. For this group of pulsators, projected rotational velocities, $v \\sin i$, around 100 km s$^{-1}$ and higher are not unusual. This makes 44~Tau an interesting and relatively simple target for asteroseismology, because the effects of rotation are weak and do not complicate mode identification and stellar modelling.  In recent years several ground-based observing campaigns of the Delta Scuti Network (DSN) were organized to retrieve photometric and spectroscopic data of 44~Tau. The first comprehensive frequency analysis of its photometric light variations was published by Antoci et al. (\\cite{antoci07}). This paper also summarizes all previous observations of the star. With two additional seasons of extensive Str\\\"omgren $vy$ photometry, Breger \\& Lenz (\\cite{breger08}) could increase the number of detected frequencies to 49, of which 15 are independent.  Garrido et al. (\\cite{garrido07}) have determined the spherical degree, $\\ell$, for the 10 dominant modes by means of photometric amplitude ratios and phase differences. Their results were later confirmed by Lenz et al. (2008). In 2004 a spectroscopic campaign was organized simultaneously to the photometric observing run. The results of the line-profile variation analyses and the determination of photospheric element abundances are discussed in detail by Zima et al. (2007). For 8 modes, the azimuthal order, $m$, could be determined. Moreover, no indications of a global magnetic field and no anomalous element abundances were found.  The spectroscopically determined $\\log g$ value of 3.6 $\\pm$ 0.1 does not allow for unambiguous determination of the evolutionary status of 44~Tau. In recent studies, both main sequence (hereafter MS) and post-MS models in the expansion phase have been considered (Garrido et al. 2007, Lenz et al. 2008). The results by Lenz et al. (\\cite{lenz08}) showed that these models fail to reproduce \\emph{all} observable parameters satisfactorily (individual frequencies, observed frequency range, fundamental stellar parameters).  The effective temperature and luminosity of MS models is significantly lower than the values derived from photometry and Hipparcos data. Moreover, the positions of predicted $\\ell$=2 modes do not match the observed frequencies (see Fig. 8 in Lenz et al. 2008).  Evolved models in the hydrogen-shell burning stage reproduce the fundamental stellar parameters well, but theory predicts many more frequencies than actually observed. While partial mode trapping may explain the position of the $\\ell$=1 modes, this mechanism should not be very effective for the observed $\\ell$=2 modes. Moreover, the frequency range of unstable modes is shifted to higher frequencies with respect to MS models. Therefore, the mode at 6.34~cd$^{-1}$ and the newly found mode at 5.30~cd$^{-1}$ are predicted to be stable by models in the post-MS expansion phase.  We reexamined asteroseismic models of 44~Tau based on the latest frequency solution of Breger \\& Lenz (\\cite{breger08}) and investigated a hitherto neglected possibility: a model in the overall contraction phase after the main sequence.  During the main sequence evolution, a $\\delta$ Scuti star transforms hydrogen into helium in its convective core. The amount of hydrogen in the core gets very low towards the end of this phase. The star increases its central temperature by an overall contraction to keep the energy production from hydrogen burning efficient. At the end of this overall contraction stage, hydrogen in the core is fully depleted and the convective core disappears\\footnote{Of course, convective overshooting also disappears, but we keep the term `model with overshooting' or `the overshooting case' for later evolutionary stages to distinguish this case from models constructed without convective overshooting.}. A nuclear hydrogen-burning shell is established outside the hydrogen-exhausted core. The stellar envelope expands again, which leads to a decrease of the effective temperature.  Finding a $\\delta$ Scuti star in the contraction phase after the TAMS is not unlikely. If we consider the same range of effective temperatures, the evolution during the overall contraction phase is approximately 10 times faster compared to the MS evolution. In the post-MS expansion phase, the evolution is approximately 16 times faster than on the MS.  The asteroseismic inferences from post-MS contraction models are presented in this paper.  The paper is organized as follows. In Sect. 2 we briefly summarize the observed parameters of 44~Tau. In Sect. 3 we discuss the computation of our seismic models. The diagnostic power of mixed modes is examined in Sect. 4. Finally, in Sect. 5 we present predictions of post-MS contraction models. ",
        "conclusions": "We constructed seismic models of the $\\delta$ Scuti star 44~Tau in the overall contraction phase after the main sequence. The two observed radial modes and an avoided crossing of a pair of dipole modes put strong constraints on pulsation models. Unlike models on the main sequence or in the post-MS expansion phase, these models successfully reproduce all 15 independent frequencies observed in 44~Tau and predict all detected modes unstable. This makes 44~Tau an example for determining the evolutionary stage by means of asteroseismology.  The observed frequency spectrum (obtained from ground-based data) can be solely explained with modes with spherical degrees of 0, 1 and 2. This result is what we expect for the given accuracy of the data.  The frequency spectrum of 44~Tau consists mainly of modes of mixed character, i.e., modes with acoustic behaviour in the envelope and gravity behaviour in the interior. These modes allow examination of the efficiency of partial element mixing in the overshoot layer around the stellar core. We find that the theoretical frequencies of several modes (e.g. the dipole $g_4$ mode at 7.79~cd$^{-1}$ and the quadrupole mode $g_5$ at 9.58~cd$^{-1}$) are in better agreement with efficient mixing in a thin overshoot layer than with less efficient mixing in an extended overshoot layer.  The main source of uncertainty in the asteroseismic modelling of 44~Tau are the stellar opacities. Determination of the size of the overshooting region above the convective core, therefore, depends on the source of opacity data and, to a smaller extent, on the choice of the element mixture. Pulsation models in the post-MS contraction phase constructed with standard OP data are more than two standard deviations outside the photometric error box in the HR diagram. Following a suggestion by Montalb{\\'a}n \\& Miglio (\\cite{mont08}), we computed a model with artificially enhanced OP opacities around the temperature region $\\log T$~=~6.0. Such a model indeed agrees better with the fundamental parameters of 44~Tau derived from photometry and Hipparcos data.  The pulsation models presented in this paper not only provide an excellent fit of all observed modes but also allow several predictions, which may be confirmed with additional observational data. A few of the theoretically predicted modes are not observed. The question arises whether the amplitudes of these modes are below the noise threshold of the current data or whether they are even excited at all. Moreover, we predict the spherical degree of already observed modes for which our mode identification techniques did not succeed because the uncertainties in amplitudes and phases are too large. The excellent fit of the individual frequencies is very promising, and we are confident that additional data will lead to even more asteroseismic inferences about this star."
    },
    "0911/0911.4755_arXiv.txt": {
        "abstract": "{We report the finding of a scaling relation among the cosmic-web anisotropy parameter $A$, the linear density rms fluctuation $\\sigma (r)$ and the linear growth factor $D(z)$. Using the tidal field derived from the Millennium Simulation on $512^{3}$ grids at $z=0,\\ 2,\\ 5$ and $127$, we calculate the largest eigenvalues $\\lambda$ of the local tidal tensor at each grid resolution and measure its distance-averaged two-point correlation function, $\\xi_{\\lambda}$, as a function of the cosines of polar angles $\\cos\\theta$ in the local principal axis frame.  We show that $\\xi_{\\lambda}$ is quite anisotropic, increasing toward the directions of minimal matter compression, and that the anisotropy of $\\xi_{\\lambda}$ increases as the redshift $z$ decreases and as the upper distance cutoff $r_{c}$ decreases.  Fitting the numerical results to an analytic fitting model $\\xi_{\\lambda}(\\cos\\theta)\\propto (1+A\\cos^{n}\\theta)^{-1}$, it is found that the best fit value of $A$, dubbed the {\\it cosmic-web anisotropy parameter}, varies systematically with $\\sigma(r_{c})$ and $D(z)$, allowing  us to determine the simple empiral scaling relation $A(r_{c},z)=0.8\\, D^{0.76}(z)\\,\\sigma (r_{c})$.} ",
        "introduction": "As confirmed by recent N-body simulations\\cite{millennium05,simulation08}, the large-scale spatial distribution of cold dark matter exhibits an anisotropic web-like pattern, which is often dubbed the {\\it cosmic web}.  The cosmic web of the dark matter distribution found in N-body simulations is also consistent with the observed large-scale filamentary distribution of galaxies in the real universe \\cite{2df,sdss}. According to the standard model of cosmic structure formation \\cite{bond-etal96}, the cosmic web originates in primordial density perturbations that are sharpened by gravitational tidal fields. The web becomes more anisotropic as the tidal fields become stronger due to the nonlinear processes during the gravitational evolution.  Various statistical tools have so far been proposed to describe the geometric properties of the cosmic web. For example, the conventional $N$-point statistics has been used to calculate the anisotropic spatial distribution of dark halos \\cite{bon-mye96}. The Minkowski functionals have been employed to determine the morphological properties of the large-scale structures embedded in the cosmic web \\cite{sch-etal99}. The $N$-dimensional skeleton approach has been found to be efficient in tracing the evolution of the cosmic web \\cite{sou-etal08,sou-etal09}.  The Multiscale Morphology Filter method has been applied to the large-scale galaxy distribution to identify the anisotropic structures of the cosmic web \\cite{ara-etal07}. A method utilizing the concept of the {\\it Local Dimension} has been suggested for the local quantification of the shapes of the galaxy neighborhood in the cosmic web \\cite{SB09}.  The ellipticity-ellipticity correlations of dark halos have been used to quantify the large-scale anisotropy of the cosmic web \\cite{lee-etal08}. Finally, tessellation techniques have been suggested to trace the evolution of the geometric structures in the universe \\cite{sha09,sha-etal09}.  Very recently, it has been pointed out by Lee, Hahn and Porciani \\cite{lee-etal09}(LHP09, hereafter) that, since the cosmic web is produced by the anisotropic compression of matter along the principal axes of the large-scale tidal fields, its anisotropic nature may be best quantified in the system of the principal axes of the tidal fields. This led them to suggest the anisotropic two-point correlations of the nonlinear traceless tidal field expressed in the principal axis frame, $\\xi_{\\lambda}({\\bf r})$, as a new statistical tool for the description of the cosmic web phenomenon.  Analyzing numerical data from high-resolution N-body simulations, they have determined $\\xi_{\\lambda}({\\bf r})$ and found that it has a much larger correlation length ($\\sim 20\\, h^{-1}$Mpc) and a higher degree of anisotropy than the density field itself. Integrating $\\xi_{\\lambda}({\\bf r})$ over distance $r$ from $0$ to a certain cut-off scale $r_{c}$, and expressing it as a function of the angle between the major principal axes and the separation vectors, they noted that the nonlinear traceless tidal field has a much higher degree of anisotropy than the nonlinear density field. Interestingly, the results of LHC09 imply that the correlations of the nonlinear traceless tidal field may have a link to the initial conditions.  Yet, the analysis of LHP09 was restricted to the present epoch and their results were obtained by setting the distance cutoff scale $r_{c}$ to the correlation length of the nonlinear tidal field, $r_{c}=20\\, h^{-1}$Mpc.  In this paper, our goal is to explore how the two-point correlations of the nonlinear traceless tidal field expressed in the principal axis system vary with redshift and distance scale by analyzing numerical data from high-resolution cosmological simulations. Throughout this paper, we assume a flat $\\Lambda$CDM cosmology. ",
        "conclusions": "To quantify the web-like pattern in the large-scale matter distribution, we have introduced the new concept of the cosmic web anisotropy parameter, which measures the degree of the anisotropy of the two-point correlations of the largest eigenvalues of the traceless tidal fields in the principal-axis frame. By analyzing the numerical data from the Millennium Simulation at different redshifts, we have empirically found a simple scaling relation among the cosmic web anisotropy parameter, the linear density rms fluctuations and the linear growth factor. Our results have allowed us to quantify how the anisotropy of the traceless tidal fields increases as the growth factor and the linear density rms fluctuations increase.  This new scaling relation thus provides a measure for the growth of anisotropy in the evolving cosmic web of the Universe.  The simple functional forms we found for the correlation and the scaling relation (eqs. [\\ref{eqn:xi_fit}] and [\\ref{eqn:ar_scale}]) suggest that there might exist a simple explanation for these functional forms in linear perturbation theory. Yet, we have not been able to develop such a theory thus far, as the analytic treatment from first principles turns out to be extremely hard since we deal here with the {\\it nonlinear} two-point correlations measured in the {\\it principal} axis frame.  A crucial implication of our result is that the anisotropy of the cosmic web may be produced by three competing effects. The trace part of the tidal field and the cosmic expansion tend to make the matter distribution more isotropic, whereas the traceless part of the tidal field stretches the matter distribution and induces large-scale anisotropy. The competition among these three effects imprints the web-like pattern in the large-scale matter distribution, just as the competition between gravity and radiation pressure has left imprints in the form of acoustic oscillations in the temperature map of the cosmic microwave background radiation \\cite{HS95}.  Since the linear growth factor $D(z)$ and the linear density rms fluctuations $\\sigma(r_{c})$ are functions of the primary cosmological parameters such as the dark energy equation of state parameter $w$ and the density parameter $\\Omega_{m}$, our results also hint that the cosmic web anisotropy parameter may be a new probe of cosmology. Given the scaling relation (eq.[\\ref{eqn:ar_scale}]), we expect that if the cosmic web anisotropy parameter $A$ is measured for two different distance cut-off scales $r_{c}$ at the same redshift, it may be possible to put a constraint on the density parameter $\\Omega_{m}$ by taking the ratio between the two values, removing the parameter degeneracy with the amplitude of the linear power spectrum $\\sigma_{8}$. Similarly, by measuring the cosmic web anisotropy parameter at two different redshifts and taking the ratio between the two values, a constraint on the dark energy equation of state parameter $w$ results.  The success of using the cosmic web anisotropy parameter as a new cosmological probe, however, is contingent upon a couple of future tests. First of all, it needs to be examined whether or not the same scaling relation also holds for different cosmologies since we have derived it assuming a $\\Lambda$CDM cosmology with a specific set of cosmological parameters. Secondly, what is readily measurable in practice is not the tidal field of the dark matter distribution but that of the galaxy distribution.  It will hence be necessary to investigate how the bias between light and matter affects the scaling relation. We investigate this question in forthcoming work, and hope to report the results elsewhere in the near future."
    },
    "0911/0911.2565_arXiv.txt": {
        "abstract": "{} {The data of sunspot numbers, sunspot areas and solar flare index during cycle 23 are analyzed to investigate the intermediate-term periodicities.} {Power spectral analysis has been performed  separately for the data of the whole disk, northern and southern hemispheres of the Sun.} {Several significant midrange periodicities ($\\sim$175, 133, 113, 104, 84, 63 days) are detected in sunspot activity. Most of the periodicities in sunspot numbers generally agree with those of sunspot areas during the solar cycle 23. The study reveals that the periodic variations in the northern and southern hemispheres of the Sun show a kind of asymmetrical behavior. Periodicities of $\\sim$175 days and $\\sim$133 days are highly significant in the sunspot data of northern hemisphere showing consistency with the findings of Lean (1990) during solar cycles 12-21. On the other hand, southern hemisphere shows a strong periodicity of about 85 days in terms of sunspot activity. The analysis of solar flare index data of the same time interval does not show any significant peak. The different periodic behavior of sunspot and flare activity can be understood in the light of hypothesis proposed by Ballester et al. (2002), which suggests that during cycle 23, the periodic emergence of magnetic flux partly takes place away from developed sunspot groups and hence may not necessarily increase the magnetic complexity of sunspot groups that leads to the generation of flares.} {} ",
        "introduction": "Sunspots are the fundamental indicators of solar activity. They exhibit a long term periodicity of 11 years, the so called solar cycle, which is known since a long time. In short term variations, the 27-day periodicity is the most prominent, which is attributed to the rotation of the Sun. In solar cycle 21, a periodicity of 154 days was discovered by Rieger et al. (1984) in the occurrence of high energy flares. Since then a number of studies have been made to search for the existence of intermediate-term or midrange periodicities between 27 days and 11 years in the various features of the active Sun (Bai, 2003).   The presence of near 155 days periodicity in flare and flare related data during cycle 21 was established by various authors (Ichimoto et al. 1985; Bogart \\& Bai 1985; Bai \\& Sturrock 1987; \\\"{O}zg\\\"{u}c and Ata\\c{c} 1989; Bai \\& Cliver 1990; Bai 2003; Joshi \\& Joshi 2005). Midrange periodicities in flare data were also investigated in solar cycles 19 and 20. Bai (1987) found a 51 days periodicity during cycle 19 in the occurrence rate of major flares. Analysis of solar flare data of cycle 20 revealed the periodicities of 78 days (Bogart \\& Bai 1985), 84 days (Bai \\& Sturrock 1991) and 127 days (Bai \\& Sturrock 1991; Kile \\& Cliver 1991). Significant peaks at 74, 77 and 83 days were reported in the flare data of solar cycle 22 (Bai 1992; \\\"{O}zg\\\"{u}c and Ata\\c{c} 1994; Joshi \\& Joshi 2005).  In several studies the sunspot data (numbers as well as areas) have been analyzed to investigate the midrange periodicities during different solar cycles (Lean \\& Brueckner 1989; Lean 1990; Carbonell \\& Ballester 1992; Oliver et al. 1992; Ballester et al. 1999; Krivova \\& Solanki 2002; Richardson \\& Cane 2005). Lean \\& Brueckner (1989) detected the near 155 days periodicity in sunspot blocking function, the 10.7 cm radio flux and sunspot number during cycles 19, 20 and 21 but this period could not be found in plage index. Lean (1990) analyzed the data of sunspot areas during cycles 12-21 and found that the periodicities in the range of 130 to 185 days occurred intermittently for the interval of 1 to 3 years during the epochs of maximum activity. Oliver et al. (1998) showed a time-frequency coincidence between the occurrence of the periodicity in both sunspot areas and high energy flares during cycle 21 which suggests a causal relationship between the two phenomena.  The aim of this paper is to detect midrange periodicities in the data of sunspot numbers, areas and solar flare index during solar cycle 23 and to investigate the causal relationship in the periodic behavior of these solar activity phenomena. The periodic variations have also been examined separately for the northern and southern hemispheres of the Sun. ",
        "conclusions": "In this paper, the daily data of sunspot numbers, areas and solar flare index are analyzed to search for midrange periodicities. The periodogram analysis of daily sunspot numbers and areas (Figures 2 and 3) reveals the existence of several significant periodicities. On the other hand, power spectra of daily solar flare index of the same interval (Figure 4) did not show any prominent peak with FAP value below 50 \\%. In table 1, the significant periodicities detected in sunspot numbers and areas are summarized.  The sunspot numbers of the whole disk (top panel of Figure 2) shows three important peaks at 84.5 days, 133.0 days and 113.5 days. A periodicity of 134.5 days is appearing in the sunspot numbers of northern hemisphere (Fig. 2 middle panel), which is very close to 133.0 days period detected in the data of whole disk. However, the northern hemisphere's sunspot numbers exhibit the highest peak at 175.4 days. Two other important periodicities at 104.3 days and 63 days are found in northern hemisphere. The sunspot numbers of the southern part of the solar disk gave rise to a highly significant peak at 84.5 days (see bottom panel of Figure 2), exactly same peak is detected in whole disk data. Similarly another peak in south at 113.5 days is also appearing in the whole disk data.  The power spectra of sunspot areas of whole disk, northern and southern hemisphere of the Sun are shown in Figure 3 (from top to bottom panel). In the whole disk data a significant peak at 112.4 days is detected which is very close to 113.5 days period present in sunspot numbers of the whole disk as well as southern hemisphere. In can be noticed that the same peak is also appearing in sunspot areas of southern hemisphere, but its significant level is low. Similarly 85.1 days periodicity present in southern hemisphere sunspot areas is very near to 84.5 days peak in the power spectra of southern hemisphere and full disk values of sunspot numbers. Both the periodicities ($\\sim$85 and 113 days) are absent is northern hemisphere data of sunspot numbers and areas. A periodicity of 175.5 days found in sunspot areas of northern hemisphere is exactly appearing in sunspot numbers of the same hemisphere and it absent in total disk as well as southern hemisphere data of both the indices. We can conclude that the near 113 and 85 days periodicities in both the indicators of sunspot activity are basically due to the activity in southern part of the solar disk, while near 175 days periodicity comes as a result of sunspot activity in northern hemisphere of the Sun. Here it is interesting to mention that a periodicity of 86 days was detected in sunspot areas of cycle 22 (Oliver \\& Ballester, 1995), which is close to 85 days periodicity appeared in  cycle 23.  Highly significant peaks at $\\sim$175 and $\\sim$133 days detected in the present analysis in sunspot numbers are consistent with the studies of Lean (1990) and Richardson \\& Cane (2005). The Periodogram analysis of sunspot areas further confirms the presence of near 175 days periodicity. Lean (1990) made an extensive study of periodic patterns in sunspot areas during cycles 12-21, and found that the ``Rieger-type\" periodicities occur intermittently in each cycle during the epochs of maximum activity, however its actual period varies from 130 to 185 days. Richardson \\& Cane (2005) applied wavelet analysis to sunspot number data of cycle 23 and reported the appearance of periodic signals of enhanced power in the range of $\\sim$120-200 days during the maximum phase of the cycle (late 1998 to 2002).  The cause for the origin of Rieger-type periodicities is suggested by several authors. Lean \\& Brueckner (1989) linked this periodicity with the magnetism of sunspots and suggested that it may be a property of emerging flux rather than the total amount of flux present on the solar disk. Carbonell \\& Ballester (1990) suggested its association with the periodicity in emergence of magnetic flux through the photosphere. Ballester et al. (2002, 2004) analyzed the photospheric magnetic flux during cycles 21-23 and found the evidence of near 160 days periodicity in cycles 21 and 23. Based on their analysis, they put forward a hypothesis that the periodicity can appear in two different ways. Periodic emergence of magnetic flux within already formed sunspot groups enhances the magnetic complexity of the active region which gives rise to similar periodic variations in solar flares. On the other hand, the periodic emergence of flux away from developed sunspot groups will not be reflected in the periodic behavior of solar flares.  The comparison of periodogram analysis of three kinds of data sets during February 1998 to December 2003, presented in this paper, suggests that the results of sunspot numbers generally agree with the results of sunspot areas. However the periodic pattern of occurrence of solar flares displayed a different behavior and we could not detected any significant periodicity in solar flare index data of same time interval. Recently Ata\\c{c} \\& \\\"{O}zg\\\"{u}\\c{c} (2006) analyzed the solar flare index data from January 1999 to December 2002 and detected a $\\sim$126 days periodicity. Their analysis also did not show any periodicity in the range of 130 to 185 days which was found in our study of sunspot data. Similar results were obtained with the analysis of soft X-ray flare index (Joshi \\& Joshi 2005). It therefore seems that periodic emergence of magnetic flux away from developed active region, during the maximum phase of current solar cycle, caused the similar periodic variations in sunspot numbers as well as areas, but did not enhance the magnetic complexity of the active regions in the same proportion. This may be a reason why the similar periodic variations could not be detected in solar flare index data."
    },
    "0911/0911.5242.txt": {
        "abstract": "I introduce an algorithm for estimating parameters from multidimensional data based on forward modelling.  It performs an iterative local search to effectively achieve a nonlinear interpolation of a template grid.  In contrast to many machine learning approaches it avoids fitting an inverse model and the problems associated with this.  The algorithm makes explicit use of the sensitivities of the data to the parameters, with the goal of better treating parameters which only have a weak impact on the data. The forward modelling approach provides uncertainty (full covariance) estimates in the predicted parameters as well as a goodness-of-fit for observations, thus providing a simple means of identifying outliers. I demonstrate the algorithm, \\ilium, with the estimation of stellar astrophysical parameters (APs) from simulations of the low resolution spectrophotometry to be obtained by Gaia. The AP accuracy is competitive with that obtained by a support vector machine. For zero extinction stars covering a wide range of metallicity, surface gravity and temperature, \\ilium\\ can estimate \\teff\\ to an accuracy of 0.3\\% at G=15 and to 4\\% for (lower signal-to-noise ratio) spectra at G=20, the Gaia limiting magnitude (mean absolute errors are quoted). \\feh\\ and \\logg\\ can be estimated to accuracies of 0.1--0.4\\,dex for stars with G\\,$\\leq18.5$, depending on the magnitude and what priors we can place on the APs.  If extinction varies a priori over a wide range (0--10\\,mag) -- which will be the case with Gaia because it is an all sky survey -- then \\logg\\ and \\feh\\ can still be estimated to 0.3 and 0.5\\,dex respectively at G=15, but much poorer at G=18.5. \\teff\\ and \\av\\ can be estimated quite accurately (3--4\\% and 0.1--0.2\\,mag respectively at G=15), but there is a strong and ubiquitous degeneracy in these parameters which limits our ability to estimate either accurately at faint magnitudes. Using the forward model we can map these degeneracies (in advance), and thus provide a complete probability distribution over solutions.  Additional information from the Gaia parallaxes, other surveys or suitable priors should help reduce these degeneracies. ",
        "introduction": "} %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%  Inferring parameters from multidimensional data is a common task in astronomy, whether this be inference of cosmological parameters from CMB experiments, photometric-redshifts of galaxies or physical properites from stellar spectra.  Important questions about the structure and evolution of stars and stellar populations require knowledge of abundances and ages, which must be obtained spectroscopically via the stellar atmospheric parameters effective temperature (\\teff), surface gravity (\\logg) and iron-peak metallicity (\\feh).  Numerous publications have presented methods for estimating such astrophysical parameters (APs) from spectra. The methods can be divided into two broad categories. The first, based on high resolution and high signal-to-noise ratio (SNR) spectra, makes use of specific line indices selected to be sensitive predominantly to the phenomena of interest.  Examples include the detection of BHB stars via Calcium and Balmer lines (e.g.\\ Brown~\\citealp{brown03}) and spectral type classification of M, L and T dwarfs via molecular band indices (e.g.\\ Hawley et al.~\\citealp{hawley02}). In each case these methods use a specific, identified phenomenon and soare generally limited to a narrow part of the AP space where the values of the other APs are reasonably well known. While simple and useful, these methods do not use all available information nor can they normally be used with low resolution data, because the effects of individual APs cannot then be separated.  The second category of methods is global pattern recognition, which try to use the full set of available data. These are generally used when we want to estimate one or more APs over a wide range of parameter space.  In such cases there is rarely a simple relation between a single feature and the physical quantity of interest. (An example is determination of surface gravity, where the relevant lines to use change with temperature.) We must therefore infer which features are relevant to which APs in which parts of the AP space. As this is generally a nonlinear, multidimensional problem, the standard approach is to use a machine learning algorithm to learn the mapping from the data space to the AP space based on labelled template spectra (spectra with known APs).  Various models have been used in astronomy for a variety of problems, including (to name just a few): neural networks for stellar parameter estimation (e.g.\\ Re Fiorentin et al.~\\citealp{rf07}) or photometric redshift estimation (e.g.\\ Firth et al.~\\citealp{firth03}); support vector machines for quasar classification (e.g.\\ Gao et al.~\\citealp{gao08}, Bailer-Jones et al.~\\citealp{cbj08}) or galaxy morphology classification (e.g.\\ Huertas-Company et al.~\\citealp{hc08}); classification trees for identifying cosmic rays in images (e.g.\\ Salzberg et al.~\\citealp{salzberg95}); linear basis function projection methods (e.g.\\ the method of Recio-Blanco et al.~\\citealp{recioblanco06} for spectral parameter estimation, which constructs the basis functions from model spectra).  More examples of machine learning methods and their use in astronomy can be found in the volume edited by Bailer-Jones~\\citep{cbj08b}.  Note that the first category of methods, line indices, is really just a special case of the second in which drastic feature selection has taken place to enable use of low (one or two) dimensional models. In both cases we must learn some relationship ${\\rm AP} = g'({\\rm Data})$ from a  model or based on some labelled templates. But this is an {\\em inverse} relation: more than one set of APs may fit a given set of data (e.g.\\ a low extinction cool star or high extinction hot star could produce the same colours). Despite this non-uniqueness we nonetheless try and fit a unique model. This causes fitting problems which become more severe the lower the quality of the data (the lower the number of independent measures) and the larger the number of APs we want to estimate from it, and could lead to poor AP estimates or biases. In contrast, the forward mapping (or {\\em generative model}), ${\\rm Data} = g({\\rm AP})$ is unique, because this is a causal, physical model (e.g.\\ a stellar atmosphere and radiative transfer model of a spectrum). %Another drawback of most inverse mapping methods is that they make a global fit to the data, normally by minimizing a global fitting error. %Consequently, the model predictions are biased to some degree by the distribution over APs in the training data (a version of the class %imbalance problem; see e.g.\\ Japkowicz \\& Stephen~\\citealp{jap02}, Bailer-Jones et al.~\\citealp{cbj08}). A further issue is that the model must learn the sensitivity of each input to each AP (and how noise affects this). Yet this information we in principle have already from the gradients of the generative model.  A pattern recognition method which tries to overcome some of these problems is $k$-nearest neighbours, which has also been applied to many problems in astronomy (e.g.\\ Katz et al.~\\citealp{katz98}, Ball et al.~\\citealp{ball08}; plus extensions thereof such as kernel density estimation, e.g.\\ Richards et al.~\\citealp{richards04}, or the method of Shkedy et al.~\\citealp{shkedy07}, which converts the distances into likelihoods and uses priors to create a full probabilistic solution.) In many ways this is the most natural way to solve the problem: we create a grid of labelled templates and find which are closest to our observation (perhaps smoothing over several neighbours). On the assumption that the APs vary smoothly with the data between the grid points, this may provide a good estimate. But for this to be accurate (and not too biased), the grid must be sufficiently dense that multiple grid points lie within the error ellipse of the observation (the covariance of these neighbours then provides a measure of the uncertainty in the estimated APs). If the SNR is high, the grid must therefore be very dense. Moreover, as the number of APs increases, the required grid density grows exponentially with it: For a stellar parametrization problem with 5 APs we might need an average of 100 samples per AP, resulting in $100^{5}=10^{10}$ templates. There is also the issue of what distance metric to use. The covariance-weighted Mahalanobis distance is often used, but it ignores the sensitivities of the inputs. (If some inputs are sometimes dominated by irrelevant cosmic scatter, this will add unmodelled noise to the distance estimate.) This is a problem when we have a mix of APs, some of which have a large and others a small impact on the variance (``strong'' and ``weak'' APs). If we just use the Mahalanobis distance we loose sensitivity to the weak APs.  We could overcome these problems if we did on-the-fly interpolation of the template grid to generate new templates as we need them. Running stellar models is far too time consuming for this, but also unnecessary because the generative model is smooth: We can instead fit a forward model to a low density grid of templates as an approximation to the generative model. As we shall see, possession of a forward model opens up opportunities not available to the inverse methods, such as direct uncertainty estimates and goodness-of-fit assessment of the solution.  In this article I introduce an algorithm for AP estimation based on this forward modelling idea and iterative interpolation.  I will demonstrate it using simulations of low resolution spectra to be obtained from the Gaia mission (e.g.\\ Lindegren et al.~\\citealp{lindegren08}).\\footnote{\\tt http://www.rssd.esa.int/Gaia} Gaia will observe more than $10^9$ stars down to 20$^{th}$ magnitude over the whole Galaxy, stars which a priori span a very wide range in several APs. This includes the line-of-sight extinction parameter, \\av, which must be estimated accurately if we want to derive intrinsic stellar luminosities from the Gaia parallaxes.  AP estimation (Bailer-Jones~\\citealp{cbj05}) is therefore an integral part of the overall Gaia data processing and comprises one of the Coordination Units in the Gaia Data Processing and Analysis Consortium (DPAC) (Mignard \\& Drimmel~\\citealp{aoresponse}, O'Mullane et al.~\\citealp{omullane07}).  I will now describe the basic algorithm (section~\\ref{algorithm}). In section~\\ref{data} I then introduce the simulated Gaia spectroscopy to which \\ilium\\ is applied, with the results and discussion thereof presented in sections \\ref{sect:tefflogg}, \\ref{sect:tefffeh} and~\\ref{sect:2d1d}. The latter section also reports on a strong and ubiquitous degeneracy between \\teff\\ and \\av.  I summarize and conclude in section~\\ref{conclusions}.  Additional plots, results and discussions can be found in a series of four Gaia technical notes (Bailer-Jones 2009a,b,c,d) available from {\\tt http://www.mpia.de/Gaia}.    %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% ",
        "conclusions": "\\label{conclusions} %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%  I have introduced an algorithm for estimating parameters from multi-dimensional data.  It uses a forward model of the data to effectively perform a nonlinear interpolation of a template grid via the Newton-Raphson method.  It is intended to overcome both the non-uniqueness issue of direct modelling of inverse problems as well as the finite grid density and metric definition issues of $k$ nearest neighbours. \\ilium\\ makes use of the sensitivity of the data to the astrophysical parameters to find an optimal solution.  This is convenient, because in principle it means we don't need to worry too much about feature selection: low sensitivity features will automatically be down weighted. %The convergence is good, with only the occasional need to impose a limit on the AP update size. An important component of the algorithm is the division of the forward model into two parts which independently model the variance of the ``strong'' APs (such as \\teff\\ and \\av) and the weak APs (such as \\logg\\ and \\feh). Good fits could be obtained using low-dimensional (1 or 2) smoothing splines.  As it is based on a forward model, \\ilium\\ naturally provides AP uncertainty estimates (actually full covariances) and a goodness-of-fit, in contrast to most inverse modelling methods.  I applied \\ilium\\ to the problem of estimating APs from simulations of the low resolution Gaia spectrophotometry, BP/RP. Results are summarized in Table~\\ref{sumres}.  When limited to zero extinction stars, \\teff\\ can be estimated at G=15 to a (mean absolute) accuracy of better than 1\\% when the metallicity and surface gravity are entirely unknown (\\logg\\,=$-0.5$ to $5.0$\\,dex; \\feh\\,=\\,$-4$ to $+1.0$\\,dex). %Limiting these ranges improves the performance. At G=18.5 and G=20.0 (the limiting Gaia magnitude) the average \\teff\\ accuracy over the full range of APs is 2\\% and 4\\% respectively. At G=15 \\logg\\ can be estimated to 0.5\\,dex if \\feh\\ is entirely unknown but to 0.15\\,dex if \\feh\\ is also modelled (full  range of the three APs). Limiting to solar metallicity stars this improves to better than 0.1\\,dex and is still 0.35\\,dex at G=18.5.  \\feh\\ for cool stars can be estimated to 0.1\\,dex at G=15 and to 0.15--0.35\\,dex at G=18.5, the better result obtained if we know the star is a dwarf.  At G=20 \\logg\\ and \\feh\\ cannot be estimated to any useful accuracy.  Of course, for population studies we can average over many stars to reduce the population estimate of metallicity, limited by the systematic errors and any correlations in the data.  If we extend the strong component of the forward model to simultaneously model \\av\\ over a very wide range (0--10\\,mag), then the performance on the weak APs is significantly degraded on account of this extra variance, such that reasonable accuracies can be reached at G=15 but not at G=18.5. \\av\\ itself can be estimated remarkably well, however, 0.07--0.2\\,mag at G=15 and 0.3--0.5\\,mag at G=18.5. The extra variance also affects the \\teff\\ determination. However, these statistical errors do not tell the whole story: I have shown that there is a strong and ubiquitous degeneracy between \\teff\\ and \\av\\ intrinsic to the BP/RP spectra. Thus in addition to reporting single optimal estimates, the Gaia catalogue will need to provide the corresponding degeneracy map, which can be built in advance using the forward model. The degeneracies could be reduced (and the weak APs also then estimated more accurately) if we can use additional information. Possible sources include the apparent magnitude and parallax measured by Gaia, external data, and/or weak priors from a very simple Galaxy and stellar evolution model.  For some problems \\ilium\\ outperformed a support vector machine in terms of smaller residuals, but in other cases the SVM was better.  This should be assessed in more detail after improvements to the \\ilium\\ algorithm have been explored. For example, the updating method is very simple -- stopping after a fixed number of iterations -- whereas a more adaptive approach may help on larger variance problems (e.g.\\ noisier data). We may also want the AP update weighting to take into account the noise (and not just the sensitivity). %More generally, we should consider the impact on the algorithm of noise correlations in the spectra. These which may be introduced by the %spectral combination in the Gaia data processing. %But sphering simply removes this  Now that we have a forward model, further possibilities for modelling open up. \\ilium\\ is just a method for locating the best APs. If we define a distance metric (such as equation~\\ref{eqn:dist}), then we can define a likelihood function, $P(D|{\\bmath \\phi})$, which when combined with a suitable prior defines a (non-analytic) posterior over the APs, $P({\\bmath \\phi} | D)$. %(This could be Gaussian in $D$, not in the APs.) We can then use one of many sampling methods, such as Markov Chain Monte Carlo, to sample this as a function of ${\\bmath \\phi}$, thus yielding a complete Bayesian probabilistic solution which varies smoothly over the APs.  This approach to parameter estimation has been used in several areas of astronomy, such as galaxy classification (e.g.\\ Heavens et al.~\\citealp{heavens00}) and inference of cosmological parameters (e.g.\\ Percival et al.~\\citealp{percival07}).  (If we dispensed with the forward model and just calculated probabilities at points in the original grid, then the principle is similar to that used by Shkedy et al.~\\citealp{shkedy07}.)  While offering some advantages, this approach is much slower than \\ilium, because for the present application it would require of order $10^4$ to $10^5$ samples and hence this many evaluations of the forward model, compared to about $10^2$ for \\ilium.  % NN method with a sufficiently dense grid for the lowest noise levels encountered % might be *very* large, perhaps 100^4 points. So finding NN and the of order 10^4 objects needed % to define the errors is probably inhibitively time consuming. In this low grid density or % low noise limit regime ilium is certainly preferable. % Yet it may anyway be preferable because of its sensitivity weighting: using noise-weighted Euclidean % distance metric shows low sensitivity to the weak APs, resulting in larger-than-necessary errors for the % weak APs. The ilium sensitivity weighting overcomes this.    %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%"
    },
    "0911/0911.5490_arXiv.txt": {
        "abstract": "{} {We present  near-infrared imaging and spectroscopic high spatial resolution observations of the SMC region N88 containing the bright, excited, extincted and compact  H II region N88A of size $\\sim$1 pc.} {To investigate its stellar content and  reddening, N88 was observed  using spectroscopy and imagery in the JHKs-  and L'-band  at a spatial resolution of $\\sim$ 0.1-0.3 $\\arcsec$, using  the VLT UT4 equipped with the NAOS adaptive optics system. In order to attempt to establish if the origin of the infra-red (IR) excess is due to bright nebulosity, circumstellar material and/or local dust, we used  Ks vs J-K colour-magnitude (CM) and JHK colour-colour (CC) diagrams, as well as L' imagery.} {Our IR-data reveal in the N88 area an IR-excess fraction of $\\geq$30 per cent of the detected stars, as well as an unprecedently detailed morphology of N88A. It consists  of  an embedded cluster of $\\sim$ 3.5$\\arcsec$ ($\\sim$ 1 pc) in diameter, of at least thirteen resolved stars  superposed with an unusual bright continuum centered on a very bright star. The four brightest stars in this cluster  lie red-ward of H-K $\\geq$ 0.45 mag, and could be classified as young stellar object (YSO) candidates. Four other probable  YSO candidates are also detected in N88 along a south-north bow-shaped thin H$_2$ filament at $\\sim$ 7$\\arcsec$ east of the young central bright star. This star, that we assume to be the main exciting source, could also be complex. At 0.2$\\arcsec$ east of this  star, a heavily embedded core is detected in the L'-band. This core with L' $\\sim$ 14 mag and L'-K $\\geq$ 4.5 mag could be a  massive class I protostar candidate. The  2.12 $\\mu$m H$_2$ image of N88A resembles a shell of diameter $\\sim$ 3$\\arcsec$ ($\\sim$ 0.9 pc)  centered on the bright star. This shell is formed of three bright components,  of  which the brightest one  superposes the ionization front. The line ratios of H$_{2}$ 2-1 S(1) and 1-0 S(0) relative to 1-0 S(1), as well as the presence of high v lines,  are indicative  of photodissociation regions, rather than shocks.} {} ",
        "introduction": "The  Small Magellanic Cloud (SMC) is rich in  H II regions and young OB associations. Because of its known and  relatively small  distance ($\\sim$ 65 kpc) (Kovacs 2000), its face-on  position relatively free from  foreground extinction  (McNamara $\\&$ Feltz 1980), and low internal extinction (Westerlund 1997), it is  well suited  to study both individual stars and  very compact objects, as well as global structures. It is an  ideal laboratory  for studying  the formation   and evolution of massive stars in a low metallicity environment. Understanding the characteristics of massive stars and their interaction with their environment is a key problem in astrophysics. We have made some progress concerning the early stages of massive star formation in the galaxy, but the current knowledge about the early stages of massive star evolution in other galaxies is mediocre at best. The main  reason is that   the earliest stages of massive star evolution are deeply enshrouded, inaccessible in the optical wavelengths. Another reason is that these stars are often members of very crowded regions.  Today, high-spatial   near-infra-red (NIR) resolution observations using NACO attached  to the VLT, are able to overcome these obstacles in  the SMC. Our search for the youngest massive stars in the Magellanic Clouds (MCs) (Heydari-Malayeri   $\\&$ Testor 1982) led to the discovery of a distinct and very rare  class of H II  regions that we called high-excitation compact H II \"blobs\" (hereafter HEBs) listed in  Testor (2001). In contrast  to the ordinary H II regions of the MCs, which are extended  structures spanning several arcminutes on the sky ($>$ 50 pc) and are powered by a large number of hot stars, HEBs are very dense small regions, usually 2$\\arcsec$ to 8$\\arcsec$ in diameter (0.8 to 3 pc), ionized  by one or a few massive stars and affected by local dust. Two other HII  regions,  MA 1796 and MG 2 (less than  1pc across) heavily extincted and ionized by a small young cluster, have been  found by Stanghellini et al. (2003) in the SMC.  In the present paper we focus on the peculiar   HEB LHA 115-N88A, hereafter labelled  N88A of diameter $\\sim$1 pc (Testor $\\&$ Pakull 1985) in the extended H II region LHA 115-N88 or N88 (Henize 1956) of diameter $\\sim$ 10 pc. N88 lies in the  Shapley Wing of the SMC and contains the young cluster HW 81 ($\\sim$ 0.6$\\arcmin$)  (Hodge $\\&$ Wright 1977). It is known that the SMC is made of four H I  layers with different velocities (McGee \\& Newton 1982). This situation complicates the study, particularly in the region of the Shapley wing. However, the available H I observations provide helpful data for the study of this region. In particular, the N88 region lying at about 35$\\arcmin$ ($\\sim$ 700 pc) west of N83/N84 is not apparently associated with the H I cloud of these H II regions (Heydari-Malayeri et al. 1988). N88A should be associated with the H I component of velocity +134 kms$^{-1}$ (McGee \\& Newton 1982).  N88A is the  brightest and  the most excited HEB in the MCs. It is also the most reddened H II region in these galaxies of low-metallicity (Heydari-Malayeri et al. 2007). Numerous optical detailed  studies have been made  on this object  (Testor $\\&$ Pakull 1985, Wilcot 1994, Heydari-Malayeri et al. 1999 (hereafter HM99), Kurt et al. 1999, Testor et al. 2003). Nevertheless, many uncertainties remain  to understand the  true nature of N88A, such as its exciting source, that still  remains   unidentified, as well as the nature of the reddening.  Israel $\\&$ Koorneef (1988, 1991) have detected the presence  of molecular hydrogen in N88 through  H$_{2}$ emission which  is either shock-excited on a small scale of 0.46$\\arcsec$ by stars embedded in the molecular cloud, or radiatively excited on a large scale (3$\\arcsec$-60$\\arcsec$). However, their low spectral (R=50) and spatial (7.5-10$\\arcsec$ aperture) resolutions did not allow discrimination of these different processes. They described N88A  as a strong NIR source dominated by nebular emission containing a strong hot dust component and noticed  that N88A has an unusual blue J-H colour.  In Testor et al. (2005), at low spatial resolution,   a pure H$_{2}$ emission is detected in N88A as well as  along a south-north diffuse long  filament at $\\sim$  6-8$\\arcsec$ to the  east. In N66 (Henize 1956) a giant HII region in the SMC, Schmeja et al. (2009) have reported that most of  the H$_{2}$ emission peaks coincide with the bright component of the ionized gas and with compact embedded young clusters where candidate YSOs have been identified. Using SEST,  Israel et al. (2003) detected a  CO  molecular cloud of 1.5$\\arcmin$ x 1.5$\\arcmin$ in the region, reporting spectra and maps of the $^{12}$CO lines J=1-0 and J=2-1. Stanimirovic et al. (2000)  found that the highest values of the  dust-to-gas mass ratio and dust temperature in the SMC are found  in N88A.  IR studies  of a  similar size young  star formation region like the Trapezium region in Orion ($\\sim$0.75 x 0.75 pc)  (Lada et al. 2000)  and the more extended 30 Doradus in the LMC (Maercker $\\&$ Burton 2005) showed that during star formation, YSOs are associated with the circumstellar  material inducing IR-excess emission, and  also that the use of JHK CC diagrams are useful tools to detect this emission.  However, for young massive stars generally found in embedded clusters,  their  surrounding material destruction time scale is short, making their observation difficult (Bik et al. 2005, 2006). The most suitable wavelength to determine the nature of the  IR-excess is the L-band,  increasing  the IR-excess and reducing  the contribution of  extended emission  from reflection nebulae and H II regions (Lada et al. 2000). IR-excess can be  useful to determine the origin of  the reddening in embedded young clusters. Martin-Hernandez et al. (2008) found in N88A,  from a  mid-IR high spatial resolution  Spitzer-IRS spectrum, a  rising dust continuum and PAH bands, typical characteristics in H II regions. Using radio observations, Indebetouw et al. (2004) found that N88A is ionized by an O5 type star.  In the  present paper we present  the results  of JHK- and L'-band  high spatial resolution observations  of  N88A  and its surroundings, using adaptive optics at the VLT. Section 2 discusses the instrumentation  employed during these observations, and the data acquisition and reduction procedures used.  Section 3 describes the results and analysis of our observations, and  Section 4 gives our conclusions. \\begin{figure*}[t] \\begin{center} \\hspace{0cm} \\includegraphics[width=14.0cm]{testor_Fig1.jpg} \\caption{JHK colour composite image of SMC N88 containing the bright HEB N88A. The image  size extracted from the S54 camera field corresponds to 50$\\arcsec$ x 50$\\arcsec$ (or $\\sim$ 14 pc x 14 pc). The lower spatial density of faint stars close to the edges of the frame delimits approximatively the cluster N88-cl. In the north-west quadrant the lack of stars is explained by the presence  of the molecular cloud (HM99). A  red filament  is visible east of N88A. J, H and K are respectively coded B, G and R.  North is up and East is left. } \\label{fig:S54} \\end{center} \\end{figure*} \\begin{figure*}[t] \\begin{center} \\hspace{0cm} \\includegraphics[width=18.0cm]{testor_Fig2.pdf} \\vspace{-9cm} \\caption{Finding chart  in the J-band   obtained with the S54 camera (Field 1).  The inset corresponds to an enlargement (factor 2) and unsaturated image of  N88A extracted from Field 1. The numbering refers to Table 3. The total field size corresponds to 48.6$\\arcsec$ x 47.5$\\arcsec$ (or $\\sim$ 14 pc x 14 pc). } \\label{fig:S54inset} \\end{center} \\end{figure*} \\begin{figure*}[t] \\begin{center} \\hspace{0cm} \\includegraphics[width=18.0cm]{testor_Fig3.pdf} \\vspace{-9cm} \\caption{Ks-band image of SMC-N88 obtained with the S27 camera (Field 2). The location of the slits labelled S1, S2 and S3  used in the spectroscopic mode is indicated by a solid line. In N88A the positions of the stars  $\\#$37, $\\#$41, $\\#$42 and $\\#$47 are indicated. The bottom left is an enlargement of 0.55$\\arcsec$ x 0.55$\\arcsec$ containing the complex object N88B. The image  size  corresponds to 25$\\arcsec$ x 25$\\arcsec$ (or $\\sim$ 7 pc x 7 pc). } \\label{fig:S27inset} \\end{center} \\end{figure*} \\renewcommand{\\arraystretch}{0.5} \\tabcolsep=0.04cm \\begin{table} \\caption{ Log of the photometric VLT/NACO observations.} \\begin{center} \\begin{tabular}{llccccc} \\noalign{\\smallskip} \\hline \\hline \\noalign{\\smallskip} Id.       &  Filter  & Expo.   &Mode   &  Date      & Seeing    & FWHM  \\\\ &          &t(s) x n &       &            &($\\arcsec$)&($\\arcsec$)      \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} N88A     &J 1.27 $\\mu$m         &20 x 30   & -         &     -     & -         & 0.35     \\\\ &H 1.66 $\\mu$m         &20 x 30   & -         &     -     &           & 0.27     \\\\ &Ks 2.18 $\\mu$m        &30 x 30   & S54       &9/10/2004  &   0.6-0.9 & 0.21     \\\\ &L' 3.8 $\\mu$m         &0.2 x 26  & S27       &10/10/2004 &   0.9-1.2 & 0.10     \\\\ &NB-2.12               &200x 20   & S54       &     -     &           & 0.19     \\\\ &NB-2.24               &150x 20   & S54       &     -     &           & 0.25     \\\\ &Ks                    &60 x 30   & S27       &11/10/2004 &           & 0.10     \\\\ \\noalign{\\smallskip} \\hline \\end{tabular} \\end{center} \\end{table} \\begin{figure*}[t] \\begin{center} \\hspace{-1cm} \\includegraphics[width=18.0cm]{testor_Fig4.pdf} \\vspace{-9cm} \\caption{Enlargements of the HEB N88A in Field 2. a) Ks-band obtained with the S27 camera. The numbering refers to  Table 4 (Field 2). b)  the DDP process  applied to the Ks-band allows to  enhance  the faint stars  of N88A-cl. A few stars not identified with DAOPHOT and small features are indicated by arrows. c) Stromgren y image obtained with the HST showing the absorption lane, the stars $\\#$1, $\\#$2 as well as  the `dark spot' of HM99. The Stars $\\#$1 and $\\#$2 correspond to our stars  $\\#$41 and  $\\#$42. d) L'-band obtained with the S27 camera, the different components are indicated. The field corresponds to 4.43$\\arcsec$ x 4.32$\\arcsec$ (or 1.27 pc x 1.27 pc) and is outlined in (Fig. \\ref{fig:S27inset}). } \\label{fig:S27zoom} \\end{center} \\end{figure*} \\renewcommand{\\arraystretch}{0.5} \\tabcolsep=0.04cm \\begin{table} \\caption{ Log of the VLT/NACO long-slit spectroscopic observations.} \\begin{center} \\begin{tabular}{lcccccc} \\noalign{\\smallskip} \\hline \\hline \\noalign{\\smallskip} Id.        &  Date      & Slit     &  Exposures & Mode     &$\\lambda$/$\\delta$$\\lambda$& Seeing       \\\\ &            & (mas)         &   t(s) x n &          &                           &($\\arcsec$)   \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} S1         &9/10/2004  &172        & 100 x 10   &S54-4-SHK &   400                     & 0.6-0.9      \\\\ S2         &10/10/2004 & -         & -          &  -       &   -                       &  -           \\\\ S3         &11/10/2004 & -         & -          &  -       &   -                       &  -           \\\\ &           &           &            &          &                           &              \\\\ hip8485    &9/10/2004  & -         & 2.5 x 4    &  -       &   -                       &   -          \\\\%274-277 hip29968   &9/10/2004  & -         & 2 x 4      &  -       &   -                       &   -          \\\\%297-300 hip103087  &11/10/2004 & -         & 1.8 x 4    &  -       &   -                       &   -          \\\\%143-146 \\noalign{\\smallskip} \\hline \\end{tabular} \\end{center} \\end{table} ",
        "conclusions": "We present high spatial resolution imaging of  FWHM $\\sim$ 0.12$\\arcsec$ - 0.25$\\arcsec$ in the JHKL'-band  of  the HEB N88A and its immediate environment, and the main results are as follows:  N88A is associated with a cluster  that contains at least  thirteen stars centered approximatively on the bright central star  $\\#$41, that could be multiple.  N88A coincides  perfectly with the 3cm radio peak and should be ionized not only by the star   $\\#$41  classified  of  type O8 V,  but also by   other low to high-mass stars.  From analysis of the JHK CC diagram  we found four possible MYSO star candidates in  the  N88A cluster, as well as three probable YSOs in the red bow east of N88A. In N88 at least 30$\\%$ of the  detected stars have an IR-excess.  From  the K-L excess we  found that the core L1-C in N88A  should be a heavily embedded, high mass protostar of Class I.  The continuum emission at the position of  $\\#$41 is very bright and represents  about 30$\\%$ of the  Br$\\gamma$ emission peak.  The H$_2$ infrared emission in N88A resembles a shell formed mainly by three peaks of which one coincides with the ionization front.  We show that the excitation mechanism may be  caused predominantly by PDRs, without excluding combination with shocks.  The morphology   of N88A could be comparable with  galactic   regions  such as the nearby Trapezium region in the Orion nebula.  Future JHK band  imaging data, using   higher spatial resolution  and  longer wavelengths, as L' and M'  provided by the NACO S13 camera  are still needed to  disentangle the IR-excess origin in N88A. Higher spectral  resolution spectra are also required to obtain a better analysis of the different spectral lines like the Br$\\gamma$  emission.  These new observations  should  allow to  investigate more  thoroughly  the HEB N88A. This object, which is the brightest, the most excited and reddened of the MCs,  presents a unique opportunity  to progress in the knowledge of newborn  massive stars in regions of low metallicity."
    },
    "0911/0911.5345_arXiv.txt": {
        "abstract": "We study the properties of SDSS galaxies with and without AGN detection as a function of the local and global environment measured via the local density, the mass of the galaxy host group (parameterised by the group luminosity) and distance to massive clusters.  {  Our results can be divided in two main subjects, the environments of galaxies and their relation to the assembly of their host haloes, and the environments of AGN.  (i)} For the full SDSS sample, we find indications that the local galaxy density is the most efficient parameter to separate galaxy populations, but we also find that galaxies at fixed local density show some remaining variation of their properties as a function of the distance to the nearest cluster of galaxies (in a range of $0$ to $10$ cluster virial radii).  These differences seem to become less significant if the galaxy samples are additionally constrained to be hosted by groups of similar total luminosity.   {  If instead of fixing the local density, the mass of the host group is held fixed at a given value, the fraction of red galaxies also increases as the distance to clusters diminishes, indicating that neither the local density or the host halo mass contain all the information on the environment}.  (ii) In AGN host galaxies, the morphology-density relation is much less noticeable when compared to the behaviour of the full SDSS sample, indicating a lack of sensitivity to the host group mass during the AGN phase probably due to the higher typical luminosities of the AGN hosts.  In order to interpret this result we analyse {  control samples constructed using galaxies with no detected AGN activity} with matching distributions of redshifts, {  stellar masses}, r-band luminosities, $g-r$ colours, concentrations, local densities, host group luminosities, {  and fractions of central and satellite galaxies; the aim in using the control sample is to detect any correlations between the AGN detection and other AGN host properties that are unrelated to the AGN selection}.  The control samples also show a similar small dependence on the local density indicating an influence from the AGN selection, but their colours are {  slightly bluer} compared to the AGN hosts regardless of local density.  {  Furthermore}, {\\it even when the local density is held fixed} at intermediate or high values, and the distance to the closest cluster of galaxies is allowed to vary, AGN control galaxies away from clusters tend to be bluer than the AGN hosts.  {  However}, AGN in bright, low concentration hosts (i.e.  disky morphologies) are bluer than galaxies in the control sample, connecting the presence of discs to AGN activity {  even under a controlled comparison between active and inactive galaxies.} ",
        "introduction": "The description of the large-scale structure in the Universe, and how it relates to the properties of galaxies in different environments has become more intricate in the last few years.  {  In order to understand the environmental dependence of galaxy properties it is now necessary to combine results where the environment comes from local density estimates, the host halo mass, the distance to cluster centres, together with a theoretical understanding of what could influence the SF activity and morphological transformations in galaxies.  This introduction will first review the assembly effects on the galaxy population of haloes of equal mass, then the results from the literature on the environment studies of galaxies, and finally the properties of the hosts of AGN, and their dependence on environment.}  \\subsection{Properties of galaxies in equal mass haloes: assembly effects}  Up to only a few years back it was believed that the mass of a dark-matter halo was the only important parameter that defined its clustering properties (see for instance Padilla \\& Baugh 2002; Sheth, Mo \\& Tormen 2001; Mo \\& White 1996) and galaxy population (e.g. Cole et al., 2000). The latter has been studied in the framework of the Halo Model (Cooray \\& Sheth 2002; Cooray 2005, 2006, Jing, Mo \\& B\\\"orner, 1998, Peacock \\& Smith, 2000), whereby the typical number of galaxies per dark-matter halo, typical central and satellite luminosities, {  star formation activity}, and possibly their colours depend solely on the host halo mass ({  Weinmann et al. 2006}; Wang et al. 2008). However, recent studies on the clustering of dark-matter haloes found that the age and assembly history of haloes of similar mass also play a role in their clustering amplitude.  The dependence on halo age was first reported by Gao, Springel \\& White (2007), and has been studied in more detail finding a general dependency on the way in which a halo is assembled by Gao \\& White (2007), Jing, Suto \\& Mo (2007), Croton, Gao \\& White (2007), Wechsler et al., (2007), and Li, Mo \\& Gao (2008). The assembly bias has also been detected in observations by Wang et al. (2008); their studies only concentrate on the clustering amplitude of different samples selected according to group colour which they associate with group age. In addition to a difference in clustering amplitude, Zapata et al. (2009) demonstrated that groups of similar mass and different assembly histories also show important differences in their population of galaxies, an aspect that can be added to current halo model and conditional luminosity functions, and that can also be used to improve the modeling and understanding of the evolutionary processes that shape galaxies in the observed Universe.  { Extending this particular problem, Wang et al. (2009) found a sub-population of red, dwarf isolated galaxies that are not satellites of larger haloes but that are concentrated around their nearest massive halo.   Such isolated dwarfs are expected to be able to form stars and therefore show blue colours.  These red, dwarf galaxies would fit within the galaxy formation scenario if they inhabit subhaloes that have been expelled from clusters and have acquired their red colours earlier, when they were under the effects that satellites suffer in Clusters such as strangulation, ram-pressure and tidal stripping, processes that quench the star formation activity. On the theoretical side, Wang, Mo \\& Jing (2009) find that some small fraction of the assembly bias effect comes from such subhaloes.}  These results however, concern only the population of galaxies inside individual dark-matter haloes (or ejected sub-haloes).  \\subsection{Measurements of environment and the properties of galaxies}  On the large-scale structure side, starting with the pioneering work of Dressler et al. (1980), the location of a galaxy has been shown to have an influence on its average properties.  Using large redshift surveys, Balogh et al. (2004) found a smooth transition from high density regions populated by red galaxies, towards blue dominated low density regions. These results pointed to the density of the region where galaxies live as being responsible for important characteristics of the galaxy population.  This was also confirmed at higher redshifts in the Sloan Digital Sky Survey (SDSS, York et al. 2000), by O'Mill, Padilla \\& Lambas (2008). Numerical simulations following the evolution of dark-matter, with galaxy populations imprinted onto them via semi-analytic models, have shown that the prevalence of red galaxies in high density regions can be naturally achieved in a $\\Lambda$CDM hierarchical cosmology (Croton et al. 2006; Bower et al. 2006; Cattaneo et al. 2006, 2008; Lagos, Cora \\& Padilla 2008; Lagos, Padilla \\& Cora 2009). A remaining question with respect to this general behaviour resided in whether this population change was due to local or global effects,{  that is, whether the dependence of properties on the local density responds mainly to the dependence on halo mass alone, or also to assembly and other effects such as ejection from hosts}.  Kauffmann et al. (2004) found little evidence of a dependence of star formation (SF) indicators on the environment measured over scales larger than $1$h$^{-1}Mpc$.  However, more recent results around voids in the SDSS by Ceccarelli, Padilla \\& Lambas (2008) suggested that the local density is not the only parameter defining a galaxy population.  They showed that at fixed local density, the fraction of blue galaxies increases towards the regions corresponding to void walls; this represented a first measurement of a large-scale modulation of star formation in galaxies. This behaviour was later confirmed in numerical simulations by Gonz\\'alez \\& Padilla (2009), who analysed semi-analytic galaxies from the SAG model (Lagos, Cora \\& Padilla 2008; Lagos, Padilla \\& Cora 2009), and found that this was also the case for galaxies of fixed local density, at increasingly larger distances from clusters of galaxies, or closer to voids.  Furthermore, they found that the reason behind this change at fixed local density, is mostly due to the effect of the mass of the host halo, with a remnant variation that can be tied to the assembly history of the haloes hosting these galaxies. The reason for a dependence on mass is related to the different physical processes taking place in haloes of different masses.  For instance, at very low halo masses, $M<10^{10}$h$^{-1}M_{\\odot}$, reionization from a UV radiation field plays an important role in the photo-ionization of baryons; this effect along with stellar feedback reduces significantly the baryon fraction available to form stars in such low mass haloes (Tassis et al. 2002). The reason for a dependence on assembly is that, for instance, more recent merger activity can have an important impact on the colours of galaxies. Therefore, even though local processes may extend out to $\\sim 1$h$^{-1}$Mpc, global processes affecting the underlying mass function and the rate of growth of haloes (assembly) help extend the variation of galaxy properties to much larger scales, up to tens of Mpc.  \\subsection{AGN host properties, and their dependence on environment}  Regarding the environmental dependence of galaxies with active galactic nuclei (AGN) there is still a debate on the  frequency of occurrence of AGN as a function of the local density.  For instance, Miller et al. (2003) found that the fraction of galaxies with AGN appeared to be constant with local density, even when dividing the sample according to the host morphology into spiral and elliptical galaxies (see also Dressler et al. 1999). Using the SDSS, Kauffmann et al. (2004) confirmed this result only for AGN with weak emission in the OIII line.  They claim that for high OIII emitters the fraction of galaxies with AGN decreases as the local density increases (see also Popesso \\& Biviano 2006). It should be noted that these studies do not reach densities as high as those in clusters of galaxies. Pasquali et al. (2009) studied the variation of the SF and AGN activity of central and satellite galaxies as a function of their host dark-matter halo mass (ranging from low mass groups to clusters of galaxies), to find a smooth and continuous transition from SF to AGN and to radio activity as the halo mass { and the stellar mass} increase, for both centrals and satellites, the latter showing lower activity than the former at all halo masses; { furthermore they show that the dependence on environment (as indicated by the halo mass) is weaker than on the stellar mass of the host galaxy}. Other studies of AGNs in clusters show a higher frequency of X-ray sources than in the field (Cappi et al. 2001; Molnar et al. 2002; Martini et al. 2002). Furthermore, using clustering analyses, Gilli et al. (2003) also found higher correlation amplitudes for X-ray selected AGNs. Additionally, Martini, Sivakoff \\& Mulchaey (2009) found that the fraction of AGN in clusters of galaxies decreases for lower redshifts; this could help obtain a smooth transition between the results from these X-ray detected AGN and those by Kauffmann et al. (2004).  There are also studies of the environment in which AGN are embedded, { which focus on the morphology and other characteristics of their hosts and their neighbors (including their local density, colours, etc)}. Choi et al. (2009) studied AGN hosts in the SDSS Data Release 5 (DR5, Adelman-McCarthy et al. 2008).  They find that late morphological types are the dominant hosts regardless of AGN power, and that bluer late types host the most powerful AGN (also in agreement with Coldwell et al. 2009), {  confirming earlier results by Kauffmann et al. (2003a), who additionally show that the young stars in the powerful AGN reside preferentially away from the nucleus, and that their hosts have suffered a starburst in the recent past.  Additionally, Kauffmann et al. (2003a) show that these results do not depend on the Seyfert type (for equal OIII luminosity).  Going into more detail on the different AGN types, Kewley et al. (2006) show that the hosts of LINERs (low-ionization narrow emission-line regions, characterised by low luminosities with respect to the Eddington value) are older, more massive, less dusty, less concentrated and have higher velocity dispersions then Seyfert galaxies (characterised by higher $L/L_{EDD}$).} Waskett et al. (2005) found that the environments of AGN at $0.4<z<0.6$ are similar to those of normal galaxies at these same redshifts, and consist of group-like environments. Gilmour et al. (2007) confirm this result finding that AGN in bright galaxies at $z=0.17$ (in a region corresponding to a supercluster) show a preference for intermediate environments, and are preferentially surrounded by blue galaxies. More recently in zCOSMOS (Lilly et al. 2007), Silverman et al. (2009) confirm that AGN tend to prefer intermediate environments such as groups, and also find AGN hosted by massive galaxies with stellar masses $M_*>2.5\\times 10^{10}M_{\\odot}$ that occur preferentially in low density environments where disruptive processes (mergers) are lessened. Results on the environment of high redshift AGN sources (detected via X-ray counterparts of ACS images) are presented by Martel et al. (2007), who find that $z\\simeq1$ AGN prefer regions which are more crowded than for the typical galaxy in the field, giving clues on the possible nature of the hosts of these AGN; however, they also warn about possible systematic effects in their results. More recently, von der Linden et al. (2009, see also Montero-Dorta et al. 2009) study the variation of the fraction of AGN hosts as a function of the distance to cluster centres in the SDSS.  They analyse the time-scale over which the AGN activity shuts down and find that this coincides roughly with the crossing time-scale of the clusters, and that this is similar to the timescale of general SFR shut-down; {  note that their analysis does not discriminate the changing local density as the distance to the cluster centres varies}.   Centering mostly on low accretion rate AGNs, these studies have helped shed light on the feeding and growing mechanisms of AGN, and on their influence on the star formation (SF) activity of the host galaxies. For instance, Koulouridis et al. (2006) studied the origin of the AGN activity in relation to their immediate environment. Their analysis of bright IRAS galaxies in comparison to that of AGN helped them conclude that close interactions drive fuel to the AGN, which produces obscured AGN activity.  They infer that only when the close neighbor moves away, the AGN becomes unobscured.  Martini et al. (2004) presents an extended review on the relation between AGN fueling and environment. These results are in good agreement with morphological studies by Pasquali, van den Bosch \\& Rix (2007) who find that the AGN activity is associated with an increase in the fraction of galaxies with stellar discs, and Hao et al. (2009) who detect an increase in the fraction of barred hosts with the presence of AGN.  Along the same lines, Donoso et al. (2009) find, via clustering studies of radio-loud AGN, that there is an influence from the gaseous content on scales of the host dark-matter halo on the probability that an AGN is radio-loud and on the jet luminosity (see Donoso, Best \\& Kauffmann 2009, for a study on the evolution of radio AGN with redshift).  \\subsection{The object of this work}  In this paper we will focus on the variation of galaxy properties at different global environments as traced simultaneously by the local density and by the distance to the nearest clusters of galaxies, {  which as mentioned above should include both a dependence on the host dark-matter halo mass, halo assembly, and other effects such as, for instance, dynamical coherence or ejection from larger haloes}. We will study galaxy colours, concentrations, and colour-magnitude relations, for normal galaxies on the one hand and for galaxies with AGN on the other, with the aim to provide further clues on the interplay between large-scale structure, assembly history, and galaxy evolution. {  The reader should bear in mind that the environmental behaviours of normal galaxies and of AGN hosts will not be compared to one another due to the different selection effects in play, but will be used to try to improve our understanding of the galaxy formation process.} {  In the case of AGN hosts,} we {  want to stress the fact that we} will focus our analysis of the environmental dependence via variations with the environment instead of studying the fraction of galaxies hosting AGNs.  In order to overcome the possible influence from selection effects associated with the detection of AGN signatures in the spectra of the galaxies we will rely on the definition of control samples that match several properties of the AGN hosts to mimic their selection function as best as possible. {  Finally, we will use the galaxy colour as an indicator of star formation activity, bearing in mind that this quantity also depends on star-formation history, metallicity and reddening.}  This paper is organised as follows, Section 2 presents the sets of observational data and definition of samples used in this work; in Section 3 we explore distributions of galaxy colours and concentrations for the full sample of galaxies and for AGN. We present our results and compare them to the literature in Section 4, which also contains our conclusions. ",
        "conclusions": "{ The aims of this paper consisted on studying the variations of the properties of galaxies with their environment, and to compare it to what is expected from galaxy formation models.  We have analysed separately these variations for normal galaxies and for galaxies with AGN.  We will now compare our results to previous studies for these two types of galaxies in the following subsections.  \\subsection{Results for the full galaxy population}  There have been numerous studies on the environmental dependence of galaxy properties, where the environment has been given for example by projected galaxy densities on different scales (e.g. Balogh et al. 2004; Baldry et al. 2005; Kauffmann et al. 2004; O'Mill, Padilla \\& Lambas 2006), the host halo mass (Weinmann et al. 2006), or the distance to cluster centres (e.g. Wang et al. 2009; Smriti \\& Raychaudhury 2009).  As these three quantities are correlated, the results obtained by different authors consistently show that galaxies located in lower density environments, or further away from clusters, or in host haloes with lower masses, are characterised by bluer colours.  Given that it is expected that the local density in the very high redshift Universe affects the final local dark-matter halo mass (e.g. Weinmann et al. 2006; Gonz\\'alez \\& Padilla 2009), and in turn, the history of mergers, it could be expected that the properties of galaxies at {\\it equal $z=0$ local densities} but different large-scale environments tracing different initial conditions, could show different properties.  Note that this would also be the case for galaxies with {\\it equal host-halo masses}, as the different merger histories, could also imprint different galaxy properties in haloes of equal mass, but different large-scale environments { and therefore, different assembly histories}. This is clearly expected in view of the assembly bias effect, now detected in simulations (starting with Gao, Springel \\& White 2005) and observations (Wang et al. 2008; Zapata et al. 2009), where equal mass haloes with different assembly histories show different clustering properties as well as different galaxy populations.  Therefore, we studied the variation of red galaxy fractions as the distance to clusters increases, but at {\\it fixed values of local densities or host-halo masses}, and found an increasing fraction of blue galaxies with a high statistical significance ($9\\sigma$ and $5\\sigma$ for fixed local densities and host-halo masses, respectively). With this result, we confirmed those found for semi-analytic galaxies by Gonz\\'alez \\& Padilla (2009).  Previously, Ceccarelli, Padilla \\& Lambas (2008) had been able to find an indication of this effect, but with much lower statistical significance, comparing galaxies in equal local density environments, but in void walls, and outside voids.  Gonz\\'alez \\& Padilla (2009) also found that even when the local densities and { DM halo} masses are constrained to fixed values, different large-scale environments can still induce changes on the modelled galaxy properties, an effect that could be attributed to  coherent motions or further assembly effects.  We repeated this analysis with our full galaxy sample, but were not able to find any significant remaining variations of colours.  Given the large errorbars involved, larger galaxy samples are needed in order to confirm this effect in observations.  \\subsection{Results on AGN hosts}  Our analysis of the AGN hosts used projected local densities, host halo masses and their distance to clusters of galaxies (and combinations of these) as proxies for their environment.  We also defined a control sample of non-AGNs such that their stellar masses, r-band luminosities, colours, local densities, concentrations, fractions of central and satellite galaxies, and host halo masses, show the same distributions as the AGN hosts.  By applying all these restrictions independently on each parameter of non-AGN galaxies, any correlations between these quantities introduced by, or allowing the, AGN activity will produce differences with the control sample, particularly when analysing different environments. In the literature, several works make comparisons between the properties of AGN hosts and of non-AGN galaxies characterised by similar properties, but still allowing for differences on several parameters. For instance, most studies on the fractions of AGN essentially make a comparison to the full galaxy population divided according to local density, host morphology, stellar mass, and nuclear OIII luminosity (Miller et al. 2003; Kauffmann et al. 2004; Popesso \\& Biviano 2006; Cappi et al. 2001; Molnar et al. 2002; Martini et al. 2002; Martini, Sivakoff \\& Mulchaey 2009; Pasquali, van den Bosch \\& Rix 2007). In other cases, the studies served to identify the environments of AGN hosts (Choi et al. 2009; Waskett et al. 2005; Gilmour et al. 2007; Silverman et al. 2009; Martel et al. 2007; von der Linden et al. 2009; Montero-Dorta et al. 2009; Koulouridis et al. 2006; Martini et al. 2004) using a single parameter such as local density, distance to cluster centres or cluster membership; or the typical AGN host properties (Kauffmann et al. 2003a; Kewley et al. 2006; Hao et al. 2009), as in Pasquali et al. (2009) who studied the hosts of AGN as a function of the host halo mass. In the process of constructing our control samples, we compared the distributions of several parameters of AGN hosts and the full galaxy population, and we were able to confirm several of these results (for example, constant fraction of AGN as a function of local density, a larger occurrence of AGN in higher mass host haloes).  Our approach is complementary to these previous works, and is designed to find the effect of changing only one parameter, the detection of AGN activity. In our study we found that as the local density or the host halo mass increase, AGN hosts become slightly redder (while occupying the green valley) and of higher luminosity, consistent with studies of the hosts of AGN.  However, even though the control galaxies also show this qualitative behaviour, they are always slightly bluer and fainter than AGN hosts, a clear difference coming from a condition related to the AGN detection, either an effect of the BH on the galaxy, or a condition that allows the BH activity.  In this sense, this could indicate that AGN in clusters are hosted by larger galaxies (a correlation in the AGN hosts not present in the control samples) in order to be able to retain gas in the intracluster conditions. Note that this result is seemingly contrary to what was found by Pasquali, Van den Bosch \\& Rix (2007, { although this paper deals only with early type galaxies}) and Hao et al. (2007), since they connect the AGN activity to star formation events (disky and barred galaxies, respectively). However, part of this signal is due to the typical hosts of AGN, and not to the detection of AGN activity in these galaxies.  Consistent with the interpretation that our control samples contain dormant black holes where the SF has regained its strength, we find that when the local density is held fixed at intermediate (or high) values and the distance to clusters varies, the hosts of AGN { show constant colours.  Thus, close to the clusters AGN are significantly bluer than the control galaxies } (however, it could also be possible that AGN near cluster centres need more available gas to remain active). In this case most of the galaxies (AGN hosts and control) are satellites of the cluster and therefore the AGN hosts under analysis would only conform to a subsample of the AGN studied by Pasquali et al. (2009), who do not divide their sample according to local density.  At low local densities, the low OIII luminosity AGN are found to extend their range of colours and concentration to lower values as the host luminosity increases.  Only part of this behaviour comes from the morphologies and characteristics of the AGN hosts, as the control sample shows a less clear trend, in slightly better agreement with the behaviour of normal galaxies (higher concentrations and colours).  This could favour the view that the AGN are not the most important factor in transforming galaxies into red and dead objects as suggested by Schawinski et al. (2007), in comparison to cluster environmental processes (Kewley et al. 2006) since at low densities low OIII luminosity AGNs can even show blue colours.  \\subsection{Perspectives}  Even though we were able to detect signatures of assembly on the properties of galaxies for our full sample, in the form of different colours for galaxies at fixed local densities and different distances from clusters, there still remain many questions regarding the effect from particular details of the development of structures in the Universe. For instance, the apparent correlation between coherent motions and the remnant variations of galaxy properties at fixed local density and host halo mass (Gonz\\'alez \\& Padilla 2009), or the contamination of the areas surrounding clusters by ejected satellite galaxies (Wang et al. 2009).  This will require considerably larger galaxy samples, which could also be useful in detecting the evolution of this phenomena with redshift.   These effects will be small but they are an expected outcome of the development of galaxies within a $\\Lambda$CDM scenario.  The same applies to the study of AGN hosts, where the comparison to a strict control sample has allowed us to find the correlations between the detection of an AGN and the local density, the distance to clusters (with respect to a control sample), and the host halo mass.  More systematic studies on the effect of each parameter included in the selection of the control sample are possible and would allow a more detailed understanding of the processes behind the detection of an AGN and its relation to the different phenomena associated to the origin and development of structure. }"
    },
    "0911/0911.0426_arXiv.txt": {
        "abstract": "No terrestrial planet formation simulation completed to date has considered the detailed chemical composition of the planets produced. While many have considered possible water contents and late veneer compositions, none have examined the bulk elemental abundances of the planets produced as an important check of formation models. Here we report on the first study of this type. Bulk elemental abundances based on disk equilibrium studies have been determined for the simulated terrestrial planets of \\cite{dave}. These abundances are in excellent agreement with observed planetary values, indicating that the models of \\cite{dave} are successfully producing planets comparable to those of the Solar System in terms of both their dynamical \\emph{and} chemical properties. Significant amounts of water are accreted in the present simulations, implying that the terrestrial planets form ``wet'' and do not need significant water delivery from other sources. Under the assumption of equilibrium controlled chemistry, the biogenic species N and C still need to be delivered to the Earth as they are not accreted in significant proportions during the formation process. Negligible solar photospheric pollution is produced by the planetary formation process. Assuming similar levels of pollution in other planetary systems, this in turn implies that the high metallicity trend observed in extrasolar planetary systems is in fact primordial. ",
        "introduction": "Terrestrial planet formation, both in terms of the dynamics and chemistry involved, is still not fully understood. Dynamically, basic planetary formation is described through the planetesimal theory (see \\cite{chambers} for a detailed review). This theory sees planetary formation occurring through three main steps. Initially, dust settles into the midplane and accretes to form planetesimals, the first solid bodies of the system. This stage is followed by the collisional accretion of planetesimals to produce planetary embryos. The growth of the embryos occurs initially via runaway growth (where an increasing geometric cross-section and gravitational field allow for the accretion of an ever increasing number of planetesimals) before transitioning to oligarchic growth (where neighboring embryos grow at similar rates). Finally, as numbers of planetesimals decrease, the interaction between embryos becomes the dominant factor as they perturb each other onto crossing orbits, thus producing accretion via violent collisions.  Many attempts have been made to simulate the third stage of terrestrial planet formation described above (e.g. \\citealt{kom1,kom2,chambers2,chambers3,raymond04}). However, these simulations have had limited success and no simulation has been able to \\emph{exactly} reproduce the terrestrial planets of the Solar System in terms of their number, masses and orbital parameters. For example, direct N-body simulations have produced systems with approximately the correct number of planets yet with an orbital excitation greater than that observed for the Solar System \\citep{chambers2,chambers3,raymond04}. Similarly, simulations incorporating tidal torques (e.g. \\citealt{kom1,kom2}) have produced too many planets but with more favorable excitation levels.  Recent developments in modeling have occurred with the incorporation of dynamical friction, the process whereby equipartitioning of energy between low mass planetesimals and larger embryos results in reduced relative velocities for the embryos, thus increasing their probability of accreting. Dynamical friction has been shown to be a viable mechanism to reduce the dynamical excitation levels of the final planets to better agree with observed values within the Solar System \\citep{levison}.  Several authors have recently performed N-body simulations with a high enough resolution to accurately model the effects of dynamical friction during terrestrial planet accretion, such as \\citet{dave} and \\citet{raymond:2006,raymond:2007,raymond:2009}.  The simulations of \\citet{dave}, which will be used in this work, represent a factor of $\\sim$5 increase in the number of gravitationally interacting bodies compared to most other previous simulations of this type (e.g. \\citealt{chambers3}), and because of the large number of small planetesimals in the simulation, dynamical friction is more accurately treated than in those previous simulations. The terrestrial planets produced in the \\citeauthor{dave} simulations are a significantly better fit to the observed properties of the terrestrial planets of the Solar System than those of previous simulations. There are fewer planets (averaging 3.25 planets per simulation) and the dynamical excitation of the planetary systems is comparable to that of the actual terrestrial planets in the Solar System. Furthermore, the accretion timescales of the planets simulated are in agreement with the accretion timescale for the Earth ($\\sim$60 Myr) as obtained from $^{182}$Hf-$^{182}$W dating \\citep{touboul}.  The simulations of \\citet{raymond:2006,raymond:2007,raymond:2009}, which have a comparable resolution to the \\citeauthor{dave} simulations, also produce a comparable improvement over previous simulations.  Thus it can be seen that these recent N-body simulations have provided us with some of the most realistic and feasible models of terrestrial planet formation completed to date, and clearly represent a significant step towards understanding the last stage of the planetary formation.  Several other types of terrestrial planet formation models have been developed recently as well. The effects of the depleting Solar nebula gas on terrestrial planet accretion, both in terms of its effects through gravitational tides and also by causing the sweeping of secular resonances through the terrestrial planet region, has been studied analytically by \\citet{nagasawa:2005}.  More recently, \\citet{thommes:2008} incorporated these effects into an N-body integrator.  That work did not incorporate a planetesimal swarm (and hence did not include dynamical friction), but found that under certain circumstances, the masses and dynamical state of the terrestrial planets in the Solar System could be reproduced. \\citet{bromley:2006} and \\citet{kenyon:2006} developed a hybrid code in which large bodies are treated via an N-body integration scheme, while small bodies are treated with a multi-annulus coagulation model that can treat the growth of bodies starting from km-scale planetesimals.  They too found that their model could reproduce many of the dynamical characteristics of the terrestrial planets in the Solar System, and while the final eccentricities of the planets were somewhat high, it is likely that a more complete treatment of the effects of fragmentation, and of tidal drag on large bodies from the nebular gas, would help to decrease them.  A full comparison of these models with the direct N-body simulations of \\citet{dave} and others is beyond the scope of this article, but we will discuss the implications of those models for the work presented here in Sec.~\\ref{sec:discussion}.  While models have advanced significantly in their ability to reproduce the terrestrial planets from a dynamical point of view, the vast majority of terrestrial planet formation studies completed to date have been limited (at best) in their simultaneous chemical composition studies. Several studies have estimated the amount of potentially hydrated material that may have accreted into the terrestrial planets (e.g. \\citealt{raymond04,dave,raymond:2007}) while \\citet{dave} and \\citet{raymond:2007} also considered the delivery of siderophile-element-rich late veneer material needed to account for the highly siderophile element budget of the Earth's mantle \\citep{chou}. However, detailed examination of the bulk chemical composition of the resulting planets as a test of dynamical simulations has never been thoroughly explored. This combination of the two approaches to produce a comprehensive model of both the chemistry and dynamics of terrestrial planet formation is essential to determine how well current numerical simulations reproduce not only the masses and dynamic state of the terrestrial planets, but also their bulk composition. The present study represents a first step towards such a model.  We have derived detailed bulk elemental abundances for all of the terrestrial planets formed in the simulations of \\cite{dave}. These predicted compositions are compared to the bulk compositions of the actual terrestrial planets as an examination of the chemical plausibility of the dynamical simulations currently being used to model terrestrial planet formation. This approach allows us to not only examine the bulk elemental composition of the final planets but to also study the compositional evolution with time. Additionally, we also investigate the delivery of hydrated material, which has important implications for the development of habitable terrestrial planets, along with the composition of the ``late veneer'' material accreted by the planets. Finally, the amount of material accreted by the Sun as ``pollution'' during the terrestrial planet formation process and the resulting changes in stellar photospheric abundances are also examined. This study is the first to attempt to produce a comprehensive, simultaneous model of both the chemistry and dynamics of terrestrial planet formation. As such, it represents a significant step forward in producing a cohesive model of terrestrial planet formation. ",
        "conclusions": "} The simulations successfully produce terrestrial planets that are in excellent agreement with the terrestrial planets of the Solar System, in terms of both their dynamics and their bulk elemental abundances. Although current simulations are unable to capture the finer details of planetary formation (such as the composition, and sometimes the amount, of the late veneer material), the success of these simulations on a broader scale provides us with increased confidence in the dynamical models of \\cite{dave}, as well as other comparably high-resolution N-body simulations such as \\citet{raymond:2006,raymond:2007,raymond:2009} (earlier, lower-resolution simulations such as \\citet{chambers2}, \\citet{chambers3}, and \\citet{raymond04} would likely also provide a reasonable fit in terms of bulk chemistry, but are unable to fully match properties of the terrestrial planets such as their masses and number, formation timescale, low dynamical excitation, etc). Furthermore, it also serves to validate the approach utilized here to combine detailed dynamical and chemical modeling together. This will allow for reliable application of this approach not only to other dynamical models but also to other planetary systems \\citep[e.g.][]{bond:thesis}.  The final bulk elemental compositions of the simulated terrestrial planets for rock forming elements are not strongly influenced by the orbital properties or evolution of Jupiter and Saturn as can be seen by the strong similarities between the CJS and EJS simulations. The only significant differences occur in the amount of water rich material accreted onto the final planets and the amount of late veneer material delivered after the last major impact. This suggests that the bulk chemical evolution of the terrestrial planets is to a large extent independent of the evolution of the giant planets. This conclusion, however, is only valid for late stage and \\textit{in situ} formation after the giant planets have formed and undertaken the vast majority of their migration. Simulations are currently running to examine the effects on planetary composition of formation occurring during giant planet migration. As Jupiter and Saturn are not believed to have undergone extensive migration \\citep{nice1,levison,nice2}, the issue of migration is not a significant one for the Solar System. However, it will be an important issue for extrasolar planetary systems where many are thought to have experienced large amounts of migration.  The delivery of water and other volatile species to the final terrestrial planets appears to be a normal outcome of terrestrial planet formation within the CJS simulations, implying that it is normal for such planets to accrete significant amounts of water. As such, it negates the need for large-scale delivery of water and other volatile species by such exotic processes as cometary impacts. The EJS simulations, on the other hand, only produce terrestrial planets with a significant amount of water for disk conditions at times of 1$\\times$10$^{6}$ years and greater. Thus in order to produce terrestrial planets with water accreted during the formation process within the EJS configuration, time-varying compositions need to be included within a single simulation, allowing for an increased hyrdous material content at later times. It should also be noted that the terrestrial planets produced in the CJS simulations are likely to always accrete more water than their EJS companions. Terrestrial planets within the CJS simulations accrete more material from the outer, more water enriched regions of the asteroid belt, thus making them inherently water-rich compared to the terrestrial planets produced in the EJS simulations.  Later planetary processing (especially within the mantle) would have resulted in dissociation of some of the primordial water, causing a significant amount of H to migrate to the core of the Earth or alternatively to be lost from the planet. This migration, however, would have left a large amount of OH in the mantle, thus increasing its redox state. This process allows us to explain the evolution of the Earth's redox state over time from the initially reduced state produced by the current simulations to the present stratified redox state. This conclusion, however, requires a more detailed experimental understanding of the efficiency of H partioning within the mantle and core of the Earth, along with the evolution of mantle processes and mixing.  It is interesting to note the order of magnitude difference between the CJS and EJS simulations in the amount of solid material added to the Sun during terrestrial planet formation. This difference in accreted material is due to stronger resonances produced when Jupiter and Saturn are in more eccentric orbits. These resonances can quickly increase the eccentricity of a body to $\\sim$1, effectively driving them into the Sun and resulting in an increased amount of accreted material. This may potentially have great bearing on extrasolar planetary host stars as many known extrasolar planets are currently in eccentric orbits. If such an orbit always results in a greater amount of solid material being accreted by the host star, then it may give more weight to the pollution hypothesis which has previously been suggested to explain the observed [Fe/H] enrichment (e.g. \\citealt{la,g6,murray}). However, no correlation has been found between known planetary eccentricity and the metallicity of the host star \\citep[e.g.][]{re,s2,fv,bond1}, suggesting that this is not a strong effect. We determined that a negligible amount of solid material was added to the Solar photosphere during terrestrial planet formation and that no observable elemental enrichment would be produced. While similar simulations are needed for extrasolar planetary systems, the current simulations support the conclusion that the observed metal enrichment in extrasolar host stars is primordial in origin, established in the giant molecular cloud from which these systems formed as has been concluded by other studies (such as \\citealt{s1,s2,s04,s05} and \\citealt{fv}). Of course, this does not rule out the possibility that extrasolar planetary host stars may have in fact accreted a giant planet during the formation and migration process and thus display higher levels of pollution.  Finally, biologically important elements are obviously of great interest, especially for the Earth. Of the six major biogenic elements (H, C, N, O, S and P), four are accreted in excess by the planets during their formation (H, O, S and P). Only C and N are not accreted during planetary formation. Both species are primarily present (under the assumption of equilibrium) as solids within the Solar nebula as clathrates (methane or ammonia trapped in a water ice lattice) and organics. As these species only form in the outermost regions of the disk considered in the current simulation (beyond $\\sim$3.6AU for disk conditions at 5$\\times$10$^{5}$ years), they could only be delivered to the final planets via cometary impacts, migration or temporal variations within the disk or some combination of all three. Therefore delivery of material from the outer regions of the disk are necessary in the current models for life to be able to develop. However, it is entirely possible that both C and N could be present at low levels in solid solutions and/or liquids within the actual Solar System. Furthermore, both the CM and CR chondrites are known to contain up to 5 wt.\\% organic C. Such compositions are not captured in our current assumption of equilibrium driven condensation, indicating the presence of another mechanism, acting to incorporate this material into carbonaceous asteroids. As such, small amounts of C and N may be delivered during the normal accretionary process, thus reducing the requirement for external delivery of substantial amounts of material.  We did not perform a comparable chemical analysis for other recent types of terrestrial planet formation simulations, such as N-body simulations incorporating the effects of tidal drag and sweeping secular resonances during gas disk dissipation \\citep{thommes:2008}, or the hybrid N-body/coagulation code of \\citet{bromley:2006} and \\citet{kenyon:2006}.  However, we expect that such simulations would also provide a reasonable fit to the terrestrial planets in terms of bulk chemistry.  As discussed in Sec.~\\ref{sec:results}, the equilibrium chemical composition of the Solar nebula is not highly zoned, and the majority of the solid mass within the disk (at least outside of $\\sim$0.75 AU) has elemental abundances that are basically solar.  While processes that are treated in these other models [eg.~secular resonance sweeping in \\citet{thommes:2008} and gas drag on planetesimals in \\citet{bromley:2006} and \\citet{kenyon:2006}] may lead to a somewhat increased degree of radial migration of material during terrestrial planet formation, that material will still have a fundamentally similar composition throughout the disk, such that the final bulk planet chemistry is unlikely to be significantly different than what we find in this work.  More subtle changes may occur, for example with regards to the amount of hydrated material delivered, or the amount and chemistry of the late veneer material, which would be interesting to explore in future work."
    },
    "0911/0911.4669_arXiv.txt": {
        "abstract": "This paper presents an overview of the radiative transfer problem of calculating the spectral line intensity and polarization that emerges from a (generally magnetized) astrophysical plasma composed of atoms and molecules whose excitation state is significantly influenced by radiative transitions produced by an anisotropic radiation field. The numerical solution of this non-LTE problem of the 2nd kind is facilitating the physical understanding of the second solar spectrum and the exploration of the complex magnetism of the extended solar atmosphere, but much more could be learned if high-sensitivity polarimeters were developed also for the present generation of night-time telescopes. Interestingly, I find that the population ratio between the levels of some resonance line transitions can be efficiently modulated by the inclination of a weak magnetic field when the anisotropy of the incident radiation is significant, something that could provide a new diagnostic tool in astrophysics. ",
        "introduction": "The non-LTE problem of the 1st kind is that described in Mihalas's (1978) book \\citep{Mihalas-1978}: it consists in calculating, at each spatial point of the astrophysical plasma model under consideration, the values of the atomic level populations that are consistent with the intensity of the radiation field generated within the medium. This requires solving jointly the radiative transfer (RT) equations for the specific intensities and the statistical equilibrium equations for the level populations. Nowadays, this non-local and non-linear problem is routinely solved via the application of ``clever'' iterative schemes, where everything goes as simply as in the $\\Lambda$-iteration method, but for which the convergence rate is extremely high \\citep[e.g., the reviews by][]{Hubeny-2003,Trujillo-Bueno-2003}. The success of this international meeting to honor Prof. Dimitri Mihalas on the occasion of his 70th birthday is a clear demonstration that the numerical solution of this multilevel radiative transfer problem has enabled many important advances to be made in 20th century astrophysics, including the modeling of the Fraunhofer spectrum and the development of the field of astrophysical spectroscopy.  In contrast, the non-LTE problem of the 2nd kind is still rather unfamiliar to astrophysicists. This term was introduced by Landi Degl'Innocenti \\citep{Landi-1987} to refer to the formidable numerical problem that implies calculating, at each spatial point of the (generally magnetized) astrophysical plasma model under consideration, the values of the $(2J+1)^2$ elements of the atomic density matrix corresponding to each atomic level of total angular momentum $J$, which quantify its overall population, the population imbalances between its magnetic sublevels, and the quantum coherences between each pair of them. The values of such density-matrix elements have to be consistent with the intensity, polarization and symmetry properties of the radiation field generated within the medium. This requires solving jointly the RT equations for the Stokes parameters and the statistical equilibrium equations for the elements of the atomic density matrix\\footnote{Note that the first Stokes parameter, $I$, is the specific intensity at a given wavelength, while Stokes $Q$ represents  the intensity difference between linear polarization parallel and perpendicular to a given reference direction (e.g., the direction perpendicular to the straight line joining the solar disk center with the observed point). Stokes $U$ would be then the intensity difference between linear polarization at $+45^{\\circ}$ and $-45^{\\circ}$ with respect to the chosen reference direction, where positive angles are measured counterclockwise by an observer facing the radiation. Stokes $V$ is the intensity difference between right-handed and left-handed circular polarization.}. Although much remains to be done, the numerical solution of the non-LTE problem of the 2nd kind is facilitating the physical understanding and modeling of the so-called {\\em second solar spectrum}, as well as the development of the field of astrophysical spectropolarimetry.  As clarified below, the second solar spectrum is the observational signature of radiatively induced population imbalances and quantum coherences in the atoms and molecules of the solar atmosphere. Magnetic fields produce partial decoherence via the Hanle effect, giving rise to fascinating observable effects in the emergent spectral line polarization. Interestingly, these effects allow us to ``see'' magnetic fields to which the Zeeman effect is blind within the limitations of the available instrumentation \\citep[e.g.,][]{Trujillo-Bueno-2004}. In fact, the Zeeman effect is of limited practical interest for the exploration of magnetic fields in hot (e.g., chromospheric and coronal) plasmas because the polarization that it induces in a spectral line scales with the ratio between the Zeeman splitting and the Doppler width, and this ratio turns out to be particularly small in such plasmas. Moreover, the polarization of the Zeeman effect as a diagnostic tool is blind to magnetic fields that are randomly oriented on scales too small to be resolved. We may thus define ``the Sun's hidden magnetism'' as all the magnetic fields of the extended solar atmosphere that are impossible to diagnose via the consideration of the Zeeman effect alone.  \\begin{figure} \\includegraphics[height=.40\\textheight]{Trujillo-Bueno-fig1} \\caption{The visible Fraunhofer spectrum (displayed as the normalized intensity of the continuum) versus the second solar spectrum (displayed as the fractional linear polarization, $Q/I$). The positive direction for the Stokes $Q$ parameter is the parallel to the observed solar limb. Extracted from \\citep{Gandorfer-2000}.} \\end{figure} ",
        "conclusions": "As mentioned in \\S1, we may define {\\em the Sun's hidden magnetism} as all the magnetic fields of the extended solar atmosphere that are impossible to diagnose via the consideration of the Zeeman effect alone. Contrary to what one might think, there are many examples that belong to this category:  \\begin{itemize}  \\item The greater part of the magnetism of the quiet solar photosphere.  \\item The magnetic fields of the solar chromosphere outside sunspots.  \\item The magnetic fields that confine the plasma of solar prominences and filaments.  \\item The magnetism of the solar transition region and corona.  \\end{itemize}  Fortunately, the Hanle and Zeeman effects can be suitably complemented for exploring magnetic fields in solar and stellar physics. To this end, it is necessary to interpret high-sensitivity spectropolarimetric observations through the numerical solution of non-LTE problems of the 2nd kind. The interested reader will find abundant information on a variety of interesting applications in solar physics in the following review papers: \\citep{Trujillo-Bueno-2006,Casini-2007,Trujillo-Bueno-2009a,Trujillo-Bueno-2009b}.  Finally, it is important to emphasize that one of the greatest challenges in astrophysics is the empirical investigation of the magnetic field vector in a variety of astrophysical systems, such as the solar corona, circumstellar envelopes, accreting systems, etc. The physical mechanisms I have considered here (anisotropic radiative pumping, atomic level polarization and the Hanle and Zeeman effects) occur in many astrophysical environments, not only in the solar atmosphere. In particular, the linear polarization signals produced by the presence of atomic level polarization are sensitive to magnetic fields in a parameter domain that ranges from very weak fields (e.g., 1 mG) to a few hundred gauss. The observation and physical interpretation of these polarization effects via the solution of the non-LTE problem of the 2nd kind provides key information, impossible to obtain via spectroscopy or conventional spectropolarimetry.   \\begin{theacknowledgments} Financial support by the Spanish Ministry of Science and Innovation through project AYA2007-63881 and by the European Commission via the SOLAIRE network (MTRN-CT-2006-035484) are gratefully acknowledged. The author is also grateful to the members of the SOC for their invitation to participate in a very good conference. \\end{theacknowledgments}"
    },
    "0911/0911.1750_arXiv.txt": {
        "abstract": "{Icy dust grains play an important role in the formation of complex inter- and circumstellar molecules. Over the past decades, laboratory studies have mainly focussed on the physical interactions and chemical pathways in ices containing rather simple molecules, such as H$_2$O, CO, CO$_2$, CH$_4$ and CH$_3$OH. Observational studies show that polycyclic aromatic hydrocarbons (PAHs) are also abundantly present in the ISM in the gas phase. It is likely that these non-volatile species freeze out onto dust grains as well and participate in the astrochemical solid-state network, but experimental PAH ice studies are largely lacking.} {The study presented here focuses on a rather small PAH, pyrene (C$_{16}$H$_{10}$), and aims to understand and quantify photochemical reactions of PAHs in interstellar ices upon vacuum ultraviolet (VUV) irradiation as a function of astronomically relevant parameters.} {Near UV/VIS spectroscopy is used to track the \\itshape in situ \\upshape VUV driven photochemistry of pyrene containing ices at temperatures ranging from 10 to 125~K.} {The main photoproducts of VUV photolyzed pyrene ices are spectroscopically identified and their band positions are listed for two host ices, \\water and CO. Pyrene ionisation is found to be most efficient in \\water ices at low temperatures. The reaction products, triplet pyrene and the 1-hydro-1-pyrenyl radical are most efficiently formed in higher temperature water ices and in low temperature CO ice. Formation routes and band strength information of the identified species are discussed. Additionally, the oscillator strengths of Py, Py$^{\\cdot+}$ and PyH$^\\cdot$ are derived and a quantitative kinetic analysis is performed by fitting a chemical reaction network to the experimental data. } {Pyrene is efficiently ionised in water ice at temperatures below 50~K. Hydrogenation reactions dominate the chemistry in  low temperature CO ice with trace amounts of water. The results are put in an astrophysical context by determining the importance of PAH ionisation in a molecular cloud. We conclude that the rate of pyrene ionisation in water ice mantles is comparable to the rate of photodesorption of \\water ice. The photoprocessing of a sample PAH in ice described in this manuscript indicates that PAH photoprocessing in the solid state should also be taken into account in astrochemical models. } ",
        "introduction": "Strong infrared emission attributed to polycyclic aromatic hydrocarbons (PAHs) is characteristic of many galactic and extragalactic objects \\citep{smith07,draine07,draine07a,tielens08}. While this emission generally originates from optically thin, diffuse regions, PAHs should also be common throughout the dense interstellar medium. There, as with most other interstellar species in molecular clouds, PAHs condense out of the gas onto cold icy grain mantles where they are expected to influence or participate in the chemistry and physics of the ice. While laboratory studies on interstellar ice analogs have shown that complex organic molecules are produced upon extended vacuum ultraviolet (VUV) photolysis \\citep[e.g.][]{briggs92,bernstein95}, the photoinduced processes at play during the irradiation of PAH containing interstellar ice analogues have not been studied in detail. In optical, {\\itshape in situ} studies on the photochemistry of naphthalene, 4-methylpyrene and quatterylene containing water ice at 20~K, \\citet{gudipati03, gudipati06a,gudipati06} and \\citet{gudipati04} showed that these PAHs are readily ionized and stabilized within the ice, suggesting that trapped ions may play important, but overlooked roles in cosmic ice processes. Beyond this, there is little information on the VUV induced, {\\itshape in situ} photochemistry and photophysics of PAH containing water rich ices.  Here, we describe a detailed study of the VUV induced photochemistry that takes place within pyrene (Py or C$_{16}$H$_{10}$) containing water ices (Py:H$_2$O=1:10,000-1:5,000). The present study is an extension of a recently published study \\citep{bouwman09} in which the focus was on the new experimental setup and where the use of PAH ice spectra was discussed to search for solid-state features of PAHs in space. In this work, the focus is on a detailed characterization of the actual chemical processes taking place upon VUV irradiation, particularly as a function of ice temperature ranging from 25 to 125~K. Additionally, measurements on Py:CO ices at 10~K have been performed to further elucidate the role of water in the reaction schemes and to clarify formation routes of identified species. {\\bf A similar study of three small PAHs is now underway to understand the general principles of PAH/ice photochemistry.} This is part of an overall experimental program at the Sackler Laboratory for Astrophysics to study the fundamental processes of inter- and circumstellar ice analogues such as thermal \\citep{acharyya07} and photodesorption \\citep{oberg07a,oberg08a}, hydogenation reactions \\citep{fuchs08,ioppolo08}, photochemistry \\citep{oberg09a} and physical interactions in interstellar ice analogues \\citep{bouwman07,oberg07,oberg09}.  \\begin{figure*}[t] \\centering \\includegraphics[width=17cm]{f1.eps} \\caption{The spectrum of a dilute pyrene:H$_2$O ice after 900~s of VUV irradiation at 125~K. The inset shows a blow-up of the pyrene photoproduct bands. Band assignments are discussed in Section~\\ref{sectionresults}. Note the broad feature ranging from about 350 to 470~nm which is indicated by a Gaussian fit. This is attributed to overlapping bands from individual pyrene photoproducts. Bands with negative optical depth indicate species destruction, those with positive optical depth show species formation. The blue bands are Gaussian profiles which co-add to the overall fit shown in red. Note the instrumental resolution indicated by the profile of the H$\\beta$ line at 486.1~nm. (This figure is available in color in electronic form)} \\label{FIGURE1} \\end{figure*}  The manuscript is organized as follows. The experimental technique is summarized in Section~\\ref{experimental}. Section~\\ref{sectionresults} describes the Py:H$_2$O and Py:CO ice photochemistry, the resulting products and their formation routes. The temperature dependent photochemistry and derived reaction dynamics are described in Section~\\ref{kineticanalysis} and astrochemical implications are discussed in Section~\\ref{sectionAstrochemical}. The main conclusions are summarized in Section~\\ref{conclusions}. ",
        "conclusions": "\\label{conclusions}  A recently constructed setup has been used to track, on a sub-second timescale, the photochemistry of a PAH in \\water and CO ices as a function of temperature. The setup used here clearly has advantages compared to the relatively slow infrared photochemical ice studies. The conclusions from this work on the PAH pyrene trapped in an interstellar ice analogue are summarized below:  \\begin{enumerate} \\item\tA set of photochemical reaction products has been assigned, both in irradiated Py:\\water and Py:CO ice experiments. The reaction products result from direct photoionisation of pyrene, or from a reaction of the parent, pyrene, with free H atoms produced in the matrix. Additionally, an absorption band is tentatively assigned to a triplet-triplet transition of pyrene. A vibrational progression assigned to HCO$^\\cdot$ is found in spectra of the VUV irradiated Py:CO ice. \\item Pyrene is easily and efficiently ionised when trapped in \\water ice. Photoionisation is a non-diffusion related reaction and hence a photon-rate of $1.2\\times10^{-17}$~cm$^2$~photon$^{-1}$, which can serve as input for astrochemical models, is derived. \\item When trapped in CO ice, pyrene ionisation is inefficient compared to water ice. \\item\tElectron-ion recombination is independent of ice temperature and is characterized as a non diffusion dominated reaction. For this process, a photon rate of $9\\times10^{-18}$~cm$^2$~photon$^{-1}$ is derived which can be directly used in astrochemical models. \\item There are two distinct reaction paths in the photochemistry of pyrene trapped in \\water ice. At low temperatures ($<50$~K) the chemistry is dominated by ion-molecule interactions and processes. At temperatures above 50~K reactions are dominated by diffusing radical species. \\item A simple model indicates that, in {\\bf dense clouds where $A_V=3$}, the rate of pyrene ionisation is comparable to the rate of photodesorption in water rich ices. Hence, chemical reactions involving pyrene, and its cation can be important in modeling grain chemistry in such environments. \\end{enumerate}"
    },
    "0911/0911.3049_arXiv.txt": {
        "abstract": "The mechanism by which outflows and plausible jets are driven from black hole systems, still remains observationally elusive. Notwithstanding, several observational evidences and deeper theoretical insights reveal that accretion and outflow/jet are strongly correlated. We model an advective disk-outflow coupled dynamics, incorporating explicitly the vertical flux. Inter-connecting dynamics of outflow and accretion essentially upholds the conservation laws. We investigate the properties of the disk-outflow surface and its strong dependence on the rotation parameter of the black hole. The energetics of the disk-outflow strongly depend on the mass, accretion rate and spin of the black holes. The model clearly shows that the outflow power extracted from the disk increases strongly with the spin of the black hole, inferring that the power of the observed astrophysical jets has a proportional correspondence with the spin of the central object. In case of blazars (BL Lacs and Flat Spectrum Radio Quasars), most of their emission are believed to be originated from their jets. It is observed that BL Lacs are relatively low luminous than Flat Spectrum Radio Quasars (FSRQs). The luminosity might be linked to the power of the jet, which in turn reflects that the nuclear regions of the BL Lac objects have a relatively low spinning black hole compared to that in the case of FSRQ. If the extreme gravity is the source to power strong outflows and jets, then the spin of the black hole, perhaps, might be the fundamental parameter to account for the observed astrophysical processes in an accretion powered system. ",
        "introduction": "High resolution observations show strong outflows and jets in black hole accreting systems, both in active galactic nuclei (AGNs) or quasars (Begelman et al. 1984; Mirabel 2003) and microquasars (SS433, GRS~1915+105) [Margon 1984; Mirabel \\& Rodriguez 1994, 1998]. Extragalactic radio sources show evidence of strong jets in the vicinity of spinning black holes (Meier et al. 2001; Meier 2002). Outflows are also observed in neutron star low mass X-ray binaries (LMXBs) (Fender et al. 2004; Migliari \\& Fender, 2006) and also in young stellar objects (Mundt 1985). It has been argued (Ghosh \\& Mukhopadhyay 2009; Ghosh et al. 2010; hereinafter GM09, G10 respectively, and references therein) that outflows and jets are more prone to emanate from strong advective accretion flows; the said paradigm is more susceptible for super-Eddington and sub-Eddington accretion flows. However, the exact (global if any) mechanism of formation of the jet, its collimation, acceleration, composition from the accretion powered systems still remain inconclusive.  Extensive works have been pursued on the origin of outflow/jet, since the pioneering work of Blandford \\& Payne (1982) (e.g. Pudritz \\& Norman 1986; Contopoulos 1995; Ostriker 1997), where the authors used a self-similar approach to demonstrate that the poloidal component of the magnetic field can be seemingly used to launch outflowing matter from the disk. The formation of the jet is directly related to the efficacy of extraction of angular momentum and energy from the accretion disk. Physical understanding of the hydromagnetic outflows from disks has been developed from magnetohydrodynamic (MHD) simulations (Shibata \\& Uchida 1986; Ustyogova et al. 1999; De Villiers et al. 2005; Hawley \\& Krolik 2006) both in non-relativistic as well as in relativistic regimes, mostly in the Keplerian paradigm. However, the large Lorentz factors as well as Faraday rotation measures suggest that the observed VLBI jets in quasars and active galaxies are in the Poynting flux regime (Homan et al. 2001; Lyutikov 2006). On the other front, strong radiation pressure can serve as a different mechanism to effuse outflows/jets. This is likely to occur when the accretion rate is super-Eddington or super-critical (Fabrika 2004; Begelman et al. 2006; GM09,G10 and references therein) and the accretion disk is precisely ``radiation trapped'' as in ultra-luminous X-ray (ULX) sources. It has been confirmed by radiation hydrodynamic simulation at super-critical accretion rate (Okuda et al. 2009) as well to explain luminosity and mass outflow rate of relativistic outflows from SS433. The under-luminous accreting sources, having high sub-critical accretion flow, were explained by an advection dominated accretion flow (ADAF) model (Narayan \\& Yi 1994). The promising outcome of this model lies in the large positive value of the Bernoulli parameter because of the small radiative energy loss. It leads to conceive that the gas in the inflowing disk is susceptible to escape, leading to strong unbounded flows in the form of outflows and jets. This also signifies that even in the absence of magnetic field and radiation pressure, outflows are plausible from strongly advective accretion flow if the system is allowed to perturb. Nevertheless, the definitive understanding of the origin of outflows/jets is sill unknown. It remains one of the most compelling problems in high energy astrophysics.  Whatever might be the reason for the origin of outflow and then jet from the disk, one aspect is however definite that strong outflows producing relativistic jets are powered by extreme gravity. Although seems paradoxical, the strength and the length scale of observed astrophysical jets vary directly with the strength of the central gravitating potential. Observationally, it is evident that strong outflows and relativistic jets are more powerful in observed AGNs and quasars, harboring supermassive black holes, compared to that in black hole X-ray binaries (XRBs). In addition, the jets observed in AGNs and quasars have greater length scale compared to that seen in stellar mass black hole systems. One of the most important signatures of relativistic gravitation is the spin of the central object. It is presumably believed that the spin (practically specific angular momentum) of the neutron star is less than that of a black hole, and thus the observed jet from black holes is much stronger and powerful than that of neutron star sources. In early, Blandford \\& Znajek (1977) demonstrated that if there is a magnetic field associated with the black hole due to threading of magnetic field lines from the disk and the angular momentum of the Kerr black hole is large enough, then the energy and the angular momentum can be extracted from the underlying black hole by a purely electromagnetic mechanism, which can thus be expected to power the jet in an AGN. These imply that the spin might play a significant role in powering jets, both in microquasars and in AGNs, rather, it can act as a fundamental parameter in an accretion powered system.  Most of the studies of accretion disk and studies of related outflow/jet have been evolved separately, assuming these two apparently to be dissimilar objects. However, several observational inferences (for details see GM09, G10 and references therein) and improved understanding of accretion flow and outflow reveal that accretion and outflow/jet are strongly correlated. The unifying scheme of disk and outflow is essentially governed by conservation laws; conservations of matter, energy and momentum. Hence, in modeling the accretion and outflow simultaneously in any accretion powered system, following aspects should be taken into account: (1) the effect of relativistic central gravitational potential including its spin, (2) proper mechanism of origin of outflow/jet from the disk, (3) appropriate hydrodynamic equations (considering that the accreting gas be treated as a continuum fluid), capturing the information about the intrinsic coupling between inflow and outflow which are governed by the conservation laws in a strongly advective paradigm.  Recently, GM09, G10 made an endeavor to explore the dynamics of the accretion-induced outflow around black holes/compact objects in a 2.5-dimensional paradigm. The authors formulated the disk-outflow coupled model in a more self-consistent way by solving a complete set of coupled partial differential hydrodynamic equations in a general advective regime through a self-similar approach in an axisymmetric, cylindrical coordinate system. They explicitly incorporated the information of the vertical flux in their model. However, they restricted their study to Newtonian regime, thus negating the requisite effect of general-relativity, especially the effect of spin of the central object. Based on the model, the authors computed the mass outflow rate and the power extracted by the outflow from the disk self-consistently, without proposing any prior relation between the inflow and outflow.  In the present paper, we propose a new model for the accretion-induced outflow by extending the work of GM09, G10, by incorporating the general relativistic effect of the central potential without limiting ourselves to a self-similar regime. As the definitive mechanism of launching of outflows/jets is still evasive, we do not embrace any specific mechanism of outflow like inclusion of the magnetic field into our model equations. Nevertheless, the importance of the magnetic field can not be, in principle, discarded to explain the launching and collimation of jets off the accretion disk. Here, owing to our inability to solve coupled partial differential MHD equations in a 2.5-dimensional advective regime, we neglect the influence of the magnetic field in our model equations. However, the implicit coupling between the inflow and outflow, dictated by the conservation equations, have been taken into account appropriately. We arrange our paper as follows. In the next section, we present the formulation of our model. The section 3 describes the computational procedure to solve the model equations of the accretion-induced outflow. In sections 4 and 5, we study the dynamics and the energetics of the flow respectively. Finally, we end up in section 6 with a summary and discussion. ",
        "conclusions": "Blandford-Znajek process  (Blandford \\& Znajek 1977) is still one of the most promising mechanisms to drive powerful jets in AGNs and XRBs. Although the exact mechanism of formation of the jet in the vicinity of the black holes is still elusive, and whatever might be the reason for the origin of jet, the said work is significant mostly due to two underlying reasons: (1) extreme gravity is indispensable to effuse strong unbounded flows in the vertical direction from the inner region of the accretion disk, (2) the spin of the black holes, which is purely a relativistic effect, is essential to power strong outflows and jets.  In the present study, we have neither laid importance to the nature of the outflow nor invoked any aspect for the origin of outflows or jets. Indeed, the distinctive or definitive understanding of the formation of strong outflows or jets is till unknown. Notwithstanding, the most distinguishable and obvious picture in this case is that outflows and jets observed in AGNs and XRBs can only originate in an accretion powered system, where the accretion of matter around relativistic gravitating objects like black holes acts as a source, and outflow and then jet takes the form of one of the possible sinks (the other sink is the central nucleus). The dynamics of the outflowing matter should then be intrinsically coupled to the accretion dynamics macroscopically through the fundamental laws of conservation (of matter, momentum and energy). The outflow is unbounded and the total energy just at the base of the outflow should be positive. The accretion flow should be bounded in the vicinity of the central object as the central potential is attractive. However, the strength of the unbounded flows in the form of jets is distinctly proportional to the attractiveness of the central gravitational potential field. This paradox is well manifested in the observed universe, as relativistic jets are more populated around extreme gravitating objects like black holes. Also noticeably, length scale of jets increases from microquasars to quasars, which are supposedly harboring stellar mass and supermassive black holes respectively.  Thus in any theoretical modeling of the accretion and outflow, it can be presumably argued that the mathematical equations governing the dynamics of the inflow and outflow should inherently be correlated and evolved self-consistently without any ad hoc proposition, as accretion and outflow should not be treated as dissimilar objects. Second, the relativistic gravitational effect of the black hole should be incorporated, as the nature of gravity is the cornerstone to both the accretion and the unbounded outflow. To capture this essential physics, in our present study to understand the connection between disk and outflow, we have incorporated the general relativistic effect of the spinning black hole through a pseudo-Newtonian approach. Although pseudo-Newtonian formalism is an approximate method to mimic the space-time geometry of the Kerr black hole, yet it captures the important salient features of the corresponding metric, and thus can be used to examine the nature of outflows from the inner region of the disk.  In \\S 2 we have described the general disk-outflow coupled hydrodynamic equations in the inviscid limit  following GM09. The necessity to simplify our model equations for an inviscid flow is discussed in \\S 2. Although GM09 explored the 2.5-dimensional accretion-induced outflow for a fully viscous system, they used a self-similar approach in order to solve the necessary partial coupled differential equations. In addition, they neglected the most indispensable effect of relativistic gravitation or,  precisely, the effect of spin of the black hole. In the present paper, we have solved the disk-outflow model equations in a more general 2.5-dimensional paradigm, while trying to limit our assumptions to the least possible extent theoretically. One of the important premises we have made is the relationship between $v_z$ and $c_s$,  which we have established empirically. We have not used the height integrated equations which are mostly valid in the circumstances where the dynamical fluid parameters are likely to be independent of $z$. Without presuming the fact that the outflow originates from the surface of the disk, as most of the authors do, we have logically constructed a disk-outflow surface with properly defined boundary conditions. One of the most important computations we have done is to evaluate the mass outflow rate and the power of the outflow extracted from the disk self-consistently in the inviscid limit, unlike the previous works (e.g. Donea \\& Biermann 1996; Blandford \\& Begelman 1999; Xie \\& Yuan 2008). We have found that the spin of the black hole plays a crucial role in determining the structure, dynamics and the energetics of the outflow coupled to the disk. With the increase of the spin, the outflow region extends further inward and then the disk-outflow region shrinks and compresses (see Fig. 2). As a result, the outflow and then any plausible jet is likely to eject out from an extreme inner region of the flow around a rapidly spinning black hole with a greater efficiency. Note that the higher spin results in the system to get more compressed with a greater outflow power. Therefore, the efficiency of outflow and jet is directly related to the disk scale-height and hence the disk-outflow surface. Previously, in a different context, Mckinney \\& Gammie (2004), while examining the electromagnetic luminosity of a Kerr black hole, assumed the ratio of the disk height to the radius ($h/r$) constant, irrespective of the black hole spin. However, as seen in the present work, the constant $h/r$ for different spin of the  black hole may not be an obvious choice.  The power extracted by the outflow from the disk not only depends directly on the mass of the black hole and the initial mass accretion rate of the flow, but also on the spin of the black hole. With our model, keeping the  black hole mass and the accretion rate the same, the power of the outflow increases with the spin of the black hole and the computed power differs in two orders of magnitude between non-rotating and maximally rotating black holes. We have restricted our study vertically up to the region where the inflow and the outflow are least coupled, i.e. the disk-outflow surface. Above this surface accretion ceased to exist and probably the outflow gets decoupled from the disk, accelerates and eventually forms relativistic jet. The modeling of the astrophysical jet is altogether a different issue and is beyond the scope of the present work. Nevertheless, it can be effectively argued that in modeling the dynamics of the jet, the computed outflow power may serve as an initial power fed to the jet. Thus, the dynamics and the energetics of the jet will eventually be related to the black hole spin.  According to the unification scenario, it is possible to device a single astrophysical scenario, which can broadly explain the observed multitude of different types of AGNs (see e.g. Antonucci 1993; Urry \\& Padovani 1995). Flat Spectrum Radio Quasars (FSRQs) and BL Lacs are probably the two most active types of AGNs which are collectively referred to blazar. It is observed that BL Lacs are relatively low luminous than FSRQs. Although both of these galaxies show high energy emissions, their spectral properties are different (Bhattacharya et al. 2009) which indicate that they are different source classes. According to the unification scheme, for FSRQs the line-of-sight is almost parallel to the jet and, hence, strong relativistic Doppler beaming of the jet emission produces highly variable and continuum dominated emission. As one moves away from the jet axis, the central continuum emission falls and the nucleus looks like a FR-II galaxy. Similarly, BL Lacs are considered to be a sub class of FR-I galaxies whose line-of-sight is almost parallel to the jet axis. The present work suggests that the total mechanical power of an outflow proportionately increases with the spin of the central supermassive black hole; higher the spin, stronger is the outflow. It is reasonable to consider that strong outflow can lead to a strong jet, and hence, one can expect to observe higher luminosity. Therefore, the work suggests that BL Lacs are slow rotators than FSRQs.  For the theoretical formulation of the disk-outflow coupling even with approximations, one needs to be very thoughtful for the proper foundation of the model which essentially needs to solve hydrodynamic or magnetohydrodynamic conservation equations in presence of strong gravity. The flow parameters vary in 2.5-dimension and are coupled to each other. Limited observational inputs put irremediable constraint on the boundary conditions as well as the fundamental scaling parameters, governing the coupled dynamics of the accretion and outflow. The inadequacy of an effective mathematical tool to handle partial coupled differential hydrodynamic equations for compressible flow motivates us to invoke approximations and assumptions. Despite of this fact, one needs to explore the possibilities to examine the accretion-induced outflow. The question then arises: what are the valid assumptions and to which extent they can be considered? In the present work, we have addressed this question in the following way.\\\\ (1) As the extreme gravity is the most important aspect to effuse jet from the disk, we have incorporated its effect through a pseudo-general-relativistic potential.\\\\ (2) The disk and outflow should not be treated as dissimilar objects, and hence their correlated-dynamics should be essentially governed by the conservation laws. The energetics of the accretion-induced outflow would then be evaluated self-consistently as is shown in our work. \\\\ (3) Any unbounded flow in the form of outflow is more plausible to emanate from a hot, puffed up region of the accretion flow (may be low/hard state of the black hole). Hence we have formulated our model in a 2.5-dimensional, strongly advective paradigm, surpassing the simplicity of height integration. \\\\ (4) We have approximated our system to an inviscid limit, whose reason has been argued in \\S 2. Nevertheless, in any future work, viscosity should be incorporated into the flow to make it more realistic. \\\\ (5) We do not account for the mechanism of formation of strong jets, like due to magnetic field or strong radiation pressure, as the definitive mechanism of the formation of jet is still unknown. Indeed, it is beyond the scope to solve magnetohydrodynamic equations in the present scenario. However, assuming that the outflow resides, and is coupled to the disk, the governing conservation equations of matter, momentum and energy should be treated accordingly. \\\\ (6) Power of the outflow and jet is expected to directly depend on the spin of the black hole. The spin, which is the signature of general relativity, can make an impact on the nature of the observed AGN classes, and thus the spin of the black has been incorporated in our study. As spin of the black hole has a direct impact on the flow parameters, we obtain different outflow power for different spin.  What can be the most defined parameter for the accretion powered system --- mass accretion rate, mass of the central star/black hole or the spin of the central star/black hole? Our analysis has shown that the power extracted by the outflow from the disk proportionately increases with the spin of the black hole. It infers that if extreme gravity is essential to power the jet, then the strength and the length-scale of the observed astrophysical jets directly depend on the spin of the black hole. If it is believed that the distant quasars harbor massive black holes of same mass scale and accrete matter with similar rate, perhaps, spin be the guiding parameter for the different observed AGN classes. We can end with the specific question: are the observed AGN classes astrophysical laboratories to measure the spin of the supermassive black holes? \\\\   This work is supported by a project, Grant No. SR/S2HEP12/2007, funded by DST, India. The authors would like to thank the referee for making important comments which helped to prepare the final version of the paper. \\\\  \\noindent{\\bf APPENDIX}  Equation (\\ref{5f}) consists of complicated terms, where ${\\cal A}, {\\cal B}$  and ${\\cal C}$ are given by following equations $$ {\\cal A} \\, = \\, 4 n+2, \\eqno(A1) $$ $$ {\\cal B} = \\biggl(F_{Grc}-\\frac{\\lambda_{c}^{2}}{r_{c}^3}\\biggr) \\biggl(\\frac{1}{2 n}+1\\biggr) \\frac{1}{v_{rc}} + \\frac{v_{rc}}{r_c} \\biggl(1-\\frac{1}{2n}\\biggr) $$ $$ -\\frac{\\imath}{z} \\biggl(\\frac{z}{r_c}\\biggr)^{\\mu} v_{r0c} \\biggl(\\frac{1}{n}+1\\biggr) + 2 \\frac{v_{rc}}{r_c} + 2 \\frac{\\imath}{z}\\biggl(\\frac{z}{r_c}\\biggr)^{\\mu} v_{rc}, \\eqno(A2) $$ and $$ {\\cal C} = \\biggl(F_{Grc}-\\frac{\\lambda_{c}^{2}}{r_{c}^3}\\biggr) \\frac{1}{2 nr_c} + \\frac{v^{2}_{rc}}{2 n r_c^2} - \\frac{\\imath}{2 n z} \\biggl(\\frac{z}{r_c}\\biggr)^{\\mu} v_{r0c} v_{rc} \\biggl[\\frac{1}{r_c} + $$ $$ \\frac{\\imath}{z} \\biggl(\\frac{z}{r_c}\\biggr)^{\\mu}\\biggr] + \\frac{1}{2n} \\biggl(\\frac{\\partial}{\\partial r} F_{Gr} \\bigg |_{c} +3 \\frac{\\lambda_{c}^{2}}{r_{c}^4}\\biggr) + \\frac{v^{2}_{rc}}{2 n r_c^2} +\\mu \\frac{\\imath}{2 n z} \\, \\frac{z^{\\mu}}{r^{\\mu +1}_c}  \\, v_{r0c} v_{rc}  \\, $$ $$ + \\frac{\\imath}{2 n z} \\biggl(\\frac{z}{r_c}\\biggr)^{\\mu} v_{rc} \\biggl[\\frac{2 n v_{r0c}}{c_{s0c}} \\,   \\biggl(\\frac{-{\\cal B}_{0}+\\sqrt{{\\cal B}_{0}^2-4{\\cal A}_{0}{\\cal C}_{0}}}{2{\\cal A}_{0}}\\biggr)  +\\frac{v_{r0c}}{r_c}\\biggr], \\eqno(A3) $$ where, \\\\ $${\\cal A}_{0} \\, = \\, 2 + 4 n, \\eqno(A4)$$ \\\\ $${\\cal B}_{0}=\\bigl(F_{Gr0c}-\\frac{\\lambda_{c}^{2}}{r_{c}^3}\\bigr) \\bigl(\\frac{1}{2 n}+1\\bigr) \\frac{1}{v_{r0c}} + \\frac{v_{r0c}}{r_c} \\bigl(1-\\frac{1}{2n}\\bigr) + 2 \\frac{v_{r0c}}{r_c}, \\eqno(A5)$$ \\\\ and \\\\ $${\\cal C}_{0}=\\bigl(F_{Gr0c}-\\frac{\\lambda_{c}^{2}}{r_{c}^3}\\bigr) \\frac{1}{2 n r_c} + \\frac{v^{2}_{r0c}}{n r^2_c} + \\frac{1}{2n} \\bigl(\\frac{\\partial}{\\partial r} F_{Gr0} \\big |_{c} +3 \\frac{\\lambda_{c}^{2}}{r_{c}^4}\\bigr). \\eqno(A6)$$"
    },
    "0911/0911.5291_arXiv.txt": {
        "abstract": "We present a new upper limit of $\\sum m_{\\nu} \\le 0.28$ ($95\\%$ CL) on the sum of the neutrino masses assuming a flat $\\mathrm{\\Lambda CDM}$ cosmology. This relaxes slightly to $\\sum m_{\\nu} \\le 0.34$ and $\\sum m_{\\nu} \\le 0.47$ when quasi non-linear scales are removed and $w \\ne -1$, respectively. These  bounds are derived from a new photometric redshift catalogue of over 700,000 Luminous Red Galaxies (MegaZ DR7) with a volume of 3.3 (Gpc $h^{-1}$)$^3$, extending over the redshift range $0.45 < z < 0.65$ and up to angular scales of $\\ell_{\\mathrm{max}} = 300$. The data are combined with WMAP 5-year CMB fluctuations,  Baryon Acoustic Oscillations (BAO), type 1a Supernovae (SNe) and an HST prior on the Hubble parameter. This is the first combined constraint from a photometric redshift catalogue with other cosmological probes. When combined with WMAP this data set proves to be as constraining as the addition of all SNe and BAO data available to date. The upper limit is one of the tightest and `cleanest' constraints on the neutrino mass from cosmology or particle physics. Furthermore, if the aforementioned bounds hold, they all predict that current-to-next generation neutrino experiments, such as KATRIN, are unlikely to obtain a detection. ",
        "introduction": " ",
        "conclusions": "\\end{table} {\\it Conclusions} -- Using the biggest ever large scale structure survey we have set bounds on the neutrino masses at $\\sum m_{\\nu} < 0.28$ eV ($\\ell_{\\mathrm{max}}=300$) and $\\sum m_{\\nu} < 0.34$ eV ($\\ell_{\\mathrm{max}}=200$) at $95\\%$ CL, when combined with WMAP5+SNe+BAO+HST data. This is the first ever determination of neutrino masses from a photometric galaxy redshift survey. Not only have we shown that photometric redshifts can be used for this problem, but also that such a galaxy survey is competitive with all currently available geometric probes (SNe+BAO) or spectroscopic clustering when added to the CMB. Our constraint is one of the tightest current bounds available without the use of data from Lyman-$\\alpha$ (e.g. \\citep{Seljak06}), which is prone to systematics, or a complicated modelling of the bias \\citep{deBernardis08}. Further, all our results show that KATRIN's \\citep{Wolf08} projected $90\\%$ sensitivity ($\\sum m_{\\nu}$ $\\lesssim 0.6$ eV) leaves an unlikely neutrino mass detection. Finally, our overall method shows great promise for the next generation of surveys, which will yield upper limits of e.g. $0.12$ eV \\citep{Lahav09} and $0.025$ eV \\citep{Abdalla07} at $95\\%$ C.L.  \\begin{figure} \\includegraphics[width=8.2cm,height=8.0cm]{./figures/3.pdf} \\caption{\\small{The 2D $68\\%$ and $95\\%$ contours and marginalised 1D distributions for 7 cosmological parameters ($\\Omega_{b} h^{2}$, $\\Omega_{c} h^{2}$, $\\Omega_{\\Lambda}$, $n_{s}$, $\\tau$, $\\mathrm{ln} (10^{10} A_{s})$ and $\\sum m_{\\nu}$) in a WMAP5 + SNe + BAO + MegaZ DR7 + HST combined constraint (red contours). The amplitude of the Sunyaev-Zeldovich fluctuations ($A_{SZ}$) and four bias parameters ($b_{1}$, $b_{2}$, $b_{3}$, $b_{4}$) are marginalised over but not plotted. The black contours are given for the WMAP-only analysis. }} \\label{fig:cmb_sn_megaz_neutrino_dr7} \\end{figure} {\\it Acknowledgements} -- ST acknowledges a STFC studentship and UCL's Institute of Origins for a Post-doctoral Fellowship. FBA and OL acknowledge the support of the Royal Society via a Royal Society URF and a Royal Society Wolfson Research Merit Award, respectively. OL and FBA acknowledge the Weizmann Institute of Science for an Erna and Jakob visiting professorship and support for a research visit, respectively."
    },
    "0911/0911.0682_arXiv.txt": {
        "abstract": "\\noindent A solenoid oscillating in vacuum will pair produce charged particles due to the Aharonov-Bohm (AB) interaction. We calculate the radiation pattern and power emitted for charged scalar particles. We extend the solenoid analysis to cosmic strings, and find enhanced radiation from cusps and kinks on loops. We argue by analogy with the electromagnetic AB interaction that cosmic strings should emit photons due to the gravitational AB interaction of fields in the conical spacetime of a cosmic string. We calculate the emission from a kink and find that it is of similar order as emission from a cusp, but kinks are vastly more numerous than cusps and may provide a more interesting observational signature. ",
        "introduction": "\\label{gaugepot}  The first step in calculating AB radiation is to find the gauge potential for a moving solenoid. This step was taken in Alford and Wilczek \\cite{Alford:1988sj}.  A moving solenoid with magnetic flux $\\Phi$ will be a current source in Maxwell's equation, and the current can be written as \\begin{equation} J^{(\\Phi )}_\\nu = \\frac{\\Phi}{2} \\epsilon_{\\mu\\nu\\alpha\\beta} \\partial^\\mu S^{\\alpha\\beta} \\end{equation} where \\begin{eqnarray} S^{\\alpha\\beta}(x) &=& \\int d\\tau d\\sigma \\sqrt{-\\gamma} \\epsilon^{ab} ~ \\partial_a  X^\\alpha \\partial_b X^\\beta \\delta^{(4)} (x-X(\\sigma,\\tau)) \\nonumber \\\\ && \\hskip -1 cm = \\int d\\tau d\\sigma ( {\\dot X}^\\alpha {X^\\beta} ' - {\\dot X}^\\beta {X^\\alpha} ' ) \\delta^{(4)} (x-X(\\sigma,\\tau)) \\end{eqnarray} where $\\sigma , \\tau$ are world-sheet coordinates, $X^\\mu (\\sigma)$ denotes the position of the string, overdots and primes denote derivatives with respect to $\\tau$ and $\\sigma$ respectively, and $\\gamma_{ab}$ is the worldsheet metric, with $\\gamma$ denoting its determinant. In what follows, we shall use $\\tau =t$, while $\\sigma =z$ for a solenoid along the $z-$axis. In applications to cosmic strings, we will also use the gauge conditions \\begin{equation} {\\dot X} \\cdot X' = 0 \\ , \\ \\ {\\dot X}^2 + {X'}^2 =0 \\label{gaugeconditions} \\end{equation}  The field strength of a static, thin solenoid satisfies \\begin{equation} \\partial^\\mu F_{\\mu\\nu} = J_\\nu \\end{equation} and the solution is \\begin{equation} F_{\\mu\\nu} = \\frac{\\Phi}{2} \\epsilon_{\\mu\\nu\\alpha\\beta} S^{\\alpha\\beta} \\label{solenoidF} \\end{equation} The corresponding gauge potential can be written in Lorenz gauge \\begin{equation} A_\\nu = \\frac{\\Phi}{2} \\epsilon_{\\mu\\nu\\alpha\\beta} \\partial^\\mu \\frac{1}{\\partial^2} S^{\\alpha\\beta} \\end{equation} where $\\partial^2$ is the D'Alembertian operator. We will only need the gauge potential in momentum space, \\begin{equation} {\\tilde A}_\\nu = - i \\frac{\\Phi}{2} \\epsilon_{\\mu\\nu\\alpha\\beta} \\frac{k^\\mu}{k^2} {\\tilde S}^{\\alpha\\beta} \\end{equation} where overtilde's denote the Fourier transformed variable \\begin{equation} {\\tilde S}^{\\alpha\\beta} (k) = \\int d^4x ~ e^{+ik\\cdot x} S^{\\alpha\\beta}(x) \\ . \\end{equation}  Note from Eq.~(\\ref{solenoidF}) that the field strength vanishes everywhere except at the location of the solenoid and so a moving solenoid does not radiate classical electromagnetic waves. ",
        "conclusions": "\\label{conclusions}  We have studied the problem of an oscillating solenoid in vacuum and determined the rate of scalar particle creation due to the Aharonov-Bohm interaction. This is novel because the electromagnetic fields of a moving solenoid vanish everywhere except at the location of the solenoid. In particular, a moving solenoid does not emit electromagnetic radiation but does emit charged particle pairs as AB radiation. We have also determined the pattern of AB radiation and find that it is pre-dominantly in the direction perpendicular to the plane of oscillation.  Cosmic string loops that have AB interactions with quantum fields will also emit AB radiation. In the thin string limit, loops with cusps and kinks emit a linearly divergent power. Imposing a cutoff due to the thickness of the string indicates that the AB radiation is significantly stronger than the naive dimensional estimate.  We point out that all quantum fields interact with cosmic strings via the gravitational AB effect, and thus there is corresponding gravitational AB radiation of all particles, including photons. Cosmic strings, even if occurring in dark matter sectors of particle physics, will emit photons. We have calculated the energy emitted in photons from loops with kinks and cusps. As in the gauge case, the radiated energy is much larger than the naive dimensional estimate. We have made a rough estimate of the energy emitted in photons and it remains to be seen if the signal is strong enough to make it an interesting observational tool."
    },
    "0911/0911.2478_arXiv.txt": {
        "abstract": "We report the discovery of an unusual source of extended X-ray emission CXOU J184846.3--013040 (`The Stem') located on the outskirts of the globular cluster GLIMPSE-C01. No point-like source falls within the extended emission which has an X-ray luminosity $L_{X}$(0.3--8 keV) $\\sim 10^{32}$ ergs s$^{-1}$ and a physical size of $\\sim 0.1$ pc at the inferred distance to the cluster. These X-ray properties are consistent with the pulsar wind nebula (PWN) of an unseen pulsar located within the 95-percent confidence error contour of unidentified \\fermi\\ $\\gamma$-ray source 0FGL J1848.6--0138. However, we cannot exclude an alternative interpretation that postulates X-ray emission associated with a bow shock produced from the interaction of the globular cluster and interstellar gas in the Galactic plane. Analysis of the X-ray data reveals that `The Stem' is most significant in the 2--5 keV band, which suggests that the emission may be dominated by non-thermal bremsstrahlung from suprathermal electrons at the bow shock. If the bow shock interpretation is correct, these observations would provide compelling evidence that GLIMPSE-C01 is shedding its intracluster gas during a galactic passage. Such a direct detection of gas stripping would help clarify a crucial step in the evolutionary history of globular clusters. Intriguingly, the data may also accommodate a new type of X-ray source. ",
        "introduction": "For decades, the Galactic plane has been a source of amazement and unexpected X-ray discoveries \\citep{giacconi,giacconi2,tanaka,voges}. The current generation of X-ray satellites \\cha\\ and {\\it XMM-Newton} are the latest to extend the census of X-ray sources along the Galactic plane to unprecedented depths \\citep{motch,champlane}. Together these experiments have revealed a rich tapestry of X-ray binaries, neutron stars, pulsar wind nebulae (PWNe), supernova remnants, coronal emitting stars, and non-thermal filaments along the Galactic plane. However, for every physical object identified with high confidence, the literature is plagued by an equal or greater amount of X-ray sources that continue to elude firm identifications \\citep{wang,muno}. Frustratingly, a combination of high extinction, excessive crowding, and the lack of direct distance indicators prevent further progress in the identification of many of these sources.  During a routine multiwavelength survey of $\\gamma$-ray error contours produced by the \\fermi\\ mission \\citep{mirabal}, we have encountered the latest puzzling source connected with the Galactic plane. The object CXOU J184846.3--013040 -- which we nickname `The Stem' -- corresponds to a source of extended emission in the outskirts of GLIMPSE-C01. The low-latitude Globular Cluster GLIMPSE-C01 is located in the Galactic plane at $(\\ell, b)=(31.\\!^{\\circ}3,-0.\\!{^\\circ}1)$ and was discovered during Galactic plane scans conducted by the {\\it Spitzer Space Observatory} \\citep{kobul} and the Two Micron All Sky Survey  \\citep[2MASS;][]{simpson}. Subsequently, superb observations with the {\\it Chandra X-ray Observatory} revealed a population of X-ray point-like sources with luminosities $L_{X} > 6 \\times 10^{31}$ ergs s$^{-1}$ around the core of the cluster \\citep{pooley}.  \\begin{figure*} \\hfil \\includegraphics[width=2.1in,angle=0.]{f1a.eps} \\includegraphics[width=2.1in,angle=0.]{f1b.eps} \\includegraphics[width=2.1in,angle=0.]{f1c.eps} \\hfil \\caption{{\\it Left panel:} 0.3--2 keV \\cha\\ image. {\\it Middle panel:}  2--5 keV \\cha\\ image. {\\it Right panel:}  5--8 keV \\cha\\ image overlain with the \\fermi\\ error contour ({\\it dashed line}) derived by \\citet{abdo1}. The X-ray images have been smoothed with a Gaussian kernel with radius $r_{k}$ = 2.5\\arcsec. The superposed arrows mark the location of the `The Stem'; as well as the  extended trail emission designated X2.  The field size is $\\sim$ 3.5\\arcmin $\\times$ 3.5\\arcmin shown in Galactic coordinates. North Galactic Pole is in the direction of the top of the figure and longitude increases toward the left. } \\label{figure1} \\end{figure*}   `The Stem' is an extended source of emission placed well outside the cluster centre and is remarkable for two completely different reasons. First, it lies within the 95-percent confidence error contour of unidentified \\fermi\\ source 0FGL J1848.6-0138 \\citep{abdo1,luque}. If powered by a compact object, `The Stem' would constitute a formidable counterpart candidate for the $\\gamma$-ray emitter. The second reason, unrelated to the $\\gamma$-ray emission, is the fact that the position of GLIMPSE-C01 along the Galactic plane indulges us in that rare geometrical arrangement where a Globular cluster is caught in the act of crossing the Galactic plane. This is the precise moment in which the cluster is believed to interact with the interstellar medium (ISM) and proceed to shed some of the intracluster medium accumulated from the continuous mass loss of individual stars within the cluster \\citep{frank,faulkner}.  In order to weigh the merits of each possibility, we analyse archival \\cha\\ images; as well as infrared observations of `The Stem' obtained with the {\\it Spitzer Space Telescope}. The organization of the paper is as follows. \\S 2  describes the observations. Spectral fits are summarized in \\S 3. In \\S 4 we discuss alternative models for the X-ray emission. Finally, conclusions and future work are presented in \\S 5. ",
        "conclusions": "In view of our discussion, two leading explanations emerge to explain the source of extended emission in the outskirts of GLIMPSE-CO1. First, based purely on its purported X-ray luminosity and physical size, `The Stem' may represent a PWN lacking a central point-like source. If so, `The Stem' becomes the most notable object in the 95-percent confidence error contour of unidentified \\fermi\\ source 0FGL J1848.6-0138. The alternative model, a potentially more intriguing explanation for the extended source of emission, is that `The Stem' traces back to a bow shock produced as the globular cluster GLIMPSE-C01 passes through the Galactic plane. If confirmed, the latter explanation would indicate that GLIMPSE-C01 is losing part of its intracluster gas in the process. More generally, it would provide evidence of the systematic stripping of intracluster gas on Galactic scales long predicted by a number of theoretical models. However, given the unusual nature of this source, it is also possible that `The Stem' represents a new type of X-ray emitter.  With these findings at hand, it is critical to determine the radial velocity and proper motion of GLIMPSE-C01 through alternative means. In addition, deeper X-ray observations of this region are needed to refine the current modelling of the X-ray emission. We envision two vital tests that may conclusively trim these models. First, the orbit of the GLIMPSE-C01 must agree with the placement of the purported bow shock. Second, radio/X-ray pulsation searches must aim for a pulsar within the extended emission and place stricter constraints on the radio component. Given the heavy crowding around this region, we believe that GLIMPSE-C01 makes an exquisite target for the newly refurbished {\\it Hubble Space Telescope} that would result in the most accurate proper motion measurements of the globular cluster."
    },
    "0911/0911.2222_arXiv.txt": {
        "abstract": "We reconsider the proposal of excited dark matter (DM) as an explanation for excess 511 keV gamma rays from positrons in the galactic center. We quantitatively compute the cross section for DM annihilation to nearby excited states, mediated by exchange of a new light gauge boson with off-diagonal couplings to the DM states. In models where both excited states must be heavy enough to decay into $e^+ e^-$ and the ground state, the predicted rate of positron production is never large enough to agree with observations, unless one makes extreme assumptions about the local circular velocity in the Milky Way, or alternatively if there exists a metastable population of DM states which can be excited through a mass gap of less than 650 keV, before decaying into electrons and positrons.\\\\  \\centerline{Dedicated to the memory of Lev Kofman} ",
        "introduction": "There is presently much discussion about whether several anomalous observations in galactic gamma ray and cosmic ray astronomy are better explained by dark matter models or by more conventional astrophysics. One such example is the 511 keV gamma ray excess from the galactic center, which has been observed for 40 years, most recently by the SPI spectrometer on the INTEGRAL satellite \\cite{spi}; for a recent review see ref.\\ \\cite{Diehl}.  Pulsars \\cite{pulsars}, gamma ray bursts \\cite{grb}, supernovae \\cite{sne}, low-mass x-ray binaries \\cite{lmxrb}, the galactic black hole \\cite{bh} and anisotropic propagation effects \\cite{prop} have been suggested as sources of positrons whose annihilation could explain the observations, but there is no consensus within the astrophysical community as to which of these might be the right explanation.  There have also been many attempts to explain the 511 keV signal using particle physics models, including annihilations of MeV scale dark matter (DM) \\cite{mevdm} or millicharged DM \\cite{milli}, emission from cosmic strings \\cite{cosmic-strings}, decays of sterile neutrinos \\cite{pp}, axinos \\cite{axinos}, moduli (or modulinos) \\cite{moduli}, WIMPs \\cite{PR,wimps}, light photinos \\cite{photini}, or emission from composite objects \\cite{composite}. Excited dark matter (XDM) \\cite{xdm} is a particularly appealing example, in which heavy DM particles in their ground state scatter into excited states, $\\chi_0\\chi_0\\to \\chi_1\\chi_1$, with mass difference $\\delta M = M_1-M_0$.  The excited states subsequently decay into $e^+ e^-\\chi_0$, where the leptons are approximately nonrelativistic due to the mass difference being close to the threshold value $\\Delta M\\gsim 2 m_e$.  If the different DM states are members of a multiplet of a new hidden gauge symmetry, this framework has the advantage of being able to explain the small mass splittings naturally through quantum loop effects \\cite{nima}; moreover the same class of theories can potentially explain excess positrons seen by higher energy experiments including ATIC \\cite{atic}, PPB-BETS \\cite{ppb-bets}, PAMELA \\cite{pamela} and the Fermi Large Area Telescope \\cite{fermi}.   The XDM proposal requires that the excitation cross section be large, in fact close to the unitarity bound, in at least a few (and possibly many) partial waves \\cite{PR}, \\cite{FSWY}.  Ref.\\ \\cite{xdm} made a first attempt to achieve such large values in a model where the excitation was mediated by exchange of a light scalar $\\phi$ with $m_\\phi$ in the MeV$-$GeV range.  It was recognized that one must resum ladder diagrams with multiple $\\phi$ exchanges when the DM particles are scattering at low velocity, but a quantitatively reliable way of doing so was not yet appreciated at the time of this work.  Subsequently ref.\\ \\cite{nima} showed how the calculation can be set up in the framework of nonrelativistic quantum mechanics of a two-state system.  It provides an example of the Sommerfeld enhancement \\cite{Sommerfeld,Slatyer}, which in the last few years has been widely studied in the context of galactic dark matter annihilations.  However, unlike the simpler one-state system where the enhancement can be approximated analytically, in the two-state system no analytic results for the excitation cross section are known (however see ref.\\ \\cite{Slatyer} for recent progress in the annihilation cross section for multi-state DM).   Ref.\\ \\cite{CCF}  made a first attempt to numerically solve the two-state system and to perform  a preliminary scan of the parameter space.  Technical challenges limited that analysis to relatively small coupling strengths of the exchanged boson to the dark matter.  One goal in the present work is to extend these results to stronger couplings and to explore more widely the range of possibilities, including the dependence on the mass of the exchanged boson.  A second is to  explore the dependence of the predicted rate  on parameters of the dark matter density and velocity profiles in the galaxy.  In section \\ref{setup} we review the nonrelativistic quantum mechanical formulation of the problem and give a classical estimate of the number of partial waves that can be expected to significantly contribute to the excitation cross section.  In section \\ref{method} we discuss the numerical  method for computing partial wave amplitudes $f_l$ as a function of velocity, and present sample results for $f_l$, as well as a  survey of the range of boost factors which arise from a broad exploration of parameter space. In section \\ref{rate_sect} we show how these results go into the  calculation of the rate of positron production in the galactic bulge; our choice of different DM density and velocity dispersion profiles is explained there.  Section \\ref{survey} presents the  results  of our scan of the XDM model's parameter space, consisting of the interaction strength and the mass of the exchanged particle, and the  mass splitting between ground and excited state.  We at first hold the DM mass $M_0$ fixed at a TeV \\cite{wimp-mass}, but then explore the dependence on  $M_0$ and galactic DM distribution parameters at some optimal values of the microscopic parameters.   Our results show that a large enough rate to match observations cannot be achieved in the most straightforward implementation of the XDM scenario, but in section \\ref{alternative} we review a modified version which can overcome this deficit.  In this version of XDM, dark matter has at least three states; the middle state is stable or metastable and can excite over a smaller gap to the highest state. In section \\ref{angular} we compare the predictions for the angular distribution of 511 keV gamma rays with the observed signal. Conclusions are given in section \\ref{conclusion}.  Appendix \\ref{born} derives the Born approximation for the excitation cross section, and \\ref{vesc_app} reviews the derivation of the radial dependence of the DM escape velocity. ",
        "conclusions": "\\label{conclusion}  In this paper we have tried to give the most quantitative treatment to date of the excited dark matter mechanism for producing positrons at the galactic center.  A main technical improvement was the numerical computation of the excitation scattering cross section, which is nonperturbative and guaranteed to satisfy unitarity constraints.  A second difference relative to previous treatments is that we included the effects of baryons on the velocity dispersion of DM in the inner kpc of the galaxy, which can significantly boost the rate of excitations.  In addition we considered NFW and Einasto profiles for the DM density, varying parameters determining the cuspiness of the distributions, as well as those affecting the velocity distributions.  Even making all the most optimistic assumptions for increasing the rate of DM excitations, we were not able to find realistic parameter values which yield a large enough rate, if the mass gap between the ground state $\\chi_0$ and excited state $\\chi_1$ of the DM is sufficient for  the subsequent decay $\\chi_1\\to \\chi_0\\, e^+ e^-$ to produce the positron.  The largest mass gap we found consistent with the observed rate was $\\delta M \\cong 650$ keV, assuming the standard values $v_c = 220$ km/s  and $R_0 = 8$ kpc, respectively, for the circular velocity of the sun around and its distance to the galactic center.  By pushing these values to  $v_c = 280$ km/s and $R_0 = 9.4$ kpc, we can barely accommodate a mass gap of $\\delta M = 2 m_e$, but these choices are respectively $4\\sigma$ and $3\\sigma$ away from the mean values.  One way to ease the tension would be to consider models where the excited states are charged and thus able to decay into single electrons or positrons; this would lower the required mass gap to just $\\delta M = m_e$.  But this requires a great deal more model-building gymnastics than the case of a neutral excited state. A theoretically more appealing alternative is the case of three DM states where $\\chi_0\\chi_0\\to \\chi_1\\chi_2$, in which the mass gap $\\delta M_{01}$ between $\\chi_0$ and $\\chi_1$ is much smaller than $\\delta M_{02} > 2 m_e$ \\cite{Cheung,CCF2}.  For example if $\\delta M_{01}=0$, then the effective  $\\delta M$ which determines the threshold velocity is half as large as in the models that produce two of the same excited state.  This possibility can arise if the hidden sector gauge group is  SU(2)$\\times$U(1) and gets completely broken, for example.  However to see if this class of models can really have a larger rate would require solving the Schr\\\"odinger equation in the three-state system, so while it is plausible that such asymmetric excitations could increase the rate, it is not proven by our analysis.  Another way of getting around our no-go result is to assume the existence of a stable or metastable population of already excited states $\\chi_1$, which need only be further excited to $\\chi_2$ through a small mass gap $\\delta M_{12}$, if $\\delta M_{02}$ is assumed to be greater than $2 m_e$ so that the decay $\\chi_2\\to \\chi_0\\, e^+e^-$ is allowed.  We refer to this as the inverted mass hierarchy scenario.  This idea comes with new complications, since it is difficult to prevent the excitation process from occuring earlier in the history of the universe, to maintain the population of $\\chi_1$ to the present day.  However there is an  intriguing possibility to overcome this by assuming the present DM particles are products of the late decay of a much heavier predecessor, and thus were initially relativistic, at a time when they would normally have been nonrelativistic had they been produced through conventional freeze-out \\cite{CCF}.  The higher velocity suppresses the Sommerfeld enhancement at early times, when it would have the undesirable effect of depleting the metastable intermediate DM states.  We hope to examine this scenario more carefully in future work.  We have not tried to impose in detail the additional constraints on the model which arise if one would like it to also account for the PAMELA and Fermi/LAT excess electrons/positrons, although we did focus on  TeV scale DM for that purpose.  If the DM explanation of the high-energy leptons is not ruled out by upcoming analyses, this would be an interesting next step.   It might also be worthwhile to investigate the effect of a nonspherically symmetric DM halo \\cite{jmr} on the rate of positron production.  \\centerline{$\\phantom{.}$} {\\bf Acknowledgment:}  We thank Gil Holder for valuable discussions about the galactic parameters.   \\appendix"
    },
    "0911/0911.3154_arXiv.txt": {
        "abstract": "\\begin{list}{ } {\\rightmargin 0.5in}{\\leftmargin 1in} \\baselineskip = 11pt \\parindent=1pc {\\small We discuss the interior structure and composition of giant planets, and how this structure changes as these planets cool and contract over time.  Here we define giant planets as those that have an observable hydrogen-helium envelope, which includes Jupiter-like planets, which are predominantly H/He gas, and Neptune-like planets which are predominantly composed of elements heavier than H/He.  We describe the equations of state of planetary materials and the construction of static structural models and thermal evolution models.  We apply these models to transiting planets close to their parent stars, as well as directly imaged planets far from their parent stars.  Mechanisms that have been postulated to inflate the radii of close-in transiting planets are discussed.  We also review knowledge gained from the study of the solar system's giant planets.  The frontiers of giant planet physics are discussed with an eye towards future planetary discoveries. \\\\~\\\\~\\\\~}%  \\end{list} ",
        "introduction": "The vast majority of planetary mass in the solar system, and indeed the galaxy, is hidden from view in the interiors of giant planets.  Beyond the simple accounting of mass, there are many reasons to understand these objects, which cut across several disciplines.  Understanding the structure of these planets gives us our best evidence as to the formation mode of giant planets, which tells us much about the planet formation process in general.  As we will see, giant planets are vast natural laboratories for simple materials under high pressure in regimes that are not yet accessible to experiment.  With the recent rise in number and stunning diversity of giant planets, it is important to understand them as a class of astronomical objects.  We would like to understand basic questions about the structure and composition of giant planets.  Are they similar in composition to stars, predominantly hydrogen and helium with only a small mass fraction ($\\sim$1\\%) of atoms more massive than helium?  If giant planets are enhanced in  ``heavy elements'' relative to stars, are the heavy elements predominantly mixed into the hydrogen-helium (H-He) envelope, or are they mainly found in a central core?  If a dense central core exists, how massive is it, what is its state (solid or liquid), and is it distinct or diluted into the above H-He envelope?  Can we understand if a planet's heavy element mass fraction depends on that of its parent star?  Giant planets are natural laboratories of hydrogen and helium in the megabar to gigabar pressure range, at temperatures on the order of 10$^4$ K. How do hydrogen and helium interact under these extreme conditions? Is the helium distribution within a planet uniform and what does this tell us about how H and He mix at high pressure?  What methods of energy transport are at work in the interiors of giant planets?  Can we explain planets' observable properties such as the luminosity and radius at a given age?  The data that we use to shape our understanding of giant planets comes from a variety of sources.  Laboratory data on the equation of state (EOS, the pressure-density-temperature relation) of hydrogen, helium, warm fluid ``ices'' such at water, ammonia, and methane, silicate rocks, and iron serve at the initial inputs into models.  Importantly, data are only avaible over a small range of phase space, so that detailed theoretical EOS calculations are critical to understanding the behavior of planetary materials at high pressure and temperature.  Within the solar system, spacecraft data on planetary gravitational fields allows us to place constraints on the interior density distribution for Jupiter, Saturn, Uranus, and Neptune.  For exoplanets, we often must make due with far simpler information, namely a planet's mass and radius only.  For these distant planets, what we lack in detailed knowledge about particular planets, we can make up for in number.  Within six years of the Voyager 2 fly-by of Neptune, the encounter that completed our detailed census of the outer solar system, came the stunning discoveries of the extrasolar giant planet 51 Peg b \\cp{Mayor95} and also the first bona fide brown dwarf, Gliese 229B \\cp{Nakajima95}.  We were not yet able to fully understand the structure and evolution of the solar system's planets before we were given a vast array of new planets to understand.  In particular the close-in orbit of 51 Peg b led to immediate questions regarding its history, structure, and fate \\cp{Guillot96,Lin96}.  Four years later, the first transiting planet, \\hd\\ \\cp{Charb00,Henry00}, was found to have an inflated radius of $\\sim$1.3 \\rj, confirming that proximity to a parent star can have dramatic effects on planetary evolution \\cp{Guillot96}.  The detections of over 50 additional transiting planets (as of August 2009) has conclusively shown that planets with masses greater than that of Saturn are composed predominantly of H/He, as expected.  However, a great number of important questions have been raised.  Much further from their parent stars, young luminous gas giant planets are being directly imaged from the ground and from space \\cp{Kalas08,Marois08,Lagrange09}.  For imaged planets, planetary thermal emission is only detected in a few bands, and a planet's mass determination rests entirely on comparisons with thermal evolution models, that aim to predict a planet's luminosity and spectrum with time.  However, the luminosity of young planets is not yet confidently understood \\cp{Marley07,Chabrier07}.  In this chapter we first outline the fundamental physics and equations for understanding the structure and thermal evolution of giant planets.  We next describe the current state of knowledge of the solar system's giant planets.  We then discuss current important issues in modeling exoplanets, and how models compare to observations of transiting planets, as well as directly imaged planets.  We close with a look at the future science of extrasolar giant planets (EGPs).   \\bigskip ",
        "conclusions": ""
    },
    "0911/0911.1198_arXiv.txt": {
        "abstract": "We report optical spectroscopy and high speed photometry and polarimetry of the INTEGRAL source IGRJ14536-5522 (=Swift J1453.4-5524). The photometry, polarimetry and spectroscopy are modulated on an orbital period of $3.1564(1)$ hours. Orbital circularly polarized modulations are seen from $\\sim 0$ to $\\sim -18$ per cent, unambiguously identifying IGRJ14536-5522 as a polar. The negative circular polarization is seen over $\\sim 95$ per cent of the orbit, which is consistent (as viewed from Earth) with a single pole accretor. We estimate some of the system parameters by modeling the polarimetric observations.  Some of the high speed photometric data show modulations that are consistent with quasi-periodic oscillations (QPOs) on the order of 5-6 minutes. Furthermore, for the first time, we detect the (5-6) minute QPOs in the circular polarimetry. We discuss the possible origins of these QPOs. In addition, we note that the source undergoes frequent changes between different accretion states.  We also include details of HIPPO, a new high-speed photo-polarimeter used for some of our observations. This instrument is capable of high-speed, multi-filtered, simultaneous all-Stokes observations. It is therefore ideal for investigating rapidly varying astronomical sources such as magnetic Cataclysmic Variables. ",
        "introduction": "The standard picture of a cataclysmic variable (CV) is a binary system consisting of a Roche lobe filling red dwarf (known as the secondary or the donor star) and an accreting white dwarf (the primary). CVs have orbital periods of typically a few hours, and mass transfer is caused by angular momentum loss - see e.g. Warner (1995) for a review of cataclysmic variables.  Approximately 20\\% of the known CVs are magnetic cataclysmic variables (mCVs), where the white dwarf has a strong magnetic field (see the catalogue of Ritter \\& Kolb 2003). These are further sub-divided into two subtypes, namely intermediate polars (IPs) and polars, depending on the strength of the magnetic field of the white dwarf and the degree of synchronism between the white dwarf spin and the binary orbit- see the reviews given by Cropper (1990) and Patterson (1994).  In polars, also known as AM Her systems, the white dwarf has a sufficiently strong magnetic field to lock the system into synchronous rotation and to prevent completely the formation of an accretion disc. Instead, the material from the secondary overflowing the Roche lobe initially falls towards the white dwarf following a ballistic trajectory until, at some distance from the white dwarf, the magnetic pressure overwhelms the ram pressure of the ballistic stream, confining the flow until it eventually reaches the surface where it forms a hot shocked region.  It has long been known that CVs are a significant source of soft ($<$ 2 keV) and medium energy (2-10 keV) X-rays (e.g. Patterson et al. 1984). Recent surveys with INTEGRAL show that CVs are also notable sources of hard ($>$ 20 keV) X-rays. A large fraction of these are made up of magnetic CVs and in particular IPs (see Barlow et al. 2006 and Revnivtsev et al 2008). Interestingly, out of the five asynchronous polars known, two are INTEGRAL sources.  IGRJ14536-5522 (=Swift J453.4-5524) was discovered as a hard X-ray source by INTEGRAL (Kuiper, Keek, Hermsen, Jonker \\& Steeghs 2006) and by Swift/BAT (Mukai et al. 2006). A pointed Swift/XRT observation led to the identification with a ROSAT all-sky survey (RASS) source 1 RXS J145341.1-552146, and hence to its optical identification (Masetti et al. 2006). Revnivtsev et al. (2008) classify it as an IP.  Follow-up spectroscopy and photometry with SALT (Southern African Large Telescope) and with the SAAO (South African Astronomical Observatory) 1.9m telescope showed that this object belongs to a rare subtype of magnetic CV: a hard X-ray bright polar or a soft intermediate polar (Mukai, Markwardt, Tueller, Buckley, Potter, Still et al. 2006). Further observations were clearly needed.  Here we report on our optical photo-polarimetric and spectroscopic observations of IGRJ14536-5522. The paper is structured as follows: In Sect. 2 we describe the overall design of a new high speed photo-polarimeter used for some of our observations. Sect. 3 gives an account of all our observations, followed by our analysis of the spectroscopic and photo-polarimetric observations in Sects. 4 and 5 respectively. We finish with a discussion and summary in Sect. 6. ",
        "conclusions": "\\subsection{Object classification and system geometry}  Our optical spectroscopy and high speed photo-polarimetry of the INTEGRAL source IGRJ14536-5522 (=Swift J453.4-5524) unambiguously confirm its identification as a polar. Negative circular polarization is seen over all of the orbit, which is consistent with a single pole accretor at a moderate inclination. We estimate some of the system parameters by modeling the polarimetric observations. The most successful model integrated the emission from an arc-shaped region extending from 170$^{\\rm o}$ to 230$^{\\rm o}$ in magnetic longitude and 10$^{\\rm o}$ from the magnetic pole. The system inclination was found to be in the range of 45-55$^{\\rm o}$ with a magnetic dipole offset angle of 10$^{\\rm o}$.  \\subsection{The QPOs and flickering}  Our high speed photo-polarimetry shows flickering on minute time-scales. Furthermore, for the first time, we detect QPOs in the photometry and circular polarimetry.     Photometric variations on the time scales of minutes are a common feature of many polars, from X-rays to infrared (e.g. Szkody \\& Margon 1980 and Watson, King \\& Williams 1987), and have been characterised as flickering, fluctuations, QPOs or erratic QPOs according to different authors. In general (although not always strictly true), QPOs describe variations that show some coherence over a period of time. They mainly cluster in time-scales in the range of a few seconds (1-5) or minutes (4-10). The (1-5s) QPOs are of low amplitude and have been observed in the optical in only a few polars (e.g. V834 Cen, AN Uma and VV Pup). The larger amplitude (4-10 min) QPOs seem to be a general feature of most polars. By eye, they can appear to show some sort of coherence, but fail to show any significant singularly persistent peaks under Fourier analysis. Instead, groupings of periods are seen and are often referred to as QPO-like. Flickering and fluctuations best describe variations that do not show any periodic behavior.  There is evidence that the 4-10 minute QPOs may consist of the superposition of regular periodic oscillations: Bonnet-Bidaud, Somova \\& Somov (1991) report on observations of AM Her during an intermediate brightness state where they find fluctuations to be nearly periodic, with a period increasing from 250s to 280s. They reason that they observed AM Her in a transition from bright-state to faint-state where a given accretion rate is able to excite stable oscillations. Alternatively, during the higher state, many oscillations may be present but masked by the superposition of different modes corresponding to different accretion tubes. Therefore, during the transition, the accretion rate decreases, reducing the number of accretion tubes until finally only one unique tube contributes to the emission.  Ultimately, the photometric oscillations are caused by variations in the accretion flow. There are several theories/models that attempt to explain the variations. King (1989) remarks that the (4-10 minute) periods are characteristic of a dynamical time at the photosphere of the companion. Irradiation of the region near the L$_{1}$ point will result in the formation of an ionisation front, which will tend to oscillate and therefore modulate the accretion. Other models involve the capture of inhomogeneities of the accretion flow at the capture radius (Kuijpers and Pringle 1982) and accretion gate mechanisms (Patterson et al. 1981).  The short period QPOs (1-5s) have been observed in $\\sim $ 6 systems (e.g. V834 Cen, AN Uma and VV Pup).  Langer, Chanmugam \\& Shaviv 1982) realised that the 1-5s oscillations were consistent with the cooling time-scales for white dwarf-radiated shock waves and thus QPOs are potentially powerful probes of the radiative shocks. Observations of VV Pup (Larson 1989) demonstrated conclusively that the QPOs arose in a region near the shock emission region on the white dwarf.   From our results presented in Figs.~\\ref{28febQPO} and \\ref{7aprQPO} it is clear that IGRJ14536-5522 exhibited photometric and polarized QPOs (5-6 minutes) on two separate occasions. Additionally, the fact that we observe the QPOs oscillations in the polarimetry, unequivocally places the emission site at the cyclotron emitting shock region. Furthermore, the QPOs were seen during phases 1.0-1.3 only. We propose that the super-position of the oscillations arising from numerous accretion tubes (similar to the mechanism of Bonnet-Bidaud, Somova \\& Somov 1991) applies to IGRJ14536-5522. However, instead of a reduced accretion state leading to a reduced number of accretion tubes, we suggest that for most of the orbit the superposition of emission from many accretion tubes leads to the observed flickering. However during the phase interval 1.0-1.3 we preferentially see the trailing edge of the accretion region as the white dwarf rotates. The leading part of the accretion region is effectively shielded from our view by its trailing edge for this phase interval. Consequently we observe a singular QPO corresponding to a singular accretion tube.  Our observations of IGRJ14536-5522 indicate that it undergoes relatively frequent changes in accretion state, making it a good candidate for capturing it in transition and testing this scenario further. If QPOs are also present in the X-rays then IGRJ14536-5522 is an ideal source for investigating the two dominating shock cooling mechanisms: i.e. bremsstrahlung and cyclotron cooling.  Finally, we would like to note the importance of optical follow-up observations of candidate CV INTEGRAL sources, in particular with photometry and/or polarimetry. It has been noted that many CVs detected by INTEGRAL are IPs, both new and re-discoveries (see e.g. Barlow et al. 2006 and Revnivtsev et al 2008). However, as Pretorius (2009) points out, it is only through follow-up observations that unequivocal identifications can be made. A case in point is IGRJ14536-5522, which had been assumed to be an IP (see e.g. Revnivtsev et al. 2008) even though a spin period had not been detected."
    },
    "0911/0911.4890_arXiv.txt": {
        "abstract": "s{ \\textit{Fermi} Large Area Telescope (LAT) has recently detected~8 gamma-ray millisecond pulsars (MSPs), providing an unprecedented opportunity to probe the magnetospheres of these low-spin-down pulsars. We performed 3D emission modeling, including various Special Relativistic effects, in the context of pair-starved polar cap (PSPC), slot gap (SG), and outer gap (OG) pulsar models. Most of the light curves are best fit by SG and OG models, surprisingly indicating the presence of narrow accelerating gaps limited by robust pair production. All model fits imply high-altitude emission, and we observe exclusive differentiation of the current gamma-ray MSP population into two sub-classes: light curve shapes and lags across wavebands impose either PSPC or SG / OG-type geometries.} ",
        "introduction": "\\label{sec:intro} Millisecond pulsars (MSPs) are rapidly-rotating neutron stars, characterized by relatively short spin periods and low surface magnetic fields. % It is thought that transfer of mass and angular momentum from a binary companion during an accretion phase, may revive their ``dead'' (radio-invisible) younger pulsar progenitors which have ``spun down'' previously.~\\cite{Alpar82} MSPs are thus believed to be ``recycled'' pulsars,~\\cite{Bhattacharya91} decending from low-mass X-ray binaries (LMXRBs).~\\cite{Archibald09}  MSPs were predicted to be visible in gamma rays.~\\cite{HUM05,Frackowiak05a,Zhang07,Venter09,Zajczyk08,Watters09,Venter09_MSPs} Observational confirmation was provided by the \\textit{Fermi} Large Area Telescope (LAT), a high-energy (HE) gamma-ray telescope with a large field of view (2.4~sr), operating in the energy range $\\sim20$~MeV to $>300$~GeV~\\cite{Atwood09}, which recently detected~8 gamma-ray MSPs.~\\cite{Abdo09_MSP}   Mainly two different models have been employed to describe HE radiation from pulsars. Polar cap (PC) models~\\cite{DH96} assume injection of primary electrons from the stellar surface. These primaries radiate HE curvature radiation (CR) or inverse Compton scattering (ICS) gamma rays which are converted into electron-positron pairs via magnetic pair production. A low-altitude pair formation front (PFF) is established close to the surface, screening the accelerating electric field.~\\cite{HM98} By allowing variation of the CR PFF altitude across the PC, a slot gap (SG) is formed along the last open magnetic field lines. This gap extends from the neutron star surface to near the so-called light cylinder, and allows acceleration of primaries up to high altitudes.~\\cite{Arons83,MH04_SG} The SG model is a possible physical realization of the two-pole caustic geometry.~\\cite{Dyks03} % Outer gap (OG) models~\\cite{CHR86a,CHR86b,Romani96} assume that HE radiation is generated by cascades of electron-positron pairs produced via photon-photon pair creation. This emission occurs close to the last open field lines, above the ``null-charge surface'', and is thus confined to the outer magnetosphere. Recently, a 3D OG solution was found~\\cite{Hirotani08} which extends toward the NS surface. For MSPs with mostly unscreened magnetospheres,  a ``pair-starved polar cap'' (PSPC) model~\\cite{MH04_PS,Venter09_MSPs} applies in which charges are accelerated up to high altitudes over the full open-field-line region. %  We study the \\textit{Fermi}-LAT MSP population~\\cite{Abdo09_MSP} using 3D emission modeling, including Special Relativistic (SR) effects of aberration and time-of-flight delays, and rotational sweepback of B-field lines,~\\cite{Morini83,Romani95,Dyks04_SR} in the context of geometric SG and OG models, as well as a CR PSPC model, and obtain fits for gamma-ray and radio light curves. Our calculations are complementary to another study~\\cite{Watters09} that focuses on younger pulsars. % More details are provided elsewhere.~\\cite{Venter09_MSPs}  \\begin{figure} \\begin{center} \\epsfig{figure=fig16.eps,height=2.5in}\\hskip0.4cm \\epsfig{figure=fig24.eps,height=2.5in} \\caption{\\textit{Left:} Phaseplots for PSR~J0030+0451. Panel~(a) is for an SG model with $(\\alpha,\\zeta)=(70^\\circ,80^\\circ)$, while panel~(b) is for an OG model with $(\\alpha,\\zeta)=(80^\\circ,70^\\circ)$. \\textit{Right:} Observed and fitted light curves for PSR~J0030+0451. In panel~(a), we show the \\textit{Fermi}-LAT data (histogram) with estimated background level (dashed line), and SG (magenta line) and OG (green line) fits. Panel~(b) shows the radio data (blue line), and the magenta and green lines correspond to the same $(\\alpha,\\zeta)$ combinations as those of the SG and OG fits in panel~(a).\\label{fig1}} \\end{center} \\end{figure} ",
        "conclusions": "\\label{sec:Con} We have compared 3D model predictions of gamma-ray and radio radiation with MSP gamma-ray data from \\textit{Fermi}-LAT in the framework of geometric SG and OG pulsar models, and also for the full-radiation PSPC model. Surprisingly, some MSPs have double-peaked light curves well fit by SG / OG models, indicating strong screening of $E_{||}$ by pair production. The larger radio beam widths of MSPs compared to those of canonical pulsars furthermore implies relatively few radio-quiet MSPs. We found exclusive differentiation between the SG / OG models on the one hand, and the PSPC model on the other hand. % Emission in \\textit{all} models considered comes from the outer magnetosphere. Our fits of $\\alpha$ and $\\zeta$ are in reasonable agreement with values inferred from MSP radio polarization measurements. For PSR~J0437-4715 and PSR~J0613-0200, the SG model predicts a small precursor to the main gamma-ray peak, but not the OG model. This may become a future model discriminator. Future phase-resolved spectroscopy made possible by the quality of \\textit{Fermi}-LAT data should challenge pulsar models."
    },
    "0911/0911.1151_arXiv.txt": {
        "abstract": "Stellar-mass black holes in the low-hard state may hold clues to jet formation and basic accretion disk physics, but the nature of the accretion flow remains uncertain.  A standard thin disk can extend close to the innermost stable circular orbit, but the inner disk may evaporate when the mass accretion rate is reduced. Blackbody-like continuum emission and dynamically-broadened iron emission lines provide independent means of probing the radial extent of the inner disk.  Here, we present an X-ray study of eight black holes in the low-hard state.  A thermal disk continuum with a colour temperature consistent with $L \\propto T^{4}$ is clearly detected in all eight sources, down to $\\approx$5$\\times10^{-4}L_{Edd}$.  In six sources, disk models exclude a truncation radius larger than 10\\rg. Iron-\\ka\\ fluorescence line emission is observed in half of the sample, down to luminosities of $\\approx$1.5$\\times10^{-3}L_{Edd}$. Detailed fits to the line profiles exclude a truncated disk in each case.  If strong evidence of truncation is defined as (1) a non-detection of a broad iron line, {\\it and} (2) an inner disk temperature much cooler than expected from the ${\\rm L} \\propto {\\rm T}^{4}$ relation, none of the spectra in this sample offer strong evidence of disk truncation.  This suggests that the inner disk may evaporate at or below $\\approx$1.5$\\times10^{-3}L_{Edd}$. ",
        "introduction": "The innermost regions of accreting black holes and the underlying physics of the accretion process can be successfully studied using the X-ray spectra of black hole binaries. Over the past years, such studies have revealed that Galactic black hole binaries (GBHBs) radiate in various distinct spectral states characterised by the relative strength of their soft and hard X-ray emission (McClintock \\& Remillard 2006).   The X-ray spectrum of black hole binary systems in their low-hard state is dominated by a powerlaw with relatively low luminosity ($\\le$0.05$L_{Edd}$; Maccarone 2003), a photon index $\\Gamma$ in the range $\\approx$1.4--2 and an exponential cut-off at about 100\\kev. This is in contrast with the high-soft state, where the spectrum is dominated by a quasi-black body component with a characteristic temperature of $\\sim$1\\kev, and the highly luminous very--high state ($L$$\\approx$$L_{Edd}$) with a powerlaw component having a flux comparable to that of the soft blackbody. In the latter cases, the powerlaw does not show evidence for a high--energy roll-over. Superimposed on the continuum is the presence of various reflection features (Ross \\& Fabian 1993; Miller 2007) with the dominant component being the Fe-\\ka\\ emission line. In the inner regions of the accretion flow the reflection spectrum appears blurred due to the combination of various relativistic effects. The degree of gravitational blurring is strongly dependent on the inner radius of the reflecting material (Fabian et al. 1989; Laor 1991).   Theoretically, the soft, quasi-black body component observed in the high-soft and very-high state is generally agreed to originate in a standard, optically thick, geometrically thin accretion disk (Shakura \\& Sunyaev 1973) extending to the innermost stable circular orbit (ISCO). For a non-rotating, Schwarzschild, black hole this radius is equal to 6\\rg\\ where \\rg{\\thinspace = $GM/c^2$}, and decreases to $\\approx$1.2\\rg\\ for a maximally-rotating, Kerr black hole (Thorne 1974). The powerlaw component dominating the X-ray spectra in the low-hard state, on the other hand, is believed to be produced by the inverse Compton scattering of soft photons in a thermal, optically thin region (Shapiro et al. 1976; Sunyaev \\& Titarchuk 1980).  However, there is no general consensus on the geometry of this optical--thin region.  In the accretion disk-corona model the hard powerlaw comes from either a patchy corona, possibly powered by magnetic flares (Di Matteo, Celotti \\& Fabian 1999; Beloborodov 1999; Merloni, Di Matteo \\& Fabian 2000; Merloni \\& Fabian 2001) or the base of a centrally located jet (Merloni \\& Fabian 2002; Markoff \\& Nowak 2004; Markoff, Nowak \\& Wilms 2005).  In both cases the thin accretion disk extends close to the ISCO. An alternative model has the thin disk truncated at large distances from the black hole (Narayan \\& Yi 1995; Esin et al. 1997) with the central region replaced by an advection-dominated accretion flow (ADAF).  In both interpretations for the geometry of the accretion disk in the low-hard state it is expected that the thermal-disk component will have a low effective temperature. The low mass accretion rates, and corresponding low luminosity, observed in systems in the low-hard state implies a peak disk temperature $<0.4$\\kev. In the case of a geometrically thin disk extending close to the ISCO, this temperature closely follows the $L\\propto T^4$ relation. The expected temperature departs dramatically from this relation when the disk is truncated at the distances predicted by the ADAF models (see e.g. McClintock et al. 2001). A thermal component has been observed in a number of sources in the low-hard state. In particular, \\xmm/\\rxte\\ observations of \\gx\\ (Miller et al. 2006; Reis et al. 2008) and Swift J1753.5-0.127 (Miller, Homan \\& Miniutti 2006; Reis et al. 2009b) as well as \\chandra\\ observation of XTE J1118+480 (Reis, Miller \\& Fabian 2009) have shown that the accretion disk in these sources are consistent with extending close to the ISCO. Evidence for the disappearance of this component has usually been the result of observations made with \\rxte\\ which lacks the low energy coverage.   As mentioned above, a further feature in the X-ray spectra of BHBs are the various reflection signatures associated with the reprocessing of hard radiation by the cool accretion disk. The most prominent of these features is often the broad, skewed Fe-\\ka\\ line observed in a number of BHBs (Miller 2007, 2009), Seyfert galaxies (Tanaka et al. 1995; Fabian et al. 2009) and accreting neutron stars (Cackett et al. 2008; Reis, Fabian \\& Young 2009).  New fits to some ultra-luminous X-ray sources (ULXs; Caballero-Garcia \\& Fabian 2009) spectra suggest that a disk reflection may also be observed, though more sensitive spectra are needed to confirm this possibility. The strength of the reflection features can be typified by the equivalent width of the iron-\\ka\\ line, $W_{K\\alpha}$ which is expected to correlate linearly with the reflection fraction $R = \\Omega/2\\pi$, where $\\Omega$ is defined as the solid angle covered by the accretion disk as viewed from the hard X-ray source. George \\& Fabian (1991) have shown that for a neutral accretion disk with solar abundances the reflection fraction closely follows $R\\approx W_{K\\alpha}/180\\ev$. For stellar mass black holes in the low-hard state, $R$ is observed to be below $\\sim 1$. This ``weak'' reflection can be interpreted as either a recession in the accretion disk, as in the ADAF interpretation (Esin et al. 1997), a highly ionised inner disk (Ross, Fabian \\& Young 1999) or mildly relativistic motion of the hot corona away from the disk (Beloborodov 1999). Using the latter model, Beloborodov showed that reflection fractions as low as $\\approx 0.3$ can be be obtained in the low-hard state without invoking disk truncation.    In the following chapters we will present a systematic analysis of a sample of eight stellar mass black hole binaries in the low-hard state with the goal of better understanding the accretion disk and flow geometry of such systems. We find that for the sources observed in this work the accretion disks are consistent with extending close to the ISCO and suggest a set of observational criteria for the support of disk truncation. We start in \\S{\\thinspace\\ref{observation}} with an introduction to the various systems investigated. This is then followed by a detailed analyses of the property of the thermal component (\\S\\S{\\thinspace\\ref{analyses_mcd}}--\\ref{innerrad_diskpn}) and the reflection component (\\S\\ref{innerrad_reflection}). Our results are summarised in \\S{\\thinspace\\ref{discussion}}. ",
        "conclusions": "In this paper we have investigated a sample of stellar mass black holes in the canonical low-hard state. By systematically analysing their X-ray spectrum we have found that in all cases the accretion disk is consistent with being at the inner-most stable circular orbit down to luminosities as low as $\\approx5\\times10^{-4}L_{Edd}$. The main points and implication of this paper are summarised as follows.  \\begin{enumerate} \\item In all sources investigated the presence of an accretion disk is required at the 5$\\sigma$ level of confidence. The temperature and flux  of this thermal component is consistent with the $L\\propto T^4$ relation and with a geometrically thin, optically thick accretion disk extending to the ISCO. \\item The presence of reflection features and  predominantly an  iron-\\ka\\ emission line with an equivalent width greater than $\\sim$70\\ev\\ is detected in half of the sample. In all these cases, the broadness of the reflection features further suggest that the accretion disk is not highly truncated. \\item Our findings suggest that transition to the low-hard state are driven by changes in the corona (perhaps related to jet formation) and not changes in the accretion disk, and \\item Jet production is not initiated at the point where the disk recedes. \\item Furthermore, we suggest the following strong and weak observation criteria for disk truncation: \\begin{itemize} \\item The data must be able to rule-out both a broad iron-\\ka\\ line with an equivalent width  $\\gtrsim$60\\ev\\ AND an effective disk temperature consistent with $L\\propto T^4$ or \\item The data must be able to rule-out either a broad iron-\\ka\\ line with an equivalent width  $\\gtrsim$60\\ev\\ OR an effective disk temperature consistent with $L\\propto T^4$. \\end{itemize} \\end{enumerate}    Whilst the number of sources presented here is small, there is a general trend that when the accretion disk is statistically required the data suggests that the disk in the low-hard state can remain at the ISCO. This result is contrary to the lore that the accretion disk in the low-hard state is truncated at hundreds of gravitational radii, as required in the strong-advection dominated interpretation."
    },
    "0911/0911.4297_arXiv.txt": {
        "abstract": "Submillimeter galaxies (SMG) represent a dust-obscured high-redshift population undergoing massive star formation activity. Their properties and space density have suggested that they may evolve into spheroidal galaxies residing in galaxy clusters. In this paper, we report the discovery of compact ($\\sim10-20\\arcsec$) galaxy overdensities centered at the position of three SMGs detected with the Max-Planck Millimeter Bolometer camera (MAMBO) in the COSMOS field. These associations are statistically significant. The photometric redshifts of galaxies in these structures are consistent with their associated SMGs; all of them are between $z=1.4-2.5$, implying projected physical sizes of $\\sim170$ kpc for the overdensities. Our results suggest that about 30\\% of the radio-identified bright SMGs in that redshift range form in galaxy density peaks in the crucial epoch when most stars formed. ",
        "introduction": "Submillimeter galaxies (SMG) are dust-obscured starburst galaxies at high-redshift \\citep{Smail1997, Hughes1998, Barger1998}. Their large dynamical \\citep{Greve2005, Tacconi2006, Tacconi2008} and stellar masses \\citep{Borys2005, Dye2008}, as well as their number densities and clustering properties \\citep{Scott2002, Blain2004, Viero2009}, suggest they could be the progenitors of present-day luminous ellipticals \\citep{Lilly1999, Swinbank2006}. Some SMGs are known to be associated with galaxy clusters at high redshifts \\citep[e.g.][]{Webb2005} and to be located in extended overdensities of LBGs and radio galaxy fields \\citep{Ivison2000, Smail2003, Stevens2003, Chapman2008, Daddi2009, Tamura2009}.  If SMGs are progenitors of massive clustered spheroids, they would likely show signs of clustering in the epoch when these galaxies have their peak in activity and luminosity ($z\\sim 1 - 3$), similar to what is observed for powerful AGN \\citep{Miley2008}. Attempts to measure the clustering of SMGs indicate that they are associated with massive dark matter halos and possibly trace the largest scale structures at high redshifts \\citep{Blain2004, Viero2009, Weiss2009}. However, current submillimeter blank-field surveys either cover small areas in the sky, yielding a few tens of sources within contiguous fields \\citep{Coppin2006, Bertoldi2007, Scott2008, Perera2008, Austermann2009}, or are limited by poor angular resolution albeit covering large regions \\citep{Devlin2009}. Hence, good quantitative studies of the small to large scale clustering of the SMG population are not feasible until large surveys comprising a few square degrees on the sky under good resolution ($\\lesssim20\\arcsec$) can be made.  Studies of the environment of SMGs are possible when utilizing the rich complementary data available for current (sub)millimeter fields. Whether SMGs are embedded in regions with an enhanced number of optical/near-IR detected high-redshift galaxies has yet not been quantified. In this paper, we investigate to what extent SMGs are located in clustered fields. For this, we make use of deep optical and near-IR imaging data in the central part of the COSMOS field to measure the density of high-redshift $BzK$ galaxies in the surroundings of SMGs. This allows us to study the relation of blank-field detected SMGs with the most prominent galaxy density peaks at high-redshift. Hereafter, we assume a cosmology with $H_0=70$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{\\Lambda}=0.7$ and $\\Omega_\\mathrm{M}=0.3$ and use all magnitudes in the AB system. ",
        "conclusions": "\\label{sect:discussion}  We find that significant overdensities of star-forming high-redshift galaxies are related to three MAMBO galaxies detected in the COSMOS field.  These groups are compact in size, and the peaks of their redshift distributions are compatible with the redshift estimated for the associated MAMBO galaxies.  \\subsection{SMGs in dense environments}  If SMGs are related to the formation of structures at high-redshift, we would expect that most SMGs are located in regions with enhanced galaxy densities. However, only a few SMGs in our sample can be associated with strong galaxy overdensities. This could be partly attributed to a selection effect, since we are comparing the overall population of SMGs at various redshifts with $K$-band selected galaxies in the range $z\\sim1.4-2.5$. According to \\citet{Chapman2005}, $\\sim55\\%$ of the radio-identified, bright SMG population lies in this redshift range. From our sample of fifteen MAMBO sources detected with a significance $>4\\sigma$, eleven were identified to have a radio counterpart. Following Chapman et al., we estimate that about six of these eleven MAMBO sources should be at $z=1.4-2.5$, while using the photometric redshifts reported by \\citet{Bertoldi2007}, seven appear to lie in this redshift range. This implies that $\\sim$30\\% of the radio-identified, significant MAMBO sources at $z=1.4-2.5$ are associated with substancial overdensities of massive galaxies at these redshifts.  Note that our study is biased in that we are selecting massive galaxies at $z=1.4-2.5$. Our $K<23$ limit roughly translates into stellar masses $\\gtrsim (2-4)\\times10^{10}\\ M_{\\sun}$ \\citep{Daddi2004}, and therefore we miss less massive galaxies that could be associated with SMGs at these redshifts. However, the fact remains that only a fraction of the SMGs are related to groups of massive galaxies in the crucial epoch of galaxy assembly.    \\begin{figure}[!t] \\centering \\includegraphics[scale=0.5]{figure5.ps} \\caption{Redshift distribution for galaxies located within 30\\arcsec \\ from the density peak of high-redshift galaxies near MAMBO sources. The dotted histogram shows the photometric redshifts for our $K$-band selected galaxies, the dark-gray histogram shows the redshifts for the $K$-band selected high-redshift $BzK$ galaxies and the black single object shows the redshift of the most likely counterpart to the MAMBO source as given by \\citet{Bertoldi2007}. For COSBO-16, the hashed entry represents the secondary solution from the photometric redshift computation. \\label{fig:photo_z}} \\end{figure}    We compared the distribution of densities at the position of $>4\\sigma$ MAMBO source detections with the distribution of densities of BzK galaxies in the field (Fig. \\ref{fig:density_map_highz}). A Kolmogorov-Smirnov (KS) test does not reject the null hypothesis that both samples follow the same distribution at a 46\\% significance level. Restricting the $>4\\sigma$ MAMBO source sample to sources at $z=1.4-2.5$, the KS test does not reject the null hypothesis at a 43\\% level. This is somewhat inconsistent with the fact that two of the brightest MAMBO galaxies are related to some of the strongest overdensities in the field, however it may merely reflect the low number of SMGs that we used to compute the KS statistic. For the $3.5-4.0\\sigma$ MAMBO source sample, the KS-test does not rule out the null hypothesis at a 23\\% level, whereas if we limit this sample to galaxies with $z=1.4-2.5$, the KS test gives a 3.4\\% significance level. This strongly suggests that the densities at the position of the fainter MAMBO sources follow a different distribution from that of BzK galaxies in the field. For both samples combined (all MAMBO sources), the KS test gives only a 9.6\\% probability that they follow the same distribution of densities of BzK galaxies in the field, which is similar to the value that we obtain if we restrict the sample to $z=1.4-2.5$, 12.6\\%, again suggesting that the distributions are different.  Overall, our results show that only a fraction (30\\%) of MAMBO sources at $z=1.4-2.5$ is located in strongly overdense regions. This suggests that only some SMGs are linked to the formation of structures at high-redshift. Although we find a hint that some SMGs could be located in environments denser than that of the general population of galaxies at $z\\sim2$, it is not possible to discern whether this is a real trend or is due to cosmic variance as our analysis is based only on a handful of SMGs.  The presence of bright SMGs in galaxy overdensities is possibly related to the fact that the density peaks of high-redshift $BzK$ galaxies associated with MAMBO galaxies (e.g. for COSBO-3 and 6) are the strongest in the whole COSBO area. Denser groups of star-forming $BzK$ galaxies are more likely to produce major mergers between galaxies and thus are prone to trigger violent star-formation activity. The merger of two gas-rich $BzK$ galaxies close to the center of a galaxy group could induce a starburst that would be seen as a bright SMG. Star-forming $BzK$ galaxies have large reservoirs of molecular gas \\citep[$\\sim10^{11} M_{\\sun}$;][]{Daddi2008}, and could easily sustain the typical star-formation rates observed in SMGs ($\\sim1000\\ M_{\\sun}$ yr$^{-1}$) for $\\lesssim100$ Myr. The densest galaxy groups are thus ideal for the formation of bright submillimeter activity.   \\subsection{Comparison with other studies}  Similar studies relating the distribution of SMGs with the large scale structures traced by optically selected galaxies at high-redshift have recently been done. In particular, \\citet{Tamura2009} found a strong overdensity of SMGs toward a massive protocluster of Lyman-$\\alpha$ emitters at $z=3.1$, reflecting a strong link between SMGs and the large scale structure. Nevertheless, this study is based on a proto-cluster field where we know that strong clustering is taking place.  Studies of the galaxy-galaxy angular correlation function in blank-fields indicate that SMGs and IR luminous galaxies at $z\\sim2$ are clustered on typical angular scales of $15-25\\arcsec$, being related to massive dark matter halos \\citep[$\\sim10^{13}\\ M_{\\sun}$;][]{Blain2004, Greve2004,Scott2006, Farrah2006, Viero2009, Weiss2009}. Our results are consistent with these results in that strong clustering between SMGs and BzK galaxies occur on similar angular scales. We note, however, that studies purely based on the angular two-point correlation function are only able to measure the average clustering properties of SMGs. They miss the important fact that not all the SMGs are located in clustered environments, as we find in this paper, and therefore only a few of them will significantly contribute to the clustering signal of the angular correlation function in small scales."
    },
    "0911/0911.2706.txt": {
        "abstract": "We present a new analysis of the Jupiter+Saturn analog system, OGLE-2006-BLG-109Lb,c, which was the first double planet system discovered with the gravitational microlensing method. This is the only multi-planet system discovered by any method with measured masses for the star and both planets. In addition to the signatures of two planets, this event also exhibits a microlensing parallax signature and finite source effects that provide a direct measure of the masses of the star and planets, and the expected brightness of the host star is confirmed by Keck AO imaging, yielding masses of $M_\\ast = 0.51{+0.05\\atop -0.04}\\,\\msun$, $M_b = 231\\pm 19\\,\\mearth$, and $M_c = 86\\pm 7\\,\\mearth$. The Saturn-analog planet in this system had a planetary light curve deviation that lasted for 11 days, and as a result, the effects of the orbital motion are visible in the microlensing light curve. We find that four of the six orbital parameters are tightly constrained and that a fifth parameter, the orbital acceleration, is weakly constrained. No orbital information is available for the Jupiter-analog planet, but its presence helps to constrain the orbital motion of the Saturn-analog planet. Assuming co-planar orbits, we find an orbital eccentricity of $\\epsilon = 0.15 {+0.17\\atop -0.10}$ and an orbital inclination of $i = 64^\\circ {+4^\\circ \\atop -7^\\circ}$. The 95\\% confidence level lower limit on the inclination of $i > 49^\\circ$ implies that this planetary system can be detected and studied via radial velocity measurements using a telescope of $\\simgt 30\\,$m aperture. ",
        "introduction": "\\label{sec-intro} The discovery of the Jupiter/Saturn analog planetary system, OGLE-2006-BLG-109Lb,c, by gravitational microlensing \\citep{gaudi-ogle109} suggests that solar systems like our own may be common, at least among systems that contain gas giant planets. This was generally assumed to be the case prior to the detection of the first extrasolar planets orbiting main sequence stars \\citep{51peg,70vir,47uma}, but these first discoveries challenged this conventional wisdom, with the discovery of hot-Jupiters and gas giants in highly eccentric orbits. However, the discovery of true solar system analogs with the radial velocity method is difficult, requiring radial velocity precision $< 3\\,{\\rm m/s}$ spanning a decade or more \\citep{jup-twin}, so it may be that the apparent paucity of systems like our own is a selection effect.  The gravitational microlensing method has very different selection effects from the radial velocity method. It is most sensitive to planets at separations similar to the Einstein radius, which in most cases is just beyond the ``snow line\" at $\\sim 2.7 M/\\msun$ \\citep{ida_lin,lecar_snowline,kennedy-searth,kennedy_snowline} where gas giant planets are expected to form, according to the core accretion theory. Microlensing does find that super-Earth and Neptune-mass planets \\citep{ogle390,ogle169,bennett-moa192} are more common at these separations than Jupiters. Of the handful of microlensing events that reveal gas giant planets \\citep{bond-moa53,ogle71}, only OGLE-2006-BLG-109 and MOA-2007-BLG-400 \\citep{dong-moa400} are highly sensitive to multiple planets of Jupiter mass or less. The fact that OGLE-2006-BLG-109 did reveal a Jupiter-Saturn-like system suggests that systems like our own are common among stars hosting gas giants near or beyond the snow line.  The OGLE-2006-BLG-109L system is the first multi-planet system found by microlensing but several other aspects of this event are also unique. The light curve reveals the microlensing parallax effect, which yields a geometric mass measurement of the planetary host star and both of the planets. In addition, the host star is five times brighter than the source in the $H$ band, and it is detected in Keck AO images. The measured $H$-band brightness confirms the microlensing parallax mass measurement.  In addition, the light curve shape is sensitive to the relative positions of the planet at the pico-arcsecond level, and the light curve signal of the Saturn-mass planet is visible for 11 days. As a result, the orbital velocity of this planet in the plane of the sky must be included to model the light curve. Moreover, these effects also provide a weak constraint on the orbital acceleration.  In this paper, we present some details of the analysis that was summarized by \\citet{gaudi-ogle109}, and a new analysis of OGLE-2006-BLG-109 that includes the orbital motion of the OGLE-2006-BLG-109L system. With the assumption of co-planar planetary orbits and orbital stability constraints, we derive limits on the full set of orbital parameters for this Saturn-mass planet.  This paper is organized as follows. In \\S~\\ref{sec-data} we discuss the image data and the photometric reductions. The modeling of the light curve is described in \\S~\\ref{sec-model}, and a new method for determining the angular radii of the source star and the Einstein Ring is presented in \\S~\\ref{sec-radius}. The follow-up observations that identify the planetary host star and their analysis are discussed in \\S~\\ref{sec-follow}, and \\S~\\ref{sec-alt} discusses possible alternative models for the light curve. The next section, \\S~\\ref{sec-orbits}, discusses the orbital motion modeling and the conversion from measured fit parameters to inferred orbital parameters, and the Bayesian analysis used to find the constraints on the physical orbital parameters is discussed in \\S~\\ref{sec-param}. Limits on additional planets in the system are discussed in \\S~\\ref{sec-limits}, and we conclude in \\S~\\ref{sec-conclude}. ",
        "conclusions": "\\label{sec-conclude}  We have presented the complete analysis of the OGLE-2006-BLG-109Lb,c planetary system that was summarized by \\citet{gaudi-ogle109}, and we have explored the physics implications of the microlensing fit parameters. We have introduced a new method to determine the angular radius of the source star, $\\theta_*$, which makes use of three color, $VIH$, light curve data to determine $\\theta_*$ to a precision of 2.8\\%. We then combine this measurement with the microlensing parallax parameter, $\\pi_E$, determined from the microlensing light curve model to yield a direct, geometric, measurement of the masses of the star and two planets in this planetary system. This is the only multiple planet system with measured masses for the stars and planets.  (The pulsar planet system PSR B1257+12 \\citep{pulsar_planets03} has measured planet masses, but the mass of the host neutron star is assumed.) The mass measurement of the host star is confirmed by Keck AO images, which detect the planetary host star, and measure its $H$-band magnitude to be $H_L = 17.17\\pm 0.05$.  These results, plus our more sophisticated modeling, including terrestrial parallax, confirm the results of \\citet{gaudi-ogle109} that this system resembles a smaller version of our own solar system with a primary about half the mass of the Sun orbited by planets slightly smaller than Jupiter and Saturn in a similar arrangement with the Saturn-like planet outside. Their semimajor axes are about half of those of Jupiter and Saturn. This similarity to Jupiter and Saturn is even greater if we consider the configuration of the system during the process of planet formation at an age of $\\sim 1\\,$million years, when nuclear fusion does not yet dominate the host star's energy production. At this time, the stellar luminosity is thought to scale as $\\sim M^2$ \\citep{burrows93,burrows97} (G. Kennedy, C. Lada, 2007, private communications). In the core accretion theory, the most massive planets are thought to form outside the ``snow line\" where it is cold enough for water-ice to form \\citep{ida_lin,lecar_snowline,kennedy-searth,kennedy_snowline}, and this predicted $\\sim M^2$ dependence of the proto-star luminosity implies that the location of the ``snow line\" should scale linearly with stellar mass. So, if our solar system's snow line was at $2.7\\,$AU, which is slightly larger than half of Jupiter's semimajor axis, then we estimate that the snow line for OGLE-2006-BLG-109L should have been at $\\sim 1.4\\,$AU, which is slightly larger than half the semimajor axis of its Jupiter-analog planet.  Due to the relatively long, 11-day, duration of the signal of the Saturn-analog planet, the light curve tightly constrains four of the six parameters that are needed to describe the planetary orbit. A fifth orbital parameter is weakly constrained. We use these light curve constraints in an MCMC analysis, along with orbital stability constraints for the combined Jupiter/Saturn-analog system, to determine the constraints on the orbital parameters implied by the light curve model. Assuming coplanar orbits, we find an orbital inclination of $i_c = 64^\\circ {+4^\\circ\\atop -8^\\circ}$, with a 2-$\\sigma$ lower limit of $49^\\circ$. Because the lens star is relatively bright and brighter than the source, this means that it should be possible to confirm at least the inner, Jupiter-analog, planet with Doppler radial velocity measurements. While, the host star, at $H_L = 17.17\\pm 0.05$ is much fainter than the host star for any known planet discovered with the radial velocity method, it seems reasonable to expect that the next generation of very large telescopes \\citep{tmt,e-elt-sci,e-elt-stat}, combined with a high-throughput, high resolution spectrograph \\citep{mayor_harps} would be able to detect the radial velocity signal of the Jupiter-analog planet with a radial velocity amplitude of $K_b = 17.4 {+1.3\\atop -1.1}\\,$m/sec. This would be a challenging measurement for these extremely large telescopes, but our prior knowledge of the planetary system will minimize the number of observations that would be required. It is likely to be at least a decade and perhaps two decades before such measurements are possible, but this is much sooner than the $\\simgt 10^6\\,$yr that we would likely have to wait before either of these planets reveals themselves again via microlensing.  With our measurements of the planetary masses and constraints on the orbital parameters, OGLE-2006-BLG-109Lb,c is arguably the best constrained multiplanet system to orbit a main sequence star, other than the Sun. None of the other known multiple exoplanet systems has measured masses, and radial velocity measurements generally do not constrain more parameters of the planetary orbits than we have constrained for OGLE-2006-BLG-109c. On the other hand, radial velocity surveys often measure the eccentricity of the planets they detect more precisely than we have done, and the eccentricity is intrinsically more interesting than the parameters we have measured describing the orientation of the orbital plane of a single planet. Nevertheless, this is certainly the first planetary system found by microlensing in which the eccentricity of a planetary orbit is constrained.  It seems likely that we will be able to extract even more information on exoplanetary systems from microlensing events discovered in the near future. There have been two major telescope hardware improvements that have occurred since the discovery of OGLE-2006-BLG-109 that should significantly improve the light curve coverage for future events. Later in the 2006 observing season, the MOA Collaboration brought its 1.8m aperture MOA-II telescope \\citep{moa2_tel} online equipped with the $2.2\\,{\\rm deg}^2$ MOAcam-3 CCD camera \\citep{sako_moacam3}. This allows an observing cadence as high as one observation every 15 minutes for the central Galactic bulge fields. OGLE has just completed an upgrade to its $1.4\\,{\\rm deg}^2$ OGLE-IV camera \\citep{udalski_bohdan} that will enable to increase its sampling rate in a similar fashion. If OGLE-IV had been in operation when OGLE-2006-BLG-109 was discovered, feature 1 in Figures~\\ref{fig-lc1} and \\ref{fig-lc3} would have been recognized much sooner, which would likely have resulted in many more observations of this feature. Similarly, the larger MOA-II telescope and higher observing cadence enabled by MOA-cam-3 would have significantly improved the quantity and quality of the MOA data. As a result, it is very likely that if a similar event were detected today, we would be able to tightly constrain the fifth orbital motion parameter ($T_{\\rm orb}$ in the parameterization we use here) and to weakly constrain the final orbital parameter. This would yield much tighter constraints on standard orbital parameters. So, for a fraction of the future exoplanets discovered by microlensing, we can expect better measurements of the orbital parameters than the ones we present here for OGLE-2006-BLG-109."
    },
    "0911/0911.4012_arXiv.txt": {
        "abstract": "{We describe a recent full-polarization radio continuum survey, performed using the Westerbork Synthesis Radio Telescope (WSRT), of several nearby galaxies in the Spitzer Infrared Nearby Galaxies Survey (SINGS) sample. The WSRT-SINGS survey has been utilized to study the polarized emission and Faraday rotation measures (RMs) in the targets, and reveals an important new observational trend. The azimuthal distribution of polarized flux seems to be intimately related to the kinematic orientation of galaxies, such that in face-on galaxies the lowest level of polarized flux is detected along the kinematic major axis. In highly inclined galaxies, the polarized flux is minimized on both ends of the major axis, and peaks near the minor axis. Using models of various three-dimensional magnetic field geometries, and including the effects of turbulent depolarization in the midplane, we are able to reproduce the qualitative distribution of polarized flux in the target galaxies, its variation with inclination, and the distribution of RMs, thereby constraining the global magnetic field structure in galaxies. Future radio telescope facilities, now being planned and constructed, will have properties making them extremely well-suited to perform vastly larger surveys of this type, and are thereby poised to significantly increase our understanding of the global structure of galactic magnetic fields. We discuss progress that can be made using surveys which will be realized with these new facilities, focusing in particular on the Aperture Tile in Focus (APERTIF) and Australian Square Kilometre Array Pathfinder (ASKAP) telescopes, both based on Focal Plane Array (FPA) designs, which are expected to be particularly useful for wide-field polarization applications.}  \\FullConference{Panoramic Radio Astronomy: Wide-field 1-2 GHz research on galaxy evolution - PRA2009\\\\ June 02 - 05 2009\\\\ Groningen, the Netherlands}  \\begin{document} ",
        "introduction": " ",
        "conclusions": ""
    },
    "0911/0911.1637_arXiv.txt": {
        "abstract": "We present high speed photometric observations of the eclipsing dwarf nova IP Peg taken with the triple-beam camera ULTRACAM mounted on the William Herschel Telescope. The primary eclipse in this system was observed twice in 2004, and then a further sixteen times over a three week period in 2005. Our observations were simultaneous in the Sloan $u'$, $g'$ and $r'$ bands. By phase-folding and averaging our data we make the first significant detection of the white dwarf ingress in this system and find the phase width $\\phi$ of the white dwarf eclipse to be $0.0935 \\pm 0.0003$, significantly higher than the previous best value of $0.0863 < \\phi < 0.0918$. The mass ratio  is found to be $q = M_2 / M_1 = 0.48 \\pm 0.01$, consistent with previous measurements, but we find the inclination to be $83.8 \\pm 0.5$ deg, significantly higher than previously reported. We find the radius of the white dwarf to be $0.0063 \\pm 0.0003$\\Rsun, implying a white dwarf mass of $1.16 \\pm 0.02$\\Msun. The donor mass is $0.55 \\pm 0.02$\\Msun. The white dwarf temperature is more difficult to determine, since the white dwarf is seen to vary significantly in flux, even between consecutive eclipses. This is seen particularly in the $u'$-band, and is probably the result of absorption by disc material. Our best estimate of the temperature is $10,000$ -- $15,000$K, which is much lower than would be expected for a CV with this period, and implies a mean accretion rate of $< 5 \\times 10^{-11} M_{\\odot}$ yr$^{-1}$, more than $40$ times lower than the expected rate. ",
        "introduction": "Cataclysmic variable stars (CVs: \\citealt{Warner95}) provide examples of white dwarfs accreting from low mass companions at rates of $\\sim$$10^{-11}$ -- $10^{-9}$\\Msun/yr. The mass transfer in CVs and similar semi-detached binary systems is driven by angular momentum loss. For long orbital period ($\\sim$$3$ h or greater) systems the angular momentum loss is thought to be driven by magnetically-coupled stellar winds. For the shorter period ($< 3$ h) systems losses due to  gravitational radiation are thought to dominate the mass transfer process. This is observationally supported to some extent: the shorter period population is dominated by lower \\Mdot \\ dwarf novae, whereas the peak of the population of high \\Mdot \\ `steady state' CVs is in the $3$ to $4$h period range \\citep{Shafter92}. However, there is no one-to-one relationship. The  distribution of steady-state, bright CVs  tails off sharply at longer periods \\citep{Ritter03}, and a number of dwarf novae  are observed to have very long periods. It is generally thought that all non-magnetic CVs are essentially the same, with only a change in \\Mdot \\ required to move between the different subclasses. While it may be the case that the high- and low-\\Mdot \\ systems at a given orbital period are different populations with a different evolutionary history, a more likely explanation is that the instantaneous \\Mdot \\ we measure is a poor indicator of the mean \\Mdot, so that even high mean \\Mdot \\ systems spend some fraction of their time as slowly accreting dwarf novae.  One independent means of inferring a longer-term ($10^3$ -- $10^5$yr) value of \\Mdot \\  is through measurement of the white dwarf temperature. The temperature is a good tracer of the long-term \\Mdot \\ since it is determined not by the accretion heating, but by the compression of the underlying white dwarf by the accreted matter \\citep{Townsley03, Townsley09}. The existing measurements agree well with the high-\\Mdot/long period, low-\\Mdot/short period consensus (figure 5 of \\citealt{Townsley09}). However, these data are likely to be biased. These white dwarf temperatures are spectroscopically determined, and the white dwarf must dominate at UV wavelengths in order for the temperature to be determined. Long period CVs are generally brighter than shorter period systems due to their higher accretion rates and more luminous donor stars, and so the white dwarfs in long period systems must be hotter in order to be measurable. An alternative strategy which reduces this bias is to make fast, multi wavelength photometric observations of CVs where the inclination is high enough for the white dwarf to be eclipsed. With sufficient time resolution the ingress and egress of the white dwarf can be identified independently of the other components, and hence the white dwarf colours can be determined.  An ideal system for a study of this nature is IP Pegasi (IP Peg hereafter). This dwarf nova was found to be eclipsing by \\citet{Goranskii85a} and has a period of $3.8$h, in the `long period' tail of the dwarf nova distribution. Following \\citet{Townsley03} this would imply a high mean \\Mdot \\ and a correspondingly high white dwarf temperature. However, the temperature of the white dwarf has yet to be determined with precision. Previous studies of IP Peg have shown that the white dwarf ingress feature is obscured by the ingress of the bright spot on the accretion disc, and it has been claimed that the white dwarf egress feature varies considerably in size and length \\citep{Wood86, Wood89}. The difficulty in detecting the white dwarf throws the system parameters of IP Peg into doubt. In an attempt to resolve this and to measure the temperature of the white dwarf we took a total of eighteen separate observations of the IP Peg eclipse with the high speed CCD camera ULTRACAM mounted on the William Herschel Telescope (WHT). In this paper we present these data, and our subsequent parameter determinations. ",
        "conclusions": "\\label{sec:conc} In 2004 and 2005 we took eighteen separate observations of the eclipse in the $3.8$h period dwarf nova IP Peg. We used the triple-band camera ULTRACAM mounted on the William Herschel Telescope, and observed simultaneously in the Sloan $u'$-, $g'$- and $r'$-bands. By phase-folding these data we identify for the first time the white dwarf ingress in this system. The phase width of the white dwarf eclipse is $0.0935 \\pm 0.0003$, significantly higher than the standard value used by previous authors. This new value implies a higher orbital inclination or a higher mass ratio or some combination of the two. Our model fits strongly favour a high orbital inclination, which we find to be $83.81 \\pm 0.45$ deg. Our determinations of the masses of the two components are higher than those reported by  previous authors.  We find the brightness of the white dwarf egress feature to be highly variable in flux and colour. We suggest this is due to photoelectric absorption as a result of obscuration of the white dwarf by material from the accretion disc. We find the duration of the egress feature to be consistent in length, and so suggest that past reports of very extended egresses were due to failures to detect the egress at all, as a result of this varying obscuration.  It is difficult to determine an accurate temperature for the white dwarf due to the obscuration, but by comparing our data with synthetic white dwarf models we find the most likely temperature to be in the range of $10,000$ - $15,000$K. This is very low for a dwarf nova above the period gap and, following \\citet{Townsley03}, implies a mean accretion rate over the past $1,000$ - $10,000$ years of $< 5 \\times 10^{-11} M_{\\odot}$ yr$^{-1}$. This is more than $40$ times lower than the expected rate. Additionally we find a mass and radius for the donor that are consistent with those of a main sequence star in thermal equilibrium, which also suggests a very low accretion rate. Unless the mass transfer in IP Peg has begun very recently, these findings imply either that CVs can sustain accretion rates well below the expected rate for very long periods of time, or that the classical picture whereby the long-term accretion rate scales with period is not correct."
    },
    "0911/0911.1895_arXiv.txt": {
        "abstract": "Most simulations of coronal mass ejections (CMEs) to date either focus on the interplanetary propagation of a giant plasma ``blob'' without paying too much attention to its origin and to the formation process or they focus on the complex evolution of the coronal magnetic field due to (sub-)photospheric motions which result in an eruption. Here, we present global simulations of CMEs where coronal motions are used to produce a realistic evolution of the coronal magnetic field and cause an eruption. We focus on active region 10069, which produced a number of eruptions in late August 2002, including the August 24, 2002 CME -- a fast ($\\sim$ 2000 km~s$^{-1}$) eruption originating from W81--, as well as a slower eruption on August 22, 2002 (originating from W62). Using a three-dimensional magneto-hydrodynamic (MHD) simulation of these ejections with the Space Weather Modeling Framework (SWMF), we show how a realistic initiation mechanism enables us to study the deflection of the CME in the corona and in the heliosphere. Reconnection of the erupting magnetic field with that of neighboring streamers and active regions modify the solar connectivity of the field lines connecting to Earth and change the expected solar energetic particle fluxes. Comparing the results at 1 AU of our simulations with {\\it in situ} observations by the {\\it ACE} spacecraft, we propose an alternate solar origin for the shock wave observed at L1 on August 26. ",
        "introduction": "\\subsection{Active Region 10069 and August Eruptions} On August 24, 2002, active region (AR) 10069 was near the western limb of the Sun (W81) when it was the source of a fast ($\\sim$ 1,900~km~s$^{-1}$) and wide coronal mass ejection (CME). This event has been well studied due to extensive observations remotely by LASCO and UVCS \\citep[]{Raymond:2003} as well in-situ by the {\\it Wind} and {\\it ACE} spacecraft. Most importantly, it was associated with an intense Solar Energetic Particle (SEP) event \\citep[e.g., see][]{Tylka:2006, Tan:2009}. The eruption was preceded 2 days earlier by a similar, but twice slower CME, which was also associated with a large SEP event but is not thought to have been associated with a shock wave at 1 AU. Particularly noteworthy is the fact that, at the start of the August 24 event, the proton flux at energies above 10 MeV was still ten times above the intensity before the August 22 CME.  Western limb events such as the August 24, 2002 CME present a number of challenges for space weather prediction. Due to the Parker spiral, the Earth is magnetically connected with regions on the solar surface at around W55 (for a 400~km~s$^{-1}$ wind). Therefore, SEP events are preferentially associated with western-limb events. Events such as the August 24, 2002 CME present additional challenges since the SEP arrival time at Earth corresponds to a particle release height of less than 5~$R_\\odot$ \\citep[]{Tylka:2006}. To explain the observations, a shock wave must be formed low in the corona at the flanks of the CME or the magnetic field line connecting the Earth to the solar surface must significantly diverge from the nominal Parker spiral. Such large differences of up to 30$^\\circ$ between flare sites and the footprint tracing of the magnetic field lines connecting the Earth to the Sun have been reported before \\citep[]{Ippolito:2005}.  Western limb events can be associated with ejecta and more often with shock waves at 1~AU \\citep[]{Cane:1988}. It is believed that the August 24, 2002 CME was associated with a shock wave at 1~AU on August 26. This fact, again, seems to imply either a very large span of the shock wave, a large deflection of the CME, or a combination of both. However, in \\citet{Lugaz:2009a}, we studied numerically the August 24, 2002 eruption, and our model could not explain the expansion and/or deflection required to associate the shock at 1~AU to the eruption on August 24. In this article, we investigate i) why the August 22 CME, which originated from the same region when it was closer to the disk center, was not associated with a shock at 1 AU, and ii) the importance of this preceding ejection on the transport of SEPs accelerated by the August 24 CME.  \\subsection{Solar Wind and Solar Eruption Models}  \\begin{figure}[t] \\begin{minipage}[]{1.0\\linewidth} \\begin{center} {\\includegraphics*[width=7.1cm]{0min_sw12.pdf}} {\\includegraphics*[width=7.1cm]{1hrfieldlines.pdf}} \\end{center} \\end{minipage}\\hfill \\caption{{\\it Top}: Magnetic topology of the solar corona in the vicinity of AR 10069 on August 22-24, 2002. {\\it Bottom}: Snapshot of the eruption at the end of the shearing phase (t = 1 hour) showing magnetic field lines color-coded with the radial velocity. The reconnection between the erupting flux from the dipole and the open magnetic field line happens at eNP and is highlighted with the white circle.} % \\end{figure}  Our simulation setup is similar to that discussed in more details in \\citet{Roussev:2007}, \\citet{Jacobs:2009} and \\citet{Lugaz:2009a}. In our numerical model, the steady-state solar corona and solar wind are constructed following the methodology of \\citet{Roussev:2003a}. The initial condition for the coronal magnetic field is calculated by means of potential field extrapolation with boundary conditions for the radial magnetic field at the Sun provided by full-disk SoHO/MDI observations. The plasma parameters are prescribed in an ad-hoc manner, through a variable polytropic index, in order to mimic the physical properties of streamers and coronal holes once a (non-potential) steady state is reached.  To the initial magnetic field constrained by MDI data, we superimpose newly emerged magnetic flux given by the dipolar magnetic field of two point charges, whose magnitude is chosen so that the peak value of the radial magnetic field at the solar surface is about 47~Gauss. To initiate the eruption, we use a similar strategy as described in \\citet{Roussev:2007}. To summarize, once the steady-state is reached at $t=0$, the two subsurface magnetic charges are moved apart quasi-steadily up to $t=t_S=60$~min with a speed which is ramped up in $t_S/3$ to V$_q$ = 33~km~s$^{-1}$; the charge motion is stopped at $t=t_S$. The two charges are moved approximatively parallel to the polarity inversion line. In order to provide sufficient buildup of magnetic energy, accompanying horizontal boundary motions are imposed as well. ",
        "conclusions": "We have performed a Sun-to-Earth simulation of the August 22, 2002 CME event with a realistic CME initiation mechanism \\citep[]{Roussev:2007}, following a similar simulation of the faster and more studied August 24, 2002 CME \\citep[]{Lugaz:2009a}. We have found that, although the August 22 CME was not detected at Earth, it is likely that the presence of its associated shock wave and ejecta in the heliosphere significantly modified the length of the field lines connected to Earth at the start of the August 24 CME. This, in turn, must be taken into account when deriving the SEP release time and the shock formation height. We have also investigated the origin of the shock wave observed in situ at 1 AU on August 26 near 14UT. Our simulations are not able to reproduce the correct transit time and hit/miss from these two eruptions: the August 24, 2002 ejection is not wide enough to hit Earth in our simulation and the August 22, 2002 CME is predicted to hit Earth but 1.5 days earlier than what was detected. It is possible that complex interaction in the heliosphere between these two eruptions and other eruptions in the period from August 22-24, 2002 modified the transit time of one of the shocks to explain the observations at Earth \\citep[as described in][]{Dasso:2009}. A disk-centered eruption on August 23 may play an important role in the formation of the detected shock.  This study has shown the advantages and the limitations of modeling to study complex events such as fast, wide CMEs from the limb. The simulation can give us an estimate of the length of the field line connected to the Earth which can be, for example, compared to the ones computed from anisotropic measurements \\citep[]{ERNE:2005} and must be used to derive SEP release time and the height at which the CME-driven shock wave was formed. Here, we have found that the length of the field line decreases by about 20$\\%$. On the other hand, the model is not able to help to explain the arrival of a shock wave 2.5 days after the August 24 CME at a location about 80$^\\circ$ east of the flare location. Although such Sun-Earth events originating from the western limb are rare, it will be important to better understand them in the future in order to have more reliable space weather prediction.     \\begin{theacknowledgments} The research for this manuscript was supported by NSF grants ATM0639335 and ATM0819653 and NASA grants NNX07AC13G and NNX08AQ16G. The SWMF was developed by the CSEM with funding support from NASA, NSF and DoD. \\end{theacknowledgments}"
    },
    "0911/0911.4154_arXiv.txt": {
        "abstract": "The discovery of Main Belt Comets (MBCs) has raised many questions regarding the origin and activation mechanism of these objects. Results of a study of the  dynamics of these bodies suggest that MBCs were formed in-situ as the remnants of the break-up of large icy asteroids. Simulations show that similar to the asteroids in the main belt, MBCs with orbital eccentricities smaller than 0.2 and inclinations lower than $25^\\circ$ have stable orbits implying that many MBCs with initially larger eccentricities and inclinations might have been scattered to other regions of the asteroid belt. Among scattered MBCs, approximately 20\\% reach the region of terrestrial planets where they might have contributed to the accumulation of water on Earth. Simulations also show that collisions among MBCs and small objects could have played an important role in triggering the cometary activity of these bodies. Such collisions might have exposed sub-surface water ice which sublimated and created thin atmospheres and tails around MBCs. This paper discusses the results of numerical studies of the dynamics of MBCs and their implications for the origin of these objects. The results of a large numerical modeling of the collisions of m-sized bodies with km-sized asteroids in the outer part of the asteroid belt are also presented and the viability of the collision-triggering activation scenario is discussed. ",
        "introduction": "The discovery of comet-like activities in four icy asteroids 7968 Elst-Pizzaro (133P/Elst-Pizzaro), 118401 (1999 ${{\\rm RE}_{70}}$, 176P/LINEAR), P/2005 U1 (Read), and P/2008 R1 (Garradd) has added a new item to the mysteries of the asteroid belt (\\cite[Hsieh \\& Jewitt 2006]{Hsieh06}; \\cite[Jewitt, Yang \\& Haghighipour 2009]{Jewitt09}). Known as Main Belt Comets (MBCs), these objects may be representatives of a new class of bodies that are dynamically asteroidal (i.e., their Tisserand parameters\\footnote{For a small object, such as an asteroid, that is subject to the gravitational attraction of a central star and the perturbation of a planetary body P, the quantity ${a_{\\rm P}}/a \\, + \\,2 {[(1-e^2)\\,a/{a_{\\rm P}}]^{1/2}}\\, \\cos i$ is defined as its Tisserand parameter . In this formula, $a$ is the semimajor axis of the object with respect to the star, $e$ is its orbital eccentricity, $i$ is its orbital inclination, and $a_{\\rm P}$ is the semimajor axis of the planet. In the solar system, the Tisserand parameter of a small body with respect to Jupiter can be used to determine the cometary or asteroidal nature of its orbit. In general, the Tisserand parameters of comets with respect to Jupiter are smaller than 3, whereas those of asteroids are mostly larger.}  are larger than 3), but have cometary appearance. As shown in Table 1, the orbits of these objects are in the outer half of the asteroid belt (Fig. 1) implying that they may contain sub-surface water ice. In fact the observation of the tail of 7968 Elst-Pizzaro by \\cite[Hsieh, Jewitt \\& Fern\\'andez (2004)]{Hsieh04} has indicated that the comet-like activity of this MBC is episodic (it is not the ejection of dust particles that were produced through an impact to this object) and is due to the dust particles that have been blown off the surface of this body by the drag force of the gas that was most likely produced by the sublimation of near-surface water ice.   \\begin{table} \\begin{center} \\caption{Orbital Elements of MBCs (\\cite[Hsieh \\& Jewitt 2006]{Hsieh06})} \\label{tab1} {\\scriptsize \\begin{tabular}{|l|c|c|c|c|c|}\\hline {\\bf MBC} & {\\bf $a$ (AU)} & {\\bf $e$} & {\\bf $i$ (deg.)} & {\\bf Tisserand} & {\\bf Diameter (km)} \\\\ \\hline (133P)/7968 Elst-Pizzaro     & 3.156  & 0.165  & 1.39  & 3.184  & 5.0  \\\\ \\hline 118401 (176P/LINEAR) & 3.196  & 0.192  & 0.24  & 3.166  & 4.4 \\\\ \\hline P/2005 U1 (Read) & 3.165  & 0.253  & 1.27  & 3.153  & 0.6\\\\ \\hline P/2008 R1 (Garradd) & 2.726 & 0.342 & 15.9 & 3.216 & 1.4 \\\\ \\hline \\end{tabular} } \\end{center} \\end{table}   The comet-like appearance of MBCs has raised questions regarding the origin of these objects. While the asteroidal orbits of these bodies, combined with the proximity of 7968 Elst-Pizzaro, 118401 (176P/LINEAR), and P/2005 U1 (Read) to the Themis and Beagle families of asteroids (Fig. 1), suggests that MBCs have formed in-situ as the remnants of collisionally broken larger objects, the cometary activities of these bodies may be taken to argue that MBCs are comets that were scattered inward from the outer regions of the solar system and were captured in their current orbits. Such a capture mechanism could not have occurred recently. Simulation of the dynamics of Kuiper belt object by \\cite[Fernandez et al. (2002)]{Fernandez02} have shown that, at the current dynamical state of the solar system, it would not be possible to scatter comets from regions outside the orbit of Neptune to the main asteroid belt. A primordial capture, on the other hand, may not be impossible. Recently \\cite[Levison et al. (2009)]{Levison09} have shown that within the context of the Nice model (\\cite[Gomes et al. 2005]{Gomes05}; \\cite[Morbidelli et al. 2005]{Morbidelli05}; \\cite[Tsiganis et al. 2005]{Tsiganis05}) many trans-Neptunian objects could  have been scattered inwards and captured in orbits in the asteroid belt as close in as 2.68 AU during the early state of the dynamics of the solar system. Whether these objects could be the source of MBCs is, however, uncertain. This paper evaluates this possibility, in particular in comparison with the in-situ formation model,  by presenting the results of a numerical study if the dynamics of the currently known MBCs, and discussing their implications for the formation and origin of these objects.  As mentioned above, tails of MBCs are generated through the interactions of dust grains on the surface of these bodies with the gas produced by the sublimation of near-surface water ice (\\cite[Hsieh, Jewitt \\& Fern\\'andez 2004]{Hsieh04}). As shown by  \\cite[Schorghofer (2008)]{Schorghofer08}, asteroids in the region between 2 AU and 3.3 AU can maintain sub-surface water ice for several billion years if their surfaces are covered by a layer of dust, even as thick as only a few meters. That implies that in order for an MBC to start its cometary activation, this dusty layer has to be removed. \\cite[Hsieh \\& Jewitt (2006)]{Hsieh06} have suggested that collisions between MBCs and objects as small as a few meter in size, can reveal the sub-surface water ice. Such collisions will result in the local exposure of ice which sublimates and creates a thin atmosphere and tail for an MBC. An activated MBC, on the other hand, may terminate its activity after sublimating all the ice at the location of its collision with a m-sized projectile. It may also start new activation if it is collided with a second m-sized object in a later time again. In other words, an MBC may be activated several times till it exhausts all its water ice, or is scattered to an unstable orbit and either leaves the asteroid belt or collides with another asteroid or a planet. It will therefore be useful to study the rate of the collisions of m-size bodies with km-size MBCs, in particular in the outer region of the asteroid belt. This paper presents the results of such simulations and discusses their implications for the activation of MBCs and the possibility of the detection of more of these objects.  \\begin{figure} \\vspace*{0.2 cm} \\begin{center} \\includegraphics[width=4in]{fig1.eps} \\caption{The four currently known MBCs and the Themis and Beagle families of asteroids. As shown here, 7968 Elst-Pizzaro and 118401 (Read) are within the Themis and Beagle families whereas P/2005 U1 is in their proximity. The MBC P/2008 R1 (Garradd) seems to be an object that was scattered out of its forming region. The locations of mean-motion resonances with Jupiter are also shown.} \\label{fig1} \\end{center} \\end{figure}    \\section {Orbital Integrations and Implications for the Origin of MBCs}  To study the long-term stability of the four known MBCs, the orbits of these objects were integrated for 1 Gyr. Integrations included all the planets and Pluto, and treated MBCs as non-interacting objects. The effects of non-gravitational forces such as Yarkovsky, and the effect of the mass-loss of MBCs due to their cometary activities were not included. Since the activation of MBCs is episodic and intensifies during the perihelion passages of these objects, which is short compared to their orbital periods, the effect of the mass-loss may not alter the dynamics of these objects significantly. Integrations were carried out with Bulirsch-Stoer and with the Second-Order Mixed-Variable Symplectic (MVS) integrators in the N-body integration package MERCURY (\\cite[Chambers 1999]{Chambers99}). The initial orbital elements of the MBCs and planets were obtained from documentation on solar system dynamics published by the Jet Propulsion Laboratory (http://ssd.jpl.nasa.gov/?bodies). The timestep of each integration was set to 9 days.  Figure 2 shows the results of the simulations. As shown here 7968 Elst-Pizzaro and 118401 (176P/LINEAR) maintain their orbits for 1 Gyr. However, P/2005 U1 (Read) and P/2008 R1 (Garradd) become unstable in approximately 20 Myr. Integrations were also carried out for different initial values of the semimajor axes and eccentricities of MBCs, changing these quantities in increments of $\\Delta a=0.0001$ AU and $\\Delta e= 0.001$ within the ranges of their observational uncertainties. Similar results were obtained. 7968 Elst-Pizzaro and 118401 were stable whereas P/2005 U1 and P/2008 R1 became unstable in all simulations with a median lifetime of $\\sim$57 Myr. For more details on the results of the simulations, in particular on the analysis of the effects of mean-motion resonances on the dynamics of these MBCs, the reader is referred to \\cite[Haghighipour (2009)]{Hagh09} and \\cite[Jewitt, Yang \\& Haghighipour (2009)]{Jewitt09}.  \\begin{figure} \\vspace*{0.2cm} \\begin{center} \\includegraphics[width=2.5in]{f2a.eps} \\includegraphics[width=2.5in]{f2b.eps} \\includegraphics[width=2.5in]{f2c.eps} \\includegraphics[width=2.5in]{f2g.eps} \\caption{Graphs of the eccentricities, semimajor axes $(a)$, perihelion $(q)$, and aphelion $(Q)$ distances of 7968 Elst-Pizzaro, 118401 (176P/LINEAR), P/2005 U1 (Read), and P/2008 R1 (Garradd) . As shown here, 7968 Elst-Pizzaro and 118401 are stable for 1 Gyr whereas P/2005 U1 (Read) and P/2008 R1 (Garradd) become unstable in approximately 20 Myr.} \\label{fig2} \\end{center} \\end{figure}   As shown by Fig. 1, the orbit of P/2008 R1 (Garradd) is close to the influence zone of the 8:3 mean-motion resonance with Jupiter. Numerical simulations by \\cite[Jewitt, Yang \\& Haghighipour (2009)]{Jewitt09} have shown that the region in the vicinity of P/2008 R1 is dynamically unstable implying that this MBC must have formed in another region of the asteroid belt and scattered to its current orbit. The orbital instability of P/2005 U1 (Read), on the other hand, may show a pathway to such scattering events. The proximity of P/2005 U1 to the Themis family and the location of the 1:2 mean-motion resonance with Jupiter suggest that this MBC was perhaps formed close to the influence zone of  the 1:2 resonance. The original proximity of P/2005 U1 to this resonance has resulted in a gradual increase in its orbital eccentricity which will eventually make its orbit unstable. Such an instability might have also happened to the orbits of other MBCs and resulted in their scattering to other regions. To study this scenario, a large number of hypothetical MBCs were considered around the region where 7968 Elst-Pizzaro, 118401 (176P/LINEAR), and P/2005 U1 (Read) exist. The semimajor axes of these objects were varied between 3.14 AU and 3.24 AU, and their initial eccentricities were taken to be between 0 and 0.4. The initial orbital inclinations of these MBCs were chosen from a range of 0 to $40^\\circ$.  The orbits of these hypothetical MBCs were integrated for 100 Myr. Figure 3 shows the results. In this figure, green circles correspond to MBCs with stable orbits whereas purple indicates instability. As shown here, 7968 Elst-Pizzaro and 118401 (176P/LINEAR) are in the stable region of the graph whereas P/2005 U1 (Read) is approaching the unstable area.  \\begin{figure} \\vspace*{0.2 cm} \\begin{center} \\includegraphics[width=4in]{fig3a.eps} \\vskip 10pt \\includegraphics[width=4in]{fig3b.eps} \\caption{Top: graph of the stability of hypothetical MBCs in terms of their inclinations. The regions of secular resonances $\\nu_5$, $\\nu_6$, and $\\nu_{16}$ corresponding to an eccentricity of 0.1 are also shown. Bottom: graph of the stability of hypothetical MBCs in terms of their eccentricities. The brown area in the top graph and solid line in the bottom graph show the region of the 2:1 MMR with Jupiter. Circles in green correspond to initial semimajor axes and eccentricities of stable MBC whereas those in purple show instability. Similar to the asteroid in the main belt, objects with inclinations larger than $\\sim25^\\circ$ and eccentricities larger than $\\sim 0.2$ are unstable.} \\label{fig3} \\end{center} \\end{figure}   An interesting result depicted by Fig. 3 is the familiar role of secular resonances in establishing the boundaries of stable zones. Similar to the asteroids in the asteroid belt, stability of an MBC depends on the values of its initial eccentricity and orbital inclination. Fig.  3 shows that for initial inclinations larger than $\\sim 25^\\circ$, the orbit of an MBC becomes unstable due to the Kozai and the $\\nu_5$, $\\nu_6$ and $\\nu_{16}$ secular resonances. For smaller values of inclination,  the apastron distance of an MBC determines its stability. Those hypothetical MBCs close to or inside the region of the 2:1 MMR with Jupiter became unstable in a short time. An analysis of the orbits of the unstable objects indicates that approximately $80\\%$ of these bodies were scattered to large distances outside the solar system. This is a familiar result that has also been reported by \\cite[O'Brien et al. (2007)]{OBrien07} and \\cite[Haghighipour \\& Scott (2008)]{Hagh08} in their simulations of the dynamical evolution of planetesimals in the outer asteroid belt. From the remaining $20\\%$ unstable MBCs, approximately $15\\%$ collided with Mars, Jupiter, or Saturn, and a small fraction $(\\sim 5\\%)$ reached the region of 1 AU implying that MBCs might have played a role in delivering water to the Earth.  The stability analysis above has direct implications for the origin of MBCs and favors the in-situ formation of these objects. In this scenario, MBCs are small asteroidal bodies that were formed as a result of the collisional break-up of their larger precursor asteroids. An alternative scenario based on the primordial capture of cometary bodies, although, as shown by \\cite[Levison et al. (2009)]{Levison09}, efficient in the inward scattering of D-type and P-type asteroids and the delivery of these objects in particular to the region of Trojans, cannot provide information on the inward scattering and distribution of C-type asteroids. That is primarily due to the fact that C-type asteroids are mainly at small semimajor axes, and the difference between their orbital distribution and that of D-type asteroids are not known. Additionally, the colorless feature of MBCs, as indicated by \\cite[Hsieh \\& Jewitt (2006)]{Hsieh06} and \\cite[Hsieh, Jewitt \\& Fern\\'andez (2008)]{Hsieh08} is not consistent with an origin model based on the inward scattering of comets from the Kuiper belt region (the latter objects are optically red).  The in-situ formation scenario is, however, consistent with MBCs orbital and spectral properties. In this scenario, the break up of the precursor asteroids could have produced many km-sized fragments, among which those with large inclinations and large eccentricities became unstable and were scattered to other regions. The remaining objects have naturally asteroidal  orbits (i.e. their Tisserand numbers are larger than 3), and similar to their parent bodies, are C-type asteroids with no specific optical color. In regard to  7968 Elst-Pizzaro, 118401 (176P/LINEAR), and P/2005 U1 (Read), this scenario points to the Themis family, and perhaps a smaller $\\sim 10$ Myr sub-family (known as Beagle) within these objects (\\cite[Nesvorn\\'y et al 2008]{Nesvorny08}), as the origin of these MBCs. This scenario also suggest that asteroid families, in particular those in the outer half of the asteroid belt and with large parent bodies capable of differentiating and forming ice-rich mantles, are the most probable places for detecting more MBCs. As indicated by the results of the dynamical simulations, some members of such families may interact with giant planets and reach orbits in other regions of the asteroid belt--a scenario that might explain the existence of P/2008 R1 (Garradd) in its current orbit. All sky surveys such as those with Pan STARRS 1 would be capable of detecting such individual MBCs, and are ideal for carrying out targeted surveys for families of these objects. ",
        "conclusions": "\\begin{itemize}  \\item Current MBCs seem to have formed through the collision and break up of bigger asteroids. The results of the simulations of the dynamic of these objects point to the Themis family as the origin of 7968 Elst-Pizzaro, 118401 (176P/LINEAR), and P/2005 U1 (Read).  \\item Interaction with giant planet might have scattered MBCs from their original orbits to other locations in the asteroid belt. P/2008 R1 (Garradd) seems to be one of such scattered MBCs.  \\item More MBCs may exist in low inclinations and low eccentricities in the vicinity of asteroid families in the outer region of the asteroid belt.  \\item Collisions with  small objects might have activated MBCs or eroded them.  \\item Many MBCs might have been active in the past and are either no longer active, or will become active if hit by a small body again.  \\item Many MBCs, with locally exposed sub-surface ice, may still be on their ways to their perihelion distances where they become active, or they may be awaiting collisions with smaller objects to get activated.  \\item All sky surveys such as Pan STARRS will be able to detect more MBC in near future.  \\end{itemize}  \\acknowledgement I gratefully acknowledge fruitful discussions with H. Hsieh, D. Jewitt, H. Levison, K. Meech, D. Nesvorny, and N. Schorghofer. This work was partially supported by the NASA Astrobiology Institute under Cooperative Agreement NNA04CC08A at the Institute for Astronomy, NASA Astrobiology Central, the office of the Chancellor of the University of Hawaii, and a Theodore Dunham J. grant administered by Funds for Astrophysics Research, Inc. I am also grateful to Newton's Institute for Mathematical Science at the Cambridge University for their great hospitality during the preparation of this manuscript."
    },
    "0911/0911.1771_arXiv.txt": {
        "abstract": "We investigate a scenario in which neutrinos are coupled to a pseudoscalar degree of freedom $\\ph$ and where decays  $\\nu_1 \\to \\nu_2+\\ph$ and inverse decays are the responsible mechanism for obtaining equilibrium. In this context we discuss the implication of the invisible neutrino decay on the neutrino-pseudoscalar coupling constant and the neutrino lifetime. Assuming the realistic scenario of a thermal background of neutrinos and pseudoscalar we update the bound on the (off-diagonal) neutrino-pseudoscalar coupling constant to $g<2.6\\times10^{-13}$ and the bound on the neutrino lifetime to $\\tau<1\\times10^{13}\\,\\rm{s}$. Furthermore we confirm analytically that kinetic equilibrium is delayed by two Lorentz $\\ga$--factors -- one for time dilation of the (decaying) neutrino lifetime and one from the opening angle. We have also confirmed this behavior numerically. ",
        "introduction": "The possibility for neutrino interactions beyond the standard model has been studied in many contexts over the years. One particularly simple possibility is that neutrinos couple to a new pseudoscalar degree of freedom, as is, for example, the case in Majoron models - see the following references for previous discussions about the dynamics of the strong neutrino-pseudoscalar coupling and its astrophysical implications \\cite{Chikashige:1980ui,Gelmini:1980re,Schechter:1981cv,Gelmini:1982rr,Kolb:1985tqa,Manohar:1987ec,Dicus:1988jh,Konoplich:1988,Berezhiani:1989za,Gelmini:1994az}.  Astrophysics provides fairly stringent constraints on such couplings. For example SN1987A provides a bound on the dimensionless coupling constant of order $10^{-7}\\lwig g \\lwig 10^{-5}$ \\cite{Choi:1989hi,Raffelt:1996wa,Farzan:2002wx,Kachelriess:2000qc} by requiring that the neutrino signal should not be significantly shortened.  In the same way there are two cosmological bounds on $g$. First, the value of $g$ should not be large enough that pseudoscalars are fully thermalized before big bang nucleosynthesis. This leads to $g \\lwig 10^{-5}$. Second, a significant value of $g$ will make neutrinos self-interacting in the late universe and prevent neutrino free-streaming. This possibility has been discussed a number of times in the literature (see \\cite{Raffelt:1987ah,Beacom:2004yd,Hannestad:2004qu,Bell:2005dr,Sawyer:2006ju}).  The effect on cosmological observables such as the cosmic microwave background (CMB) spectrum were studied in \\cite{Beacom:2004yd,Hannestad:2004qu,Bell:2005dr,Sawyer:2006ju,Cirelli:2006kt,% Friedland:2007vv,Hannestad:2005ex,Basboll:2008fx}); particularly it was found that although models with no neutrino free-streaming can mimic the matter power spectrum of $\\Lambda$CDM models, they produce a distinct signature in the CMB spectrum which is much harder to reproduce. The feature arises because neutrinos act as a source term for photon perturbations. If there is no free-streaming, the source term is stronger and consequently the CMB anisotropy is increased for all scales inside the particle horizon at recombination. On the other hand, there is no effect on larger scales.  This distinct signature has been used to constrain models without neutrino free-streaming and in \\cite{Hannestad:2005ex,Basboll:2008fx} it was used to constrain the corresponding neutrino-pseudoscalar coupling parameters.  However, for decays and inverse decays the interaction was treated in a somewhat simplified manner in the sense that the momentum equilibration rate was assumed to be roughly $\\Gamma^* \\sim 1/(\\ga^2 \\ta)$, where $\\ta$ is the rest-frame lifetime and $\\gamma \\sim E_\\nu/m_\\nu$ is the Lorentz boost factor.  In this paper we wish to check this assumption in an explicit way with a realistic setup. The possible departure from the simple relation $\\Gamma^* \\sim 1/(\\ga^2\\ta)$ is something which is highly relevant for parameter estimations - such as placing bounds on the neutrino-pseudoscalar coupling and hence also for constraining the neutrino lifetime. Furthermore it is something which needs to be taken into account in numerical studies in which we allow for nonstandard neutrino interactions. One particular area where detailed knowledge of the interaction would be useful is the search for the cosmic neutrino background \\cite{Hannestad:2009xu,de Bernardis:2007bu,Ringwald:2009bg,DeBernardis:2008ys}.  One comment is in order here: The motivation for looking at decays and inverse decays rather than various scattering processes involving neutrinos and pseudoscalars ($\\nu\\nu\\to\\ph\\ph$,$\\,\\,\\ph\\ph\\to\\nu\\nu$,$\\,\\,\\nu\\ph\\to\\nu\\ph$) is that the probabilities of these scattering processes are proportional to $g^4$, where $g$ is the neutrino-pseudoscalar coupling constant. The probability of the decay $\\nu_1 \\to \\nu_2+\\ph$, on the other hand, is proportional to $g^2$. Consequently, at small values of $g$ the decay actually dominates over the scattering processes and allows us to put severe constraints on $g$.  This paper is organized as follows: In Sec.~\\ref{setup} we look at the setup with a gas consisting of two neutrino species and a pseudoscalar -- a gas that only has decays and inverse decays to obtain equilibrium. We argue for the $\\frac{1}{\\tau\\gamma^2}$ in the decay rate. In Sec.~\\ref{nobackg} we look at an initial situation of a standing wave of the heavy neutrinos and no other particles. In Sec.~\\ref{thermback} we introduce thermal distributions of the light neutrino and of pseudoscalars into a thermal background while keeping the initial conditions for the heavy neutrino. Furthermore we discuss the implication of the decay on the neutrino-pseudoscalar coupling and on the neutrino lifetime. We present numerical results in Sec.~\\ref{numerics}. Finally we have a conclusion and an appendix concerning calculations for the numerical implementation of the system. ",
        "conclusions": "We have investigated a neutrino-pseudoscalar gas with only decays  $\\nu_1 \\to \\nu_2+\\ph$ and inverse decays to obtain equilibrium. We started with an anisotropic distribution of $\\nu_1$ and confirmed that kinetic equilibrium is delayed by two Lorentz $\\ga$--factors -- one for time dilation of the heavy neutrino lifetime and one from the opening angle ie. from the transformation of the isotropic distribution of the decay products in the rest frame of the decaying particle back to the (cosmic) laboratory frame. We found explicit analytical expressions for the rate of creation of transverse momentum -- both in a case with no background of the decay products and in the case of thermal backgrounds and the ultrarelativistic limits hereof. We have confirmed this behavior in numerical simulations as well -- though we had to make a thermal smear of the initial anisotropy, making the analytical and numerical results open to a qualitative, but not quantitative, comparison.  Furthermore we have obtained updated bounds on the neutrino-pseudoscalar coupling constant as well as on the neutrino lifetime in the realistic case of a thermal background of neutrinos and pseudoscalars."
    },
    "0911/0911.4376_arXiv.txt": {
        "abstract": "We report the detection, from observations using the James Clerk Maxwell Telescope, of CO J\\,$=$\\,3$\\rightarrow$\\,2 transition lines in six  carbon stars, selected as members of the Galactic Halo and having similar infrared colors.  Just one Halo star had been detected in CO before this work.  Infrared observations show that these stars are red (J-K $>$3), due to the presence of large dusty circumstellar envelopes.  Radiative transfer models indicates that these stars are losing mass with rather large dust mass-loss rates in the range 1--3.3 $\\times$$10^{-8}$M$_{\\odot}$yr$^{-1}$, similar to what can be observed in the Galactic disc.  We show that two of these stars are effectively in the Halo, one is likely linked to the stream of the Sagittarius Dwarf Spheroidal galaxy (Sgr dSph), and the other three stars certainly belong to the thick disc.  The wind expansion velocities of the observed stars are low compared to carbon stars in the thin disc and are lower for the stars in the Halo and the Sgr dSph stream than in the thick disc. We discuss the possibility that the low expansion velocities result from the low metallicity of the Halo carbon stars. This implies that metal-poor carbon stars lose mass at a rate similar to metal-rich carbon stars, but with lower expansion velocities, as predicted by recent theoretical models.  This result implies that the current estimates of mass-loss rates from carbon stars in Local Group galaxies will have to be reconsidered. ",
        "introduction": "Stars with initial masses in the range  0.8--8~M$_{\\odot}$ end their life with a phase of catastrophic mass-loss. During the asymptotic giant branch (AGB) phase, they develop a superwind leading to mass-loss rates up to 10$^{-4}$M$_{\\odot}$\\,yr$^{-1}$. This superwind enriches the ISM with newly synthesized elements.  The mass-loss mechanism of AGB stars is likely due to pulsations from the star and radiation pressure on dust grains.  Shocks due to pulsation extend the atmosphere, so that the material ejected by the star becomes dense and cold enough for dust to form.  Due to its opacity, dust absorbs the radiation from the star and is driven away by radiation pressure, carrying the gas along through friction. Theoretical models (Winters et al.\\ 2000), show that the mass loss evolves from a pulsation driven regime characterized by a low mass-loss rate and a slow expansion velocity to a dust-driven regime with a high mass-loss rate and a high expansion velocity ($>$5km.s$^{-1}$).  Studying the effect of metallicity on the mass-loss process is important to understand the formation of dust around AGB stars in the early Universe.  At low metallicity, less seeds are present for dust formation, so one might expect dust formation to be less efficient and thus the mass-loss rates to be lower.  Theoretical work by Bowen \\& Willson (1991) predicts that for metallicities below [Fe/H]$=-1$ dust-driven winds fail, and the wind must stay pulsation-driven.  However, observational evidence for any metallicity dependence is still very limited (Zijlstra 2004).  More recent observational (Groenewegen et al.\\ 2007) and theoretical works (Wachter et al.\\ 2008; Mattsson et al. 2008) indicate that the mass-loss rates from carbon stars in metal-poor environments are similar to our Galaxy.  To obtain the first observational evidence on mass-loss rates at low metallicity, we have carried out several surveys with the {\\it Spitzer Space Telescope} of stars in nearby dwarf galaxies.  These show significant mass-loss rates down to Z=1/25 Z${_\\odot}$ (Lagadec et al.\\ 2007b; Matsuura et al.\\ 2007; Sloan et al.\\ 2009), but only for carbon-rich stars. The current evidence indicates that oxygen-rich stars have lower mass-loss rates at lower metallicities.  For carbon stars, no evidence for a dependency of mass-loss rate on metallicity has yet been uncovered.  Consequently, (Lagadec \\& Zijlstra 2008) have proposed that the carbon-rich dust plays an important role in triggering the AGB superwind.  The main uncertainty arises from the unknown expansion velocity.  This parameter is needed to convert the density distribution to a mass-loss rate.  The expansion velocity is also in itself a powerful tool.  Hydrodynamical simulations (Winters et al.\\ 2000), have shown that radiation-driven winds have expansion velocity in excess of 5\\,km\\,s$^{-1}$, while pulsation-driven winds are slower.  There is some evidence that expansion velocities are lower at low metallicity, from measurement of OH masers (Marshall et al.\\ 2004).  However, OH masers are found only in oxygen-rich stars, which appear to have suppressed mass-loss at low metallicities.  We lack equivalent measurements for metal-poor carbon-rich stars, which do reach substantial mass-loss rates.  For these, the only available velocity tracer is CO.  Currently, extra-galactic stars are too distant for CO measurements.  However, there are a number of carbon stars in the Galactic Halo, which are believed to have similarly low metallicity.  One CO measurement exists for a Galactic Halo star:  its expansion velocity has been estimated to be $\\sim$3.2 km\\,s$^{-1}$ through the $^{12}$CO $J=2 \\rightarrow 1$ transition (Groenewegen et al.\\ 1997).  This is much lower than typical expansion velocities, which are in the range 10-40 km\\,s$^{-1}$ for stars with similar colours, and thus optical depths (Fig.\\ref{histo_vexp}).  A large number of metal-poor carbon stars have recently been discovered in the Galactic Halo (Totten \\& Irwin 1998; Mauron et al.\\ 2004, 2005, 2007).  These stars may have been stripped from the Sagittarius Dwarf Spheroidal galaxy (Sgr dSph), which has a metallicity of [Fe/H]$\\sim-1.1$ (Van de Bergh 2000).  These stars are the closest metal-poor carbon stars known, and they are bright enough to be detected in CO using ground-based millimeter telescopes.  We have carried out observations of six Halo carbon stars in the CO J $= 3 \\rightarrow 2$ transition.  Here, we report the results of these observations.  \\begin{figure} \\begin{center} \\includegraphics[width=9cm]{halo_plot.ps} \\caption{\\label{histo_vexp} Distribution of the expansion velocity for the observed Halo carbon stars compared with stars from the disc with similar J$-$K colors.} \\end{center} \\end{figure} ",
        "conclusions": "We have detected the CO J\\,$=$\\,3$\\rightarrow$\\,2 in six carbon stars selected as Halo stars.  Only one carbon star had been detected in CO previously.  Comparison of the infrared observations and radiative transfer models indicates that these stars are losing mass and producing dust.  Their mass-loss rates are larger than their nuclear burning rates, so their final evolution will be driven by this mass-loss phenomenon.  We show that three of the observed stars are certainly members of the thick disc, while one is in the Sgr dSph stream and two are in the Halo.  The CO observation of the Sgr dSph stream star are thus the first identified millimetre observations of an extragalactic AGB star.  The expansion velocity we determined from our CO observations are lower than those of carbon stars in the thin disc with similar near-infrared colours.  The observed carbon stars with the lowest expansion velocities are Halo or Sgr dSph stream carbon stars.  There is a strong indication that the expansion winds are lower in metal-poor environments, which agrees with recent theoretical models (Wachter et al.\\ 2008).  So far, the effect of metallicity on the mass-loss from carbon-rich AGB stars has studied primarily from infrared observations of AGB stars in Local Group galaxies, mostly with the {\\it Spitzer Space Telescope}.  Infrared observations measure the infrared excess, which can be converted to a mass-loss rates assuming an expansion velocity for the circumstellar material.  So far all of the mass-loss rates have been estimated using the assumptions that the expansion velocity is independent of the metallicity.  The results presented here show that this assumption needs to be reconsidered.  We have recently obtained spectra of some of the present sample with the Infrared Spectrograph on {\\it Spitzer}. Combining our CO observations and these spectra together and comparing them to {\\it Spitzer} spectra from carbon stars in other galaxies in the Local Group will allow us to quantitatively study the mass-loss from carbon-rich AGB stars in the Local Group and its dependence on metallicity. This will be the subject of a forthcoming paper."
    },
    "0911/0911.4695_arXiv.txt": {
        "abstract": "We report the detection of very-high-energy (VHE) gamma-ray emission from supernova remnant (SNR) G106.3+2.7. Observations performed in 2008 with the VERITAS atmospheric Cherenkov gamma-ray telescope resolve extended emission overlapping the elongated radio SNR. The 7.3$\\sigma$ (pre-trials) detection has a full angular extent of roughly $0.6^\\circ$ by $0.4^\\circ$.  Most notably, the centroid of the VHE emission is centered near the peak of the coincident $^{12}$CO (J=1-0) emission, $0.4^\\circ$ away from the pulsar PSR J2229+6114, situated at the northern end of the SNR. Evidently the current-epoch particles from the pulsar wind nebula are not participating in the gamma-ray production. The VHE energy spectrum measured with VERITAS is well characterized by a power-law $dN/dE = N_0(E/3\\;TeV)^{-\\Gamma}$ with a differential index of $\\Gamma = 2.29 \\pm 0.33_{stat} \\pm 0.30_{sys}$ and a flux of $N_0 = (1.15 \\pm 0.27_{stat} \\pm 0.35_{sys}) \\times 10^{-13}$ cm$^{-2}$s$^{-1}$TeV$^{-1}$. The integral flux above 1 TeV corresponds to $\\sim 5$ percent of the steady Crab Nebula emission above the same energy.  We describe the observations and analysis of the object and briefly discuss the implications of the detection in a multiwavelength context. ",
        "introduction": "The supernova remnant (SNR) G106.3+2.7 was first identified as a faint extended radio source in a northern Galactic Plane radio survey with the Dominion Radio Astrophysical Observatory (DRAO) by \\citet{1990AAS...82..113J}.  Working with the earlier radio data of \\citet{1980AAS...42..227K}, those authors determined that the object was a previously undetected supernova remnant, with a radio spectral index of $0.45$ and a flux density at 1~GHz of 6~Jy.  Subsequent work by \\cite{2000AJ....120.3218P} confirmed the object as a supernova remnant, with an estimated age and distance of 1.3~Myr, and 12~kpc.  The object can be described as cometary in shape, with a compact radio-bright head to the north, and a dimmer, extended tail to the southwest.  Located at the northern edge of the remnant's head is the pulsar PSR~J2229+6114 and its associated boomerang-shaped radio and x-ray-emitting wind nebula, G106.6+2.9 \\citep{2001ApJ...547..323H}.  The pulsar, which has a period of 51.6~ms, was discovered in a search of the error box of the EGRET source 3EG~J2227+6122 by \\cite{2001ApJ...552L.125H} and is notable for being one of the most energetic pulsars known, with a spin-down power $\\dot{E} = 2.2 \\times 10^{37}$ erg~s$^{-1}$, despite having a low radio luminosity \\citep{2006ApJ...638..225K}.  Based on polarization studies and velocity maps of HI and CO emission, \\cite{2001ApJ...560..236K} associate the pulsar wind nebula and the supernova remnant with the same progenitor event and invoke a scenario of shock-wave evolution into a complex molecular environment to explain the unusual configuration of the constituent components.  These authors derive a distance to the remnant of 0.8~kpc, and adopt an age of 10~kyr, inferred from the pulsar data \\citep{2001ApJ...552L.125H}.  At GeV energies, the EGRET source 3EG~J2227+6122, with relatively large error contours, is compatible with the pulsar position, as well as with the main bulk of the radio remnant \\citep{1999ApJS..123...79H}.  The newly-released Large Area Telescope Bright Source List of the Fermi Gamma-ray Space Telescope (FGST) contains the object 0FGL J2229.0+6114, the position of which is consistent with the pulsar position \\citep{2009arXiv0902.1559A}.  The AGILE source 1AGL J2231+6109 \\citep{web:Agile} is also located nearby.  At TeV energies, the most stringent flux limits come from the MAGIC Collaboration \\citep{EsterThesis}, who quote a point-source upper limit of $\\sim10$\\% of the Crab Nebula flux above 220~GeV at the pulsar position.  Less constraining point-source limits have also been presented by the Whipple \\citep{2005ApJ...624..638F} and VERITAS \\citep{2007arXiv0709.3975K} Collaborations.  Additionally, the Milagro Collaboration has recently reported a new analysis of their data for this region, resulting in the detection of gamma rays over a broad $\\sim 1^\\circ$ area encompassing both the pulsar position and the main bulk of the remnant \\citep{2009arXiv0904.1018A}.  The detection corresponds to the location of their previously published ``source candidate C4\" \\citep{2007ApJ...664L..91A}, but lacks the angular resolution to provide a definitive association with a particular region within the SNR/pulsar complex. The statistical significance of the detection at the pulsar location is $6.6\\sigma$ and the median energy of the detected gamma rays is 35~TeV.  The quoted differential flux is $(70.9 \\pm 10.8_{stat} \\pm 35.5_{sys}) \\times 10^{-17}$ TeV$^{-1}$ cm$^{-2}$ s$^{-1}$, though this depends on assumptions about the source spectral shape. ",
        "conclusions": "The VERITAS observatory has made a high-confidence detection of TeV gamma rays from the supernova remnant G106.3+2.7.  This gamma-ray emission is significantly extended, and seems well-correlated with the radio extent of the remnant, in particular, with the location of known nearby molecular clouds.  The emission seen at the location of the associated pulsar PSR J2229+6114 is non-zero, but only of marginal statistical significance.  The total flux from the remnant above 1 TeV is about $\\sim 5\\%$ of the Crab Nebula, and can be described by a power-law in energy with differential index 2.29.  With the information available, it is not yet possible to make a firm determination of whether the TeV gamma-ray emission is of leptonic or hadronic origin, though a solid association with the Milagro source could favor the latter.  With additional exposure and a refined observation strategy (\\textit{i.e.}, wobbling around the location of the emission peak rather than the pulsar) a more in-depth study of the object will be achievable, including an investigation of possible energy dependence in the source morphology.  Such an effort will benefit from the enhanced sensitivity and angular resolution of the reconfigured VERITAS array, in which the T1 telescope (see Section \\ref{sec:anal}) is relocated to obtain a more optimal baseline.  This upgrade will be completed in autumn, 2009."
    },
    "0911/0911.4189_arXiv.txt": {
        "abstract": "Extended high-energy($\\ga100$ MeV) gamma-ray emission that lasts much longer than the prompt sub-MeV emission has been detected from quite a few gamma-ray bursts (GRBs) by Fermi Large Area Telescope (LAT) recently. A plausible scenario is that this emission is the afterglow synchrotron emission produced by electrons accelerated in the forward shocks. In this scenario, the electrons that produce synchrotron high-energy emission also undergo inverse-Compton (IC) loss and the IC scattering with the synchrotron photons should be in the Klein-Nishina regime. Here we study effects of the Klein-Nishina scattering on the high-energy synchrotron afterglow emission. We find that, at early times the Klein-Nishina suppression effect on those electrons that produce the high-energy emission is usually strong and therefore their inverse-Compton loss is small with a Compton parameter $Y\\la {\\rm a ~few}$  for a wide range of parameter space. This leads to a relatively bright  synchrotron afterglow at high energies that can be detected by Fermi LAT. As the Klein-Nishina suppression effect weakens with time, the inverse-Compton loss increases and could dominate over the synchrotron loss in some parameter space. This will lead to a faster temporal decay of the high-energy synchrotron emission than what is predicted by the standard synchrotron model, which may explain the observed rapid decay of the early high-energy gamma-ray emission in GRB090510 and GRB090902B. ",
        "introduction": "With the launch of Fermi satellite, one of the new features of gamma-ray bursts at high-energies  has been established, i.e. GRBs show extended high-energy ($\\ga100$ MeV) emission which lasts much longer than the prompt phase. This extended emission has been seen in both long and short GRBs and the flux usually decays with time after the initial peak. In some cases (e.g. GRB090510 and GRB090902B), the temporal decay is a simple power-law decay with a slope ranging from $-1.3$ to $-1.5$ (Abdo et al. 2009a,b; Ghirlanda, Ghisellini, Nava 2009; De Pasquale et al. 2009). One of the proposed models for such emission is the hadronic cascade emission model, in which the high-energy photons produced by the accelerated ultra-high energy protons can not escape the soft photon field and a cascade is induced (Abdo et al. 2009b). This model has been applied to the extended emission of GRB941017 (Dermer \\& Atoyan 2004), but whether it can explain the simple power-law decay of the Fermi LAT bursts is unknown. The long-lived behavior and not very rapid decay of the high-energy emission from GRB090510 and GRB090902B can not be easily explained by  the reverse shock emission model either (Wang et al. 2001a,b){\\footnote{The decay of the high-energy emission in GRB080825C is  steeper than $t^{-1.7}$ (Abdo et al. 2009c), which is consistent with the synchrotron self-inverse Compton scattering emission from the reverse shock or cross IC scatterings between the reverse shock and forward shock (Wang et al. 2001a,b; Granot \\& Guetta 2003; Pe'er \\& Waxman 2004; Wang et al. 2005).}}. On the other hand, the simple power-law decay with a modestly large slope is reminiscent of the afterglow emission. The self inverse-Compton (IC) emission of the afterglow has been long thought to produce a high-energy component (e.g. Zhang \\& \\Meszaros 2001; Sari \\& Esin 2001; Fan et al. 2008; Gou \\& \\Meszaros 2007 ), but the light curve is expected to rise initially and start to decay minutes to hours after the burst. Kumar \\& Barniol Duran (2009a) proposed that the extended high-energy emission from GRB080916C is due to afterglow synchrotron emission. This mechanism has also been proposed to explain the extended high-energy emission from GRB090510 (Gao et al. 2009; Ghirlanda, Ghisellini, Nava 2009; Ghisellini, Ghirlanda, Nava 2009, De Pasquale et al. 2009) and GRB090902B (Kumar \\& Barniol Duran 2009b).  In the latter synchrotron afterglow scenario, the high-energy extended emission is produced by the electrons  in the forward shock via synchrotron emission. The shock-accelerated electrons are usually assumed to have a power-law form in energy distribution, i.e. $dN_e/d\\gamma_e\\propto \\gamma_e^{-p}$, where $\\gamma_e$ is the Lorentz factor of electrons. These electrons also suffer IC loss by scattering synchrotron photons. Due to that the scattering between large $\\gamma_e$ electrons and the synchrotron photons could enter the Klein-Nishina (KN) scattering regime, higher energy electrons may suffer smaller IC loss and as a consequence, their synchrotron emission is stronger. Since the Lorentz factor of the electrons producing the high-energy afterglow emission are usually large, the KN scattering effect must be taken into account when one calculate the synchrotron high-energy afterglow emission. In this paper, we study the effect of the KN scattering on the high-energy afterglow emission of GRBs. Recently, Wang et al.(2009) studied the KN effect on the prompt emission spectrum of GRBs and Nakar et al. (2009) studied the KN effect on optically thin synchrotron and synchrotron self-Compton spectrum in general. In this paper, we focus on the early high-energy afterglow emission and confront the theoretical results  with the high-energy afterglow observations by Fermi LAT.  The paper is organized as follows. First we study how the KN scattering affects the electron distribution in the forward shock in $\\S$ 2. Then in $\\S$ 3 we calculate the Compton parameters for the high-energy electrons that produce high-energy gamma-ray emission and study their evolution with time. With the Compton parameters known, we calculate the light curves of high-energy afterglows in $\\S$ 3. Finally, we summarize  our findings in In $\\S$ 4. ",
        "conclusions": "We have studied the KN effect on early high-energy  afterglow emission of GRBs. Our findings are summarized as:  i) The IC scatterings between high-energy electrons that produce early-time high-energy ($\\ga100$ MeV) afterglow emission and synchrotron peak photons of the afterglow  are generally in the deep KN scattering regime. As a result, the IC loss of these electrons is small and  the synchrotron luminosity at $h\\nu\\ga100$ MeV is high, which is favorable for the detection of high-energy gamma-ray emission from the early afterglow by Fermi LAT.  ii)The high-energy gamma-ray emission at the early afterglow phase is dominated by the synchrotron emission. The SSC emission becomes dominated only at energies above the maximum synchrotron photon energy of shock-accelerated electrons, but at such energies the SSC flux is usually too weak to be detectable by Fermi LAT.  iii)The KN suppression effect of high-energy electrons weakens with time, so that the IC loss increases with time. In the parameter space where the Compton parameter is $Y(\\gamma_*)>1$, the increasing IC loss leads to a faster temporal decay of the synchrotron afterglow emission at high frequency. The decay slope could be steeper by a factor of $\\Delta \\alpha=0.5$ at most under favorable conditions. This may explain the rapid decay of the early-time high-energy emission observed in GRB090510 and GRB090902B.  {"
    },
    "0911/0911.0413_arXiv.txt": {
        "abstract": "We present a detailed analysis of the relation between infrared luminosity and molecular line luminosity, for a variety of molecular transitions, using a sample of 34 nearby galaxies spanning a broad range of infrared luminosities ($10^{10} < L_{IR} < 10^{12.5} L_{\\sun}$). We show that the power-law index of the relation is sensitive to the critical density of the molecular gas tracer used, and that the dominant driver in observed molecular line ratios in galaxies is the gas density.  As most nearby ultraluminous infrared galaxies (ULIRGs) exhibit strong signatures of active galactic nuclei (AGN) in their center, we revisit previous claims questioning the reliability of HCN as a probe of the dense gas responsible for star formation in the presence of AGN. We find that the enhanced HCN(1-0)/CO(1-0) luminosity ratio observed in ULIRGs can be successfully reproduced using numerical models with fixed chemical abundances and without AGN-induced chemistry effects.  We extend this analysis to a total of ten molecular line ratios by combining the following transitions: \\coOne, \\hcoOne, \\hcoThr, \\hcnOne, and \\hcnThr. Our results suggest that AGNs reside in systems with higher dense gas fraction, and that chemistry or other effects associated with their hard radiation field may not dominate (NGC~1068 is one exception). Galaxy merger could be the underlying cause of increased dense gas fraction and the evolutionary stage of such mergers may be another determinant of the HCN/CO luminosity ratio. ",
        "introduction": "\\label{sec:intro}  Over the past 50 years, there have been a number of studies relating galaxy star formation rate (SFR) and the amount of molecular gas available.  \\citet{1959ApJ...129..243S} proposed a power-law relationship between star formation rate volume density and molecular gas volume density.  \\citet{1998ApJ...498..541K} \\citep[also see][]{1989ApJ...344..685K} framed the problem in terms of surface densities, yielding the Kennicutt-Schmidt relation: $\\Sigma_{SFR} \\propto \\Sigma_{gas}^{1.4 \\pm 0.15}$.  While a power-law index around 1.4 or 1.5 holds for the molecular gas traced by \\coOne, a power-law index of unity was found when relating SFR to the molecular gas traced by \\hcnOne\\ \\citep[hereafter GS04b]{1992ApJ...387L..55S,  2004ApJ...606..271G}.  This difference in index is interpreted as a consequence of the different critical densities\\footnote{$n_{crit} \\equiv A_{ul}/\\Gamma_{ul}$, where $A_{ul}$ is the Eintein A coefficient and $\\Gamma_{ul}$ is the collision rate for a transition from upper (u) to lower (l) energy level.} of the molecular tracers used.  Having a lower critical density, \\coOne\\ traces the global molecular gas content, some of which may not be involved in the star formation process, whereas \\hcnOne\\ traces denser molecular gas ($n > 10^{4}~cm^{-3}$) more closely linked to star formation. Since the \\citet{1992ApJ...387L..55S} and \\citet[hereafter GS04a]{2004ApJS..152...63G} observations, several interpretations regarding the linear $L_{\\rm IR}$-\\hcnOne \\ relation in galaxies have been put forth by both observational arguments and numerical models.  For example, by extending the linear relationship between the total infrared luminosity, $L_{\\rm IR} [\\equiv L$(8-1100$\\mu$m)], and \\hcnOne \\ to Galactic cloud cores, \\citet{2005ApJ...635L.173W} added to the original GS04b interpretation by framing the relationship in terms of individual star-forming dense-gas ``units''. As one increases the number of dense star-forming clumps in a galaxy, $L_{\\rm IR}$ and \\hcnOne \\ both increase in lock-step.  In this view, the only difference between an extreme starburst galaxy and Galactic star-forming region is the number of dense star-forming units emitting HCN and infrared emission.  Recent observations may challenge the existence of the tight, linear correlation between $L_{\\rm IR}$ and \\hcnOne \\ luminosity\\footnote{We use the notation L' for molecular line luminosities as they are expressed in K km s$^{-1}$ pc$^2$, whereas other luminosities are expressed in L$_{\\sun}$.} reported by GS04b. For instance, \\citet[hereafter GC08]{2008A&A...479..703G} obtained new HCN and HCO$^{+}$ measurements for 17 luminous and ultra-luminous infrared galaxies (LIRGs: $10^{11} < L_{IR} < 10^{12}~L_{\\sun}$ and ULIRGs: $L_{IR} > 10^{12}~L_{\\sun}$).  They found conflicting \\hcnOne\\ measurements for several galaxies that overlap with the \\citet{1992ApJ...387L..55S} sample. They attribute the differences to calibration errors in the older survey, and proceed to compare their revised $L_{IR}-$\\Lphcn\\ relation with the original one.  They find a larger and more significant increase of $L_{IR}-$\\Lphcn\\ with $L_{IR}$, meaning that the $L_{IR}-$\\Lphcn\\ power-law would be steeper than an index of unity.  These authors interpreted the higher $L_{\\rm IR}$/\\Lphcn\\ ratio as an enhanced star formation efficiency in ULIRGs.  A large fraction of nearby LIRGs and ULIRGs appear to be mergers of gas-rich progenitor galaxies and therefore have a high concentration of dense molecular gas in their center which feeds strong starburst and AGN activity \\citep[see][and references therein]{1996ARA&A..34..749S}. It has been noted that sources with brighter IR luminosities are both more likely to host an AGN \\citep[e.g.][]{1995ApJS...98..171V, 1998ApJ...505L.103L, 1999ApJ...522..113V}, as well as show higher \\hcnOne/\\coOne\\ molecular line luminosity ratios \\citep[GS04b,][]{2006ApJ...640L.135G}. It is tempting to associate the increased \\hcnOne/\\coOne \\ luminosity ratio with the presence of AGN. Indeed, \\citet{2008A&A...479..703G} report evidence for an enhanced abundance of HCN in ULIRGs, and caution that care must be taken when converting HCN line luminosities into dense gas masses.  Moreover, some theoretical work suggests that AGN-induced processes could potentially increase the abundance of HCN \\citep[e.g.][]{2006ApJ...646L..37L}, causing the observed $L_{\\rm IR}$-\\Lphcn \\ relation to have a shallower slope than the $L_{\\rm IR}$-\\Lpco \\ relation. In this picture, X-ray emission resulting from AGN buried in ULIRGs may enhance the abundance of HCN molecules via an increased availability of free electrons that facilitates combination with ions (e.g., HCNH$^+$ + e$^-$ $\\rightarrow$ HCN + H). Some high-resolution HCN observations of nearby galaxies have shown stronger emission of \\hcnOne\\ relative to \\coOne\\ in the center of a few LIRGs and ULIRGs known to host an AGN \\citep[e.g.][]{2006ApJ...640L.135G,2007A&A...468L..63K}. These authors have interpreted this higher molecular line ratio in terms of an increased abundance of HCN.  Others interpret high molecular line ratios in terms of the mid-plane pressure, in the sense that larger ambient pressure could trap or maintain dense gas \\citep{2008ApJ...673..183L, 2006ApJ...650..933B}.  In recent years, new theoretical models suggest a different physical explanation for the observed $L_{\\rm IR}$-molecular line luminosity relationships.  Specifically, models that couple non-local thermodynamic equilibrium (LTE) radiative transfer calculations with models of star-forming giant molecular clouds (GMCs) \\citep{2007ApJ...669.289K} and hydrodynamic simulations of disk galaxies and galaxy mergers \\citep{2008ApJ...684..996N} find that the driving force controlling the SFR-\\lmol \\ relation in galaxies is the fraction of thermalized molecular gas at a given molecular transition. In the \\citet{2007ApJ...669.289K} and \\citet{2008ApJ...684..996N} picture, the SFR is controlled by the relation SFR$\\propto n^{1.5}$, where $n$ is the number density of molecular hydrogen.  When the observed molecular line traces the bulk of the molecular gas in the galaxy (e.g. \\coOne), the SFR-\\lmol \\ power-law index (hereafter ``slope'' as this relationship is typically considered in log-log space) is equivalent to the Kennicutt-Schmidt index (i.e. close to 1.5). On the other hand, higher critical density molecular line tracers (e.g. \\hcnOne) trace an increasingly smaller fraction of the gas in a galaxy, and so the index in the observed SFR-\\lmol\\ power-law relation decreases. These theories were tested by \\citet{2008ApJ...681L..73B}, who observed \\hcnThr \\ --- a higher critical density tracer than \\hcnOne --- in the GS04a,b sample of galaxies, and found a SFR-\\hcnThr \\ index less than unity ($0.72 \\pm 0.08$).  In this paper, we expand on the Bussmann et al. work by studying the relationship between five molecular line tracers and infrared luminosity. In particular, we emphasize the importance of considering the critical density of each molecular line tracer. We present support that \\hcnOne\\ is a valid probe of dense molecular gas and we revisit the interpretation of the enhanced \\hcnOne/\\coOne\\ luminosity ratio in high IR luminosity systems. Our results suggest that this feature is primarily driven by the increased dense gas fraction in galaxies undergoing a merger and/or a strong episode of starburst activity; we do not require X-ray driven chemical abundance effects to explain the observed molecular line ratios. Overall, we build a coherent picture where the governing parameter is the density distribution of the molecular gas.  We describe our sample of 34 galaxies in \\S\\ref{sec:method}, along with the measurements that we use.  Section~\\ref{sec:results} contains our results regarding power-law relationships between infrared luminosity and various molecular line luminosities (\\S\\ref{sec:lum}), as well as the study of molecular line ratios as a function of infrared luminosity (\\S\\ref{sec:lineratios}). The interpretation is partially based on the theoretical models of \\citet{2008ApJ...684..996N}, which we apply to this study in \\S\\ref{sec:model}. We include a brief discussion of chemistry effects in \\S\\ref{sec:Chem} and present our conclusions in \\S\\ref{sec:conclusions}.  Throughout this paper we assume a flat $\\Lambda$CDM cosmological model with ($\\Omega_{M}, \\Omega_{\\Lambda}$, H$_0$) = (0.3, 0.7, 70~km s$^{-1}$ Mpc$^{-1}$). ",
        "conclusions": "\\label{sec:conclusions}  Using a sample of 34 nearby infrared luminous galaxies ($10^{10} < L_{IR} < 10^{12.5}~L_{\\sun}$), we characterized the infrared luminosity dependence of various molecular gas tracers.  \\begin{itemize}  \\item{The presence of AGN was assessed using the optical {\\it BPT} diagram. In agreement with previous publications, we find a more frequent occurrence of AGN in more IR-luminous galaxies. This may be related to the availability of larger amounts of dense gas during the mergers of gas-rich galaxies. }  \\item{The molecular transitions used, \\coOne, \\hcoOne, \\hcnOne, \\hcoThr\\ and \\hcnThr, span 4 orders of magnitude in critical density.  We find that the relationship between (F)IR luminosity and molecular line luminosity $L_{mol}$ is shallower for transitions with higher $n_{crit}$.  This trend is in agreement with theoretical models of N08 and KT07 and can be explained by the varying degree of thermalization of the gas giving rise to molecular line emission. }  \\item{Trends of molecular line ratios with $L_{IR}$ are consistent with an increased molecular gas density in more IR-bright galaxies. This result agrees with the picture presented by \\citet{2004ApJ...606..271G} and \\citet{2005ApJ...635L.173W}. When comparing high-density and low-density molecular gas tracers, we observe an increase in their luminosity ratio ($L_{HD}/L_{LD}$) with increasing infrared luminosity.  Interestingly, this trend vanishes when comparing two tracers with nearly equal critical density (\\hcnOne and \\hcoThr).  We infer that the main driver of the enhanced HD/LD luminosity ratio is the molecular gas density distribution, in agreement with the observations that ULIRGs host large reservoirs of dense molecular gas in their central regions. }  \\item{We compare our observed molecular line ratios with theoretical values obtained from a set of galaxy models.  We consider SPH simulations of two galaxy mergers as well as individual gas-rich galaxies.  With constant Galactic abundances, the models successfully produce an enhanced HD/LD luminosity ratios at brighter infrared luminosity.  This provides additional support for a higher molecular gas density in galaxies that are extremely gas-rich or undergoing gas-rich mergers. Indeed, this result suggests that AGN-induced chemistry (or other) effects may not be necessary to reproduce the observations, but rather that AGNs are more likely to reside in galaxies that are very gas-rich and/or experiencing a merger. The simulations also demonstrate important variations in molecular luminosity ratio with the evolutionary stage of the mergers. These variations may be connected with the dynamics of the molecular gas during mergers (inflows, outflows, gas compression, etc.). }  \\item{We investigate possible chemistry effects on the well-known \\hcnOne/\\coOne\\ luminosity ratio using common optical emission line diagnostics of the ISM ionization parameter and metallicity. With the exception of one of two outliers, we do not observe significant correlations between the molecular line ratio and these properties.  Decoupling between the dense molecular gas traced by \\coOne\\ and \\hcnOne\\ and the more diffuse gas that traces the optical nebular line may be responsible for the absence of correlation.  However, we note that NGC~1068 shows extreme conditions in the sense of having the largest optically-measured ionization parameter as well as the largest MIR-excess.  This outlier is thus more subject to chemistry and/or MIR-pumping effects quoted in the literature.  We emphasize that NGC~1068 has properties that are very distinct from the rest of our sample, and thus may not be representative of its class. }  \\item{More high-density molecular line observations would be extremely beneficial to confirm the trends outlined in this work. Only half of our combined sample of 34 galaxies have data for all five molecular lines used here. In particular, future observations should target galaxies with fainter IR luminosities to test whether the results presented here extend into the lower luminosity regime. Higher-resolution studies of high-density tracers may complement this work by allowing for a more detailed analysis of the density and excitation structure of the molecular gas within galaxies. Such studies could provide information on the \\emph{local} influence of AGN, whereas in this work we probe the \\emph{global} influence of AGN. }  \\item{Overall, our results support that \\hcnOne\\ is a valid tracer of dense molecular gas in galaxies even in the presence of AGN. We expect scatter in the relationship due to variations in temperature and possible radiative (de-)excitation of this transition.  Also, the dense-to-total molecular gas fraction is expected to differ from galaxy to galaxy, especially in systems undergoing significant mergers. }  \\end{itemize}"
    },
    "0911/0911.0887_arXiv.txt": {
        "abstract": "The pulse shapes detected during multiple outbursts of \\saxj are analyzed in order to constrain the neutron star's mass and radius. We use a hot-spot model with a small scattered-light component to jointly fit data from two different epochs, under the restriction that the star's mass and radius and the binary's inclination do not change from epoch to epoch. All other parameters describing the spot location, emissivity, and relative fractions of blackbody to Comptonized radiation are allowed to vary with time. The joint fit of data from the 1998 ``slow decay'' and the 2002 ``end of outburst maximum'' epochs using the constraint $i<90^\\circ$ leads to the $3\\sigma$ confidence constraint on the neutron star mass $ 0.8 M_\\odot < M < 1.7 M_\\odot$ and equatorial radius $5 \\km < R < 13 \\km$.  Inclinations as low as $41^\\circ$ are allowed. The best-fit models with $M > 1.0 M_\\odot$ from joint fits of the 1998 data with data from other epochs of the 2002 and 2005 outbursts also fall within the same $3\\sigma$ confidence region. This $3\\sigma$ confidence region allows a wide variety of hadronic equations of state, in contrast with an earlier analysis \\citep{LMC08} of only the 1998 outburst data that only allowed for extremely small stars. ",
        "introduction": "\\label{s:intro}  The discovery \\citep{WvdK98} of the first accretion powered millisecond pulsar, SAX J1808.4 3658 (hereafter SAX J1808)), confirms the proposal that neutron stars in low-mass X-ray binaries are the progenitors of millisecond-period pulsars. SAX J1808 is now one of eleven known accreting millisecond X-ray pulsars (see \\cite{W06} and \\cite{P06} for reviews). The X-ray pulsations are most likely produced from accretion onto the neutron star's magnetic poles (see, e.g., Fig. 12 of \\citet{GDB02}). Spectral models \\citep{GDB02} provide strong evidence that the X-rays correspond to blackbody emission from a spot on the star which is then Compton scattered by electrons above the hot spot. Since the pulsed light is emitted from (or very close to) the neutron star's surface, accreting millisecond X-ray pulsars are excellent targets for light-curve fitting in order to constrain the neutron star equation of state (EOS). The X-ray light curve depends on the intrinsic properties of the emission spot (size, shape, location, and emissivity), as well as the neutron star's properties (mass, radius, and spin). Constraints on the star's mass and radius lead to constraints on the EOS of supernuclear density material.  The first pulse shape analysis (\\citet{PG03}, hereafter PG03) for SAX J1808 provided interesting constraints on the neutron star's mass and radius. However, this analysis did not take into account two effects that are important for rapidly rotating neutron stars: variable time delays due to light travel time across the star \\citep{CLM05} and the oblate shape of the star \\citep{CMLC07}. The pulse shape analysis including these effects was done by \\cite{LMC08}, resulting in smaller mass and radius than found in PG03. Since the pulse shapes of SAX J1808 were variable during the time period analyzed by PG03 and during subsequent outbursts by SAX J1808 (see the comprehensive analysis by \\citet{Har08}), it is important include the time variability in the analysis. In effect, the results of an analysis of a time-averaged pulse profile and an analysis of individually different pulse profiles are expected to be different. Here we carry out an analysis of a set of different pulse profiles from different time periods from SAX J1808. In addition, fitting multiple pulse profiles with the same model should provide a much more stringent test of the validity of the model than fitting of a single pulse profile.  The outline of our paper is as follows. In section 2, we introduce the different pulse shapes observed from SAX J1808, which come from different phases of the 1998, 2002 and 2005 outbursts. In section 3, the hot spot model is described and we note that an additional scattered-light component is required in order to fit all of the pulse shapes. In section 4, the results are given for joint fits of the pulse shape data, including resulting mass, radius and inclination values. We conclude with a general discussion of the results, comparing with previous work. ",
        "conclusions": "In this paper we analyze multiple pulse shapes resulting from the different epochs of the 1998, 2002 and 2005 outbursts of SAX J1808. We jointly fit data from 1998 and several other epochs in order to find consistent models for the neutron star's mass, radius and inclination.  In each joint fit to the data from two different epochs, the mass, equatorial radius and the inclination angle are kept fixed between epochs, while the spot location, anisotropy and blackbody-to-compton ratios are allowed to change from epoch to epoch. We use a spectral model, motivated by \\citet{GDB02} and \\citet{PG03}, that includes blackbody radiation and Comptonized radiation in the low-energy (3-4 keV) band and only Comptonized radiation in the high-energy (9-20 keV) band. We find that an extra scattered-light component is required to fit all of the pulse shapes.   The inclusion of the scattered radiation model (Section \\ref{s:disk}), together with a hot spot that wanders a little in its co-latitude, allows a consistent fit for all different epochs with a consistent mass. Increased radius or increased spot co-latitude increases the second harmonic content in the pulse shape; scattered radiation reduces it; and increased light bending (smaller star at a given mass) can increase or decrease second harmonic content depending on viewing geometry. Without including all these physical effects, the fits to some epochs require a small star  whereas other epochs require a larger star. As we have shown, by including all effects, we can consistently fit different epochs with a consistent mass.      The joint fit of data from the 1998B4 and 2002B3 epochs using the constraint $i<90^\\circ$ leads to the $3\\sigma$ confidence constraint on the neutron star mass $0.81 M_\\odot < M < 1.72 M_\\odot$ and equatorial radius $5.0 \\km < R < 12.7 \\km$, as summarized in Figure~\\ref{fig:mr90}.  Inclinations as low as $41^\\circ$ are allowed.  If the inclination is further constrained to $i<70^\\circ$, as suggested by many separate observations, the $3\\sigma$ limits are $ 0.83 M_\\odot < M < 1.50 M_\\odot$ and $ 7 \\km < R < 13 \\km$, as summarized in Figure~\\ref{fig:mr70}. Inclinations as low as $35^\\circ$ are allowed in this case.    The joint fits between 1998B4 and the other epochs (2002B4, 2005B1, 2005B2 and 2005B4) yield independent best-fit models summarized in Tables~\\ref{tab:joint98-02B4}-\\ref{tab:joint98-05B4}. The best-fit models from these epochs with $M > 1.0 M_\\odot$ have mass and radius values that lie within the $3\\sigma$ confidence region for the 1998B4-2002B3 joint fits. This indicates that the various joint fits between the different epochs give consistent results.  The $3\\sigma$ confidence region for the 1998B4-2002B3 joint fits allows a much wider range of masses and radii for the neutron star than a similar fit that only made use of data from the 1998 outburst \\citep{LMC08}. In the previous analysis by \\citet{LMC08} a similar type of analysis using 2 narrow energy bands constrained the neutron star mass and radius to very small values with $R<7$ km. The inclusion of a scattered-light component and use of joint fits with data from other epochs (2002 and 2005 outbursts) results in stars with larger masses and radii than are given by fits to the 1998 data alone without the scattered-light component.  In their analysis of the 2002 data, \\citet{IP09} do not attempt to find a best-fit mass and radius for SAX J1808. For some of their models they make use of a fixed mass of 1.4 $M_\\odot$ and a radius of either 10 or 12 km. However, the results of their modelling (such as the magnetic moment, visibility of an antipodal spot) do not strongly depend on the assumed values of mass and radius. A mass of 1.4 $M_\\odot$ and a radius in the 10-12 km range is consistent with our final results.   Observations of SAX J1808 in quiescence by \\citet{Heinke07,Heinke09} have set limits on the luminosity. The low luminosity during quiescence implies that rapid neutrino cooling takes place in the inner core, such as the direct URCA process. \\citet{YP04} show that direct URCA can take place if the core density is larger than approximately $10^{15}$ g/cm$^3$. For a standard hadronic EOS such as APR \\citep{APR}, stars with masses larger than about $1.5 M_\\odot$ (and spinning at 400 Hz) have central densities larger than this critical density. Softer EOS will have lower mass stars for the same central density. This suggests that many of the models allowed by the $3\\sigma$ confidence region shown in Figure \\ref{fig:mr90} could have cores dense enough to allow cooling through the direct URCA process.     A pulse-shape analysis of the accreting ms-period X-ray pulsar XTE 1814-338 \\citep{LMCC09} using the similar methods as used in this paper implied that the mass and radius of XTE 1814-338 must be quite large. Although the $3\\sigma$ regions for XTE 1814-338 and SAX J1808 do not overlap, the results are not necessarily inconsistent. It is only necessary for an EOS mass-radius curve to pass through the regions for both stars. For example, an EOS mass-radius curve that is slightly to the right of APR in Figure~\\ref{fig:mr90} would still be allowed by the SAX J1808 data, and would be stiff enough to be allowed by the XTE 1814 data, and also allowed by the constraints on the slow X-ray pulsar Hercules X-1 \\citep{Leahy04}."
    },
    "0911/0911.5693_arXiv.txt": {
        "abstract": "{We analyze the dependence of the membership probabilities obtained from kinematical variables on the radius of the field of view around open clusters (the sampling radius, $R_s$). From simulated data, we show that the best discrimination between cluster members and non-members is obtained when the sampling radius is very close to the cluster radius. At higher $R_s$ values more field stars tend to be erroneously assigned as cluster members. From real data of two open clusters (NGC~2323 and NGC~2311) we obtain that the number of identified cluster members always increases with increasing $R_s$. However, there is a threshold $R_s$ value above which the identified cluster members are severely contaminated by field stars and the effectiveness of membership determination is relatively small. This optimal sampling radius is $\\simeq 14$ arcmin for NGC~2323 and $\\simeq 13$ arcmin for NGC~2311. We discuss the reasons for such behavior and the relationship between cluster radius and optimal sampling radius. We suggest that, independently of the method used to estimate membership probabilities, several tests using different sampling radius should be performed in order to evaluate the existence of possible biases.} ",
        "introduction": "The large astrometric catalogues derived from surveys covering very wide areas of the sky are allowing the systematic searching of new star systems \\citep[see, for example,][and references therein]{Lop98,Hoo99,Kaz02,Myu03,Cab08b,Zha09}. The searching process is based on the detection of well defined structures in some subsets of the phase space. The presence of both spatial density peaks and proper motion peaks indicates the existence of star clusters; peaks visible only in the proper motion distributions suggest the existence of moving groups; whereas more spread and less dense velocity-position correlated structures could be associated to stellar streams. Once these structures have been detected, the next step is to search for identify possible members of the star system. For the particular case of open clusters, the most often used procedure to select possible cluster members is the algorithm designed by \\citet{San71}. This algorithm is based on a former model proposed by \\citet{Vas58} for the proper motion distribution. The model assumes that cluster members and field stars are distributed according to circular and elliptical bivariate normal distributions, respectively. The Sanders' algorithm, or some variation or refinement of it, has been and still is widely used to estimate cluster memberships either as the only method or as part of a more complet treatment that includes, for example, spatial and/or photometric criteria. Some recent representative references are \\citet{Wu02,Jil03,Bal04,Dia06,Kra07,Wir09}.  With the advent of large catalogues and databases available via internet and future surveys such as the forthcoming {\\it Gaia} mission of ESA, the interest in developing and applying fully automated techniques is increasing among the astronomical community. However, special care must be taken to avoid obtaining biased results. In this work we will show that the results yielded when using the Sanders' algorithm significantly depend on the choice of the size of the field of view surrounding the cluster. So, once detected a possible open cluster, it is natural to ask what area of the sky should be sampled in order to get the most reliable membership determinations. It is equally important to ask about the robustness of used methodology, i.e. how does the solution change when the sampled area is varied? Here we explore these subjects by using both simulated and real data. In Section~\\ref{metodo} we briefly present the method used to determine memberships and describe the simulations that we performed to analyze the expected behavior. The results of applying the Sanders' algorithm on the simulated data are discussed in Section~\\ref{resultados}. After this, in Section~\\ref{cumulos} we use real astrometric data of two open clusters (NGC~2323 and NGC~2311) to evaluate the performance of the algorithm. We discuss strategies to estimate the optimal sampling radius, i.e. the maximum radius beyond which the identified cluster members are expected to be severely contaminated by field stars. The main results of the present work are summarized in Section~\\ref{conclusiones}. ",
        "conclusions": "\\label{conclusiones}  We have evaluated the performance of the commonly used Sanders' method \\citep{Vas58,San71,Cab85} in the determination of star cluster memberships. In general, the results depend on the radius of the field containing the sampled cluster (the sampling radius, $R_s$). The main reason for this dependence lies in the differences between the assumed gaussian and the true underlying proper motion distributions. The contamination of cluster members by field stars increases as the sampling radius increases. The rate at which this effect occurs depends on the intrinsic characteristics of the data set. There is a threshold value of $R_s$ above which the identified cluster members are highly contaminated by field stars and the effectiveness of membership determination is relatively small. Thus, care must be taken when applying the Sanders' method (just by itself or as part of a more extensive procedure) especially when we do not have reliable information about the real cluster radius and/or when the sampling radius is larger the cluster radius. If this type of effects is not taken into consideration in automated data analysis then significant biases may arise in the derived cluster parameters. The optimal sampling radius can be estimated by plotting the number of cluster members and/or the fraction of members as a function of the sampling radius. Moreover, this type of analysis can also be used as an objective procedure that can be applied systematically to determine cluster radii."
    },
    "0911/0911.0763_arXiv.txt": {
        "abstract": "We examine deep \\emph{XMM-Newton} Reflection Grating Spectrometer (RGS) observations of the X-ray luminous galaxy cluster A1835. For the first time in a galaxy cluster we place direct limits on turbulent broadening of the emission lines. This is possible because the coolest X-ray emitting gas in the cluster, which is responsible for the lines, occupies a small region within the core.  The most conservative determination of the 90 per cent upper limit on line-of-sight, non-thermal, velocity broadening is $274\\kmps$, measured from the emission lines originating within 30~kpc radius. The ratio of turbulent to thermal energy density in the core is therefore less than 13 per cent.  There are no emission lines in the spectrum showing evidence for gas below $\\sim 3.5$~keV. We examine the quantity of gas as a function of temperature and place a limit of $140 \\Msunpyr$ (90 per cent) for gas cooling radiatively below 3.85~keV. ",
        "introduction": "Little direct information is known about the dynamical state or microphysics of the intracluster medium (ICM). Simulations of galaxy clusters within a cosmological environment predict that as matter accretes along filaments, or as clusters merge, turbulent motions should be set up in the ICM. This energy should then cascade from large to small scales where it can be dissipated. Simulations have predicted that even in the most relaxed clusters there are substantial flows of ICM (e.g. \\citealt{Evrard90}). In these simulations the fraction of pressure support in gas motions is typically 5 to 15 per cent. Most simulations, however, have not included physical processes which may be relevant, such as viscosity or magnetic fields.  In addition, it is observed that the active galactic nuclei (AGN) in cluster cores inject jets and bubbles of relativistic plasma into their surroundings. The feedback of energy into the cluster is thought to combat the high radiative losses of energy which would otherwise lead to substantial cooling rates (for reviews see \\citealt{PetersonFabian06} and \\citealt{McNamaraNulsen07}).  It is predicted that feedback could induce gas motions in the core in the range $500-1000\\kmps$ \\citep{Bruggen05}. Some models also require turbulence in order to heat the ICM using the energy in inflated cavities (e.g. \\citealt{Scannapieco08}). Strong turbulent motions would provide significant non-thermal pressure support. This would bias the determination of cluster mass profiles measured under the assumption of hydrostatic equilibrium. The measurement of turbulence is thus important for mass determination, as well as AGN feedback.  Evidence for resonance scattering has also been used to place upper limits on turbulence in elliptical galaxies. In NGC~4636, \\cite{Xu02} examined the strength of the resonantly scattered 15{\\AA} Fe~\\textsc{xvii} line relative to the line at {17.1\\AA}, inferring scattering. As turbulence broadens emission lines and prevents scattering, they obtained an upper limit on the turbulent velocity dispersion of 10 per cent of the sound speed. For the same object, \\cite{Werner09} calculated a limit of 25 per cent of the sound speed, with an implied maximum of 5 per cent of energy in turbulent motions.  The Perseus cluster is an interesting test case. The tens of kpc long quasi-linear filaments \\citep{FabianPerFilament03} imply laminar flow, significant viscosity in the ICM and low levels of turbulence on the scale of the filaments.  However, the lack of evidence for resonant scattering in the strongest 6.7~keV line of He-like iron in Perseus argues for velocity spread of at least half the sound velocity \\citep{Churazov04}. \\cite{RebuscoDiff05} also require motions of $300\\kmps$ on scales of $20\\kpc$ to account for the overall metallicity profile in Perseus.  Comparing the gravitational mass profiles derived from optical and X-ray data in M87 and NGC~1399, \\cite{Churazov08} obtained combined nonthermal pressures of $\\sim 10$ per cent of the thermal gas pressure.  \\cite{Reynolds05} simulated the evolution of cavities in the ICM, requiring viscous processes in order for the cavities to retain the observed spherical cavity shapes and prevent them from being shredded by instabilities \\citep{Reynolds05,Dong09}.  A \\emph{lower} limit of 10 per cent of the total pressure in turbulence was found in the Coma cluster by \\cite{Schuecker04}. The Coma cluster is unrelaxed with two central galaxies. It is unlikely that it can be directly compared to relaxed clusters or ellipticals. In summary, the evidence for turbulence in relaxed clusters is uncertain.  A1835 (at $z=0.2523$) is the most luminous cluster in the BCS sample \\citep{Ebeling98}, with an inferred cooling rate of $\\sim 1000 \\Msunpyr$ in the absence of heating \\citep{Allen96}.  The optical-spectroscopic-derived rate of star formation in the central galaxy is 40--70\\Msunpyr \\citep{Crawford99}. The IR luminosity of the object is $\\sim 7 \\times 10^{11} \\Lsun$ \\citep{Egami06}, implying a star formation rate of $\\sim 125 \\Msunpyr$. In addition there is $\\sim10^{11} \\Msun$ of molecular gas in the cluster \\citep{Edge01}.  Despite the evidence for star formation, cool gas and the high X-ray luminosity, little X-ray emitting gas was seen in this cluster below 1--2~keV when it was observed using the \\emph{XMM-Newton} RGS instruments \\citep{Peterson01}. Less than $200\\Msunpyr$ cools below 2.7~keV (90 per cent confidence).  AGN feedback, at a level of $1.6\\times 10^{45} \\ergps$, prevents most of the X-ray gas from cooling.  Later observations confirmed the lack of X-ray cool gas is a common feature of clusters (e.g. \\citealt{Peterson03}).  In this paper we examine deep RGS observations of A1835 totalling 254~ks, placing limits on velocities in the gas and the amount of cool gas detectable in X-rays. We use a Hubble constant of $70 \\kmpspMpc$ and the relative Solar metallicities of \\cite{AndersGrevesse89}. 1 arcsec corresponds to 3.94~kpc at $z=0.2523$. Errors quoted are $1\\sigma$ unless stated otherwise. ",
        "conclusions": "For the first time we measure meaningful constraints of the ICM turbulent velocity from the width of the X-ray emission lines. We place a limit of 13 per cent on the turbulent energy density as a fraction of the thermal energy density. We also constrain the cooling rate of X-ray emitting gas in A1835 to a value comparable to the star formation rate."
    },
    "0911/0911.1659_arXiv.txt": {
        "abstract": "\\grb\\ is a long gamma-ray burst (GRB) discovered by \\swift, featuring a bright X-ray afterglow as well as a faint infrared transient with very red and peculiar colors. The burst attracted a large interest because of a possible quasi-periodicity at $P=8.1$ s in the prompt emission, suggesting that it could have a different origin with respect to standard, long GRBs. In order to understand the nature of this burst, we obtained a target of opportunity observation with \\xmm. X-ray spectroscopy, based on \\xmm\\ and \\swift\\ data, allowed us to model the significant excess in photoelectric absorption with respect to the Galactic value as due to a large column density ($\\sim$$6.5\\times10^{22}$ cm$^{-2}$) in the GRB host, located at $z\\sim4.2$. Such a picture is also consistent with the infrared transient's properties. Re-analysis of the prompt emission, based on \\emph{International Gamma-Ray Astrophysics Laboratory} and on \\swift\\ data, excludes any significant modulation at $P=8.1$ s. Thus, we conclude that \\grb\\ is a distant, standard, long GRB. ",
        "introduction": "The bright, long gamma-ray burst (GRB) \\grb\\ was discovered by the Burst Alert Telescope (BAT) onboard \\swift\\ on 2009 July 9 at $T_0=07$:38:34.59 UT \\citep{morris09}. The prompt emission had a complex structure (see Fig.~\\ref{lcbat}), with a multi-peaked light curve lasting $t_{90}\\sim89$ s. The peak flux was $7.8\\pm0.3$ ph cm$^{-2}$ s$^{-1}$ at $T_0+21$ s and the fluence was $(2.57\\pm0.03)\\times10^{-5}$ erg cm$^{-2}$ in the 15--350 keV energy range \\citep{sakamoto09}. A bright X-ray afterglow was observed by the X-ray Telescope (XRT) onboard \\swift\\, starting as soon as 77 s after the trigger \\citep{morris09}. Analysis of XRT data up to $T_0+0.3$ days yielded evidence for a break in the decay at $T_0+0.1$ days. The XRT spectrum was described by a power law with an absorbing column exceeding by a factor $\\sim$3 the Galactic value in the direction of \\grb\\ \\citep{rowlinson09}.  On the optical/infrared side, several follow-up observations of the field of \\grb\\ were performed. A possible, faint transient with very red colors was detected in early observations in the near infrared, performed within a few minutes from $T_0$ by automated instruments such as the PAIRITEL Telescope \\citep*{morgan09}, the Palomar Observatory's 60-inch telescope \\citep{cenko09} as well as the Faulkes North Telescope \\citep{guidorzi09}. The peculiar colors of such a transient prompted to suggest a very high redshift ($z\\sim10$) for such GRB. The same infrared source was possibly detected in a Subaru image collected $\\sim$2 hours after the trigger \\citep{aoki09}. A very deep observation performed with the 10.4 m Gran Telescopio CANARIAS telescope $\\sim$41 hours after the GRB detected no source at the same position, down to a $3\\sigma$ upper limit $I>25.5$ \\citep{castrotirado09}. \\begin{figure} \\centering \\resizebox{\\hsize}{!}{\\includegraphics[angle=-90]{lcbat.eps}} \\caption{\\label{lcbat} Light curve of \\grb\\ obtained with the \\swift/BAT in the 15--350 keV. The time binning is 0.5 s. } \\end{figure}\\\\ \\indent \\grb\\ attracted much interest because of a very peculiar timing phenomenology. \\citet{markwardt09}  reported evidence for quasi-periodical pulsations at $\\sim$8 s in the prompt emission, based on analysis of \\swift/BAT data in the 15--350 keV energy range. Such a quasi-periodicity was then confirmed by the analysis of Konus-Wind and Konus-RF data \\citep{golenetskii09}, as well as by analysis of \\igr/ACS data \\citep{gotz09}. Taken at face value, such a result is tantalizing. This would be the first detection of periodic activity in the prompt gamma-ray emission of a GRB, which could yield rare information on the behavior of the central engine.\\\\ \\indent Indeed, the peculiar quasi-periodicity immediately prompted several authors (e.g. \\citealt{markwardt09} and \\citealt{guidorzi09}) to suggest that \\grb\\ could be different in origin with respect to standard long GRBs, and possibly related to activity of a magnetar, either within the Galaxy, or extragalactic. However, a search for pulsations in early afterglow data collected by \\swift/XRT yielded null results \\citep{mirabal09}.\\\\ \\indent To shed light on the nature of the peculiar \\grb\\, we asked for an \\xmm\\ Target of Opportunity (ToO) observation, aimed at characterizing the afterglow emission at a later stage (about 2 days after the event), when the target is too faint to be studied with \\swift/XRT. Here we report on the outcome of the \\xmm\\ ToO, which, coupled to a re-analysis of \\swift\\ and \\igr\\ data, allowed us to obtain a comprehensive description of the high-energy phenomenology of \\grb. ",
        "conclusions": "\\label{disc} \\subsection{The afterglow of a distant, standard, long GRB} Prompted by early reports of the detection of a strong quasi-periodical signal in the prompt emission of \\grb, we have performed a ToO observation with \\xmm.\\\\ \\indent The afterglow of \\grb\\, as seen in soft X-rays with the EPIC instrument, is fully similar to other cases of typical, well-behaving, long GRBs. No pulsations are seen in the X-ray emission $\\sim$48 hr after the trigger, with a $3\\sigma$ upper limit of 15--25\\% on the pulsed fraction, assuming a sinusoidal pulse shape. The energy spectrum has a typical power-law shape, with a remarkable excess in photoelectric absorption (by a factor $\\sim$4) with respect to the Galactic $N_{\\rm H}$ in the direction of \\grb.\\\\ \\indent Such a picture of a ``standard'' X-ray afterglow is completed by \\swift/XRT observations, which started as soon as 77 s after the burst, and extended up to 20 days after the trigger. \\grb\\ turns out to have a rich phenomenology, featuring -  in addition to the intrinsic photoeletric absorption - a clear spectral evolution (the photon index $\\Gamma$ steepens from  1.7 to 2.0 in $\\sim$2 days), as well as a break in the flux decay, with a $\\sim0.3$ variation in the power law decay index occurring $\\sim$0.26 days after the burst.\\\\ \\indent We will focus here on the extra absorption only, the evidence for which is fully confirmed - and strenghtened - by \\swift\\ data. Indeed, adopting a simple, redshifted neutral absorption model, all datasets yield consistent values for both the redshift and the redshifted absorbing column, as well as of the spectral steepening as a function of time. \\xmm\\ and \\swift/XRT data point  to a rather high redshift ($z\\sim4.2$) for \\grb\\ as well as to a huge column density in the host galaxy ($N_{\\rm H}\\sim6\\times10^{22}$ cm$^{-2}$). Although a word of caution is required in considering such results (see Sect.~\\ref{spectra}), to our knowledge, this is the largest intrinsic absorption ever observed in an X-ray afterglow spectrum \\citep[see e.g.][]{deluca05,campana06short,grupe07}.\\\\ \\indent This high absorption is also on the high side among expected column densities for sources located inside molecular clouds \\citep{reichart02}. We note that part of the absorption could be due to intervening systems along the line of sight rather than being local to the host galaxy, as seen in the optical band in QSO and GRB studies \\citep[the so-called Damped Lyman-$\\alpha$ Absorbers;][]{reichart02,wolfe05}. Lack of optical spectroscopy in the case of \\grb\\ prevents from drawing firm conclusions.\\\\ \\indent Such a picture suggests an explanation of the peculiar, very red colors of the infrared transient likely associated to \\grb\\ \\citep{morgan09} as due to  reddening in the host galaxy. Indeed, if the GRB spectrum has no breaks in the X-ray to optical range and the redshift and intrinsic column derived from X-ray spectroscopy are correct, an extinction A$_V\\sim3$ in the GRB host galaxy (at $z=4.2$) would fit the infrared data. This would point to a dust-to-gas ratio $\\sim$10\\% of that found in the Milky Way (assuming Solar abundances), similar to the findings of other investigations \\citep[e.g. ][]{galama01,hjorth03short}. Such an interpretation of the infrared data is possibly supported by the time decay of the infrared transient. The fading between the PAIRITEL and the Subaru observations is consistent with the $t^{-1.15}$ law describing the X-ray afterglow decay before the break at $T_0+0.26$ days, which suggests a common origin for the X-ray and infrared emission.\\\\ \\indent In any case, X-ray spectroscopy yields a robust indication that \\grb\\ was a very distant event.  \\subsection{On the prompt temporal variability}  Coming back to the temporal properties of the prompt emission, our reanalysis, based on both \\swift/BAT and \\igr\\ SPI/ACS data, could not confirm the presence of any significant signal at $\\sim$8.1 s. In any case, a complex multi-peaked light curve is clearly apparent, with a peak-to-peak separation of order $\\sim$8--10 s, and a peak at 0.125\\,Hz is present in the power spectrum derived from BAT data, just below the 3$\\sigma$ detection threshold. It is premature to comment on the apparent discrepancy between our results and those reported in a circular by \\citet{markwardt09}, in view of the lack of detailed information concerning the analysis performed by these authors. Indeed, several issues affect estimation of the significance of a signal when dealing with de-trending algorithms and Fourier transform techniques. Among the most important, we remember the presence of non-Poissonian noise in the power spectrum and the change in the statistical properties of a time series induced by operations such as subtraction and division. The former, if not taken into account properly, may result in an overestimation of the statistical significance of any peak sitting on a non-Poissonian underlying power spectrum continuum (see \\citealt{israel96}). The latter affects the statistical properties of the original time-series. In particular, the division of a Possonian variable (the time series) with a (model-dependent) de-trending algorithm might result in a non-Poissonian distribuited time series with effects on the (unknown) statistical properties of the power spectrum noise.\\\\ \\indent The claimed period of about 8 s is precisely in the range of periods of soft gamma-ray repeaters (SGRs) and anomalous X-ray pulsars (AXPs). These sources are thought to be magnetars: isolated neutron stars powered by strong (10$^{14}$--10$^{15}$ G) magnetic fields (see \\citealt{mereghetti08} for a review). Thus, association of \\grb\\ with the activity of a magnetar was suggested, also possibly based on some similarity of the light curve to the one of a magnetar giant flare, featuring a sort of ``pulsating tail''.\\\\ \\indent Based on high energy properties, we can rule out such an interpretation. First, we could find no unambiguous, coherent pulsations in the $\\gamma$-ray emission. Then, the spectrum of \\grb\\ is much harder than the typical spectra observed in the pulsating tails of the giant flares from SGRs. These are well described by thermal bremsstrahlung emission with $kT\\sim15$--30 keV,  while the spectrum of \\grb\\  is typical for a GRB.  Giant flares are also characterized by a short ($<$0.5 s), very bright and spectrally hard initial spike that was clearly absent in \\grb\\  (although this feature was lacking in some ``intermediate flares'' from SGRs, all of them had a very sharp initial rise, contrary to the light curve of \\grb ). The location at high Galactic latitude ($b\\sim20\\degr$) would also be very unusual for a Galactic SGR. Moreover, a Galactic origin may be safely ruled out by the observation of an absorbing column exceeding by a factor $\\sim$4 the Galactic value in the direction of \\grb. A giant flare from an extragalactic SGR can also be excluded, based on the absence of a visible nearby galaxy at the location of \\grb. The observed fluence would imply a distance not larger than 600 kpc (assuming an energetic similar to the one of other observed magnetar giant flares). The enormous energy requirement implied by a cosmological distance rules out the SGR giant flare hypothesis.\\\\ \\indent As already discussed in the previous section, we conclude that \\grb\\ is a standard, long GRB, with a multi-peak structure in the prompt emission. We will not go into pure speculations by considering physical processes that could produce a true quasi-periodic signal.\\\\ \\indent The variable prompt emission could be the signature of non-stationary processes in the GRB inner engine.  As a likely  possibility -- discussed by \\citet{beskin09} who observed a similar phenomenology in the optical emission from the ``naked-eye'' GRB\\,080319B (but with a possibly stronger evidence for quasi-periodicity) -- the peculiar variability could be related to cyclic accretion by the central newborn compact object. Such a phenomenon could be due to the fragmentation of an accretion disc due to some kind of instability. For instance, \\citet{masada07} explain short-time variability in the prompt emission by GRBs as due to magneto-rotational instabilities developing in a massive, hot hyperaccreting disc surrounding a central black hole of a few stellar masses. At a redshift $z\\sim4.2$, the variability time scale in the source frame would be of $\\sim1.5$ s. Thus, adopting the model of \\citet{masada07}, the observed properties of \\grb\\ could fit into such a scenario by assuming a beaming factor $\\sim$100, for a $\\sim$1 $M_{\\odot}$ accretion disc with inner radius of $\\sim$30 gravitational radii, surrounding a $\\sim$$4M_{\\odot}$ central black hole."
    },
    "0911/0911.1003_arXiv.txt": {
        "abstract": "We present a spectral variability study of the \\xmm\\ and \\suzaku\\ observations of one of the most extreme Narrow Line Seyfert 1 galaxies, \\iras.  The X-ray spectrum is characterized by two main peculiar features, i) a strong soft excess with a steep rise below about 1.3 keV and ii) a deep drop in flux above 8.2 keV. Although absorption--based interpretations may be able to explain these features by a suitable combination of ionization, covering factors, column densities and outflowing velocities, we focus here on a reflection--based interpretation which interprets both features, as well as the large soft excess, in terms of partially ionized reflection off the inner accretion disc. We show that the two peculiar spectral features mentioned above can be reproduced by two relativistic emission lines due to Fe K and Fe L. The lines are produced in the inner accretion disc and independently yield consistent disc parameters. We argue that the high L/K intensity ratio is broadly consistent with expectations from an ionized accretion disc reflection, indicating that they belong to a single ionized reflection component. The spectral shape, X--ray flux, and variability properties are very similar in the \\xmm\\ and \\suzaku\\ observations, performed about 5 years apart. The overall X--ray spectrum and variability can be described by a simple two--component model comprising a steep power law continuum plus its ionised reflection off the inner accretion disc. In this model, a rapidly rotating Kerr black hole and a steep emissivity profile are required to describe the data.  The simultaneous detection of broad relativistic Fe L and K lines in IRAS 13224-3809 follows that in another extreme NLS1 galaxy, 1H~0707--495. Although the data quality for \\iras\\ does not allow us to rule out competing models as in 1H~0707--495, we show here that our reflection-based interpretation describes in a self--consistent manner the available data and points towards \\iras\\ being a very close relative of 1H~0707--495 in terms of both spectral and variability properties. These results, together with those based on pure broad Fe K detections, are starting to unveil the processes taking place in the immediate vicinity of accreting radiatively efficient black holes. ",
        "introduction": "There is little doubt that narrow-line Seyfert 1 (NLS1) galaxies provide an extreme view of the AGN phenomenon (e.g. Boller \\etal\\ 1996, 2002; Brandt \\etal\\ 1997; Fabian \\etal\\ 2002, 2004; Gallo 2006; Gierlinski \\etal\\ 2008).  A strong soft excess below $\\sim$2 keV and significant X-ray variability are often seen, and with \\xmm\\ came the discovery of high-energy spectral drop and curvature in some objects (e.g. Boller \\etal\\ 2002, 2003).  The physical interpretation is a topic of debate, and the suggestions are as extreme as the objects themselves involving relativistically blurred reflection (e.g. Fabian \\etal\\ 2002, 2004; Ponti et al. 2006) or ionised absorption (Gierlinski \\& Done 2004; Middleton \\etal\\ 2007), and partial covering (Tanaka \\etal\\ 2004, 2005) that may be outflowing (Gallo \\etal\\ 2004a).  A breakthrough was the discovery of broad iron (Fe) L$\\alpha$ emission and reverberation-like time lags in a long observation of 1H~0707--495 (Fabian \\etal\\ 2009).  Confirmation of this type of behaviour in other objects would bolster the importance of NLS1s in understanding the innermost regions of AGN.  \\iras\\ ($z=0.0667$) along with 1H~0707--495 are perhaps the most remarkable members of this class, displaying the above properties in outstanding fashion.  In many ways \\iras\\ and 1H~0707--495 are very similar objects, especially in the X--ray band.  They both show a high energy Fe K drop around 7-8 keV implying either an outflowing partial covering absorber or a large emission line (Boller et al. 2002; Fabian et al. 2004; Boller et al. 2003). In both cases the material producing the structure must be iron overabundant with respect to the Sun. They also have large soft excesses (Ponti PhD thesis) with sharp drops around 1 keV (Leighly 1999). Moreover they show a high degree of variability indicating low black hole masses and high accretion rates. %  The motivation of this present work is to examine \\iras\\ in the light of the discoveries from the new observations of 1H~0707--495 (Fabian \\etal\\ 2009). To do so we make use of a previously analysed \\xmm\\ observation (e.g. Boller \\etal\\ 2003; Gallo \\etal\\ 2004b; Ponti et al. 2006) along with a new \\suzaku\\ observation.  The purpose here is not to duplicate previous efforts by fitting the spectra with partial-covering (or other absorption models) and reflection models and comparing quality of fits. Rather, our goal is to test the robustness of the reflection model and show it to be a valid interpretation, similar to that used for 1H0707-495. ",
        "conclusions": "\\iras\\ is a remarkable source, with one of the strongest soft excesses observed, and two prominent features at $\\sim$1.2 and $\\sim$8.2 keV.  \\begin{itemize} \\item{} These features can be reproduced by 2 broad lines, in particular: i) the best fit energies of the lines are consistent with emission from Fe~L and Fe~K; ii) the same relativistic profile can explain the shape of both lines; iii) the ratio between the intensities of the lines is broadly in agreement with expectations. \\item{} A single ionized disc reflection model (produced in the inner accretion disc) can reproduce the main features of the spectrum, implying a iron abundance of $\\sim$5 times solar. \\item{} The main spectral variations can be reproduced by a steep power law varying in normalisation and a reflection component from the inner accretion disc varying in a correlated way, but with a residual component, dominating at low fluxes. \\item{} The spectral features may, in an alternative interpretation, be associated with absorption components. Nevertheless in this scenario the two main features require at least two physically different absorbing components in highly relativistic outflow and without a clear physical link. A high iron abundance is also required. Moreover, absorption-dominated models over--predict Fe M absorption in the soft X--rays, and they still require a likley unphysical additional blackbody component to describe the soft excess, casting further doubts on the overall interpretation.  \\end{itemize}"
    },
    "0911/0911.2000_arXiv.txt": {
        "abstract": "We present the results of a photometric and spectroscopic survey of 321 Lyman break galaxies (LBGs) at \\z3 to investigate systematically the relationship between \\lya emission and stellar populations. \\lya equivalent widths ($W_{\\rm Ly \\alpha}$) were calculated from rest-frame UV spectroscopy and optical/near-infrared/\\emph{Spitzer} photometry was used in population synthesis modeling to derive the key properties of age, dust extinction, star formation rate (SFR), and stellar mass. We directly compare the stellar populations of LBGs with and without strong \\lya emission, where we designate the former group ($W_{\\rm Ly \\alpha}$ $\\ge$ 20 \\AA) as Ly$\\alpha$-emitters (LAEs) and the latter group ($W_{\\rm Ly \\alpha}$ $<$ 20 \\AA) as non-LAEs. This controlled method of comparing objects from the same UV luminosity distribution represents an improvement over previous studies in which the stellar populations of LBGs and narrowband-selected LAEs were contrasted, where the latter were often intrinsically fainter in broadband filters by an order of magnitude simply due to different selection criteria. Using a variety of statistical tests, we find that \\lya equivalent width and age, SFR, and dust extinction, respectively, are significantly correlated in the sense that objects with strong \\lya emission also tend to be older, lower in star formation rate, and less dusty than objects with weak \\lya emission, or the line in absorption. We accordingly conclude that, within the LBG sample, objects with strong \\lya emission represent a later stage of galaxy evolution in which supernovae-induced outflows have reduced the dust covering fraction. We also examined the hypothesis that the attenuation of \\lya photons is lower than that of the continuum, as proposed by some, but found no evidence to support this picture. ",
        "introduction": "An increasing number of high-redshift galaxies have been found in the last two decades using selection techniques reliant on either color cuts around the Lyman limit at 912 \\AA\\ in the rest frame \\citep[e.g.,][]{steidel1996a,steidel1999} or strong \\lya line emission \\citep[e.g.,][]{cowie1998,rhoads2000,gawiser2006}. These two methods, which preferentially select Lyman break galaxies (LBGs) and Ly$\\alpha$-emitters (LAEs), respectively, have successfully isolated galaxies at redshifts up to $z$ = 7 \\citep{iye2006,bouwens2008}. Extensive data sets of LBGs and LAEs have afforded detailed studies of galactic clustering \\citep[e.g.,][]{adelberger1998,giavalisco2001}, the universal star formation history \\citep[e.g.,][]{madau1996,steidel1999}, and the galaxy luminosity function \\citep[e.g.,][]{reddy2008,mclure2009}. While the nature of LBGs and LAEs have been studied at a range of redshifts \\citep[e.g.,][]{shapley2001,gawiser2006,verma2007,pentericci2007,nilsson2009,ouchi2008,finkelstein2009}, the epoch around \\z3 is particularly well-suited to investigation of these objects' detailed physical properties. At this redshift, the prominent H\\Rmnum{1} \\lya line ($\\lambda_{\\rm rest}$ = 1216 \\AA), present in all LAE spectra and a significant fraction of LBG spectra, is shifted into the optical, where current imaging and spectroscopic instrumentation is optimized. Consequently, there are large existing data sets of spectroscopically-confirmed \\z3 LBGs \\citep[e.g.,][]{steidel2003} and LAEs \\citep[e.g.,][]{lai2008}, where extensive multiwavelength surveys often complement the former and, less frequently, the latter.  The mechanism responsible for LAEs' large \\lya equivalent widths is not fully understood, although several physical pictures have been proposed \\citep[e.g.,][]{dayal2009,kobayashi2010}. As \\lya emission is easily quenched by dust, one explanation for LAEs is that they are young, chemically pristine galaxies experiencing their initial bursts of star formation \\citep[e.g.,][]{hu1996,nilsson2007}. Conversely, LAEs have also been proposed to be older, more evolved galaxies with interstellar media in which dust is segregated to lie in clumps of neutral hydrogen  surrounded by a tenuous, ionized dust-free medium \\citep{neufeld1991,hansen2006,finkelstein2009}. In this picture, \\lya photons are resonantly scattered near the surface of these dusty clouds and rarely encounter dust grains. Continuum photons, on the other hand, readily penetrate through the dusty clouds and are accordingly scattered or absorbed. This scenario preferentially attenuates continuum photons and enables resonant \\lya photons to escape relatively unimpeded, producing a larger \\lya equivalent width than expected given the underlying stellar population. To date, the distribution of dust in the interstellar medium has only been investigated using relatively small samples \\citep[e.g.,][]{verhamme2008,atek2009,finkelstein2009}.  Given the different selection techniques used to isolate LBGs and LAEs, understanding the relationship between the stellar populations of these objects has been an important goal of extragalactic research. Recent work by \\citet{gawiser2006} has suggested that LAEs are less massive and less dusty than LBGs, prompting these authors to propose that LAEs may represent the beginning of an evolutionary sequence in which galaxies increase in mass and dust content through successive mergers and star formation episodes \\citep{gawiser2007}. The high specific star formation rate -- defined as star formation rate (SFR) per unit mass -- of LAEs \\citep[$\\sim$ 7 $\\times$ 10$^{-9}$ yr$^{-1}$;][]{lai2008} relative to LBGs \\citep[$\\sim$ 3 $\\times$ 10$^{-9}$ yr$^{-1}$;][]{shapley2001} illustrates that LAEs are building up stellar mass at a rate exceeding that of continuum-selected galaxies at \\z3. This rapid growth in mass is consistent with the idea that LAEs represent the beginning of an evolutionary sequence of galaxy formation. However, results from \\citet{finkelstein2009} cast doubt on this simple picture of LAEs as primordial objects, given that these authors find a range of dust extinctions ($A_{\\rm 1200}$ = 0.30--4.50) in a sample of 14 LAEs at \\z4.5. \\citet{nilsson2009} also find that \\z2.25 LAEs occupy a wide swath of color space, additional evidence that not all LAEs are young, dust-free objects. Furthermore, the assertion that LAEs are pristine galaxies undergoing their first burst of star formation is called into question by the results of \\citet{lai2008}. These authors present a sample of 70 \\z3.1 LAEs, $\\sim$ 30\\% of which are detected in the 3.6 $\\mu$m band of the \\emph{Spitzer} Infrared Array Camera \\citep[IRAC;][]{fazio2004}. These IRAC-detected LAEs are significantly older and more massive ($\\langle t_{\\star} \\rangle$ $\\sim$ 1.6 Gyr, $\\langle M \\rangle$ $\\sim$ 9 $\\times$ 10$^{9}$) than the IRAC-undetected sample ($\\langle t_{\\star} \\rangle$ $\\sim$ 200 Myr, $\\langle M \\rangle$ $\\sim$ 3 $\\times$ 10$^8$ $M_{\\odot}$); \\citet{lai2008} suggest that the IRAC-detected LAEs may therefore be a lower-mass extension of the LBG population. Narrowband-selected LAEs are clearly marked by heterogeneity, and the relationship between these objects and LBGs continues to motivate new studies.  When comparing the stellar populations of LBGs and LAEs, it is important to take into account the selection biases that result from isolating these objects with broadband color cuts and line flux/equivalent width requirements, respectively. By virtue of selection techniques that rely on broadband fluxes and colors, LBGs generally have brighter continua than LAEs. Spectroscopic samples of LBGs typically have an apparent magnitude limit of ${\\cal R}$ $\\le$ 25.5 \\citep[0.4$L^*$ at \\z3;][]{steidel2003} while LAEs have a median apparent magnitude of $R$ $\\sim$ 27 \\citep{gawiser2006}, where ${\\cal R}$ and $R$ magnitudes are comparable. Even though the majority of LBGs studied to date are an order of magnitude more luminous in the continuum than typical LAEs, both populations have similar rest-frame UV colors \\citep{gronwall2007}. Therefore, LAEs fainter than ${\\cal R}$ = 25.5 are excluded from LBG spectroscopic surveys not because of their colors, but rather because of their continuum faintness. Given the significant discrepancy in absolute magnitude between LBGs and LAEs, understanding the relationship between these objects can be fraught with bias. An preferable approach to comparing these populations is to investigate how the strength of \\lya emission is correlated with galaxy parameters, for a controlled sample of objects at similar redshifts \\emph{drawn from the same parent UV luminosity distribution}.  Several authors have looked at the question of the origin of \\lya emission in UV flux-limited samples \\citep[e.g.,][]{shapley2001,erb2006a,reddy2008,pentericci2007,verma2007}. \\citet{shapley2001} analyzed 74 LBGs at \\z3 and constructed rest-frame UV composite spectra from two samples of ``young\" ($t_{\\star}$ $\\leq$ 35 Myr) and ``old\" ($t_{\\star}$ $\\geq$ 1 Gyr) galaxies, respectively.  These authors found that younger objects exhibited weaker \\lya emission than older galaxies; \\citet{shapley2001} attributed the difference in emission strength to younger LBGs being significantly dustier than their more evolved counterparts. On the other hand, \\citet{erb2006a} examined a sample of 87 star-forming galaxies at \\z2 and found that objects with lower stellar mass had stronger \\lya emission features, on average, than more massive objects. In a sample of 139 UV-selected galaxies at \\z2--3, \\citet{reddy2008} isolated 14 objects with \\lya equivalent widths $\\ge$ 20 \\AA\\ and noted no significant difference in the stellar populations of strong Ly$\\alpha$-emitters relative to the rest of the sample. \\citet{pentericci2007} examined 47 LBGs at \\z4 and found that younger galaxies generally showed \\lya in emission while \\lya in absorption was associated with older galaxies (in contrast to the \\citet{shapley2003} results). Probing even earlier epochs, \\citet{verma2007} examined a sample of 21 LBGs at \\z5 and found no correlation between \\lya equivalent width and age, stellar mass, or SFR. These authors noted, however, that only 6/21 of the brightest LBGs had corresponding spectroscopy from which equivalent widths were estimated. Therefore, the lack of a correlation between \\lya equivalent width and stellar populations may, in this case, have been masked by a small sample that was biased towards the brightest objects. These aforementioned investigations have shown that there does not yet exist a clear picture relating stellar populations to \\lya emission.  \\begin{figure}[t] \\centering \\includegraphics[trim=0in 0in 0in 0in,clip,width=3in]{f1.eps} \\caption{Redshift distribution of the sample, where $\\langle z \\rangle$ = 2.99 $\\pm$ 0.19.} \\label{fig: redshift} % \\end{figure}  In this paper, we present a precise, systematic investigation of the relationship between \\lya emission and stellar populations using our large photometric and spectroscopic data set of \\z3 observations. As an improvement over previous studies, we approach the data analysis from multiple aspects: we compare not only the stellar population parameters derived from population synthesis modeling, but also examine the objects' best-fit SEDs and photometry. Furthermore, all analysis is conducted on objects drawn from the same parent sample of continuum-bright (${\\cal R}$ $\\le$ 25.5) LBGs. By controlling for continuum magnitude, we avoid the biases of comparing objects with significantly different luminosities while still retaining the ability to comment on the nature of strong Ly$\\alpha$-emitting galaxies within the LBG sample. Our conclusions are applicable to both LBGs and bright ({$\\cal R$} $\\le$ 25.5) narrowband-selected LAEs. While we are unable to make inferences about the population of faint LAEs, our study is complete with respect to bright LAEs given these objects' similar colors to LBGs in the rest-frame UV.  We are motivated by the following questions: how do the stellar populations of Ly$\\alpha$-emitting LBGs differ from those of other LBGs at \\z3 where the \\lya emission line is weaker (or absent altogether)? To what degree are galactic parameters such as dust extinction, SFR, age, and stellar mass correlated  with \\lya line strength? What do the relative escape fractions of \\lya and continuum photons reveal about the distributions of gas and dust in these objects' interstellar media?  This paper is organized as follows. In \\S \\ref{sec:sample}, we present details of the observations and data reduction, including a description of the systematic technique used to calculate \\lya equivalent widths. Stellar population modeling is discussed in \\S \\ref{sec:pops}. The properties of objects with and without strong \\lya emission are presented in \\S \\ref{sec: results} and we discuss how our data can be used to address several of the outstanding questions pertaining to the physical nature of LBGs and LAEs in \\S \\ref{sec: disc}. A summary and our conclusions appear in \\S \\ref{sec: sum}. We assume a standard $\\Lambda$CDM cosmology throughout with \\emph{H}$_{\\rm 0}$ = 70 \\kms\\ Mpc$^{-1}$, $\\Omega_{\\rm M}$ = 0.3, and $\\Omega_{\\Lambda}$ = 0.7. All magnitudes are based on the AB system \\citep{oke1983}\\footnote{AB magnitude and $f_{\\nu}$, the flux density in units of ergs s$^{-1}$ cm$^{-2}$ Hz$^{-1}$, are related by \\emph{m}$_{\\rm AB}$ = $-$2.5~log$_{10}~f_{\\nu}-$48.6. Conversions between AB and Vega magnitudes for the near-infrared passbands are as follows: $K_s$(AB) = $K_s$(Vega) + 1.82; $J$(AB) = $J$(Vega) + 0.90.}, with the exception of the infrared passbands which are in the Vega system. ",
        "conclusions": "\\label{sec: sum}  We have analyzed a sample of 321 optically-selected \\z3 Lyman break galaxies in order to investigate the relationship between stellar populations and \\lya emission. The equivalent width of the \\lya feature was robustly estimated from rest-frame UV spectroscopy and broadband $G{\\cal R}JK_s$ + $Spitzer$ IRAC photometry was used to conduct stellar population modeling to derive the key properties of age, extinction, star formation rate, and stellar mass for all objects with photometric coverage in at least one near- or mid-infrared passband ($\\sim$ 80\\% of the sample). The limited luminosity range of LBGs enabled a controlled investigation of the nature of \\lya emission \\emph{within a sample of objects spanning similar luminosities.} Given the luminosity-dependent trends in \\lya equivalent width, stellar populations, and clustering \\citep[e.g.,][]{gronwall2007,lai2008,adelberger2005a}, this controlled study represents an improvement over previous work. We used a variety of statistical tests to analyze the correlation between \\lya emission and stellar populations in LBGs from the standpoint of comparing spectroscopic observations, photometric data, and best-fit stellar population parameters. The relative \\lya escape fraction was used to probe the relative distributions of gas and dust in the objects' interstellar media. Below, we summarize our results.  \\begin{itemize} \\item \\lya equivalent width is correlated with age, SFR, and E(B--V), respectively, in the sense that galaxies with strong \\lya emission are older, more quiescent, and less dusty than their counterparts with weak or absent \\lya emission. Taking into account the uncertainties on the best-fit population parameters, the probability of a null hypothesis (i.e., no correlation) was excluded at the 3--5$\\sigma$ level for these relationships. We found that stellar mass was not significantly correlated with \\ewlya\\ -- the null hypothesis could be excluded at only the $\\sim$ 1$\\sigma$ level. These results were consistently derived from several different analyses of the data, including correlation tests, composite spectra, and the binary division of the sample into LAEs ($W_{\\rm Ly \\alpha}$ $\\ge$ 20 \\AA) and non-LAEs ($W_{\\rm Ly \\alpha}$ $<$ 20 \\AA).  \\item Analysis of the relative escape fraction of Ly$\\alpha$ is consistent with \\lya photons experiencing more attenuation than non-resonance continuum photons. We also find that the \\lya escape fraction is strongly correlated with E(B--V), where galaxies with more dust attenuation also have lower escape fractions. \\end{itemize}  The observed correlations between \\lya emission and stellar populations are consistent with the physical picture proposed by \\citet{shapley2001}, in which young, dusty LBGs experience vigorous outflows from supernovae and massive star winds. \\citet{shapley2003} reported evidence that both gas and dust are entrained in the outflows, evidence that these ``superwinds\" could tenably decrease a galaxy's dust and gas covering fraction over several tens of Myr. While the more mature LBGs may have as much or more overall dust content than younger galaxies, the fact that these older objects ($t_{\\star}$ $\\gtrsim$ 100 Myr) are also more likely to exhibit \\lya emission is an important point: this trend demonstrates that it is the dust covering fraction -- not the total amount of dust -- that modulates \\lya emission. Independent evidence for a lower covering fraction of dusty gas in older objects is provided by a trend toward weaker low-ionization interstellar absorption lines with increasing age in the composite spectra presented in \\S \\ref{sec: bootstrap}.  While this physical picture has been put forth previously, our current results are more strongly supported due to the wide range of statistical tests and analysis methods that we employed. We compared our larger sample of data in different forms -- spectroscopic, photometric, and best-fit parameters -- using a battery of statistical methods including correlation and K--S tests. \\lya emission strength was furthermore treated as both an independent and dependent variable. These analyses represent a systematic, powerful approach to elucidating trends.  Results of both the escape fraction and the relative escape fraction can be used to further constrain the physical picture of LBGs and bright narrowband-selected LAEs. While the observed inverse relationship between E(B--V) and $f_{esc}$ is consistent with earlier results \\citep[e.g.,][]{verhamme2008,atek2009}, the scatter present in our larger sample highlights the myriad factors modulating the \\lya escape fraction (e.g., galaxy kinematics, dust, and outflow geometry). Further analysis of this scatter will be instrumental in constraining how the \\lya escape fraction varies. The observed relative escape fraction of \\lya emission -- where the majority of LBGs have $f_{esc,rel}$ below unity -- is indicative of a physical picture in which LBGs contain gas and dust that is well mixed. In this geometry, \\lya photons experience more attenuation than continuum photons do because of the former's increased path lengths from resonant scattering. This is in contrast to the scenario proposed by \\citet{finkelstein2009}, in which dust is segregated in neutral gas clumps surrounded by an ionized, dust-free medium. We caution, however, that the relative escape fraction of \\lya is likely modified by a variety of factors including galaxy kinematics, dust, and outflow geometry. As such, additional observations are necessary in order to more fully understand the distribution (and evolution) of gas and dust in LBGs.  While we hope that this work has illuminated the relationship between \\lya emission and stellar populations in \\z3 LBGs, future systematic studies at other redshifts -- ideally over larger dynamic ranges in UV luminosity -- are necessary in order to address the discrepancies between trends reported here and those at \\z2 or \\z4 \\citep[e.g.,][]{erb2006a,pentericci2007}. The nature of complex \\lya emission morphologies -- such as the double-peaked objects that comprise $\\sim$ 4\\% of the sample discussed here -- has also yet to be explored in light of understanding if these objects' stellar populations and gas distributions differ from those of more typical Ly$\\alpha$-emitters or absorbers. Similarly, an analysis of the geometry of dust and gas in the interstellar media of \\z3 LBGs is vital to understanding what kind of dust attenuation law -- the \\citet{calzetti2000} relation, a SMC-like law, or another relation completely -- most accurately describes these objects. Furthermore, a basic property critical to placing galaxies in an evolutionary context -- metallicity -- remains poorly constrained for the \\z3 LBG population \\citep[but see][]{maiolino2008}. While metallicity has been investigated in the rare cases of gravitationally lensed objects \\citep[e.g.,][]{teplitz2000,pettini2002,hainline2009,quider2009}, a systematic study of elemental abundances in a larger sample of \\z3 LBGs would be illuminating \\citep[akin to the \\z2 sample discussed in][]{erb2006a}. Given anticipated advances in both telescopes and instrumentation, we hope that these questions are illustrative of many others that will help to guide future LBG research."
    },
    "0911/0911.0005_arXiv.txt": {
        "abstract": "The star formation rate (SFR) and black hole accretion rate (BHAR) functions are measured to be proportional to each other at $z \\la 3$. This close correspondence between SF and BHA would naturally yield a BH mass--galaxy mass correlation, whereas a BH mass--bulge mass correlation is observed. To explore this apparent contradiction we study the SF in spheroid-dominated galaxies between $z=1$ and the present day. We use 903 galaxies from the COMBO-17 survey with M$_\\ast >2\\times 10^{10}$\\,M$_\\odot$, ultraviolet and infrared-derived SFRs from Spitzer and GALEX, and morphologies from GEMS HST/ACS imaging. Using stacking techniques, we find that $<$25\\% of all SF occurs in spheroid-dominated galaxies (S\\'ersic index $n>$2.5), while the BHAR that we would expect if the global scalings held is three times higher. This rules out the simplest picture of co-evolution, in which SF and BHA trace each other at all times. These results could be explained if SF and BHA occur in the same events, but offset in time, for example at different stages of a merger event. However, one would then expect to see the corresponding star formation activity in early-stage mergers, in conflict with observations.  We conclude that the major episodes of SF and BHA occur in different events, with the bulk of SF happening in isolated disks and most BHA occurring in major mergers. The apparent global co-evolution results from the regulation of the BH growth by the potential well of the galactic spheroid, which includes a major contribution from disrupted disk stars. ",
        "introduction": "The last decade has seen the discovery and characterization of an unexpectedly tight correlation between the mass of supermassive black holes (SMBH; $M_{\\rm BH}$) and the mass ($M_{\\rm bulge}$) or velocity dispersion of their host galaxy's bulge \\citep{Magorrian98,Ferrarese00,Gebhardt00,Marconi03,Haring04}.  The black hole mass appears to be most strongly correlated with the bulge mass, not the total mass \\citep{Kormendy01,Ho07}; the scatter in this relation between bulge mass and black hole mass is estimated to be less than a factor of two \\citep{Haring04}.  This relationship indicates that galaxy and black hole formation and evolution are interconnected.\\footnote{\\citet{Peng07}, noting that multiple generations of mergers would reduce the scatter of a weak or non-existent $M_{\\rm bulge}-M_{\\rm BH}$ relation essentially through the central limit theorem, found that more than 10 generations of major mergers would be required to imprint a tight $M_{\\rm bulge}-M_{\\rm BH}$ relation in a population that was initially uncorrelated.  Such a large number of major mergers is highly unlikely, arguing that while mergers may tighten this relation much of the correlation must be imprinted through a strong interconnection between $M_{\\rm bulge}$ and $M_{\\rm BH}$.} In its weak form, such an interconnection could result if the black hole growth is limited by the wider galaxy environment (`dog wagging the tail'). Stronger forms of interconnection are also possible. The energy released by a growing SMBH is sufficient, if it couples effectively with its surroundings, to have dramatic consequences on the galaxy (`tail wagging the dog'). For example, the injection of energy into hot halo gas in galaxy clusters is indicated by the inflation of bubbles and the propagation of sound waves \\citep{McNamara00,McNamara05,Fabian03,Rafferty06}. Furthermore, it is possible that SMBHs with high accretion rates drive powerful outflows that remove (at least some) cold gas from galaxies \\citep{Chartas03,Crenshaw03,Pounds03,Tremonti07}. In fact, feedback from an accreting SMBH/AGN may help address many of the most difficult issues affecting models of galaxy evolution in a cosmological contest: e.g., overcooling in massive halos and late star formation in elliptical galaxies \\citep{Binney95,Croton06,Bower06,Monaco07,Somerville08}.   A diversity of models has been constructed to explore the interrelationship between bulge mass and black hole mass, featuring both weak and strong aspects of the possible interconnections \\citep[][and references therein]{Somerville08}.  However, the details of the physical processes governing BH formation and growth, and the physical mechanisms whereby SMBHs couple with their host galaxies and surroundings remain poorly understood.  It will be impossible to model the full range of interconnected processes from theoretical first principles for some time, as the physics of the formation of SMBHs and galaxy evolution spans at least $\\simeq$ 7 orders of magnitude in scale. Empirical and phenomenological constraints are therefore extremely important.  A common feature of most ``unified'' models of SMBH and galaxy formation is the concept of co-evolution, whereby the SMBHs and their galaxies evolve under each other's influence. The term co-evolution means many things to many people \\citep[e.g.,][]{Silk98,Kauffmann00,Wyithe03,Granato04,Merloni04,Haiman04,Cattaneo05,Fontanot06}. For the purpose of having strawmen to take aim at, we outline three possible scenarios:  \\begin{itemize} \\item {\\bf Strong Co-evolution: } In its strongest sense, co-evolution implies that bulges and black holes grow {\\it together}, i.e., in the same objects and at the same time. In this case, BH growth could be thought of as a sort of a `tax': whenever star formation occurs, the SMBH accretes some fraction of the mass involved. This scenario in its simplest form can already be ruled out --- there are many examples of star-forming bulges without AGN activity and bulges with AGN activity and no SF (e.g., M87). \\item {\\bf Time Offset: } One could imagine a scenario in which bulge stars and SMBHs are formed in the same event, but on different timescales and/or stages of the event. For example, star formation might occur predominantly in the early stages of a merger, while BH growth might occur in the later stages \\citep[e.g.,][]{DiMatteo05,Hopkins05,Hopkins06}. Such a picture would predict little evolution in the BH-bulge mass correlation for an unbiased ensemble, but galaxies in the throes of their bulge building events (or black holes in their most rapid growth phase) could exhibit significant deviations from the scaling relations; this hypothesis is considerably more challenging to rule out. \\item {\\bf Regulated Growth: } A third possibility is that the major episodes of SF and BH accretion occur in different events, but the growth of one component regulates that of the other. For example, if black hole growth is regulated by feedback from the AGN itself, as suggested by many workers \\citep[e.g.,][]{Silk98,Murray05,DiMatteo05}, the potential well of the pre-existing bulge determines how much energy is needed to stop further accretion and halt the growth of the black hole, and hence the final black hole mass. It is possible, even likely, that the bulge stars that dominate this potential well are formed mostly in previous star formation ``events'', i.e. in progenitor disks which have subsequently merged. In this picture, one would expect a larger possible disconnect between observable episodes of SF vs. BH accretion {\\it activity}, even while the endpoints of this activity (bulges and black holes) are expected to end up being tightly related. \\end{itemize}  There are a number of ways to attack the problem of co-evolution observationally.  Obviously, measuring the BH-bulge mass correlation at the present day at extremely low or high mass will provide powerful insights into the processes linking galaxy and SMBH evolution \\citep{Barth05,Lauer07}.  The study of the BH-bulge mass correlation during BH or bulge-building events is also an important avenue of approach \\citep[e.g.,][]{Barth05,Greene06,Treu04,Treu07,Woo06,Peng06,Borys05}. The most commonly-used technique in the latter approach is studying the masses of actively-accreting Type I AGN for which BH masses can be estimated using an empirically-calibrated combination of luminosity and broad line region line width \\citep[][and references therein]{Wandel99,Baskin05,Kaspi05,Vestergaard06}; current results are controversial but tentatively support larger BH to bulge mass ratios at early times than those seen today \\citep[][see also, e.g., \\citealt{Alexander05} for the smaller ratios of black hole to bulge mass in rapidly-star-forming galaxies detected in the submm]{Treu04,Woo06,Peng06,Shields06}.  With current generations of wide and deep cosmological surveys, another means has opened up: comparison of evolution of the BH accretion rate (BHAR) and that of the SFR of galaxies \\citep[\\citealt{Silverman09}; see \\citealt{Boyle98} and][for early efforts]{Franceschini99}. The history of BH accretion has been explored using primarily optical and X-ray techniques, with important insights being gained from observations in the radio and infrared \\citep{Miyaji00,Ueda03,Hasinger05,LaFranca05,Barger05,Donley05,Brown06,Richards06a}. Furthermore, a decade of intensive multi-wavelength and redshift surveying has placed powerful and interesting constraints on the volume-averaged SF history of the Universe \\citep{HB06,PerezGonzalez05,LeFloch05,Schiminovich05}. In both cases, the integrated BH accretion and SF histories give values that match to within a factor of 2 or 3 of the present-day BH mass and stellar mass\\footnote{In the case of stellar mass density, recycling of nearly 1/2 of the initial stellar mass back into interstellar gas must be accounted for to avoid dramatically overproducing the present-day stellar mass density.} densities \\citep{Yu02,Marconi04,Borch06,Fardal07, Wilkins08}; we will discuss in detail uncertainties and possible inconsistencies between the integral of the SFR and the present-day stellar mass in \\S \\ref{sec:overall}.  The observational maturity of current datasets at $z<1$ allows the use of a novel angle of attack on the problem of co-evolution.  Recent multi-wavelength/HST surveys have covered enough volume that SFR density and BHAR density can be estimated in well-defined subsamples of galaxies for the first time, offering potentially decisive insight into not only the statistical evolution of these quantities, but into whether black holes and bulges are actually growing {\\it in the same objects and at the same time}.  This is the approach we will adopt in this paper.  To do this, we first examine the global SFR and BHAR (\\S \\ref{sec:overall}), then the multiplicity functions\\footnote{The comoving number density of sources with a given SFR or BHAR --- e.g., a luminosity function is a multiplicity function} of SFR and BHAR as a function of redshift (\\S \\ref{sec:multifuncs}).  In \\S \\ref{sec:host}, we present new measurements of SF activity in mass and morphology-limited samples of galaxies, and examine accretion activity driven by massive BHs in massive spheroids under assumptions widely used in literature.  In \\S \\ref{sec:summ}, we discuss the results in the context of co-evolution and summarize our conclusions.  Throughout the paper we assume a cosmology with H$_0$\\,=\\,70\\,km\\,s$^{-1}$\\,Mpc$^{-1}$, $\\Omega_{\\rm M}$\\,=\\,0.3 and $\\Omega_{\\Lambda}$\\,=\\,0.7. ",
        "conclusions": "\\label{sec:summ}  We have shown that the global SF history is proportional to the global BH accretion up to (at least) $z\\sim 1$, and the proportionality factor (2000) is that expected if the BHAR is a fixed fraction of the star formation that ends up in galactic bulges. This parallel between the cosmic SF history and the cosmic BH accretion history is an important constraint on the co-evolution between SMBHs and galaxies. It suggests that {\\it overall} accretion-driven BH growth and {\\it overall} SF-driven galaxy growth trace each other, and that the globally averaged BH mass to stellar mass ratio remains constant. Furthermore, we find that the SFR and BHAR {\\it multiplicity functions} are also scaled versions of one another at $z<1$, suggesting that the intensity of star formation and BH accretion are at least statistically linked.  We then refined our analysis to test the hypothesis that star formation and BH accretion do not just trace each other in a statistical fashion, but that {\\it the SFR and BHAR are related by a fixed factor} $f_{\\rm co} \\equiv f_{\\rm recycle} \\, f_{\\rm bg-BH} \\simeq 1300$ {\\it in every bulge/BH-building event}. We focused our analysis on massive galaxies with $M_* > 2 \\times 10^{10}$\\,M$_{\\sun}$, for which our observational sample is complete to $z\\sim 1$. Assuming that the local relationship between BH mass and bulge mass remains constant, spheroids with $M_* > 2 \\times 10^{10}$\\,M$_{\\sun}$ should host BH with mass $M_{\\rm BH}\\ge 2.6\\times 10^{7}$\\,M$_\\odot$. Assuming an average Eddington ratio $L_{\\rm bol}/L_{\\rm Edd}=0.25$ \\citep{Kollmeier06}, and $\\epsilon =0.1$, these BH, when active, power AGN with $BHAR \\geq 0.13 \\sfrunit$ or bolometric luminosity $\\log L_{\\rm bol}\\ge 44.9$. We computed the BHAR contributed by AGN with bolometric luminosities above this limit, and compared it with the SFR contributed by galaxies with $M_* > 2 \\times 10^{10}$\\,M$_{\\sun}$. We found that the SFR and BHAR again trace each other  but are offset by approximately the factor of $f_{\\rm co} \\simeq 1300$ expected for massive spheroid-dominated galaxies. This result is nearly unchanged if we compute the SFR from galaxies limited in both stellar mass and SFR (SFR\\,$\\ge 3\\sfrunit$), as these ``high-intensity'' SF events dominate the SFR budget in massive galaxies. The massive galaxies are simply dominating the total SFR, and the big accretion systems are dominating the accretion history (although with a little more evolution). Just as before when we compared the overall SFRD with the overall BHAR we got a match, we should get a match for the top two panels.  Further examination of the nature of the galaxies hosting star formation and BH activity showed that only about 30 \\% of the SF ``budget'' needed to match the BHAR is observed in massive {\\it spheroid dominated} galaxies. Furthermore, the distribution of S\\'ersic indices for star forming galaxies and that for AGN detected in the soft X-ray indicate that these two classes of activity tend not to reside in the same types of host galaxies. The bulk of star formation occurs in {\\it isolated disks} while the majority of luminous AGN hosts are intermediate- and early-type galaxies \\citep[see also][]{Grogin05,Nandra07,Georgakakis08,Alonso08,Watson09,Gabor09}. These results clearly rule out the simplest picture of object-by-object co-evolution, in which BH and bulges grow proportionally and simultaneously. This is also supported by the detection of a small fraction of luminous X-ray detected AGN in spirals, indicating that BH accretion may happen in non-bulge-building events \\citep[e.g.,][]{Georgakakis09}.  There are both observational claims and theoretical expectations that the relationship between BH mass and spheroid mass might have been larger in the past. While this evolution would impact our calculation by shifting the mass of the BH (and therefore the cutoff luminosity for AGN) to include in our ``matched'' BHAR, this would only cause a minor change in the slope of the BHAR and cannot remove the mismatch between SF and BHA in spheroid dominated galaxies. Therefore, this would not change any of our conclusions.  There are two possible ways to reconcile this apparent paradox.  One possibility (which shall now call Merger Scenario 1) is that bulges and BH grow in proportion to one another, and in the same {\\it events}, but with an offset in time.  For example, in simulations of galaxy-galaxy mergers including BH growth, the luminous AGN phase is delayed by about a dynamical time (a few hundred Myr) relative to the peak of the starburst activity \\citep{DiMatteo05,Hopkins05,Hopkins06}. In this picture, we would expect that most AGN hosts would resemble late-stage mergers, i.e. objects with spheroidal morphology but young stellar populations, in plausible agreement with our observational results. Moreover, in this picture, we would expect that a small fraction of AGN to be detected in early stage mergers, and the star formation associated with this BH growth should be observed in early-stage mergers, i.e. close pairs or highly morphologically disturbed objects.  However, other studies have found that only a small fraction of the global star formation at $z<1$ is occurring in galaxy mergers \\citep{Bell05,Wolf05,Jogee09,Robaina09}: the strongest limit is set by Robaina et al., who find that $<$10\\% of the star formation in massive galaxies ($m_* > 10^{10} M_\\odot$) is directly triggered by {\\it all identifiable phases} (from close pairs to morphologically disturbed remnants) of major galaxy merging. This is strong evidence against the Merger Scenario 1. Therefore, while a small fraction of the star formation associated with the BH growth that we witness at $z\\la 1$ may be occurring in the same merger events that feed the BH, it seems that this does not represent the majority of the SF activity needed to ``match'' the observed BH accretion.  The second possibility (Merger Scenario 2) is that most star formation takes place in isolated disks. When these disks merge, the pre-existing stars are scrambled into a dynamically hot spheroidal remnant\\footnote{In what follows, much of what we argue for merger scenario 2 would also apply if the main route for forming low-mass spheroids were secular evolution.  The key point for this paper is that the bulk of the stars that end up in the spheroid are formed long before the epoch of actual bulge creation, which applies equally in Merger Scenario 2 and secular evolution.}. The presence of even a modest amount of gas leads to the formation of a remnant with higher phase-space density and hence a deeper gravitational potential well \\citep{Dekel-cox06,Robertson06,Cox06}. If the depth of the potential well in the central parts of the galaxy determines how large the BH can grow, as in the picture of self-regulated BH growth outlined by \\citet{Hopkins07a}, then the post-merger BH will end up on the BH mass-bulge mass relation, {\\it despite the fact that the SFR and BHAR during the merger event may not necessarily obey the universal proportionality factor}.  One can easily see that in this scenario, in many cases the BH accretion and the star formation will not trace each other during the merger event. For example, in a major merger of two pure disk progenitors, the scaled BHAR will be larger than the SFR associated with bulge growth. Most of the bulge mass will be contributed by the pre-existing stars in the progenitor disks, and the BH will grow to ``catch up'' with the bulge. In the opposite extreme, in a merger of two galaxies with pre-existing massive BH, there may be very little BH accretion because the AGN will almost immediately ``shut itself off''. In this case, the SF in the merger might be larger than the scaled BH accretion (if the progenitors contain enough gas).  Because what we see is more BH accretion than we would expect based on the amount of star formation in spheroids, we conclude that most of the mass in these spheroids was contributed by the pre-existing stars in the progenitors rather than by new stars formed in the merger-associated bursts. Indeed, we have independent lines of evidence that this must be the case. We know that the colors of elliptical galaxies and bulges at $z\\la 1$ are characteristic of fairly old stellar populations \\citep[][and references therein]{Gallazzi06,Renzini06}, and similarly, fossil evidence from line strengths in nearby ellipticals also precludes the very recent formation of a large fraction of the stellar mass \\citep{Trager00,Thomas05}. In addition, this picture is consistent with the relatively small fraction of star formation associated with ongoing mergers at $z<1$: in cosmological models that incorporate a scenario like Merger Scenario 2 \\citep[e.g.][]{Somerville08}, the predicted fraction of star formation triggered by major mergers at $z<1$ is about 7\\%, in excellent agreement with the observational estimates (see Robaina et al. 2009 for details of the comparison).  A slightly modified version of Merger Scenario 2 can also be accommodated in a Universe in which the BH-spheroid mass ratio has decreased with time. This is in fact expected in the theoretical picture that has emerged from the analysis of hydrodynamic simulations of galaxy mergers with self-regulated BH growth \\citep[e.g.][]{Robertson06b,Hopkins07a}. These studies find that the final BH-spheroid mass ratio depends on the gas fraction of the progenitors. This is because gas-rich progenitors can dissipate energy during the merger, and all else equal, form more compact remnants than gas-poor mergers \\citep{Dekel-cox06,Robertson06,Cox06,Covington08}. In this deep potential well, the BH can grow to a larger mass (relative to the spheroid) before the pressure-driven outflow halts further accretion.  If gas fractions in merging galaxies were higher at high redshift, as is generally expected, this would then lead to the prediction that BH were more massive relative to their spheroids at high redshift. A significant fraction of mergers at $z<1$ are expected to involve at least one galaxy that already contains a massive BH. If the pre-existing BH is already larger than the ``critical mass'' set by the new remnant, further BH growth will be stifled. Theoretical models predict that this expected evolution is relatively mild: less than a factor of two since $z\\sim1$ \\citep{Hopkins07a}.  \\citet{Borch06} have shown that the stellar mass in red (predominantly spheroidal) galaxies has increased by a factor of 2--3 since $z\\sim1$, while the stellar mass in blue (disk-dominated) galaxies has remained about the same over this time period. At the same time, we know that the bulk of new star formation is occurring in blue (disk-dominated) galaxies, implying that mass in blue (disk) galaxies must be transformed into red (spheroidal) galaxies at approximately the same rate that new stars are forming \\citep{Bell07}. Given that most of the red galaxies are spheroid-dominated, which is well established at least at $z<1$ \\citep{Bell04b}, the mass on the red sequence cannot grow simply via the quenching of star formation in disks --- a strong dynamical process like merging is required.   Taken together, this evidence strongly favors a picture in which, at least at $z<1$, the bulk of the growth of mass in spheroids is due to merging of pre-existing, already old disk stars, and BHs preferentially grow to ``catch up'' to their newly assembled bulges following a merger event.  In this case, the parallel evolution of the global SFR and BHAR may simply reflect the apparent coincidence that, at $z<1$, the rate of formation of new stars is approximately equal to the rate that stellar mass is transferred from disks into bulges via mergers."
    },
    "0911/0911.2236_arXiv.txt": {
        "abstract": "We measure the average mass properties of a sample of 41 strong gravitational lenses at moderate redshift (z $\\sim$ 0.4 -- 0.9), and present the lens redshift for 6 of these galaxies for the first time. Using the techniques of strong and weak gravitational lensing on archival data obtained from the Hubble Space Telescope, we determine that the average mass overdensity profile of the lenses can be fit with a power-law profile ($\\Delta\\Sigma \\propto R^{-0.86 \\pm 0.16}$) that is within 1-$\\sigma$ of an isothermal profile ($\\Delta\\Sigma \\propto R^{-1}$) with velocity dispersion $\\sigma_v$ = 260 $\\pm$ 20 km~s$^{-1}$.  Additionally, we use a two-component de Vaucouleurs+NFW model to disentangle the total mass profile into separate luminous and dark matter components, and determine the relative fraction of each component. We measure the average rest frame V-band stellar mass-to-light ratio ($\\Upsilon_V = 4.0 \\pm 0.6~h~ M_{\\odot}/L_{\\odot}$) and virial mass-to-light ratio ($\\tau_V = 300 \\pm 90~h~M_{\\odot}/L_{\\odot}$) for our sample, resulting in a virial-to-stellar mass ratio of $M_{vir}/M_{*} = 75 \\pm 25$.  Relaxing the NFW assumption, we estimate that changing the inner slope of the dark matter profile by $\\sim$20\\% yields a $\\sim$30\\% change in stellar mass-to-light ratio.  Finally, we compare our results to a previous study using low redshift lenses, to understand how galaxy mass profiles evolve over time.  We investigate the evolution of $M_{vir}/M_{*}(z) = \\alpha(1+z)^{\\beta}$, and find best fit parameters of $\\alpha = 51 \\pm 36$ and $\\beta = 0.9 \\pm 1.8$, constraining the growth of virial to stellar mass ratio over the last $\\sim$7 Gigayears.  We note that, by using a sample of strong lenses, we are able to constrain the growth of $M_{vir}/M_{*}(z)$ without making any assumptions about the IMF of the stellar population. ",
        "introduction": "Observations over the last few decades have suggested that galaxies are embedded in an extended, diffuse halo of dark matter \\citep[e.g.][]{pet78,whi78,rub79,blu84}, which contributes up to 95 percent of the total mass \\citep{hoe05,jia07,avi08}.  Because of the overwhelming fraction of dark matter, numerical simulations dealing with the formation and assembly of galaxy mass often focus on dark matter alone \\citep[e.g., the Via Lactea Simulation;][]{kuh08}, and these simulations have made very specific predictions about the overall profile of galaxy-sized mass distributions \\citep[e.g.,][]{nav97,moo98,jin00,sto06,die07,sch08}.  However, when compared to observations, these simulations often fail to accurately recreate the observed shape of galaxy mass profiles \\citep{sal01,gen04,sim05,gen07,kuz08,sei08,san08}. This disagreement is attributed largely to the complex physics of baryon interaction (e.g. frictional dissipation and radiative cooling) which is still not well understood and is difficult to model.  There are many possible methods capable of measuring the distribution of matter on small ($<$ 1 Mpc) scales, and techniques used for characterizing a radial mass profile include measuring the rotation curves from planetary nebulae \\citep{rom03,arn04,mer06,del08,nap09}, kinematics of stellar populations \\citep{ber94,cap06,van08} and \\ion{H}{1} gas \\citep{uso03,jac04,and06,mat08}, or the temperature of X-ray emitting gas \\citep{hum06,chu08}.  In this paper, we use a powerful alternative technique: gravitational lensing \\citep[e.g.,][]{bra96,fis00,wil01,kle06,par07}.  The power of lensing is due to the fact that the technique is able to directly trace the total mass (luminous + non-luminous) enclosed within a given radius without needing to make any assumptions about the dynamical state of the mass in question.  In addition, lensing does not rely on the presence of kinematic tracers, which are often only visible in very low-redshift galaxies, and even then are typically not present in the outer halo regions where dark matter dominates.  Furthermore, since we analyze a sample of strong lenses, we are able to combine the mass constraints determined from both strong lensing and weak lensing.  This allows us to probe the total mass distribution over a much wider range of physical distances than by using weak lensing alone, and also allows us to constrain the properties of stellar mass distribution without selecting a specific stellar initial mass function (IMF).  Instead, we rely on directly-derived properties: lensing-inferred mass and total luminosity (although we do make an assumption about the shape of the dark matter halo).  For early-type galaxies, several lensing-based studies have shown the typical shape of the mass density profile to be consistent with an isothermal ($\\rho(r) \\propto r^{-2}$) model, given the uncertainties on the measurements, between distances of $\\sim 50-300 ~h^{-1}$ kpc \\citep{she04,man06}.  However, due to the low signal-to-noise ratio (SNR) on the lensing signal obtained from an individual galaxy, these studies rely on stacking a large number of galaxies into a single sample to determine the average profile of the entire stack.  For this work, we utilize deep imaging from the Hubble Space Telescope (HST). In general, space-based data will have a much higher density of background galaxies ($n_s$) when compared to data taken from the ground over an identical exposure time.  This increased background density count not only enables us to detect a weak lensing signal much closer to the center of the lensing galaxy, but also allows for a significant detection with fewer stacked galaxies ($\\rm N_{lens}$), because the expected weak lensing SNR scales as $\\sqrt{{\\rm N_{lens}}~n_s}$.  The ability to work on a small sample size is particularly useful, given the relative paucity of known strong lenses at all redshifts.  In this work, we are able to extend the results of \\citet[hereafter G07]{gav07}, who conducted a similar mass profile study on a small (22 lens) sample of low redshift (z $\\sim$ 0.2) early-type strong gravitational lenses collected from the Sloan Lens ACS Survey \\citep[SLACS; see e.g.][]{bol06,bol08} over a wide range of physical distances ($\\sim 3 - 300 ~h^{-1}$ kpc).  From this work, it was shown that, like their non lensing counterparts, the total mass density profile of strong lensing ellipticals could be described by a roughly isothermal model (although they note that the comparison between the data and the model was done without a formal fit).  This was thought to be due to the combination of baryonic and dark matter, the so-called Bulge-Halo Conspiracy, in which a relatively steep luminous matter profile and a shallower dark matter profile combine in such a way that the total mass distribution of the galaxy can be well approximated by an isothermal model \\citep{tre06}.  By comparing this data set to our moderate redshift sample, we thus explore the evolutionary trends of massive early-type galaxies.  This paper is organized as follows: In Section 2, we describe our lens sample, the lens selection criteria, and the data reduction pipeline used for this work.  At this point, the interested reader may turn to the appendix, where we describe our techniques for analyzing weak lensing shear, paying specific attention to signal to noise optimization and the removal of systematic errors.  In Section 3, we use a two-component model to determine the contributions of the stellar and dark matter components to the total profile.  We present all of our results in Section 4, the reader interested in the science may skip directly to this section.  We discuss the results and compare them to the SLACS results in Section 5, looking for evolutionary differences between the samples.  Finally, we summarize our results in Section 6.  Throughout this work, we assume a flat cosmology with $H_0 = 100 ~h$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_M = 0.3$, and $\\Omega_\\Lambda = 0.7$. All magnitudes presented in this paper are AB magnitudes. ",
        "conclusions": "Taking advantage of the large density of background galaxies obtained by relatively short ($\\sim 1800$ sec) ACS exposures, we present the results of a joint weak and strong lensing analysis of a sample of 41 massive elliptical galaxies at moderate ($z \\sim 0.4 - 0.9$) redshift, determining the mass overdensity profile over nearly 3 orders of magnitude in physical radius.  Using a 2-component de Vaucouleurs + NFW profile model, we are able to determine stellar and virial $M/L$ ratios for the sample, and disentangle the relative contributions of stellar and dark matter from the total mass budget.  Furthermore, we compare all of our results to those obtained from a subset of the lower redshift SLACS sample ($\\overline{z} \\sim 0.2$), placing constraints on the evolution of these mass properties over cosmic time.  Our results can be summarized as follows:  \\begin{itemize}  \\item We present new redshift information for 7 COSMOS lenses: we find lens redshifts for J095930.93+023427.7, 0038+4133, J100140.12+020040.9, 0211+1139, and 0216+2955, a source redshift for 0056+1226, and both lens and source redshifts for 0254+1430.  \\item The mass overdensity profile of our sample has a best-fit power-law profile of $\\gamma = -0.86 \\pm 0.16$ and is consistent with an isothermal model between $\\sim$ 10 $h^{-1}$ kpc and $\\sim$ 1000 $h^{-1}$ kpc.  \\item The best-fit stellar and virial $M/L$ ratios for our lens sample are given by $\\Upsilon_V = 4.0 \\pm 0.6 ~^{+1.1}_{-1.8} ~h~M_\\odot/L_\\odot$ and $\\tau_V = 300 \\pm 90 \\pm 75 ~h~M_\\odot/L_\\odot$ respectively.  \\item Given our average sample luminosity of $\\overline{L_V} = 4.5 \\times 10^{10}~h^{-2}~ L_\\odot$, we find an average virial mass of $\\overline{M_{vir}} = (1.4 \\pm 0.4 \\pm 0.3) \\times 10^{13}~h^{-1}~ M_\\odot$ and an average stellar mass of $\\overline{M_*} = (1.8 \\pm 0.3 ~^{+0.5}_{-0.8}) \\times 10^{11}~h^{-1}~ M_\\odot$.  \\item We find an average mass ratio of $M_{vir}/M_* = 75 \\pm 25 ~^{+36}_{-27}$, which is consistent with the lower-redshift SLACS sample ratio of $M_{vir}/M_* = 54 \\pm 28$.  \\item Using our assumed cosmology, we convert mass fraction into a stellar baryon fraction, finding that $f_* = 0.075 \\pm 0.030 ~^{+0.030}_{-0.040}$.  This result indicates a low stellar formation efficiency in early-type galaxies, in agreement with previous results \\citep{man06,hey06}.  \\item By comparing our mass fraction data to the lower-redshift SLACS sample, we are able to place constraints on the evolution of the virial to stellar mass fraction of massive early-type galaxies.  We find that the quantity evolves as $M_{vir}(z)$/$M_{*}(z)$ = $(51 \\pm 36)(1 + z)^{(0.9 \\pm 1.8)}$ over the last $\\sim$ 7 Gigayears.  This is the first such study to place constraints on the evolution of $M_{vir}$/$M_*$ without invoking a specific IMF.  \\end{itemize}  Overall, these results show that our method is a promising new technique: with only a small sample of lenses (41 systems), were are able to significantly measure the average mass overdensity profile and mass-to-light ratios of the sample, allowing us to place constraints on the mass evolution of early-type galaxies from $z \\sim 0.8$ to $z \\sim 0.2$.  As always though, this work can be strengthened by increasing the sample size.  While galaxy-scale strong lenses are relatively rare today ($\\sim$ 100 systems), future large-scale observational surveys such as LSST ($\\sim$ 10000 new lenses) and JDEM \\citep[$\\sim$ 10000 new lenses;][]{mar05} are expected to dramatically increase the number of known strong lenses.  Including these new lenses in future work will greatly reduce the size of our present statistical uncertainties.  In the near future however, we plan to increase our lens sample by including other known deep space-based lens data from CASTLES that have been imaged using the smaller HST Wide Field / Planetary Camera 2 (WFPC2).  Though the smaller area will limit our ability to probe out to the large radii made available with the ACS sample, the longer exposure times for the WFPC2 data yield background densities that are comparable to those in the ACS sample, allowing us to place better constraints on the inner mass overdensity profile by incorporating these galaxies into our current lens sample.  We will also take advantage of the HAGGLeS strong-lens search of bright red galaxies \\citep{mar09} which is expected to explore the whole of the HST ACS archive, discovering $\\sim$ 10 lenses per square degree of coverage. We plan to further augment this study by combining our ACS sample with the full ($\\sim$ 40) long-exposure SLACS sample, which will allow us to probe the differences in the mass properties of strong lenses across a wider variety of samples and categories, such as morphology, luminosity, and finer redshift slices.  We are particularly interested in segregating the lens sample by total luminosity, since the SLACS lenses appear to be more luminous than their moderate-redshift counterparts when passively evolved to $z = 0$ (Figure \\ref{fig:mlum}), suggesting that the SLACS sample may have evolved from a slightly different population of galaxies.  Of course, by increasing the sample size of lenses and thus reducing the statistical errors associated with these measurements, correcting systematic uncertainties will become much more important.  In future studies we hope to improve our control of systematics in several ways. Foremost is the ability to measure accurate redshifts for both the lens and source galaxies of strong lens systems, as these measurements are crucial for accurate lens modeling.  While some future survey instruments should provide redshift information (either spectroscopically or photometrically) as part of their normal operation, it is still important to improve the redshift information for currently known lenses, and we plan to obtain more spectroscopic and photometric redshift data in the near future.  Additional systematics controls could come in the form of improved profile modeling, relaxing the strict de Vaucoulers + NFW mass profiles in favor of a more general free-index S\\'{e}rsic profile for the stellar mass and a freely-varying gNFW profile for the dark matter. Similarly, improved strong lens mass modeling should be utilized in the future, increasing our ability to distinguish between various mass profiles at very small scales, as well as improving the precision of our stellar mass fraction measurements.  Finally, we could improve galaxy selection and weak lensing measurements by using HST data obtained through other filters, allowing us to determine photometric redshifts for many of the background source galaxies.  We note however that this should not be a dominant source of systematic error, as we have shown previously that altering the background galaxy redshift distribution does not significantly change the results of any of our fitted parameters."
    },
    "0911/0911.3752.txt": {
        "abstract": "We describe the procedure of the construction of a sample of distant ($z>0.3$) radio galaxies  using the NED, CATS and SDSS databases for further use in various statistical tests. We believe the sample to be free of objects with quasar properties. This paper is the third part of the description of the radio galaxies catalog that we plan to use for cosmological tests. We report the results of the sample of angular sizes for the NVSS survey list objects, and its preliminary statistical analysis. Three-parameter diagrams ``angular size--redshift--flux density'' and ``angular size--redshift--spectral index'', and their two-parameter projections are constructed. Three subsamples of radio galaxies are separated in the ``source size--flux density'' diagram. \\mbox{\\hspace{1cm}}\\\\ PACS: 98.54.Gr, 98.62.Py, 98.70.Dk, 98.80.Es %\\pacs{98.54.Gr, 98.62.Py, 98.70.Dk, 98.80.Es} ",
        "introduction": "The present work is the third part of the catalog of radio galaxies located within $0.3<z<5.2$ redshifts. The first two parts \\cite{rg_listI:Verkhodanov3_n_en,rg_listII:Verkhodanov3_n_en} (further, Papers I and II) describe the way  the sample was constructed, based on the NED\\footnote{\\tt http://nedwww.ipac.caltech.edu}, CATS\\footnote{\\tt http://cats.sao.ru}, and SDSS\\footnote{\\tt http://sdss.org} data, including an analysis of other works with the aim of excluding the objects with quasar properties. Paper~I lists 2442 objects with their equatorial coordinates for epoch J2000.0, spectroscopic redshifts, and spectral indices, deduced according to the cross-identification results of the selected radio galaxies with the CATS radio catalogs, and computed at 325, 1400 and 4850\\,MHz. The list of objects in Paper~I also contains the names of the initial catalogs from where we adopted the data for the objects (and the names of radio sources in case of the 3C and 4C surveys). Paper~II presents photometric data from the NED and SDSS archives. We constructed the ``magnitude--redshift'' Hubble diagrams for 11 filters. For the SDSS survey \\cite{sdss:Verkhodanov3_n_en} (g, i, r, z passbands) we determined the parameters of linear regression on the Hubble diagram. The preparation of the catalog presented in the three papers is the first step for making cosmological tests on the observed physical parameters of radio galaxies at different epochs.  An  objective of our catalog was a manifold increase of the comparatively uniform sample of radio galaxies, which would allow increasing the accuracy of the determination of cosmological parameters in an independent set of tests, apart from the cosmic microwave background and a large-scale structure. In the last part of the catalog, which is supplemented with this paper, we give the list of objects with their angular sizes from the NVSS survey\\footnote{\\tt http://www.cv.nrao.edu/nvss/} \\cite{nvss:Verkhodanov3_n_en} and the flux density measurements data at 1400\\,MHz.  The radio source characteristics collected in this list can be used in the ``radio luminosity--redshift'' and ``angular size--redshift'' tests. Both tests are connected with the galaxy characteristics that greatly evolve. Hence, the evaluation of the cosmological parameters, based on them, has a large dispersion. Nevertheless, the ``angular size--redshift'' (or \\mbox{``$\\theta-z$'')} diagram is widely used in the papers, that consider both the minimal size of point sources \\cite{gurvits:Verkhodanov3_n_en}, and the average angular size of a radio galaxy at a given redshift, corrected for evolutionary effects \\cite{guerra:Verkhodanov3_n_en}. The angular size of a radio source depending on its redshift can be used in a combined test with the spot sizes in the cosmic microwave background ~\\cite{jackson_janneta:Verkhodanov3_n_en}, where both the scale of the CMB fluctuations, and the characteristic sizes of millisecond sources are used.  Note once more that all the methods of parameter finding from the relation ``$\\theta-z$'' are aggravated with uncertainties linked with evolutionary effects, and the   \\begin{figure*}[tbp] \\centerline{\\hbox{ \\psfig{figure=fig1_map_sources.ps,width=14cm} }} \\caption{Sky positions of selected radio sources in Galactic coordinates. Objects location is determined by the simultaneous presence of observational data in optical and radio ranges. White circles indicate SDSS objects and gray crosses mark all other sources.} \\label{f1:Verkhodanov3_n_en} \\end{figure*}    \\noindent extended objects, in addition, with the task of determining the outer edge of an object. Nevertheless, using the source size as a standard ruler is justified. If radio galaxies have a fundamental plane, determined by such physical parameters as, e.g., the mass of the radio galaxy or its central black hole, which was already mentioned in Paper\\,II, then we have reason to link the angular size of the source with the history of formation of the active nucleus, described by a set of standard measurable parameters, like the accretion velocity, angular momentum, magnetic field magnitude, and mass of the central black hole. This gives us grounds to expect that the proposed in the literature approaches of parameter estimation of the cosmological model will work as well. Note that the highest certainty of parameter estimation may be reached via a combination of independent tests (for radio galaxies, see, e.g., \\cite{vo_par1:Verkhodanov3_n_en,vo_par2:Verkhodanov3_n_en,vo_par3:Verkhodanov3_n_en}).  Describing the catalog in Papers I and II, we made a preliminary statistical analysis, and constructed the Hubble diagrams, describing the behavior of different parameters with the change in redshift. In the present work, devoted to the construction of a sample, and a statistical analysis of the radio galaxy catalog, we collected the objects with known spectroscopic redshifts and with measured angular sizes. To make the sample uniform, we adopted the angular sizes from the NVSS catalog \\cite{nvss:Verkhodanov3_n_en}, obtained on the Very Large Array (VLA) with the beam of 45\\arcsec. In further studies we shall as well use the angular size data from other surveys, like, for example, the FIRST survey \\cite{first:Verkhodanov3_n_en} data with the beam of 5\\arcsec. Below we describe our sample, make its statistical analysis and construct diagrams for various physical parameters depending on redshifts. ",
        "conclusions": "We report the final part of the radio galaxies catalog with redshifts $z>0.3$, which contains complete information on the angular sizes of the objects, and flux densities at 1400\\,MHz. The catalog is based on lists of objects from the NED and CATS databases. Our complete sample includes the SDSS objects, and the objects from the ``Big Trio'' program \\cite{par_big3a:Verkhodanov3_n_en,par_big3b:Verkhodanov3_n_en}. We perform a preliminary statistical analysis of the angular size data, flux density and spectral index calculations. On the diagram ``source size--flux density'' we separated three radio galaxy subsamples, varying in angular sizes and flux densities. We approximate the boundaries of these subsamples based on analytic functions. The bottom object separation on the ``source size--flux density'' diagram is preconditioned by the selection effect defined by the size of the VLA directivity pattern in the NVSS survey. In order to use the sources from this area in the cosmological tests, connected with the source size, some better resolution data are needed.  %"
    }
}